{
    "1112/1112.5249.txt": {
        "abstract": "The 3D structure of the solar wind and its evolution in time is needed for heliospheric modeling and interpretation of energetic neutral atoms observations. We present a model to retrieve the solar wind structure in heliolatitude and time using all available and complementary data sources. We determine the heliolatitude structure of solar wind speed on a yearly time grid over the past 1.5 solar cycles based on remote-sensing observations of interplanetary scintillations, \\textit{in~situ} out-of-ecliptic measurements from \\textit{Ulysses}, and \\textit{in~situ} in-ecliptic measurements from the OMNI-2 database. Since the \\textit{in~situ} information on the solar wind density structure out of ecliptic is not available apart from the \\textit{Ulysses} data, we derive correlation formulae between the solar wind speed and density and use the information on the solar wind speed from interplanetary scintillation observations to retrieve the 3D structure of solar wind density. With the variations of solar wind density and speed in time and heliolatitude available we calculate variations in solar wind flux, dynamic pressure and charge exchange rate in the approximation of stationary H atoms. ",
        "introduction": "The goal of this paper is to retrieve the solar wind structure at 1~AU as a function of time and heliolatitude based on available \\textit{in~situ} data sources and interplanetary scintillation observations for the time interval since 1990 until present that covers 1.5 solar activity cycles.  The existence of the solar wind was predicted on a theoretical basis by \\citet{parker:57a} and discovered experimentally by \\textit{Lunnik}~II and \\textit{Mariner}~2 at the very beginning of the space age \\citep{gringauz_etal:60a, neugebauer_snyder:62a}. Regular measurements of its parameters began in the early 1960s and data from many spacecraft are now available, obtained using various techniques of observations and data processing \\citep[see for review][]{bzowski_etal:12b}. Shortly after the discovery of the solar wind (SW hereafter), a question of whether or not it is spherically symmetric was put forward. Most spacecraft with instruments to measure the SW parameters are at orbits close to the ecliptic plane and the information about the latitudinal structure of the SW is hard to obtain. There are a few sources of data on the out of ecliptic SW parameters, but only one of them from \\textit{in~situ} measurements, namely from \\textit{Ulysses}.  While direct observations of the SW in the ecliptic plane have been collected for many years, information on its latitudinal structure had been available only from indirect observations of the cometary ion tails \\citep{brandt_etal:75a}, until radio-astronomy observations of interplanetary scintillation \\citep{hewish_etal:64a, coles_maagoe:72a} and spaceborne measurements of the Lyman-$\\alpha$ helioglow \\citep{lallement_etal:85a, bertaux_etal:95} became available. To our knowledge these two techniques remain the only source of global, time-resolved information on the the solar wind structure. The launch of the \\textit{Ulysses} spacecraft \\citep{wenzel_etal:89a} improved our understanding of the 3D behavior of the solar wind by offering direct \\textit{in~situ} observations and a very high resolution in latitude, but a poor resolution in time\\footnote{The same latitudes were visited only a few times during the $\\sim 20$-year mission.}.  The solar wind structure varies in latitude with the solar activity cycle. Knowledge of its evolution is needed to construct credible models of the heliosphere and its boundary regions. With the history of the SW evolution based on a homogeneous series of data and retrieved using homogeneous analysis method, one obtains a tool to interpret both present heliospheric observations, such as the ongoing measurements of Energetic Neutral Atoms (ENAs) by the NASA \\textit{Interstellar Boundary Explorer}~\\citep[IBEX,][]{mccomas_etal:09a} and \\textit{in~situ} measurements of the heliospheric environment by the Voyagers, and to compare them with the results from past and current long-lived experiments such as \\textit{Solar Wind ANisotropy}~(SWAN) onboard \\textit{SOlar and Heliospheric Observatory}~(SOHO) etc.  We assume that the solar wind expansion is purely radial, its speed does not change with solar distance, and density drops down quadratically with distance from the Sun. These assumptions are valid for close distances from Sun $\\left( r < 10~AU \\right)$. Their validity at farther distances will be considered in the Discussion section.  In the following text ``CR-averaged'' data mean Carrington rotation (CR hereafter) averaged values, where Carrington rotation period is the synodic period of solar rotation equal to 27.2753 days \\citep{franz_harper:02a}. Throughout the paper, ``adjusted'' means scaled from the value $x_0$ measured at a heliocentric distance $r_\\mathrm{0}$ to the distance $r_\\mathrm{E} = 1$~AU, where the value specific for $r_\\mathrm{E}$ is calculated as $x_\\mathrm{E} = x_0 \\left( r_0 / r_\\mathrm{E} \\right)^2$. ",
        "conclusions": "To check the credibility of our results we compare them with the \\textit{Ulysses} data from all scans. The \\textit{Ulysses} data are prepared by splitting the hourly data into full Carrington rotations and calculating average values. Next we linearly interpolate our model values to the times and heliolatitudes corresponding to the Carrington rotation-averaged \\textit{Ulysses} data.  A comparison is presented in Figure~\\ref{figCompareModelWithUlysses}. The agreement is quite good in the ecliptic parts of the \\textit{Ulysses} orbit and satisfactory for higher latitudes. The model retrieves the fast solar wind speed, but some discrepancies exist for the slow and variable solar plasma flows. We have a better agreement in solar wind speed than in density, which is understandable, since the density values were derived from the already approximate speed values using approximate statistically-derived solar wind density-speed relations.  The worst agreement between \\textit{Ulysses} density measurements and our model is during the third slow scan (during descending phase of solar activity), when \\textit{Ulysses} \\textit{in~situ} measurements for this solar cycle phase give highly variable values, nearly $50\\%$ greater than our model predictions. The source of the disagreement might be connected with the density-speed correlation formula we adopted for this time interval: for solar maximum we use an average formula from the two fast scans during solar minimum assuming that the correlation does not change with the solar cycle.  The overall agreement, however, is much better. The residuals of speed are typically $10\\%$, not exceeding $30\\%$, and typical residuals of density are $20-30\\%$, not exceeding $60\\%$. The sign of the residuals varies, which suggests there is no systematic global bias in our method. Given all the uncertainties in the absolute calibration of both \\textit{in~situ} measurements and IPS observations and the relative simplicity of our approach, we believe such a level of agreement between the model and the measurements is quite good and hard to improve on without an additional data source.  Such a source of additional information might result from an inversion of photometric maps of the Lyman-$\\alpha$ helioglow obtained from SWAN/SOHO observations \\citep{lallement_etal:10a}, aimed at calculation of the total ionization rate of neutral interstellar hydrogen in the inner heliosphere as a function of heliolatitude and Carrington rotation. With this information, one might be able to follow the idea presented by \\citet{bzowski_etal:12b} and independently calculate the profiles of solar wind density.  In this analysis we assumed that the solar wind parameters obtained from OMNI-2, \\textit{Ulysses} and IPS are directly comparable, i.e. that there is no systematic change in solar wind speed with the solar distance between the region from the solar wind acceleration region (a dozen solar radii) to 1~AU, from which the IPS solar wind speeds are retrieved, and \\textit{Ulysses}, which measured between $\\sim 1.4$ and $\\sim 5.5$~AU. Also, we assume there is no distance-related change in the density other than the simple $1/r^2$ scaling.  The assumption of purely radial expansion of solar wind implies that the latitude structure we inferred in this paper does not change until the termination shock. This may not be exactly true, as pointed out by \\citep{fahr_scherer:04b}, who argue that the pickup ions (PUIs) pressure induces nonradial flows at large  heliocentric distances. Such flows would cause changes both in the radial component of solar wind speed and in the local density. A more thorough study of this effect requires, in our opinion, using a multifluid, 3D and time dependent MHD modeling with turbulence and boundary/initial conditions taken from observations. Such a model, to our knowledge, is still in development \\citep{usmanov_etal:11a}. Our results seem to be well suited as the boundary conditions for such modeling.  The assumption of a $1/r^2$ drop in density and constancy of speed is increasingly invalid with an increase of solar distance because of the interaction with neutral interstellar gas, which results in the creation of pickup ions and slowdown and heating of the distant solar wind. These phenomena were extensively discussed by \\citet{fahr_rucinski:99,fahr_rucinski:01a,fahr_rucinski:02a,fahr:07a,lee_etal:09a,richardson_etal:95a, richardson_etal:08a,richardson_etal:08b}.  However, for the global modeling of the heliosphere and calculation of survival probabilities it is the total flux of solar wind, being a sum of the core solar wind and pickup ions, that counts most. In this respect, the total solar wind flux quite exactly follows the $1/r^2$ scaling due to the continuity conditions, as discussed by \\citet{bzowski_etal:12b}. This is because a vast majority of PUIs are created due to charge exchange and thus are not new members of the solar wind population.  \\begin{figure*}[!h] \\centering \\begin{tabular}{cc} \\includegraphics[scale=0.35]{figCompareModelWithUlysses_speed1s} & \t\\includegraphics[scale=0.35]{figCompareModelWithUlysses_dens1s}\\\\ \\includegraphics[scale=0.35]{figCompareModelWithUlysses_speed1f} & \t\\includegraphics[scale=0.35]{figCompareModelWithUlysses_dens1f}\\\\ \\includegraphics[scale=0.35]{figCompareModelWithUlysses_speed2s} & \t\\includegraphics[scale=0.35]{figCompareModelWithUlysses_dens2s}\\\\ \\includegraphics[scale=0.35]{figCompareModelWithUlysses_speed2f} & \t\\includegraphics[scale=0.35]{figCompareModelWithUlysses_dens2f}\\\\ \\includegraphics[scale=0.35]{figCompareModelWithUlysses_speed3s} & \t\\includegraphics[scale=0.35]{figCompareModelWithUlysses_dens3s}\\\\ \\includegraphics[scale=0.35]{figCompareModelWithUlysses_speed3f} & \t\\includegraphics[scale=0.35]{figCompareModelWithUlysses_dens3f}\\\\ \\includegraphics[scale=0.35]{figCompareModelWithUlysses_speed4s} & \t\\includegraphics[scale=0.35]{figCompareModelWithUlysses_dens4s}\\\\ \\end{tabular} \\caption{Comparison of \\textit{Ulysses} \\textit{in~situ} measurements for slow and fast scans with model results for solar wind speed and adjusted density. Presented are Carrington rotation averaged data. Blue - model, red - \\textit{Ulysses}.} \\label{figCompareModelWithUlysses} \\end{figure*} \\clearpage"
    },
    "1112/1112.4841_arXiv.txt": {
        "abstract": "Thermal instability (TI) can strongly affect the structure and dynamics of the interstellar medium (ISM) in the Milky Way and other disk galaxies. Thermal conduction plays an important role in the TI by stabilizing small scales and limiting the size of the smallest condensates.  In the magnetized ISM, however, heat is conducted anisotropically (primarily along magnetic field lines).  We investigate the effects of anisotropic thermal conduction on the nonlinear regime of the TI by performing two-dimensional magnetohydrodynamic simulations.  We present models with magnetic fields of different initial geometries and strengths, and compare them to hydrodynamic models with isotropic conduction. We find anisotropic conduction does not significantly alter the overall density and temperature statistics in the saturated state of the TI. However, it can strongly affect the shapes and sizes of cold clouds formed by the TI.  For example, for uniform initial fields long filaments of cold gas are produced that are reminiscent of some observed HI clouds. For initially tangled fields, such filaments are not produced.  We also show that anisotropic conduction suppresses turbulence generated by evaporative flows from the surfaces of cold blobs, which may have implications for mechanisms for driving turbulence in the ISM. ",
        "introduction": "\\label{sec:Intro}  The thermal instability (TI) plays an important role in controlling several different aspects of the interstellar medium (ISM) and star formation. For example, it has been invoked to explain the observed multiphase structure of the ISM \\citep{par53,spi58,fie65,fie69}. The linear stage of the TI in astrophysical gases was first studied in detail by \\citet{fie65}. He identified three unstable modes: an isobaric and two isentropic modes. In the nonlinear regime, the isobaric mode produces condensations that are fundamental to the classical two-phase model of the ISM \\citep{fie69}, as well as the extension to a three-phase model by \\citet{cox74,mck77}. The TI also regulates the mass flow between the different components of the ISM, and therefore affects the star formation rate \\citep{chi87}. Thus, it is important to investigate the role of TI in determining the distribution of density, temperature, and other physical variables in the multiphase ISM. For this reason, a variety of authors have studied the linear and nonlinear stages of the TI using numerical hydrodynamic simulations \\citep{hen99,kri02a,kri02b,kri04,vaz00,gaz01,san02,pio04,pio05,koy04,vaz07,kim08,ino08}.  Thermal conduction is important to include in studies of the TI for two reasons.  First, it suppresses the growth rate at small scales, in fact isobaric perturbations with wavelength smaller than the Field length $\\lambda_F$ \\citep{fie65} do not grow at all.  Second, it produces evaporation from the surfaces of cold dense fragments, and the interaction of the evaporative flows can induce turbulence.  Including thermal conduction is essential for numerical studies of the TI. Without explicit thermal conduction, perturbations at the grid scale grow fastest, and may eventually come to dominate the dynamics. For this reason, \\citet{koy04} concluded that numerical studies of the TI must satisfy a \"Field criterion\", that is the Field length must be resolved by at least a few cells to prevent artificial fragmentation at the grid scale, and to avoid the results being dominated by grid noise.  Satisfying the Field criterion requires including explicit thermal conduction \\citep{pio04,pio05,koy04,bra07}.  To highlight one example, \\citet{pio04,pio05} studied the interaction of the TI and the magnetorotational instability (MRI) in disks including isotropic thermal conduction; they found the MRI could drive turbulence and fragmentation in the diffuse ISM at amplitudes consistent with HI observations.  To date most studies of the TI with conduction have assumed the conductivity is isotropic. However, in a magnetized plasma, electrons can flow more freely along magnetic field lines than across them, leading to anisotropic transport coefficients \\citep{spi62}.  The degree of anisotropy is measured by the ratio of the electron gyro radius to the mean free path between collisions.  For the warm medium, where typically $T=1500$ K, $n=2$ cm$^3$, and $B=1 \\mu$G, the Columb mean free path is $\\lambda_{\\rm mfp} \\sim 10^{10}$ cm, while the electron gyro radius is $r_g=10^{6}$ cm.  Thus, in this medium the thermal conduction should be highly anisotropic, and primarily along magnetic field lines.  The implications of anisotropic transport terms on the dynamics of astrophysical plasmas has begun to be explored recently \\citep{bal01,sha03,par05}.  For example, in stratified atmospheres, anisotropic conduction can result in the magnetothermal and heat-flux buoyancy instabilities in the intracluster medium \\citep{par08,sha09}. In addition, it can have effects on the evolution of supernova remnants \\citep{bal08} and on magnetized spherical accretion flows \\citep{sha08}. Recently, \\citet{sha10} studied the TI in the intracluster medium, including heating by cosmic rays.  Their primary result is that with anisotropic thermal conduction, the TI could produce filaments of cold gas orientated along magnetic field lines.  In this study, we perform two-dimensional numerical hydrodynamic and magnetohydrodynamic simulations of TI in the ISM that include the anisotropic heat conduction.  The purpose of this paper is to investigate the effects of anisotropic thermal conduction on the structure and dynamics of interstellar medium, and to compare the results to models which include only isotropic conduction. Since the geometry and strength of magnetic field in different regions of the ISM may vary, we perform simulations of the TI with various magnetic field strengths and two different initial geometries for the field.  Ideally, the numerical studies of the thermal instability should include both isotropic and anisotropic conduction with the temperature dependency since the thermal instability develops density distribution with distinct peaks in the cold/dense phase dominated by isotropic conduction and in the diffuse/hot phase with higher ionization fraction and thus more affected by anisotropic conduction. However, including both isotropic and anisotropic conduction with a realistic temperature dependency is challenging, since it would decrease the Field length, and require higher resolution. In this study, we restrict our exploration to the simulations only with isotropic conduction or ones only with anisotropic conduction to better understand the effects of anisotropic conduction.  This paper is organized as follows: our numerical methods and code tests are summarized in \\S~\\ref{sec:Model}. In \\S~\\ref{sec:iso}, we first discuss the effects of the conduction rate and resolution on the TI in hydrodynamical simulations, and then use these results to choose a specific set of model parameters for our simulations. Results from calculations of the nonlinear regime of the TI with anisotropic conduction are presented in \\S~\\ref{sec:aniso}. We summarize and discuss our results in \\S~\\ref{sec:Summary}. ",
        "conclusions": "\\label{sec:Summary}  In this work, we have studied the nonlinear regime of the thermal instability (TI) and the effect of anisotropic thermal conduction by performing two-dimensional hydrodynamical and magnetohydrodynamic simulations incorporating radiative cooling and heating.  Our main results can be summarized as the following:  1.  As found in previous studies, it is crucial to include explicit thermal conduction in numerical studies of the TI so that the Field length is resolved, in order to prevent artificial fragmentation driven by numerical noise at the grid scale (e.g. Figure~\\ref{K_density}).  2. The amplitude of the thermal conductivity controls the rate of evaporation from the surfaces of dense clouds formed by the TI, and therefore strongly affects the amplitude of turbulent motions induced by the TI (e.g. Figure~\\ref{K_KEmean}).  3. Even when the Field length is resolved, explicit viscosity must be included to obtain numerical convergence of some quantities, for example the amplitude of turbulent motions driven by evaporative flows (e.g. Figure~\\ref{R_study_sigma}).  4.  Although the statistics of the density and temperature are not strongly affected by anisotropic conduction, the geometry of structures formed by the TI are quite different.   With anisotropic conduction and a uniform magnetic field, the TI saturates as long thin filaments of dense gas aligned with the field.  In a tangled field, spherical clouds are formed in regions of closed field lines (e.g. Figure~\\ref{fig_logd}).  5.  The combination of anisotropic conduction and MHD strongly suppresses the rate of evaporation of cold gas from the surfaces of dense structure in regions where the field is parallel to the interface. This reduces the amplitude of turbulence driven by the TI (e.g. Figure~\\ref{mhd_velo}).  These results have a number of implications for observations of cold neutral gas in the ISM.  In particular,  the thin filaments along the magnetic field in the weak uniform magnetic field case agree well with recent observations of the Riegel-Crutcher cloud conducted by \\citet{mcc06}. This neutral hydrogen (HI) cloud lies on the edge of the Local Bubble \\citep{cru84} filled with a hot and diffuse gas where the anisotropic conduction can be expected. \\citet{mcc06} found a network of dozens of hairlike filaments of cold hydrogen with widths of less than $\\sim 0.1$ pc and up to 17 pc long. They also have found that the filaments are aligned with the magnetic field of the cloud which agrees well with our results. They also calculated the magnetic field strength by using the Chandrasekhar-Fermi method, finding $\\sim 60 \\mu$G.  We find that compression and twisting of field during the formation of filaments by the TI with anisotropic conduction can amplify it by up to a factor of 100 in two-dimensions.  In three-dimensions, the amplification is likely to be larger. This may be enough to explain the observed field.  Recently, \\citet{sha10} have reported a study of the TI with anisotropic thermal conduction in the hot X-ray emitting plasma in clusters of galaxies. The heating and cooling processes in this regime are very different than those in the ISM studied here (equations \\ref{eq:E_heat} and \\ref{eq:E_cool} respectively), in particular there are no stable phases in the cooling curve they adopt, so that magnetic pressure sets the only limit on how cold and dense the gas becomes.  Nonetheless, the structure of density condensations in this case are very similar to our results: filaments of cold gas along magnetic field lines.  Our results show that in the case of anisotropic conduction, the geometry of the magnetic field with respect to the interface between cold and hot phases is very important.  Only if the field is normal to the interface can thermal conduction drive evaporation and outflows. The inclusion of anisotropic conduction could have important implications for the structure of interfaces, and the interpretation of observations of these regions \\citep{ino06,sto09,sto10}.  There are a number of limitations to our work that should be addressed in future investigations.  Firstly, the ISM is highly turbulent \\citep{hei03}, with typical turbulent velocities approximately 7 $\\rm km s^{-1}$ \\citep{hei03,moh04}, i.e. more than 100 times larger than those produced by the TI with isotropic conduction and viscosity. In the traditional picture of the ISM, supernova driven turbulence leads to a hot, diffuse third phase \\citep{cox74,mck77}.  More recently, \\citep{pio04} have considered the interaction of turbulence driven by the MRI and TI. We have not considered the effect of externally forced turbulence on the TI in this work, however this would be a productive direction for study in the future.  Secondly, in this work we have adopted a constant conductivity $\\mathcal{K}$.  However, in reality the conductivity is a function of the temperature \\citet{par53}, so that the rate of thermal conduction decreases as the gas cools. The amplitude of the conductivity can vary by 2 orders of magnitude between warm and cold phases at the end of our typical simulation at $t\\sim 100$~$\\rm Myr$.  As discussed in Section \\ref{sec:coneffect}, the value of the conductivity can affect the rate of evaporation from dense clouds, and therefore the kinetic energy of turbulence driven by the TI, thus assuming a constant value may alter the result.  Using  a realistic temperature dependent conductivity is challenging, since it would decrease the minimum value of Field length and require much higher resolution.  Since the geometry of the magnetic field limits the amount of evaporation to a very small surface area at the ends of the filaments in the case of anisotropic conduction, we do not expect the use of a more realistic conductivity to produce qualitative changes in our result.  Nonetheless, more realistic studies which use temperature dependent conductivities would be fruitful.  Finally, this study has considered flows in only two dimensions. In full three dimensions, the amplification of the magnetic field due to geometrical compression into filaments may be larger, and the turbulence driven by evaporative flows may be of a different character. Fully three-dimensional simulations of the TI with anisotropic conduction would also be interesting for future studies."
    },
    "1112/1112.1463_arXiv.txt": {
        "abstract": "A wide field of view Cherenkov/fluorescence telescope array is one of the main components of the Large High Altitude Air Shower Observatory project. To serve as Cherenkov and fluorescence detectors, a flexible and mobile design is adopted for easy reconfiguring of the telescope array. Two prototype telescopes have been constructed and successfully run at the site of the ARGO-YBJ experiment in Tibet. The features and performance of the telescopes are presented. ",
        "introduction": "The energy spectrum of primary cosmic rays spans almost 12 orders of magnitude, from $10^{9}$~eV to $10^{21}$~eV, and can be well fitted by a simple power law except in several small energy regions. A region called the ``knee\"  of the spectrum existing at around $10^{15}$~eV is one of these regions where the spectrum becomes steeper at higher energy side. Many experiments have observed this phenomenon; however, controversial arguments on its origin persist because of limited discrimination power on the primary cosmic ray composition and ambiguities in nucleus-nucleus interaction modeling. These two aspects are closely related to each other. Modern balloon borne experiments, such as ATIC~\\cite{ATIC} and CREAM~\\cite{CREAM}, have efficiently measured the energy spectra of individual elements at the top of the atmosphere. The energy spectra for all nuclei are measured up to $\\sim$100~TeV which is not far from the ``knee\". Because the detector area is constrained by the payload, the spectrum measurement has to be extended to a higher energy using a ground based air shower detector array. The spectrum should initially be measured well below 100~TeV to create an overlap with the balloon experiments which serve as absolute calibrations for the ground-based techniques. Identifying the individual components of cosmic rays continues to be a major challenge in ground-based experiments.  Multiple parameter measurements on an air shower seem to be a plausible approach.  The ultimate goal is to separate individual species out of the total observed-event samples and measure a clear individual ``knee\" for every single species, enabling the discovery of the origin of the ``knee\". As one of the major scientific goals of the Large High Altitude Air Shower Observatory (LHAASO) project~\\cite{LHAASO-caozh,LHAASO-hhh}, the energy spectrum for a separated composition will be measured at energies above dozens of ~TeV. To tag each primary particle that causes an air shower, the atmospheric depth of the shower maximum should be measured as one of the important parameters.  The wide field of view Cherenkov/fluorescence telescope array (WFCTA), one of the components of the LHAASO project, is designed to accomplish this goal.  A portable design of WFCTA telescopes is adopted to maximize the flexibility of changing the configuration of the array of telescopes.  The elevations, pointing directions, and locations of the telescopes are then easily reconfigured. This is one way of using the same telescopes to serve as both fluorescence and Cherenkov detectors.  In the fluorescence detector, the telescopes are tilted down to a horizontal position. In such an operational mode, which is analogous to the HiRes experiment~\\cite{Hires-detector}, most of the Cherenkov photons are avoided except those that are scattered onto the field of views (FOVs) of the telescope, such as in the fluorescence detector of the telescope array experiment~\\cite{TA-exp} and the fluorescence detector of the Pierre Auger experiment~\\cite{auger-fluorescence, auger-spectrum}.  Only the fluorescence light from the shower is collected together with the scattered Cherenkov light to trigger the telescopes. This requires showers having much higher energy, usually above 100 PeV, such as in the HiRes prototype experiment~\\cite{Hires-hybrid}, because the fluorescence light by a single electron is considerably weaker and  isotropic. In the Cherenkov detector, the telescopes run in high elevation mode to directly measure Cherenkov light from the showers, similar to what was done in the Dice experiment~\\cite{dice-exp}. A Cherenkov light radiation provides considerably more photons along the shower axis that are useful for lowering the shower energy.  In 2007, two prototype Cherenkov telescopes~\\cite{CRTNT-caozh,CRTNT-hhh} were deployed at Yangbaijing (YBJ) Cosmic Ray Observatory near the ARGO-YBJ experiment~\\cite{ARGO-detector}. Moreover, two WFCTA telescopes have been successfully running in Cherenkov mode beginning August 2008.  To date, millions of cosmic ray events that simultaneously trigger the telescopes and the ARGO-YBJ detector carpet array have been collected. An analysis of these events is carried out to study the performance of the telescopes.  Detailed descriptions of the telescopes and the analysis of the findings are presented in this paper.  Several details about the apparatus are presented in Section 2. The detector calibration is then discussed in Section 3. The test-run of the two telescopes and results are reported in Section 4 including summaries on the detector performance. The conclusions drawn are provided in the last section. ",
        "conclusions": "The telescopes were successfully run at YBJ from August 2008 up to July 2009. Millions of coincidence events with the ARGO-YBJ experiment have been collected. The performance of the telescopes was studied using these events. The trigger rate is about 0.5 Hz in stereo mode. Moreover, the mode energy of the telescope is 40 TeV when a pure proton composition is assumed.  The features of the two WFCTA prototype telescopes are summarized as follows:  \\begin{itemize} \\item a 4.7~$m^2$ spherical mirror, \\item a 16$\\times$16 PMT array covers an FOV of $14^{\\circ}\\times16^{\\circ}$ with $1^{\\circ}$ pixels, \\item dual gain system for a dynamic range to 3.5 orders of magnitude, \\item DC coupling and modulized design for electronics, \\item three-level online trigger logic: single channel trigger based on S/N ratio, telescope trigger based on pattern recognition, and event trigger for stereoscopic observation, \\item maximized mobility and the telescope can be uplifted from $0^{\\circ}$ to $60^{\\circ}$ in elevation. \\end{itemize}  The absolute gains of the telescopes are calibrated using calibrated LEDs mounted at the centers of the mirrors. The systematic uncertainty of the calibration constant is about 7\\%. The pixel gains are monitored on a daily basis."
    },
    "1112/1112.4070_arXiv.txt": {
        "abstract": "Radiative feedback is among the most important consequences of clustered star formation inside molecular clouds. At the onset of star formation, radiation from massive stars heats the surrounding gas, which suppresses the formation of many low-mass stars. When simulating pre-main-sequence stars, their stellar properties must be defined by a prestellar model. Different approaches to prestellar modeling may yield quantitatively different results. In this paper, we compare two existing prestellar models under identical initial conditions to gauge whether the choice of model has any significant effects on the final population of stars. The first model treats stellar radii and luminosities with a ZAMS model, while separately estimating the accretion luminosity by interpolating to published prestellar tracks. The second, more accurate prestellar model self-consistently evolves the radius and luminosity of each star under highly variable accretion conditions. Each is coupled to a raytracing-based radiative feedback code that also treats ionization. The impact of the self-consistent model is less ionizing radiation and less heating during the early stages of star formation. This may affect final mass distributions. We noted a peak stellar mass reduced by 8\\% from 47.3$\\Msun$ to 43.5$\\Msun$ in the evolutionary model, relative to the track-fit model. Also, the difference in mass between the two largest stars in each case is reduced from 14$\\Msun$ to 7.5$\\Msun$. The HII regions produced by these massive stars were also seen to flicker on timescales down to the limit imposed by our timestep ($<$ 560 years), rapidly changing in size and shape, confirming previous cluster simulations using ZAMS-based estimates for prestellar ionizing flux. ",
        "introduction": "The conversion of molecular gas into fully-formed stars is complex, involving several diverse processes. These different processes are linked to each other through feedback mechanisms that make isolating and understanding the contribution of each process a difficult task. A key point in this regards is that stars also rarely form in isolation, but instead are seen to be forming in clusters and subclusters within molecular clouds \\citep{Clarke+2000,Testi+2000}. In the cluster environment, the formation of a sufficiently massive star can affect all the others through the energy it radiates back into the cloud. Numerical simulations of star formation have made it very clear that the effects of stellar radiation cannot be neglected. Simulations including some form of radiative transfer show a dramatic reduction in the production of brown dwarfs and other low-mass stars \\citep{Offner2009}, due to an increase in gas temperatures reducing fragmentation \\citep{Krumholz+2007,Peters2010c}. More of the available gas mass ends up being accreted by the fewer, larger stars formed, and the fragmentation that does occur takes place in optically thick self-shielding discs \\citep{Krumholz+2007,Peters2010a,Peters2010c}. The fact that radiation affects the mass spectrum in simulations of molecular cloud clumps has obvious implications for the shape of the initial mass function, for example the suppression of excessive brown dwarf formation \\citep{Bate2009,Krumholz+2010,Peters2010c}.  Massive stars also emit prodigious amounts of UV radiation \\citep{Hoare+2007,Beuther+2007} creating expanding HII regions. The hot ($10^4$ K) gas expands into the colder ($10^2$ K) surrounding low-pressure gas, creating another feedback mechanism and ionized region that may contribute to the destruction of molecular clouds \\citep{Keto2002,Keto2003,Keto2007,Matzner2002,Peters2010a,Peters2010c}. HII regions can be observed by their radio continuum emission \\citep{MezgerHenderson1967}, or by their recombination lines (e.g.\\ \\citet{WoodChurchwell1989} use the H76$\\alpha$ line). More recently, observations have shown time variability in HII regions \\citep{Franco-HernandezRodriguez2004,Rodriguez+2007,Galvan-Madrid+2008,gomezetal08}. \\citet{Franco-HernandezRodriguez2004} have suggested that such observed time-variability may be due to the changes occurring in the source of the ionizing radiation, though it may also be due to increased absorption in the rapidly-evolving core of the nebula. \\citet{Peters2010b} present a technique for using synthetic radio maps to study the time-evolution of stars forming in a cluster environment and variability in the morphology and size of HII regions. Analysis of these simulations by \\citet{Peters2010a} and \\citet{Galvan-Madrid+2011} confirmed variability in the flux and size measurements of HII regions, which in a few cases might be observable on timescales of $\\sim 10$  years. They also noted that positive changes were more likely to occur than negative changes, i.e.\\ that most of the flux variations were increases rather than decreases.  To further explore the impact of radiative feedback and the possible variability in HII regions, simulations must be equipped with good protostellar models. These have been investigated by \\citet{PallaStahler1991}, \\citet{PallaStahler1992}, \\citet{Nakano2000}, \\citet{McKeeTan2003}, \\citet{Offner2009} and \\citet{HosokawaOmukai2009}, among others. It is clear from these models, that the evolution of a protostar depends heavily on the mass accretion rate. Among other things, they show that the radius of the protostar may grow or contract depending on the stellar evolutionary stage. With a radius that can change significantly during the pre-main-sequence lifetime of the star, the effective temperature can also be expected to vary significantly. To study this, we simulate the formation of a cluster of stars inside a molecular cloud. We equip the stars with one of two prestellar models based on the ones described in \\citet{Peters2010a} and \\citet{Offner2009}, each with its own characteristics. The \\citet{Offner2009} model has already been used to study star cluster formation in \\citet{Krumholz+2011}, though with different initial conditions. Ours is the first simulation with the protostellar model to also include the effects of ionizing radiation and HII region formation. We connect the model to a radiative transfer method that computes the heating and ionization due to radiation from the stars formed in the simulation.  The differences between the two models is explained in \\ref{sec:protmod}, but the key difference is that the \\citet{Offner2009} model treats the evolution of the radius and luminosity self-consistently. The choice of stellar model affects the early evolution of stars in a cluster, and may have repercusions for the final mass spectrum. Though not entirely conclusive, we find that reduced heating and ionization in the early stages of star formation when using the \\citet{Offner2009} model resulted in a more equitable mass distribution. With the \\citet{Peters2010a} model, the cluster came to be dominated more by a single star about $14\\Msun$ more massive than the next largest, compared to a $\\sim7.5\\Msun$ gap in the \\citet{Offner2009} model simulations.  Other effects of the self-consistent prestellar modeling are delayed ionization of the cluster gas by 3\\% of a freefall time (17.7 kyr), and delayed heating of the cluster gas by 1\\% of a freefall time (5.9 kyr).  Our numerical approach is described in Section \\ref{sec:numerical_methods}. In Section \\ref{sec:results} we list our results for the early evolution of star clusters with massive stars. Our assessment of the impact of protostellar modeling we discuss in Section \\ref{sec:discussion} and summarize our findings in Section \\ref{sec:conclusion} with a view to future simulations. ",
        "conclusions": "\\label{sec:conclusion}  Stars begin to affect their birth environments as soon as they are born through radiative feedback. We have considered the impact that pre-main-sequence modeling can have on a star cluster by comparing two different prestellar models already described in the literature. We did this by repeating the simulation of \\citet{Peters2010a}. We then upgraded the FLASH code to include a protostellar evolution module based on the one described in the appendices of \\citet{Offner2009}.  Each model works by equipping the stars in the simulation (``sink particles'') with a stellar radius and luminosity. The greatest difference between the two models was self-consistency. The \\citet{Peters2010a} model calculated approximate stellar parameters on-the-fly, while the evolving protostar model evolved the stellar parameters self-consistently through the simulation as the stars grew and accreted mass.  In terms of the overall gas structure, HII regions, temperature structure, mean ionization fraction, or stellar binarity, the two models produced qualitatively the same results. This is because a cluster comes to be dominated by its most massive stars, which are evolved, main-sequence, highly luminous stars, regardless of the choice of stellar model. These one or two massive stars control the overall dynamics.  The differences exist in the early phase of star formation. Major ionization of the gas in the evolving protostar model lagged the \\citet{Peters2010a} model by about 3\\% of a freefall time, or about 17.7 kyr. Major heating of the gas lagged by about 1\\% of a freefall time, or about 5.9 kyr. The difference in heating and ionization was due to the fact in \\citet{Peters2010a}, the stellar radius was underestimated (a ZAMS-equivalent value was taken), when protostars have radii an order of magnitude larger than a zero-age main-sequence star of equal mass. The correspondingly higher surface temperatures resulted in excess heating and ionization in this model. When both models had stars converging onto the main sequence, the differences between the two models diminished.  It is possible that these initial differences could have had a lasting effect on the stellar population. The most massive star at the end of each simulation was 43.5$\\Msun$ in the evolving protostar model, and 47.3$\\Msun$ in the \\citet{Peters2010a} case---a difference of 8\\%. The differences in mass between the most massive star and the next largest star was 7.5$\\Msun$ in the evolving protostar case and 14$\\Msun$ in the \\citet{Peters2010a} case. It would require further simulations, varying the initial conditions, to confirm that this is always the case.  The cluster of stars is embedded in a rotating disc of gas approximately 0.2 pc in size. The expanding HII regions above and below the disc are rapidly changing in shape and size on timescales shorter than 570 years. The physical size of these HII regions in our simulation is at most about 0.2 pc. This flickering is observed regardless of the prestellar model used.  Future simulations will have initial conditions including turbulence to model molecular clouds as realistically as possible. The stars will no longer be forming within a global disc, but rather along sheets and filaments in diverse parts of the cloud. With star formation thus spread out more in space and time, we expect the influence of individual young stellar objects on their environments to be more significant than when all stars form in a central cluster. It will be important to have the radiative feedback accurately modeled in these cases."
    },
    "1112/1112.1655_arXiv.txt": {
        "abstract": "We present the results of the optical, X-ray and gamma-ray analysis of some recent GRBs. The data were obtained by the automated P60 telescope and the Swift telescope (UVOT, XRT and BAT).  We present some example fits for the lightcurves.  The data reduction and the investigations were made by the Konkoly Observatory HEART group (http://www.konkoly.hu/HEART/index.html). ",
        "introduction": "The UVOT photometry was done using NASA's HEASOFT software package. This package provides a complete assistance for doing photometry for measurements done by various telescopes (e.g. Swift, CGRO, INTEGRAL).  For photometry, contrary to the suggestions of the software manual, we explored various aperture sizes between $1$\" and $10$\" (instead of keeping $5$\" at all times) to get the most usable data. Then we chose those aperture sizes, which provided the most accurate results.  For a given filter band we used the best aperture size, meaning that in some cases when obtaining photometry for a given GRB we used various apertures depending on the filter band.  All UVOT magnitudes are in the Standard UVOT Photometric System \\citep{uvot_phot}.  Magnitudes obtained by ground based telescopes were taken from GCNs (in detail see at references \\cite{gcn}).  When converting magnitudes  to fluxes, we used the methods described in \\cite{bessel,p60,sdss}.  Magnitudes are not corrected for galactic extinction. X-ray data is taken from the XRT observations available from \\cite{evans}. According to our results we suggest 19 magnitude as acceptable faintest limit in the UVOT photometry system.   We used the \\textit{mpfit} package \\cite{mpfit} for fitting broken power-law functions to the lightcurves with the $F_{\\nu} \\propto t^{\\alpha}$ convention. ",
        "conclusions": "\\ In this work we aimed  to produce photometric data with the highest possible accuracy.  Our sample consisted of relatively bright and well observed bursts in order to achieve a reliable data set with our reduction method.  We excluded those large error measurements which we had to disregard during our fitting procedure. According to our results we suggest 19 magnitude as the faintest limit when using the UVOT system.   \\begin{theacknowledgments} Thanks to Bob Wiegand for the help us to make HEASOFT working correctly. This work was partially supported by OTKA grant K077795 (P.V.). \\end{theacknowledgments}"
    },
    "1112/1112.0379_arXiv.txt": {
        "abstract": "\\himass\\ is one of the most \\HI-massive galaxies in the local ($z<0.1$) Universe. Not only are such galaxies extremely rare, but this ``coelacanth\" galaxy exhibits characteristics -- in particular its active, inside-out stellar disk-building -- that appear more typical of past ($z \\sim 1$) star formation, when large gas fractions were more common.  Unlike most local giant \\HI\\ galaxies, it is actively star forming. Moreover, the strong infrared emission is not induced by a merger event or AGN, as is commonly found in other local LIRGs. The galaxy is suggestive of a scaled-up version of local spiral galaxies; its extended star formation activity likely being fueled by its large gas reservoir and, as such, can aid our understanding of star formation in systems expected to dominate at higher redshifts. The multi-wavelength imaging and spectroscopic observations that have led to these deductions will be presented. These include NIR ($JHK$) and MIR (Spitzer; $3-24\\micron$) imaging and photometry, MIR spectroscopy, ATCA \\HI-interferometry and Mopra CO line emission observations. But no optical data, as the galaxy is heavily obscured due to its location in Vela behind the Milky Way. ",
        "introduction": "The galaxy \\himass, originally discovered in the Parkes deep Multibeam \\HI-survey of the Zone of Avoidance (ZOA; e.g. Kraan-Korteweg et al. 2005), is a rapidly rotating disk galaxy ($v_{\\rm rot}=630$\\kms; $D_{\\rm HI} = 130$kpc) with an \\HI-mass of $M_{\\rm HI} = 7.5 \\cdot 10^{10}\\msun$ and a dynamical mass of $M_{\\rm tot} = 1.4 \\cdot 10^{12}\\msun$ (Donley et al. 2006). It has a prominent bulge and smallish stellar disk ($B/D \\sim 0.8$) of evolved stars of $M^* = 4.4 \\cdot 10^{10}\\msun$ embedded in the five times larger HI-disk, while its current hearty star forming activity makes it a luminous infrared galaxy (LIRG; next section).  \\himass\\ is the most \\HI-rich galaxy known. Such galaxies are extremely rare ($\\sim 3 \\cdot 10^{-8}$/Mpc$^3$) according to the current best determined local galaxy HI-mass function (HIMF; Zwaan et al. 2005) and should not have been found at all in the explored volume. Interestingly, the recent, larger volume ALFALFA survey 40\\%-data release (Haynes et al. 2011) identify several, similarly extreme \\HI-massive galaxies in excess of the predicted HIMF number density. Independently, galaxies with such vast reservoirs of gas were more common in the past, as were LIRGs. Hence this enigmatic, relatively nearby ($v=10\\,689$\\kms) galaxy is an ideal {\\sl local} probe that enables detailed studies of evolutionary processes and the transformation of gas into stars. ",
        "conclusions": ""
    },
    "1112/1112.5240_arXiv.txt": {
        "abstract": "We discuss a possible physical connection between   helium-rich ($Y \\ge 0.35$) stellar populations of massive globular clusters (GCs) and the ultraviolet (UV) upturn of galactic spheroids by using analytical and numerical models.  In our model, all stars are initially formed as bound or unbound star clusters (SCs) formed from giant molecular clouds (GMCs) and the  SCs can finally  become GCs, open clusters, and field stars depending on physical properties of their host GMCs. An essential ingredient of the model is that helium-rich stars are formed almost purely from gas ejected from  massive asymptotic giant branch (AGB) stars. The helium-rich star formation is assumed to occur within massive SCs if the masses of the progenitor  GMCs are larger than a threshold mass ($M_{\\rm thres}$). These  massive SCs can finally become  either massive GCs or helium-rich field stars depending on whether they are disintegrated or not. Using this model, we show that if the initial mass functions (IMFs) in galactic spheroids are mildly top-heavy, then the mass fractions  of helium-rich main-sequence stars ($F_{\\rm He}$) can be as large as $\\sim 0.1$ for $M_{\\rm thres}=10^7 {\\rm M}_{\\odot}$. $F_{\\rm He}$ is found to depend on IMFs and $M_{\\rm thres}$ such that it can be larger for shallower IMFs and smaller $M_{\\rm thres}$. The inner regions of galactic spheroids show larger $F_{\\rm He}$ in almost all  models. Based on these results, we suggest that if the UV upturn of elliptical galaxies is due to the larger fractions of helium-rich stars, then the origin can be closely associated with top-heavy IMFs in the galaxies. ",
        "introduction": "Recent observational and theoretical studies of the Galactic GCs have suggested that some of the massive GCs ($\\omega$ Cen and NGC 2808) have significant fractions of helium-rich stars (e.g., Bedin et al. 2004; Norris 2004; Lee et al. 2005a; Piotto et al. 2005, 2007; Renzini 2008). The observed unusual level of helium enhancement ($\\Delta Y/\\Delta Z \\approx 70$; Piotto et al. 2005) and sizable fractions of helium-rich stars in these GCs have driven many theoretical studies to make great efforts in understanding where and how helium-rich stars can be formed (e.g., Bekki \\& Norris 2006; D'Antona \\& Ventura  2007; D'Ercole et al. 2008). As discussed by Renzini (2008),  helium-rich stars can be formed from gaseous ejecta from massive AGB stars with initial masses ($m_{\\rm I}$) ranging from $3{\\rm M}_{\\odot}$ to $8{\\rm M}_{\\odot}$, though the ejecta should not be diluted by interstellar gas (ISM) to keep the original high helium abundances of the ejecta.  Code \\& Welch (1979) investigated the integrated light of seven elliptical and S0 galaxies for a wavelength range of  1550-4250 \\AA \\, and found that some of them show an increase in energy at the shortest wavelengths. Since this discovery of the UV upturn, the origin of the UV upturn has been extensively discussed both in observational and theoretical studies (e.g., Bertola et al. 1980; Burstein et al. 1988; Greggio \\& Renzini 1990; Horch et al. 1992; Dorman et al. 1995; Brown et al. 1997; O'Connell 1999; Yi et al. 1997, 1998, 2005). Although there can be a number of stellar candidates responsible for the UV upturn in galaxies, one of the most promising ones  is old  horizontal-branch stars (e.g., Yi 2008 for a recent review). It is  suggested that enhanced helium abundances can play an important role in the formation of hot stars responsible for the UV upturn (Yi 2008).  Brown et al. (2003) investigated the UV emission in eight early-type galaxies at $z=0.33$ (a look-back time of 3.9 Gyr) and found that the UV emission in these galaxies is significantly weaker than it is in the current epoch. Observational studies by {\\it Galaxy Evolution Explorer (GALEX)} investigated the UV properties  of bright cluster galaxies (BCGs) in clusters at $z<0.2$ and compared them with those of nearby giant elliptical galaxies (e.g., Lee et al. 2005b; Ree et al. 2007). Ree et al. (2007)  concluded that the observed evolution of FUV $-$ $V$ color with a model in which the dominant FUV source is hot horizontal-branch stars (e.g., Ree et al. 2007). Yi et al. (2011) investigated correlations between the strength of the UV upturn and global galactic  properties (e.g., luminosity) in BCGs and did not find any remarkable correlations (see also Loubser \\& S\\'anchez-Bl\\'azquez 2011 for  similar results). Yi et al. (2011) therefore concluded that the helium sedimentation scenario proposed by Peng \\& Nagai (2009) can not be supported by their  observational results.    These observational and theoretical studies appear to suggest that there can be helium-rich stellar populations in diverse objects with dramatically different masses and sizes ranging from GCs to BCGs. Although it would be likely that different astrophysical objects with helium-rich stars have different origins, it would be  possible that they can have a common origin. Previous theoretical studies, however, have not yet provided  an unified picture for the origins of helium-rich stars in GCs, galactic bulges, and elliptical galaxies with the UV-upturn. It is therefore worthwhile to construct a  model to discuss the origins of helium-rich stars observed in GCs and galactic spheroids in a self-consistent manner.      The purpose of this paper is to present a new  model which can provide a possible explanation for the origins of helium-rich stars with $Y \\ge 0.35$ in GCs and galactic spheroids in a self-consistent manner. Based on the model, we mainly discuss what types of IMFs are required to explain the observed possible fraction of helium-rich stars in galactic spheroids with the UV upturn (e.g., Chung et al. 2011). The most important ingredient of the  model is that helium-rich stars are formed almost purely from gas ejected from massive AGB stars with $3 {\\rm M}_{\\odot} \\le m_{\\rm I} \\le 8 {\\rm M}_{\\odot}$ with no/little dilution of the gas with ISM. This formation of new stars from AGB ejecta have been included also in recent theoretical models of GC formation (e.g., D'Ercole et al. 2008).   In the present study, we assume that such formation of helium-rich stars almost directly from AGB ejecta can be possible only within SCs formed from massive GMCs. By considering that the vast majority of stars are observed to be initially formed as bound or unbound SCs (e.g., Lada \\& Lada 2003), we assume that all stars form as SCs from GMCs in the model. Some massive SCs can retain gaseous ejecta from AGB stars and consequently have new helium-rich stars formed form the ejecta, and the helium-rich stars can become field stars of their host galaxies  when the SCs become disintegrated. Such massive SCs can finally become GCs or nuclear SC (or stellar galactic nuclei) depending on their birth places, if they are not destroyed by their host galaxies. The IMFs and GMCMFs are key parameter that can determine mass fractions of helium-rich stars in galaxies. By using this new model, we discuss the origin of physical properties of the Galactic GCs and the UV upturn in galactic bulges and elliptical galaxies.  Although it remains less clear how much fractions of stars should be helium-rich stars in galactic spheroids with  the UV upturn (Yi 2008), Chung et al (2011) have recently suggested that if about 11\\% of stellar populations in galactic spheroids are helium-rich stars, then they can show the UV upturn for a given IMF. In their models, they considered that the major source of far-UV flux originates from metal-poor and helium-enhances hot horizontal branch stars. We therefore consider that if the fractions of helium-rich stars in galactic spheroids are  less than  $\\sim 0.1$, the spheroids are unlikely to show the UV upturn in the present models. We investigate in what physical conditions (e.g., IMFs) the fractions of helium-rich stars can be  above 0.1 for galactic spheroids to show clearly the UV upturn.   A possible physical link between helium-rich stellar populations and the UV upturn of galactic spheroids has been already pointed out by a number of authors (e.g., Yi 2008). Thus the main point of the present study  is  {\\it not} to propose the importance of helium-rich stellar populations in the origin of the UV upturn {\\it but to discuss how galactic spheroids can have significant fractions  of helium-rich stellar populations.} It should be also stressed that the present model is idealized and less realistic in some points (e.g., no chemical evolution of galaxies): the model should be regarded as a first step toward better understanding the origin of helium-rich stars in GCs and galaxies. More sophisticated numerical simulations including SC formation processes in galaxies will need to be done in our future studies to address the UV upturn problem in a much more quantitative way.     The plan of the paper is as follows: in the next section, we describe a  model which enables us to estimate (i) the mass fraction of helium-rich stars in a single GC and (ii) the mass fraction of helium-rich main-sequence (MS) stars in a galaxy for a given IMF and GMCMF. In \\S 3, we present the results on the number/mass fraction of massive GCs with helium-rich stars in the Galaxy and the mass fractions of helium-rich MS stars in galaxies with different IMFs and GMCMFs In \\S 4, we provide important implications of the present results in terms of the origin of helium-rich stars in GCs, the Galactic bulge, and elliptical galaxies. We also discuss the origin of the radial gradients of helium-rich stars in galactic spheroids in this section. We summarize our  conclusions in \\S 5. ",
        "conclusions": "We have adopted a model in which (i) all stars are formed as bound or unbound SCs from GMCs and (ii) HRSs with $Y\\ge 0.35$ can be formed from gas ejected from AGB stars only in  massive SCs with $m_{\\rm gmc} \\ge M_{\\rm thres}$. Some SCs can become field stars in their host galaxies after disintegration and others can become massive GCs or nuclear star clusters (or stellar nuclei) in galaxies. If massive SCs with HRSs are destroyed by their host galaxies, then HRSs become helium-rich field stars in the galaxies. We have investigated mass fractions of MS HRSs ($f_{\\rm He}$ and $F_{\\rm He}$) in GCs and galactic spheroids and their dependences on IMFs, $M_{\\rm thres}$, and GMCMFs based on the  CDS (``cluster disintegration scenario''). We have also investigated the radial gradients of $F_{\\rm He}$ in galaxies with mass-size relations of GMCs by using N-body simulations on dynamical evolution of multiple cluster systems. We summarize our principle results as follows. \\\\    (1) The mass ratios  ($f_{\\rm m, agb}$) of gaseous ejecta from  massive AGB stars with $3{\\rm M}_{\\odot} \\le m_{\\rm I} \\le 8 {\\rm M}_{\\odot}$ to initial total masses of SCs are typically $\\sim 0.08$ for a canonical IMF (${\\alpha}_1 = 2.35$) with the reasonable lower and upper cut-off masses. If these gas can be efficiently converted into new stars with enhanced helium abundances (i.e., HRSs) in SCs, then the mass fractions of MS HRSs  among all MS stars ($f_{\\rm He}$) are $\\sim 0.08$  for canonical IMFs and a SC age of 12 Gyr. Only if original SCs have top-heavy IMFs (e.g., ${\\alpha}_1 \\sim 1.5$), then $f_{\\rm He}$ can be larger than  0.3 for $\\alpha_2 \\ge  1.85$. \\\\    (2) The initial number and mass fractions of massive GCs (MGCs) with  HRSs ($F_{\\rm n, mgc}$ and $F_{\\rm m, ngc}$, respectively) among the Galactic GCs formed from GMCs with $m_{\\rm gmc} \\ge 10^6 {\\rm M}_{\\odot}$ is $\\sim 0.16$ and $0.66$, respectively,  for $M_{\\rm th}=10^7 {\\rm M}_{\\odot}$ and $\\beta =1.7$. Both $F_{\\rm n, mgc}$ and $F_{\\rm m, mgc}$ are lager for smaller $M_{\\rm th}$ and smaller $\\beta$. MGCs with $m_{\\rm gc}= 5 \\times 10^6 {\\rm M}_{\\odot}$ can sink into the Galactic center owing to dynamical friction to disappear from the halo region within 13 Gyr, if the initial radii from the Galactic center are less than 3.4 kpc. About 37\\% (49\\%)  of the Galactic GCs with $m_{\\rm gc} \\ge 5 \\times 10^6 {\\rm M}_{\\odot}$ ($m_{\\rm gc} \\ge  10^7 {\\rm M}_{\\odot}$) have already been lost from the halo owing to efficient dynamical friction of these massive GCs. Thus, the Galactic halo initially could have a larger fraction of  MGCs. \\\\     (3) The mass fractions of HRSs among all stars ($F_{\\rm He, t}$) in galactic spheroids are $\\sim 0.05$ for $M_{\\rm thres}=10^7 {\\rm M}_{\\odot}$,   $\\beta =1.7$, $f_{\\rm m, agb}=0.1$, and ${\\epsilon}_{\\rm bsc}=1$.  $F_{\\rm He, t}$ of galactic spheroids are  larger for smaller $M_{\\rm th}$ and smaller $\\beta$ for a given IMF. The mass fractions of MS HRSs among all MS stars ($F_{\\rm He}$) in galactic spheroids are 0.06 for canonical IMFs ($\\alpha_1=\\alpha_2=2.35$), $M_{\\rm th}=10^7 {\\rm M}_{\\odot}$, $\\beta=1.7$, and a galaxy age of $\\sim 12$ Gyr. $F_{\\rm He}$ in galactic spheroids can be  larger for smaller $M_{\\rm th}$ and smaller $\\beta$ for a given IMF,  and it can be significant ($>0.1$), if original SCs (i.e., building blocks of the spheroids) have top-heavy IMF (e.g. $\\alpha_1 <  2$). These results suggest that if the observed UV upturn in bright elliptical galaxies  is due to the larger fractions ($\\sim 10$\\%) of HRSs  within them, the IMFs need to be top-heavy. If ${\\epsilon}_{\\rm bsc}$ is as small as $\\sim 0.1$, then bright elliptical galaxies are unlikely to show the UV upturn owing to the small fractions of HRSs. \\\\  (4) The Galactic bulge can have a larger $F_{\\rm He}$ if it had a top-heavy IMF at its formation in the CDS. Given a number of suggestions on the top-heavy IMF of the bulge by chemical evolution models of the bulge, the observed possible presence of helium-rich stars with $Y \\sim 0.35$ in the bulge (e.g., Nataf et al. 2011) can be understood in the context of  the CDS. HRSs in  MGCs that were initially located in the Galactic halo and had  sunk into the bulge can be currently observed as field HRSs  in the bulge, though the contribution of such stars to the entire helium-rich population is rather minor. It is doubtlessly worthwhile for observational studies to investigate whether or not the bulge has a larger fraction of HRSs in the inner regions. \\\\  (5) If $M_{\\rm th}$ and $\\beta$ are constant in elliptical galaxies with different physical properties, then elliptical galaxies with shallower (or more top-heavy) IMFs can show larger fraction of HRSs. Therefore, if more massive/luminous elliptical galaxies have shallower IMFs, then they can have larger fraction of HRSs. More massive elliptical galaxies can retain more efficiently the ejecta of supernovae so that they can finally have higher metallicities of their stars (e.g., Arimoto \\& Yoshii 1987). Therefore, if the Burstein relation (a correlation between UV flux and Mg index; Burstein et al. 1988) in bright elliptical galaxies is due to the correlation between $F_{\\rm He}$  and Mg indices, then the relation  can be understood in terms of shallower IMF slopes ($\\alpha_1$) in more massive/luminous elliptical galaxies. \\\\   (6) Galactic spheroids  can show negative radial gradients of $F_{\\rm He,t}$ and $F_{\\rm He}$  (i.e., higher in inner parts) for a reasonable set of model parameters. This is mainly because massive SCs with HRSs can sink rapidly into the inner regions of galaxies owing to dynamical friction and disintegrate there. The HRSs initially in the  SCs can be dispersed to become field HRSs there after disintegration of the host SCs. Therefore they can show the stronger UV upturn in their inner regions, if $F_{\\rm He}$ can determine the strengths of the UV upturn. The observed spatial distribution of the FUV excess in elliptical galaxies can be understood in the context of formation and evolution of MGCs with HRSs. The central regions of galactic spheroids can be composed of two different stellar populations with normal helium abundances and enhanced ones (i.e., HNSs and HRSs).\\\\       Thus, an advantage of the CDS is that both  (i) the absence or presence of the UV upturn in galactic spheroids  and (ii) correlations between the strength of the UV upturn  and galactic properties (e.g., Mg$_2$ index) can be discussed in the context of IMF slopes of galaxies in a self-consistent manner.  A disadvantage  of the CDS is that it is unclear whether and in what physical conditions ${\\epsilon}_{\\rm bsc}$ can be $\\sim 1$ for massive GMCs in galaxies. The formation of HRSs is possible, if helium-rich gas from AGB stars is converted into new stars without dilution of the gas by  ISM with normal helium abundances. In the CDS,  the inner regions of massive SCs are assumed to be the production sites of HRSs. It would be possible that AGB ejecta can be converted into new cold gas clouds outside SCs and consequently into new stars in deep potential wells of galactic spheroids, if there is no/little cold gas with normal helium abundances there. We need to investigate where and how AGB ejecta from field stars and SCs can be converted into new stars in galaxies using a more sophisticated chemodynamical  simulation in our future studies."
    },
    "1112/1112.2710_arXiv.txt": {
        "abstract": "Pristine stars with masses between $\\sim$140 and 260 $M_{\\odot}$ are theoretically predicted to die as pair-instability supernovae.  These very massive progenitors could come from Pop III stars in the early universe.  We model the light curves and spectra of pair-instability supernovae over a range of masses and envelope structures.  At redshifts of reionization $z\\geq 6$, we calculate the rates and detectability of pair-instability and core collapse supernovae, and show that with the \\emph{James Webb Space Telescope}, it is possible to determine the contribution of Pop III and Pop II stars toward reionization by constraining the stellar initial mass function at that epoch using these supernovae.  We also find the rates of Type Ia supernovae, and show that they are not rare during reionization, and can be used to probe the mass function at 4-8 $M_{\\odot}$.  If the budget of ionizing photons was dominated by contributions from top-heavy Pop III stars, we predict that the bright end of the galaxy luminosity function will be contaminated by pair-instability supernovae. ",
        "introduction": "The life of a massive star ends in a supernova (SN).  The detection of neutrinos from SN 1987A verified the idea that some SNe are set off by the gravitational collapse of the iron core of their progenitor star \\citep{Krauss1987}.  However, theory predicts that very massive stars with helium cores between $\\sim$64 and 133 $M_{\\odot}$ could find another way to blow up, through the thermonuclear explosion of oxygen via the pair-production instability \\citep{Rakavy1967, Barkat1967, Heger2002}.  The production of electron/positron pairs in the core softens the equation of state, leading to collapse and the ignition of explosive oxygen burning.  The subsequent thermonuclear runaway reverses the collapse and ejects the entire star, leaving no remnant behind.  Unlike iron core-collapse supernovae (CCSNe), which involves poorly constrained physical processes such as turbulence, pulsations, perhaps rotation and magnetic fields, the physics involved in pair-instability supernovae (PISNe) is fairly well understood and can be modeled with fewer uncertainties \\citep{Langer2009}.  Due to the extremely large stellar mass required, the progenitors of PISNe are expected to be rare, and may only form under unusual conditions.  One such condition existed in the early universe, when metal-free Population III stars were born \\citep{Loeb2010}.  In star formation, it is the accretion process that ultimately sets the final mass of a star.  From dimensional arguments, the mass growth rate is simply given by the Jeans mass $M_{J} \\sim c_{s}^3 G^{-\\frac{3}{2}} \\rho^{-\\frac{1}{2}}$ over the free-fall time $t_{ff} \\sim 1/\\sqrt{G\\rho}$, implying $dM/dt \\propto c_{s}^3/G \\propto T^{\\frac{3}{2}}$, where the sound speed $c_{s} \\sim \\sqrt{kT/m_{p}}$. In present day star-forming regions, heavy elements radiatively cool the gas to a temperature as low as $T\\sim 10$K.  However, in primordial clouds, the primary coolant at low temperatures is molecular hydrogen, which can only cool the gas to $T\\sim 200-300$K, implying an accretion rate higher than present day by two orders of magnitude.  Hence, theoretical studies suggest that the initial mass function (IMF) of Pop III stars might have been biased toward masses much higher than today, e.g. several hundred $M_{\\odot}$ \\citep{Bromm2004}.  The nucleosynthesis imprints of this top heavy IMF have been seen in globular clusters and damped Lyman alpha systems \\citep{Cooke2011, Puzia2006}.  Moreover, massive stars have strong winds driven by radiation pressure through spectral lines, with a mass loss rate scaling with stellar metallicity $\\dot{M}\\propto Z^{0.5\\sim 0.7}$ \\citep{Vink2001, Kudritzki2002}.  Most PISNe should therefore be from Pop III stars, which have weak radiation-driven winds due to their extremely low metallicities, and retain enough of their initial masses at the end of their lives to undergo a pair-instability explosion.  Naturally, studies of the rates and detectability of PISNe focused on high redshifts before reionization.  \\citet{Mackey2003} found the PISNe rate to be $\\sim 50$ deg$^{-2}$ yr$^{-1}$ at $z > 15$, while \\citet{Weinmann2005}, using more conservative assumptions for the number of PISNe produced per unit Pop III stellar mass formed, found the PISNe rate to be $\\sim 4$ deg$^{-2}$ yr$^{-1}$ at similar redshifts.  Assuming that only one supermassive Pop III star forms in unenriched minihalos, and that none form in protogalaxies, \\citet{Wise2005} found the PISNe rate be $\\sim 0.34$ deg$^{-2}$ yr$^{-1}$ at $z\\sim 20$.  After our paper was submitted, \\citet{Hummel2011} presented a complimentary analysis of the source density of PISNe from pristine minihalos, and determined the observability of such events with the \\emph{James Webb Space Telescope} (JWST), finding approximately $\\sim 0.4$ PISNe visible per JWST field of view at any given time.  PISNe after the epoch of reionization were also considered; \\citet{Scannapieco2005a} calculated a suite of PISNe model light curves with blackbody spectra, and analyzed the detectability and rates of PISNe from Pop III stars formed from leftover pristine gas at $z \\lesssim 6$.  As for CCSNe during reionization, \\citet{Mesinger2006} presented detailed predictions for the number of core collapse SNe that JWST could observe as a function of different survey parameters.  In this paper, we present light curves and spectral time series for PISNe from our multi-wavelength radiation-hydrodynamics simulations. As the stellar population responsible for reionization is currently unknown, instead of predicting a fixed SNe rate, we normalize the star formation rate by requiring that enough ionizing photons must be produced by either Pop III or Pop II stars in protogalaxies to complete reionization by $z \\sim 6$, and calculate the rates of pair-instability, core-collapse, and Type Ia SNe and their detectability with JWST for these two scenarios; the actual SNe rates will be in between these limiting cases.  We show that using the observed rates of these SNe, it is possible to distinguish the contribution of Pop III and Pop II stars toward reionization by characterizing the IMF at that time. ",
        "conclusions": "We analyzed simulated light curves and spectra of pair-instability supernovae for a variety of progenitor masses and envelope types, and found that the supernovae from the more massive progenitors are super-luminous and have extended light curves,  traits that would help  photometrically distinguish pair-instability supernovae from other types of supernovae using repeated snapshots.  We calculated the rates and detectability of pair-instability, core collapse, and Type Ia supernovae during the redshifts of reionization, and showed that it is possible to constrain the initial mass function of stars at that time, and identify the stellar population responsible for reionization.  If Pop III stars made the dominant contribution of ionizing photons during reionization, the bright end of the galaxy luminosity function will be contaminated by pair-instability supernovae.  \\bigskip"
    },
    "1112/1112.0009_arXiv.txt": {
        "abstract": "We carry out adaptive mesh refinement (AMR) cosmological simulations of Milky-Way mass halos in order to investigate the formation of disk-like galaxies in a $\\Lambda$-dominated Cold Dark Matter model.  We evolve a suite of five halos to $z=0$ and find gaseous-disk formation in all; however, in agreement with previous SPH simulations (that did not include a subgrid feedback model), the rotation curves of all halos are centrally peaked due to a massive spheroidal component.  Our standard model includes radiative cooling and star formation, but no feedback.  We further investigate this angular momentum problem by systematically modifying various simulation parameters including: (i) spatial resolution, ranging from 1700 to 212 pc; (ii) an additional pressure component to ensure that the Jeans length is always resolved; (iii) low star formation efficiency, going down to 0.1\\%; (iv) fixed physical resolution as opposed to comoving resolution; (v) a supernova feedback model which injects thermal energy to the local cell; and (vi) a subgrid feedback model which suppresses cooling in the immediate vicinity of a star formation event.  Of all of these, we find that only the last (cooling suppression) has any impact on the massive spheroidal component.  In particular, a simulation with cooling suppression and feedback results in a rotation curve that, while still peaked, is considerably reduced from our standard runs. ",
        "introduction": "\\label{intro}  In cosmologies dominated by Cold Dark Matter (CDM), galaxy rotation is produced by gravitational tidal torques arising from the hierarchical collapse of structure.   Analytic models and N-body simulations have shown that this can produce enough angular momentum to explain the observed sizes of disk galaxies \\citep{Fall:1980p1033, Mo:1998p1034}.  However, computational models including gas dynamics have struggled to reproduce realistic disk galaxies.  Such models initially produced undersized disks with low angular momentum \\citep{Navarro:1991p1002, Navarro:1994p1004}.  Later work did generate disks with approximately the correct extent \\citep{Steinmetz:2002p1035, Abadi:2003p641}, but these had oversized stellar spheroidal components and therefore unnaturally large core circular velocities \\citep{Steinmetz:1999p635} far in excess of those found observationally.  This \\emph{angular momentum problem} remains one of the major shortcomings of the CDM paradigm.  The origin of the angular momentum problem is not entirely understood, but it probably stems from the fact that such simulations do not achieve sufficient spatial and mass resolution to correctly model the appropriate physical processes.  For example, insufficient spatial resolution leads to the spurious mixture of hot and cold components of the ISM, producing an artificial warm component which is very efficient at radiating away energy \\citep{Katz:1992p1037, Katz:1996p1031, Steinmetz:1995p1005}.  Thus, underresolved gas in such simulations can cool very quickly, and in cooling it loses its pressure support and collapses into dense knots of material.  These knots interact with the galactic dark matter component through dynamical friction processes, and much of the angular momentum of the gas knots is transferred to the dark matter halo of a galaxy \\citep{Donghia:2006p1032}.  Consequently, the knots tumble into the center of the galaxy to produce a dense cusp of material in the core of the simulated galaxy.  Typically the buildup of these massive cores with cold gas and stars occurs rapidly, and even nascent simulated galaxies exhibit evidence of these cusps as early as $z \\sim 4$-5.  Energetic feedback from star formation events has the ability, in theory, to alleviate this problem either by heating and ``puffing\" up the collapsing knots so they are more easily disrupted before losing their angular momentum \\citep{Weil:1998p1030}, and/or by preferentially ejecting low angular momentum gas \\citep[e.g.,][]{2010arXiv1010.1004B}.  However, existing simulations are unable to resolve the detailed structures of star-forming regions.  Individual star-forming and stellar feedback events occur on parsec scales.  Stellar feedback processes can be resolved in simulations confined to local regions of the ISM \\citep[e.g.,][]{Joung:2006p1038} but not in cosmological simulations, which need to accurately co-evolve a galaxy's environment on scales of $ > 10$ Mpc.  Therefore, cosmological simulations simplify and parameterize star formation and stellar feedback on scales more easily met by current computational resources (i.e. scales of $10^2 - 10^3$ pc).  A variety of techniques have been suggested in order to achieve this puffiness within the confines of low-resolution models.  In the most primitive prescription, stellar feedback is simply the return of energy from newly created star particles to their surrounding gas, usually in the form of thermal energy.  Star ``particles'' in these simulations typically represent clusters of stars of mass $10^{2} - 10^{5}$ \\msun, so the thermal feedback is justified as the energy output from the most massive stars in the cluster becoming type II supernovae shortly after the creation of the star particle.  In nature, this supernovae heating produces a small component of very hot gas surrounding the stellar population, so hot that its cooling time is very long.  Unfortunately at the resolution currently achievable in cosmological simulations, this primitive thermal prescription for feedback deposits the same SN energy into a much larger reservoir of gas, which does not reach the same high temperature as it should.  This now-warm component of gas can easily radiate the excess energy away, cool further and proceed with star-formation runaway thus defeating the purpose of the feedback \\citep[e.g.,][]{Steinmetz:1999p635}.  Building from these failures, a number of research groups artificially turn off cooling in a gas parcel for a period of time ($t \\sim 10^7$ yr) after a cluster of stars has formed out of it \\citep[e.g.,][]{Gerritsen:1997p1039, Thacker:2000p1040, SommerLarsen:2003p1116, Stinson:2006p1023, Governato:2007p1022, Agertz:2010p461, Colin:2010p1053, Piontek:2011p1041, Guedes:2011p1080}.  This method is justified as an application of the Sedov-Taylor blast wave solution for a Type II SN \\citep{Taylor:1950p1075, Sedov:1959p1074}, which blows out any cold media from the immediate environment of a star formation event.  Using this prescription, any gas in a galaxy which starts to collapse into knots will reach the star formation criteria, form a star, and then heat up without any allowed cooling, thus preventing further collapse.  Not surprisingly, these research groups have found some success with this method, yielding simulated galaxies with reduced inner rotation curves due to less massive bulge components; however, gas parcel masses and sizes in cosmological simulations of this sort are typically too large for the Sedov-Taylor solution to apply (see Section \\ref{code_coolingsuppression}).  Thus despite the successes of the cooling suppression feedback model, the community continues to search for other more physically-motivated solutions.  Another subgrid model for feedback (i.e. on scales smaller than the true resolution of the simulation) is to inject kinetic energy directly into the gas; this can alleviate the problem of thermal energy being radiated away.  For example, some studies \\citep[e.g.][]{Springel:2003p1044, Scannapieco:2006p1118, Oppenheimer:2008bu} using Smoothed Particle Hydrodynamics (SPH) give some of the SN energy to individual gas particles in the form of momentum.  This method can result in significant mass outflows (by design), but at the cost of decoupling wind particles from hydrodynamic interaction for a period of time.  An alternate approach, to keep wind particles coupled to the disk gas was explored by \\citet{Schaye:2008p1045}.  Both approaches help but, by themselves, do not appear to generate realistic rotation curves.  In addition, \\citet{Truelove:1997p1046} showed that insufficient resolution in a simulation can lead to artificial fragmentation of the gas, perhaps resulting in a further overproduction of stars.  One way to prevent artificial fragmentation is to add additional (numerical) pressure in high-density, low-temperature regions to ensure that the Jeans length is always resolved \\citep{Machacek:2001p1047, Robertson:2008p1017}.  This can be achieved by modifying the equation of state (EOS) itself, making it stiffer in order to provide an additional source of pressure to gas in denser regions \\citep{Schaye:2008p1045, Ceverino:2009p1014, Agertz:2010p461}.  A polytropic EOS ($P \\propto \\rho^{\\Gamma}$) with $\\Gamma = 4/3$ will keep the ratio of Jeans length to resolution length constant (assuming Lagrangian resolution such that the resolution length decreases as $\\rho^{-1/3}$ -- for fixed resolution $\\Gamma = 2$ is required), but even stiffer relations have been used.  For example, \\citet{Agertz:2010p461} ran simulations with such an equation of state, where in low-density regions it behaved as an ideal gas, but in high-density (star forming) regions it followed a polytropic equation of state with $\\Gamma = 2$.  In this paper, we undertake an investigation of galaxy formation using an Adaptive Mesh Refinement (AMR) hydrodynamics code.  The majority of work in this field has used SPH codes, and so this allows us to investigate the problem from a new angle.  Although there has been some work with AMR codes \\citep{Joung:2009p1010, Ceverino:2009p1014, Agertz:2010p461, Colin:2010p1053}, there has not been a clear demonstration that an equivalent AMR calculation (i.e. one without a subgrid feedback model) actually does reproduce the classic SPH result.  We begin by simulating a set of five halos without any feedback or subgrid model (except a minimum pressure support to prevent artificial fragmentation).  We find, in agreement with SPH codes that a large, concentrated bulge is produced, resulting in a rotation curve that rises to $\\sim 500$ km/s at $r \\sim 1$ kpc.  We then vary a number of numerical and physical parameters in order to understand how sensitive the result is to our a choice of parameters.  The paper is organized as follows.  Section \\ref{method} describes the details of our hydrodynamics code, our initial conditions and the relevant parameters for this study.  In Section \\ref{results}, we present the results of our simulations including the five canonical runs, our resolution study, and our modified runs.  Section \\ref{discussion} is a discussion of our results and their implications.  Finally, Section \\ref{summary} summarizes our conclusions and makes predictions for future solutions to the angular momentum problem. ",
        "conclusions": "\\label{discussion}  These results demonstrate that it is challenging to generate disk systems with the correct mass distribution.  Without effective feedback, the default outcome is for dense clumps to lose angular momentum and result in centrally-cusped rotation curves.  In particular, the simulation with a very low star formation efficiency nicely demonstrates that this result is fundamentally a dynamical one, and does not depend on whether the clumps are primarily gas, or mostly stars.  As long as they are concentrated, they will lose angular momentum and hence rotational support.  Although this result is not new, and there is a long history of SPH simulations which found this result much earlier, we show it here clearly and systematically using a completely different numerical method (AMR).  Therefore, the result is quite general.  In addition, we went on to systematically vary our numerical parameters and investigate a range of resolution and feedback methods.  We confirmed that only cooling-suppression feedback models are capable of significantly changing the mass distribution and hence the rotation curve.  Cooling suppression models do effectively enhance feedback, although it is unclear how physically meaningful this technique is (see Section \\ref{code_coolingsuppression}), and a better approach might be to use high-resolution local models to generate subgrid models \\citep[see][for some attempts in this direction]{Yepes:1997p1093, Tasker:2006p1072, Ceverino:2009p1014, Joung:2009p1094}.  Another constraint is the baryon content of galactic halos: a variety of techniques have been used to infer that the baryon-to-dark-matter ratio in galaxies is much smaller than the cosmic mean \\citep[e.g.][]{Moster:2010p1095, Behroozi:2010p1096}, implying that a significant amount of mass has been ejected from galactic systems (or never accreted in the first place).  For example, Milky-Way massed halos only appear to host 20\\% of their baryons, with the fraction decreasing rapidly for smaller-mass systems \\citep{Behroozi:2010p1096}.  We find that all of our simulations result in very high disk baryon fractions; even the run with feedback and cooling suppression has a baryon fraction of about 63\\% of the cosmic mean.  This appears to be a general issue with cosmological galaxy simulations \\citep[see also the discussion in][]{AvilaReese:2011p1097}.  We demonstrate that simulations must be run to $z\\sim0$ in order to gauge the efficacy of model parameters at minimizing the effects of the angular momentum problem.  Many past studies \\citep[e.g.][]{Ceverino:2009p1014, Joung:2009p1010} have traded off simulation run time for increased resolution in their simulations.  While there are some early indicators of the angular momentum problem like bursts of star formation and peaked rotation curves at redshifts as early as $z=5$, successful results during this epoch do not guarantee successful results at $z\\sim0$.  We specifically demonstrate this in the case of the low star formation efficiency run H26SPML, which staved off early bursts of star formation but eventually succumbed to the same fate as runs with a normal star formation efficiency.  \\subsection{Comparison to Previous Work}  In agreement with previous work using SPH, we find that unless we include efficient feedback, the resulting systems are dominated by a too-large spheroidal component, and the resulting rotation curve is peaked in the center \\citep[e.g.][]{Navarro:1991p1002, Navarro:1997p1098, Weil:1998p1030, Abadi:2003p641, Donghia:2006p1032, Zavala:2008p1099, Piontek:2011p1041}.   \\begin{figure*} \\begin{center} \\plotone{figure14a.png} \\plotone{figure14b.png} \\end{center} \\caption{We render faceon and edgeon views for the last four of our modified halos at $z = 0$.  The top two rows show each galaxy in gas, whereas the bottom two rows display each galaxy's stellar component.  Each postage stamp has a width of 25 kpc.  For the gas, edgeon views show a volume rendering with isocontours of the gas density, whereas faceon views are column density bricks.  For the stars, we use the column density of stars in the faceon and edgeon views respectively.  The color of a star particle is representative of its age where blue represents young stars and red represents old stars according to the color bars on the right.  Notably, these modified runs produce young stellar disks and have a much higher density of gas in their cores than the canonical runs.} \\label{fig:modified_runs2_renderings} \\end{figure*}  A number of reasons have been suggested for this in the past, including purely numerical causes, such as angular momentum transfer between the disk and the hot halo \\citep[e.g.][]{Okamoto:2005p1101}, or between cold gas clumps and the hot halo \\citep[e.g.][]{Thacker:2000p1040}.  The concern was particularly that SPH simulations might be susceptible to this issue because of smoothing between hot and cold phases.  However, the fact that AMR simulations -- which use a completely different numerical method to solve the fluid equations -- find the same result, is an indication that these effects do not dominate.  Another suggested source of numerical angular momentum loss is in the form of gravitational instabilities which arise from inadequate resolution of the Jeans length in the disk \\citep{Truelove:1997p1046, Robertson:2004p1102}.  We tested this idea by running with and without an additional numerical (``Jeans\") pressure designed specifically to ensure that the Jeans length was adequately resolved, finding no difference in our results.  A third numerical reason is the lack of resolution \\citep{Governato:2004p1024, Kaufmann:2007p1103}; however, we specifically test this over the computational range available for us, and find no significant difference from 200-1700 pc.  This is in agreement with the SPH simulations of \\citet{Piontek:2011p1041}, who also varied their numerical resolution over a wide range, and found no difference.  There have been a number of recent cosmological AMR simulations which we can compare to.  \\citet{Colin:2010p1053} used the \\emph{ART} code \\citep{Kravtsov:1997p1104} to simulate a halo which is smaller by an order of magnitude (about $10^{11}$ M$_\\odot$), finding peaked rotation curves (decreasing as the star formation density criterion was reduced).  Although the halo masses are quite different, the essential result seems to be in agreement.    \\citet{Agertz:2010p461} used the \\emph{RAMSES} AMR code \\citep{Teyssier:2002p1105} to simulate a Milky-Way mass galaxy with similar resolution to that found here.  They argued that a low star formation efficiency by itself was enough to produce nearly flat rotation curves (and that feedback was only efficient if extreme amounts of energy were injected).  We have not been able to confirm the first suggestion -- using a range of low efficiencies for star formation, we find that clumps lose angular momentum at high-redshift, and generate steep rotation curves, whether they are in gas or stellar form.  The efficiency controls when the gas is converted to stars, but has little impact on the distribution of the material.  In this, we are in agreement with previous SPH work \\citep[e.g.][]{Weil:1998p1030, Donghia:2006p1032, Piontek:2011p1041}.  Finally, we have found that the only way to significantly decrease the peak of the rotation curve was to introduce a sub-grid model which enhanced the efficiency of stellar feedback.  We briefly explored the cooling suppression model and found this to be effective.  This agrees with a substantial number of SPH simulations which adopt this mechanism \\citep{Gerritsen:1997p1039, Thacker:2000p1040, SommerLarsen:2003p1116, Stinson:2006p1023, Governato:2007p1022, Agertz:2010p461, Colin:2010p1053, Piontek:2011p1041, Guedes:2011p1080}.  In addition, \\citet{AvilaReese:2011p1097} used AMR simulations (with the \\emph{ART} code) and also found cooling suppression to be effective in obtaining approximately flat rotation curves. \\citet{Ceverino:2009p1014} also used the \\emph{ART} code, but with a different sub-grid model, arguing that a model in which stars are born with significant velocities relative to the nascent gas will feed energy into low-density regions, producing efficient feedback and flat rotation curves, although the simulation is only run to $z \\sim 3$."
    },
    "1112/1112.1419_arXiv.txt": {
        "abstract": "We obtained $\\rm{U_{330}}$ and B band images of the M31 nucleus using the High Resolution Camera of the Advanced Camera for Surveys on board the Hubble Space Telescope (HST). The spatial resolution in the $\\rm{U_{330}}$-band, $0\\asec03$ FWHM, or 0.1 pc at M31, is sufficient to resolve the outskirts of the compact cluster (P3) of UV-bright stars surrounding the M31 black hole. The center of the cluster is marked by an extended source that is both brighter and redder than the other point sources within P3; it is likely to be a blend of several bright stars. We hypothesize that it marks the location of the M31 black hole. Both stellar photometry and a surface brightness fluctuation analysis, show that the P3 stellar population is consistent with early-type main sequence stars formed in a $\\sim100~-~200$ Myr old starburst population. Evolutionary tracks of post early asymptotic giant-branch stars, associated with late-stage evolution of an old population, also traverse the U and ${\\rm U-B}$ domain occupied by the P3 stars; but we argue that only a few stars could be accounted for that way. PEAGB evolution is very rapid, and there is no progenitor population of red giants associated with P3. The result that P3 comprises young stars is consistent with inferences from earlier HST observations of the integrated light of the cluster. Like the Milky Way, M31 harbors a black hole closely surrounded by apparently young stars. ",
        "introduction": "There is a highly-compact cluster of blue stars \\citep{l98, brown} at the heart of the M31 double nuclear-cluster of old stars. \\citet{kb} showed that this blue cluster in turn hosts the $\\sim10^8M_\\odot$ super-massive black hole also known to reside in the nucleus \\citep{d84, dr, k88, r90}.  \\citet{bend} argue that the integrated spectrum and spectral-energy distribution of the stars are consistent with their formation in a burst 200 Myr ago. Young stars are also closely bound to the nuclear black hole in our galaxy \\citep{forrest, allen, krabbe}, suggesting that this may not be a rare phenomenon. Understanding how apparently young stars were formed deeply interior to the tidal field of a supermassive black hole probes several issues of how stars and gas interact within such an extreme environment (see \\citealt{alex} for a review). Study of the blue cluster in M31 provides an additional context in which to test theories for the formation of the unusual population surrounding the Milky Way black hole.  There is a diversity of ideas for explaining the origin of the M31 blue cluster. \\citet{chang}, for example, argue that the cluster is the expected consequence of the non-axisymmetric structure of the surrounding M31 nucleus funneling gas from stellar mass-loss into orbit around the black hole, where it periodically reaches surface densities sufficient to induce collapse and star formation. Likewise, \\citep{levin}, \\citet{br}, and \\citet{ward} have advanced the formation of compact disks of young stars around the Milky Way black hole by massive central accretion of gas. In contrast, \\citet{demar} argue that the blue stars in M31 result from stellar collisions and tidal stripping acting on an old population of stars interacting within the high-velocity orbits and strong tidal-field associated with the black hole. We present new observations of the blue cluster obtained at the maximum angular resolution offered by the {\\it Hubble Space Telescope} to provide better constraints on the origin of the stars residing in this unusual environment.  \\subsection{The Discovery of the Blue Cluster}  The discovery of the blue cluster proceeded in stages. \\citet{king} imaged the nucleus of M31 with the HST Faint Object Camera at 1750\\AA, seeing a double structure similar to that discovered in the optical by \\citet{l93}. By the simple expedient of checking the position-angle of their images, however, they discovered that the optically-dimmer peak, designated P2 by Lauer et al., was actually brighter in the UV than P1. \\citet{king} lacked $S/N$ sufficient to resolve the spatial structure of the UV source at P2, noting only that it was highly-compact and possibly consistent with a point source, such as low-level emission associated with the weak AGN detected in the radio \\citep{crane}, or emission from a single PAGB star, such as those seen further out from the nucleus \\citep{k92,bert}.  In retrospect, the discovery of a UV-bright nuclear source explains the observations of \\citet{nieto}, who found the nucleus to have color gradient becoming {\\it bluer} towards its center, using CFHT images obtained at 3750\\AA. Discovery of the cluster may also have been presaged by an enigmatic reference in \\citet{red} to third-party unpublished ``Mt. Wilson material'' purporting to show that ``the spectrum of the nucleus is of a peculiar dwarf A type,'' a result supported by no other work prior to \\citet{bend} to the best of our literature research.\\footnote{We have also been unable to find photographic spectra of the M31 nucleus prior to 1937 in catalogues of Mt.\\ Wilson observations.  We thank Dr. Christopher Burns of the Carnegie Observatories for searching for this material.}  Later HST observations by \\citet{l98} and \\citet{brown} were able to resolve the UV-source; both papers argued that it is a cluster of stars. WFPC2/PC images at 3000\\AA\\ obtained by \\citet{l98} showed that the half-power radius of the blue source is $\\sim0.2$ pc. \\citet{l98} further combined their U and V band fluxes of the source with the \\citet{king} fluxes at 1750\\AA\\ to conclude that the source is consistent with an A-star spectral energy distribution.  \\citet{bend} refined this picture further by obtaining HST STIS spectroscopy of the nucleus over $\\lambda\\approx3600-5100$\\AA, and reanalyzing the F300W images presented in \\citet{l98}. The spectroscopy shows strong Balmer lines, including a strong Balmer break, consistent with either A0 main-sequence or giant stars, but not a population of white dwarfs of the same photospheric temperature.  The best match to the spectrum is provided by a $200\\pm50$ Myr population formed in a single burst. Populations of less than half this age would exhibit too much UV-flux to the blue side of the Balmer break.  \\citet{bend} further show that the stars are distributed in a flat disk with an exponential profile of scale-length of $0.37\\pm0.04$ pc in rapid rotation about the black hole.\\footnote{We use 770 Kpc as the distance to M31 \\citep{dm31}.} The rotational broadening of the STIS spectrum implies a black hole mass of $M_\\bullet=1.4\\times10^8M_\\odot,$ a significant upwards revision from earlier determinations.  \\subsection{The Blue Cluster in the Context of the M31 Nucleus}  The M31 black hole and its cluster of blue stars mark the center of a much larger nuclear star-cluster of complex structure. The surface brightness of the nuclear stellar system  begins to rise above the underlying M31 bulge interior to $r\\sim3''$ or $\\sim10$ pc from the center. The cluster becomes increasingly elliptical at smaller radii, with $1-b/a >0.3$ for $r<1\\asec7,$ but its isophotes remain concentric with the photo-center of the bulge for $r>1\\asec4$ \\citep{l93}. At yet smaller radii the nuclear cluster exhibits the double-peaked structure discovered by \\citet{l93}. The optically-brighter peak, P1, is offset by $0\\asec49$ from the secondary brightness peak, P2, and a slightly lesser amount from the M31 bulge photo-center.  The best explanation for the double morphology is that both peaks result from the line-of-sight projection of an eccentric disk of stars roughly $\\sim2$ pc in diameter bound to the black hole \\citep{trem}. The black hole strongly dominates the potential, thus the stars in the disk follow Keplerian orbits, spending most of their time at the orbital apo-center, creating an enhancement of light at P1. Significantly, the P1---P2 line twists away from the major axis of the outer nucleus by $\\sim20^\\circ.$ P1 and P2 are both redder in $V-I$ than the surrounding nucleus \\citep{l98}. High spatial-resolution spectroscopy shows that both P1 and P2 comprise old stellar populations with characteristics more like each other, than the underlying bulge \\citep{kb}. The ``Tremaine Disk'' thus appears to be a distinct component embedded in a much larger nuclear-cluster. The initial model of the disk provided by \\citet{trem} has been refined by the more detailed analysis presented by \\citet{pt03}. An interesting feature of the refined model is that the disk requires a central hole in order to generate the apparent minimum of stellar emission that occurs between P1 and P2.  The cluster of blue stars occurs at still smaller radii, $r<0\\asec1,$ and must be closely bound to the black hole, which was first inferred to exist from ground-based spectroscopy \\citep{d84, dr, k88, r90}. \\citet{kb} carefully registered their high spatial-resolution spectroscopy to the WFPC2 imagery of \\citet{l98}, concluding that the M31 black hole was coincident with the cluster and that both were within $0\\asec07$ of the M31 bulge photo-center. The conclusion that the cluster hosts the black hole was further confirmed by \\citet{bend}, who used STIS spectroscopy to show that the central velocity dispersion in the cluster rises to $1183\\pm200~\\rm{km~s^{-1}},$ with organized rotation peaking at $618\\pm81~\\rm{km~s^{-1}}.$ Building on the analysis of \\citet{kb}, they reduced the offset of the black hole from the bulge center to $0\\asec033$ in the ``anti-P1'' direction along the P1-P2 line. They further demonstrated that this slight offset is balanced by the asymmetry of the Tremaine disk --- the center of mass of the complete nuclear system is coincident with the bulge photo-center to $<0\\asec01.$  \\citet{bend} also designated the cluster as a third component, P3, of the inner nucleus. This designation includes a subtle redefinition of P2 as it previously had been introduced by \\citet{l93} and used by \\citet{l98}. The center of the cluster is coincident with the peak of a shallow surface-brightness cusp seen in V-band that falls within $0\\asec1$ of the bulge center; it was this location that the two Lauer et al.\\ papers adopted as the center of P2. \\citet{bend}, however, use P2 to denote the elongated and more diffuse concentration of older stars that extends from the cluster on the side opposite from P1.  In this schema, P1 and P2 explicitly correspond to the apo- and peri-center apses of the Tremaine disk, which has a mean peri-center that is {\\it not} coincident with the cluster.  The present environment of P3 appears to be largely quiescent, despite the immediate proximity of the black hole \\citep[see the extensive discussion of this topic by][]{li}. \\citet{crane} find a weak radio source coincident with P3, and \\citet{garcia, garcia2} find low-level X-ray emission associated with P3, as well \\citep[however, see][]{li}. On the other hand, a high spatial-resolution map of emission from ionized gas within in the nucleus finds a few sources within an arcsecond of P1, but nothing coincident with P2 or P3 \\citep{delb}. While central accretion of gas may have created the P3 stars in the past, there is no evidence that any significant reservoir of cold gas exists there at present.  \\subsection{A Closer Look\\dots}  Under the \\citet{bend} description of P3, the cluster is a disk comprising only a few hundred A-stars.  Given the scale length of the disk, we concluded that it should begin to show resolution into discrete sources in images with slightly higher angular resolution than were obtained by \\citet{l98}, such as could be provided by the HST ACS/HRC in the blue. The photometry of the brighter sources might be obtained directly, but in any case, the full image of the cluster could also be analyzed by surface-brightness fluctuation (SBF) analysis \\citep{sbf} to constrain the P3 stellar population. If the cluster indeed comprises main sequence A stars, it would look ``lumpy'' in HRC images. SBF could quantify the amplitude of the lumpiness, and in turn the typical luminosity of the stars that the cluster comprises. ",
        "conclusions": "We began the paper with a history of earlier work on P3, which already built a strong case that it is a cluster of young stars.  We have now partially resolved the cluster into individual stars, and characterized its form with significantly higher spatial resolution than was provided by the WFPC2/PC observations of \\citet{l98}; however, in broad detail the conclusions of that paper, and the later work of \\citet{bend} remain unaltered. Both a color-magnitude diagram (CMD) obtained of the resolved stars, and an SBF analysis performed on the total image, show that P3 plausibly comprises a population formed in a burst 100 to 200 Myr ago, consistent with the P3 SED and the spectrum obtained by \\citet{bend}. The result may also be supported by the \\citet{sag} spectroscopy observation of the inner bulge, which suggest that a burst of star formation occurred within the nucleus about $\\sim100$ Myr ago.  One small modification of the earlier results provided by the present work is that that P3 may be slightly younger than the 200 Myr age derived by \\citet{bend}, but that work really only ruled out an age much younger than 100 Myr.  Another new result is that the color of S11 suggests that a single-age starburst model for the cluster does not capture the full story of P3.  At some level, the large luminosity of P3 blends with the much bluer stars around it to generate an overall A-like SED.  It is intriguing that S11 provides the only visible trace of P3 in the V-band.   If it does mark the location of the M31 black hole, then it suggests that there may yet be another change in the stellar population of the nucleus within $\\sim0\\asec03$ or $\\sim0.1$ pc of the black hole.  Although we have concluded that the P3 stars are young, \\citet{king} and \\citet{brown} argue that the UV-bright stars in the bulge of M31, which of course surrounds the nucleus, represent the final stages of stellar evolution in an old metal-rich stellar population.  If there are truly two different populations contributing UV-bright stars to the bulge and nucleus of M31, it still begs the question of whether or not old remnants may be present in P3, as well as how far out the putative young stars in P3 may be traced out into the surrounding nucleus or even bulge.  We show that PEAGB tracks traverse the U vs.\\ ${\\rm U-B}$ domain defined by the P3 stars, thus it is possible that a few of the P3 stars are indeed old-evolved stars, rather than being young massive main sequence stars. However, since this is a very rapid evolutionary phase, it is extremely unlikely that the entire P3 cluster could be accounted for by PEAGB stars.  A strong factor in constraining the population of P3 is the compactness of P3, itself. The UV-stars outside the nucleus do trace the bulge-light at long wavelengths, but yet are not strongly concentrated within the outer nucleus; their abundance by total light or number is but a small faction of the bulge. Requiring any sort of HP-HB stars as the dominant P3 population requires an associated massive population of red giants also closely bound to the M31 black hole to continually generate the short-lived PEAGB stars.  There is no evidence for such a population in the V and I nuclear images of \\citet{l98}.  This same consideration would appear to also apply to the hypothesis of \\citet{demar}, who suggest that the P3 UV-bright stars may be the remnant cores of giants stripped by close encounters with the black hole. The high velocities associated with the black hole are such that stars passing close to the black hole from the bulge, or even the outer nucleus will not dwell in the vicinity of P3. If the P3 stars are really processed giants, they still must have been initially closely bound to the black hole, which again begs the question of where the progenitor giants are. The only way that this mechanism might work is if the stripping of the P3 stars takes place just as they begin their first ascent up the red giant branch. This might produce long-lived UV-bright stripped-cores, and would account for lack of progenitor giants as well.  Further investigation of the origin of P3 might be linked to understanding the early-type stars closely bound to the black hole in our own galaxy.  The existence of P3 in a second Local Group galaxy suggests that this is not a rare phenomenon.  The present observations offer an additional site to test mechanisms that can form stars within the strong tidal field of a supermassive black hole.  We finish by noting that as with the galactic center, it may be possible to measure the proper motions of stars closely bound to the M31 black hole.  Figure \\ref{fig:stars} shows that circular velocities in the plane of the sky around the M31 black hole will exceed $1000~{\\rm km~s^{-1}}$ over the entire extent of P3.  It is possible that proper motions may be detected in about a decade after the present observations were made."
    },
    "1112/1112.4716_arXiv.txt": {
        "abstract": "We show that, even for extremely stiff systems, explicit integration  may compete in both accuracy and speed with  implicit methods if algebraic methods are used to stabilize the numerical integration. The required stabilizing algebra depends on whether the system is well-removed from equilibrium or near equilibrium. This paper introduces a quantitative distinction between these two regimes and addresses the former case in depth, presenting explicit asymptotic methods appropriate when the system is extremely stiff but only weakly equilibrated. A second paper \\cite{guidQSS} examines quasi-steady-state methods as an alternative to asymptotic methods in systems well away from equilibrium and a third paper \\cite{guidPE} extends these methods to equilibrium conditions in extremely stiff systems using partial equilibrium methods. All three papers present systematic evidence for timesteps competitive with implicit methods. Because an explicit method can execute a timestep faster than an implicit method, algebraically-stabilized explicit algorithms might permit integration of larger networks than have been feasible before in various disciplines. ",
        "introduction": "In many scientific and technical contexts one encounters phenomena that may be modeled by fluxes transferring population between sources and sinks for various species.  Examples include kinetic processes that modify abundances and transfer energy in atomic, molecular, and nuclear systems;  geochemical, climate, and other environmental systems; electrical circuits; economic models; and population dynamics. Terminology varies  but let us refer generically to these sources and sinks as {\\em boxes,} and term the resulting systems of boxes connected by fluxes {\\em reaction networks.} Such systems are commonly modeled by a coupled set of differential equations that describe a continuous flow of population through the boxes.  The reaction network is often classified as a {\\em stiff system}, which we shall define to be a system of equations containing multiple timescales ranging over many orders of magnitude \\cite{gear71,lamb91,press92,oran05}. Most  physical systems involve important processes operating on very different timescales, so realistic problems tend to be at least moderately stiff. Some, such as those encountered in many astrophysics applications, are extremely stiff, with fastest and slowest timescales in the problem differing by as much as 10--20 orders of magnitude. In stiff systems the timestep constraints are set by numerical stability requirements rather than accuracy considerations. Hence, explicit numerical integration of stiff systems is usually impractical because the maximum stable timestep is far too small for efficient solutions (see, for example, Refs.\\ \\cite{press92,oran05}).  This is commonly addressed by employing implicit or semi-implicit stiff solvers that are stable, but that require time-consuming iterative matrix solutions.  A given box in a reaction network often is connected strongly only to a few other boxes.  For example, the explosive burning conditions encountered in astrophysical novae, X-ray bursts, or supernovae may require reaction networks with hundreds to thousands of nuclear isotopes.  Yet individual isotopes are typically connected directly to other isotopes through (at most) $\\sim$10 reactions of consequence, and under many conditions no more than 2--3 reactions are important for a given isotope. Such restrictions on the direct box reaction coupling imply that the matrices appearing in the iterative implicit solution are {\\em sparse.} Although various methods are available to deal with sparse matrices, in practice many codes for solving large reaction networks have not exploited sparseness in particularly effective ways.  For example, in astrophysical calculations with implicit solvers in large networks (say $\\sim 150$ species or more), one finds often that greater than 90\\% of the processor time is consumed in matrix operations \\cite{timmes,hix05}. Efficient algorithms exist for the required matrix algebra (with incremental improvements in them over time), but the matrix nature of the core problem implies that the time required for implicit solution grows non-linearly with the size of the network. In typical working codes for large-scale applications, increasing the size of the network increases the time for solution, often quadratically, sometimes as much as cubically, until there are enough boxes in the network to justify the overhead of sparse-matrix methods with more favorable scaling. In applications in thermonuclear networks, for example, it is often found that the overhead required to implement sparse-matrix iterative solutions is not justified until there are several hundred boxes in the network. Thus, many present implicit stiff-network algorithms do not scale very gracefully to larger networks.   We are primarily interested in the most ambitious applications of large networks, where the reaction network is only a portion of a larger problem.  Let us take as representative astrophysical thermonuclear reaction networks, where a proper description of the overall problem typically requires multi-dimensional hydrodynamics or radiation hydrodynamics coupled tightly to a large thermonuclear reaction network.  The hydrodynamical evolution controls the conditions in the network such as temperature and density, and the network influences the hydrodynamic evolution strongly through energy production and modification of composition variables. As a consequence of the limitations discussed in the preceding paragraphs, the solution of large networks by the usual approaches is time-consuming and few calculations have attempted to couple the element and energy production strongly to the hydrodynamics with a network of realistic complexity.  The most ambitious approaches use very small networks, perhaps tuned empirically to get critical quantities like energy production correct on average, coupled to the hydrodynamical simulation.  In many calculations even this is not done and the network is replaced entirely by parameterization.  Then a more complete network is run in a separate ``post-processing'' step, where fixed hydrodynamical profiles computed in the hydrodynamical simulation with the small network are used to specify the variation of thermodynamic variables such as temperature and density with time.  Astrophysical thermonuclear networks have been used for illustration, but many problems of scientific and technical interest  exhibit similar complexity. Examples include astrochemical kinetics, where one must model large chemical evolution networks in contracting molecular clouds, or combustion chemistry, where chemical burning networks are strongly coupled to simulations of the dynamics of the air and fuel mixture. Physically-realistic networks in such contexts would often be quite large. In combustion of larger hydrocarbon molecules or studies of soot formation, hundreds to thousands of reacting species undergoing as many as 10,000 reactions may be encountered \\cite{oran05}, and in supernova explosions hundreds to thousands of nuclear isotopes with tens of thousands of reaction couplings make non-zero contributions \\cite{hix05}. For such cases one finds that current techniques do not allow for a coupling of realistic reaction networks to the full dynamics of the problem and often severely truncated or highly schematic reaction networks have been used in even the most realistic simulations. ",
        "conclusions": "Explicit numerical integration can compute a timestep faster then implicit methods, and the time to compute a network explicitly scales linearly and therefore more favorably with network size than for implicit codes. Nevertheless, previous discussions of numerical integration for very stiff systems have concluded rather uniformly that explicit methods are not competitive with implicit methods for stiff networks because they are unable to take large enough stable timesteps. To quote {\\em Numerical Recipes} \\cite{press92}, ``For stiff problems we {\\em must} use an implicit method if we want to avoid having tiny stepsizes.''  Improvements in explicit methods based on using asymptotic and steady-state limiting solutions to remove stiffness from the network have had some success for systems of moderate stiffness such as various chemical kinetics problems. However, it has been concluded in the previous literature that such methods are not competitive, failing even to give correct results, with timesteps that are far too short to be useful even if they gave correct results, for extremely stiff networks such as those encountered commonly in astrophysical thermonuclear networks \\cite{oran05,mott99}. This paper has presented evidence strongly challenging all of these conclusions.   We have cleanly identified three fundamentally different sources of stiffness in large networks, only the first of which is commonly emphasized in the literature: \\begin{enumerate} \\item Situations where small populations can become negative if the explicit timestep is too large, with the propagation of this anomalous negative population leading to exponentially growing terms that destabilize the network.  \\item Situations where the right sides of the differential equations expressed as $dY = F = \\fplus{} - \\fminus{}$ approach a constant derived from the difference of two large numbers (the total flux in $\\fplus{}$ and total flux out $\\fminus{}$), and numerical errors in taking this difference destabilize the network if the timestep is too large.  \\item Situations where on the right sides of the differential equations expressed in the form of Eq.~(\\ref{equilDecomposition}) the net flux in specific forward-reverse reaction pairs $(f^+_i - f^-_i)$ tends to zero as the system approaches equilibrium, leading to large errors if the timestep is too large because the net flux is derived from the difference of two large numbers and the timescale equilibrating the populations is short compared with the desired numerical timestep.  \\end{enumerate} We have shown that these distinctions are important because different sources of stiffness require different approximations for their removal in an algebraically-stabilized explicit integration.  Using the extremely stiff systems characteristic of astrophysical thermonuclear networks as a stringent test, we have shown that asymptotic methods are very successful at removing the first two types of stiffness, and give correct results, even for the stiffest of thermonuclear networks, provided that adequate attention is paid to conservation of probability in the network.  Furthermore, we have shown various examples of stable and accurate timestepping with these methods in extremely stiff systems that are competitive with that of standard implicit codes, demonstrating in some simple but physically-important networks timesteps that are as much as 20 orders of magnitude larger than the maximum timestep that would be stable in a standard explicit method.  Asymptotic methods are adept at removing the first two types of stiffness listed above, permitting explicit numerical timesteps that are competitive with implicit methods even in the stiffest networks.  However,  we have also shown that such methods give correct results but fail to exhibit competitive timestepping when the system approaches microscopic equilibrium and the third type of stiffness instability begins to dominate. In a following paper \\cite{guidPE}, we shall provide evidence for competitive timestepping, even in the approach to equilibrium, if the explicit asymptotic method is supplemented by partial equilibrium approximations designed specifically to deal with the third type of stiffness instability.  Taken together, this paper and the following ones on quasi-steady-state methods \\cite{guidQSS} and partial equilibrium methods \\cite{guidPE} present compelling evidence that algebraically-stabilized explicit integration methods are capable of timesteps competitive with implicit integration methods for a variety of highly-stiff reaction networks. Since explicit methods can execute a timestep faster than an implicit method in a large network, our results suggest that algebraically-stabilized explicit algorithms may be capable of performing as well as, or even substantially outperforming, implicit integration in a variety of moderate to extremely stiff applications. Because of the highly-favorable linear scaling for explicit methods, this fundamentally new view of the efficacy of explicit integration for stiff equations may be particularly important for applications in any field where it is imperative that more realistic---and therefore much larger---networks be used in complex physical simulations.   \\begin{ack} We thank Tony Mezzacappa for useful discussions, Austin Harris for help with some of the calculations, Eric Lingerfelt for programming assistance, and Christian Cardall for a careful reading of the manuscript. Research was sponsored by the  Office of Nuclear Physics, U.S. Department of Energy.  \\end{ack}  \\clearpage"
    },
    "1112/1112.3300_arXiv.txt": {
        "abstract": "We measure the gas-phase oxygen abundances of $\\sim3000$ star-forming galaxies at $z=0.05-0.75$ using optical spectrophotometry from the AGN and Galaxy Evolution Survey (AGES), a spectroscopic survey of $I_{\\rm AB}<20.45$ galaxies over $7.9$~deg$^{2}$ in the NOAO Deep Wide Field Survey (NDWFS) \\bootes{} field.  We use state-of-the-art techniques to measure the nebular emission lines and stellar masses, and explore and quantify several potential sources of systematic error, including the choice of metallicity diagnostic, aperture bias, and contamination from unidentified active galactic nuclei (AGN).  Combining volume-limited AGES samples in six independent redshift bins and $\\sim75,000$ star-forming galaxies with $r_{\\rm AB}<17.6$ at $z=0.05-0.2$ selected from the Sloan Digital Sky Survey (SDSS) that we analyze in the identical manner, we measure the evolution of the stellar mass-metallicity (\\mz) between $z=0.05$ and $z=0.75$.  We find that at fixed stellar mass galaxies at $z\\sim0.7$ have just $30\\%-60\\%$ the metal content of galaxies at the present epoch, where the uncertainty is dominated by the strong-line method used to measure the metallicity.  Moreover, we find no statistically significant evidence that the \\mz{} relation evolves in a mass-dependent way for $\\mass\\simeq10^{9.8}-10^{11}$~\\msun{} star-forming galaxies.  Thus, for this range of redshifts and stellar masses the \\mz{} relation simply shifts toward lower metallicity with increasing redshift without changing its shape. ",
        "introduction": "\\label{sec:intro}  Like stellar mass, the gas-phase metallicity of a galaxy is a sensitive observational diagnostic of its past star formation history and present-day evolutionary state for the simple reason that in a closed system metallicity increases monotonically with each successive generation of massive stars.\\footnote{Throughout this paper we use the terms {\\em metallicity} and {\\em abundance} interchangeably to mean the heavy-element content of the warm ($T\\approx10^{4}$~K) interstellar medium of galaxies.  In particular, we make the reasonable assumption that the nebular oxygen abundance, historically written as \\logoh, traces the total gas-phase metallicity, $Z_{\\rm gas}$.}  In detail, however, galaxies are not closed systems: infall of cold, metal-poor gas from the intergalactic medium, rapid gas accretion via minor and major mergers, and supernova-driven winds of gas and metals can modulate the metallicity of individual galaxies according to their large-scale environment, gas supply, and assembly history.  Therefore, accurate abundance measurements provide valuable insight into the interplay between many fundamental processes in galaxy evolution, including star formation, gas accretion, and supernova-driven feedback across cosmic time \\citep{tinsley80a, pagel97a, pettini04b}.  Nebular emission lines---ubiquitous in the rest-frame optical spectra of star-forming galaxies---provide a particularly powerful way to study the chemical abundances of both nearby and distant galaxies. Among the most commonly observed lines are the Balmer \\halam, \\hblam, and \\hglam{} hydrogen recombination lines, and the collisionally excited \\oiidoublet, \\neonlam, \\oiiidoublet, \\niidoublet, and \\siidoublet{} forbidden lines.  Because these lines originate principally in star-forming (\\hii) regions, they trace the physical conditions in the gas from which the current generation of massive stars is forming.  The metal lines in particular are the principal coolants in \\hii{} regions, making them sensitive to the total abundance of oxygen and other heavy elements in the interstellar medium.  Moreover, the relative line-strengths can be used to infer the interstellar pressure, density, temperature, ionizing radiation field strength, dust reddening, and the presence of an active galactic nucleus (AGN).  Finally, from an observational standpoint, many of the nebular lines are intrinsically strong, making them measurable in even relatively low signal-to-noise spectra over a broad range of redshifts.  Following the discovery of the correlation between dynamical mass and oxygen abundance in dwarf irregular galaxies (\\citealt{lequeux79a}; see also \\citealt{garnett87a, oey93a}), numerous subsequent studies demonstrated that star-forming galaxies ranging from the lowest-luminosity dwarfs to massive disk galaxies obey a well-defined luminosity-metallicity (\\lz) correlation spanning several orders of magnitude in $B$-band luminosity \\citep{garnett87a, skillman89a, zaritsky94a, pilyugin04a, lama04a, salzer05a, lee06b, moustakas10a}. Metallicity was also found to correlate with morphological type \\citep{edmunds84a}, surface mass density \\citep{mccall82a, ryder95a, garnett97a}, and maximum rotational velocity \\citep{zaritsky94a, garnett02a, dalcanton07a}.  However, because these global properties all correlate with one another, the underlying physical driver of the \\lz{} relation remained elusive.  \\citet{tremonti04a} was the first to leverage the tremendous statistical power of the Sloan Digital Sky Survey \\citep[SDSS;][]{york00a} spectroscopic database to show that the gas-phase metallicity of a galaxy correlates best with stellar mass, \\mass, now known as the stellar mass-metallicity (\\mz) relation.  The \\mz{} relation reveals that the gas-phase metallicity of star-forming galaxies increases monotonically with stellar mass and then remains relatively constant above $\\mass\\approx3\\times10^{10}~\\msun$. Although subsequent studies have found weak residual correlations from the \\mz{} relation with stellar mass density \\citep{tremonti04a, liang10a}, size \\citep{ellison08a}, large- and small-scale environment \\citep{mouhcine07a, cooper08a, ellison08b, ellison09a, peeples09a}, star formation rate \\citep[SFR;][]{mannucci10a, lara-lopez10a, yates11a, cresci11a}, and the existence of bars \\citep{ellison11a}, the intrinsic dispersion in the \\mz{} relation is $\\lesssim0.1$~dex, making it among the tightest empirical correlations known.  The physical origin of the \\mz{} relation remains under debate.  One possibility is that low-mass galaxies started forming stars later than massive galaxies \\citep[i.e., they are ``younger'';][]{noeske07b, leitner11b} and have been less efficient at synthesizing metals via star formation \\citep{brooks07a, mouhcine08a, calura09a}.  This interpretation is qualitatively consistent with the measured low gas fractions in massive galaxies relative to lower-mass galaxies \\citep{mcgaugh97a, geha06a, garcia-appadoo09a}, and with the observed correlation between stellar mass and SFR \\citep[the {\\em star formation sequence};][]{brinchmann04a, salim07a, elbaz07a, noeske07a}.  Another popular interpretation is that supernova-driven galactic winds preferentially expel metals from low-mass galaxies \\citep{larson74a, garnett02a, tremonti04a, dalcanton07a}.  Indeed, state-of-the-art hydrodynamic and semianalytic theoretical models \\emph{require} metal-enriched outflows to match many observed galaxy properties, including the \\mz{} relation \\citep{kobayashi07a, finlator08a, dutton09a, peeples11a, dave11b}.  A related outstanding question is whether gas and metals ejected in previous star formation episodes are re-accreted onto galaxies \\citep{delucia04a, oppenheimer10a}, or whether galaxies are predominantly fed by infall of cold, metal-poor gas from the intergalactic medium \\citep{koppen99a, faucher11a}.  Alternatively, \\citet{koppen07a} show that a SFR-dependent, and therefore mass-dependent, stellar initial mass function (IMF) naturally explains the observed \\mz{} relation without needing to invoke metal-enriched outflows from galaxies.  One way of gaining insight into the complex and multifaceted interrelationship between chemical abundance measurements, star formation, gas accretion, and the role of supernova feedback during galaxy growth is to measure the evolution of the \\mz{} relation.  The time since $z=1$, spanning $\\sim60\\%$ of the age of the Universe, is important for many reasons.  Measurements of the star formation sequence \\citep{noeske07a, wuyts11a} suggest that star-forming galaxies over this redshift range evolve smoothly as a population, driven by continuous, secular processes like gas consumption and {\\em in situ} star formation \\citep{bell05a, noeske09a, leitner11b}. Bolstered by measurements of their morphological distribution \\citep{bell05a, konishi11a} and clustering properties \\citep{coil08a, zehavi11a}, these results suggest that we can draw direct evolutionary connections linking star-forming galaxies across this redshift range.  In addition, the metallicity-sensitive \\pagel{} parameter (defined in \\S\\ref{sec:oh}) can be measured from ground-based optical spectroscopy of galaxies at least to $z\\sim0.8$, ensuring that metallicities can be derived using a single consistent abundance calibration.  Unfortunately, previous studies have reported widely varying results on the evolution of both the shape and normalization of the \\mz{} relation since $z=1$.  \\citet{savaglio05a} and \\citet{zahid11a} find that the mean metallicity of galaxies at $z\\sim0.8$ with $\\mass\\sim10^{10.3}$~\\msun{} differs by just $\\sim0.05$~dex (factor of $\\sim1.12$) relative to similarly massive star-forming galaxies at $z\\sim0.1$, whereas they find that galaxies with $\\mass\\sim10^{9.5}$~\\msun{} undergo a factor of $2-3$ more chemical evolution over the same redshift interval.  On the other hand, \\citet{liang06a} and \\citet{cowie08a} report $0.2-0.3$~dex (factor of $1.6-2$) of chemical evolution for galaxies with $\\mass\\sim10^{10.3}$~\\msun{} since $z\\sim0.7$ and no statistically significant evidence for evolution in the shape of the \\mz{} relation \\citep[see also][]{rodrigues08a}.  \\citet{lama09a} and \\citet{perez-montero09a} find an even greater amount of metallicity evolution, $\\sim0.35$~dex (factor of $2.2$) since $z=0.6-1$, and a \\emph{flattening} of the \\mz{} relation at intermediate redshift. Other studies based on the evolution of the $B$-band \\lz{} relation report that at fixed luminosity star-forming galaxies at $z=0.5-1$ are $0.1-0.7$~dex (factor of $1.25-5$) more metal-poor than local star-forming galaxies \\citep{kobulnicky99b, carollo01a, lilly03a, kobulnicky03a, kobulnicky04a, maier04a, maier05a, maier06a, lama09a}.  However, the significant amount of luminosity evolution experienced by blue, star-forming galaxies since $z=1$ \\citep{blanton06a, faber07a, cool12a} renders the interpretation of the \\lz{} relation less straightforward.  The origin of these widely varying results on the evolution of the \\mz{} relation can be attributed to various issues, including cosmic variance, small sample size, heterogenous selection criteria, systematic differences in the methods used to infer nebular abundances, an inconsistent analysis of the local \\mz{} relation, spectroscopy with insufficient spectral coverage, signal-to-noise (S/N) ratio, or instrumental resolution, and uncertain stellar mass estimates due to limited broadband photometric coverage. Consequently, the efficacy of existing observational metallicity constraints on state-of-the-art theoretical models of galaxy formation \\citep{delucia04a, brooks07a, dutton09a, dave11b} has been fairly limited.  From an observational standpoint, there exists a clear need for a large, homogeneously selected sample of galaxies with high-quality spectroscopy and reliable stellar mass estimates to better constrain the evolution of the \\mz{} relation at intermediate redshift.  To address this need, we measure the evolution of the \\mz{} and $B$-band \\lz{} relations at intermediate redshift using oxygen abundances of $\\sim3000$ star-forming galaxies at $z=0.05-0.75$ observed as part of the AGN and Galaxy Evolution Survey \\citep[AGES;][]{kochanek11a}, and $\\sim75,000$ galaxies at $z=0.05-0.2$ selected from the SDSS.  The AGES main galaxy survey consists of optical spectrophotometry for $\\sim12,000$ galaxies in the $\\sim9$~deg$^{2}$ NOAO Deep Wide Field Survey \\citep[NDWFS;][]{jannuzi99a, brown03a, brown07a, brown08a} \\bootes{} field at a median redshift of $z\\sim0.3$.  This sample is statistically complete over $7.9$~deg$^{2}$ for galaxies with $I_{\\rm AB}<20.45$ \\citep{kochanek11a}, reaching $\\sim2$~mag deeper than the SDSS main galaxy sample \\citep{strauss02a} over a considerably larger area than other recent or ongoing surveys of intermediate-redshift galaxies such as DEEP2 \\citep{davis03a}, AEGIS \\citep{davis07a}, VVDS \\citep{lefevre04a, lefevre05a, garilli08a}, and zCOSMOS \\citep{lilly09a}.\\footnote{For galaxies at $z\\lesssim0.5$, the Galaxy and Mass Assembly (GAMA) survey will ultimately supplant all these surveys by obtaining optical spectroscopy for several hundred thousand galaxies brighter than $r_{\\rm AB}\\approx19.8$ over $360$~deg$^{2}$ \\citep{driver11a}.}  The statistical completeness, large sample size, and availability of high-quality optical spectroscopy and deep optical and near-infrared photometry enables us to investigate the evolution of both the \\mz{} and optical \\lz{} relations of star-forming galaxies at $z=0.05-0.75$ using a single, self-consistent abundance diagnostic over the entire redshift range.  The plan of the paper is as follows.  In \\S\\ref{sec:data} we present the ground-based optical and near-infrared imaging we use, summarize the AGES observations, and describe how we measure the nebular emission lines for the galaxies in our sample.  In \\S\\ref{sec:sed} we describe the methodology used to derive rest-frame luminosities, colors, and stellar masses, and in \\S\\ref{sec:sample} we select a subset of the AGES galaxies for chemical abundance analysis.  We present the methods we use to derive oxygen abundances in \\S\\ref{sec:oh}, and our principal results in \\S\\ref{sec:evol}, where we quantify the mass-dependent evolution of the \\mz{} relation for star-forming galaxies since $z=0.75$.  In \\S\\ref{sec:syseffects} we address the effect of various potential sources of systematic uncertainty on our results, and in \\S\\ref{sec:discussion} we discuss recent theoretical work on the origin and evolution of the \\mz{} relation.  Finally, we summarize our principal conclusions in \\S\\ref{sec:summary}.  We adopt a concordance cosmology with $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $h_{70}\\equiv H_{0}/100=0.7$, the AB magnitude system \\citep{oke83a}, and the \\citet{chabrier03a} initial mass function (IMF) from $0.1-100~\\mathcal{M}_{\\sun}$ unless otherwise indicated.  For reference, the conversion from Vega to AB for the $I$-band filter used to select AGES targets is $+0.45$~mag, that is $I_{\\rm AB}=I_{\\rm Vega}+0.45$. ",
        "conclusions": ""
    },
    "1112/1112.1844_arXiv.txt": {
        "abstract": "We present the result of a systematic study of pileup phenomena seen in the X-ray Imaging Spectrometer, an X-ray CCD instrument, onboard the Suzaku observatory. Using a data set of observed sources in a wide range of brightness and spectral hardness, we characterized the pileup fraction, spectral hardening, and grade migration as a function of observed count rate in a frame per pixel. Using the pileup fraction as a measure of the degree of pileup, we found that the relative spectral hardening (the hardness ratio normalized to the intrinsic spectral hardness), branching ratio of split events, and that of detached events increase monotonically as the pileup fraction increases, despite the variety of brightness and hardness of the sample sources. We derived the pileup fraction as a function of radius used for event extraction. Upon practical considerations, we found that events outside of the radius with a pileup fraction of  1\\% or 3\\% are useful for spectral analysis. We present relevant figures, tables, and software for the convenience of users who wish to apply our method for their data reduction of piled-up sources. ",
        "introduction": "\\label{section:1}  Charge coupled devices (CCDs)  have been playing an important role in the modern X-ray astronomy (Lumb et al. 1991) since the SXT (Soft X-ray Telescope;  \\cite{Tsuneta1991}) mounted on YOHKOH. As a photon-counting detector, the first successful use of it in space was realized with the SIS (Solid-state Imaging Spectrometer; Burke et al. 1991) on board the ASCA satellite.  The fifth Japanese X-ray satellite Suzaku (\\cite{Mitsuda2007}) also carries the CCD devices, the X-ray Imaging Spectrometer (XIS; \\cite{Koyama2007}), located at the foci of the X-ray Telescope (XRT; \\cite{Serlemitsos2007}) modules, as well as a non-imaging hard X-ray instrument, the Hard X-ray Detector (HXD; \\cite{Takahashi2007}). The XIS covers soft X-rays from 0.2 to 12 keV, while the HXD covers from 10 to 600 keV. The wide-band coverage by a combination of these instruments is a key competence of Suzaku.  Sources that have rich statistics in the HXD are usually bright in the XIS as well. In observations of such a bright source, more than one photon can strike the same CCD pixel, or one of its immediate neighbors during the exposure time per frame. It  causes a complicated and irreversible phenomenon called ``pileup''. For example, when two photons with energies of $E_1$ and $E_2$ fall in the same pixel within an integration period, it is impossible to distinguish them from one with an energy of $E_1 + E_2$. It affects obtained spectra, images, and light curves. Incorrect evaluation of pileup can lead to wrong interpretations of the data. This is especially the case for bright sources because their small statistical uncertainty can be overwhelmed by systematic uncertainties brought by pileup. Therefore, it is important to understand pileup effects in order to maximize the wide-band capability of Suzaku.  In general, when pileup occurs, flux and spectra become lower and harder, respectively. The pileup effects of the SISs of ASCA, have been studied by \\citet{Ebisawa1996}, \\citet{Ueda1997}, and Kotani et al. (2000). Recently, a detailed study for the CCD detector of Chandra, the ACIS, as well as a generalized formula of the pileup, has been compiled by Davis (2001). However, the telescopes used for the XIS and the ACIS are entirely different in the shape and size of the point spread function (PSF), which adds non-linear complications to the effect. Therefore, it is difficult to simply apply methods developed for the ACIS to the XIS. Likewise, the way of estimating pileup\\footnote{http://xmm.esac.esa.int/sas/current/documentation/ \\\\ threads/epatplot.shtml} for the CCD detector of XMM-Newton, the EPIC, can not be directly applied to the XIS. Thus, studies on the pileup effects should be performed specifically for the combined system of XISs and XRTs.   One of the practical solutions to mitigate the pileup in the XIS is the \"core-exclusion\" method, in which the central region of the PSF is excluded to avoid the CCD region severely affected by pileup. This method was adopted in the analysis of 4U~1630$-$472 (Kubota et al. 2007), Cygnus X-1 (Makishima et al. 2008), GX 339$-$4 (Yamada et al. 2009), and others. Based on these experiences, we released a recipe to apply the method for XIS data as an official document\\footnote{http://www.astro.isas.ac.jp/suzaku/analysis/xis/pileup/ \\\\ HowToCheckPileup\\_v1.pdf}. In the document, it is poorly stated how to estimate the size of the core to be excluded. This is often difficult to do by using data that are affected by pileup.  In this paper, we present the results of a systematic analysis for several bright sources affected by pileup in the XIS. A brief summary on pileup characteristic for phenomenological specific to the XIS+XRT system is presented in section 2. After describing target selection and data processing in section 3, systematic studies on pileup effects are presented in section 4. In section 5, we discuss a criterion that utilizes pileup fraction and show concrete examples for the application. The main results of the paper are summarized in section 6. ",
        "conclusions": "We have comprehensively studied pileup phenomena of the XIS, by analyzing the actually observed data for several bright point sources. Using the surface brightness, the pileup fraction as the primary measure, we characterized grade-branching ratios of split events and detached events. We examined the effects on spectral analysis when excluding the central regions within $R_{3\\%}$ and $R_{1\\%}$. Our results are summarized as follows:  \\begin{itemize} \\item Pileup causes spectral hardening in FI and BI CCDs, as well as changes in the sharp spectral features such as atomic edges of instrumental response. The changes in hardness is not dependent on intrinsic spectral hardness. \\item Pileup increases the ratio of split events by $\\sim$ 20\\% from a pileup fraction of $\\sim$ 1\\% to $\\sim$ 10\\% in both FI and BI CCDs. On the other hand, the ratio of the detached events increases proportionally with the pileup fraction by $\\sim$ several orders of magnitude from $\\sim$ 0.1 and $\\sim$ 10 \\% of pileup fraction. \\item When surface brightness, spectral hardness, and split event ratio begin to increase, the pileup fraction starts to exceed $\\sim$ 1\\%, which illustrate that the pileup fraction can be a good measure for estimating pileup. \\item The relation between the measured surface brightness and the radius corresponding to pileup fraction of 3\\% and 1\\%, is useful to estimate the extent of pileup in a practical manner. \\item Based on the observation of Cygnus X-1, the differences between the spectra accumulated outside $R_{3\\%}$ and $R_{1\\%}$ are less than 5\\% in the 2--10 keV and $\\sim$ 20 \\%  in 0.5--2 keV. The latter might be caused by uncertainties of contamination modeling in the XIS. \\end{itemize}  Thus, we have phenomenologically compiled the pileup effects specific to the XIS + XRT systems and found a practical way of estimating pileup effects, which would be useful for those who analyze the bright sources taken with Suzaku, and even with similar CCD sensors to the XIS+XRT system, such as Soft X-ray Imager on board ASTRO-H (Tsunemi et al. 2010).  \\vspace{0.5cm}  The authors would like to express their thanks to Suzaku and the XIS team members, and anonymous referee for helpful comments and suggestions. The present work was supported by Grant-in-Aid for JSPS Fellows and RIKEN."
    },
    "1112/1112.4466_arXiv.txt": {
        "abstract": "In this paper we investigate scaling relations between star formation rates and molecular gas masses for both local Galactic clouds and a sample of external galaxies. We  specifically consider relations between the star formation rates and measurements of dense, as well as total, molecular gas masses.  We argue that there is a fundamental empirical scaling relation that directly connects the local star formation process with that operating globally within galaxies. Specifically, the total star formation rate in a molecular cloud or galaxy is linearly proportional to the mass of dense gas within the cloud or galaxy. This simple  relation, first documented in previous studies,  holds over a span of mass covering nearly nine orders of magnitude and indicates that the rate of star formation is directly controlled by the amount of dense molecular gas that can be assembled within a star formation complex. We further show that the star formation rates and total molecular masses, characterizing both local clouds and galaxies, are correlated over similarly large scales of mass and can be described by a family of linear star formation scaling laws, parameterized by $f_{DG}$, the fraction of dense gas contained within the clouds or galaxies. That is, the underlying star formation scaling law is always linear for clouds and galaxies with the same dense gas fraction. These considerations provide a single unified framework for understanding the relation between the standard (non-linear) extragalactic Schmidt-Kennicutt scaling law, that is typically derived from CO observations of the gas, and the linear star formation scaling law derived  from HCN observations of the dense gas. ",
        "introduction": "Knowledge of the physical factors that control the conversion of interstellar gas into stars is of fundamental importance for both developing a predictive physical theory of star formation and understanding the evolution of galaxies from the earliest epochs of cosmic history to the present time. An essential first step to obtaining such knowledge is to establish empirically the underlying relation or relationships that most directly connect the rate of star formation in a galaxy to some general physical property of the interstellar gas from which stars form. A little more than a half-century ago, Schmidt (1959) conjectured that this might take the form of a scaling relation between the rate of star formation and some power, n,  of the surface density of atomic (HI) gas. From evaluation of  the distributions of local HI gas and stars orthogonal to the Galactic plane, he suggested that n$\\approx$ 2.  Subsequent studies comparing the surface densities of OB stars and HII regions with those of atomic gas within nearby external galaxies produced scaling laws with similar, super-linear, power-law indices (e.g., Sanduleak 1969; Hamajima \\& Tosa 1975).  By the 1980s it became clear that molecular, not atomic, clouds were the sites of star formation in galaxies.  The ability to make sensitive CO molecular-line observations enabled,  for the first time, the measurement of total  gas surface densities ($\\Sigma_{HI + H_2}$) in external galaxies while advancements in infrared and ultraviolet observations  led to significant improvements in the measurements of star formation rates. Significant effort was then expended by a number of researchers to systematically measure star formation rates and total gas surface densities in increasingly large samples of galaxies (e.g., Kennicutt 1989 and references therein). These efforts culminated in the study of Kennicutt (1998a) who compiled galaxy averaged measurements of star formation rates and gas surface densities for a large sample of galaxies including normal spirals and starbursts. He derived a scaling relation between the star formation rate surface density ($\\Sigma_{SFR}$) and total gas surface density ($\\Sigma_{HI + H_2}$) that was characterized by a power-law index of n $\\approx$ 1.4. This value was shallower than that Schmidt and others found for individual galaxies using only atomic gas but still super-linear.  Wong and Blitz ( 2002), employing spatially resolved observations of seven nearby, molecular rich, spiral galaxies, showed that the star formation rate was better correlated with the molecular hydrogen surface density, $\\Sigma_{H_2}$, than with the atomic surface density, but still obtained n $\\approx$ 1.4. More recently, Bigiel et al. (2008) analyzed spatially resolved observations of 18 nearby galaxies containing both atomic rich and molecular rich objects and confirmed that $\\Sigma_{SFR}$ was better correlated with $\\Sigma_{H_2}$ than $\\Sigma_{HI}$, but they determined that n $=$ 1.0 ($\\pm$ 0.2)  for the $\\Sigma_{SFR}$ -- $\\Sigma_{H_2}$ relation.  However, recent observations of M 101 and  M 81 have suggested that the index of the scaling law can vary within a galaxy with values of n ranging between 1 and 2 (Suzuki et al. 2010).  Among the more interesting investigations of the extragalactic scaling laws for star formation was that of Gao and Solomon (2004) who used molecular-line emission from HCN, rather than CO, to trace the molecular gas. They found a linear (n $=$ 1) correlation between the total far-infrared luminosities and the HCN molecular-line luminosities of a large sample of star forming galaxies including normal spirals and luminous and ultra-luminous infrared galaxies.  Since the total infrared luminosity is a good proxy for the total star formation rate (SFR) and the HCN luminosity a good proxy for the total amount of dense (i.e., n(H$_2$) $\\geq$ 3 $\\times$ 10$^4$ cm$^{-3}$) gas in a galaxy, this also implied a linear correlation between the SFR and the mass of dense molecular gas. %  The various determinations of differing power-law indices for the extragalactic star formation scaling relations present a somewhat confused and problematic picture. Particularly since the difference between a linear and non-linear scaling relation can have significant consequences for the theoretical understanding of the star formation process in galaxies.  Therefore it is important to understand the nature of such differences. Are the different scaling relations consistent with each other? Are the differences due to such effects as the choice of the samples studied (e.g., normal spirals vs starbursts, CO rich vs. HI rich galaxies, distant vs. nearby systems, etc.) or the different quantities actually measured (e.g., SFR vs.  $\\Sigma_{SFR}$, $\\Sigma_{HI + H_2}$ vs. $\\Sigma_{H_2}$, or CO vs. HCN, etc.), or the systematic uncertainties in the quantities measured (e.g., observational tracers or IMFs adopted for SFR determinations, conversion factor for transforming CO measurements into H$_2$ masses, etc.), or some linear combination of all these effects? Do any of these scaling relations represent the fundamental underlying physical relationship that most directly connects star formation activity with interstellar gas?  Schmidt's original scaling law was determined from observations of the local region of the Galaxy. Since our knowledge of the local Milky Way has improved profoundly over the last half century, it would seem that important insights into the relation between star formation and interstellar gas could and should be derived from observations of local star formation activity. In a previous paper (Lada et al. 2010; hereafter Paper I) we presented a study of the star formation activity in a sample of local (d $<$ 0.5 kpc) molecular clouds with total masses between 10$^3$ and 10$^5$ \\msun. We employed infrared extinction measurements derived from wide-field surveys to determine accurate cloud masses and mass surface densities, and compiled from the literature both ground and space-based infrared surveys of young stellar objects to construct complete inventories of star formation within the clouds of our local sample. We found the specific star formation rates (i.e., the star formation rates per unit cloud mass) in these clouds to vary by an order of magnitude, independent of total cloud mass. However, we also found the dispersion in the specific star formation rate,  to be minimized (and reduced by a factor of 2-3) if one considers only the mass of molecular gas characterized by high extinction in calculating the specific star formation rates. As a result we showed that the (total) star formation rate  in local clouds is linearly proportional to the cloud mass contained above an extinction threshold of A$_K$ $\\geq$ 0.8 magnitudes, corresponding to a gas surface density threshold of $\\Sigma_{H_2}$ $\\approx$ 116 \\msun pc$^{-2}$. Similar surface density thresholds for star formation in local clouds have been suggested in other recent studies (e.g., Goldsmith et al 2008; Heiderman et al. 2010). Given the density stratification of molecular clouds, we argued that such surface density thresholds also correspond to volume density thresholds of n(H$_2$) $\\approx$ 10$^4$ cm$^{-3}$. These findings are consistent with and reinforce those of Wu et al. (2005) who had already demonstrated a linear correlation between far-infrared luminosity and HCN luminosity (i.e., between SFR and dense gas mass) for more massive and distant star formation regions in the Milky Way.  The correspondence between these results and those obtained by Gao and Solomon (2004) for external galaxies is intriguing and especially striking because the scalings of the Galactic and extragalactic power-law relations, that together span more than nine orders of magnitude in cloud mass, agree to within a factor of 2-3. This suggested to us that the close relationship between the star formation rates and the {\\it dense} gas masses of molecular clouds could be the underlying physical relation that connects star formation activity with interstellar gas over vast spatial scales from the immediate vicinity of the sun to the most distant galaxies.  However, if this is so, how does one understand these observations in the context of the classical Schmidt-Kennicutt scaling relations based on CO observations?  These classical relations are often super linear and moreover, as Heiderman et al. (2010) point out, they  under predict the $\\Sigma_{SFR}$ in local regions by factors of 17 - 50 (see also Evans et al. 2009). In this paper we  attempt to address this issue by re-examining the extinction observations of local clouds to include low extinction material and re-examining the CO observations of the clouds studied by Gao and Solomon. We show that all the observations can be understood within a self-consistent framework in which the differences are primarily due to the dense gas fractions that characterize the molecular gas being observed, supporting a hypothesis originally put forward  by Gao and Solomon (2004). ",
        "conclusions": "The SFR-Molecular Mass diagram of Figure 2 provides a physical context for understanding the star formation scaling laws over spatial scales ranging from those of local molecular clouds to those of entire galaxies. The  close correlation of the star formation rate with the mass of dense gas over these immense scales has been established in previous studies (Wu et al. 2005, Paper I).  Here we find that a close relation also appears to hold between the SFR and the total molecular mass over a similarly large range, 8-9 orders of magnitude in both quantities.  Both the local clouds and galaxies appear to scatter around the linear relation given by Equation 1 for $f_{DG}$ $=$ 0.1 and $M_G = M_{TG}$.  From extrapolation of the results for local clouds we suggest that this particular line corresponds to the case where 10\\% of the measured gas mass is in the form of dense, star forming material for the galaxies as well as for the local clouds. The smaller scatter of the galaxies around this relation compared to that of the local clouds is likely the result of the fact that the galaxy measurements are averages over entire systems.  These results indicate that, similar to the situation for dense gas, the star formation scaling law for total (H$_2$ + He) gas mass is likely linear across all scales for molecular clouds with similar dense gas fractions.  This notion is reinforced by the recent observations of Daddi et al. (2010) who studied infrared-selected BzK galaxies at $z\\sim1.5$ and found evidence for unusually high gas fractions and extended molecular reservoirs in these distant systems. Using the star formation rates and CO gas masses provided by Daddi et al. (2010), we plot these six galaxies (open triangles) on Figure 2 and find that the BzK galaxies occupy an area in the SFR-Molecular Mass plot that is close to the linear relation described by Equation 1, consistent with the locations of Gao and Solomon galaxies and the extrapolation of the local Galactic cloud sample.  These results lead us to the conclusion that there is a basic and universal physical process of star formation that presently operates in our Milky Way galaxy and is also responsible for the bulk of star forming activity occurring in external galaxies both in the present epoch (z $\\approx$ 0; GS04) and perhaps at much earlier (z $\\approx 1-2$; Daddi et al. 2010) cosmic times. It is a process in which the rate of star formation is simply and directly controlled by the amount of dense molecular gas that can be assembled within a star forming complex. In most situations massive molecular clouds appear to be able to convert only about 10\\% or less of their total mass into a sufficiently dense (n(H$_2$) $\\ge$ 10$^4$ \\cc) form to actively produce stars. This may be considered as the normal process of star formation in GMCs.  Closer inspection of Figure 2 suggests that for starburst galaxies, particularly the ULIRGS, this standard scenario may be modified. As the SFRs for starbursts (i.e., LIRGs and ULIRGs in Figure 2) increase with gas mass, the open symbols (CO derived gas masses) appear to approach and then merge with the filled symbols (HCN derived gas masses), almost  overlapping at the highest SFRs. As originally hypothesized by Gao and Solomon (2004), we interpret this to indicate that these galaxies are characterized by an increasingly  high dense gas fraction and consequently, the CO observations begin to trace nearly the same material as the HCN observations. Nonetheless, the star formation rate is still dictated by the amount of dense gas within the galaxies. This interpretation is also favored by Heiderman et al. (2010) who suggested that the maximal starburst activity occurs when $f_{DG} = \\rm 1$ which they posit to happen when the mass surface density exceeds values $\\sim$ 10$^4$ \\msun\\ pc$^{-2}$. ULIRGS (e.g.,  Arp 200)  are believed to be experiencing major mergers and we suggest that this extreme process likely produces conditions (e.g., high pressures) that could increase the dense gas fractions of the molecular clouds within these systems (e.g., Blitz \\& Rosolowsky 2006).  In contrast the BzK galaxies studied by Daddi et al. (2010) have similarly high SFRs but lower dense gas fractions. Their high star formation rates appear to result from high global molecular gas mass fractions (i.e., M$_{H_2}$/M$_*$), as might be expected for very young galaxies.  We note that a linear relation in the SFR-Mass plane should transform to a linear relation in the $\\Sigma_{SFR}$-$\\Sigma_{g}$ plane (provided the surface densities for the galaxies are global averages) and we can express our star formation scaling law in this latter plane as:  \\begin{equation} \\Sigma_{SFR} \\propto f_{DG}\\Sigma_{g} \\end{equation}  \\noindent where $\\Sigma_{g}$ refers to the H$_2$ gas mass. Moreover, the Spitzer study of Galactic clouds by Heiderman et al. (2010) suggested a linear star formation law in the $\\Sigma_{SFR}$-$\\Sigma_{g}$ plane that holds for gas above a threshold surface density of $\\sim$ 130 \\msun\\ pc$^{-2}$ (i.e., A$_K$ $>$ 0.9 mag) and extrapolates smoothly to the GS04 galaxies.  Our result  is apparently not consistent with the standard Schmidt-Kennicutt, super-linear, scaling law (Kennicutt 1998a). Both are based on valid empirical relations. However, here we argue that the underlying scaling law for star formation is linear over all scales for all clouds and galaxies, provided they are characterized by the same dense gas fraction. Kennicutt (1998a) uses total (HI $+$ H$_2$) gas mass surface densities with CO derived  molecular masses and combines results for both normal star-forming disk galaxies and starburst galaxies to derive his star formation scaling law. Note that for these latter galaxies the total gas surface densities are dominated by the molecular component. The starbursts dominate his relation for $\\Sigma_{gas}$ $>$ 100 \\msun pc$^{-2}$. It is possible that the fit of a single relation to the combined sample with CO determined masses is inappropriate and skewed by the starbursts because $f_{DG}$ for starbursts is higher than that for normal star forming spirals. Indeed, Gao and Solomon (2004) showed that using the masses calculated from the CO observations  produced a super-linear (n $\\approx$ 1.7) scaling law (in the SFR vs M$_G$ plane) for a sample that included their galaxies and an additional number of luminous starbursts drawn form the literature. Using gas masses derived solely from HCN observations, however, produces a linear star formation law connecting both normal star forming galaxies and starbursts. The standard Schmidt-Kennicutt relation may also be skewed at low mass surface densities. For galaxies in this portion of the diagram, the HI surface density is a large fraction of the total gas surface density and thus the measured total gas surface density, $\\Sigma_{HI + H_2}$, contains a large component of inert, non-star forming, (HI) gas; this dilutes and lowers the SFR corresponding to a fixed mass surface density, resulting in a steepening of the slope of the $\\Sigma_{SFR}$ vs $\\Sigma_{gas}$ relation.  These two effects, the increasing dense gas fraction for the starbursts and the dilution of the dense gas fraction by HI gas at low gas surface densities, which act together to steepen the slope of the Schmidt-Kennicutt relation, can also explain the finding of Heiderman et al. (2010) and Evans et al. (2009) that the extrapolation of the extragalactic scaling relations to local scales (i.e., low mass surface densities) lies below the data for Galactic clouds.  It can also be shown that our scaling law (Equation 1) is consistent with a volumetric scaling law, $\\dot{\\rho_*}  \\propto \\rho_G^n$ if and only if $n = 1$ and $\\rho_G \\geq \\rho_{thres}$, where  $\\rho_{thres}/\\mu$ corresponds to the threshold volumetric number density for star formation  for a mean particle mass given by $\\mu$ (i.e., n(H$_2$) $\\geq$ 10$^4$ cm$^{-3}$).  As discussed earlier, taking the empirical, linear star-forming scaling relations at face value leads to a simple interpretation of the observations in Figures 1 \\& 2. Namely, that the total rate of star formation, $\\dot{M}_*$, is directly proportional to the mass of dense molecular gas above a threshold density, $M_{DG} = \\int_{\\rho_{thres}}^\\infty M(\\rho) d\\rho$.  Moreover, once the gas has reached this threshold density, the SFR does not depend on the exact value of the density but only on the total mass of gas whose density has exceeded the threshold. This interpretation of the observations differs from those that explain the observed non-linear index of the Schmidt-Kennicutt law as resulting from star formation timescales dictated by the free-fall time, e.g.,  SFR $\\sim$ M/$\\tau_{ff}$ $\\sim$ $\\rho/ \\rho^{-0.5}$ $\\sim$ $\\rho^{1.5}$ since $\\tau_{ff} \\sim \\rho^{-0.5}$ (e.g., Elmegreen 1994; Krumholz \\& Thompson 2007; Narayanan et al. 2008). A recent variant of such a model has been studied by Krumholz et al. (2011). They propose that the underlying physical law governing the relation between star formation rates and cloud properties is given by  $\\dot{\\rho_*}  \\propto \\rho_G/ \\tau_{ff}$. They find that the standard Schmidt-Kennicutt law can be linearized if the data are plotted in the $\\Sigma_{SFR}-\\Sigma_G/\\tau_{ff}$ plane as long as the free-fall times are measured using the typically higher densities of the star forming gas rather than those derived from the mean densities averaged over kpc scales. Their interpretation differs from the one in this paper in that Krumholz et al. (2011) posit that the positions of galaxies in the standard $\\Sigma_{SFR}$-$\\Sigma_{G}$ plane are a consequence of both their gas surface densities and their local free-fall times, while here we posit that the locations of these galaxies instead depend on their gas surface densities and their dense gas fractions. Although both interpretations are consistent with the observations, they appear not to be consistent with each other. However, we point out that Figure 1 empirically demonstrates that the locations of Galactic clouds in the SFR-Mass diagram are in fact a result of their dense gas fractions. Therefore it seems reasonable to infer that the locations of galaxies in the diagram are due to the same cause.  Finally, we reiterate  our point that the linear scaling law of Equation 1 implies that the process of star formation across entire galaxies as well as individual local clouds is governed by a very similar and simple physical principle: the rate at which molecular gas is turned into stars depends on the mass of dense gas within a molecular cloud or cloud population. The underlying star formation scaling law is linear over all scales for all clouds and galaxies characterized by  the same dense gas fraction. The star formation rate appears therefore to be controlled by local processes and not by global, galactic scale mechanisms, except to the extent that such mechanisms can alter the dense gas fractions in the molecular gas. If this interpretation is correct, then the key problem that needs to be addressed in future studies is that of the origin of the dense gas component of molecular clouds."
    },
    "1112/1112.6149_arXiv.txt": {
        "abstract": "We consider inflationary models in which vector fields are responsible for part or eventually all of the primordial curvature perturbation $\\zeta$.  Such models are phenomenologically interesting since they naturally introduce anisotropies in the probability distribution function of the primordial fluctuations that can leave a measurable imprint in the cosmic microwave background. Assuming that non-Gaussianity is generated due to the superhorizon evolution, we use the $\\delta N$ formalism to do a complete tree level calculation of the % non-Gaussianity parameters $f_{\\rm NL}$ and $\\tau_{\\rm NL}$ in the presence of vector fields. % We isolate the isotropic pieces of the non-Gaussianity parameters, which anyway have contributions from the vector fields, and show that they obey the Suyama-Yamaguchi consistency relation $\\tau^{\\rm iso}_{\\rm NL}\\geqslant(\\frac{6}{5}f^{\\rm iso}_{\\rm NL})^2$. Other ways of defining the non-Gaussianity parameters, which could be observationally relevant, are stated and the respective Suyama-Yamaguchi-like consistency relations are obtained. ",
        "introduction": "The study of non-Gaussianities in the primordial curvature perturbation $\\zeta$ is a subject of major interest in modern cosmology because the evaluation of the non-Gaussianity (NG) parameters provide criteria to discriminate among the many models proposed to explain the origin of the large-scale structure that we observe today\\cite{obs,obs1,obs2}. Particular attention has been paid to the study of a consistency relation between the $f_{\\rm NL}$ and $\\tau_{\\rm NL}$ parameters\\cite{cr10,cr11,cr12,cr13,cr14,cr20,cr21}, the Suyama-Yamaguchi (SY) consistency relation ($\\tau_{\\rm NL} \\geqslant (\\frac{6}{5} f_{\\rm NL})^2$ at least at tree level), because its violation would rule out many of the popular models of inflation. So far, it has been shown that this consistency relation works, in principle, on  models that only include scalar fields\\cite{cr10,cr11,cr12,cr13,cr14} under just a few assumptions, but there is no any conclusive result for models that include other type of fields, for instance, vector fields (VF)\\footnote{VF naturally generate scale-dependence and, therefore, the proofs in Ref. \\cite{cr20,cr21} do not apply to those scenarios where VF contribute to the generation of the primordial curvature perturbation.}.  Vector field models are particularly interesting because they are suitable candidates to explain the apparent violation of statistical isotropy observed in recent analysis of data from the WMAP satellite\\cite{gr1,gr2}, given that VF define inherently a preferred direction for the expansion, for the distribution of primordial fluctuations, or for both\\cite{vi1,vi2,vi3,vi4,vi5,vi6,ValenzuelaToledo:2011fj,dklr}. Because of this reason, it is pertinent to study consistency relations in models that include scalar fields as well as vector fields as these not only include NG parameters but also the amount of primordial statistical anisotropy $g_\\zeta$ \\cite{varios1,varios2,varios3,varios4}.   The main purpose of this paper is to find out how the well known SY consistency relation, valid for models involving only scalar fields under just a few assumptions, is modified when including VF. To this end, we do a complete tree level calculation of the $f_{\\rm NL}$ and $\\tau_{\\rm NL}$ parameters based on the $\\delta N$ formalism\\cite{dklr}, and concentrate on the isotropic pieces of these parameters since these are the ones that are currently constrained by observations.  Under almost the same few assumptions made in the scalar fields case, and although the VF do contribute to the isotropic pieces of the NG parameters, we prove that the SY consistency relation is still obeyed, at least at tree level. Finally, the modified SY consistency relations for other ways of defining the NG parameters, which could be observationally relevant when taking into account the level of statistical anisotropy, are obtained. ",
        "conclusions": "In this paper, we have studied under which conditions new kinds of the well known Sayama-Yamaguchi consistency relation are allowed  when vector fields are included in the inflationary dynamics. In particular we have shown that the isotropic pieces of the non-gaussianity parameters obey the Suyama-Yamaguchi consistency relation $\\tau^{\\rm iso}_{\\rm NL}\\geqslant(\\frac{6}{5}f^{\\rm iso}_{\\rm NL})^2$ when its calculation is performed at tree level.  We have also derived a modified consistency relation for the complete non-gaussianity parameters $f_{\\rm NL}$ and $\\tau_{\\rm NL}$ and for a set of new parameters $f'_{\\rm NL}$ and $\\tau'_{\\rm NL}$ that could be observationally relevant; these relations are given by the Eqs. \\eqref{cr1} and \\eqref{cr2}. Since the latter depend on the level of statistical anisotropy and the configuration of the wavevectors, the naive relations $\\tau_{\\rm NL}\\geqslant(\\frac{6}{5}f_{\\rm NL})^2$ and $\\tau'_{\\rm NL}\\geqslant(\\frac{6}{5}f'_{\\rm NL})^2$ could eventually be violated in some particular inflationary model that includes vector fields. On this basis, we expect that in a near future the derived relations could provide a criteria to discriminate among the different models for the generation of the primordial curvature perturbation when vector fields are involved."
    },
    "1112/1112.2076_arXiv.txt": {
        "abstract": "\\begin{center} {\\bf Abstract} \\end{center} We calculate the $^4$He abundance in a universe of Bianchi type I whose cosmic anisotropy is dynamically generated by a fluid with anisotropic equation of state. Requiring that the relative variation of mass fraction of $^4$He is less than $4\\%$ with respect to the standard isotropic case to be consistent with astrophysical data, we constrain the parameter of cosmic anisotropy, the shear $\\Sigma$, as $|\\Sigma(T_f)| \\lesssim 0.4$, where $T_f$ is the freeze-out temperature of the weak interactions that interconvert neutrons and protons. Anisotropic fluids, whose energy density is subdominant with respect to the energy content of the Universe during inflation and radiation era, generate much smaller shears at the time of freeze-out and then do not appreciably affect the standard $^4$He production. This is the case of anisotropic dark energy, and of a uniform magnetic field with energy density much smaller than about $1.25$ times the energy density of neutrinos. ",
        "introduction": "} \\renewcommand{\\thesection}{\\arabic{section}}  The high level of isotropy of the cosmic microwave background (CMB) radiation is the most convincing justification of the Cosmological Principle: the Universe is homogeneous and isotropic at large cosmological scales~\\cite{Weinberg}. However, tiny deviations from perfect isotropy are not excluded by present CMB data. Indeed, a particular anisotropic cosmological model of Bianchi type I, known as {\\it ellipsoidal universe}~\\cite{ellipsoidal1,ellipsoidal2}, can better match CMB data and solve the so-called ``quadrupole problem,'' namely the lack of CMB power detected on large angular scales.  Various mechanisms could give rise to an ellipsoidal universe, such as a uniform cosmological magnetic field~\\cite{Barrow1,ellipsoidal1,ellipsoidal2}, topological defects (e.g. cosmic stings, domain walls)~\\cite{Barrow1}, or a dark energy fluid with anisotropic equation of state~\\cite{Barrow1,Koivisto-Mota}. Independently on the nature of the mechanism, however, a modification of the standard picture of primordial nucleosynthesis can occur if universe anisotropization takes place during the early Universe~\\cite{R1,R2,R3,R4}.  The aim of this paper is, indeed, to constrain the level of cosmic anisotropy, so as to be consistent with observational bounds on primeval $^4$He abundance. ",
        "conclusions": "} \\renewcommand{\\thesection}{\\arabic{section}}  In this paper, we have analyzed the effects caused by cosmic anisotropy on the primordial production of $^4$He. We worked in the context of a cosmological model of Bianchi type I where the anisotropy of spatial geometry, the shear $\\Sigma$, is generated by a fluid with anisotropic equation of state.  We found that in such an anisotropic universe there is an overproduction of $^4$He with respect to the standard isotropic case. Imposing that the relative increase of $^4$He abundance is below the $4\\%$ to be consistent with observational data, we constrained the absolute value of the shear to be less than $0.4$ at the time of freeze-out. This limit does not depend on the equation(s) of state of the anisotropic fluid and has been obtained assuming that the energy density of the anisotropic fluid is small compared to that of radiation.  Moreover, we showed  that anisotropic fluids generated at inflation, such as dark energy with anisotropic equation of state and a uniform magnetic field, create anisotropies much smaller than the above limit if their energy densities are subdominant with respect to that of the Universe during inflation and radiation era. In particular, the existence of a uniform magnetic field at the time of nucleosynthesis is compatible with astrophysical data if its energy density is much smaller than about $1.25$ times the energy density of neutrinos."
    },
    "1112/1112.0303_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}  \\subsection{Background and Motivation}  The recent discovery of the accelerating expansion of the Universe \\cite{Perlmutter1999, Riess1998} has prompted many theoretical speculations about the underlying mechanism. The most likely mechanism is a cosmological constant, which is the simplest model and is in good agreement with observational data \\cite{Blake2011}.  More complicated models involve new dynamical sources of gravity that act as dark energy, and/or modifications to general relativity on large scales. A plethora of models have been postulated and explored in recent years, including Quintessence, K-essence \\cite{Armendariz-Picon2000, Chiba2000}, Ghost Condensates \\cite{Arkani-Hamed2004}, DGP gravity \\cite{Dvali2000}, and $f(R)$ gravity, to name but a few.  See Refs.\\ \\cite{Peebles2003, Nobbenhuis2006, Copeland2006, Caldwell:2009ix,Silvestri:2009hh,Skordis2011,Amendola2010} for detailed reviews of these and other models.  A common feature of the majority of dark energy and modified gravity models is that in the low energy limit, they are equivalent to general relativity coupled to one or more scalar fields, often called quintessence fields. Therefore it is useful to try to construct very general low energy effective quantum field theories of general relativity coupled to light scalar fields, in order to encompass broad classes of dark energy models. Considering dark energy models as quantum field theories is useful, even though the dynamics of dark energy is likely in a classical regime, because it facilitates discriminating against theories which are theoretically inconsistent or require fine tuning.    A similar situation occurs in the study of models of inflation, where it is useful to construct generic theories using effective field theory. Cheung \\textit{et al.} \\cite{Cheung2008} constructed a general effective field theory for gravity and a single inflaton field, for perturbations about a background Friedman-Robertson-Walker cosmology in unitary gauge.  This work was later generalized in multiple directions \\cite{Senatore:2010wk,2011arXiv1106.2189F} and has been very useful. An alternative approach to single field inflationary models was taken by Weinberg \\cite{Weinberg2008}, who constructed an effective field theory to describe both the background cosmology and the perturbations.  This theory consisted at leading order of a standard single field inflationary model with a potential, together with higher order terms in a covariant derivative expansion up to four derivatives.  More detailed discussions of this type of effective field theory were given by Burgess, Lee and Trott \\cite{Burgess:2009ea}.   When one turns from inflationary effective field theories to quintessence effective field theories, the essential physics is very similar, but there are three important differences that arise:  \\begin{itemize}  \\item First, the hierarchy of scales is vastly more extreme in quintessence models.  The Hubble parameter $H$ is typically several orders of magnitude below the Planck scale $\\mpl \\sim 10^{28}$ eV in inflationary models, whereas for quintessence models $H_0 \\sim 10^{-33} \\, $eV is $\\sim 60$ orders of magnitude below the Planck scale.  Quintessence fields must have a mass that is smaller than or on the order of $H_0$.  It is a well-known, generic challenge for quintessence models to ensure that loop effects do not give rise to a mass much larger than $H_0$.  Because of the disparity of scales, this issue is more extreme for quintessence models than inflationary models.  \\item In most inflationary models, it is assumed that the dynamics of the Universe are dominated by gravity and the scalar field (at least until reheating).  By contrast, for quintessence models in the regime of low redshifts relevant to observations, we know that cold dark matter gives an $O(1)$ contribution to the energy density.  Therefore there are additional possible couplings and terms that must be included in an effective field theory.  \\item For any effective field theory, it is possible to pass outside the domain of validity of the theory even at energies $E$ low compared to the theory's cutoff $\\Lambda$, if the mode occupation numbers $N$ are sufficiently large (see Sec.\\ \\ref{sec:validity} below for more details).  This corresponds to a breakdown of the classical derivative expansion. For quintessence theories, mode occupation numbers today can be as large as $N \\sim (\\mpl/H_0)^2$ and it is possible to pass outside the domain of validity of the theory.  By contrast in inflationary models, this is less likely to occur since mode occupation numbers for the perturbations are not large before modes exit the horizon. Thus, the effective field theory framework is less all-encompassing for quintessence models than for inflation models. This issue seems not to have been appreciated in the literature and we discuss it in Sec. \\ref{sec:validity} below.  \\end{itemize}     Several studies have been made of generic effective field theories of dark energy.  Creminelli, D'Amico, Nore\\~{n}a and Vernizzi \\cite{Creminelli2009} constructed a the general effective theory of single-field quintessence for perturbations about an arbitrary FRW background, paralleling the similar construction for inflation \\cite{Cheung2008}.  Park, Watson and Zurek constructed an effective theory for describing both the background cosmology and the perturbations, following the approach of Weinberg \\cite{Weinberg2008} but generalizing it to include couplings to matter \\cite{Watson2010}.  The two approaches to effective field theories of quintessence -- specialization to perturbations about a specific background, and maintaining covariance and the ability to describe the dynamics of a variety of backgrounds -- are complementary to one another. The dynamics of the cosmological background FRW solution can be addressed in the covariant approach of Weinberg, but not in the background specific approach of Creminelli et al., which restricts attention to the dynamics of perturbations about a given, fixed background. On the other hand, a background specific approach can describe a larger set of dynamical theories for the perturbations than can a covariant derivative expansion\\footnote{To see this, consider for example a term in the Lagrangian of the form $f(\\phi) (\\nabla \\phi)^{2 n}$, where $\\phi$ is the quintessence field.  Such a term would be omitted in the covariant derivative expansion for sufficiently large $n$.  However, upon expanding this term using $\\phi = \\phi_0 + \\delta \\phi$, where $\\phi_0$ is the background solution, one finds terms $\\sim (\\nabla \\phi_0)^{2n-2} (\\nabla \\delta \\phi)^2$ which are included in the Creminelli approach of applying standard effective field theory methods to the perturbations.}.    \\subsection{Approach and Assumptions} \\label{sec:approach}  The purpose of this paper is to revisit, generalize and correct slightly the covariant effective field theory analysis of Park, Watson and Zurek \\cite{Watson2010}. Following Weinberg and Park et al., we restrict attention to theories where the only dynamical degrees of freedom are a graviton and a single scalar.  We allow couplings to an arbitrary matter sector, but we assume the validity of the weak equivalence principle, motivated by the strong experimental evidence for this principle. We assume that the theory consists of a standard quintessence theory coupled to matter at leading order in a derivative expansion, with an action of the form \\begin{align} S[g_{\\alpha\\beta},\\phi,\\psi_{\\rm m}] = \\int d^4 x \\sqrt{-g} \\left\\{ \\frac{\\mpl^2}{2} R - \\frac{1}{2} (\\nabla \\phi)^2 - U(\\phi) \\right\\} + S_{\\rm m} \\left[ e^{\\alpha(\\phi)} g_{\\mu \\nu},  \\psi_{\\rm m}\\right]. \\label{eq:intro:leadingorder} \\end{align} Here $\\psi_{\\rm m}$ denotes a set of matter fields, and $\\mpl$ is the Planck mass. The factor $e^{\\alpha(\\phi)}$ in the matter action provides a leading-order non-minimal coupling of the quintessence field to matter, in a manner similar to Brans-Dicke models in the Einstein frame.  Our analysis then consists of a series of steps:  \\begin{enumerate}  \\item We add to the action all possible terms involving the scalar field and metric, in a covariant derivative expansion up to four derivatives. We truncate the expansion at four derivatives, as this is sufficient to yield the leading corrections to the action (\\ref{eq:intro:leadingorder}).  As described by Weinberg \\cite{Weinberg2008} there are ten possible terms, with coefficients that can be arbitrary functions of $\\phi$ [see Eq.\\ (\\ref{eq:S1}) below]. Section \\ref{sec:scaling} below describes one possible justification of this covariant derivative expansion from an effective field theory viewpoint, starting from a set of ultralight pseudo Nambu-Goldstone bosons (PNGBs).  It is likely that the same expansion can be obtained from other, more general starting points.   \\item We allow for corrections to the coupling to matter by adding to the metric that appears in the matter action all possible terms involving the metric and $\\phi$ allowed by the derivative expansion, that is, up to two derivatives.  There are six such terms [see Eq.\\ (\\ref{eq:JordanMetric}) below.] We also add to the action terms involving the stress energy tensor $T_{\\mu\\nu}$ of the matter fields, up to the order allowed by the derivative expansion using $T_{\\mu\\nu} \\sim \\mpl^2 G_{\\mu\\nu}$ [see Eq.\\ (\\ref{eq:S1}) below]. Including such terms in the action seems poorly motivated, since {\\it a priori} there is no reason to expect that the resulting theory would respect the weak equivalence principle (see Appendix \\ref{app:WEP}). However we show in Appendix \\ref{app:WEP} that the weak equivalence principle is actually satisfied, to the order we are working to in the derivative expansion.  In addition, all the terms in the action involving $T_{\\mu\\nu}$ can be shown to have equivalent representations not involving the stress energy tensor, using field redefinitions (see Appendix \\ref{app:WEP}).  \\item The various correction terms are not all independent because of the freedom to perform field redefinitions involving $\\phi$, $g_{\\mu\\nu}$ and the matter fields, again in a derivative expansion.  In Sec.\\ \\ref{sec:techs} we explore the space of such field redefinitions, finding eleven independent transformations and tabulating their effects on the coefficients in the action (see Table \\ref{tab:transformations} below).   \\item Several of the correction terms that are obtained from the derivative expansion are ``higher derivative'' terms, by which we mean that they give contributions to the equations of motion which involve third-order or higher-order time derivatives of the fields\\footnote{The precise definition of higher derivative that we use, which is covariant, is that an equation will be said not to contain any higher derivative terms if there exists a choice of foliation of spacetime for which any third-order or higher-order derivatives contain at most two time derivatives. Theories which are higher derivative in this sense are generically associated with instabilities (Ostragradski's theorem) \\cite{Woodard2007}, although the instabilities can be evaded in special cases, for example $f(R)$ gravity.  For most of this paper (except for the Chern-Simons term), a simpler definition of higher derivative would be sufficient: a term in the action is ``higher derivative'' if it gives rise to terms in the equation of motion that involve any third-order or higher order derivatives.}. Normally, such higher derivative terms give rise to additional degrees of freedom. However, if they are treated perturbatively (consistent with our derivative expansion) additional degrees of freedom do not arise.  Specifically, one can perform a {\\it reduction of order} procedure on the equations of motion \\cite{1971ctf..book.....L,Parker:1993dk,Flanagan:1996gw}, substituting the zeroth-order equations of motion into the higher derivative terms in the equations of motion to eliminate the higher derivatives\\footnote{This is more general than requiring the solutions of the equation of motion to be analytic in the expansion parameter, as advocated by Simon \\protect{\\cite{Simon1990}}; see Ref.\\ \\protect{\\cite{Flanagan:1996gw}}.}. We actually use a slightly different but equivalent procedure of eliminating the higher derivative terms directly in the action using field redefinitions\\footnote{This procedure is counterintuitive since normally field redefinitions do not change the physical content of a theory; here however they do because the field redefinitions themselves involve higher derivatives.} (see Appendix \\ref{app:backsubs}).   Weinberg \\cite{Weinberg2008} and Park {\\it et al.} \\cite{Watson2010} use a slightly different method, consisting of substituting the leading order equations of motion directly into the higher derivative terms in the action.  This method is not generally valid, but it is valid up to field redefinitions that do not involve higher derivatives, and so it suffices for the purpose of attempting to classify general theories of dark energy (see Appendix \\ref{app:backsubs}).    \\item Another issue that arises with respect to the higher derivative terms is the following.  Is it really necessary to include such terms in an action when trying to write down the most general theory of gravity and a scalar field, in a derivative expansion? Weinberg \\cite{Weinberg2008} suggested that  perhaps a more general class of theories is generated by including these terms and performing a reduction of order procedure on them, rather than by omitting them. However, since it is ultimately possible to obtain a theory that is perturbatively equivalent to the higher derivative theory, and which has second order equations of motion, it should be possible just to write down the action for this reduced theory.  In other words, an equivalent class of theories should be obtained simply by omitting all the higher derivative terms from the start.  We show explicitly in Sec.\\ \\ref{sec:details} that this is the case for the class of theories considered here.   \\item We fix the remaining field redefinition freedom by choosing a ``gauge'' in field space, thus fixing the action uniquely (see Sec. \\ref{sec:details:final}).     \\end{enumerate}  \\subsection{Results and Implications} Our final action is [Eq.\\ (\\ref{eq:details:finalaction}) below] \\begin{align} S =&  \\int d^4 x \\sqrt{-g} \\left\\{ \\frac{\\mpl^2}{2} R - \\frac{1}{2} (\\nabla \\phi)^2 - U(\\phi)\\right\\} + S_{\\rm m}[e^{\\alpha(\\phi)} g_{\\alpha\\beta},\\psi_{\\rm m}] \\nonumber \\\\ & + \\epsilon \\int d^4 x \\sqrt{-g}  \\bigg\\{a_1 (\\nabla \\phi)^4 +  b_2 T (\\nabla \\phi)^2 +  c_1 G^{\\mu \\nu} \\nabla_\\mu \\phi \\nabla_\\nu \\phi \\nonumber \\\\ & +  d_3 \\left( R^2 - 4 R^{\\mu \\nu} R_{\\mu \\nu} + R_{\\mu \\nu \\sigma \\rho} R^{\\mu \\nu \\sigma \\rho} \\right) +  d_4 \\epsilon^{\\mu \\nu \\lambda \\rho} \\tensor{C}{_{\\mu \\nu}^{\\alpha \\beta}} C_{\\lambda \\rho \\alpha \\beta} + e_1 T^{\\mu \\nu} T_{\\mu \\nu} + e_2 T^2 + \\ldots \\bigg\\}. \\label{eq:details:finalaction0} \\end{align} Here the coefficients $a_1$, $b_2$ etc.\\ of the next-to-leading order terms in the derivative expansion are arbitrary functions of $\\phi$, and the ellipsis $\\ldots$ refers to higher order terms with more than four derivatives. The corresponding equations of motion do not contain any higher derivative terms.  This result generalizes that of Weinberg \\cite{Weinberg2008} to include couplings to matter.     \\begin{figure}[h!] \\centering \\ifx\\dofigures\\undefined \\else \\includegraphics[width=0.55\\textwidth]{graph3.pdf} \\fi \\caption{\\sl The parameter space of fractional density perturbation $\\delta \\rho/\\rho$ for perturbations to the quintessence field, and cutoff scale $M$ for the effective field theory, illustrating the constraint (\\ref{eq:ddd}) on the domain of validity.  Near the boundary of the domain of validity the higher derivative terms in the action are potentially observable, this is labeled the ``interesting regime''.  Further away from the boundary the higher derivative terms are negligible and the theory reduces to a standard quintessence model with a matter coupling.} \\label{fig:scales0} \\end{figure}   We can summarize our key results as follows:  \\begin{itemize}  \\item The most general action contains nine free functions of $\\phi$: $U, \\alpha, a_1, b_2, c_1, d_3, d_4, e_1, e_2$, as compared to the four functions that are needed when matter is not present \\cite{Weinberg2008}.  \\item There are a variety of different forms of the final theory that can be obtained using field redefinitions.  In particular some of the matter-coupling terms in the action can be re-expressed as terms that involve only the quintessence field and metric. Specifically, the term $T (\\nabla \\phi)^2$ term could be eliminated in favor of $\\square \\phi (\\nabla \\phi)^2$, the $(\\nabla \\phi)^4$ could be eliminated in favor of a term $T^{\\mu\\nu} \\nabla_\\mu \\phi \\nabla_\\nu \\phi$, or the $G^{\\mu\\nu} \\nabla_\\mu \\phi \\nabla_\\nu \\phi$ term could be eliminated in favor of a term $T^{\\mu\\nu} \\nabla_\\mu \\phi \\nabla_\\nu \\phi$ (see Sec.\\ \\ref{sec:details:final}).   \\item As mentioned above, one obtains the correct final action if one excludes throughout the calculation all higher derivative terms.   \\item The final theory does contain terms involving the matter stress-energy tensor.  Nevertheless, the weak equivalence principle is still satisfied (see Appendix \\ref{app:WEP}).  It is possible to eliminate the stress-energy terms, but only if we allow higher derivative terms in the action (where it is assumed that the reduction of order procedure will be applied to these higher derivative terms).  Thus, for a fully general theory, one must have either stress-energy terms or higher derivative terms; one cannot eliminate both (see Sec.\\ \\ref{sec:details:final}).    \\item We can estimate how all the coefficients $a_1$ etc. scale with respect to a cutoff scale $M$ for an effective field theory as follows (see Sec.\\ \\ref{sec:scaling}).  We assume that several ultralight scalar fields of mass $\\sim H_0$ arise as pseudo Nambu-Goldstone bosons (PNGBs) from some high energy theory \\cite{1995PhRvL..75.2077F,2003JCAP...07..003A}, and are described by a nonlinear sigma model at low energies.  We then suppose that all but one of the these PNGB fields have masses $M$ that are somewhat larger than $\\sim H_0$, and integrate them out.  This will give rise to a theory of the form discussed above for the single light scalar, where the higher derivative terms are suppressed by powers of $M$.  The scalings for each of the coefficients in the action are summarized in Table \\ref{tab:scalings}.  We find that the fractional corrections to the cosmological dynamics due to the higher derivative terms scale as $H_0^2/M^2$, as one would expect.  \\item Finally, we can use these scalings to estimate the domain of validity of the effective field theory (see Sec.\\ \\ref{sec:validity}).  We find that cosmological perturbations with a density perturbation $\\delta \\rho$ in the quintessence field must have a fractional density perturbation that satisfies \\begin{align} \\frac{\\delta \\rho}{\\rho} \\ll \\frac{M^2}{H_0^2}. \\label{eq:ddd} \\end{align} Thus perturbations can become nonlinear, but only modestly so, if $M$ is close to $H_0$.  The parameter space of fractional density perturbation $\\delta \\rho/\\rho$ and cutoff scale $M$ is illustrated in Fig.\\ \\ref{fig:scales0}.  In addition there is the standard constraint for derivative expansions \\begin{align} E \\ll M \\label{eq:ddd1} \\end{align} where $E^{-1}$ is the length-scale or time-scale for some process. We show in Fig.\\ \\ref{fig:scales1} the two constraints (\\ref{eq:ddd}) and (\\ref{eq:ddd1}) on the two dimensional parameter space of energy $E$ and mode occupation number $N$.   \\end{itemize}          \\begin{figure}[h!] \\centering \\ifx\\dofigures\\undefined \\else \\includegraphics[width=0.80\\textwidth]{graph2.pdf} \\fi \\caption{\\sl The domain of validity of the effective field theory in the two dimensional parameter space of energy $E$ per quantum of a mode of the quintessence field, and mode occupation number $N$. The cutoff scale $M$ must be larger than the Hubble parameter $H_0$ in order that the background cosmology lie within the domain of validity. Perturbation modes on length-scales that are small compared to $H_0^{-1}$ but large compared to $M^{-1}$ can be described, but only if the mode occupation number and fractional density perturbation are sufficiently small.  See Sec.\\ \\ref{sec:validity} for details.} \\label{fig:scales1} \\end{figure}   Finally, in Appendix \\ref{app:compare} we compare our analysis to that of Park, Watson and Zurek \\cite{Watson2010}, who perform a similar computation but in the Jordan frame rather than the Einstein frame (see also Ref.\\ \\cite{Jimenez:2011nn}). The main difference between our analysis and theirs is that they use a different method to estimate the scalings of the coefficients, and as a result their final action differs from ours, being parameterized by three free functions rather than nine.   \\ifx\\doejects\\undefined \\else \\eject \\fi ",
        "conclusions": "\\label{sec:conclusion}   In this paper, we have investigated effective field theory models of cosmic acceleration involving a metric and a single scalar field. The set of theories we considered consists of a standard quintessence model with matter coupling, together with a general covariant derivative expansion, truncated at four derivatives.  We showed that this class of theories can be obtained from a PNGB scenario, where one of the PNGB fields is lighter than all the others, and the heavier fields are integrated out.  We showed that in constructing this class of theories, including higher derivative terms in the action, as suggested by Weinberg \\cite{Weinberg2008}, does not give any increased generality.  We also showed that complete generality requires one to include terms in the action that depend on the stress-energy tensor of the matter fields.  We now turn to a discussion of some of the advantages and shortcomings of the approach adopted here to describe models of dark energy.  Some of the shortcomings are:   \\begin{itemize}  \\item By construction, our approach excludes theories where nonlinear kinetic terms in the action give an order unity contribution to the dynamics, such as K-essence, ghost condensates etc., since such theories do not arise from the PNGB construction used in this paper. On the other hand, such theories are less natural than the class of theories considered here, from the point of view of loop corrections: they require very nontrivial physics at the scale $\\sim H_0$, instead of at the scale $\\sim \\sqrt{H_0 \\mpl}$ required in the PNGB approach. The most general class of theories of this kind is that of Horndeski \\cite{Horndeski1974}, which contains four free functions of $\\phi$ and $(\\nabla \\phi)^2$ \\cite{Skordis2011}, and which is the most general class of theories of a metric and a scalar field for which the equations of motion are second order. As discussed in the Introduction, these theories are included in the alternative, background-dependent approach to effective field theories of quintessence of Creminelli et al. \\cite{Creminelli2009}.    \\item Our class of theories will be observationally distinguishable from vanilla quintessence theories only if the cutoff $M$ is near the Hubble scale $H_0$.  In this regime, our framework cannot be used to analyze Solar System tests of general relativity, since they are outside the domain of validity of the effective field theory. Also, when the background cosmology is evolved backwards in time it passes outside the domain of validity at fairly low redshifts. (This is not a serious disadvantage since dark energy dominates only at low redshifts.)  \\item We have restricted attention to theories with a metric and a single scalar field, with the only symmetry being general covariance.  Thus, our analysis does not include models with several scalar fields, vector fields etc.  In addition, our analysis excludes an interesting class of models that one obtains by imposing that the action be invariant under $\\phi \\to f(\\phi)$, where $f$ is any monotonic function, as such a symmetry cannot be realized with our derivative expansion.  This class of models includes Horava-Lifshitz gravity and has the same number of physical degrees of freedom as general relativity \\cite{Skordis2011,Blas:2010hb}. It would be interesting to explore the most general dark energy models of this kind.  \\end{itemize}  Some of the advantages of the approach used here are:  \\begin{itemize}  \\item Our class of theories is generic within the PNGB construction, which itself is a well motivated way to obtain the ultralight fields needed for cosmic acceleration.  The theories are fairly simple and it should be straightforward to confront them with observational data.  \\item Our class of theories allow for a unified treatment of the cosmological background and perturbations, unlike the background-dependent approach of Ref.\\ \\cite{Creminelli2009}.   \\end{itemize}  Finally, we list some possible directions in which the approach used here could be extended:  \\begin{itemize}  \\item It would be interesting to compute the relation between the nine free functions used in our theories to the free functions of the post-Friedmannian approach to parameterizing dark energy models \\cite{Skordis2011}.   \\item It would be interesting to explore the phenomenology of the various higher order terms in our action, for the cosmological background evolution and perturbations.  Many of the terms have already been explored in detail, see for example Refs.\\ \\cite{Deffayet:2010qz,Charmousis2011}.  \\item Either by using the post-Friedmannian approach, or more directly, it would be useful to compute the current observational constraints on the free functions in the action.  \\item An interesting open question is the extent to which our final action is generic.  That is, is there a class of theories more general than nonlinear sigma model PNGB theories for which our action is obtained by integrating out some of the fields?  \\end{itemize}         \\subsubsection*"
    },
    "1112/1112.2901_arXiv.txt": {
        "abstract": "{Supernova remnants (SNRs) have emerged as one of the largest source classes in very-high-energy (VHE; E$>$0.1\\,TeV) astronomy. Many of the now known VHE $\\gamma$-ray emitting SNRs have been discovered by the H.E.S.S. imaging Cherenkov telescope array, thanks to its unique access to the inner galaxy. Statistically-significant emission of VHE $\\gamma$-rays has now been detected from the direction of the supernova remnant \\g. While the centroids of the H.E.S.S. source and the shell-type SNR are compatible, the VHE morphology suggests a center-dominated source at TeV energies, something which is at odds with the shell-like morphology observed at radio frequencies. This suggests that H.E.S.S. may be observing TeV emission from a previously unknown pulsar wind nebula (PWN) located within the boundaries of the radio shell. If this interpretation is correct, \\g\\ would in fact be a composite SNR, the first case in which an SNR is identified as a composite on the basis of VHE $\\gamma$-ray observations. Archival data from MAGPIS gives exciting hints that there is radio emission from the central parts of the remnant, giving support to this hypothesis. Unfortunately, image artefacts from a nearby strong radio source produce considerable uncertainties in the radio analysis. Additional observations in both the radio and X-ray are needed to confirm the composite nature of \\g\\ suggested by H.E.S.S.} ",
        "introduction": "Very-high-energy (VHE; E$>$0.1\\,TeV) $\\gamma$-ray emission has now been detected from more than 60 Galactic sources, and the vast majority of these are in some way connected to supernova remnants (SNRs) \\cite{Gast2011}. TeV $\\gamma$-rays have been observed from SNR shells themselves and also from the pulsar wind nebulae (PWNe) at the center of many SNRs. While the former gives rise to VHE $\\gamma$-ray emission tracing the shock-front of the SNR, the latter shows a centre-filled morphology, i.e. a central nebula, generally visible from low-energy radio wavelengths up to $\\gamma$-ray photon energies of tens of TeV in the most extreme cases. Young PWNe typically have a rather symmetric morphology and are centrally located inside the host SNR. Such PWNe show a distinct spectral fingerprint with a roughly flat radio spectral index and a much steeper spectrum at higher energies \\cite{pwnevolution}. The pulsar powering the nebula is usually found close to its birthplace at the centre of the SNR in such systems. However, only in a few cases has the pulsar actually been detected. One prominent example of a young PWN is found in the SNR G\\,0.9$+$0.1 which was first classified as a composite SNR on the basis of observations in radio \\cite{helfland1987} then subsequently detected in TeV $\\gamma$-rays \\cite{hessg0.9}, before ultimately the energetic pulsar was found in a deep radio observation of the source \\cite{radiog0.9}.   The SNR \\g\\ is a poorly studied object which was initially discovered in a 90-cm survey of the inner Galaxy conducted by the Very Large Array (VLA) \\cite{brogan2006}. The SNR was reported to have a shell-like morphology with a size about $14^{\\prime} \\times 15^{\\prime}$ and an average spectral index $\\alpha = -0.6 \\pm 0.2$, indicating that the radio emission is dominated by non-thermal synchrotron emission from the shell. For the reasons outlined above, this source is a prime candidate for H.E.S.S. observations. Furthermore, since \\g\\ is larger than the H.E.S.S. point-spread-function (PSF; $\\sim6^{\\prime}$), morphological studies with H.E.S.S. have the capability to distinguish between a shell-type SNR or PWN-like origin of the VHE $\\gamma$-ray emission, thus shedding more light on this little known SNR. ",
        "conclusions": "Since a chance positional coincidence of the centroid of the VHE $\\gamma$-ray emission with the center of the radio shell of \\g\\ seems highly unlikely, an association between \\hessj\\ and the SNR is assumed here. The question is then whether the VHE emission originates from the shell of the SNR or from a putative PWN embedded inside the SNR. However, since the extension of the H.E.S.S. source is significantly smaller than the extension of the SNR as seen in the archival VLA data, it is clear that the bulk of the VHE emission originates from a yet unknown source located within the shell. The discovery of a radio or X-ray PWN counterpart to \\hessj\\ would confirm this scenario. High-quality X-ray observations of the region are unfortunately lacking, but preliminary analyses of 20-cm and 90-cm radio observations suggest that the SNR is composed of two distinct regions with different spectra. The central region is found to be characterised by a relatively flat radio spectrum, characteristic of the central PWN of a composite SNR \\cite{pwnevolution}, while the shell is characterised by a steeper spectrum, indicative of the non-thermal synchrotron emission expected from an SNR shell. \\g\\ thus appears to be an example of a composite SNR, similar to G\\,0.9$+$0.1, the first composite SNR to be detected at TeV energies. If the nature of \\g\\ can be confirmed, \\hessj\\ would be the first case in which a SNR is identified as a composite SNR on the basis of VHE $\\gamma$-ray observations.  Additional supporting evidence for a PWN origin of the TeV emission would be the detection of a pulsar within the central regions of the SNR. Unfortunately, no pulsar is currently known, in radio or X-rays, inside the shell of \\g, nor are there any known hard X-ray point sources. However, a typical young radio pulsar located at $\\sim$9\\,kpc with an assumed luminosity L$_{21\\mathrm{\\,cm}} \\sim 56$\\,mJy\\,kpc$^2$ (i.e. similar to the median of high-luminosity, young, rotation-powered pulsars estimated by \\cite{camilio}) would have a flux density of S$_{21\\mathrm{\\,cm}} \\sim 0.08$\\,mJy, considerably lower than the sensitivity of the current VLA data. Beaming effects could also make the pulsar virtually undetectable in radio. If future observations of \\g\\ are able to detect the putative central pulsar associated with \\hessj, it is expected to have a very high spin-down luminosity ($\\dot{\\mathrm{E}} \\sim 4 \\times 10^{37}$\\,erg\\,s, as seen in other TeV $\\gamma$-ray emitting PWNe \\cite{radiog0.9}.  Little is known about the distance and age of \\g. Using the uncertain $\\Sigma-D$ relationship \\cite{sigmaD}, a distance of $10 \\pm 3$\\,kpc is derived. Adopting 10\\,kpc as the nominal distance of the SNR would imply a Sedov age of $\\sim$9\\,kyr, assuming an ambient ISM density of 1\\,cm$^{-3}$. However, varying the density by $\\pm$50\\% and taking into account uncertainties in the $\\Sigma-D$ relationship, an age as low as 4\\,kyr can be derived. These age estimates would make SNR \\g\\ older than similar systems previously detected in VHE $\\gamma$-rays."
    },
    "1112/1112.5197_arXiv.txt": {
        "abstract": "A plausible scenario for the gamma ray and the hard x-ray burst in a strongly magnetized plasma, based on the collective plasma maser instability, is proposed. The physical parameters with which this scenario becomes relevant are estimated. The attractive feature of this scenario over the conventional cyclotron radiation theory is discussed. ",
        "introduction": "According to a recent study, % the thermal electron gyro-motion in a strongly magnetized plasma could lead to instabilities of  electromagnetic (E\\&M) waves~\\citep{sonmaxwell}. This instability, so-called the maser instability, has significant implications on strongly magnetized plasmas~\\citep{Stone,Biskamp, Innes,Fishman, Nakar, Horse}. The goal of this paper is to estimate its relevance to the gamma ray or hard x-ray burst in the strongly magnetized astrophysical plasmas~\\citep{Fishman, Nakar}. The gamma ray burst in the absence of the magnetic field~\\citep{Mendon\u00e7a} is not relevant to our analysis. Our analysis shows that the photons of 10 keV to 1 MeV can be emitted from a relativistic plasma with the electron density $ 10^{19} - 10^{26}~\\mathrm{cm^{-3}} $ and the magnetic field of order of $10^{8}$ to $10^{13}$ gauss. In contrast to the conventional cyclotron radiation where the photons are emitted rather uniformly, the angular distribution of the radiated photons could be sharply concentrated. Advantages of this possible phenomenon over the ones based on the incoherent cyclotron radiation are discussed.  The rest of the paper is organized as follows. The instability growth rate is presented in Sec. 2, based on the Landau damping theory, which is analyzed when the photon wave vector is parallel to the magnetic field (Sec. 3) and when the photon wave vector is not parallel to the magnetic field (Sec. 4). The gamma ray burst is discussed in Sec. 5, and the difference between this theory and the conventional theories is presented Sec. 6, which is followed by summary (Sec. 7). ",
        "conclusions": "A scenario of the gamma ray burst, based on the recent radiation theory~\\citep{sonmaxwell}, is proposed and examined. The previous analysis~\\citep{sonmaxwell} is generalized to an arbitrary angle. The estimation shows  that  the coherent burst of 10 keV to 1 MeV photons is plausible in a relativistic plasma when $\\gamma_0 = 10\\sim  1000$ and the electron density is higher than some critical value, $n_e > 10^{18} \\sim 10^{26}  \\ \\mathrm{cm}^{-3}$. It is shown that a rather compact dense object with less available energy could cause the short gamma ray burst and that the observed gamma rays would be coherent rather than incoherent. In addition, the coherent photons of lower frequency comparable to the plasma frequency might be observed due to the Weibel instability. This is particularly relevant to the case when $\\theta = \\pi/2$ because the low frequency coherent photons (high frequency coherent photons) from the Weibel instability (from the instability studied here) can be observed simultaneously. The above conjecture  might be useful in verifying whether the scenario proposed here would account for some of the short gamma ray burst events observed in the satellites~\\citep{Nakar}.  While our estimation is rather focused on the gamma and hard x-rays, a similar mechanism would manifest in generating soft x-rays in the inertial confinement fusion plasma~\\citep{tabak}. The electron beam of $\\gamma > 10 \\sim 100$ and the magnetic field of $10^8$ gauss can be readily generated in laboratories. Even a magnetic field of $10^9$ gauss might be possible~\\citep{sonprl}. Then, the photon generated from the instability may have energy between 10 eV and 1 keV. Complications would be the electron quantum diffraction effect and the degeneracy~\\citep{sonpla, sonprl, sonlandau}. The plausibility study is in progress."
    },
    "1112/1112.3122_arXiv.txt": {
        "abstract": "The initial shear field, characterized by a primordial perturbation potential, plays a crucial role in the formation of large scale structures. Hence, considerable analytic work has been based on the joint distribution of its eigenvalues, associated with Gaussian statistics. In addition, directly related morphological quantities such as ellipticity or prolateness are essential tools in understanding the formation and structural properties of halos, voids, sheets and filaments, their relation with the local environment, and the geometrical and dynamical classification of the cosmic web. To date, most analytic work has been focused on Doroshkevich's unconditional formulae for the eigenvalues of the linear tidal field, which neglect the fact that halos (voids) may correspond to maxima (minima) of the density field. I present here new formulae for the constrained eigenvalues of the initial shear field associated with Gaussian statistics, which include the fact that those eigenvalues are related to regions where the source of the displacement is positive (negative): this is achieved by requiring the Hessian matrix of the displacement field to be positive (negative) definite. The new conditional formulae naturally reduce to Doroshkevich's unconditional relations, in the limit of no correlation between the potential and the density fields. As a direct application, I derive the individual conditional distributions of eigenvalues and point out the connection with previous literature. Finally, I outline other possible theoretically- or observationally-oriented uses, ranging from studies of halo and void triaxial formation, development of structure-finding algorithms for the morphology and topology of the cosmic web, till an accurate mapping of the gravitational potential environment of galaxies from current and future generation galaxy redshift surveys. ",
        "introduction": "The large-scale spatial distribution of dark matter, as revealed from numerical simulations, shows a characteristic anisotropic web-like structure. This \\textit{cosmic web}, which arises through the gravitational clustering of matter, is mainly due to the effects of the tidal field: in fact, the competition between cosmic expansion, the trace, and the traceless part of the tidal field imprints anisotropies in the large-scale matter distribution in much the same way that gravity and radiation pressure imprints baryonic acoustic oscillations (BAO) on the Cosmic Microwave Background (CMB) sky (Hu \\& Sugiyama 1995; Lee \\& Springel 2010). Hence, the initial shear field plays a crucial role in the formation of large scale structures, and a number of studies in the literature have been devoted to this subject -- among the plethora of papers, see for example the classic works by Zeldovich (1970), Icke (1973), Peebles (1980), White (1984), Bardeen et al. (1986), Kaiser (1986), Bertschinger (1987), Bond \\& Myers (1996), Bond, Kofman \\& Pogosyan (1996) and van de Weygaert \\& Bertschinger (1996). In addition, if the cosmic web originates from primordial tidal effects and its degree of anisotropy increases with the evolution of the Universe, then studying this initial field is crucial in understanding the subsequent nonlinear evolution of cosmic structures (Springel et al. 2005; Shandarin et al. 2006; Desjacques 2008; Desjacques \\& Smith 2008; Pogosyan et al. 2009), the alignment of shape and angular momentum of halos (West 1989; Catelan et al. 2001; Lee \\& Springel 2010; Rossi, Sheth \\& Tormen 2011), the statistical properties of voids (Lee \\& Park 2006; Platen, van de Weygaert \\& Jones 2008), and more generally for characterizing the geometry and morphology of the cosmic web (Shen et al. 2006;  van de Weygaert \\& Bond 2008; Forero-Romero et al. 2009; Aragon-Calvo et al. 2010a,b; Shandarin et al. 2010).  The basic theory for the formation and evolution of structures is now well-understood, thanks to the pioneering work of Doroshkevich \\& Zeldovich (1964),  Doroshkevich (1970), Zeldovich (1970), and Sunyaev \\& Zeldovich (1972) -- the latter in the context of galaxy formation. In particular, Doroshkevich (1970) derived the joint probability distribution of an ordered set of eigenvalues in the tidal field matrix -- at random positions -- given the variance of the density field, corresponding to a Gaussian potential; we will refer to it as the \\textit{unconditional} probability distribution of eigenvalues. In addition, Zeldovich (1970) provided the fundamental understanding of anisotropic collapse on cosmological scales, and recognized the key role of the large scale tidal force in shaping the cosmic web. Subsequently, Doroshkevich and Shandarin (1978) calculated some statistical properties of the maxima of the largest eigenvalue of the shear tensor. Their study proved that the most probable formation process starts first with a one-dimensional collapse (cosmic pancake formation). The directions (orientations) for the one-dimensional collapsed sheets are determined by the largest eigenvalue of the deformation tensor, which can be attributed to the initial linear density perturbations; moreover, the probability that two or even three of the initial eigenvalues are identical or nearly equal is extremely small, indicating that the collapse is triaxial. Later on, following a constrained field approach pioneered by Bertschinger (1987), van de Weygaert \\& Bertschinger (1996) developed an algorithm for setting up tailor-made initial conditions for cosmological simulations, which addressed the role of the tidal fields in shaping large-scale structures. In the same period, Bond, Kofman \\& Pogosyan (1996) developed a cosmic web theory which naturally explains the filamentary structure present in the Cold Dark Matter (CDM) cosmology, due to the coherent nature of the primordial tidal field. They realized that an ``embryonic'' cosmic web is already present in the primordial density field, and explained why in overdense regions sheet-like membranes are only marginal features. Since then, because of the correspondence between structures in the evolved density field and local properties of the linear tidal field pointed out by the same authors, the statistics of the shear has received more attention in the literature. For example, Lee \\& Shandarin (1998) computed some probability distributions for individual shear eigenvalues and obtained an analytic approximation to the halo mass function, and Catelan \\& Porciani (2001) explored the two-point correlation of the tidal shear components.  However, a variety of studies in the physics of halo formation and cosmic web classification are based on Doroshkevich's unconditional formulae for the ordered eigenvalues of the initial shear field associated with Gaussian statistics (Doroshkevich 1970), but those formulas cannot differentiate between random positions and peak/dips as they neglect the fact that halos (voids) may correspond to maxima (minima) of the density field. According to Bardeen et al. (1986), if one assumes the cosmological density fluctuations to be Gaussian random fields, the local maxima of such fields are plausible sites for the formation of nonlinear structures. Hence, the statistical properties of the peaks can be used to predict the abundances and clustering properties of objects of various types, and in studies of the non-spherical formation of large-scale structures. The study of Bertschinger (1987) goes in this direction, by generalizing the treatment of Bardeen et al. (1986) and proposing a path integral method for sampling constrained Gaussian random fields. The method allows one to study the density field around peaks or other constrained regions in the biased galaxy formation scenario, but unfortunately it is too elaborate and inefficient in its implementation. Instead, by applying the prescription of Hoffman \\& Ribak (1991) to construct constrained random fields, van de Weygaert \\& Bertschinger (1996) were able to show that it is possible to generate efficiently initial Gaussian random density and velocity fields, and specify the presence and characteristics of one or more peaks and dips at arbitrary locations -- with the gravity and tidal fields at the site of the peaks having the required strength and orientation. Some other studies on constrained initial conditions were also pursued by van Haarlem \\& van de Weygaert (1993), and by van de Weygaert \\& Babul (1994). The idea has been expanded in Bond \\& Myers (1996), who presented a \\textit{peak-patch} picture of structure formation as an accurate model of the dynamics of peaks in the density field. Their approach goes further, as it involves the explicit formalism for identifying objects in a multiscale field (i.e. it does not restrict to a single scale). The model is even more precise for void patches, the equivalent framework for studying voids (Sahni et al. 1994; Sheth \\& van de Weygaert 2004; Novikov, Colombi \\& Dore 2006; Colberg et al. 2008). In general, density peaks define a well-behaved point-process which can account for the discrete nature of dark matter halos and galaxies, and on asymptotically large scales are linearly biased tracers of the dark matter field (Desjacques \\& Sheth 2010).  Therefore, it would be desirable to incorporate the peak (or dip) constraint in the statistical description of the initial shear field, in order to characterize more realistically the geometry and dynamics of the cosmic web. The main goal of this paper is to do so, by providing a set of analytic expressions which extend the work of Doroshkevich (1970) and Bardeen et al. (1986), and are akin in philosophy to that of van de Weygaert \\& Bertschinger (1996). This is achieved by constraining the Hessian of the displacement field (the matrix of the second derivatives) to be positive (negative) definite, which is the case in the vicinity of minima (maxima) of the source of the displacement field. The new \\textit{conditional} probability distributions derived in this study include the correlation between the potential and density fields through a reduced parameter $r$, and naturally recover Doroshkevich's (1970) \\textit{unconditional} formulae in the absence of correlation.  Hence, the main focus of this work is to derive explicitly the joint probability distribution of the eigenvalues of the shear field, given the fact that positions are peaks or dips in the corresponding density field -- and not random locations. In this sense the field is termed \\textit{constrained} (i.e. it has \\textit{constrained} eigenvalues, which are the result of looking only at peak/dip regions), and it is statistically described by a \\textit{conditional} probability distribution. Note that even though the formalism is essentially restricted to one scale, the formulae derived here are useful in a variety of applications -- some of which will be discussed at the end of this paper.   The layout is organized as follows. Section \\ref{new_formula} provides the derivation of the new analytic expressions for the conditional distributions of eigenvalues of the initial shear field. In particular, Section \\ref{notation} illustrates the basic notation adopted; Section \\ref{joint_eigen_new} contains the derivation of the joint distribution of eigenvalues in the peak/dip picture, while Section \\ref{bbks_connection} shows the reverse conditional distribution function. As a direct application of the new formulae, Section \\ref{individual_eigenvalues} presents the individual distributions of eigenvalues subjected to the extremum constraint, along with some other related conditional probabilities. This part extends previous work by Lee \\& Shandarin (1998), which is briefly summarized in Appendix \\ref{Lee_Shandarin}, and complements the study of van de Weygaert \\& Bertschinger (1996). Finally, Section \\ref{conclusion} highlights the main results and discusses ongoing and future applications, which will be presented in forthcoming publications. ",
        "conclusions": "\\label{conclusion}  Since the initial shear field associated with Gaussian statistics plays a major role in the formation of large scale structures, considerable analytic work has been based on the joint distribution of its eigenvalues -- i.e. Doroshkevich's formulae (Equation \\ref{doro_standard}). However, Doroshkevich's equations neglect the fact that halos (voids) may correspond to maxima (minima) of the density field. The main goal of this work was to provide new analytic expressions, in the context of the peak/dip picture (Bardeen et al. 1986; Bond et al. 1991), which include the fact that the eigenvalues of the linear shear field are related to regions where the source of the displacement is positive (negative). These new conditional probabilities, derived in Section \\ref{new_formula}, are Equations (\\ref{doro_inter_extended}), (\\ref{doro_eigen}), (\\ref{bbks_inter_extended}) and (\\ref{bbks_eigen}): they represent the main results of this paper. Written in Doroshkevich-like format, they naturally reduce to Doroshkevich's (1970) unconditional relations in the limit of no correlation between the potential and the density fields (i.e. when $r=0$). As a first direct application of (\\ref{doro_eigen}), the individual conditional distributions of eigenvalues were obtained in Section \\ref{individual_eigenvalues}; these relations extend some previous work by Lee \\& Shandarin (1998) -- see Appendix \\ref{Lee_Shandarin} for a compact summary of their results. Much more analytic work can be carried out using these new formulae, especially in connection with the statistics of peaks developed by Bardeen et al. (1986): their calculations can be extended within this framework (see Section \\ref{bbks_connection}),  and results of this extension will be presented in a forthcoming publication -- along with more insights on the equations derived in Section \\ref{joint_eigen_new}.  To obtain the conditional distribution $p(\\tilde{T}|r,\\tilde{H}>0)$ (i.e. Equation \\ref{what_we_want}), in principle one needs to sample numerically the probability (\\ref{cond}) as done in Lavaux \\& Wandelt (2010). However, the new analytic results of this paper suggest a simpler generalized excursion set algorithm, which will be also presented in a forthcoming study. The algorithm allows for a fast sampling of (\\ref{cond}), and permits to test the new relations derived here (i.e. Sections \\ref{new_formula} and \\ref{individual_eigenvalues}, and Appendix \\ref{Lee_Shandarin}) against mock data. In addition, along these lines it is possible to extend the main shape distributions involved in the triaxial formation of nonlinear structures (i.e. ellipticity $e$, prolateness $p$, axis ratios $\\nu$ and $\\mu$, etc.). Halos and voids are in fact triaxial rather than spherical (i.e. Rossi, Sheth \\& Tormen 2011), and the initial shear field plays a crucial role in their formation. For example, in the ellipsoidal collapse framework the virialization condition depends on ellipticity and prolateness, which are directly related to the eigenvalues of the external tidal field. Starting from $p(\\zeta_1,\\zeta_2,\\zeta_3|r)$, it is then straightforward to turn this joint conditional distribution into $p(e,p,\\delta|r)$ -- to include the peak constrain, -- and then to characterize $p(e|r)$, $p(p|r)$, $p(\\mu|r)$, $p(\\nu|r)$ and so forth. Applications to halos and voids will be also presented next, including implications for the skeleton of the cosmic web. The extension to non-Gaussian fields, along the lines of Lam et al. (2009), is a natural follow-up of this work and will be presented in a forthcoming publication as well. In the context of cosmic voids, other interesting uses of the new formulae (\\ref{doro_inter_extended}), (\\ref{doro_eigen}), (\\ref{bbks_inter_extended}) and (\\ref{bbks_eigen}) involve the Monge-Amp\\`{e}re-Kantorovitch reconstruction procedure -- see for example Lavaux \\& Wandelt (2010).  The analytic framework described here can be also useful in several observationally-oriented applications, and in particular for developing algorithms to find and classify structures in the cosmic web. For example, Bond, Strauss \\& Cen (2010) presented an algorithm that uses the eigenvectors of the Hessian matrix of the smoothed galaxy distribution to identify individual filamentary structures. They used the distribution of the Hessian eigenvalues of the smoothed density field on a grid to study clumps, filaments and walls. Other possibilities include a web classification based on the multiscale analysis of the Hessian matrix of the density field (Aragon-Calvo et al. 2007), the skeleton analysis (Novikov et al. 2006), as well as a morphological (Zeldovich-based) classification (for instance Klypin \\& Shandarin 1983; Forero-Romero et al. 2009). More generally, the fact that the eigenvalues of the Hessian matrix can be used to discriminate different types of structure in a particle distribution is fundamental to a number of structure-finding algorithms (Forero-Romero et al. 2009), shape-finders algorithms (Sahni et al. 1998), and structure reconstruction on the basis of tessellations (Schaap \\& van de Weygaert 2000; Romano-Diaz \\& van de Weygaert 2007), etc. In addition, the classification of different environments should provide a framework for studying the environmental dependence of galaxy formation (see for example Blanton et al. 2005). To this end, recently Park, Kim and Park (2010) extended the concept of galaxy environment to the gravitational potential and its functions -- as the shear tensor. They studied how to accurately estimate the gravitational potential from an observational sample finite in volume, biased due to galaxy biasing, and subject to redshift space distortions, by inspecting the dependence of dark matter halo properties on environmental parameters (i.e. local density, gravitational potential, ellipticity and prolateness of the shear tensor). It would be interesting to interpret their results with the theoretical formalism developed here, and ultimately to study the gravitational potential directly from a real dataset such as the SDSS Main Galaxy sample within this framework.  The formalism presented in this paper is restricted to one scale (i.e. peaks and dips in the density field, as in Bardeen et al. 1986), but the extension to a multiscale \\textit{peak-patch} approach along the lines of Bond \\& Myers (1996) is duable and subject of ongoing work. This will allow to account for the role of the peculiar gravity field itself, an important aspect not considered here but discussed for example in van de Weygaert \\& Bertschinger (1996). In fact, these authors introduced the peak constraints to describe the density field in the immediate surroundings of a peak, and then addressed the constraints on the gravitational potential perturbations; in particular, they constrained the peculiar gravitational acceleration at the position of the peak itself, in addition to characterizing the tidal field around the peak. Including all these effects in our formalism is ongoing effort. Finally, the interesting and more complex question of the local expected density field alignment/orientation distribution as a function of the local field value (or the other way around -- see Bond 1987; Lee \\& Pen 2002; Porciani et al. 2002;  Lee 2011) can be addressed within this framework, and is left to future studies."
    },
    "1112/1112.1916_arXiv.txt": {
        "abstract": "The new multi-color $BVRI$ photometric light curves of the short-period eclipsing binary \\astrobj{GSC 3576-0170} were obtained on two consecutive nights (October 5 and 6, 2009). With the 2003 version of Wilson-Devinney program, the precise photometric solutions are derived for the first time. The result shows that \\astrobj{GSC 3576-0170} is a semi-detached binary system with a large temperature difference of approximately 1490~K. The light-curve distortions are further explained by a hot spot on the secondary component through mass transfer via a stream hitting the facing surface of the secondary component. By analyzing all available light minimum times, we also derived an update ephemeris and found for the first time a possible periodic oscillation with an amplitude of 0.0038 days and a period of 4.3 years. The periodic oscillation could be explained either by the light-time effect due to a presumed third component or by magnetic activity cycle of the system. ",
        "introduction": "\\label{intro} \\astrobj{GSC 3576-0170}(P$_{orb}$=$0^{d}$.405, G1 V) is a new near-contact solar-type eclipsing binary, which displays asymmetric light curves (two different light maxima) (Nelson et al. 2006). Therefore, it is a very intriguing target for understanding the property of the system.\\\\ \\indent \\astrobj{GSC 3576-0170} was discovered to be variable by Nelson et al. (2006). They attempted a preliminary light-curve synthesis with Wilson-Devinney modeling and preferred a detached model with a mass ratio in the range $0.15-0.35$ and an orbital inclination of 65-$70^{\\circ}$. For the period variation, they invoked a quadratic function to fit the values of the observational times of light minimum - calculational times of light minimum $(O-C)$. Subsequently, new minimum times of GSC 3576-0170 were published by several astronomers (Nelson et al. 2006; H$\\ddot{u}$bscher \\& Walter 2007; H$\\ddot{u}$bscher et al. 2006; Br\\'{a}t 2007; etc.).\\\\ \\indent In this paper, we present our new $B, V, R,$ and $I$ LCs of \\astrobj{GSC 3576-0170} which were analyzed using the 2003 version of Wilson-Devinney code (Wilson \\& Devinney 1971; Wilson 1979, 1990, 1994; Wilson \\& Van Hamme 2004). We accumulated all available times of light minimum and discussed the period change.\\\\ \\begin{table} \\caption{Minimum times of \\astrobj{GSC 3576-0170}.} \\tabcolsep 0.30truecm \\renewcommand\\arraystretch{0.6} \\begin{tabular}{lrrc} \\hline \\hline \\multicolumn{1}{l}{JD(Hel.)} & \\multicolumn{1}{c}{Cycle} & \\multicolumn{1}{c}{(O-C)}  & \\multicolumn{1}{c}{References} \\\\ \\hline 2452794.863     & $-2.5$ & $-0.0001$ & 1\\\\ 2452795.8716     & $0.0$ &$-0.0038$  & 1\\\\ 2452799.9230     & $10.0$ & $-0.0024$  & 1 \\\\ 2452802.5542       &$16.5$ &$-0.0039$  &1,4\\\\ 2452806.8076     &$27.0$ &$-0.0029$   &1\\\\ 2452807.8210&$29.5$&$-0.0020$& 1\\\\ 2452812.4781&$41.0$&$-0.0026$& 1,4\\\\ 2452826.8600&$76.5$&$0.0017$&1\\\\ 2452829.4887&$83.0$&$-0.0021$&4\\\\ 2452831.5151&$88.0$&$-0.0007$&4\\\\ 2452863.5105&$167.0$&$-0.0007$&4\\\\ 2452864.5304&169.5&0.0067&4\\\\ 2452867.5607&177.0&$-0.0006$&4\\\\ 2452868.3701&179.0&$-0.0012$&4\\\\ 2452946.3385&371.5&$0.0037$&4\\\\ 2453215.4640&1036.0&0.0034&4\\\\ 2453216.4767&1038.5&0.0036&4\\\\ 2453217.4888&1041.0&$0.0032$&4\\\\ 2453221.5370&1051.0&0.0013&4\\\\ 2453263.8659&$1155.5$&0.0072&1 \\\\ 2453264.6735&$1157.5$&$0.0048$&1\\\\ 2453305.7787 &$1259.0$&0.0020&1\\\\ 2453612.3645&2016.0&$-0.0011$&4\\\\ 2453612.5733&2016.5&0.0052&4\\\\ 2453621.4824&2038.0&0.0042&4\\\\ 2453837.9506&$2573.0$&$-0.0028$&1\\\\ 2453852.9370&$2610.0$&$-0.0016$&1\\\\ 2453900.7278&$2728.0$&$-0.0014$&1\\\\ 2453941.8337&$2829.0$&$-0.0035$&1\\\\ 2453943.8605&$2834.5$&$-0.0017$&1\\\\ 2454073.2580&$3154.0$&$-0.0033$&2\\\\ 2454270.4947&$3641.0$&$-0.0041$&3\\\\ 2455109.0636&5711.5&$0.0019$&5\\\\ 2455110.0739&5714.0&$-0.0004$&5\\\\ \\hline \\end{tabular} \\small \\renewcommand\\arraystretch{0.1} \\\\References - (1) Nelson et al. 2006; (2) H$\\ddot{u}$bscher \\& Walter 2007; (3) Br\\'{a}t 2007; (4) H$\\ddot{u}$bscher et al. 2006; (5) Present paper. \\end{table} ",
        "conclusions": "Our new LCs and the period variation of \\astrobj{GSC 3576-0170} are included in present investigation. \\subsection{Cyclic magnetic activity or the light-time effect due to the third Body?} For \\astrobj{GSC 3576-0170}, there seems to be the sinusoidal oscillation of the O-C residuals. The oscillating characteristic may be caused by the light-time effect due to the existence of the third body orbiting the eclipsing binary or a possible result of magnetic activity cycles of the system.\\\\ \\indent On one hand, the period oscillation may be explained by the light-time effect. Assuming that the orbital of the presumed third body is circular, we can obtain the mass function for the third body f($M_{3}$)=0.016($\\pm$0.001) M$\\odot$ using the following equation: \\begin{equation} f(m) =\\frac{4\\pi^{2}}{G{T}^{2}}\\times(a\\prime_{12}\\sin{i^{\\prime}})^{3}=\\frac{(M_{3}\\sin{i_3^{\\prime}})^{3}} {(M_{1}+M_{2}+M_{3})^{2}}. \\end{equation} where $M_{1}$, $M_{2}$ and $M_{3}$ are the masses of the binary, and the third body, respectively. The mass of the third component can also be calculated with the above equation, which depends on the orbital inclination. The minimal mass $M_{3,min}$ is 0.38 M$\\odot$ when $i^{\\prime}$=$90^{\\circ}$. \\\\%Because the third mass is much \\indent On the other hand, the period oscillation may be accounted by magnetic cycle of the system (Applegate 1992; Lanza et al. 1998; Lanza \\& Rodon$\\grave{o}$ 1999). The orbital period change corresponding to a variation of the quadrupole moment is given by $\\Delta Q$~=~-($\\Delta p$/p)$\\times$($Ma^{2}$/9) (Applegate 1992) where a is the semi-major axis of the binary orbit and M is the mass of the active component. Since no spectroscopic solutions are available in literature \\astrobj{GSC 3576-0170}, its absolute parameters can not be directly determined. We estimated the primary mass 1.024 as M$\\odot$, and the radius as 1.06 R$\\odot$ by assuming the primary component to be a normal and main-sequence GIV star. Based on our photometric solutions, the mass of the secondary component is $M_{2}$ = 0.478 M$\\odot$ and the separation between the two components is a~=~2.37 R$\\odot$. Therefore, the quadrupole moment $\\Delta Q$ is calculated to be 0.5$\\times$$10^{50}$ g $cm ^{2}$ for the primary.% This value is similar to the typical values for active binaries (Lanza \\& Rodono 1999). Therefore, the magnetic activity cycle is a possible mechanism to explain the variation of period.\\\\ \\subsection{Long-term change and their evolutionary status} For our observations, the two light maxima are basically equal. However, the light curves show that the first quadrature is brighter than the second one in 2003 (Nelson et al. 2006). These indicate a variability on a time scale of about six years. This variation might be attributed to surface magnetic activity of one or both stars. More importantly, GSC 3576-0170 is perhaps associated with possibly thick deep convection according to their spectral types G1 V. Thus, it confirms that magnetic activities may produce the sinusoidal variation of the period.\\\\   \\indent Our result showed that \\astrobj{GSC 3576-0170} is a marginal-contact binary system with a large temperature difference of about 1490~K between the two components, which belongs to the subclass of V1010 Oph binaries (Shaw 1994). It is similar to several semi-detached systems with a possible third body, such as \\astrobj{GW Tau} (Zhu \\& Qian 2006). They may be potential candidates of binaries that will become overcontacted or remain in the broken-contact in accordance with the thermal relaxation oscillation theory of overcontact binaries (Lucy 1976; Flannery 1976; Robertson \\& Eggleton 1977; etc.).\\\\ Of course, our solutions are based on photometric observations only. For better understanding of the properties and the evolutionary state of \\astrobj{GSC 3576-0170}, spectroscopic observations are needed. Moreover, the O-C residual of \\astrobj{GSC 3576-0170} is about seven years and there is an observational gap. Therefore, it is too early to decide the character of the periodic variation, and it might take another 10-20 years are needed for confirmation.\\\\ {\\bf Acknowledgements} The authors gratefully acknowledge the assistance provided for the 85 cm telescope at Xinglong station. We also thank Profs. X., Zhou, A. Y. Zhou, X.J. Jiang, and Y. H. Zhao, for allocation of time for observation and their kind helps during the visit to NAOC. This work was partly supported by GuiZhou University under Grant No. 2008036, GuiZhou Natural Science Foundation 20092263, Shandong Natural Science Foundation (ZR2009AM021) and the Joint Fund of Astronomy of the National Natural Science Foundation of China and the Chinese Academy of Sciences Grant No. 10978010, Natural and Scientific funding supported by Department of Education of Guizhou Province No. 20090130, and Dezhou University Foundation (402811)."
    },
    "1112/1112.5915_arXiv.txt": {
        "abstract": "We report on the first simultaneous observation of an H$\\alpha$ Moreton wave, the corresponding EUV fast coronal waves, and a slow and bright EUV wave (typical EIT wave). Associated with an X6.9 flare that occurred on 2011 August 9 at the active region NOAA 11263, we observed a Moreton wave in the H$\\alpha$ images taken by the Solar Magnetic Activity Research Telescope (SMART) at Hida Observatory of Kyoto University. In the EUV images obtained by the Atmospheric Imaging Assembly (AIA) on board the {\\it Solar Dynamic Observatory} ({\\it SDO}) we found not only the corresponding EUV fast ``bright'' coronal wave, but also the EUV fast ``faint'' wave that is not associated with the H$\\alpha$ Moreton wave. We also found a slow EUV wave, which corresponds to a typical EIT wave. Furthermore, we observed, for the first time, the oscillations of a prominence and a filament, simultaneously, both in the H$\\alpha$ and EUV images. To trigger the oscillations by the flare-associated coronal disturbance, we expect a coronal wave as fast as the fast-mode MHD wave with the velocity of about 570 -- 800~km~s~$^{-1}$. These velocities are consistent with those of the observed Moreton wave and the EUV fast coronal wave. ",
        "introduction": "Moreton waves, flare-associated waves seen in H$\\alpha$, have been observed \\citep{mor1960,smi1971,shi2011} to propagate in restricted angles with the velocity of about 500 -- 1500~km~s$^{-1}$. They sometimes show arc-shaped fronts, and are often associated with type-II radio bursts \\citep{kai1970}. They are transient, and appear only for about 10 minutes. Associated with flares, remote filaments and prominence are sometimes activated or excited to oscillate. These ``winking filaments'' are also thought to be caused by flare-associated waves, and are called as invisible Moreton waves \\citep{smi1971,tri2009,her2011}. After the findings, \\citet{uchi1968} suggested that Moreton waves are the intersection of the fast-mode magnetohydrodynamic (MHD) shock propagating in the corona with chromosphere. This model has been widely accepted, and the coronal counterparts have been surveyed for a few decades. Moreton waves are rare to be observed even for large flares \\citep{shi2011}.  After the launch of the {\\it Solar and Heliospheric Observatory} ({\\it SOHO}), the EUV Imaging Telescope (EIT) found wavelike phenomena associated with flares, which are called ``EIT waves'' \\citep{thom1999,thom2000}. Although EIT waves were expected to be the coronal counterpart of Moreton waves, they show different physical characteristics from those of Moreton waves: the propagating velocity is much slower than that of Moreton wave and is about 200 -- 400~km~s$^{-1}$, the lifetime is much longer and is about 45 -- 60 minutes, they can show isotropic propagation, while Moreton waves propagate with restricted angles \\citep{kla2000,war2007,thom2009}. There have been, therefore, remained a question whether EIT waves are really coronal counterparts of Moreton waves or no. As for searching for a coronal counterpart of Moreton waves, the Soft X-ray Telescope (SXT) on board {\\it Yohkoh} found wavelike phenomena in soft X-rays, called X-ray waves \\citep{khan2000,khan2002}. X-ray waves are confirmed to be a real counterpart of Moreton waves by simultaneous observations of X-ray waves and Moreton waves \\citep{naru2002,naru2004}.  Then, we come to an issue what EIT waves are. \\citet{eto2002} clearly showed that an EIT wave is different from a Moreton wave, based on simultaneous observations of them. On the other hand, \\citet{war2004a,war2004b} argue that the velocity discrepancy of EIT and Moreton waves can be explained by the deceleration of coronal waves. The mechanism of EIT waves remains, therefore, very controversial \\citep{war2007,will2009,gall2010}. \\citet{del1999} and \\citet{chen2002,chen2005} proposed the field-line stretching model for EIT waves. They suggested that EIT bright fronts were not ``waves'' at all, but instead plasma compression at stable flux boundaries due to rapid magnetic field expansion. This model can also resolve the puzzle why EIT waves often stop at magnetic separatrices.  Recently, by the Atmospheric Imaging Assembly (AIA; Title \\& AIA team 2006, Lemen et al. 2011) on board the {\\it Solar Dynamic Observatory} ({\\it SDO}), fast coronal waves have been observed associated with flares \\citep[e.g.][]{liu2010,ma2011}. These waves (hereafter called ``EUV fast coronal waves'') are thought to be the fast-mode MHD waves. Coronal X-ray waves and EUV fast coronal waves have been also observed spectroscopically with the EUV Imaging Spectrometer (EIS) on board {\\it Hinode} \\citep{asa2008,harra2011}.  \\citet{chen2011} found two different coexisting coronal waves, slow coronal wave (i.e. EIT wave) and fast coronal wave, from EUV observations taken by {\\it SDO}/AIA. Although the fast coronal wave seems to be the coronal counterpart of a Moreton wave, they used only EUV images, and it remained to be confirmed whether it is a classical H$\\alpha$ Moreton wave. This letter presents the first simultaneous observation of EUV waves and a Moreton wave by using EUV and H$\\alpha$ images with high spatial and temporal resolutions. Moreover, we found not only a winking filament on the disk, but also an oscillating prominence on the limb, triggered by the coronal wave (Moreton wave). ",
        "conclusions": "We simultaneously observed the H$\\alpha$ Moreton wave and the corresponding EUV fast coronal wave. The Moreton wave front was well consistent with the fast-bright-sharp EUV wave front (F2b). Even after the Moreton wave disappeared, we identified the propagation of the fast-faint EUV wave (F2f). Even along the Line 3 (Figure~\\ref{shock}a), which is the direction without the Moreton wave, we found that the fast EUV waves (F3b, F3f). The fast EUV waves (F1, F2f, F3b, and F3f) are thought to be the fast-mode MHD waves (coronal waves). The EUV fast coronal waves appear more frequently than Moreton wave, although they are very faint and have not been observed until the launch of {\\it SDO}. Especially, only when the shock strongly contacts with the chromosphere, the intersection is observed as the Moreton wave (F2b). The temporal evolutions of the H$\\alpha$ Moreton wave and the EUV waves are summarized in Figure~\\ref{hei_tim}(b).  In Figure~\\ref{hei_tim} we also show the temporal evolutions of a type-II radio burst associated with the flare. In the metric radio spectrogram (25 -- 2500~MHz) observed with the Hiraiso Radio Spectrograph (HiRAS; Kondo et al. 1995), we identify the type-II radio burst from 08:02:40 to 08:06:30~UT. Assuming the coronal density model proposed by \\citet{new1961} and \\citet{mann1999}, we derived the propagation velocity of about 850~km~s$^{-1}$. The observed type-II radio burst seems to be consistent with the Moreton wave and fast bright EUV wave, which also supports the interpretation of the wave as a fast-mode MHD shock.  We also found oscillations of a prominence and a filament.% To trigger the oscillations by a flare-associated coronal disturbance, we expect a coronal wave as fast as the fast-mode MHD wave with the velocity of about 570 -- 800 km~s~$^{-1}$. These velocities are consistent with the propagation velocities of the observed Moreton wave and the EUV fast coronal wave. An invisible Moreton wave could be such an EUV fast-faint coronal wave. It is known that a typical slow EIT wave sometimes causes filament oscillations \\citep{oka2004}, or such filament oscillations triggered by an EIT wave are expected to be stronger than those by an invisible Moreton wave (P. F. Chen, private communication). In the current case, however, the role of the EIT wave on the filament/prominence oscillations is unclear.  Along the Line 2 and 3, we identified slow-bright EUV waves (S2b, S3b) behind the fast-faint EUV waves. From the propagating features (i.e. the velocity and the isotropic propagation; Figure~\\ref{euv}), we think it is a typical EIT wave. We simultaneously observed the EIT wave and the EUV fast wave (fast-mode MHD wave) as reported by \\citet{chen2011}. This means that the EIT wave is different from a fast-mode MHD wave, which supports the field-line stretching model \\citep{chen2002,chen2005}. It is, however, difficult to clearly identify both features separately in the very initial phase of the flare, and it is unclear the relation between the fast-bright (F3f) and the slow-bright (S3b) EUV waves along the Line 3.  There is an alternative possibility that we observed a single coronal disturbance, and the two different waves (F2f/F3f and S2b/S3b) correspond to the fronts of two different heights (faster ones are higher) of the disturbance \\citep{vero2008,war2011}.  We, however, do not think the possibility due to the following reasons: First, the slow bright waves (S2b, S3b) stopped propagating at small active regions. This conflicts with the features of a fast-mode wave, % while this is possibly reconsiled by considering the stopping front as CME flanks as reported by \\cite{pat2009}. Second, we observed the prominence/filament oscillations. Since they are located low in the corona, we can derive the velocity of the shocks/waves there, from the distances and the times of the oscillations. In the current case we need waves with velocities of 570~km~s$^{-1}$ or more. Especially, the direction of the filament is close to the Line 3, and the required velocity is much faster than the bright slow wave (S3b), while it is consistent with the fast faint wave (F3f)."
    },
    "1112/1112.2018_arXiv.txt": {
        "abstract": " ",
        "introduction": "The outer crust of a neutron star, below densities of $\\rho \\sim 4 \\times 10^{11}$ g cm$^{-3}$ consists of matter in a state not too far removed from that found in white dwarfs: a lattice of nuclei permeated by a relativistic, degenerate electron gas which gives the dominant contribution to the pressure of the matter \\cite{BPS1971}. As pressure increases with depth, equilibrium with respect to weak interactions drives the nuclei to become more neutron rich. There comes a point when the intra-nuclear forces can no longer bind all the neutrons, and \\emph{neutron drip} occurs. Above $\\rho \\sim 4 \\times 10^{11}$ g cm$^{-3}$, a new regime is entered in which the nuclear lattice is bathed in a fluid of (`dripped') neutrons. These neutrons are delocalized much like conduction band electrons in metals. From this density inwards, the \\textbf{equation of state} (\\textbf{EOS}) is dominated by pressure arising from nucleon-nucleon interactions.  \\begin{figure}[!tb]\\label{fig:1} \\begin{center} \\includegraphics[width=14cm,height=8cm]{Fig1.pdf} \\caption{Cartoon representation of the inner neutron star crust.} \\end{center} \\end{figure}  A cartoon representation of the crustal layers below the outer crust is shown in Fig.~1. Models predict two distinct layers. (1) The \\textbf{inner crust} between densities $\\rho \\sim 4 \\times 10^{11}$ g cm$^{-3}$ and $\\rho \\sim 10^{14}$ g cm$^{-3}$ is an elastic solid consisting of a lattice of heavy, neutron rich nuclei surrounded by fluid neutrons, with the nuclei increasing in size and mass with density while the inter-nuclear spacing and nuclear proton fraction decrease \\cite{BBP1971}.  (The presence of a background electron gas will be taken as given from now on). The dripped neutrons are expected to become superfluid shortly after the neutron star is formed as it rapidly cools below temperatures of $10^8 - 10^9$ K \\cite{Dean2003}. (2) The \\textbf{mantle} \\cite{Gusakov2004} between $\\rho \\sim 10^{14}$ g cm$^{-3}$ and the crust-core transition density consists of frustrated matter: the competition between the nuclear surface energy and the nuclear and lattice Coulomb energies over similar length scales drives the formation of exotic nuclear geometries termed nuclear `pasta'  \\cite{Ravenhall1983, Oyamatsu1984} which proceed through a canonical sequence of phases: cylindrical (spaghetti) $\\to$ slab (lasagna) $\\to$ cylindrical bubble $\\to$ spherical bubble. The latter three (`bubble') phases are distinguished by the delocalization of the \\textbf{charged nuclear component} of the matter (containing the protons) in one or more dimensions, and corresponding localization of the \\textbf{charge-neutral nuclear component} (fluid neutrons). Similar microscopic structures are observed in terrestrial soft condensed matter systems such as surfactants \\cite{Jones2002,Watanabe2005}; by analogy, we can expect rich mechanical properties, intermediate between liquid and elastic solid, to emerge in the mantle \\cite{Pethick1998}. Some crust models predict the absence of the mantle \\cite{Douchin2001}; its presence depends sensitively on the nuclear microphysics of the crust. Fig.~1 gives a range of widths for the inner crust and mantle taken from the model presented in this chapter and encompassing a range of neutron star masses from 1-2 $M_{\\odot}$ and a range of equations of state as discussed later.  To describe the states of matter in a neutron star, one needs a model for the nucleon-nucleon interactions as they are manifested in a many-nucleon context. A useful concept that bridges the gap between \\emph{ab initio} nucleon-nucleon calculations, nuclear experimental observables, and neutron star matter is that of uniform \\textbf{nuclear matter (NM)}. This is an idealized system, homogeneous and infinite in extent, of neutrons and protons interacting solely via the strong force. The energy per particle of such a system at a density $\\rho$ and proton fraction $x$, $E(\\rho, x)$, is referred to as the nuclear matter equation of state (\\textbf{NM EOS}). In the regions of the neutron star core where protons and neutrons exist, the NM EOS can be combined with the electron energy, and under conditions of charge neutrality and beta-equilibrium gives an EOS for the core. In the inner crust, the NM EOS can be used to describe the dripped neutrons ($x=0$) and the bulk matter in the nuclear clusters. A consistent model for the EOS of crust and core necessarily uses a unique NM EOS, and one should expect parameters characteristic of a given NM EOS to correlate with both crust and core properties.  Nuclear matter with equal numbers of neutrons and protons ($x=0.5$) is referred to as \\textbf{symmetric nuclear matter (SNM)}; nuclear matter with $x=0.0$ is naturally referred to as \\textbf{pure neutron matter (PNM)}. Nuclei on Earth contain closely symmetric nuclear matter at densities close to nuclear saturation density $\\rho_0 \\approx 2.7 \\times 10^{14}$ g cm$^{-3} \\equiv 0.16$ fm$^{-3} = n_0$, where we use $n$ to refer to baryon number density. Thus experiment has constrained the properties of $E(\\sim n_0,\\sim0.5)$ to within relatively tight ranges, but the properties of PNM remain uncertain from an experimental standpoint. In the past decade, much experimental activity has been devoted to extending our knowledge of nuclear interactions to more neutron-rich systems and to higher and lower densities. Although we cannot produce pure neutron matter in the laboratory, we can produce matter with proton fractions as low as $x \\approx 0.3$ in certain neutron rich isotopes and in the products of heavy ion collisions. This allows us to obtain information on how $E(\\sim n_0, x)$ changes as $x$ decreases.  By expanding $E(n, x)$ about $x = 0.5$ using the isospin asymmetry variable $\\delta = 1-2x$, we can define a useful quantity called the \\emph{symmetry energy} S(n), \\be\\label{eq:eos1} E(n,\\delta) = E_{\\rm 0}(n) + S(n)\\delta^2 + ...; \\;\\;\\;\\;\\;\\;\\;\\; S(n) = {1 \\over 2}{\\partial^2 E(n,\\delta) \\over \\partial \\delta^2}\\bigg|_{\\delta=0}, \\ee \\noindent which encodes the change in the energy per particle of NM as one moves away from isospin symmetry. This allows extrapolation to the highly isospin asymmetric conditions in neutron stars. The simplest such extrapolation, referred to as the \\textbf{parabolic approximation (PA)}, truncates the expansion to second order, giving \\be \\label{PAapp} E_{\\rm PNM}(n) \\equiv E(n, \\delta=1) \\approx E_{\\rm 0}(n) + S(n) \\ee for the PNM EOS. Expanding the symmetry energy about $\\chi=0$ where $\\chi = \\frac{n-n_{\\rm 0}}{3n_{\\rm 0}}$ we obtain \\be\\label{eq:eos3} S(n) = J + L \\chi + \\half K_{\\rm sym} \\chi^{2} + ..., % \\ee \\noindent where $J$, $L$ and $K_{\\rm sym}$ are the symmetry energy, its slope and its curvature at saturation density.  Since neutron star matter contains a low fraction of protons, many inner crust and global stellar properties are sensitive to the symmetry energy parameters $J$,$L$, etc. To give a simple example, the pressure of PNM at saturation density is given by $P_{\\rm PNM}(n_0)$=$n_0L/3$. The pressure in the inner crust and outer core is dominated by neutron pressure so a strong correlation exists between the pressure in neutron stars near saturation density and $L$. Neutron star EOSs which have higher pressures are often referred to as `stiff'; lower pressure EOSs are referred to as `soft'. Thus, around 1 to 2$n_0$, `stiff' EOSs are associated with high values of $L$ and `soft' EOSs with low values of $L$.  That large uncertainties exist in characteristic NM parameters such as $J$ and $L$ is one reason why many model predictions of potential neutron star observables span such wide ranges. On the other hand, observations of neutron stars offer the opportunity to obtain constraints on NM parameters and hence on the underlying models of the nucleon-nucleon interactions \\cite{Lattimer2001}.  The following is a non-exhaustive list of (potentially) observable neutron star phenomena whose precise properties depend on the properties of the inner crust and mantle and of the star as a whole; some of these will be described in detail elsewhere in this book.  \\vspace{\\baselineskip}  $\\bullet$ \\hspace{5pt} \\textbf{Pulsar glitches.} Young pulsars spin down under the action of magnetic torque, their rotational energy powering the radiation beam. Many are observed to undergo occasional, sudden, spin-ups called glitches \\cite{Espinoza2011}. Proposed mechanisms include crust-cracking as the star attempts to adjust its shape to become more spherical \\cite{Baym1969} and angular momentum transfer from one internal component to another such as some part of the crust superfluid neutrons to the rigid part of the crust \\cite{Anderson1975, Link1999}, or a combination of both \\cite{Ruderman1998}. In such models, the size, frequency and post-glitch relaxation of the spin period depend on, among other microscopic properties, the crust and core sizes, moments of inertia and composition.  $\\bullet$ \\hspace{5pt} \\textbf{Free precession.} Certain pulsars exhibit long timescale periodic variation in their timing residuals, with periods of order years, suggestive of free precession of the star \\cite{Stairs2000,Shabanova2001}. Free precession is can arise from mechanically or magnetically supported crustal deformation \\cite{Cutler2002,Wasserman2003,Cutler2003}, and the period depends also on details of crust-core coupling, notably through the properties of the crust and core superfluid \\cite{Jones2001,Link2003,Link2006,Glampedakis2009}.  $\\bullet$ \\hspace{5pt} \\textbf{QPOs from SGR giant flares.} Quasi-periodic oscillations (QPOs) in the tails of light curves of giant flares from soft gamma-ray repeaters (SGRs) have been observed  \\cite{Israel2005, Watts2006, Strohmayer2005, Strohmayer2006}, and their frequencies lie in the range of possible torsional vibrations of the crust. The crust thickness, composition (through, e.g., the shear modulus) and the stellar size all affect the frequencies of such modes \\cite{SteinerWatts2009,Samuelsson2007,Andersson2009,Gearheart2011,Sotani2011}.  $\\bullet$ \\hspace{5pt} \\textbf{Neutron star cooling.} The crust thermalization timescale depends on the crust thickness as well as the thermal conductivity and specific heats arising from the heat transport mechanisms operating in the crust \\cite{Lattimer1994, Gnedin2001}. One intriguing possibility is the operation of the direct Urca process in the bubble phases of the mantle, where the delocalization of the protons may allow it \\cite{Gusakov2004}. The thickness of this layer plays an important role in determining how effective a cooling mechanism this might be.  $\\bullet$ \\hspace{5pt} \\textbf{Gravitational waves (GWs) from neutron stars.} A rich array of stellar oscillation modes are possible, some of which might generate GWs detectable on Earth \\cite{Andersson2011,Abbott2010}, or lead to other observational signatures such as limiting neutron star spin-up \\cite{Bildsten1998,Andersson1999}. Stability of modes can depend sensitively on the physics at the crust core interface and crust thickness, \\cite{Bildsten00,Andersson00,Lindblom00,Rieutord01,Peralta06,Glampedakis06,Wen11}. GWs can also be generated by a quadrupole deformation in the stellar shape supported, among other possibilities, by the elastic crust \\cite{Ushomirsky2000, Haskell2006}. Whether the crust is strong enough to support a large enough deformation to produce detectable gravitational waves depends on the shear modulus throughout the crust, and thus its composition (especially in the mantle where the mechanical properties are particularly uncertain) and the crust thickness and stellar size \\cite{Gearheart2011}. \\vspace{\\baselineskip}  In this chapter we will review the dependences of the composition, thickness of the crust and the mantle, and certain global neutron star properties on the symmetry energy parameters $J$ and $L$. The interplay of such relationships in modeling neutron star observables and the potential for obtaining astrophysical constraints on $J$ and $L$ will be illustrated by constructing consistent crust and core models based on a model of uniform nuclear matter whose symmetry energy parameters can be smoothly varied. We shall use the \\textbf{compressible liquid drop model (CLDM)} of crustal matter for its expediency; an outline of more sophisticated models will be given in the discussion at the end of the chapter. ",
        "conclusions": "We have constructed sets of neutron star EOSs that consistently encompass the inner crust and core, and include the crust composition and transition densities, using the compressible liquid drop model for the crust. These have been used to demonstrate the effect of the magnitude $J$ and slope $L$  of the symmetry energy at nuclear saturation density on microscopic and global crustal properties, and potential neutron star observables. The crust-core transition density and pressure, crustal composition, stellar and crustal mass, thickness and moment of inertia, torsional crust oscillation frequencies and maximum crust deformation all depend sensitively on both $J$ and $L$ within their experimentally constrained ranges. One of the dominant neutron star model dependences is therefore the correlation between $J$ and $L$ which constrains their possible values in $J-L$ space. Experimental and theoretical information about $L$ and $J$ and their correlation will continue to improve; in order to add neutron star observations to this investigation, consistent explorations of neutron star properties over the constrained ranges is of great importance. The sets of EOSs used in this paper are available to interested parties \\cite{mywebsite}.  The simplicity of the CLDM allows useful exploration of the dependence of composition and transition density on $J$ and $L$, but it possesses several drawbacks. Firstly shell effects within the nuclei, or arising from scattering of dripped neutrons off of nuclear clusters, are ignored. Such shell effects can dominate the determination of nuclear geometry, the equilibirium size of the nuclear clusters (which will proceed in discrete jumps corresponding to changes in the nuclear `magic' numbers with density) and the ordering of the pasta phases, as well as transport properties such as contributions to heat transport from nuclear components \\cite{Sandulescu2007} and entrainment of dripped neutrons by clusters \\cite{Chamel2005b}. Secondly, the WS approximation is expected to break down at when the nuclear separation becomes comparable to the cell size \\cite{Chamel2007}, which occurs in the mantle. Thirdly, effects that act over ranges greater than the unit cell are not consistently accommodated in the CLDM; longer range electron screening, larger scale self-organization of pasta phases and long wavelength transport effects are all unaccounted for.   \\begin{figure}[!t]\\label{fig:12} \\begin{minipage}{6.6cm} \\begin{center} \\includegraphics[width=3.3cm,height=3cm]{Fig10a.pdf}\\includegraphics[width=3.3cm,height=3cm]{Fig10b.pdf}\\\\ \\includegraphics[width=3.3cm,height=3cm]{Fig10c.pdf}\\includegraphics[width=3.3cm,height=3cm]{Fig10d.pdf} \\end{center} \\caption{Integrated neutron density in a cubic cell at densities of $n$=0.04, 0.06, 0.08 and 0.1 fm$^{-3}$ (top left to bottom right) calculated using the 3DHF method \\cite{Newton2011} with the SLy4 Skyrme parameterization.} \\end{minipage} \\hfill \\begin{minipage}{6.6cm} \\begin{center} \\includegraphics[width=6.6cm,height=6cm]{Fig10e.pdf} \\end{center} \\caption{Transition densities obtained using the 3DHF method compared to those obtained using the dynamical method of \\cite{Xu2009} for the SLy4, SII, SkM* and SkMp Skyrme parameterizations.} \\end{minipage} \\vspace{10pt} \\end{figure}  Some of these effects can be taken into account by more sophisticated crust models. The Thomas-Fermi and Extended Thomas-Fermi methods (e.g \\cite{Buchler1971,Oyamatsu1993,Cheng1997}) are semi-classical models employing the local density approximation which allows the nuclear surface energy to be calculated self-consistently with the bulk nuclear energy. Shell corrections can be self-consistently added using the Strutinsky integral method \\cite{Onsi2008}. The 1D-Hartree-Fock method is fully microscopic and accounts for surface and shell energies self-consistently (e.g. \\cite{Sandulescu2007,Negele1973,Montani2004,Baldo2005}), but it describes only spherically symmetric configurations and is thus constrained by the spherical Wigner-Seitz approximation. This restriction is lifted in the more computationally demanding 3D-Hartree-Fock (3DHF) method \\cite{Magierski2002,Gogelein2008}, allowing a self-consistent probe of the shape-phase-space of pasta layers with shell effects included. As example, some results of a 3DHF model \\cite{Newton2009} are shown in Figs.~10 and ~11. The former shows the local neutron density is plotted over the unit cell at four different densities encompassing the crust-core transition, showing the evolution of nuclear shape from spherical through bubble to uniform matter. The latter shows the resulting transition densities for four different parameterizations of the Skyrme interaction, compared to the dynamical method of \\cite{Xu2009}. The Hartree-Fock method is naturally extended to include pairing effects self-consistently (Hartree-Fock-Bugoliubov). Longer range effects can be simulated via the semi-classical quantum molecular dynamics method \\cite{Maruyama1998,Horowitz2004,Watanabe2001,Sonoda2007}. Between these models, a complete physical description of the crust can be built up. Many of the methods mentioned are much more time-consuming than the CLDM, making a wide-ranging survey of nuclear matter parameters unwieldy. However, if we know the quantities for which the CLDM provides a reasonable estimate, and to what densities it remains reasonable, one can use the CLDM as a useful guide for more realistic calculations. Ultimately, the goal should be to have all relevant microscopic inputs to neutron star models calculated consistently with the nuclear matter EOS; much work still needs to be done in this direction."
    },
    "1112/1112.5312_arXiv.txt": {
        "abstract": "Recent $\\gamma$-ray observations suggest that the $\\gamma$-ray millisecond pulsar (MSP) population is separated into two sub-classes with respect to the pair multiplicity. Here, we calculate the cosmic ray electron/positron spectra from MSPs. Based on the assumption of the equipartition in the pulsar wind region the typical energy of electrons/positrons ejected by a MSP with the pair multiplicity of order unity is $\\sim50$ TeV. In this case, we find that a large peak at 10 - 50 TeV energy range would be observed in the cosmic ray electron/positron spectrum. Even if the fraction of pair starved MSPs is 10\\%, the large peak would be detectable in the future observations. We also calculate the contribution from MSPs with high pair multiplicity to the electron/positron spectrum. We suggest that if the multiplicity of dominant MSP population is $\\sim 10^3$, electrons/positrons from them may contribute to the observed excess from the background electron/positron flux and positron fraction. ",
        "introduction": "The {\\it Fermi} Gamma-Ray Space Telescope has detected $\\gamma$-ray pulsed emissions from more than twenty millisecond pulsars (MSPs) \\citep{Ab11}, which have a rotation angular frequency $\\Omega\\sim 10^3$s$^{-1}$ and a stellar surface magnetic field $B_s\\sim 10^{8.5}$G. The detection of the GeV emissions from a pulsar magnetosphere means that electrons and positrons are accelerated to more than $\\sim$ TeV by the electric field parallel to the magnetic field, which arises in a depleted region of the Goldreich-Julian (GJ) charge density \\citep{GJ69}. The $\\gamma$-ray light curve is an important tool for probing the particle acceleration process in the pulsar magnetosphere. Therefore, the $\\gamma$-ray emission region has been explored by comparing theoretical models such as polar cap \\citep{DH96}, outer gap \\citep{CHR86} and slot gap models \\citep{MH04} with the observed light curve (e.g., Venter, Harding \\& Guillemot 2009; Romani \\& Watters 2010; Kisaka \\& Kojima 2011).  \\citet{VHG09} fitted the pulse profiles of the {\\it Fermi} detected MSPs with the geometries of $\\gamma$-ray emission region predicted by different theoretical models. They found that the pulse profiles of six of eight MSPs could be fitted by the geometries of either the outer gap or the slot gap model, as was the case of canonical pulsars. They interpreted that copious pairs are produced in the magnetosphere of these MSPs.  However, \\citet{VHG09} also found that the pulse profiles of remaining two MSPs show the unusual behavior in the $\\gamma$-ray light curves and could not be fitted by the geometry of either the outer gap or the slot gap models. They proposed that these unusual light curves could be fitted by the pair starved polar cap model \\citep{MH04b}, in which the multiplicity of the pairs is not high enough to completely screen the electric field above the polar cap, and the particles are continuously accelerated up to high altitude over the entire open field line region. Thus, from the model fitting of the $\\gamma$-ray light curves, \\citet{VHG09} suggested that the $\\gamma$-ray MSP population is separated into two sub-classes.  The important fact is that radio pulsed emission is also detected from all currently detected $\\gamma$-ray MSPs and remarkably similar to that from canonical pulsars. The pulsar radio emission is a highly coherent process because the brightness temperature is extremely high. In the theoretical models of the radio emission mechanisms, some authors have believed the conditions that there are a highly relativistic primary beam with the large Lorentz factor ($\\sim 10^7$) and the number density nearly equal to the GJ density, and the secondary electron/positron plasma with relatively small bulk streaming Lorentz factor ($\\sim 10$ - $10^3$) and the large pair multiplicity ($\\sim 10^3$ - $10^5$) in the radio emitting region (e.g., Melrose 1995; Lyutikov, Blandford \\& Machabeli 1999; Gedalin, Gruman \\& Melrose 2002). However, the existence of pair starved MSPs suggests that the radio emission mechanisms should be insensitive to the particle number density down to sub-GJ number density. The pulsar radio emission mechanism is still poorly understood, so that the observationally-based constraints are valuable \\citep{M95}. Therefore, another verification for the extent of the MSP multiplicity, especially the existence of the pair starved MSPs is important for the pulsar radio emission mechanisms.  Recently, HESS has discovered a new TeV source, which is located in the close vicinity of the globular cluster Terzan 5 \\citep{HE11}. Several globular clusters, including Terzan 5 also emit GeV $\\gamma$-ray \\citep{Ab09b, KHC10, Ab10b, Ta11}, which may plausibly be due to a number of MSPs residing in these clusters \\citep{HUM05, VDC09}. Thus, inverse Compton scattering by the high-energy particles ejected from MSPs are proposed for the origin of the observed TeV emission \\citep{BS07, VDC09}. The high-energy electron/positron spectrum ejected from MSPs would be a useful probe for the multiplicity of the MSPs. However, only from the TeV spectra, we cannot distinguish two models \\citep{BS07, VDC09}, which assume different pair multiplicities \\citep{HE11}. Another way to investigate the electron/positron spectrum ejected from MSPs is its direct measurement. Since high-energy electrons/positrons can propagate only about a few kpc due to the energy losses by the synchrotron and the inverse Compton emission, the direct detection of the electrons/positrons ejected from MSPs in the globular clusters is unlikely. However, for the following reasons, we may detect those from nearby MSPs.  MSPs have much lower spin-down luminosity than canonical pulsars. \\citet{BVD08} investigated the possible contribution of the nearby MSP, PSR J0437-4715 to the cosmic ray electron/positron spectrum. They concluded that unlike canonical pulsars such as Geminga pulsar, the contribution from a MSP to the observed electron/positron flux is negligible. However, since the lifetime of MSPs is much longer than canonical pulsars ($>10^{10}$ yr), there should be much more nearby active MSPs. Furthermore, Kashiyama, Ioka \\& Kawanaka (2011; hereafter KIK11) pointed out that since white dwarf pulsars have long lifetime and continue to inject the electrons/positrons after the nebulae stop expanding, the adiabatic energy losses of electrons/positrons in the pulsar wind nebula region are negligible. Also the synchrotron cooling of electrons/positrons is so small and the high-energy electrons/positrons can escape the nebulae without losing much energy. Although they consider the case of the white dwarfs, their results are also applicable to MSPs. Therefore MSPs could potentially contribute to the observed high energy cosmic ray electrons/positrons and will be detectable by the next generation experiments, such as CALET \\citep{To08} and CTA \\citep{CTA}.  In this paper, we investigate the contribution of electrons/positrons ejected from the MSPs to the observed cosmic ray spectrum. In section 2, we apply KIK11 model to the case of MSPs. We estimate the typical energy of electrons/positrons from the MSPs and show that during the propagation in pulsar wind nebulae, the adiabatic losses and radiative cooling of electrons/positrons are not so large. We also describe the propagation in the interstellar medium (ISM). In section 3, we calculate the energy spectrum of cosmic ray electrons/positrons from the MSPs and show the possibility that the electrons/positrons from these MSPs are detectable for the future observations. ",
        "conclusions": "First, we calculate the cosmic ray electron/positron spectra from the pair starved MSPs. We set the pair multiplicity $\\kappa= 1$, the lifetime $\\tau= 5\\times 10^{10}$yr, the total energy $E_{\\rm rot}= 10^{52}$ erg, the local birth rate $R = 3\\times 10^{-9}$ yr$^{-1}$ kpc$^{-2}$ and the fraction of the lost energy due to synchrotron emission 30\\% for each MSP. We assume that each MSP has the same value of the parameters ($B_0=10^{8.5}{\\rm G}, \\Omega=10^3{\\rm s}^{-1}, R=10^6{\\rm cm}$), because most MSPs have the almost same spin-down luminosity. For the injection distribution function, we assume mono-energetic distribution eq.(\\ref{mono}) with the energy $\\varepsilon_e=50$ TeV. Even if we consider the power-law distribution eq.(\\ref{pl}), the energy range of the distribution is small because the allowed range of the cutoff energy should be only 50 - 80 TeV due to the observed constraint by KASKADE/GRAPES/CASA-MIA \\citep{KY09}. Thus the distribution should be nearly mono-energetic distribution. Note that our local MSP birth rate is based on the MSP local surface density of 38$\\pm$ 16 pulsars kpc$^{-2}$ for 430 MHz luminosity above 1 mJy kpc$^2$ \\citep{L08}. Actually, now 23 MSPs are detected within 1kpc \\citep{ATNF}. We can only detect MSPs that have the radio beam directed toward us and the radio flux larger than the threshold of detectors. However, the cosmic ray electrons/positrons ejected from MSPs are distributed isotropically because of the effect of Galactic magnetic field, so that a large number of MSPs will contribute to the observed electron/positron spectrum. Therefore, our local MSP birth rate corresponds to the lower limit for the current radio observations.  In figure \\ref{fig:3.1}, the electron/positron flux from multiple pair starved MSPs is shown. Thick solid, dashed and dash-dotted lines show the total electron/positron spectra if the fraction of the pair starved MSPs is 100\\%, 25\\% and 10\\%, respectively. Thin solid line shows the contribution of electrons/positrons from the pair starved MSPs when the fraction is 100\\%. The background flux is shown as a thin dashed line. We adopt the background model of an exponentially cutoff power law with an index of -3.0 and a cutoff at 1.5 TeV, which is similar to that shown in \\citet{Ah08} and reproduces the data in $\\sim$10 GeV-1 TeV well. It is very interesting that there is a large peak at 10-50 TeV energy range. The existence of this peak cannot be ruled out from the current observations. The high-energy component is more enhanced for the long-duration sources \\citep{AAV95, KIN10}. This is because the longer the duration of injection is, the larger fraction of fresh electrons/positrons are expected to reach the Earth without losing their energy so much during the propagation. MSPs are the continuously injecting sources with the duration as long as the cosmic age, so that the spectrum from them has nearly the same shape with the injection spectrum with the soft energy tail component. This is the difference from other sources such as young pulsars, whose typical duration is only $\\sim 10^4$-$10^5$ yrs. Fitting result of \\citet{VHG09} showed that the fraction of the pair starved MSPs is 25\\%, although they have only eight samples. Even if the fraction is 10\\%, the flux is $\\sim 20$m$^{-2}$ s$^{-1}$ sr$^{-1}$ GeV$^2$ at 10 TeV. In this case, we can detect the electron/positron flux with near future missions such as CALET (we assume the geometrical factor times the observation time $\\sim 5$ yrs as $\\sim 220$ m$^2$ sr days) because the predicted electron/positron flux is sufficiently large. It was considered that the number of astrophysical sources contributing to the above several TeV energy range is quite small according to the birth rate of the supernovae and the canonical pulsars in the vicinity of the Earth \\citep{Ko04, KIN10, KIOK11}. However, we find that it is possible for multiple pair starved MSPs to contribute to the 10 TeV energy range in the electron/positron spectrum. Therefore, if the anisotropy of the observed electrons/positrons are weak in the 10 TeV energy range, we suggest that the pair starved MSPs may contribute to the spectrum significantly.  Next, we also investigate the contribution of MSPs with high pair multiplicity to the observed cosmic ray electrons/ positrons. In this model, we assume that the injection function of these MSPs is power-law distribution with index $\\alpha=2$, the cutoff energy $\\varepsilon_{e, {\\rm cut}}=1$TeV and the minimum energy $\\varepsilon_{e, \\min}=1$GeV. In this case, the pair multiplicity is $\\sim 2000$. Other parameters are the same values as in the case of the pair starved MSPs. We assume that the fraction of MSPs with high multiplicity is 100\\%. The results are shown in figure \\ref{fig:3.2} for the electron/positron spectrum and in figure \\ref{fig:3.3} for the positron fraction. Both figures show that electrons/positrons ejected from MSPs with high multiplicity partially contribute to the excess from background flux observed by PAMELA, HESS and {\\it Fermi}. In this energy range, the other sources such as canonical pulsars \\citep{S70, AAV95, CCY96, ZC01, Ko04, G07, Bu08, P08, HBS09, YKS09, MCG09, Gr09, KIN10, HGH10} would also contribute to the observed electron/positron spectrum. Note that even if other sources are dominant for the observed excess, the total spectrum added to the contribution of MSPs with high multiplicity does not significantly exceed the observed electron/positron spectrum."
    },
    "1112/1112.5548_arXiv.txt": {
        "abstract": "{Rotation is known to have a strong impact on the nucleosynthesis of light elements in massive stars, mainly by inducing mixing in radiative zones. In particular, rotation boosts the primary nitrogen production, and models of rotating stars are able to reproduce the nitrogen observed in low-metallicity halo stars. } {Here we present the first grid of stellar models for rotating massive stars at low metallicity, where a full s-process network is used to study the impact of rotation-induced mixing on the neutron capture nucleosynthesis of heavy elements. } {We used the Geneva stellar evolution code that includes an enlarged reaction network with nuclear species up to bismuth to calculate 25~$M_\\odot$ models at three different metallicities ($Z = 10^{-3}, 10^{-5}$, and $10^{-7}$) and with different initial rotation rates. } {First, we confirm that rotation-induced mixing (shear) between the convective H-shell and He-core leads to a large production of primary $^{22}$Ne (0.1 to 1\\% in mass fraction), which is the main neutron source for the s process in massive stars. Therefore rotation boosts the s process in massive stars at all metallicities. Second, the neutron-to-seed ratio increases with decreasing $Z$ in models including rotation, which leads to the complete consumption of all iron seeds at metallicities below $Z = 10^{-3}$ by the end of core He-burning. Thus at low $Z$, the iron seeds are the main limitation for this boosted s process. Third, as the metallicity decreases, the production of elements up to the Ba peak increases at the expense of the elements of the Sr peak. We studied the impact of the initial rotation rate and of the highly uncertain $^{17}$O$(\\alpha,\\gamma)$ rate (which strongly affects the strength of $^{16}$O as a neutron poison) on our results. This study shows that rotating models can produce significant amounts of elements up to Ba over a wide range of $Z$, which has important consequences for our understanding of the formation of these elements in low-metallicity environments like the halo of our galaxy and globular clusters. Fourth, compared to the He-core, the primary $^{22}$Ne production induced by rotation in the He-shell is even higher (greater than 1\\% in mass fraction at all metallicities), which could open the door for an explosive neutron capture nucleosynthesis in the He-shell, with a primary neutron source. } {} ",
        "introduction": "Massive stars in the early universe are different from the ones observed today. Their low metal content makes them more compact and rotate faster than their equivalents found in the Milky Way. This view is supported by observations of an increasing Be/B-type star ratio with decreasing metallicity \\citep{2007A&A...472..577M} and by more rapidly rotating, massive stars in the SMC compared to the Milky Way \\citep{2008A&A...479..541H}. Rapidly rotating stellar models at low $Z$ have been calculated by \\citet{2006A&A...447..623M} and \\citet{2007A&A...461..571H}. In these models, primary nitrogen yields are much larger than in non-rotating models. When yields from these rotating models are used as input in chemical evolution models, a nice fit of the N/O in very metal poor halo stars \\citep[see e.g.][]{2005A&A...430..655S} can be obtained \\citep{2006A&A...449L..27C}. This provides a strong support for the occurrence of rotation-induced mixing at low $Z$ and for our models. The primary nitrogen production in rotating low-$Z$ stellar models is accompanied by the primary production of other isotopes like $^{13}$C, and especially $^{22}$Ne, which is the neutron source for s process in massive stars \\citep[e.g.][and references therein]{2011RvMP...83..157K}.  A first attempt to assess the impact of rotation on the s-process nucleosynthesis in low-$Z$ massive rotating stars was made by \\citet{2008ApJ...687L..95P}, who investigated the impact of primary $^{22}$Ne in a parametrised way. The recent observation of large s-process enhancements in one of the oldest globular clusters in the bulge of our galaxy supports the view that massive stars could indeed also be important sources for these elements \\citep{2011Natur.472..454C}, highlighting the need for comprehensive calculations of the s process in low-$Z$ massive rotating stars. In this letter, we present the first such low-$Z$, massive rotating star models including a full s-process network.  We describe our models in Sect.~\\ref{sec:mixing}. The impact of rotation on the s process in massive 25~$M_\\odot$ stars at different metallicities is discussed in Sect.~\\ref{sec:sproc}. We give our conclusions in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} We have calculated the first complete stellar models including a full s-process network to study the effects of rotation on the s process in massive stars in low-metallicity environments. Our models confirm that rotation-induced mixing in radiative zones leads to a primary production of $^{22}$Ne in the He core over the wide $Z$ range studied ($\\mathrm{[Fe/H]}=-5.8$, $-3.8$, and $-1.8$). From $0.1$ to $1$\\% of $^{22}$Ne in mass fraction is produced in central He-burning. The large primary production of $^{22}$Ne also occurs in the He-shell in all the rotating models calculated, i.e. 1-2\\% in mass fraction at all $Z$. This can have a strong impact on the supernova nucleosynthesis in He-burning shell, since it could provide a primary neutron source for neutron capture nucleosynthesis at all values of Z. We investigated the impact of the $^{17}$O$(\\alpha,\\gamma)$-rate on the strength of $^{16}$O as a neutron poison and find the impact of current uncertainties still significant, thus deserving further experimental studies.  The s process in massive stars with rotation contributes significantly to the production of elements up to Sr above $\\mathrm{[Fe/H]}\\approx-2$, the main limiting factor being the iron seeds at low $Z$. The production of elements up to Ba is sensitive to the initial rotation rate, and we generally obtain $\\mathrm{[Sr/Ba]}>0$ from low-$Z$ massive rotating star s process. Pb is only produced in very extreme cases, but with $\\mathrm{[Pb/Sr]}\\apprle-1$ massive stars cannot be responsible for CEMP-s stars. A similar spread for [Sr/Ba] to the one used in \\citet{2011Natur.472..454C} was obtained to explain the unique abundance of several stars in one of the oldest globular clusters in the galactic bulge, thus supporting the conclusions of that study. Further stellar and galactic chemical evolution models will assess the full impact of this boosted s~process."
    },
    "1112/1112.4778_arXiv.txt": {
        "abstract": "In contrast to the prevailing view in the literature, it is shown that even extremely stiff sets of ordinary differential equations may be solved efficiently by explicit methods if limiting algebraic solutions are used to stabilize the numerical integration. The stabilizing algebra differs essentially for systems well-removed from equilibrium and those near equilibrium.  Explicit asymptotic and quasi-steady-state methods that are appropriate when the system is only weakly equilibrated are examined first. These methods are then extended to the case of close approach to equilibrium  through a new implementation of partial equilibrium approximations. Using stringent tests with astrophysical thermonuclear networks, evidence is provided that these methods can deal with the stiffest networks, even in the approach to equilibrium, with accuracy and integration timestepping comparable to that of implicit methods. Because explicit methods can execute a timestep faster and scale more favorably with network size than implicit algorithms, our results suggest that algebraically-stabilized explicit methods might enable integration of larger reaction networks coupled to fluid dynamics than has been feasible previously for a variety of disciplines. ",
        "introduction": "Introduction}  Problems from many disciplines require solving large coupled reaction networks. Representative examples include  reaction networks in combustion chemistry \\cite{oran05}, geochemical cycling of elements \\cite{magick}, and thermonuclear reaction networks in astrophysics \\cite{hix05,timmes}.  The differential equations used to model these  networks usually exhibit stiffness, which arises from multiple timescales in the problem that differ by many orders of magnitude \\cite{oran05,gear71,lamb91,press92}. Sufficiently complex physical systems often involve important processes operating on widely-separated timescales, so realistic problems tend to be at least moderately stiff. Some, such as astrophysical thermonuclear networks, are extremely stiff, with 10--20 orders of magnitude between the fastest and slowest timescales in the problem. Our concern here is with stiffness as a numerical issue,  but we remark that stiffness can have important physical implications because complex processes often function as they do precisely because of the coupling of very slow and very fast scales within the same system.  Books on numerical and computational methods routinely state \\cite{oran05,lamb91,press92} that stiff systems cannot be integrated efficiently using explicit finite-difference methods because of stability issues: for an explicit algorithm, the maximum stable timestep is set by the fastest timescales, even if those timescales are peripheral to the main phenomena of interest. The standard resolution of the stiffness problem uses implicit integration, which is stable for stiff systems but entails substantial computational overhead because it requires the inversion of matrices at each integration step. Because of the matrix inversions, implicit algorithms  tend to scale from quadratically to cubically with network size unless favorable matrix structure can be exploited.  Thus, implicit methods can be expensive  for large networks.  A simple but instructive example of stiffness is provided by the CNO cycle for conversion of hydrogen to helium, which powers main-sequence stars more massive than the Sun (Fig.~\\ref{fig:cnoCycle}). In the CNO cycle the fastest rates under characteristic stellar conditions are $\\beta$-decays with half-lives $\\sim 100$ seconds, but to track the complete evolution of main-sequence hydrogen burning may require integration of the network for hydrogen burning over timescales as large as billions of years ($\\sim 10^{16}$ seconds). If one tries to implement this integration using explicit forward differencing, the largest stable integration timestep will be set by the fastest rates and will be of order $10^2$ seconds.  Thus $10^{14}$ or more explicit integration steps could be required to integrate the CNO cycle to hydrogen depletion. Conversely, typical implicit integration schemes can take stable and accurate timesteps equal to 1-10\\% of the local time over most of the integration range, and would  compute the above numerical integration in a few hundred implicit steps. By virtue of examples such as this, it is broadly accepted that explicit methods are not viable for stiff networks. To quote the authoritative reference {\\em Numerical Recipes} \\cite{press92}, ``For stiff problems we {\\em must} use an implicit method if we want to avoid having tiny stepsizes.'' \\putfig {cnoCycle} {0pt} {\\figdn} {0.90} {fig1.eps} {The CNO (carbon--nitrogen--oxygen) cycle. On the left side the main branch of the cycle is illustrated with solid arrows and a side branch is illustrated with dashed arrows. On the right side, the main branch of the CNO cycle is illustrated in more detail.}  Our main interest lies in simulations where the reaction network is one part of a broader problem.  Let us take as representative astrophysical thermonuclear networks coupled to multidimensional hydrodynamics. The hydrodynamical evolution controls network conditions (temperature, density, \\ldots), and the network influences the hydrodynamics through energy production and modification of composition. Solution of large networks by the usual means is costly in this context and even ambitious simulations  use only small networks, or replace the network entirely by parameterization. Then a more realistic network is used in a separate ``post-processing'' step, where fixed hydrodynamical profiles computed in the original simulation specify the variation of temperature and density with time. Such approximations are especially at issue for problems like Type Ia supernovae, which are 3D asymmetric explosions powered by a complex reaction network releasing energy greater than that of a large galaxy on a timescale of order 1 s.   Astrophysical reaction networks have been used to illustrate, but problems in various fields exhibit a similar complexity.  For example, in astrochemical kinetics large chemical evolution networks must be modeled in dynamical environments such as contracting molecular clouds, or in combustion chemistry the burning networks are strongly coupled to dynamical simulations of the air and fuel mixture. Realistic networks in all such applications may be quite large. Modeling combustion of larger hydrocarbon molecules or soot formation can require hundreds to thousands of reacting species with up to 10,000 reactions \\cite{oran05}, and realistic networks for supernova explosions imply hundreds to thousands of nuclear isotopes with tens of thousands of reaction couplings \\cite{hix05}. In all such cases  current techniques do not permit the coupling of realistic reaction networks to the full dynamics and highly-schematic reaction networks are used in even the most realistic contemporary simulations. ",
        "conclusions": "Conclusions}  This paper demonstrates that algebraically-stabilized explicit integration is capable of timesteps competitive with those of implicit methods for various extremely-stiff reaction networks. Since explicit methods can execute a timestep faster than an implicit method in a large network, our results suggest that algebraically-stabilized explicit algorithms may be capable of performing as well as, or even substantially outperforming, implicit integration in a variety of moderate to extremely stiff applications. Because of the linear scaling with reaction network size for explicit methods, this fundamentally new view of explicit integration for stiff equations is particularly important for applications in fields where more realistic---and therefore larger---reaction networks are required for physical simulations. Arguably, this means almost all scientific and technical disciplines, since the sizes of reaction networks being used in simulations to this point have been dictated more often by what was feasible than by what was physical.  Of particular significance is that these new explicit methods might permit coupling of more physically-realistic reaction kinetics to fluid dynamics simulations in a variety of disciplines.  \\begin{ack} I thank Jay Billings, Reuben Budiardja, Austin Harris, Elisha Feger, and Raph Hix for help with some of the calculations.  Discussions with Raph Hix, Bronson Messer, Kenny Roche, Jay Billings, Brad Meyer, Friedel Thielemann, Michael Smith, and Tony Mezzacappa have been useful in formulating the ideas presented here, and I thank Bronson Messer for a careful reading of the manuscript. Research was sponsored by the Office of Nuclear Physics, U.S. Department of Energy. \\end{ack}  \\vfill  \\clearpage  \\setcounter{section}{0} \\appendix"
    },
    "1112/1112.1688_arXiv.txt": {
        "abstract": "Astronomy has long had a working network of archives supporting the curation of publications and data. The discipline has already created many of the features which perplex other areas of science: (1) data repositories: (supra)national institutes, dedicated to large projects; a culture of user-contributed data; practical experience of long-term data preservation; (2) dataset identifiers: the community has already piloted experiments, knows what can undermine these efforts, and is participating in the development of next-generation standards; (3) citation of datasets in papers: the community has an innovative and expanding infrastructure for the curation of data and bibliographic resources, and through them a community of authors and editors familiar with such electronic publication efforts; as well, it has experimented with next-generation web standards (e.g. the Semantic Web); (4) publisher buy-in: publishers in this area have been willing to innovate within the constraints of their commercial imperatives. What can possibly be missing?  Why don't we have an integrated framework for the publication and preservation of all data products already?  Are there technical barriers?  We don't believe so. Are there cultural or commercial forces inhibiting this?  We aren't aware of any. This Birds of a Feather session (BoF) attempted to identify existing barriers to the creation of such a framework, and attempted to identify the parties or groups which can contribute to the creation of a VO-powered data-publishing framework. ",
        "introduction": "This BoF session provided a forum for data providers, publishers, librarians and scientists to explore the issues surrounding the preservation, identification and citation of data products in astronomy.  These are small but critical steps towards the ultimate goal of identifying the right practices, resources, and incentives to ensure that the entire research lifecycle in astronomy is properly captured and described in the coming era of data-intensive astronomical research. The authors of this paper (and BoF panelists) were chosen to represent varied perspectives from different corners of our community, including: the ADS, the primary literature portal in astronomy; the CDS, one of the largest data curation hubs in the Virtual Obsevatory (VO); the AAS, the largest publisher of scholarly literature in astronomy; and the University of Glasgow, an institution involved in data management and resource discovery for the VO. The following sections highlight some of the topics covered in the BoF. ",
        "conclusions": "That there is a need for curated datasets, preserved indefinitely for scholarly purposes, is axiomatic by now. The details of a fully-fledged data preservation environment are difficult to discern at this point, but there are sensible places to begin, and from such places the mechanisms can evolve safely. The scientific and scholarly communities are vigorously pursuing possibilities, and there are many fruitful opportunities, such as this session at ADASS, for interested parties to engage in discussions about the promises and the challenges of broadly-scaled data preservation."
    },
    "1112/1112.3191_arXiv.txt": {
        "abstract": "We announce the identification of a proper motion companion to the star HII~1348, a K5\\,V member of the Pleiades open cluster.  The existence of a faint point source 1$\\farcs$1 away from HII~1348 was previously known from adaptive optics imaging by Bouvier et al.  However, because of a high likelihood of background star contamination and in the absence of follow-up astrometry, Bouvier et al.\\ tentatively concluded that the candidate companion was not physically associated with HII~1348. We establish the proper motion association of the pair from adaptive optics imaging with the Palomar 5~m telescope.  Adaptive optics spectroscopy with the integral field spectrograph OSIRIS on the Keck 10~m telescope reveals that the companion has a spectral type of M8$\\pm$1.  According to substellar evolution models, the M8 spectral type resides within the substellar mass regime at the age of the Pleiades.  The primary itself is a known double-lined spectroscopic binary, which makes the resolved companion, HII~1348B, the least massive and widest component of this hierarchical triple system and the first substellar companion to a stellar primary in the Pleiades. ",
        "introduction": "As one of the nearest young \\citep[125~Myr;][]{stauffer1998} open clusters, the Pleiades have long been recognized as an important astrophysical laboratory for studying stellar evolution and the dynamics of stellar associations. Multiplicity studies of the Pleiades have focused both on stellar \\citep[e.g.,][]{stauffer1984, mermilliod1992, bouvier1997} and substellar \\citep[e.g.,][]{martin2000, bouy2006} members of the cluster.  However, mixed star--brown dwarf systems have not been identified. One of the most extensive studies is the adaptive optics imaging survey of \\citet{bouvier1997}. Conducted in the near-IR, it covered 144 G and K stars to a relatively shallow depth. As a result, its sensitivity encompassed only stellar and massive substellar companions. A systematic high-contrast imaging survey of the Pleiades on a similar scale but at a higher sensitivity has not been performed since. Consequently, no substellar companions are known to $>$0.2~\\Msun\\ stars in the Pleiades. With the frequency of wide substellar companions to field-aged Sun-like stars now estimated at $\\approx$\\,3\\,\\% \\citep[0.012\\,--\\,0.072~\\Msun\\ brown dwarfs in 28\\,--\\,1590\\,AU orbits;][]{metchev2009}, the frequency of brown dwarf secondaries around Sun-like stars in the Pleiades is expected to be comparable.   In the present paper, we announce the identification of a low mass companion to the Pleiad HII~1348, a K5\\,V double-lined spectroscopic binary \\citep[hereafter refered to as the primary or HII~1348A]{queloz1998}. The faint companion, HII~1348B, was already detected by \\citet{bouvier1997}.  However, without follow-up astrometric observations and due to a non-negligible probability of background star contamination, \\citet{bouvier1997} conservatively assumed that the candidate companion was an unrelated background star.  The astrometric measurements confirm the proper motion association of the pair, and AO spectra obtained at Palomar and Keck reveal that the companion has a spectral type of M8. % ",
        "conclusions": "HII~1348B is a new M8 brown dwarf member of the Pleiades, and the first substellar companion discovered around a Pleiades star.  Given that no other substellar companions were discovered in the \\citet{bouvier1997} survey at similar or wider separations, it is worth considering whether HII~1348B may be unusually weakly bound, compared to other binary systems in the Pleiades or in the field.  \\citet{bouvier1997} found a total of 28 stellar binaries in the Pleiades in their CHFT AO survey.  HST surveys of very low mass stars and brown dwarfs conducted by \\citet{martin2003} and \\citet{bouy2006} revealed three additional binaries. In Figure \\ref{fig:Ebind} we compare the binding energies of these systems, and those of field A--M binaries \\citep{close1990, close2003, close2007}, to the binding energy of HII~1348A/B.  As can be seen, HII~1348A/B sits in the middle of the locus for stellar binaries, and is comparably or more tightly bound even than the three very low mass Pleiades binaries.  What is more, substellar companions up to 10 times further away from their primaries would still be well above the minimum stellar binding energy in the Pleiades.  The dearth of known brown dwarf companions to stars in the Pleiades may thus be attributable to the lack of a follow-up sensitive and comprehensive high-contrast imaging survey of the cluster.  Small samples of Pleiades stars have since been observed in deep surveys by \\citet[23 stars;][]{metchev2009} and \\citet[14 stars;][]{tanner2007}, with no new brown dwarf companion detections.  However, a much more comprehensive survey is needed to reveal brown dwarf companions with any statistical confidence \\citep{metchev2009}.  Current AO systems at Keck or Gemini North should allow the detection of companions with masses down to $\\sim$0.03\\,$\\Msun$ at separations larger than $\\sim$60AU \\citep[0$\\farcs$5;][]{lafreniere2007, metchev2009b}. Using the Gemini North AO system Altair in combination with NIRI, \\citet{lafreniere2007} showed that companions up to 9.5\\,mag fainter can be detected at separations of $>$0$\\farcs$5 from the primary. Accounting for the narrow band filter used during the observation, brown dwarfs down to 0.03\\,$\\Msun$ should be detectable around bright (V$<$12\\,mag) Pleiades stars. Likewise, using angular differential imaging \\citep[ADI,][]{marois2006} in combination with the Keck AO system, \\citet{metchev2009b} demonstrated that companions up to $\\Delta$H=10.5\\,mag fainter than the primary can be detected at separations larger than 0$\\farcs$5. Thus, at the distance of the Pleiades brown dwarfs down to 0.02\\,$\\Msun$ should be detectable at separations larger than $\\sim$60AU. Upcoming instruments like the Gemini planet imager \\citep{macintosh2008, mcbride2011} and the extreme AO system at Palomar \\citep[PALM3000,][]{bouchez2008} will enable detections of brown dwarf companions at even smaller separation down to 0$\\farcs$2.  Given the importance of the Pleiades as an age and a dynamical benchmark for stellar and substellar evolution, a deep AO survey could have a high impact in delivering substellar objects much fainter and cooler than the ones presently known from wide-area surveys. While wide-area surveys are more efficient in discovering large numbers of substellar candidate members, the limited precision of seeing-limited astrometry, and the increasingly challenging radial velocity measurements at fainter magnitudes, prevent the unequivocal confirmation of the lower mass candidates as bona-fide members.  The factor of $\\sim$100 higher precision attainable over narrow angles with AO leaves little doubt about the astrometric association of candidate binaries, and is thus also an excellent tool for confirming the cluster membership of even the lowest-mass, faintest companions."
    },
    "1112/1112.3969_arXiv.txt": {
        "abstract": "{ During the last decade, the FORS1 instrument of the ESO Very Large Telescope has been extensively used to study stellar magnetism. A number of interesting discoveries of magnetic fields in several classes of stars have been announced, many of which obtained at a $\\sim 3\\,\\sigma$ level; some of the discoveries are confirmed by measurements obtained with other instruments, some are not. } { We investigate the reasons for the discrepancies between the results obtained with FORS1 and those obtained with other instruments. } { Using the ESO FORS pipeline, we have developed a semi-automatic procedure for magnetic field determination. We have applied this procedure to the full content of circular spectropolarimetric measurements of the FORS1 archive (except for most of the observations obtained in multi-object spectropolarimetric mode). We have devised and applied a number of consistency checks to our field determinations, and we have compared our results to those previously published in the literature. } { We find that for high signal-to-noise ratio measurements, photon noise does not account for the full error bars. We discuss how field measurements depend on the specific algorithm adopted for data reduction, and we show that very small instrument flexures, negligible in most of the instrument applications, may be responsible for some spurious field detections in the null profiles. Finally, we find that we are unable to reproduce some results previously published in the literature. Consequently, we do not confirm some important discoveries of magnetic fields obtained with FORS1 and reported in previous publications. } { Our revised field measurements show that there is no contradiction between the results obtained with the low-resolution spectropolarimeter FORS1 and those obtained with high-resolution spectropolarimeters. FORS1 is an instrument capable of performing reliable magnetic field measurements, provided that the various source of uncertainties are properly taken into account. } ",
        "introduction": "During a full decade of operations, the FORS1 instrument of the ESO Very Large Telescope has collected a large number of magnetic field measurements in various kinds of stars. Together with the ESPaDOnS instrument of the Canada-France-Hawaii Telescope, and with the MuSiCoS and NARVAL instruments of the 2\\,m Telescope Bernard Lyot of the Pic-du-Midi Observatory, FORS1 has been one of the workhorse instruments for the observational studies of stellar magnetism.  Several important detections obtained with FORS1 have led to the conclusion that magnetic fields are quite common in a variety of stars across the Hertzsprung-Russell diagram, including for instance central stars of planetary nebulae \\citep{Joretal05}, hot subdwarfs \\citep{Otoetal05}, $\\beta$\\,Cephei and slowly pulsating B stars \\citep{Hubetal09a}, B stars with emission lines \\citep{Hubetal09b}, and normal O-type stars \\citep{Hubetal08b}.  In practice, a close inspection to the published results shows a number of problems:  \\noindent \\textit{i)} Inconsistencies between field measurements obtained with FORS1 and field measurements obtained with other instruments. For instance, \\citet{Hubetal04a} reported the discovery of a magnetic field in the Herbig Ae/Be star HD\\,139614, while repeated ESPaDOnS measurements failed to confirm the magnetic nature of that star \\citep{Wadetal05}. Similarly, \\citet{Siletal09} failed to confirm the presence of a magnetic field in several of the $\\beta$\\,Cep and SPB stars observed by \\citet{Hubetal06a} and \\citet{Hubetal09a}.  \\noindent \\textit{ii)} Inconsistencies between the analysis of the same FORS dataset performed by different authors. For instance, \\citet{McSwain08} observed the normal B stars NGC 3766 MG 111 and NGC 3766 MG 176, and the Be star NGC 3766 MG 200, and reported no field detection. Using the same data, \\citet{Hubetal09b} reported new field detections for all three stars.   \\noindent \\textit{iii)} Inconsistencies between the analysis of the same FORS dataset performed by the same authors but at different epochs. For instance, \\citet{Wadetal05} reported a possible detection in the young Ap star HD\\,72106A, which was not confirmed by the later analysis of the same data performed by the same group \\citep{Wadetal07a}. Note that the magnetic nature of that star was established with data independently obtained with ESPaDOnS \\citep[see][]{Foletal08}.  \\noindent \\textit{iv)} Inconsistencies between field measurements obtained from different subsets of an observing series of frames. Magnetic field measurements are often obtained by combining a number of pairs of frames obtained at two different position angles of the retarder waveplate. This redundancy is mainly motivated by the need to reach a very high signal-to-noise ratio. In some rare cases, a magnetic field may be firmly detected in a pair of frames, but not in the remaining pairs.  This is for instance the case of the measurements of HD\\,139614 reported by \\citet{Hubetal04a}, where the magnetic field is detected only in a subset of frames, and in a couple of H Balmer lines.  \\noindent \\textit{v)} Finally, there are some global inconsistencies of the full FORS dataset, revealed for instance by the high incidence of field detection in the null profiles; this kind of problem was not previously reported in the literature, and will be discussed in Sect.~\\ref{Sect_Internal_Tests}.  Issues \\textit{ii)} and \\textit{iii)} point to possible glitches in the data reduction method, while issues \\textit{i)}, \\textit{iv)}, and \\textit{v)} might be the symptom of a wider, possibly instrumental, problem.  The release of the FORS pipeline for spectropolarimetric data \\citep{Izzetal10,Izzetal11} gave us the opportunity to develop an accurate and efficient reduction method. The FORS pipeline is a software tool specifically designed for that instrument, which makes it easier to handle some characteristics that are specific to the data obtained with FORS1. Furthermore, the FORS pipeline allows a high degree of automatization in the data reduction process. Complemented with a suite of software tools for data pre-processing (for frame classification and quality check) and data post-processing (for the computation of the magnetic field), we were able to build up a nearly automatic tool for data analysis which allowed us to treat the entire FORS archive data in a homogeneous way. Using our tool suite, the difference in terms of effort required to perform the analysis of a single series of observations compared to performing the analysis of the entire archive of FORS1 data consists mainly in the amount of the necessary disk space and CPU time.  Compared to the reduction of raw data coming from individual observing runs, the mass-production of reduced spectra offers the possibility to perform a quality check of the final products on a very large scale. The aim of this paper is to present: (1) a preliminary discussion of the methods, (2) the results of our quality checks, and (3) to caution the reader about the lack of robustness of some of the results previously published in the literature. A deeper analysis and a comprehensive and homogeneous catalogue of FORS1 magnetic field measurements will be published in a forthcoming paper. ",
        "conclusions": ""
    },
    "1112/1112.4408_arXiv.txt": {
        "abstract": "We investigate the observability of cold accretion streams at redshift 3 via Lyman-alpha (\\lya{}) emission and the feasibility of cold accretion as the main driver of \\lya{} blobs (LABs). We run cosmological zoom simulations focusing on 3 halos spanning almost two orders of magnitude in mass, roughly from $10^{11}$ to $10^{13}$ solar masses. We use a version of the \\ramses{} code that includes radiative transfer of ultraviolet (UV) photons, and we employ a refinement strategy that allows us to resolve accretion streams in their natural environment to an unprecedented level. For the first time in a simulation, we self-consistently model self-shielding in the cold streams from the cosmological UV background, which enables us to predict their temperatures, ionization states and \\lya{} luminosities with improved accuracy. We find the efficiency of gravitational heating in cold streams in a $\\sim10^{11}$ solar mass halo to be around 10-20\\% throughout most of the halo but reaching much higher values close to the center. As a result most of the \\lya{} luminosity comes from gas which is concentrated at the central 20\\% of the halo radius, leading to \\lya{} emission which is not extended. In more massive halos, of $\\ga 10^{12}$ solar masses, cold accretion is complex and disrupted, and gravitational heating does not happen as a steady process. Ignoring the factors of \\lya{} scattering, local UV enhancement, and SNe feedback, the cold `messy' accretion alone in these massive halos can produce LABs that largely agree with observations in terms of morphology, extent, and luminosity. Our simulations slightly and systematically over-predict LAB abundances, perhaps hinting that the interplay of these ignored factors may have a negative net effect on extent and luminosity.  We predict that a factor of a few increase in sensitivity from current observational limits should unambiguously reveal continuum-free accretion streams around massive galaxies at $z=3$. ",
        "introduction": "\\label{Intro.sec} The last decade has seen a shift in the way galaxies are thought to have assembled. In the classic theory \\citep{Rees:1977p4388, Silk:1977p4383,White:1978p4389}, galaxies collect their baryons via so-called hot mode accretion where diffuse gas symmetrically falls into dark matter (DM) halos and is shock-heated as it hits the gas residing in them. Depending on the mass of the halo, the gas may or may not eventually settle into the galaxy. However, it has become increasingly apparent through theoretical work and simulations that at high redshift ($z\\ga 2$), galaxies get their baryons primarily via accretion of relatively dense, cold ($10^4$ K) and pristine gas which penetrates in the form of \\textit{streams} through the diffuse shock-heated medium \\citep{Fardal:2001p3736, Birnboim:2003p3602, Keres:2005p3601, Dekel:2006p4450, Birnboim:2007p4448, Ocvirk:2008p2688, Dekel:2009p1318, Brooks:2009p3604, vandeVoort:2011p5673, FaucherGiguere:2011p5611, vandeVoort:2011p5669}. Simulations consistently show these streams to exist and peak in activity around redshift 3, though it appears that their widths are still dictated mostly by resolution.  The problem is that cold accretion streams have never been directly observed, though we are starting to see some hints, both in emission \\citep{Rauch:2011p5439} and absorption \\citep{Ribaudo:2011p5454}.  Is this lack of observational evidence consistent with the existence of cold accretion streams? Do we not observe them because they're not easily observable or simply because they don't exist?  \\cite{FaucherGiguere:2011p3606} showed that the streams are hard to detect directly via absorption due to their small covering factor and surrounding galactic winds that overwhelm their signature. \\cite{Kimm:2011p4491} came to the same conclusion, adding that the low metallicity in streams ($\\leq 10^{-3}$ solar) further inhibits their detection via metal line absorption. Even so, \\cite{Fumagalli:2011p2943} and \\cite{vandeVoort:2011p5667} have argued that a large fraction of observed metal-poor Lyman-limit systems (LLSs) make up for indirect detections of cold streams.  Furthermore, we may possibly have been directly observing the tips of these streams during the last decade in the form of Lyman-alpha blobs (LABs).   \\vsk LABs are extremely bright ($\\ga 10^{43} \\; \\mathrm{erg \\; s^{-1}}$) and extended ($\\ga 30 \\;\\mathrm{kpc}$ in diameter) \\lya{} nebulae \\citep[e.g.][]{Francis:1996p4544, Keel:1999p4529, Steidel:2000p2213, Matsuda:2004p3081, Palunas:2004p4550, Nilsson:2006p3525, Smith:2007p4610, Prescott:2009p3951, Yang:2010p3447, Erb:2011p5386}.  They have a slight tendency to be filamentary in structure \\citep[][hereafter M11]{Matsuda:2011p5426}, and often have short limbs protruding from the main body. They often coincide with galactic sources that give hints about their physical origin but the mechanism by which the emission becomes so strong and extended is a matter of debate. A subset of LABs however have no apparent coinciding galactic sources \\citep[e.g.][]{Steidel:2000p2213, Weijmans:2010p5527, Erb:2011p5386}. Up until now about two hundred LABs have been discovered, including about fifteen giant ones ($>100$ kpc). Smaller extended \\lya{} emitters exist in large quantities over a continuous range of sizes down to point sources. LABs appear to be specific to the high-redshift Universe \\citep{Keel:2009p4857} and most of them have been detected at $2<z\\la3$.  The physical nature of LABs is still a matter of debate, but by most accounts they are powered by a combination of some or all of the following processes: \\textit{(a)} Cold stream accretion is a natural explanation, where the fuel source is the dissipation of gravitational potential, also termed gravitational heating \\citep[e.g.][]{Steidel:2000p2213, Haiman:2000p4632, Fardal:2001p3736, Dijkstra:2006p4697, Dijkstra:2009p3780}.  \\textit{(b)} Photo-fluorescence by near-lying sources, such as active galactic nuclei (AGN) or starbursts \\citep[e.g.][]{Haiman:2001p4742, Cantalupo:2005p4317, Kollmeier:2010p3256}, \\textit{(c)} \\lya{} scattering, also fuelled by neighbouring star-forming regions \\citep[e.g.][]{Laursen:2007p4741, Zheng:2011p5486}. \\textit{(d)} Cooling radiation in galactic outflows, fuelled by AGN or supernovae \\citep[e.g.][]{Taniguchi:2000p4771, Ohyama:2003p4783, Mori:2004p4829}.  \\cite{Furlanetto:2005p3744} used cosmological simulations to look at the contributions of each of these processes, and found that star-forming regions can in principle power all but the largest LABs via photo-fluorescence and \\lya{} scattering, but that cold accretion alone cannot, except under very optimistic assumptions.  They however pointed out that the \\lya{} emissivity of their simulated gas is highly uncertain due to the lack of modelling of self-shielding from UV radiation: The self-shielding state of the gas affects both the temperature and ionization state, which sensitively dictates the \\lya{} emissivity.  They also pointed out that the efficiency of star-formation in powering LABs is very dependent on the presence of dust. As pointed out by \\cite{Cen:2011p5588}, massive galaxies tend to have large dust content which makes them very efficiently transform their UV (and \\lya{}) output into infrared radiation. Thus it appears problematic to associate the largest and most luminous LABs to star-formation in the most massive halos in the Universe.  \\subsection {Recent work on gravitationally driven \\lya{} emission} Notably, two recent simulation papers have studied gravitational heating as the driver of LABs, but have reached conflicting",
        "conclusions": "We have in this work addressed the questions of whether gravitational heating may be the main driver of LABs, and how close we are to making direct and unambiguous detections of cold accretion streams via their \\lya{} emission.  To this purpose we have run and analyzed cosmological RHD simulations specifically tailored to accurately predict \\lya{} emission from extended structures. These simulations are idealized in the sense that the effects of stellar feedback and \\lya{} scattering are ignored such as to isolate the efficiency of gravitational heating in generating \\lya{} photons. Our analysis is focused on redshift 3, which corresponds to most LAB observations.  Our approach improves upon previous works in the following ways: (a) Using \\ramsesrt{}, our newly developed RHD version of the \\ramses{} code, we include on-the-fly propagation of UV photons, which allows us to consistently and accurately model the self-shielding state in accretion streams and their resulting temperatures and ionization fractions, which are all very important to accurately predict their \\lya{} emissivity.  (b) We apply a novel refinement strategy that allows us to optimally resolve accretion streams to an unprecedented degree and on much larger scales than previously. This allows us to spatially resolve the competition between gravitational heating and radiative cooling in those streams.  (c) We post-process our simulation outputs with very small timesteps to ensure we also resolve said competition temporally. Failing to do this leads to a dramatic underestimate of \\lya{} emissivity of gas due to the commonly utilized numerical method of operator splitting, and previous works may have been marked by this problem.  (d) We simulate more massive halos than hitherto done, based on the growing consensus that LABs are hosted by the most massive halos in the Universe.  \\vsk There are nontheless issues regarding uncertainties that potentially affect our results. One is the likely presence of artificial overcooling in shocks. As pointed out by \\cite{Creasey:2011p5926}, shocks -- or the mean-free paths of particles inside them -- are almost exclusively under-resolved in cosmological simulations. The artificially broadened shocks can prevent the creation of hot and diffuse gas phase and instead allow it to efficiently cool and remain at temperatures where \\lya{} emission is the most effective cooling channel. We may thus over-predict the \\lya{} emissivities of shock regions in our simulations. However, this effect should be most severe in regions where gas is shocking on to galactic disks, and should thus be mainly constrained to CGM regions, and to the densest gas under consideration, i.e. $n_H\\sim 1 \\ \\cci{}$. Weaker shocks may also exist at the boundaries of the disrupted streams in and around our more massive halos, but it seems unlikely that numerical overcooling is a big issue here, due to the high resolution, large volumes, low densities, and the fact that the \\lya{} emissivity is not particularly concentrated at the stream boundaries.  It is an unavoidable fact that the denser the gas in our simulations, or in any simulations for that matter, the larger the uncertainty in its \\lya{} emissivity. In particular, at densities $\\ga 0.1 \\ \\cci$, gas may cool down to $<< 10^4$ K via molecular or metal-line cooling, neither of which is included in our simulations.  The densest gas is also in general the most \\lya{} luminous in our simulations: What we term CGM gas ($n_H>0.3\\ \\cci$) consistently accounts for $40\\%$ of the total \\lya{} luminosities of our halos, so we can estimate the total \\lya{} luminosities to be uncertain by (very) roughly $50\\%$, and even more if we exclude still more diffuse gas than the CGM.  We have however shown that our results and conclusions regarding LAB areas are not sensitive to the density threshold applied (i.e. above which densities we ignore \\lya{} emissivity).  \\vsk Our main results are the following: \\begin{itemize} \\item{ Cold accretion streams in halos more massive than $\\sim 10^{12}$~M$_\\odot$ produces extended and luminous \\lya{} nebulae which are by large compatible with LABs observed at $z\\sim 3$, in terms of morphology, luminosity and extent. Gravity alone provides most of the energy, and we find that extra sources such as UV fluorescence, \\lya{} scattering or superwinds are not necessary. This clearly doesn't rule out these other processes though, as they are likely all significant in the case of LABs, and further work is needed to study their complex interplay.}  \\item{In our simulations, LAB area and luminosity are reasonably well-behaved functions of halo mass. We use these relations to compute the cumulative luminosity and area distributions, and find that they are in reasonable agreement with observations given the relatively large uncertainties. This comparison however suggests that the combined effects of SN feedback, \\lya{} scattering and an enhanced local UV field may possibly have a negative impact on the luminosity and extent of simulated LABs, when conjoined with cold accretion. }  \\item{ The model of gravitational heating as a driver of extended \\lya{} emission works according to our results, but we need to alter our notion of \\textit{how} it works: It is inefficient in the classic sense where gas accretion is smooth. Rather the accretion is messy and disrupted in massive halos and probably involves some mass loss to the surrounding hot diffuse medium.  }  \\item{Our examination of maximum photofluorescence hints that in extreme cases local UV enhancement, e.g. near quasars, can boost the \\lya{} luminosity of LABs and to a lesser degree their extent. As demonstrated by \\citet[][]{Cantalupo:2005p4317} and \\citet{Kollmeier:2010p3256}, this means that large accretion flows may be more easily observed in the proximity of quasars than elsewhere.}  \\item{We find that cold accretion streams should be unambiguously observable via direct \\lya{} emission for the first time in the near future, on upcoming instruments such as MUSE and K-CWI which will allow to probe emission at surface brightnesses as low as $\\sim 10^{-19}\\IS$.}  \\end{itemize}  Although we have significantly improved on previous work, a large number of theoretical issues remain to be addressed. In forthcoming papers, we plan to investigate the effects of \\lya{} scattering SNe-driven winds and local UV enhancement from star formation."
    },
    "1112/1112.5455_arXiv.txt": {
        "abstract": "Existing photometry for NGC 2264 tied to the Johnson \\& Morgan (1953) {\\it UBV} system is reexamined and, in the case of the original observations by Walker (1956), reanalyzed in order to generate a homogeneous data set for cluster stars. Color terms and a Balmer discontinuity effect in Walker's observations were detected and corrected, and the homogenized data were used in a new assessment of the cluster reddening, distance, and age. Average values of $E_{B-V}= 0.075 \\pm0.003$ s.e. and $V_0${\\it --M}$_V=9.45 \\pm 0.03$ s.e. ($d = 777 \\pm12$ pc) are obtained, in conjunction with an inferred cluster age of $\\sim 5.5 \\times 10^6$ yr from pre-main-sequence members and the location of the evolved, luminous, O7 V((f)) dwarf S Mon relative to the ZAMS. The cluster main sequence also contains gaps that may have a dynamical origin. The dust responsible for the initial reddening towards NGC 2264 is no more than 465 pc distant, and there are numerous, reddened and unreddened, late-type stars along the line of sight that are difficult to separate from cluster members by standard techniques, except for a small subset of stars on the far side of the cluster embedded in its gas and dust and background B-type ZAMS members of Mon OB2. A compilation of likely NGC 2264 members is presented. Only 3 of the 4 stars recently examined by asteroseismology appear to be likely cluster members. NGC 2264 is also noted to be a double cluster, which has not been mentioned previously in the literature. ",
        "introduction": "} ``How often have I said to you that when you have eliminated the impossible, whatever remains, however improbable, must be the truth?'' Sherlock Holmes to Dr. Watson (Conan~Doyle 1890).  In the study of open clusters a similar totalogy could be paraphrased as ``when you have eliminated the effects of extinction and differential reddening for cluster stars, the resulting cluster color-magnitude diagram, however unusual, must represent an accurate picture of the temperatures and luminosities of cluster stars.'' That philosophy was demonstrated to be the case for relatively young (10$^7$--10$^8$ yr old) open clusters by Turner (1996), as well as for the extremely young ($3.5 \\times 10^6$ yr old) cluster IC 1590 (Guetter \\& Turner 1997). The last study demonstrated how closely pre-main-sequence members of IC 1590 matched model isochrone predictions by Palla \\& Stahler (1993) converted to the Johnson \\& Morgan (1953) {\\it UBV} system.  The situation for other young clusters is more complicated, primarily because most of the original {\\it UBV} studies of the brighter members of the class (NGC 2264, NGC 6530, IC 1546, and NGC 6611) were made by Merle Walker (Walker 1956, 1957, 1959, 1961, respectively) using a detector system that was not an ideal match to the Johnson system, although that was not realized at the time. For example, later studies of NGC 6611 (Hiltner \\& Morgan 1969) and NGC 2264 (e.g., Mendoza \\& G\\'{o}mez 1980) noted differences between Walker's photometry and observations tied more closely to the {\\it UBV} system. The origin of such differences can be explained by the work of Moffat \\& Vogt (1977) and Guti\\'{e}rrez-Moreno, Moreno \\& Cort\\'{e}z (1981), who noted that systematic errors, specifically in {\\it U--B} measures, can arise from the manner in which the Balmer discontinuity in the continua of early-type stars is sampled by non-standard telescope/filter/detector systems, as well as by the treatment of atmospheric extinction (see Cousins \\& Caldwell 2001). The differences in the case of NGC 6611 are fairly extreme, amounting to offsets of $0.10$ in {\\it U--B} and $0.03$ in {\\it B--V} (Hiltner \\& Morgan 1969).  The purpose of the present study is to redo Walker's original study of NGC 2264 (Walker 1956) in order to generate a new reddening-free and extinction-corrected color-magnitude diagram for the cluster. The cluster has been studied many times previously, but never for the purpose of improving upon Walker's results. Recent detections of non-radial pulsation via asteroseismology in many of the pre-main-sequence members of NGC 2264 (e.g., Zwintz 2008; Kallinger, Zwintz \\& Weiss 2008; Guenther et al. 2009) make it imperative to have a clear picture of the evolutionary status and exact H-R diagram location of cluster stars. ",
        "conclusions": "} One advantage of removing the effects of interstellar reddening and extinction from cluster color-magnitude diagrams is that the resulting distribution of data points must be closely linked to the true variations in effective temperature and luminosity for cluster members. In the case of NGC 2264, the color-magnitude diagram presented here in Fig.~\\ref{fig5} provides a number of important insights into the evolutionary status of cluster stars. The main sequence gaps, for example, imply an advanced dynamical state, provided they arise from close binary mergers, as argued by Turner (1996). And the close coincidence in the cluster age inferred from both pre-main-sequence stars and the slightly-evolved S Mon imply that the formation of cluster stars was not spread out greatly in time, the creation of cluster stars probably consuming no more than a few hundred thousand years.    \\subsection*{{\\rm \\scriptsize ACKNOWLEDGEMENTS}} \\scriptsize{The present study has used the open cluster data compilation maintained by WEBDA, operated at the Institute for Astronomy of the University of Vienna.}"
    },
    "1112/1112.0046_arXiv.txt": {
        "abstract": " ",
        "introduction": "Small bodies in the Solar System are usually divided into stable populations (like the main asteroid belt, Jupiter Trojans or the Kuiper belt) and the unstable, planet-crossing populations (like the near-Earth asteroids, comets or the Centaurs). Stable populations are seen as fossilized remnants from the formation of the Solar System, while the planet-crossing small bodies are short-lived and are constantly replenished from these stable reservoirs. There are some exceptions to this classification, like the Trans-Neptunian Scattered Disk, which is very long-lived but evolves appreciably over the age of the Solar System \\citep{vol08}. The small Hungaria asteroid group, on low-eccentricity, high-inclination orbits just beyond Mars, has also been suggested as a constantly eroding population \\citep{mce10}, but the full implications of the Hungarias' long-term decline have not been explored yet.  The issue of small body stability and instability is important for determining the impact history of planets and satellites. The Moon is the best studied body in the Solar System after Earth, and reveals a record of the system's early history that has since been erased on our planet. The lunar samples returned by the Apollo missions firmly established that the Moon was subject to intense bombardment at 3.85 Gyr ago (Gya) \\citep{ter74}. The nature of lunar bombardment before that time is controversial, as is the source of the impactors at 3.85 Gya. The terms \"Lunar cataclysm\" (LC) and \"Late Heavy Bombardment\" (LHB) have been introduced to describe this bombardment, and the two terms are now used almost interchangeably. A non-primordial source of the lunar cataclysm at 3.85 Gya is based on two lines of argument. One argument is that the decay of lunar bombardment at the time of the large Imbrium impact at about 3.85 Gya is too rapid to be consistent with a slowly decaying primordial population \\citet{ccg07}. Independently, dynamical calculations show that the primordial small bodies should have been all but extinct by 3.85 Gya \\citep{bot07}. This leads to a conclusion that the lunar cataclysm impactor population must have become unstable after a period of stability.  The only plausible mechanism for a delayed instability of a numerous small body population would be a late rearrangements of the planets. This could be a local event, involving an additional terrestrial planet \\citep{cha07}, or global, affecting the giant planets and, indirectly, the rest of the system \\citep{lev01, tsi05, gom05}. While there is independent evidence for an unstable episode in the Solar System's history \\citep{mor05, nes07, mor09, mor10}, there is no independent evidence about the timing of this instability; it is most often assumed to coincide with the lunar cataclysm \\citep{gom05}.  In \\citet{cuk10} we argue that the record of the cataclysm does not match the predictions of the planetary migration-based models. Youngest basins (Imbrium and Orientale) seem to have formed much later than the older basins (see section 3), and were produced by a non-asteroidal impactor population \\citep{cuk10}, while planetary migration predicts the late impactors to be derived from main-belt asteroids \\citep{gom05}. More generally, the necessity of a connection between the planetary instability and bombardment at 3.85 Gya hinges on the lack of other plausible explanations for the late delivery of lunar impactors. In this paper, we will show that there should still be a significant number of primordial Mars-crossers present at 3.85 Gya, so that this population can indirectly trigger a relatively short-lived lunar cataclysm, in full agreement with the available evidence. ",
        "conclusions": "Here we will detail a new scenario of the chronology and sources of early lunar bombardment. This scenario is tentative and will need further improvements if it is to explain the full complexity of the lunar record. However, we believe that by proposing a number of testable predictions we can help advance lunar science.  On the basis of magnetic and cratering data, we conclude that there were two populations of ancient lunar impactors, a pre-Imbrian one and an Imbrian one. These groups are largely equivalent to Populations I and II from \\citet{str05}, although there is some disagreement about the nature of Population II. Here we find that the most likely source of Population I were the primordial Mars crossers, which would have decayed with the half-life of 80 Myr and may have been able in producing $D > 900$~km basins until 4.0~Gya. The size-distribution of the Population I impactors was very similar to present-day main belt asteroids (MBAs), but there is no reason to think that other (now extinct) inner solar system small body populations did not have the same SFD. It is reasonable to expect that MBAs, Mars-crossers, the \"E-belt\" \\citep{cuk08, bot10} and the leftovers from Earth's formation \\citep{bot07} would all have had similar size-distributions.    The above discussion implies that the size-distribution and composition of Population I impactors are unlikely to be specific to their source region. However, chronology of their impacts should be indicative of their dynamical evolution. If the Nectaris basin is older than about 4.2 Gyr, it could plausibly be the product of a primordial impactor population, as proposed here. If, on the other hand, it is younger than 4.1 Gyr ago, it is likely to be a part of a cataclysm-type event \\citep{bot10}. The absolute age of Nectaris basin may be impossible to determine from the existing Apollo 16 samples, which appear dominated by Imbrium ejecta \\cite{kor87, has98, war03, nor10}. The 3.9~Gya Apollo 17 ages attributed to Serenitatis may actually reflect the formation of Imbrium, with Serenitatis being much older \\citep{has98, spu11}.  Unfortunately, the factors that complicate dating of the nearside Nectarian basins may similarly affect future samples from other locations on the Moon. Moscoviense may be the only lower Nectarian basin sufficiently distant from both Imbrium and Orientale impacts that it was likely not contaminated much by their ejecta, and may be (in theory) the best candidate for sample return among Nectarian basins. Dating of the largest and oldest South Pole-Aitken (SPA) basin would put a upper limit on the age of lunar surface features and confirm or rule out the total resurfacing scenario \\citep{ter74}, but would not give us any information about the chronology of much later impactors. An age of about 4.4~Gyr for SPA would allow both \"medium\" \\citep{bot10} and minimal (as proposed here) versions of the lunar cataclysm, so at least one Nectarian basin would need to be dated in order to distinguish between these two tentative chronologies.  In contrast, the timing of (lower Imbrian) Population II impactors is well established, as the Imbrium basin formed at 3.85-3.87 Gya, and the most Imbrian craters formed shortly afterwards. Since the size distribution of Imbrian impactors appears to be different from that of Nectarian ones \\citep{wil78, cuk10}, and different from asteroids, their source is unknown. We propose a collisional disruption of a Vesta-sized Mars-crosser as their source \\citep{wet75} on the basis of theoretical considerations, tentative composition of the Imbrium impactor \\citep{puc08}, and possible connection to mesosiderites.   A minimal cataclysm from a Mars-crossing collisional disruption also affected Mars, but the number of superposed craters indicate that the largest impact features on Mars like Argyre and Hellas are likely to predate the lunar cataclysm \\citep{tan86, wer08}. In any case, amount of mass in 200 km bodies necessary to form these basins is not consistent with the breakup of one Vesta-sized body. Our hypothesis predicts that none of the largest basins formed at Mars at 3.85 Gya, and, in general, implies much older ages for Noachian features than the models that assume a larger-scale bombardment at 3.85 Gya \\citep{wer08}. Despite the limited extent of the spike at 3.9 Gya, Mars would have experienced substantial bombardment prior to 3.6 Gya due to Mars-crossers and this bombardment was the likely the source of shocking found in the ancient martian meteorite ALH84001 \\citep{ash96}.  Shock ages of HED meteorites and H-chondrites are also sometimes used to constrain the early Solar System bombardment \\citep{bog03, swi09}. Interestingly, both HED meteorites and H-chondrites record a much longer bombardment episode than the Moon, making it hard to explain both records with the same impactor population \\citep{har03}. We think that it is significant that the HED and H-chondrite parent bodies \\citep[Vesta, and, most likely, Hebe ]{mcs10, gaf98, bot10b}  are both located in the inner main belt and would be likely to suffer significant number of collisions with late surviving Mars-crossers/proto-Hungarias. The shock resetting of meteorite Ar-Ar ages appears to be dependent both on the size of the body, parameters of impacts and the family history \\citep{nes09}. Therefore, while the number of impacts by Mars-crossers onto Vesta and Hebe from 3.5 to 4.1 Gya did not dwarf their collisions with other asteroids since then (H-chondrites also show more recent Ar-Ar ages, while HEDs do not), it is likely that the shock ages of HED meteorites and H-chondrites do have a connection to Mars-crossers. Even more than meteorite shock ages, impact events recorded in lunar meteorites \\citep{coh00, coh05} seem at odds the lunar impact history inferred from Apollo samples \\citep{har03, har07}. Events postdating 3 Gya are recorded in lunar meteorites, long after the heavy bombardment ended. Our hypothesis does not offer an explanation for this discrepancy, and we suspect that the solution lies in the physics of lunar regolith rather than the dynamics of impactors.  How can our hypotheses be tested? First, new crater counts (down to about $D=8$~km) on all non-magnetized basins (Imbrium, Orientale, Schr\\\" odinger and Hertzsprung) would confirm or falsify the existence of a non-asteroidal Imbrian size distribution, without having to resort to stratigraphic or morphological identification of craters. Otherwise, the study of mesosiderites and Hungarias could probably test their relationship with the lunar cataclysm faster than new lunar samples can be collected. Our hypothesis predicts that the mesosiderite parent body (or bodies) are likely located among Hungarias (although innermost main belt and Mars Trojans are also plausible refuges).  Remote sensing of asteroid composition has advanced considerably in the last few decades, and it is possible that we may be able to identify the intermediate mesosiderite parent body, if it is big and bright enough (small Hungarias are brighter than small MBAs). Radar observations may also in principle be able to break the degeneracy among X-type asteroids among Hungarias \\citep{ock10}. Additionally, if the Dawn spacecraft identifies a mesosiderite unit exposed on Vesta (or if mesosiderites are traced to any other large asteroid deep in the main belt), our hypothesis about the connection between the mesosiderites and the cataclysm would be falsified. While the collisional disruption hypothesis can be formulated independently from mesosiderites \\citep{wet75}, the case becomes weaker once mesosiderites are excluded from consideration."
    },
    "1112/1112.4386_arXiv.txt": {
        "abstract": "We comment on arXiv:1112.1320 and point out that baryonic oscillations of the matter power spectrum, while predicted by theories that do not incorporate collisionless cold dark matter, are strongly suppressed by the statistical window function that is used to process finite-sized galaxy samples. We assert that with present-day data sets, the slope of the matter power spectrum is a much stronger indicator of a theory's validity. We also argue that MOND should not be used as a strawman theory as it is not in general representative of modified gravity theories; some theories, notably our scalar-vector-tensor MOdified Gravity (MOG), offer much more successful predictions of cosmological observations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.1106_arXiv.txt": {
        "abstract": "Despite the success of Maxwell's electromagnetism in the description of the electromagnetic interactions on small scales, we know very little about the behaviour of electromagnetic fields on cosmological distances. Thus, it has been suggested recently that the problems of dark energy and the origin of cosmic magnetic fields could be pointing to a modification of Maxwell's theory on large scales. Here, we review such a proposal  in which the scalar state which is usually eliminated be means of the Lorenz condition is allowed to propagate. On super-Hubble scales,  the new mode is essentially given by the temporal component of the electromagnetic potential  and contributes as an effective cosmological constant to the energy-momentum tensor. The new state can be generated from quantum fluctuations during inflation and it is shown that the predicted value for the cosmological constant agrees with observations provided inflation took place at the electroweak scale. We also consider more general theories including non-minimal couplings to the space-time curvature in the presence of the temporal electromagnetic background. We show that both in the minimal and non-minimal cases, the modified Maxwell's equations include  new effective current terms which can generate magnetic fields from sub-galactic scales up to the present Hubble horizon. The corresponding amplitudes could be enough to seed a galactic dynamo or even to account for observations just by collapse and differential rotation in the protogalactic cloud. ",
        "introduction": "Out of the four fundamental interactions, gravity and electromagnetism are the only long-range forces in nature. The standard theories for these two interactions, namely General Relativity (GR) and Maxwell's electromagnetism have received a vast experimental support in a huge range of scales. GR is able to explain the gravitational phenomena from sub-millimeters distances up to Solar System scales and standard electromagnetism has been probed from the tiny scales involved in high-energy colliders up to distances of order  1.3 A.U. corresponding to the coherence lengths of the magnetic fields dragged by the solar wind{\\cite{limit}. However, despite the enormous success of these two theories, both of them present some unresolved problems when large scales are involved. For GR, cosmological observations require the presence of exotic components in the universe, i.e. dark matter and dark energy. Indeed, some attempts to account for such observations without invoking dark components include infrared modifications of GR. On the other hand,  the presence of cosmic magnetic fields found  in galaxies, clusters\\cite{galactic1,galactic2,galactic3} and, very recently\\cite{extragalactic1,extragalactic2,extragalactic3,extragalactic4}, also in the voids cannot be accommodated within Maxwell's theory. This could also be signaling that  a more careful analysis of the behavior of electromagnetic fields in cosmological contexts is needed.  A particularly interesting aspect is the quantization of gauge theories in non-trivial spacetimes. The usual quantization procedures of the electromagnetic field in flat spacetime rely on some type of subsidiary condition on the physical states to get rid of the unphysical gauge modes. Although all the approaches are equivalent in flat spacetime, such an equivalence has not been proven in curved background. In fact, the Gupta-Bleuler formalism for the covariant quantization has been shown to present difficulties in a time-dependent spacetime \\cite{Parker,EM2}. The BRST method also exhibits similar pathologies in certain spacetimes\\cite{ghosts1,ghosts2}.  Here, we shall review a recent proposal of an extended theory of electromagnetism in which we can avoid the aforementioned difficulties. This extension is based on allowing the propagation of the state that is usually eliminated from the physical Hilbert space by means of the subsidiary condition. We shall show how this theory can be consistently quantized in an expanding universe with {\\it three} physical states comprising the two polarizations of the usual photons plus one additional scalar state. The remarkable feature of the resulting theory is that the additional state can be generated from quantum fluctuations during inflation and give rise to and effective cosmological constant on large scales, whereas on sub-Hubble scales it leads to the generation of cosmic magnetic fields. ",
        "conclusions": "We have reviewed the extended electromagnetic theory in which a gauge-fixing term is promoted into a physical contribution  in the fundamental action. The quantization of the free theory can be performed without having to impose any subsidiary condition at the price of introducing an additional degree of freedom. This new mode can be gravitationally produced from quantum fluctuations during inflation and its amplitude will be determined precisely by the scale at which inflation takes place.  In the subsequent cosmological evolution, this new mode has two effects. On super-Hubble scales, it behaves as an effective cosmological constant whose observed value can be explained if inflation occurred at the electroweak scale. On sub-Hubble scales, the additional mode gives rise to a stochastic background  of longitudinal electric waves that generate magnetic fields from sub-galactic scales up to the present Hubble radius. On the other hand, we have also considered the effects of non-minimal couplings in the presence of the temporal electromagnetic mode and shown that in such a case, neutral massive objects can act as sources of electromagnetic fields thus implementing the old conjecture by Schuster, Einstein and Blackett of gravitational magnetism.  The theory studied in this work shows how a modification of  electromagnetism which does not require the introduction of new fields, dimensional parameters or potential terms  could provide a simple explanation for the tiny value of the cosmological constant and, at the same time, a mechanism for the generation of magnetic fields on cosmological scales. Some open questions still remain to be studied such as the inclusion of electromagnetic interactions in the theory, since all the analysis performed so far are limited to the free theory. Also, it would be interesting to know the behaviour of the new mode in more general background space-times in order to determine the viability of the model from the theoretical and phenomenological point of views. Finally,  the possibility of detecting  the longitudinal electric wave background generated during inflation could provide a clear signal of the modification of electromagnetism on cosmological scales.   \\vspace{0.2cm}  {\\em Acknowledgments:} This work has been  supported by MICINN (Spain) project numbers FIS 2008-01323 and FPA 2008-00592, CAM/UCM 910309 and MICINN Consolider-Ingenio MULTIDARK CSD2009-00064. JBJ is also supported by the Ministerio de Educaci\\'on under the postdoctoral contract EX2009-0305. \\vspace{0.2cm}"
    },
    "1112/1112.2513_arXiv.txt": {
        "abstract": "{The quiet Sun magnetic fields produce ubiquitous bright points (BPs) that cover a significant fraction of the solar surface. Their contribution to the total solar irradiance (TSI) is so-far unknown.} {To measure the center-to-limb variation (CLV) of the fraction of solar surface covered by quiet Sun magnetic bright points. The fraction is referred to as {\\em fraction of covered surface}, or FCS. } {Counting of the area covered by BPs in $G$-band images obtained at various heliocentric angles with the 1-m Swedish Solar Telescope on La Palma. Through restoration, the images are close to the diffraction limit of the instrument ($\\sim 0\\farcs 1$). } { The FCS  is largest at disk center ($\\simeq$ 1\\,\\%), and then drops down to become $\\simeq$ 0.2\\,\\% at $\\mu\\simeq 0.3$ (with $\\mu$ the cosine of the heliocentric angle). The relationship has large scatter, which we evaluate comparing different subfields within our FOVs. We work out a toy-model to describe the observed CLV, which considers the BPs to be depressions in the mean solar photosphere characterized by a depth, a width, and a spread of inclinations. Although the model is poorly constrained by observations, it shows the BPs to be shallow structures (depth$\\,<\\,$width) with a large range of inclinations. We also estimate how different parts of the solar disk may contribute to TSI variations, finding that 90\\,\\% is contributed by BPs having $\\mu > 0.5$, and half of it is due to BPs with $\\mu > 0.8$. } {} ",
        "introduction": "Our understanding of the quiet Sun magnetic fields has drastically improved during the last decade \\citep[for recent reviews, see, e.g.,][]{2009SSRv..144..275D,2011ASPC..437..451S}. We have gone from magnetic signals present only at the network boundaries \\citep[e.g.,][]{bec77b,sol93}, to ubiquitous polarization signals created through Hanle effect \\citep[e.g.,][]{fau93,tru04} and Zeeman effect \\citep[e.g.,][]{lin99,2000ApJ...532.1215S,2003ApJ...582L..55D,har07,lit08}. The wealth of quiet Sun magnetic structures makes them  potentially important to understand the global magnetic properties of the Sun \\citep[][]{2003ApJ...585..536S,tru04}, and also makes it unlikely that the quiet Sun magnetism results from the decay of active regions \\citep[e.g.,][]{san03b}. Theoretical arguments, corroborated by numerical experiments, favor a different production mechanism \\citep{pet93,cat99b,vog07,pie10,2011ApJ...736...36M}. An efficient turbulent dynamo transforms into magnetic fields part of the kinetic energy of the granular convection. It generates a complex magnetic field which evolves in short time scales (a few min) and has small characteristic length-scales ($<\\,1\\,$Mm).  In agreement with the turbulent dynamo scenario, quiet Sun magnetic fields come with strengths in the full range covering from almost zero to 2\\,kG \\citep[][]{2000ApJ...532.1215S,dom06,2008A&A...477..953M, 2009A&A...506.1415B,2011A&A...526A..60V}. Even if they only fill a small fraction of the quiet photosphere, the part having strong kG fields may be particularly important for a number of reasons. Firstly, the magnetic flux and energy increase with field strength, therefore, the energy and flux provided by kGs may surpass the contribution of the more  common but weaker fields \\citep[][]{san04}. The need to consider kGs is illustrated by the numerical experiments set up by \\citet{2011arXiv1108.1155C}, where realistic granular convection redistribute initial hG fields so that daG, hG and kG field strengths have the same energy despite their very different area covering. Magnetic concentrations with kG fields may also be  important because  buoyancy makes them vertical \\citep[e.g.,][]{sch86} and so, they naturally provide  a mechanical connection between the photosphere and the upper atmosphere \\citep[e.g.,][]{bal98,sch03b,goo04,jen06,san08}. They can function as guides that sustain magneto-acustic wave propagation, or be physical channels connecting plasmas of different atmospheric layers. Finally, kG magnetic concentrations are expected to be particularly bright due to the so-called hot-wall effect\\footnote{The magnetic pressure suffices to maintain kG structures in mechanical balance within the photosphere, therefore, they are evacuated and transparent, allowing us to look through into the sub-photosphere, which is generally hotter and so brighter.}\\citep[]{spr77,car04,kel04}. They produce bright points (BPs) which, depending on the variation across the solar disk, and during the solar cycle, may even contribute to the Total Solar Irradiance (TSI) variations as network and plage magnetic fields do \\citep[e.g.,][]{1997ARA&A..35...33L, 2004A&ARv..12..273F, 2011SSRv..tmp..133F}.  The finding of BPs in the quiet Sun was immediately identified as the kG magnetic concentrations inferred from polarization measurements \\citep{2004ApJ...609L..91S}. Their basic properties and their ubiquitous presence have been confirmed by a number of researchers \\citep[][]{dew05,dew08,bov08,san08, 2009ApJ...700L.145V,2010ApJ...723..787V, 2010ApJ...714L..31G}. They vary on time scales of minutes similar to that of granulation, and many BPs are at the resolution limit of the current instrumentation (some 0.1~Mm). They are extremely common, filling at least 1\\% of the solar surface and outnumbering granules \\citep{2010ApJ...715L..26S}.  \\citet{2011A&A...532A.136S} have recently estimated the excess of brightness produced by the quiet magnetic fields at the  disk center, turning out to be of the order of 0.15\\,\\%.  Their contribution is larger than the 0.08\\,\\%~TSI variations associated with cycle, however, determining the impact of these kG fields on TSI demands knowing their center-to-limb variation (CLV), as well as the variation of the quiet Sun magnetic fields with the solar cycle. These two properties, central to assess the role of quiet Sun fields on TSI, are  poorly known. As far as the variation along the cycle is concerned, the claims in the literature are for little variation  if any \\citep{san03d,shc03,har10}. Unfortunately, the  uncertainty of such claims can be as large as a factor two \\citep{fau01}. There are several works on the CLV % signals associated with the quiet Sun magnetic fields \\citep[e.g.,][]{2008A&A...477..953M,lit08,% 2011ApJ...737...52L} however, to the best of our knowledge, nothing is known on the CLV % of the quiet Sun magnetic BPs. This is precisely the subject of our work, i.e., providing a first observational description of the CLV % of the quiet Sun BPs. To be more exact, we evaluate how the area covered by quiet Sun BPs varies with the position on the disk.   Two main difficulties hinder the analysis. First, one has to use images with enough spatial resolution and of uniform quality, since the number of BPs depends critically  on the spatial resolution of the observations \\citep{1996ApJ...463..797T,2010ApJ...715L..26S}. This issue is sorted out using data from the 1-m Swedish Solar Telescope \\citep[SST,][]{2003SPIE.4853..341S} obtained in a single day during moments of excellent seeing, and then restored to get images with uniform resolution close to the diffraction limit of the instrument \\citep{2005SoPh..228..191V}. Second, the BPs can be misidentified with granule borders and other structures \\citep[e.g.,][]{2007SoPh..243..121B}, a problem particularly severe in near-limb images. This second problem is addressed resorting to  the cumbersome method of eye-ball identification which, however, has been proven to be reliable for our purpose \\cite[][and references therein]{2010ApJ...715L..26S}.  The work is organized as follows; the observations are presented in \\S~\\ref{observation}. The actual CLV measurements are described in \\S~\\ref{sec_ctl}. In principle, such measurements has to be interpreted using realistic models of magneto-convection (such as those used for plages by \\citeauthor{car04}~\\citeyear{car04} or \\citeauthor{kel04}~\\citeyear{kel04}), where the 3-dimensional geometry of the photosphere is considered consistently. It would require repeating the analyses carried out on the observed images using synthetic images from the simulations. Then the CLV coming from a comprehensive battery of numerical simulations should be compared with the observations. This detailed realistic approach clearly exceeds the scope of the work. However, we attempt a toy-modeling which, despite its simplicity, considers the key ingredients that within the hot-wall paradigm determine the CLV  (\\S~\\ref{interpretation}). The model is compared with the observed CLV in \\S~\\ref{results}, which constraints some of the properties of the magnetic structures. The implications for TSI variations are analyzed in \\S~\\ref{tsi_var}. These results are discussed and put into context in  \\S~\\ref{conclusions}. ",
        "conclusions": "We have measured the center-to-limb variation (CLV) of the area covered by $G$-band bright points (BPs) in the quiet Sun (fraction of covered surface or FCS). It is a parameter difficult to determine since the detection of BPs critically depends on the angular resolution of the observation (\\S~\\ref{intro}). We employ several time series taken in two hours during moments of excellent seeing with the SST. The images were post processed using MOMFBD (see \\S~\\ref{observation}) which provides a homogeneous set of images adequate for these subtle measurements. They were restored to provide an angular resolution close to the diffraction limit of the instrument at the working wavelength (some 0\\farcs 1 at the $G$-band).  We find the FCS to be largest at disk center ($\\simeq$ 1\\,\\%), and then it drops down to become $\\simeq$ 0.2\\,\\% at $\\mu\\simeq 0.3$. The relationship has large scatter, which we managed to estimate comparing different subfields within our FOVs (see the error bars in Fig.~\\ref{curvas8ptos} and Table~\\ref{ff_results}). The value obtained at the disk center agrees with previous estimates based on data of similar quality.  We work out a toy model to describe the observed CLV. It assumes the magnetic bright points to be depressions in the mean solar photosphere, characterized by a depth, a width, and with a spread of inclinations. It is only an exploratory modeling which, however, includes the physical ingredients that seems to be responsible for a kG magnetic concentration to show up as a bright feature on the solar disk. The solutions offered by our toy-model are poorly constrained, but they seem to show the BPs to be shallow structures\\footnote{This result is not inconsistent with the magnetic concentrations producing BPs being modeled as a compactly-packed ensemble of narrow magnetic concentrations arranged in a micro-structured magnetic atmosphere MISMA fashion \\citep[][]{1996ApJ...466..537S, 2000ApJ...544.1135S,2001ApJ...555..978S}. The CLV of the FCS is sensitive to the global scale of the ensemble, whereas the MISMA accounts for the smallest scales responsible, among others, for the asymmetries of the spectral lines formed in the atmosphere. } (ratio depth to width $\\simeq 0.7\\pm 0.2$) with a large spread of inclinations  ($\\simeq \\pm 70^\\circ$).   Among others, the FCS is of interest because it determines the impact that quiet Sun magnetic fields may have on TSI variations, an influence so far unknown. Since the measured FCS is so peaked towards disk center, any role that quiet Sun magnetic fields may have to play on TSI will be due to the center of the disk. According to our estimate, 90\\,\\% of the TSI is contributed by BPs having $\\mu > 0.5$,  and half of it is due to BPs with $\\mu > 0.8$ (\\S~\\ref{tsi_var}). This estimate is based on assuming the CLV of the BP (wavelength integrated) flux to scale as the observed $G$-band intensity. It is an ad-hoc assumption  adopted for lacking of a better one which, however, has allowed us to provide a first constrain on the effect of quiet Sun BPs on TSI.   In this sense, we have to stress that the FCS has been measured in the $G$-band, because the magnetic concentrations are particularly conspicuous at this wavelengths.Then we implicitly assume throughout the work that the FCS is the same at all wavelengths, i.e., that the area covered by magnetic concentration does not change with wavelength. (The $G$-band BPs are brighter but not bigger.) This assumption remains to be proven however, for the time being, it seems to be a reasonable working hypothesis.   Our toy-model suggests the BPs to be shallow structures with varied inclinations. One could test these predictions observing polarization signals of quiet Sun BPs at disk center. On the one hand, the polarization signals provide magnetic field inclinations through Stokes inversion \\citep[e.g.,][]{2008ApJ...674..596S}. On the other hand, assuming the magnetic concentrations to be in mechanical balance, one can infer their depths \\citep[e.g.,][]{2000ApJ...532.1215S}. These inferences require low-noise spectro-polarimetry with an angular resolution similar to that of our $G-$band images. They represent a technical challenge, but such observations seem to be doable in the near future \\citep[see, e.g.,][]{2010ApJ...723L.164L}. As we mentioned in the introduction, the  existing numerical simulations of magneto-convection predict the BPs to be depressed with respect to the mean photosphere \\citep[e.g.,][]{vog05,vog07}. Whether the faint BPs observed in quiet Sun are predicted to be superficial or deep remains to be worked out. However, the existing simulations of plage magnetic concentrations suggest the continuum intensity to be formed in a rather shallow region \\citep{kel04,car04}.  -------------------------------"
    },
    "1112/1112.5105_arXiv.txt": {
        "abstract": "We discuss  a fast cross-Wigner transform based technique for detecting gravitational wave bursts, and estimating the direction of arrival, using a network of (three) non co-located interferometric detectors. The performances of the detector as a function of signal strength and source location, and the accuracy of the direction of arrival estimation are investigated by numerical simulations. The robustness of the method against instrumental glitches is illustrated. ",
        "introduction": "\\label{sect:intro}  The next generation of  interferometric detectors, of gravitational waves (henceforth GW) including AdLIGO \\cite{AdLIGO}, AdVirgo \\cite{AdVirgo} and GEO-HF \\cite{superGEO}, hopefully to be followed soon by LCGT \\cite{LCGT} and ACIGA \\cite{ACIGA}, and eventually by ET \\cite{ET}, is expected to observe tens of events per year, opening the way to gravitational wave astronomy \\cite{GWBs}. Identifying the direction of arrival (henceforth DOA) of the signals, and retrieving their shapes, will be a primary task in reconstructing the physics of the sources and their environments.\\\\ The possibility of retrieving the DOA from {\\it independent} estimates of the signal arrival time at {\\it each} detector was first suggested in \\cite{Boulanger}, and further discussed in Saulson seminal book \\cite{Saulson_book}. It was shown that three-interferometers are sufficient to retrieve the DOA up to a mirror-image  ambiguity which can be solved in principle from knowledge of the detectors' directional responses. This method, often referred to as {\\it triangulation} was further elaborated by Sylvestre \\cite{Sylvestre_03}, Cavalier et al. \\cite{Cavalier_06}, and Merkovitz et al. \\cite{Zanolin_08}. In \\cite{Cavalier_06} a Gaussian distribution was assumed for the (independent) arrival time estimation errors, and a $\\chi^2$ minimization algorithm was accordingly proposed for retrieving the DOA, in the maximum likelihood spirit. In \\cite{Zanolin_08} it was shown that this method is affected by a systematic bias in the estimated DOA, a possible technique for removing the bias was discussed, and amplitude consistency tests for removing the mirror-image ambiguity were suggested. Fairhurst developed a similar analysis of the effect of arrival time estimation errors on the DOA estimation accuracy, for the special case of chirping signals, including waveform and calibration errors \\cite{Fairh_09}, \\cite{Fairh_11}.\\\\ DOA estimation algorithms are already implemented in the {\\it coherent} LIGO-Virgo pipelines for GW burst (henceforth GWB) detection \\cite{Xpipe}, \\cite{Klimenko_11}. DOA estimation in {\\it coherent} network data analysis, was studied  first by Krolak and Jaranowski \\cite{Kro_Jar_94}, and then by Pai et al. \\cite{Pai_Dhur}, as part of  the waveform parameter estimation problem,  with specific reference to chirping waveforms from coalescing binaries, in a Gaussian noise background. The conceptual foundations of coherent data analysis for unmodeled waveforms were laid out by Flanagan and Hughes \\cite{Flan_Hug}, and further developed by Klimenko et al. \\cite{Klim1}-\\cite{Rakh}. G\\\"{u}rsel and Tinto \\cite{GurTin} first suggested the possibility of retrieving the DOA for un modeled signals using null-streams. This concept was  analyzed in depth by Schutz and Wen \\cite{Wen_Schutz_05}, and further exploited by Chatterji et al. \\cite{Shurov}. A Fisher-matrix based analysis of arrival time estimation error in coherent network detection of modeled as well as {\\it unmodeled} signals was made in Wen et al. \\cite{Wen_Fan_Chen}.\\\\ In this paper we capitalize on the time-shift and localization properties of the cross-Wigner-Ville (henceforth XWV) transform to introduce a new and conceptually simple GWB detection and DOA reconstruction algorithm, using a network of non co-located interferometric detectors.\\\\ The Wigner-Ville transform is a well known powerful tool for the analysis of non-stationary signals \\cite{XWV1}, whose potential in GW data analysis, has been highlighted by several Authors, under different perspectives \\cite{Feo_etal}-\\cite{Chass}. Here we suggest its possible use as an effective tool for detecting GWBs, and estimating their DOA, which offers nice features in terms of performance, robustness against spurious instrumental/environmental transients (glitches).\\\\ Instead of using {\\it independent} estimates of the arrival times at {\\it each} detector, our DOA estimator uses data from (all) detector {\\it pairs} to estimate the needed propagation delays. In addition, it also provides an effective detection statistic, combining the data from {\\it all} detectors in the network, at a remarkably light computational cost.\\\\ DOA reconstruction from arrival-time delay estimation in a network of sensors is a well known problem in the technical Literature on Acoustics and Radar  (see, e.g.,  \\cite{Berdugo} for a broad review). The standard method for time-delay estimation in Gaussian noise is (generalized) cross-correlation \\cite{Knapp}, which is known to  perform reasonably well for relatively large signal to noise ratios \\cite{Fert_Sjo}. Remarkably, the correlation-based estimator offers worse performances compared to the XWV in the present context, as shown in Sect. 5.2\\\\ This paper is  organized as follows. In Sect.  2 we introduce the XWV transform, and recall its time-shift properties, which are illustrated  for the simplest case of sine-Gaussian (henceforth SG) GWBs. In the same section we recall the relationship between arrival time delays and DOA. In Sect. 3, we  illustrate the proposed XWV transform based DOA reconstruction algorithm. In Sect. 4 we discuss the effect of noise in the data, and the related DOA reconstruction uncertainties. In Sect. 5 we present the results of extensive numerical simulations, aimed at characterizing the performance of our XWV based algorithm both as a detector and as a DOA estimator. The simulations are based on SG-GWBs, but the case of more realistic waveforms (including Dimmelmeier and binary merger waveforms) is also discussed. In Sect. 6  we include a short discussion  of the robustness of the proposed algorithm against instrumental/environmental transients (glitches). Conclusions follow under Sect. 7. ",
        "conclusions": "We presented a simple, computationally light and fast algorithm for detecting short unmodeled GWB in a network of three interferometric GW detectors, and estimating the related DOA, based on XWV spectra. The algorithm is reasonably performant, and nicely robust against spurious transients (glitches) of instrumental origin corrupting the (otherwise Gaussian) detectors noise floor.\\\\ It  does not provide waveform reconstruction; this latter, however can be accomplished in principle off-line, once the DOA has been estimated.\\\\ Generalization to larger networks, and other potentially interesting waveforms (e.g., chirps) is relatively straightforward. Such extensions will be explored in a forthcoming paper.\\\\ Based on the above preliminary results, we suggest that the proposed algorithm may be used as a quick-and-(not-so)-dirty  on-line data sieving tool.\\\\ A quantitative comparison  with existing GWB detection/DOA estimation algorithms in terms of  efficiency and computational burden  will be the subject of future investigation."
    },
    "1112/1112.0516_arXiv.txt": {
        "abstract": "{Massive stars have a profound effect on the surrounding interstellar medium. They ionize and heat the neutral gas, and due to their strong winds, they swept  the gas up forming large \\hi\\ shells. In this way, they generate a dense shell where the  physical conditions for the formation of new stars are given.} {The aim of this study is to analyze the origin and evolution of the large \\hi\\ shell \\gs\\, and its role in triggering star forming processes.}{To characterize the shell and its environs, we carry out a multi-wavelength study. We analyze the \\hi\\ 21 cm line, the radio continuum, and infrared emission distributions.  }{The analysis of the \\hi\\ data shows an expanding  shell structure  centred at ($l, b$) = (100\\fdg6, --2\\fdg04) in the velocity range from  --29 to --51.7  \\kms. Taking into account non circular motions, we infer for \\gs\\, a  kinematical distance of 2.8 $\\pm$ 0.6 kpc. Several massive stars belonging to Cep\\,OB1 are located in projection within the  large \\hi\\, shell boundaries. The analysis of the radio continuum and infrared data reveal that there is no continuum counterpart of the \\hi\\ shell. On the other hand, three slightly extended radio continuum sources are observed in projection onto the dense \\hi\\ shell. From their flux density determinations we infer that they are thermal in nature. An analysis of the \\hi\\ emission distribution in the environs of  these sources shows, for each of them,  a region of low emissivity  having a good morphological correlation with the ionized gas in a velocity range similar to the one where \\gs\\, is detected.} {Based on an  energetic analysis, we conclude that the origin of \\gs\\, could have been mainly due to the action of the Cep\\,OB1 massive stars located inside the \\hi\\ shell. The obtained age difference between  the \\hi\\ shell and the \\hii\\ regions, together with their relative location, led us to conclude that the ionizing stars could have been created as a consequence of the shell evolution. } ",
        "introduction": "The presence in the interstellar medium (ISM) of the Milky Way, when viewed in the $\\lambda\\sim$ 21 cm line emission of the neutral hydrogen (\\hi), of giant structures having linear dimensions of a few hundred parsecs in diameter is a widely known and well observed phenomena first noticed by \\citet{hei79}. These structures are usually detected as huge shells or arc-like features of enhanced  \\hi\\, emission surrounding regions of low \\hi\\,  emissivity, receiving the generic name of \\hi\\, supershells. These features may even be the dominant structure in the interstellar medium, taking up a large fraction of the volume of the galactic disk. The \\hi\\, structures could be also observed at infrared wavelengths. Based on the 60 and 100 $\\mu$m IRAS databases, \\citet{kon07} have performed an all-sky survey of loop- and arc-like structures.  Similar structures have also been observed in nearby spiral galaxies  \\citep{sta07,cha11}. In the Milky Way, these structures were initially catalogued by  \\citet{hei79,hei84}. Though a large number of \\hi\\, features likely to be classified as either large \\hi\\, shells or \\hi\\, supershells have been catalogued in the outer (90$^\\circ \\leq l \\leq$ 270$^\\circ$) part of the Galaxy \\citep{ehl05}, only a small number of them have been studied in some detail. The later \\citep{jun96, sti01, uya02, mcc02, caz03, arn07, cic11} have galactocentric distances ranging from 9.7 to 16.6 kpc, diameters from 120 to $\\sim$ 840 pc, expansion velocities between $\\sim$10 and $\\sim$ 20 \\kms, and kinetic energies from $\\sim$ 1 $\\times$ 10$^{50}$ up to $\\sim$ 6 $\\times$ 10$^{51}$ erg. Among them only two (GS\\,305+01-24 and the feature studied by \\citet{caz03})  have an OB-association as their likely powering source, and other three (GSH\\,91.5+2--114, GS\\,234--02 and GS\\,263--02+45) show evidence of having induced the formation of new generation of stars.     The general consensus is that those structures whose kinetic energy is of the order of, or less than, a few times 10$^{51}$ erg, very likely may have been created by the joint action of stellar winds and supernova explosions. A large number of examples are reported in the literature \\citep[e.g.][]{uya02,arn07,cic11}. On the other hand, for expanding \\hi\\, structures having kinetic energies in excess $\\sim 10^{52}$ erg, termed {\\it supershells}, the above mechanism may not be adequate to create them because one would need a stellar grouping (either an open cluster or an OB-association) with many more stars than the average found in the Milky Way. In these cases alternative mechanisms like the infalling of high velocity clouds \\citep{ten81} or gamma-ray bursts \\citep{per00} may be at work.   Along its expansion, these structures (either a shell or a supershell) may became gravitationally unstable, forming clouds that later on may lead to the formation of stars along the periphery of these \\hi\\, structures, or else the expanding structure may hit and compress  pre-existing {\\rm ISM} molecular clouds from one side. During this process a high density perturbation may move into the molecular clouds, which may eventually collapse into denser cores in which star formation may occur. A thorough review of observations and theory related to trigger star formation is given by \\citet{elm98}.               A new large scale study aim at detecting  in the outer part of the galaxy structures likely to be either large \\hi\\, shells or supershells is being carried out by one of the authors (L. A. Suad) as part of her PhD Thesis.  In this paper we analyze a new large \\hi\\, shell  observed at (\\textit{l, b}) $\\sim$ ($100^\\circ$, $-2^\\circ$), with the purpose of elucidating both its origin and  its interaction with the surrounding ISM. We also look for signs of recent star formation activity likely to be related to this shell. ",
        "conclusions": "The large \\hi\\, shell \\gs\\, has been analyzed to study the interaction of massive stars with the interstellar medium and, in particular, the process of triggered star formation. From this analysis we conclude the following:  \\begin{enumerate}  \\item \\gs\\, is a large shell of a radius of about 102 pc located at a distance of 2.8 $\\pm\\,$ 0.6 kpc. The swept up mass in the shell is ($1.5 \\pm\\, 0.7$) $ \\times\\, 10^5$ M$_\\odot$ and the shell density $n_{sh} = 2.5 \\pm\\, 0.4$ cm$^{-3}$. The shell is expanding at a velocity of 11 $\\pm\\, 2$ \\kms\\, and its kinetic energy is ($1.8 \\pm\\, 0.8$) $\\times\\, 10^{50}$ erg.  \\item Several evolved massive stars members of Cep\\,OB1 are projected inside the large shell. The distance to the OB association is compatible with the kinematical distance of \\gs\\, when non-circular motions are considered. An energetic analysis suggests that the wind energy provided during the main sequence phase of the stars could explain the origin of the shell. However, taking into account the SN rate in OB associations, the energy contribution of a SN explosion as well as of its massive progenitor can not be discarded.  \\item From the 2695 MHz radio continuum and 60 $\\mu$m infrared images, we found three slightly extended sources, labelled  G98, G100, and G103 projected onto the borders of \\gs. From the radio flux densities estimated at different wavelengths, the thermal nature of the sources was confirmed by the estimation of their spectral indexes. In addition, dust temperatures were estimated and found to be typical of \\hii\\, regions.  \\item An inspection of the 1\\am\\, CGPS \\hi\\, data reveals \\hi\\, minima  having a good morphological correlation with the \\hii\\, regions at velocity ranges compatible with the velocity spanning by \\gs.  This leads to the conclusion that G98, G100, and G103 are located at the same distance than \\gs.  \\item  From O and OB star catalogues, the massive star candidates to be responsible for the ionized gas were identified.  \\item The obtained age difference among the \\hii\\, regions and the shell, together with their relative location leads us to the conclusion that G98, G100, and G103 may have been created as a consequence of the action of a strong shock produced by the expansion of \\gs\\, into the surrounding gas.  \\end{enumerate}"
    },
    "1112/1112.0666_arXiv.txt": {
        "abstract": "The measurement of the anisotropies of cosmic ray arrival direction provides important informations on the propagation mechanisms and on the identification of their sources. In this paper we report the observation of anisotropy regions at different angular scales. In particular, the observation of a possible anisotropy on scales between $\\sim$10$^{\\circ}$ and $\\sim$30$^{\\circ}$ suggests the presence of unknown features of the magnetic fields the charged cosmic rays propagate through, as well as potential contributions of nearby sources to the total flux of cosmic rays. Evidence of new weaker few-degree excesses throughout the sky region $195^{\\circ}\\leq$ R.A. $\\leq 315^{\\circ}$ is reported for the first time. ",
        "introduction": "As cosmic rays (CRs) are mostly charged nuclei, their arrival direction is deflected and highly isotropized by the action of galactic magnetic field (GMF) they propagate through before reaching the Earth atmosphere. The GMF is the superposition of regular field lines and chaotic contributions. Altough the strength of the non-regular component is still under debate, the local total intensity is supposed to be $B=2\\div 4\\textrm{ $\\mu$G}$. In such a field, the gyroradius of CRs is given by $r_{a.u.}=100\\,R_{\\textrm{\\scriptsize{TV}}}$, where $r_{a.u.}$ is in astronomic units  and $R_{\\textrm{\\scriptsize{TV}}}$ is in TeraVolt.  However, different experiments \\cite{nagashima,kam07,tibet06,milagro09,eastop09,icecube11} observed an energy-dependent \\emph{\"large scale\"} anisotropy in the sidereal time frame with an amplitude of about 10$^{-4}$ - 10$^{-3}$, suggesting the existence of two distint broad regions, one showing an excess of CRs (called \"tail-in\"), distributed around 40$^{\\circ}$ to 90$^{\\circ}$ in Right Ascension (R.A.). The other a deficit (the \"loss cone\"), distributed around 150$^{\\circ}$ to 240$^{\\circ}$ in R.A.. The origin of these anisotropies is still unknown. Some authors claim that it can be explained within the diffusion approximation taking into account the role of the few most nearby and recent sources \\cite{blasi,erlykin06}. Other studies suggest that the observations may be due to a combined effect of the regular and turbolent GMF \\cite{battaner09}, or to local uni- and bi-dimensional inflows \\cite{amenomori10}.  In the last years the Tibet AS$\\gamma$ \\cite{amenomori07} and Milagro \\cite{milagro08} Collaborations reported evidence of the existence of a medium angular scale anisotropy contained in the tail-in region. The observation of similar small scale anisotropies has been recently claimed by the Icecube experiment in the Southern hemisphere \\cite{icecube11}. So far, no theory of CRs in the Galaxy exists which is able to explain few degrees anisotropies in the rigidity region 1-10 TV leaving the standard model of CRs and that of the local GMF unchanged at the same time. More beamed the anisotropies and lower their energy, more difficult to fit the standard model of CRs and GMF to experimental results.  The observation of anisotropy effects at a level of 10$^{-4}$ with an air shower array is a tricky job. A wrong estimation of the exposure may affect the significance and the relative intensity sky maps, even create artifacts (i.e. fake excesses or deficit regions). In fact, drifts in detector response and atmospheric effects on air shower development are quite hard to be modeled to sufficient accuracy. The envisaged solution is to use the data to estimate the detector exposure. But data contain either signal and background events, so that some distortions could be present in the results. The shape and the size of possible artifacts depend on the characteristic angle and time scale over which all the aspects of the data acquisition vary less than the effect to be observed. Therefore, if anisotropies of the order 10$^{-4}$ are looked for, operating conditions must be kept (or made up) stable down to this level, all across the field of view and during all the acquisition time. In addition, as it is well known, the partial sky coverage carried out by a given experiment will bias the observations of structure on largest scales. Finally, the observation of a possible small angular scale anisotropy region contained inside a larger one rely on the capability for suppressing the anisotropic structures at larger scales without, simultaneuosly, introducing effects of the analysis on smaller scales.  The ARGO-YBJ experiment, located at the YangBaJing Cosmic Ray Laboratory (Tibet, P.R. China, 4300 m a.s.l., 606 g/cm$^2$), is an air shower array able to detect the cosmic radiation at an energy threshold of a few hundred GeV. The full detector is in stable data taking since November 2007 with a duty cycle greater than 85\\%. The trigger rate at the threshold is 3.6 kHz. The detector characteristics and performance are described in \\cite{moon11}, the main results are reported in \\cite{demitri11}. In this paper the observation of CR anisotropy at different angular scales with ARGO-YBJ is reported as a function of the primary energy. ",
        "conclusions": "In this paper the observation of CR anisotropy at different angular scales with ARGO-YBJ is reported as a function of the primary energy. The large scale CR anisotropy has been clearly observed up to about 25 TeV. Evidence of existence of different few-degree excesses in the Northern sky (the strongest ones positionally coincident with the regions detected by Milagro in 2008) is reported. The study of the CR anisotropy with air shower arrays is challenging, therefore deeper analysis is under way to investigate possible artifacts produced in the background calculations. However, a joint analysis of concurrent data recorded by different experiments in both hemispheres, as well as a correlation with other observables like the interstellar energetic neutral atoms distribution \\cite{ibex09,ibex11}, should be a high priority to clarify the observations."
    },
    "1112/1112.2349_arXiv.txt": {
        "abstract": "{Dust particles are observed at mm sizes in outer regions of the disk, although theoretically, radial drift does not allow dust particles to form pebbles.} {In order to explain  grain growth to mm sized particles and their retention in outer regions of protoplanetary disks, as it is observed at sub-mm and mm wavelengths, we investigate if strong inhomogeneities in the  gas density profiles can slow down excessive radial drift and can help dust particles to grow.} {We use coagulation/fragmentation and disk-structure models, to simulate the evolution of dust in a bumpy surface density profile which we mimic with a sinusoidal disturbance. For different values of the amplitude and length scale of the bumps, we investigate the ability of this model to produce and retain large particles on million years time scales. In addition,  we introduced a comparison between the pressure   inhomogeneities  considered in this work and the pressure profiles that come from magnetorotational instability. Using the Common Astronomy Software Applications ALMA simulator, we study if   there are observational signatures of these pressure inhomogeneities that can be seen with ALMA.} {We present the favorable conditions to trap dust particles and the corresponding calculations predicting the spectral slope in the mm-wavelength range, to compare with current observations. Finally we present simulated images using different antenna configurations of ALMA  at different frequencies, to show that  the ring structures will be detectable at the distances of the Taurus Auriga or Ophiucus star forming regions.} {} ",
        "introduction": "\\label{sec1}  The study of planet formation is an important field in astronomy with an increasing research since the middle of the twentieth century,  however there are still countless unanswered questions. One of these questions is the observed grain growth to mm sized particles in the outer disk regions (\\cite{1991ApJ...381..250B};  \\cite{2000ApJ...534L.101W, 2005ApJ...626L.109W}; \\cite{2001ApJ...554.1087T,  2003A&A...403..323T}; \\cite{2005ApJ...631.1134A}; \\cite{2006A&A...446..211R}; \\cite{2007prpl.conf..767N}; \\cite{2009ApJ...701..260I}; \\cite{2009A&A...495..869L}; \\cite{2010A&A...512A..15R, 2011A&A...525A..81R}; \\cite{2011A&A...529A.105G}) that suggests a mechanism operating that prevents the rapid inward drift  (\\cite{1997Icar..128..213K}, \\cite{2007A&A...469.1169B}, \\cite{2007Natur.448.1022J}). Different efforts are aimed to explain theoretically the growth from small dust particles to planetesimals, which have led to the development of different numerical models, e.g. \\cite{1981Icar...45..517N},  \\cite{2005A&A...434..971D}, \\cite{2008A&A...480..859B}, \\cite{2008A&A...489..931Z}, \\cite{2009ApJ...698.1122O}, \\cite{2010A&A...513A..79B}.  Due to the fact that circumstellar disks exhibit a wide range of temperatures, they radiate from micron wavelengths to millimeter  wavelengths, which is why  they can be observed with infrared and radio telescopes. With the construction of different kinds of these  telescopes, e.g. Spitzer, Herschel, SMA, EVLA or  ALMA,  astronomers can observe with more details the material inside accretion disks around young stars. The parallel development of theory and observations have allowed astrophysicists  to study the different stages of  planet formation, making this topic  one of the most active fields in astronomy today.  In the first stage of planet formation, the growth from sub-micron sized particles to larger objects is a complex process that contains many physical challenges. In the case of smooth disk with a radial gas pressure profile that is monotonically decreasing, the dust particles drift inwards owing to the fact that the gas moves with sub-keplerian velocity due to the gas pressure gradient. Before a large object can be formed, the radial drift  causes dust pebbles to move towards the star. Moreover, the high relative velocities due to turbulence and radial drift cause  the solid particles to reach velocities  that lead to fragmentation collisions which do not allow dust particles to form larger bodies \\citep{1977MNRAS.180...57W, 2008A&A...480..859B, 2010A&A...516L..14B}. The combination of these two problems is called ``meter-size barrier'' because on timescales as short as 100 years, a one meter sized object at 1 AU moves towards the star due to the radial drift, preventing that any larger object could be formed.  The observations in the inner regions of the disk, where planets like Earth should be formed,  are very difficult because these regions are so small on the sky that few telescopes can spatially resolve them. Also, these regions are optically thick. However, what amounts to the meter-size barrier in the inner few AU is a ``millimeter-size barrier'' in the outer regions of the disk. These outer regions ($\\gtrsim 50$ AU)  are much easier to spatially resolve and are optically thin. Moreover, one can use millimeter observations, which probe precisely the grain size range of the millimeter-size. Therefore, the study of dust growth in the outer disk regions may teach us something about the formation of planets in the inner disk regions.  Observations of protoplanetary disks at sub-millimeter and mm wavelengths show that the disks remain dust-rich during several million years with large particles in the outer regions \\citep{2007prpl.conf..767N, 2010A&A...512A..15R}. However, it is still unclear how to prevent the inward drift and how to explain theoretically that mm-sized particles  are observed in the outer regions of the disk.  Different mechanisms of planetesimal formation have been proposed to resist the rapid inwards drift like: gravitational instabilities \\citep{2002ApJ...580..494Y}, the presence of zonal flows\\citep{2009ApJ...697.1269J, 2011A&A...529A..62J, 2011ApJ...736...85U} or dead zones of viscously accumulated gas which form vortices \\citep{2006A&A...446L..13V}.  With the model presented here, we want to imitate mechanisms that allow to have long-lived pressure inhomogeneities in protoplanetary disks, by artificially adding pressure bumps onto a smooth density profile.  \\begin{table} \\caption{Parameters of the model} \\label{table1} \\centering \\begin{tabular}{c c } \\hline\\hline Parameter & Values \\\\ \\hline $A$ & $\\{0.1; 0.3; 0.5; 0.7\\}$  \\\\ $f$ & $\\{0.3; 0.7; 1.0; 3.0\\}$  \\\\ $\\alpha$ & $10^{-3}$  \\\\ $R_{\\star} [R_\\odot]$ & $2.5$  \\\\ $T_{\\star} [K]$ & $4300$  \\\\ $M_{disk} [M_\\odot]$ & $0.05$  \\\\ $\\rho_s[g/cm^3]$& $1.2$  \\\\ $v_f[m/s]$& $10$  \\\\ \\hline \\end{tabular} \\end{table}  To confront the millimeter-size barrier, it is necessary to stop the radial drift considering a radial gas pressure profile that is not monotonically decreasing with radius. Instead, we take a pressure profile with local maxima adding a sinusoidal perturbation of the density profile. These  perturbations influence directly the pressure, following a simple equation of state for the pressure in the disk. Depending on the size of the particle, the dust grains will be nearly perfectly trapped in the pressure peaks, because a positive pressure gradient can cause those dust particles move outwards. On the other hand, turbulence can mix part of the dust particles out of the bumps, so that overall there may still be some net radial inward drift. More importantly, dust fragmentation may convert part of the large particles into micron size dust particles, which are less easily trapped and thus drift more readily inward.  In the work of \\cite{2010A&A...516L..14B}, they compared the observed fluxes and mm spectral indices from Taurus \\citep{2010A&A...512A..15R} and Ophiucus \\citep{2010A&A...521A..66R} star-forming regions with predicted fluxes and spectral indices at mm wavelengths. They neglected the radial drift, forcing the dust particles to stay  in the outer disk regions. They aimed to keep the spectral index at low values, which implies that the dust particles could acquire millimeter sizes \\citep {1991ApJ...381..250B}. However, they found over-predictions of the fluxes. As an extension of their work, the purpose of this paper is to model the combination of three processes:  the radial drift,  the radial turbulent mixing and  the dust coagulation/fragmentation cycle in a bumpy surface density profile. Our principal aim is to find out how the presence of pressure bumps  can help explain the retainment of dust pebbles in the outer regions of protoplanetary disks, while still allowing for moderate drift and thus obtaning a better match with the observed fluxes and mm spectral indeces. In addition, we show simulated images using different antenna configurations of the complete stage of ALMA, to study if it is possible to detect these kind of inhomogeneities  with future ALMA observations.  This paper is ordered as follows: Sect. \\ref{sec2} will describe the coagulation/fragmentation model and the sinusoidal perturbation that we take for the initial condition of the gas surface density. Section \\ref{sec3}  will describe the results of these simulations, the comparison between existing mm observations of young forming disks and the results from our model.  We discuss  if the type of structures generated by our model can be detectable with  future  ALMA observations. In Sect. \\ref{sec4}, we explore  the relation of our model with  predictions of current simulations of the magnetorotational instability (MRI) \\citep{1991ApJ...376..214B}. Finally, Sect. \\ref{sec5} will summarize our results and the conclusions of this work.   \\begin{figure*} \\centering \\includegraphics[width=18cm]{Figure1.pdf} \\caption{Comparison between: The gas surface density (left plot) taken in this work (Eq. \\ref{eq1}) for two different values of the amplitude and constant width (dashed and dot-dashed lines). The Rossby wave instability \\citep{2011arXiv1109.6177R}, and the presence of zonal flows due to MHD instabilities  \\citep{2011ApJ...736...85U}. Right plot shows the pressure gradient for each of the gas surface density profiles.} \\label{comparacion} \\end{figure*} ",
        "conclusions": "\\label{sec5}  Theoretical models of dust evolution in protoplanetary disks show that the growth   from sub-micron sized particles to larger objects is prevented basically by two phenomena: radial drift and fragmentation. Nevertheless, infrared and radio observations show that millimeter sized particles can survive under those circumstances in the outer regions of disks. Therefore, various theoretical efforts have been focused on explaining the survival of those bodies.  Taking into account strong inhomogeneities expected to be in the gas density profile e.g. zonal flows,  and using the coagulation/fragmentation and disk-structure models by  \\cite{2010A&A...513A..79B}, we have investigated how the presence of pressure bumps can cause the reduction of radial drift, allowing the existence of millimeter sized grains in agreement with observations. In this work, we assumed a sinusoidal function for the gas surface density  to simulate pressure bumps. The amplitude and wavelength disturbances are chosen considering the necessary conditions to have outward angular momentum transport in an $\\alpha$-turbulent type disk, outward radial drift of dust and reasonable values compare to predictions from the recent work of zonal flows \\citep{2011ApJ...736...85U}.  The results presented here suggest that the presence of pressure bumps with a width of the order of the disk scale-height and an amplitude of $30\\%$ of the gas surface density of the disk, provide the necessary physical conditions for the survival of larger grains in a disk with properties summarized in Table ~\\ref{table1}.  Comparisons between the observed fluxes of the Taurus, Ophiucus and Orion Nebula Cluster star forming regions with the results of the models ratify that the effect of the radial drift is reduced allowing particles to grow. Figure \\ref{Fig3}  shows how models with these kind of disturbances  reproduce much better mm-observations than models with full or without radial drift.  In addition, we presented a comparison between the bumpy density profile assumed in this work and 3D MHD models of zonal flows that can cause long lived bumps in protoplanetary disks. We showed that the pressure bumps cause by zonal flows of  \\citep{2011ApJ...736...85U} are in agreement with the amplitudes and wavelengths used in this work. Therefore, taking those bumps, the survival of dust particles is possible in the outer regions after some Myr.  The simulated images using CASA ALMA simulator (version 3.2.0) show that, with different antenna configuration of the final ALMA stage, the ring structures, due to the presence of the pressure bumps, should be detectable. Future ALMA observations will have an important impact for understanding the first stages of planet formation and it will  be very important to investigate if the  grain growth and retetion can be explained with the presence of these kind of inhomogeneities in the gas density profile."
    },
    "1112/1112.2692.txt": {
        "abstract": "% Context: {Recent work identified a growth barrier for dust coagulation that originates in the electric repulsion between colliding particles. Depending on its charge state, dust material may have the potential to control key processes towards planet formation such as MHD (magnetohydrodynamic) turbulence and grain growth which are coupled in a two-way process.} %Aims: {We quantify the grain charging at different stages of disc evolution and differentiate between two very extreme cases: compact spherical grains and aggregates with fractal dimension $D_f = 2$.} %Methods: {Applying a simple chemical network that accounts for collisional charging of grains, we provide a semi-analytical solution. This allowed us to calculate the equilibrium population of grain charges and the ionisation fraction efficiently. The grain charging was evaluated for different dynamical environments ranging from static to non-stationary disc configurations.} %Results {The results show that the adsorption/desorption of neutral gas-phase heavy metals, such as magnesium, effects the charging state of grains. The greater the difference between the thermal velocities of the metal and the dominant molecular ion, the greater the change in the mean grain charge. Agglomerates have more negative excess charge on average than compact spherical particles of the same mass. The rise in the mean grain charge is proportional to $N^{1/6}$ in the ion-dust limit. We find that grain charging in a non-stationary disc environment is expected to lead to similar results.} % Conclusions {The results indicate that the dust growth and settling in regions where the dust growth is limited by the so-called ''electro-static barrier'' do not prevent the dust material from remaining the dominant charge carrier.} ",
        "introduction": "\\label{sec1} Small dust particles are regarded as the key ingredient to control the planet formation process in protoplanetary discs. At early stages of planet formation, submicron-sized compact particles are thought to grow towards larger but fluffy agglomerates. At mm to cm sizes these agglomerates are supposed to be compacted, whereby they loose their fractal structure. Because of the rise in the mass-to-surface ratio associated with the compaction, agglomerates tend to be more concentrated locally. However, there are a number of potential processes that may control grain growth, one of which is disc turbulence. The magnetorotational instability - MRI - (Balbus \\& Hawley 1991; Hawley \\& Balbus 1991) has been shown to be robust in generating turbulence in Keplerian discs. MHD turbulence provides the internal stress required for mass accretion, the rate of which is constrained by observations. Observations of young stars indicate that discs usually show signatures of active gas accretion onto the central star with mass flow rates of about $10^{-8\\pm1} \\ \\rm M_{\\odot} yr^{-1}$ (e.g., Sicilia-Aguilar et al. 2004). Recent observations from discs around young stars also set constraints on the turbulent linewidth and provide support for subsonic turbulence in discs (Hughes et al. 2011). However, numerical studies (e.g. Fleming \\& Stone 2003) on MRI-driven MHD turbulence have identified locations within the planet forming regions, the so-called ''dead zones'', where the coupling between the gas and the magnetic field is not sufficient to maintain MHD turbulence. Fleming \\& Stone also showed that a low Reynolds stress can be maintained in the dead zone, such that low levels of accretion are sustained there. That is why dead zones have been considered to advance the planet formation process. Dead zones provide a safe environment for grain growth and for planetesimals (Gressel et al. 2011). However, the presence of small grains in the weakly ionized dead zones leads to significant changes in the abundances of the charge carriers in the gas phase and therefore affects the MRI. In other words, MHD turbulence and grain growth are coupled in a two-way process.\\\\ \\indent Grain charging may also effect the dust growth: Okuzumi (2009) pointed out that the so-called ''electro-static barrier'' may inhibit dust coagulation in planet forming regions. In a subsequent study, Okuzumi et al. (2011a) extended the analysis including the dust size distribution. They found that under certain conditions the dust undergoes bimodal growth. While the small aggregates sweep up free electrons, the large aggregates continue to grow. However, the growing stalls if the electrostatic repulsion becomes strong and the collisions are driven by Brownian motion. The results of Okuzumi et al. (2011b) indicate that under minimum mass solar nebula conditions the dust growth halts at aggregate sizes beyond $[4 \\cdot 10^{-5}, 3\\cdot 10^{-2}] \\ \\rm cm$ depending on the radial position. Conventional chemical models studying the dust-grain chemistry in protoplanetary discs account for up to $|Z| = 3$ grain charges (e.g. Sano et al. 2000). More recently, the range of grain charges was extended considerably for various purposes. Focusing on surface layers of T-Tauri and transitional discs, Perez-Becker  \\& Chiang (2011) examined the charge distribution on grains covering grain charges $[Z_{\\rm min}, Z_{\\rm max}] = [-200, 200]$. Okuzumi (2009) calculated the charge distribution on dust aggregates. He demonstrated that aggregates made of $N=10^{10}$ monomers can carry excess charges of about $Z_{\\rm min} \\sim - 10^5$. To account for higher grain charges we therefore improved our simple chemical network introduced in Ilgner \\& Nelson (2006a) allowing $[Z_{\\rm min}, Z_{\\rm max}] $ to become free of choice and appropriate to the conditions applied.\\\\ % \\indent In this paper we examine the grain charging for various stages of disc evolution ranging from static to non-stationary disc environment. Linking the grain charging with the gas-dust dynamics provides a natural environment to mimic the variation of the dust-to-gas ratio $\\Sigma_{\\rm d}/\\Sigma_{\\rm g}$. The grain charging is assumed to originate in collisional charging processes between grains, electrons, and gas-phase ions. We include X-ray ionisation from the central star as the primary source for ionisation and consider, to some extent, the contributions from cosmic ray particles. We furthermore assume an equilibrium distribution of dust material in vertical direction balancing turbulent stirring and sedimentation. In particular, we apply the model of Takeuchi \\& Lin (2005) to simulate the dynamics of the gas-dust disc. One of our primary goals is to compare the grain charging obtained for compact spheres and dust agglomerates and the mean grain charge in particular.\\\\ % \\indent We have investigated the effect that thermal adsorption/desorption of metals has on grain charging using two different chemical networks. In comparison with the values obtained by switching off the thermal adsorption of metals, the modified Oppenheimer-Dalgarno model produces lower, the Umebayashi-Nakano model higher values for the mean grain charge. This originates in differences in the thermal velocities of the dominant gas-phase ion: $v_{\\rm HCO^+} < v_{\\rm Mg^+} < v_{\\rm H_2^+}$. In line with expectations, we find that dust agglomerates have a higher charge-to-mass ratio carrying more charges than the corresponding compact spherical particles. Considering stationary disc configurations, we observe that grain charging on dust agglomerates is always associated with the so-called \"ion-dust regime\",  where the charge balance is mainly maintained by negatively charged grains and positively charged gas-phase ions. (The ion-electron regime is characterised by the balance between gas-phase ions and electrons.) In particular, we show that the fractional abundances of the charge carriers are independent of the number $N$ of constituent monomers. Concerning compact spherical dust particles, we identified regions $z/h_{\\rm g} > 0$ associated with the ion-electron regime. Another conclusion of our work is that  - for the parameter range considered (i.e., compact radii $a_{\\ast} \\le 10^{-3} \\ \\rm cm$ and agglomerates with characteristic radii $a_{\\rm c} \\le 10^{-2} \\ \\rm cm$ with $N \\le 10^6$, respectively\\footnote{The characteristic and compact radius is defined in Eqs. (\\ref{eq16:sec2}) and (\\ref{eq18:sec2}), respectively.}) - grain charging in a non-stationary disc environment is expected to lead to similar results. We also present semi-analytical solutions for both agglomerates and compact spheres that allow us to determine the equilibrium distribution of  $<\\!\\!Z\\!\\!>$, $\\sqrt{<\\!\\!\\Delta Z^2\\!\\!>}$, and the fractional abundances of the charge carriers a priori.\\\\ % \\indent The paper is organised as follows. In Sect.2  we discuss the disc model we apply. In Sect. 3 we introduce the chemical network that we used to compare our results with the reaction network Okuzumi (2009) applied. Key problems concerning the ionisation rates are discussed in Sect. 4 while the numerical methods applied are described in Sect. 5. In Sect. 6 we present the results of our model, and finally in Sect. 7 we summarise the key findings of our study.  % ================================================================= % New chapter: % ================================================================= ",
        "conclusions": "\\label{sec7} We have presented calculations of the grain charging under conditions that mimic different stages of protoplanetary disc evolution. Instead of parametrising the dust-to-gas ratio, we inferred the  value for $\\Sigma_{\\rm d}(R)/\\Sigma_{\\rm g}(R)$ from the evolution of the gas-dust disc. In particular, we applied the disc model of Takeuchi \\& Lin (2002, 2005) and content ourself with order-of-magnitude estimates. We considered collisional charging as the dominant process to determine the charge state of grains. For that purpose, we generalised the modified Oppenheimer-Dalgarno model introduced in Ilgner \\& Nelson (2006a) to account for higher grain charges. Based on that simple chemical network, we examined the grain charging for two different types of grain topology: compact spherical grains and fractal agglomerates of $D_{\\rm f} = 2$. Our main conclusions are:\\\\ \\begin{enumerate} \\item The effect that thermal adsorption/desorption of metals has on grain charging depends on the mass of the dominant molecular gas-phase ion. If its mass is heavier than that of metal, the inclusion of the thermal adsorption of metals results in high negative excess charges on grains on average. Less negative charge excess on average is observed for molecular ions lighter than metals. \\item  We extended the semi-analytical method proposed by Okuzumi (2009), which allowed us to determine steady-state solutions for the mean grain charge, electron and ion abundances associated with the modified Oppenheimer-Dalgarno model. The semi-analytical solutions were derived for both grain topologies: compact grains and fractal aggregates. \\item The results obtained confirm that dust agglomerates have a higher charge-to-mass ratio than the corresponding compact spheres. \\item We found that reducing the number $N$ of constituent monomers causes a drop in the mean value $<\\!\\!Z\\!\\!>$ of the grain charge and the standard deviation $<\\!\\!\\sqrt{\\Delta Z^2}\\!\\!>$. Under conditions valid for the ion-dust plasma limit (i.e., $x_{\\rm i} \\gg x_{\\rm e}$) the profiles for $<\\!\\!Z\\!\\!>x_{\\rm d}$, $x_{\\rm e}$, and $x_{\\rm i}$ remain unchanged with varying $N$. This is because the cumulative projected surface area of all BCCA agglomerates is independent of $N$. \\item The results obtained by switching from one type of grain topology (fractal agglomerates) to another (compact spheres) reveal that for compact spherical grains $<\\!\\!Z\\!\\!>$ is shifted towards more negative values while $<\\!\\!\\sqrt{\\Delta Z^2}\\!\\!>$ decreases. \\item The results obtained for agglomerate sizes $a_{\\rm c} = [10^{-4}, 10^{-2}] \\ \\rm cm$ indicate that the grain charging of BCCA agglomerates is governed by the ion-dust plasma limit. Transitions from the ion-dust to the ion-electron plasma limit are observed for compact spheres depending on altitude $z/h_{\\rm g}$. \\item Another conclusion is that grain charging in a non-stationary disc environment is expected to lead to similar results as long as the effective ionisation rate and the temperature of the gas match the stationary values. \\end{enumerate} We regard the drastically simplified description of the temperatures of the gas and the dust mentioned above as a potentially serious omission of our model. Working out more realistic conditions may result in significant changes in the local disc structure and might be a potential problem for future investigations.  \\begin{acknowledgement} \\noindent I appreciate the discussions with Hiroshi Kobayashi, Satoshi Okuzumi, and Taku Takeuchi very much indeed. \\end{acknowledgement}  % ================================================================= % References % ================================================================="
    },
    "1112/1112.6040_arXiv.txt": {
        "abstract": "One possibility for explaining the apparent accelerating expansion of the universe is that we live in the center of a spherically inhomogeneous universe. Although current observations cannot fully distinguish $\\Lambda$CDM and these inhomogeneous models, direct measurement of the acceleration of the universe can be a powerful tool in probing them. We have shown that, if $\\Lambda$CDM is the correct model, DECIGO/BBO would be able to detect the positive redshift drift (which is the time evolution of the source redshift $z$) in 3--5 year gravitational wave (GW) observations from neutron-star binaries, which enables us to rule out any Lema\\^itre-Tolman-Bondi (LTB) void model with monotonically increasing density profile. We may even be able to rule out any LTB model unless we allow unrealistically steep density profile at $z\\sim 0$. This test can be performed with GW observations alone, without any reference to electromagnetic observations, and is more powerful than the redshift drift measurement using Lyman $\\alpha$ forest. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.3948_arXiv.txt": {
        "abstract": "It is well known that the long-term evolution of the photospheric magnetic field plays an important role in building up free energy to power solar eruptions. Observations, despite being controversial, have also revealed a rapid and permanent variation of the photospheric magnetic field in response to the coronal magnetic field restructuring during the eruption. The \\hmi\\ instrument (HMI) on board the newly launched \\Sdo\\ (\\sdo) produces seeing-free full-disk vector magnetograms at consistently high resolution and high cadence, which finally makes possible an unambiguous and comprehensive study of this important back-reaction process. In this study, we present a near disk-center, \\goes-class X2.2 flare, which occurred in NOAA AR 11158 on 2011 February 15. Using the magnetic field measurements made by HMI, we obtained the first solid evidence of a rapid (in about 30 minutes) and irreversible enhancement in the horizontal magnetic field at the flaring magnetic polarity inversion line (PIL) by a magnitude of $\\sim$30\\%. It is also shown that the photospheric field becomes more sheared and more inclined. This field evolution is unequivocally associated with the flare occurrence in this sigmoidal active region, with the enhancement area located in between the two chromospheric flare ribbons and the initial conjugate hard X-ray footpoints. These results strongly corroborate our previous conjecture that the photospheric magnetic field near the PIL must become more horizontal after eruptions, which could be related to the newly formed low-lying fields resulted from the tether-cutting reconnection. ",
        "introduction": "Almost two decades ago, we discovered rapid and permanent changes of vector magnetic fields associated with flares \\citep{wang92,wang94}. Specifically, the transverse field near the flaring magnetic polarity inversion line (PIL) is found to enhance substantially and irreversibly across the time duration of the flare, which is also often accompanied by an increase of magnetic shear. Similar trend indicating a more horizontal orientation of the photospheric magnetic field after flares and coronal mass ejections (CMEs) has continued to be observed later on in many observations \\citep{wang02b,wang04,wang+liu05,liu05,wang07a,jing08,li09,chang11}, and shows some agreement with recent model predictions \\citep{li10}. Nevertheless, a majority of such studies are unavoidably hampered by the obvious limitations, ground-based observations (e.g., seeing variation and the limited number of observing spectral positions), probably because of which mixed results were also reported \\citep{ambastha93,hagyard99,chen94,li00a,li00b}.  On the other hand, flare-related variations in the line-of-sight (LOS) component of photospheric magnetic field have been clearly recognized \\citep[e.g.,][]{wang02b,spirock02,yurchyshyn04,sudol05,wang06,wang10,petrie10}. In particular, the feature of unbalanced flux evolution of the opposite polarities could provide an indirect evidence for the more horizontal orientation of photospheric fields after flares/CMEs \\citep{wang10}. However, it is noted that the changes of the LOS field alone cannot provide complete understanding of the field restructuring \\citep{hudson11}.  It is notable that vector magnetic field data has been made available with the \\hmi\\ (HMI) instrument \\citep{schou11} on board the newly launched \\Sdo\\ (\\sdo). Its unprecedented observing capabilities give a favorable opportunity to finally resolve any uncertainties regarding the evolution of photospheric magnetic field in relation to flares/CMEs.  In this study, we investigate a near disk-center X2.2 flare on 2011 February 15, which provides the first solid evidence of the enhancement in the horizontal field at the flaring PIL using the seeing-free HMI data. We will discuss the implications of such a change in the context of magnetic reconnection model for flares. ",
        "conclusions": "We have used the unprecedented SDO/HMI vector field observations to analyze the changes of the photospheric magnetic field associated with the first X-class flare in the solar cycle 24, with the aid of images of flare emissions in multiple wavelengths. Main results are as follows.  \\begin{enumerate} \\item A compact region R along the flaring PIL shows a rapid and permanent enhancement of $\\langle B_h \\rangle$ by 400~G (\\sm30\\% of the preflare magnitude) within about 30~minutes, which has a close temporal relationship with the flare HXR emission. Meanwhile, the nonpotentiality represented by magnetic shear also exhibits a pronounced increase near the surface.  \\item The initial HXR sources FP1 and FP2 as well as the double J-shaped flare ribbons are at the two ends of the region R lying at the central of this sigmoidal active region. Two additional flare footpoints FP3 and FP4 are clearly seen in the hot 94~\\AA\\ channel, located at the far ends of the sigmoid. We suggest that the tether-cutting reconnection \\citep{moore01} between the loops FP3--FP2 and FP4--FP1 produces the short and low-lying loops FP1--FP2, which could explain the enhanced $B_h$ as well as $\\tilde{S}$ and $\\mathring{S}$ at the region R \\citep{melrose97}. The detected enhancement of nonpotentiality on the surface could also be due to the newly emerging, sheared fields \\citep{jing08}. \\end{enumerate}  In summary, the HMI observations presented in this study constitute the first solid evidence of flare-induced changes of the photospheric magnetic field, which strongly endorses our previous results using ground-based vector magnetograms subject to seeing variation \\citep[][and references therein]{wang10}. The unambiguously observed enhancement of horizontal field on the surface strongly suggests that the photospheric magnetic field could respond to the coronal field restructuring by tilting toward the surface (i.e., toward a more horizontal state) near the PIL, and that this development may be due to the tether-cutting reconnection producing the flare. This view is also well in line with the recent theoretical development \\citep{hudson08,fisher10}, where the back reaction on the solar surface resulting from the coronal field evolution as required by the energy release is quantitatively assessed. Further systematic studies of flare-related magnetic field change, especially when aided with extrapolation models, are promising to provide further insight into the relationship between the surface field change and coronal magnetic reconnection \\citep[e.g.,][]{sun11,chang11}."
    },
    "1112/1112.3386_arXiv.txt": {
        "abstract": "We present 15-20 \\mum\\, spectral maps towards the reflection nebula NGC\\,2023 obtained with the Infrared Spectrograph in short-wavelength, high-resolution mode on board the Spitzer Space Telescope. These spectra reveal emission from PAHs, C$_{60}$, and H$_2$ superposed on a dust continuum. These emission components exhibit distinct spatial distributions: with increasing distance from the illuminating star, we observe the PAH emission followed by the dust continuum emission and the H$_2$ emission. The C$_{60}$ emission is located closest to the illuminating star in the south while in the north, it seems to be associated with the H/H$_2$ transition. Emission from PAHs and PAH-related species produce features at 15.8, 16.4, 17.4, and 17.8 \\mum\\, and the 15-18 \\mum\\, plateau. These different PAH features show distinct spatial distributions. The 15.8 \\mum\\, band and 15-18 \\mum\\, plateau correlate with the 11.2 \\mum\\, PAH band and with each other, and are attributed to large, neutral PAHs. Conversely, the 16.4 \\mum\\, feature correlates with the 12.7 \\mum\\, PAH band, suggesting that both arise from species that are favored by the same conditions that favor PAH cations. The PAH contribution to the 17.4 \\mum\\, band is displaced towards the illuminating star with respect to the 11.2 and 12.7 \\mum\\, emission and is assigned to doubly ionized PAHs and/or a subset of cationic PAHs. The spatial distribution of the 17.8 \\mum\\, band suggests it arises from both neutral and cationic PAHs. In contrast to their intensities, the profiles of the PAH bands and the 15-18 \\mum\\, plateau do not vary spatially. Consequently, we conclude that the carrier of the 15-18 \\mum\\, plateau is distinct from that of the PAH bands. ",
        "introduction": "\\label{intro}  Strong emission bands at 3.3, 6.2, 7.7 and 11.3 \\mum\\,- the so-called unidentified infrared (UIR) bands - dominate the mid-IR spectra of almost all objects, including reflection nebulae, planetary nebulae, the interstellar medium and \\HII\\, regions. These bands are generally attributed to IR fluorescence of a family of Polycyclic Aromatic Hydrocarbon molecules (PAHs) pumped by the UV radiation field. A key result in the observational studies of PAHs is that their mid-IR bands show clear variations in peak positions, shapes and (relative) intensities, not only between sources, but also spatially within extended sources \\citep[e.g.][]{Hony:oops:01, Peeters:prof6:02, Brandl:06, SmithJD:07, Galliano:08}. The variability in intensity ratio of the main PAH bands has been interpreted in terms of the charge state and excitation levels of the PAHs \\citep[e.g.][]{Allamandola:modelobs:99, Galliano:08} while the variation in the profiles of the main PAH bands \\citep[classified as class A, B and C profiles;][]{Peeters:prof6:02, vanDiedenhoven:chvscc:03} is attributed to the environment \\citep[CSM versus ISM;][]{Peeters:prof6:02, VanKerckhoven:phd:02, Boersma:08} and the degree of UV processing \\citep{Sloan:07, Tielens:08, Boersma:08}.   In addition to the well-studied main PAH bands, a plethora of weaker PAH bands are present, for example in the 15-20 \\mum\\, region. PAH emission in this region was first observed with the Infrared Space Observatory (ISO): new features were reported at 15.8, 16.4, 17.4 and 17.9 \\mum\\, with the 16.4 and then 17.4 \\mum\\, PAH bands being the most prominent \\citep{Beintema:pahs:96, Moutou:16.4:00, Sturm:swsgal:00, VanKerckhoven:plat:00}. In addition, \\citet{VanKerckhoven:plat:00} presented evidence for a variable, broad plateau from 15 to 20 \\mum\\, which seems to be present solely in \\HII\\, regions \\citep{VanKerckhoven:plat:00, Peeters:plat:04, Peeters:spitzer:04}. Based upon observations with the Spitzer Space Telescope, additional broad components are reported at 16.6, 17.0 and 17.2 \\mum\\, \\citep{SmithJD:04, SmithJD:07, Sellgren:07}. But it is really the high sensitivity of Spitzer that revealed the omnipresence of these bands within the Milky Way and in other galaxies \\citep[e.g.][]{Werner:04, SmithJD:04, Brandl:04, Sellgren:07, SmithJD:07, Tappe:06}, allowing - for the first time -  systematic investigations of these weaker bands.  The 15-20 \\mum\\, emission is attributed to C-C-C bending vibrations \\citep[e.g.][]{ATB, Moutou:firempahs:96, VanKerckhoven:plat:00, Peeters:plat:04, Mattioda:09, Boersma:10, Ricca:10}. As a consequence, the bands in the 15-20 \\mum\\, region are expected to reveal more about the overall molecular structure of the carriers than do the major bands below 15 \\mum. An intriguing result of the Spitzer observations is the remarkable similarity (to first order) of the 15-20 \\mum\\, PAH emission spectra from regions spanning a large range of physical conditions. This suggests that the astronomical PAH population may be dominated by a handful of stable, molecular structures (cf. survival of the fittest) rather than comprised of a large number of PAHs with widely varying structures \\citep{Boersma:10}. Hence, a systematic study of a large sample of spectra showing these weaker PAH bands will provide a complementary view on the characteristics of the emitting PAH population.  This paper reports such a study.  Here we analyze Spitzer-IRS spectral maps of NGC\\,2023 in the 15-20 \\mum\\, region. In Section \\ref{source},  we describe the reflection nebula NGC\\,2023, while the observations and data reduction are discussed in Section \\ref{data}. The data analysis is presented in Section \\ref{analysis} and discussed in Section \\ref{discussion}. We end with a short summary in Section \\ref{conclusion}.   \\begin{figure}[t!] \\centering \\resizebox{14cm}{!}{% \\includegraphics{fig1.eps}} \\caption{The IRAC [8.0] image of NGC\\,2023 with the SH FOV shown (white) for both the north and south positions studied here. The star HD37903 is indicated by a black circle, sources A, C and D are from \\citet{Sellgren:83}, and the white circles indicate 2MASS point sources located inside the SH apertures. S refers to the south ridge, SSE to the south-southeastern ridge, N to the north ridge, and NW to the northwestern ridge. Maps shown in this paper use the orientation denoted by the vectors (X, Y).} \\label{fov} \\end{figure} ",
        "conclusions": "\\label{conclusion}  We studied the emission in the 15-20 \\mum\\, region towards the reflection nebula NGC\\,2023 by analysing two spectral maps obtained with Spitzer's Infrared Spectrograph (IRS), short-wavelength high-resolution mode. We observed PAH emission bands at 15.8, 16.4, 17.4, and 17.8 \\mum, a broad PAH plateau between 15-18 \\mum, C$_{60}$ emission at 17.4 and 18.9 \\mum, and H$_2$ emission at 17.0 \\mum\\, superposed on a dust continuum. \\\\  We found distinct spatial distributions for these emission components and found tight intensity correlations between some emission components. Based on these results, we collected the emission components in seven groups and discussed their specific characteristics.  {\\it Group 1: the 11.2, 15.8 \\mum\\, PAH bands and the 15-18 \\mum\\, PAH plateau.} We attributed this group to large, neutral PAHs.  {\\it Group 2: the 12.7 and 16.4 \\mum\\, PAH bands.} Compared to Group 1, these bands are displaced towards the illuminating star. We concluded that they must arise from species that are co-spatial with the cationic portion of the emitting PAH population.  {\\it Group 3: the 17.4 \\mum\\, diffuse component.} We estimated the PAH contribution to the 17.4 \\mum\\, band and found that its spatial distribution is displaced towards the illuminating star compared to Groups 1 and 2. We assigned this PAH band to doubly ionized PAHs and/or a subset of the cationic PAH population, such as for example dehydrogenated PAHs.  {\\it Group 4: the 17.8 \\mum\\, PAH band.} This band has a spatial distribution in between that of Groups 1 and 2. We suggested this band arises from both neutral and ionised PAHs.  {\\it Group 5: the 18.9 and (some of) the 17.4 \\mum\\, features.} The C$_{60}$ emission shows a spatial distribution distinct from that of the PAH emission, consistent with \\citet{Sellgren:10}.  Moreover, the spatial behavior of C$_{60}$ towards NGC2023 is different in the south and north. Specifically, the C$_{60}$ emission is located closer to the illuminating star compared to the PAH and H$_2$ emission in the south while in the north, it seems to be associated to the H/H$_2$ transition.  {\\it Group 6: the underlying continuum producing the emission measured here at 14.9 and 19.2 \\mum.} The continuum emission is strongest deeper into the PDR compared to the PAH emission. This is consistent with the VSG component attributed to PAH clusters by, for example, \\citet{Rapacioli:05} and \\citet{Berne:07}.  {\\it Group 7: the H$_2$ emission.} The H$_2$ emission is located furthest into the PDR.\\\\  We also investigated the profiles of the PAH bands and 15-18 \\mum\\, PAH plateau. All PAH bands have a symmetric band profile except for the 16.4 \\mum\\, band, which is slightly asymmetric. We did not observe any changes in the band profiles spatially across the maps. Likewise, we found that the profile of the 15-18 \\mum\\, PAH plateau does not vary.  We concluded that the carrier of the underlying 15-18 \\mum\\, plateau is distinct from the individual PAH bands located on top of it."
    },
    "1112/1112.4790_arXiv.txt": {
        "abstract": "In a power spectrum of the Sun's acoustic waves, the $p$ modes have distinctly skewed frequency profiles. Furthermore, the asymmetry is observed to have the opposite sign in power spectra made from line-of-sight velocity and continuum intensity. The asymmetry and its reversal in sign has previously been explained using a combination of mechanisms that involve the acoustic source. A localized acoustic source within an acoustic cavity naturally generates asymmetric profiles through wave interference; however, the sign of the asymmetry due to this mechanism is identical for all observables. The reversal of the asymmetry between velocity and intensity observations has been attributed to the visibility of the source itself (i.e., ``correlated noise\"). In this paper, I will show that asymmetry generated by a localized source can be interpreted as either a wave interference effect in physical space, or a mode interference effect in spectral space. I advocate a new mode-fitting procedure based on this new interpretation, whereby the complex phases of all the modes determine the mode asymmetries. Further, I suggest that information about the acoustic source function is encapsulated in the amplitude of each mode, and present a scheme by which the source function can be obtained from measured mode amplitudes by standard helioseismic inversion techniques. ",
        "introduction": "\\label{sec:introduction} \\setcounter{equation}{0}  In power spectra of solar acoustic oscillations the frequency profiles of the $p$ modes are skewed and markedly non-Lorentzian. This property was first discovered using Ca II K observations made from the geographic South Pole \\citep{Duvall:1993} and were later confirmed \\citep{Nigam:1998a} with modern spacecraft-based telescopes using data from the Michelson Doppler Imager (MDI) onboard the Solar and Heliospheric Observatory (SOHO). In both sets of observations, a puzzling fact was revealed: for modes of the same frequency and harmonic degree, the asymmetry or skewness of the line profile has opposite signs for observations made in the continuum intensity and line-of-sight Doppler velocity. For both types of observations, the lowest-frequency modes are the most severely skewed and the asymmetry appears to be only a weak function of harmonic degree; however, for modes of high harmonic degree ($l \\gtrsim 150$) this statement is difficult to verify because the mode peaks are highly blended with the spatial leaks from modes of nearby degree, forming a broad ridge of power. Figures~\\ref{fig:GONGpower_l=190} and \\ref{fig:GONGpower_l=20} illustrate the asymmetry for both high and low degree. From the figures it is evident that the frequency profile of each individual mode is asymmetric and, at high degree, the blended ridge is also asymmetric.  The physical mechanism for this skewness depends somewhat on whether we are considering the asymmetry of each mode profile or the asymmetry of a blended ridge. In the former case, the asymmetry of each mode is now thought to arise from a combination of effects due to the nature and location of the acoustic wave source. At high harmonic degree, the asymmetry of the $p$ mode ridge is the result of the skewness of each component mode profile as well as an uneven distribution of power amongst the spatial leaks. The wave source can introduce skewness in a variety of ways. The first mechanism that I discuss here relies in detail on the location of the acoustic source. The primary source of acoustic waves in the Sun is granular convection. Short-lived convective events of relatively high Mach number prolifically emit sound waves.  Since granulation is confined to a narrow band just below the photosphere, the acoustic source is highly localized in depth. Waves generated in this layer propagate both upwards and downwards. Those that propagate downwards are eventually refracted back upwards at the bottom of the acoustic cavity by the sound speed gradient. When this initially downward propagating wave returns to the excitation layer, it can destructively interfere with the wave that was initially upward propagating. At particular frequencies, complete destructive interference occurs and zero wave amplitude is observed above the excitation layer (which is where the spectral lines used in actual observations are formed). The frequencies at which the interference is complete are asymmetrically placed between the mode ridges, thereby leading to asymmetric mode profiles \\citep{Gabriel:1992, Gabriel:1995, Roxburgh:1995, Abrams:1996, Rast:1998}.  The second mechanism due to the source arises because the convective events that generate sound waves are themselves regions of intense temperature fluctuation and high velocity. The sources therefore provide a contribution to the power in either intensity or velocity observations. Since the acoustic waves are correlated with the source that generates them, the visibility of the source can lead to power that varies significantly over the line profile and cause skewness. This mechanism has been dubbed ``correlated noise\" and has been suggested as the reason why the asymmetry is opposite for intensity and velocity spectra \\citep{Roxburgh:1997, Nigam:1998a, Nigam:1998b}. The current understanding is that in order to reproduce the asymmetry of both intensity and velocity power spectra and to explain the rapid variation with frequency of the phase of the cross-spectra between intensity and velocity \\citep{Oliviero:1999}, correlated noise must provide significant power to both types of spectra \\citep[e.g.,][]{Nigam:1999, Skartlien:2000, Jefferies:2003}.  Numerical simulations of granulation have provided a counter explanation \\citep{Georgobiani:2003}. In these simulations, the velocity and intensity fluctuations have power spectra with asymmetric mode profiles. However, if the two quantities are sampled at the same geometric height the sign of the asymmetries are identical. Only after artificial Dopplergrams and intensity images have been constructed using radiative transfer do the two types of power spectra exhibit opposite asymmetries. Therefore, these simulations suggest that details of the radiative transfer, such as the different heights at which the two observables are formed and the nonlinear interaction between wave fluctuations and the radiation field, are responsible for the opposite skewnesses.  Unfortunately, many mode-fitting schemes used to measure $p$ mode frequencies assume that the line profile is symmetric, a Lorentzian in fact. \\cite{Duvall:1993} estimate that failing to account for the asymmetry in the mode's frequency profile can lead to mismeasurement of mode frequencies by as much as one part in $10^4$, a significant error for the purposes of helioseismic inversions. This is probably the source of the systematic offset ($\\sim 0.1~\\mu$Hz) that was observed between frequencies obtained from intensity and velocity spectra \\citep{Toutain:1997}. To account for such systematic shifts, researchers have suggested a variety of asymmetric model profiles with which to fit the observations \\citep[e.g.,][]{Duvall:1993, Nigam:1998b, Rosenthal:1998}. Due to its practicality, the formulation of \\cite{Nigam:1998b} is probably the most widely used, having been adapted for both the measurement of global mode frequencies \\citep{Toutain:1998, Reiter:2002, Korzennik:2005} and the fitting of ring-analysis spectra \\citep{Basu:1999, Basu:2001, Tripathy:2009}. In essence, the profile of \\cite{Nigam:1998b} is a Lorentzian times a polynomial in frequency, where the polynomial has been calculated by performing a low order expansion about the mode frequency $\\omega_0$,  \\begin{eqnarray} P(\\omega) &=& \\frac{A}{1+x^2} \\, \\left[s^2 + (1 + sx)^2\\right] \\; , \\label{eqn:Nigam} \\\\ x &\\equiv& (\\omega-\\omega_0)/\\gamma \\; . \\end{eqnarray}  \\noindent In this formula, $\\gamma$ is the linewidth, $A$ is the power amplitude, and $s$ is an asymmetry parameter.  This expression is formally valid under fairly limited conditions. For example, it has been derived as a small argument expansion in terms of $sx$, which requires that asymmetry is weak and more restrictively that we confine our attention to the core of the line and avoid the wings. One clear sign of this restriction is that the profile isn't integrable, and does not contain a finite amount of energy. Second, it ignores the presence of nearby modes which may contribute to the power. This is a particularly onerous assumption for high degree modes for which the line widths are sufficiently large that the mode and the neighboring spatial leaks blend into a ridge. Figure~\\ref{fig:GONGpower_l=190} illustrates this later point quite well. Each ridge is a blending of mode peaks and the ridges are wide enough that they have begun to overlap.  Our goal here is to generate a fitting function that takes the skewness of the $p$-mode frequency profiles into account while maintaining its validity over a wide range of conditions. I do so by using a modal expansion of the wavefield to represent the power spectra using a small number of free parameters. This expansion is accomplished in a general manner that doesn't depend on details of the atmosphere or the mechanisms for wave damping and excitation. In the next section, Section~\\ref{sec:driven}, I discuss the driving and damping of acoustic $p$ mode oscillations in broad terms. In Section~\\ref{sec:ModeDecomp}, I derive the modal expansion of the wave field and compute the resulting power spectrum. In Section~\\ref{sec:ModeFitting}, I promulgate a fitting profile of the spectrum that treats the asymmetry as a mode interference phenomenon. Finally, in Section~\\ref{sec:discussion}, I discuss my findings and suggests ways in which properties of the acoustic source could be deduced through inversion of the fitted complex mode amplitudes. ",
        "conclusions": "\\label{sec:discussion} \\setcounter{equation}{0}  I have suggested an alternate formula---Equation~\\eqnref{eqn:FitFunc}---for fitting $p$-mode frequency profiles that is based on an eigenfunction expansion of the Green's function. This formulation takes into account the skewness of the line profiles that is inherently produced by a localized acoustic source by including the effects of mode interference in spectral space. In addition, the fitting formula has the salutory features that it is valid for all frequencies across the mode profile (not just in the core of a line) and each mode contains a finite amount of energy.  The suggested scheme can be applied to the fitting of global modes of relatively low harmonic degree ($l \\lesssim 200$) directly as envisioned. The asymmetry in a mode's profile is primarily the result of interference with nearby spatial leaks, which are most likely from the same mode order or ridge (except at very low harmonic degree $l \\lesssim 10$). On the other hand, the proposed fitting model is not immediately relevant to the fitting of ring-analysis spectra. The asymmetry observed in the $p$-mode ridges at large harmonic degree ($l \\gtrsim 200$) are not produced by the same mechanisms that generate the asymmetry for global modes. Instead, the asymmetry is largely a result of the spatial window function (the ring-analysis version of the blending of a mode and the nearby leaks).  Ring-analysis spectra are manufactured by Fourier transforming the wavefield observed on a small patch on the surface of the Sun. The small spatial domain translates into a broad spread of power in horizontal wavenumber. This widens the ridges and produces an effective line width that is much larger than one would expect based just on the mode lifetime. Furthermore, the resulting line profile depends on the window function applied during the Fourier analysis, the slope of the $p$-mode ridge in the dispersion diagram ($d\\omega_n/dk$), and the variation in power along the ridge. Unlike in the global mode analyses where the effects of leakage have been considered in detail for over a decade, leakage has not been treated self-consistently in any ring-analysis studies. While the fitting formula presented here could be used to model the resulting ring-analysis spectra quite well with the asymmetry being produced by the interference between the ridges themselves, the underlying physical mechanism would be wrong. I leave it to a subsequent paper to include the effects of the leakage function on mode profiles.   \\subsection{Measuring the Acoustic Source Function} \\label{subsec:SourceFunc}  One novel property of the fitting function given by Equation~\\eqnref{eqn:FitFunc} is that the effect of the source is directly represented by the complex mode amplitudes. Therefore, the additional parameters have a physical significance that can be used to ferret out information about acoustic sources and the convection that engenders them. This can be seen most easily if I relax the assumption that the source is proportional to a delta function in height. If instead the source is distributed, as usual the wavefield is obtained through a convolution of the Green's function with the source,   \\begin{eqnarray} \\psi(z;\\bvec{k},\\omega) &=& \\sum_{n \\neq 0} \\frac{a_n(\\bvec{k}, \\omega)}{\\omega-\\hat{\\omega}_n(k)} \\; , \\\\ &=& \\sum_{n \\neq 0} \\intall dz' \\; \\frac{\\psi_n(z';k) \\, \\psi_n(z;k)}{2\\hat{\\omega}_n(k)}\\; \\frac{S(z';\\bvec{k},\\omega)}{\\omega-\\omega_n(k)} \\; . \\end{eqnarray}  \\noindent Assuming that the source function is isotropic with a frequency variation that is slow compared to the damping rate, the complex mode amplitude is therefore given by an integral over the source,  \\begin{eqnarray} a_n(k) &=& \\intall dz' \\; {\\cal K}_n(z';k) \\, S(z';k,\\omega_n) \\; , \\label{eqn:Aintegrals} \\\\ {\\cal K}_n(z';k) &\\equiv& \\frac{\\psi_n(z;k) \\, \\psi_n(z';k)}{2\\hat{\\omega}_n(k)} \\; , \\end{eqnarray}  \\noindent where the set of ${\\cal K}_n$ are sensitivity kernels. Given a set of measured amplitudes, the coupled integral equations~\\eqnref{eqn:Aintegrals} could be inverted using standard helioseismic inversion procedures to obtain the source as a function of height in the atmosphere.  \\subsection{Effects of Correlated Noise} \\label{subsec:CorrNoise}  The theory and procedures that have been developed in this paper have so far ignored the existence of correlated noise. Previous studies have shown that correlated noise is a necessary ingredient to explain the reversal of asymmetry between Doppler-velocity spectra and continuum-intensity spectra \\citep{Roxburgh:1997, Nigam:1998a, Nigam:1998b}, as well as the rapid phase variation in velocity--intensity phase difference spectra \\citep{Skartlien:2000}. The modal expansion presented here can easily be adapted to include the effects of correlated noise \\citep[e.g.,][]{Skartlien:2000}; one simply adds a complex correlated noise component, $\\psi_{\\rm corr}$, to the wavefield before computing the power,  \\begin{equation} P_{\\rm model}(\\omega) = B(\\omega) + \\left| \\frac{a_j}{\\omega-\\hat{\\omega}_j + i\\Gamma_j} + \\sum_{j'} c_{jj'} \\, \\frac{a_{j'}}{\\omega-\\hat{\\omega}_{j'} + i\\Gamma_{j'}} + \\psi_{\\rm corr} \\right|^2 \\; , \\label{eqn:PowerCorrNoise} \\end{equation}  \\noindent Assuming that the correlated noise is due to the visibility of the source itself, one would expect that the phase and amplitude of the correlated noise $\\psi_{\\rm corr}$ would be a slowly varying function of frequency. Therefore, if one is fitting the power in a narrow frequency band, it can safely be assumed that the noise is a complex constant. Thus, in any fit to data the correlated noise adds only two additional parameters that must be fit.  If one were fitting a broad range of frequencies, it might be necessary to parameterize the frequency dependence of the correlated noise. Since, the frequency variation should be rather weak, almost any smooth function with a only a few free parameters (polynomial, power law, gaussian, etc.) would probably work."
    },
    "1112/1112.4926_arXiv.txt": {
        "abstract": "Nanoflares, the basic unit of impulsive energy release may produce much of the solar background emission. Extrapolation of the energy frequency distribution of observed microflares, which follows a power law to lower energies can give an estimation of the importance of nanoflares for heating the solar corona. If the power law index is greater than 2, then the nanoflare contribution is dominant.  We model time series of extreme ultraviolet emission radiance, as random flares with a power law exponent of the flare  event distribution. The model is based on three key parameters, the flare rate, the flare duration and the power law exponent of the flare intensity frequency distribution. We use this model to simulate emission line radiance detected in 171 \\AA, observed by STEREO/EUVI and SDO/AIA. The Observed light curves are matched  with simulated light curves using an Artificial Neural Network and parameter values are determined across regions of active region, quiet sun, and coronal hole. The damping rate of nanoflares is compared with radiative losses cooling time. The effect of background emission, data cadence, and network sensitivity on the key parameters of model is studied.  Most of the observed light curves have a power law exponent, $\\alpha$,  greater than the critical value 2. At these sites nanoflare heating could be significant. ",
        "introduction": "The mechanism of forming nanoflares is the dissipation of current sheets arisen from tangential discontinuities in the continuously evolving corona (Levine 1974, Parker 1988). Parker presumes that the change in the magnetic field across the current sheet, $\\Delta B$, is critical for the onset. When the strength $|\\Delta{B}|$ of the discontinuity exceeds some threshold, there is a runaway dynamical instability leading to an explosive reconnection phase. This is similar to the sand pile model that has been used to explain the comparison is also made between results of different methods. Avalanches of magnetic reconnection (Lu \\& Hamilton 1991), pointed out that the magnetic field of the corona is in a state of self organized criticality. To determine whether nanoflares are the main source of heat input to the corona or not, it is necessary to measure the energy frequency distribution of the smallest observed flares.  Hudson (1991) pointed out that the flare occurrence follows a power law distribution such as, $dN\\sim E^{-\\alpha}dE$, in which $dN$ is the number of flares per energy interval $E$ and $E+dE$  (Lu \\& Hamilton (1991) obtained this distribution for complex systems which are in a self organized critical state). The power law index, $\\alpha$, is a critical value for determining whether more weight is given to small scale events (nanoflares) or larger ones (flares). Determination of the power law index of the flare frequency distribution is a scientific challenge. Many authors have attempted to find this index. Benz \\& Krucker (1998), using Yohkoh and SoHO, shown that the power law index at the  microflare frequency distribution is 2.5. Parnell \\& Jupp (2000), used TRACE observations and concluded that the energy of nanoflares is insufficient to heat the quiet sun corona. Aschwanden \\& Parnell (2002) also get some other power law distribution by combining scaling law and fractal geometry for different observations. Aschwanden \\& Charbonneau (2002), gathered the statistics of solar flares, microflares, and nanoflares and concluded that, the power law index  falls bellow the critical value. An extended review on the simulation and observational results is given by ( Klimchuk 2006 and Klimchuk et al. 2009).   {To determine the contribution of small scale events (nanoflares) on the solar coronal heating, some applicable models has been investigated (e. g., Vekstein \\& Katsukawa 2000, Sakamoto et al. 2009, Terzo et al. 2011). Here,} we use a model to simulate the observed Extreme Ultra Violet (EUV) emitted radiation from STEREO and SDO which has been applied successfully to the UV radiance fluctuation in the quiet Sun (Pauluhn \\& Solanki 2007) and later to SUMER observations of the corona in an active region (Bazarghan et al. 2008).  This paper is organized as follow: STEREO/EUVI and SDO/AIA data analysis are described in Section \\ref{data}. The nanoflare model is treated in Section \\ref{model}.  A kind of Artificial Neural Networks (Probabilistic Neural Network) is briefly discussed in Section 4. In sections \\ref{method} and \\ref{result} the method and results are presented.  The conclusions are given in Section \\ref{conc}. ",
        "conclusions": "\\label{conc}  The basic aim of the model, as mentioned earlier, is to determine the coronal heat input by numerous randomly distributed small scale events specially nanoflares by finding the power law index. This will decides  whether the small scale event contribution is of importance or not. The problem of the too power law index indicated by  researchers may have been just a bias due to neglecting the overlapping nanoflares.   Here, a simple nanoflare model based on three key parameters (the flare rate, the flare decay time, and the power law exponent of the flare energy frequency distribution) is used to simulate emission line radiances from the STEREO/EUVI and SDO/AIA in the corona. The simulation code ran to generate more than 22000 light curves (train set) for each combination of $\\alpha$, $\\tau_d$, and $p_f$. For each of the marked regions on the full solar disk images, more than 15000 light curves were generated for average intensities. This large number of perfect light curves which enable us in statistical description was focus of the present-paper.  Light curve pattern recognition by a Probabilistic Neural Network (PNN) was employed to determine values of the key parameters. We found that more than 85\\% of the observed light curves have a power law index greater than 2. Since the network's sensitivity depends on the training set in which the network must see all the possible patterns during the training session, we may have some errors in recognizing the correct patterns.  Empirically, the network is sensitive to steps $\\Delta \\tau_d\\approx1$, $\\Delta p_f\\approx0.1$, and $\\Delta\\alpha\\approx0.1$ for three key parameters. This means that, for shorter steps the light curves are too like, so that the network is not able to classify.   The results can be summarized as follows: \\begin{description} \\item[-] A physical picture of how the model's parameters affect the simulated light curves is discussed.   Decreases in both  average and variance of the light curves are the function of increasing  power law index (greater $\\alpha$ value corresponds to greater number of small events). The higher moments (skewness  and kurtosis) values  of the  time series are accompanied with increasing $\\alpha$, $\\tau_d$, and $p_f$ values (because of the lognormal shape of the distributions). Both the skewness and kurtosis values are  positive numbers. {The distributions of both simulated and observed light curves are asymmetric (Terzo et al. 2011). } \\item[-] The average $\\alpha$ and $ p_f$ did not change with changes in data cadences but the average  $\\tau_d$ is a sensitive function of decrease in cadence. \\item[-]  With regard to average dimensionless range of $\\tau_d$=40-50s and by multiplying it by cadence of $2.5$ min (for STEREO/EUVI, 13 June 2007) and 1.5 min (for SDO/AIA, 22 August 2010) we obtain values of $100-125$ min and  $100-75$ min, respectively. Assuming that plasma cooling through the narrow band filter is dominated by radiative cooling  we find that the ranges are consistent with previous results. \\item[-] The effect of the background emission on the flare rate, $p_f$ is studied. In the coronal hole regions with  less background majority  of events has low flare rates. \\end{description}   The next logical step is to determine the actual flare energies and the total energy input to the corona which is still a problem for researchers to be solved in the future."
    },
    "1112/1112.3447_arXiv.txt": {
        "abstract": "We study gaseous outflows from disk galaxies driven by the combined effects of ram pressure on cold gas clouds and radiation pressure on dust grains. Taking into account the gravity due to disk, bulge and dark matter halo, and assuming continuous star formation in the disk, we show that radiation or ram pressure alone is not sufficient to drive escaping winds from disk galaxies, and that both processes contribute. We show that in the parameter space of star formation rate (SFR) and rotation speed of galaxies, the wind speed in galaxies with rotation speed $v_c\\le 200$ km s$^{-1}$  and SFR $\\le 100$ M$_{\\odot}$ yr$^{-1}$, has a larger contribution from ram pressure, and that in high mass galaxies with large SFR, radiation from the disk has a greater role in driving galactic winds. The ratio of wind speed to circular speed can be approximated as ${v_w \\over v_c} \\sim 10^{0.7} \\, \\left[{\\rm SFR\\over 50 \\, {\\rm M}_{\\odot} \\, {\\rm yr}^{-1}}\\right] ^{0.4} \\ \\left[{v_c\\over 120\\, km/s}\\right]^ {-1.25}$. We show that this conclusion is borne out by observations of galactic winds at low and high redshift and also of circumgalactic gas. We also estimate the mass loading factors under the combined effect of ram and radiation pressure, and show that the ratio of mass loss rate to SFR scales roughly as $v_c^{-1} \\Sigma_g^{-1}$, where $\\Sigma_g$ is the gas column density in the disk. ",
        "introduction": "Galactic winds have been observed at different wavelengths in galaxies of various masses and in a range of redshifts. Galaxies, especially with star formation rates ($\\Sigma_{SFR}$) $\\ge 10^{-1}$ M$_{\\odot}$ yr$^{-1}$ kpc$^{-2}$, often show large outflow of hot gas that emits X-rays and in which cold clouds are found to be embedded, which are observed with H$\\alpha$ or NaD  lines \\citep{hec00,mar05}. The speed of the clouds in the wind range from a few tens to several hundred km s$^{-1}$, and the total mass loss rate can be several times the star formation rate \\citep{vei05}.  These outflows play a crucial role in the evolution of galaxies by expunging gas, and thereby suppressing the star formation. The attempts to understand galactic evolution in the cosmological context have since long encountered the so-called 'cooling catastrophe' problem, since left to its own device the baryonic gas would cool and form stars more rapidly than observed. It is generally believed that a feedback loop inhibits this, and that the process of star formation excites an outflow and quenches itself. The observed mass-metallicity relation in galaxies also indicate that galactic outflows play a major role in the chemical evolution in galaxies. Furthermore, these outflows enrich the intergalactic medium with metals.  The standard model to understand galactic outflows involves a heated interstellar medium (ISM) under the influence of supernovae (SN), and the hot gas  being driven by thermal pressure \\citep{che85, hec02}. The expansion speed of this hot gas can be large enough to eject it out of  the galaxy \\citep{lar74,sai79,dek86}. The observations of cold gas in these outflows \\citep{hec00} led to the proposal that the cold gas entrained in the hot gas moved due to ram pressure. The wind speed was however not found to correlate with galaxy mass \\citep{hec00,mar99}, and it was argued that the supernovae rate increased with SFR and hence the wind velocity might correlate with SFR.  Simulations  also supported this scenario \\citep{suc94,str00}. However, there is a limiting cloud speed implicit in this process since ram pressure acts on the cold gas until the cold gas velocity becomes equal to that of hot gas.  This scenario, however, has met with problems from new observations of cold component which show that the terminal outflow speed  depends on galactic properties like rotation speed \\citep{mar05,rup05}.% It has been proposed that these observations can be explained  by radiation pressure driving the outflow \\cite{mqt05, mar05, sha11}. It has also been pointed out that a natural course of events leading from a starburst would be a radiation pressure driven wind in the beginning, and ram pressure acting on it after a period of $\\sim 3\\hbox{--}5$ Myr, the life time of massive stars \\citep{nath08, mur11}. This scenario also naturally explains the puzzling fact that cold clouds are observed at large distances although their survival time-scales in the hot gas would have inhibited them from being pushed out to such distances.    In the face of two processes leading to outflows, one wonders if both processes contribute equally, or if there are regimes in which one of these two processes dominate over the other. In this paper we present an analytical calculation for the dynamics of cold clouds taking into account both ram and radiation pressure and all sources of gravity, and compare our results with observations. ",
        "conclusions": "The most important result of our calculation is that galactic outflows require both ram and radiation pressure, especially for high mass and high SFR cases. Our calculation has a number of ingredients from stellar physics and disk and halo parameters, and apart from the value of the hot wind speed $v_h$, there is no free parameter in this calculation. It is therefore interesting to note that our theoretical results are consistent with most data of outflows when studied in the parameter space of $v_c$ and SFR. It is also interesting that a recent simulation with ram and radiation pressure driven outflows has concluded that these two processes are important in different mass regimes, although it is not clear where the dividing line between the two regimes lies \\citep{hopkins11,sch11}. Cold cloud outflows from galaxies on the left of the contours in figure \\ref{figcie} are unlikely to escape into the IGM and likely get trapped in the circumgalactic region as observed by \\cite{tum11} (data shown by blue cross) or fall back \\citep{opp08}.  Although strictly speaking our calculation refers to cold clouds being driven out along the pole of the disk galaxies, and we cannot infer the mass loss rate without doing a 2-D calculation, but we can speculate on the scaling of the mass loss rate with galactic mass by making some simple assumption. Let us assume that the dynamics of cold clouds beyond the polar regions are similar to that along the pole. Assuming a one-dimensional mass flow, the mass loss rate from the disk is approximately $\\dot{M}_w\\propto v_w [\\Sigma_g \\pi r_d^2]$, where $\\Sigma_g$ is the gas column density  and $r_d$ is the scale length of the disk. We note that in the prescription of MMW98, one has $v_c \\propto r_d$. We therefore have, $\\dot{M}_w \\propto v_c^{2-0.25} \\dot{M}_\\ast ^{0.4} \\Sigma_g$, where we have used eqn \\ref{eq:fit}, after multiplying both sides by $v_c$. The ratio of mass outflow rate to the SFR is therefore $\\dot{M}_w/\\dot{M}_\\ast \\propto v_c^{1.75} \\Sigma_g \\dot{M}_\\ast^{-1.4}$. Using  Kennicutt's law of star formation, which gives $\\dot{M}_\\ast \\propto \\Sigma_g^{1.4} r_d^2 \\propto \\Sigma_g ^{1.4} v_c^2$, we have finally, $ {\\dot{M}_w \\over \\dot{M}_\\ast} \\propto v_c^{1.05} \\Sigma_g^{-0.96} $. We can therefore conclude that roughly, \\begin{equation} {\\dot{M}_w \\over \\dot{M}_\\ast} \\propto v_c^{-1} \\Sigma_g ^{-1} \\,. \\end{equation} Interestingly, similar power law dependence has also been found in simulations \\citep{hopkins11}.  We note that our results assumed a value of $v_h\\sim 800$ km s$^{-1}$, and a column density of cold clouds of $\\sim 10^{21}$ cm$^{-2}$. If we assume a larger value of $v_h$ ($\\sim 1000$ km s$^{-1}$), then the contour for only ram pressure will be able to explain the winds in ULIGs with large SFR and high mass. A similar result will follow from larger values of $\\kappa$ for the radiation pressure case. %  It is interesting to note that the contour for only radiation pressure can explain the ULIG region of Figure \\ref{figcie} (top right corner). Extending to larger SFR, our results indicate that radiation pressure will also be important for HLIGs (Hyperluminous Infra-red galaxies) \\citep{row00}. Lastly, although it may appear that the role of radiation pressure in galaxies other than ULIGs is less dominant than ram pressure as far as energetics is concerned, radiation pressure may still play an important role in lifting the clouds to a large height before it is embedded in the hot wind to help it survive long \\citep{nath08, mur11}."
    },
    "1112/1112.1168_arXiv.txt": {
        "abstract": " ",
        "introduction": "The Australian SKA Pathfinder (ASKAP) \\citep{Johnston08, Deboer09} is a new radio telescope being built on the Australian candidate Square Kilometre Array (SKA) site in Western Australia. ASKAP will consist of 36 12-metre antennas spread over a region 6 km in diameter. Although the array of antennas is no larger than many existing radio telescopes, each antenna will be equipped with a phased-array feed of 96 dual-polarisation pixels, giving it a 30 deg$^2$ field of view and a very fast survey speed. The Evolutionary Map of the Universe (EMU) project \\citep{Norris11} is a wide-field radio continuum survey planned for ASKAP. The primary goal of EMU is to make a deep (rms $\\sim$ 10$\\mu$Jy/bm) radio continuum survey of the entire Southern Sky at 1.3 GHz, extending as far North as $+30^{\\circ}$ declination, with a 10 arcsec resolution. EMU is expected to detect and catalogue about 70 million galaxies, including typical star-forming galaxies up to z=1, powerful starbursts to even greater redshifts, and AGNs to the edge of the visible Universe. The amount of data involved with ASKAP ($\\sim$2.5 GB/s, or 100 PB/year) requires that the source detection and measurement is fast, robust and highly automated.  Source detection and measurement is a problem common to all astronomical imaging surveys and projects, and numerous software tools have been developed in order perform this initial step in the analysis of imaging data. With the advent of large area surveys the automation of this process is clearly crucial. This has led to a variety of survey-specific source-finders being developed, each optimised to address the specific issues associated with the imaging technology of each survey, and its corresponding image properties and artifacts. What has not developed in parallel is an analysis of the common steps in source-identification and measurement in order to assess the optimal approaches or algorithms that should be used to maximise the robustness and scientific utility of the resulting source catalogues. This is partially a consequence of images obtained using different telescope or imaging technologies (UV/optical/NIR/FIR, compared to radio, or X-ray, or gamma-ray) having very different characteristics. Consequently, the assumptions being used in source-finders developed for images at one wavelength or technology are not typically applicable for others.  The analysis presented here looks at the first steps taken in the source-finding process, those of background estimation and thresholding. This is done explicitly in the context of radio interferometric imaging, although the expectation is that the conclusions should be more broadly applicable. Furthermore, our investigation has a focus on identifying an optimal approach for the source-identification and measurement to be implemented in the ASKAP image analysis software pipeline. This analysis complements that of Hancock et al.\\ 2011 (in prep.), which presents a detailed exploration of existing source-finding tools, their underlying algorithms, and how they perform on simulated radio interferometer images. Together, these analyses comprise the first stage of a thorough investigation of each stage of the source-identification and measurement process that is being pursued as part of the Design Study for EMU.  Several radio source identification and measurement tools are in common use. These include the {\\sc Miriad}/AIPS Gaussian fitting routines IMSAD, SAD and VSAD, {\\sc sfind} \\citep{hopkins2002}, and Duchamp (Whiting 2008) as well as SExtractor \\citep{bertin1996}. There are also a variety of survey-specific tools, such as HAPPY (a modified version of SAD used in the FIRST survey, \\citealp{white1997}), a machine-learning back-end to VSAD used to construct the SUMSS catalogue \\citep{mauch2003}, BDSM (used for the LOFAR source-finding, N. Mohan, in prep), and the floodfill algorithm being used in the Australia Telescope Large Area Survey (ATLAS) Data Release 2 (\\citealp{murphy2007}; Hales et al.\\ 2011 in prep.).  As standalone tools, none of these is adequate for EMU, due to limitations evident in the treatment of background estimation, approaches to treating a varying noise level across an image, and importantly in the numbers of artifacts incorrectly identified as sources. Here we aim to test the approaches to both background estimation and thresholding, in order to identify an optimal approach for the EMU survey. All existing tools have implemented the complete sequence of steps, from background and noise level estimation, through threshold setting, and ``source-pixel\" identification, to source-measurement. It is thus challenging to extract robust information about each independent step in the source-identification process. We do this here for a selection of tools through a judicious choice of parameters in the tasks we investigate, and interpret our results cautiously as a consequence. We emphasise that we are focussing here on two-dimensional data (radio continuum images), with our results expected to be applicable generically to two-dimensional image source-identification.  We present details of the simulated data used in our analysis in \\S\\,\\ref{simdata}, and the algorithms being tested in \\S\\,\\ref{algorithms}. The background estimation and local rms noise estimates are discussed in \\S\\,\\ref{bkg}, with the reliability and completeness statistics being used as a metric to compare the different approaches in \\S\\,\\ref{relcom}. Our results are summarised in \\S\\,\\ref{conc}. ",
        "conclusions": "\\label{conc}  We have tested SExtractor, Selavy and {\\sc sfind} to explore the effects of background and noise estimation, along with two approaches to thresholding, a simple $n\\sigma$ level compared to the false-discovery rate method, on source extraction. The tests were performed on two sets of simulations, the ASKAP simulations which are based on SKADS input source catalogue and include instrumental artefacts, and the Hancock et al.\\ simulation which has only Gaussian noise. The Hancock et al.\\ simulation is an idealised case that is useful for testing the algorithms in ``perfect\" conditions.  The first step in source extraction is background subtraction and noise estimation. We have confirmed the results from the Planck team (Planck Collaboration 2011) and previous deep continuum radio surveys \\citep{huynh2005,schinnerer2007,schinnerer2010} that mesh-sizes of 10 to 20 PSFs or beamwidths produce satisfactory background and noise images. We find that SExtractor background mesh-sizes of 10 and 20 beamwidths produce similar results in terms of reliability. A visual inspection shows the background and noise images are still affected by local bright sources for a mesh-size of 10 beamwidths, but combining a mesh-size of 10 beamwidths with smoothing of 3 meshes produces the highest quality background and noise images. The reliability of the catalogues resulting from 10 and 20 beamwidth mesh-sizes, however, does not differ significantly.  The fact that the background estimation step, so crucial in all the subsequent stages of source-identification and measurement, still needs to be manually tuned in most existing source-detection software, is a major concern. The optimum mesh-size for background subtraction and noise estimation is likely to be image specific, and to vary perhaps substantially depending on the distribution of sources within the image. Developing an automated process for setting the mesh-size when implementing the background and noise properties is clearly a priority, and will need to be developed as part of an automated pipeline for radio telescopes of the future such as the Australian Square Kilometre Array Pathfinder (ASKAP) and the Square Kilometre Array (SKA). One possible method could be to use a tree-based approach, identifying rms noise levels for the whole image and for progressively smaller regions, so that each pixel can be associated with a ``tree\" of rms noise values on each scale. Identifying the scale for which the rms noise plateaus for each pixel could be a suitable approach, and will be tested as part of the EMU Design Study.  The thresholding comparison was limited for the ASKAP simulation due to a ceiling in the completeness values, resulting from bright input sources with small separations that were not deblended well by any of the algorithms, and by the low peak flux density values for extended sources. Nevertheless, in this simulation we find that {\\sc sfind} with $\\alpha$ values of 1 to 5 results in similar completeness to SExtractor run with thresholds of 4 to 5$\\sigma$. The reliability of the {\\sc sfind} sources is higher however, $\\sim$92\\% for {\\sc sfind} compared to 80 -- 88\\% for the SExtractor. Selavy results in higher completeness than SExtractor or {\\sc sfind} for the ASKAP simulations, but at the cost of lower reliability. In regions with significant artefacts such as sidelobes from bright sources {\\sc sfind} and Selavy perform much better than SExtractor in rejecting spurious detections.  In the case of the Hancock et al.\\ simulation, where noise is Gaussian and the sources more well-separated, SExtractor is 98\\% reliable with thresholds of $4\\,\\sigma$ or greater. For the Hancock et al.\\ simulation we find that $n\\,\\sigma$ threshold approach of SExtractor and the FDR approach of {\\sc sfind} perform similarly, although the FDR thresholding of {\\sc sfind} seems to give somewhat better reliability at thresholds that produce comparable levels of completeness. In this idealised simulation SExtractor and Selavy gives similar completeness but the Selavy sources are 15\\% to 30\\% less reliable.  There is a trade-off between completeness and reliability in source extraction algorithms: parameters which give high completeness result in lower reliability. Overall, the false-discovery rate method, as tested with {\\sc sfind}, results in more reliable sources than SExtractor or Selavy, for parameters that give similar completeness levels. While more fine-tuning of Selavy, the prototype source finder for EMU, is required, our analysis suggests that the FDR approach is worthwhile pursuing and we recommend this be implemented in Selavy.  This analysis demonstrates that existing approaches to background and noise estimation seem to be limited not by the specific algorithms but rather the requirement to select appropriate mesh-sizes over which to calculate a ``local\" background and noise estimate. Future work will require development of an automated background mesh-size estimation process for the ASKAP software pipeline, in particular for the EMU images. In addition, a complementary analysis of the source-fitting and parameter measurement approaches is also underway (Hancock et al.\\ in prep.) which will establish the optimum approaches to these latter stages of source extraction."
    },
    "1112/1112.0028_arXiv.txt": {
        "abstract": "Recently published {\\it Spitzer Space Telescope} observations of the classical Cepheid archetype $\\delta$~Cephei revealed an extended dusty nebula surrounding this star and its hot companion HD~213307. At far infrared wavelengths, the emission resembles a bow shock aligned with the direction of space motion of the star, indicating that $\\delta$~Cephei is undergoing mass-loss through a stellar wind.  Here we report \\HI\\ 21-cm line observations with the Very Large Array (VLA) to search for neutral atomic hydrogen associated with this wind. Our VLA data reveal a spatially extended \\HI\\ nebula ($\\sim13'$ or 1~pc across) surrounding the position of \\DC. The nebula has a head-tail morphology, consistent with circumstellar ejecta shaped by the interaction between a stellar wind and the interstellar medium (ISM). We directly measure a mass of circumstellar atomic hydrogen $M_{\\rm HI}\\approx0.07~M_{\\odot}$, although the total \\HI\\ mass may be larger, depending on the fraction of circumstellar material that is hidden by Galactic contamination within our band or that is present on angular scales too large to be detected by the VLA. It appears that the bulk of the circumstellar gas has originated directly from the star, although it may be augmented by material swept from the surrounding ISM. The \\HI\\ data are consistent with a stellar wind with an outflow velocity $V_{\\rm o}=35.6\\pm1.2$~\\kms\\ and a mass-loss  rate of ${\\dot M}\\approx(1.0\\pm0.8)\\times 10^{-6}~M_{\\odot}$ yr$^{-1}$. We have computed theoretical evolutionary tracks that include mass loss across the instability strip and show that a mass-loss rate of this magnitude, sustained over the preceding Cepheid lifetime of \\DC, could be sufficient to resolve a significant fraction of the discrepancy between the pulsation and evolutionary masses for this star. ",
        "introduction": "} Owing to the tight coupling between the pulsation properties of Cepheid variables  and their fundamental stellar parameters,  these stars have long provided crucial tests for stellar evolution models (Gautschy \\& Saio 1996) and play a key role in the calibration of the extragalactic distance scale (Feast \\& Walker 1987; Freedman et al. 2001). Nonetheless, for more than 40 years, a puzzle has persisted concerning the lingering discrepancy between stellar masses derived for Cepheids using different methods (e.g., Cox 1980). In spite of improved radiative opacity calculations (Iglesias et al. 1990; Seaton et al. 1994), advances in evolutionary modeling (Bono et al. 2002), and the consideration of metallicity effects (Keller \\& Wood 2006), masses derived from stellar evolutionary models are found to be systematically $\\sim$10-15\\% higher than those derived from stellar pulsation models (Caputo et al. 2005) or orbital dynamics, when available (Evans 2009).  As described by Neilson et al. (2011), one of the most promising solutions to the so-called ``Cepheid mass discrepancy'' is likely to be convective core overshooting (Chiosi et al. 1992) coupled with {\\em mass-loss} during the Cepheid phase of evolution (see also Bono et al. 2006). Mass-loss from Cepheids has long been postulated, not only as a means of resolving the Cepheid mass discrepancy, but as a natural consequence of stellar evolution and pulsation models (e.g., Iben 1974; Willson \\& Bowen 1984; Neilson \\& Lester 2008). The observational confirmation of mass loss from Cepheids would impact not only our understanding of Cepheid evolution, but would also have implications for the use of Cepheids as distance indicators. For example, mass-loss can non-negligibly affect the structure and scatter in the infrared period-luminosity relation (Neilson et al. 2009), and the presence of circumstellar material can impact Cepheid distance determinations made using the interferometric Baade-Wesselink method (M\\'erand et al. 2007).  A number of previous authors have attempted to identify direct evidence of past or ongoing mass-loss from Cepheids using observations ranging from ultraviolet to radio wavelengths (e.g., Deasy \\& Butler 1986; Deasy 1988; Welch \\& Duric 1988; Bohm-Vitense \\& Love 1994; M\\'erand et al. 2007; Kervella et al. 2006, 2009; Neilson et al. 2009). However, the results have been inconclusive, as the mass-loss rates (or upper limits) derived from these studies (${\\dot M}\\approx 10^{-12}$ to $10^{-5}~M_{\\odot}$~yr$^{-1}$) span many orders of magnitude and are typically rather uncertain owing to their reliance on trace atomic and molecular species (which may not be optimal tracers of past and ongoing mass-loss), their inability to sample very extended spatial scales, and finally, the strong dependencies of derived mass-loss rates on  a variety of underlying assumptions (e.g., wind structure, gas-to-dust ratios, ionization fractions).  With the goal of obtaining new empirical constraints on Cepheid mass-loss, Marengo et al. (2010a) recently used the {\\it Spitzer Space Telescope} to survey a sample of 29 nearby Cepheids. While IR excesses were not directly detected, ruling out the presence of a large amount of warm ($\\sim$500~K) dust in close proximity to the stars, the {\\it Spitzer} images revealed extended emission around several targets (Marengo et al. 2010a; Barmby et al. 2011). Perhaps the most intriguing was the discovery of a prominent nebula around the archetype Cepheid variable, \\DC\\ (Marengo et al. 2010b).  Some basic stellar properties of \\DC\\ are summarized in Table~1. \\DC\\ is part of a wide binary (Benedict et al. 2002), with a hot companion HD~213307 ($T_{\\rm eff}=$8800~K; Cenarro et al. 2007) separated by a projected distance of $40''$. The IR emission surrounding \\DC\\ (and HD~213307) is visible in multiple {\\it Spitzer} bands and shows a roughly parabolic structure with an extent of $\\sim5'$ ($\\approx2\\times10^{4}$~AU). The symmetry axis of the parabola is aligned with the direction of the star's motion through the interstellar medium (ISM; Marengo et al. 2010b), where $V_{\\rm space}\\approx 10.3$~\\kms\\ and P.A.=\\ad{58}{3}.\\footnote{For the present work, we have recomputed the space motion vector for \\DC\\ using the updated solar constants from Sch\\\"onrich et al. 2010.} The IR emission therefore appears to trace a bow shock structure, as arises when a moving, mass-losing star interacts with the local ISM (e.g., Wilkin 1996). Bow shocks have been seen previously in FIR images of a number of mass-losing stars, including the supergiants $\\alpha$~Cam (Van Buren \\& McCray 1988) and $\\alpha$~Ori (Noriega-Crespo et al. 1997; Ueta et al. 2008) and the asymptotic giant branch (AGB) stars R~Hya (Ueta et al. 2006; Wareing et al. 2006), Mira (Ueta 2008), R~Cas (Ueta et al. 2010), IRC+10216 (Ladjal et al. 2010), TX~Psc, and X~Her (Jorissen et al. 2011). The emission mechanism responsible for the FIR emission from these circumstellar bow shocks is uncertain, but is likely to be mainly thermal emission from dust, with possible contributions from low-excitation atomic emission lines (Ueta et al. 2008; Marengo et al. 2010b).  While the companion of \\DC, HD~213307, cannot strictly be excluded as the source of the IR emitting material seen by {\\it Spitzer}, this appears quite unlikely based on the B7-B8~III-V spectral type of this star (see Marengo et al. 2010b for discussion). Thus the {\\it Spitzer} observations of Marengo et al. (2010b) provide strong evidence that \\DC\\ is undergoing mass loss through a stellar wind. To further characterize the nature of this wind and its interaction with the ISM, we have now used the Very Large Array (VLA) to search for \\HI\\ 21-cm line emission in the circumstellar environment of \\DC.  \\HI\\ observations have been  used previously to trace extended circumstellar emission surrounding a number of mass-losing AGB stars (e.g., G\\'erard \\& Le~Bertre 2006; Matthews \\& Reid 2007; Libert et al. 2008; Matthews et al. 2008, 2011). Frequently, the \\HI\\ emission associated with AGB stars is highly extended (up to $\\sim$1~pc) and shows signatures of interaction with the ISM in the form of trailing \\HI\\ wakes, velocity gradients caused by ram pressure effects, and/or density enhancements that demarcate the interstellar-circumstellar interaction zone. Furthermore, several of the evolved stars detected in \\HI\\ are among those with FIR (and/or FUV) emitting bow shocks, including Mira, R~Cas, and IRC+10216, underscoring the complementarity of these tracers for probing the chemistry and kinematics of stellar outflows and their interaction with their environments. This paper represents the first extension of this approach to the study of Cepheids.  \\begin{deluxetable}{lcc} \\tabletypesize{\\scriptsize} \\tablewidth{0pc} \\tablenum{1} \\tablecaption{Coordinates and Stellar Properties of $\\delta$~Cephei} \\tablehead{\\colhead{Parameter} & \\colhead{Value} & \\colhead{Ref.}} \\startdata  $\\alpha$ (J2000.0) & 22 29 10.2 & 1\\\\  $\\delta$ (J2000.0) & +58 24 54.7 & 1 \\\\  $l$ & \\ad{105}{19} & 1 \\\\  $b$ & \\ad{+0}{53} & 1\\\\  Distance (pc) & 273$\\pm$0.011 & 2\\\\  Spectral Type & F5Ib-G1Ib & 3\\\\  Pulsation period (days) & 5.366341 & 3\\\\  Mean $T_{\\rm eff}$ (K) & 5910 & 4\\\\  Mass (pulsation)$^{b}$ ($M_{\\odot}$) & 4.5$\\pm$0.3 & 5\\\\  Mass (evolutionary)$^{c}$ ($M_{\\odot}$) & 5.7$\\pm$0.5 & 5\\\\  Mean $M_{V}$ & $-$3.47$\\pm$0.10 & 2 \\\\  Mean Luminosity ($L_{\\odot}$)& $\\sim$2000 & ...\\\\  Mean radius ($R_{\\odot}$) & 44.5$^{a}$ & 6\\\\  $V_{\\rm LSR}^{d}$ & $-4.7$~\\kms\\ & 7 \\\\  \\enddata  \\tablenotetext{a}{Based on the mean, limb darkened angular diameter $\\phi_{\\rm LD}$=1.520~mas and the distance adopted here.}  \\tablenotetext{b}{Based on $V-K$ colors.}  \\tablenotetext{c}{Assuming a canonical model with no overshoot; see \\S~\\ref{resolution}).}  \\tablenotetext{d}{Derived from the heliocentric radial velocity $V_{\\rm h}=-16.8$~\\kms.}  \\tablecomments{Units of right ascension are hours, minutes, and seconds, and units of declination are degrees, arcminutes, and arcseconds. All quantities have been scaled to the distance adopted in this paper.}  \\tablerefs{(1) SIMBAD database; (2) Benedict et al. 2002; (3) Samus et al. 2011;  (4) Andrievsky et al. 2005; (5) Caputo et al. 2005; (6) Armstrong et al. 2001; (7) Wilson 1953 }  \\end{deluxetable} ",
        "conclusions": "We have used the VLA to search for \\HI\\ 21-cm line emission in the circumstellar environment of the archetype of Cepheid variables, \\DC. We have detected an extended ($13'$, or $\\sim$1~pc) nebula at the position of the star.  The nebula exhibits a head-tail morphology, with the head of the structure aligning closely with the infrared-emitting nebula and bow shock previously detected by Marengo et al. (2010b), while the tail appears to be a turbulent structure that trails the motion of the star through the ISM. We measure an \\HI\\ mass for the nebula of $M_{\\rm HI}\\approx0.07~M_{\\odot}$, although its total \\HI\\ mass could be $\\sim$2-3 times larger, depending on the fraction of emission that is hidden by the strong Galactic emission that contaminates a portion of our observing band. Additional material may also be present on large angular scales ($\\gsim15'$) to which the VLA is insensitive. We interpret the bulk of the \\HI\\ nebula surrounding \\DC\\ as arising from a stellar wind with a mass-loss rate of ${\\dot M}\\approx(1.0\\pm0.8)\\times10^{-6}~M_{\\odot}$~yr$^{-1}$ and an outflow velocity $V_{\\rm o}\\approx35$~\\kms. By computing evolutionary models that include mass loss across the instability strip, we show that a mass-loss rate of this magnitude, sustained over the preceding Cepheid lifetime of \\DC, could be sufficient to resolve a significant fraction of the longstanding discrepancy between the masses of the star derived from stellar pulsation versus stellar evolutionary models."
    },
    "1112/1112.2507_arXiv.txt": {
        "abstract": "The newly discovered Be/X-ray binary in the Small Magellanic Cloud, \\src, provides the first example of a robust association with a supernova remnant (SNR). The short age estimated for the SNR qualifies \\src\\ as the youngest known source in its class, $\\tau\\approx 10^4\\ \\textrm{yr}$. As such, it allows to test current models of magneto-rotational evolution of neutron stars in a still unexplored regime. Here we discuss possible evolutionary scenarios for \\src\\ in the attempt to reconcile its long spin period, $P=1062\\ \\textrm{s}$, and short age. Although several options can be considered, like an anomalously long initial period or the presence of a fossil disc, our results indicate that \\src\\ may host a neutron star born with a  large initial magnetic field, typically in excess of $\\sim 10^{14}\\, {\\rm G}$, which then decayed to $\\sim 10^{13}\\, {\\rm G}$. ",
        "introduction": "\\label{intro}  Be/X-ray binaries (or BeXBs for short) form a subclass of the high-mass X-ray binaries (HMXBs) in which the neutron star (NS) companion is a Be star, a spectral class B giant/subgiant with emission lines and large IR flux. These peculiar properties are explained in terms of an equatorial disc, formed by matter lost by the rapidly rotating Be star. X-ray emission is believed to be powered by accretion of material in the equatorial disc onto the NS \\cite[see e.g.][for a recent review]{reig11}.  BeXBs are both transient and persistent X-ray sources. Transient systems are characterized by type I-II outbursts during which their flux increases by a factor 10--$10^4$ over the quiescent level. They typically contain a not-too-slow NS ($P\\la 100\\, {\\rm s}$) on a moderately eccentric orbit, $\\po\\la 100\\, {\\rm d}$, $e\\ga 0.2$. On the other hand, persistent BeXBs exhibit a rather flat lightcurve, lower X-ray luminosity ($L\\approx 10^{34}$--$10^{35}\\, \\lum$), longer spin and orbital periods \\cite[$P\\ga 200\\, {\\rm s}$, $\\po\\ga 200\\, {\\rm d}$; see again][]{reig11}.  There are presently about 30 well-established BeXBs in the Galaxy, plus $\\sim 40$ candidates. In addition, $\\sim 50$ sources (plus $\\sim 20$ candidates) are known in the Small Magellanic Cloud (SMC)\\footnote{Figures from \\cite{reig11} and the BeXB online catalogue at http://xray.sai.msu.ru/\\textasciitilde raguzova/BeXcat; see also \\cite{ragpo05}.}. Very recently \\cite{henbru11} reported the discovery of a new BeXB in the Wing of the SMC. The new source (\\src) has the typical properties of a persistent BeXB: a B0-0.5(III)e+ donor, $L\\sim 6\\times 10^{35}\\, \\lum$, $P\\sim 1062\\, {\\rm s}$, and an orbital period $\\po\\sim 300\\, {\\rm d}$, as estimated from the Corbet diagram \\cite[][]{corb09}. What makes \\src\\ to stand out amongst its kinship is its (likely) association with a supernova remnant (SNR). BeXB-SNR associations have been already reported (again in the SMC) but in all previous cases they appear uncertain \\cite[][]{hugsm94,coe00}. In the case of \\src\\ the association looks robust and allows for the first time to estimate the NS age in a BeXB from that of the parent SNR, $\\tau\\sim 2$--$4\\times 10^4\\, {\\rm yr}$ \\cite[][]{henbru11}. The suggested association of \\src\\ with a SNR has been further strengthened by a reanalysis of the same \\xmm\\ datasets, supplemented with optical and radio observations, by \\cite{hab11}. Their estimate for the SNR age, $1.6\\times 10^4$ yr, is even shorter than, albeit fully compatible with, that of \\cite{henbru11}. \\cite{hab11} were also able to measure the source period evolution, which results in a positive (i.e. spin-down) rate $\\dot P\\sim 3\\times 10^{-6}\\, {\\rm s\\,s}^{-1}\\sim 95\\, {\\rm s\\,yr}^{-1}$.  The long spin periods ($P \\ga 1000\\, {\\rm s}$) of some persistent BeXBs have been for a long time a major issue. According to the standard picture, there are four stages in the spin evolution of a neutron star embedded in a medium: ejector, propeller, accretor and georotator\\footnote{The last one is of no concern here, since it occurs under specific conditions, hardly met in HMXBs.} \\cite[e.g.][]{lip92}. Once the NS entered the accretor stage after a short propeller phase, its spin period quickly settles at an equilibrium value, $P_{eq}$. In the conventional model, based on \\cite{davpri81} results, the star dipole field should be $B\\ga 10^{14}\\, {\\rm G}$ to have $P_{eq}>10^3 \\, {\\rm s}$, unless the accretion rate is orders of magnitude below what needed to account for the observed X-ray luminosity. Since observations support the presence of a NS with standard magnetic field (at least in some BeXBs), the subsonic propeller stage, which can delay the onset of accretion until a much longer period is reached \\cite[][]{iks07}, has been invoked to explain ultra-slow BeXBs \\cite[see][]{reig11}. More recent investigations of wind-fed accretion onto magnetized NSs indicate, however, that the equilibrium period can be as high as $\\sim 1000\\, {\\rm s}$ even for $B\\approx 10^{12}$--$10^{13}\\, {\\rm G}$ and  $\\dot M\\approx 10^{16}\\, {\\rm g\\,s}^{-1}$, as expected in BeXBs \\cite[][]{shak11}.  Whatever the details of the braking torques, the previous argument implicitly relies on the assumption that the present age of the source is long enough for the NS to have entered the propeller stage. Unless the accretion rate is way above typical, a NS in a BeXB with $B\\ga 10^{12}\\, {\\rm G}$ starts its evolution in the ejector (or pulsar) phase. Its duration can be roughly estimated as $\\tau_{ej}\\ga 10^6 (B/10^{12}\\, {\\rm G})^{-1}(\\dot M/10^{15}\\, {\\rm g\\,s}^{-1})^{-1/2}\\, {\\rm yr}$ (see Section~\\ref{spindown}). This is comfortably below (by a factor $\\approx 10$) the lifetime of the Be companion, so there is ample room for the binary to start an accretion-powered X-ray stage. In the case of \\src, however, it would be impossible for the NS to enter the propeller stage (and hence to become an accretor) in a time as short as a few $\\times  10^4\\, {\\rm yr}$, the estimated SNR age, for typical values of $B$ and $\\dot M$. The accretion rate in \\src\\ is $\\dot M=L/\\eta c^2\\sim 6\\times 10^{15} \\, {\\rm g\\,s}^{-1}$ for an efficiency $\\eta=0.1$, so this points to a highly magnetized NS, with an initial magnetic field substantially above $10^{12}\\, {\\rm G}$. ",
        "conclusions": "The quite short age implied by the association of the newly discovered BeXB in the SMC \\src\\ with a SNR \\cite[][]{henbru11, hab11} raises a number of questions about the properties of the neutron star and its evolutionary status. Normally, one expects that accreting X-ray pulsars spin close to their equilibrium period. However, for a such a young system this appears far from granted. \\cite{hab11} reported a large spin-down rate for \\src, which may suggest that $P_{eq}$ has not been reached yet. According to the standard evolutionary scenario \\cite[][]{lip92}, the (maximum) spin-down rate in the accretor stage is $\\dot P\\sim 2\\pi B^2R_{NS}^6/(GMI)$ which implies $B\\approx 3\\times 10^{14}\\, {\\rm G}$ for $\\dot P\\approx 100\\, {\\rm s\\,yr}^{-1}$. On the other hand, if Eq. (\\ref{shaktorque}) is used to estimate the magnetic field, a much lower value is obtained, $B\\approx 10^{13}\\, {\\rm G}$, very close to what is predicted assuming that the source spins at the equilibrium period (see Eq. [\\ref{bshak}]). This supports a picture in which the NS actually rotates close to $P_{eq}$. A further argument in favor of this is the very short duration of the spin-down phase in the accretor stage, $P/\\dot P\\la 100\\, {\\rm yr}$, which makes it very unlikely to catch the source in this state. Our conclusion is that both the young age and the large spin-down rate of \\src\\ argue in favor of an initially highly magnetic NS which experienced field decay.  Details of evolution past the ejector stage are uncertain, so a fine tuning of the model discussed in Sec. \\ref{spindown} is entirely premature. Still, some general considerations can be made. The main one is that the timescale required for reaching $P_{eq}$ after the NS left the ejector stage is very short. A further point to be stressed is that the zero, or negative, period derivative at $P=P_{eq}$ (see Figure \\ref{perplot}) does not account for variations in $\\dot M$ due to orbital motion in the BeXB and irregularities in the wind. This can result in quite large values of $\\dot P$ and rapid changes in the period derivative which are indeed measured in wind-fed X-ray binaries \\cite[see e.g.][]{bild97}. In this respect the spin-down rate in \\src\\ is large but not exceptionally so. For example, GX~301-2 is known to have a larger value of $\\dot P$ \\cite[][]{koh97}. \\cite{hab11} measured the average period derivative over a time-span less than a month, $\\sim$ one tenth of the orbital period. Fluctuations in $\\dot P$ can occur on a timescale of days, so a more complete dataset is definitely needed to assess the real nature of the period variation.  Alternative scenarios to explain the long period and short age of \\src\\ can be envisaged. For instance, \\cite{hab11} suggested that the NS could have been born with an initial period much longer than $\\sim 0.01\\, {\\rm s}$. The value of $P_0$ can be evaluated by requiring that the transition period given by Eq. (\\ref{pejec}) is reached in less than the source age, and it turns out to be $\\sim 1\\, {\\rm s}$ for the $B$-field corresponding to $P_{eq}$ [Eq. (\\ref{bshak})]. If this is the case no field decay is required. Although possible, no compelling observational evidence for such long initial periods in NSs exists. There are hints that some central compact objects (CCOs) in SNRs may have a current period very close to the initial one, $P_0\\approx 0.1\\, {\\rm s}$. These sources, however, are suspected to host weakly magnetized NSs \\cite[$B\\approx 10^{11}\\, {\\rm G}$; e.g.][]{gotthalp10}. Whether the low field and the long period are related, depending on the conditions under which the NS formed, is still unclear. Still, the estimated value of $B$ in \\src\\ is much higher and the initial period required in the present case is more than one order of magnitude longer. Very recently \\cite{knig11} presented evidence that the population of BeBXs is bimodal, with the two sub-populations having distinct typical spin and orbital periods, and eccentricities. They suggest that the NSs in the two groups are produced through different channels, iron-core-collapse and electron capture supernovae, with the long spin period sources associated to the former channel. Whether this may result in a population of long $P_0$ NSs is a still open question.  Another possibility is that the NS in \\src\\ could, at least in the early stages following its formation, have been surrounded by a debris (or fossil) disc, left by the parent supernova explosion. The existence of such a disc could also lead to rapid spin-down and large period. This, in addition to a super-strong field, was suggested to explain the ultra-long period ($\\sim 6.67\\, {\\rm hr}$) in the enigmatic source RCW~103 \\cite[][]{deluca06,li07}. Although this remains a possibility worth of further investigations, preliminary calculations, based on the model by \\citet[][see also \\citealt{esp11}]{li07}, indicate that a large initial field ($B\\ga 10^{14}\\ {\\rm G}$) is still required, unless the disc is quite massive. For $M_{disc}(0)\\sim 10^{-2}\\ M_\\odot$, the NS can enter the propeller stage earlier than $\\sim 10^4\\ {\\rm yr}$ also for $B\\sim 10^{13}\\ {\\rm G}$.  If \\src\\ indeed contains an initially strongly magnetized neutron star, then studies of this system can shed light on the origin of magnetars. In the standard scenario \\cite[][]{dt92} super-strong fields are generated via dynamo processes. This requires rapid rotation of the proto neutron star. Primary components in high-mass binary systems are not expected to form rapidly rotating cores at the end of their lives \\cite[][]{pp06, bp09}. Up to now no strongly magnetized compact objects have been identified in binary systems with certainty. \\cite{chashkina11}  have recently derived estimates of the $B$-field in HMXBs using \\cite{shak11} formula and no evidence of ultra-high fields was found. It has been suggested that supergiant fast X-ray transients \\cite[SFXTs, see e.g.][for a review]{sid11} may host a magnetar \\cite[][]{boz08}. More recently, \\cite{tor11} reported magnetar-like behaviour from the peculiar binary LS I +61 303. However, no definite confirmation has been given yet, and, in this respect, \\src\\ may be a unique example. The existence of a (evolved) magnetar in a high-mass binary system will pose new challenges on the origin of such neutron stars."
    },
    "1112/1112.5111_arXiv.txt": {
        "abstract": "A complete sample of bright {\\it Swift} Gamma--ray Bursts (GRBs) has been recently selected by Salvaterra et al. (2011). The sample has a high level of completeness in redshift ({\\bf 90}$\\%$). We derive here the intrinsic absorbing X--ray column densities of these GRBs making use of the {\\it Swift} X--ray Telescope data. This distribution has a mean value of $\\log(N_H/{\\rm cm^{-2}})=21.7\\pm0.5$. This value is consistent with the distribution of the column densities derived from the total sample of GRBs with redshift. We find a mild increase of the intrinsic column density with redshift. This can be interpreted as due to the contribution of intervening systems along the line of sight. Making use of the spectral index connecting optical and X--ray fluxes at 11 hr ($\\beta_{OX}$), we investigate the relation of the intrinsic column density and the GRB `darkness'. We find that there is a very tight correlation between dark GRBs and high X--ray column densities. This clearly indicates that the dark GRBs are formed in a metal-rich environment where dust must be present. ",
        "introduction": "Long duration Gamma--ray bursts (GRBs) are thought to originate from the collapse of massive stars. Several lines of evidence points toward this conclusion, ranging from the association to type Ib/c supernovae (Woosley \\& Bloom 2006 and references therein), to the occurrence of GRBs in the most luminous part of their host galaxies (Svensson et al. 2010). The ambient medium in which GRBs explode is expected to be denser than the interstellar medium and typical of star forming regions. The values of the absorbing column densities as measured in X--rays are high (Galama \\& Wijers 2001; Stratta et al. 2004; Campana et al. 2006; Watson et al. 2007). An analysis of the intrinsic column densities of all {\\it Swift} GRBs observed within 1,000 s and with a known redshift has been carried out by Campana et al. (2010). The selected sample consisting of 93 GRBs was biased. The distribution of the intrinsic X--ray absorption column density is consistent with a lognormal distribution with mean $\\log N_H(z)=21.9\\pm0.1$ ($90\\%$ confidence level). This distribution is in agreement with the expected distribution for GRBs occurring randomly in giant molecular clouds similar to those within the Milky Way (Campana et al. 2006; Reichart \\& Price 2002). Looking at the distribution of X--ray column densities vs. redshift, there is a lack of non-absorbed GRBs at high redshift and a lack of heavily absorbed GRBs at low redshift. This might be the outcome of biases present in the sample. Looking at the distribution of X--ray column densities versus redshift  a lack of non-absorbed GRBs at high redshift and a lack of heavily absorbed GRBs at low redshift were found in previous studies (Campana et al. 2010). The former might be explained in terms of more compact and dense star formation regions in the young Universe (or to a sizable contribution from intervening systems). The latter might be interpreted as due to a change in the dust properties with redshift, with GRBs at redshift $z\\lsim 2-3$ having a higher dust to gas ratio for the same X--ray column density (e.g. different grain size or composition). This will naturally provide a lack of heavily (X--ray) absorbed GRBs at small redshifts.  In the optical the presence of a large amount of absorbing material  is much less clear, since a large number of GRB afterglows are not affected by absorption (Kann, Klose \\& Zeh 2006; Schady et al. 2010; Zafar et al. 2011; but see Greiner et al. 2011 and Covino et al. 2011, in preparation). In this respect photoionisation of dust grains can play an important role (Lazzati, Perna \\& Ghisellini 2001; Lazzati, Covino \\& Ghisellini; Draine \\& Hao 2002; Campana et al. 2007). Moreover, the absorbing column densities measured in the optical based on damped Lyman-$\\alpha$ absorption are a factor of $\\sim 10$ lower than those measured in the X--ray band (Campana et al. 2010; Fynbo et al. 2009). This has been interpreted as due to photoionization of the surrounding medium by GRB photons (Campana et al. 2006, 2007; Watson et al. 2007; Campana et al. 2010; Schady et al. 2011).  The presence of a large amount of material is also testified by the existence of `dark' GRBs. There are several definitions of dark GRBs. The easiest is that they do not show an optical counterpart (Fynbo et al. 2001). Since this definition is clearly related to the sensitivity (and availability) of the instruments used for the follow-up a more general definition is needed. Based on the predictions of the fireball model (M\\'esz\\'aros \\& Rees 1997) one can require that the optical to X--ray spectral index $\\beta_{OX}$ (i.e. the slope between the fluxes in the $R$-band and at 3 keV at 11 hr after the burst) should be lower than 0.5 (Jakobsson et al. 2004). This will individuate optically sub-luminous bursts, i.e. fainter than expected from the fireball model. Alternatively, with the advent of {\\it Swift}, X--ray spectral slopes were more easily available and a somewhat different definition was put forward by van der Horst et al. (2009) for which $\\beta_{OX}$ is shallower than $\\beta_X - 0.5$.  The darkness of a GRB can have different causes: it can be due to intrinsically optically faint GRBs, it can be due to absorption by intervening material within the host galaxy or it can be due to high redshift GRBs, thus being absorbed by the intergalactic medium. Several works have addressed this topic in the {\\it Swift} era when a number of facilities allowed a quick follow-up of the afterglows. The fraction of dark bursts has been estimated to be $\\sim 25-50\\%$ according to Jakobsson's definition (Melandri et al. 2008; Roming et al. 2009; Cenko et al. 2009; Greiner et al. 2011; Melandri et al. 2011). It is now believed that the faint optical afterglow emission of dark bursts might be due to a moderate intrinsic extinction at moderate redshifts. Greiner et al. (2011) estimated a $\\sim 20\\%$ contribution from high redshift ($z\\gsim 4-5$) GRBs to the dark population only.  Salvaterra et al. (2011, see also Nava et al. 2011) selected a complete sample of bright GRBs based on optical observability (Jakobsson et al. 2006) and {\\it Swift} BAT peak flux $P\\ge 2.6$ ph s$^{-1}$ cm$^{-2}$. The sample consists of 58 GRBs and it is complete in spectroscopic redshift at $90\\%$ (with $95\\%$ of GRBs having some constraints on the redshift). This sample offers the opportunity to study in an unbiased way the distribution of the X--ray column densities and its relation to GRB darkness.  The paper is organised as follows. In section 2 we derive the X--ray absorbing column densities for the Salvaterra's sample and briefly describe how the slope $\\beta_{OX}$ has been computed for each GRB of the sample. In section 3 we discuss our findings and in section 4 we draw our conclusions. ",
        "conclusions": "Salvaterra et al. (2011) selected a complete sample of bright GRBs with a high degree of completeness in redshift. In a series of papers we investigate the impact of this sample on GRB studies. Here we focus on the properties of the sample with respect to the intrinsic X--ray absorption. We found that the intrinsic column density distribution of our complete sample is consistent with the total distribution of column densities presented in Campana et al. (2010). The mean of the two distributions are in fact $21.7\\pm0.5$ and $21.9\\pm0.1$, respectively. This likely indicates that the GRB brightness, as well as any other bias present in the total sample of GRBs with redshift (e.g. dust), does not heavily influence the total distribution of intrinsic column densities.  At variance with the total distribution presented in Campana et al. (2010), we see in the complete sample presented here that the region at high column densities and low redshift is now more populated by GRBs. This clearly reveals a bias present in the non-complete sample, where this region is heavily underpopulated due to the lack of  a redshift determination of dark bursts.  Even if not statistically compelling there is an increase of the intrinsic column density with redshift  (this is more apparent in the full sample of GRBs with redshift, Campana et al. 2010). We evaluate the mean contribution to $N_H(z)$ due to the intervening systems along the GRB line of sight. We find that, if we take into account the larger number of observed systems affecting the line of sight of GRBs with respect to the quasar one (Vergani et al. 2009), the population of sub-Damped Lyman-$\\alpha$ and Damped Lyman-$\\alpha$ systems can account for the increase with redshift of $N_H(z)$. It would be interesting to confirm this directly through the study of high-$z$ GRB lines of sight. Unfortunately this effect plays a significant role at very high redshift, where the number of GRB afterglow spectra is very low. It is indeed difficult to measure absorption from Lyman-$\\alpha$. However, the column density of neutral gas can still be traced by weakly ionised metal lines (e.g. Zn II, Si II), which in fact is a more logical method of comparing absorption in X--rays and the optical, given that the X--rays are absorbed by metals and not neutral hydrogen (e.g. Schady et al. 2011). The X-shooter instrument mounted at the ESO/VLT offers the best opportunities for these studies.  Making use of the $\\beta_{OX}$ computed by Melandri et al. (2011), we found a strong correlation between GRB darkness and X--ray absorbing column densities. Since metals are a key ingredient for dust production (Draine 2003), our findings are consistent with a picture in which the darkness of a GRB is in most cases due to absorption by circumburst material."
    },
    "1112/1112.6218_arXiv.txt": {
        "abstract": "{Solar cycles vary in their amplitude and shape. There are several empirical relations between various parameters that link cycle's shape and amplitude, the foremost the Waldmeier relations.} {The solar cycle is believed to be a result of the solar dynamo action, therefore these relations require an explanation in the framework of this theory, which  we aim to present here.} {We related the cycle-to-cycle variability of solar activity to fluctuations of solar dynamo drivers and primarily to fluctuations in the parameter responsible for the recovery of the poloidal magnetic field from the toroidal one. To be specific, we developed a model in the framework of the mean-field dynamo based on the differential rotation and $\\alpha$-effect.} {We demonstrate that the mean-field dynamo model, which is based on a realistic rotation profile and on nonlinearity that is associated with the magnetic helicity balance, reproduces both qualitatively and quantitatively the Waldmeier relations observed in sunspot data since 1750. The model also reproduces more or less successfully other relations between the parameters under discussion, in particular, the link between odd and even cycles (Gnevyshev-Ohl rule).} {We conclude that the contemporary solar dynamo theory provides a way to explain the cycle-to-cycle variability of solar activity as recorded in sunspots. { We discuss the importance of the model for stellar activity cycles which, as known from the data of the Mount Wilson HK project, which measures the Ca H and K line index for other stars, demonstrate the cycle-to-cycle variability similar to solar cycles.}} ",
        "introduction": "}  Solar activity has a periodic nature, but the cycle amplitude and shape vary from one cycle to the other. { This challenges the prognostic abilities of solar activity models. The sunspot activity can be quantified by using various tracers derived from observations. These tracers show an interrelation among each other. The indices characterizing the tracers can be employed to predict the future evolution of solar activity.} Waldmeier (\\citeyear{w35}) first suggested this option (an inverse correlation between the length of the ascending phase of a cycle and the peak sunspot number of that cycle) and applied it (\\citealp{w36}) to predict the subsequent cycle. Later, other relations of this kind were proposed and called Waldmeier relations (see for a review \\citealp{vetal86} and \\citealp{hathaw02}). The nature of physical processes underlying the Waldmeier relations is not clear \\citep[see discussion, e.g., in][]{camschu07,detal08,karak11}. We note,  that these statistical properties of the magnetic activity also exist for other tracers related to the sunspot activity (e.g., sunspot group number and area, see \\citealp{vetal86,hathaw02,karak11}), and even for the other kind of solar and stellar activity indices, e.g., for the Ca II index \\citep{soon94}. The Waldmeier  relations are considered as a valuable test for dynamo models \\citep{karak11,pk11}.  Clarifying the physics underlying the Waldmeier relations is particularly attractive to support the prognostic abilities concerning solar cycle. It is believed that the cyclic solar activity is driven by a dynamo, i.e. a mechanism that transforms the kinetic energy of hydrodynamic motions into a magnetic one. Many modern solar dynamo models \\citep[see, e.g.,][]{stix:02}  assume that the toroidal magnetic field that emerges on the surface and forms sunspots is generated near the bottom of the convection zone, in the tachocline or just beneath it in a convection overshoot layer. This kind of dynamo can be approximated by the Parker dynamo waves \\citep{park93}. The direction of the dynamo wave propagation in the framework of the $\\alpha\\Omega$-dynamo is defined by the Parker-Yoshimura rule \\citep{par55,yosh1975}, according to which  the wave propagates along iso-surfaces of the angular velocity. The propagation can be affected by the turbulent transport (associated with the mean drift of magnetic activity in the turbulent media by means of the turbulent mechanisms), by the anisotropic turbulent diffusivity \\citep{k02}, and by meridional circulation \\citep{yosh1975,choud95}. An alternative to the Parker's surface dynamo waves is the distributed dynamo with subsurface shear \\citep[e.g.,][]{b05}, where the dynamo wave propagates along the radius in the main part of the solar convection zone \\citep{k02}. Near-surface activity is determined by the subsurface shear. Another popular option is the flux-transport dynamo \\citep[e.g.,][]{choud95,dc99}.  In the context of dynamo theory, the Waldmeier relations can be explained by invoking physical mechanisms of the solar magnetic field generation and a mechanism that drives variations of the amplitude and shape of the activity cycle. For example, \\cite{pk11}, hereafter PK11, showed that variations of the $\\alpha$-effect amplitude may explain the correlation between the cycle rise rate and the cycle amplitude and other types of the Waldmeier relations as well. It was suggested \\citep{choud92,h93} that the fluctuations of the $\\alpha$-effect (associated with kinetic helicity fluctuations) are likely to be one of the natural sources of the cycle parameter variations.  In addition to the statistical relations between the cycle parameters within a separate cycle there are correlations relating the parameters in subsequent cycles, for example, the odd-even cycle and the last cycle period-amplitude effects. These effects are closely related to the memory of the dynamo processes and to the strength of the saturation processes, which damp deviations of the cycle parameters from the reference state characterizing the cycle attractor \\citep{oss-h96a,oss-h96b}.   It was argued \\citep{choud92,h93}, that a dozen percent is a reasonable estimate for the noise component of the $\\alpha$-effect. Previous calculations (see the above cited papers) showed that a straightforward application of the idea with the vortex turnover time and the vortex size as the correlation time and length for the $\\alpha$-fluctuations needs fluctuations much stronger than the mean $\\alpha$. On the other hand, the results of direct  numerical simulations \\citep[e.g,][]{bs02} and results of current helicity (related to $\\alpha$) observations in solar active regions \\citep[e.g.,][]{zetal10} {suggest that the correlation time for $\\alpha$-fluctuations can be comparable to the cycle length and the correlation length comparable to the extent of the latitudinal belts}. Using these results, \\citet{moss-sok08} and \\citet{uetal09}  showed that an $\\alpha$-noise on the order of few dozen percents is sufficient to explain the Grand minima of solar activity. The aim of this paper is to examine the result of $\\alpha$-fluctuations in the statistical properties of the solar cycle including the Waldmeier relations and the odd-even cycle effect.  We chose a particular model for the solar cycle in which $\\alpha$-fluctuations are introduced. Of course, it is impractical to try all available models to learn which one is better to obtain the relations under discussion, but we select below the model among a relative wide choice of the models that gave better results in the preliminary simulations \\citep{ps11}. ",
        "conclusions": "We have studied the impact of low-amplitude Gaussian fluctuations of the $\\alpha$-effect on the statistical properties of the magnetic dynamo cycle, such as the Waldmeier relations and the Gnevyshev-Ohl rules. The dynamo model includes long-term fluctuations of the $\\alpha$-effect and employs two types of a nonlinear feedback of the mean-field on the $\\alpha$-effect, including algebraic quenching and  dynamic quenching due to the magnetic helicity generation. The general properties of the dynamo, such as the direction of the toroidal magnetic field drift, the polar magnetic field sign reversal at the maximum of a cycle, etc., are consistent with observations.  Our model does not include the meridional circulation effect, which is advocated by the Babcock-Leighton and the flux-transport type dynamo models \\citep[e.g.,][]{choud95,dc99,detal08}. It was shown that a part of the Waldmeier relations can be possibly explained by a specially tuned flux-transport model that considers fluctuations of the meridional circulation speed \\citep{karak11}. Still, observational constraints on the distribution of the meridional circulation inside the convection zone are not  very strong because we have measurements for the surface.  The angular momentum balance in mean-field models supports for the circulation pattern, which has a deep circulation stagnation point and a strong concentration of the velocity speed towards the bottom and the top boundaries of the solar convection zone (e.g., \\citealp{kit11}). Yet, most of the flux-transport models (including \\citealp{karak11}) use a very different circulation pattern. Following to this argumentation, we postpone a more complete study of the effects of the meridional circulation fluctuations to the future.  We showed, confirming the previous findings of \\cite{pk11} and \\cite{ps11}, that variations of the $\\alpha$-effect amplitude result in variations of the cycle amplitude and period. Taking into account random fluctuations  of the $\\alpha$-effect, we calculated statistical properties relating the cycle amplitude, the cycle shape, the rise time, etc., on the basis of the simulated SN data set covering period of more than 10000 years. Our results agree well  with observations for the Waldmeier relations and the Gnevyshev-Ohl rules.  From the qualitative point of view these results were anticipated from the earlier analysis of the helicity fluctuation effect in the dynamo given by \\cite{choud92} and \\cite{h93} (see, also \\citealp{oss-h96a,oss-h96b,moss-sok08,uetal09}). Our results presented in Figure~\\ref{minsurf} about the correlation of the polar dipole field and the cycle amplitude and as the results for the Gnevyshev-Ohl rules suggest that the Waldmeier relations can be understood by considering the general properties of the magnetic field generation processes, which are involved in the dynamo.  Our model shows a good correlation (with low variance) between the strength of the polar dipole magnetic field in the cycle minimum and the amplitude of the subsequent cycle. This results from the deterministic process of the toroidal magnetic field generation by the differential rotation from the large-scale poloidal magnetic field. This correlation is often used for the cycle prediction \\citep{hath09} by Babcock-Leighton type dynamo models and it is for the first time demonstrated in the mean-field dynamo. The rise rate of the sunspot cycle depends on the differential rotation and the amplitude of the poloidal field. Therefore, the correlation between the rise rate and amplitude of the cycle is a derivative property and is a consequence  of the link  between the polar dipole magnetic field in the cycle minimum and the strength of the toroidal field in the subsequent cycle.  Furthermore, following the general idea of \\cite{GM78} (cf, \\citealp{h93,chetal07}), we can interpret  the  Gnevyshev-Ohl rule as evidence that the solar cycle is a nonlinear self-excited oscillation that tends to preserve the property of the attractor under random perturbations. The amplitude and phase of the subsequent cycles are related by the so-called Zaslavsky map. The strength of the link between the parameters of the subsequent cycles is controlled by the fluctuation amplitude and by the perturbation's decrement. The latter strongly depends on the nonlinear mechanisms involved in the dynamo. To examine this idea we made additional simulations with a lower helicity dissipation rate (high $R_{\\chi}$) and found that the correlation coefficients between the parameters of the subsequent cycle increase with increasing $R_{\\chi}$. Therefore  the link between the odd and even cycles, and  the period to amplitude correlation in subsequent cycles can be considered  as evidence for the fluctuation impact on the dynamo and evidence for nonlinear damping of these perturbations in the dynamo. This conclusion needs to be investigated in more detail especially by comparing the results of the $\\alpha$-effect and meridional circulation fluctuations.  Long-term variations of the magnetic cycle in the dynamo can be induced in different ways. Two main mechanisms can be identified: nonlinear deterministic chaos and an effect of fluctuations of the turbulent parameters involved in the dynamo process. Generally, we anticipate that statistical properties of long-term cycle variations are depended on the force that drives the long-term variations. We examined weakly nonlinear models with the amplitude of the $\\alpha$-effect close to the threshold. In our models, the typical ratio between the energy of the large-scale toroidal field and the kinetic energy of convective flows does not exceed $0.3$. As a result, the chaotic regime in the model is not as evident as the impact of the $\\alpha$-effect fluctuations.  Figure~\\ref{prob} illustrates the difficulty to obtain extended episodes of low magnetic activity in this case, while these episodes are common in the solar dynamo \\citep{setal04}. To amplify the chaotic regime in the model, we have tried additional possibilities and included an angular momentum balance into the dynamo problem to take into account the nonlinear feedback of the magnetic field on the differential rotation. The model was described earlier by \\cite{p99,p04}. Figure~\\ref{ndifr} shows the simulated SN for the model involving the nonlinear effect of the magnetic field on the angular momentum balance in the solar convection zone. This model shows higher intermittency in the cycle variations than that in Figure~\\ref{waldr}, and indeed it has a similar probability of the occurrence of lowactivity episodes as that in the reconstruction data set. \\begin{figure} \\begin{centering} \\includegraphics[width=0.95\\columnwidth]{fig9} \\par\\end{centering} \\caption{Time series of the simulated sunspot number for the extended 2D1$\\alpha$ model involving the fluctuations of the alpha effect and the magnetic feedback on the differential rotation.}\\label{ndifr} \\end{figure}  It is natural to expect that at least  stellar magnetic cycles of solar-like stars should demonstrate a variability similar to the solar one, including relations comparable with the Waldmeier relations. Available observations of stellar activity provide some hints that support this expectation.  { Stellar activity data of the Mount Wilson HK project measuring the Ca H and K line index for other stars \\citep{bal} are available for two activity cycles.} The wavelet analysis of the data for several stars \\citep{balfrick} demonstrated that the subsequent cycles for a given star can differ in their cycle amplitudes. We note that a monitoring of stellar activity of solar-like stars to obtain relations similar to the Waldmeier ones could establish our prognostic abilities of solar activity based on these relations much better.   Summarizing the results of the paper, we conclude that the mean-field solar dynamo theory provides  a way to explain the cycle-to-cycle variability of solar activity as recorded in sunspots. The results given in the literature and the results obtained in the paper suggest that the Waldmeier relations can be explained invoking very different kinds of dynamo models. More work is necessary to study the relations between the statistical properties of the dynamo cycle, and the dynamo mechanisms involved in the magnetic activity will help to obtain more insight into the processes operating in the stellar and solar dynamo."
    },
    "1112/1112.1438_arXiv.txt": {
        "abstract": "We present the results of a systematic Green Bank Telescope (GBT) and Giant Metrewave Radio Telescope (GMRT) survey for 21-cm absorption in a sample of 10 damped Lyman-$\\alpha$ (DLA) systems at $2\\le$ \\zabs$\\le 3.4$. Analysis of L-band Very Long Baseline Array (VLBA) images of the background QSOs are also presented.  We detect 21-cm absorption in only one DLA (at \\zabs\\ = 3.1745 towards J1337+3152). Thus the detection rate of 21-cm absorption is  $\\sim 10\\%$ when no limit on the integrated optical depth ($\\int \\tau(v) dv$) is imposed and $\\sim13$\\% for a 3$\\sigma$ limit of 0.4~\\kms. Combining our data with the data from the literature (a sample of 28 DLAs) and assuming the measured core fraction at milliarcsecond scale to represent the gas covering factor, we find that the H~{\\sc i} gas in DLAs at $z\\ge 2$ is predominantly constituted by warm neutral medium. The detection rate of 21-cm absorption seems to be higher for systems with higher $N$(H~{\\sc i}) or metallicity. However, no clear correlation is found between the integrated 21-cm optical depth (or the spin-temperature, $T_{\\rm S}$) and either  $N$(H~{\\sc i}), metallicity or velocity spread of the low ionization species. There are 13 DLAs in our sample for which high resolution optical spectra covering the expected wavelength range of \\h2 absorption are available. We report the detection of \\h2 molecules in the \\zabs\\ ~=~3.3871 21-cm absorber towards J0203+1134 (PKS 0201+113). In 8 cases, neither \\h2 (with molecular fraction $f$($H_2$)$ \\le 10^{-6}$) nor 21-cm absorption (with $T_{\\rm S}/f_{\\rm c} \\ge$~700~K) are detected. The lack of 21-cm and H$_2$ absorption in these systems can be explained if most of the H~{\\sc i} in these DLAs originate from low density high temperature gas. {In one case we have a DLA with 21-cm absorption not showing \\h2 absorption.} In two cases, both species are detected but do not originate from the same velocity component. In the remaining 2 cases 21-cm absorption is not detected despite the presence of \\h2 with evidence for the presence of cold gas. All this is consistent with the idea that the \\h2 components seen in DLAs are compact (with sizes of $\\le$~15~pc) and contain only a small fraction (i.e typically $\\le 10\\%$) of the total $N$(H~{\\sc i}) measured in the DLAs. This implies that the molecular fractions $f$(\\h2) reported from the \\h2 surveys should be considered as conservative lower limits for the \\h2 components. ",
        "introduction": "The Galactic interstellar medium (ISM) has a multiphase structure with neutral hydrogen being distributed between the cold neutral (CNM), warm neutral (WNM) and warm ionized (WIM) media. A large fraction of the gas is also found in diffuse, translucent and dense molecular clouds. Newly formed stars are associated with these dense molecular clouds and strongly influence the physical state of  the rest of the gas in different forms through radiative and mechanical inputs. The physical conditions in the multiphase ISM depend on the UV background radiation field, metallicities, dust content and the density of cosmic rays \\citep[see  Figs.~5, 6 and 7 in][]{Wolfire95}. In addition, the filling factor of the different phases depends sensitively on the supernova rate \\citep{deavillez2004}. Therefore, detecting and studying the multiphase ISM in external galaxies has great importance for our understanding of galaxy evolution.  Damped Lyman-$\\alpha$ systems (DLAs) are the highest H~{\\sc i} column density  absorbers seen in QSO spectra, with $N$(H~{\\sc i})$\\ge 2\\times 10^{20}$~cm$^{-2}$. These absorbers trace the bulk of the neutral hydrogen at $2\\le z\\le 3$ \\citep [][]{Prochaska05,Noterdaeme09dla} and have long been identified as revealing the interstellar medium of the high-redshift precursors of present day galaxies \\citep[for a review see,][]{wolfe05}.  The typical dust-to-gas ratio of DLAs, is less than one tenth of that observed in the local ISM, and only a small fraction ($<\u223c10$\\%) of DLAs show detectable amounts of molecular hydrogen \\citep{Petitjean00,Ledoux03,Noterdaeme08} with the detection rate being correlated to the dust content of the gas \\citep{Petitjean06}. The estimated temperature and molecular fraction in these systems are consistent with them originating from the CNM \\citep{Srianand05}. It has been shown recently that strong C~{\\sc i} absorbers detected in low-resolution Sloan Digital Sky Survey (SDSS) spectra are good candidates for H$_2$ bearing systems. Indeed these absorbers have yielded the first detections of CO molecules in high-$z$ DLAs \\citep{Srianand08,Noterdaeme09co,Noterdaeme10co,Noterdaeme11}. The properties of these absorbers are similar to those of translucent molecular clouds. The fact that no DLA is found to be associated with a dense molecular cloud, a fundamental ingredient of star-formation, is most certainly related to the large extinction that these clouds are expected to produce and/or the small size of such regions \\citep{Zwaan06} making detections difficult.  Thus, most DLAs detected in optical spectroscopic surveys seem to probe the diffuse H~{\\sc i} gas \\citep{Petitjean00}. However, about 50\\% of the DLAs show detectable C~{\\sc ii}$^*$ absorption \\citep{Wolfe08}, and \\citet{Wolfe03b} argued that a considerable fraction of the C~{\\sc ii}$^*$ absorption in DLAs originates from CNM gas \\citep[see however][]{Srianand05}. { Detection of 21-cm absorption is the best way to estimate the CNM fraction of DLAs as it is sensitive to both $N$(H~{\\sc i}) and thermal state of the gas \\citep{Kulkarni88}.} \\begin{figure*} \\centerline{ \\vbox{ \\hbox{ \\includegraphics[trim= 25.0mm 50.0mm 90.0 40.0mm, clip, scale=0.33,angle=90.0]{J0733+2721c.ps} \\includegraphics[trim= 25.0mm 50.0mm 90.0 50.0mm, clip, scale=0.33,angle=90.0]{J0801+4725c.ps} \\includegraphics[trim= 25.0mm 50.0mm 90.0 50.0mm, clip, scale=0.33,angle=90.0]{J0816+4823c.ps} } \\hbox{ \\includegraphics[trim= 25.0mm 50.0mm 70.0 40.0mm, clip, scale=0.33,angle=90.0]{J0839+2002c.ps} \\includegraphics[trim= 25.0mm 50.0mm 70.0 50.0mm, clip, scale=0.33,angle=90.0]{J0852+2431c.ps} \\includegraphics[trim= 25.0mm 50.0mm 70.0 50.0mm, clip, scale=0.33,angle=90.0]{J1017+6116c.ps} } \\hbox{ \\includegraphics[trim= 25.0mm 50.0mm 70.0 40.0mm, clip, scale=0.33,angle=90.0]{J1223+5037c.ps} \\includegraphics[trim= 25.0mm 50.0mm 70.0 50.0mm, clip, scale=0.33,angle=90.0]{J1237+4708c.ps} \\includegraphics[trim= 25.0mm 50.0mm 70.0 50.0mm, clip, scale=0.33,angle=90.0]{J1242+3720c.ps} } \\hbox{ \\includegraphics[trim= 25.0mm 50.0mm 70.0 40.0mm, clip, scale=0.33,angle=90.0]{J1406+3433c.ps} \\includegraphics[trim= 25.0mm 50.0mm 90.0 50.0mm, clip, scale=0.33,angle=90.0]{J1413+4505c.ps} \\includegraphics[trim= 25.0mm 50.0mm 90.0 50.0mm, clip, scale=0.33,angle=90.0]{J1435+5434c.ps} } }} \\caption[]{ SDSS spectra showing the \\lya lines for 12 DLAs in our sample. The best fitted Voigt profiles (solid curves) together with the associated 1$\\sigma$ errors (shaded regions) are over-plotted. The dotted curve gives the best fitted continuum (in some cases the continuum fit includes the emission lines also). We have used VLT UVES spectra to get $N$(H~{\\sc i}) in the case of \\zabs = 3.1745 system towards J1337+3152 \\citep[see][]{Srianand10} and \\zabs = 2.595 and  2.622 systems towards J0407$-$4410 \\citep[CTS 247, see][]{Ledoux03}. } \\label{dlafits} \\end{figure*} This is why it is important to search for 21-cm absorption in DLAs over a wide redshift range. While a good fraction of DLAs/sub-DLAs preselected through Mg~{\\sc ii} absorption seems to show 21-cm absorption at $z\\sim 1.3$ \\citep[see for example][]{Gupta09, Kanekar09mg2}, searches for 21-cm absorption in DLAs at $z_{\\rm abs}\\ge2$ have mostly resulted in null detections \\citep[see][]{Kanekar03,Curran10} with only four detections reported till now \\citep[see][]{Wolfe85,Kanekar06,Kanekar07,York07}. { The low detection rate of 21-cm absorption in high-$z$ DLAs can be related to either the gas being warm (i.e high spin temperature, $T_{\\rm S}$, as suggested by \\citet{Kanekar03}) and/or the low value of covering factor ($f_c$) through high-z geometric effects \\citep{Curran06}.  The best way to address the covering factor issue is to perform milliarcsecond scale spectroscopy in the redshifted 21-cm line using very long baseline interferometry (VLBI) to measure the extent of absorbing gas \\citep{Lane00}. Unfortunately due to limited frequency coverage and sensitivity of the receivers available with VLBI such studies cannot be extended to high redshift DLAs. Alternatively, the core fraction measured in the  milliarcsecond scale images can be used to get an estimate of the covering factor \\citep[see][] {Briggs89, Kanekar09vlba}. Here one assumes that the absorbing gas completely covers at least the emission from the milliarcsec scale core. Therefore, to address this issue, one needs, not only to increase the number of systems searched for 21-cm absorption but also to perform milliarcsecond scale imaging of the background radio sources.  We report here the results of a search for 21-cm  absorption in 10 DLAs at $z>2$ we have carried out using GBT and GMRT, complemented by L-band VLBA images of the background QSOs. \\begin{table} \\caption{Log of GBT and GMRT observations to search for 21-cm absorption} \\begin{center} \\begin{tabular}{lccccc} \\hline \\hline \\multicolumn{1}{c}{Source}& Tele-    & Date       & Time    &  BW      &    Ch.         \\\\ \\multicolumn{1}{c}{name}  & scope    &            &         &          &                \\\\ &          & yy-mm-dd           & (hr)    & (MHz)    &                \\\\ \\multicolumn{1}{c}{(1)} & (2)      &  (3)       &  (4)    &   (5)    &    (6)         \\\\ \\hline  J0407$-$4410 &  GBT  &  2006-10-20  &   4.7   &  0.625   &    512         \\\\ (CTS247)  &       &  2007-01-05  &         &          &                \\\\ &       &  2007-01-06  &         &          &                \\\\ &       &  2007-01-08  &         &          &                \\\\ J0733+2721  &  GBT  &  2007-12-05  &  10.7   &  1.25    &    512         \\\\ &       &  2007-12-06  &         &          &                \\\\ J0801+4725  &  GMRT &  2006-12-22  &  10.8   &  1       &    128         \\\\ &       &  2006-11-23  &         &          &                \\\\ J0852+2431  &  GBT  &  2009-08-07  &   5.6   &  1.25    &    1024        \\\\ &       &  2009-08-08  &         &          &                \\\\ &       &  2009-09-06  &         &          &                \\\\ J1017+6116  &  GBT  &  2008-10-15  &   4.5   &  1.25    &    1024        \\\\ &       &  2008-10-16  &         &          &                \\\\ &       &  2008-10-18  &         &          &                \\\\ J1242+3720  &  GMRT &  2007-06-08  &   6.1   &  0.5     &    128         \\\\ &       &  2007-06-10  &         &          &                \\\\ J1337+3152  &  GMRT &  2009-01-13  &   6.2   &  1       &    128         \\\\ &       &  2009-03-17  &   7.8   &  0.25    &    128         \\\\ J1406+3433  &  GBT  &  2009-03-05  &   8.0   &  1.25    &    1024        \\\\ &       &  2009-03-06  &         &          &                \\\\ &       &  2009-03-07  &         &          &                \\\\ &       &  2009-03-08  &         &          &                \\\\ &       &  2009-04-07  &         &          &                \\\\ J1435+5435  &  GMRT &  2007-06-10  &   6.7   &  1       &    128         \\\\ &       &                 &         &          &                \\\\ \\hline \\end{tabular} \\end{center} \\begin{flushleft} Column 1: Source name. Column 2: Telescope used for 21-cm absorption search. Column 3: Date of observations. Column 4: Total time on source i.e. after excluding telescope set-up time and calibration overheads. Columns 5 and 6: Spectral setup for the observations i.e. bandwidth (BW) and number of spectral channels respectively. \\end{flushleft} \\label{obslog} \\end{table} This survey has resulted in the detection of 21-cm absorption in the \\zabs = 3.1745 DLA towards J1337+3152. A detailed analysis of this system and two sub-DLAs close to this system are presented in \\citet{Srianand10}. Section~\\ref{samp} presents the details of our sample. In Section 3 we present the details of  GBT and GMRT spectroscopic observations, VLBA continuum observations, and data reduction. The detection rate of 21-cm absorption in DLAs is discussed in Section~\\ref{detect}. In Section~\\ref{gencor} we study the correlations between the parameters derived from 21-cm observations, $N$(H~{\\sc i}), metallicity and redshift. In Section~\\ref{mole} we study the relation between 21-cm and \\h2 absorption. The results  are summarized in Section~\\ref{results}. In this work we assume a flat Universe with $H_0$~=~71~\\kms~Mpc$^{-1}$, $\\Omega_{\\rm m}$~=~0.27 and $\\Omega_\\Lambda$~=~0.73. ",
        "conclusions": "\\label{results}  We have carried out a systematic search for 21-cm absorption in 10 DLAs at \\zabs$>2$ using GMRT and GBT. We detect 21-cm absorption in only one of them. From our sample we find the 21-cm detection rate is 13\\% for a $\\int \\tau dv$ limit of 0.4 km/s (the detection limit reached in the case of J1337+3152). We also obtained 1420~MHz VLBI images for the  sources in our sample.   The 21-cm detection at $z\\ge 2$ seems to favour systems with high metallicity  and/or  high $N$(H~{\\sc i}) \\citep[see also][]{Kanekar09ts, Curran10}. This basically means that the probability of detecting cold components that can produce detectable 21-cm absorption is higher in systems with high values of $N$(H~{\\sc i}) and Z.  However, we do not find any correlation between the integrated optical depth (or T$_{\\rm S}$/$f_{\\rm c}$) and $N$(H~{\\sc i}) or metallicity.   It is important to address the covering factor issue before drawing any conclusions on $T_S$. Ideally one should do high spatial resolution VLBA spectroscopy for this purpose \\citep[see for example][]{Lane00}. However, this is not possible at present specially for $z\\ge2$ absorbers. Therefore, we proceed by  assuming that the core fraction  found in the VLBA images as the covering factor of the absorbing gas \\citep[as in the case of][]{Kanekar09vlba}. We find that more than 50\\% of DLAs have weighted mean spin  temperature ($T_{\\rm S}$) in excess of 700 K. For the  assumed temperature of the CNM gas $T_{\\rm S}^C = 200$ K (as seen in \\h2 components in high-z DLAs) we find that more than 73\\% of H~{\\sc i} in such systems is originating from WNM. The median value CNM fraction (i.e $f$(CNM)) obtained for the detections and the median value of upper limits in the case of non-detections are in the range 0.2 to 0.25.  We study the connection between 21-cm and \\h2 absorption in a sub-sample of 13 DLAs where both these species can be searched for. We report the detection and detailed analysis of \\h2 molecules in the \\zabs=3.3871 DLA system towards J0203+1134 where 21-cm absorption is also detected. For a $b$ parameter in the range 1-5~\\kms\\  we find 14.57$\\le$log~$N$(\\h2)$\\le$16.03. The inferred kinetic temperature is in the range 48-108~K based on $T_{01}$ of H$_2$. However no 21-cm absorption is detected at the very position of this \\h2\\  component. This suggests that the H~{\\sc i} column density associated with this component is $\\le$ 10$^{19}$~cm$^{-2}$. However, the lack of proper coincidence between 21-cm and any of the strong UV absorption components may also mean that the radio and optical sight lines probe different volumes of the gas.  In the case of 8 DLAs, neither 21-cm nor H$_2$ are detected. Typical upper limits on the molecular fraction  ($f_{\\rm H_2}$) in these systems are $\\le 10^{-6}$. The lack of \\h2 in DLAs can be explained if the H~{\\sc i} gas originates from low density regions photoionized by the metagalactic UV \\citep[see for example,][]{Petitjean92,Petitjean00,Hirashita05}. This also indicates that the volume filling  factor of \\h2 in DLAs is small  \\citep{Zwaan06}.  Typical limits obtained for  $T_{\\rm S}$ in these systems are consistent with only a small fraction of the H~{\\sc i} gas originating from the CNM phase as suggested by the lack of \\h2 absorption.  In two cases strong \\h2 absorption is detected and kinetic temperatures are in the range 100-200~K, but 21-cm absorption is not detected.  Even in two cases where both the species are  detected they do not originate from the same velocity component. The lack of 21-cm absorption directly associated with \\h2\\  indicates that only a small fraction (typically $\\le$ 10\\%) of the neutral hydrogen seen in the DLA is associated with the \\h2 components \\citep[see also][]{Noterdaeme10co}. This implies that the molecular fractions $f$(\\h2) reported from the \\h2 surveys should be considered as conservative lower limits for the \\h2 components.  For two of the \\h2-bearing DLAs with density measurements based on C~{\\sc i} fine-structure excitation we derive an upper limit on the line of sight thickness of $\\le 15$~pc. { This is consistent with the size estimate for the H$_2$-bearing gas in \\zabs = 2.2377 DLA towards Q1232+082 based on partial coverage \\citep{Balashev11}.}  In principle, the presence of \\h2 and 21-cm absorptions in a single component provides a unique combination to simultaneously constrain the variation of the fine-structure constant ($\\alpha$), the electron-to-proton mass ratio ($\\mu$) and the proton G-factor. As shown here, DLAs with 21-cm and \\h2 detections are rare. Even in these cases the presence of multiphase structure at parsec scale is evident, introducing velocity shifts between the different absorption components that will affect the constraints on the variation of constants."
    },
    "1112/1112.3860_arXiv.txt": {
        "abstract": "{The clumpy density structure of photon-dominated regions is well established, but the physical properties of the clumps and of the surrounding interclump medium are only approximately known.} {The aim of this paper is to constrain the physical and chemical conditions in the Orion Bar, a prototypical nearby photon-dominated region.} {We present observations of the HF $J$=1--0 line, which appears in emission toward the Orion Bar, and compare the brightness of the line to non-LTE radiative transfer calculations. } {The large width of the HF line suggests an origin of the emission in the interclump gas, but collisional excitation by \\hh\\ in the interclump gas underpredicts the observed line intensity by factors of 3--5. In contrast, an origin of the line in the dense clumps requires a density of $\\sim$10$^9$\\ccm, 10--100 times higher than previous estimates, which is unlikely. However, electron impact excitation reproduces our observations for $T$=100\\,K and $n_e$=10\\,\\ccm, as expected for the interclump gas.} {We conclude that HF emission is a signpost of molecular gas with a high electron density. Similar conditions may apply to active galactic nuclei where HF also appears in emission.} ",
        "introduction": "\\label{s:intro}   Photon-dominated regions (PDRs) are the surface regions of molecular clouds, where ultraviolet radiation with photon energies between a few and 13.6\\,eV drives the thermal and chemical balance of the gas \\citep{hollenbach1999}. This situation occurs in regions of high-mass star formation, but also in protoplanetary disks and in the nuclei of active galaxies. Studying the structure of PDRs therefore has a wide astrophysical application. In PDRs, gas heating proceeds by photo-electric emission from dust grains, while the main cooling channels are the fine structure lines of C$^+$ and O and the rotational lines of CO \\citep{kaufman1999}. Absorption of the impinging ultraviolet radiation by dust and gas in the PDR creates a layered structure, where chemical transitions such as H$^+$ $\\to$ H $\\to$ \\hh\\ and C$^+$ $\\to$ C $\\to$ CO occur.  The Orion Bar is a prototypical PDR, located between the Orion molecular cloud and the Orion Nebula, the H$^+$ region surrounding the Trapezium stars, at a distance of 414\\,pc \\citep{menten2007}. Observations at infrared and submillimeter wavelengths indicate a geometry for the Bar where the PDR is wrapped around the Orion Nebula, and changes from a face-on to an edge-on view where the molecular emission peaks (\\citealt{hogerheijde1995}; \\citealt{walmsley2000}). The mean temperature of the molecular gas in the Bar is 85\\,K, while the temperature rises to several 100\\,K toward the ionization front, where the emission from PAH particles and vibrationally excited \\hh\\ peaks.  While the temperature structure of the Orion Bar is reasonably well understood, the same cannot be said about the density structure. The mean density of the molecular gas is 10$^5$\\,\\ccm, but single-dish observations already indicate the presence of small-scale density variations without apparent pattern, usually called ''clumps'' \\citep{hogerheijde1995}, that are also seen toward other PDRs \\citep{stutzki1988,wang1993}. While interferometric observations have confirmed the presence of clumps \\citep{youngowl2000}, the densities of both the clumps and the interclump medium are somewhat uncertain. The interclump medium probably has a density between a few 10$^4$ and \\pow{2}{5}\\,\\ccm\\ \\citep{simon1997}, while estimates of the clump density range from \\pow{1.5}{6} to \\pow{6}{6}\\,\\ccm\\ \\citep{lis2003}.  This Letter presents observations of the HF $J$=1--0 line, which appears in emission toward the Orion Bar. The HF molecule is expected to be the dominant carrier of gas-phase fluorine, because the reaction F + \\hh\\ $\\to$ HF + H is exothermic. For diffuse clouds where the effect of depletion on grains should be unimportant, models by \\citet{neufeld2009} predict an HF abundance of $\\sim$\\pow{3.6}{-8} relative to \\hh. Recent observations of the HF 1--0 line confirm this prediction: the line is seen in absorption toward several background sources, indicating abundances of $\\sim$\\pow{1--2}{-8} \\citep{neufeld2010}. Toward dense clouds, the abundance is measured to be $\\sim$100 times lower \\citep{phillips2010}, suggesting that depletion of F on grain surfaces or excitation effects play a role.  Like CO, HF is a linear rotor with a regular line spectrum, where $\\nu = 2B(J+1)$ and $A_{ij} \\propto \\nu^3$. Unlike CO, HF has a small reduced mass and a large dipole moment, so that the lines have high frequencies and radiative decay is rapid. In particular, the HF 1--0 line has a frequency of 1232\\,GHz and an Einstein~A coefficient of \\pow{2.422}{-2}\\,s$^{-1}$. Thermal excitation of the $J$=1 level of HF thus requires extremely high gas densities, which is why the line usually appears in absorption \\citep{sonnentrucker2010,monje2011}. ",
        "conclusions": "\\label{s:concl}  We have presented observations of the HF $J$=1--0 line toward the Orion Bar, which is the first time that this line is seen in emission from the Galactic interstellar medium. We present calculations and arguments why this emission cannot be caused by collisional excitation of HF by \\hh, either from dense clumps or from the inter-clump medium. However, collisional excitation by electrons in the interclump gas explains our observations. The appearance of HF in emission therefore seems to be a signpost of molecular gas with a high electron density, which in the case of the Orion Bar is the combined effect of the high gas density and the strong ultraviolet radiation field.  The physical conditions of the Orion Bar may be similar to those in the nucleus of the active galaxy Mrk 231, where the radiation field is strong and where the HF line also appears in emission \\citep{vdwerf2010}. Other active galactic nuclei where HF appears in absorption such as Arp 220 \\citep{rangwala2011} and the Cloverleaf \\citep{monje2011clover} probably have lower electron densities, which may indicate a softer radiation field.  Support for an origin of the HF emission in the interclump medium comes from the detection of CF$^+$ toward the Orion Bar \\citep{neufeld2006}. Additional support for this model would come from observations of H$_2$F$^+$, for which line frequencies are known \\citep{fujimori2011}. However, our data do not show any features at the predicted line frequencies down to an rms of $\\approx$50\\,mK in $T_A^*$, which is a factor of $\\sim$10 above the expected signal.  In the future, maps of HF emission with Herschel-HIFI will help to clarify the spatial distribution of dense gas clumps in the Orion Bar and other PDRs. Another useful test would be the observation of HF $J$=2--1 or higher-$J$ lines in the $v$=0 or $v$=1 states. Unfortunately, the $J$=2--1 line at 2463\\,GHz lies very close (0.13\\,\\mic) to the N$^+$ $^3P_2$--$^3P_1$ fine-structure line at 121.8\\,\\mic, so that the lines are blended at the resolution of the PACS instrument onboard Herschel. Heterodyne observations in this frequency range may become possible with the future STO telescope, the SOFIA airborne observatory, or the Millimetron mission."
    },
    "1112/1112.4995_arXiv.txt": {
        "abstract": "The universe evolution during the radiation-dominated epoch in the $R^2$--extended gravity theory is considered. The equations of motion for $R$ and $H$ are solved analytically and numerically. The particle production rate by the oscillating curvature is calculated in one-loop approximation and the back-reaction of particle production on the evolution of $R$ is taken into account. Possible implications of the model for cosmological creation of non-thermal dark matter is discussed. ",
        "introduction": "During the last two decades the accumulated astronomical data have unambiguously proved that the universe expansion is accelerated. The data include the observation of the large scale structure of the universe, the measurements of the angular fluctuations of the cosmic microwave background radiation, the determination of the universe age (for a review see~\\cite{cosm-prmtr}), and especially the discovery of the dimming of distant Supernovae~\\cite{Nobel_2011}.   The driving force behind this accelerated expansion in unknown. Among possible explanations, the most popular is probably the assumption of a new (unknown) form of cosmological energy density with large negative pressure, $ P < -\\rho/3$, the so-called dark energy, for a review see e.g.~\\cite{DE_Peebles_Ratra}. %  A competing mechanism to describe the accelerated expansion is represented by gravity modifications at small curvature, the so-called $f(R)$-gravity theories, as suggested in ref.~\\cite{grav-mdf}. In these theories the standard Einstein-Hilbert Lagrangian density, proportional to the scalar curvature $R$, is replaced by a function $f(R)$, so the usual action of General Relativity acquires an additional term: \\be S = -\\frac{m_{Pl}^2}{16\\pi} \\int d^4 x \\sqrt{-g}\\,f(R)+S_m= -\\frac{m_{Pl}^2}{16\\pi} \\int d^4 x \\sqrt{-g}\\,\\left[R+F(R)\\right]+S_m\\, , \\label{A1} \\ee where $m_{Pl}= 1. 2 2\\cdot 10^{19}$ GeV is the Planck mass and $S_m$ is the matter action.  The original version of such models suffers from a strong instability in presence of gravitating bodies~\\cite{DolgKaw} and because of that more complicated functions $ F(R) $ have been proposed~\\cite{Starob, HuSaw, ApplBatt,Noj-Odin-2007}, which are free from the mentioned exponential instability. It was claimed in an earlier paper~\\cite{odin-instab} that the instability could be eliminated by a mere addition of terms of the type $R^q/m_q^{2(q-1)}$ to the action, where $m_q$ is a constant parameter with dimension of mass. To some extent such terms may suppress the instability for systems with relatively high mass density, $\\rho > 1 {\\rm g/cm^{3}}$, but for less dense systems e.g. for those with $\\rho \\sim 10^{-24} {\\rm g/cm{^3}}$ this mechanism demands too large coefficients in front of $R^q$ (small $m_q$), which are at odds with big bang nucleosynthesis.   Though free of instability~\\cite{DolgKaw}, the models proposed in~\\cite{Starob, HuSaw, ApplBatt} possess another troublesome feature, namely in a cosmological situation they should evolve from a singular state in the past~\\cite{appl-bat-08}. Moreover, it was found in refs.~\\cite{frolov, Arb_Dolgov} that in presence of matter, a singularity may arise in the future if the matter density rises with time; such future singularity is unavoidable, regardless of the initial conditions, and is reached in a time which is much shorter than the cosmological one. In the standard Friedmann cosmology, in which the energy density decreases with time, a future singularity may appear in several scenarios (phantom cosmology, quintessence models, \\dots), but not in models~\\cite{Starob, HuSaw,ApplBatt}, as shown in ref.~\\cite{Arb_Dolgov}. This statement is in straight disagreement with earlier papers~\\cite{odin-future-sing}.  The subject of modified gravity includes about $10^3$ papers at the present time and it is impossible to quote many relevant works. For a review one may look into references~\\cite{App-Bat-Star,NojOd,clifton}.  The aforementioned problems can be cured by adding to the action an $R^2$-term, which prevents from the singular behavior both in the past and in the future. In the present work we study the cosmological evolution of the Universe in a theory with only an additional $R^2$ term in the action, neglecting other terms which have been introduced to generate the accelerated expansion in the contemporary universe. The impact of such terms is negligible in the limit of sufficiently large curvature, $|R|\\gg |R_0|$, where $R_0$ is the cosmological curvature at the present time.  In other words, we study here the cosmological evolution of the early and not so early universe in the model with the action: \\be S = -\\frac{m_{Pl}^2}{16\\pi} \\int d^4 x \\sqrt{-g} \\left(R-\\frac{R^2}{6m^2}\\right)+S_m\\,, \\label{A-R2} \\ee where $m$ is a constant parameter with dimension of mass.\\\\ Cosmological models with an action quadratic in the curvature tensors were pioneered in ref.~\\cite{Gur-Star}. Such higher-order terms appear as a result of radiative corrections to the usual Einstein-Hilbert action after taking the expectation value of the energy-momentum tensor of matter in a curved background. In such models, like for instance the Starobinsky model~\\cite{Starobinsky_1980}, the universe may have experienced an exponential (inflationary) expansion without invoking phase transitions in the very early universe. This model has a graceful exit to matter-dominated stage which is induced by the new scalar degree of freedom, the \\textit{scalaron} (curvature scalar), which becomes a dynamical field in $R^2$-theory.  The reheating process, due to gravitational particle production from scalaron (curvature scalar) oscillations, leads to a transition to a Friedmann-like universe. These features of the model are thoroughly discussed, for instance, in refs.~\\cite{Vilenkin_1985,Mijic_Morris_Suen_1986,Suen_Anderson}. Cosmological dynamics of fourth-order gravity were investigated in several works, see e.g.\\cite{Carloni_Dunsby_Troisi_2009} and references therein.  A somewhat similar study was performed in ref.~\\cite{davidson} where a version of massive Brans-Dicke (BD) theory without kinetic term (i.e. with BD parameter $\\omega =0$) was considered. The Hubble parameter and curvature demonstrate oscillating behavior which resembles the one found in our paper (and earlier in many others), but quantitative features are very much different.   Beside $R^2$-terms, terms containing the Ricci tensor squared $ R_{\\mu\\nu} R^{\\mu\\nu}$ are induced by radiative corrections as well, and with similar magnitude. However, the natural magnitude of such radiatively induced terms is quite small. The characteristic mass parameter, in fact, is of the order of the Planck mass in both cases, which makes this situation non-interesting for applications discussed below. On the other hand, $R^2$ cosmology (without  $ R_{\\mu\\nu} R^{\\mu\\nu} $) has been considered in the literature with much larger magnitude of $R^2$ than the natural value from radiative corrections. The assumption of large $R^2$ terms is made \\textit{ad hoc} to formulate a model which could, for instance, cure singularities. We follow the spirit of those works. However, it could be worthwhile to study the consequences of more complicate models with both $R^2$ and $ R_{\\mu\\nu} R^{\\mu\\nu} $ terms. It may be a subject for future investigation.  The paper is organized as follows. In sec.~\\ref{sec:eqs_of_motion} the cosmological equations modified by the presence of $R^2$-term are presented and solved both analytically and numerically in the case of a radiation-dominated (RD) universe.  We study two equivalent sets of equations, for the time-time component of the (modified) Einstein equations and for their trace. In sec.~\\ref{s-part-prod} particle production by the external oscillating gravitational field is studied. First, we derive the equation of motion for the evolution of $R$ with the account of the back-reaction from particle production. This leads to an exponential damping of the oscillating part of $R$, while the non-oscillating \"Friedmann\" part remains practically undisturbed. The particle production influx into the cosmological plasma is estimated in the case of a massless, minimally-coupled scalar field. In the conclusions we discuss possible implications of the scenario, in particular for heavy supersymmetric dark matter. ",
        "conclusions": "The characteristic decay time of the oscillating curvature is \\be \\tau_R = \\frac{1}{2\\Gamma_R} =  \\frac{24 m^2_{Pl}}{m^3} \\simeq 2 \\left(\\frac{10^5\\text{ GeV}}{m}\\right)^3 \\text{ seconds}\\,. \\label{tau-R} \\ee The contribution of the produced particles into the total cosmological energy density reaches its maximum value at approximately this time. The ratio of the energy density of the newly produced energetic particles and that of those already existing in the plasma, according to eq.~(\\ref{rho-long-time}), is: \\be \\frac{\\rho_{hi}}{\\rho_{therm}} = \\frac{8\\beta^2 N_{eff}}{\\kappa(2\\alpha_1-1)}\\,\\frac{1-(2\\Gamma_Rt_{in})^{2\\alpha_1-1}}{(2\\Gamma_R t_{in})^{2\\alpha_1-2}}\\,. \\label{ratio} \\ee If we take $t_{in} \\simeq 1/m$, then $t_{in} \\Gamma_R \\simeq m^2/m_{Pl}^2 \\ll 1$ and the effects of non-thermalized matter may be negligible. However, for sufficiently large $\\beta$ and possibly small $\\kappa$ the non-thermal particles may play a significant role in the cosmological history.\\\\ The influx of energetic protons and antiprotons could have an impact on BBN. Thus this would either allow to obtain some bounds on $m$ or even to improve the agreement between the theoretical predictions for BBN and the measurements of primordial light nuclei abundances.  The oscillating curvature might also be a source of dark matter in the form of heavy supersymmetric (SUSY) particles. Since the expected light SUSY particles have not yet been discovered at LHC, to some people supersymmetry somewhat lost its attractiveness. The contribution of the stable lightest SUSY particle into the cosmological energy is proportional to \\be \\Omega \\sim{ m^2_{SUSY} }/m_{Pl} \\label{Omega-SUSY} \\ee and for $m_{SUSY} $ in the range $100-1000$ GeV the cosmological fraction of these particles would be of order unity. It is exactly what is necessary for dark matter. However, it excludes thermally produced LSP's if they are much heavier. If LSP's came from the decay of $R$ and their mass is larger than the ``mass'' of $R$, i.e. $m$, the LSP production could be sufficiently suppressed to make a reasonable contribution to dark matter.  These and other manifestations of the considered scenario will be discussed elsewhere."
    },
    "1112/1112.2057_arXiv.txt": {
        "abstract": "We consider the construction of generic spherically symmetric thin-shell traversable wormhole spacetimes in standard general relativity. By using the cut-and-paste procedure, we comprehensively analyze the stability of arbitrary spherically symmetric thin-shell wormholes to linearized spherically symmetric perturbations around static solutions. While a number of special cases have previously been dealt with in scattered parts of the literature, herein we take considerable effort to make the analysis as general and unified as practicable. We demonstrate in full generality that stability of the wormhole is equivalent to choosing suitable properties for the exotic material residing on the wormhole throat. ",
        "introduction": "Traversable wormholes are hypothetical tunnels in spacetime, through which observers may freely travel~\\cite{Morris, 269623}. These geometries are supported by ``exotic matter'', involving a stress-energy tensor violating the null energy condition (NEC).  That is, there exists at least one null vector $k^\\mu$ such that  $T_{\\mu\\nu}k^{\\mu}k^{\\nu}<0$ on or in the immediate vicinity of the wormhole throat (see, in particular,~\\cite{Visser, gr-qc/9704082, gr-qc/9710001}). In fact, wormhole geometries violate all the standard pointwise energy conditions, and all the standard averaged energy conditions~\\cite{Visser}. Although (most) classical forms of matter are believed to obey (most of) the standard energy conditions~\\cite{hawkingellis}, it is a well-known fact that they are violated by certain quantum effects, amongst which we may refer to the Casimir effect, Hawking evaporation, and vacuum polarization. (See~\\cite{Visser}, and more recently~\\cite{gr-qc/0205066}, for a review. For some technical details see~\\cite{gvp}.) It is interesting to note that the known violations of the pointwise energy conditions led researchers to consider the possibility of averaging of the energy conditions over timelike or null geodesics~\\cite{Tipler-AEC}. For instance, the averaged weak energy condition (AWEC) states that the integral of the energy density measured by a geodesic observer is non-negative. (That is, $\\int T_{\\mu\\nu}U^\\mu U^\\nu \\,d\\tau \\geq 0$, where $\\tau$ is the observer's proper time.) Thus, the averaged energy conditions are weaker than the pointwise energy conditions, they permit localized violations of the energy conditions, as long as they hold when suitably averaged along a null or timelike geodesic~\\cite{Tipler-AEC}.  As the theme of exotic matter is a problematic issue, it is useful to minimize its usage. In fact, it is important to emphasize that the theorems which guarantee the energy condition violation are remarkably silent when it comes to making quantitative statements regarding the ``total amount'' of energy condition violating matter in the spacetime. In this context, a suitable measure for quantifying this notion was developed in \\cite{VKD}, see also~\\cite{gr-qc/0405103}, where it was shown that wormhole geometries may in principle be supported by arbitrarily small quantities of exotic matter (an interesting application of the quantification of the total amount of energy condition violating matter in warp drive spacetimes was considered in \\cite{gr-qc/0406083}). In the context of minimizing the usage of exotic matter, it was also found that for specific models of stationary and axially symmetric traversable wormholes the exotic matter is confined to certain regions around the wormhole throat, so that certain classes of geodesics traversing the wormhole need not encounter any energy condition violating matter \\cite{Teo:1998dp}. For dynamic wormholes the null energy condition, more precisely the averaged null energy condition, can be avoided in certain regions \\cite{hochvisserPRL98, hochvisserPRD98, gr-qc/9901020}. Evolving wormhole geometries were also found which exhibit ``flashes'' of weak energy condition (WEC) violation, where the matter threading the wormhole violates the energy conditions for small intervals of time \\cite{Kar}. In the context of nonlinear electrodynamics, it was found that certain dynamic wormhole solutions obey (suitably defined versions of) the WEC \\cite{Arellano:2006ex}.  It is interesting to note that in modified theories of gravity, more specifically in $f(R)$ gravity, the matter threading the wormhole throat can be forced to obey (suitably defined versions of) all the energy conditions, and it is the higher-order curvature terms that are responsible for supporting these wormhole geometries \\cite{modgravity}. (Related issues arise when scalar fields are conformally coupled to gravity~\\cite{gr-qc/0003025, gr-qc/9908029}.) In braneworlds, it was found that it is a combination of the local high-energy bulk effects, and the nonlocal corrections from the Weyl curvature in the bulk, that may induce an effective NEC violating signature on the brane, while the ``physical''/``observable'' stress-energy tensor confined to the brane, threading the wormhole throat, nevertheless satisfies the energy conditions~\\cite{Lobo:2007qi}.  An interesting and efficient manner to minimize the violation of the null energy condition, extensively analyzed in the literature, is to construct thin-shell wormholes using the thin-shell formalism \\cite{Visser,thinform,Israel} and the cut-and-paste procedure as described in~\\cite{Visser,VisserPRD,VisserNPB,VisserPLB,Poisson}. Motivated in minimizing the usage of exotic matter, the thin-shell construction was generalized to nonspherically symmetric cases~\\cite{Visser,VisserPRD}, and in particular, it was found that a traveler may traverse through such a wormhole without encountering regions of exotic matter. In the context of a (limited) stability analysis, in \\cite{VisserNPB}, two Schwarzschild spacetimes were surgically grafted together in such a way that no event horizon is permitted to form. This surgery concentrates a nonzero stress energy on the boundary layer between the two asymptotically flat regions and a dynamical stability analysis {(with respect to spherically symmetric perturbations)} was explored. In the latter stability analysis, constraints were found on the equation of state of the exotic matter that comprises the throat of the wormhole. Indeed, the stability of the latter thin-shell wormholes was considered for certain specially chosen equations of state \\cite{Visser,VisserNPB}, where the analysis addressed the issue of stability in the sense of proving bounded motion for the wormhole throat. This dynamical analysis was generalized to the stability of spherically symmetric thin-shell wormholes by considering linearized radial perturbations around some assumed static solution of the Einstein field equations, without the need to specify an equation of state \\cite{Poisson}. This linearized stability analysis around a static solution was soon generalized to the presence of charge \\cite{Eiroa}, and of a cosmological constant \\cite{LoboCraw}, and was subsequently extended to a plethora of individual scenarios \\cite{thinshellwh,Ishak}, some of them rather ad hoc.  The key point of the present paper is to develop an extremely general, flexible, and robust framework that can quickly be adapted to general spherically symmetric traversable wormholes in 3+1 dimensions. We shall consider standard general relativity, with traversable wormholes that are spherically symmetric, with all of the exotic material confined to a thin shell. The bulk spacetimes on either side of the wormhole throat will be spherically symmetric and static but otherwise arbitrary (so the formalism is simultaneously capable of dealing with wormholes embedded in Schwarzschild, Reissner--Nordstr\\\"om, Kottler, or de~Sitter spacetimes, or even ``stringy'' black hole spacetimes).  The thin shell (wormhole throat), while constrained by spherical symmetry,  will otherwise be permitted to move freely in the bulk spacetimes, permitting a fully dynamic analysis. This will then allow us to perform a general stability analysis against spherically symmetric perturbations, where wormhole stability is related to the properties of the exotic matter residing on the wormhole throat. We particularly emphasize that our analysis can deal with geometrically imposed ``external forces'', (to be more fully explained below), a feature that has so far been missing from the published literature. Additionally we emphasize the derivation of rather explicit and very general rules relating the internal structure of the wormhole throat to a ``potential'' that drives the motion of the throat.  This paper is organized in the following manner: In Sec. \\ref{secII} we outline in detail the general formalism of generic dynamic spherically symmetric thin-shell wormholes, and provide a novel approach to the linearized stability analysis around a static solution. In Sec. \\ref{secIII}, we provide specific examples by applying the generic linearized stability formalism outlined in the previous section. In Sec. \\ref{conclusion}, we draw some general  conclusions. ",
        "conclusions": "In this work, we have developed an extremely general, flexible and robust framework, leading to the linearized stability analysis of spherically symmetric thin shells. The analysis is well-adapted to general spherically symmetric thin shell traversable wormholes and, in this context, the construction confines the exotic material to the thin-shell. The latter, while constrained by spherical symmetry is allowed to move freely within the bulk spacetimes, which permits a fully dynamic analysis. To this effect, we have considered in great detail the presence of a flux term, which has been widely ignored in the literature. This flux term corresponds to the net discontinuity in the conservation law of the surface stresses of the bulk momentum flux, and is physically interpreted as the work done by external forces on the thin shell.  Relative to the linearized stability analysis, we have reversed the logic flow typically considered in the literature, and introduced a novel approach to the analysis. We recall that the standard procedure extensively used in the literature is to define a parametrization of the stability of equilibrium, so as not to specify an equation of state on the boundary surface \\cite{Poisson,Eiroa,LoboCraw}. More specifically, the parameter $\\eta(\\sigma)= d{\\cal P}/d\\sigma$ is usually defined, and the standard physical interpretation of $\\eta$ is that of the speed of sound. In this work, rather than adopt the latter approach, we considered that the stability of the wormhole is fundamentally linked to the behaviour of the surface mass $m_s(a)$ of the thin shell of exotic matter, residing on the wormhole throat, via a pair of stability inequalities. More specifically, we have considered the surface mass as a function of the potential. This novel procedure implicitly makes demands on the equation of state of the matter residing on the transition layer, and demonstrates in full generality that the stability of thin shell wormholes is equivalent to choosing suitable properties for the material residing on the thin shell.    We have applied the latter stability formalism to a number of specific examples of particular importance: some presented to emphasize the features specific to possible asymmetry between the two universes used in traversable wormhole construction; some to emphasize the importance of NEC non-violation in the bulk; and some to assess the simplifications due to symmetry between the two asymptotic regions. In particular, we have considered the case of borderline NEC non-violation in the bulk. This is motivated by the knowledge that, on extremely general grounds, the NEC must be violated somewhere in the spacetime of a generic traversable wormhole, and if this were to happen in the bulk region, this would be equivalent to imposing $\\Phi'_\\pm(r)< 0$ in the bulk.  Thus to minimize NEC violations for thin-shell wormholes, we have considered the specific case of $\\Phi_\\pm=0$, which is particularly interesting for its mathematical simplicity, and for its physical interest as it corresponds to the constraint that the bulk regions on either side of the wormhole throat be on the verge of violating the NEC. We have also considered the simplification when the two bulk regions are identical, and analysed the stability regions of asymmetric thin-shell Schwarzschild wormholes and thin-shell Reissner-Nordstr\\\"{o}m wormholes in great detail. It was instructive to consider explicit cases which violate some of the energy conditions in the bulk. For instance, we considered thin-shell variants of the Ellis wormhole and two new toy models, and explored the linearized stability analysis in the presence of zero momentum flux and non-zero external forces. Finally, we analyzed thin-shell dilatonic wormholes, where the exterior spacetime solutions corresponded to a combined gravitational-electromagnetic-dilaton system.   In conclusion, by considering the matching of two generic static spherically symmetric spacetimes using the cut-and-paste procedure, we have analyzed the stability of spherically symmetric dynamic thin-shell traversable wormholes --- stability to linearized spherically symmetric perturbations around static solutions. The analysis provides a \\emph{general and unified framework} for simultaneously addressing a large number of wormhole models scattered throughout the literature.  As such we hope it will serve to bring some cohesion and focus to what is otherwise a rather disorganized and disparate collection of results. A key feature of the current analysis is that we have been able to include ``external forces'' in the form of non-zero values for the metric functions $\\Phi_\\pm(r)$. (This feature is absent in much of the extant literature.) Another key aspect of the current analysis is the focus on $m_s(a)$, the ``mass'' of the thin shell of exotic matter residing on the wormhole throat --- and the realization that stability of the wormhole is fundamentally linked to the behaviour of this exotic matter via a pair of simple and relatively tractable inequalities."
    },
    "1112/1112.0863_arXiv.txt": {
        "abstract": "Peculiar A stars are so named because they exhibit abundance peculiarities in their atmospheres.  It is believed that these arise as a result of differentiation of chemical species in large magnetic spots in which convective mixing is inhibited: there might be just two antipodal spots, whose axis is inclined to the axis of rotation.  Many of the Ap stars that are rotating slowly also pulsate, with periods substantially shorter than the period of the fundamental radial mode. The pulsations appear to be nonradial, but axisymmetric, with their common axis usually aligned with the axis of the spots.  In this lecture I shall first discuss the magnetic suppression of convection in the spots, and then I shall try to explain the pulsation phenomenon, reviewing some of the suggestions that have been made to explain the alignment and the excitation mechanism, and finally raising some issues that need to be addressed. ",
        "introduction": "It was my privilege to be the first research student to be registered by the University of Cambridge under Roger Tayler's supervision. That did not make me Roger's first student in practice, however, because Roger was a very kindly man and would supervise other astrophysics students in the department who needed more help than their registered supervisors provided; and so there were already several students under Roger's wing when I arrived. Consequently Roger was at that time intellectually active in a variety of fields, embracing stellar structure and evolution, the physics of thermonuclear reactions, magnetohydrodynamics -- particularly stability theory, with which Roger had a great deal of experience -- and stellar pulsation. It is therefore appropriate that I devote this lecture to a subject that combines many of these fields. I have chosen to address some of the issues that arise from trying to understand the phenomenon of rapid oscillations in peculiar A stars. Roger did not himself work in this area, but, recognizing his interests, one can easily imagine that he might well have done so had his attention not been drawn more strongly to other matters.  My own interest in rapidly oscillating Ap (roAp) stars was triggered by Donald Lynden-Bell, when in 1981, if I recall correctly, he told me on his return from a visit to the South African Astronomical Observatory of Don Kurtz's recent, and yet unpublished, discovery. What intrigued me was that the stars were spotted and that the oscillations were apparently dipolar, with axes aligned with the spots. There were only a few examples, but nevertheless they must surely have at least hinted very strongly that the oscillations are always (almost) aligned with the spots, and therefore do not wander far off under the influence of Coriolis precession. Soon after that time, N\\\"{o}el Dolez visited me from Paris, and we decided to carry out some simple calculations that might address the most obvious questions: why are only dipolar oscillations excited; why are their axes aligned with the spots; and why are only rapid oscillations observed, with frequencies much higher than that of the fundamental dipole mode? I address unashamedly in this lecture the picture that is emerging at least in my mind from that early work, and which has been substantially elaborated upon in recent years; I cannot here also review in detail the alternative suggestions that have been propounded. And I emphasize here that the line of argument that I follow is strongly influenced by what I learnt at Roger's feet when I was a student. ",
        "conclusions": "It is my opinion that a theory of the roAp-star phenomenon is beginning to emerge.  The stars are slow rotators.  They are spotted as a result of there being a large-scale predominantly dipole magnetic field pervading the star, inclined from the axis of rotation, which suppresses subsurface convection to form two antipodal spots where the field is almost vertical.  This permits the establishment of chemical abundance anomalies brought about a combination of radiative levitation and gravitational settling against diffusion.  There might also be a stellar wind from the spot, where the magnetic field lines are effectively open. This would modify the chemical abundances, as also might transport by fingering in regions where the mean molecular mass increases upwards, although this process, unfortunately, is not well understood. The oscillations are low-degree $\\rm{p}$ models of high order, which are probably axisymmetric with axes more-or-less aligned with the spots. They are excited by the $\\kappa$ mechanism in the spots, and damped by convection elsewhere.  Only a few high-order modes are overstable, in accord with observation. Some calculations have found overstability also in some relatively-low-order modes, with much lower growth rates; perhaps such oscillations are present too, but at amplitudes too low yet to have been observed.  An important matter that has never been seriously addressed is an explanation of the amplitudes of the modes.  If the modes are overstable, their growth must be limited by some nonlinear process. Is that process similar to that which limits the growth of the oscillations of Cepheids and RR Lyrae stars?  If so, why does it operate effectively at so low and amplitude?  Only when the spots are sufficiently large, and the rotation of the star is sufficiently small, can axisymmetric oscillation modes be dynamically locked with their axes close enough to the spots for excitation to dominate over damping, causing overstability; nonaxisymmetric modes precess, due to Coriolis effects, and cannot be firmly locked. Because the oscillation amplitude is greatest on the axis of symmetry, overstability is most likely to result when the modes are aligned with the spots. Of course, the larger the spots, the less precise need be the alignment, and in the fullness of time a robust theoretical criterion for overstability might be found: a relationship between spot size and  inclination of the oscillation axis from the spots.  Nonaxisymmetric modes have lower growth rates, and are overstable only if the spots are very large; that they appear not to be observed would set an upper limit to the size of the spots, once a robust theory is available.  In some circumstances the stellar rotation might be so large that the principal axis of pulsation is more-or-less aligned with the rotation axis, and not with the spots; in that case, it might be that other modes of the same degree precess slowly across the spot and become overstable only when they are appropriately orientated, giving the impression of dynamical alignment.  The theoretical models are at present extremely primitive. The influence of the spots on the pulsations has been accommodated in a piecemeal way, and must in future be properly incorporated into the pulsation dynamics to produce a plausible quantitative model.  There has been substantial progress in studying the influence of the Lorentz forces on the pulsations, but only in stellar models without spots. Lateral inhomogeneity of the background state, convection-pulsation interactions, Lorentz forces and nonadiabaticity must all be combined in a consistent way.  And our understanding of the dynamics of the oscillations in the atmosphere, where nonlinearity in the dynamics can be important, and where the processes of amplitude limitation are probably the most effectual, must be improved.  Only then can one reliably use future observations of mode frequencies and amplitudes, intensity-velocity phase relations and nonlinearities in the light-curves to draw reliable inferences about the structure of the stars.  We have many (but not all) of the ingredients in hand. Therefore a preliminary theory might be almost in sight."
    },
    "1112/1112.2782_arXiv.txt": {
        "abstract": "We demonstrate that the integrated gravitational wave signal of Type Ia supernovae (SNe Ia) in the single-degenerate channel out to cosmological distances gives rise to a continuous background to spaceborne gravitational wave detectors, including the Big Bang Observer (BBO) and Deci-Hertz Interferometer Gravitational wave Observatory (DECIGO) planned missions. This gravitational wave background from SNe Ia acts as a noise background in the frequency range 0.1 - 10 Hz, which heretofore was thought to be relatively free from astrophysical sources apart from neutron star binaries, and therefore a key window in which to study primordial gravitational waves generated by inflation.  While inflationary energy scales of $\\gtrsim 10^{16}$ GeV yield inflationary gravitational wave backgrounds in excess of our range of predicted backgrounds, for lower energy scales of $\\sim10^{15}$ GeV, the inflationary gravitational wave background becomes comparable to the noise background from SNe Ia. ",
        "introduction": "Type Ia supernovae (SNe Ia) are powerful thermonuclear explosions of a white dwarf star that release a kinetic energy of $\\sim10^{51}$ ergs, and serve as standard candles for cosmology \\cite {phillips93, riessetal98, perlmutteretal99}.  While the precise nature of the SNe Ia progenitor remains a topic of active investigation, current theoretical models   \\cite {townsleyetal07, jordanetal09, hoflichetal10} of  the detonation of a white dwarf star brought to a mass near the Chandrasekhar limit by a companion red giant or main sequence star are consistent with both the energetics and nucleosynthetic yield of the explosion, as well as its observed asymmetry \\cite {maedaetal10}. We refer to this particular mechanism of SNe Ia as the single degenerate channel (SD)\\citep{whelanetal73,nomotoetal84}, as opposed to the double degenerate (DD) channel, which consists of the merger and subsequent detonation of two degenerate white dwarfs \\citep {webbink84, ibentutukov84}.  Two models for detonations in the SD channel have been studied extensively -- the deflagration-to-detonation transition (DDT) \\citep {khokhlov91, blondinetal11}, and the  gravitationally-confined detonation (GCD). Utilizing detailed three-dimensional (3D) hydrodynamical simulations of the GCD mechanism of the SD channel of SNe Ia, \\cite{faltaetal11} demonstrated that the intrinsic asymmetry of the explosion gives rise to a gravitational wave (GW) signature. Based on the GCD model prediction, a nearby individual SNe Ia event may possibly be detected in the Milky Way and nearby galaxies by third-generation GW detectors in the $\\sim 2$ Hz frequency range. While the focus of the earlier paper \\citep {faltaetal11} was upon the GCD model for detonation in the SD channel,  both the DDT and DD models are likely to yield off-centered detonations \\citep {ropkeetal07, vankerkwijketal10}.  Recent work has demonstrated that the explosion energies and nucleosynthetic yields are similar in off-centered detonations in which only a small fraction of the total energy release occurs prior to detonation \\citep {chamulaketal11}. Since the great majority of the GW energy is radiated during the detonation itself, we expect both the DDT and DD models are likely to yield off-centered detonations whose gravitational wave signature is qualitatively similar to the GCD model prediction.  Previous work explored the stochastic gravitational wave background due to the distribution of  unresolved ``chirping'' binary white dwarfs, and demonstrated that this background obscures the inflationary signal between frequencies of 10$^{-5}$ Hz and 0.1 Hz \\cite {farmerandphinney03}. The frequency range between 0.1 - 10 Hz was previously thought to be relatively free from gravitational wave emission from astrophysical sources, with the primary contribution believed to originate from neutron star binaries \\cite {schneider01}, and possibly also core-collapse supernovae at cosmological distances, and Pop III (e.g., first-generation) stars \\cite {buonannoetal05}. In this paper, we consider an additional stochastic GW background source : SNe Ia at cosmological distances. If sufficiently strong, the SNe Ia stochastic background may be a source of noise to future detectors in the 0.1 - 10 Hz frequency range. Thus, the presence of a stochastic gravitational wave background due to SNe Ia may further impose additional challenges in extracting a possible inflationary GW background in this frequency range.  The paper is organized as follows. In section II, we present the formalism underlying the computation of the cosmological gravitational wave background of stochastic sources. In section III, we discuss the cosmic rate of SNe Ia. In section IV, we review the gravitational wave signature of a single SNe Ia event predicted by the GCD mechanism. Finally, in sections V and VI, we present our results and conclude. ",
        "conclusions": "A population of purely SD GCD SNe Ia at cosmological distances will generate a continuous GW stochastic background within the 1 - 2 Hz frequency range that could be a source of noise to third-generation detectors seeking the inflationary GW signal. The SNe Ia GW stochastic signal strength is relatively insensitive to the DTD, with the prompt DTD producing a slightly larger signal than the $\\beta = -1.1$ power law. The SNe Ia background could be a significant source of noise for the ultimate DECIGO detector if the inflationary signal is at a level less than $\\Omega_{GW} h_0 \\simeq 10^{-19}$. However assuming a slow roll inflation that can be constrained by the CMB/LSS, the inflationary signal could be as strong as $\\Omega_{GW} h_0 \\lesssim 6\\times 10^{-15}$, in which case the SNe Ia background would be too weak to be a source of noise. If the SNe Ia GW signal strength is significantly stronger, such that the background is closer to the upper bound estimate, it could also be a source of noise to the correlated BBO.    Previous work on the primordial inflationary gravitational wave spectrum has primarily focused upon neutron-star binaries as a source of noise in the frequency range of 0.1 - 10 Hz \\cite {schneider01}. The gravitational wave signal generated by these binaries is relatively clean. Therefore it seems that it may be possible to subtract the neutron-star binary spectrum to reveal the underlying primordial gravitational wave spectrum. In contrast, the physics underlying the SNe Ia is far more complex, involving both hydrodynamics and nuclear combustion. Furthermore, though supernova rates are relatively well-constrained from optical observations at low redshift, they are poorly constrained at the high-redshifts responsible for the background.  In this respect, the numerous uncertainties surrounding the stochastic background of SNe Ia are similar to those of core-collapse SNe \\cite {buonannoetal05}.  If the inflationary energy scale does lie close to $10^{15}$ GeV, the GW emission of SNe Ia will likely prove to be an additional source of noise to the primordial gravitational wave signal -- one which will be very challenging to extract.         The authors acknowledge useful discussions with Gaurav Khanna. DF acknowledges support from NSF grant numbers PHY-0902026. RF acknowledges research support from NSF grant CNS-0959382 and AFOSR DURIP grant FA9550-10-1-0354. The software used in this work was in part developed by the DOE-supported ASC / Alliance Flash Center at the University of Chicago. This research was supported in part by the National Science Foundation through TeraGrid resources provided by the  Louisiana Optical Network Initiative under grant number TG-AST100038.    \\newcommand{\\aj}{Astronomical Journal} \\newcommand{\\apjl}{Astrophysical Journal Letters} \\newcommand{\\apjs}{Astrophysical Journal Supplement} \\newcommand{\\aap}{Astronomy and Astrophysics} \\newcommand{\\araa}{Annual Review of Astronomy and Astrophysics} \\newcommand{\\mnras}{Monthly Notices of the Royal Astronomical Society} \\newcommand{\\jcap}{Journal of Cosmology and Astroparticle Physics}"
    },
    "1112/1112.0114_arXiv.txt": {
        "abstract": "We present multi-epoch Very Long Baseline Array (VLBA) observations of V773~Tau~A, the 51-day binary subsystem in the multiple young stellar system V773~Tau. Combined with previous interferometric and radial velocity measurements, these new data enable us to improve the characterization of the physical orbit of the A subsystem. In particular, we infer updated dynamical masses for the primary and the secondary components of 1.55 $\\pm$ 0.11 \\Msun, and 1.293 $\\pm$ 0.068 \\Msun, respectively, and an updated {\\it orbital parallax} distance to the system of 135.7 $\\pm$ 3.2 pc, all consistent with previous estimates. Using the improved orbit, we can calculate the absolute coordinates of the barycenter of V773~Tau~A at each epoch of our VLBA observations, and fit for its {\\it trigonometric parallax} and proper motion. This provides a direct measurement of the distance to the system almost entirely independent of the orbit modeling. The best fit yields a distance of 129.9 $\\pm$ 3.2 pc, in good agreement (i.e.\\ within 1$\\sigma$) with the distance estimate based on the orbital fit. Taking the mean value of the {\\it orbital} and trigonometric parallaxes, we conclude that V773~Tau is located at $d$ = 132.8 $\\pm$ 2.3 pc. The accuracy of this determination is nearly one order of magnitude better than that of previous estimates.  In projection, V773~Tau and two other young stars (Hubble~4 and HDE~283572) recently observed with the VLBA are located toward the dark cloud Lynds 1495, in the central region of Taurus. These three stars appear to have similar trigonometric parallaxes, radial velocities, and proper motions, and we argue that the weighted mean and dispersion of their distances ($d$ = 131.4 pc and $\\sigma_d$ = 2.4 pc) provide a good estimate of the distance to and depth of Lynds~1495 and its associated stellar population. The radio emission from the two sources in V773~Tau~A is largely of gyrosynchrotron origin. Interestingly, both sources are observed to become typically five times brighter near periastron than near apastron (presumably because of increased flaring activity), and the separation between the radio sources near periastron appears to be systematically smaller than the separation between the stars. While this clearly indicates some interaction between the individual magnetospheres, the exact mechanisms at play are unclear because even at periastron, the separation between the stars ($\\sim$ 30 $R_*$) remain much larger than the radius of the magnetospheres around these low-mass young stars ($\\sim$ 6 $R_*$). ",
        "introduction": "In spite of significant progress in recent years, the formation and early evolution of stars are still not fully understood (e.g.\\ Hillenbrand \\& White 2004, Mathieu et al.\\ 2007). One way to foster progress is to measure as accurately as possible the intrinsic characteristics of individual young stars (luminosity, effective temperature, mass, disk properties, etc.), and compare them with the predictions of detailed theoretical models. Young binary systems are particularly interesting in this respect, because tracking their orbital motions provides a direct means to estimate their dynamical mass. In particular, if astrometric and spectroscopic data are combined, the physical orbit and the individual masses of the system members can be determined.  A recurrent obstacle to the accurate determination of the intrinsic properties of young stars have been fairly large uncertainties (typically 20 to 50\\%) in distance estimates to even the nearest star-forming regions. Significant progress has been possible in recent years thanks to direct trigonometric parallax measurements obtained using Very Long Baseline Interferometry (VLBI) multi-epoch observations. Such results have been reported, in particular, in the previous papers in this series (Loinard et al.\\ 2005, 2007, 2008; Torres et al.\\ 2007, 2009; Dzib et al.\\ 2010, 2011). They provide uncertainties of a few percent or better, that typically surpass the accuracy of previous determinations by one order of magnitude.  V773~Tau (HD~283447, HBC~367) is a young stellar system located toward the dark cloud Lynds~1495 in Taurus. V773~Tau was first identified as a T~Tauri star by Rydgren et al.\\ (1976), and was established as a visual double (components designated A and B) with an apparent separation of roughly 150 mas in high-angular resolution studies independently by Ghez et al.\\ (1993) and Leinert et al.\\ (1993). Roughly contemporaneously, Martin et al.\\ (1994) and Welty (1995) established the brighter (A) visual component as a short-period (51 days) double-lined spectroscopic binary. Duch\\^ene et al.\\ (2003; hereafter D2003) and Woitas (2003) independently identified a third visual component (designated C; note that D2003 use an alternate component notation) in the system, indicating that V773~Tau is (at least) a quadruple system.\\footnote{It should be noted that multiple nomenclatures have been used to describe the components of V773~Tau. D2003 designate the spectroscopic binary V773~Tau~A/B, and the optical and infrared companions V773~Tau~C and V773~Tau~D, respectively. Alternately White \\& Ghez (2001), Woitas (2003), and Massi et al.\\ (2008) among others, designate the spectroscopic binary as V773~Tau~A (containing stars Aa and Ab), and the two companions as B and C, respectively. A third nomenclature is used in the Double Star Catalog of Mason et al.\\ (2001), which designates the spectroscopic binary Aa, containing stars Aa1 and Aa2. In this work we will follow the notation established by White \\& Ghez (2001), reflecting the hierarchical nature of the V773~Tau architecture.}  V773~Tau~A has long been known to be a strong radio source (Kutner et al.\\ 1986).\\footnote{The two companions V773~Tau~B and C, on the other hand, are not detected at radio wavelengths at the level of sensitivity of existing observations.} Indeed, it was the strongest source in the 5 GHz VLA survey of WTTS in the Taurus-Auriga molecular cloud complex by O'Neal et al.\\ (1990). From detailed multi-frequency observations, Feigelson et al.\\ (1994) concluded that the radiation was most likely of non-thermal origin. This was confirmed by Phillips et al.\\ (1996; hereafter P1996) who obtained VLBI observations, and resolved the radio emission into a clear double source, most likely corresponding to the two components of the spectroscopic binary. More recently, Massi et al.\\ (2002, 2006) showed that the radio emission exhibits periodic variations with a period corresponding to the 51 day orbital period of the spectroscopic binary. This variability is due to an increase in the flaring activity near periastron and might reflect interactions between extended magnetic structures associated with the two stars when they get close to one another. Finally, Boden et al.\\ (2007; hereafter B2007) and Massi et al.\\ (2008) also resolved the radio emission from V773~Tau~A into two components, which they associate with the two stars in the spectroscopic binary.  The dynamics in V773~Tau has been studied by a number of authors. Relative orbital motion among the A, B, C components has been monitored by D2003, B2007 and Boden et al.\\ (2011; a companion paper to this one -- hereafter B2011). Following Welty (1995) and P1996, B2007 used radio and near-IR interferometry and spectroscopic radial velocity (RV) datasets to estimate the A subsystem physical (three-dimensional) orbit. They obtain dynamical mass estimates of 1.54 and 1.33 \\Msun\\ for the Aa and Ab components, respectively. Further, B2007 estimated the A subsystem distance by means of ``orbital parallax'' (comparing the angular and physical orbit size), yielding 136.2 $\\pm$ 3.7 pc. There is also a direct trigonometric parallax measurement based on multi-epoch VLBI observations for this source (Lestrade et al.\\ 1999). This VLBI-based distance measurement ($d$ = 148.4$^{+5.7}_{-5.3}$ pc) is roughly consistent (at the 2--3 $\\sigma$ level) with the value obtained from the A subsystem orbit modeling.  In this paper we present new VLBI observations of V773~Tau~A which resolve the A subsystem, and yield additional insights on its structure. These data are described in \\S 2, and modeled jointly with earlier observations from B2007 in \\S3 to update the A subsystem physical orbit and resulting physical parameters (component dynamical masses, orbital parallax). Then the updated A orbit model is used in conjunction with the VLBA global astrometry to compute a new trigonometric parallax to V773~Tau~A. The relevance of these new results for the distance to the dark cloud Lynds~1495 is discussed in \\S~3.3. Several interesting features of the V773~Tau~A radio emission are apparent in these new observations; the insights that they yield on the magnetospheric physics of the A subsystem components are discussed in \\S~4.2.  {\\scriptsize \\begin{deluxetable}{clcr} \\tablecaption{Project code and date for each observation.\\label{table_observations}} \\tablehead{ \\multicolumn{1}{c}{Epoch}& \\multicolumn{1}{c}{Project Code}& \\multicolumn{1}{c}{Mean UT date}& \\multicolumn{1}{c}{Julian Day}\\\\ {}& {}& {[yyyy.mm.dd ~ hh:mm]}& {}}% \\startdata 01 & BM 198 A  & 2004.03.11 ~ 20:12 & 2453076.3417 \\\\ 02 & BM 198 B  & 2004.03.12 ~ 20:12 & 2453077.3417 \\\\ 03 & BM 198 C  & 2004.03.13 ~ 20:12 & 2453078.3417 \\\\ 04 & BM 198 D  & 2004.03.14 ~ 20:12 & 2453079.3417 \\\\ 05 & BM 198 E  & 2004.03.15 ~ 20:11 & 2453080.3414 \\\\ 06 & BM 198 F  & 2004.03.16 ~ 20:11 & 2453081.3414 \\\\ 07 & BM 198 G  & 2004.03.17 ~ 20:11 & 2453082.3414 \\\\ 08 & BL 128 AA & 2005.09.08 ~ 12:01 & 2453622.0013 \\\\ 09 & BL 128 AB & 2005.11.15 ~ 07:31 & 2453689.8136 \\\\ 10 & BL 128 AC & 2006.01.21 ~ 03:11 & 2453756.6327 \\\\ 11 & BL 128 AD & 2006.04.01 ~ 22:31 & 2453827.4386 \\\\ 12 & BL 128 AE & 2006.06.12 ~ 17:48 & 2453899.2423 \\\\ 13 & BL 128 AF & 2006.09.05 ~ 12:14 & 2453984.0102 \\\\ 14 & BL 146 B  & 2007.08.23 ~ 13:06 & 2454336.0461 \\\\ 15 & BL 146 C  & 2007.08.29 ~ 12:42 & 2454342.0295 \\\\ 16 & BL 146 D  & 2007.09.05 ~ 12:15 & 2454349.0106 \\\\ 17 & BL 146 E  & 2007.09.11 ~ 11:51 & 2454354.9943 \\\\ 18 & BL 146 F  & 2007.09.16 ~ 11:32 & 2454359.9806 \\\\ 19 & BL 146 G  & 2007.09.21 ~ 11:12 & 2454364.9669 \\\\ 20 & BL 146 H  & 2007.09.27 ~ 10:48 & 2454370.9504 \\\\ 21 & BL 146 I  & 2007.10.03 ~ 10:25 & 2454376.9342 \\\\ 22 & BL 146 J  & 2007.10.09 ~ 10:01 & 2454382.9177 \\\\ 23 & BL 146 K  & 2007.10.17 ~ 09:30 & 2454390.8960 \\\\ 24 & BL 146 L  & 2007.10.23 ~ 09:06 & 2454396.8794 \\\\ 25 & BL 146 M  & 2007.10.27 ~ 08:50 & 2454400.8684 \\\\ 26 & BL 146 N  & 2007.11.17 ~ 07:28 & 2454421.8114 \\\\ 27 & BM 306    & 2009.09.27 ~ 08:07 & 2455101.8386 \\\\ \\enddata \\end{deluxetable} } ",
        "conclusions": ""
    },
    "1112/1112.0004_arXiv.txt": {
        "abstract": "Microwave Kinetic Inductance Detectors, or MKIDs, have proven to be a powerful cryogenic detector technology due to their sensitivity and the ease with which they can be multiplexed into large arrays.  A MKID is an energy sensor based on a photon-variable superconducting inductance in a lithographed microresonator, and is capable of functioning as a photon detector across the electromagnetic spectrum as well as a particle detector.  Here we describe the first successful effort to create a photon-counting, energy-resolving ultraviolet, optical, and near infrared MKID focal plane array.  These new Optical Lumped Element (OLE) MKID arrays have significant advantages over semiconductor detectors like charge coupled devices (CCDs).  They can count individual photons with essentially no false counts and determine the energy and arrival time of every photon with good quantum efficiency.  Their physical pixel size and maximum count rate is well matched with large telescopes.  These capabilities enable powerful new astrophysical instruments usable from the ground and space.  MKIDs could eventually supplant semiconductor detectors for most astronomical instrumentation, and will be useful for other disciplines such as quantum optics and biological imaging. ",
        "introduction": "Cryogenic detectors are currently the preferred technology for astronomical observations over most of the electromagnetic spectrum, notably in the far infrared through millimeter (0.1--3~mm)~\\cite{Bintley:2010p4108,Niemack:2008p4173,Carlstrom:2011p4239}, X-ray~\\cite{Kelley:2009p4337}, and gamma-ray~\\cite{Doriese:2007p4284} wavelength ranges.  In the important ultraviolet, optical, and near infrared (0.1--5~$\\mu$m) wavelength range a variety of detector technologies based on semiconductors, backed by large investment from both consumer and military customers, has resulted in detectors for astronomy with large formats, high quantum efficiency, and low readout noise.  However, these detectors are fundamentally limited by the band gap of the semiconductor (1.1 eV for silicon) and thermal noise sources from their high ($\\sim$100~K) operating temperatures~\\cite{Eisaman:2011p6692}.  Cryogenic detectors, with operating temperatures on the order of 100~mK, allow the use of superconductors with gap parameters over 1000 times lower than typical semiconductors.  This difference allows new capabilities.  A superconducting detector can count single photons with no false counts while determining the energy (to several percent or better) and arrival time (to a microsecond) of the photon.  It can also have much broader wavelength coverage since the photon energy is always much greater than the gap energy.  While a CCD is limited to about 0.3--1~$\\mu$m, the new arrays described here are sensitive from 0.1~$\\mu$m in the UV to greater than 5~$\\mu$m in the mid-IR, enabling observations at infrared wavelengths vital to understanding the high redshift universe.  This approach has been pursued in the past with two technologies, Superconducting Tunnel Junctions (STJs)~\\cite{Martin:2006p4412,Hijmering:2008p4501} and Transition Edge Sensors (TESs)~\\cite{Romani:2001p1716,Burney:2006p4521}.  While both of these technologies produced functional detectors, they are limited to single pixels or small arrays due to the lack of a credible strategy for wiring and multiplexing large numbers of detectors, although recently there have been proposals for larger TES multiplexers~\\cite{Niemack:2010p4657}.  Microwave Kinetic Inductance Detectors, or MKIDs\\cite{Day03}, are an alternative cryogenic detector technology that has proven important for millimeter wave astrophysics\\cite{Schlaerth:2010p3957,Roesch:2010p4655} due to their sensitivity and the ease with which they can be multiplexed into large arrays.  MKIDs use frequency domain multiplexing~\\cite{Mazin:2006p5} that allows thousands of pixels to be read out over a single microwave cable. While the largest STJ array is 120 pixels~\\cite{Verhoeve:2006p3383} and the largest optical TES array is 36 pixels~\\cite{Burney:2006p4521}, the MKID arrays described below are 1024 pixels, with a clear path to Megapixel arrays.  The ability to easily reach large formats is the primary advantage of MKID arrays.  In this paper we describe the first photon-counting, energy-resolving ultraviolet, optical, and near infrared MKID focal plane array.  These Optical Lumped Element (OLE) MKID arrays have significant advantages over semiconductor detectors like charge coupled devices (CCDs)~\\cite{Smith:2011p4039}.  They can count individual photons with essentially no false counts and determine the energy and arrival time of every photon with good quantum efficiency.  Their physical pixel size and maximum count rate is well matched with large telescopes.  These capabilities enable powerful new astrophysical instruments usable from the ground and space. ",
        "conclusions": ""
    },
    "1112/1112.2237_arXiv.txt": {
        "abstract": "We have discovered two compact sources of shocked H$_2$ 2.12-$\\mu$m emission coincident with \\object{Mol 160} (\\object{IRAS 23385+6053}), a massive star-forming core thought to be a precursor to an ultracompact \\ion{H}{2} region. The 2.12-$\\mu$m sources lie within 2$^{\\prime\\prime}$ (0.05 pc) of a millimeter-wavelength continuum peak where the column density is $\\ge$ 10$^{24}$ cm$^{-2}$. We estimate that the ratio of molecular hydrogen luminosity to bolometric luminosity is $>$ 0.2\\%, indicating a high ratio of mechanical to radiant luminosity. CS J=2$\\rightarrow$1 and HCO$^+$ J=1$\\rightarrow$0 observations with CARMA indicate that the protostellar molecular core has a peculiar velocity of $\\sim$ 2 km s$^{-1}$ with respect to its parent molecular cloud. We also observed 95 GHz CH$_3$OH J=8$\\rightarrow$7 Class I maser emission from several locations within the core. Comparison with previous observations of 44-GHz CH$_3$OH maser emission shows the maser sources have a high mean ratio of 95-GHz to 44-GHz intensity. Our observations strengthen the case that \\object{Mol 160} (\\object{IRAS 23385+6053}) is a rapidly accreting massive protostellar system in a very early phase of its evolution. ",
        "introduction": "Studying the earliest evolutionary phases of massive protostars is difficult, because they are rare and their formation times are brief. Mol 160/IRAS 23385+6053 (hereafter Mol 160) is one of a number of sources identified by Molinari et al. (1996, 1998a, 2000, 2002) as possible precursors to ultracompact (UC)  \\ion{H}{2} regions (Wood \\& Churchwell 1989). They were selected on the basis of their cool IRAS colors, high infrared luminosities, absence of centimeter-wavelength continuum emission, H$_2$O maser activity, ammonia line emission, millimeter continuum observations, and presence of energetic molecular outflows. Based on follow-up observations with the Owens Valley Radio Observatory (OVRO) and the Infrared Satellite Observatory (ISO), Molinari et al. (1998b, hereafter M98) suggested that Mol 160 might be the first bona fide example of a massive Class 0 protostar, an object still accreting matter at a high rate from a massive envelope.  It is generally accepted that outflows from low-mass young stellar objects (YSOs) are driven by magnetic stresses in accretion disks that eject some of the inflowing matter and carry off angular momentum (Pudritz \\& Norman 1983, 1986; Lovelace et al. 1987; K\\\"onigl 1989; Pelletier \\& Pudritz 1992; Wardle \\& K\\\"onigl 1993; Safier 1993; Paatz \\& Camenzind 1996; Shu et al. 2000). Optical and near-infrared observations often reveal jets and Herbig-Haro objects that show in detail the interaction of the underlying winds with the ambient medium.  If massive stars are built up through similar accretion processes, one might expect them to have outflows that are scaled-up versions of their low-mass counterparts. However, observationally linking jets and outflows with their sources is challenging, even in regions forming intermediate-mass stars (e.g., Wolf-Chase et al. 2003). The problems are compounded for massive stars, which often form in tight clusters, lie in heavily obscured regions at distances greater than several kiloparsecs, and may be evolving on dynamical time scales less than a few thousand years.  Observations of molecular hydrogen lines in the near infrared are particularly valuable, since they directly measure kinetic energy deposited by high-speed shocks and can be observed with large ground-based telescopes at angular resolutions similar to those that can be achieved with millimeter wavelength interferometry. In this paper, we present narrowband, 0.7$^{\\prime\\prime}$-resolution, near-infrared images of Mol 160 made with the Apache Point Observatory (APO) 3.5-m telescope and 2$^{\\prime\\prime}$-resolution 3-mm continuum and line observations with the Combined Array for Research in Millimeter-wave Astronomy (CARMA). The near-infrared measurements were made as part of an emission-line survey for shocked molecular hydrogen gas associated with massive protostars (Wolf-Chase et al. 2011, in preparation). In the Mol 160 region, we detected a number of compact molecular hydrogen emission-line sources. Surprisingly, two of the brightest lie within 2$^{\\prime\\prime}$ of a millimeter-wavelength continuum peak where the hydrogen column density is $\\ge$ 10$^{24}$ cm$^{-2}$. Our CARMA measurements include CS, a tracer of dense gas, HCO$^+$, a tracer of cloud kinematics, and a Class I maser transition of CH$_3$OH, a tracer of shock-excited gas. Our observations support the hypothesis that Mol 160 is a massive protostellar system in a very early phase of evolution, but call into question some of the specific conclusions drawn from earlier outflow studies of this object (see \\S 4.3.1).  Prior work on Mol 160 includes observations covering a wide range of wavelengths and angular resolutions, from the arcminute scales of the IRAS measurements to the arcsecond scales of interferometric observations. This can potentially lead to ambiguities in the association of source names with specific physical systems. In this paper, we follow the nomenclature generally used in the literature (e.g., Wang et al. 2011 and references therein). We use ``clumps'' to refer to parsec-scale molecular cloud structures and``cores'' to the smaller structures ($\\sim$ 0.1 pc) that may further collapse and coalesce into stars or groups of stars. We will use the term ``Mol 160'' to narrowly denote the protostar (or protostellar cluster) and the dense molecular core from which it is forming (diameter $\\le$ 6$^{\\prime\\prime}$) and ``Mol 160 region'' to refer to the arcminute-scale area that also encompasses potentially related structures including larger molecular clouds and two \\ion{H}{2} regions.  \\section {OBSERVATIONS AND DATA REDUCTION}  \\subsection{Apache Point Observatory: NICFPS}  We obtained narrowband, near-infrared images using the Near-Infrared Camera and Fabry-Perot Spectrometer (NICFPS) on the Astrophysical Research Consortium (ARC) 3.5-m telescope at APO in Sunspot, NM. NICFPS uses a 1024 $\\times$ 1024 Rockwell Hawaii HgCdTe detector with 18-$\\micron$ pixels. On the APO 3.5 m, this yields a pixel scale of 0.273$\\arcsec$ pixel$^{-1}$ and gives a field of view of 4.6$\\arcmin \\times 4.6\\arcmin$ (Vincent et al. 2003). We used the following filters: K$_s$ (500 s), H$_2$ 2.12-$\\micron$ (1800 s), H$_2$ 2.25-$\\micron$ (1200 s), and a narrow-band H$_2$-continuum filter centered on 2.13 $\\micron$ (1200 s). Because of weather and scheduling constraints, the H$_2$ 2.12-$\\micron$ and K$_s$ data were recorded on 2005 Jun 17 and the remaining data were obtained on 2007 Nov 21. Seeing conditions were good the first night ($\\sim0.7\\arcsec$ at 2.2 $\\micron$) and average the second night ($ \\sim1.3\\arcsec$ at 2.2 $\\micron$).  All images were processed using IRAF and standard reduction techniques (e.g. Joyce 1992). We dithered the telescope between each exposure in a five-point pattern, and used the median of the resulting images to subtract the sky background from each image. To produce final images in each filter, we used stars to register the reduced images and then took the median of the aligned frames.  We produced a continuum-subtracted H$_2$ 2.25-$\\micron$ image by taking the difference between H$_2$ 2.25-$\\micron$ and H$_2$-continuum images. Ideally, we would have done the same for the H$_2$ 2.12-$\\micron$ data, but it was not possible to obtain H$_2$-continuum filter data during our 2005 run. Additionally, the instrument electronics were upgraded between the two runs, which resulted in different read noise and gain levels on the chip that complicated the comparison of 2005 and 2007 data. For these reasons, we chose to use K$_s$ data taken on the same night and thus with the same electronics and seeing conditions as the H$_2$ 2.12-$\\micron$ data, scaling the K$_s$ image so that the stars disappeared when we subtracted the K$_s$ image from the H$_2$ 2.12-$\\micron$ image. Both continuum subtraction techniques worked well on all but the brightest stars, which show small residual artifacts.  We performed aperture photometry on compact sources using the DAOPHOT package in IRAF. To avoid source confusion while still maintaining good signal-to-noise, we chose an aperture radius of 8 pixels (2.2$\\arcsec$), which is roughly 3 times the FWHM of the stars in the H$_2$ 2.12-$\\micron$ images.  The data were calibrated by comparing the instrumental magnitudes of several dozen stars to their K$_s$ magnitudes in the 2MASS point source catalog. For each filter, we calculated a zero point correction between our instrumental magnitudes and the 2MASS K$_s$ magnitudes, then computed fluxes using:  $$ F_{Jy} = F_0 \\times 10^{-Kmag/2.5}  \\eqno{(1)}  $$ $$ F_{line} = F_{Jy} \\times 10^{-26} \\times d\\nu \\eqno{(2)} $$   \\noindent where F$_0 =$ 666.7 Jy and $d\\nu = c  d\\lambda/\\lambda^2$.  For the H$_2$ 2.12-$\\micron$ filter, $d\\lambda = 6.93$ nm and for the H$_2$ 2.25-$\\micron$ filter, $d\\lambda = 7.6$ nm (Apache Point Observatory 2010).  WCS coordinates were calibrated using 2MASS point sources. The 2MASS positions are accurate to within 0.1$^{\\prime\\prime}$ for K$_s$ $<$ 14 (Skrutskie et al. 2006). Given the seeing conditions ($\\sim$ 0.7$^{\\prime\\prime}$) during acquisition of the NICFPS data  and pixel size of 0.27$^{\\prime\\prime}$, we estimate that our positions are accurate to within $\\pm$ 0.3$^{\\prime\\prime}$.  \\subsection{CARMA}  We observed a field around the Mol 160 core approximately 1.5$^{\\prime}$ across with CARMA on 2010 Mar 23 and 24 during three tracks. The interferometer was in the C configuration with baselines ranging from 26-370 meters. We simultaneously observed the continuum and 3 spectral lines. At the time, CARMA was configured to observe in both upper and lower sidebands of three spectral bands, and a fourth was being integrated into the system.  For the first track, two bands were configured for maximum continuum bandwidth, covering 468 MHz each, for a total of 1.872 GHz (including both sidebands). A third observed CH$_3$OH J=8$\\rightarrow$7 at 95.169 GHz in the upper side band with a 31-MHz bandwidth and 63 channels, resulting in a velocity resolution of 1.538 km s$^{-1}$. The fourth band observed HCO$^+$ J=1$\\rightarrow$0 at 89.189 GHz in the lower side band with 62-MHz bandwidth and 383 channels, giving a resolution of 0.547 km s$^{-1}$. There were calibration problems with this band; the upper side band data were corrupted, so the fourth band was not used during the other tracks.  For the second and third tracks, we used only three bands. One observed continuum, for a total of 936-MHz bandwidth. The band observing CH$_3$OH was unchanged. The third band observed the HCO$^+$ line in the lower side band and CS J=2$\\rightarrow$1 in the upper side band at 97.981 GHz. It was set to a 62-MHz bandwidth with 63 channels, yielding 3.283 km s$^{-1}$ resolution for HCO$^+$ and 2.988 km s$^{-1}$ resolution for CS.  The passband calibrator was 3C84 for the first track, 1927$+$730 on the second track, and 3C454.3 on the third. The flux calibrator was Uranus for the first and second track and MWC349 on the third. The phase calibrator was 0102$+$584 for all tracks.  In our analysis, we combined CS and HCO$^+$ data from the second and third tracks and CH$_3$OH and continuum data from all three tracks. The total integration time on source was 2.35 hours for the second and third tracks and 4.75 hours for all three tracks. The size (FWHM) of the synthesized beam was 1.982$^{\\prime\\prime}\\times$1.811$^{\\prime\\prime}$. We reduced and calibrated the data using the MIRIAD data reduction package (Sault et al. 1995). We used an intermediate Brigg's visibility factor to balance sensitivity to compact structure in the Mol 160 core with suppression of sidelobes that might otherwise distort clump-scale emission. Coverage of the u-v plane over the multiple tracks contains numerous shorter baselines, so we expect good qualitative rendering of structures $\\le$30$^{\\prime\\prime}$ in size, though some fraction of the flux will be missing.  \\section {RESULTS}  \\subsection{H$_2$ Emission in the Mol 160 Region}  Our H$_2$ 2.12-$\\micron$ and continuum-subtracted H$_2$ 2.12-$\\micron$ images are shown in Figures 1 \\& 2, respectively. Positions and designations of the 24 $\\micron$ point sources (crosses) presented in Table 1 of Molinari et al. (2008a, hereafter M08) are plotted in Figure 2. Source ``A'' lies close to the Mol 160 molecular core. Some of the pure line emission in Figure 2 arises from fluorescently excited gas in photo-dissociation regions (PDRs) associated with two \\ion{H}{2} regions (indicated by the large ovals in Figure 1) and some from shock-excited gas.  In a catalog hosted by the Joint Astronomy Centre in Hawaii, Davis et al. (2010) have established a numbering scheme for ``molecular hydrogen emission-line objects'' or ``MHOs'', compact emission-line sources thought to be associated with outflows. We have identified 10 candidate MHOs and numbered them according to this scheme. Their positions are plotted with open diamonds in Figure 2 and identified by the last two digits of the MHO numbers.  Table 1 lists the MHO designations (column 1), their positions (columns 2 \\& 3), and their 2.12-$\\micron$ line fluxes (column 4). None of the MHOs were detected in the H$_2$ 2.25-$\\micron$ filter. Using equations (1) \\& (2) to calibrate the faintest stars in the H$_2$ 2.25-$\\micron$ image against stars also present in the 2MASS catalog, we derive a 3$\\sigma$ upper limit for H$_2$ 2.25-$\\micron$ line intensity of $9.49 \\times 10^{-19}$ W m$^{-2}$. Column 5 lists the ratios of the 2.12-$\\micron$ fluxes to the 2.25-$\\micron$ upper limit. The difference in extinction at 2.12 $\\micron$ and 2.25 $\\micron$ amounts to only a tenth of a magnitude (assuming A$_{\\lambda} \\sim {\\lambda}^{-1.7}$), with a corresponding brightness ratio of $\\sim$1.1, which would change the numbers in column 5 only slightly. The expected 2.12-$\\micron$/2.25-$\\micron$ line ratio for UV excitation (as in PDRs) is $\\sim$ 1.9, and $\\sim$ 7.7 for shocks (Black \\& van Dishoeck 1987; Gredel \\& Dalgarno 1995). Shock excitation is consistent with the observed ratios for all the MHOs and is very likely for the brightest.  Star formation is clearly occurring throughout the Mol 160 region. MHOs 2921 \\& 2922 are of particular interest because of their proximity to the Mol 160 3.2-mm continuum peak.  At the present time, there is no evidence linking any of the other MHOs to the Mol 160 core itself. MHO 2922 is extremely compact, appearing almost stellar in the 2.12-$\\micron$ image (see Figure 3a). Comparison with stellar images of comparable peak intensity yields a deconvolved source size of 0.7$^{\\prime\\prime} \\times 0.6^{\\prime\\prime}$. Assuming a distance of 4.9 kpc to Mol 160 (M98)\\footnote{Zhang et al. (2005) used a distance of 6.9 kpc, and M08 suggest the distance may be closer to 8 kpc.}, this corresponds to a physical size of $\\sim$ 3430 AU. MHO 2921 is slightly larger and more irregular in appearance, particularly at low intensity levels. It is displaced to the south of the peak position of the 3.2-mm continuum source, and there is a very dark region symmetrically displaced to the north, with a low-intensity arc of H$_2$ 2.12-$\\micron$ emission at its northern edge, as can be seen in the continuum-subtracted 2.12-$\\micron$ image (Figure 3b). Fitting a two-dimensional Gaussian to the bright central region of MHO 2921 yields a deconvolved size of  1.4$^{\\prime\\prime} \\times 0.7^{\\prime\\prime}$.  Subtracting the Gaussian fits to MHO 2921 and MHO 2922 from the continuum-subtracted H$_2$ image of Figure 3b results in the image shown in Figure 3c. The residual emission has the form of a ring of approximately the same diameter as the outer contours of the  -53.7 km s$^{-1}$ CS and -52.5 km s$^{-1}$ HCO$^+$ peaks seen in Figure 4. The ring also coincides with a number of CH$_3$OH maser sources. The only feature of comparable intensity within the ring is a patch of emission near the positions of Masers 3 and 4 (see \\S 3.2.2). In Figure 3, the relative intensities at the peaks along the ring and the peaks at MHO 2921 \\& MHO 2922 are roughly 1:4:10. Note that the intensities of the dark area and arc north of the main continuum source are the same in Figures 3b and 3c. That is, they are not artifacts of the subtraction.  \\subsection{CARMA Results}  \\subsubsection{HCO$^+$, CS, and 3.2-mm Continuum Emission}  The 3.2-mm (93.6 GHz) continuum emission peaks at the position 23$^h$40$^m$54.5$^s$, +61$^{\\circ}$10$^{\\prime}$28.1$^{\\prime\\prime}$. The beam size was 2.0$^{\\prime\\prime}$ $\\times$ 1.8$^{\\prime\\prime}$, with a position angle of 40$^{\\circ}$. A two-dimensional Gaussian fit to our continuum data yields a size of 3.3$^{\\prime\\prime}$ $\\times$ 2.6$^{\\prime\\prime}$, with a position angle of 40$^{\\circ}$. The deconvolved source size is 2.7$^{\\prime\\prime}$ $\\times$ 1.9$^{\\prime\\prime}$ with a position angle of 35$^{\\circ}$. We measured a peak intensity of 10.8 mJy beam$^{-1}$. Integrating under the Gaussian fit yields a total flux of 21.6 mJy. Our position lies within 0.4$^{\\prime\\prime}$ of the positions measured by M98 and Fontani et al. (2004, hereafter F04) at 3.4mm (88 GHz) and 3 mm ($\\sim$100 GHz), respectively. M98 measured a peak intensity of 9 mJy beam$^{-1}$ in a 4.1$^{\\prime\\prime}$ $\\times$ 3.5$^{\\prime\\prime}$ beam, a total flux of 19 mJy, and derived a deconvolved source size of 4.5$^{\\prime\\prime}$ $\\times$ 3.6$^{\\prime\\prime}$. F04 measured a peak flux of 8 mJy beam$^{-1}$ in a 2.3$^{\\prime\\prime}$ $\\times$ 1.9$^{\\prime\\prime}$ beam, a total flux of 12.4 mJy (integrated to the 3$\\sigma$ contour), and a deconvolved source size of 1.5$^{\\prime\\prime}$. Our measured flux at 3.2 mm is consistent with the SED for Mol 160 presented in Fig. 4 of M08, and does not change their conclusions regarding envelope properties or integrated bolometric luminosity (3170 L$_{\\odot}$, assuming a distance of 4.9 kpc).  Maps of the three CS and three HCO$^+$ channels containing significant signal are shown in Figure 4 and Figure 5. The bandwidths of each channel are $\\sim$ 3 km s$^{-1}$ (3.283 km s$^{-1}$ for HCO$^+$ and 2.988 km s$^{-1}$ for CS). The central frequencies of the HCO$^+$ channels lie near the boundaries of the CS channels. Figure 4 shows the data at full angular resolution. In Figure 5, the data have been smoothed using a Gaussian kernel with a radius of 4$^{\\prime\\prime}$, to emphasize structure with scale lengths between those of the full-resolution interferometer data and measurements made with single-dish telescopes. Contours of velocity-integrated CS emission are plotted on the H$_2$ image in Figure 1.  Although the interferometer data do not recover all the flux at larger scales, the data have been smoothed on scales to which CARMA is still sensitive, so the structures seen in the smoothed maps are either real emission, or real emission contaminated by side-lobe artifacts.  To test for the latter, we compared the smoothed, velocity-integrated CS emission with the CARMA PSF smoothed to the same radius. We found no correlation of CS emission features with side-lobe features, and therefore conclude that the emission features seen in the smoothed maps are real. Qualitatively, Figure 5 is more sensitive to parsec-scale ``clumps'' and Figure 4 to sub-parsec-scale ``cores''. Furthermore, we note that the CS clumps in the smoothed maps correlate well with depressions in the extended H$_2$ emission (Figure 1), as would be expected for dense clumps in the foreground of a bright PDR background.  In the smoothed maps, clump-scale emission is most evident in the CS channel centered at -50.7 km s$^{-1}$ and the HCO$^+$ channels centered at -49.2 and 52.5 km s$^{-1}$. The most prominent clump is the one containing Mol 160. Near the position of the 3.2-mm continuum peak, the velocity distribution is broader and ``bluer'' than at other positions in the clump, and the compact nature of the Mol 160 core is most evident in the CS channel centered at -53.7 km s$^{-1}$ and the HCO$^+$ channel centered at -52.5 km s$^{-1}$.  \\subsubsection{CH$_3$OH Masers}  Kurtz, Hofner, \\& \\'Alvarez (2004, hereafter KHA) detected 44-GHz CH$_3$OH Class I maser emission at four positions near the Mol 160 core with a synthesized beam of $2.01^{\\prime\\prime} \\times 1.25^{\\prime\\prime}$ with the VLA in its D configuration. We found 95.169-GHz CH$_3$OH J=8$\\rightarrow$7 emission at all the positions they reported and also at a position south of the Mol 160 core continuum peak. We obtained a copy of the 44-GHz  data (S. Kurtz, private communication), and compared it with our CARMA observations. Examination of the individual channel maps revealed 44-GHz emission at the position of our southern 95-GHz source and at slightly lower but statistically significant intensities at several other positions. The channel maps also show that the KHA Source 3 position lies between two partially resolved sources, one to the north and one to the east of the Mol 160 core continuum peak. At the 0.166 km s$^{-1}$ resolution of the 44-GHz data, many of the sources shift in position from channel to channel, particularly Source 2, which extends in velocity across 6 channels and traces a path several arcseconds long. In Table 2 we list the peak positions and velocities of all the 44-GHz sources. We have extended the nomenclature used by KHA to include the three additional masers with intensities below the cut of their survey data and to reflect the resolution of their Source 3 into two spatial components.  To more directly compare the intensities of the sources at 95 GHz and 44 GHz, we velocity-binned the KHA data to the channel resolution of the 95-GHz data and smoothed the higher angular resolution VLA data to match the CARMA beam. The resulting maps are shown in Figure 6. Table 2 compares the properties of the 95-GHz and 44-GHz sources. The positions listed in Table 2 correspond to positions of peak single-channel intensity of the 44-GHz observations, which display more structure than the 95-GHz observations because of their higher angular resolution. The large diamonds in the panels identify the four KHA positions. Small diamonds indicate the positions of the sources with intensities lying below the KHA cutoff (Sources 5, 6, and 7) and the two partially-resolved components of Source 3 (3N and 3E). The crosses in the panels show the positions of peak single-channel 95-GHz intensity. Column 7 in Table 2 lists the peak single-channel 44-GHz intensities. Columns 8 and 9 comprise, respectively, the 44-GHz intensities integrated over a velocity range corresponding to the 95-GHz channel bandwidth and the 95-GHz peak intensities at the positions of Column 1 and 2. Column 10 lists the ratios of the 95-GHz intensities to the smoothed, velocity-integrated 44-GHz intensities. ",
        "conclusions": "In their review of high mass star formation, Zinnecker and Yorke (2007) suggest four evolutionary phases. (1) Compression: the formation of dense, cold, gravitationally bound cores with masses $\\ge$ 100 M$_{\\odot}$ from a larger molecular cloud by gravo-turbulent cloud fragmentation. (2) Collapse: gravitational collapse of portions of the core into one or more quasi-hydrostatic protostellar embryos, accretion disks, and infall envelopes. (3) Accretion: accretion of matter onto protostellar objects, eventually leading to one or more massive stars. (4) Disruption: disruption of the natal core by stellar winds, outflows, ionizing radiation, or supernovae.  In practice, linking details of these proposed evolutionary categories to observations of massive star-forming regions is complicated by many factors. The total radiant luminosity of a massive protostellar core may arise from a combination of Kelvin-Helmholtz emission from quasi-static protostellar embryos, protostellar and disk accretion shocks, and shocks associated with outflows or relative motions of core components. During the earliest evolutionary phases, there could also be significant contributions from external heating sources. The relative importance of each could vary dramatically and perhaps erratically with time, depending on the detailed dynamics of what may be a complex and chaotically evolving system. The expected conversion efficiency of gas to stars is sufficiently small that the mass of a core may not change dramatically until it enters the disruption phase (see, e.g., the theoretical tracks of core mass vs. luminosity in Molinari et al. 2008b). However, its temperature, density structure, chemistry, and kinematics may undergo significant alterations throughout the compression, assembly, and accretion phases. The ratio of mechanical luminosity from outflows to total radiant luminosity may also vary in a complicated way. In a general sense, high accretion rates probably correlate with high outflow rates, but kinetic energy can be stored and released asynchronously. Additionally, outflows in massive star-forming regions may be produced through different mechanisms.  Two distinctly different types of outflows have been observed in some of the nearest massive star-forming regions. Whereas the outflow from IRAS 20126$+$4104 appears to be a scaled-up version of the disk-mediated accretion outflow scenario associated with low-mass YSOs (Caratti o Garatti et al. 2008), outflow from the Orion BN/KL complex has been attributed to the explosive disintegration of a massive star cluster about 500 years ago (e.g., Rosenthal et al. 2000; Zapata et al. 2009; Bally et al. 2011). Since most massive star-forming regions (including Mol 160) are too distant to fully resolve outflow structure with current instruments, the frequency of occurrence of these fundamentally different types of outflow events, and their effects on evolving protoclusters, remains to be seen. For example, Mol 160 is approximately ten times farther away than OMC-1. At this distance, the entire BN/KL complex, including both the ``18 km s$^{-1}$'' SiO outflow and the higher velocity gas associated with the shocked H$_2$ ``bullets'' or ``fingers''  (Allen and Burton 1993; Rosenthal et al. 2000; Colgan et al. 2007), would span $\\le$6$^{\\prime\\prime}$, similar to the size of the Mol 160 core.  Theoretical models can offer some insights but are poorly constrained by actual observations. The observational problem at hand is how detailed studies of Mol 160 and other young objects can illuminate the phenomenology of massive star formation and place constraints on possible models.  \\subsection{Environment of Mol 160}  The molecular cloud containing Mol 160 was studied by F04. They observed C$^{18}$O J=1$\\rightarrow$0, C$^{17}$O J=1$\\rightarrow$0, and C$^{17}$O J=2$\\rightarrow$1 with the IRAM 30-m telescope with 23$^{\\prime\\prime}$, 22$^{\\prime\\prime}$, and 11$^{\\prime\\prime}$ beams, respectively. At the position of the 3-mm continuum peak, they found there were two distinct kinematic components, centered at -50.5 and -47.8 km s$^{-1}$. Maps over a 40$^{\\prime\\prime} \\times 40^{\\prime\\prime}$ area showed that the -47.8 km s$^{-1}$ emission peaked south of the position of the Mol 160 continuum peak and that the peak of the -50.5 km s$^{-1}$ emission was extended in a N-S direction and displaced about 4$^{\\prime\\prime}$ to the west. They suggested that the lack of a prominent feature at the position of the Mol 160 core might be explained by depletion of molecular gas on grains in the cool, high-density core. They also mapped NH$_3$ (1,1) over the same region with the VLA, using a 4.3$^{\\prime\\prime} \\times 3.7^{\\prime\\prime}$ beam. They found the emission was extended in a N-S ridge parallel to the CO peak, but displaced only 1$^{\\prime\\prime}$-2$^{\\prime\\prime}$ to the west of Mol 160. At the declination of Mol 160, there was actually a depression in the ridge, with the peak intensity occurring $\\sim$ 5$^{\\prime\\prime}$ to the south, again suggesting depletion in the core. In NH$_3$, they saw only the -50.5 km s$^{-1}$ component, and they concluded that the interferometer was resolving out the -47.8 km s$^{-1}$ component seen in single-dish observations with the Effelsberg 100-m telescope. Their CO maps also showed that the -47.8 km s$^{-1}$ emission was more extended E-W than the -50.5 km s$^{-1}$ component. Henceforth, we will refer to these components as the -50.5 and -47.8 km s$^{-1}$ clouds.  On a larger scale, Molinari et al. (2002) observed 3.6-cm continuum emission from two extended \\ion{H}{2} regions lying to the west and east of Mol 160. They named them VLA 1 and VLA 2, respectively. Their ionization fronts almost intersect at a position just to the east of the Mol 160 core (see the large ovals on Figure 1). They appear to be bounded along portions of their ionization fronts by dense molecular clouds. In such cases, the high optical depths and densities toward the cloud result in high-surface-brightness H$_2$, polycyclic aromatic hydrocarbon (PAH), and dust emission from layers parallel to the ionization front, leading to particularly intense emission along tangent lines of sight. For the purposes of this discussion, we will broaden the conventional definition of the term PDR to include not only the zone just outside the ionization front but also the high-density zone inside the front where grains flowing into the \\ion{H}{2} region from the molecular cloud are exposed to high radiation densities and emit strongly at mid-infrared wavelengths.  Contours of velocity-integrated CS emission are plotted on the H$_2$ image in Figure 1. We smoothed the CS data using a Gaussian kernel with a 4$^{\\prime\\prime}$ radius in order to highlight emission from parsec-scale structures generally referred to as ``clumps'' (Wang et al. 2011 and references therein). Most of the flux seen in the integrated map is also present in the single velocity channel centered at -50.7 km s$^{-1}$. In the H$_2$ images, one can see dark nebulosity overlying a brighter background of fluorescent emission from PDRs excited by the stars ionizing the \\ion{H}{2} regions. A similar morphology can be seen in the 15-$\\micron$ ISO image of M98. The molecular clump containing the Mol 160 core lies in a particularly dark band whose E-W extent is comparable to the width of the -50.5 km s$^{-1}$ CO cloud. It seems likely that the -50.5 km s$^{-1}$ cloud and the Mol 160 core are associated with this dark nebulosity and that they lie on the near side of the \\ion{H}{2} regions and their associated PDRs. The CS clumps lie along the eastern side of the dark band, to the east of the CO ridge and just to the west of the intense arc of PDR emission at the western boundary of VLA 2.  M08 observed Mol 160 with MIPS on Spitzer at 24 $\\micron$ and 70 $\\micron$. The prominent arcs of emission in their 24-$\\micron$ map all have analogs in our 2.12-$\\micron$ image. The extended 24-$\\micron$ and fluorescent H$_2$ radiation probably arise from contiguous regions in the PDR, within the first few optical depths at FUV wavelengths. Our arcsecond-resolution H$_2$ data suggest that much of the flux that they attribute to embedded 24-$\\micron$ point sources (the positions for which are shown in Figure 2) could arise from knots of intense PDR emission. Separating truly point-like mid-infrared sources will require significantly better angular resolution than possible with Spitzer. The exciting stars of VLA 1 and VLA 2 may as yet be undetected.  It is likely that much of the 70-$\\micron$ emission also arises from the PDRs. The overall morphologies of the Spitzer 24-$\\micron$ and 70-$\\micron$ maps made by M08 are similar. One can plausibly ascribe most of the differences to the fact that the hotter grains that dominate the 24-$\\micron$ emission lie close to or within the \\ion{H}{2} regions, where the stellar radiation is most intense. The principal exception is the far infrared peak near the position of Mol 160. Even in that case, the relative importance of internal and external heating is uncertain. The 70-$\\micron$ flux peaks $\\sim$ 5$^{\\prime\\prime}$ northeast of the Mol 160 continuum source. The discrepancy could result from positional uncertainties, as they suggest, but it is perhaps as likely that the shift is real and a significant contribution to the peak intensity in the 22$^{\\prime\\prime}$ Spitzer beam comes from nearby PDR emission, as is the case in arcminute-scale observations of the Trapezium/BNKL region in Orion (Harper 1974). Because the column density of the  Mol 160 core is sufficiently high to absorb mid-infrared flux from the PDR, there could also be a significant external contribution to the heating of the core. For these reasons, the value M08 derive for the radiant luminosity of the Mol 160 core should be considered an upper limit.  \\subsection{The Mol 160 Core}  F04 have summarized estimates of the mass, column density, volume density, and temperature of the Mol 160 core derived from both line and continuum data taken with single dish telescopes and interferometers at angular resolutions between 1.5$^{\\prime\\prime}$ and 18$^{\\prime\\prime}$. For a region with diameter of $\\sim$4$^{\\prime\\prime}$-8$^{\\prime\\prime}$ and an assumed distance of 4.9 kpc, they find  a kinetic temperature of 26 K, masses of $\\sim$100-400 M$_{\\odot}$, H$_2$ column densities of $\\sim$ 1.5-4 $\\times$10$^{24}$ cm$^{-2}$, and space densities of $\\sim$ 3.4-16 $\\times$10$^{6}$ cm$^{-3}$, with higher numbers from continuum data than from molecular lines. They note that there is evidence that molecules are depleted onto grains in the densest part of the core and suggest that the higher values should be preferred. For a region 1.3$^{\\prime\\prime}$-1.9$^{\\prime\\prime}$ in diameter ($\\sim$ 0.03 pc), they derive a kinetic temperature of 42 K, masses of 15-150 M$_{\\odot}$, H$_2$ column densities of $\\sim$3.3-15.0 $\\times$ 10$^{24}$ cm$^{-2}$, and space densities of $\\sim$ 1.7-6.0 $\\times$10$^{7}$ cm$^{-3}$, again with the highest numbers derived from continuum data. From the source SED, they estimate a temperature of 40 K. M08 derived a similar temperature of 37 K by using the automatic SED fitting tool provided by Robitaille et al. (2007), and a bolometric luminosity of 3170 L$_{\\odot}$.  They showed that these values place Mol 160 firmly into the protostar category on the evolutionary diagram developed by Molinari et al. (2008b). We note that their estimate of the bolometric luminosity of Mol 160 should be considered an upper limit, due to emission from the nearby, overlapping PDRs. Taken together, these observations and models indicate Mol 160 is in an evolutionary stage that is a precursor to a hot core.  In spite of the uncertainties, it seems reasonable to conclude that the mass of the core is sufficient to form one or more massive stars. Furthermore, the high column densities suggest that massive stars {\\it should} form once the core begins to collapse. Recent theoretical work predicts a mass column density threshold $\\ge$ 1 g cm$^{-2}$ for massive star formation (Krumholz \\& McKee 2008; Krumholz et al. 2010). Regions where intermediate mass stars are thought to be forming typically have mass column densities $\\sim$ 0.1 - 0.5 g cm$^{-2}$ (e.g., Arvidsson et al. 2010; Wolf-Chase et al. 2003; Wolf-Chase, Walker, \\& Lada 1995). In the Arvidsson et al. (2010) study, estimates were made from peak column densities derived from single-antenna millimeter-wave observations of regions at different distances, implying a `clump' rather than `core' scale for many of these regions; nevertheless, the derived values are significantly smaller than those obtained for UC \\ion{H}{2} regions at comparable distances. In the Wolf-Chase et al. (1995, 2003) studies, the 2264 S1 core in the Mon OB1 dark cloud has a mass column density of  $\\sim$ 0.27 g cm$^{-2}$ on $\\sim$ 0.1 pc scale. Spitzer observations indicate 2264 S1 contains at least 10 low-mass protostars (Young et al. 2006; Teixeira et al. 2006). Taking 4.0 $\\times$ 10$^{24}$ cm$^{-2}$ as a lower limit for the column density of the compact Mol 160 core (see Tables 5 \\& 6 in F04), the mass column density is  $\\ge$ 16 g cm$^{-2}$.  M98 noted that the Mol 160 core was slightly elongated to the southeast and to the north at 3.4 mm. Using a $0.94^{\\prime\\prime} \\times 0.76^{\\prime\\prime}$ beam at a wavelength of 1.3 mm, F04 resolved the millimeter continuum source into two components. The principal component was centered at the position previously determined by M98. The flux density per beam and integrated flux density of the smaller source were, respectively, 50\\% and 8\\% of those of the primary.  Our 3.2-mm continuum observations do not resolve these two components, but are slightly extended in the direction of the secondary peak (see contours on Figure 3c). Also, the position of the secondary source coincides with the position of MHO 2922. If the mass within a 1.5$^{\\prime\\prime}$ diameter region centered on the primary continuum peak is $\\sim$ 150 M$_{\\odot}$, as F04 suggest, and if the mass scales with the 1.3-mm continuum flux, the mass associated with the secondary peak would be $\\sim$ 12 M$_{\\odot}$.  The Mol 160 region was not covered in the Spitzer GLIMPSE survey (Benjamin et al. 2003; Churchwell et al. 2009), but has been observed by WISE. Figure 7 is a 3-color WISE image (3.4, 4.6, \\& 12-$\\micron$ bands) of the Mol 160 region. As in the color scheme chosen for the GLIMPSE survey, objects that are unusually bright in the 4.6-$\\mu$m band appear to be green. In the GLIMPSE data, ``extended green objects'' or ``EGOs'' have been linked to shocked H$_2$ emission from massive outflow candidates (Cyganowski et al. 2008). MHOs 2921 \\& 2922 are the only MHOs listed in Table 1 that are detected as EGOs in the WISE image. Furthermore, Mol 160 is {\\it only} detected in the WISE 4.6-$\\mu$m band and was not seen in the continuum by ISO at 15 $\\mu$m (M98). Hence it is likely, as for the EGOs, that the 4.6-$\\mu$m emission is dominated by spectral lines from shocked molecular gas.  \\subsection{Outflow from Mol 160}  Observations of millimeter-wave, H$_2$, and CH$_3$OH maser emission independently trace outflowing gas from Mol 160. Taken together, these observations point to the extreme youth of Mol 160 and suggest new avenues to explore in establishing an evolutionary sequence for massive star formation.  \\subsubsection{Millimeter-wave Spectroscopy}  M98 and Molinari et al. (2002) presented evidence for energetic molecular outflows in Mol 160. The wings of their SiO spectrum at the position of Mol 160 form a broad plateau in velocity space. The wings have a relatively large amplitude compared to the intensity in the central frequency channel, and the intensity drops sharply at the edges of the plateau. The HCO$^+$ wings are much smaller, compared to the central peak. The blue wing tapers down smoothly from the systemic peak within $\\sim$ ~10 km s$^{-1}$. The red wing is broader, with a dip at -46 km s$^{-1}$ and a peak at -43 km s$^{-1}$. The blue and red HCO$^+$ outflow maps of M98 peak $\\sim$ 1$^{\\prime\\prime}$ to the east of the continuum peak, and the blue peak (integrated from -63 km s$^{-1}$ to -53 km s$^{-1}$) is much stronger and broader than the barely detected red peak (integrated from -47 km s$^{-1}$ to -37 km s$^{-1}$). The frequency range over which they integrated to make their ``blue'' outflow map overlaps a third of our -52.5 HCO$^+$ channel and nearly all of our 53.7 km s$^{-1}$ CS channel, so it is likely that their map is dominated not by outflow gas but by the blue wing of the systemic core gas. On the other hand, the morphologies of their maps of the SiO wings are similar to each other. The peak of the blue lobe lies $\\sim$ 0.5$^{\\prime\\prime}$ south of the continuum peak, and the red peak lies $\\sim$ 1$^{\\prime\\prime}$ to the north. Based on the small angular separation between the lobes, M98 argued that the outflow is essentially parallel to the line of sight, ruling out the possibility that their failure to detect 15-$\\mu$m emission with ISO was caused by geometrical effects such as an edge-on dust disk and thus implying that the flow is completely contained within the optically thick core.  In our HCO$^+$ data, we do see a NE-SW velocity gradient across the core (cf. Figure 4 and the color image in Figure 8). The channel maps of the 1.5 km s$^{-1}$ data from Track 1 show the same trend, but the instrumental problems with the correlator prevent a more precise estimate of the velocity dispersion. One can set a rough upper limit of $\\sim$6 km s$^{-1}$ from the span between the central velocities of our -49.2 and -55.8 km s$^{-1}$ HCO$^+$ channels. This is much less than the $ >$20 km s$^{-1}$ extent of the SiO plateau and could plausibly result from systematic motions established during the compression/collapse phase of core evolution rather than protostellar outflows. We suggest that kinematics of the HCO$^+$ and CS gas can best be understood in terms of a model in which the compact, high-surface-brightness molecular core associated with the Mol 160 continuum source is blueshifted by  $\\sim$2 km s$^{-1}$ from a parent clump with systemic velocity of -50.5 km s$^{-1}$. Most of the flux from the compact core lies within the -52.5 km s$^{-1}$ HCO$^+$ channel and -53.7 km s$^{-1}$ CS channel. The source brightness distributions in these two channels are similar and their outer contours are roughly congruent with the low-intensity 2.12-$\\micron$ halo around MHO 2921 and MHO 2922 seen in Figure 3. The observed blueshift of the core gas with respect to the -50.5 km s$^{-1}$ cloud would be consistent with compression by a background \\ion{H}{2} region.  Our observations are consistent with the assertion by M98 that the extent of the SiO outflow is small compared to the size of the core but suggest an alternative possibility for the geometry. The symmetrical displacements of MHO 2921 to the south of the continuum peak position and an opposing dark area to the north is reminiscent of the simulated near-infrared images of wide-angle outflows constructed by Zhang \\& Tan (2011). See, for example, the example with an outflow inclination of 60$^{\\circ}$ in their Figure 9. Their models assumed the source of the near-infrared radiation from the arc was continuum radiation from the protostar but would also be consistent with an intense source of line emission located close to the center of the core. Such a scenario would imply that H$_2$ emission associated with the red SiO outflow lobe is hidden by dust in the waist of the accretion envelope (and/or a disk or pseudodisk).  From their SiO observations, M98 deduced an outflow mass of $\\sim$ 20 M$_{\\odot}$, a dynamical timescale of $\\le 6 \\times 10^3$ yr, a mass loss rate of 3.5 $\\times 10^{-3}$ M$_{\\odot}$ yr$^{-1}$ and a mechanical luminosity of 22 L$_{\\odot}$. They based these estimates on the assumptions that the outflow angle was 30$^{\\circ}$ with respect to the line of sight and that the extent of the outflow was equal to the radius of the molecular core. In light of the small angular displacement of MHO 2921 from the continuum peak, the mechanical luminosity and mass loss rate could be larger and the time scale shorter, indicating an even younger, more powerful outflow.  \\subsubsection{H$_2$ Luminosity}  Until spectra of MHOs 2921 \\& 2922 are available, we can only estimate the effects of extinction. Assuming that the 2.12-$\\micron$ flux is $\\sim$ 5\\% of the total ro-vibrational H$_2$ emission (e.g., Caratti o Garatti 2006, 2008), the combined H$_2$ luminosity for the two MHOs is $3.4 \\times 10^{-3} \\times$ (D/4.9 kpc)$^2 \\times 10^{0.4 A_{2.12}}$ L$_{\\odot}$, where D is the distance in kpc and $A_{2.12}$ is the extinction at 2.12 $\\micron$. Because of extinction, the actual luminosities of the 2.12-$\\micron$ sources are probably much greater than the observed luminosities. The average column densities, N(H$_2$),  of the large -50.5 km s$^{-1}$ cloud inferred from large-beam measurements with single-dish telescopes range from 1$\\times 10^{23}$ to 3.6$\\times 10^{23}$ cm$^{-2}$. The values derived from interferometric measurements at higher angular resolution are larger by an order of magnitude or more. Conservatively assuming a column density of N(H$_2$) = 5$\\times 10^{22}$ cm$^{-2}$, an amount equal to half the lowest value derived from large-beam measurements, leads to a 2.12-$\\micron$ extinction of 5.1 and a combined H$_2$ luminosity of 7.4 L$_{\\odot}$ for MHO 2921 and MHO 2922. This is 0.23\\% of the bolometric luminosity inferred from the SED of the continuum source (L$_{bol} = 3170$ L$_{\\odot}$, M08) and is comparable to the mechanical luminosity inferred by M98 from the high-velocity SiO outflow.  Caratti o Garatti (2006) found a relationship between total H$_2$ luminosity and L$_{bol}$  for low-mass YSOs. Caratti o Garatti (2008) suggested that this relationship could tentatively be extended to high-mass YSOs based on their results for IRAS 20126$+$4104. They further suggested that this relationship might be taken as evidence that outflows from massive YSOs are similar in nature to those from lower-mass YSOs (i.e., produced by disk-mediated accretion). However, the Orion BN/KL complex, which has a bolometric luminosity of 10$^5$ L$_{\\odot}$, and a total H$_2$ luminosity of 120 L$_{\\odot}$ (Rosenthal et al. 2000), would also fit their relationship. In this kind of impulsive event there is no reason to expect strict proportionality between the bolometric luminosity of the source and either the kinetic energy of the outflow or its rate of dissipation in shocks. Therefore, caution must be exercised in attempting to draw conclusions about the nature of outflows from massive YSOs from the L$_{H_2}$/L$_{bol}$ relationship alone. We note that the L$_{H_2}$/L$_{bol}$ ratio for Mol 160 conservatively places Mol 160 a factor of two above the L$_{H_2}$/L$_{bol}$ relationship determined by Caratti o Garatti (2006, 2008) and the discrepancy could be larger if (a) the internal heating is smaller; (b)  the extinction has been underestimated; or (c) the ratio of total H$_2$ to H$_{2.12}$ luminosity has been underestimated (e.g., if there is significant emission from cooler shocked hydrogen, cf. Caratti o Garatti 2006, 2008).  As noted in \\S 4.3.1, MHO 2922 is coincident with a secondary continuum peak (F04). It is possible that MHO 2922 emission arises from a distinct outflow from the central source, a secondary accretion center, or an object ejected from the core. Present observations do not allow us to distinguish between the possibilities; however, we note that this does not change our conclusion that Mol 160 has a very high L$_{H_2}$/L$_{bol}$ ratio, since this ratio was computed based on the estimated bolometric luminosity for the entire core, and our computed ratio is a conservative lower limit.  \\subsubsection{CH$_3$OH Masers}  Class I CH$_3$OH masers are frequently and perhaps exclusively found near young sources. They are well-correlated with molecular outflows in massive star forming regions and are thought to be collisionally pumped (Cyganowski et al. 2009 and references therein; Fontani et al. 2010). Schnee \\& Carpenter (2009) found a strong correlation between the presence of compact 3-mm continuum emission and 95-GHz Class I CH$_3$OH maser emission. In contrast, they detected no 3-mm continuum emission toward UC \\ion{H}{2} regions lacking maser emission, suggesting that the masers are signposts of an early stage in the evolution of a massive protostar before an expanding UC \\ion{H}{2} region has destroyed the accretion disk.  DeBuizer et al. (2009) note that the relatively low velocities of CH$_3$OH masers, combined with their occasional locations slightly offset from the outflow axis, suggest they arise in ``systemic'' gas in outflow cavity walls. The velocities of the Mol 160 masers vary from -49.0 to 52.2 km s$^{-1}$, and the most intense emission comes from sources nearest the continuum peaks and at velocities near $\\sim$ 52 km s$^{-1}$. The maser positions seem to lie along two principal axes (see Figure 8), one approximately north-south through the main continuum source and one roughly along a line between the main and secondary continuum sources. It is tempting, but probably premature, to associate them with the walls of a wide-angle outflow cavity. The source could be significantly more complex than a single central source and a single bipolar outflow. The maser sources are as yet unresolved, spectrally or spatially, but they are bright enough to permit observations at higher spectral and spatial resolution that could significantly improve our understanding of the structure, kinematics, and excitation of the outflows.  In their survey of CH$_3$OH masers, Fontani et al. (2010) found that sources for which both 44-GHz and 95-GHz Class I masers were observed, had similar spectra, confirming a common physical origin. They also concluded that the 95-GHz line is intrinsically fainter, based on their detection rates. The mean value of the ratio of 95-GHz to 44-GHz line intensity in Table 2 is 4.4. Some of the individual sources studied by Fontani et al. (2010) also have intensity ratios this large, suggesting that their detection statistics must result not from the intrinsic line formation process but from the probability of finding sources with conditions favoring the 95-GHz transition. Perhaps the ratio will prove to be a useful diagnostic for the shortest-lived (e.g., very early) phases of massive star formation.  \\subsection{Mol 160: Current Status \\& Future Explorations}  Taken together, the core and outflow properties of Mol 160 strongly suggest that this object is in an early phase of the 3rd stage (accretion) outlined for massive star formation at the beginning of this section (Zinnecker \\& Yorke 2007). The temperature, mass, mass column density, and luminosity of Mol 160 are consistent with physical properties expected for a precursor to a hot core, before winds and radiation from embedded protostars have removed a significant amount of the accreting envelope, and while the protostar mass and luminosity are still increasing. The presence of a very compact outflow, which is independently confirmed via millimeter-wave spectroscopy, H$_2$ shocks, and CH$_3$OH masers, also points to the extreme youth of one or more embedded, accreting, massive objects.  Because we expect the youngest massive protostars to be rare, more refined selection methods are highly desirable. A high value of L$_{H_2}$/L$_{bol}$ may be a useful discriminant, but the example of Mol 160 also highlights some of the difficulties. Near- and mid-infrared spectroscopy are needed for more accurate estimates of extinction and L$_{H_2}$. It is also likely that massive accreting objects will be located in or near regions containing other products of massive star formation such as bright PDRs. For these sources, bolometric luminosities can easily be overestimated. High-angular-resolution, multicolor, far-infrared observations with Herschel or the Stratospheric Observatory for Infrared Astronomy (SOFIA) will help reduce confusion in complex regions. Systematic searches in infrared dark clouds more clearly separated from PDRs (e.g., for compact Herschel sources and GLIMPSE ``green objects'' more pointlike than EGOs) may also be helpful. The Milky Way Project is producing a new catalog of PDRs in the galactic plane using citizen scientist identifications from Spitzer GLIMPSE/MIPSGAL images (Simpson et al. 2011) that should be useful in this regard. High-angular-resolution observations of CH$_3$OH Class I maser line ratios may also prove to be a useful diagnostic for physical conditions characteristic of very young sources, one that will become even more valuable with the deployment of advanced receivers on CARMA and as the Atacama Large Millimeter/submillimeter Array  (ALMA) comes on line."
    },
    "1112/1112.1658_arXiv.txt": {
        "abstract": "We present a new version of our code for modeling the atmospheric circulation on gaseous exoplanets, now employing a ``double-gray\" radiative transfer scheme, which self-consistently solves for fluxes and heating throughout the atmosphere, including the emerging (observable) infrared flux.  We separate the radiation into infrared and optical components, each with its own absorption coefficient, and solve standard two-stream radiative transfer equations.  We use a constant optical absorption coefficient, while the infrared coefficient can scale as a powerlaw with pressure.  Here we describe our new code in detail and demonstrate its utility by presenting a generic hot Jupiter model.  We discuss issues related to modeling the deepest pressures of the atmosphere and describe our use of the diffusion approximation for radiative fluxes at high optical depths.  In addition, we present new models using a simple form for magnetic drag on the atmosphere.  We calculate emitted thermal phase curves and find that our drag-free model has the brightest region of the atmosphere offset by $\\sim$12\\degrees~from the substellar point and a minimum flux that is 17\\% of the maximum, while the model with the strongest magnetic drag has an offset of only $\\sim$2\\degrees~and a ratio of 13\\%.  Finally, we calculate rates of numerical loss of kinetic energy at $\\sim$15\\% for every model except for our strong-drag model, where there is no measurable loss; we speculate that this is due to the much decreased wind speeds in that model. ",
        "introduction": "It has been almost one decade since the first atmospheric measurement of a hot Jupiter \\citep{Charbonneau2002} and yet this class of exotic exoplanet still provides us with many mysteries waiting to be solved.  These include the culprit responsible for the stratospheric temperature inversions inferred for many hot Jupiters \\citep{Hubeny2003,Fortney2008,Spiegel2009,Zahnle2009,Knutson2010}, the implication of a super-solar carbon-to-oxygen ratio in at least one hot Jupiter \\citep{Madhu2011}, the possible influence of magnetic drag on the atmospheric circulation \\citep{Perna2010a}, and a peculiar planet where dawn seems to be much hotter than noon \\citep{Crossfield2010}.  In the midst of a growing collection of observations, some with unexpected or surprising results, is the continuing development of atmospheric models trying to interpret and understand the measurements.  One set of atmospheric models are numerical ones that simulate the global circulation patterns on close-in gas giants.  With masses and radii comparable to Jupiter, but subject to incident stellar fluxes 10,000 times stronger than what Jupiter receives from the Sun, and expected to be tidally locked into synchronous orbits, hot Jupiters exist in an atmospheric regime unlike anything in our solar system, and models of their atmospheric circulation are expanding into uncharted territory.  In order to try to understand how atmospheres work in this new regime, the models that have been developed so far represent a range of complexities, various approaches, and the use of different sets of assumptions.  Some of the simplest models solve the shallow water or equivalent barotropic equations and use various schemes to include the effect of radiative heating \\citep{Cho2003,Cho2008,Langton2007}.  Others solve the primitive equations of meteorology, either using a Newtonian relaxation scheme for the radiative forcing \\citep{Showman2002,Cooper2005,Showman2008,MR09,RM10}, a dual-band radiative transfer scheme \\citep[][similar to the one we present here]{Heng2011}, or more complex non-gray radiative transfer \\citep{Showman2009}.  Finally, there are also models that solve the full set of fluid equations, using dual-band flux-limited diffusion for the radiative transfer \\citep{DobbsDixon2008,DobbsDixon2010}.  Here we present an updated version of our previous general circulation model that now includes a ``double-gray\" radiative transfer scheme.  Fluxes throughout the atmosphere are separated into optical and infrared components, each with its own absorption coefficient, which is constant for the optical band and can scale as a powerlaw with pressure for the infrared band.  We use standard two-stream radiative transfer equations to solve for the vertical fluxes and calculate heating rates from those.  In addition to producing self-consistent radiative fluxes and heating rates, including the observable infrared flux emerging from the top boundary, this new code has the advantage of maintaining only a moderate level of complexity.  This facilitates comparison between simulated results and analytic profiles \\citep[][see also Hansen 2008]{Guillot2010} and makes it easier to clearly identify the effects due to changes in opacity.  Our new code is described in detail in Section~\\ref{sec:code}.  We demonstrate the functionality of our code by presenting a model of a generic hot Jupiter (Section~\\ref{sec:models}) and we investigate issues related the atmosphere's behavior at deep pressures (Section~\\ref{sec:deep}).  We then study the effect of magnetic drag on atmospheric circulation by applying a simplified drag scheme to our model, with drag strengths ranging from weak to strong (Section~\\ref{sec:drag}).  In Section~\\ref{sec:conc} we summarize our results. ",
        "conclusions": "\\label{sec:conc}  We have presented a new radiative transfer scheme for our atmospheric circulation code.  It divides all radiation into optical and infrared wavelengths.  The optical flux is incident on the top boundary and its attenuation is controlled by a constant optical absorption coefficient.  The infrared flux is absorbed and emitted at each level, as governed by an infrared absorption coefficient that goes as a powerlaw with pressure (and can be constant).  Below the infrared photosphere the infrared fluxes are calculated using the flux-limited diffusion approximation, which we found was necessary in order to reproduce analytic temperature profiles correctly.  Using this new code, we present a fiducial model for a generic hot Jupiter and find that the temperature and wind structures agree well with previously published models.  We also show vertical velocities, radiative fluxes, and heating/cooling rates as a function of pressure and region of the atmosphere.  Our radiative scheme allows us to self-consistently predict the infrared flux emitted by the planet and we map this as a function of location on the globe, reproducing the standard shift of the brightest region eastward of the substellar point.  We show temperature profiles for models that: 1) did not use the flux diffusion scheme, or 2) assumed zero flux from the interior.  In both cases the deepest pressure levels are mostly isothermal, instead of increasing in temperature toward the inner convective zone, significantly changing the ability of the atmosphere to exchange energy and momentum with the interior.  However, the upper atmosphere (including and above the infrared photosphere) are relatively insensitive to the behavior in the deepest levels and have similar observable properties to our fiducial model.  We present models that use our new radiative transfer scheme in combination with a simplified form of magnetic drag, as an improvement on the models presented in \\citet{Perna2010a}.  As in those earlier models, we find that as the drag strength is increased, the structure of the atmosphere changes; the strongest level of drag is able to significantly alter the atmosphere from its drag-free state.  As a new result, we are able to measure the amount by which the brightest region of the atmosphere is offset from the substellar point, as a function of the strength of the magnetic drag used in each model.  The weak- and medium-drag models have orbital thermal phase curves similar to the drag-free model.  The winds are slowed enough in the strong-drag model that the brightest region of the atmosphere is closely aligned with the substellar point and, related to this effect, the day-night flux amplitude in the strong-drag model is slightly larger than for the drag-free model.  We estimate the amount of ohmic heating that would be produced for each of our models and find that the medium-drag model should have a significantly inflated planetary radius, while the amount of heating in the strong-drag model is so high that it could perhaps lead to planet evaporation.  A better analysis of magnetic effects on hot Jupiter atmospheres must wait for improvements in our code, however, using a more realistic form for the drag and explicitly including the effect of ohmic heating.  Finally, we estimate the rate of numerical loss of kinetic energy in each of our models and find it to be at the level of $\\sim$15\\%, except for our model with strong magnetic drag, which we calculate to have losses consistent with zero at the percent level.  Our strong-drag model has much slower winds than the other models presented here and we speculate that strong numerical loss is directly related to those high speed winds."
    },
    "1112/1112.3970_arXiv.txt": {
        "abstract": "We use a sample of active galaxies from the Cosmic Evolution Survey to show that host galaxy morphology is tied to the accretion rate and X-ray obscuration of its active galactic nucleus (AGN). Unobscured and rapidly accreting broad-line AGNs are more likely to be in spheroid-dominated hosts than weak or obscured AGNs, and obscured AGNs are more likely to have disturbed host galaxies.  Much of the disagreement in previous work on the AGN-merger connection is likely due to each study probing AGNs with different obscuration and accretion properties.  Instead, only obscured AGNs seem to merger-driven, while weak AGNs are fed by stochastic processes in disks and rapidly-accreting broad-line AGNs require massive bulges. Our observed ``unified model'' for AGN hosts fits with theoretical models for merger-driven AGN evolution, but is also consistent with steady-state AGN activity. ",
        "introduction": "The most popular theoretical framework for coevolving active galactic nuclei (AGNs) and galaxies invokes major mergers to fuel both starbursts and quasars \\citep{san88,hop06}.  But there also exist secular processes which can grow AGNs in disks, through stochastic gas accretion and stellar mass loss \\citep{hh06} or cold streams and disk instabilities \\citep{bou11}.  The relative dominance of these fueling modes may evolve: perhaps secular processes are efficient in the gas-rich $z>0.5$ universe but mergers are required to funnel gas to galaxy nuclei at $z \\sim 0$, in a way analogous to star-forming galaxies \\citep[e.g.][]{gen10}.  Observations of AGN host galaxy morphologies have the potential to distinguish secular processes from merger fueling.  Figure \\ref{fig:mergerfrac} shows examples of host morphology studies from X-ray AGNs at three different redshifts: $z \\sim 0.03$ \\citep{koss10}, $z \\sim 0.75$ \\citep{cis11}, and $z \\sim 1.8$ \\citep{koc11}.  In each survey, AGNs and mass-matched inactive galaxies are blindly classified by visual inspection.  Naively comparing the three results suggests evolution in the importance of merger fueling for AGNs, but each study is biased to selecting different kinds of AGNs.  \\begin{figure}[t] \\begin{center} \\scalebox{0.8} {\\plotone{jrtrump_fig1.eps}} \\end{center} \\caption{The fraction of AGNs and inactive galaxies which are disturbed, measured from X-ray selected samples at three different redshifts: $z \\sim 0.03$ \\citep{koss10}, $z \\sim 0.75$ \\citep{cis11}, and $z \\sim 1.8$ \\citep{koc11}.  At first glance, it appears that the AGN-merger connection evolves, with mergers driving AGNs locally but not in the past.  However the apparent trend is probably caused by differences in the AGN samples rather than true evolution. \\label{fig:mergerfrac}} \\end{figure}  Contrary to the classical unified model \\citep{ant93}, recent work has shown that rapidly accreting quasars have very different fueling and feedback modes from weak AGNs \\citep{ho08,tru11,ant11}.  In particular, \\citet{tru11} shows that broad-line quasars have strong radiative winds and luminous accretion disks, while weakly accreting narrow-line AGNs have radiatively inefficient accretion flows and radio-mode feedback.  Understanding AGN-galaxy coevolution then requires studying host galaxy type across the two axes of AGN accretion rate and obscuration.  Yet previous AGN host studies (including those shown in Figure \\ref{fig:mergerfrac}) have generally not distinguished between weakly accreting LINERs, Compton-thick AGNs, and powerful quasars.  Here we study host galaxy morphologies for 70 X-ray selected AGNs from the Cosmic Evolution Survey \\citep[COSMOS,][]{sco07}.  These AGNs span three orders of magnitude in each of accretion rate and column density, and their observed data demonstrate that host galaxy type is connected to AGN properties. ",
        "conclusions": ""
    },
    "1112/1112.0138_arXiv.txt": {
        "abstract": "{The distance to an isolated dark globule is often unknown and yet crucial for understanding its properties, in particular its mass. A new approach to this problem is discussed} {The purpose of the present paper is to investigate how well the distances of more or less reddened field stars can be determined by using multi-colour imaging.} {We observed a test globule, B\\,335 in U, B, g, r, and I, and together with the 2MASS survey, this data set gives a well-defined spectral energy distribution (SED) of a large number of stars. The SED of each star depends on the interstellar extinction, the distance to the star, and its intrinsic SED. As we had good reasons to suspect that the wavelength dependence of the extinction (the reddening) changes from the outskirts of the globule to the central parts, we did not assume any specific reddening law. Instead, we use a scheme that allows independent determination of the extinction in each line of sight as determined by groups of adjacent stars. The method is based on the use of stellar atmospheric models to represent the intrinsic SEDs of the stars. Formally, it is then possible to determine the spectral class of each star and thereby its distance. For some of the stars we have optical spectra,  allowing us to compare the photometric classification to the spectrometric. } {As expected, the main problem is that there are few stars found within each distance bin for the small field size defining a typical dark globule. However, the characterisation of the extinction and photometric classification give consistent results and we can identify one star at the front side of the globule. It has a photometric distance of 90\\,pc. The closest star behind the B\\,335 globule has a distance of only $\\approx 120$\\,pc and we therefore determine the distance to B\\,335 as 90--120\\,pc.  Our deep U image shows a relatively bright south-western rim of the globule, and we investigate whether it might be due to a local enhancement of the radiation field. A candidate source, located 1.5 arcminutes outside our field, would be the field star, HD\\,184982. This star has an entry in the Hipparcos Catalogue and its distance is 140--200\\,pc. However, we come to the conclusion that the bright SW rim is more likely due to the wing of the point-spread-function (PSF) of this star.} {} ",
        "introduction": "It is in principle easy to determine the distance to a dark cloud: it is just a matter of finding a star  just in front of the cloud and another  closely behind it. Or, even better, a star close enough to the cloud to give rise to a reflection nebula. The closest cloud complexes, like that in Taurus/Aurigae, cover large regions on the sky and thanks to the Hipparcos mission (\\cite{Hipparcos}) there are a number of stars with known distances within the field. The remaining task is that of placing the stars in front of or behind the cloud. The traditional way is to look for the onset of interstellar reddening (see Lombardi et al. \\cite{Lombardi} for a recent example). As an attractive alternative, the onset of polarisation can be determined (Alves \\& Franco \\cite{Alves}). For small isolated clouds, however, there are too few (if any)  Hipparcos stars within the field, and the distances to the stars must be determined by spectroscopic/photometric means.  The reason why it is important to determine the distances to such clouds, including dark globules, is their role of hosting isolated star formation. Thus,  important properties as the luminosity of the forming star and the mass and the radius of the cloud cannot be determined without knowing the distance.  Maheswar \\& Bhatt (\\cite{Maheswar}) have estimated the distances to nine dark globules using V, R and I photometry in combination with the J, H and K photometry from the 2MASS catalogue (Skrutskie et al. \\cite{2MASS}). By assuming a standard interstellar extinction law (valid for \\emph {diffuse} clouds) and adjusting visual extinction, A$_V$, until the SED fits that of a \\emph{main sequence} (MS) star, these authors determine the extinction and the distance for each star. There are two uncertain assumptions in their approach, which may have strong impact on the estimates of the distances. The first is the assumption of a standard extinction law, while it is known that the $R_{V}$ value ($R_{V}$=A$_{V}$/E(B-V)) changes from typically $R_{V}$=3.1 in the diffuse ISM to $R_{V}\\approx 5$ in molecular clouds (see e g Ossenkopf \\& Henning \\cite{Ossenkopf}, Strafella et al \\cite{Strafella}). The second is the assumption that all the stars are MS stars. In the present paper we discuss these problems and present a more elaborate method which is based on imaging in eight filters.  As a test case we use the well studied dark globule B\\,335. It hosts a protostar heavily obscured by a flattened core structure, including an accretion disk seen edge-on (G\\aa lfalk \\& Olofsson \\cite{Gaalfalk}) . In most of the numerous investigations of this source, a distance of 250\\,pc as given by Tomita et al. (\\cite{Tomita}) is assumed.  Stutz et al. (\\cite{Stutz}) have recently revised the distance estimate by using archive photometry of stars within a radius of 20\\arcmin. Their revised distance is 150\\,pc (60--200\\,pc). The diameter of B\\,335 as seen in the optical is just $\\approx$ 4\\arcmin\\,. CO observations (Frerking  et al. \\cite{Frerking}), however, show that it extends far to the NE, including stars within such a large radius as 20\\arcmin.   \\begin{table} \\begin{minipage}[t]{\\columnwidth} \\caption{Observation log} \\label{obslog} \\renewcommand{\\footnoterule}{}  % \\begin{tabular}{llllll} \\hline \\hline object &camera &filter &center &exp time&exp time\\\\ & & &  wavel   &per frame &total\\\\ & & &$ [nm]$&[s]&[s]\\\\ B\\,335 &EMMI-blue &U602  &353.96 &720 &26600\\\\ B\\,335 &EMMI-red  &Bb605  &413.22 &720 &7900\\\\ B\\,335 &EMMI-red  &g772  &508.86 &480 &4700\\\\ B\\,335 &EMMI-red  &r773  &673.29 &300 &2100\\\\ B\\,335 &EMMI-red  &I610  &798.50 &300 &2400\\\\ B\\,335 &NOTCAM\\footnote{co-added image provided by M. G{\\aa}lfalk\\\\} &Ks &2144. &300  &2800\\\\ B\\,335ref &EMMI-blue &U602 &353.96 &720 &2200\\\\ B\\,335ref &EMMI-red &Bb605 &413.22 &720 &800\\\\ B\\,335ref&EMMI-red &g772 &508.86 &480 &500\\\\ B\\,335ref &EMMI-red &r773 &673.29 &300 &300\\\\ B\\,335ref &EMMI-red &I610 &798.50 &300 &300\\\\ \\hline\\\\ \\end{tabular} \\end{minipage} \\end{table} ",
        "conclusions": "\\begin{enumerate} \\item The fundamental limitation when it comes to determining the distance to a small dark cloud is the space density of stars and as the cool MS stars are the most abundant, any method should be able to sort out these stars and reliably separate them from cool giant stars. \\item It turns out that multi-colour imaging in combination with synthetic colours can provide reliable spectral classification and extinction, and we find that the U band is particularly important, both for the spectral classification and for tracing the extent of the globule. \\item We have tested a new SED-fitting method which allows for different R$_{V}$-values across the field. For the test case, B\\,335, we get a distance of 90--120\\,pc. \\item Our deep U image of B\\,335 shows a bright halo due to scattered star light, and the south-western rim is much brighter. We find that a possible explanation is that the globule is illuminated by a star, HD\\,184982, located 1.5 arcminutes  to the south-west of our field and a few tenths of a parcec behind the cloud. The distance to this star is 140--200\\,pc according to the Hipparcos Catalogue. However, the surface brightness can also be explained as due to the PSF wing of the same star, and we judge that this is in fact a more likely explanation. \\end{enumerate}"
    },
    "1112/1112.0656.txt": {
        "abstract": "Observed in cool chromospheric lines such as H$\\alpha$ or \\ion{Ca}{2} H, coronal rain corresponds to cool and dense plasma falling from coronal heights. Considered rather as a peculiar sporadic phenomenon of active regions, it has not received much attention since its discovery more than 40 years ago. Yet, it has been shown recently that a close relationship exists between this phenomenon and the coronal heating mechanism. Indeed, numerical simulations have shown that this phenomenon is most likely due to a loss of thermal equilibrium ensuing from a heating mechanism acting mostly towards the footpoints of loops. We present here one of the first high resolution spectroscopic observations of coronal rain, performed with the CRISP instrument at the Swedish Solar Telescope. This work constitutes the first attempt to assess the importance of coronal rain in the understanding of the coronal magnetic field in active regions. With the present resolution, coronal rain is observed to literally invade the entire field of view. A large statistical set is obtained in which dynamics (total velocities and accelerations), shapes (lengths and widths), trajectories (angles of fall of the blobs) and thermodynamic properties (temperatures) of the condensations are derived. Specifically, we find that coronal rain is composed of small and dense chromospheric cores with average widths and lengths of $\\sim370$~km and $\\sim1500$~km respectively, average temperatures below 7000~K, displaying a broad distribution of falling speeds with an average of $\\sim70$~km~s$^{-1}$ and accelerations largely below the effective gravity along loops. Through estimates of the ion-neutral coupling in the blobs we show that coronal rain acts as a tracer of the coronal magnetic field, thus supporting the multi-strand loop scenario, and acts as a probe of the local thermodynamic conditions in loops. We further elucidate its potential in coronal heating. We find that the cooling in neighboring strands occurs simultaneously in general suggesting a similar thermodynamic evolution among strands, which can be explained by a common footpoint heating process. Constraints for coronal heating models of loops are thus provided. Estimates of the fraction of coronal volume with coronal rain give values between 10~\\% and 40~\\%. Estimates of the occurrence time of the phenomenon in loops set times between 4 and 15~hours, implying that coronal rain may be a common phenomenon, in agreement with the frequent observations of cool downflows in EUV lines. The coronal mass drain rate in the form of coronal rain is estimated to be on the order of $10^{10}$~g~s$^{-1}$, a significant quantity compared to the estimate of mass flux into the corona from spicules. ",
        "introduction": "Solar observations over the past 30 years have presented a very dynamic picture of the Sun. Loops, closed magnetic field entities encountered basically everywhere in the Sun at different spatial scales (of a few Mm and up to a few hundred Mm) are perfect examples of this dynamic character. Especially in active region coronae, studies of such entities reveal a large range of lifetimes, dynamics, and thermodynamic properties which have led to a main trifold classification as hot, warm and cool loops \\citep[see the extensive review by][]{Reale_2010LRSP....7....5R}. Several debates exist over the characteristics of the heating leading to such observed properties. For instance, it is not known whether the heating has mainly an impulsive or a uniform character in time. It is not known whether the heating is generally localized at specific atmospheric heights or is rather uniform over the loop. And, maybe most importantly, the main agent converting the magnetic energy into thermal energy is not known \\citep[generally known as AC or DC mechanism depending on its timescale,][]{Klimchuk_2006SoPh..234...41K}.  Several observational studies have shown, however, that a significant fraction of loops in active regions are out of hydrostatic equilibrium \\citep{Aschwanden_2001ApJ...550.1036A, Aschwanden_2001ApJ...559L.171A, Winebarger_etal_2003ApJ...587..439W, Schmieder_etal_2004ApJ...601..530S, Aschwanden_2009ApJ...695...12A}. This family of loops shows in general temperatures below $2-3\\times10^{6}~$K and show significant changes over relatively short timescales as compared to hotter loops.  Indeed, heating and cooling processes seem to occur continuously in these loops \\citep{Kjeldseth_Brekke_1998SoPh..182...73K, UgarteUrra_etal_2006ApJ...643.1245U, Warren_2007PASJ...59S.675W, Ugarte-Urra_etal_2009ApJ...695..642U, Landi_etal_2009ApJ...695..221L, Dudik_etal_2011AA...531A.115D}. A recurrent finding in these and other active region studies is the presence of cool downflows along the legs of the loops \\citep{Foukal_1976ApJ...210..575F, Foukal_1978ApJ...223.1046F, Schrijver_2001SoPh..198..325S, Oshea_etal_2007AA...475L..25O, Tripathi_etal_2009ApJ...694.1256T}.  This non-equilibrium state is generally explained through footpoint heating, a scenario in which the loops have their main heating source located towards the footpoints.  Footpoint heating in active region loops has received further observational support. Upflows often appear correlated with nonthermal velocities at the footpoints of these loops \\citep{Doschek_etal_2007ApJ...667L.109D, Hara_2008ApJ...678L..67H, Nishizuka_Hara_2011ApJ...737L..43N}. Recently, \\citet{Hansteen_2010ApJ...718.1070H} and \\citet{DePontieu_etal_2011Sci...331...55D}  have shown that a considerable part of the hot coronal plasma could be heated at low spicular heights, thus explaining the fading character of the ubiquitous type II spicules \\citep{Rutten_2006ASPC..354..276R, DePontieu_etal_2007PASJ...59S.655D, Rouppe_etal_2009ApJ...705..272R}. But perhaps one of the clearest evidences for footpoint heating is put forward by the presence of cool structures in the active region coronae, such as filaments/prominences and coronal rain.  Coronal rain is not a newly observed phenomenon. Although observed for the first time almost 40 years ago \\citep{Kawaguchi_1970PASJ...22..405K, Leroy_1972SoPh...25..413L} few observational work exists on the subject. Furthermore, the term `coronal rain' is often erroneously attributed to basically any cold material falling down from coronal heights, as has been the case for describing prominence material falling down following a prominence eruption. Coronal rain corresponds to cool and dense blob-like material forming in the hot coronal environment in a timescale of minutes, which subsequently falls down to the surface along loop-like paths. The temperatures of the plasma composing coronal rain seem to range from transition region to chromospheric, according to the lines in which it has been observed: in \\ion{Ne}{7} 465~\\AA\\ and \\ion{O}{6} 1032~\\AA\\ with {\\it Skylab} \\citep{Levine_Withbroe_1977SoPh...51...83L}, in \\ion{O}{5} 629~\\AA\\ with {\\it SoHO}/CDS \\citep{Kjeldseth_Brekke_1998SoPh..182...73K}, in the 1600~\\AA\\ channel of {\\it TRACE} \\citep{Schrijver_2001SoPh..198..325S}, in the 304~\\AA\\ channel of {\\it SoHO}/EIT \\citep{DeGroof_2004AA...415.1141D, DeGroof05, Stenborg_etal_2008ApJ...674.1201S}, in the 304~\\AA\\ channel of {\\it SDO}/AIA \\citep{Kamio_etal_2011AA...532A..96K}, in \\ion{Ca}{2}~H with {\\it Hinode}/SOT \\citep{Antolin_2010ApJ...716..154A, Antolin_Verwichte_2011ApJ...736..121A}, in H$\\alpha$ \\citep[at Okayama Astrophysical Observatory, the Observatoire du Pic du Midi, Big Bear Observatory, the Swedish Vacuum Solar Telescope (SVST):][respectively]{Kawaguchi_1970PASJ...22..405K, DeGroof_2004AA...415.1141D, Leroy_1972SoPh...25..413L, Muller_2005ESASP.596E..37M}.  \\citet{Schrijver_2001SoPh..198..325S} followed the various stages of coronal rain formation in loops with the channels of {\\it TRACE}. The hot loops rapidly cool down through thermal conduction and radiation until becoming thermally unstable. These radiative instabilities occur locally in the corona and lead to the rapid formation of blob-like condensations. These first appear in the 1216~\\AA\\ channel and then in the 1600~\\AA\\ channel suggesting continuous cooling below transition region temperatures. \\citet{DeGroof05} analyzed coronal rain simultaneously with {\\it SOHO}/EIT in the 304~\\AA\\ channel and in H$\\alpha$ images from Big Bear Observatory. The blobs first appear in EIT before becoming bright in H$\\alpha$, supporting the progressive cooling scenario. Furthermore, \\citet{Schrijver_2001SoPh..198..325S} shows that the loops hosting the condensations bright in the 1216~\\AA\\ were still visible in the 171~\\AA\\ channel, suggesting a multi-temperature structure in the loops. Recently, \\citet{Kamio_etal_2011AA...532A..96K} in observations with {\\it SDO}/AIA, similarly report cool condensations bright in the 304~\\AA\\ channel that appear to be surrounded by diffuse hotter plasma in the 171~\\AA\\ channel.  Numerical simulations have pinpointed the now generally accepted mechanism behind coronal rain. Termed thermal non-equilibrium (or catastrophic cooling) the mechanism is based on the fact that a coronal loop whose heating input is mostly concentrated towards the footpoints will undergo a thermal instability locally in the corona. Due to the footpoint heating source the loop gets rapidly hot and dense. Thermal conduction acts efficiently for slowing the temperature increase over the loop. The heating source together with chromospheric evaporation will gradually increase the density in the loop bringing the heating per unit mass down and with it the average coronal temperature. Eventually, in a time scale of an hour, the temperature is sufficiently low to allow recombination of atoms locally, decreasing the temperature and the pressure locally in a timescale of minutes. The pressure drop accretes mass from the surroundings and produces upflows from the footpoints leading to the condensations. The dense clumps cool rapidly through radiation until the plasma becomes optically thicker and the cooling slows down \\citep{Goldsmith_1971SoPh...19...86G, Hildner_1974SoPh...35..123H, Mok_etal_1990ApJ...359..228M, Antiochos_Klimchuk_1991ApJ...378..372A, Antiochos_1999ApJ...512..985A, Karpen_etal_2001ApJ...553L..85K, Muller_2003AA...411..605M, Muller_2004AA...424..289M, Muller_2005AA...436.1067M, Mendozabriceno_2005ApJ...624.1080M, Mok_etal_2008ApJ...679L.161M, Antolin_2010ApJ...716..154A, Xia_etal_2011ApJ...737...27X}. It is still not clear how cool can coronal rain become from this process (before falling down to the solar surface).  Since they have so far been observed in the same spectrum lines, coronal rain is often associated with prominences, and probably for both the thermal non-equilibrium mechanism is equally important \\citep{Dahlburg_etal_1998ApJ...495..485D, Karpen_etal_2001ApJ...553L..85K, Karpen_etal_2005ApJ...635.1319K, Karpen_etal_2006ApJ...637..531K, Liu_etal_11SPD....42.2119L}. However, one fundamental difference seems to be the magnetic field structure in which they appear, most likely determining the time scale of the phenomenon. While a prominence can live for days or even weeks, coronal rain can only exist during the time it takes for the blobs to fall to the surface (on the order of half an hour or less). The thermodynamic properties of the plasma are very likely linked to the underlying magnetic field topology. Coronal rain is restrained to coronal loops, where the magnetic field is basically bipolar and magnetic field lines are not far from being potential. On the other hand, the magnetic field topology in prominences seems to present complicated geometries such as `dips' or `peaks'  \\citep[see][and references therein]{Lin_2010SSRv..tmp..112L}. Helical magnetic fields conforming flux ropes have also been proposed, in which apart from offering support from the Lorentz force to balance gravity, thermal isolation from the surrounding hot corona is also achieved due to the cross-field thermal conduction inhibition \\citep{Kippenhahn_Schluter_1957ZA.....43...36K, Low_Hundhausen_1995ApJ...443..818L}. The longer timescales involved in prominences may produce different element population leading to different ionization fraction, and consequently mechanisms such as ambipolar diffusion may play a different role. Another type of prominence model has been suggested in which the material is continuously renewed through flows between the prominence and the chromosphere, giving the impression of long-lived material \\citep{Priest_Smith_1979SoPh...64..267P, Tandberg-Hansen_1995ASSL..199.....T, Zirker_etal_1998Natur.396..440Z, Liu_etal_11SPD....42.2119L, Karpen_etal_2001ApJ...553L..85K, Karpen_etal_2005ApJ...635.1319K}. Since this case does not involve long-term gravitational support the required magnetic field may be more simple and features such as `dips' are not necessary: conditions that are more alike to coronal rain formation.  High resolution observations have revealed that prominences are often composed of a myriad of fine threads, outlining a fine-scale structure of the magnetic field and the presence of flows along the threads \\citep{Heinzel_Anzer_2006ApJ...643L..65H, Lin_2005SoPh..226..239L, Lin_2008ASPC..383..235L, Lin_2010SSRv..tmp..112L, Martin_2008SoPh..250...31M}. Observations of coronal rain with {\\it Hinode}/SOT have also revealed a thread-like character in coronal loops, due to the separation and elongation of the blobs down to very small sizes \\citep{Antolin_2010ApJ...716..154A, Antolin_Verwichte_2011ApJ...736..121A}. The observations with the {\\it SST} presented here allow a closer look into this thread-like character of loops, suggesting the actual tracing of strands by the blobs.  Due to the strong footpoint heating and the subsequent very high densities that follow a flare, thermal non-equilibrium leading to the formation of condensations is also proposed as the mechanism behind the observed cool downflows in post-flare loops  \\citep{Foukal_1978ApJ...223.1046F, Schmieder_etal_1995SoPh..156..337S, Shimojo_etal_2002SoPh..206..133S, Hara_etal_2006ApJ...648..712H}. Despite the very different thermodynamic conditions in post-flare loops with respect to warm loops, the coronal rain observed in both cases appears similar. Although a rigorous comparison awaits, this could be easily explained since radiative cooling increases with the square of the density, and therefore achieving the same low temperatures in a similar timescale. It is possible however that the element population of the obtained blobs in post-flare loops is different, leading to different properties at short length scales. It is thus important to make the distinction between flaring and non-flaring active regions for coronal rain studies.  Coronal rain is often observed to follow loop-like paths, suggesting a strong coupling between neutrals coming from recombination and the local ion population. Since the thermodynamical state and ionization fraction of such clumps are still poorly known it is not clear how strong this coupling is and, consequently, how well coronal rain blobs can trace the coronal magnetic field. Assuming that a blob does not alter the magnetic field lines, by studying its dynamics, morphology changes and thermodynamics we can learn about the local conditions in coronal loops. This is an idea that is exploited in this work. Supported by theoretical estimates of coupling processes based on the obtained results we show that the very small sizes of blobs allow them to be magnetic field tracers and probes of the local environment.  In all of the observational papers concerning coronal rain it has been shown that the falling speeds are smaller than free fall, leading to downward accelerations considerably lower than that of the solar gravity at the surface. The gas pressure gradient inside the loops has often been suggested as the main agent regulating the speeds of the blobs. This is also supported by numerical simulations of thermal non-equilibrium \\citep{Mackay_2001SoPh..198..289M, Muller_2003AA...411..605M}. Magnetic pressure exerted by transverse MHD waves have also been suggested as factors determining the kinematics of blobs. In \\ion{Ca}{2}~H observations with {\\it Hinode}/SOT, \\citet{Antolin_Verwichte_2011ApJ...736..121A} detected such waves through coronal rain tracking. It is shown that the pressure force from the waves can explain the low falling speeds. Further investigation on the effect of waves on blobs such as those of coronal rain is being carried out \\citep{Verwichte_etal_2011}. Whether blobs are capable of producing such waves is also being investigated in that work.  Based on the {\\it TRACE} observations, \\cite{Schrijver_2001SoPh..198..325S} estimated the occurrence time of coronal rain in a coronal active region loop to be at most once every two days stressing the sporadic character of coronal rain in active regions, and thus suggesting that thermal non-equilibrium may not play a significant role in coronal heating. These estimates are constrained by the resolution of {\\it TRACE}. Here we show that in order to have a clear picture about the phenomenon a resolution at least twice that of {\\it TRACE} is needed. New estimates of the occurrence time of coronal rain and the coronal volume involved are presented, as well as the coronal mass drain rate.  The importance of thermal non-equilibrium (catastrophic cooling) in coronal heating is a matter of significant debate. So far, observations have pointed out five characteristics of warm loops ($\\sim1~$MK) in active regions, which any coronal heating model must explain. Namely, an excess density, a flat temperature profile, a super-hydrostatic scale height, an unstructured intensity profile, and a $1000-5000~$s lifetime. \\cite{Klimchuk_2010ApJ...714.1239K} have analyzed through one-dimensional simulations the case of thermal non-equilibrium in this context, considering both monolithic and multi-stranded loops. The simulations failed in general to satisfactorily reproduce the observations. For instance, monolithic models end up with far too much intensity structure, while multi-strand models are either too structured or too long-lived. On the other hand, 3D-MHD models of an active region with a footpoint heating function leading to thermal non-equilibrium have provided satisfactory results \\citep{Mok_etal_2008ApJ...679L.161M}. The resulting evolution of the light-curves, the variation of temperature along the loops, the density profile, and the absence of small scale structures in the intensity profiles are all compatible with real loops, thus supporting the thermal non-equilibrium scenario \\citep{Lionello_etal_2010AGUFMSH31C1811L}.  Impulsive versus steady heating of loops is a common debate in studies of coronal heating \\citep[see][and references therein]{Klimchuk_2006SoPh..234...41K}. The scenario in which loops are heated impulsively to hot temperatures and subsequently cool down to EUV or lower temperatures is often proposed in observational studies of the corona \\citep{Warren_etal_2002ApJ...579L..41W, Winebarger_etal_2003ApJ...593.1164W, Winebarger_etal_2003ApJ...593.1174W, Viall_Klimchuk_2011ApJ...738...24V} but neither this nor the steady heating scenario are fully compatible with some observations \\citep{UgarteUrra_etal_2006ApJ...643.1245U}. Indeed, in some cooling events the {\\it TRACE} or {\\it SDO}/AIA filters do not seem to brighten according to their temperature. A possible solution to this has been put forward by \\citet{Kamio_etal_2011AA...532A..96K}, where, apart from investigating the upward propagation of quasi-periodic brightness fluctuations observed with {\\it SDO}/AIA, they also analyze sporadic cool downflows that they interpret as coronal rain. The peculiar order of the brightening in the AIA channels in time seem to be due to these cool downflows, particularly, because of the contribution of cool lines to hot channels in the AIA temperature response functions. Throughout their lifetimes the condensations progressively cool and the densities increase (in order to keep the local pressure roughly constant), producing a considerable increase in the cool line emission. The authors thus favor the thermal non-equilibrium mechanism to explain the observational results.  Being the observational signature of thermal non-equilibrium (and hence to footpoint heating), coronal rain is deeply linked with coronal heating. As shown by the numerical simulations of the works cited previously thermal non-equilibrium is very sensitive to parameters such as the heating scale height and the loop length. By studying coronal rain we can then learn about the heating mechanisms. \\citet{Antolin_2010ApJ...716..154A} showed that Alfv\\'en wave heating, a strong coronal heating candidate, is not a predominant heating mechanism in loops with coronal rain. When propagating from the photosphere into the corona, Alfv\\'en waves can nonlinearly convert to longitudinal modes due to density fluctuations, wave-to-wave interaction, and deformation of the wave shape during propagation \\citep{Vasheghani_2011AA...526A..80V}. These modes subsequently steepen into shocks and heat the plasma uniformly along the loop \\citep{Moriyasu_2004ApJ...601L.107M, Antolin_2010ApJ...712..494A}, thus avoiding the loss of thermal equilibrium in the corona. In this work we deepen the link between coronal rain and coronal heating by studying the occurrence character of thermal non-equilibrium in neighboring strands.  This paper is organized as follows. In section~\\ref{observations} we present the {\\it SST}/CRISP observations. In section~\\ref{results} the results of the statistical analysis are presented. These comprise dynamics, shapes, trajectories and thermodynamic properties. Discussion is presented in section~\\ref{discussion}, where we try to assess the role of coronal rain in the understanding of the coronal magnetic field. Conclusions are presented in section~\\ref{conclusions}. ",
        "conclusions": "In this paper we have analyzed H$\\alpha$ observations of coronal rain with the CRISP spectropolarimeter of the {\\it SST}. The phenomenon is observed above an active region at the limb. Kinematics (total velocities and accelerations), shapes (widths and lengths), trajectories (angles of falling blobs) and thermodynamic properties (temperatures) were derived for the condensations. On-disc blobs dark in absorption were also analyzed and shown to have the same dynamical, morphological and thermodynamical properties as the off-limb blobs. We have further shown that the profile of off-limb blobs becomes an absorption profile when passing against a bright background,  from which we have concluded that on-disc blobs correspond to the same phenomenon.  Being one of the first spectroscopic studies on coronal rain we have been able to measure the total velocities and shown that the falling velocities are lower than free fall, as has been suggested in previous works on the subject \\citep{Schrijver_2001SoPh..198..325S, DeGroof_2004AA...415.1141D, Muller_2005ESASP.596E..37M}. The corresponding downward accelerations are lower than the effective gravity along elliptic loops, suggesting the presence of other forces than gravity, possibly gas pressure as suggested by most numerical simulations. Transverse magnetic waves could not be clearly detected due to seeing effects so we cannot discard their presence and influence on the dynamics of the blobs. The blobs could often be traced to chromospheric heights, close to what appears to be the footpoints of the loops that host them. Strong decelerations were observed in some cases, suggesting the increase of gas pressure at those heights expected from the higher local densities.  Through orders of magnitude estimates we have shown that the coronal rain observed in H$\\alpha$ is expected to come from a neutral Hydrogen population strongly coupled to the protons. At least on length scales on the order of the blob sizes (a few hundred km) no diffusion across the magnetic field is expected. Coronal rain can thus be considered as a tracer of the coronal magnetic field. Combining the projected and Doppler velocities we have shown that it is possible to retrieve the angles of fall of the blobs, allowing a reconstruction of the coronal magnetic field. No evidence for twisting or braiding of strands in loops was found, but we do not discard their existence at lower length scales (under 100 km), where turbulence may also be important. The tracing of strands by the blobs further suggest a constant area cross-section in the corona for these loops with no significant expansion down to chromospheric heights. The expected flux tube expansion may happen at lower (photospheric?) heights.  One of our main results concerns the sizes of the observed blobs. Lengths and widths on the order of 1.5~Mm and 0.37~Mm, respectively, were found, stressing the need of high spatial (and temporal) resolution to fully observe this phenomenon. The blobs further present average temperatures of 7000~K, and possibly even lower if turbulence and the Stark effect are important. Coronal rain thus present cool and compact chromospheric cores of a few 100~km in width. Some evidence of progressive cooling over time was found, as suggested from previous work. We further found that coronal rain occurs frequently in neighboring strands in a simultaneous way, forming groups of condensations which, if close enough together, are seen as large clumps. We have termed these sporadic events as `showers', and they can have widths up to a few Mm. This is probably what has been observed in the past with instruments of lower spatial resolution such as {\\it TRACE} and {\\it SOHO}/EIT. Since we did not observe here any evidence of cool blobs nor showers in the co-temporal {\\it Hinode}/EIS observation we suspect that loops with coronal rain are multi-thermal, as suggested previously by \\cite{Schrijver_2001SoPh..198..325S}. A scenario in which the cool coronal rain cores are surrounded by diffuse hotter plasma needs to be further investigated.  Our observations support the multi-stranded loop scenario and suggest that a significant fraction of strands in a loop do not have an independent thermodynamic evolution. Indeed, neighboring strands often display a coherent cooling (in the form of coronal rain) of a significant number of strands in a loop, suggesting a heating mechanism which acts in roughly the same way on neighboring strands. A footpoint heating mechanism imparting a similar heating scale height over all or part of the strands in a loop is a simple scenario that can explain these observations. Since we expect an expansion of the loops at low (possibly photospheric) heights spatially small heating events may suffice for such purpose.  A coherent evolution of strands in a loop throughout their lifetime (from heating to cooling, from being dense to being evacuated) is a puzzling scenario that needs further observational testing. In case of being a common scenario it  may pose several constraints on coronal heating models, especially on one-dimensional models. For instance, one constraint may be that a loop cannot be modeled by a collection of independently evolving one-dimensional strands. This may explain the disagreement regarding the importance of thermal non-equilibrium in the corona between the 1D \\citep{Klimchuk_2010ApJ...714.1239K} and the 3D simulations \\citep{Klimchuk_2010ApJ...714.1239K, Lionello_etal_2010AGUFMSH31C1811L}.  The present observations suggest that coronal rain may not be a sporadic phenomenon of active regions as previously thought, but a rather common phenomenon deeply linked to the heating mechanisms of coronal loops. We have estimated the fraction of coronal volume with coronal rain to be between 10~\\% and 40~\\%. This is however strongly dependent on the magnetic field filling factor and on the projection effects. We have further estimated the occurrence time of this phenomenon in a loop to be between 4 and 15~hours. Longer datasets with larger fields of view are needed to clearly answer the question on how common coronal rain is. The obtained mass drain rate in the form of coronal rain is roughly $10^{10}$~g~s$^{-1}$ taking a mass density of $10^{-13}$~g~cm$^{-3}$ for the condensations. This number is on the same order of magnitude as the estimated mass flux into the corona from spicules, reinforcing the idea that coronal rain is an important phenomenon. This suggests a scenario in which the hot spicular material injected into the corona falls back cool, `raining' down.  Irrespectively of the occurrence time of coronal rain, we have shown in this paper the large potential that this phenomenon beholds in the understanding of the coronal magnetic field. In \\citet{Antolin_2010ApJ...716..154A} we have shown that coronal rain can act as a marker of the coronal heating mechanisms. Here, besides further elucidating its link to coronal heating, we have shown that it constitutes a tracer of the internal structure of loops and that it can be a probe of the local thermodynamic conditions in loops."
    },
    "1112/1112.0886_arXiv.txt": {
        "abstract": "The Kilo Degree Survey (KiDS) is a 1500 square degree optical imaging survey with the recently commissioned OmegaCAM wide-field imager on the VLT Survey Telescope (VST). A suite of data products will be delivered to ESO and the community by the KiDS survey team. Spread over Europe, the KiDS team uses \\textsf{Astro-WISE} to collaborate efficiently and pool hardware resources. In  \\textsf{Astro-WISE} the team shares, calibrates and archives all survey data. The data-centric architectural design realizes a dynamic 'live archive' in which new KiDS survey products of improved quality can be shared with the team and eventually the full astronomical community in a flexible and controllable manner. ",
        "introduction": "One of the radical advances that optical astronomy has seen in recent years is the advent of wide-field CCD-based surveys. On Paranal, ESO has recently started operating two dedicated survey telescopes: VISTA in the infra-red wavelength region and the VLT Survey Telescope (VST) in the optical. The lion's share of the observing time on both survey telescopes will be invested in a set of Public Surveys.  The largest of the optical surveys is the Kilo-Degree Survey (KiDS), which will image 1500 square degrees in four filters ($u$,$g$,$r$,$i$) over a period of 3--4 years. Combined with one of the VISTA surveys, VIKING, which will observe the same area in ZYJHK, this will provide a sensitive, 9-band multi-colour survey.  \\begin{figure}[ht] \\includegraphics[width=\\textwidth]{P157_fig1.ps} \\caption{Lay-out of the KiDS-North (top) and KIDS-South (bottom) fields, shown by the hatched areas. Also shown are the areas where 2DF spectra are available, indicated by the large dots, and the area covered by DR7 of the SDSS survey, indicated by the small dots. The CFHTLS-W2 field and the DS/COSMOS deep field are overplotted on the top panel.} \\label{fig:areas} \\end{figure}  {\\bf Observational set-up.} KiDS will cover 1500 square degrees, which is approximately 7\\% of the extragalactic sky. It consists of two patches that ensure that observations can take place year-round. The Northern patch lies on the celestial equator, while the Southern area straddles the South Galactic Pole (Fig. \\ref{fig:areas}). These specific areas were chosen because they were the target of massive spectroscopic galaxy surveys already: the 2dF redshift survey \\citep{2dfgrs} covers almost the same area, and KiDS-North overlaps with the SDSS spectroscopic and imaging survey \\citep[SDSS, ][]{sdssdr8}. The exposure times for KiDS and VIKING have been chosen to yield a median galaxy redshift of 0.8, so that the evolution of the galaxy population and matter distribution over the last $\\sim$ half of the age of the universe can be studied. They are also well-matched to the natural exposure times for efficient VST and VISTA operations, and balanced over the astro-climate conditions on Paranal (seeing and moon phase) so that all bands can be observed at the same average rate. This strategy makes optimal use of the fact that all observations are queue-scheduled, allowing the best seeing time to be used for deep $r$-band exposures, for example, and the worst seeing for $u$.  {\\bf Science drivers.} The main scientific objective of KiDS and VIKING is to map the matter distribution in the universe through weak gravitational lensing and photometric redshift measurements. The large numbers of galaxies that KiDS will detect, with accurate photometric redshifts up to $z\\simeq1.2$ will allow the Baryonic Accoustic Oscillations, an important cosmological standard candle, to be measured over a large redshift range, and thus unveil its evolution. Galaxy-galaxy lensing (GGL) studies into the structure of galaxy halos for various redshifts and galaxy types, will exploit the excellent image quality of the OmegaCAM wide-field camera and the VST on the one hand and the shear size of the KiDS data set on the other. The deep photometry and accurate photometric redshifts also will ensure that KiDS data will be a powerful tool to study the evolution of galaxies and clusters out to redshifts of $z\\simeq1.5$. Additionally, the extensive data set that KiDS will deliver, will be useful in a broad range of rsearch areas in astronomy, for example the study of stellar streams and the Galactic halo.  {\\bf Survey data products.} Being a Public Survey, all KiDS data will be made publicly available. The KiDS catalogue will contain some 100,000 sources per square degree (150 million sources over the full survey area), and for each square degree there will be 10 GB of final image data, 15 TB for the whole survey. A set of basic data products will be made public, both through ESO and through the \\textsf{Astro-WISE} database: calibrated coadded images, weight maps, calibration images, single-band and multi-band catalogues. In the long-term, we intend to provide more advanced data products, for example images with gaussianized point-spread-functions, or morphological parameters of all detected sources. ",
        "conclusions": ""
    },
    "1112/1112.0248_arXiv.txt": {
        "abstract": "In cosmic-ray physics, large field of view experiments are triggered by a number of signals laying on different angular scales: point-like and extended gamma-ray sources, diffuse emissions, as well as large and intermediate scale cosmic-ray anisotropies. The separation of all these contributions is crucial, mostly when they overlap with each other. Needlets are a form of spherical wavelets that have recently drawn a lot of attention in the cosmological literature, especially in connection with the analysis of CMB data. Needlets enjoy a number of important statistical and numerical properties which suggest that they can be very effective in handling cosmic-ray and gamma-ray data analysis. An application of needlets to astroparticle physics is shown here. In particular, light will be thrown on how useful they might be for estimating background and foreground contributions. Since such an estimation is expected to be optimal or nearly-optimal in a well-defined mathematical sense, needlets turn out to be a powerful method for unbiased point-source detections. In this paper needlets were applied to two distinct simulated datasets, for satellite and EAS array experiments, both large field of view telescopes. Results will be compared to those achievable with standard analysis tecniques in any of these cases. ",
        "introduction": "\\label{sec:needlets} Needlets enjoy quite a few important properties that make them very suitable for spherical data analysis. Indeed, needlets do not rely on any tangent plane approximation and they are perfectly adapted to standard packages; their main property is the capability of unfolding directional data in the harmonic space, so that every order $j$ of needlets contains power only from within a certain range in the $l$ space. In the real space for any fixed angular distance the tail of the needlets decay faster than any polynomial, i.e. quasi-exponentially as the frequency increases. It is worth noticing that for needlets the exact reconstruction property holds, i.e. back-transform exists giving the input signal within numerical accuracy.  We do not give details about needlets construction here (on this purpose, see \\cite{npw1,bookmarinucci}). Nonetheless, we recall that the needlet idea has been extended by Geller and Mayeli \\cite{gm1} with the construction of so called Mexican needlets, see  also \\cite{scodeller} for numerical analysis and implementation in a  cosmological framework. Mexican needlets have localization properties in real space even better than standard needlets, that is why they were applied to get the results presented here.  Numerical codes have been developed by the author to compute the needlet estimators $\\widehat{\\beta}_{jk}$ for any sky map containing density information. Some other codes are publicly available \\cite{needatool} and results reported here have been suitably cross-checked. The index $j$ refers to the needlet order, while $k$ runs along the pixels into which the sphere has been discretized (the Healpix package is used here \\cite{Healpix}, with the ``ring'' numbering scheme).  When the needlet transform is numerically computed, three parameters are important, i.e. $j_0$ (the first order computed) and $n_j$ (how many needlet orders are computed) and $B>1$ which is a number used to pass from the $j$-th order to the $(j+1)$-th. A golden rule to be used is that the order $j$ contains power mostly from multipoles in the range $[B^{j-1},B^{j+1}]$\\footnote{That is, if $B=1.6$, then $j=6$ corresponds to $l\\sim1.6^{6}=17$, or $\\psi\\sim10^{\\circ}$.}. If mexican needlets are used, one more parameter is important, $p$, which describes how fast needlets go to zero in the harmonic space.  We shall now consider the analysis of data from cosmic rays observatories, following classical approaches to wavelet-based density estimation. An idea that we shall pursue is to implement \\emph{thresholding} estimates, as discussed for instance by \\cite{bkmpAoSb}. More precisely, we can consider the nonlinear estimate $\\widehat{f}_{n}^{\\ast }(x)$ obtained by suitably weighting the needlet coefficients in the back-transform. The weight function $w_{jk}(\\widehat{\\beta }_{jk})$ is some nonlinear function that ``shrinks'' beta towards zero. Such estimates can be shown to be robust and nearly optimal over a wide class of density functions and different loss functions, i.e. figures of merit by which to measure when the estimate is \"close\" to the density to be estimated.  Hereafter, we will name {\\em source-map} the sky-map as it comes from the experiment; {\\em beta-maps} those containing the coefficient estimators $\\widehat\\beta_{jk}$; {\\em reconstructed maps} or {\\em back-transformed maps}, those obtained after the application of the inverse needlet transform. ",
        "conclusions": "Recently, needlets drew the attention of the scientific community for their important applications in data-analysis of cosmological data as a new form of spherical wavelets. We presented here some applications of the needlet transform to high energy cosmic ray physics, showing the very good properties of localization.The needlet transform is sensitive in the whole harmonic domain, provided that enough orders of needlets are computed. In particular, the application to a simulation of the ARGO-YBJ and the Fermi-LAT data-sets found again the well-known intermediate scale anisotropy at low orders and the point gamma-ray sources at higher orders, thus showing the possiblity of decting sources directly in the needlet space. The significance of the beta-maps is carried out quite easily if the background distribution is known and the variance of the beta coefficients may be used to threshold the signal in the beta space, which turned out to be a very promising technique."
    },
    "1112/1112.4141_arXiv.txt": {
        "abstract": "Rastall's theory is based on the non-conservation of the energy-momentum tensor. We show that, in this theory, if we introduce a two-fluid model, one component representing vacuum energy whereas the other pressure-less matter (e.g. baryons plus cold dark matter), the cosmological scenario is the same as for the $\\Lambda$CDM model, both at background and linear perturbation levels, except for one aspect: now dark energy may cluster. We speculate that this can lead to a possibility of distinguishing the models at the non-linear perturbation level. ",
        "introduction": "Since the formulation of General Relativity, about one hundred years ago, many alternative geometric theories have been proposed in order to explain gravitation phenomena (e.g. \\cite{Cartan1922, Cartan1923, Brans1961, Rosen1971, Rastall1973, Rastall1976, Moffat:1994hv, Bekenstein:2004ne}). Some of these touch one important aspect of General Relativity: the conservation law. To our knowledge, the first non-conservative theory of gravity was the steady-state model \\cite{Bondi1948, Hoyle1948}, following some ideas already presented in the end of the forties by Jordan \\cite{Jordan:1949zz}. In the beginning of the seventies, Rastall proposed one new version of a non-conservative theory of gravity, following the remark that the conservation law ${T^{\\mu\\nu}}_{;\\mu} = 0$ may not hold true in a curved space-time \\cite{Rastall1973, Rastall1976}. Hence, he argued that new gravitational equations can be obtained considering a modification of the conservation law such that \\begin{eqnarray} \\label{mot1} {T^{\\mu\\nu}}_{;\\mu} = \\kappa R^{;\\nu}, \\end{eqnarray} where $T^{\\mu\\nu}$ is the energy-momentum tensor, $\\kappa$ is a coupling constant and $R$ is the Ricci scalar curvature. Hence, in the weak field limit, the usual expressions are preserved. Since, generally, the Ricci scalar curvature is connected with the trace of the energy-momentum tensor, Eq.~(\\ref{mot1}) can be re-written as \\begin{eqnarray} {T^{\\mu\\nu}}_{;\\mu} = \\bar\\kappa T^{;\\nu}\\;, \\end{eqnarray} where $\\bar\\kappa$ is a new constant and $T$ is the trace of the energy-momentum tensor. It is curious to remark that the phenomenon of particle creation in cosmology \\cite{Gibbons:1977mu, Parker:1971pt, Ford:1986sy} also leads to a violation of the classical conservation laws and, in this sense, Rastall's idea may be viewed as a kind of classical formulation of that quantum phenomenon, since the violation of the energy-momentum conservation is connected with the curvature.  More in detail, Rastall's modification to Einstein equations take the following form ($c = 1$ units): \\begin{eqnarray} \\label{eq1R} R_{\\mu\\nu} - \\frac{1}{2}g_{\\mu\\nu}R &=& 8\\pi G\\left(T_{\\mu\\nu} - \\frac{\\gamma- 1}{2}g_{\\mu\\nu}T\\right)\\;,\\\\ \\label{eq2R} {T^{\\mu\\nu}}_{;\\mu} &=& \\frac{\\gamma - 1}{2}T^{;\\nu}\\;, \\end{eqnarray} where $\\gamma$ is a parameter (the choice $\\gamma = 1$ restores General Relativity). Note that it seems possible to have a Lagrangian formulation from which the above equations are deduced \\cite{Smalley1984}. Since for a radiative fluid $T = 0$, implying $R = 0$, we can expect that the cosmological evolution during the radiative phase is the same as in the standard cosmological scenario. At same time, a single fluid inflationary model, described by a cosmological constant, is the same as it would be in the General Relativity case. Hence, Rastall cosmologies may have an important departure from the standard cosmological model from the beginning of the matter dominated phase on \\cite{Batista:2010nq, Fabris:2011rm, Fabris:2011wz, Capone:2009xm}. ",
        "conclusions": "We have shown that there is a subset of cosmological scenarios based on Rastall's energy-momentum tensor non-conservation that are equivalent to the $\\Lambda$CDM cosmology, except for one aspect: vacuum energy may agglomerate. Two effects could allow to discriminate between the two models: non-linear effects in the matter power spectrum and the transfer function for cosmological perturbations. Vacuum energy is negligible in the past, hence the impact of its fluctuation in the evolution of the perturbations in the primordial periods of the evolution of the universe may not be relevant. However, according to the scenario described above, dark energy must be present in virialized systems, like galaxies and cluster of galaxies, and the effect of agglomeration of dark energy must be relevant at this level. Therefore, Rastall's cosmology and the $\\Lambda$CDM model seem to be distinguishable only at the non-linear regime of the evolution of cosmic perturbations. Though we have given some hints of the latter possibility addressing the simple case of a spherical top-hat collapse, a deeper analysis shall be the subject of future investigation.  \\subsection*{Acknowledgements}  The authors thank the anonymous referee for his kind and useful suggestions. J. C. F. thanks CNPq (Brasil) for partial financial support."
    },
    "1112/1112.1880_arXiv.txt": {
        "abstract": "The thermal evolution of isothermal neutron stars is studied with matter both in the hadronic phase as well as in the mixed phase of hadronic matter and strange quark matter. In our models, the dominant early-stage cooling process is neutrino emission via the direct Urca process. As a consequence, the cooling curves fall too fast compared to observations. However, when superfluidity is included, the cooling of the neutron stars is significantly slowed down. Furthermore, we find that the cooling curves are not very sensitive to the precise details of the mixing between the hadronic phase and the quark phase and also of the pairing that leads to superfluidity. ",
        "introduction": "Neutron stars are natural laboratories for physics under extreme conditions with their extremely high densities, powerful energy emission, large magnetic fields, and millisecond rotation periods. At the densities near the surface of such a star, atoms break apart into nuclei and electrons. At higher densities, the electrons neutralize with the protons in the nuclei to form neutrons. These stars thus consist of a large fraction of neutrons and are supported from gravitational collapse by the neutron degeneracy pressure, from which the neutron star derives its name.  However, at high densities the existence of more exotic particles is expected. These particles are generated by processes which produce strangeness, such as \\begin{equation}\\label{lambda} n+n\\rightarrow n+\\Lambda^0+K^0, \\end{equation} where $n$ is the neutron, $\\Lambda^0$ the Lambda hyperon, and $K^0$ the strange meson. The strange meson can decay via various weak processes and the final products, usually photons and neutrinos, leak out of the star. Therefore, the reverse process is reduced and some net amount of strangeness survives in the dense core of the star \\cite{Glendenning:1985}. It is generally believed that these strangeness carrying particles, called hyperons, can exist in the center of neutron stars. They coexist with the nucleons as well as the leptons $e^-$ and $\\mu^-$. The interactions between them are dominated by the complicated nuclear force whose carriers are the mesons. We refer to such a system as the hadronic phase of matter. The system with only nucleons and leptons, i.e., without hyperons, is referred to as the nuclear phase of matter.  At even higher densities, as a consequence of asymptotic freedom, quarks become deconfined from the hadrons. Therefore, strange quark matter, which consists of $u$, $d$ and $s$ quarks, may also exist in the neutron star core. There may also be a phase-separated mixture of strange quark matter and hadronic matter in a certain range of densities \\cite{Glendenning:1992}, which we call the mixed phase. At present, there is still a lot unknown about the deconfinement phase transition of quarks. The contemporary knowledge on matter at high densities and temperatures is shown by the schematic QCD phase diagram in Fig.~\\ref{QCD}. Since the baryon chemical potential $\\mu_B$ is a monotonously increasing function from the surface to the center of the neutron star, the $\\mu_B$ axis can be mapped to the stellar radius and the different phases in the neutron star are explicitly indicated. \\begin{figure} \\resizebox{1\\linewidth}{!}{\\includegraphics*{QCD.eps}} \\caption{The schematic QCD phase diagram as a function of temperature $T$ and baryon chemical potential $\\mu_B$. The solid lines denote first- and second-order phase transitions, whereas the dashed lines denote smooth crossovers. According to the typical temperature of a proto-neutron star of $T \\sim 10^{10}$ K $\\sim 1$ MeV, the state of matter in the neutron star should be close to the bottom line, as indicated by the grey thick line.}\\label{QCD} \\end{figure}  Because of the limited knowledge on the state of matter at high densities and the complexity of the interactions between the particles, many effective models for matter inside neutron stars have been constructed. In general, these models supply an equation of state, which determines the particle composition of the neutron star. The observation of neutron stars, in turn, constrain these models. For example, if an equation of state is too soft it is incapable of supporting a very large stellar mass, such that some models become disfavored whenever the data on the heaviest neutron star is updated. The stellar mass is not the only constraint on the models. The cooling of neutron stars can be studied from their luminosity as a function of time. A proto-neutron star is born with a typical temperature larger than $10^{10}\\textrm{ K}$, after which it mainly loses its energy by two processes, namely by neutrino emission everywhere inside the star and by photon radiation at the surface. In the early stages of the thermal evolution of the star, neutrino emission is the dominant cooling effect after which photon radiation ultimately takes over \\cite{Yakovlev:2001}. Since neutrino emission occurs everywhere in the neutron star, it provides a probe for studying the state of matter inside the star.  The most efficient neutrino-emission process is called the direct Urca (DUrca) process \\begin{equation}\\label{DUrca} n\\rightarrow p+e^-+\\bar{\\nu}_e,\\quad p+e^-\\rightarrow n+\\nu_e. \\end{equation} This process is only possible if the proton fraction is more than a certain threshold in order for it to satisfy energy and momentum conservation \\cite{Lattimer:1991}. Historically, the proton abundance in neutron stars was underestimated. In that scenario, it is reasonable to take into account also the modified Urca (MUrca) process, where a bystander helps the momentum conservation, for example \\begin{align}\\label{MUrca} n+n\\rightarrow p+n+e^-+\\bar{\\nu}_e,&\\ p+n+e^-\\rightarrow n+n+\\nu_e,\\nonumber\\\\ n+p\\rightarrow p+p+e^-+\\bar{\\nu}_e,&\\ p+p+e^-\\rightarrow n+p+\\nu_e. \\end{align} Some other processes can have a neutrino emissivity which is smaller or comparable to the modified Urca process, such as neutrino bremsstrahlung and plasmon decay \\cite{Yakovlev:2001}. However, all of them are negligible whenever the direct Urca channel is open. In the hadronic phase, the direct Urca processes are quite rich and can be summarized as \\begin{equation}\\label{DUrcabaryon} b_1\\rightarrow b_2+l^-+\\bar{\\nu}_l,\\quad b_2+l^-\\rightarrow b_1+\\nu_l, \\end{equation} where $b_1$ and $b_2$ denote two different baryons and $l^-$ represents one of the leptons, $e^-$ or $\\mu^-$.  For strange quark matter, the direct Urca processes are \\begin{align}\\label{DUrcaquark} d\\rightarrow u+l^-+\\bar{\\nu}_l,\\quad u+l^-\\rightarrow d+\\nu_l,\\nonumber\\\\ s\\rightarrow u+l^-+\\bar{\\nu}_l,\\quad u+l^-\\rightarrow s+\\nu_l, \\end{align} which are simply the direct Urca processes of the baryons at the quark level. Inside the neutron star, there is no threshold for the direct Urca process in quark matter, since the Fermi momenta of the two quarks and lepton can always satisfy momentum conservation. As mentioned, the neutrino emissivity depends strongly on the type of matter contributing to the process, such that the thermal evolution of a star directly probes its composition. The main goal of this paper is to study the cooling of neutron stars with different types of equations of state and try to constrain them by comparing with observational data.  A large amount of research on the cooling process of neutron stars has been carried out in the last several decades, with a strong focus on the nuclear phase of matter. To the best of our knowledge, a unifying model covering the low-density nuclear phase to the deconfined quark phase, consistent with the QCD phase diagram, has been less thoroughly explored. The strange quark star composed of strange matter, either with or without the nuclear crust, has been considered in some papers \\cite{Schaab:1996,Page:2002}. However, the cooling behavior of neutron stars containing also all hyperons and possibly also a mixed phase of strange quark matter and hadronic matter, has not been extensively studied. Two possible reasons for this can be given. First, the hadronic and mixed equations of state are rather soft and are continuously being challenged by new data on heavy neutron stars, such as PSR J1903+0327 which has $M=1.67\\pm0.01\\textrm{ M}_\\odot$ \\cite{Freire:2009} and PSR J1614$-$2230 which even has $M=1.97\\pm0.04\\textrm{ M}_\\odot$ \\cite{Demorest:2010}. Second, the existence of a mixed phase is being questioned in view of screening and surface effects \\cite{Voskresensky:dual}. With respect to the first concern, we note that most of the observed neutron star masses still lie below the maximum mass a hadronic or mixed equation of state can allow for. According to Ref.~\\cite{Timmers:1996}, most nearby young neutron stars \\cite{Popov:2003} have masses no bigger than $1.4\\textrm{ M}_{\\odot}$ \\cite{Page:2006}. Therefore, at least for the study of the thermal evolution of these stars, the hadronic or mixed phase can still be of great importance. In fact, because of the uncertainties in the interaction between particles at extremely high density, it is hard to exactly determine the equation of state at high density. Although the equations of state used in this paper cannot support neutron stars as massive as those reported above, we can still use them for a discussion on medium-mass neutron stars with more complicated compositions. As for the stability of the mixed phase, the arguments are still indecisive. For example, many details of the surface tension, which strongly influences the stability calculation, are still uncertain. Although it is expected that screening and surface effects diminish the mixed-phase regime, it is far from certain that its existence can be excluded \\cite{Endo:2011}. The mixed phase can in particular have significant effects on the cooling behavior, especially for the heat transport inside the star. However, we show below that its effect will be less important after the star has become isothermal, when the thermal evolution is determined by the heat capacity and neutrino emissivity integrated over the whole volume, and the inner thermal conductivity no longer plays a role.  The temperature of a neutron star is generally much smaller than the typical Fermi energy as a consequence of the very high densities in the star. Therefore, superfluidity may play an important role. According to BCS theory, fermions can form Cooper pairs at low temperatures via an attractive interaction and thereby lower the energy of the system. The resulting pairs, which obey Bose statistics, can form a Bose-Einstein condensate and the system becomes superfluid. Pair formation changes the single-particle dispersion around the Fermi surface and consequently the heat capacity and neutrino emissivity will be influenced. We find that, without superfluidity, the cooling of neutron stars is too fast compared with observations. However, by including the effects of superfluidity we obtain a more realistic cooling behavior.  The quark phase is usually referred to as an exotic phase of extremely dense matter, in contrast to the nuclear phase or the hadronic phase. Other exotic phases have also been proposed, among which the pion and kaon condensates have attracted much attention \\cite{Ohnishi:2009,Brown:2008}. However, the existence of such phases inside the neutron star is still an open question. On the one hand, with such condensates, the equation of state is further softened and thus the corresponding maximum star mass is reduced \\cite{Potekhin:2011}, which makes such phases less favored when compared to the observational data of massive stars, as mentioned above. On the other hand, for the cooling process, it was reported that a meson condensate can increase the neutrino emissivity over the typical modified Urca emissivity by several orders of magnitude \\cite{Yakovlev:2001}, but it is still much less efficient than the direct Urca process. Since in our calculation the direct Urca process is always present, such enhancement from the meson condensate has a negligible effect. Besides, the kaon condensation may even reduce the pressure and cause the star to collapse into a black hole \\cite{Brown:2008}. Therefore, considering all the above arguments, we do not include such meson condensates in this paper.  The paper is organized as follows, we first introduce in Sec. II our theoretical framework which includes the relativistic stellar structure, the unified mean-field model describing the state of matter inside the star, and the thermal evolution equations. In Sec. III, we then present the numerical results for the cooling curves of neutron stars with the nuclear, hadronic, and mixed equations of state. Due to the efficiency of the direct Urca process the cooling is seen to be too fast compared to observations. Therefore superfluidity is included in Sec. IV and, as a consequence, the cooling is slowed down and we find a much better agreement with observations. ",
        "conclusions": "We have carefully compared the thermal evolution of different types of neutron stars, namely with the hadronic, the mixed phase of hadronic and strange quark matter, and the nuclear equation of state. We find that the direct Urca process is open in all of these cases and thus results in relatively fast cooling behavior. Although the details concerning the heat capacity and neutrino emissivity can be rather different in these cases, the cooling curves are quite similar after the stars become isothermal. However, the behavior in the early stages before the stars become isothermal could be significantly different, but such a study requires the knowledge of the thermal conductivity of these complex systems. The geometrical structure of the mixed phase is also expected to play an important role, although no decisive conclusion has yet been drawn.  As we have seen, the fast cooling in the non-superfluid case did not agree with observations. In order to remedy this discrepancy, superfluidity was introduced, which significantly reduces the efficiency of the direct Urca process as well as the heat capacity. The resulting cooling curve is much closer to the observational data. We also found that the particular pairing type of the superfluid baryons is not very important to the thermal evolution of the star. The thermal evolution after the star becomes isothermal is not strongly dependent on the initial temperature of the star or on the critical temperature related to superfluidity, as long as they are within reasonable ranges. The robustness of the results with superfluidity is quite helpful in order to remove the uncertainties concerning baryon superfluidity at high density. Note that the cooling process is almost completely determined by the direct Urca process and photon radiation even after including the effects due to superfluidity. The cooling curves of the neutron stars with the three equations of state are still quite similar when superfluidity is included. We expect that a calculation including the magnetic field and rotation gives an even better agreement with the observational data. Besides, since our nuclear model is geared originally towards nuclear matter near the saturation point, a further improvement of it may play an important role in getting a better agreement with the observations. By incorporating the properties at higher densities, e.g., the coupling constants for hyperons, it is possible to get cooling curves covering most of the observational data, as in Ref. \\cite{Page:2006}."
    },
    "1112/1112.4539_arXiv.txt": {
        "abstract": "The $k$d-tree is a fundamental tool in computer science. Among other applications, the application of $k$d-tree search (by the tree method) to the fast evaluation of particle interactions and neighbor search is highly important, since the computational complexity of these problems is reduced from $O(N^2)$ for a brute force method to $O(N \\log N)$ for the tree method, where $N$ is the number of particles. In this paper, we present a parallel implementation of the tree method running on a graphics processing unit (GPU). We present a detailed description of how we have implemented the tree method on a Cypress GPU. An optimization that we found important is localized particle ordering to effectively utilize cache memory. We present a number of test results and performance measurements. Our results show that the execution of the tree traversal in a force calculation on a GPU is practical and efficient. ",
        "introduction": "A technique for gravitational many-body simulations is a fundamental tool in astrophysical simulations because the gravitational force drives structure formation in the universe. The length scales that arise in structure formation range from less than 1 cm for the aggregation of dust to more than $10^{24}$ cm for the formation of cosmological structures. At all scales, gravity is a key physical process for the understanding of structure formation. The reason behind this is the long-range nature of gravity.  Suppose we simulate structure formation with $N$ particles. The flow of the many-body simulation is as follows. First, we calculate the mutual gravitational forces between the $N$ particles, then integrate the orbits for the $N$ particles, and repeat this process as necessary. Although it is simple in principle, the force calculation is a challenging task from the point of view of computer science. A simple, exact method for the force calculation requires $O(N^2)$ computational complexity, which is prohibitively compute-intensive for large $N$. An exact force calculation is necessary in some types of simulations, such as few-body problems, the numerical integration of planets orbiting around a star (e.g., the Solar System), and the evolution of dense star clusters. For simulations that do not require exact forces, however, several approximation techniques have been proposed \\citep{Hockney_1981, Barnes_1986, Greengard_1987}. The particle--mesh/particle--particle--mesh method \\citep{Hockney_1981} and the tree method \\citep{Barnes_1986} reduce the computational complexity of the force calculation to $O(N \\log N)$. The fast multipole method (FMM) reduces it further to $O(N)$ \\citep{Greengard_1987}. Of these methods, the tree method has been used extensively in astrophysical simulations, since its adaptive nature is essential for dealing with clumpy structure in the universe \\citep[e.g.,][]{Bouchet_1988}.  Despite the $O(N \\log N)$ complexity, computational optimization of the tree method by techniques such as vectorization and parallelization is necessary to accommodate demands for simulations with larger and larger $N$. \\cite{Hernquist_1990}, \\cite{Makino_1990}, and \\cite{Barnes_1990} have reported various techniques to vectorize the force calculation with the tree method. \\cite{Warren_1992}, \\cite{Dubinski_1996}, and \\cite{Yahagi_1999} have reported a parallel tree method for massively parallel processors (MPPs). In a recent publication \\citep{Springel_2005}, a simulation of large-scale structure formation in the universe with more than ten billion particles, using a parallel tree code running on an MPP, has been reported. Another computational technique to speed up the tree method utilizes the GRAPE special-purpose computer \\citep{Sugimoto_1990,Makino_1998}. Using a combination of vectorization techniques for the tree method, the tree method can be executed efficiently on a GRAPE system \\citep{Makino_1991}.  Cosmological simulations are a ``grand challenge\" problem. The Gordon Bell prizes have been awarded many times for cosmological simulations \\citep{Warren_1992a, Fukushige_1996, Warren_1997, Warren_1998, Kawai_1999, Hamada_2009}. In those simulations, both parallel tree codes \\citep{Warren_1992a, Warren_1997, Warren_1998} and a tree code running on a GRAPE system \\citep{Fukushige_1996, Kawai_1999} and a graphics processing unit (GPU) \\citep{Hamada_2009} were used to perform cosmological simulations.  In the present paper, we describe our implementation of the tree method on a GPU. The rise of the GPU forces us to rethink our way of doing parallel computing, since the performance of recent GPUs has reached the impressive level of $> 1$ Tflops. Acceleration techniques for many-body simulations with a GPU have already been reported \\citep[e.g.,][]{Nyland_2007, Portegies_2007, Hamada_2007,Belleman_2008}; however, these techniques have implemented an exact, brute force method with $O(N^2)$ complexity. It is apparent, however, that for applications that do not require exact forces, it is possible to do much more efficient computation by the tree method. We have directly implemented the tree method on a GPU so that we can enjoy the speed of an $O(N \\log N)$ algorithm. For small $N < 30,000$, the brute force method on a GPU is faster than the tree method owing to extra work concerning the tree data structure. However, our results show that the tree method significantly outperforms the brute force method on a GPU for $N \\gg 10,000$, which is the standard size in current astrophysical simulations is. Our code is simple and easy to extend to other numerical algorithms that require a neighbor list or a short-range force, such as algorithms for the smoothed particle hydrodynamics (SPH) method \\citep{Gingold_1977, Lucy_1977}. ",
        "conclusions": "In this paper, we have described our implementation of the tree method on a Cypress GPU. By transforming a recursive tree traversal into an iterative procedure, we have shown that the execution of a tree traversal together with a force calculation on a GPU can be practical and efficient. In addition, our implementation shows performance comparable to that of a recently reported FMM code on a GPU.  We can expect to get further performance gains by fully utilizing the four-vector SIMD operations of the SCs of the GPU. Moreover, since 10--20\\% of $T_{\\rm total}$ is spent on the tree construction, parallelization of this part on a multicore CPU will be an effective way to boost the total performance. Provided that we can easily extend our code to implement a force calculation for short-range interactions by a method such as the SPH method, we believe that a future extended version of our code will enable us to do a realistic astrophysical simulation that involves baryon physics with $N > 1,000,000$ very rapidly. It is fairly easy to incorporate higher-order multipole expansion terms into our method, and it would be a natural extension to the present work. Another good application of the tree method on a GPU would be to simulations that adopt a symmetrized Plummer potential \\citep{Saitoh_2010}. We believe that our method is the best for implementing that proposal, and hence that we shall certainly obtain better accuracy and good performance in simulating galaxy evolution and formation with different mass resolutions."
    },
    "1112/1112.3319_arXiv.txt": {
        "abstract": "We derive bounds on the mixing between the Standard Model (``active'') neutrinos and their right-chiral (``sterile'') counterparts in the see-saw models, by combining neutrino oscillation data and results of direct experimental searches.  We demonstrate that the mixing of sterile neutrinos with any active flavour can be significantly suppressed for the values of the angle $\\theta_{13}$ measured recently by Daya Bay and RENO experiments. We reinterpret the results of searches for sterile neutrinos by the PS191 and CHARM experiments, considering not only charged current but also neutral current-mediated decays, as applicable in the case of see-saw models.  The resulting \\emph{lower bounds} on sterile neutrino lifetime are up to an order of magnitude \\emph{stronger} than previously discussed in the literature.  % Combination of these results with the upper bound on the lifetime coming from primordial nucleosynthesis rule out the possibility that two sterile neutrinos with the masses between 10 MeV and the pion mass are solely responsible for neutrino flavour oscillations. We discuss the implications of our results for the Neutrino Minimal Standard Model (the $\\nu$MSM). ",
        "introduction": "\\label{sec:introduction}    Transitions between neutrinos of different flavours (see e.g.~\\cite{Strumia:06} for a review and Refs.~\\cite{Schwetz:08a,Schwetz:2011zk,Fogli:2011qn} for the recent update of experimental values) are among the few firmly established phenomena \\emph{beyond the Standard Model of elementary particles}. The simplest explanation is provided by the ``neutrino flavour oscillations'' -- non-diagonal matrix of neutrino propagation eigenstates in the weak charge basis. While the absolute scale of neutrino masses is not established, particle physics measurements put the sum of their masses below $2\\ev$~\\cite{PDG:11} while from the cosmological data one can infer an upper bound of $0.58\\ev$ at 95\\% CL ~\\cite{WMAP7}.  A traditional explanation of the smallness of neutrino masses is provided by the \\emph{see-saw mechanism}~\\cite{Minkowski:77,Ramond:79,Mohapatra:79,Yanagida:80}.  It assumes the existence of several \\emph{right-handed neutrinos}, coupled to their Standard Model (SM) counterparts via the Yukawa interaction, providing the Dirac masses, $M_\\dir$, for neutrinos. The Yukawa interaction terms dictate the SM charges of the right-handed particles: they turn out to carry no electric, weak and strong charges; therefore they are often termed ``singlet,'' or ``sterile'' fermions.  Sterile neutrinos can thus have Majorana masses, $M_\\m$, consistent with the gauge symmetries of the Standard Model. If the Majorana masses are much larger than the Dirac ones, the \\emph{type I seesaw formula}~\\cite{Minkowski:77,Ramond:79,Mohapatra:79,Yanagida:80} holds, expressing the mass matrix of observed neutrinos $(\\mathcal{M}_\\nu)$ via \\begin{equation} \\mathcal{M}_\\nu = - M_\\dir M_\\m^{-1} M_\\dir^T, \\label{see-saw} \\end{equation} where $\\mathcal{M}_\\nu$ is a $3\\times 3$ matrix of active neutrino masses, mixings, and (possible) CP-violating phases. The masses of sterile neutrinos are given by the eigenvalues of their Majorana mass matrix (with the corrections of the order $M_\\dir^2/M^2_\\m$). They are much heavier than the active neutrino masses as a consequence of~(\\ref{see-saw}).  Numerous searches for sterile neutrinos in the mass range up to $\\sim 100\\gev$ had been performed (see the corresponding section in Particle Data Group~\\cite{PDG:11},\\footnote{\\url{http://pdglive.lbl.gov/Rsummary.brl?nodein=S077&inscript=Y}} see also~\\cite{Atre:09} and refs. therein).  These searches provided upper bounds on the strength of interaction of these neutral leptons with the SM neutrinos of different flavours -- active-sterile neutrino \\emph{mixing angles} $\\vartheta_\\alpha^2 \\propto \\left|\\frac{M_{\\dir,\\alpha}}{M_\\m}\\right|^2$ for sterile neutrino with the mass $M_\\m$.\\footnote{Here and below we use the letter $\\vartheta$ for \\emph{active-sterile mixing angles} (defined by Eq.~(\\ref{eq:19}) below) while reserving $\\theta_{12},\\theta_{13}$ and $\\theta_{23}$ for the measured parameters of the active neutrinos matrix $\\mathcal{M}_\\nu$.  These quantities $\\vartheta_\\alpha$ are often denoted $|V_{4\\alpha}|^2$ or $|U_{x\\alpha}|^2$ in the experimental papers, to which we refer. Here and below the Greek letters $\\alpha,\\beta$ are flavour index $e,\\mu,\\tau$ and $i,j=1,2,3$ denote active neutrino mass eigenstates.}  These bounds then can be interpreted as \\emph{lower bounds} on the lifetime of sterile neutrinos $\\tau_s$ via \\begin{equation} \\label{eq:lifetime_expression} \\tau_s^{-1} = \\frac{G_F^2 M_\\m^5}{96\\pi^3}\\sum_X  \\vartheta_\\alpha^2 B_{X}^{(\\alpha)}, \\end{equation} where the sum runs over various particles to which sterile neutrino can decay, depending on their mass ($\\nu$, $e^\\pm$,$\\mu^\\pm$,$\\tau^\\pm$, $\\pi$, $K$, heavier mesons and baryons) and dimensionless quantities $B_X^\\alpha$ depend { on the branching ratios (see Appendix \\ref{app:lifetime} for details).  The lower bound on the lifetime $\\tau_s$ is usually dominated by the least constrained mixing angle, $\\vartheta_\\tau^2$ (as will be shown later).  This bound can be made \\emph{stronger} if one assumes that the same particles are also responsible for the neutrino oscillations. The see-saw formula~(\\ref{see-saw}) limits (at least partially) possible values of ratios of the mixing angles $\\vartheta_\\alpha^2/\\vartheta_\\beta^2$. In the simplest case when only two sterile neutrinos are present (the minimal number, required to explain two observed neutrino mass differences) the ratios of mixing angles varies within a limited range, see e.g.~\\cite{Gorbunov:07a,Shaposhnikov:08a}. While this range can be several orders of magnitude large (owing to our ignorance of certain oscillation parameters, such as e.g. a CP-violating phase~\\cite{Gorbunov:07a,Shaposhnikov:08a}), the implied (lower) bounds on the lifetime become much stronger, essentially being determined by the \\emph{strongest}, rather than the weakest direct bound on $\\vartheta_\\alpha$.} When confronted with the upper bound from Big Bang Nucleosynthesis~\\cite{Dolgov:00a,Dolgov:00b}, they seem to close the window of parameters for sterile neutrinos with the mass lighter than about $150\\mev$~\\cite{Gorbunov:07a,Boyarsky:09a}. It was argued in~\\cite{Asaka:2011pb}, however, that in the case of normal hierarchy there can be a small allowed window of parameters of sterile neutrinos with the mass below the pion mass.  In this paper we reanalyze restrictions on sterile neutrino lifetime in view of the recent results of the Daya Bay~\\cite{An:2012eh} and RENO~\\cite{Ahn:2012nd} collaborations, that measured a non-zero mixing angle $\\theta_{13}$ (see also~\\cite{Adamson:2011qu,Abe:2011sj}).  We demonstrate that in the case when there are only two sterile neutrinos, responsible for the observed neutrino oscillations, the oscillation data allow for such a choice of the active-sterile Yukawa couplings that the mixing of sterile neutrinos with any given flavour can be strongly suppressed. This happens \\emph{only} for a non-zero values of $\\theta_{13}$, in the range consistent with the current measurements~\\cite{Schwetz:08a,Fogli:2011qn,An:2012eh,Ahn:2012nd}.  We also confront our results with the recently reanalyzed bounds from primordial nucleosynthesis~\\cite{Ruchayskiy:2012si} and show that the window in the parameter space of sterile neutrinos with masses $10\\MeV \\lesssim M_s \\lesssim 140\\MeV$ discussed in previous works~\\cite{Asaka:2011pb} (see also~\\cite{Gorbunov:07a}) is closed. For larger masses the window remains open. The results of this paper partially overlap with~\\cite{Asaka:2011pb} (also~\\cite{Gorkavenko:2009vd}), and we compare in the corresponding places to the previous works.  The paper is organized as follows: in Section~\\ref{sec:numsm} we briefly describe the type I see-saw model we use. We then investigate the relations between different mixing angles imposed by the see-saw mechanism and demonstrate that the mixing with any flavour $\\vartheta_\\alpha^2$ can become suppressed (Section~\\ref{sec:solution-see-saw}). Section~\\ref{sec:experiments} is devoted to the overview of the experiments, searching for sterile neutrinos with the masses below $2\\GeV$, and the way one should interpret their results to apply to the see-saw models that we study. Section~\\ref{sec:discussion} summarizes our revised bounds on mixing angles and translates them into the resulting constraints on sterile neutrino lifetime (Figs.~\\ref{fig:lifetime_result}). We conclude in Section~\\ref{sec:conclusion}, discussing implications of our results and confronting them with the bounds from primordial nucleosynthesis. ",
        "conclusions": "In this Section we summarize our results: the upper bound on the (combination of) mixing angles of sterile and active neutrinos in the see-saw models~(\\ref{eq:lifetime_expression}) in the range 10~MeV -- 2~GeV and the lower bound on sterile neutrino lifetime, obtained in combination of these bounds with constraints, coming from neutrino oscillation data.      \\begin{figure}[htp] \\centering % \\subfloat [Normal hierarchy, mass range $10\\mev - 2\\gev$] {\\includegraphics[width=.5\\textwidth]{figures/lifetime_NH+MI+NC-largemass-revDaya}}% ~\\subfloat [Normal hierarchy, zoom at the mass range $10\\mev - 140\\mev$] % {\\includegraphics[width=.5\\textwidth]{figures/lifetime_NH+MI+NC-smallmass-revDaya}}\\\\ \\subfloat% [Inverted hierarchy, mass range $10\\mev- 2\\gev$] { \\includegraphics[width=.5\\textwidth]{figures/lifetime_IH+MI+NC-largemass-revDaya} }% \\subfloat [Inverted hierarchy, zoom at the mass range $10\\mev - 140\\mev$] { \\includegraphics[width=.5\\textwidth]{figures/lifetime_IH+MI+NC-smallmass-revDaya}% } \\caption{The resulting lower bounds on sterile neutrino lifetime $\\tau_s$ as a function of their mass, obtained by requiring that two Majorana sterile neutrinos are responsible for neutrino oscillations and their parameters do not contradict the negative results of direct experimental searches. In all figures the upper curve comes from using of the best fit neutrino oscillation parameters, the middle one -- from their variation within the $3\\sigma$ limits, and the lower one does not take into account neutrino oscillation data and puts all three mixing angles equal to their direct experimental bounds. The dashed line for NH corresponds to the best-fit values of PMNS parameters with $\\theta_{13} = 0$ and shows how much the bounds on the lifetime relax for non-zero value of $\\theta_{13}$ (see text, Section~\\protect\\ref{sec:discussion} for discussion). } \\label{fig:lifetime_result} \\end{figure}    \\subsection{Bounds on the mixing angles of sterile neutrinos} \\label{sec:bounds-mixing}  For the models~(\\ref{eq:nuMSM_Lagrangian}) (two Majorana sterile neutrinos, interacting through both charged and neutral interactions), the compilation of constraints on various combinations of active-sterile mixing angles ($\\vartheta_e^2$, $\\vartheta_\\mu^2$, $\\vartheta_e \\sqrt{\\sum c_\\alpha \\vartheta_\\alpha^2}$, $\\vartheta_\\mu \\sqrt{\\sum c_\\alpha \\vartheta_\\alpha^2}$) that we used in this work are plotted in Figs.~\\ref{fig:PS191-bounds} and~\\ref{fig:mixing_constraints}.\\footnote{Notice that in the published results of the PS191 experiment~\\cite{Bernardi:1987ek} bounds are given up to $M_\\m=400\\MeV$. We extend these bounds up to $450\\MeV$, using the PhD Thesis of J.-M. Levy \\cite{Levy}.}    \\subsection{The lower bound on the lifetime of sterile neutrinos} \\label{sec:bounds-lifetime}  The result of the Sections~\\ref{sec:normal-hierarchy}--\\ref{sec:inverted-hierarchy}, combined with these experimental bounds can be translated into the \\emph{lower} limits on the lifetime of sterile neutrinos. These results are presented in Figs.~\\vref{fig:lifetime_result}.  Additionally, we plot the lifetime bounds for the best-fit values of the PMNS parameters yet with $\\theta_{13}=0$ (as used e.g. in~\\cite{Gorbunov:07a,Canetti:10a}). For normal hierarchy we see that our bounds with $\\theta_{13} \\neq 0$ are relaxed by as much as the order of magnitude at some masses, compared to $\\theta_{13}=0$ case. The difference for IH is not so pronounced. \\emph{Notice}, that the bounds of~\\cite{Gorbunov:07a,Canetti:10a,Asaka:2011pb} were different from what we show as dashed line in Fig.~\\ref{fig:lifetime_result} because of ignoring the neutral current contributions to the results of PS191 experiment (for details see discussion in Section~\\ref{sec:experiments} and Figs.~\\ref{fig:lifetime-interpretations-peaks},~\\ref{fig:lifetime-CC+CCNC-prev-studies})."
    },
    "1112/1112.3543_arXiv.txt": {
        "abstract": "{The earliest stages of high-mass star formation are still poorly characterized. Densities, temperatures and kinematics are crucial parameters for simulations of high-mass star formation. It is also unknown whether the initial conditions vary with environment.} {We want to investigate the youngest massive gas clumps in the environment of extremely active star formation. } {We selected the IRDC\\,18454 complex, directly associated with the W43 Galactic mini-starburst, and observed it in the continuum emission between 70\\,$\\mu$m and 1.2\\,mm with Herschel, APEX and the 30\\,m telescope, and in spectral line emission of N$_2$H$^+$ and $^{13}$CO with the Nobeyama 45\\,m, the IRAM 30\\,m and the Plateau de Bute Interferometer.} {The multi-wavelength continuum study allows us to identify clumps that are infrared dark even at 70\\,$\\mu$m and hence the best candidates to be genuine high-mass starless gas clumps. The spectral energy distributions reveal elevated temperatures and luminosities compared to more quiescent environments. Furthermore, we identify a temperature gradient from the W43 mini-starburst toward the starless clumps. We discuss whether the radiation impact of the nearby mini-starburst changes the fragmentation properties of the gas clumps and by that maybe favors more high-mass star formation in such an environment. The spectral line data reveal two different velocity components of the gas at 100 and 50\\,km\\,s$^{-1}$. While chance projection is a possibility to explain these components, the projected associations of the emission sources as well as the prominent location at the Galactic bar -- spiral arm interface also allow the possibility that these two components may be spatially associated and even interacting.} {High-mass starless gas clumps can exist in the close environment of very active star formation without being destroyed. The impact of the active star formation sites may even allow for more high-mass stars to form in these 2nd generation gas clumps. This particular region near the Galactic bar -- spiral arm interface has a broad distribution of gas velocities, and cloud interactions may be possible.} ",
        "introduction": "\\label{intro}  The initial conditions required to form massive stars is one of the major topics in high-mass star formation research today. The advent of new observational capabilities often opened the window to earlier and earlier evolutionary stages in that field. While the IRAS satellite combined with cm continuum surveys of the galactic plane revealed the populations of ultracompact H{\\sc ii} regions and high-mass protostellar objects (HMPOs, e.g., \\citealt{wc89,wc1989b,kurtz1994,molinari1996,sridha}), the advent of the mid-infrared satellites ISO, MSX and Spitzer allowed to access even earlier evolutionary stages, namely the infrared dark clouds (IRDCs, e.g., \\citealt{perault1996,egan1998,carey2000,sridharan2005,simon2006,peretto2009}). However, it turned out that most of the so far studied IRDCs host weak 24\\,$\\mu$m emission sources and drive molecular outflows, both strong indicators for star formation activity (e.g., \\citealt{beuther2005d,beuther2007a,beuther2007g,rathborne2006,motte2007,cyganowski2008}). Therefore, we are still lacking a thorough understanding of the initial conditions for high-mass star formation prior to any star formation activity. The advent of the Herschel far-infrared satellite \\citep{A&ASpecialIssue-HERSCHEL} now offers the unique opportunity to identify targets for these initial conditions and to study their physical and chemical properties in detail. Major characteristics of these earliest evolutionary stages are, that the massive gas cores with masses in excess of several 100\\,M$_{\\odot}$ observed at (sub)mm wavelengths are still dark at 70\\,$\\mu$m, hence they have not formed any protostellar objects to Herschel sensitivity limits yet.  The Herschel guaranteed time key project EPOS (Early Stages of Star Formation, P.I.~O.~Krause) is a dedicated observation campaign targeting 45 IRDCs that are promising candidates to search for and to study these initial conditions. Early results from the first data revealed already exciting results, e.g., we identified candidate high-mass starless cores (HMSCs) as well as very weak 70\\,$\\mu$m sources within IRDCs that may be the very first manifestation of high-mass star formation \\citep{beuther2010b,henning2010,linz2010}.  \\begin{figure*}[htb] \\includegraphics[width=0.99\\textwidth]{17850f1.jpg} \\caption{Compilation of the continuum data from 70\\,$\\mu$m to 1.2\\,mm wavelength as labeled in each panel. The top-row presents the data on a smaller spatial scale as marked in the bottom-left panel. The color-scale is chosen in each image individually to highlight the most important features. Contour levels of the 870\\,$\\mu$m data start at the 3$\\sigma$ levels of 0.27\\,mJy\\,beam$^{-1}$ and continue in in 3$\\sigma$ steps to 2.7\\,Jy\\,beam$^{-1}$, from where they continue in 2.7\\,Jy\\,beam$^{-1}$ steps. The 1.2\\,mm data are contoured in 3$\\sigma$ levels of 30\\,mJy\\,beam$^{-1}$. The numbers in the top-left panel label the sub-cores, the white small box there marks the region shown in Fig.~\\ref{18454_spitzer}, and a scale-bar for a distance of 5.5\\,kpc is shown as well. The crosses in the bottom-left panel mark the positions of a double-component radio recombination line H{\\sc ii} region from \\citet{anderson2011}, the W43-MM1 position from \\citet{motte2003} as well as the approximate center of the Wolf-Rayet/OB cluster \\citep{blum1999,motte2003}.} \\label{spire_pacs} \\end{figure*}  \\paragraph{The IRDC\\,18454 close to W43:} One particularly interesting region within the sample are the IRDCs associated with the IRAS source IRAS\\,18454-0158. Henceforth, we will call the region IRDC\\,18454. While the IRAS source was first studied in the framework of the HMPO survey by \\citet{sridha} and \\citet{beuther2002a}, it was soon recognized that several of the sub-sources are indeed infrared dark \\citep{sridharan2005}. Another curiosity of that region is that it harbors two distinctly different velocity components, one around 50 and one around 100\\,km\\,s$^{-1}$ \\citep{sridharan2005,beuther2007g}. It is also interesting to note that only the 100\\,km\\,s$^{-1}$ component shows SiO emission, indicative of a molecular outflow, whereas we do not detect any SiO from the 50\\,km\\,s$^{-1}$ component \\citep{beuther2007g}. Furthermore, the measured H$^{13}$CO$^+$(1--0) line-width from the 100\\,km\\,s$^{-1}$ component is also broader than that at 50\\,km\\,s$^{-1}$ (2.7 versus 1.7\\,km\\,s$^{-1}$, \\citealt{beuther2007g}), also indicating that the 50\\,km\\,s$^{-1}$ component is in a less turbulent state. The kinematic distances derived by \\citet{beuther2007g} for the 50 and 100\\,km\\,s$^{-1}$ components, using the \\citet{brand1993} rotation curve were 3.5 and 6.4\\,kpc, respectively. Applying the new rotation curve by \\citet{reid2009}, we now get kinematic distances of 3.3 and 5.5\\,kpc for both components, respectively. The 5.5\\,kpc distance corresponds well to the distance derived for the large H{\\sc ii} region W43 by \\citet{wilson1970}. A similar distance of $\\sim 6$\\,kpc was also recently derived for the nearby red supergiant cluster RSGC3 \\citep{negueruela2011}). We will discuss later whether these are really two different components just chance-projected together on the plane of the sky, or whether it may be interacting cloud components at similar distances.  In the direct vicinity of that region we find the large Galactic mini-starburst region W43 (Fig.~\\ref{spire_pacs}), which was already studied in detail in the past with a broad wavelength coverage (e.g., \\citealt{smith1978,lester1985,blum1999,motte2003,nguyen2011}) including also early Herschel results \\citep{bally2010,elia2010}. The IRDC\\,18454 is located at the northeastern edge of the Z-shaped filament first discussed in \\citet{motte2003} (see also Fig.~\\ref{spire_pacs}). This region is very luminous with $L\\sim 3\\times 10^6$\\,L$_{\\odot}$, and it contains sources at different evolutionary stages from a Wolf-Rayet cluster to active young star-forming cores (e.g., \\citealt{blum1999,motte2003}). \\citet{nguyen2011} recently discussed several broad velocity components in that regions and suggested that this whole complex might have formed via converging gas flows due to its special location at the interface of the Galactic bar with the inner Scutum spiral arm \\citep{benjamin2005,lopez2007,rodriguez2008}. Different velocity components were also identified in atomic HI emission and absorption \\citep{liszt1993}. However, our target region IRDC\\,18454 which hosts supposedly the earliest evolutionary stages, has not been subject to a detailed investigation yet.  Among others, we like to address the following questions: Do we identify bona-fide high-mass starless cores at all? If yes, what are their physical properties, e.g., their temperatures and turbulent line-widths? Are the two velocity components spatially associated or are they rather chance alignments along the line of sight? Does the mini-starburst W43 impose any impact on the earliest evolutionary stages or are the young cores evolving independently of that? Do we find any signature of triggering? ",
        "conclusions": "\\label{discussion}  \\subsection{Luminous starless clumps} \\label{luminous}  The measured high luminosities from the starless clump candidates (Fig.\\ref{seds}), in particular clumps 1, 3 and 8a are intriguing since -- down to our detection limits -- they do not have an internal radiation sources.  Also other sources like clump 2 are interesting because this region is still dark at 24\\,$\\mu$m and only becomes an emission source from 70\\,$\\mu$m onwards. Hence, it has to be at a very young evolutionary stage as well. This implies that the bolometric luminosities of these regions must be largely caused by external radiation sources of which the prime candidates are the diverse sources comprising the neighboring mini-starburst W43 (see Introduction or, e.g., \\citealt{motte2003,nguyen2011}). Do we find differences between IRDCs in the vicinity of very active massive star-forming regions and more isolated IRDCs?  For comparison, we select the IRDC\\,18223 where luminosities and temperatures were measured in a similar fashion via SED fits of Herschel data \\citep{beuther2010b}. The derived temperatures for the clumps in IRDC\\,18454 are slightly elevated compared to IRDC\\,18223 (between 19 and 20\\,K for IRDC\\,18454, and between 16 and 18\\,K in IRDC\\,18223), and the measured bolometric luminosities in IRDC\\,18454 are factors between 4 and 10 larger than found in IRDC\\,18223. While the observed aperture size and spatial scale directly translates into different luminosities (Fig.~\\ref{seds}), the smaller aperture of $19.2''$ (corresponding to $\\sim$0.5\\,pc) used for IRDC\\,18454 is even smaller than that used for IRDC\\,18223 ($\\sim$0.64\\,pc). Therefore, spatial scale arguments cannot account for that difference. Another important parameter for the total bolometric luminosity is the available gas mass. While clumps 1 and 3 are factors of a few more massive than IRDC\\,18223-s1 and IRDC\\,18223-s2, clump 8 presented here and IRDC\\,18223-s2 have approximately the same mass (140 and 170\\,M$_{\\odot}$, respectively).  Nevertheless, the luminosity of clump 8 is about a factor 4 larger than that of IRDC\\,18223-s2.  It is no surprise that the temperature differences between IRDC\\,18223 and IRDC\\,18454 are less strong than the luminosity differences because the luminosity has a strong exponential dependence on the temperature (Stefan-Boltzmann law $L \\propto T^4$).  While all these parameters have to be taken a bit cautiously since the associated errors are large (e.g., for the SEDs the uncertain background subtraction may contribute significantly), in principle the errors should be comparable for the different regions because we applied the same techniques. Therefore, it is tempting to speculate what may cause the measured luminosity and temperature differences between infrared dark clumps in the vicinity of the active star-forming region W43 compared to the more quiescent region IRDC\\,18223. It could be possible that the neighboring mini-starburst injects considerable energy even into the starless clumps that the total bolometric luminosities are increased above their expected more quiescent values. To first order, with such a luminosity increase one also expects a temperature increase within the gas clumps. Although the measured temperature difference between the regions is less strong than the luminosity increase, nevertheless, we see elevated temperatures in IRDC\\,18454 compared to IRDC\\,18223. Considering the large apertures used for the SEDs, the derived average temperatures may still be dominated by the average temperature of the general ISM that is in the same range (e.g., \\citealt{reach1995}).  We can also quantitatively estimate the influence of the total bolometric luminosity of the Wolf-Rayet/OB cluster ($\\sim 3.5\\times 10^6$\\,L$_{\\odot}$) on the gas and dust environment at the typical projected separations from the target IRDCs (Fig.~\\ref{spire_pacs}). Assuming that all radiation of the central cluster is reprocessed as far-infrared radiation, and using an approximate projected separation of the infrared dark clumps of 11\\,pc, as well as Gaussian source FWHM of $36.6''$ or $19.2''$, respectively (Sec.~\\ref{sedsection} and Fig.~\\ref{seds}), the received luminosity at the location of the IRDCs are on the order of 2500 and 700\\,L$_{\\odot}$, respectively. Comparing these values with the measured luminosities in the same given apertures (Fig.~\\ref{seds}), elevated luminosities caused by the nearby luminous cluster appear clearly feasible. Furthermore, one can estimate the approximate required interstallar radiation field to account for the observed luminosity. A clump with a surface corresponding to a radius $19.2''/2$ has to intercept an interstellar radiation field of $\\sim$306 (in units of the Habing field $G_0=1.6\\times 10^{-3}$\\,erg\\,s$^{-1}$cm$^{-2}$, \\citealt{habing1968}) to exhibit a luminoisty of 1000\\,L$_{\\odot}$.  In comparison, if one just accounts for geometric dillution, the $\\sim 3.5\\times 10^6$\\,L$_{\\odot}$ of the Wolf-Rayet/OB cluster correspond at a distance of 11\\,pc to $\\sim$588 times the interstellar radiation field. Subtracting additional scattering and blockade by interfering gas clumps, these two values agree well with each other. Based on kinematic arguments, \\citet{motte2003} also discussed the influence of the luminous cluster on the environment, and they similarly infer that there should be considerable influence from that first generation of high-mass star formation on its environment.  In  addition  to  the   individual   SED fits   presented  in  section \\ref{sedsection} and Figure \\ref{seds},  one can also smooth  all maps between 70  and  870\\,$\\mu$m to  the  same  spatial  resolution of the 500\\,$\\mu$m   data  and   fit   the SEDs pixel  by    pixel  (see also \\citealt{stutz2010} for  details  of the  fitting).  This  large-scale fitting approach allows us to infer the dust temperature structure and gradients across the whole field. Figure \\ref{t_herschel} presents the results.  As expected, we identify  a clear temperature gradient  from more than 35\\,K toward  the W43 main  complex down to low values below 20\\,K in the east  of the region. In  particular, we see that  many of the  submm continuum   structures   correlate  well  with the   lowest temperatures.  This is true  for the dark   clumps of this analysis as well as for the filament  going approximately in northsouth  direction at the eastern edge of  the map. The large-scale east-west temperature gradient throughout   the   region is  consistent   with   significant influence of the   W43 complex  on  the here  discussed infrared  dark high-mass starless clump candidates.  \\begin{figure}[htb] \\includegraphics[width=0.49\\textwidth]{17850f11.jpg} \\caption{The color-scale shows a temperature map in K derived for the whole region from the combined Herschel and ATLASGAL data between 70 and 870\\,$\\mu$m, all smoothed to the spatial resolution of the 500\\,$\\mu$m data. Contour levels of the overlaid 870\\,$\\mu$m data start at the 3$\\sigma$ levels of 0.27\\,mJy\\,beam$^{-1}$ and continue in in 3$\\sigma$ steps to 2.7\\,Jy\\,beam$^{-1}$, from where they continue in 2.7\\,Jy\\,beam$^{-1}$ steps. Our target clumps are marked.} \\label{t_herschel} \\end{figure}  Since the Jeans length and Jeans mass as indicators of the fragmentation properties of star-forming gas clumps are proportional to the temperature ($\\lambda_J \\propto T^{1/2}$ and $M_J \\propto M^{3/2}$), increases in luminosity and temperature of young gas clumps could change their fragmentation properties to preferentially form larger and more massive fragments. Although still speculative, such a scenario could imply that nearby active star formation does not only have a destructive character destroying pristine gas clumps, but that the energy input of close-by active regions could even elevate the ability to form massive stars in a triggered fashion in the immediate environment of the earlier generations of stars. While the energy input from a first generation of low-mass stars appears to be too little to halt fragmentation \\citep{longmore2011}, the influence of the nearby luminous high-mass star-forming region may be more important.  \\subsection{Colliding clouds or chance alignment?} \\label{converge}  As indicated in sections \\ref{kinematic1} \\& \\ref{kinematic2}, the two different velocity components at 100 and 50\\,km\\,s$^{-1}$ may either be chance projections along the line of sight, or they could be associated and potentially interacting cloud complexes. Based mainly on large-scale $^{13}$CO(1--0) data from the Galactic Ring survey (GRS, \\citealt{jackson2006}), \\citet{nguyen2011} suggest that the velocities associated directly with W43 range between 60 and 120\\,km\\,s$^{-1}$, whereas lower velocity components are unlikely part of the W43 complex.  \\begin{figure}[htb] \\includegraphics[width=0.49\\textwidth]{17850f12.jpg} \\caption{IRAM\\,30m $^{13}$CO(2--1) observations toward a larger part of W43 and IRDC\\,18454. The color-scale shows the $^{13}$CO emission around 53\\,km\\,s$^{-1}$ whereas the contours present gas at 93\\,km\\,s$^{-1}$. The stars mark the positions of our clumps 1 to 8, a scale-bar is shown at the bottom-left.} \\label{13co} \\end{figure}  In contrast to this, the spatial structure of the two velocity components as presented in Figs.~\\ref{n2hplus} and \\ref{18454-1_100_50} indicates that the two components may also be spatially associated and interacting. Also the fact that the mm fluxes of clump 1 and 2 are so similar could be taken as circumstantial support for this idea.  To further investigate whether we find additional signatures of associated clouds in that region, we investigated the $^{13}$CO(2--1) data of a recent IRAM 30\\,m large program to study the entire W43 complex (PI F.~Motte). The whole project, the data and more detailed analysis will be presented in forthcoming papers by Nguyen-Luong et al.~and Carlhoff et al., here we only show a small sub-set where two selected velocity components are compared. Figure \\ref{13co} presents an overlay of the $^{13}$CO(2--1) emission around 50 and 100\\,km\\,s$^{-1}$ over a larger area but also encompassing our region of interest. The two velocity components appear in projection as a layered structure where one velocity component mainly emits at locations that are devoid of emission from the other velocity component. Unfortunately, this is also no unambiguous proof for associated or interacting clouds.  However, the spatial alignment of different velocity components for molecules with varying critical densities over a broad range of spatial scales is suggestive for the idea of connected and interacting structures.  Additional interesting information about different velocity components in this region is found for other tracers and wavelengths as well. For example, CH$_3$OH Class {\\sc ii} maser were found in this region at both velocities (e.g., \\citealt{menten1991}). Furthermore, the recent H{\\sc ii} region radio recombination line survey of \\citet{anderson2011} detected toward 21 out of 23 targeted position within 0.5 degrees from W43 two velocity components as well (one of the positions is marked in Fig.~\\ref{spire_pacs}). While projection effects of molecular and atomic gas along the line of sight through the Galaxy are well known, having such a large number of multiple-velocity recombination line spectra is more surprising since one does not expect that many projected H{\\sc ii} regions along each line of sight. Hence interacting clouds and H{\\sc ii} regions are possible. Furthermore, it is interesting to note that the largest number of H{\\sc ii} regions in our Galaxy are also found toward Galactic longitudes around $l\\approx \\pm30$\\,degrees \\citep{wilson1970,lockman1979,anderson2011}, hence again including the W43/IRDC\\,18454 complex. All this evidence suggests special environmental physical properties in this region that may affect the star formation processes.  How could such interacting velocity components be explained? The whole W43-IRDC\\,18454 complex is located at approximately a Galactic longitude of 30.85 degrees where the Galactic bar ends and the Scutum spiral arm starts (for a recent description of the different Galactic components, see e.g., \\citealt{churchwell2009,nguyen2011}). Hence, one could fathom interaction between Galactic bar and spiral arm structures, e.g., streaming motions between both Galactic components. Following the Galaxy modeling by \\citet{bissantz2003}, the inner spiral arms connecting to the bar and the bar should co-rotate.  In this framework, the 100\\,km\\,s$^{-1}$ component should be associated with the beginning of the Scutum arm (see also \\citealt{vallee2008}). Spiral arm velocities for the Sagittarius and Perseus arm are, according to \\citet{vallee2008}, around 70 and 25\\,km\\,s$^{-1}$, respectively.  However, there is no clear association with a 50\\,km\\,s$^{-1}$ gas component. Inspecting the Galactic plane CO survey by \\citet{dame2001}, at a Galactic longitude of $\\sim 30.85$ degrees one also find strong emission at 100\\,km\\,s$^{-1}$ and only less prominent features at 50\\,km\\,s$^{-1}$. Therefore, from the spiral arm structure of our Galaxy, it is not obvious to what distance one can associate the 50\\,km\\,s$^{-1}$ component. Furthermore, \\citet{nguyen2011} associate even 60\\,km\\,s$^{-1}$ gas with W43 and not so much with the Sagittarius arm as one would expect from the Galaxy models (e.g., \\citealt{vallee2008}). Although we cannot firmly establish whether the 100 and 50\\,km\\,s$^{-1}$ gas components are spatially associated or just chance alignments, the close projected association of them as well as the extremely busy and active nature of that part of the Galaxy at the interface of the Galactic bar with the inner spiral arm, where streaming motions between both components are likely to occur, makes spatial association and interaction of different cloud components here a fair possibility. See also \\citet{anderson2011} for a similar discussion of the origin of the ionized gas.  One way to distinguish the different explanation is to derive accurate distances to the different components. And the best way to do that today is parallax measurements of maser sources (e.g., \\citealt{reid2009}). The VLBI now has a very large project to do that for many maser sources throughout the Galactic plane (PI M.~Reid), and the masers toward the W43 complex are part of that program (K.~Menten, priv.~comm.). Therefore, we expect to get accurate distances for the different velocity components in the coming years. That should then allow us to answer the question whether these different velocity components are indeed spatially associated or whether they are just chance projections along the lines of sight."
    },
    "1112/1112.4363_arXiv.txt": {
        "abstract": "Using the gamma-ray bursts simultaneously detected by \\textit{Swift}/BAT and \\textit{Fermi}/GBM we performed a joint spectral and temporal analysis of the prompt emission data and confirm the rough correlation between the BAT-band photon index $\\Gamma^{\\rm BAT}$ and the peak spectral energy $E_{\\rm peak}$. With the redshift known sub-sample, we derived the isotropic gamma-ray energy $E_{\\rm \\gamma,iso}$ and also confirm the $E_{\\rm \\gamma,iso}-E_{\\rm peak,rest}$ relation, with a larger scatter than the Amati sample but consistent with GBM team analyses. We also compare the $T_{90}$ values derived in the GBM band with those derived in the BAT band and find that for long GRBs the BAT $T_{90}$ is usually longer than the GBM $T_{90}$, while for short GRBs the trend reverses. This is consistent with the soft/hard nature of long/short GRBs and suggests the importance of an energy-dependent temporal analysis of GRBs. ",
        "introduction": "Since the launch of the \\textit{Fermi} satellite in 2008, a significant number of gamma-ray bursts (GRBs) have been observed by both the \\fermi Gamma-ray Burst Monitor (GBM, 8 keV-40 MeV) and \\swift \\cite{gehrels04}. Since the energy band of \\textit{Swift} BAT is narrow (15-150 keV) and the peak of the spectrum is often outside the observed energy window \\cite{sakamoto09}, the joint BAT-GBM sample is valuable for \\textit{Swift} science by providing important information about the prompt emission, especially the characteristic peak energy of the $\\nu F_{\\nu}$ spectrum, $E_{\\rm peak}$. Using the \\textit{Fermi}/GBM-\\textit{Swift}/BAT joint sample we performed a two-part analysis on the prompt emission data.  First, a joint spectral analysis where we fit the data from both detectors and investigated several empirical correlations.  Second,  a temporal analysis where we rigorously derived $T_{90}$ in the 8-1000 keV band and compared it to the published BAT values.  A variety of empirical relations have been discussed in the literature \\citep{amati02,amati03,sakamoto04,sakamoto06,ghirlanda04,liangzhang05,yonetoku04}, many of which utilize $E_{\\rm peak}$. Since $E_{\\rm peak}$ cannot be well measured for most \\swift GRBs, efforts have previously been made to look for indicators of $E_{\\rm peak}$ based on the available observed quantities. In particular, the effective photon power law index in the BAT band, $\\Gamma^{\\rm BAT}$, has been found to be broadly correlated with $E_{\\rm peak}$. Two versions of this relation are found in the literature: one presented in Sakamoto et al. (2009) using simulations of BAT spectra and $E_{\\rm peak}$ measurements from broad-band detectors (e.g. Konus-Wind), and another presented in Zhang et al. (2007b) with $E_{\\rm peak}$ estimates based on the hardness ratio in the BAT band itself (Zhang et al. 2007a; see also Cui et al. 2005). In \\S3.2 we re-calibrate this correlation with the GBM-BAT joint sample. The $E_{\\rm \\gamma,iso} - E_{\\rm peak,rest}$ relation \\citep[or Amati relation;][]{amati02,amati03} is one of the most widely studied, and hotly debated, empirical correlations that may connect to GRB physics  \\citep[e.g.][]{band05,nakar05,kocevski11}. In \\S3.3 we test the Amati relation with the redshift-known sub-sample of the GBM-BAT joint sample.  Finally, traditionally long and short GRBs were defined in the BATSE band with a separation of about 2 seconds in the observed frame (Kouveliotou et al. 1993).  When applied to \\swift GRBs and combined with multi-wavelength afterglow data,  this definition creates confusion when used for GRB classification (e.g. Zhang et al. 2007b, 2009). Since the energy band of GBM is similar to that of BATSE, a comparison between the measured $T_{90}$ durations in the two detectors is of great interest and is performed in \\S4. ",
        "conclusions": "Using the data available from \\fermi and \\swift we performed an analysis of the jointly detected GRB sample. The spectral analysis indicates that fitting the joint spectra shows consistency with values reported in the literature based solely on GBM data  \\cite{nava11,bissaldi11}, and is able to reproduce a variety of empirical relations presented in the literature, including the $\\Gamma^{\\rm BAT}-E_{\\rm peak}$ and Amati relations.  The joint fits do not, in general, give substantially different values for the spectral parameters from the GBM-only analyses.  The updated sample of peak energies shows agreement with the relation between the PL slope of the BAT spectrum and $E_{\\rm peak}$ derived previously by Sakamoto et al. (2009) and Zhang et al. (2007b) and indicates the relation's ability to estimate the value of the observed $E_{\\rm peak}$ from a fit to the rather narrow \\swift band.  The relation shows a fair amount of scatter, however, and we caution against using it for robust measurements of $E_{\\rm peak}$, especially when compared to spectra from \\fermi or other missions that often contain the spectral peak within the energy range of the detectors.  It may be suited, for example, to give a rough estimate of the values of $E_{\\rm peak}$ for \\swift bursts that have no additional observations.  The redshift known subsample of bursts reproduces the general relationship in the $E_{\\rm iso}-E_{\\rm peak,rest}$ plane and further confirms the conclusions of Gruber et al. (2011) that show generally harder spectra and a larger scatter around the relation than previously believed.  Li (2007) proposed the argument that the $E_{\\rm iso}-E_{\\rm peak,rest}$ relation might evolve with redshift.  We believe that the sample of bursts with known redshift, especially the subsample of joint bursts, is still too small to test such a claim.  The necessary binning to test the hypothesis would thin out the sample sufficiently to make firm conclusions tenuous at best.  Finally, the interesting discrepancy between the $T_{90}$ derived in the GBM and BAT detectors confirms the soft/hard nature of long/short GRBs, respectively, and suggests the need for a detailed energy-dependent temporal analysis of GRBs."
    },
    "1112/1112.1400.txt": {
        "abstract": "It was recently shown that our peculiar velocity $\\boldsymbol{\\beta}$ with respect to the CMB induces mixing among multipoles and off-diagonal correlations at all scales which can be used as a measurement of $\\boldsymbol{\\beta}$, which is independent of the standard measurement using the CMB temperature dipole. The proposed techniques rely however on a perturbative expansion which breaks down for $\\,\\ell \\gtrsim 1/\\beta \\approx 800$. Here we propose a technique which consists of deboosting the CMB temperature in the time-ordered data and show that it extends the validity of the perturbation analysis multipoles  up to $\\,\\ell \\sim 10000$. We also obtain accurate fitting functions for the mixing between multipoles valid in a full non-linear treatment. Finally we forecast the achievable precision with which these correlations can be measured in a number of current and future CMB missions. We show that Planck could measure the velocity with a precision of around $60$ km/s, ACTPol in 4 years  around $40$ km/s, while proposed future experiments could further shrink this error bar by over a factor of around 2. ",
        "introduction": "\\label{sec:intro}   The dipole of the CMB is measured to be much larger than the other multipoles and this is usually attributed to a Doppler effect due to our peculiar velocity $\\boldsymbol{\\beta}$ with respect to the CMB rest frame. Under this assumption we can infer, by combining the measured temperatures of WMAP dipole~\\cite{Hinshaw:2008kr} with the COBE monopole~\\cite{Lineweaver:1996xa,Mather:1998gm}, its direction to be $\\, l = 263.99^\\circ \u00b1 0.14^\\circ$, $b= 48.26^\\circ \u00b1 0.03^\\circ\\,$ in galactic coordinates, and its modulus to be $\\beta \\equiv |\\boldsymbol{\\beta}| = (1.231 \\pm 0.003) \\times 10^{-3}$. These very precise numbers rely however entirely on the above assumption, but generically the CMB dipole is not necessarily due only to a relative velocity. One way to test this assumption was proposed in~\\cite{Burles:2006xf} where it was shown that our peculiar velocity could also be measured using the asymmetry in the location of the peaks of the power spectrum between forward and backward hemispheres, achieving a possible detection of $\\beta$ at $2 - 3 \\sigma$ for resolution $\\ell = 1000 - 1500$. Recently it has also been pointed out~\\cite{Challinor:2002zh,Kosowsky:2010jm,Amendola:2010ty} that {\\it all} the CMB multipoles $a_{\\ell m}$ have a correction due to our local peculiar velocity because the primordial anisotropies are distorted by the Doppler and aberration effects.  This shows up as a correlation between different multipoles $\\ell$ and  could be used as an alternative way of measuring our velocity, as an independent consistency check. This fact might offer also a way to test the isotropy of the Universe on very large scales: it is known in fact that the CMB sky and other observations seem to exhibit some anomalies on the very large scales and this effect offers an observational handle which could either confirm and make more robust the standard assumptions or perhaps point to global anisotropies of the Universe.  The analysis of these CMB correlations has been performed~\\cite{Challinor:2002zh,Kosowsky:2010jm,Pereira:2010dn,Amendola:2010ty} relying upon a Taylor expansion in orders of $\\beta$ of boost effects (Doppler and aberration) on the multipole coefficients $a^{X}_{\\ell m}$, where $X$ stands for $T$ (temperature) or $E, B$ (the $E$ and $B$ modes of polarization). It turns out, however, that each of the series coefficients brings together ever higher powers of $\\,\\ell$, effectively transforming the expansion into one of powers of the product $\\,\\beta\\,\\ell$. Since $\\beta = 1.231 \\times 10^{-3}$ this means that for $\\ell \\gtrsim 800$ the series can no longer be relied upon. We shall therefore distinguish between two regimes of angular scales where the the above correlations show up. In the first regime on large scales, $\\ell \\ll 1/\\beta \\simeq 800$, we can use a perturbative approach and therefore precisely predict that there are  correlations of ${\\cal O} (\\beta \\ell)$ between neighbor $\\ell$'s and look for a signal in the CMB in the form of correlators $\\,a_{\\ell m}^* a_{\\ell+1 m'}$.  Moreover it was shown in~\\cite{Amendola:2010ty} that one can non-trivially and very efficiently define three different estimators (for $m'-m=0,\\pm 1$) in order to measure directly the three cartesian components of $\\boldsymbol{\\beta}$, without having to scan all the possible directions in the sky (i.e., without the need to compute and minimize a numerical $\\chi^2$ for the correlations for each direction in a grid of $\\{\\theta,\\phi\\}$ coordinates).  The multipoles in the complementary regime ($\\ell \\gtrsim 1/\\beta$) cannot be so easily treated but nevertheless also carry information about our peculiar velocity in the form of similar correlations, and it turns out that such information is needed in order to reach a signal-to-noise ratio larger than 1. In this regime the deviation angle $\\theta-\\theta'$ due to aberration is larger than the angular scale of interest $1/\\ell$, so things are much more complicated. There are in fact correlations which are large (${\\cal O} (1)$!)  and nonzero  also for distant $\\ell$'s given by a transformation $ a^{[A]}_{\\ell m} \\;=\\; \\sum_{\\ell'=0}^\\infty K_{\\ell' \\, \\ell\\, m}\\, a^{[P]}_{\\ell' m}\\,,$ between the aberrated $[A]$ frame and the primordial $[P]$ correlations. One in principle would have to compute the matrix elements $\\,K_{\\ell' \\ell m}\\,$, sometimes referred to as the \\emph{aberration kernel}, which are integrals of spherical harmonics with different arguments, plug the matrix elements into all the possible two-point correlation functions and compare with the data. In fact, as we will show in Section~\\ref{sec:nonlinear} the correlations in non-neighboring multipoles ($a_{\\ell m}^* a_{\\ell\\pm n\\, m'}$, $n>1$) also carry a measurable signal, which should be taken into account to measure the velocity with reasonable precision. While in principle straightforward this procedure has some disadvantages: \\emph{(i)} this can be a heavy and delicate numerical task, because of the highly oscillating integrands and also because of the huge number of correlators that one would have to consider (future experiments propose to measure all multipoles as far as $\\ell \\sim 3000+$ for both temperature and polarization); \\emph{(ii)} it is not obvious to understand in this case whether the three simple estimators can be written explicitly for the three cartesian components of the velocity for any $n$ in $a_{\\ell m}^* a_{\\ell\\pm n\\, m'}$, so perhaps the procedure would have to be carried out scanning the sky in all possible directions which would probably make this approach even more expensive in terms of computational time.   Even though we explore this approach further and propose a solution (in Section~\\ref{sec:nonlinear}) to the first of the two challenges  listed above, the main point of this paper is to suggest a very simple trick which may be used to overcome all of these problems and measure directly the three components of $\\boldsymbol{\\beta}$ from a map in a much faster and straightforward way. Using this result we are able to predict with which precision we can measure $\\boldsymbol{\\beta}$ for several future experiments. This trick relies on the fact that we may already {\\it assume} to know the direction and modulus of the velocity $\\boldsymbol{\\beta}$ from the CMB dipole and that we can the use the other multipoles with their correlations as an independent consistency check, in order to confirm (or not) the assumed value of $\\boldsymbol{\\beta}$ up to some precision.  The trick (described in more detail in Section~\\ref{sec:pre-deboost}) is as follows. Given the central assumed value of the velocity $\\,\\boldsymbol{\\beta}_{\\rm dip} = \\boldsymbol{\\beta}_{\\rm dip}^{\\rm fit} \\pm \\delta\\boldsymbol{\\beta}_{\\rm dip}$  indicated by the dipole, we may take the CMB map $T(\\theta,\\phi)$ and ``deboost'' it by a Lorentz transformation  into the frame with opposite velocity $\\,-\\boldsymbol{\\beta}_{\\rm dip}^{\\rm fit}$. As discussed in~\\cite{Menzies:2004vr} such transformation should be carried in the CMB time-ordered data (TOD), i.e., before data treatment to extract the harmonic multipole coefficients \\alm\\ and thus before one constructs the CMB temperature maps. Thus we call this technique ``pre-deboosting'' the CMB data. In the new (pre-deboosted) frame the residual velocity  $\\,\\boldsymbol{\\beta}_{\\rm res}$ compared to the CMB frame is expected to be given just by the residual error $\\,\\delta\\boldsymbol{\\beta}_{\\rm dip}\\,$ on the experimental determination of $\\,\\boldsymbol{\\beta}_{\\rm dip}$. As quoted above, such error is currently approximately $\\,|\\delta\\boldsymbol{\\beta}_{\\rm dip}|  \\,\\simeq\\, 3\\,\\times\\,10^{-6}$. As a consequence the correlations due to Doppler and aberration in this new frame are expected to be of the order $\\,{\\cal O}(\\beta_{\\rm res} \\; \\ell)\\,$ for $\\,\\ell \\ll 1/(\\beta_{res})\\simeq 3\\times10^5\\,$ and so we could safely use the first order perturbative equations to compute the correlation functions $\\,a^*_{\\ell m} a_{\\ell+1 m'}\\,$ up to, say, $\\,\\ell \\simeq 10000\\,$ which is more than enough for all future experiments. %Now, in this new frame we should be able to check that the correlations functions are small and measured the residual velocity $\\beta_{\\rm res}$ with some error $\\delta\\beta_{\\rm res}$, %consistent with zero (if the subtraction has been made to very high precision) %just evaluating the linear estimators for the velocity as described in~\\cite{Amendola:2010ty}. Note that this allows us to use directly again the three efficient define estimators (for $m'-m=0,\\pm 1$) for measuring direction and modulus of the residual velocity.   An unexpected but very interesting prospect is the case in which the measured velocity turns out to be different from the expected $\\beta_{\\rm res}$. This would be a clear indication that the CMB dipole is not completely (nor around $99\\%$, as sometimes stated) due to our peculiar velocity alone. Therefore, this would imply that there are other contributions to the CMB dipole, and it would be interesting to understand whether such correlations may distinguish even the nature of such contributions: adiabatic, isocurvature perturbations, dipolar lensing or other more exotic contributions. %contributions due to vector fields.  The ability to measure exotic contributions to the dipole is of great interest to test Cosmology on very large scales, which could hide non-trivial phenomena, as suggested by some reported anomalies on the low-$\\ell$ CMB multipoles itself~\\cite{Copi:2010na, Bennett:2010jb}. For instance, it could provide valuable information about some proposed tilted cosmological models in which the dipole arises partly due to primordial superhorizon-scale isocurvature fluctuations~\\cite{Turner:1991dn,Grishchuk:1992wv,Langlois:1995ca,Erickcek:2008sm,Zibin:2008fe}, which could provide a possible explanation some recent controversial claims of a high galaxy cluster~\\cite{Kashlinsky:2008ut,Kashlinsky:2009dw,AtrioBarandela:2010wy} and galaxy~\\cite{Watkins:2008hf} bulk flow on large scales and it could be used as a test of non-standard cosmological models. In any case, if $\\beta_{\\rm res}\\neq 0$ we would also be able to measure the direction of our velocity with an error $\\delta\\theta=\\delta\\beta/\\beta_{\\rm res}$ (see~\\cite{Amendola:2010ty}).   This paper is organized as follows. We start by discussing in Section~\\ref{sec:nonlinear} the full non-linear approach to estimate our velocity, and derive a very accurate fitting function for the oscillating integrals. We then discuss in Section~\\ref{sec:pre-deboost} the pre-deboost technique originally proposed by~\\cite{Menzies:2004vr}. In Section~\\ref{sec:applications} we forecast the sensitivity expected from present and future experiments. Finally, we draw our conclusions and summarize our results. ",
        "conclusions": "In this work we proposed a method to overcome the challenges posed to the measurement of our peculiar velocity $\\boldsymbol{\\beta}$ through its aberration and Doppler  effect in the CMB at very small scales, due to the breakdown of the perturbative calculations when $\\ell \\gg 1/\\beta$. This technique consists of pre-deboosting the CMB (at the level of the TOD) and subsequently evaluate the estimators originally constructed in~\\cite{Amendola:2010ty}, which are based on a linear order Taylor expansion in $\\beta$. The expected small residual velocity after such deboost justifies the use of only the linear order terms, and greatly simplifies the analysis. In particular, the pre-deboost method validates analysis of the aberration effect constrained to the correlations between $\\{\\ell,\\,\\ell+1\\}$ only.  Making use of the estimator of the velocity derived in~\\cite{Amendola:2010ty} for such correlations, we investigated the precision with which many of the current and future CMB experiments might be able to measure our velocity $v$. We find that Planck will put an error bar $\\,\\delta v \\simeq 60$ km/s, similar to what ACTPol could achieve in only 2 years with a survey covering $40\\%$ of the sky (4 years of observation would result in $\\,\\delta v \\simeq 40$ km/s). SPTPol, could also be competitive if it carried out a similar wide survey. Proposed future space experiments such as EPIC and Core could put a limit $\\,\\delta v < 30$ km/s. Even more precise experiments, able to measure the CMB signals (and get systematics under control) all the way to $\\ell \\simeq 5000$ could in principle achieve  $\\,\\delta v < 10$ km/s.  Since one expects from the CMB dipole that $v \\simeq 370$~km/s, with current (near future) experiments one would be able to detect a residual term of primordial origin which contributes to a fraction larger that $20\\%$ ($10\\%$) of the dipole. This is of course true unless the primordial effect itself also induces a dipolar kernel leading to correlations, with exactly the same coefficients as the peculiar velocity we analyze here, on the $\\ell>1$ multipoles. Apart from velocity, other possible components to the dipole and/or a dipolar kernel include intrinsic adiabatic, isocurvature and lensing effects, as well as more exotic effects, some of which we now briefly discuss.  An example would be given by some fundamental vector field or perhaps magnetic fields coherent over the horizon scale, which could single out a preferred direction and therefore induce a dipole. Some inhomogeneous and anisotropic cosmological models are also expected to contribute to the dipole. For instance inhomogeneous gigaparsec-scale void models, embodied by the Lema\\^itre-Tolman-Bondi metric, are sometimes constructed to serve as a candidate for dark energy, and have recently drawn considerable attention from the literature (see, for instance,~\\cite{Marra:2011ct}). In these spherical symmetric models, the observer is sometimes assumed to be in the center for simplicity, but any off-center displacement induces a dipole contribution. This in fact happens to be the tightest constrain in this off-center distance, if one assumes the CMB dipole to be entirely due to this displacement and that any additional peculiar velocity is actually zero~\\cite{Quartin:2009xr}. Metrics which exhibit vorticity (e.g., Bianchi V and VII) also can contain a preferred direction and sense, and so also introduce a dipole~\\cite{Barrow1985}. On the other hand, homogenous anisotropic metrics without vorticity, such as Kantowsky-Sachs, Bianchi I and III, exhibit a preferred direction but no preferred sense (i.e., a preferred axis, but not an arrow), so they contribute to the CMB angular power spectrum at most at the quadrupole level~\\cite{Graham:2010hh,Koivisto:2010dr}. Finally, it is also possible that more generically large-scale vector perturbations in FLRW with unusually large amplitude may induce such dipoles, as well as a primordial power spectrum which explicitly depends on the direction of the Fourier mode $\\mathbf{k}$ instead of just $k=|\\mathbf{k}|$~\\cite{Hanson:2009gu}.   %,BlancoPillado:2010uw  %, such as non-trivial topologies,  \\emph{et cetera}.  It is an interesting question, which deserves future investigation,  to check what kind of primordial perturbations or other effects such as the ones discussed above could induce similar correlations between different $\\ell$'s and if the produced correlations are proportional to the produced Doppler effect. Since it is natural to expect that in some cases there is no such proportionality, aberration can provide in principle an observational handle to distinguish between these different contributions.    While the pre-deboost is probably the simplest way to proceeed, we have also proposed very precise fitting functions for the complicated integrals~\\eqref{eq:non-linear-coef}, valid also when $\\ell \\gg 1/\\beta $. %which allows not only to easily compute the correlation coefficients but to quickly their dependence on $\\ell$, $m$ and $\\beta$. Although another method was recently proposed in~\\cite{Chluba:2011zh} to compute this kernel elements, it relies on recursion relations between some of its elements and is not as simple to implement as just using the simple Bessel functions~\\eqref{eq:non-linear-fit-general} and~\\eqref{eq:non-linear-fit-n0}.  %Still concerning the integrals~\\eqref{eq:non-linear-fit-general}, the very high precision and breadth of applicability of an arguably simple fitting function~\\eqref{eq:non-linear-fit-general} (which holds for any $n$ and either temperature or polarization) hints to the existence of an analytical solution involving these Bessel functions. Such solution could perhaps only be obtained in some special limit, such as for instance in the so-called flat-sky approximation, which should be a reasonable approximation for small-scales (although see~\\cite{hu:2000}). An analytical derivation of~\\eqref{eq:non-linear-fit-general} or similar equation, if it exists, could shed some light into the mathematics of this problem and into the limits of validity of the approximation needed to make such a derivation. Nevertheless, such analysis is beyond the scope of this work.  Apart from introducing a correlation between different \\alm 's, it was showed in~\\cite{Challinor:2002zh} that our peculiar velocity would also cause a bias in the measured CMB \\emph{intensity} power spectrum proportional to $4\\beta^2$. Here we showed instead that for the \\emph{temperature} power spectrum (the one which is physically of more interest) the bias is exactly \\emph{zero} at ${\\cal O}\\big(\\beta^2\\big)$. We also showed that the small cross-correlation between $E$ and $B$ modes computed in~\\cite{Challinor:2002zh} is instead also \\emph{zero} -- at least up to ${\\cal O}\\big(\\beta^6\\big)$ -- when one corrects the aberration kernel to transform as temperature instead of as intensity.  With some recent claims of the existence of possible unexpected bulk flows~\\cite{Watkins:2008hf,Kashlinsky:2008ut,Kashlinsky:2009dw,AtrioBarandela:2010wy} it is a very interesting prospect that we will soon have a complementary measurement of our peculiar velocity with respect to the CMB. A dipolar effect has already been measured outside the CMB: in supernovae~\\cite{Weyant:2011hs}, in the cosmic infrared background~\\cite{Fixsen:2011qk} and also marginally in X-rays~\\cite{Boughn:2002bs}; it could also be measured in the future using the cosmic parallax technique~\\cite{Quercellini:2008ty,Quartin:2009xr,Quercellini:2010zr}, and in the optical, radio and gamma bands. It would therefore be interesting to explore the interconnections and compare the future precision that these different approaches can achieve."
    },
    "1112/1112.5686_arXiv.txt": {
        "abstract": "We examine the claim made by Hara et al.\\cite{Hara} in 1969 of the observation of a $10^{19}$eV cosmic ray extensive air shower using the air fluorescence technique.  We find that it is likely that fluorescence light was observed, confirming this as the first such observation.  The work of Hara et al. and their friendly competitors at Cornell University paved the way for modern experiments like the Pierre Auger Observatory and the Telescope Array. ",
        "introduction": "Investigations into the feasibility of detecting air fluorescence light from extensive air showers were conducted in the 1960's by groups led by Suga in Japan and Greisen in the United States.  Results from the Japanese experiment, reported by Hara et al.\\cite{Hara} in 1969, are reviewed here.  In that report, the authors say ``One event is very likely due to the atmospheric scintillation [fluorescence] light from an air shower whose primary energy and distance are about $10^{19}$eV and 3\\,km, respectively''.  This was the first reported observation of fluorescence light from an air shower.  The purpose of the present short note is to the review this observation in the light of our modern understanding of fluorescence detection.  The Japanese experiment ran at the Dodaira Observatory (altitude 876\\,m) for a period of 5 months from December 1968.  The fluorescence telescope consisted of a 1.6\\,m diameter Fresnel lens focussing light onto a camera of 24 PMTs, each of which imaged a 4.5$^\\circ$ degree portion of the sky.  For the observation described here, the telescope with its field of view of $23^\\circ \\times 20^\\circ$ was centered at an elevation of 30$^\\circ$.  The design was similar to Greisen's Cornell telescope \\cite{Cornell}.  However the Japanese design had the advantage of a larger Fresnel lens, and faster electronics.  The rise and fall-times of pulses on the cathode-ray tube displays were 0.12$\\mu$s and 0.2$\\mu$s respectively.  The potential fluorescence observation (event \\#12 in Fig 3 of \\cite{Hara}) triggered 8 PMTs with an angular track length of $18.4^\\circ$ and a duration of 1.9$\\mu$s.  In the next section we review the event geometry before considering the shape of the light profile received at the telescope. ",
        "conclusions": "It appears very likely that the signal detected by Hara et al. was one dominated by fluorescence light, and it is reasonable that they lay claim to the first such detection.  While the timing information of the event is not sufficient to reconstruct a unique shower axis geometry, the event's light profile is consistent with a fluorescence profile, and quite unlike the peaked profile expected for a Cherenkov dominated signal.  Of course some doubt does remain, and the first {\\em unambiguous} detection of fluorescence light from an air shower was with the Utah Fly's Eye prototype operated in coincidence with the Volcano Ranch surface array in 1977 \\cite{FE}.  The $5\\times10^{18}$eV shower simulated in this note produced a light profile with a peak flux of about $5\\times10^{10}$\\,photons/m$^2$/s, the flux observed in \\cite{Hara}.  This energy is a factor of two lower than that claimed by Hara et al., but this can be explained by their assumption of a fluorescence yield of 1.9 photons per metre of electron track, which is approximately a factor of two lower than the modern accepted value.  One remaining uncertainty in this energy is the absolute light calibration of the experiment.  Hara et al. discuss an inconsistency between the observed light intensity spectrum and that expected from simulations.  The inconsistency led to them increasing the observed light fluxes by a factor of 5.  The origin of the inconsistency is unclear, but it would probably be safe to say that $5\\times10^{18}$eV is an upper limit to the energy of the observed air shower."
    },
    "1112/1112.5479_arXiv.txt": {
        "abstract": "Dark matter-dominated cluster-scale halos act as an important cosmological probe and provide a key testing ground for structure formation theory. Focusing on their mass profiles, we have carried out (gravity-only) simulations of the concordance $\\Lambda$CDM cosmology, covering a mass range of $2 \\times 10^{12}-2\\times 10^{15} h^{-1}$M$_\\odot$ and a redshift range of $z=0-2$, while satisfying the associated requirements of resolution and statistical control. When fitting to the Navarro-Frenk-White profile, our concentration-mass ($c-M$) relation differs in normalization and shape in comparison to previous studies that have limited statistics in the upper end of the mass range. We show that the flattening of the $c-M$ relation with redshift is naturally expressed if $c$ is viewed as a function of the peak height parameter, $\\nu$. Unlike the $c-M$ relation, the slope of the $c-\\nu$ relation is effectively constant over the redshift range $z=0-2$, while the amplitude varies by $\\sim 30\\%$ for massive clusters. This relation is, however, not universal: Using a simulation suite covering the allowed $w$CDM parameter space, we show that the $c-\\nu$ relation varies by about $\\pm$ 20\\% as cosmological parameters are varied. At fixed mass, the $c(M)$ distribution is well-fit by a Gaussian with $\\sigma_c/\\langle c\\rangle \\simeq 0.33$, independent of the radius at which the concentration is defined, the halo dynamical state, and the underlying cosmology. We compare the $\\Lambda$CDM predictions with observations of halo concentrations from strong lensing, weak lensing, galaxy kinematics, and X-ray data, finding good agreement for massive clusters ($M_{vir} > 4\\times 10^{14} \\mau$), but with some disagreements at lower masses. Because of uncertainty in observational systematics and modeling of baryonic physics, the significance of these discrepancies remains unclear. ",
        "introduction": "\\label{section:introduction}  According to the current cosmological model, structure forms in the Universe primarily by the amplification of primordial fluctuations driven by the gravitational Jeans instability. The process of nonlinear structure formation is hierarchical and complex, the initial perturbations evolving eventually into a `cosmic web' network consisting of voids, filaments, and clumps. The clumps, termed halos in cosmological parlance, are dark matter dominated localized mass overdensities with their own complex substructure. Observed baryonic systems such as galaxies and hot gas reside in these halos. Although the dark matter within halos cannot be observed directly, its presence can be inferred by dynamical arguments, and much more directly, through gravitational lensing of background sources.  The notion of the dark matter dominated halo is one of the fundamental building blocks in studies of the formation of individual galaxies, galaxy groups, and galaxy clusters (for an overview, see \\citealt{mo_book}). The structure of halos has been extensively studied using N-body simulations over a wide range of halo masses. Even though individual halos can be, and are, dynamically and morphologically complex, it was shown by \\cite{nfw1, nfw2} (NFW) that the spherically averaged density profiles of `relaxed' halos formed in cold dark matter (CDM) simulations can be described by a roughly universal functional form -- the NFW profile -- independent of their mass, the spectrum of initial fluctuations, and cosmological parameters. The NFW profile has a fixed shape, albeit with two scale parameters; as applied to individual halos it has been remarkably successful and is often applied to all halos, regardless of their dynamical state. (When applied to stacked or average halos, this profile is somewhat less succesful, as discussed later below.)  The two parameters of the NFW profile are a mass and a scale radius. The scale radius, $r_s$, specifies the point where the logarithmic slope of the profile equals -2 (at small radii, the profile $\\sim 1/r$, while at large radii, it asymptotes to $\\sim 1/r^3$). Instead of $r_s$, one often uses the concentration, which is the radial scale set by the halo mass divided by $r_s$. In cluster cosmology, the usual key observable is the halo mass, rather than the profile per se. The cluster mass function (cluster abundance, more generally), is a sensitive probe of dark energy, since clusters form very late, during the epoch of dark energy dominance. However, measuring the concentration parameter, the simplest first measurement of a profile, can also be very useful.  First, as shown originally by NFW, the concentration of a halo, $c$, has a strong correlation to its mass, $M$, therefore measuring the $c-M$ relation observationally is a direct test of the CDM paradigm. In fact, combining cluster $c-M$ predictions and measurements, and the measured gas mass fraction, one can aim to constrain $\\Omega_m$ and $\\sigma_8$ \\citep{ettori11}. As another example, lensing shear peak counts, a proposed weak lensing survey cosmological probe, is very sensitive to the form of the $c-M$ relation \\citep{king11}. Finally, future measurements of the weak lensing power spectrum will probe small enough spatial scales that results will be sensitive to baryonic effects on the halo profile, i.e., modifications to the gravity-only $c-M$ relation \\citep{mw04, zhan04}. We will return to these points in more detail below.  The correlation of halo concentration with mass is based on the idea -- as first explicated by NFW -- that the concentration is determined by the mean density of the universe when the halo is assembled, with higher concentrations corresponding to higher densities. Thus cluster mass halos, which are assembling today, should have a lower concentration than halos of lower mass that were built up at an earlier epoch, where the mean density was higher. Furthermore, one may expect this trend to flatten out (sufficiently) beyond the nonlinear mass scale $M_*$, and therefore, since $M_*$ falls rapidly with redshift, flatten out over an extended range in mass as redshift increases. Although the general arguments are plausible and are broadly consistent with simulation results, a predictive theory for the mean of the $c-M$ relation, and its scatter, does not exist. Several simple heuristic models tuned to simulations have been suggested (NFW; \\citealt{bullock99}; \\citealt{eke01}; \\citealt{zhao09}) but their predictive status cannot be considered satisfactory, especially at the higher end of halo masses (see, e.g., \\citealt{gao07}; \\citealt{hayashi07}; \\citealt{maccio08}; \\citealt{zhao09}). Indeed there is sufficient uncertainty even when comparing simulation results from different groups, that the general problem is still open. However, as the mass resolution in large-volume simulations continues to improve, we may expect this situation to be merely temporary.  On the observational front, cluster (and group) halo profiles can be studied using both strong and weak gravitational lensing, individually, and in combination (see, e.g., \\citealt{comerford07}; \\citealt{broadhurst08}; \\citealt{mandelbaum08}; \\citealt{okabe09}; \\citealt{oguri11}; \\citealt{zitrin11}; \\citealt{coe12}), projected gas density and temperature profiles from X-ray observations (see, e.g., \\citealt{vikhlinin05}; \\citealt{buote06}; \\citealt{schmidt06}; \\citealt{gastaldello07}; \\citealt{vikhlinin09}; \\citealt{sun09}; \\citealt{ettori11}), and galaxy kinematics (\\citealt{diaferio05}; \\citealt{rines06}; \\citealt{wojtak10} and references therein). Results from these observations have generally shown qualitative agreement with the $c-M$ relation obtained from simulations, although there have been difficulties with matching the shape and normalization. Additionally, there are discrepancies between different sets of observations, presumably because the underlying systematics are not fully understood and modeled.  The purpose of this paper is to present a set of predictions for the NFW mass profile targeted primarily towards massive clusters. To do so, however, a fairly large mass range must be considered in order to obtain a sufficiently well-defined $c-M$ relation. Our simulations cover three orders of magnitude in mass ($\\sim 10^{12}-\\sim 10^{15} h^{-1}$M$_\\odot$) with very good control of statistics over the entire range. The high dynamic range and excellent statistics enable us to derive a new set of results for the mass profile, including profile evolution and probability distribution functions (PDFs) for the concentration as a function of mass. We compare our results with previous simulations and with a set of recent observations of the cluster mass profile.  The paper is organized as follows. In Section~\\ref{section:halo} we discuss general features of the $c-M$ relation in the simulation context focusing on the role of differing definitions and analyses. In Section~\\ref{section:sims}, we describe the main features of the simulation runs. We present our results for the $c-M$ relation and its redshift evolution in Section~\\ref{section:results}. This is followed (Section~\\ref{sec:wcdm}) by a presentation of results from a suite of $w$CDM cosmologies in order to further study how the concentration depends on cosmology. Next, in Section~\\ref{sec:comp}, we provide a detailed comparison with recent observations, noting areas of agreement and disagreement. Finally, Section~\\ref{section:disc} is devoted to a summary of the results and further discussion. An Appendix discusses various systematic issues that need to be considered when deriving concentrations from simulation results. A number of tests are used to illustrate these points and to verify the robustness of the numerical procedures carried out in this paper. ",
        "conclusions": "\\label{section:disc}  We presented results for the concentrations of dark matter halos using a set of large-volume simulations. With a total volume roughly 1-2 orders of magnitude larger compared to previous simulations, we focused on studying the $c-M$ relation for massive clusters. As shown in the past, at the high mass end, the $c-M$ relation becomes flatter at $z=0$ and the flattening becomes more significant at higher $z$. The mean concentration of the sample when expressed in terms of the peak height parameter, $\\nu(M,z)=\\delta_c/\\sigma(M,z)$, shows a roughly uniform slope at all redshifts. Indeed, the slope of the $c-\\nu$ relation does not change with redshift. The amplitude of the $c-\\nu$ relation evolves by about $30\\%$ at the high mass end from $z=0-2$. The $z$-evolution is consistent with the results of \\cite{gao07}, although the overall amplitude of the concentration differs because of the different choices of $\\sigma_8$. We do not observe a rise in concentration at higher masses as reported by \\cite{klypin10} and \\cite{prada11} (the Appendix includes further discussion). Because of our large halo sample, we can study the distribution of the concentration in individual mass bins; we find that the distribution of concentrations is well-described by a Gaussian PDF \\citep{lukic09,reed11}. Thus the halo profile shape can be described by two parameters -- the mean concentration and its standard deviation. By comparing results across a number of $w$CDM cosmologies, we find that the standard deviation is roughly universal, $\\sigma_c=0.33 c$, and does not change with redshift, halo dynamical state, or cosmological parameters.  We investigated how the concentration changes as the cosmological parameters are varied using a set of 18 runs spanning the $w$CDM cosmology parameter space. The parameter range covers the 2$\\sigma$ variation around the best fit WMAP5 cosmology. We find that over this parameter range, the $c-M$ relation varies by $\\sim \\pm$ 20\\%, although the standard deviation $\\sigma_c$ follows the relation $\\sigma_c=0.33 c$. As suggested by our work on the $w$CDM models, and also previous studies of redshift evolution, the halo formation epoch, and hence the concentration, depends on the matter fluctuations, slope of the power spectrum and the growth factor. Thus calibrating the $c-M$ relation as a function of cosmology is important for a wide variety of problems, ranging from galaxy formation, the weak lensing shear power spectrum, to the case of assembly bias in clusters. We will address the cosmology dependence of the $c-M$ relation in detail using more simulations and analytical models elsewhere.  The simulation predictions are in good agreement with observations from strong lensing, weak lensing, galaxy kinematics, and X-ray data for massive clusters with masses $M_{vir} > 4\\times 10^{14} \\mau$. At lower masses, different observations suffer from different sources of systematic error. For example, the lensing data need to account for bias due to the triaxiality of halos while the X-ray data typically ignore the non-thermal pressure component in galaxy clusters which can lead to a systematic underestimate of the cluster mass \\citep{lau09}. The simulations, on the other hand also need to account for baryonic effects which play a bigger role as the halo mass decreases. As a result, due to cooling, gravity-only simulations may predict $20-30\\%$ lower concentration for clusters with masses $M_{vir} < 4\\times 10^{14} \\mau$. The fact that most recent observations are in agreement with the simulation results (and amongst themselves) to better than $20\\%$ for massive clusters ($M_{vir} > 4\\times 10^{14} \\mau$) indicates that baryonic effects influencing the cluster mass profile are indeed small and that the individual observational systematics are under some level of control.  \\appendix"
    },
    "1112/1112.0540_arXiv.txt": {
        "abstract": "We present results from an analysis of the broad-band, 0.3--195\\,keV, X-ray spectra of 48 Seyfert 1--1.5 sources detected in the very hard X-rays with the Swift Burst Alert Telescope (BAT).  This sample is selected in an all-sky survey conducted in the 14--195\\,keV band.  Therefore, our sources are largely unbiased towards both obscuration and host galaxy properties.  Our detailed and uniform model fits to Suzaku/BAT and XMM-Newton/BAT spectra include the neutral absorption, direct power-law, reflected emission, soft excess, warm absorption, and narrow \\ion{Fe}{1} K$\\alpha$ emission properties for the entire sample. We significantly detect O VII and O VIII edges in 52\\% of our sample. The strength of these detections are strongly correlated with the neutral column density measured in the spectrum. Among the strongest detections, X-ray grating and UV observations, where available, indicate outflowing material. The ionized column densities of sources with O VII and O VIII detections are clustered in a narrow range with N$_{\\rm warm} \\sim 10^{21}$\\,cm$^{-2}$, while sources without strong detections have column densities of ionized gas an order of magnitude lower.  Therefore, we note that sources without strong detections likely have warm ionized outflows present but at low column densities that are not easily probed with current X-ray observations. Sources with strong complex absorption have a strong soft excess, which may or may not be due to difficulties in modeling the complex spectra of these sources. Still, the detection of a flat $\\Gamma \\sim 1$ and a strong soft excess may allow us to infer the presence of strong absorption in low signal-to-noise AGN spectra.  Additionally, we include a useful correction from the Swift BAT luminosity to bolometric luminosity, based on a comparison of our spectral fitting results with published spectral energy distribution fits from 33 of our sources. ",
        "introduction": "The Swift Burst Alert Telescope (BAT) is conducting the first all-sky, very hard X-ray survey in thirty years.  With hundreds of detections in a harder X-ray band than the previous survey by HEAO-1 \\citep{1982ApJ...253..485P}, the Swift survey presents an unprecedented sample of active galactic nuclei (AGN) sources.  Due to their detection in the 14--195\\,keV, very hard X-ray band, the Swift sources are unbiased to all but the highest levels of obscuration.  Therefore, the Swift BAT-detected AGN are an important sample for determining the global properties of AGN.  In \\citet{2009ApJ...690.1322W}, the 0.3--10\\,keV X-ray properties for the 153 sources detected in the 9-month Swift BAT catalog \\citep{2008ApJ...681..113T} are presented.  The 9-month catalog includes all sources with very hard X-ray detections of F$_{14 - 195 {\\rm keV}} > 10^{-11}$\\flux.  These sources are local, with an average redshift of 0.03, and are bright IR/optical/UV/X-ray sources.  The X-ray analysis of the 9-month sources uncovered many results of the properties of local AGNs, relying on simple models.  However, since both broader X-ray band and higher signal-to-noise spectra are available for a majority of these sources, more detailed analyses are now possible.  Our main goal in this paper is the detection and characterization of ``warm absorbers'', signatures of ionized gas potentially from outflows, in the Type 1 AGN.  Previous studies searching for warm absorbers were biased -- selecting all the bright sources with archived X-ray or UV data \\citep{1997MNRAS.286..513R,1998ApJS..114...73G,1999ApJ...516..750C,2007AJ....134.1061D}. Therefore, the results of these studies may also be biased.  Mass outflows or more specifically AGN feedback mechanisms are important in galaxy formation \\citep{1998AA...331L...1S,2003ARA&A..41..117C,2005ARA&A..43..769V} and a potential cause of the well-known relationship between galaxies and their black holes (i.e., the M-$\\sigma$ relation first described in \\citealt{2000ApJ...539L...9F,2000ApJ...539L..13G}). Thus, it is important to understand the AGN warm absorber properties in an unbiased sample.  Since X-ray warm absorbers are most evident in the soft emission ($< 2$\\,keV), where more absorbed sources exhibit low flux levels, we rely on a study of the Type 1 sources detected by Swift's BAT. In this paper, we present the detailed broad-band spectral properties of 48/51 Seyfert 1--1.5 sources, with $| b | > 15$\\degr and high signal-to-noise X-ray CCD spectra available, which were detected in the Swift BAT 9-month survey.  Our analysis relies preferentially on X-ray spectra from Suzaku, which provides simultaneous coverage from $\\sim 0.2$--$50$\\,keV, and time-averaged spectra from Swift's BAT in the 14--195\\,keV band.  With a broad bandpass, we obtain tight constraints on the continuum emission, which is vital in determining the warm absorber properties.  Where Suzaku spectra were not available, we use data from XMM-Newton, which covers the 0.3--10\\,keV band, with the time-averaged Swift BAT spectra.  We describe our data reduction in \\S~\\ref{sect-data}.  Our spectral analysis is detailed in \\S~\\ref{sect-spectra}.  Discussion of our results is found in \\S~\\ref{sect-discussion}, followed by conclusions in \\S~\\ref{sect-conclusion}. ",
        "conclusions": "\\label{sect-discussion} We performed detailed, uniform spectral fitting on a sample of 48/51 Seyfert 1--1.5 sources selected in the 14-195\\,keV band with the Swift BAT.  In this section, we discuss our findings, based on general spectral properties, including the continuum shape, luminosity, high energy emission lines (e.g., \\ion{Fe}{1} K), and soft excess.  Additionally, we specifically highlight results of an analysis of the warm absorption properties and their implications for AGN feedback.  \\subsection{Basic Spectral Properties}   \\subsubsection{\\bf Intrinsic Neutral Column Density and Redshift}\\label{sect-nh} In the X-rays, type 1 AGN typically have low columns of obscuring material in the line of sight to the central source.  Since the emission is largely unobscured, we view directly the central regions surrounding the black hole. In the optical, type 1 AGN exhibit broad permitted emission lines, particularly in H$\\alpha$ and H$\\beta$.  The distinction between sub-types of Sy 1s relies on the strength of the narrow H$\\beta$ emission component, where a Sy 1.5 is intermediate between a Sy 1 and Sy 2 with a clearly distinguishable narrow emission-line component \\citep{1977ApJ...215..733O}.  Since optical Sy 1.5 galaxies are in an intermediate state, in a simple unified model of AGN (e.g., \\citealt{1993ARA&A..31..473A}) we expect the Sy 1.5s to have intermediate column densities between the low column Sy 1s and the higher column Sy 2s.  Our sample includes 22 Sy 1s, 9 Sy 1.2s, and 17 Sy 1.5s.  Among the type 1s, there are several broad line radio galaxies (BLRGs), including 1H 0419-577, 3C 120, 3C 382, 3C 390.3, 4C +74.26, 2MASX J21140128+8204483, and MR 2251-178.  Average values for the logarithm of neutral column density (using an upper limit of $10^{19}$\\,cm$^{-2}$ for spectra with no clearly detected neutral absorption component) correspond to $19.65 \\pm 0.53$ for Sy 1s, $20.12 \\pm 0.63$ for Sy 1.2s, and $20.61 \\pm 1.14$ for Sy 1.5s.  Broad line radio galaxies, as a subset of Sy 1s, have an average neutral column density measured at $20.26 \\pm 0.74$.  Therefore, Sy 1s do have the lowest column densities, while Sy 1.5s are more obscured in the X-rays, on average.  However, if we exclude the four Sy 1.5s with columns higher than $10^{22}$\\,cm$^{-2}$, a K-S test shows a smaller probability of the Sy 1s and Sy 1.5s being drawn from different distributions ($D = 0.359$ and $P = 0.264$).  % Further, the optical Sy 1--1.5 sources are, as expected, less obscured in the X-rays than the X-ray obscured sources (with typical values of \\nh\\,$> 10^{22}$\\,cm$^{-2}$, see  \\citealt{2009ApJ...690.1322W}).  In Figure~\\ref{fig-nhz}, we plot the measured column densities for our sample versus redshift.  The highest redshift sources in our sample are the broad line radio galaxies (black squares), with $\\langle z \\rangle = 0.068$.  The Sy 1.5s are the lowest redshift sources ($\\langle z \\rangle = 0.012$), while the Sy 1s have an average redshift of 0.026 and the Sy 1.2s are intermediary with a value of 0.017.  In Figure~\\ref{fig-nhz}, we find that there are no Sy 1s at the lowest redshifts. From the plot, we find that there are 5 Sy 1.5s and 2 Sy 1.2s at $z < 0.008$.  Since Sy 1s are relatively unobscured, we expect that Swift's BAT has no bias against selecting Sy 1s in this redshift range.  The lack of Sy 1.5 detections at higher redshift may indicate that the Sy 1.5s are less luminous, since obscured sources in the BAT sample have lower luminosities than unobscured sources \\citep{2009ApJ...690.1322W}.  However, at the low column densities measured from our broad-band X-ray fits (from 0.3--195\\,keV), we expect the obscuring material in Sy 1.5s to have little affect at the high energies to which BAT is sensitive.  In the following section, we discuss the derived X-ray luminosities, with respect to optical classifications.  While we find no Sy 1s at low redshift, it is difficult to make statistical claims on the paucity of low redshift Sy 1s due to the small number of objects in this redshift range.  However, this result is intriguing since the X-ray analyses of the Swift sample reveal that the type 2 sources are less luminous than the type 1 objects \\citep{2009ApJ...690.1322W} and the fraction of obscured/unobscured AGN also changes with redshift (e.g., \\citealt{koss-thesis}).  We will test this further with optical and X-ray follow-ups of the 600 AGNs detected in the 58-month Swift catalog (Baumgartner et al., submitted).  If the same result of few Sy 1s at low-redshift is found, one possible explanation is that there is a physical difference in the broad line region/accretion state between Sy 1.5s and Sy 1s.  \\citet{2009ApJ...701L..91E} posit that in the disk wind model, where the broad line region and torus are part of a wind from the accretion disk, that at low accretion rates the broad line region disappears.  If the Sy 1.5s are accreting at a lower rate (see \\S~\\ref{lum}), the difference in the appearance of the optical broad lines could be correlated with a different accretion state.  In the current study, there is the possibility, however, of our sources being mis-classified (i.e., between Sy 1 sub-types). This is a concern since the NED classifications we use come from heterogeneous sources.  We will test this further in a follow-up paper examining our uniform optical spectra obtained with the Apache Point Observatory 3.5-m and CTIO SOAR 4-m telescopes.  Using the optical follow-ups, we will carefully classify the optical spectra of our sample to determine whether the NED classifications are correct.  We confirm that the optical follow-ups show broad Balmer lines in all of these sources, however, we will include more detailed analyses in the optical paper (Winter et al., in prep).   \\subsubsection{\\bf Luminosity, Mass, and Accretion Rate}\\label{lum}  The luminosity, black hole mass, and accretion rate are among the most basic properties of an AGN.  Since the Swift BAT bandpass is at high enough energies (14--195\\,keV) to be unaffected by all but the highest levels of obscuration, it is a good proxy for the bolometric luminosity of the AGN.  Our measurements of the BAT band luminosity are computed using the 14--195\\,keV flux listed in Table~\\ref{tbl-flux}.  We find that the BLRGs are the most luminous sources in our sample, while the Sy 1.5s are least luminous.  Average values and standard deviations for $\\log$\\,L$_{14-195\\,{\\rm keV}}$ are: $44.8 \\pm 0.2$ (BLRGs), $43.8 \\pm 0.4$ (Sy 1s), $43.4 \\pm 0.6$ (Sy 1.2s), and $43.3 \\pm 0.7$ (Sy 1.5s).  The standard deviations demonstrate that there is no statistical difference between the luminosity of Sy 1s, 1.2s, or 1.5s.  However, the BLRGs are significantly more luminous -- approximately 4 times more luminous than the Seyfert sources.  The mass determinations recorded in Table~\\ref{tbl-flux} are derived from 2MASS bulge photometry of the host galaxies (see \\citealt{2008ApJ...684L..65M} and \\citealt{2009ApJ...690.1322W}) and are in good agreement with the independent analysis of 33 of our sources in  \\citet{2007MNRAS.381.1235V} and  \\citet{2009MNRAS.392.1124V}.  Comparison of the 2MASS-derived masses with both reverberation mapping and H-$\\beta$ FWHM derived masses were discussed in \\citet{2010ApJ...710..503W}.  We found masses derived from each of these methods to be well-correlated, with $\\log M_{2MASS} = (0.91 \\pm 0.14) \\times \\log M_{H\\beta} + (1.07 \\pm 1.13)$ and the H$\\beta$-derived masses equivalent to the reverberation results.  The masses recorded in Table~\\ref{tbl-flux} are the 2MASS values corrected to match the more accurate H-$\\beta$ derived masses, with the exceptions of NGC 4051, NGC 4593, Mrk 279, and NGC 7469, where we use the reverberation mapping derived masses (in each of these cases there was a significant difference between the masses from the alternate method).  Additionally, we note that the mass of the Sy 1.5 2MASX J11454045-1827149, derived from the 2MASS K-band photometry, is low ($\\log M/M_{\\sun} = 6.2$), requiring a very high accretion rate ($L_{\\rm bol}/L_{\\rm Edd} = 3.5$).  We have obtained optical spectroscopy of the H-$\\beta$ region of this source from the CTIO SOAR telescope and are in the process of reducing the data.  An alternative mass estimate will be included in Winter et al., in prep.  Comparing the average masses for our sample, we find $\\log M/M_{\\sun}$ of $8.5 \\pm 0.2$ (BLRGs), $7.9 \\pm 0.6$ (Sy 1s), $7.9 \\pm 0.7$ (Sy 1.2s), and $7.8 \\pm 0.7$ (Sy 1.5s).  We find that within the standard deviation the black hole masses are similar between all of our sources.  This is also true for the obscured type 2 sources, shown in \\citet{2009ApJ...690.1322W}.   To determine the accretion rate of our sources, we must first determine the bolometric luminosity.  A reliable method for determining the bolometric luminosity is through broad-band fitting of the spectral energy distribution (e.g., \\citealt{2002ApJ...579..530W}).  For 33 sources in our sample, bolometric luminosities and accretion rates were obtained through fitting simultaneous optical/UV/X-ray data from XMM-Newton \\citep{2007MNRAS.381.1235V} or Swift \\citep{2009MNRAS.392.1124V}.  In Figure~\\ref{fig-comparebol}, we compare the Swift BAT band luminosities obtained from our spectral fits with the bolometric luminosities calculated from SED fitting in  \\citet{2007MNRAS.381.1235V} and  \\citet{2009MNRAS.392.1124V}.  We expect that the 14--195\\,keV luminosity is the direct unobscured signature from the AGN for our Seyfert 1 sources and our figure confirms this.  Fitting an ordinary least squares line to the data, we find that $\\log L_{\\rm bol} = (1.1157 \\pm 0.1172) \\log L_{14-195\\,{\\rm keV}} + (-4.2280 \\pm 5.1376)$.  This fit is very significant, with a correlation coefficient of $R^2 = 0.82$.  Using this relationship, we find that the average bolometric luminosities of our sources are $5.7 \\times 10^{45}$\\,\\lum (BLRGs), $4.4 \\times 10^{44}$\\,\\lum (Sy 1s), $1.6 \\times 10^{44}$\\,\\lum (Sy 1.2s), and $1.2 \\times 10^{44}$\\,\\lum (Sy 1.5s).  The bolometric luminosities are included in Table~\\ref{tbl-flux}.  The mass accretion rate, $\\dot{M} = \\frac{L}{\\eta c^2}$, ranges roughly from 1\\,M$_{\\sun}\\,{\\rm yr}^{-1}$ for the BLRGs and 0.02\\,M$_{\\sun}\\,{\\rm yr}^{-1}$ in the Sy 1.5s (assuming $\\eta = 0.1$).  It is useful to parameterize the accretion rate relative to the Eddington limit.  The Eddington ratio is defined as the ratio of the bolometric luminosity to the Eddington luminosity, the luminosity where the gravitational and radiation pressure balance ($L_{\\rm Edd} = 1.3 \\times 10^{38} \\times M/M_{\\sun}$\\,ergs\\,s$^{-1}$), or equivalently the ratio of the mass accretion rate to the Eddington accretion rate ($\\dot{M}_{\\rm Edd} = \\frac{L_{\\rm Edd}}{\\eta c^2}$).  We include the Eddington ratio for our sample in Table~\\ref{tbl-flux}.  We find values of $L_{\\rm bol}/L_{Edd}$ corresponding to $0.14 \\pm 0.08$ (BLRGs), $0.05 \\pm 0.14$ (Sy 1s), $0.03 \\pm 0.05$ (Sy 1.2s), and $0.02 \\pm 0.09$ (Sy 1.5s).  The main uncertainties in these measurements are in the estimates of the mass.  For one source,   2MASX J11454045-1827149, which does not have a measurement from alternative methods, the Eddington ratio is suspiciously high ($> 1$).  As discussed above, we will determine a mass estimate based on the H-$\\beta$ width (Winter et al., in prep).  The general picture, however, emerging from our estimates shows that local AGN selected in the very hard X-rays tend to have accretion rates from $10^{-3} - 0.5$ times the Eddington rate.   The neutral column densities measured in the X-ray spectroscopy are a combination of the effects of absorption features of metals, which impose a soft X-ray cut-off in the spectrum.  This material may be associated with intergalactic material, gas local to the host galaxy, and/or intrinsic to the AGN.  In  \\citet{2009ApJ...690.1322W}, we tested whether the obscuration was associated with the accretion rate of the AGN (Figure 10a) and the inclination of the host galaxy (Figure 16).  We found that the bulk of the obscuring material is not associated with the host galaxy (highly inclined hosts have slightly higher neutral column densities, but not enough to account for all of the obscuration measured) and that there is no correlation between the column and accretion rate.  In Figure~\\ref{fig-lumlledd}, we test the dependence of the measured column on both the luminosity and Eddington ratio measured from our broad-band spectral fits.  There is no correlation found, which is consistent with the unified model, since our line of sight viewing angle with respect to the torus torus, a potential source of the neutral obscuration, is expected to only be an effect of our line of sight to the AGN (while this is true for low luminosity sources like the BAT-detected AGN, we note that this is not the case for higher luminosity sources).   \\subsubsection{\\bf Direct and Reflected Continuum} The direct emission from our sources was modeled as a cutoff power-law.  The power-law index, $\\Gamma$, measured values correspond to $1.90 \\pm 0.16$ (BLRGs), $1.78 \\pm 0.34$ (Sy 1s), $1.84 \\pm 0.18$ (Sy 1.2s), and $1.52 \\pm 0.61$ (Sy 1.5s).  The low average photon index for the Sy 1.5s is a result of several outliers with $\\Gamma \\sim 1.0$, including NGC 526A, Mrk 6, NGC 3227, and NGC 3516, with NGC 4151 having $\\Gamma << 1.0$. Among these, NGC 3227, NGC 3516, and NGC 4151 are poorly fit with the base model (reduced $\\Delta\\chi^2 > 2.0$).  Two Sy 1s also have low values: EXO 055620--3820.2 and NGC 3783.  For these sources, we attribute the unusual $\\Gamma$ measurements to difficulty modeling the absorbing components. Warm absorber signatures, measured through the \\ion{O}{7}/\\ion{O}{8} edges (see Table~\\ref{tbl-warmabs}), are the strongest in our sample for all but EXO 055620--3820.2.  While we do not measure strong absorbers in EXO 055620--3820.2, analysis of the archived X-ray data by \\citet{2009MNRAS.394L...1L} reveals obscuring clouds and complex structure (which we also detect in this observation).  Removing these values, we find that the mean $\\Gamma$ value for Sy 1s and Sy 1.5s agree -- $1.90 \\pm 0.17$ and $1.90 \\pm 0.23$, respectively -- and a K-S test reveals a high probability ($P = 0.947$) that both are drawn from the same distribution.  Further, these average values are consistent with the distribution of $\\Gamma$ for the BLRGs, demonstrating that there is no statistical indication of spectral hardening with luminosity/accretion rate in the BAT sample (see Figure~\\ref{fig-gammalum}).  %  In Figure~\\ref{fig-gammalum}, we also plot the average power-law index binned by luminosity/accretion rate (right plots).  We divided the range of luminosities/accretion rates for our sources into equally-spaced bins and computed the average power-law index in each of these bins.  The error-bars indicate the difference between the average value and the maximum/minimum in each bin, with average values of $\\Gamma = 1.07$, and are plotted as the logarithm of this value in order to illustrate the average values more clearly.  In  \\citet{2009ApJ...690.1322W}, we found that there was no correlation between power-law index and luminosity/accretion rate, consistent with results from previous low-redshift samples \\citep{2000ApJ...531...52G}.  Since our study consisted of all of the Swift-BAT detected AGN (e.g., including both absorbed and lower luminosity sources), we concluded that we found no correlation because the properties of the Swift sources were more diverse than the higher-redshift, more luminous sources which showed these correlations \\citep{2004ApJ...605...45D,2008AJ....135.1505S}.  We did, however, see correlations for individual sources where we had multiple X-ray observations \\citep{2008ApJ...674..686W}.  As shown in Figure~\\ref{fig-gammalum}, in the current study we find that the photon-index is well-correlated with luminosity and accretion rate, when the values are binned.  We find that: \\begin{equation} \\Gamma = (0.16 \\pm 0.03) \\log L_{\\rm bol} + (-5.31 \\pm 1.17), \\end{equation} with a correlation coefficient of $R^2 = 0.86$, and \\begin{equation} \\Gamma = (0.23 \\pm 0.03) \\log L_{\\rm bol}/L_{\\rm Edd} + (2.08 \\pm 0.04), \\end{equation} with a coefficient of determination (the square of the Pearson linear correlation coefficient) of $R^2 = 0.83$ (the probability for both corresponds to $P < 0.001$ for 48 degrees of freedom).  Our linear correlation between binned $\\Gamma$-L$_{\\rm bol}$ is consistent with the rest-frame $\\Gamma$-(2-10\\,keV) luminosity relation found in \\citep{2004ApJ...605...45D} ($\\Gamma = (0.13 \\pm 0.04) L_{2-10\\,{\\rm keV}} + (-4.1 \\pm 1.7)$ for $0.3 \\la z \\la 0.96$ AGN).  This suggests that the X-ray power-law index and luminosity/accretion rate are related for the Sy 1s and that a similar process is at work as with the higher redshift radio quiet quasars.  However, the fact that a similar relation is not found for the entire Swift BAT-detected sample \\citep{2009ApJ...690.1322W} enforces the result that the lowest luminosity sources (absorbed sources whose optical emission line properties classify them as \\ion{H}{2} galaxies/composites/LINERs) are in a different accretion state  \\citep{2010ApJ...710..503W}.  Cutoff energies of the power-law component are poorly constrained for the majority of our sample (40/48).  For many of the sources, the parameter was best-fit with the upper limit of $E_{cutoff} = 500$\\,keV.  Of the sources with constrained cutoffs, EXO 055620--3820.2 and Mrk 79 have cutoffs between 20--60\\,keV, and the remaining five sources have cutoffs from the range $\\sim 60$--$400$\\,keV.  By contrast, \\citet{2009MNRAS.399.1293M} find cutoff energies between 50--150\\,keV in their sample of Sy 1s, derived from the INTEGRAL survey.  As with our analysis, \\citet{2009MNRAS.399.1293M} use broad-band X-ray spectral fits and use a very similar model, including reflection with the {\\tt pexrav} model.  Of the overlapping sources between both the INTEGRAL and our Swift study, we find general agreement in the cutoff values, however, our error bars on the cutoff energies are much larger than those in the  \\citet{2009MNRAS.399.1293M} analysis.  It is unclear why this is the case, particularly since our study includes, on average, a higher number of degrees of freedom from using the joint Suzaku/BAT or XMM-Newton/BAT data points.  A reflected spectrum is found to be significant in 37/48 (77\\%) of the sample, as determined by $\\Delta\\chi^2$ on adding this parameter (see Table~\\ref{tbl-fits}).  We find good agreement with measured reflection parameters between our study and that of  \\citet{2009MNRAS.399.1293M}, of which 10 sources overlap.  The average and standard deviation for the measured reflection parameter ($R \\sim \\Omega/2\\pi$) correspond to $1.56 \\pm 1.45$ (BLRGs), $2.23 \\pm 1.47$ (Sy 1s), $1.96 \\pm 1.22$ (Sy 1.2s), and $1.75 \\pm 1.61$ (Sy 1.5s).  There is no statistical difference in the reflection parameter between sources of different Sy type.  In Figure~\\ref{fig-reflection}, we look for correlations between the reflection parameter and both AGN luminosity and power-law index, $\\Gamma$.  Clearly, we find no correlation between $R$ and L$_{\\rm bol}$.  A correlation does appear, however, between $R$ and $\\Gamma$, where higher $\\Gamma$ values coincide with higher reflection parameters (note that $R \\sim 0.0$ corresponds to no reflection, while more negative values correspond to larger reflection).  Excluding sources with $\\Gamma < 1.5$, which include sources with complex absorbers, we find that $R = (-8.58 \\pm 1.84) + (5.74 \\pm 0.97) \\times \\Gamma$.  This correlation is not strong, with a correlation coefficient of $R^2 = 0.49$.  Correlations in $R - \\Gamma$ were previously seen in a number of studies, including those of \\citet{1999MNRAS.303L..11Z},  \\citet{2007ApJ...664..101M},  and \\citet{2009MNRAS.399.1293M}.  Since both parameters are linked in the fitting process, this trend may not be physical.  If the correlation is physical, \\citet{1999MNRAS.303L..11Z} interpret the $R - \\Gamma$ relation as follows.  The reflection parameter is proportional to the angle subtended by the reflector.  This ``cold'' reflector emits soft photons, which irradiate the X-ray source (i.e., the corona) and act as the seed for Compton upscattering.  If $R$ is large, there is more reflecting material and therefore a stronger flux from the soft photons.  This leads to stronger cooling of the plasma, a smaller Comptonization parameter ($y$), and a softer spectrum (or larger $\\Gamma$).  \\citet{2001ApJ...556..716P} show that with the {\\tt pexrav} model larger reflection parameters also correspond to higher corona temperatures ($kT_e \\propto E_{cutoff}$) and lower optical depths towards inverse Compton scattering, as well as softer spectra (higher $\\Gamma$).  For sources in our sample with well-constrained $R$ and $E_{cutoff}$, there is no correlation between cutoff energy, thereby corona temperature, and reflection parameter/$\\Gamma$.   We conclude that the relationship between the direct and reflected emission is not easy to interpret in our sample.  Particularly, degeneracies between the fitted values from the {\\tt pexrav} model make it impossible to determine whether the noted correlation between $R - \\Gamma$ is physical.  Sources with low $\\Gamma$ tend to have higher N$_H$, low $R$, and a range of luminosities.  Sources with high $\\Gamma$ tend to have low N$_H$ and high $R$.  Cutoff energies, which are related to the Comptonization parameter ($y$), corona temperature, and optical depth towards Compton scattering, are not well-constrained for the majority of the sample.  Where they are constrained, there is no correlation with $\\Gamma$.  In the future, we will compare the {\\tt pexrav} model results with those from alternative reflection models, such as the dusty torus model MYTORUS \\citep{2009MNRAS.397.1549M}, but such analysis is beyond the scope of the present paper.   \\subsubsection{\\bf Soft Excess}~\\label{softexcess} An X-ray soft excess, which we modeled with a simple blackbody, is statistically significant in 45/48 (94\\%) of our sources.  Of the three sources without a significant blackbody component, the spectrum of Mrk 926 is poorly fit and the spectrum of 2MASX J11454045$-$1827149 has a relatively short exposure time with XMM-Newton (9\\,ks).  Further, a soft excess is detected in a long XMM-Newton spectrum of the final source, NGC 5548  \\citep{2003MNRAS.341..953P}.  Therefore, a soft excess is likely present in all local AGN.  This contrasts with our previous analyses of the Swift BAT sources, which found a soft excess in $40-50$\\% of AGN \\citep{2008ApJ...674..686W,2009ApJ...690.1322W}.  The higher detection rate in the current study is likely due to the higher signal-to-noise observations used in this study.  In  \\citet{2009ApJ...690.1322W}, we used data from a variety of sources, with the majority of spectra from ASCA and lower signal-to-noise Swift XRT spectra.  This made it difficult to detect fainter blackbody components.  The average and standard deviation for the best-fit blackbody temperature ($kT$) corresponds to $0.11 \\pm 0.06$\\,keV (BLRGs), $0.11 \\pm 0.04$\\,keV (Sy 1s), $0.32 \\pm 0.60$ (Sy 1.2s), and $0.15 \\pm 0.22$ (Sy 1.5s).  The average value for the Sy 1.2s is skewed by one point  (MCG -01-13-025 has a best-fit $kT = 2.0$\\,keV, with large error bars) and if this point is disregarded the average and standard deviation ($0.11 \\pm 0.04$\\,keV) are in line with the other sources in our sample.  Both the measured blackbody temperature and the small amount of scatter in our sample are consistent with measurements of sources in the Lockman Hole from the XMM-Newton survey \\citep{2005AA...444...79M}, PG-selected QSOs \\citep{2004AA...422...85P, 2005AA...432...15P,2010ApJ...725.1848T}, and our previous analyses of the Swift-selected sources \\citep{2008ApJ...674..686W,2009ApJ...690.1322W}.  The origin of the soft excess in AGN is as yet unknown.  It may arise from thermal emission from star formation (particularly as seen in ULIRGs or AGN hosted in galaxies with strong nuclear starbursts), a population of near nuclear X-ray binaries/ULXs  \\citep{2010MNRAS.407.2399M}, blurred reflected emission \\citep{2003AA...412..317C, 2005MNRAS.358..211R}, or blurred absorption \\citep{2004MNRAS.349L...7G}.  In Figure~\\ref{fig-bbodytemp}, we test whether a thermal model is plausible.  If the emission is thermal and associated with the black hole, we expect the blackbody temperature to be proportional to $M^{-1/4} L/L_{Edd}^{1/4}$, where $M$ is the black hole mass and $L/L_{Edd}$ is the accretion rate.  We find no correlation between the blackbody temperature and either of these parameters, in agreement with the results of our previous analysis of the X-ray spectra of the Swift BAT-selected AGN \\citep{2009ApJ...690.1322W}.  However, we do find a slight correlation between the flux in the power-law and the flux in the blackbody component.  In Figure~\\ref{fig-ktnorm}, we plot the blackbody normalization versus the power-law normalization.  The correlation ($\\log kT\\,{\\rm norm} = (-1.64 \\pm 0.41) + (1.16 \\pm 0.19) \\log \\Gamma\\,{\\rm norm}$; excluding the Sy 1.5s, which have more uncertainty in determining the soft excess parameters due to complex absorption) is weak, with $R^2 = 0.28$, but it indicates the same direct correlation shown from our more careful comparison of the luminosity in the blackbody and power-law components in  \\citet{2009ApJ...690.1322W} (i.e., $L_{power-law} \\propto L_{soft excess}$).  Similarly, this correlation is also seen in the PG QSOs  \\citep{2010ApJ...725.1848T}.  The relationship between the power-law and soft excess shows that for the Sy 1s, the soft excess is either created by or affected by the direct AGN emission and not from thermal emission from star formation.  Further clues to the origin of the soft excess are seen in a comparison with the absorbing gas.  Namely, we find that sources with the largest soft excess also have the highest neutral column densities and the strongest absorber signatures (measured through the optical depth of the O\\,VII absorption edge).  These results are clearly shown in Figure~\\ref{fig-nhoviiktnorm}.  Binning these relationships, shows that the correlations are very strong.  We find that \\begin{equation} \\log A_{\\rm kT} = (1.74 \\pm 0.26) \\log N_{\\rm H} + (-38.77 \\pm 5.50), \\end{equation} with a correlation coefficient of $R^2 = 0.88$.  We also find that \\begin{equation} \\log A_{\\rm kT} = (1.62 \\pm 0.19) \\log \\tau_{\\rm O\\,VII} + (-0.58 \\pm 0.35), \\end{equation} with $R^2 = 0.92$.  Therefore, the soft excess, neutral column density, and warm ionized gas are connected.   Since the soft excess is also correlated with the direct emission, these processes must all be related (i.e., the direct power-law emission, soft excess, absorption).  %   \\subsubsection{\\bf \\ion{Fe}{1} K$\\alpha$ Emission} The Fe K$\\alpha$ emission feature is typically the most prominent feature in AGN spectra.  Recent work shows that the hard X-ray region surrounding the Fe K$\\alpha$ band can be quite complex, with both narrow and broad features from \\ion{Fe}{1} K$\\alpha$, additional narrow features from, e.g., \\ion{Fe}{25} K$\\alpha$, \\ion{Fe}{1} K$\\beta$, and \\ion{Ni}{1} K$\\alpha$, iron K-edges (7.11\\,keV), and absorption features from outflowing highly ionized gas.  We focus only on the narrow \\ion{Fe}{1} K$\\alpha$ line in this work, and defer further analysis of the hard X-ray emission/absorption to a future study.  We note, however, that many of our sources are included in more in-depth studies, such as the Chandra HETG study in \\citet{2004ApJ...604...63Y} and the XMM-Newton study in \\citet{2007MNRAS.382..194N}.  Details of the best-fit parameters for the \\ion{Fe}{1} K$\\alpha$ emission are included in Table~\\ref{tbl-fek}.  For many of our sources (40\\% or 19/48), we could not constrain the energy or the width of the emission line.  Instead, we fixed these values to default values of 6.41\\,keV and 0.01\\,keV, respectively, to determine limits on the equivalent width.  We find that for the sources with well-constrained energies and widths, the average and standard deviation correspond to $E = 6.41 \\pm 0.04$\\,keV and a range in width from $\\sigma = 0.06 - 0.14$\\,keV (corresponding to ranges from $FWHM \\approx 2780 - 9180$\\,km\\,s$^{-1}$). Therefore, while the energy of this line is very similar for our sources, we find that the width varies between sources.  However, the average value is still below the resolution of the XIS at 6.4\\,keV, which is $\\sim 130$\\,eV. Still, our best-fit values from the Suzaku and XMM-Newton spectroscopy are in good agreement with Chandra grating results for 14 Seyferts, which found $E = 6.404 \\pm 0.005$\\,keV and a $FWHM = 2380 \\pm 760$\\,km\\,s$^{-1}$ \\citep{2004ApJ...604...63Y}.  Estimates of the $EW$ were made for all of the sources in the sample.  We find average and standard deviations of the equivalent width of the \\ion{Fe}{1} K$\\alpha$ emission of $108.96 \\pm 190.34$\\,eV.  Therefore, there is also a large spread in the $EW$, as well as the width, of the emission feature.  The source with the strongest $EW$ line in our sample, with $EW \\approx 1.3$\\,keV, is EXO 055620-3820.2, whose 2006 observation corresponds to a low flux, possibly Compton-thick phase (see \\citealt{2009MNRAS.394L...1L} for a comparison of this observation with previous data showing a high flux state).  Disregarding this outlier, we find $EW = 83.10 \\pm 70.08$\\,eV.  While past studies found an anti-correlation between the equivalent width and luminosity of AGN (the X-ray Baldwin effect \\citealt{1993ApJ...413L..15I}), we find no correlation -- in agreement with our earlier results on the Swift sources \\citep{2009ApJ...690.1322W}.  In Figure~\\ref{fig-fekew}, we show the relationship between the \\ion{Fe}{1} K$\\alpha$ equivalent width and both the bolometric luminosity and Eddington ratio.  Linear correlations are not seen with either parameter, with correlation coefficients of $R^2 = 0.25$ and $0.19$, respectively.  We previously found this in our \\citet{2009ApJ...690.1322W} study, determining that a correlation exists for multiple observations of individual sources or when the luminosity/accretion rate proxy are binned.  To test this further, we looked for evidence of the X-ray Baldwin effect, by binning the data by luminosity and accretion rate.  As shown in Figure~\\ref{fig-fekew}, there is no strong correlation between the \\ion{Fe}{1} K$\\alpha$ equivalent width and the bolometric luminosity.  We find that $EW \\propto L_{\\rm bol}^{-0.27 \\pm 0.08}$, but with a correlation coefficient of $R^2 = 0.36$ ($P \\sim 0.1$).  This correlation is consistent, however, with the relationship we found ($EW \\propto L^{corr}_{2-10\\,{\\rm keV}}$) for the entire 9-month sample \\citep{2009ApJ...690.1322W}.  We find a strong correlation between the \\ion{Fe}{1} K$\\alpha$ equivalent width and accretion rate.  The linear relationship is parameterized as \\begin{equation} \\log EW = (-0.38 \\pm 0.07) \\log L_{\\rm bol}/L_{\\rm Edd} + (1.35 \\pm 0.08), \\end{equation} with $R^2 = 0.80$.  This relationship is consistent with the\\\\ $EW \\propto (L^{corr}_{2-10\\,{\\rm keV}}/L_{\\rm Edd})^{(-0.26 \\pm 0.03)}$ relationship found for the entire Swift sample \\citep{2009ApJ...690.1322W}.  This confirms that the primary driver for the observed X-ray Baldwin effect is the correlation between the $EW$ and accretion rate, as suggested by \\citet{2006ApJ...644..725J}.   \\subsection{Warm Absorbers} Using the base continuum model described in the preceding section, we constrained the properties of potential warm absorber signatures in our sample.  The simplest model for detecting potential outflows in X-ray CCD data is with the addition of absorption edge models fit to the 0.73\\,keV \\ion{O}{7} and 0.87\\,keV \\ion{O}{8} features.  We describe details of these fits in \\S~\\ref{o7}.  We then fit the spectra with more sophisticated models using the {\\tt warmabs} model, which determines the ionization state and column density of the warm absorbing gas, described in \\S~\\ref{warmabs}.  We discuss our conclusions on the warm absorption fits in \\S~\\ref{subsect-conclusions}.  \\subsubsection{\\ion{O}{7} and \\ion{O}{8} Absorbers}\\label{o7} The results of adding absorption edges to detect \\ion{O}{7} and \\ion{O}{8} edges are included in Table~\\ref{tbl-warmabs}.  Among our sources, 25/48 (52\\%) have clear detections of the edge features.  This fraction is higher than our original report of 18/44 (41\\%) in \\citet{2010ApJ...725L.126W}, since we added analysis of 4 sources with recently available spectra (NGC 985, 2MASX J11454045$-$1827149, NGC 6814, and 2MASX J21140128+8204483) and re-analyzed longer exposures available for 3 additional sources (IC 4329A, NGC 5548, and Mrk 926).  As in \\citet{2010ApJ...725L.126W}, we classify a detection based on $\\Delta\\chi^2 \\ga 13.39$, which corresponds to a probability of $P = 0.01$ for the four additional degrees of freedom added.  The detection rate for sources with Suzaku spectra is 18/34 (53\\%), while the detection rate for XMM-Newton spectra is 7/14 (50\\%).  We find no relationship between $\\Delta\\chi^2$, on adding the edges, and the total number of counts in the spectrum (i.e., there are spectra with few counts and strong detections of absorption edges, as well as spectra with many counts and low $\\Delta\\chi^2$ on adding the absorption edges).  Likewise, there is no relationship between the measured optical depth of the absorption features and the number of counts in the spectrum (e.g., the four sources with the highest $\\tau$ in the \\ion{O}{7} edge have between $\\sim 1.5 - 4.2 \\times 10^{5}$\\,counts, while spectra in our sample range from total counts of $\\sim 5 \\times 10^{5} - 2 \\times 10^{6}$).  Therefore, we find no evidence for biases in our detection rates or measured edge strengths with the total counts in the spectra.  In the companion paper, \\citet{2010ApJ...725L.126W}, we showed that the detection of \\ion{O}{7} and \\ion{O}{8} is dependent on luminosity, accretion rate, and column density.  In particular, we found that detection rates are higher in sources with larger neutral column densities and lower in the more luminous sources (see Table 2 and Figure 2 in \\citealt{2010ApJ...725L.126W}).  Additionally, we find that the strength of the warm absorber features are also dependent upon column density.  This is illustrated in Figure~\\ref{fig-eddingtonlimit}, where the optical depth in the \\ion{O}{7} edge is conveyed by the size of the symbols (i.e., the largest symbols correspond to the highest optical depths).  In the plot, the neutral column densities are from Table~\\ref{tbl-fits1}, with the exception of NGC 3783.  Determinations of the neutral column density of NGC 3783 are available from several observations from 2001--2007 utilizing XMM-Newton, Chandra, and Suzaku.  The published measurements range from $\\sim 5 \\times 10^{21} - 10^{22}$\\,cm$^{-2}$ \\citep{2004ApJ...602..648R,2007MNRAS.379.1359M,2009PASJ...61.1331M}.  The source is variable in the X-rays and it is unclear whether our low measurement in the 2009 Suzaku observation represents a change in column density from the earlier epochs or is the result of difficulties measuring the column in such a complex spectrum.  In the plot, we adopt N$_{\\rm H} = 5 \\times 10^{21}$\\,cm$^{-2}$ for NGC 3783.  The N$_{\\rm H}$-Eddington ratio plot is a useful illustration of the effects of radiation pressure on dusty gas surrounding the AGN (see for instance \\citealt{2006MNRAS.373L..16F}, \\citealt{2008MNRAS.385L..43F}, and \\citealt{2010MNRAS.402.1081V}).  In Figure~\\ref{fig-eddingtonlimit}, the blue line represents the {\\it effective} Eddington limit, where gravitational pressure balances with radiation pressure on dusty gas with solar abundances for the dust grains, as computed in  \\citet{2006MNRAS.373L..16F}.  The radiation pressure on dusty gas is enhanced relative to that of ionized gas; such that the the Eddington limit for dusty gas is lower.  Outflows driven by radiation pressure in a dusty medium are expected when both the Eddington ratio and the neutral column density are high (i.e., where the effective Eddington limit is exceeded).  Sources in this region are believed to be short-lived, as the radiation pressure will eventually drive away the obscuring dusty material, leaving unobscured, low Eddington ratio sources.  The sources with the strongest absorption features, i.e., highest optical depth for \\ion{O}{7} absorption, clearly fit with the Fabian et al., model.  They have high column densities and Eddington ratios near or above the effective Eddington limit for dusty gas.  The estimates of the Eddington ratios for 4/6 of these sources are secure, since the mass estimates are based on reverberation mapping \\citep{2004ApJ...613..682P}.  Therefore, a wind driven by radiation pressure on the dusty gas is a likely mechanism for creating the absorption features seen in these sources.  Assuming that the absorption edges trace outflowing ionized gas, the sources with the strongest outflow detections are NGC 3516, NGC 4151, Mrk 6, NGC 3227, NGC 526A, and NGC 3783.  For NGC 3516, NGC 4151, NGC 3227, and NGC 3783, UV and X-ray grating results confirm that the warm ionized gas is indeed outflowing, with measured absorption lines blue-shifted by $\\sim 100 - 600$\\km~(see, for instance,  \\citealt{1999ApJ...516..750C,2001ApJ...555..633C} for analysis of UV Hubble spectroscopy; also note that the low-ionization lines of NGC 3227 are red-shifted, whereas the higher ionization lines are blue-shifted).  Neither Mrk 6 or NGC 526A have available UV spectroscopy, as these are highly reddened sources, and X-ray grating observations of these sources do not have sufficient signal-to-noise to confirm that the gas is outflowing.  Since the strong outflow sources also have complicated spectra, we searched the literature for confirmation that the high neutral column densities we measured were also detected in alternative spectral fits.   We already described above the column density measurements for NGC 3783, which indicate N$_H \\ga 5 \\times 10^{21}$\\,cm$^{-2}$.  \\citet{2011ApJ...731...21M} find that there is a variable absorber in Mrk 6, with N$_{\\rm H} \\sim 10^{21} - 10^{23}$\\,cm$^{-2}$ over the past 6 years of available X-ray spectroscopy.  The ASCA and Chandra analyses of NGC 4151 also reveal variable high column density gas (e.g., \\citealt{2010ApJ...714.1497W}).  For this source, we note that the soft X-rays are dominated by emission features, with a weak and heavily absorbed continuum, and that the grating observations reveal that the edges from \\ion{O}{7}, \\ion{O}{8}, \\ion{Ne}{9}, and \\ion{Ne}{10} are blended together \\citep{2005ApJ...633..693K}.  Simple fits to the Suzaku data of NGC 4151 indicate that there is a strong component of photoionized emission, as in the earlier observations, but a more comprehensive analysis of the complex spectrum is beyond the scope of the current paper.  Despite the complexity of NGC 4151's X-ray spectra, the grating observations clearly indicated a high column density ($\\sim 10^{22}$\\,cm$^{-2}$) of absorbed gas, consistent with this study, particularly when NGC 4151 is in a lower flux state (as is the case in the Suzaku spectrum presented here).  Since the soft continuum is very weak and dominated by strong emission, we can not constrain the warm absorption properties of NGC 4151.  We can say that it is heavily absorbed and that the complexity of this source is unique in our sample, as no other source shows as weak a continuum or as strong of emission features.  An independent analysis of the XMM-Newton spectrum of NGC 526A agrees with our measured column of $\\sim 10^{22}$\\,cm$^{-2}$ \\citep{2011MNRAS.413.1206B}.  NGC 3516 also is known to exhibit high column density ($\\ga 10^{22}$\\,cm$^{-2}$) gas in its X-ray spectrum \\citep{2008PASJ...60S.277M}. The XMM-Newton analysis by \\citet{2009ApJ...691..922M} of NGC 3227 uncovers a lower neutral component, $9 \\times 10^{20}$\\,cm$^{-2}$, than our measured value of $\\sim 1.6 \\times 10^{22}$\\,cm$^{-2}$, but since the source's column density has been observed to vary in the past and there is no published analysis of the Suzaku spectrum we analyzed, it is unclear how our results compare.  As a whole, our literature search supports our conclusion that the sources with the strongest outflow detections are associated with higher column densities of gas.  Our spectral fits allow insight into the spectral properties of sources with strong outflows.  In particular, we identify signs of strong absorbers that can be used in lower quality data, for instance, for low flux or higher redshift sources where X-ray grating spectroscopy is infeasible.  Two clear signs of sources with strong outflows are a flat power-law index and a strong soft excess. We find that among the sources with the strongest warm absorber detections, all of these sources have measured values for $\\Gamma \\approx 1.0$.  Two other sources with $\\Gamma < 1.5$ also have complex absorption in their spectra: EXO 055620-3820.2 and NGC 5548.  These low measured values are likely due to difficulties in modeling the spectra, which are absorbed by multiple components of dusty and ionized gas.  Likely, the ``true'' direct emission of these sources is a $\\Gamma \\sim 1.9$ power-law, as found for the majority of our sources, but the absorption distorts the soft X-ray spectrum, flattening the fit.  Even with the extended hard X-ray spectra from Suzaku HXD and Swift BAT, the intrinsic power-law continuum is difficult to uncover from the fitting process.  While a flat spectrum is not physically representative of the direct emission, it is an easy diagnostic for identifying sources with complex absorbed spectra, both for Seyfert 1s and heavily absorbed Compton-thick spectra, in lower signal-to-noise data.  Additionally, we find that sources with strongly detected outflows (confirmed through the archived grating and UV data) also have the strongest soft excesses (see also \\S~\\ref{softexcess}). As with the flat power-law index, a strong soft excess can be an effect of strong absorption features distorting the spectrum.  Alternatively, the soft excess could be related to the warm ionized absorbing gas.  If the soft excess is a feature of the outflow/ionized absorption, this is in agreement with the soft excess as complex absorption model \\citep{2004MNRAS.349L...7G,2007MNRAS.374..150S}.   From our spectral fitting of  \\ion{O}{7} and \\ion{O}{8} absorption edges, we find that the 0.73\\,keV \\ion{O}{7} edge tends to have the higher optical depth.  For the sources classified as not having strong warm absorber detections, based on $\\Delta\\chi^2$, only 10 sources have an optical depth of zero in the \\ion{O}{7} feature.  This suggests that warm ionized absorbing gas is present in at least 80\\% of the sample.  However, it is unclear whether this material is outflowing, as high signal-to-noise UV or X-ray grating observations are needed to measure the velocity shifts in individual absorption lines.  Among the sources with zero optical depth, three sources are BLRGs, four are Sy 1s, two are Sy 1.2s, and one is a Sy 1.5.  All of these sources have low neutral column densities (N$_{\\rm H} \\la 10^{20}$\\,cm$^{-2}$).   In \\citet{2010ApJ...725L.126W}, we discussed how the published observations of each of these three BLRGs reveal ionized absorption, with ionization parameters higher or lower than are expected to create the \\ion{O}{7} and \\ion{O}{8} edges.  It is possible that this is the case for the remaining seven sources, as well, in which case the covering fraction of an outflow is $\\Omega \\sim 1$.  In the following section, we explore this through a discussion of our analytic warm absorber fits.  With these models, we determined both the column density and ionization parameter of ionized gas in our entire sample, including both the sources with and without clear detections of \\ion{O}{7} and \\ion{O}{8} absorption.    \\subsubsection{Analytic Warm Absorption Models}\\label{warmabs} In \\citet{2010ApJ...725L.126W}, we claimed that while the detection rate of \\ion{O}{7} and \\ion{O}{8} edges was low in the highest luminosity sources (33\\%) and high in the low luminosity sources (60\\%), more detailed spectral fits reveal the presence of ionized gas that in many cases is either more or less ionized than required to produce strong \\ion{O}{7} and \\ion{O}{8} features.  As a follow-up to this work, we fit analytic models to the X-ray CCD spectra.  Results of these fits are included in Tables~\\ref{tbl-warmabs-outflow} and \\ref{tbl-warmabs-nonoutflow}.  The X-ray CCD data does not have the energy resolution to accurately determine the velocity of outflow components.  However, we were able to determine both the column density of warm ionized gas (N$_{\\rm warm}$) and the ionization parameter ($\\xi =  {\\rm L}_{\\rm ion}/(n_e R^2)$; where L$_{\\rm ion}$ is the ionizing luminosity, $n_e$ is the electron density, and $R$ is the distance of the ionized gas from the central ionizing source) for our entire sample.  For the sources with clear detections of absorbers through \\ion{O}{7} and \\ion{O}{8} edges, we fit the spectra with two ionized components. We add only one warm absorber component for the sources without strong detections of the edge features, to determine limits on ionized absorbers present.  For the majority of sources in Table~\\ref{tbl-warmabs-outflow}, the two component analytic model improves the fits significantly, as indicated by the $\\Delta\\chi^2$ values (e.g., out of the 25 sources fit with a two-component model, only 8 show a second component to have low significance with $\\Delta\\chi^2 < 20$ upon adding the second warm absorber).  \\subsubsection*{Caveats on the Analysis and Detection of Ionized Absorbers in CCD Data} While ideally we would compare our results to those from an X-ray grating analysis, the complication of X-ray variability makes such comparisons difficult.  For instance, while 8 sources overlap between our analysis and the uniform Chandra grating analysis in \\citet{2007MNRAS.379.1359M}, a comparison of the observed soft fluxes from our observations and the Chandra observations reveals that 6 of the sources vary considerably.  We find, for example, that NGC 4051 is ten times brighter in the Suzaku observation, while NGC 4593 is four times fainter.  Therefore, there is no consistent way to compare our results unless we know that a source does not vary or we analyze simultaneous grating and CCD data.  In \\citet{2010ApJ...719..737W}, such an analysis is presented using the 2009 XMM-Newton CCD and grating spectrum of NGC 6860, along with a Suzaku observation, taken a year earlier, which showed no variability from the XMM observation.  In this analysis, the best-fit neutral and ionized column densities were consistent within the error-bars between the Suzaku CCD, XMM-Newton CCD, and XMM-Newton grating observations.  We find, however, that the ionization parameter, consistent in both CCD analyses ($\\log \\xi \\sim 2.0$), is higher in the grating observation ($\\log \\xi = 2.4$).  The second warm absorber fit had best-fit values that varied considerably between the CCD and grating measurements.  Since multiple components are present in AGN spectra and the quality of the grating spectrum from NGC 6860 was relatively low signal-to-noise, it is unclear how these results reflect on our entire sample.  However, it is likely that the measured column densities are accurate, particularly for significant warm absorption components.  %  It is also important to note that since we do not have velocity information on the detected ionized absorbers, due to the low resolution of the CCD data, we can not confirm that the absorbers originate in an AGN driven outflow (i.e., due to the low velocity resolution, we fixed the velocity of the ionized absorbers to the systemic velocity of the AGN).  We have shown that for the sources with the strongest detections, outflows are likely present, since UV/X-ray grating observations confirm blue-shifted absorption features.  For the remaining sources, the absorbing gas could be intrinsic to the AGN or a feature from ionized gas in our own Galaxy, the host galaxy, or intervening systems (signatures of the warm hot ionized medium).  A Chandra grating study by \\citet{2004ApJ...617..232M}, for instance, detected oxygen absorption in half of a sample of 15 type 1 AGN sources that was identified as hot gas from local interstellar structures and potential intergalactic medium features.  In order to test whether the ionized gas is outflowing from the AGN, high signal-to-noise X-ray grating observations and/or ultraviolet spectra are needed to measure the velocity of the absorption features.  In future work, we will present an analysis of the X-ray grating and ultraviolet spectra of our sample.  \\subsubsection*{Results of the Analytic Warm Absorption Spectral Fits} For the six sources with the strongest \\ion{O}{7} and \\ion{O}{8} absorption features, we find average column densities of warm ionized gas of a few $10^{21}$\\,cm$^{-2}$ in each of the two absorption components fit to the spectra.  All but NGC 3227 have a warm ionized component with an ionization parameter near $\\xi \\approx 100$\\,ergs\\,s$^{-1}$ (NGC 3227's highest measured ionization component is $\\xi \\approx 28$\\,ergs\\,s$^{-1}$).  Out of the entire sample of sources with detected \\ion{O}{7} and \\ion{O}{8} edges, we find that the average column and ionization parameters, for the component with the highest warm ionized column density, are $N_{\\rm  warm} = (3.8 \\pm 9.1) \\times 10^{21}$\\,cm$^{-2}$ and $\\log \\xi = 2.04 \\pm 1.27$.  We find $N_{\\rm warm} = (3.9 \\pm 4.1) \\times 10^{20}$\\,cm$^{-2}$ and $\\log \\xi = 1.59 \\pm 1.69$ for the second component.  These values are in line with the range in column density (N$_{\\rm warm} \\sim 10^{20} - 10^{23}$\\,cm$^{-2}$) and ionization parameter ($\\xi \\sim 10^{0-4}$\\,ergs\\,s$^{-1}$) found from studies of X-ray grating observations (see \\citealt{2005AA...431..111B} and \\citealt{2007MNRAS.379.1359M}).  For the sources without strong detections of \\ion{O}{7} and \\ion{O}{8}, we find $N_{\\rm warm} = (1.8 \\pm 2.2) \\times 10^{20}$\\,cm$^{-2}$ and $\\log \\xi = 1.09 \\pm 2.04$.  Of these, three have significant warm absorption detections with the analytic model, including, 1H 0419-577 (low column and low ionization parameter), EXO 055620-3820.2 (low column and low ionization parameter), and 3C 382 (with low column and a higher ionization parameter).  To test whether the analytic model result of a very low ionization parameter in the spectrum of 1H 0419-577 (the source with the highest significance with the analytic model in Table~\\ref{tbl-warmabs-nonoutflow} and lowest measured ionization parameter) was plausible, we alternatively fit the spectrum with an edge model for lowly ionized oxygen (\\ion{O}{1}/\\ion{O}{2}) at 0.545\\,keV.  We found that the lowly ionized edge was very significant, with $\\Delta\\chi^2 = 73.3$, and find an optical depth of $\\tau = 0.15^{+0.03}_{-0.03}$.   Our results, then, show that sources without strong detections of \\ion{O}{7} and \\ion{O}{8} absorption edges tend to have lower column densities and lower ionization parameters, on average, than sources with strong detections.  In particular, the column densities of warm ionized gas are an order of magnitude lower in the non-detection sources.  The ionization parameters, however, do show a large range of values, including both lowly ionized and highly ionized gas.  In Figure~\\ref{fig-nwarm1}, we compare the results on the properties of the warm ionized gas in sources with strong \\ion{O}{7} and \\ion{O}{8} absorption edges.  Component 1 is chosen as the model component with the highest column density from Table~\\ref{tbl-warmabs-outflow}.  In the top plots, we compare the ionization parameters to the bolometric luminosity.  We find that the ionization parameter, $\\xi$, has no dependence on the luminosity of the AGN.  Instead, we find the majority of values clustered near $\\xi = 100$\\,ergs\\,s$^{-1}$ for component 1 and a broader range in $\\xi$ for component 2.  This method is not particularly sensitive to very high $\\xi$ ($>3.5$) or very low  $\\xi$ ($<1.0$) absorbers, since those absorber components are driven by fits to a small number of highly ionized lines (e.g. \\ion{Si}{14} Lya) or Fe UTAs, which require grating spectral resolution to fit properly. The Fe UTAs in particular may show up in the \\ion{O}{7} and \\ion{O}{8} edge fits.  The most likely scenario to account for the different components is that the absorbing region consists of a warm ionized medium with higher density blobs or filaments embedded within it \\citep{2001ApJ...561..684K}.  In this case, the sources without the strong detections of \\ion{O}{7} and \\ion{O}{8} have similar ionization parameters with the lower density, ionized medium also detected in the sources with 'strong' absorber detections.               \\subsubsection{Conclusions on Warm Absorber Properties}\\label{subsect-conclusions} The most comparable previous studies to the work we present on the warm absorber properties in the Swift-detected Seyfert 1s are the ASCA studies of  \\citet{1997MNRAS.286..513R} and \\citet{1998ApJS..114...73G}.  In \\citet{2010ApJ...725L.126W}, we discussed the overlap in our sample with the  \\citet{1997MNRAS.286..513R} study.  Out of the 24 sources in the ASCA study, 18 are also in our sample.  Between publication of \\citet{2010ApJ...725L.126W} and our present results, we revised the classification of two sources (IC 4329A and NGC 5548) to have warm absorption detections.  Thus, we find good agreement in our classification scheme with the \\citet{1997MNRAS.286..513R} study, with 16/18 sources classified accordingly.  Both 3C 382 and 3C 390.3 were classified by  \\citet{1997MNRAS.286..513R} as exhibiting \\ion{O}{7} and \\ion{O}{8} features in the ASCA data, while our analysis shows that these sources do not have significant detections.  Grating spectroscopy of these sources reveal that ionized outflows are present, but at higher ionization parameters \\citep{2010MNRAS.401L..10T,2009ApJ...700.1473S}.  Our detection rate for warm absorption features, revealed by the significance of \\ion{O}{7} and \\ion{O}{8} absorption edge features, is 52\\%.  This is in agreement both with the more biased samples presented in the ASCA study of  \\citet{1997MNRAS.286..513R} and the UV outflow studies of \\citet{1999ApJ...516..750C}, but much lower than the $\\sim 70$\\% detection rate reported in \\citet{1998ApJS..114...73G}.  While our detection rate, assuming that the warm absorbers are produced in outflowing gas, initially suggests a covering fraction of AGN outflows of $\\Omega \\sim 0.5$, there is clearly more that needs to be considered in this simple picture.  In \\citet{2010ApJ...725L.126W}, we pointed out that the detection rate of \\ion{O}{7} and \\ion{O}{8} edges was related to the luminosity of the AGN.  There are fewer detections in the most luminous ($\\sim 30$\\%) and higher detection rates for the least luminous ($\\sim 60$\\%) sources.  Grating observations, however, reveal that the detection rate is higher -- consistent with outflows in at least 80\\% of the most luminous sources.  In the current study, we find that the main difference between sources with and without \\ion{O}{7} and \\ion{O}{8} detections is the measured column density of potentially outflowing ionized gas.  We find that sources with strong detections have an order of magnitude or higher ionized column densities than those without detections.  Therefore, we can not rule out that sources without \\ion{O}{7} and \\ion{O}{8} detections do not have warm absorbers present.  Our results support the hypothesis that they do have ionized gas, just at lower column densities than are easily detected in the X-ray data, which is most sensitive to absorption with N$_{\\rm H} > 10^{20}$\\,cm$^{-2}$.  This is still orders of magnitude higher than the ionized gas detected in the UV, which has column densities from N$_{\\rm ion} \\sim 10^{12} - 10^{14}$\\,cm$^{-2}$.  Through this work, we suggest a change in the paradigm of assuming a 50\\% covering fraction for AGN warm absorbers (and by extension outflows, since grating observations reveal that this gas is typically outflowing) to $\\sim$100\\%.  Our analysis of an unbiased sample of type 1 AGNs selected in the very hard X-rays with Swift shows that while the covering fraction is likely unity, there are clear distinctions between the higher column density warm absorbers, which have neutral and ionized columns  $> 10^{20}$\\,cm$^{-2}$, and lower column density warm absorbers.  The higher column density sources make up half of the sample.  Among these sources with strong detections of warm absorbers/outflows, we suggest that the strongest outflows are driven by radiation pressure on dusty gas, as claimed by  \\citet{2006MNRAS.373L..16F}.  A main goal of AGN outflow study work is to determine how much mass and energy is entrained in the outflow.  While work such as the grating study of \\citet{2005AA...431..111B} places estimates on the energy in the outflows (e.g., \\citealt{2005AA...431..111B} estimate that outflows account for less than 1\\% of the bolometric AGN luminosity), many assumptions, particularly relating to the geometry of the outflow, are necessary to make these estimates (see, for instance, \\citealt{2007MNRAS.379.1359M} for a discussion of these assumptions).  Therefore, we do not include estimates of the mass outflow rate or the total energy in the current paper.  In future work, we will follow-up the current study with a uniform analysis of the archived grating observations, available for the majority of our sample but without as uniformly high a signal-to-noise ratio as in the CCD data presented in this work.   Additionally, estimates of the total outflow rate/energy must rely on multi-wavelength data, characterizing the AGN outflow contribution from neutral through the most highly ionized gas (e.g., \\citealt{2011arXiv1109.2882T} detect Fe K-shell absorption lines in $> 35$\\% of their sample of 42 low redshift, radio-quiet AGN), and include studies of the obscured AGN (X-ray type 2) sources, whose UV and soft X-ray ionized absorber properties are difficult to determine.  We will explore both of these lines of research in our future work."
    },
    "1112/1112.2918_arXiv.txt": {
        "abstract": "Helium abundances and atmospheric parameters have been determined from high resolution spectra for a new sample of 46 bright hot subdwarf B (sdB) stars. The helium abundances have been measured with high accuracy. We confirm the correlation of helium abundance with temperature and the existence of two distinct sequences in helium abundance found previously. We focused on isotopic shifts of helium lines and found $^{3}$He to be strongly enriched in 8 of our programme stars. Most of these stars cluster in a small temperature range between $27\\,000\\,{\\rm K}$ and $31\\,000\\,{\\rm K}$ very similar to the known $^{3}$He-rich main sequence B stars, which cluster at somewhat lower temperatures. This phenomenon is most probably related to diffusion processes in the atmosphere, but poses a challenge to diffusion models. ",
        "introduction": "The atmospheric helium abundances of sdBs are poorly understood. From the typical abundance patterns of these stars, which show a depletion of light elements as well as an enrichment of heavy metals, it has been concluded that diffusion play an important role in their atmospheres. However, diffusion models predict an almost total depletion of helium in contrast to what is observed. The helium abundances of sdBs range from slightly above solar down to $\\log{y}<-4$. Mass loss caused by stellar winds as well as extra mixing in the atmosphere have been invoked to counteract gravitational settling and explain the observed helium abundances \\citep[see ][ and references therein]{hu11}.  \\citet{edelmann03} found a correlation of helium abundance with temperature. The hotter the sdB, the more helium is present in its atmosphere. Similar correlations have been found by other groups (see e.g. Vennes et al. these proceedings). However, \\citet{edelmann03} also reported the discovery of two distinct sequences showing a similar correlation with temperature, the ''lower sequence'' being offset by about $2\\,{\\rm dex}$ from the upper sequence. The majority of stars lie on the ''upper sequence''. Those sequences could not be clearly identified in other datasets so far \\citep[e.g. ][]{lisker05, geier11}. \\citet{otoole08} combined the then published data sets and found the stars of the upper sequence to lie near the Extreme Horizontal Branch (EHB) band in the $T_{\\rm eff}-\\log{g}$-plane, as expected, whereas the lower-sequence stars lie in a much more dispersed area \\citep[see Figs.~2,3 in][]{otoole08}.  Gravitational settling can also lead to isotopic anomalies in stellar atmospheres. In the case of helium the light isotope $^{3}$He can be enriched with respect to the usually much more abundant $^{4}$He. Such an enrichment has initially been found in main sequence B stars with subsolar helium abundance \\citep{hartoog79}. However, \\citet{heber91} detected strong line shifts in the sdB star SB\\,290 and the blue horizontal branch star PHL\\,25 indicating that basically the whole helium content of the atmosphere consists of $^{3}$He. Later on \\citet{edelmann01} and \\citet{heber04} found another three sdBs, where $^{3}$He is enriched in the atmosphere.  Here we present the results of a quantitative spectral analysis of a sample of 46 sdB stars from high resolution spectra.  \\begin{figure} \\begin{center} \\includegraphics[width=10cm]{nhevsteff.eps} \\caption{Helium abundance $\\log{y}$ plotted against effective temperature. The filled symbols mark the results from our study. Filled red diamonds mark objects where isotopic shifts due to an enrichment of $^{3}$He were detected, filled circles objects with atmospheres dominated by $^{4}$He. Upper limits are marked with triangles. The solid horizontal line is drawn at solar helium abundance. The two dotted lines are regression lines for the two distinct helium sequences taken from \\citet{edelmann03}. Results taken from the literature are plotted as grey symbols \\citep{saffer94, maxted01, edelmann03, morales03, lisker05, geier11, vennes11, oestensen10}.} \\label{fig:abun} \\end{center} \\end{figure} ",
        "conclusions": "The distribution of these stars in the $T_{\\rm eff}$-$\\log{g}$-diagram is shown in Fig.~\\ref{fig:tefflogg} including the three sdBs with isotopic shifts taken from literature. It can be clearly seen that they cluster in a narrow temperature range between $27\\,000\\,{\\rm K}$ and $31\\,000\\,{\\rm K}$ with BD+48\\,2721 ($T_{\\rm eff}=24\\,800\\,{\\rm K}$) being the only exception. Given the uncertainties, this $^{3}$He-strip may be pure.  Most stars show clear shifts of the He\\,{\\sc i} line at $6678\\,{\\rm \\AA}$ indicating that almost all helium in the atmosphere is $^{3}$He. BD+48\\,2721, EC\\,12234$-$2607 and PG\\,1519+640 show strong lines of $^{3}$He blended with weak components of $^{4}$He. These three stars cover the whole $^{3}$He temperature strip. The isotope ratio is therefore not correlated to the effective temperature.  An $^{3}$He isotope anomaly has first been found for chemically peculiar main sequence stars of spectral type B. The $^{3}$He-stars were found at effective temperatures between $18\\,000\\,{\\rm K}$ and $21\\,000\\,{\\rm K}$ separating helium-poor stars at lower $T_{\\rm eff}$ from helium-rich stars at higher $T_{\\rm eff}$ \\citep{hartoog79}. In Fig.~\\ref{fig:tefflogg} a similar pattern can be seen for the sdBs. The stars enriched in $^{3}$He occupy a small strip in $T_{\\rm eff}$, while the helium abundance decreases towards lower temperatures and rises towards higher temperatures. \\citet{michaud11} carried out diffusion calculations and predict a mild enrichment of $^{3}$He, but due to gravitational settling of the heavier isotope this should be the case in all sdBs. Hence, the $^{3}$He strip stars lacks an explanation.  \\citet{hartoog79} argued that diffusion is responsible for this effect. At low temperatures the radiation pressure is not strong enough to support helium in the atmosphere. As soon as the temperature reaches a certain threshold value, the less massive $^{3}$He can be supported, but not the more abundant $^{4}$He. This leads to an enrichment of $^{3}$He in the atmosphere. At even higher temperatures both isotopes are enriched and the isotopic anomaly vanishes as the helium abundance rises."
    },
    "1112/1112.2403_arXiv.txt": {
        "abstract": "{ The stellar metallicity is a direct measure of the amount of metals present in a galaxy, as a large part of the metals lie in its stars. In this paper we investigate new stellar metallicity indicators suitable for high-z galaxies studying the stellar photospheric absorption lines in the rest frame ultraviolet, hence sampling predominantly young hot stars. We defined these new indicators based on the equivalent widths (EW) of selected features using theoretical spectra created with the evolutionary population synthesis code {\\it Starburts99}. We used them to compute the stellar metallicity for a sample  of UV-selected galaxies at $z>3$ from the AMAZE survey using very deep (37h per object) VLT/FORS spectra. Moreover, we applied the new metallicity indicators to  eight additional high redshift galaxies found in literature. We then compared stellar and gas-phase metallicities measured from the emission lines for all these galaxies, finding that within the errors the two estimates are in good agreement, with possible tendency to have stellar metallicities lower than the gas phase ones. For the first time, we are able to study the stellar mass-metallicity relation at $z>3$. We find that the metallicity of young, hot stars in galaxies at $z\\sim3$ have similar values of the aged stars in local SDSS galaxies, contrary to what observed  for the gas phase metallicity. } ",
        "introduction": "Metallicity is one of the important properties of galaxies, and its study is able to shed light on the details of galaxy evolution. It is, in fact, an integrated property, related to the whole past history of the galaxies. In particular, metallicity is sensitive to whole star formation history, and so to evolutionary stage of the galaxy. Moreover, it is affected by presence of infalls and outflows, i.e. by feedback processes and by the interplay between the forming galaxy and the intergalactic medium (see e.g. Erb et al. 2008, Mannucci et al. 2009, Cresci et al. 2010). As consequence, it has become an important test of galaxy evolution (e.g. Nagamine et al. 2001, Spitoni et al. 2010, Dav\\`e et al. 2011).\\\\ Local galaxies show a clear correlation between mass and metallicity (MZR), for which the galaxies with larger stellar mass have higher metallicities, and this correlation appears to hold both in term of gas-phase metallicity (e.g. Tremonti et al. 2004) and stellar metallicity (Gallazzi et al. 2005, Panter et al. 2008).\\\\  At high-redshift the gas-phase metallicity of the ISM of star-forming galaxies has been measured using primarily oxygen abundances. The most common techniques to determine the gas phase metallicity are based either on theoretical calibrations (see Kewley \\& Ellison 2002, and Kewley \\& Ellison 2008) or on empirical metallicity calibrations, the so-called ``strong line diagnostics'', which are based on the ratios of collisionally excited forbidden lines to hydrogen recombination lines. Previous studies have shown that the mass-gas phase metallicity relation presents evidence of strong redshift evolution. Among others, Savaglio et al. (2005) and Zahid et al. (2011) studied  star forming galaxies at redshift  $z\\sim0.7$ and demonstrated that, at given mass, these galaxies shown lower metallicity than the SDSS sample at $z\\sim0.1$. Erb et al. (2006) reported a more significant decrease of metallicity in galaxies at $z\\sim2.2$. Two projects were specifically designed to extend the investigation of MZR at  $z>3$: LSD (Lyman-break Stellar population and Dynamic) and AMAZE (Assessing the Mass-Abundance redshift Evolution). With these projects, Maiolino et al. (2008) and Mannucci et al. (2009) showed  for the first time the evolution of the mass-metallicity relation at $z>3$. However, the redshift evolution of the gas phase  metallicity in galaxies have been questioned recently by Mannucci et al. (2010). They discovered that metallicity depends not only on the mass, but also from the Star Formation Rate (SFR): for a given stellar mass, galaxies with higher SFRs systematically show lower metallicities. This is the so-called ``Fundamental Metallicity Relation (FMR)'', i.e., a tight relation between stellar mass, gas-phase metallicity, and star formation rate (SFR). Local SDSS galaxies show very small residuals around this relation, of the order of 0.05dex. Yates et al. (2011) found a similar relation, with some differences due to the metallicity calibration adopted. According to Mannucci et al. (2010), the FMR does not appear to evolve with redshift up to z$\\sim$2.5, with the high redshift galaxies following the same FMR defined by the local SDSS galaxies. This suggests that the observed evolution of the mass-metallicity relation is due to selection effects and to the increase of the average SFR with redshift. In fact, the measured metallicity in several additional samples of high-z galaxies results to be in agreement with the predictions of the FMR given the mass and SFR of the galaxies: galaxies having lower SFRs than the general population at their redshifts also have higher metallcities (e.g., Richard et al. 2011, Nakajina et al. 2011), and galaxies with higher SFRs also have lower metallicities (e.g., Erb et al. 2010, Contini et al. 2011, Sanders et al. 2011), so that all these galaxies follow the FMR. Also, the FMR allowed Mannucci et al. 2011 and Campisi et al. 2011 to show that the hosts of the long- GRBs have the same metallicity properties of the other star-forming galaxies. However, they found some metallicity evolution of the FMR at $z\\sim3.3$, where galaxies tend to have lower metallicities.\\\\    All the observational studies mentioned in the previous  paragraph refer to the gas-phase metallicity, as measured by emission lines. Gallazzi et al. (2005) presented the local mass-stellar metallicity relation based on $\\sim170000$ SDSS galaxies (Sloan Digital Sky Survey Data Release Two). The stellar metallicities were derived using the Lick system of spectral indices in the optical region which are sensitive to the overall metallicity of the stellar population, primarily dominated by intermediate/old stars. They adopted a Bayesian statistical approach and  derived the stellar metallicities by comparing the observed spectrum of each galaxy with a comprehensive library of model spectra corresponding to different star formation histories. They found that at  low  masses, the stellar metallicity increase with  mass, while above $\\sim3\\times10M_{\\odot}$ the relation flattens out. In addiction they noted that gas-phase metallicity is best determined for star-forming galaxies, whereas stellar metallicity is best determined for early-type galaxies, and found that the stellar metallicity is generally lower than the gas-phase metallicity (by \u223c0.5 dex). More recently Panter et al. (2008) inferred the stellar metallicity history of SDSS galaxies and determined their stellar mass-metallicity relation. They used a different approach respect to Gallazzi et al. (2005), but they found similar results. Moreover,  considering only the younger  population of galaxies ($\\leq$1 Gyr) they found good agreement also with the gas phase metallicities. \\\\  A very  limited work has instead been done on stellar metallicity at high redshift, see Shapley (2011) for a recent review. As in local {\\it starbursts}, the strongest features in the rest-frame UV spectrum of distant galaxies are interstellar and photospheric absorption lines of C, N, O, Si, and Fe, produced by  hot, young O-B stars  (see Shapley et al. 2003). One advantage in using these spectra to measure the metallicity at $z\\sim3$ is that the UV rest frame is shifted into the optical spectral region, which is easier to observe from the ground-based telescope. However, very high signal to noise on the stellar continuum is required to study the relevant absorption features for metallicity measurements, and therefore such studies were obtained, until now, mainly for gravitationally lensed galaxies(Rix et al. 2004, Quider et al. 2009, Dessauges-Zavadsky et al. 2010) or for co-added star forming galaxy spectra (Halliday et al. 2008).\\\\  Several authors have presented calibrations for stellar metallicity based on UV absorption lines. Leitherer et al. (2001) investigated the influence of metallicity on the spectra of star forming galaxies. In particular, they investigated the existence of some blended photospheric lines whose strengths depend on metallicities only. They found that the two blends of lines near $\\lambda$1370 and $\\lambda$1425 (which they attributed to OV $\\lambda$1371 and FeV $\\lambda$1360-$\\lambda$1380 and to $\\lambda$SiIII 1417, CIII $\\lambda$1427  and FeV $\\lambda$1430 respectively) have equivalent widths that increase steadily with metallicity and do not depend on other stellar parameters, as age and IMF. Rix et al. (2004) using the {\\it Starburst99} plus their non-LTE model atmosphere code WM-basic, supported the conclusions that these lines are useful metallicity indicators, and suggested a new indicator at $\\lambda$ 1978. Rix et al. (2004) applied the new indicators to measure the stellar metallicity  of two lensed galaxies, MS 1512-cB58 at z=2.73, and Q1307-BM1163 at z=1.411, finding good agreement with the gas phase metallicity.\\\\   On the basis of these works, Mehelert et al. (2006) calculated the EWs of the $\\lambda$1370  $\\lambda$1425  for 12 galaxies with $2.37<z<3.40$ in the FORS deep field. They investigated the evolution of the EWs and the metallicity with the redshift, and found that the abundances of heavy elements is increasing between $z\\sim0.5$ and   $z\\sim2.5$. Halliday et al. (2008), instead,  computed the stellar metallicity using the  $\\lambda$1978 indicator for a co-added unlensed star forming galaxies spectra at $z\\sim2$. Comparing their results with the gas phase metallicity of galaxies with similar mass, they found that their stellar metallicities were lower by a factor of $\\sim0.25dex$. Quider et al. (2009) computed the stellar metallicity for the gravitationally lensed galaxies Cosmic Horseshoe at $z=2.38$ by using the $\\lambda$1425  metallicity indicator, found good agreement between stellar and gas phase metallicity. More recently, Quider et al. (2010) carried on the same work for the lensed galaxies Cosmic Eye at $z=3.07$. Because of the presence a strong sky lines in the region of the stellar metallicity indicators, they do not give the exact value of the stellar metallicity, but just an indication by comparing the observed spectra with the model. Finally, Dessauges-Zavadsky et al. (2010) studied the stellar metallicty for the lensed galaxy 8 o'clock arc at  $z=2.73$. Using the  $\\lambda$1425 and  $\\lambda$1978 photospheric indices, they found comparable metallicity for the gas and the stellar component.\\\\    All these studies provide important constraints to the models of galaxy formation. In fact, many models study  the processes of gas infall and outflows, and  the properties of feedback and galactic winds (Dekel \\& Woo 2003, Veilleux, Cecil \\& Bland-Hawthorn 2005, Murray, Quataert \\& Thompson 2005, Dav\\'e, Finlator \\& Oppenheimer 2007, de Rossi, Tissera \\& Scannapieco 2007, Brooks et al. 2007, Tornatore et al. 2007, Somerville et al. 2008, Oppenheimer \\& Dav\\'e 2008, Finlator \\& Dav\\'e 2008, Mouchine et al. 2008, Dayal et al. 2009, Tassis, Kravtsov \\& Gnedin 2008, Kobayashi et al 2007, Calura et al. 2009, Spitoni et al. 2010, Salvaterra et al. 2010, Sakstein et al. 2011). The no evolution of the FMR, the associated evolution of the mass-metallicity relation, and the properties of the effective yields put strong constraint on them. In this picture, the stellar metallicity can be used as independent measurement to test the predicted metallicity evolution (De Rossi et al. 2007, Dav\\'e \\& Oppenheimer 2007).\\\\  In this work we study the stellar metallicity of a sample of  galaxies at $z>3$.  In the next section we review the different indices used for stellar metallicity estimates at high-z, and define new rest frame UV feature suitable for such measurements  to enlarge the number of available metallicity indicators. In Section 3 we present our sample of high redshift galaxies and we compute their stellar metallicity applying the calibrations found. In Section 4 we calculate the stellar and gas phase metallicity for some lensed galaxies found in literature, and we compare the two. Finally we present the first stellar mass-metallicity relation obtained at $z>3$, followed by the conclusions.     \\begin{figure*} \\centering \\includegraphics[width=12cm,angle=270]{Fig1.ps} \\caption{ Variation of the line indices with stellar population age and  metallicity. In each plot, the EW is shown in left vertical axis, while the right axis shows the average fractional depth below the continuum. For each feature we draw the models constructed with   {\\it Starburts99} at five value of metallicity as shown in different colors. For each metallicity, the color regions represent the error due to the dependence on the IMF assumed. In this way we can highlight the dependence of the models and robustness of the indices from the IMF and stellar population age assumed. } \\label{linvar} \\end{figure*} ",
        "conclusions": "In this paper, we have investigated for the first time the stellar mass-metallicity relation at high redshift, $z\\sim3$.\\\\ Using the theoretical spectra created with the population synthesis code  {\\it Starburts99}, we looked for photospheric absorption lines to be used as indicators of stellar metallicity.\\\\ First we tested the line indices proposed by Leitherer et al. (2001) and Rix et al. (2004), the  F1370, F1425 and F1978 using the last version of {\\it Starburts99}, although the F1978 shows a strong dependence from the resolution and the IMF.\\\\ Then we defined two new photospheric lines, F1460 and F1501, and we found that these lines are sensitive to the metallicity and almost independent of the age and the IMF, and therefore useful stellar metallicity indicators. The F1501 index seems to be the most promising because it is defined on the narrowest wavelength range and less affected by the uncertainties on the continuum definition.\\\\ We provided the metallicity calibrations, see Fig.~\\ref{cal1}, with two different definitions of the continuum: the first relations are referred to the real continuum, that we suggest to use in case of spectra with low signal-to-noise ratio, the others were  obtained using the definition of the ``pseudo-continuum'' provided by Rix et al. (2004), that we recommend in case of high signal-to-noise spectra, see Sec.~\\ref{mc}.\\\\ We applied the relations on one galaxy and a composite spectra comprised of three additional galaxies of the AMAZE sample at $z\\sim3.3$, for which the gas phase metallicity and the galaxy masses were already know.\\\\ We took from  the literature the spectra of eight additional galaxies, and we recompute their stellar metallicity using the new calibrations.\\\\ At the end we compared the results found with the gas phase metallicity for each object, see Fig.~\\ref{dif}. The main conclusion  of this work is that within the  errors, the stellar and the gas phase metallicity are consistent, although there seems to be a tendency to find stellar metallicity lower than the gas phase one by  $\\sim0.1dex$, as already found by Halliday et al. (2008). This result supports the low metal content derived for the gas phase of high-z galaxies from optical strong line ratios, as well as an evolution of the Fundamental Metallicity Relation at $z>3$ as found by Mannucci et al. (2010). \\\\ For the first time, we obtained the stellar mass-metallicity at redshift $z>2.5$,  see Fig.~\\ref{MZ}. We notice that the stellar metallicities found at high redshift is comparable with  those found by Panter et al. (2008)  for local galaxies, although the two are not straightforwardly comparable as in high redshift galaxies the stellar metallicity are computed for hot, young stars, while in the local galaxies for cold, older stellar population.\\\\  In summary, the rest-frame UV is rich in metallicity dependent features, which are able to provide a measure of stellar metallicity in high redshift galaxies. This represent an independent measure of the chemical abundances in galaxies with respect to the more widespread gas phase metallicities, which can provide important constraints to the star formation histories of galaxies in the early Universe. Although this technique is currently limited to very bright or lensed galaxies by the high S/N required, the advent of next generation of telescopes will give us much higher quality spectra for high redshift galaxies, and the stellar metallicity indicators will play a more important role in chemical abundances studies at high redshift."
    },
    "1112/1112.2479.txt": {
        "abstract": "{We use cosmological hydrodynamical simulations to show that a significant  fraction of the gas  in high redshift rare massive halos falls nearly radially to their very centre on extremely short timescales.  This process results in the formation of very compact bulges with  specific angular momentum a  factor $5-30$ smaller than the average  angular momentum of the baryons  in the whole halo. Such low angular momentum originates both from segregation and effective cancellation when the gas flows to the centre of the halo  along  well defined cold filamentary streams. These filaments penetrate deep inside the halo and connect to the bulge from  multiple rapidly changing  directions.  Structures falling in  along  the filaments (satellite galaxies) or formed by gravitational instabilities triggered by the inflow (star clusters) further reduce the  angular momentum of the gas in the bulge. Finally, the fraction of gas radially falling to the centre  appears to increase  with the mass of the halo; we argue that this is most likely due to an enhanced cancellation of angular momentum in rarer halos which are  fed by more isotropically distributed cold streams. Such an increasingly efficient funneling of low-angular momentum gas to the centre of very massive halos at high redshift may account for the rapid pace at which the most massive super massive black holes grow to reach observed masses around $10^9$M$_\\odot$ at an epoch when the Universe is barely 1 Gyr old.} ",
        "introduction": "Supermassive black holes (BH) have been established to be ubiquitous  at the centre of local galactic bulges and their mass has been shown to correlate well with the stellar mass and perhaps even more strongly  with the stellar velocity dispersion of these bulges \\citep{magorrianetal98,  tremaineetal02, haring&rix04}. As the accretion of gas onto BHs can drive strong feedback from Active Galactic Nuclei (AGN) through spherical winds or collimated jets, and can potentially self-regulate the growth of the BH along with the cold baryon content of galaxies \\citep{silk&rees98, haehneltetal98, king03} AGN feedback is a natural candidate to explain the observed correlations. Indeed, this picture has been further substantiated by several numerical implementations of such feedback using either semi-analytical models \\citep{boweretal06, crotonetal06, somervilleetal08} or hydrodynamical simulations \\citep{dimatteoetal05, booth&schaye09, duboisetal12}.  The issue is further complicated by the fact that the most supermassive BHs (several $10^9\\, \\rm M_\\odot$) seem to be already in place  at  $z\\approx 6-7$,  i.e. less than a Gyr  after the Big-Bang \\citep{willottetal03, fanetal06,jiangetal09,mortlocketal11}. Growth to such large masses in such a short time scale constitutes  a significant challenge for any model of BH evolution, especially those with strong AGN feedback, as it requires sustained feeding at close to the Eddington  accretion rate \\citep{haiman04, sijackietal09}. However, considering that a substantial fraction  of the sky had to be surveyed to discover the luminous QSOs powered by these very massive high-redshift BHs ($10^{-9} \\, \\rm Mpc^{-3}$ comoving number density of high-redshift quasars according to~\\citealp{fanetal06}), it appears likely that the halos hosting these BHs  are the most  massive formed  at these early times and thus are very rare objects. Obviously, the main caveat of this 'rareness' argument is that since these young BHs are thought to be growing fast because of the presence of large amounts of gas in their surroundings, a larger number of their quasar counterparts could possibly be obscured in optical wavebands \\citep{alexanderetal03, treisteretal11, willott11} and only be visible  in X-rays~\\citep[e.g.][]{daddietal07} and the Infrared. This selection effect could increase the {\\it true} number of very bright quasars/super massive BHs at high redshift considerably.  Nevertheless, even when adopting the view that these objects are rare, one still needs to propose an effective mechanism to funnel low angular momentum gas all the way down to the sphere of influence of the BHs at the very centre of  their very massive halos hosts. Hydrodynamical simulations suggest that high-redshift quasars are accreting gas at rates close to their Eddington (or super-Eddington) limit because large amounts of cold gas are indeed trapped in the centre of galaxies, and that the fraction of Eddington-limited growth of BHs diminishes with time as this gas reservoir is depleted~\\citep{sijackietal07, dimatteoetal08, duboisetal12}.   \\cite{dimatteoetal12} have emphasized  that filamentary infall of cold gas is the major mode of gas  supply for the continuous feeding of BHs similar to what is discussed for the build-up of galaxies at intermediate redshift \\citep{keresetal05, agertzetal09}. Note that this is  a very different mode of gas supply from that due to the often invoked more episodic feeding of BHs  due to large amounts of gas driven to  galactic nuclei by  (major) mergers \\citep{mihos&hernquist96, kauffmann&haehnelt00, hopkinsetal06, mayeretal10}  \\cite{pichonetal11} and  \\cite{kimmetal11}  have recently used hydro cosmological Adaptive Mesh Refinement (AMR)  simulations  to investigate  the angular momentum properties and the infall of gas in Milky Way class halos at low and intermediate redshifts with a view to study the formation and evolution of their central galaxy discs. With a similar technique~\\cite{bournaudetal11} have performed simulations of isolated galaxy discs with parsec resolution to  show that Toomre instabilities drive strong inflows of gas to the centre of these discs.  We build here on this work, by investigating the more massive and rarer halos expected  to host the most massive supermassive BHs at  high redshift. We therefore concentrate our exploratory study on the gas with the lowest angular momentum falling to the centre of these rare, massive  halos.  %--------------------------------------------------------------------------------------- \\begin{figure*} \\includegraphics[width=\\columnwidth]{fig/nice_SH_z6.eps} \\includegraphics[width=\\columnwidth]{fig/nice_LH_z6_zoom64.eps} \\includegraphics[width=\\columnwidth]{fig/nice_SH_z6_20pc_zoom1024.eps} \\includegraphics[width=\\columnwidth]{fig/nice_stars_ugr_zoom4096.eps} \\caption{{\\sl Top panels:} Polychromatic view of the two halos in the SH and the LH simulation at $z=6$. Gas density is color coded in green, gas temperature in red, and tracer particles in blue. The circles in the two top panels corresponds to  the virial radius $r_{\\rm vir}$. The {\\sl bottom left panel} displays the inner region of the top left panel simulated with higher resolution (SHhr), showing the transient  disc and its connected filaments as discussed in the text. The  circle corresponds to 0.1 $r_{\\rm vir}$ (see table~\\ref{tab:feature} for numerical values). The {\\sl bottom right panel } displays the stellar emission in the inner region of the bottom left panel  as it would be observed  in  ugr filter bands. Note the compact bulge (the saturated spherical region at the centre) and the significant number of satellites  flowing in along the cold streams  on their way to merge with the central galaxy. Finally, we point out that even though in projection the hot gas phase seems to dominate, this halo is still mainly accreting cold gas. } \\label{fig:visual} \\end{figure*} %---------------------------------------------------------------------------------------  Whilst \\cite{dimatteoetal12} have recently investigated  this problem using a simulation with a large enough volume ($533\\, h^{-1}\\cdot\\rm Mpc$) to capture halos sufficiently  massive to plausibly host these $\\sim 10^9\\, \\rm M_\\odot$ BHs, we take a different approach here and  focus our study on two individual very massive halos at high redshift which we follow using a high-resolution cosmological hydrodynamics re-simulation technique. The main aim here is to better resolve the dynamics and angular momentum history of the inflowing cold gas as it falls to the centre of halos, forms a galactic bulge and potentially feeds the growth of the central supermassive BH.  %-------------------------------------------------------------------------------------- \\begin{figure*} %   \\includegraphics[width=\\columnwidth]{fig/vrot_vcirc_vdisp.ps} \\includegraphics[width=\\columnwidth]{fig/vrot_vcirc3.ps} %   \\includegraphics[width=\\columnwidth]{fig/vrot_vcirc4.ps} \\includegraphics[width=\\columnwidth]{fig/vrot_vcirc_vdisp_LH.ps} \\caption{Circular velocities (black lines and inset), radial velocity dispersion (blue lines) and rotational velocities (red lines) of the DM (dotted lines), the stars (dashed lines), and the gas (dot-dashed lines) for the SHhr simulation (left panel), and the LH simulation (right panel) at $z=6$. The blue line corresponds to the SH run. The bulge radius for the SHhr simulation is shown as a vertical arrow. At this  redshift, these galaxies have the rotation curve of  very compact ellipticals. } \\label{fig:vrot_vcirc} \\end{figure*} %--------------------------------------------------------------------------------------  The paper is organized as follows. Section~\\ref{section:numerics} describes the numerical setup of our simulations. Section~\\ref{sec:mass} investigates the angular momentum evolution and the typical trajectories of the baryonic material forming the bulge. Section~\\ref{sec:discussion} discusses the implications for BH build up at high redshift while Section~\\ref{sec:conclusion}  gives our conclusions.  %=============================================== % Simulation setup %=============================================== ",
        "conclusions": "\\label{sec:conclusion}  We have investigated here the formation of a compact central bulge in two rare very massive halos at redshift $z=6$.  Our main results are as follows.  \\begin{itemize}  \\item{The typical path for a parcel of fluid entering a halo and eventually ending up in the bulge, is to stream out of a void into the surrounding wall, from there into the filaments defining the intersection of the cosmic walls/sheets, and then along the filaments which penetrate deep into the halo.}  \\item{ Early on, most  of the  baryons ending up in the bulge stream nearly radially along the cold filamentary infall directly into a very compact bulge. As time progresses an increasing fraction of the baryons settles first into a compact gravitationally unstable transient central disc from where they  progress rapidly into the bulge. By redshift $z=6$ about half the baryons forming the bulge move first through this surrounding disc before reaching the bulge.}  \\item{In the outer part of the halo  the angular distribution of the baryonic inflowing matter is stable and reflects the angular distribution of the cosmic web, while well inside the virial radius the angular distribution becomes much more random and changes rapidly.}  \\item{A small fraction ($\\sim$ 10 \\%) of the baryons in the halos falls directly from the structures of the cosmic web into the centre of the halos in the form of proto-galactic clumps; these clumps are carried along the filamentary streams, that represent more than half of the total accretion of gas, and join the compact central bulge as minor mergers. }  \\item{In our simulations most of the baryons in the bulge turn into stars with a remaining gas fraction of 5-30\\%. The gas fraction  is, however,  strongly  dependent on the resolution  and the star formation criterion  of the simulation. The gas fraction should also be strongly affected by feedback both from  the stars but even more by feedback from the accreting central supermassive BH once the BH  has become sufficiently massive.  Note again that the simulations used here  do  not include any feedback.}  \\item{The baryon dominated bulge has angular momentum  a factor 10-30 lower than the average specific angular momentum of the  baryons in the two simulated halos. The baryonic material which is mainly in the form of stars is  extremely strongly bound and compact with a peak of the circular velocity approaching 1000 km/s at a radius $< 100$pc. }  \\item{The low value of the angular momentum is due to three effects. First the bulge is preferentially made from material which already has a somewhat lower specific angular momentum than average as it enters the halo along the  filaments of the cosmic web. Second the baryonic matter entering the halo from the cosmic environment does so with  a broad angular momentum distribution. As the gas streams in cold the lower angular momentum material falls quickly to the centre of the halo due to the lack of pressure support where the randomly oriented angular momentum  effectively cancels. Third there is additional angular momentum loss due to dynamical friction on inspiralling proto-galactic clumps, exchange between spin and orbital angular momentum and gravitational instabilities in the disc.}  \\item{The fraction of gas  streaming  nearly radially to the centre of the halo making up the low-angular momentum bulge appears to strongly increase for rare halos. We argue that  this is due to the fact that isolated rare halos are connected to the cosmic web more isotropically and with a larger  number of filaments.  This leads both to lower specific angular momentum when the baryonic matter enters the halo and to more complete angular momentum cancellation when the low angular momentum material streams cold  into the  compact bulge. }  \\end{itemize}  Numerical simulations of galaxy formation are still subject to a  range of uncertainties both numerical and in terms of the implemented physics in particularly with respect to feedback due to stars and AGN. The filamentary cold inflow of gas appears nevertheless  to emerge as the key ingredient in the formation of the observed compact discs and bulges at high redshift. As we have discussed here it may also be key to  the efficient  build up of the several billion solar mass black holes inferred to exist as early  as $z\\sim 6-7$.  %=============================================== % Acknowledgments %==============================================="
    },
    "1112/1112.0360_arXiv.txt": {
        "abstract": "The k=0 Friedmann Lemaitre Robertson Walker model with a positive cosmological constant and a massless scalar field is analyzed in detail. If one uses the scalar field as relational time, new features arise already in the Hamiltonian framework of classical general relativity: In a \\emph{finite} interval of relational time, the universe expands out to infinite proper time and zero matter density. In the deparameterized quantum theory, the true Hamiltonian now fails to be essentially self-adjoint both in the Wheeler DeWitt (\\WDW) approach and in LQC. Irrespective of the choice of the self-adjoint extension, the big bang singularity persists in the \\WDW theory while it is resolved and replaced by a big bounce in loop quantum cosmology (LQC). Furthermore, the quantum evolution is surprisingly insensitive to the choice of the self-adjoint extension. This may be a special case of an yet to be discovered general property of a certain class of symmetric operators that fail to be essentially self-adjoint. ",
        "introduction": "\\label{s1}  Loop quantum cosmology (LQC) of the k=0, $\\Lambda=0$ Friedmann Lemaitre Robertson Walker (FLRW) model with a massless scalar field was discussed in detail in \\cite{aps3}. The scalar field serves as a viable internal time variable both in the classical and the quantum theory, with respect to which relational observables such as the matter density and curvature evolve \\cite{aps1,aps2}. This makes it possible to explicitly construct the physical Hilbert space and introduce relational Dirac observables to unravel physics of the Planck regime in a large number of cosmological models \\cite{as}, and a scheme has been sketched even for full general relativity \\cite{warsaw-full}. Using this setup it was rigorously established that, while the big bang singularity persists in the \\WDW theory of the k=0, $\\Lambda=0$ model, it is resolved due to the quantum geometry effects of loop quantum gravity (LQG) \\cite{aps3}.% \\footnote{Recently, this result has been conceptually sharpened using the consistent histories framework in which one can calculate probabilities for the occurrence of certain histories without recourse to external measurement devices or interaction with environment. Using appropriate coarse grained histories which completely decohere, it was shown that the probability of encountering a singularity in the distant past or future is 1 in the \\WDW theory and 0 in LQC for any state (which is in the domain of operators used to construct coarse-grained histories) \\cite{consistent2,consistent3,as}.}  An appendix in \\cite{aps3} also outlined how the cosmological constant $\\Lambda$ with either sign can be incorporated. A subsequent, detailed discussion of the $\\Lambda<0$ case appeared in \\cite{bp}. It firmly established that, as in the $\\Lambda=0$ case, in LQC the big bang singularity is replaced by a quantum bounce which occurs when the total energy density $\\rho_{\\rm tot}$ reaches its maximum value $\\rcr$. Furthermore the numerical value of $\\rcr$ is the same as in the $\\Lambda=0$ case, $\\rcr \\approx 0.41 \\rho_{\\rm Pl}$, although now $\\rho_{\\rm tot}$ includes a contribution from the cosmological constant in addition to the matter density $\\rho$;\\, $\\rho_{\\rm tot} = \\rho + \\Lambda/8\\pi G$. It turns out that, by a suitable choice of time variable (or lapse function), the $\\Lambda=0$ model can be solved exactly \\cite{acs}. This is not the case for $\\Lambda\\not=0$. Therefore, results of \\cite{bp} for the $\\Lambda <0$ case are conceptually important also because they demonstrate that the LQC bounce and the qualitative features of the resulting Planck scale physics are not tied to exact solvability. Finally, although the situation with the bounce is the same, the presence of the cosmological constant does alter the underlying mathematical structure in non-trivial ways. In particular, in the deparameterized picture, while the spectrum of the true Hamiltonian is continuous in the $\\Lambda=0$ case, it is purely discrete in the $\\Lambda <0$ case.  The goal of this paper is to present an analogous, detailed account of the $\\Lambda >0$ case. Even though we will again consider a massless scalar field, rather surprisingly, the flip of the sign of the cosmological constant changes the underlying mathematical and conceptual structure significantly. Let us begin with the classical theory. If one again uses the scalar field $\\phi$ for internal time, in contrast to the $\\Lambda =0$ and $\\Lambda <0$ cases \\cite{aps3,bp}, the Hamiltonian vector field on the phase space is now incomplete. As a result, volume of any compact co-moving region becomes infinite and the matter density vanishes at a finite instant $\\phi_o$ of internal time $\\phi$. This situation is qualitatively similar to that in the case of a non-relativistic particle in a steep negative potential whose dynamical trajectories reach infinity in a finite time. In such situations, typically, the Hamiltonian operator in Schr\\\"odinger quantum mechanics is symmetric but not essentially self-adjoint. Each self-adjoint extension then yields a unitary evolution but evolutions obtained from distinct operators are both mathematically and physically inequivalent. In the present case, one again finds that the true Hamiltonian operator generating evolution in the scalar field time is symmetric but not essentially self-adjoint. However, rather surprisingly, this ambiguity has negligible effect on states of physical interest: those that start out being peaked at a classical solution in a low curvature region. In particular, all these states undergo a quantum bounce and the total density $\\rho_{\\rm tot}$ at the bounce is again universal. Furthermore, while the evolution of expectation values of physical observables does depend on the choice of self-adjoint extension, the dependence is extremely weak. This robustness may be related to the fact that, on the classical phase space, one can extend both the evolution equations and the solutions simply by analytical continuation, without having to introduce specific boundary conditions at infinity. Our analysis raises the possibility that there may well be a general pattern and new results could be found on properties of certain sub-classes of operators that fail to be essentially self-adjoint.  The paper is organized as follows. In Sec. \\ref{s2} we discuss the Hamiltonian framework for the k=0, $\\Lambda>0$ model. Sec. \\ref{s3} is devoted to the \\WDW quantum theory and Sec. \\ref{s4} to LQC. We conclude in \\ref{s5} with a brief summary and discussion. Because the numerical simulations in this paper were completed soon after the initial analysis in \\cite{aps3}, they use an older value of LQC area gap which turned out to be half the value that is relevant for states used in LQC \\cite{awe1}. In the main text we use this more recent value. Therefore, unfortunately, in Sec. \\ref{s4} there is an occasional mismatch of factors of two between the text and the figures. ",
        "conclusions": "\\label{s5}  In this paper, we analyzed in detail the \\WDW theory and LQC of the k=0, $\\Lambda>0$ FLRW model along the lines of the treatment of the $\\Lambda <0$ case of \\cite{bp}, thereby completing the program outlined in an Appendix of \\cite{aps3}. As in the $\\Lambda \\le 0$ cases, the scalar field can be used as a global clock, providing us with a natural notion of relational time both in the classical and quantum theories. However, interestingly, there is a key difference in the physically most interesting case, that of $\\Lambda >0$: In classical general relativity, solutions that start with infinite matter density $\\rho$ at the big bang at time $\\phi = -\\infty$ expand out and now achieve $\\rho=0$ at some \\emph{finite} value $\\phi_o$ of internal time (when the volume $v$ of the fiducial cell $\\mathcal{C}$ becomes infinite). Thus, in the $\\rho$-$\\phi$ plane each of these dynamical trajectories starts out at $\\phi = -\\infty$ but ends at $\\phi=\\phi_o$. But it can be \\emph{analytically} extended beyond $\\phi=\\phi_o$ and the extension represents a universe which starts out with zero matter density at $\\phi=\\phi_o$ but contracts, ending in a big crunch singularity at $\\phi=\\infty$. >From the relational time perspective, then, one is led to regard the two branches as providing a single dynamical trajectory because it is artificial to simply end dynamics at a finite value of time. Now, in non-relativistic mechanics if the potential is negative and steep the particle may roll off to infinity in a finite amount of time. In that case, one has to choose from a one (or more) parameter family of boundary conditions at infinity to continue dynamics beyond that time. In the present case, by contrast, we did not have to resort to making a choice because the most interesting Dirac observable, $\\rho|_\\phi$, is analytic in $\\phi$.  Nonetheless, the fact that the universe expands out to $\\rho=0$ at a finite value of the relational time introduces ambiguities in the \\emph{quantum} evolution: The operator $\\Theta_\\Lambda$ which generates dynamics with respect to $\\phi$ now fails to be essentially self-adjoint. Before discussing this point in detail, let us first note two aspects of this phenomenon. First, it is not a peculiarity of LQC; it occurs also in the \\WDW theory. Second, in both cases, the lack of essential self-adjointness is related directly with the behavior of the system at large $v$ and low matter density $\\rho$; its origin does \\emph{not} lie in the Planck scale physics. In both quantum theories, the dynamical operator admits a one parameter family of self-adjoint extensions. For a general system, different choices of extensions can give rise to very different dynamics. However, in this model the results are surprisingly robust with respect to this choice. Not only is the qualitative behavior of dynamics the same, but the differences in the dynamics of the expectation values of the most interesting Dirac observables in theories resulting from two different extensions are smaller than their dispersions in any one theory, even for general states. Furthermore, numerical simulations show that, irrespective of the choice of extension, quantum states which are semi-classical in the low curvature (or low total density $\\rho_{\\rm tot}$) regime remain sharply peaked at the \\emph{extended} classical trajectory in the low curvature regime both in the \\WDW theory and LQC. It is tempting to conjecture that this robustness of quantum dynamics is related to the fact that we did not have to choose a boundary condition at $\\phi=\\phi_o$ to extend the classical $\\rho-\\phi$ trajectory. In the remainder of this section, most of our discussion on the behavior of wave functions will refer to these states.  As in the $\\Lambda \\le 0$ cases, there is a \\emph{pronounced} difference between the quantum dynamics of the two theories in the Planck regime. In the \\WDW theory, the wave function simply follows the extended classical trajectory into the big-bang and the big-crunch singularities. In LQC, by contrast, while these states remain peaked at the classical trajectory so long as the curvature (or $\\rho_{\\rm tot}$) is low compared to the Planck scale, there is a dramatic departure in the Planck regime. There is again a new repulsive force with origin in the quantum geometry that overwhelms classical gravity and cases a quantum bounce. Again, the numerical simulations show that, although the force is so strong in the Planck regime, it dies very quickly and becomes negligible once $\\rho_{\\rm tot}$ falls below  $10^{-2}- 10^{-3}$ Planck density. In LQC then, even though we are in the k=0 case, we are led to a scenario that is approximately cyclic. As in the k=1 LQC models, the quantum evolution spans an infinite number of epochs. In each epoch the universe begins with a quantum bounce where $\\rho_{\\rm tot} \\approx 0.41 \\rho_{\\rm Pl}$, expands out till $\\rho_{\\rm tot} = \\Lambda/8\\pi G$ and then undergoes a collapse till it reaches another quantum bounce. For states under consideration, dynamics is nearly periodic.  How does this dynamics appear in the space-time picture? Let us begin with the classical theory and consider a solution in which the universe starts out with a big-bang at $\\phi=-\\infty$. It expands out to $\\scri^+$ ---which is space-like for $\\Lambda>0$--- where the matter density vanishes and $\\rho_{\\rm tot} = \\Lambda/8\\pi G$. In terms of the physical metric, this space-time is future complete. However, the extended phase space trajectory analytically continues the space-time geometry across $\\scri^+$, effectively gluing it with $\\scri^-$ of a contracting solution.% \\footnote{The detailed gluing procedure will involve a conformal completion along the lines of \\cite{ar} where the normal component to $\\scri$ of the metric is rescaled by a different power of the conformal factor than the tangential one.} Quantum states under consideration remain peaked at these extended space-time geometries across $\\scri$. An extension is but to be expected both in the \\WDW theory and in LQC: quantum evolution in the internal time $\\phi$ is \\emph{unitary} and $\\phi$ achieves a finite value $\\phi_o$ at $\\scri$ of the given classical solution, unitary evolution could not just stop there. What is interesting is that, \\emph{irrespective of the choice of the self-adjoint extension}, the state remains sharply peaked on the analytically extended geometry.  This extension, and the ensuing nearly cyclic scenario has some similarities with Penrose's recent proposal of a cyclic conformal cosmology \\cite{rp}. However, there are also key differences. In our case, $\\scri^+$ of the expanding branch is glued to the $\\scri^-$ of the contracting branch; not to the big-bang singularity of the next `aeon'. More importantly, quantum geometry effects are crucial in LQC. In particular $\\hbar$ appears in the denominator of the expression of the maximum density $\\rcr$ whence, as one would expect, $\\rcr$ would diverge in the classical limit $\\hbar \\to 0$. Therefore, quantum effects and a non-zero $\\hbar$ play an essential role in the approximately cyclic scenario of LQC. By contrast, a central feature of the cyclic conformal cosmology paradigm is that, although one does have unboundedly large curvatures, $\\hbar$ plays no role at all in this regime.  Use of the scalar field as a relational time variable played a key role throughout our analysis, both in the classical and quantum theory. What would have happened if we had made some other choice? In a recent analysis \\cite{hp-qg,hp-qc} non-rotating dust has been used in place of the scalar field. In this case, the expression (\\ref{eq:constr-quant}) of the gravitational part $\\Theta_\\Lambda$ of the Hamiltonian constraint is modified because the lapse is now tailored to proper time. In particular, the coefficient of $\\Lambda$ is now \\emph{linear} rather than \\emph{quadratic} in $v$. Consequently, the analog of $\\Theta_\\Lambda$ is now essentially self-adjoint and the LQC evolution resembles that in the $\\Lambda=0$ case \\cite{aps3}: The universe starts out with infinite volume in the distant past, collapses, undergoes a quantum bounce and then expands out to infinite volume. However, in the Planck regime, quantum matter should be described using quantum field theory and for all standard quantum fields the kinetic term in the Hamiltonian is quadratic in momenta. Therefore the specific feature that simplifies the mathematics in the case of dust is no longer available and the overall situation is then the same as that in the case of the scalar field.  But what if we return to using scalar field as matter source but let a geometric variable be the relational time? An obvious choice is volume. But in LQC, volume fails to be single valued making it difficult to introduce the `time-dependent' relational Dirac observables ---such as the matter density operator $\\hat\\rho|_v$--- especially near the bounce. On the other hand, the conjugate variable $b$ is single valued both on classical and effective trajectories and, as the form of the Hamiltonian constraint suggests, the function $y$ it determines is a possible candidate for relational time. However, it appears that the evolution would then be unambiguously unitary only for $|y| >y_o$. Furthermore, defining the physical state would require specification of the initial data at $y=y_o$ but the theory does not provide any selection principle for this task. These difficulties with $v$ and $b$ could well be surmountable with new ideas and more careful analysis. But as of now they seem to be more serious handicaps than the complications associated with the use of the scalar field as relational time we encountered in this paper.  Could one perhaps retain scalar field as the matter source but set lapse $N=1$ in the quantum constraint and use a `timeless framework' in LQC? Results of \\cite{klp1,klp2} imply that we would have been led to a theory that is mathematically free of ambiguities associated with self-adjoint extensions. However, as we now explain, this theory is difficult to interpret and it is unclear whether it is physically viable. We will conclude our discussion with a detailed elaboration of this point.  With lapse $N=1$, the quantum Hamiltonian constraint has the form \\, $\\hat{C}= B(v) \\otimes \\partial_\\phi^2 + \\hat{C}_{\\gr}\\otimes \\id$ where $C_{\\gr}$ is the gravitational part of the Hamiltonian constraint and $B(v) \\sim 1/v$ for large values of $v$ (see, e.g., \\cite{aps3}). This total constraint operator $\\hat{C}$ has been shown to be essentially self-adjoint \\cite{kp-posL} and for simplicity we will denote its self-adjoint extension also by $\\hat{C}$. One can therefore use it directly for group averaging and construct the physical Hilbert space $\\Hil^{\\phy}$ in the `timeless framework' without recourse to deparametrization. Then, although states can be represented as wave functions $\\Psi(v,\\phi)$, the physical scalar product is no longer given by an integral over $v$ at a fixed value of $\\phi$. What is the relation between this $\\Hil^{\\phy}$ and the Hilbert spaces $\\Hil^{\\phy}_{\\beta}$ associated with self-adjoint extensions $\\Theta_{\\Lambda,\\beta}$ we constructed in this paper? It turns out \\cite{klp2} that $\\Hil^{\\phy}$ is huge; it is given by the \\emph{direct integral} of all $\\Hil^{\\phy}_{\\beta}$;\\,\\, $\\Hil^{\\phy}= \\int_{\\oplus\\, I} \\Hil^{\\phy}_{\\beta} \\, \\rd\\beta$, where $I$ is the interval $(0, \\pi)$.  We will briefly discuss a simple example ---due to Wojciech Kami\\'nski in \\cite{k-pres}--- to illustrate the relation between these Hilbert spaces. Consider a 2-dimensional strip, $M= \\mathbb{R} \\times [0,1]$, with coordinates $\\phi \\in \\mathbb{R}$ and $x\\in [0,1]$ and a constraint thereon in the form of the Schr\\\"odinger equation $(-i\\partial_\\phi + i\\partial_x)\\Psi(x,\\phi) =0$ (with a first-order Hamiltonian). As in LQC, $\\phi$ plays the role of time while $x$ is to be thought of as the analog of the compactified volume coordinate $\\theta$. The operator $\\Theta_x := -i\\partial_x$ fails to be essentially self-adjoint on the closed interval $[0,1]$; it admits a 1-parameter family of self-adjoint extensions $\\Theta_{x,\\beta}$, labeled by $\\beta \\in [0, 2\\pi)$ (with domain $\\Dom_\\beta$ given by wave functions $\\psi(x)$ satisfying $\\psi(1) = e^{i\\beta}\\psi(0)$). For each extension, we can construct a physical Hilbert space $\\Hil^{\\phy}_\\beta$ in the standard manner from solutions to the quantum constraint: $\\Psi(x,\\phi) \\in \\Hil^{\\phy}_\\beta$ if and only if $-i\\partial_\\phi \\Psi(x,\\phi) = \\Theta_{x,\\beta}\\, \\Psi(x,\\phi)$. Next, let us define $\\Hil := \\int_{\\oplus\\, I} \\Hil^{\\phy}_{\\beta} \\, \\rd\\beta$, with $I= (0,2\\pi)$. ($\\Hil$ is analogous to $\\Hil^{\\phy}$ obtained by group averaging in \\cite{klp2}). Every $\\Psi_\\beta(x,\\phi) \\in \\Hil^{\\phy}_\\beta$ can be expanded as \\be \\Psi_\\beta(x,\\phi) = \\sum_{n=-\\infty}^{\\infty} \\tilde{\\psi}_{\\beta, n}\\,\\, e^{i(2\\pi n+\\beta) x}\\,\\, e^{ik\\phi} \\ee where $k:=2\\pi n+\\beta$. On the other hand $\\Psi(x,\\phi) \\in \\Hil$ has the form \\begin{equation}\\begin{split} \\Psi(x,\\phi) &= \\int_0^{2\\pi}\\!\\! \\rd\\beta\\, \\Psi_\\beta(x,\\phi) = \\sum_{n=-\\infty}^{\\infty} \\, \\int_0^{2\\pi}\\!\\rd \\beta\\, \\tilde{\\psi}_{\\beta,n}\\, e^{i(2\\pi n +\\beta) x}\\, e^{ik\\phi}  \\\\ &= \\int_{-\\infty}^{\\infty}\\!\\! \\rd k\\, \\tilde\\Psi(k)\\, e^{ikx + ik\\phi}\\, , \\end{split}\\end{equation} where we have set $k = 2\\pi n +\\beta$ as above. The norms in $\\Hil^{\\phy}$ are given by: \\be \\|\\Psi(x,\\phi)\\|^2 = \\int_{-\\infty}^{\\infty}\\!\\! \\rd k\\, |\\tilde{\\Psi}(k)|^2\\, . \\ee Note that $\\tilde{\\Psi}(k)$ has support on the entire $k$-axis and furthermore, the norm of $\\Psi(x,\\phi)$ also involves the integral of $|\\tilde{\\Psi}(k)|^2$ over the entire $k$ axis. Therefore, while elements $\\Psi_\\beta(x,\\phi)$ of any one $\\Hil^{\\phy}_\\beta$ are restricted to have support only on the physical configuration space $M = \\mathbb{R}\\times [0,1]$, $\\Psi(x,\\phi)$ in $\\Hil^{\\phy}$ are allowed to be non-zero all along the $x$-axis, even though points outside the $x$-interval $[0,1]$ have no physical interpretation in the model under consideration. In other words, although we restricted ourselves only to a `strip' $M= \\mathbb{R}\\times [0,1]$ of the Minkowski space in defining the classical system, in effect $\\Hil^{\\phy}$ describes a system on the entire 2-dimensional Minkowski space $\\mathbb{R}^2$. \\footnote{Note that, since the basis functions $e^{i(2\\pi n+\\beta)x}$ can be extended analytically to the entire $x$-axis, the quantum constraint can be uniquely extended to the full Minkowski space $\\mathbb{R}^2$.} To summarize, for the particle on the strip $M$, one would physically expect that the quantum theory should be formulated just on $M$ since values of $x$ outside $[0,1]$ have no physical meaning. This expectation is borne out if one works with any one self-adjoint extension $\\Theta_{x,\\beta}$ of $\\Theta_{x}$ but not if one works with the direct integral $\\Hil^{\\phy}$ of all the resulting Hilbert spaces.  In LQC, the situation is analogous. Working with a specific self-adjoint extension allows us to remain in the interval $[0, \\infty]$ of the $|v|$-axis, introduce Dirac observables and track their evolution in the internal time $\\phi$ and compare it with classical trajectories. On the other hand, working in the timeless framework in effect requires us to extend the $v$-axis beyond $v=\\infty$ and this extension is difficult to interpret physically.% \\footnote{Furthermore, unlike in the simple example discussed above, here one cannot extend the constraint operator uniquely even if one expresses the basis vectors as functions of the `compactified volume' coordinate $\\theta$ of \\eqref{eq:p-theta}, as these functions are not analytic in $\\theta=\\pi/2$.} For the same reason, while one can introduce Dirac observables also in the timeless framework, we cannot ask for their `evolution' and it is difficult to compare predictions of the quantum theory with those of the classical. In particular, while the choice $N=1$ of the lapse yields evolution in proper time in the classical theory, unfortunately this interpretation does not extend to the quantum theory in a simple way. Thus, while at first it seems mathematically natural to work with the lapse $N=1$ because the full constraint is then essentially self-adjoint, the resulting Hilbert space $\\Hil^{\\phy}$ appears to be simply too large to be physically viable in the above context.  Finally, our analysis brought out an unexpected robustness of the quantum evolution with respect to the choice of self-adjoint extensions $\\Theta_{\\Lambda,\\beta}$. As we noted above, this may be related to the fact that, in the physical sector associated with \\emph{any} self-adjoint extension, under unitary evolution quantum states of interest follow the natural and unambiguous analytic extension of classical trajectories. Is this perhaps a special case of as yet unknown general result? Is the lack of sensitivity of the quantum evolution on the choice of analytic extensions have its origin in some special features of the classical evolution?"
    },
    "1112/1112.2365_arXiv.txt": {
        "abstract": " ",
        "introduction": "Red supergiants are long--period variables with semiregular light variations on a timescale $\\gtrsim 10^2$ day. The period--luminosity relation (Glass 1979; Feast et al. 1980) and the linear theory of adiabatic oscillations (Stothers 1969, 1972) allow us to suppose that such a type of variability is due to radial stellar pulsations. At the same time together with semiregular variability some red supergiants exhibit superimposed irregular light variations (Kiss et al. 2006). The secondary stochastic variability is thought to be due to the large--scale convection in the outer subphotospheric layers (Stothers and Leung 1971; Schwarzschild 1975; Stothers 2010). Red supergiants are also remarkable due to intensive mass loss revealed through a large infrared excess indicating dust production in the stellar wind (Verhoelst et al. 2009).  The period--luminosity relation of radially pulsating red supergiants is used for determination of extra--galactic distances and in comparison with Cepheids the red supergiants allow us to substantially extend the distance scale due to their higher luminosities (Pierce et al. 2000; Jurcevic et al. 2000). Application of the theory of stellar pulsation to the analysis of observed variability of red supergiants allows us to verify some conclusions of the stellar evolution theory in a way similar to that employed earlier for Cepheids. It should be noted also that the growing bulk of recent observations indicate that the strong stellar wind of massive late--type supergiants is due to nonlinear stellar oscillations (van Loon et al. 2008)  The nature of radial oscillations in red supergiants is still not completely clear yet. The linear analysis of pulsational instability of red supergiants with masses $15M_\\odot\\le M\\le 30M_\\odot$ was performed by Li and Gong (1994) and Guo and Li (2002). According to their calculations radial oscillations of red supergiants are due to instability of the fundamental mode and, perhaps, the first overtone. However the theoretical period--luminosity relation agrees only with fundamental mode oscillations.  Nonlinear radial oscillations of red supergiants were considered only in two studies. In the first one (Heger et al. 1997) the authors investigated radial oscillations of the red supergiant with mass $M = 15M_\\odot$ at the final stage of the core helium burning. In the second work (Yoon and Cantiello 2010) the authors investigated pulsational instability of red supergiants with masses $15M_\\odot\\le M \\le 40M_\\odot$. It should be noted that in both these studies the self--exciting stellar oscillations were treated with modified methods of stellar evolution calculation and effects of interaction between pulsation motions and turbulent convection were not taken into account.  Below we present results of investigation of nonlinear pulsations of red supergiants obtained from the self--consistent solution of the equations of radiation hydrodynamics and turbulent convection. The need for such an approach is due to the significant length and mass of the outer convection zone involved in pulsation motions. The treatment of convective heat transport uses the solution of the diffusion--type equations for the enthalpy and the mean turbulent energy obtained by Kuhfu\\ss (1986) for spherically--symmetric gas flows from the Navier--Stokes equation. Thus, the results presented below deal with modelling the semiregular variability and the secondary stochastic variability is not considered because this problem is beyond the approximation of spherical geometry. We consider the stars with masses at the zero--age main sequence $8M_\\odot\\le\\mzams\\le 20M_\\odot$ and initial fractional mass abundunces of hydrogen and elements heavier than helium $X=0.7$ and $Z=0.02$. ",
        "conclusions": "Given in the previous section estimates of masses of seven galactic red supergiants allow us to conclude that the theory of stellar evolution is in an agreement with observational estimates of stellar radii. To compare more stars with the theoretical period--luminosity relation one should consider pulsational instability of red supergiants in the wider interval of initial masses $\\mzams$.  For more detailed theoretical period--luminosity relation one should consider the role of some parameters used in evolutionary computations. One of them is the overshooting parameter. In the present study the evolutionary computations were done for the ratio of the overshooting distance to the pressure scale height $l_\\mathrm{ov}/\\hp = 0.15$. The need to know the role of this parameter is due to the dependence of the mass--luminosity relation of helium burning stars on convective overshooting.  In stars with masses $M\\ge 20M_\\odot$ effects of mass loss during the red supergiant evolutionary stage become significant. In the present study the evolutionary calculations were done with mass loss rates by Nieuwenhuijzen and de Jager (1990) however determination of the mass loss rate $\\dot M$ as a function of fundamental stellar parameters remains disputable (Mauron and Josselin, 2011) Therefore, one should employ parametrization of the expression for $\\dot M$ and consider the mass--luminosity and period--luminosity relations as a function of this parameter.  Another parameter which significantly affects the period--luminosity relation of red supergiants is the mass fraction abundance of heavy elements $Z$. Of special interest is the period--luminosity relation for $Z=0.008$ which is typical for the Large and Small Magellanic Clouds.   \\newpage"
    },
    "1112/1112.0010_arXiv.txt": {
        "abstract": "In this work I investigate the statistical properties of a huge catalog of closely interacting pairs of dark matter haloes, extracted from the Millennium Simulation database.  Only haloes that reach a minimum mass $\\geq 8.6 \\times 10^{10} M_{\\odot}\\, h^{-1}$ (corresponding to 100 particles) are considered.  Close pairs are selected if they come within a critical distance $d_{\\rm crit}$.  I explore the effects of replacing $d_{\\rm crit}=1\\, {\\rm Mpc}\\, h^{-1} \\rightarrow 200\\, {\\rm kpc}\\, h^{-1}$ on the evolution of separations, lifetimes, total masses and mass ratios of these pairs. ",
        "introduction": "Mergers of galaxies play a fundamental role in essentially all modern theories of galaxy formation.  They are believed to determine the morphology of a galaxy, drive its star formation and even activate its nuclear supermassive black hole.  This has motivated the development of very detailed numerical simulations of merging galaxies (often involving only two galaxies in isolation).  Unfortunately, with very few exceptions \\cite[e.g.,][]{tonnesen11}, these simulations typically concentrate primarily on the post-merger aftermath, often neglecting the early stages of interaction.  Observations, on the other hand, tell us a different story.  Surveys like the SDSS and zCOSMOS have confirmed, in large numbers, that galaxies in pairs tend to be bluer, have their star formation enhanced, and are more likely to be active \\citep[e.g.,][]{ellison11,silverman11}.  Moreover, the discovery rate of binary quasars has accelerated to unprecedented levels in the last few years \\citep[e.g.,][]{liu11}.  For this reason, it is vital to re-focus our attention to the early stages of galactic interactions. ",
        "conclusions": "Using the Millennium Simulation, I analyze the behavior of a very large set of halo pairs selected by a proximity criterion.  As this critical distance is reduced, (1) more mergers are pre-selected, (2) pairs end up with more elongated orbits, (3) interactions tend to be shorter, and (4) non-mergers split more violently.  Using a smaller $d_{\\rm crit}$ probes smaller scales, enhancing the importance of pairs involving a massive member and the prevalence of massive neighbors capable of splitting pairs.  With these results, I wish to highlight the following important lessons: \\begin{itemize} \\item {\\bf Extra care must be taken when using close galaxy pairs as proxies for mergers. This subtlety is important even in physical three-dimensional space!}  \\item {\\bf Systems with two merging galaxies in isolation are just an approximation.  In reality, the universe can nurture these duos, or split them altogether!}  \\end{itemize}  Next I will use this catalog to explore the symbiosis between interacting galaxies and binary quasars.  Ultimately, this will set up realistic initial conditions for black hole mergers.  This work is just a puzzle piece of an ambitious long-term research program centered on the evolution of galaxies, supermassive black holes, and their environment."
    },
    "1112/1112.3338_arXiv.txt": {
        "abstract": "{\\normalsize The search for classically stable Type IIA de-Sitter vacua typically starts with an ansatz that gives  Anti-de-Sitter supersymmetric vacua and then raises the cosmological constant by modifying the compactification. As one raises the cosmological constant, the couplings typically destabilize the classically stable vacuum, so the probability that this approach will lead to a classically stable de-Sitter vacuum is Gaussianly suppressed. This suggests that classically stable de-Sitter vacua in string theory (at least in the Type IIA region), especially those with relatively high cosmological constants, are very rare. The probability that a typical de-Sitter extremum is classically stable (i.e., tachyon-free) is argued to be Gaussianly suppressed as a function of the number of moduli.  }  \\vspace{3cm}    \\begin{flushleft} \\today \\end{flushleft}  \\end{center} \\end{titlepage}   \\setcounter{page}{1} \\setcounter{footnote}{0}  \\tableofcontents  \\parskip=5pt ",
        "introduction": "Recent cosmological data strongly suggests that our universe is sitting at a vacuum state with a very small positive vacuum energy density, or cosmological constant. Further cosmological data also suggests that our universe went through an inflationary epoch in its very early stage. This epoch follows from the presence of a vacuum energy density much bigger than today's value. So it is very likely that our universe started with a relatively large vacuum energy density during the inflationary epoch; it subsequently moved down a ``potential landscape\" and landed at the present small value before nucleosynthesis time. A question of great interest is whether this picture is compatible with our present understanding of string theory. In particular, we would like to examine whether the stringy cosmic landscape has features that may suggest the above cosmological scenario.   Although recent studies of flux compactification in string theory suggest that there are numerous solutions to the string/supergravity equations of motion with different vacuum energies \\cite{Bousso:2000xa,Kachru:2003aw,Susskind:2003kw}, we believe that most of them are only extrema of the resulting effective potential. In fact, explicit model building shows that a meta-stable (i.e., classically stable but may have a finite decay time due to quantum tunneling) de-Sitter ($dS$) vacuum is hard to come by. In particular, the search of a single $dS$ minimum (i.e., with only semi-positive scalar field mass-squares) in Type IIA models has so far come up empty. This is somewhat discouraging as the search includes a collection of exponentially many extrema (by varying the fluxes) in Type IIA vacua. On the other hand, this result is really not that surprising from the properties of multidimensional potentials \\cite{Aazami:2005jf,Easther:2005zr}.  Consider a flux compactification with $N$ moduli. An extremum will be (meta-)stable if all the scalar mass-squares are semi-positive. The axionic component of each modulus presumably has an oscillating behavior, so it will hit a minimum half the time. If we hit a maximum, we expect a nearby minimum to which the wavefunction of the universe can easily move. To simplify the discussion, we may assume that it is easy to reach an axionic minimum and we focus mostly only on the real moduli. Now (with canonical kinetic terms), a typical stringy effective potential $V(\\phi_j)$ ($j=1,2, \\cdots ,N$) of $N$ (real) moduli have non-trivial behavior. Since none (or almost none) of the moduli takes a constant value, they are expected to have some extrema too. For a complicated potential $V(\\phi_j)$ sitting at a minimum, the Hessian (i.e., the $N \\times N$ mass-squared (symmetric) matrix, or simply mass matrix) must have only semi-positive eigenvalues. However, this likelihood is very small, as first pointed out by \\cite{Aazami:2005jf}. Let ${\\cal P}$ be the probability that a given de-Sitter solution (an extremum of a positive $V(\\phi_j)$) turns out to be a meta-stable $dS$ minimum (that is, the $dS$ vacuum is tachyon-free). To avoid de-compactification, we consider only the meta-stable vacua within the finite ranges of the moduli. Suppose all the real entries in the Hessian is random, then the probability ${\\cal P}$ that it has only positive eigenvalues is roughly given by \\begin{equation} \\label{intro1} {\\cal P}  \\sim   e^{- \\frac{\\ln 3}{4} (N+ 0.7)^2 } \\end{equation} where $\\ln (3)/4= 0.275$ is obtained in \\cite{Dean:2006wk, Dean2008}. For a relatively large $N$, ${\\cal P}$ is Gaussian-suppressed. Even if the moduli do not couple to each other so the Hessian is diagonal, the probability ${\\cal P}$ for large $N$ is still exponentially suppressed, \\begin{equation} \\label{intro2} {\\cal P}  \\sim   \\left( \\frac{1}{2} \\right)^N = e^{- N {\\ln 2}} \\end{equation} where $\\ln 2= 0.693$. These probabilities apply to searches via trial and errors only.  One may argue that a generic potential $V(\\phi_j)$ must hit some minima somewhere. This is certainly true. In the Type IIA cases, we see that all known (meta-stable) minima happen to have zero (Minkowski) or negative (Anti-de Sitter) vacuum energy densities, where supersymmetry as well as other symmetries help to guide the search for minima. In fact, some searches start from an $AdS$ minimum with only positive mass squares and then lift it to de-Sitter space. However, upon lifting to de-Sitter space, tachyon generically appears. The couplings among the moduli introduces off-diagonal terms in the Hessian. As the cosmological constant increases, the magnitudes of the off-diagonal terms increase as well, and that tend to cause instability.   To see the impact on the stability of the vacuum due to the increase of the vacuum energy, let us start with the diagonal positive mass-square matrix $A$ for an $AdS$ vacuum. Given the masses, we can determine the variance $\\sigma_A$ of $A$. As we lift the cosmological constant, the Hessian at the extremum becomes $A+B$ where $B$ may be treated as a random matrix for a complicated generic $V(\\phi_j)$. The matrix $B$ has variance $\\sigma_B$. This allows us to define the size of the average magnitude of the off-diagonal terms relative to the diagonal mass-squared terms in terms of $y= \\sigma_B/\\sigma_A$. The parameter $y$ essentially describes how the uplifting potential impact on the stability of the vacuum. Now let \\begin{equation} {\\cal P} = a\\, e^{-b N^2 - c N} \\end{equation} where the Gaussian suppression dominates the exponential suppression when $bN/c >1$. Numerically we find that, for small $y$, \\begin{equation} \\begin{split} b&= 0.000395 y + 1.05 y^2 - 2.39 y^3,\\\\ {b\\over c}&= 0.0120+ 2.99 y - 12.2 y^2 + 1650 y^3. \\end{split} \\label{p suppressed off-diagonal comp} \\end{equation} Here we see that Gaussian suppression becomes dominant when $y \\ge 0.0241$ for $N=10$ (where $b/c \\ge 0.1$ and $Nb/c \\ge 1$) and $y \\ge 0.00269$ for $N=50$. For fixed $N$, Gaussian suppression becomes more dominant as $y$ increases.  As an example, we look at a concrete search for a $dS$ vacuum starting from an $AdS$ vacuum undertaken in \\cite{Caviezel:2008tf}. In this $SU(2) \\times SU(2)$ model, there are 14 moduli (7 complex moduli), i.e., $N=14$. At the extremum with positive vacuum energy, $y \\simeq 0.274$, so $bN/c \\gg 1$. This indicates that such a search has a Gaussianly small probability of success. Sure enough, tachyon appears at this extremum.  On the other hand, some generic argument suggests that $dS$ vacua  exist in Type IIB models, especially when non-perturbative effects are turned on. For example, a KKLT vacuum may be obtained by uplifting an $AdS$ minimum with  non-perturbative effects. Attempts to construct $dS$ vacua in Type IIA models studied so far do not include non-perturbative effects. This is because the no-scale structure is present at tree level in Type IIB, while not in Type IIA. The no-scale structure may help to have a hierarchical structure in Type IIB with sub-leading corrections. We shall discuss Type IIB models in a separate paper. Even without going into the details here, we shall use the simple belief that there are metastable $dS$ vacua present in string theory, even though they may be very rare.   Suppose all potentials of $dS$ vacua may be treated as ``uplifts\" of $AdS$ vacua with semi-positive mass-squares. A small uplifting will introduce relatively small off-diagonal terms into the Hessian so the chance of being a $dS$ minimum is relatively good. As we increase the uplifting to higher vacuum energy densities, the off-diagonal terms in the Hessian increase accordingly and the Hessian becomes complicated. The off-diagonal terms tend to push the lowest eigenvalues to negative values. So the chance of this being a $dS$ minimum becomes Gaussianly small (i.e., (\\ref{intro1})) as we go to higher cosmological constants. This leads to the conjecture that there are essentially no $dS$ minima in the relatively higher CC regions in the Type IIA cosmic landscape. The message leads to the following proposal : \\begin{quote} {\\it Raising the cosmological constant destabilizes the classically stable vacua}. \\end{quote} This suggests that, as the universe evolves down the potential, it encounters no $dS$ minima along its way for relatively large cosmological constant (CC) values. Presumably, this happened during the inflationary epoch. Towards the end of inflation, when the universe reaches regions with sufficiently small CC, there may be some $dS$ minima around for the universe to be trapped in one of them. That is, the percolation probability is of order unity for high CC, but decreases substantially by the time the universe reaches the small CC region.  It is very likely that there are many more $AdS$ vacua than $dS$ minima around, but since the universe starts from a relatively high CC region (to generate enough inflation), it has to go through the small positive CC region before reaching the negative CC region. It is not unreasonable that it becomes trapped in a low $dS$ minimum on its way towards negative CC region if the probability of finding a $dS$ minimum with a small CC is not too suppressed. As we shall see, it is reasonable to expect that, for the small positive CC region, some of the moduli with large masses essentially decouple, thus reducing $N$ to a smaller effective value, so the probability of finding a low $dS$ minimum may not be exponentially suppressed compared to the number of $AdS$ vacua nearby.  Since no classically stable Type IIA $dS$ vacuum has been found so far, the above proposal cannot be (non-trivially) checked at the moment. However, we do know some Type IIB solutions; so presumably such a check can be performed for Type IIB regions of the landscape. This study (which involves some subtleties) is under way.  This paper is organized as follows. In section \\ref{sec:prob-estim-rand}, we review and discuss the properties of a real symmetric random matrix as a typical example of a Hessian. We then discuss a Hessian that is the sum of a random diagonal positive matrix plus a random symmetric matrix. This mimics the Hessian that generically appears in the search for a $dS$ minimum in Type IIA models. In Sec. 3, we consider a few examples to illustrate the main point of this paper. Sec. 4 contains some discussions. In particular, we shall comment on an earlier estimate of the probability of obtaining $dS$ vacua  in \\cite{Denef:2004cf}. Some details are relegated to the appendix. ",
        "conclusions": "In this work, we use random matrix theory to estimate the probability of metastable $dS$ vacua in IIA string theory. Our purpose here is to quantify how unlikely we can stabilize all of the K\\\"ahler moduli and the complex structure moduli simultaneously at a $dS$ minimum of a tree level potential. Including  quantum effects will simply yield a more complicated potential, which in turn tends to improve the randomness approximation. Assuming the possibility of introducing some hierarchies in the mass matrix in this setup, we expect different probabilities for different situations. Relatively large off-diagonal terms tend to make the probability ${\\cal P}$ of locating a meta-stable $dS$ vacuum more Gaussianly suppressed than  exponentially suppressed. In the IIA models, we explained why tachyons are ubiquitous in meta-stable de-Sitter extrema in terms of the appearance of the off-diagonal terms in the Hessian. An increase in the cosmological constant typically destabilizes the vacuum as a result of the increase of off-diagonal components.  Note that, in the search for meta-stable $dS$ vacua, the potential used (see (\\ref{potentials in two moduli}) for example) in the search typically introduces multiple free parameters. In attempts to find meta-stable $dS$ vacua, the search allows any appropriate values for these parameters. In reality, these parameters and other hidden ones must be dynamically determined. In the present searches, it is assumed that the modes associated with these parameters (there can be a number of modes associated with each parameter) are heavy and have been dynamically stabilized already. Furthermore, they have enough multiple stable solutions so that their values needed for $dS$ vacuum stability will be among the possible stabilized solutions. This may be a reasonable assumption for low energy scales. However, as the vacuum energy scale is increased to scales comparable to some of lighter masses of these modes, we can no longer ignore the dynamics of these modes. For example, consider the supergravity potential (\\ref{SUGRAV}) (for simplicity without the D term). Suppose one is allowed to increase the energy scale simply by adding a positive constant to $K$, without changing the spectrum. By doing so, we may have to include the lighter part of heavy moduli into the potential that have not been included in $K$ and $W$ so far. The resulting relevant Hessian with increased number of moduli will not be simple in general. So the typical relevant Hessian will grow in size and complication as the vacuum energy is increased. Even though changing the overall scale does not alter the hierarchy between the diagonal and off-diagonal entries in the original mass matrix (which is a subset of the entire mass matrix), the increase of the number of relevant moduli increases the likelihood of instability as we argued in section \\ref{sec:random-matrix-with}. The implication of this viewpoint is as follows: As the vacuum energy scale increases, the number of moduli to be included in the analysis of the relevant mass-squared matrix increases. As a result, the unavoidably complicated form of the matrix will imply that the probability of having a meta-stable $dS$ vacuum becomes increasingly Gaussianly unlikely. This reinforces our view that: \\begin{quote} {\\it The probability of finding a meta-stable de-Sitter vacuum in the string landscape becomes asymptotically Gaussianly unlikely as the vacuum energy is increased.} \\end{quote} Since not a single classically stable Type IIA $dS$ vacuum has been found so far, this proposal/conjecture is compatible with this fact, but at the same time, we cannot check the proposal/conjecture in a non-trivial way. On the other hand, we do know some Type IIB approximate solutions; so the above proposal may be checked (to some extent) for the Type IIB regions of the landscape. However, all known meta-stable Type IIB vacua involve non-perturbative contributions to the 4-dimensional effective potential, thus the analysis is somewhat more subtle.  The probability of obtaining a $dS$ vacuum in string theory was previously considered in \\cite{Denef:2004cf}. In this seminal work, the authors estimated such probability in the context of ${\\cal N}=1$ supergravity (and for simplicity, they considered cases where the D-terms are absent). To maintain stability, the authors required the off-diagonal terms to be suppressed relative to the positive diagonal terms. Since this potential is a function of the superpotential $W$ and its derivatives, such off-diagonal suppression is achieved if we satisfy the condition $|D_A D_B D_C W \\bar{\\psi}_1^A \\bar{\\psi}_1^B \\bar{\\psi}_1^C| < {\\cal O}(DW)$ where the vector $\\psi_1$ specifies a direction for extrema with the smallest eigenvalue, under the assumption of $|DW| \\sim |W|$ (in Planck unit). Since a term proportional to $D_A D_B D_C W$ shows up in the off-diagonal components of the Hermitian mass matrix, such a constraint provides a hierarchy between the diagonal and off-diagonal entries. Now,  we have shown that the probability of stability is Gaussianly suppressed (as a function of the number of moduli) when the off-diagonal terms are within an order of magnitude of the diagonal terms; and  we typically expect one of the lightest modes to turn tachyonic first. This implies that the above condition will be Gaussianly unlikely to satisfy as the number of moduli increases.    Since the Hessian is built  from $W$ and its derivatives, the elements in the Hessian are not totally random. However, for a complicated system with a large number of moduli, we believe that some form of the central limit theorem should  hold and our approach is valid. While this paper was in preparation, we were informed that a more detailed model, based on \\cite{Denef:2004cf} using some random distributions of $W$ and its derivatives has recently been performed \\cite{Marsh:2011aa}.     Furthermore, in our setup, we had in mind a much wider class of constructions, some of which cannot be written  in the form of an ${\\cal N}=1$ supergravity potential without $D$ terms; in these cases, our approach should provide a good estimates of the probability ${\\cal P}$ when the hierarchy between the diagonal and the off-diagonal components of the Hessian can be approximately estimated. We will report on our findings for the probability of metastable $dS$ vacua in IIB string theory in a forthcoming paper."
    },
    "1112/1112.1084_arXiv.txt": {
        "abstract": "We have obtained long-slit spectra of 174 star-forming galaxies with stellar masses greater than $10^{10} M_{\\odot}$  from the {\\it GALEX} Arecibo SDSS (GASS) survey. These galaxies have both HI and H$_2$ mass measurements. The average metallicity profile is strikingly flat out to $R_{90}$, the radius enclosing 90\\% of the $r$-band light. Metallicity profiles which decline steadily with radius are found primarily for galaxies in our sample with low stellar mass (Log$(M_*)<10.2$), concentration, and/or mean stellar mass density. Beyond $\\sim R_{90}$, however, around 10 percent of the galaxies in our sample  exhibit a sharp downturn in metallicity. Remarkably, we find that the magnitude of the outer metallicity drop is well correlated with the {\\it total} HI content of the galaxy (measured as $f_{HI}=M_{HI}/M_*$). We examine the radial profiles of stellar population ages and star formation rate densities, and conclude that the galaxies with largest outer metallicity drops are actively growing their stellar disks, with mass doubling times across the whole disk only one third as long as a typical GASS galaxy. We also describe a correlation between {\\em local} stellar mass density and metallicity, which is valid across all galaxies in our sample. We argue that much of the recent stellar mass growth at the edges of these galaxies can be linked to the accretion or radial transport of relatively pristine gas from beyond the galaxies' stellar disks. ",
        "introduction": "A proper characterization of the ages and metallicities of the stars  in spiral galaxies, as well as the radial dependence of the metallicity of stars and gas in these systems, has long been recognized as a key stepping-stone to unravelling disk galaxy formation processes, including the roles of gas accretion, supernovae-driven outflows and the radial migration of stars \\citep{quirk73, tinsley80, lacey85, wyse89, kauffmann96, chiappini97, schonrich09}.  There have, however, been rather few systematic studies of how radial star formation and metal abundance gradients vary across {\\em populations} of disk galaxies. \\citet{vilacostas92} carried out an analysis of abundance gradients in a sample of 30 spiral galaxies with spectroscopy from the literature. Barred galaxies were found to have flatter abundance gradients than un-barred galaxies. When the analysis was restricted to un-barred galaxies, gradients were found to be stronger in late-type galaxies (see also \\citep[see also][]{oey93}. Vila-Costas \\& Edmunds also found the the central metallicities of spiral galaxies were correlated with both the total stellar mass of the galaxy, and the local central surface density of stars. Zaritsky, Kennicutt \\& Huchra (1994) obtained uniform HII-region based abundance gradients for a sample of 39 nearby disk galaxies. In this study, the slopes of radial abundance gradients, when expressed in units of dex/isophotal radius, did not exhibit any significant correlation with luminosity or Hubble type, but it was later noted by \\citet{garnett97} that gradients expressed in dex/kpc steepen with decreasing luminosity. Further observations by \\citet{vanzee98} focused on increasing the number of measurements in the outer regions of galaxies; they identified several galaxies where the fitted gradient changed significantly with the addition of higher-radius data. Recently, \\citet{moustakas10} presented resolved metallicity measurements for a sample of nearby galaxies from the {\\it Spitzer} Nearby Galaxies Survey \\citep[SINGS,][]{kennicutt03a}, and pointed out that derived metallicity gradients can be sensitive to the methodology used to calibrate the strong-line abundance estimates.  A complete, systematic and carefully-executed survey of metallicity gradients in a large sample of nearby disk galaxies is urgently needed to help put into context some recent studies that have reported unexpected or peculiar metallicity profiles in a number of individual galaxies. A number of these \\citep[e.g.,][]{werk11, rupke10, kewley10, bresolin09} find surprisingly flat metallicity gradients, even out to extreme distances, both for interacting galaxies and peculiarly gas-rich objects, in contrast with earlier work documenting low metallicities or large drops in the extreme outskirts of galaxies \\citep[e.g.,][]{ferguson98, kennicutt03}. The more recent papers argue that these differences reflect the important role of gas inflows and outflows, as well as mixing, in determining how metallicity changes as a function of radius in galaxies.  In order to learn more about the relations between cold gas and star formation in nearby galaxies, we are carrying out the {\\it GALEX} Arecibo SDSS Survey (GASS)\\footnote{http://www.mpa-garching.mpg.de/GASS} \\citep[][hereafter C10]{catinella10}. GASS is designed to measure the neutral hydrogen content of a representative sample of $\\sim 1000$ galaxies uniformly selected from the Sloan Digital Sky Survey \\citep[SDSS,][]{york00} and {\\it Galaxy Evolution Explorer} ({\\it GALEX}, Martin et al. 2005) imaging survey, with stellar masses in the range $10^{10}-10^{11.5} {\\rm M}_{\\odot}$ and redshifts in the range $0.025<z<0.05$. GASS observations are designed to detect HI down to a gas-fraction limit of 1.5-5\\%, so the full GASS sample will be the first HI survey able to place meaningful, unbiased constraints on the atomic gas reservoirs that may contribute to future growth in massive galaxies. We are also pursuing a companion project on the IRAM 30m telescope, COLD GASS\\footnote{http://www.mpa-garching.mpg.de/COLD$\\textunderscore$GASS}, which has obtained accurate and homogeneous molecular gas masses for a subset of 350 galaxies from the GASS sample \\citep{saintonge11}. These data will allow us to characterize the balance between atomic and molecular gas in the galaxies in our sample, and understand the  physical processes that determine how the condensed baryons are partitioned into stars, HI and H$_2$ in the local Universe.  The third component of the GASS survey, and the subject of the current paper, is a follow-up campaign to obtain long-slit spectroscopy for the COLD GASS subset of 350 galaxies. Such spectra allow us to link the measured gas contents of galaxies to their rotational dynamics, metal abundance gradients and resolved star formation histories. In this first paper in our series on  the GASS long-slit data, we focus on studying the gas-phase metallicities of star-forming regions as a function of radius for a partially-complete sample of 174 GASS galaxies.  In \\S2 \\& 3 below, we will describe our observations, data reduction pipeline, and our methodology for  measurement of metallicities. In \\S4, we will show how metallicity  gradients vary as a function of both global and localized galaxy properties. In \\S5, we will present a remarkably tight relation between the total HI content and metallicity in the outer regions of galactic disks, and examine in detail a subset of galaxies exhibiting steep metallicity drops in their outer disks. In the following, we adopt a standard $\\Lambda$CDM cosmology with H$_0=70$km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_\\Lambda=0.7$. Stellar masses and star formation rates are calculated assuming a Kroupa (2001) initial mass function (IMF). ",
        "conclusions": "\\subsection{The Relation Between Gas Content and Metallicity}  \\begin{figure*} \\includegraphics[width=0.66\\columnwidth]{f9a.eps} \\includegraphics[width=0.66\\columnwidth]{f9b.eps} \\includegraphics[width=0.66\\columnwidth]{f9c.eps} \\\\ \\includegraphics[width=0.66\\columnwidth]{f9d.eps} \\includegraphics[width=0.66\\columnwidth]{f9e.eps} \\includegraphics[width=0.66\\columnwidth]{f9f.eps}  \\caption{\\label{global_plots} Outer metallicity, calculated as in the left panel of Figure~8, plotted as a function of, from top left to bottom right, ratio of molecular to atomic gas ($f_{H2}/f_{HI}$), the combination of NUV--{\\it r} and full-galaxy $\\mu_*$ found to best predict HI content by C10, NUV--{\\it r} color alone, molecular gas fraction ($f_{H2}$), residuals from the C10 HI plane, and stellar mass ($M_*$). None predict outer metallicity as well as the HI fraction ($f_{HI}$) alone.} \\end{figure*}   In C10, we found that the HI content of a given GASS galaxy was most accurately predicted by an fundamental-plane style ``HI plane'' constructed from whole-galaxy measurements of $\\mu_*$ and NUV$-r$ color (which is a good proxy for sSFR). Since we have seen in the above section that local $\\mu_*$ and sSFR are also tied closely to the local gas-phase metallicity, it is natural to wonder if we can identify a direct link between metallicity and gas content.  Remarkably, as can be seen in Figure~\\ref{FP}, we have found that there is a surprisingly tight relation between the global HI fraction ($f_{HI}=M_{HI}/M_*$) and the lowest metallicity point that we measure ($R>0.7R_{90}$, O3N2 scale) for each galaxy. We measure a correlation coefficient $|\\rho|=0.53$, significant at the $5\\sigma$ level. We include upper limits as if they were detections for this calculation, but excluding these limits---either in $f_{HI}$, metallicity, or both---does not significantly change the result.  The C10 HI plane can be used to predict $f_{HI}$ with rms accuracy of $\\sim0.3$~dex, and for comparison the scatter in our outer metallicity-$f_{HI}$ relation is 0.4~dex in $f_{HI}$. We note, however, that outer metallicities that are near solar provide little predictive power for HI. In contrast, low metallicity points {\\it do} seem to be universally associated with gas-rich galaxies. Whether or not Figure~\\ref{FP} reflects a {\\it continuous} relation between outer metallicity and $f_{HI}$, or whether it is more properly described as a {\\it threshold} in $f_{HI}$, above which outer metallicities are suppressed in proportion to the gas content, is unclear from the current data. We hope to revisit this issue with the full GASS data-set once it is in hand.  The quantity plotted in the left panel of Figure~\\ref{FP}, being defined as the lowest metallicity we measure in a galaxy's outskirts, could be subject to a number of biases due to the imprecise definition. To check whether the way we define the metallicity value is driving this relation in any way, we also calculated metallicity values by coadding all spectroscopic flux from $R>R_{90}$ into a single high-S/N spectrum for each galaxy. Metallicities measured from these spectra, being integrated over a much larger area of the galaxy, often do not reach such low values as our more finely sampled points. Even so, when plotting these metallicity values against $f_{HI}$, in Figure~\\ref{FP}, right panel, we find a nearly identical relation as before, with a marginally higher $|\\rho|=0.57$ and $6\\sigma$ significance. We note, also, that plotting the size of the outer metallicity {\\it drop}, rather than the outer metallicity directly, gives largely the same results. Considering that the inner metallicities of most GASS galaxies are nearly solar everywhere, this is not surprising.  To understand what might be driving this peculiar relation between a very {\\it local} quantity, the outer metallicity, and the {\\it global} gas content, we must examine how outer metallicity relates to other global galaxy characteristics. In Figure~\\ref{global_plots}, we plot the minimum outer metallicity on the O3N2 scale (as in the left panel of Figure~\\ref{FP}), as a function of several such quantities.  The first quantity of interest is the molecular gas fraction, $f_{H2}$. As can be seen in the lower-left panel, $f_{H2}$ by itself does not correlate with outer metallicity. Since H$_2$ is generally more centrally concentrated in galaxies than HI \\citep[e.g.,][]{leroy09}, this result may not be surprising. In fact, we do detect a weak correlation ($|\\rho|=0.39$, $3.9\\sigma$) between $f_{H2}$ and {\\it central} metallicity. The outer metallicity may simply be more closely associated with the size of a galaxy's gas {\\it reservoir}, as traced by HI, rather than by the amount of currently star-forming gas, traced by H$_2$.  Interestingly, the {\\it ratio} of molecular to atomic gas, shown in the top left panel, {\\it does} correlate with outer metallicity, in the sense that galaxies more dominated by their HI exhibit lower outer metallicity. However, this correlation is weaker ($|\\rho|=0.44$) and less statistically significant ($3.7\\sigma$) than the HI relation in Figure~\\ref{FP}, and so we may only be seeing here a reflection of the underlying relation between HI and outer metallicity.  \\begin{figure*}[t] \\includegraphics[width=\\columnwidth]{f10a.eps} \\includegraphics[width=\\columnwidth]{f10b.eps} \\caption{\\label{metal_gradients} Left: Metallicity versus $R/R_{90}$ for the 13 galaxies with the largest outer metal drops. Right: Metallicity versus $R/R_{90}$ for the matched comparison galaxies. In both panels, solid lines connect adjacent measurements for each individual galaxy, and both halves of each galaxy are plotted separately, `folded' onto the same $R/R_{90}$ scale. Points and lines are color coded for galaxy HI content, according to the color bar at top, which indicates $Log(f_{HI})$ associated with each shade. Galaxies that do not yet have an HI measurement are plotted in black.} \\end{figure*}  In the upper middle panel of Figure~\\ref{global_plots}, we plot outer metallicity as a function of the same combination of global NUV--{\\it r} and $\\mu_*$ that goes into the C10 HI plane--i.e., the x-axis can be thought of as the $f_{HI}$ that is predicted given the galaxy's NUV--{\\it r} and $\\mu_*$. We do this to evaluate the possibility that the $f_{HI}$--metallicity relation is simply an induced correlation: if it is really the star formation rates and stellar mass densities that set metallicity, as our results above seem to suggest is true locally, it may be that the correlation with HI content is not fundamental, but instead driven by the dependence of HI on the same underlying parameters (sSFR and $\\mu_*$, but galaxy-averaged now) that drive metallicity.  What the upper middle panel of Figure~\\ref{global_plots} shows, however, is that the correlation of outer metallicity with global $\\mu_*$ and NUV--{\\it r} is relatively weak compared to Figure~\\ref{FP}, with $|\\rho|=0.37$ and a significance of $4\\sigma$. Since C10 found the relation between $f_{HI}$ and NUV--{\\it r} color alone was nearly as tight as the one including $\\mu_*$, we also plot in the upper right panel outer metallicity versus NUV--{\\it r}. The correlation here is more significant, at 4.8$\\sigma$, but again the correlation coefficient is not quite as strong ($|\\rho|=0.44$) as we saw in Figure~\\ref{FP}.  The {\\it residuals} from the C10 HI plane are useful for identifying galaxies that have either abnormally high or low atomic gas content, compared to that expected for galaxies with their specific combination of $\\mu_*$ and NUV--{\\it r}. In the lower middle panel of Figure~\\ref{global_plots}, we plot these residuals versus outer metallicity. The correlation is quite low and only marginally significant ($2.8\\sigma$), suggesting that it is the absolute level of HI content that matters, rather than the {\\it excess} or {\\it deficit} of HI compared to other similar galaxies.  Finally, we checked for relations between metallicity and other global galaxy properties such as concentration, $\\mu_*$, and $M_*$. None show any significant correlation; the lower right panel of Figure~\\ref{global_plots} shows metallicity versus Log$(M_*)$, which is representative of the others. Outer metallicity thus appears to depend hardly at all on the structural properties of galaxies in GASS.  \\begin{figure*} \\centering \\includegraphics[width=2\\columnwidth]{f11a.eps} \\medskip  \\includegraphics[width=2\\columnwidth]{f11b.eps} \\caption{\\label{low_stamps}SDSS postage stamp images of galaxies with the steepest outer metallicity drops (top), and the matched control sample selected without regard for metallicity (bottom). Images are $80\\arcsec\\times80\\arcsec$, and are arranged so that the match to a galaxy in the top panel is in the same relative position in the lower panel. White bars in each image indicate a projected physical length of 10kpc at the redshift of that galaxy.} \\end{figure*}  In short, we have found that the relation between $f_{HI}$, a global property of the galaxy, and outer metallicity $12+Log(O/H)$, a very local one, is tighter and therefore likely {\\it more} fundamental than any relation between global SFR/$\\mu_*$ and metallicity. The question, then, is why this should be true. One possibility is that metallicity at the outer edge of the star-forming disk is simply a sensitive `thermometer' of sorts for measuring the amount of new gas accreting onto the disk of the galaxy, or perhaps for the rate that existing gas is transported inward. Lower metallicities would simply be reflecting a higher proportion of pristine or relatively unenriched gas residing in or flowing through the outer stellar disk.  If this is correct, then varying rates of gas accretion/flow may leave other signatures in, for example, the stellar populations and star formation rates of galaxies, or on their radial gradients. To evaluate this possibility, in the following section we will examine in detail the subset of galaxies exhibiting the strongest metallicity drops in our sample, and compare to a control sample selected without regard to outer metallicity.  \\subsection{The Galaxies with Steepest Metallicity Drops}  In order to determine whether a low outer metallicity is associated with any distinctive features in the star formation rates and histories of our galaxies, either at their outer edges or across their disks, we select for detailed study a small sample containing only those galaxies with the steepest drops in outer metallicity. We will refer to this as the  `low metallicity' or `large drop' sample, and we will compare it to a `control' sample selected without regard to metallicity, but matched in global characteristics one-for-one with galaxies in the large-drop sample. We note that both samples are selected from the subset of 119 galaxies that exhibit measurable star formation in their outskirts at $R>0.7R_{90}$, which ensures that all the galaxies in the control sample have significant star formation and measurable metallicities at the same large radii as the large-drop galaxies.  In total, we identify 13 galaxies with outer-disk metallicities $12+Log(O/H)<8.4$, a threshold we chose both because no measured point in the inner region of any galaxy reaches this low, and because it is significantly below solar at the $>3\\sigma$ level (0.25~dex). Such galaxies make up about 10\\% of the total sample, and so even though the number identified is small, they represent a significant proportion of all massive star-forming galaxies. For the control sample, we select each matched counterpart by requiring that it be within 0.2~dex in stellar mass, 0.4~dex in global $\\mu_*$, and 0.3 in NUV-{\\it r} color, similar to the procedure in \\citet{wang11a}. These limits were chosen to ensure that each large-drop galaxy has at least one match within our spectroscopic sample, and for cases with more than one match we select randomly among the choices.  In Figure~\\ref{metal_gradients}, in the left-hand panel we plot the radial metallicity profiles of all 13 of our low-metallicity galaxies, as a function of $R/R_{90}$. Solid lines connect adjacent points to more easily follow the two folded halves of each profile (i.e., from both sides of each galaxy). Points are color coded for galaxy HI content, as illustrated by the color bar. It is clear from this plot that even those galaxies with the strongest metallicity drops show typically flat profiles in their inner regions (though some variation can, indeed, be seen). The median difference between the inner metallicity and the lowest measured point is $\\sim0.3$ dex for these objects. In the right-hand panel of Figure~\\ref{metal_gradients}, we now plot the metallicity profiles of the control sample. Again, we see quite flat inner profiles, though there may be a hint that even these galaxies show subtle drops in metallicity at the highest measured points. Consistent with the overall correlation shown in Figure~\\ref{FP}, galaxies with large metallicity drops have on average higher $f_{HI}$, which can be seen by noting the quite different range of colors between the two panels.   \\subsubsection{Morphologies} Our low-metallicity and control samples of galaxies by design have certain features in common, in particular their stellar masses and colors, as well as widespread star-formation extending out to $R>R_{90}$. Yet the disparity in the magnitude of the metallicity drop, as well as the wide variation in HI content, suggests that it is worthwhile to examine the images of both samples directly, to check for any subtle structural differences between the two that are not apparent when considering just the broadest characteristics.  In Figure~\\ref{low_stamps}, we show SDSS postage stamp images of these galaxies, with large-drop galaxies in the top panel, and the control galaxies at bottom. Images are arranged so that each matched pair of galaxies is in the same relative location in the top panel and bottom. We remind the reader that galaxy pairs were not matched in redshift (beyond the normal GASS constraint of $0.025<z<0.05$), so differences in apparent size are to be expected. For aid in comparison, white bars beneath each image indicate 10kpc projected distance at the redshift of each galaxy.  The first striking feature of Figure~\\ref{low_stamps} is the large number of relatively edge-on galaxies in the low-metallicity sample. There are two possible explanations for this trend, both related to observing geometry. First, if the lowest metallicity regions are distributed unevenly, in a stochastic manner around the outer edges of galaxy disks, then an edge-on galaxy may afford us a higher probability of actually observing one, as the light collected by our slit has passed through a larger proportion of the galaxy. As can be seen in the Appendix, where individual profiles of large-drop galaxies are shown, some of these galaxies (but not all) show a metallicity drop on one side only, which might argue in favor of this scenario.  Alternatively, surface brightness enhancement due to viewing the galaxy edge-on could simply enable our fixed-integration time observations to reach further out in an edge-on galaxy, where such low-metallicity star-forming regions are perhaps more likely to reside. We do note that the maximum radius reached for edge-on galaxies is somewhat higher than for other objects (median of $1.2R_{90}$ for $b/a>0.5$ vs $1.4R_{90}$ for $b/a<0.5$), but it is unclear if this difference is large enough to cause the observed bias. The median radius reached for the control sample galaxies is only slightly lower than the large drop galaxies: $1.2R_{90}$ and $1.3R_{90}$, respectively.  A few galaxies in Figure~\\ref{low_stamps} seem to harbor close-in tidally disturbed companions, which could suggest that new gas accreted or cannibalized from these companions is the source of the low-metallicity material. However, the majority of the large-drop galaxies do not have obvious companions, and the frequency of companions is not obviously larger than that seen in the control sample images. Indeed, the one large-drop galaxy that we have previously studied in detail \\citep{moran10}, shows remarkably little evidence for any sort of dynamical disturbance, such as one would expect in the case of an accreted companion. Though beyond the scope of this paper, rotation curves are available for all of these galaxies, and we expect to address the relation between gas content and galaxy dynamics and/or mergers in a future paper.  We measured a number of quantitative morphological parameters for both sets of galaxies to search for any differences. These included asymmetry, bar fraction, and an asymmetry variant weighted to the outer disk \\citep{wang11b}. Though such small samples make it difficult to assess any statistical difference between the two sets of galaxies, we find no clear evidence that the populations are different in any of the key morphological measures. This remains true whether or not we exclude the most edge-on galaxies (which could have peculiar values of these morphological parameters).  Besides the obvious difference in orientation, then, we find no dramatic differences in appearance between the large-drop and control sample galaxies in Figure~\\ref{low_stamps}. Below, we will examine the radial profiles of quantities detailing the star formation histories and current star formation rates as a function of radius across each galaxy type.   \\subsubsection{Star Formation Histories}  So far, then, our large-drop sample of galaxies appears to differ from the control sample only in two important properties: metallicity drop and HI content. But the key question remains: why does a high HI content appear to drive low-metallicity star formation in galaxies' outskirts? To help answer this question, we can look to the star formation rates and histories of our two groups of galaxies. Specifically, we can examine the radial profiles of a number of spectroscopic diagnostics of both past-averaged and present-day star formation.  \\begin{figure*}[t] \\includegraphics[width=0.66\\columnwidth]{f12a.eps} \\includegraphics[width=0.66\\columnwidth]{f12b.eps} \\includegraphics[width=0.66\\columnwidth]{f12c.eps} \\\\ \\includegraphics[width=0.66\\columnwidth]{f12d.eps} \\includegraphics[width=0.66\\columnwidth]{f12e.eps} \\includegraphics[width=0.66\\columnwidth]{f12f.eps} \\caption{\\label{other_gradients} From left to right, $D4000_n$, SFR density, and sSFR gradients as a function of $R/R_{90}$, for galaxies in the large-metallicity-drop sample (top row), and those in the control sample (bottom row). Grey points show gradients with adjacent points connected by lines, as in Figure~10. Thick black points with error bars show the mean and 1-sigma values for bins of $\\sim25$ points each. Light grey shading shows the $\\pm1-\\sigma$ region occupied by the full GASS sample with extended star formation, divided into radial bins of $\\sim150$ points each, and including only star-forming regions for $D4000_n$. We plot $D4000_n$ only for points with continuum S/N greater than 5 per angstrom, and uncertainty in $D4000_n$ less than 0.1. Typical uncertainties are less than 0.05.} \\end{figure*}  In Figure~\\ref{other_gradients}, we plot, from left to right, the $D4000_n$ index, SFR density (in $M_\\odot$~yr$^{-1}$~kpc$^{-2}$), and sSFR (yr$^{-1}$) as a function of $R/R_{90}$, for both the galaxies with large metallicity drops (top), and the control sample (bottom). As in Figure~\\ref{metal_gradients}, grey lines and dots connect the points for individual galaxies. Black dots with error bars show the average and scatter in bins of radius, for all the galaxies in each subsample. Light grey shading denotes the $\\pm1\\sigma$ region occupied by our {\\it entire} 119-galaxy GASS sample with extended star formation, as a function of radius. We note that $D4000_n$ averages for the full sample are calculated only from points with detected star-formation, to ensure the curves reflect the same set of galaxies (and individual points) as those in the other two panels.  The top panels of Figure~\\ref{other_gradients} show that, for the large-drop galaxies, the radial profiles of all three quantities show rather similar trends from galaxy to galaxy. $D4000_n$ generally decreases monotonically from center to outskirts (signifying decreasing stellar population ages), specific star formation rate {\\it increases} monotonically towards the outside, and SFR density seems to exhibit largely flat radial profiles for most galaxies (though the precise level of SFR density varies considerably).  Moreover, it is clear by comparing to the underlying shaded region marking the full GASS sample, that, while the trends for decreasing age and increasing sSFR with radius are typical of all star-forming galaxies in this mass range, large-drop galaxies appear overall younger and are building up stellar mass faster than the typical GASS galaxy. Large-drop galaxies have $D4000_n$ shifted an average of 0.16 lower than the full-sample mean, and sSFR is fully 0.5~dex higher, which translates to mass-doubling times three times shorter---as low as 1~Gyr at the galaxy edges in the most extreme cases. Both of these differences are statistically significant in all radial bins. SFR density is also higher by an average of 0.27~dex, though the difference in the innermost bin is not significant.  We note that all three plots for the large-drop galaxies seem remarkably similar to the detailed properties of UGC8802 described in \\citet{moran10}, where we argued that the combination of flat SFR density profile and declining $D4000_n$ profile could easily be replicated by a toy model featuring a recent episode of constant star formation spread evenly across the galaxy, on top of an older stellar population that built up most of the pre-existing stellar mass at an earlier time ($>1-2$ Gyr ago). The fact that $D4000_n$ indicates stellar populations that are everywhere younger than the full-sample average, while sSFR is everywhere higher than the full-sample average, lends support to this scenario.  One of the most striking features in Figure~\\ref{other_gradients} is the large difference in scatter between the large-drop and control samples, visible in all three measured quantities. In contrast to the fairly uniform profiles of the large-drop galaxies, those of the control sample are quite heterogeneous. For each of $D4000_n$, SFR density, and sSFR, the formal rms scatter at every radius bin in the large-drop sample is very nearly half what we measure in the corresponding control sample bin (with the precise ratio ranging from 50\\% to 65\\%). The level of heterogeneity displayed by the control sample is essentially indistinguishable from that of the full GASS sample, and the two samples have the same mean values as well, with one exception discussed below. Thus, the spread in star formation rates and stellar population ages of the control sample can be considered typical for star-forming galaxies in this mass range, while those for the large-drop sample are abnormally uniform and offset.  Only in the outermost bin (or two bins, for sSFR) do the mean properties of the control sample of galaxies differ significantly from those of the full GASS sample. In each measured parameter, the mean value of this outermost bin (or two) is instead consistent with that of the large-drop sample. In other words, the high sSFRs and young ages at the very edges of our big-drop galaxies are indistinguishable from those of other galaxies selected to have the same global NUV-{\\it r}, $M_*$, and $\\mu_*$, {\\it even though the metallicities are dramatically different.} What this means is that we cannot definitively link the high sSFRs and young ages to the high HI content and/or low metallicities at these locations, because we cannot exclude the possibility that these features are generic for galaxies with this combination of global properties. In contrast, the uniformly younger and more vigorously star-forming profiles at lower radii can perhaps {\\it only} be attributed to the high HI and/or low outer metallicities, since all the other parameters are identical between the two samples.  Can these two statements, seemingly at odds, be tied together into a uniform picture of what is happening in the high HI, large-metallicity drop galaxies? We believe the following scenario is plausible:  \\bigskip  In galaxies with a large HI content, it has been shown that much of the gas often resides beyond the optical disk of the galaxy \\citep[e.g.,][]{bigiel08}. During the process of galaxy growth, much of this gas must eventually be transported inward, where it will form molecular gas and then stars. At the same time, it has been known for a long time that radial flows of gas can naturally lead to metallicity gradients \\citep[e.g.,][]{lacey85}, and that an `inside--out' buildup of galaxy disks---like we see in the sSFR profiles of big drop galaxies---may be required to explain the observed strengths of gradients \\citep{boissier00}. We speculate that it is this transport of gas inward that sets the level of metallicity suppression at the outer edges of our galaxies: star-forming gas at the optical edge of the galaxy is diluted in simple proportion to the total amount of gas residing in the extended reservoir, and the continuous flow of such gas serves to {\\it keep} the metallicity low.  Under this scenario, the homogeneity of the radial profiles in the `large-drop' sample might arise because the star formation across the entire galaxy becomes dominated by the dynamics of this inwardly-transported gas. A dense flow of gas, by providing ample fuel to all corners of the galaxy, could well act to suppress the normal spatial variations in star formation rate seen in more typical galaxies, and at the same time cause a period of intense disk-building that elevates sSFRs everywhere, including at the outskirts where the disk is building fastest. In the control sample, small amounts of gas may also be building up the disk outskirts, but with a more-nearly-complete metal enrichment (or, rather, a less effective dilution of metals). In these systems, however, the lower quantities of gas involved are not enough to homogenize the star formation rates or stellar population age gradients along the lines seen for the large-drop objects.  Let us return briefly to the question of flat inner metallicity gradients, because it is puzzling why gradients should `saturate' and flatten out so effectively in our high-mass galaxies. Generically, our observational result implies that the new metals produced by star formation are everywhere (except at the outskirts) precisely counterbalanced by either the net inflow of gas, or outflowing winds, or both. Since inner metallicity profiles are flat for galaxies with both high and low $f_{HI}$ (Figure~\\ref{metal_gradients}), this balance presumably holds over a range of gas densities. Such an equilibrium in metallicity may seem implausible, but on the other hand could just be the latest of many such observational `conspiracies' in the properties of galaxies. In any case, it appears that models of the Milky Way with varying assumptions about cosmological gas infall and the details of `inside--out' formation can produce a range of flat to sloping inner gradients \\citep[][and references therein]{colavitti09}. A full comparison to models will not be undertaken here, but it is important to recognize that at least some models {\\it can} reproduce both the flat inner gradients and the steep outer drops we see.  To recap, we have clearly shown that galaxies with large outer metal drops not only have high HI content, but show evidence that this gas is currently involved in a substantial and widespread disk-building phase. Though the fates of these galaxies as their gas depletes and metallicities rise is still unclear, it does seem to be the case that a low outer metallicity is a sensitive signpost or thermometer for the presence of large amounts of active gas in a galaxy. Given that 10\\% of our overall GASS sample appears to be in this phase actively building disks, even though GASS probes a stellar mass range where such activity is thought to be {\\it decreasing}, it will be interesting to see whether the abundance of these systems is compatible with theoretical models of galaxy formation and evolution."
    },
    "1112/1112.4497_arXiv.txt": {
        "abstract": "Using the Inamori Magellan Areal Camera and Spectrograph (IMACS) integral-field unit (IFU) on the 6.5\\,m Magellan telescope, we have designed the first statistically significant investigation of the two-dimensional distribution and kinematics of ionised gas and stars in the central kiloparsec regions of a well-matched sample of Seyfert and inactive control galaxies selected from the Sloan Digital Sky Survey. The goals of the project are to use the fine spatial sampling (0.2 arcsec~pixel$^{-1}$) and large wavelength coverage (4000--7000\\AA) of the IMACS-IFU to search for dynamical triggers of nuclear activity in the central region where active galactic nucleus (AGN) activity and dynamical timescales become comparable, to identify and assess the impact of AGN-driven outflows on the host galaxy and to provide a definitive sample of local galaxy kinematics for comparison with future three-dimensional kinematic studies of high-redshift systems. In this paper, we provide the first detailed description of the procedure to reduce and calibrate data from the IMACS-IFU in `long-mode' to obtain two-dimensional maps of the distribution and kinematics of ionised gas and stars. The sample selection criteria are presented, observing strategy described and resulting maps of the sample galaxies presented along with a description of the observed properties of each galaxy and the overall observed properties of the sample. ",
        "introduction": "Supermassive black holes are widely accepted to lie at the centres of all bulge-dominated galaxies, with the observed correlation between black hole mass and bulge velocity dispersion evident in both active and inactive galaxies confirming a causal link between previous phases of accretion-driven nuclear activity in the galactic life cycle \\citep{geb00,merr01}. Although, such active galactic nuclei (AGN) are recognised to be integral to galaxy formation and evolution, the AGN triggering and fuelling mechanisms remain unknown. Further, the transportation of fuel to the AGN and any corresponding ejection or feedback are key components of cosmological models and are tuned to deliver observed properties of today's galaxies (e.g., \\citealt{spring05}). Merger-driven galaxy evolution and nuclear activity is central to these models but is unlikely to be a valid mechanism in the local Universe, where the merger rate has declined and the $\\sim$20\\% of galaxies currently exhibiting nuclear activity \\citep{gould10} are commonly found in early type spiral galaxies in non-cluster environments and do not show signatures of strong interactions \\citep{west07,gabor09}. Identifying the triggering and fuelling mechanism and the origin and transportation of fuel in AGN therefore remain an important goal in astronomy.  Optical/IR {\\em imaging} studies have remained ambiguous in identifying a single fuelling mechanism \\citep{86bus,fuentes88,deRob98,hunt04,tang08,kuo08}. External perturbations, such as tidal interactions or minor mergers, and internal instabilities in non-axisymmetric potentials such as bars or $m=1$ spirals have long been suggested as viable mechanisms for removing angular momentum from host galaxy gas to allow it to move closer to the AGN \\citep{wada04,witold97}, and more recently hydrodynamic turbulence in the nuclear interstellar medium has been suggested to contribute to low-level accretion \\citep{alig11}, but direct evidence of such physical mechanisms has been hard to obtain (e.g., \\citealt{cgm99,martini01}). The availability of large numbers of uniformly derived galaxy properties from the Sloan Digital Sky Survey (SDSS; \\citealt{york00}) has allowed statistical comparisons of active and inactive galaxies to be performed and evolutionary trends to be investigated in an attempt to explain the location of active and inactive galaxies across the so-called blue and red sequences in galaxy color-magnitude space (\\citealt{bald04,west07,schaw09}). These studies build on earlier work on smaller samples, such as that by \\cite{hunt04}, who suggested an evolutionary sequence driven by an underlying, but unidentified, dynamical instability to explain structural differences observed in a sample of 250 active and inactive host galaxies imaged with NICMOS on the \\textit{Hubble Space Telescope} (\\textit{HST}). Direct determination of the dynamical properties of active and inactive host galaxies is therefore vital.  Radio interferometric spectroscopy provides valuable information on the large-scale gaseous environment and gaseous structure and kinematics of the host galaxy disks to large radii (e.g., \\citealt{cgm07,haan08,haan09}) but surface brightness sensitivity constraints of current interferometers limits achievable angular resolutions to $\\sim5$\\as\\ at best (\\citealt{mund99,things08}), but more typically $\\sim$18\\arcsec. Millimetric observations can sample dense gas, such as CO, closer to the nucleus, but this gas may be clumpy and discontinuous, limiting the ease with which the velocity fields can be interpreted (e.g., \\citealt{dumas10}). Optical observations of galactic kinematics offer the capability of sampling the velocity field close to the centre of a galaxy; in Seyferts, such kinematic observations were traditionally performed with a combination of narrowband imaging to identify line-emitting regions and long-slit spectroscopy in a preferred position angle (e.g., \\citealt{ekers83,tad89a,wilson94}), but disentangling the kinematics of the host galaxy from non-gravitational AGN-driven outflows is difficult and model dependent.  With recent developments in integral-field unit (IFU) technology on large telescopes, which provide 2D imaging spectroscopy, it is now feasible to study the distribution and kinematics of gas {\\em and} stars on scales ever closer to the galactic nucleus (e.g., \\citealt{fathi06,barb06,riffel08}). Spurious misinterpretation of one-dimensional velocity fields is less likely when full two-dimensional velocity fields are sampled; the comparison of stellar and gaseous velocity fields is vital to identify kinematics associated with the galactic gravitational potential rather than non-circular gaseous streaming motion that may be erroneously interpreted in long-slit data e.g., mistakenly implying the presence of black holes offset from their galactic centres \\citep{wilson85,ferruit04}.  The two-dimensional approach was particularly effective in the comparison of stellar and gaseous kinematics in a small matched sample of active and inactive galaxies by \\cite{dumas07}, who used the SAURON IFU \\citep{bacon01} on the 4.2-m William Herschel Telescope to tentatively identify a kinematic difference between Seyfert and inactive galaxy hosts, after removal of non-gravitational AGN dynamics. The SAURON IFU provides a large field of view (FOV, 33\\as $\\times$ 41\\as) but at the cost of pixel sampling ($\\sim 1$\\as) and the sample was selected before SDSS data were available. We, therefore, carefully selected a larger, distance-limited and well-matched sample of active and inactive control galaxies from the SDSS (see \\citealt{west07}) and observed them with the IMACS-IFU on the 6.5-m Magellan telescope \\citep{bigelow98,schmoll04}. The IMACS-IFU is particularly well-suited to the study of galactic nuclei as it has small pixels (\\das02) to sample its (4\\as $\\times$ 5\\as) field of view and an unusually large wavelength coverage ($\\sim 4000 - 7000$~\\AA), which provides access to all major emission lines from \\hb\\ to \\ha\\ as well as the underlying stellar continuum and stellar absorption lines such as Mg\\,$b$ and \\feii. Table~\\ref{tab:IFUs} shows a comparison of selected integral-field spectrographs.   IMACS is most commonly used as a multi-object spectrograph (MOS) and published documentation concentrates on this MOS mode. In this paper, we describe the full procedure for obtaining stellar and gaseous distribution and kinematic maps from Magellan IMACS-IFU data. In particular we present, for the first time, the detailed procedure for reducing and calibrating IMACS in IFU `long-mode'. We present the resultant gaseous and stellar maps in catalogue form for our SDSS-selected IMACS-IFU sample of active and matched-inactive galaxies. The detailed dynamical analysis, interpretation and comparison will be presented in Paper II. ",
        "conclusions": "We have presented the first two-dimensional maps of ionised gas and stars in an SDSS-selected sample of active and inactive galaxies obtained using the IMACS-IFU on the Magellan-I 6.5m telescope. This study required the development of a full reduction pipeline for IMACS-IFU, which was previously unavailable. In this paper we have provided a detailed description of the reduction process and analysis of IMACS-IFU data, from observing strategy to final extracted maps, with particular emphasis on aspects that affect IMACS-IFU such as telescope flexure. Following extensive testing of the IMACS-IFU pipeline, two-dimensional stellar and gas kinematics were derived for 26 of the 28 galaxies in the sample thus far.  In contrast to other IFUs, the unusually large wavelength coverage ($\\sim 4000-7000$ \\AA) provided by the IMACS-IFU, coupled with the fine pixel sampling across the \\das415 $\\times$ \\das500 FOV has allowed the extraction of gaseous and stellar kinematics that probe both the AGN-related regions and the host galaxy. In particular, simultaneous observation of \\oiii\\ and \\ha\\ emission lines provide independent probes of galaxy rotation and AGN-driven outflow.  Evidence of rotation was found in the stellar velocity fields of 11 out of the 26 galaxies analysed, while the S/N of the remaining galaxies was deemed insufficient to make any conclusions. Fifteen of twenty-six galaxies showed clear evidence of rotation in the \\ha\\ velocity fields. Seven Seyfert galaxies show possible evidence of an \\oiii\\ outflow component, the most extreme of which exceeds a blue-shifted line-of-sight velocity of 200\\,\\kms\\ and extends over 2\\as\\ ($\\sim$1.5~kpc).  We have demonstrated the value of large simultaneous wavelength coverage, which allows for the derivation of the underlying host-galaxy dynamics from \\ha\\ kinematics, regardless of the perturbations revealed in \\oiii\\ kinematics. Full kinematics, modelling and its interpretation will be presented in Paper II (P. B. Westoby et al., in preparation)."
    },
    "1112/1112.0561_arXiv.txt": {
        "abstract": "{Future cosmological surveys will provide 3D large scale structure maps with large sky coverage, for which a 3D Spherical Fourier-Bessel (SFB) analysis in spherical coordinates is natural. Wavelets are particularly well-suited to the analysis and denoising of cosmological data, but a spherical 3D isotropic wavelet transform does not currently exist to analyse spherical 3D data.}{The aim of this paper is to present a new formalism for a spherical 3D isotropic wavelet, i.e. one based on the SFB decomposition of a 3D field and accompany the formalism with a public code to perform wavelet transforms.}{We describe a new 3D isotropic spherical wavelet decomposition based on the undecimated wavelet transform (UWT) described in Starck et al. 2006. We also present a new fast Discrete Spherical Fourier-Bessel Transform (DSFBT) based on both a discrete Bessel Transform and the HEALPIX angular pixelisation scheme. We test the 3D wavelet transform and as a toy-application, apply a denoising algorithm in wavelet space to the Virgo large box cosmological simulations and find we can successfully remove noise without much loss to the large scale structure.}{We have described a new spherical 3D isotropic wavelet transform, ideally suited to analyse and denoise future 3D spherical cosmological surveys, which uses a novel Discrete Spherical Fourier-Bessel Transform. We illustrate its potential use for denoising using a toy model. All the algorithms presented in this paper are available for download as a public code called {\\tt MRS3D} at \\url{http://jstarck.free.fr/mrs3d.html}}{} ",
        "introduction": "\\subsection*{Challenges in Modern Cosmology}  The wealth of cosmological data in the last few decades \\citep{wmap7,Schrabback:2010,Percival:2007} has led to the establishment of a standard model of cosmology, which describes the Universe as composed today of approximately 4\\% baryons, 22\\% dark matter and 74\\% dark energy. The main challenges in modern cosmology are to understand the nature of both dark energy and dark matter, as well as the initial conditions of the Universe \\citep{DETF,WGFC}. A thorough understanding of these three topics may lead to a revision of Einstein's theory of General Relativity and our view of the early Universe.   New surveys are planned who aim to answer these important questions e.g. Planck for the CMB \\citep{Planck}, DES \\citep[Dark Energy Survey,][]{DES:2005}, BOSS \\citep[Baryon Oscillation Spectroscopic Survey,][]{Schlegel:2007}, LSST \\citep[Large Synoptic Survey Telescope,][]{LSST} and Euclid \\citep{Euclid:2011} for weak lensing and the study of large scale structure with galaxy surveys.  The challenge with these upcoming large data-sets is to extract the cosmological information in the most suitable manner in order to test the cosmological paradigm. Depending on the signal one wishes to extract, and/or survey geometry, different bases may be more or less suitable (e.g., Fourier, Spherical Harmonic, Configuration or Wavelet Space). Moreover, future surveys may be in 2D (e.g. Planck) or in 3D (e.g. galaxy or weak lensing surveys). Where 3D data is available, a tomographic analysis is possible (also known as 2D 1/2), or a full 3D analysis can be done. For data in spherical coordinates, this corresponds to a Spherical Fourier-Bessel (SFB) decomposition \\citep{Heavens:1995,Fisher:1995,weak3d,Erdogdu:2005wi,Erdogdu:2006dv,Leistedt2011,Rassat:2011bao}.   \\subsection*{Wavelet Transform on the Sphere}  Wavelets are particularly well suited to the analysis of cosmological data \\citep{aw:martinez93,starck:sta05_2}, since cosmological data can often be sparsely represented in wavelet space.  2D Wavelets have been used  in many astrophysical studies \\citep{starck:book06} for a broad range of applications such as denoising, deconvolution, detection, etc. CMB studies have motivated the development of 2D spherical wavelet decompositions.  Continuous  wavelet transforms on the sphere~\\citep{wave:antoine99,wave:tenerio99,wave:cayon01,wave:holschneider96} have been proposed, mainly for non Gaussianity studies. In \\citet{starck:sta05_2},  an invertible isotropic undecimated wavelet transform (UWT) on the sphere based on spherical harmonics was described, that can be also used for other applications such as  deconvolution, component separation \\citep{starck:yassir05,bobin-gmca-cmb,delabrouille08}, inpainting \\citep{inpainting:abrial06,starck:abrial08}, or Poisson denoising \\citep{schmitt2010}. A similar wavelet construction has been published  in \\citep{marinucci08,fay08a,fay08} using so-called ``needlet filters\", and in \\citet{wiaux08},  an algorithm was proposed which allows us to reconstruct an image from its steerable wavelet transform. Other multiscale  transforms on the sphere such as ridgelets and curvelets have been developed  \\citep{starck:sta05_2}, and are well adapted to detect anisotropic features. Other multiscale transforms on the sphere, such as ridgelets and curvelets, have been developed \\citep{starck:sta05_2} and are well adapted to detect anisotropic features. An extension of this UWT has also been developed for polarised CMB data  in \\citet{starck:pola09}.    In this paper, we describe a new 3D isotropic spherical wavelet decomposition, which is reversible, and could therefore be useful for many different applications. It is based on the UWT proposed by \\cite{starck:sta05_2} and extended into 3D. The 3D-UWT proposed here can be used to analyse 3D data in spherical coordinates, such as a 3D galaxy or weak lensing survey with large (but not necessarily full) sky coverage. ",
        "conclusions": "Modern cosmology requires the analysis of 3D fields on large areas of the sky, i.e. where the field is best viewed in spherical coordinates. In this configuration, a Spherical Fourier Bessel (SFB) transform is the most natural way to statistically analyse the field. Wavelet transforms have been shown to be ideally suited for cosmological fields, which tend to be sparse in wavelet space. The wavelet transform can be used e.g. for denoising, but there is yet no 3D wavelet transform in spherical coordinates.  We present in this paper a new 3D spherical wavelet transform, based on the undecimated wavelet transform (UWT) described in \\citep{starck:sta05_2}. In order to perform operations on the wavelet transforms (such as denoising), we require a discrete version of the SFB transform for both the direct and inverse transforms. We show a novel way to perform such a fast Discrete Spherical Fourier-Bessel Transform (DSFBT) based on both a discrete Bessel transform and the HEALPIX angular pixelisation scheme.  Using the 3D wavelet transform and the DSFBT, both introduced in this paper, we denoise a test large scale structure data set, taken from the Virgo large box simulations\\footnote{\\url{http://www.mpa-garching.mpg.de/Virgo/VLS.html}}. We find we can satisfactorily remove artificially added Gaussian noise without much loss to the large scale structure. All the algorithms presented in this paper are available for download as a public code called {\\tt MRS3D} at \\url{http://jstarck.free.fr/mrs3d.html}."
    },
    "1112/1112.2087_arXiv.txt": {
        "abstract": "{Transit and radial velocity observations continuously discover an increasing number of exoplanets. However, when it comes to the composition of the observed planets the data are compatible with several interior structure models. Thus, a planetary parameter sensitive to the planet's density distribution could help constrain this large number of possible models even further.} {We aim to investigate to what extent an exoplanet's interior can be constrained in terms of core mass and envelope metallicity by taking the tidal Love number $k_2$ into account as an additional, possibly observable parameter.} {Because it is the only planet with an observationally determined $k_2$, we constructed interior models for the Hot Jupiter exoplanet HAT-P-13b by solving the equations of hydrostatic equilibrium and mass conservation for different boundary conditions. In particular, we varied the surface temperature and the outer temperature profile, as well as the envelope metallicity within the widest possible parameter range. We also considered atmospheric conditions that are consistent with nongray atmosphere models. For all these models we calculated the Love number $k_2$ and compared it to the allowed range of $k_2$ values that could be obtained from eccentricity measurements of HAT-P-13b.} {We use the example of HAT-P-13b to show the general relationships between the quantities temperature, envelope metallicity, core mass, and Love number of a planet. For any given $k_2$ value a maximum possible core mass can be determined. For HAT-P-13b we find $M_\\mathrm{core}<27$\\,\\ME, based on the latest eccentricity measurement. We favor models that are consistent with our model atmosphere, which gives us the temperature of the isothermal region as $\\sim2100$\\,K. With this external boundary condition and our new $k_2$-interval we are able to constrain both the envelope and bulk metallicity of HAT-P-13b to 1 -- 11 times stellar metallicity and the extension of the isothermal layer in the planet's atmosphere to 3 -- 44\\,bar. Assuming equilibrium tidal theory, we find lower limits on the tidal $Q$ consistent with $10^3 - 10^5$.} {Our analysis shows that the tidal Love number $k_2$ is a very useful parameter for studying the interior of exoplanets. It allows one to place limits on the core mass and estimate the metallicity of a planet's envelope.} ",
        "introduction": "Today more than 180 transiting exoplanets have been discovered. Within the exoplanet family planets that are found to transit their host star are especially important. The knowledge of mass \\emph{and} radius enables us to infer the density and bulk composition of a planet. Further, transit and secondary eclipse observations reveal information about an exoplanet's atmosphere.  Despite all the information given, we are still far away from knowing the composition of the deep interior of exoplanets. Characteristics like the existence and mass of a potential core or the amount and distribution of heavy elements are quite ambiguous. Yet, this information is highly desired in order to understand planet formation.  For the solar system planets, the ambiguity of interior models can be reduced by taking into account information from the gravity field, which is quantified by the gravitational moments $J_2$, $J_4$, and $J_6$. In our solar system these parameters are measureable and improved methods continuously increase their accuracy. However, for an extrasolar planet gravitational moments cannot be determined. Hence, we need a similar parameter that is accessible and will also provide us with information about the interior density distribution of the planet. This can be accomplished with the tidal Love number $k_2$ \\citep{Love11,Gavrilovetal75,GavrilovZharkov77,ZharkovTrubitsyn78}. As it is equivalent to $J_2$ for the solar system planets \\citep{Hubbard84}, it is promising that $k_2$ will help us to further constrain the interior models of exoplanets. Recent studies have shown that $k_2$ is itself in general a degenerate quantity if considered in a three-layer model \\citep{Krammetal11}.  In this study, we focus on the transiting Hot Jupiter HAT-P-13b \\citep{Bakosetal09}. This planet is of great interest because it is the only planet so far which has an observationally determined value for the Love number $k_2$ \\citep{Batyginetal09}. Hence, it gives us the chance to investigate how much information about the interior structure can be inferred from an actually measured $k_2$ and serves as a case study in this work. We especially evaluate the impact of the new eccentricity measurement for HAT-P-13b by \\citet{Winnetal10} as it gives new information about allowed $k_2$ values.  In \\S\\ref{sec:Methods} we describe our modeling procedure, the equation of state (EOS) used, and our calculation of the Love number. The known observational constraints of the examined planet HAT-P-13b are given in \\S\\ref{sec:inputmodelassumptions} where we also discuss the effect of different eccentricity measurements, and describe the setup of our interior models. The results are explained in \\S\\ref{sec:Results}, and in \\S\\ref{sec:Discussion} we compare with previous results. The applicability of the underlying theory is discussed in the Appendix. A summary is given in \\S\\ref{sec:Summary}. ",
        "conclusions": "\\label{sec:Summary}  In this work, we presented new interior models of the transiting Hot Jupiter HAT-P-13b with the aim of showing to what extent interior models of extrasolar giant planets can be constrained by using the tidal Love number $k_2$ in addition to the known observables mass and radius. We also varied the envelope temperature and metallicity in order to demonstrate the uncertainties imposed on the inferred interior models.  One main result of our work is that based on the Love number $k_2$ one cannot draw a conclusion on the precise core mass of the planet. Only a \\emph{maximum possible} core mass can be inferred which is given by adiabatic zero-envelope-metallicity models. Taking into account the allowed $k_2$ interval ($0.116<k_2<0.425$) derived from \\citet{Batyginetal09}, we find a maximum value for HAT-P-13b's core mass of 87\\,\\ME. Further radial velocity observations from \\citet{Winnetal10} gave a new value for the planet's eccentricity, allowing us to construct a \\emph{new $k_2$ interval} of $0.265 < k_2 < 0.379$. With this smaller interval we were able to constrain the core mass of HAT-P-13b further down to 27\\,\\ME.  We also showed the influence of the envelope's temperature and metallicity. The coldest possible model defines a maximum possible value for $k_2$ as this is the most homogeneous model. For HAT-P-13b we find $k_{2,\\mathrm{max}}=0.379$. This is well below the upper $k_2$ limit from \\citet{Batyginetal09}.  As our analysis showed, the temperature and metallicity of the planet's envelope are important parameters whose knowledge would significantly increase the effectiveness of constraining the planet's interior with the Love number $k_2$. We therefore calculated a model series based on a theoretical model atmosphere from \\citet{Fortneyetal07}. Models of this kind have $T_\\mathrm{1\\,bar}=2080$\\,K and fall in our \\emph{new $k_2$ interval} if they have an outer isothermal layer with $3\\,\\mathrm{bar} < P_\\mathrm{ad} < 44\\,\\mathrm{bar}$. With this result we also estimated lower limits of the tidal $Q$ value of HAT-P-13b.  Based on our analysis and the applicability of the tidal fixed point theory we can rule out the majority of the calculated models. Assuming our new $k_2$-interval, we note that hot models with $T_{1\\,\\mathrm{kbar}}>6500$\\,K can be ruled out (compare Fig.~\\ref{fig:HAT-P-13b_k2Z}). On the other hand, the very cold models with $T_\\mathrm{1\\,kbar}<4000\\,$K are unlikely because they have a metallicity that is smaller than stellar. We favor models with $T_\\mathrm{1\\,bar}=2080$\\,K, an isothermal outer layer of 3 -- 44\\,bar, and an envelope metallicity of $1 < Z/Z_* < 11$ because they (i) are placed in our \\emph{new $k_2$ interval}, (ii) are consistent with the model atmospheres from \\citet{Fortneyetal07}, and (iii) have an envelope metallicity that is above the stellar metallicity. Assuming these conditions really apply for HAT-P-13b the planet would have a core mass of 27\\,\\ME~or less. Given the calculated tidal $Q$, if HAT-P-13b turns out to have a small core and an envelope enrichment similar to Jupiter, it could represent a Jupiter-like extrasolar planet.  By comparing with the previous results from \\citet{Batyginetal09} we have seen that the core EOS affects the core mass and Love number, resulting into a difference of about 28\\% in the prediction of a maximum possible core mass of HAT-P-13b.  As pointed out in \\S\\ref{subsec:e}, the eccentricity of the inner planet is a crucial parameter because it determines the Love number $k_2$. The value of the eccentricity is still uncertain. We analyzed the consequences of an $e_b = 0.021\\pm0.009$ \\citep{Bakosetal09} and $e_b = 0.0133\\pm0.0041$ \\citep{Winnetal10}.  Finally, despite the uncertainty of the inner planet's eccentricity it is still unclear whether the tidal fixed point theory from \\citet{Mardling07} can really be applied to HAT-P-13b. Further observations are necessary to prove the co-planarity and apsidal alignment configuration of the system. Also, TTV measurements could provide clues about the existence of other long-period companions in the system."
    },
    "1112/1112.5458_arXiv.txt": {
        "abstract": "We report the first detection of the intrinsic velocity dispersion of the Arches cluster - a young ($\\sim$~2 Myr), massive ($10^4 M_{\\odot}$) starburst cluster located only 26 pc in projection from the Galactic center. This was accomplished using proper motion measurements within the central $10'' \\times 10''$~of the cluster, obtained with the laser guide star adaptive optics system at Keck Observatory over a 3 year time baseline (2006-2009).  This uniform dataset results in proper motion measurements that are improved by a factor $\\sim$~5 over previous measurements from heterogeneous instruments.  By careful, simultaneous accounting of the cluster and field contaminant distributions as well as the possible sources of measurement uncertainties, we estimate the internal velocity dispersion to be $0.15 \\pm 0.01$~mas yr$^{-1}$, which corresponds to $5.4 \\pm 0.4$~km s$^{-1}$~at a distance of 8.4 kpc.  Projecting a simple model for the cluster onto the sky to compare with our proper motion dataset, in conjunction with surface density data, we estimate the total {\\it present-day}~mass of the cluster to be \\besttot. The mass in stars observed within a cylinder of radius $R$~(for comparison to photometric estimates) is found to be \\bestmass~at formal $3\\sigma$~confidence. This mass measurement is free from assumptions about the mass function of the cluster, and thus may be used to check mass estimates from photometry and simulation. Photometric mass estimates assuming an initially Salpeter mass function ($\\IMF = 1.35$, or $\\PDMF\\sim 1.0$~at present, where $dN/d(\\log M) \\propto M^{\\PDMF}$)~suggest a total cluster mass $M_{cl} \\sim (4-6) \\times 10^4 M_{\\odot}$~and projected mass ($\\sim 2 \\le M(R<0.4~{\\rm pc}) \\le 3$)~$\\times 10^4 M_{\\odot}$. Photometric mass estimates assuming a globally top-heavy or strongly truncated present-day mass function (PDMF, with $\\PDMF \\sim 0.6$) yield mass estimates closer to $\\MProjObs \\sim 1-1.2 \\times 10^4 M_{\\odot}$. Consequently, our results support a PDMF that is either top-heavy or truncated at low-mass, or both.  Collateral benefits of our data and analysis include: (i) cluster membership probabilities, which may be used to extract a clean cluster sample for future photometric work; (ii) a refined estimate of the bulk motion of the Arches cluster with respect to the field, which we find to be 172 $\\pm$~15 km s$^{-1}$, which is slightly slower than suggested by previous VLT-Keck measurements; and (iii) a velocity dispersion estimate for the field itself, which is likely dominated by the inner galactic bulge and the nuclear disk. ",
        "introduction": "The spectrum of masses produced during the star formation process (the Initial Mass Function, or IMF) is a key prediction of the star formation process as it indirectly links to the observable Present-Day Mass Function (PDMF) of the population \\citep[for example, see][for review]{millerscalo79, mckee07, bastian10}. Because star formation depends on collapse by self-gravity out of a turbulent medium threaded with a magnetic field, there is some expectation that the physical conditions in the parent cloud should affect the slope of the IMF, its minimum mass, or both \\citep[e.g.][]{morris93}. Numerical modeling provides some support for environment-dependent IMF variations, particularly in the unusual environment of the Galactic center \\citep[e.g.][]{bonnell04, klessen07, krumholz08, bonnell08}. However, the resulting IMF variations may be so small as to be observable only in extreme environments \\citep[e.g. ][]{elmegreen08}. There is some observational support for a varying IMF and turn-over mass in the extreme environments of the young ($\\sim$~few Myr), massive ($\\sim 10^4 M_{\\odot}$) starburst clusters NGC 3603 \\citep{harayama08} and for the stellar cluster at the Galactic center itself \\citep{bartko10}.  The young, massive cluster (YMC) known as the Arches Cluster \\citep[e.g.][]{nagata95, cotera96}\\footnote{Throughout this communication, ``the Arches'' refers to the star cluster, not the arched radio filaments \\citep{yusef84, morris89}, against which the cluster is projected, and with which it is physically associated \\citep{lang04}.} is a particularly well-studied example of an extreme environment for star formation. It is massive \\citep[Total mass $M_{cl}\\sim (2-7) \\times 10^{4} M_{\\odot}$;][]{figer99,figer02} dense \\citep[$\\rho_c \\sim 10^5 M_{\\odot}$~pc$^{-3}$][]{espinoza09} and young \\citep[$\\sim 2-2.5$~My;][]{najarro04, martins08}. It contains a substantial number of massive stars \\citep{serabyn98} which both contribute to and heat the surrounding medium \\citep[e.g.][]{figer02, yusef02, lang04}.  The Arches cluster is located only 26 pc in projection from the Galactic center~(hereafter the GC). It therefore likely formed in an environment characterised by high gas pressure and velocity dispersion in the parent cloud, and high ambient temperature, particularly when compared to the relatively more benign environment of NGC 3603. As these parameters are thought to impact the IMF \\citep{morris93, klessen07}, the Arches cluster is expected to be an excellent candidate for observing a non-canonical IMF, whether in its mass function exponent, low-mass turnover, a low-mass cut-off, or all three \\citep[e.g.][]{stolte02, klessen07}. It is also young enough that the most massive main-sequence stars are still present, making it one of the few clusters in which the upper mass-limit to the star formation process may be observationally tested \\citep{figer05, crowther10}. It has thus received substantial observational attention, with efforts focused particularly on estimates of its IMF. Indeed, the Arches was originally the prototypical object for a non-standard IMF \\citep{figer99}, with an observed present-day luminosity function indicating an overabundance of massive stars compared to the canonical Salpeter IMF (parameterized as $dN/d(\\log M) \\propto M^{-\\Gamma}$, with $\\Gamma = 1.35$; see \\citealt{bastian10} for a review).  However, a number of effects conspire to obscure the true IMF from observation, complicating the interpretation of the PDMF, and indeed the present consensus seems to be that the Arches began with an IMF that is consistent with the canonical Salpeter IMF found in most environments. Photometric efforts to chart the present-day luminosity function of the cluster suffer from two important limitations. Firstly, the observations are difficult; strong crowding and high, spatially-variable extinction are observed across the field of view, so that the photometric completeness is challenging to estimate for masses lower than a few $M_{\\odot}$. There is evidence for mass segregation in the cluster \\citep{figer99, stolte05}, seen as a steepening of the present-day luminosity function towards the cluster center, implying strong spatial selection effects when attempting to constrain the IMF.  Secondly, the relationship of the PDMF to the IMF is not trivial to evaluate. Stellar evolution must be taken into account when relating the PDMF to the IMF, requiring a prescription for mass-loss from high-mass objects \\citep[e.g.,][]{espinoza09}. In addition, mass segregation and tidal stripping are both likely to have been important for the evolution of the Arches. Present-day mass segregation need not be primordial, since the Arches is likely already in a post-collapse phase \\citep[e.g.][]{pz07, allison09}. The Arches cluster sits in a strong tidal field, such that as much as half of its stars may already have been stripped into the field over the $\\sim 2.5$~My of its history \\citep{kim00, pz02}. Mass segregation and tidal stripping together imply that the true IMF of the cluster may differ from the IMF drawn from the subset of stars that have remained within the Arches cluster to the present day. We review the literature mass determinations in Section \\ref{ss_masscomp}.  A {\\it kinematic} mass measurement provides a direct test of the PDMF of the Arches cluster, because its selection effects are somewhat less stringent. Stars below the typical photometric completeness limit of $\\sim 1-2 M_{\\odot}$~are observable through their contribution to the total cluster mass. \\citet{figer02}~were the first to attempt this, by estimating the radial velocity dispersion of a sample of emission-line stars and assuming the cluster is spherically symmetric and in virial equilibrium. However, the estimate is complicated by the difficulty of resolving the blended lines, their high width, and intrinsic line-profile variation among the sample, so that the resulting mass estimate is strongly dependent on atmosphere models. Mass estimates using the velocity dispersion derived from proper motions are independent of the details of the atmospheres of the tracer stars, and in principle allow for the mass distribution to be derived in a more assumption-free manner \\citep{lm89}.  The advent of adaptive optics on large telescopes in the near-infrared has enabled the measurement of precise proper motions of a large number of stars in the Galactic center clusters. In a pioneering proper motion study of the Arches, \\citet{stolte08} used one epoch each of VLT/NACO and Keck/NIRC2 separated by 4.3 years to measure the motion of the cluster. However, differential distortion between the cameras limited the proper motion precision to $\\sim 0.7$~mas yr$^{-1}$, somewhat too coarse to measure the internal velocity dispersion\\footnote{We use the term ``velocity dispersion'' to refer to both the dispersion in mas yr$^{-1}$~and km s$^{-1}$~throughout.}, for which the expected order of magnitude is about $\\sim 0.2$~mas yr$^{-1}$.  We have observed the central $10'' \\times 10''$~of the Arches across five epochs in three years (2006-2009) with a uniform observational setup (PI Morris). Using NIRC2 on Keck-2, behind the LGS Adaptive Optics facility \\citep{ghez05, wizinowich06}, these cross-instrument systematics encountered by \\citet{stolte08} are not present in our observations, and we are able to attain proper motion measurements with error lower than the expected velocity dispersion. We report here on our results, which provide the first kinematic mass estimate of the Arches cluster from proper motions.  This communication is organised as follows. Section \\ref{s_obs} describes the observations and positional measurement technique, while Section \\ref{s_analyse} describes the process of proper motion measurement and error assignment. Section \\ref{s_res} describes the techniques used to fit the cluster membership probabilities and kinematic parameters. Section \\ref{s_discuss} provides our mass measurement and new bulk motion measurement for the Arches, and briefly discusses the implications. ",
        "conclusions": "With uniform observational setup over a sufficient time baseline and careful accounting for a number of sources of proper motion error, we have measured the internal velocity dispersion of the Arches cluster for the first time, finding $\\sigma = 0.15 \\pm 0.01$~mas yr$^{-1}$, which corresponds to $5.4 \\pm 0.4$~km s$^{-1}$~at a distance of 8.4 kpc.  We have used this dispersion to test the photometric estimates of the present-day mass function (PDMF) of the Arches cluster. The total mass is likely in the range \\besttot, but this is only weakly constrained by kinematic data and is consistent with nearly all suggestions of total cluster mass from the modeling literature. The projected mass (i.e., mass contained within a cylinder of radius $R$~on the sky) is rather better constrained; we find \\bestmass~at formal $3\\sigma$~confidence. The upper bound of this range is $3\\sigma$~below the photometric estimate for \\MProjObs~estimated by \\citet{espinoza09}~under the assumptions of a non-top-heavy mass function without low-mass truncation. If a substantial contribution from massive binaries were unknowingly included in our measurement, this would strengthen our conclusion because the upper cluster mass bound would accordingly be reduced. This is the first mass estimate for the Arches based on proper motion velocity-dispersion.  We have also revised the bulk motion of the Arches slightly downward. Our updated motion of $172 \\pm 15$~km s$^{-1}$~is only slightly lower than the $212 \\pm 29$~km s$^{-1}$~determined previously (Stolte et al. 2008). Taken at face value, this supports the previous conclusion that the Arches Cluster is unlikely to pass within 10~pc of the GC.  Finally, we have provided the first estimate (to our knowledge) of the velocity dispersion of the Bulge along such a close sight-line to the Galactic center; this estimate is (103, 64)~$\\pm$~(7.7,5.0)~km s$^{-1}$, with the major axis coincident with the Galactic plane, to within the uncertainties.   \\begin{figure} \\begin{center} \\includegraphics[width=14cm]{f1.jpg} \\caption{NIRC2 $K'$~mosaic of the core field of the Arches in May 2009. This is our best map in terms of both resolution (51 mas) and sensitivity ($K'_{lim} = 20.59$~mag; Table 1). All stars on which we report in this paper fall within the field of view indicated here. The scale-bar is two arcseconds in length. Stars used as PSF reference-stars are indicated by circles. When stellar membership probabilities are reported, positions are reported as offsets from the reference star indicated by the square in this figure.} \\label{fig_findingchart} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=12cm]{f2.png} \\caption{Positional errors as measured for an example epoch (2008.50). Positions are those in the image-stack with centroiding errors assessed as the rms of measurements within an epoch (Section \\ref{ss_centroiding}). Top row: centroiding errors along detector-X and detector-Y (top-left and top-middle respectively), and the average of the two as a function of distance from the field center (top-right). Bottom row: alignment errors along X and Y (bottom-left and bottom-middle; Section \\ref{ss_align}). The magnitude histogram is given in the bottom-right panel.} \\label{fig_additive_det} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[height=120mm]{f3.jpg} \\caption{Five example positional time-series. Left columns show motion along X, right columns along Y. Object IDs and $K'$~magnitudes are indicated in the right top and bottom corners respectively. Note that each vertical axis is scaled to accommodate the motion of the star and is in units of pixels in the \\treference~reference frame. The best-fit straight line to the motions are indicated in each case, as are 1$\\sigma$~positional error curves. Object 154 is likely a field object, as indicated by its large proper motion relative to the reference frame.} \\label{f_example_curves} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=15cm]{f4.png} \\caption{The distribution of adopted proper motion precision (Section \\ref{ss_errors} and Table \\ref{tab_budget}), for all objects with five position-measurements. Outliers due to likely mismatches are indicated by squares and were removed from the analysis. An object qualifies as an outlier if the rms in either co-ordinate falls obviously outside the sequence defined by most of the points.} \\label{f_motion_errors} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[height=18cm]{f5.png} \\caption{Vector Point Diagrams for the overlapping magnitude-ranges of Section \\ref{ss_fit} and Table \\ref{tab_kinparams}. Shaded ellipses give the $1\\sigma$~contours for the two-dimensional gaussian components fit to the field and cluster components. Within each ellipse, the lines indicate the length and direction of the semimajor (thick red line) and semiminor (thin red line) axes.} \\label{f_VPD_compilation} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[height=10cm]{f6.jpg} \\caption{Observed velocity dispersion in each coordinate for stars with $14.0 \\le K' < 17.0$~compared to a gaussian of width equal to the mean measurement error over this range (Section \\ref{ss_errors} and Table \\ref{tab_budget}). Panels correspond to detector-X (Top) and detector-Y (bottom). This figure was constructed after removing likely field objects (Section \\ref{ss_membprob})} \\label{f_compare_erro} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=15cm]{f7.png} \\caption{Color-magnitude diagram (CMD) and Vector Point Diagram (VPD) for all objects with proper motion error $< 0.5$~mas yr$^{-1}$~and five epochs of measurement. The CMD presented here was constructed by matching $K'$~measurements to photometry taken in $H$-band in 2006 May with Keck-2/NIRC2-LGS \\citep{mccrady11}, which limits the depth in the CMD. Objects with $P_{\\rm{cluster}} > 0.995$~are shown in black, all other objects denoted with open circles. Red objects in the VPD correspond to the stars within the red dot-dashed box in the CMD, and represent well-measured objects with a possible $H-K'$~excess. Of these objects, those with $P_{\\rm cluster} > 0.995$~are shown with a red circle; their field counterparts are shown with red squares. See Section \\ref{ss_excess}~and \\citet{stolte10}~for more information on these objects.} \\label{fig_CMD} \\end{center} \\end{figure}  \\begin{figure} \\includegraphics[width=8cm]{f8a.png} \\includegraphics[width=8cm]{f8b.png} \\caption{Views of the $\\Delta \\chi^2_{full} < 7.82$~region when both kinematic and surface density data (for stars of mass $10 \\le M/M_{\\odot} \\le 30$; Espinoza et al. 2009) are included in the assessment. Axes are: $R_c, R_t, M_{cl}$, with total cluster mass $M_{cl}$~vertical in each case. Limits shown are: $0.05 \\le R_c \\le 0.8$~pc; $1.0 \\le R_t \\le 30$~pc; $0.5 \\le M_{cl} \\le 6.0 \\times 10^4 M_{\\odot}$.} \\label{fig_chibubble_mtot_10-30} \\end{figure}   \\begin{figure} \\includegraphics[width=8cm]{f9a.png} \\includegraphics[width=8cm]{f9b.png} \\caption{As Figure \\ref{fig_chibubble_mtot_10-30}, but with \\MProjObs~along the vertical axis. Limits shown are: $0.05 \\le R_c \\le 0.8$~pc; $1.0 \\le R_t \\le 30$~pc; $0.5 \\le \\MProjObs \\le 2.5 \\times 10^4 M_{\\odot}$.} \\label{fig_chibubble_mproj_10-30} \\end{figure}   \\begin{figure} \\includegraphics[width=16cm]{f10.png} \\caption{Radial profiles corresponding to parameter-sets within the $\\Delta \\chi^2_{full} < 7.82$~surface; both our own kinematic data and surface density (by number) \\SurfdensN~data are used (corresponding to stars of mass $10 \\le M/M_{\\odot} \\le 30$; Espinoza et al. 2009). Top-left and top-middle panels show radial and tangential velocity dispersions from proper motions (points) over the projected profiles corresponding to model parameters (lines). Top-right panel shows the \\SurfdensN~dataset with model predictions. Bottom-left and bottom-middle panels show the total mass within cylindrical radius $R$~on the sky, with $R=0.4$~pc indicated by the vertical dashed line. Bottom-right panel shows the mass enclosed within a sphere of radius $r$~pc from the cluster center. See also Table \\ref{table_massest_10-30_full}.} \\label{fig_radprofs_finer_10-30} \\end{figure}    \\begin{figure} \\begin{center} \\includegraphics[width=16cm]{f11.png} \\caption{Photometric mass estimates for the Arches cluster from the literature. Symbols give the directly-observed photometric mass (circles) and extrapolated mass (where reported; squares and pentagons) depending on the low-mass cut-off assumed. The citation for each estimate is shown inside the symbols (using the numbering of Table \\ref{table_all_mass_estimates}). The horizontal bands show our model-dependent mass estimate using our dispersion data (at 1$\\sigma$, 2$\\sigma$~and 3$\\sigma$), and the \\SurfdensN~dataset of \\citet{espinoza09}; see Section \\ref{ss_massmodel} for more detail on the mass modeling used. All masses are reported as \\MProjObs.} \\label{fig_literature_estimates} \\end{center} \\end{figure}    \\begin{table} \\small \\begin{center} \\begin{tabular}{lcccccccc} % Epoch & (t$_{int} \\times N_{coadd}$) & $N_{images}$ & $N_{use}$ & FWHM & Strehl & $N_{\\ast}$ & $N_{\\ast, {\\rm uncrowd}}$ & $K'_{lim}$ \\\\ % &         (s)                &             &           & (mas) &        &       &      & (mag)    \\\\ % \\hline 2006 May 21 &   3.00 $\\times$ 10 & 15 & 15 &  61.05 &  0.261 & 660 & 649 & 19.43 \\\\ % 2006 Jul 18 &   3.00 $\\times$ 10 & 52 & 38 &  56.95 &  0.349 & 657 & 642 & 19.89 \\\\ % 2008 May 13 &   3.00 $\\times$ 10 & 146 & 72 &  66.66 &  0.219 & 556 & 536 & 19.74 \\\\ % 2008 Jun 01 &   3.00 $\\times$ 10 & 89 & 83 &  54.96 &  0.373 & 845 & 810 & 20.42 \\\\ % 2009 May 02 &   2.80 $\\times$ 10 & 119 & 108 &  51.47 &  0.442 & 968 & 917 & 20.59 \\\\ % \\end{tabular} \\end{center} \\normalsize \\caption{Summary of observations. Reading left-right, the columns are: Epoch of observation, the total integration time for each image, the number of images observed, the number of images used, the median FWHM and Strehl ratio over the set of accepted images $N_{use}$, the number of stars measured within the mean image stack in each epoch; the number surviving the cut on proximity to a known neighbour, and finally the magnitude $K'_{lim}$~at which the cumulative distribution function of the observed $K'$~magnitudes reaches 90$\\%$~of the total number of stars in the sample at each epoch.} \\label{tab_obsns} \\end{table}  \\scriptsize \\begin{deluxetable}{lcccccccc} \\tablewidth{0pt} \\tablecaption{PSF stars. Reading left-right, columns are: sequential star number in the master table of membership probabilities, estimated brightness, and finally the position of the star expressed as an offset in arcseconds from the reference star along the (E-W) and (S-N) directions. See Figure 1 for the locations of these stars on the field of view.} \\tablehead{ Row & $K'$ & $\\Delta x$ & $\\Delta y$ \\\\ & (mag) & ($\"$) & ($\"$)  \\\\ } \\startdata 1 &   10.24 &    2.736 &   -3.943 \\\\ 2 &   10.48 &    2.063 &   -1.193 \\\\ 3 &   10.49 &    0.791 &    0.755 \\\\ 4 &   10.66 &    3.150 &   -2.899 \\\\ 5 &   11.08 &   -0.633 &   -4.252 \\\\ 7 &   11.22 &   -1.650 &    1.730 \\\\ 8 &   11.25 &    1.385 &   -2.334 \\\\ 16 &   12.18 &    1.012 &   -5.199 \\\\ 24 &   12.49 &  0 &  0 \\\\ 25 &   12.50 &    5.407 &   -0.218 \\\\ 32 &   12.88 &    2.499 &   -5.402 \\\\ \\enddata \\label{tab_psfstars} \\end{deluxetable} \\normalsize   \\begin{sidewaystable} \\centering \\scriptsize \\begin{tabular}{l|r||l|l|l|l|l||} \\multicolumn{2}{l||}{Parameter} & 2006.38 & 2006.54 & 2008.37 & 2008.50 & 2009.33 \\\\ \\hline \\hline \\multicolumn{2}{l||}{$N_{ref}$} & 238  & 239  & 235  & 241  & 233  \\\\ \\multicolumn{2}{l||}{$N_{<4}$} & 9     & 10    & 21    & 11    & 6     \\\\ \\hline \\hline $\\Delta$ & $x'$ & ~-83.94 $\\pm$~$6.03 \\times 10^{-3}$ & ~9.26 $\\pm$~$4.01 \\times 10^{-3}$ & ~4.16 $\\pm$~$3.71 \\times 10^{-3}$ & (-4.06 $\\pm$~17.716)$\\times 10^{-4}$ & ~6.53 $\\pm$~$3.52 \\times 10^{-3}$\\\\(pix)& $y'$ & ~-16.88 $\\pm$~$4.74 \\times 10^{-3}$ & ~5.69 $\\pm$~$4.39 \\times 10^{-3}$ & ~3.35 $\\pm$~$5.64 \\times 10^{-3}$ & (3.66 $\\pm$~25.421)$\\times 10^{-4}$ & ~3.51 $\\pm$~$3.66 \\times 10^{-3}$\\\\ \\hline $x$ & $x'$ & ~1.00005 $\\pm$~$2.11 \\times 10^{-5}$ & ~0.99994 $\\pm$~$9.12 \\times 10^{-6}$ & ~1.00025 $\\pm$~$1.34 \\times 10^{-5}$ & ~0.99999 $\\pm$~$6.18 \\times 10^{-6}$ & ~1.00002 $\\pm$~$8.24 \\times 10^{-6}$\\\\()& $y'$ & (6.64 $\\pm$~0.185)$\\times 10^{-4}$ & (7.73 $\\pm$~0.079)$\\times 10^{-4}$ & (1.33 $\\pm$~1.425)$\\times 10^{-5}$ & (2.32 $\\pm$~5.858)$\\times 10^{-6}$ & (2.75 $\\pm$~0.090)$\\times 10^{-4}$\\\\ \\hline $y$ & $x'$ & (-3.15 $\\pm$~0.177)$\\times 10^{-4}$ & (-3.53 $\\pm$~0.089)$\\times 10^{-4}$ & (1.23 $\\pm$~0.105)$\\times 10^{-4}$ & (-2.40 $\\pm$~4.766)$\\times 10^{-6}$ & (-2.54 $\\pm$~0.072)$\\times 10^{-4}$\\\\()& $y'$ & ~1.00046 $\\pm$~$1.97 \\times 10^{-5}$ & ~1.00028 $\\pm$~$9.62 \\times 10^{-6}$ & ~1.00014 $\\pm$~$1.34 \\times 10^{-5}$ & ~1.00000 $\\pm$~$6.66 \\times 10^{-6}$ & ~1.00009 $\\pm$~$7.99 \\times 10^{-6}$\\\\ \\hline $x^2$ & $x'$ & ~28.90 $\\pm$~5.946 & ~20.61 $\\pm$~3.463 & ~-18.24 $\\pm$~3.989 & ~-0.27 $\\pm$~1.986 & ~4.63 $\\pm$~2.818\\\\($\\times 10^{-8}$ pix$^{-1}$)& $y'$ & ~11.78 $\\pm$~4.751 & ~15.66 $\\pm$~3.266 & ~23.10 $\\pm$~4.404 & ~0.48 $\\pm$~1.937 & ~20.65 $\\pm$~2.938\\\\ \\hline $xy$ & $x'$ & ~1.80 $\\pm$~5.818 & ~20.39 $\\pm$~3.772 & ~58.96 $\\pm$~4.594 & ~-0.28 $\\pm$~2.040 & ~-0.58 $\\pm$~2.973\\\\($\\times 10^{-8}$ pix$^{-1}$)& $y'$ & ~-5.05 $\\pm$~6.935 & ~-10.19 $\\pm$~3.985 & ~-28.61 $\\pm$~5.203 & ~-1.52 $\\pm$~2.680 & ~-20.03 $\\pm$~3.335\\\\ \\hline $y^2$ & $x'$ & ~-3.04 $\\pm$~5.917 & ~5.55 $\\pm$~3.048 & ~1.09 $\\pm$~3.540 & ~1.71 $\\pm$~1.686 & ~-0.59 $\\pm$~2.959\\\\($\\times 10^{-8}$ pix$^{-1}$)& $y'$ & ~33.22 $\\pm$~6.228 & ~52.07 $\\pm$~3.664 & ~75.11 $\\pm$~5.451 & ~-1.13 $\\pm$~2.488 & ~-0.47 $\\pm$~2.633\\\\ \\hline \\hline \\multicolumn{2}{l||}{$1.0$-$M$} & ~-2.55 $\\pm$~0.150 & ~-1.14 $\\pm$~0.064 & ~-1.95 $\\pm$~0.095 & (1.97 $\\pm$~4.303)$\\times 10^{-2}$ & ~-0.58 $\\pm$~0.062\\\\ \\multicolumn{2}{l||}{($\\times 10^{-4}$ )} & & & & &\\\\ \\hline \\multicolumn{2}{l||}{$1.0$-$M_y/M_x$} & ~-4.08 $\\pm$~0.275 & ~-3.34 $\\pm$~0.135 & ~1.10 $\\pm$~0.189 & (-5.72 $\\pm$~9.564)$\\times 10^{-2}$ & ~-0.71 $\\pm$~0.104\\\\ \\multicolumn{2}{l||}{($\\times 10^{-4}$ )} & & & & &\\\\ \\hline \\multicolumn{2}{l||}{$\\theta_{rot}$} & ~-10.11 $\\pm$~0.250 & ~-11.62 $\\pm$~0.129 & ~1.13 $\\pm$~0.182 & (-4.88 $\\pm$~7.925)$\\times 10^{-2}$ & ~-5.46 $\\pm$~0.116\\\\ \\multicolumn{2}{l||}{($''$)} & & & & &\\\\ \\hline \\multicolumn{2}{l||}{$\\theta_{skew}$} & ~7.20 $\\pm$~0.553 & ~8.66 $\\pm$~0.234 & ~2.82 $\\pm$~0.364 & -0.002 $\\pm$~0.153 & ~0.43 $\\pm$~0.245\\\\ \\multicolumn{2}{l||}{($''$)} & & & & &\\\\ \\hline \\end{tabular} \\normalsize \\caption{Transformation parameters taking the starlist in each epoch into the reference frame $t_{ref}$. {\\it Top row}: Number of reference stars $N_{ref}$~used in the mapping, along with the number of reference stars $N_{<4}$~that appear in fewer than four epochs. {\\it Next six rows:} Coefficients of the polynomial fits $x' = f(x,y)$~and $y' = g(x,y)$~(top and bottom rows respectively in each pair). For reference, a quadratic term of size $10.0\\times10^{-8}$~pix$^{-1}$~would introduce displacement 0.25 mas at the edges of the detector, comparable to the centroiding error for bright ($K' < 16$)~objects (Table \\ref{tab_budget}). {\\it Bottom four rows:} The linear parts of the transformations re-expressed as a global scaling $M$, nonuniform magnification $M_y/M_x$, rotation $\\theta_{rot}$~and departure from perpendicular axes $\\theta_{skew}$. (Global shifts $\\Delta$~appear in the polynomial fits and are not repeated.)} \\label{tab_transform} \\end{sidewaystable}   \\begin{table} \\scriptsize \\begin{center} \\begin{tabular}{l|l||l|l|l|l|l||c||l||l||} \\multicolumn{2}{l||}{$K'$} & \\multicolumn{5}{|l||}{Centroiding,  Alignment (mas)} & Additive & Confusion & Motion \\\\ \\multicolumn{2}{l||}{}  & 2006.39  & 2006.54  & 2008.37  & 2008.50  & 2009.33  & (mas) & bias (mas) & (mas/y)   \\\\ \\hline \\hline \\hline \\multirow{2}{*}{$10-16$}  & $x$  & 0.25, 0.09 & 0.10, 0.05 & 0.25, 0.06 & 0.06, 0.04 & 0.08, 0.03 & 0.16 $\\pm$~0.02& 0.0 - 0.03 & {\\bf 0.076 } \\\\ & $y$  & 0.23, 0.08 & 0.07, 0.05 & 0.19, 0.08 & 0.11, 0.04 & 0.09, 0.05 & 0.15 $\\pm$~0.02 &   & {\\bf 0.074 } \\\\ \\hline \\multirow{2}{*}{$16-18$}  & $x$  & 0.41, 0.11 & 0.23, 0.05 & 0.40, 0.08 & 0.14, 0.05 & 0.14, 0.04 & 0.24 $\\pm$~0.02& 0.03 - 0.10 & {\\bf 0.130 } \\\\ & $y$  & 0.42, 0.10 & 0.20, 0.05 & 0.43, 0.09 & 0.17, 0.05 & 0.18, 0.05 & 0.30 $\\pm$~0.03 &   & {\\bf 0.153 } \\\\ \\hline \\multirow{2}{*}{$18-20$}  & $x$  & 1.10, 0.11 & 0.92, 0.05 & 1.03, 0.08 & 0.60, 0.05 & 0.59, 0.04 & 0.59 $\\pm$~0.06& 0.1- 1.0 & {\\bf 0.378 } \\\\ & $y$  & 1.38, 0.10 & 1.05, 0.05 & 1.35, 0.09 & 0.81, 0.05 & 0.77, 0.06 & 0.71 $\\pm$~0.08 &   & {\\bf 0.478 } \\\\ \\end{tabular} \\caption{Astrometric error budget. For each magnitude bin, the top (bottom) row gives errors in X (Y). For each star, centroiding, alignment and additive error describe random variation between epochs. The effect of confusion bias on motions depends on its variation between epochs; random variation is already included in the additive error, while linear trends masquerading as spurious motions are expected to be $\\lesssim 10\\%$~of the confusion bias across the epochs for all objects (Section \\ref{ss_additive}).} \\label{tab_budget} \\end{center} \\normalsize\\end{table}   \\scriptsize \\begin{deluxetable}{lcccccccc} \\tablewidth{0pt} \\tablecaption{Membership table for objects in the Arches Central field. Reading left-right, columns are: Sequential star number, estimated brightness, offset from reference star (E-W and S-N), proper motion and error, and the formal probability that the object is associated with the cluster and field, respectively.$^\\dagger$} \\tablehead{ Row & $K'$ & $\\Delta x$ & $\\Delta y$ & $\\mu_x$ & $\\mu_y$ & $P({\\rm cluster})$ & $P({\\rm field})$ \\\\ & (mag) & ($\"$) & ($\"$) & (mas yr$^{-1}$) & (mas yr$^{-1}$) &  &   \\\\ } \\startdata 1$^{\\ast}$ &   10.24 &    2.736 &   -3.943 &    0.14 $\\pm$  0.07 &    0.13 $\\pm$  0.08 &  0.999 &  $8.54 \\times 10^{-4}$ \\\\ 2$^{\\ast}$ &   10.48 &    2.063 &   -1.193 &    0.08 $\\pm$  0.07 &    0.08 $\\pm$  0.07 &  0.999 &  $5.85 \\times 10^{-4}$ \\\\ 3$^{\\ast}$ &   10.49 &    0.791 &    0.755 &    0.15 $\\pm$  0.07 &   -0.09 $\\pm$  0.07 &  0.999 &  $1.06 \\times 10^{-3}$ \\\\ 4$^{\\ast}$ &   10.66 &    3.150 &   -2.899 &   -0.01 $\\pm$  0.07 &    0.20 $\\pm$  0.07 &  0.999 &  $8.91 \\times 10^{-4}$ \\\\ 5$^{\\ast}$ &   11.08 &   -0.633 &   -4.252 &    0.03 $\\pm$  0.07 &   -0.11 $\\pm$  0.08 &  0.999 &  $6.86 \\times 10^{-4}$ \\\\ 6 &   11.16 &    4.603 &    1.092 &    0.00 $\\pm$  0.08 &    0.07 $\\pm$  0.08 &  1.000 &  $4.88 \\times 10^{-4}$ \\\\ 7$^{\\ast}$ &   11.22 &   -1.650 &    1.730 &    0.21 $\\pm$  0.09 &    0.05 $\\pm$  0.08 &  0.999 &  $1.11 \\times 10^{-3}$ \\\\ 8$^{\\ast}$ &   11.25 &    1.385 &   -2.334 &   -0.28 $\\pm$  0.07 &    0.06 $\\pm$  0.07 &  0.999 &  $1.34 \\times 10^{-3}$ \\\\ 9 &   11.63 &   -1.758 &   -1.287 &   -0.00 $\\pm$  0.07 &    0.05 $\\pm$  0.08 &  1.000 &  $4.72 \\times 10^{-4}$ \\\\ 10 &   11.67 &    2.038 &    0.445 &    0.05 $\\pm$  0.07 &    0.03 $\\pm$  0.07 &  0.999 &  $5.04 \\times 10^{-4}$ \\\\ 11 &   11.81 &   -2.337 &   -0.540 &   -0.28 $\\pm$  0.08 &   -0.20 $\\pm$  0.08 &  0.997 &  $2.59 \\times 10^{-3}$ \\\\ 12 &   11.88 &    5.528 &   -3.874 &   -0.04 $\\pm$  0.08 &    0.04 $\\pm$  0.11 &  1.000 &  $4.66 \\times 10^{-4}$ \\\\ 13 &   11.89 &    0.285 &   -1.191 &    0.04 $\\pm$  0.07 &    0.01 $\\pm$  0.07 &  1.000 &  $4.98 \\times 10^{-4}$ \\\\ 14 &   12.00 &   -0.158 &   -3.382 &   -0.06 $\\pm$  0.07 &   -0.08 $\\pm$  0.07 &  0.999 &  $5.54 \\times 10^{-4}$ \\\\ 15 &   12.18 &    5.362 &    1.667 &   -0.02 $\\pm$  0.08 &    0.09 $\\pm$  0.08 &  0.999 &  $5.02 \\times 10^{-4}$ \\\\ 16$^{\\ast}$ &   12.18 &    1.012 &   -5.199 &    0.02 $\\pm$  0.07 &   -0.13 $\\pm$  0.08 &  0.999 &  $7.41 \\times 10^{-4}$ \\\\ 17 &   12.19 &   -1.490 &    0.681 &    0.06 $\\pm$  0.08 &   -0.07 $\\pm$  0.07 &  0.999 &  $6.26 \\times 10^{-4}$ \\\\ \\enddata \\tablenotetext{\\ast}{PSF Star} \\tablenotetext{\\dagger}{This Table will be published in its entirety in the electronic edition of the Astrophysical Journal, A portion is shown here for guidance regarding its form and content. Until publication, an electronic copy of this table is available from the first author.} \\label{table_memprob} \\end{deluxetable} \\normalsize   \\begin{table} \\scriptsize \\begin{center} \\begin{tabular}{c||c|c|c|c|c} $K'$  & 14.0-16.0  & 15.0-17.0  & 16.0-18.0  & 17.0-19.0  & 18.0-20.0  \\\\ &  &  &  & \\\\ \\hline \\hline $N$   &   75  &  105  &  135  &  165  &  135  \\\\ &  &  &  & \\\\ \\hline $\\pi_{cl}$  &  0.80 $\\pm$~0.036  &  0.72 $\\pm$~0.035  &  0.72 $\\pm$~0.030  &  0.59 $\\pm$~0.029  &  0.52 $\\pm$~0.032  \\\\ &  &  &  & \\\\ \\hline $\\Delta \\mu$  &  4.28 $\\pm$~0.526  &  4.54 $\\pm$~0.430  &  5.15 $\\pm$~0.356  &  3.68 $\\pm$~0.293  &  3.51 $\\pm$~0.320  \\\\ (mas yr$^{-1}$)  &  &  &  & \\\\ \\hline $\\phi_{f}$  &  30.9 $\\pm$~4.8  &  30.5 $\\pm$~3.9  &  37.1 $\\pm$~3.1  &  36.6 $\\pm$~2.5  &  32.5 $\\pm$~2.7  \\\\ ($^o$)  &  &  &  & \\\\ \\hline $\\sigma_{a,f}$  &  2.21 $\\pm$~0.338  &  2.56 $\\pm$~0.280  &  2.87 $\\pm$~0.240  &  2.89 $\\pm$~0.198  &  2.80 $\\pm$~0.216  \\\\ (mas yr$^{-1}$)  &  &  &  & \\\\ \\hline $\\sigma_{b,f}$  &  1.50 $\\pm$~0.231  &  1.64 $\\pm$~0.193  &  1.85 $\\pm$~0.159  &  1.81 $\\pm$~0.130  &  1.64 $\\pm$~0.137  \\\\ (mas yr$^{-1}$)  &  &  &  & \\\\ \\hline $\\sigma_{a,cl}$  &  0.15 $\\pm$~0.013  &  0.17 $\\pm$~0.012  &  0.16 $\\pm$~0.014  &  0.24 $\\pm$~0.022  &  0.45 $\\pm$~0.034  \\\\ (mas yr$^{-1}$)  &  &  &  & \\\\ \\hline $\\sigma_{b,cl}$  &  0.12 $\\pm$~0.010  &  0.16 $\\pm$~0.010  &  0.16 $\\pm$~0.012  &  0.16 $\\pm$~0.019  &  0.17 $\\pm$~0.029  \\\\ (mas yr$^{-1}$)  &  &  &  & \\\\ \\hline $\\theta_{f}$  &  33.9 $\\pm$~17.3  &  27.8 $\\pm$~14.2  &  35.0 $\\pm$~11.8  &  28.5 $\\pm$~8.8  &  26.7 $\\pm$~10.3  \\\\ ($^o$)  &  &  &  & \\\\ \\hline $\\theta_{cl}$  &  70.2 $\\pm$~21.7  &  78.6 $\\pm$~30.3  &  67.5 $\\pm$~48.3  &  117.1 $\\pm$~63.5  &  114.0 $\\pm$~66.4  \\\\ ($^o$)  &  &  &  & \\\\ \\hline $\\sigma_{b,cl} / \\sigma_{a,cl}$  &  0.83 $\\pm$~0.086  &  0.91 $\\pm$~0.078  &  0.96 $\\pm$~0.069  &  0.69 $\\pm$~0.075  &  0.37 $\\pm$~0.085  \\\\ &  &  &  & \\\\ \\hline $\\sigma_{b,f} / \\sigma_{a,f}$ &  0.68 $\\pm$~0.130  &  0.64 $\\pm$~0.113  &  0.65 $\\pm$~0.099  &  0.63 $\\pm$~0.083  &  0.58 $\\pm$~0.090  \\\\ &  &  &  & \\\\ \\hline \\end{tabular} \\end{center} \\caption{Fitted kinematic parameters of cluster and field. For each magnitude range, rows give: the cluster fraction, the separation between cluster and field centers in the Vector Point Diagram (VPD), the orientation of the separation vector from the cluster center to the field center in the VPD, the semimajor and minor axes of the field component, the semimajor and minor axes of the cluster component, the orientation of the semimajor axis of the field component, the orientation of the semimajor axis of the cluster component, and finally the axis ratio (minor/major) of the cluster and field components. Errors are estimated from Monte Carlo simulations: populations in the VPD are simulated under the intrinsic kinematic parameters estimated from observation, perturbed by the measured errors for stars in each magnitude range, and re-fitted. Orientations are position-angles reported in degrees East of North.} \\label{tab_kinparams} \\normalsize \\end{table}   \\begin{table} \\begin{center} \\begin{tabular}{l|c||c|c||c|c} $K'$ & $N$ & $\\sigma_x$ & $\\sigma_y$ & $\\sigma_x$ & $\\sigma_y$ \\\\ & & (mas yr$^{-1}$) & (mas yr$^{-1}$) & (km s$^{-1})$ & (km s$^{-1}$) \\\\ \\hline 10.0-14.0 & 67 & 0.130 $\\pm$~0.017 & 0.123 $\\pm$~0.016  & 4.912 $\\pm$~0.639 & 4.680 $\\pm$~0.593 \\\\ 14.0-16.0 & 72 & 0.161 $\\pm$~0.019 & 0.129 $\\pm$~0.016  & 6.088 $\\pm$~0.739 & 4.878 $\\pm$~0.606 \\\\ 16.0-18.0 & 107 & 0.177 $\\pm$~0.027 & 0.180 $\\pm$~0.030  & 6.721 $\\pm$~1.034 & 6.839 $\\pm$~1.142 \\\\ 18.0-20.0 & 97 & 0.224 $\\pm$~0.039 & 0.148 $\\pm$~0.046  & 8.508 $\\pm$~1.498 & 5.629 $\\pm$~1.753 \\\\ \\end{tabular} \\end{center} \\caption{Arches velocity dispersion in each co-ordinate. Reading left-right, columns are: Magnitude range of interest, number of cluster stars in this magnitude range, intrinsic velocity dispersion and error in each coordinate, first in mas yr$^{-1}$~and then km s$^{-1}$~assuming the Arches is at 8.4~kpc.} \\label{tab_dispersions} \\end{table}   \\include{table8}  \\begin{table} \\begin{tabular}{r|c|c|c} $\\Delta \\chi^2_{full}$ & 3.50  & 7.82  & 13.93  \\\\ Confidence & 68\\% & 95\\% & 99.7\\% \\\\ & \"$1\\sigma$\" & \"$2\\sigma$\" & \"$3\\sigma$\" \\\\ \\hline $M(R< 0.40~{\\rm pc})$ & 0.69 - 1.10 & 0.62 - 1.20 & 0.55 - 1.30\\\\ ($10^4~M_{\\odot}$) &   &   &   \\\\ \\hline $M(r < 1.0~{\\rm pc})$ & 1.16 - 1.88 & 1.04 - 2.06 & 0.91 - 2.24\\\\ ($10^4~M_{\\odot}$) &   &   &   \\\\ \\hline $\\rho_0$ & 0.45 - 1.66 & 0.30 - 2.34 & 0.20 - 3.19\\\\ ($10^5 M_{\\odot}~{\\rm pc^{-3}}$) &   &   &   \\\\ \\hline $R_c$ & 0.18 - 0.31 & 0.15 - 0.44 & 0.13 - 0.80\\\\ (pc) &   &   &   \\\\ \\hline $R_t$ & 2.00 - 30.00 & 2.00 - 30.00 & 1.00 - 30.00\\\\ (pc) &   &   &   \\\\ \\hline $M_{cluster}$ & 1.64 - 4.29 & 1.45 - 4.86 & 1.07 - 5.62\\\\ ($10^4~M_{\\odot}$) &   &   &   \\\\ \\hline $1000 \\times \\Sigma_{N,0} / \\rho_0$ & 0.00 - 0.07 & 0.00 - 0.18 & 0.00 - 15.62 \\\\ (stars pc$^{-2} / M_{\\odot}~{\\rm pc}^{-3}$) &   &   &   \\\\ \\hline \\end{tabular} \\caption{Significance regions for isotropic King modeling of the Arches cluster. Ranges of each parameter corresponding to the stated significance level are given, when $R_c, R_t, M_{cluster}$~are all allowed to vary. The quantity $\\chi^2_{full}$~denotes the badness-of-fit when comparing model predictions to both the Arches kinematic dataset and the surface density dataset of Espinoza et al. (2009), over the mass range ($10 \\le M \\le 30$)~$M_{\\odot}$.} \\label{table_massest_10-30_full} \\end{table}   \\begin{sidewaystable} \\scriptsize \\begin{tabular}{lcccccccccl} Ref& $R_{in}$& $R_{out}$& $R_{ext}$&  $\\Gamma$& $M_{low}$& $M_{obs}$& $M_{calc}$& $\\delta (M_{obs})$& $\\delta(M_{calc})$& Notes \\\\ &      (pc) &    (pc)     &  (pc)    &           & $(M_{\\odot})$ &  $(\\times 10^4 M_{\\odot})$ &  $(\\times 10^4 M_{\\odot})$ &  $(\\times 10^4 M_{\\odot})$ & & \\\\ \\tableline \\tableline 1 & -   &  1.15 & -   & -     & -     & 0.08  &  0.24 &   -  &  -    &  Lower limit on total mass \\\\ 2 & -   &  0.35 & -   &  1.35 &  2.0  & 0.50  &  1.5  &  0.1 &  -    &  PDMF Salpeter; \\MProjRad \\\\ 2 & -   &  0.35 & -   &  1.35 &  0.1  & 0.50  &  6.0  &  0.1 &  -    &  PDMF Salpeter; \\MProjRad \\\\ 3 & 0.12 & 0.35 & -   &  0.6  &\t  1.0 & 0.51  &  1.08 &    - &  -    &  PDMF Top-heavy; \\MProjRad \\\\ 3 & 0.12 & 0.35 & -   &  0.6  &\t  0.1 & 0.51  &  1.20 &    - &  -    &  PDMF Top-heavy; \\MProjRad \\\\ 4 & -    & 0.40 & -   &  -    &   2.0 & 0.63  &  -    &\t   - &  -    &  Rough limit on \\MProjRad~reported \\\\ 5 & -\t & 0.23 & -   &  -    &  -    & 7.0   &  -    &    - &  -    &  Upper limit on \\MProjRad~from radial velocities \\\\ 6 & - \t & 0.40 & -   &  1.1  &\t  1.0 & 0.557 & 2.0 &    - &  0.6 &  low-mass truncation; \\MProjRad \\\\ 6 & -\t & 0.40 & -   &  1.35 &  0.08 & 0.557 & 3.1  &    - &  0.6 &  Kroupa PDMF; no low-mass truncation; \\MProjRad \\\\ \\tableline 7 & -   & 2.50 & -    & 0.5   &\t 1.0  & -     &  1.60 &    - & \t-    & N-body; IMF top-heavy; $M_0$~reported \\\\ 7 & -   & 2.50 & -    & 0.75  & 1.0   & -     & 2.00  &    - &\t-    &\t  ``        `` \\\\ 7 & -  \t& 2.50 & -    & 1.0   &\t 1.0  & -     & 2.80  &    - &\t-    &\t  ``        `` \\\\ 8 & -\t& 1.26 & 3.0  & 2.8   & -     & -     & 4.00  &    - &\t-    & N-body; Multi-component IMF \\\\ 9 & -\t& 0.35 & 2.5  & 0.9   &\t 1.3  & -     & 4.00  &    - &\t-    & N-body; Representative $M_0$~reported in Figure caption \\\\ 10 & -\t& 0.35 & 2.5  & 0.9   &\t 0.1  & -     & 14.50 &    - &\t-    & Turbulent-fragmentation calculation \\\\ 11 & -\t& 0.35 & 2.4  & 1.1  & 0.9    &\t-     & 5.90  &    - &\t-    & N-body; Observations reported in Kim et al. (2006) \\\\ 12  & -\t& 0.40 & -    & 1.35 & 0.5    & -     & 1.80  &    - &\t-    & N-body, Salpeter IMF;  Present-day \\MProjObs \\\\ 12  & -\t& 0.40 & 1.0  & 1.35 & 0.5    & -     & 3.60  &    - &\t-    & Present-day simulated mass within projected radius $R = 1.0$~pc \\\\ 12  & -\t& 0.40 & 2.8  & 1.35 & 0.5    & -     & 4.90  &    - &\t0.8  & Total initial cluster mass $M_0$~with low-mass truncation \\\\ 12  & -\t& 0.40 & 2.8  & 1.35 & 1.0    & -     & 3.60  &    - &\t0.6  & Total initial cluster mass $M_0$~with moderate-mass truncation \\\\ 12  & -\t& 0.40 & 2.8  & 1.35 & 4.0    & -     & 1.90  &    - &\t0.3  & Total initial cluster mass $M_0$~with moderate-mass truncation \\\\ \\tableline \\tableline \\end{tabular} \\caption{All Arches literature mass estimates of which the authors are aware, current as of September 2011. Observational estimates are listed first, followed by cluster mass estimates from models. Reading left-right, columns give: 1: Reference cited. 2,3: Inner and outer radii within which stars were observed $R_{in}, R_{out}$. 4: Radius to which mass function has been extrapolated $R_{ext}$. 5: Mass function slope $\\Gamma$. 6, lower stellar mass used for IMF. 7:Total mass of stars directly observed, $M_{obs}$. 8: Extrapolated mass $M_{calc}$~where appropriate. 9,10: errors (where given) in the observed and extrapolated masses; 11. Brief description. References are: 1. \\citet{cotera96}; 2. \\citet{serabyn98}; 3. \\citet{figer99}; 4. \\citet{stolte02}; 5. \\citet{figer02}; 6. \\citet{espinoza09}. 7. \\citet{kim00}; 8. \\citet{pz02}; 9. \\citet{kim06}; 10. \\citet{dib07}; 11. \\citet{chatterjee09}; 12. \\citet{harfst10}} \\label{table_all_mass_estimates} \\end{sidewaystable}"
    },
    "1112/1112.2752_arXiv.txt": {
        "abstract": "The spatial distribution of matter in clusters of galaxies is mainly determined by the dominant dark matter component, however, physical processes involving baryonic matter are able to modify it significantly. We analyse a set of 500 pc resolution cosmological simulations of a cluster of galaxies with mass comparable to Virgo, performed with the AMR code RAMSES. We compare the mass density profiles of the dark, stellar and gaseous matter components of the cluster that result from different assumptions for the subgrid baryonic physics and galaxy formation processes. First, the prediction of a gravity only N-body simulation is compared to that of a hydrodynamical simulation with standard galaxy formation recipes, then all results are compared to a hydrodynamical simulation which includes thermal AGN feedback from Super Massive Black Holes (SMBH). We find the usual effects of overcooling and adiabatic contraction in the run with standard galaxy formation physics, but very different results are found when implementing SMBHs and AGN feedback. Star formation is strongly quenched, producing lower stellar densities throughout the cluster, and much less cold gas is available for star formation at low redshifts. At redshift $z=0$ we find a flat density core of radius 10 kpc in both of the dark and stellar matter density  profiles. We speculate on the possible formation mechanisms able to produce such cores and we conclude that they can be produced through the coupling of different processes: (I) dynamical friction from the decay of black hole orbits during galaxy mergers; (II) AGN driven gas outflows producing fluctuations of the gravitational potential causing the removal of collisionless matter from the central region of the cluster; (III) adiabatic expansion in response to the slow expulsion of gas from the central region of the cluster during the quiescent mode of AGN activity. ",
        "introduction": "Clusters of galaxies are the most massive virialised structures observed in the Universe and provide a wonderful laboratory to test astrophysical theories. In the $\\Lambda$CDM cosmological scenario, clusters are assembled via a hierarchy of mergers of less massive structures like galaxies and groups of galaxies. Many physical processes play a role during the formation of a cluster. When satellite galaxies are accreted into a cluster, their properties can be changed by tidal and ram pressure stripping, leading to the formation of a wide variety of galaxy morphologies. Furthermore, clusters are known to be dark matter dominated structures, with most of the baryonic matter residing in a hot diffuse X-ray emitting gaseous phase, the Intracluster Medium (ICM). The stellar mass is less significant and mainly contained in the massive central elliptical galaxy.  Since they are dominated by dark matter, this mass component determines the global properties of the mass distribution in the cluster. However, from the theoretical side, it is well known that baryonic processes can produce significant differences in the distribution of matter in collapsed structures with respect to models including only collisionless cold dark matter. For example, baryons are known to condense the centre of dark matter halos due to dissipative processes, producing adiabatic contraction of the total mass distribution  \\citep{Gnedin:2004p569}. Several models including baryonic physics have been invoked to solve the so--called cusp/core problem in the $\\Lambda$CDM cosmological framework, i.e. the discrepancy between the centrally cuspy dark matter profiles observed in dark matter halos in numerical N-body simulations and the centrally cored dark matter profiles inferred by observations in dwarf galaxies and low surface brightness galaxies (see the recent review by \\citealp{2010AdAst2010E...5D} and references therein). The study of baryon physics induced modifications in the mass distribution in collapsed structures, and clusters of galaxies in particular, is still a field with many open issues.  The present paper is dedicated to the study of the effects of baryonic processes on the mass distribution in clusters of galaxies. In particular, we use a set of cosmological hydrodynamical simulations performed using the AMR code RAMSES \\citep{Teyssier:2002p451} to study the effect of different models for baryons and galaxy formation physics on the mass density profile of a cluster of galaxies comparable to the Virgo cluster. This work can be considered as an extension of the analysis performed by \\cite{2011MNRAS.414..195T} on the same simulations, and is complementary to the analysis presented by \\cite{2012MNRAS.420.2859M}. Here, we focus on the peculiar properties produced in the mass density profile when including Super Massive Black Holes (SMBHs) and the related AGN feedback in the recipes for galaxy formation physics. We stress that AGN feedback was initially introduced to solve the so--called \"overcooling problem\", namely the fact that too much stellar mass is produced in massive collapsed structure in hydrodynamical simulations with respect to what is observed in the real Universe \\citep{Borgani:2009p728}. The strong quenching of star formation produced by processes that couple AGN activity with the gas is expected to improve the match between simulations and observations \\citep{Tabor:1993p1080, Ciotti:1997p1087, Silk:1998p941}. Strong evidence for the existence of AGN feedback is provided by observations of X-ray cavities and radio blobs in galaxy clusters, typically interpreted as buoyantly rising bubbles of high entropy material injected in the central region of clusters by jets of relativistic particles. In this paper, we show that by including SMBH physics and AGN feedback it is possible to obtain interesting predictions on the modifications they can induce on the mass distribution in massive dark matter halos.  The paper is organized as follows: the first section is dedicated to the numerical methods and the galaxy formation recipes adopted for our simulations; we show our main results and provide our interpretation in the second section; the last section is left for a short summary of our results and to a discussion. ",
        "conclusions": "\\label{sec:summary}  We used the results of a set of cosmological simulations to study the effect of baryons in massive dark matter halos, focusing on phenomena involving SMBHs and AGN feedback. We use the zoom-in technique to simulate the formation of a galaxy cluster of mass comparable to Virgo at high spatial and mass resolution. In this work we analyse three different simulations: the first is a gravity only run with no baryons (DMO), the second is a hydrodynamical run with standard galaxy formation physics (AGN-OFF), the third is a hydrodynamical run with standard galaxy formation physics plus SMBHs and AGN feedeback (AGN-ON). We adopt a modified version of the \\cite{Booth:2009p501} model to implement thermal AGN feedback in the AGN-ON run. This model allows to reproduce the self-regulated growth of black holes throughout the cosmic ages and it is tuned to reproduce the observed $M_{\\rm BH}-\\sigma$ relation.  Our analysis focuses on the evolution of the 3D mass density profiles of the cluster from redshift $z=5$ to redshift $z=0$. We find a number of interesting differences between our models that highlight the importance of accounting for all the physical processes taking place during the formation and evolution of the most massive bound structures in the Universe. Our results can be summarized in the following points:  \\begin{itemize} \\item At redshift $z=0$ the dark matter density profiles of the three simulations are consistent with each other at distances $r>10$ kpc from the centre, but they differ significantly in the central region. The DMO profile is consistent with a standard Einasto model. The AGN-OFF profile shows a mass excess with respect to an Einasto fit due to adiabatic contraction of the baryons at the centre of the cluster. The AGN-ON profile presents a flat core of radius 10 kpc, i.e. a mass deficiency with respect to an Einasto fit. The dark matter core forms gradually from $z=5$ to $z=0$. We stress that dark matter cores in dark matter profiles have been claimed to be observed in several clusters \\citep{2004ApJ...604...88S, 2008ApJ...674..711S, 2009ApJ...706.1078N, 2011ApJ...728L..39N, 2011A&A...531A.119R}.  \\item The stellar density profile at $z=0$ is also very different between the AGN-OFF and AGN-ON runs. The AGN-OFF profile is very peaked at the centre, while the AGN-ON profile has a core of the same size as that of the dark one. Unlike the dark core, the stellar one forms between $z=1$ and $z=0$. Note that, cores in the surface brightness profiles of massive ellipticals and cluster central galaxies have been observed by several authors \\citep{1999ASPC..182..124K, 2000ApJS..128...85Q, 2003AJ....125..478L, 2004AJ....127.1917T, 2005AJ....129.2138L, 2007ApJ...671.1456C,  2009ApJS..182..216K, 2011arXiv1108.0997G}, so our model is not in disagreement with observations and it {\\it predicts} the existence of dark matter cores associated to stellar cores in massive ellipticals. Furthermore, in the AGN-OFF run the stellar density is larger than in the AGN-ON run at all radii. This is a clear effect of overcooling of gas leading to enhanced (and too efficient) star formation. In the AGN-ON run, AGN feedback strongly quenches star formation and the overcooling problem is avoided \\citep{2011MNRAS.414..195T,2012MNRAS.420.2859M}. The result is also that less cool gas is available at the centre of the AGN-ON cluster with respect to the AGN-OFF run: gas is heated by AGN feedback and carried away via convective motions and slow adiabatic expansion.  \\item The core in the dark matter end stellar mass density profiles is a peculiar feature of our simulation that includes the physics associated with active galactic nuclei. In this paper, we suggest that the coupling of several mechanisms is responsible for their formation. First, SMBHs transfer part of their orbital energy to collisionless matter via dynamical friction during dry galaxy mergers \\citep{2010ApJ...725.1707G}, especially at redshift $z<1$. Second, AGN driven gas outflows modify the gravitational potential in regions close to SMBHs with resulting ejection of collisionless matter from the central region of the cluster; subsequent gas outflows followed by the central 'revirialisation' of the central material are expected to produce cores \\citep{2011arXiv1106.0499P}; due to the stronger AGN feedback at $z>1$, this mechanism is more effective at high redshift, but preserves an important role at low redshift. Third, the central hot gas slowly cools radiatively, falling onto the SMBHs in convective flows and is subsequently ejected impulsively; the slow loss of mass from the central region will produce an expansion of the inner mass distribution. \\end{itemize}  Despite assertions that AGN feedback can affect central cluster dynamics, in fact our results seem to be complementary to the those obtained by other authors for less massive dark matter halos \\citep{Governato:2010p1442, 2011arXiv1106.0499P, 2011arXiv1111.5620M}.  Related processes involving baryons seem to be active at the low and high mass end of the galaxy mass function, especially mechanisms involving gas outflows produced by feedback.  Given the fact that we perform the analysis on zoom simulations of one cluster, we need to stress that our result needs support from a suite of dedicated cosmological simulation aimed at exploring the effect of SMBHs, AGN feedback and baryon physics in general on the most massive clusters in the universe. Furthermore, a number of numerical experiments is also needed to explore in detail the core formation processes that have been considered in our discussion, with particular attention on their coupling. {The role of baryon mass outflows in galaxy clusters has been recently studied by \\cite{2012arXiv1202.1527R} using collisionless matter simulations where gas is modeled as a time varying external contribution to the gravitational potential; their results provide important support to ours.} To conclude, we are convinced that our results are robust enough to assert that the implementation of AGN feedback and SMBHs in cosmological hydrodynamical simulations is an important ingredient for modeling massive clusters of galaxies."
    },
    "1112/1112.4338_arXiv.txt": {
        "abstract": "Stellar wind-emission features in the spectrum of eta Carinae have decreased by factors of 1.5--3 relative to the continuum within the last 10 years. We investigate a large data set from several instruments (STIS, GMOS, UVES) obtained between 1998 and 2011 and we analyze the progression of spectral changes in the direct view of the star, in the reflected polar-on spectra at FOS4, and at the Weigelt knots. We find that the spectral changes occurred gradually on a time scale of about 10 years and that they are dependent on the viewing angle. The line strengths declined most in our direct view of the star. About a decade ago, broad stellar wind-emission features were much stronger in our line-of-sight view of the star than at FOS4. After the 2009 event, the wind-emission line strengths are now very similar at both locations. High-excitation \\ion{He}{1} and \\ion{N}{2} absorption lines in direct view of the star strengthened gradually. The terminal velocity of Balmer P Cyg absorption lines now appears to be less latitude-dependent and the absorption strength may have weakened at FOS4. Latitude-dependent alterations in the mass-loss rate and the ionization structure of eta Carinae's wind are likely explanations for the observed spectral changes. ",
        "introduction": "\\label{sec:intro}  Eta Carinae, one of the most massive and most luminous stars in our Galaxy, is famous for its Great Eruption about 170 years ago. Its recovery has been unsteady with unexplained photometric and spectral changes in the 1890s and 1940s (\\citealt{2008AJ....135.1249H}, and references therein). The spectral changes described in this paper may represent another rapid step in \\ec's recovery from its Great Eruption.  Eta Car has a complex spectroscopic cycle, most likely regulated by a companion star in an eccentric orbit (\\citealt{1997NewA....2..107D}, and many references in \\citealt{2005ASPC..332.....H} and \\citealt{2012eta}). So-called spectroscopic events occur every 5.54 years since 1948 \\citep{2001MNRAS.322..741F,1996ApJ...460L..49D,2008MNRAS.384.1649D}. The events are characterized by drastic changes in \\ec's spectrum and photometry, e.g., high-excitation emission lines disappear for a few months (e.g., \\citealt{1953ApJ...118..234G,1984A&A...137...79Z}) and light curves at all wavelength regions show significant variations (e.g., \\citealt{1994MNRAS.270..364W,1997Natur.390..587C,2001MNRAS.322..741F,2006JAD....12....3V,2009A&A...493.1093F,2004AJ....127.2352M}).   In a previous paper \\citep{2010ApJ...717L..22M} we compared spectra at corresponding phases of successive spectroscopic cycles and found dramatic changes in observations after the 2009 event.\\footnote{ We define ``phase'' by $P = 2023.0$ days and $t_0 =$ MJD 50814.0 = J1998.00, consistent with the Eta Carinae Treasury Program Archive (http://etacar.umn.edu/). Phases 0.00, 1.00, and 2.00 mark the 1998.0, 2003.5, and 2009.0 spectroscopic events. } Major stellar-wind emission features in the spectrum of \\ec\\ had decreased by factors of order 2 relative to the continuum within 10 years and helium P Cyg absorption had become stronger. Most of the broad emission lines in \\ec's spectrum originate in the primary star's wind, see many papers and refs.\\ in \\citet{2005ASPC..332.....H}, and the simplest explanation for the observed spectral changes is a decrease in \\ec's wind density,   by a factor of 2 or more. The early exit from \\ec's 2009 X-ray minimum and the observed decrease of the 2--10 keV photons over the last two cycles are consistent with this interpretation \\citep{2009ApJ...701L..59K,2010ApJ...725.1528C,2011ApJ...740...80M}.  In this paper we analyze spectra obtained between 1998 and 2011 with several instruments to investigate in detail spectral changes in \\ec's wind.  We are not concerned here with the temporary spectral changes observed during the events -- the spectral changes discussed are of secular nature. In \\citet{2010ApJ...717L..22M} we noted only a few examples;  here we explore a wider range of effects, and whether or not they have developed gradually as opposed to sporadically. Section {\\ref{sec:obs}} describes the observations. In Section {\\ref{sec:change}} we confirm the observations made by \\citet{2010ApJ...717L..22M} and show that the broad stellar wind features were still weak in {\\it HST\\/} STIS data obtained several months after our initial discovery in 2010 March data. The temporal progression of spectral changes and the dependence on the viewing direction is discussed. High-excitation emission and continuum from the nearby Weigelt knots, which are thought to be photoionized by a hot companion star, reveal additional information. In Section {\\ref{sec:massloss}} we discuss the implications of these observations and estimate the decrease in mass-loss rate over the last 10 years. In Section {\\ref{sec:discussion}} we give a short summary. ",
        "conclusions": "\\label{sec:discussion}  In this paper we analyzed spectral data obtained with several instruments between 1998 and 2012. We confirmed the spectral changes in the wind emission lines first reported  by \\citet{2010ApJ...717L..22M};  {\\it HST} STIS spectra obtained in 2010 August, $\\sim 170$ days after the first discovery, are comparable to the observations in 2010 March. Furthermore, we analyzed the long-term development of spectral changes in our direct line-of-sight view of  the star, at FOS4, and the Weigelt knots.  Eta Car's recent spectral changes involve both emission and absorption lines: \\begin{enumerate} \\item Broad stellar wind-emission features in our line-of-sight to the star have decreased by factors of 1.5--3 relative to the continuum within the last 10 years. These changes occurred gradually and are dependent on the viewing angle; spectra at higher stellar latitudes and from the outlying ejecta show smaller changes. The simplest explanation  is a decrease in $\\eta$ Car's primary wind density. However, the decrease in wind density appears to be latitude dependent, with emission features showing much less change at higher latitudes. After the 2009 event, emission line strengths are now very similar in our direct line-of-sight view and in the reflected polar-on spectrum at FOS4 suggesting a more spherical wind and/or a more uniform distribution of circumstellar extinction. \\item High-excitation \\ion{He}{1} and \\ion{N}{2} absorption lines strengthened gradually over the last decade indicating a change in \\ec's wind ionization structure. Hydrogen P Cyg absorption at FOS4 might have weakened after the 2009 event.  The terminal velocity of hydrogen P Cyg lines was found to be similar at all stellar latitudes. Those findings provide additional clues for a more spherical wind. \\end{enumerate}  The observational results presented here are difficult to reconcile with a decrease in mass-loss rate primarily at lower stellar latitudes since it is generally assumed that \\ec's wind had higher densities at the poles \\citep{2003ApJ...586..432S}. Our observations may be more readily reconciled with alternative explanations for latitude-dependent spectral features, such as a complex ionization structure of  \\ec's wind modulated by the secondary star's UV radiation \\citep{2010AJ....139.1534R} or the presence of a wind cavity in the primary wind caused by the secondary star \\citep{2010ApJ...716L.223G,2011arXiv1111.2280M}.    Using H$\\alpha$ emission and the method by \\citet{1988ApJ...326..356L} we found that \\ec's mass-loss rate decreased by a factor of 2--3 between 1999 and 2010. A decrease in mass-loss rate on the order of 2--3 is consistent with changes in the  X-ray light curve \\citep{2009ApJ...701L..59K,2010ApJ...725.1528C}. We did not attempt to derive the absolute value with any accuracy because there are too many unknown factors, such as latitudinal dependence and clumping of the wind. New theoretical models updating \\cite{2001ApJ...553..837H} are needed.  Observations in 2012 and 2013 will be extremely valuable to further analyze the nature of the spectral changes in \\ec's wind. It is of great importance to monitor the star consistently since spectral changes may occur on time scales of only weeks to months. For the long-term recovery of \\ec\\ it is important to investigate if the wind will further decline or if it will stabilize or even recover to its former strength. But by mid-2013, the onset of the  next event will dominate the spectrum, so observations in 2012 are needed. The last three events all differed from each other and considering the long-term spectral changes described in this paper we can expect many interesting new results from \\ec's 2014.5 event.    {\\it Acknowledgement} We thank the staff and observers of the Gemini-South Observatory in La Serena for their help in preparing and conducting the observations, and Beth Perriello at STScI for assistance with {\\it HST\\/} observing plans. We also thank Otmar Stahl and Kerstin Weis for their effort in planning and obtaining the UVES spectra. AM was co-funded under the Marie Curie Actions of the European Commission (FP7-COFUND). MTR received partial support from Center for Astrophysics FONDAP and PB06 CATA (CONICYT)."
    },
    "1112/1112.3342_arXiv.txt": {
        "abstract": "We present results from a suite of axisymmetric, core-collapse supernova simulations in which hydrodynamic recoil from an asymmetric explosion produces large proto-neutron star (PNS) velocities.  We use the adaptive-mesh refinement code CASTRO to self-consistently follow core-collapse, the formation of the PNS and its subsequent acceleration.  We obtain recoil velocities of up to 620 km\\,s$^{-1}$ at $\\sim$1 s after bounce.  These velocities are consistent with the observed distribution of pulsar kicks and with PNS velocities obtained in other theoretical calculations.  Our PNSs are still accelerating at several hundred km\\,s$^{-1}$ at the end of our calculations, suggesting that even the highest velocity pulsars may be explained by hydrodynamic recoil in generic, core-collapse supernovae. ",
        "introduction": "At birth, pulsars achieve velocities well above those of their progenitor population \\citep{Gunn:1970ys,Lyne:1994rr}.  These pulsar ``kicks\" typically range from 200 km\\,s$^{-1}$ to 400 $\\rm km$\\,$\\rm s^{-1}$, with the fastest neutron stars achieving velocities near, or in excess of, 1000 km\\,s$^{-1}$ \\citep{Lyne:1994rr,Chatterjee:2005uq,Hobbs:2005fk}.  The current distribution of observed pulsar velocities is Maxwellian, hinting at a common acceleration mechanism \\citep{Hansen:1997dq,Hobbs:2005fk,Zou:2005yq,Faucher-Giguere:2006lr}.  Many scenarios have been proposed for the origin of pulsar recoil and neutron star kicks.  Popular mechanisms often require strongly magnetized systems, exotic neutrino physics, and/or rapid rotation to produce substantial kicks.  For example, in the presence of strong magnetic fields, neutrino-matter interactions can generate neutron star velocities on the order of a few hundred km\\,s$^{-1}$ by producing $\\sim$1\\% dipole asymmetries \\citep{Lai:1998tg,Nardi:2001hc,Lai:2001ij,Kusenko:1999kx,Lambiase:2005yq,Barkovich:2004vn,Fuller:2003yq,Kishimoto:2011fk}. Many of these scenarios require magnetic fields in the magnetar range (i.e.~10$^{14-16}$ G) and may not produce substantial kicks in typical core-collapse supernovae.  Other scenarios involve jet/counter-jet misalignment launched near the proto-neutron star (PNS).  In such situations, the jets accompany an associated gamma-ray burst (GRB) or form through magneto-rotational processes during core collapse \\citep{Cen:1998oq,Khokhlov:1999qy,Sawai:2008kx,Papish:2011qy}.  These scenarios require rapid rotation and therefore, may only manifest in a small subset of core-collapse events.  If neutron star kicks are a generic outcome of core collapse, then a natural explanation is recoil during an asymmetric supernova explosion.  In the current, most sophisticated simulations, the bounce shock, launched when the equation of state stiffens at nuclear densities, stalls due to thermal energy losses from neutrino emission and dissociation of nuclei into nucleons.  The stalled shock itself is subject to hydrodynamic and neutrino-driven instabilities, which manifest as prominent low-order $\\ell$-mode oscillations in axisymmetric simulations of non-rotating progenitors \\citep{Blondin:2003ul,Scheck:2004rt,Buras:2006qf,Scheck:2006vn,Burrows:2007eu,Blondin:2007mz,Ott:2008cr,Fernandez:2009uq,Nordhaus:2010fr,Nordhaus:2010kx,Fernandez:2010ly,Brandt:2011fj}. At the onset of shock revival, the PNS may recoil if large-scale asymmetries are present during the ensuing supernova explosion.  While the mechanism by which core-collapse supernova progenitors explode is not fully understood, the most probable scenario involves absorption of neutrinos in the post-shock ``gain region\" \\citep{Bethe:1985fr} and likely requires the development of multi-dimensional instabilities in fully three-dimensional radiation-hydrodynamic simulations to succeed \\citep{Nordhaus:2010fr}.  Nonetheless, recoil from an asymmetric neutrino-driven explosion presents a natural mechanism by which PNSs achieve high velocities \\citep{Burrows:1996th,Scheck:2004rt,Scheck:2006vn,Nordhaus:2010kx,Wongwathanarat:2010yq} and appears to be supported by recent X-ray observations of the Cassiopeia A supernova remnant \\citep{Hwang:2011lr}.  Computational studies of recoil require multi-dimensional, radiation-hydrodynamics calculations which start at the onset of collapse and follow the dynamics self-consistently.  This includes the formation of the PNS, the explosion, the propagation of the shock front through the stellar envelope and eventually, decoupling of the PNS from the surrounding material.  Such an approach is computationally challenging and as such, various techniques have been adopted to make the computations tractable.  One popular approach is to commence the calculations after bounce by mapping spherically symmetric solutions onto a multi-dimensional grid and excising the PNS from the computational domain \\citep{Scheck:2004rt,Scheck:2006vn,Wongwathanarat:2010yq}.  This requires one to infer a PNS kick through a rigid, impenetrable inner boundary, but allows one to track the supernova explosion for several seconds and to distances greater than 10,000 km.  While this approach is appealing, it must be checked by simulations which include the PNS in the computational domain.  Recently, we have carried out the first axisymmetric, radiation-hydrodynamic simulation of recoil with the multi-group, arbitrary, Lagrangian-Eulerian code {\\sc VULCAN/2D} \\citep{Nordhaus:2010kx}.  By transitioning from a spherical-polar mesh to a pseudo-Cartesian mesh at the center of the domain, we self-consistently tracked the PNS's formation and off-center motion. This calculation was computationally expensive, as it employed multi-group flux limited diffusion neutrino transport and followed the supernova explosion until the shock reached the 5,000 km radial outer boundary of the domain.  At that time, the PNS had reached a velocity of $\\sim$150 km\\,s$^{-1}$ but had yet to fully decouple from the ejecta.  The PNS recoil was due almost entirely to hydrodynamical processes and was consistent with previous excised-core calculations \\citep{Scheck:2004rt,Scheck:2006vn}.  Asymmetric neutrino emission contributed $\\sim$2\\% to the overall kick magnitude.  This suggests that neutrinos play no significant role in accelerating neutron stars to high velocities during typical core-collapse supernovae.  At the end of the calculation, significant acceleration ($\\sim$350 km\\,s$^{-2}$), in addition to the degree of asymmetry in the ejecta, suggested that higher PNS velocities were possible.  Verifying these estimates requires tracking the supernova shock to greater radial distances and later times.  To expand upon the work of \\cite{Nordhaus:2010kx}, we carry out a suite of axisymmetric collapse calculations with the adaptive-mesh-refinement (AMR), radiation-hydrodynamics code, CASTRO \\citep{Almgren:2010fk,Zhang:2011lr}.  By employing AMR and a simplified form of transport (see Sec. \\ref{2}), we expand the domain to a radial distance of 10,000 km and perform multiple calculations. Our PNSs achieve recoil velocities ranging from tens of km\\,s$^{-1}$ up to $\\sim$620 km\\,s$^{-1}$.  In general, the magnitude of the recoil depends on the degree of asymmetry at the time of explosion and the energy of the explosion itself \\citep{Burrows:2007ly}.  In Sec.~3, we discuss the physical processes that accelerate the PNS.  In Sec. \\ref{4}, we compare our results with previous work before concluding and discussing future work in Sec. \\ref{5}. ",
        "conclusions": "} We have carried out a suite of axisymmetric simulations of the collapse of a massive star's core with the AMR, radiation-hydrodynamic code {\\sc CASTRO}.  For each calculation, we follow the core collapse, PNS formation, ensuing neutrino-driven supernova explosion and PNS recoil.  By incorporating the effects of neutrino heating and cooling in place of more detailed and computationally expensive neutrino transport, we are able to perform multiple calculations that simultaneously follow the evolution of the PNS and the global explosion for $\\sim$1 second and to distances of $\\sim$10,000 km.  The PNSs in our simulations achieved recoil velocities comparable to the those of observed pulsars.  After $\\sim$1 second of post-bounce evolution, the highest PNS velocity obtained was 620 km\\,s$^{-1}$ (model $\\rm L_{2.3}$).  After $\\sim$0.6 seconds of post-bounce evolution, this acceleration was supplied primarily by the gravitational pull of slow-moving ejecta in front of the PNS.  This gravitational effect dominates the late-time PNS acceleration in all of our calculations.  While our PNSs have started to decouple from the surrounding fluid (see Fig.~\\ref{fig:accel_3panel}), the substantial and ongoing gravitational acceleration suggests that higher velocities may ultimately be achievable.  Our results suggest that hydrodynamic recoil during an asymmetric supernova explosion provides a natural explanation for pulsar kicks. After the bounce shock stalls, hydrodynamic instabilities deform the shocked material and ensure that the ensuing explosion is asymmetric. Recoil during the supernova explosion and gravitational interaction with the expanding ejecta subsequently accelerate the PNS to high velocities.  The results presented in this work are consistent with the findings of \\cite{Nordhaus:2010kx} and previous axisymmetric simulations that excised the PNS from the computational domain \\citep{Scheck:2004rt,Scheck:2006vn}.  Taken together, these studies strongly suggest that generic core-collapse supernovae can accelerate neutron stars to the high velocities observed in the pulsar population.  Additionally, these studies demonstrate that hydrodynamic processes, and not asymmetric neutrino emission, are responsible for this acceleration \\citep{Scheck:2006vn,Burrows:2007ly,Nordhaus:2010kx}.  In fact, recent simulations of neutron star kicks in three dimensions suggest that velocities comparable to those from axisymmetric calculations are achievable \\citep{Wongwathanarat:2010yq}.  In this work, we have provided substantial numerical support to the hydrodynamic mechanism of pulsar kicks. Recoil due to a neutrino-driven, core-collapse supernova explosion provides a natural explanation for pulsar kicks without appealing to more exotic scenarios.  As computational methods and resources improve, self-consistent three-dimensional studies will enable a full comparison of theoretical models to observed distributions of pulsar kicks and spins."
    },
    "1112/1112.2728_arXiv.txt": {
        "abstract": "A dedicated mission to investigate exoplanetary atmospheres represents a major milestone in our quest to understand our place in the universe by placing our Solar System in context and by addressing the suitability of planets for the presence of life.  EChO -- the Exoplanet Characterisation Observatory -- is a mission concept specifically geared for  this purpose.  EChO will provide simultaneous, multi-wavelength spectroscopic observations on a stable platform that will allow very long exposures. The use of passive cooling, few moving parts and well established technology gives a low-risk and potentially long-lived mission. EChO will build on observations by Hubble, Spitzer and ground-based telescopes, which discovered the first molecules and atoms in exoplanetary atmospheres. However, EChO's configuration and specifications are designed to study a number of systems in a consistent manner that will eliminate the ambiguities affecting prior observations.  EChO will simultaneously observe a broad enough spectral region -- from the visible to the mid-infrared -- to constrain from one single spectrum the temperature structure of the atmosphere, the abundances of the major carbon and oxygen bearing species, the expected photochemically-produced species and magnetospheric signatures. The spectral range and resolution are tailored to separate bands belonging to up to 30 molecules and retrieve the composition and temperature structure of planetary atmospheres.  The target list for EChO includes planets ranging from Jupiter-sized with equilibrium temperatures $T_{eq}$ up to 2000 K, to those of a few Earth masses, with $T_{eq}$ $\\sim$300~K.  The list will include planets with no Solar System analog, such as the recently discovered planets GJ1214b, whose density lies between that of terrestrial and gaseous planets, or the rocky-iron planet 55 Cnc~e, with day-side temperature close to 3000~K. As the number of detected exoplanets is growing rapidly each year, and the mass and radius of those detected steadily decreases, the target list will be constantly adjusted to include the most interesting systems.  We have baselined a dispersive spectrograph design covering continuously the 0.4--16\\,$\\mu$m spectral range in 6 channels (1 in the visible, 5 in the InfraRed), which allows the spectral resolution to be adapted  from several tens to several hundreds, depending on the target brightness.  The instrument will be mounted behind a 1.5 m class telescope, passively cooled to 50~K, with the instrument structure and optics passively cooled to $\\sim$45~K. EChO will be placed in a grand halo orbit around L2.  This orbit, in combination with an optimised thermal shield design, provides a highly stable thermal environment and a high degree of visibility of the sky   to  observe repeatedly several tens of targets over the year. Both the baseline and alternative designs have been evaluated and no critical items with Technology Readiness Level (TRL) less than 4 to 5 have been identified.  We have also undertaken a first-order cost and development plan analysis and find that EChO is easily compatible with the ESA M-class mission framework. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.0025_arXiv.txt": {
        "abstract": "{Observations of the $\\gamma$-ray sky with {\\textit Fermi} led to significant advances  towards understanding blazars, the most extreme class of Active Galactic Nuclei. A large fraction of the population detected by {\\textit Fermi} is formed by BL~Lacertae  (BL~Lac)  objects, whose sample has always suffered from a severe redshift incompleteness due to the quasi-featureless optical spectra.} {Our goal is to provide a significant increase of the number of confirmed high-redshift BL~Lac objects contained in the 2~LAC {\\textit Fermi}/LAT cataloge.} {For 103 {\\textit Fermi}/LAT blazars, photometric redshifts using spectral energy distribution fitting have been obtained. The photometry includes 13 broad-band filters from the far ultraviolet to the near-IR observed with {\\it Swift}/UVOT and the multi-channel imager GROND at the MPG/ESO 2.2m telescope. Data have been taken quasi-simultaneously and the remaining source-intrinsic variability has been corrected for.} {We release the UV-to-near-IR 13-band photometry for all 103 sources and provide redshift constraints for 75 sources without previously known redshift. Out of those, eight have reliable photometric redshifts at $z\\gtrsim1.3$, while for the other 67 sources we provide upper limits. Six of the former eight are BL~Lac objects, which quadruples the sample of confirmed high-redshift BL~Lac. This includes three sources with redshifts higher than the previous record for BL~Lac, including CRATES~J0402-2615, with the best-fit solution at $z\\approx1.9$.} {} ",
        "introduction": "Since  its  launch in  2008,  the  {\\it  Fermi} Space  Laboratory  has dramatically extended  our view of the high-energy  sky.  The recently released  24-month catalog  (2LAC)  of Active  Galactic Nuclei  (AGN) detected  by the Large  Area Telescope  \\citep[LAT; ][]{Atwood:2009aa} revealed  885 high-significance  sources, the  large majority  of them being  blazars  \\citep{Ackermann:2011lr}.  The  latter  form the  most extreme class of AGN with their observational characteristics governed by the small angle between  their relativistic jets and the observer's sight line  \\citep{Blandford:1978lr}.  The resulting  Doppler boosting makes blazars  exceptionally bright sources at  nearly all wavelengths and therefore visible out to high redshift.  The  scientific  relevance of  blazars  is  very  broad, ranging  from laboratories for  the physics and structure of  relativistic jets (and thus the  extraction of  energy from the  central massive  black hole) \\citep[e.g.,][]{Abdo:2010ab}   to   probes   of   the   extra-galactic background  light (EBL)  through attenuation  of  $\\gamma$-ray photons \\citep[e.g.,][]{Abdo:2010ac}.  One of the crucial parameters for these applications is  the distance to  the objects, which  unfortunately is not easy to obtain in most cases.  Two classes  of blazars  dominate the 2LAC  population. These  are the flat-spectrum radio  quasars (FSRQs,  310 sources) and  BL~Lac objects (395),      named       after      the      prototype      BL~Lacertae \\citep[][]{Hoffmeister:1929lr}.    While  for   the   former  redshift measurements are routinely performed using their strong emission lines at UV-optical wavelengths,  the featureless, power-law optical spectra of    BL~Lac    objects    have    proven   to    be    a    challenge \\citep[e.g.,][]{Shaw:2009fk}. Indeed,  220 of  the 395 BL~Lacs  in the 2LAC (55\\,\\%)  lack redshift estimates.  Until  this incompleteness is resolved,   conclusions   about   the   EBL,   the   blazar   sequence \\citep[e.g.,][]{Fossati:1998lr,Ghisellini:1998fk},   and   the  blazar population in general remain tentative at best.  Several methods  have been exploited  to increase the  BL~Lac redshift sample.   At low  distance,  one can  utilize  the remarkably  uniform absolute  brightness  of the  giant  elliptical  BL~Lac host  galaxies \\citep{Sbarufatti:2005qy,Meisner:2010uq} and very-high signal-to-noise optical spectroscopy to help  to identify weak emission or absorption features  in  a  few  other  cases  \\citep[e.g.,][]{Shaw:2009fk}.   An alternative  method, applicable  to the  more distant  sources,  is the photometric  redshift technique, which consists of fitting spectral energy distribution (SED) templates to multi-band photometry.  Neutral hydrogen  along the  line of sight to the blazar will  imprint a clear attenuation signature at the Lyman limit and thus allows an accurate estimate of the redshift of the absorber.  Even  though the absorber  will be located  somewhere along the line of sight and its redshift will thus not necessarily correpond to that of  the blazar, photometric redshifts, $z_{\\rm phot}$, will  provide a reliable lower limit on the blazar redshift.  In  this  paper  we  explore  the  use  of  ultra-violet  to  near-infrared quasi-simultaneous  photometry to  measure the  redshift of  103 2LAC blazars, 86 of them without previous redshift constraints.  The sample selection is  described in Sect.~\\ref{sec:sample}  while the observations and   data   reduction    are   detailed   in   Sect.~\\ref{sec:obs}   and Sect.~\\ref{sec:data}. The  results and their discussion  are presented in Sect.~\\ref{sec:results} and Sect.~\\ref{sec:discussion}, respectively. ",
        "conclusions": "\\label{sec:discussion}  In this paper  we presented redshift constraints for  103 blazars from the  2LAC catalog using  UV-to-near-IR multi-band  photometry obtained quasi-simultaneously with {\\it Swift}/UVOT and GROND.  We provided the first reliable  redshift measurements for eight sources  and new upper limits  for  an additional  66  targets.  Of  the eight  sources  with reliable  redshift, seven are  located at  $z_{\\rm phot}>1.3$.  Six of those belong  to the  BL Lac population.   For comparison, out  of the total 395 BL  Lac in the 2LAC sample only  two sources have previously been  known to  lie at  $z>1.3$ (Fig.~\\ref{fig:zDis}).   Redshifts for these  two,   \\object{PKS~0332-403}  and  \\object{CRATES~J1312-2156},   have  also  been confirmed with our photometric observations.  \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth]{zDis.eps} \\caption{Redshift distribution for the nine BL Lac objects with reliable photometric redshifts (blue filled histogram) together with the distribution for the BL~Lac in the 2LAC catalog \\citep[empty histogram;][]{Ackermann:2011lr}. Sources with $z<0.5$ have been omitted for clarity.} \\label{fig:zDis} \\end{figure}    The six new BL Lac  redshifts at $z>1.3$ represent a dramatic (from two to eight)  increase of the confirmed high-z  {\\it Fermi} sample and thus demonstrate the  opportunity that the SED template fitting technique  holds for  obtaining  photometric redshifts  for BL  Lac sources.    For  our  sample,   this  was   possible  due   to  the densely-covered wide spectral  range (160--2200\\,nm), necessary for a  reliable  constraint  of   the  spectral  slope,  the  excellent ultra-violet   coverage  to  measure   the  Lyman-limit,   and  the quasi-simultaneity  of the  observations, important  for minimizing the impact of source-intrinsic  variability.  The method applied in this  work  overcomes  significant  challenges  inherent  in  other redshift techniques, namely the  simplicity of the optical emission of BL Lacs manifested  as the power-law-shaped synchrotron spectral component.   This complicates spectroscopic  redshift measurements, which are  in most cases  limited to the optical  wavelength regime and  thus  insensitive to  the  Lyman-limit. Instead,  spectroscopy relies  on the detection  of very  faint emission  features, mainly from the  underlying host galaxy.  It therefore requires  very high signal-to-noise, often at  the cost of long exposure  times. On the other hand,  the photometric  redshift method for  BL Lac  does not require significant  telescope time,  in particular when  data from efficient multi-channel instruments  like GROND are available.  Its weakness, however, is that the  method relies on the detection of a particular spectral  feature, which,  albeit strong, is  located in the  rest-frame far-ultra-violet.  For $z\\lesssim3$,  this  is only accessible with space-based  ultra-violet telescopes, and even then usually limited to redshifts above $z\\approx1.2$, as shown above.   The sensitivity of this method to high-redshift sources also allows the placement of upper limits  for those SEDs that do not show the  imprint of  the Lyman-limit.  Only for  three targets  is this boundary above $z=2$. In one case, \\object{CRATES~J0250+1708} ($z<3.1$), the optical counterpart  is too faint  ($r^\\prime\\approx20.7$\\,mag) and only  upper limits in  all six  UVOT bands  could be  obtained. The counterpart of \\object{CRATES~J0705-4847} is  only detected in the GROND $J$ ($\\approx19.7$\\,mag$_{\\rm  AB}$)  and $H$  bands  and  its SED  and redshift  are   thus  poorly  constrained.  Finally,   the  SED  of \\object{ATG20~J0124-0625}   ($z<2.46$)   shows   a  significant   break   at $z\\sim1.9\\pm0.5$. However, as the redshift probability distribution is broad, and the $\\chi^2$ of  the power law fit comparable to that of a $z\\approx0.4$ solution  for a galaxy template, the photometric redshift is considered to be unreliable. Except for those three, no other source in  our sample has a best-fit  photometric redshift at $z>2$.  Thus,  \\object{CRATES~J0402-2615} can  now  be  considered the  most distant known BL Lac with a measured redshift of $z\\approx1.92$.  Fig.~\\ref{fig:zDis}  indicates  that  our  high-redshift findings  are  the natural extension  of the  existing 2LAC redshift  sample. However, due to the incompletenesses of both samples, we refrain from drawing any   quantitative   conclusions   at   this  stage.  We note, however, the our result is general agreement with the theoretical predictions from \\cite{Giommi:2011lr}.  A  more detailed physical interpretation  of the  establishment of  an increased  fraction of high-z BL Lacs in the 2LAC sample will be reported separately."
    },
    "1112/1112.0355_arXiv.txt": {
        "abstract": "In accretion systems, outflows may have significant influence on the luminosity fluctuations. In this paper, following the Lyubarskii's general scheme, we revisit the power spectral density of luminosity fluctuations by taking into account the role of outflows. Our analysis is based on the assumption that the coupling between the local outflow and inflow is weak on the accretion rate fluctuations. We find that, for the inflow mass accretion rate $\\dot M \\propto r^{s}$, the power spectrum of flicker noise component will present a power-law distribution $p(f) \\propto f^{-(1+4s/3)}$ for advection-dominated flows. We also obtain descriptions of $p(f)$ for both standard thin discs and neutrino-cooled discs, which show that the power-law index of a neutrino-cooled disc is generally larger than that of a photon-cooled disc. Furthermore, the obtained relationship between $p(f)$ and $s$ indicates the possibility of evaluating the strength of outflows by the power spectrum in X-ray binaries and gamma-ray bursts. In addition, we discuss the possible influence of the outflow-inflow coupling on our results. ",
        "introduction": "The emission of Galactic Black Hole Binaries (BHBs) and active galactic nuclei (AGN) displays a significant aperiodic variability on a broad range of time-scales. The Power Spectral Density (PSD) of such variability is generally modeled with a power law, $p(f) \\propto f^{-\\beta}$, where $p(f)$ is the power at frequency $f$, and the power-law index ${\\beta}$ keeps a constant in a certain range of $f$, but changes among different ranges. At high frequencies, the PSDs of both BHBs and AGN present a steep slope with $\\beta \\sim 2$. On the contrary, below a break frequency, typically at a few Hz for BHBs, they flatten to a slope with $\\beta \\sim 1$, representing the flicker noise (see \\citealp{King04} and references therein).  Several models have been proposed in order to understand this nearly featureless character of power spectra. The so-called ``shot noise models\" (\\citealp{Terrell72}) attempted to describe the light curves as a series of independent overlapping shots with specific time-scales, amplitudes, and occurence rates. Due to lack of physical picture in this scenario, various physically motivated ideas have been put forward subsequently, such as the fluctuations of hydrodynamic or magnetohydrodynamic turbulence (\\citealp{Nowak95}; \\citealp{Hawley01}), magnetic flares or density fluctuation in the corona (\\citealp{Galeev79}; \\citealp{Poutanen99}; \\citealp{Goosmann06}; \\citealp{Kawanaka08}), and Lyubarskii's general scheme (\\citealp{Lyubarskii97}; \\citealp{King04}). In the Lyubarskii's scheme, it was noted that any variation of accretion rate, which is caused by small amplitude variations in the viscosity, would induce a variation in the accretion rate at the inner radius of the disc, where most of the energy is released. Moreover, observations showed that the variability is non-linear and the rms variability is proportional to the average flux over a wide range of time-scales (e.g., \\citealp{Uttley01}; \\citealp{Uttley05}; \\citealp{Gleissner04}). It indicates that the short time-scale variations are modulated by the longer time-scales, which favors the Lyubarskii's scheme.  However, the observed power spectra often deviate from the form $f^{-1}$. For example, the PSD of Cyg X-1 is well described with the form $f^{-1}$ in the soft state, however, exhibits the form $f^{-1.3}$ in the hard state (\\citealp{Gilfanov10}). In particular, it is shown that the power-law index is around $0.8-1.3$ both in the soft state of BHBs and in narrow-line Seyfert 1 galaxies (\\citealp{Janiuk07}). Such a dispersion of the power-law index reveals that there must exist some other mechanism. A radius-dependent amplitudes of $\\alpha$ fluctuations may help to alleviate the discrepancy between theories and observations. However, it remains unclear why the fluctuations of $\\alpha$ should have a strong radius-dependent form.  In the present paper, we will take into account another mechanism, outflows, which is a popular phenomenon in accretion systems and has strong observational evidence. One of the best examples comes from Sgr A*, whose center harbors a supermassive black hole surrounded by an accretion flow that is likely to be in the form of the advection-dominated accretion flow (ADAF, \\citealp{Narayan94}). Radio polarization observations constrain the accretion rate in the innermost region is nearly two orders of magnitude lower than that measured at the Bondi radius (e.g., \\citealp{Marrone06}), which indicates that intense outflows may present in this system. Besides, the absorption lines from highly ionized elements, which have been detected in the X-ray spectrum of some microquasars such as GRO J1655-40 (\\citealp{Ueda98}; \\citealp{Yamaoka01}; \\citealp{Miller06}), GRS 1915+105 (\\citealp{Kotani00}; \\citealp{Lee02}) and Atoll sources (see e.g. the review by \\citealp{Diaz Trigo06} and references therein), also indicate the existence of outflows. On the other hand, \\citet{Jiao11} found that outflows generally exist in accretion discs no matter that the flow is advection-dominated such as the slim disc (\\citealp{Abramowicz88}) and the ADAF, or is radiation-dominated such as the standard thin disc (\\citealp{Shakura73}). In particular, for the three types of advection-dominated flows: ADAFs (gas internal energy dominant), slim discs (trapped photon energy dominant), and hyper-accretion discs (trapped neutrino energy dominant), outflows may be significantly strong due to positive Bernoulli parameters (e.g., \\citealp{Narayan97}; \\citealp{Liu11}) or the large radiation pressure (e.g., \\citealp{Gu07}). Furthermore, outflows have generally been found in many simulation works (e.g., \\citealp{Ohsuga11} and references therein).  In the Lyubarskii's scheme, the power spectrum of luminosity fluctuations is sensitive to the varying mass accretion rate, thus we expect that outflows may have significant effects on the power-law index. The paper is organized as follows. In Section 2, we investigate the fluctuation power spectrum under a radius-dependent accretion rate following the method of Lyubarskii. The potential application of our results to observations is discussed in Section 3. ",
        "conclusions": "In this paper, we evaluate the effects of outflows on the luminosity fluctuations with the Lyubarskii's general scheme. With a radius-dependent accretion rate $\\dot M \\propto r^{s}$, the power spectrum of the luminosity fluctuations is $p(f) \\propto f^{-\\beta}$, where the value of $\\beta$ varies with the disc structure. By assuming that the coupling between the local outflow and inflow is weak on the accretion rate fluctuations, we obtain the following explicit expressions of $\\beta$ for different disc models: $\\beta = 1 + 4s/3$ for advection-dominated discs, $\\beta = 1 + 40s/(25-6s)$ for the outer region of standard thin discs, $\\beta = 1 + {10s}/(7-2s)$ for the middle region, $\\beta = 1 + {4}{s}/({7-4s})$ for the inner region, and $\\beta = 1 + 5s/2$ for NDAFs. The above expressions imply that $\\beta$ in a GRB is generally larger than that in a BHB for comparable $s$. The expressions of $\\beta$ indicate the possibility of evaluating the strength of outflows by the power spectrum in X-ray binaries and GRBs. In addition, if the coupling is not negligible, the value of $\\beta$ will probably be located between unity and the value presented in the above expressions.  In both BHBs and AGN, ADAFs are usually adopted to describe the quiescent state, the low/hard state, and the corona which lies above a cold disc. ADAFs may produce significant outflows, and therefore the power spectrum can deviate from $f^{-1}$ based on the present analysis. The exact value of ${s}$ is, however, difficult to estimate from the theoretical point of view, except for the general constraint $0 < s < 1$ (\\citealp{Narayan08}). On the other hand, some observations indicate $s \\sim 0.3$ (\\citealp{Yuan03}; \\citealp{Zhang10}). Taking this value, we obtain $\\beta=1.4$ for ADAFs, which is close to $1.3$, the power-law index of PSDs presented in the low mass X-ray binary systems (\\citealp{Gilfanov05}). The quantitative difference may be relevant to the coupling between the outflow and inflow as discussed in \\S2.4.  If there is only outflow that operates in the accreting system, $s$ should be positive. However, $s$ can also be negative due to the evaporation mechanism of a cold disc. For the accreting black hole in BHBs, the observed power-law components in the X-ray spectra are generally attributed to hot, tenuous plasmas, namely accretion disc coronae. Due to the high temperature in the corona, the interaction between the disc and corona would lead to mass evaporating from the disc to the corona (\\citealp{Meyer00}; \\citealp{Spruit02}). In this case, the value of $s$ for the corona should be negative if outflows are not strong, and therefore it is quite possible for $\\beta$ to be less than unity. Consequently, in this scenario $s$ for the underneath cold disc should be positive."
    },
    "1112/1112.2350_arXiv.txt": {
        "abstract": "We extend the holographic Ricci dark energy model to include some direct, non-gravitational interaction between dark energy and dark matter. We consider three phenomenological forms for the interaction term $Q$ in the model, namely, $Q$ is taken proportional to the Hubble expansion rate and the energy densities of dark sectors (taken to be $\\rho_{\\rm de}$, $\\rho_{\\rm m}$, and $\\rho_{\\rm de}+\\rho_{\\rm m}$, respectively). We obtain a uniform analytical solution to the three interacting models. Furthermore, we constrain the models by using the latest observational data, including the 557 Union2 type Ia supernovae data, the cosmic microwave background anisotropy data from the 7-yr WMAP, and the baryon acoustic oscillation data from the SDSS. We show that in the interacting models of the holographic Ricci dark energy, a more reasonable value of $\\Omega_{\\rm m0}$ will be obtained, and the observations favor a rather strong coupling between dark energy and dark matter. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.0952_arXiv.txt": {
        "abstract": "We present a brief overview of the theory of stellar winds with a strong emphasis on the radiation-driven outflows from massive stars. The resulting implications for the evolution and fate of massive stars are also discussed. Furthermore, we relate the effects of mass loss to the angular momentum evolution, which is particularly relevant for the production of long and soft gamma-ray bursts. Mass-loss rates are not only a function of the metallicity, but are also found to depend on temperature, particularly in the region of the bi-stability jump at 21 000 Kelvin. We highlight the role of the bi-stability jump for Luminous Blue Variable (LBV) stars, and discuss suggestions that LBVs might be direct progenitors of supernovae. We emphasize that radiation-driven wind studies rely heavily on the input opacity data and linelists, and that these are thus of fundamental importance to both the mass-loss predictions themselves, as well as to our overall understanding of the lives and deaths of massive stars. ",
        "introduction": "\\label{s:intro}  Despite their relative rarity, massive stars are dominantly dangerous for their environments. This is a result of their mass and energy input via stellar winds and subsequent core-collapse supernovae (SN). A overarching parameter for the life expectancy of a massive star concerns its strong mass outflow which is driven by radiative forces on millions of ionic spectral line transitions, providing it with its name ``line-driven wind''.  On the one hand, mass loss is thought to be a key agent in revealing chemically processed material at the stellar surface, making it responsible for evolutionary scenarios such as the O $\\rightarrow$ Luminous Blue Variable (LBV) $\\rightarrow$ Wolf-Rayet (WR) star $\\rightarrow$ SN sequence (e.g. Conti 1976, Chiosi \\& Maeder 1986, Langer et al. 1994). Furthermore, it determines the stellar mass before collapse and is thus relevant for the type of compact remnant that is left behind (i.e. neutron star or black hole). On the other hand, the role of mass loss may be equally relevant for the loss of angular momentum (e.g. Meynet \\& Maeder 2003).  With respect to the latter, it has been suggested that low metallicity (actually low ``iron'' contents; Vink \\& de Koter 2005) leads to less mass and angular momentum loss in low metallicity environments, perhaps resulting in a preference of long gamma-ray bursts (GRB) in the early Universe, but however interesting the metallicity dependence and the long GRB puzzle may be, the temperature dependence of stellar winds and its role in the angular momentum evolution of massive stars has been highlighted more recently with respect to the possibility of {\\it bi-stability braking} (Vink et al. 2010).  Given the crucial role that mass loss plays for massive star evolution, we discuss the theory of massive star mass loss and its implications, with a focus on the metallicity ($Z$) and effective temperature ($T_{\\rm eff}$) dependence. We will see that the {\\it iron line opacity} plays a dominant role in both cases. ",
        "conclusions": ""
    },
    "1112/1112.1165_arXiv.txt": {
        "abstract": "Using the {\\hf Yunnan} evolutionary population synthesis (EPS) models with and without binary interactions, we present the luminosity of $\\rm H\\alpha$ recombination line ($L_{\\rm H\\alpha}$), the luminosity of [OII]$\\lambda$3727 forbidden-line doublet ($L_{\\rm [OII]}$), the ultraviolet (UV) fluxes at 1500 and 2800\\,$\\rm \\AA$ ($L_{i, {\\rm UV}}$) and far-infrared flux ($L_{\\rm FIR}$) for Burst, S0, Sa-Sd and Irr galaxies, and present the calibrations of star formation rate (SFR) in terms of these diagnostics.    By comparison, we find that binary interactions lower the SFR.vs.$L_{\\rm H\\alpha}$ and SFR.vs.$L_{\\rm [OII]}$ conversion factors by $\\sim$0.2\\,dex. The main reason is that binary interactions raise the UV flux (shortward of the Lyman limit) of the stellar population (SP) in the age range 6.7$<$ log$t$/yr $<$8.4 and thus more ionizing photons are present in the nebula. Moreover, binary interactions do not significantly vary the calibrations of SFR in terms of $L_{i, {\\rm UV}}$. This is because binary interactions raise the flux at 1500\\,$\\rm \\AA$ of the SP in the range 8.75 $<$ log$t$/yr $<$ 9.2 and the maximal difference is about 1\\,dex. In addition, binary interactions have little effect on the flux at 2800\\,$\\rm \\AA$. At last, the calibration of SFR from $L_{\\rm FIR}$ is almost unaffected by binary interactions. This is caused by the fact that binary interactions almost do not affect the bolometric magnitudes of SPs.    We also discuss the effects of initial mass function (IMF), gas-recycle assumption and EPS models (including {\\hf GISSEL98, BC03, STARBURST99, PopSTAR and P\\'{E}GASE} models) on these SFR calibrations. Comparing the results by using Salpeter (S55) IMF with those by using Miller \\& Scalo (MS79) IMF, we find that the SFR.vs.$L_{\\rm H\\alpha}$ and SFR.vs.$L_{\\rm [OII]}$ conversion factors by using S55 IMF are greater by 0.4 and 0.2\\,dex than those by using MS79 IMF for the {\\hf Yunnan} models with and without binary interactions, respectively. The SFR.vs.$L_{i,{\\rm UV}}$ and SFR.vs.$L_{\\rm FIR}$ conversion factors by using S55 IMF are larger by an amount of 0.2\\,dex than the corresponding ones by using MS79 IMF. The inclusion of gas-recycle assumption only lowers these SFR calibrations at faint SFR. Moreover, comparing the results when using different EPS models, we find that the differences in the SFR.vs.$L_{\\rm H\\alpha}$ and SFR.vs.$L_{\\rm [OII]}$ conversion factors reach $\\sim$ 0.7 and 0.9\\,dex, the difference in the SFR.vs.$L_{\\rm FIR}$ conversion factor reaches 0.4 and 0.8\\,dex, and the differences in the SFR.vs.$L_{i, {\\rm UV}}$ conversion factors reach 0.3 and 0.2\\,dex when using S55 and NON-S55 IMF (including Cha03, K01, K93' and MS79 IMFs, partly caused by the difference in the IMF), respectively. At last, we give the conversion coefficients between SFR and these diagnostics for all models. ",
        "introduction": "One of the most recognizable features of galaxies along the Hubble sequence (loose definition, including not only morphological type but also gas content, mass, bar structure and dynamical environment) is the wide range in young stellar content and star formation activity. Understanding its physical nature and origin of the variation in stellar content are fundamental to understand evolution of galaxies \\citep[][hereafter K98]{ken98}. Star formation rate (SFR) can be used to compare with those distant galaxies at cosmological lookback times, and extrapolate the future timescales for star formation in galaxies by combining with HI and CO measurements \\citep{ken94}. Moreover, star formation activity is usually correlated with cold gas and stars in galaxies: stars continuously produce mass, energy and metals during their evolution processes, and return them to galactic medium (gas), affecting the status of the next generation of stars. SFR carries the information on the evolution of galaxies. Therefore, it is important to determine SFR and its variation with Hubble type (loose definition) and environment, which can help us to understand the evolution of galaxies.    The commonly used SFR tracers include the flux of H$\\alpha$ nebular recombination line, the ultraviolet (UV) continuum flux, the flux of [OII]$\\lambda$\\,3727 forbidden-line doublet and far infra-red (FIR) continuum flux (K98). These SFR tracers are more or less correlated with the UV passband.  \\begin{itemize} \\item First, the UV flux is directly tied to the photospheric emission of the young stellar population (SP).  \\item Second, the integrated luminosity of galaxy shortward of the Lyman limit (Far-UV) can ionize the hydrogen in the nebula and produce the recombination lines (such as H$\\alpha$, H$\\beta$, and so on). Thus the luminosities of these lines can be used to trace SFR.  \\item Third, the luminosity of the strong [OII]$\\lambda$\\,3727 forbidden-line doublet is often empirically obtained through the H$\\alpha$ luminosity, although it is not coupled to the ionizing luminosity and the excitation of this line is sensitive to abundance and the ionization state of the gas.  \\item Finally, the last SFR trace, the FIR luminosity, is also correlated with the UV passband. The interstellar dust can absorb the bolometric luminosity of galaxy and re-emit it in the thermal IR passband. The absorption cross section of the dust peaks in the UV passband, and, since the UV flux is considered as a tracer of young SP, the FIR luminosity can also diagnose SFR. \\end{itemize}   Furthermore, the last three diagnostics and the corresponding calibrations of SFR are correlated with the UV flux. Since in our previous studies, we found that the inclusion of binary interactions in evolutionary population synthesis (EPS) models can raise the UV flux by $\\sim$ 2-3 magnitudes for SP at an age of $\\sim 1$\\,Gyr \\citep{zha04,zha05}, in this study we will discuss the effect of binary interactions on these calibrations.    The outline of the paper is as follows. In Section 2 we describe the used EPS models and algorithm. In Section 3 we overview the previous results about SFR calibrations, and the advantages and disadvantages of these SFR tracers. In Section 4 we give the effect of binary interactions on these SFR calibrations. In Section 5 we discuss the effects of initial mass function (IMF), gas-recycle assumption and the EPS models on these SFR calibrations, and give the conversion coefficients between SFR and these tracers for all models. Finally we present a summary and conclusions in Section 6. ",
        "conclusions": "We use the {\\hf Yunnan} EPS models with and without binary interactions to present the luminosities of the $\\rm H\\alpha$ recombination line, the [OII]$\\lambda$3727 forbidden-line doublet, the UV (at $1500$ and $2800$\\,$\\rm \\AA$) and the FIR continuum for Burst, E, S0, Sa-Sd and Irr galaxies, and present the calibrations of SFR in terms of these diagnostics.   By comparison, we find that binary interactions lower the SFR.vs.$L_{\\rm H\\alpha}$ and SFR.vs.$L_{\\rm [OII]}$ conversion factors by $\\sim $0.2\\,dex, and do not significantly vary the SFR.vs.$L_{i,{\\rm UV}}$ (at 1500 and 2800 $\\rm \\AA$) and SFR.vs.$L_{\\rm FIR}$ calibrations.    We also consider the effects of IMF, the gas-recycle assumption and EPS models on these calibrations. By comparison, we find that the SFR.vs.$L_{\\rm H\\alpha}$ and SFR.vs.$L_{\\rm [OII]}$ conversion factors of Models A/B-S55 are larger by 0.4 and 0.2\\,dex than the corresponding ones of Models A/B-MS79, and that the SFR.vs.$L_{i,{\\rm UV}}$ and SFR.vs.$L_{\\rm FIR}$ conversion factors are larger by 0.2\\,dex. By comparing the results between Models C and Cr, we find that the inclusion of gas-recycle assumption only lowers the SFR calibrations at faint SFR. Also we use the other EPS models ({\\hf BC03}, {\\hf GISSEL98}, {\\hf PopSTAR}, {\\hf P\\'{E}GASE} and {\\hf STARBURST99}) to obtain these SFR calibrations. By comparison, we find that the differences in the SFR($L_{\\rm H\\alpha}$) and SFR($L_{\\rm [OII]}$) calibrations reach $\\sim$ 0.7 and 0.9\\,dex, the difference in the SFR($L_{\\rm FIR}$) calibration reaches 0.4 and 0.8\\,dex, and the differences in the SFR($L_{i,{\\rm UV}}$) calibration reach 0.3 and 0.2\\,dex when using S55 and NON-S55 (partly caused by the difference in the IMF) IMFs, respectively.  At last, in this paper we give the conversion coefficients between $SFR$ and these diagnostics for all models. In this paper we have only considered the effects of binary interactions for solar metallicity galaxies - more detailed studies will be given."
    },
    "1112/1112.1954_arXiv.txt": {
        "abstract": "The current dynamical structure of the Kuiper belt was shaped by the orbital evolution of the giant planets, especially Neptune, during the era following planet formation, when the giant planets may have undergone planet-planet scattering and/or planetesimal-driven migration. Numerical simulations of this process, while reproducing many properties of the belt, fail to generate the high inclinations and eccentricities observed for some objects while maintaining the observed dynamically ``cold\" population. We present the first of a three-part parameter study of how different dynamical histories of Neptune sculpt the planetesimal disk. Here we identify which dynamical histories allow an \\emph{in situ} planetesimal disk to remain dynamically cold, becoming today's cold Kuiper belt population. We find that if Neptune undergoes a period of elevated eccentricity and/or inclination, it secularly excites the eccentricities and inclinations of the planetesimal disk. We demonstrate that there are several well-defined regimes for this secular excitation, depending on the relative timescales of Neptune's migration, the damping of Neptune's orbital inclination and/or eccentricity, and the secular evolution of the planetesimals. We model this secular excitation analytically in each regime, allowing for a thorough exploration of parameter space. Neptune's eccentricity and inclination can remain high for a limited amount of time without disrupting the cold classical belt. In the regime of slow damping and slow migration, if Neptune is located (for example) at 20 AU, then its eccentricity must stay below 0.18 and its inclination below 6$^\\circ$. ",
        "introduction": "\\label{sec:intro}  The solar system is often used as a case study for the formation of planetary systems from proto-planetary disks. The current configuration of Kuiper belt objects (KBOs) provides a map for how the dynamical evolution of the giant planets in our solar system sculpted the disk of planetesimals.  Therefore it is possible to use the orbital properties of the Kuiper belt to constrain how the orbits of the giant planets evolved in the early solar system, particularly for Neptune, the primary sculptor of the Kuiper belt. Models employing N-body integrations to trace the effects of the giant planets' migration and orbital eccentricity evolution on the planetesimal disk have enjoyed substantial success in reproducing the dynamical populations of KBOs observed today \\citep[e.g.][]{1993M,1995M,1999H,2003G,2005H,2008L,2008M}. These populations include objects near orbital resonance with Neptune, objects scattering off Neptune, and ``classical\" objects decoupled from Neptune. \\citep[See][for definitions of the dynamical classes.]{2008G} Yet substantial discrepancies still exist between the simulations and observations, particularly for the classical population. The bias-corrected inclination distribution of observed classical KBOs is bimodal \\citep{2001B,2010G,2011V}. The low-inclination, dynamically ``cold\" component and the high-inclination, dynamically ``hot\" component also have distinct physical properties, including colors \\citep{2000T,2003T,2008P}, sizes \\citep{2001L,2010F}, albedos \\citep{2009B}, and binary fractions \\citep{2006S,2008N}. To date no simulations have been able to produce both the high and low inclination classical objects while qualitatively matching their observed eccentricity distribution.  Several theories of the dynamical history of the giant planets in our solar system have been proposed, inspired by the dynamical populations within the Kuiper belt.  One of the most widely accepted models, the Nice Model, stems from a postulated large scale instability in the early solar system \\citep[e.g.][]{1999T,2003G,2008M}. Inspired by the Nice Model, \\citet{2008L} proposes a scenario in which Neptune is scattered outward onto an eccentric orbit to near its current location from a formation location closer to the sun \\citep[e.g.][]{1999T,2002T}. In this scenario, Neptune's eccentricity subsequently damps due to dynamical friction with a remnant disk of planetesimals. This model reproduces the observed resonant population, scattered population, and the hot classical population, but, we argue, does not satisfactorily match the low eccentricities of the observed population of dynamically cold objects in the classical region. Furthermore, it does not produce a sufficient number of high-inclination objects. Other incarnations of the Nice model \\citep[e.g.][]{2003G,2008M} include an \\emph{in situ} population of cold objects, but, over the course of Neptune's evolution, these objects become excited to higher eccentricities. Whether or not the particular history described by the Nice Model is correct, the prevalence of large eccentricities among extrasolar giant planets suggests that large-scale orbital excitation is common during the formation of planetary systems.  In contrast to the upheaval of the Nice Model, another model \\citep{1993M} proposes a period of extensive, smooth migration of the giant planets. Planetesimal-driven migration \\citep{1984F} is likely important in shaping the architecture of many planetary systems.  \\citet{1993M,1995M,1999H} demonstrated that the migration of the giant planets in our solar system on flat, circular orbits would have perturbed the disk of planetesimals, scattering some to more eccentric orbits and capturing others into resonance. They also found that migration results in a large number of planetesimals in the location known today as the ``Scattered Disk,\" which encompasses those objects that have had gravitational interactions with one or more of the giant planets and have high orbital eccentricities and inclinations. The key inconsistency between this model and current observations is that it cannot produce, without additional processes such as stochasticity \\citep[e.g][]{2003L,2006M}, both the cold and hot classical populations from a single set of initial disk conditions or account for the differences in their physical properties.  Because detailed N-body simulations are computationally expensive, previous works have investigated a limited number of solar system planetary histories. Currently, no comprehensive parameter study has been done exploring the evolution of Neptune's orbit given constraints from the sculpting of the Kuiper belt. In addition, although planet-planet scattering often produces large mutual inclinations \\citep{2011C}, as in observed in the Upsilon Andromedae system \\citep{2010M}, no serious, detailed treatment has included the possibility that Neptune underwent a period of high orbital inclination, a conceivable outcome of the orbital instability in the early solar system. Here we consider a general model that can encompass the two detailed models described above -- the Nice-Model-inspired scenario in which Neptune is scattered to a high eccentricity \\citep{2008L} and the scenario of extensive migration of Neptune on a low-eccentricity, low-inclination orbit \\citep[e.g.][]{1995M}  -- as well as potential scenarios in which Neptune undergoes a period of high inclination. In this generalized model, Neptune undergoes some combination of migration and/or evolution of its eccentricity and/or inclination. Here we take a step toward understanding the qualitative differences in the dynamics of the Kuiper belt generated by a wide range of planetary histories. In a series of three papers, we perform a parameter study of the effects on a disk of planetesimals of the migration and the eccentricity and inclination evolution of Neptune in the early solar system. No component of this parameterization is new; rather, our goal is to comprehensively consider all possible parameters for Neptune's history.  In this first paper, we introduce our method for computationally and analytically modeling Neptune's orbital evolution and its effects on the planetesimal population. Then we demonstrate several concepts that will allow us to thoroughly explore the parameter space of Neptune's dynamical history: \\begin{itemize} \\item If Neptune undergoes a period of elevated orbital eccentricity, it will secularly excite the eccentricities and inclinations of an \\emph{in situ} planetesimal population. We can analytically model this secular excitation using a simple expression. \\item As Neptune's eccentricity and inclination damp, the planetesimals evolve to final eccentricities and inclinations that depend on Neptune's initial eccentricity and inclination and damping timescales. \\item The effects of Neptune's eccentricity and inclination evolution on the planetesimals can be treated separately to first order. \\item The migration rate sets Neptune's effective location for the secular evolution of the planetesimals. \\end{itemize} Having demonstrated these points, we can place robust constraints on Neptune's semi-major axis, eccentricity, and inclination during its late evolution (after any period of planet-planet scattering), and on its migration, eccentricity damping, and inclination damping rates, identifying which parameters are consistent with maintaining the low eccentricities and inclinations of the cold classicals. In the second paper \\citep{2012D}, we place even stronger constraints by incorporating additional effects, including the effects of other planets on Neptune's orbital evolution and the effects of proximity to mean-motion resonance with Neptune on the secular excitation of the planetesimals, as well as additional constraints from the hot classical population. A third paper (Dawson and Murray-Clay 2012b, in prep) will focus on the inclinations of the classicals.  In Section \\ref{sec:nep}, we present our parameterization of Neptune's orbital evolution. In Section \\ref{sec:small}, we discuss the observational constraints that today's cold classical KBOs place on past sculpting of the planetesimal disk and describe our computational and analytical models of the excitation of the planetesimal disk by an inclined and eccentric Neptune. We present the results of the parameter study in Section \\ref{sec:results}. In Section \\ref{sec:conclusions}, we present our conclusions and describe how the companion papers will expand upon this work. ",
        "conclusions": "\\label{sec:conclusions}  As a first step in a comprehensive study of the impact on the Kuiper belt of a wide range of possible dynamical histories of the outer solar system, we have performed a suite of numerical integrations probing the impact of the orbital evolution of a single a planet on a disk of planetesimals.  We have presented the observational evidence for a population of dynamically cold objects in the Kuiper belt in the region from 42.5 to 45 AU that are confined to $e < 0.1$ and $i < 6^\\circ$. We argued that recent models of Kuiper belt sculpting -- which explain many of the observed dynamical and physical properties of the Kuiper Belt -- do not generate or preserve sufficiently low eccentricities for the cold classicals \\citep[e.g.][]{2003G,2008L,2008M}. Our results have revealed several principles key to constraining which orbital histories of Neptune are consistent with preserving the \\emph{in situ} planetesimal population at the low eccentricities and inclinations required by the observations:  \\begin{itemize} \\item If Neptune is scattered onto an eccentric and/or inclined orbit, it will secularly excite the eccentricities and inclinations of an \\emph{in situ} planetesimal population. \\item Planetesimals starting with $e = i = 0$ reach eccentricities and inclinations up to twice their forced eccentricity and inclination, on timescales given by Eqn. (\\ref{eqn:sece}) and (\\ref{eqn:seci}). \\item As Neptune's eccentricity and inclination damp, the planetesimals evolve to their final eccentricities and inclinations. If the damping timescales $\\tau_e$ and $\\tau_i$ of Neptune's orbit are slow compared to the secular evolution time, a planetesimal evolves to its initial free eccentricity and inclination, which are set by Neptune's initial eccentricity and inclination. If the damping is fast compared to the secular evolution time, the planetesimal's $e$ and $i$ effectively freeze at the values they reach after one damping time, reaching final values of $\\efree \\sin(\\gkbo \\tau_e)$ and $\\eforced \\sin(\\gkbo \\tau_i)$ respectively. See Table \\ref{tab:constrain} for constraints on Neptune's eccentricity and inclination in the two damping regimes. \\item The effects of Neptune's: 1) eccentricity evolution, and 2) inclination evolution, on the planetesimals can be treated separately to first order. \\item At a given location in the planetesimal disk, the secular excitation timescales and forced eccentricity (but not the forced inclination) depend on Neptune's location, which is altered by Neptune's migration. When Neptune's migration is slow relative to the damping time, the secular evolution effectively takes place at Neptune's initial location. When Neptune's migration is fast relative to the damping time, the secular evolution effectively takes place at Neptune's final location. \\end{itemize}  From these principles, it is evident that the three models described in \\citet{2008L} would not be able to retain a cold classical population. Neptune begins with an eccentricity of 0.3 and a semi-major axis of 27.5 AU (run A and C) or 28.9 AU (run B) and damps on a timescale of 1 Myr (run A and B) or 3 Myr (run C), in the slow migration regime. When Neptune is at 27.5 AU, a planetesimal at 42.5 AU has a forced eccentricity of $e = 0.76 e_N = 0.23$ and a secular evolution timescale of 6 Myr. The final eccentricity of the planetesimal evolves to is reduced by a factor of $\\sin(\\gkbo \\tau_e) = 0.85$ for a damping timescale of 1 Myr (fast damping) and is not reduced for a damping timescale of 3 Myr (slow damping). Thus the final eccentricity of the planetesimal is 0.19 (run A) or 0.23 (run C), well above the observational limit. When Neptune is at 28.9 AU (run B), a planetesimal at 42.5 AU has a forced eccentricity of $e = 0.79 e_N = 0.24$ and a secular evolution timescale of 5 Myr. The final eccentricity the planetesimal evolves to is reduced by a factor of $\\sin(\\gkbo \\tau_e) = 0.97$ for a damping timescale of 1 Myr. Thus the final eccentricity of the planetesimal is 0.23, well above the observational limit. According to the constraints established in our paper, any of these three initial conditions for Neptune could retain the cold classicals if Neptune's eccentricity were to damp more quickly.  Based on these principles, we can place robust constraints on Neptune's dynamical history. For example, in the regime of slow damping and slow migration, if Neptune's initial semi-major axis is 20 AU, then its eccentricity must stay below 0.18 and inclination below 6 degrees. If Neptune's initial semi-major axis is 30 AU -- or if it migrates quickly, relative to the damping time, from its initial location to 30 AU -- then its eccentricity must stay below 0.12. In the case of fast damping -- on a timescale shorter than a planetesimal's secular excitation time -- the initial eccentricity and inclination of Neptune can be even higher.  Having established these principles, we complete our parameter study in two companion papers \\citet{2012D} and Dawson and Murray-Clay (2012b), in prep. In the first companion paper, we consider more generally the constraints on Neptune's dynamical history from the eccentricity distribution of the classical KBOs and incorporate several other important effects in the constraints from the cold classicals: \\begin{itemize} \\item A more accurate model for secular excitation that includes higher-order terms. \\item The greatly increased secular frequency near Neptune's mean-motion resonances, which places strong constraints on Neptune's semi-major axis when its eccentricity is high. \\item The effects of the other giant planets, including precession of Neptune, which can reduce a planetesimal's forced eccentricity \\citep{2011B}, and oscillations in Neptune's semi-major axis, which can create a chaotic sea in the Kuiper belt region. \\end{itemize} Then we combine these constraints for retaining the cold classicals with constraints for creating the hot classical population and identify which regions of parameter space of Neptune's dynamical history can produce the hot classical population without disrupting the cold population. In the second companion paper, we place constraints on Neptune's inclination and inclination damping time.  Our parameterization of Neptune's orbital history allows us to constrain which orbital histories are consistent with maintaining the cold classical population. By combining the observational constraints (Sec. \\ref{sec:obs}) with the principles derived in this paper, we can immediately check whether a particular set of initial conditions for Neptune's semi-major axis, eccentricity, and inclination, and the rates of its migration and eccentricity and inclination damping, are consistent with maintaining the observed dynamically cold population, without performing computationally expensive integrations. A picture, both qualitative and quantitative, is emerging of what dynamical histories of Neptune allow a promising general scenario -- Neptune's delivery of the hot classicals from the inner disk to classical region, where the cold population has formed \\emph{in situ} -- to be consistent with observed low eccentricities and inclinations of the cold classical population. Barring fast precession of Neptune's orbit, which we do not consider in this work, the existence of the cold classical population implies that Neptune could have spent only a limited time at high eccentricity and/or inclination during its dynamical history."
    },
    "1112/1112.4789_arXiv.txt": {
        "abstract": "We determine the observational signatures of protostellar cores by coupling  two-dimensional radiative transfer calculations with numerical hydrodynamical simulations that predict accretion rates that both decline with time and feature short-term variability and episodic bursts caused by disk gravitational instability and fragmentation.  We calculate the radiative transfer of the collapsing cores throughout the full duration of the collapse, using as inputs the core, disk, and protostellar masses, radii, and mass accretion rates predicted by the hydrodynamical simulations. From the resulting spectral energy distributions, we calculate standard observational signatures (\\lbol, \\tbol, \\lbolsmm) to directly compare to observations. We show that the accretion process predicted by these models reproduces the full spread of observed protostars in both \\lbol\\ $-$ \\tbol\\ and \\lbol\\ $-$ \\mcore\\ space, including very low luminosity objects, provides a reasonable match to the observed protostellar luminosity distribution, and resolves the long-standing luminosity problem.  These models predict an embedded phase duration shorter than recent observationally determined estimates (0.12 Myr vs.~0.44 Myr), and a fraction of total time spent in Stage 0 of 23\\%, consistent with the range of values determined by observations. On average, the models spend 1.3\\% of their total time in accretion bursts, during which 5.3\\% of the final stellar mass accretes, with maximum values being 11.8\\% and 35.5\\% for the total time and accreted stellar mass, respectively.  Time-averaged models that filter out the accretion variability and bursts do not provide as good of a match to the observed luminosity problem, suggesting that the bursts are required. ",
        "introduction": "Low-mass stars form from the gravitational collapse of dense cores of gas and dust (e.g., Beichman et al.~1986; Di Francesco et al.~2007; Ward-Thompson et al.~2007a).  In the simplest model, the collapse of a singular isothermal sphere initially at rest as first considered by Shu (1977) and later extended by Terebey, Shu, \\& Cassen (1984; TSC84) to include rotation (often called the ``standard model'' of star formation), collapse proceeds in an ``inside-out'' fashion, beginning in the center of the core, moving outward at the sound speed, and giving rise to a constant mass accretion rate of $\\sim 2 \\times 10^{-6}$ \\msun\\ yr$^{-1}$.  Many modifications to this model have been explored, including non-zero initial inward motions (Larson 1969; Penston 1969; Hunter 1977; Fatuzzao, Adams, \\& Myers 2004), magnetic fields (Galli \\& Shu 1993a, 1993b; Li \\& Shu 1997; Basu 1997), isothermal spheres that are \\emph{not} singular but feature flattened density profiles at small radii (Foster \\& Chevalier 1993; Henriksen, \\andre, \\& Bontemps 1997), and a finite outer boundary (Henriksen, \\andre, \\& Bontemps 1997; Vorobyov \\& Basu 2005a). All but the latter (finite outer boundary) generally increase the accretion rate over that predicted by the standard model.  A significant shortcoming of the standard model is the classic ``luminosity problem,'' whereby accretion at the above rate produces accretion luminosities ($L_{acc} \\propto M_* \\dot{M}$) factors of $10-100$ higher than typically observed for embedded protostars.  First noticed by Kenyon et al.~(1990) and further investigated by Kenyon et al.~(1994) and Kenyon \\& Hartmann (1995), this problem has recently been emphasized by studies presenting results from the \\emph{Spitzer Space Telescope} ``From Molecular Cores to Planet Forming Disks'' (cores to disks, hereafter c2d; Evans et al.~2003) Legacy Program. One of the first results to come from the c2d project was the discovery of Very Low Luminosity Objects (VeLLOs), objects embedded within dense cores with \\lint\\footnote{The internal luminosity, \\lint, is the luminosity of the central source and excludes luminosity arising from external heating.} $\\leq$ 0.1 \\lsun, most in cores previously classified as starless (Young et al.~2004; Bourke et al.~2006; Dunham et al.~2006, 2008, 2010a; Di Francesco et al.~2007; Terebey et al.~2009; Lee et al.~2009).  The discovery of such low luminosity protostars only exacerbated the existing luminosity problem.  Furthermore, Dunham et al.~(2008), Enoch et al.~(2009a), and Evans et al.~(2009) all examined protostellar luminosities from the c2d survey and showed that the protostellar luminosity distribution spans more than three orders of magnitude, is strongly skewed towards low luminosities (greater than 50\\% of protostars feature luminosities indicating $\\dot{M} \\la 10^{-6}$ \\msun\\ yr$^{-1}$), and is inconsistent with the standard model as well as with the modifications described above, which tend to increase the mass accretion rate and thus make the problem worse.  One possible resolution to the luminosity problem is the idea that mass accretion is not constant.  As noted by Kenyon et al.~(1990), either accretion rates that decline with time or accretion rates that are very low most of the time but occasionally very high could resolve the luminosity problem. The latter process, commonly referred to as episodic accretion, features prolonged periods of lower-than-average accretion punctuated by short bursts of higher-than-average accretion, a scenario already invoked to explain the luminosity flares seen in FU Orionis objects (Hartmann \\& Kenyon 1985).  First proposed by Kenyon et al.~(1990), such a solution was also suggested by Dunham et al.~(2008), Enoch et al.~(2009a), and Evans et al.~(2009) as a plausible explanation for both the large spread in observed luminosities and the significant population of sources at low luminosities.  Theoretical studies have provided several mechanisms for such a process in the embedded protostellar phase.  Hydrodynamical and MHD simulations have demonstrated that material accreting from a core can pile up in a circumstellar disk until the disk becomes gravitationally unstable and fragments into spiral structure and dense clumps, which are then driven onto the protostar in short-lived accretion bursts generated through the gravitational torques associated with the spiral structure (Vorobyov \\& Basu 2005b, 2006, 2010; Machida et al.~2011).  Numerical hydrodynamic simulations without self-consistent disk-core interaction but with gravitationally overstable massive disks have also shown quick migration of dense clumps into the disk inner region and probably onto the star (Boss 2002, Cha \\& Nayakshin 2011).   Alternatively, several authors have explored a scenario where a similar process involving gravitational instabilities in the outer disk triggers rapid accretion into the inner disk; the subsequent heating of the inner disk by this accretion activates magnetorotational instabilities (MRI) that then drive material onto the protostar in short accretion bursts (Armitage et al.~2001; Zhu et al.~2009a, 2009b, 2010). Other possible mechanisms have been proposed as well, including quasi-periodic magnetically driven outflows in the envelope causing magnetically controlled accretion bursts (Tassis \\& Mouschovias 2005), decay and regrowth of MRI turbulence (Simon et al.~2011), close interaction in binary systems or in dense stellar clusters (Bonnel \\& Bastein 1992; Pfalzner et al.~2008), and disk-planet interaction (Lodato \\& Clarke 2004; Nayakshin \\& Lodato 2011).  Direct observational evidence for episodic mass accretion bursts in the embedded phase is less clear.  Several Class I sources, including V1647 Ori (e.g., \\'{A}brah\\'{a}m et al.~2004; Acosta-Pulido et al.~2007; Fedele et al.~2007; Aspin et al.~2009), OO Serpentis (K\\'{o}sp\\'{a}l et al.~2007), [CTF93]216-2 (Caratti o Garatti et al.~2011), and VSX J205126.1 (Covey et al.~2011; K{\\'o}sp{\\'a}l et al.~2011), have undergone recent flares attributed to accretion bursts.  However, these sources only flared in \\lbol\\ by factors of $\\sim$ $2 - 10$, and all appear to be very late Class I sources near the end of the embedded phase with little remaining envelope material. On the other hand, the detection of silicate features and CO$_2$ ice bands in absorption in several known FU-Ori-like objects suggests the presence of massive envelopes (Quanz et al.~2007) and one of these objects, RHO~1B, has recently brightened by a factor of 1000 (Staude \\& Neckel 1991). Episodic mass ejection is seen in jets ejected from some protostellar systems, suggesting an underlying variability in the mass accretion, although the combination of jet velocities and spacing between knots often suggest shorter periods of episodicity than found by the above theoretical studies (e.g., Lee et al.~2007; Devine et al.~2009).  Additionally, several Class 0 sources drive strong outflows implying higher average mass accretion rates than expected from their current luminosities (\\andre\\ et al.~1999; Dunham et al.~2006, 2010a; Lee et al.~2010) and Watson et al.~(2007) showed a mismatch between the accretion rates onto the disk and protostar of NGC 1333-IRAS 4B, although their results are currently under debate (Herczeg et al.~2011). Finally, White \\& Hillenbrand (2004) showed that a sample of optically visible Class I sources in Taurus have very low accretion rates comparable to values observed for typical Class II sources rather than those expected for accreting Class I objects, although it is unclear if this result is simply a consequence of the Class I objects in their study being optically visible and thus very near the end of the embedded stage.  In this paper we test the hypothesis that the accretion process predicted by Vorobyov \\& Basu (2005b, 2006, 2010), which includes both accretion rates that decline with time and episodic accretion bursts, can resolve the luminosity problem and match the observed properties of embedded protostars.  We couple the hydrodynamic simulations presented by Vorobyov \\& Basu with radiative transfer models to calculate the observational signatures of cores collapsing to form protostars in the manner predicted by these simulations, and use these results to directly compare to observations.  This work is a direct follow-up to two previous studies.  In the first, Young \\& Evans (2005; hereafter Paper I) used a one-dimensional dust radiative transfer code to calculate the observational signatures of cores undergoing inside-out collapse following Shu (1977).  They followed three different cores with initial masses of 0.3, 1, and 3 \\msun\\ and showed that such models match only the upper end of the protostellar luminosity distribution, reconfirming the luminosity problem.  In the second study, Dunham et al.~(2010b; hereafter Paper II) revisited the models from Paper I with a two-dimensional radiative transfer code and showed that including improved dust opacities, a circumstellar disk and rotationally flattened envelope structure, mass-loss, and outflow cavities improved the match to observations but did not fully resolve the luminosity problem, whereas a toy-model representation of episodic accretion could in fact resolve the luminosity problem.  However, the latter conclusion is only tentative since the episodic accretion was included in a very simple manner where all mass accreted from the core builds up in the disk until the disk reaches 20\\% of the protostellar mass, at which point the accretion rate from the disk onto the star jumps from 0 \\msun\\ yr$^{-1}$ to $10^{-4}$ \\msun\\ yr$^{-1}$ until the disk is fully drained of its mass.  The accretion rate from the disk onto the star then drops back down to 0 \\msun\\ yr$^{-1}$ and the cycle begins anew.  This simplification likely exaggerates the fraction of total mass accreted in bursts, as pointed out by Offner \\& McKee (2011) (see also \\S \\ref{sec_discussion_bursts}).  Here we will revisit the Paper II results using accurate predictions for the evolution of collapsing cores from hydrodynamical simulations.  This paper is complementary to several other recent studies.  Myers et al.~(1998) presented simple radiative transfer calculations predicting observational signatures of collapsing cores, including an exponentially declining accretion rate with time and the effects of mass-loss, and showed that such models could generally reproduce the median observed protostellar luminosities but not the full range of values.  However, their model evolution is not based on a fully self-consistent analytic or numerical model.  Lee (2007) modified the Paper I model to include episodic accretion in a very simple manner in order to study the effects such a process has on the chemical evolution of collapsing cores.  Vorobyov (2009b) compared the distribution of mass accretion rates in their simulations featuring episodic accretion (Vorobyov \\& Basu 2005b, 2006) to those inferred from the luminosities of protostars in the Perseus, Serpens, and Ophiuchus molecular clouds compiled by Enoch et al.~(2009a) and showed that their simulations reproduced some of the basic features of the observed distribution of mass accretion rates.  However, the observed distribution of accretion rates is quite uncertain since it is calculated from the observed protostellar luminosity distribution with assumed values for the protostellar masses and radii and with the assumption that all observed luminosity is accretion luminosity.  Finally, Offer \\& McKee (2011) presented analytic derivations of the protostellar luminosity distribution for different models and concluded that models that tend toward a constant accretion time rather than constant accretion rate produce a greater spread in luminosities and are in better agreement with observations, similar to the result obtained by Myers (2010) that analytic models with accretion rates that increase with mass can at least partially resolve the luminosity problem.  However, both Offner \\& McKee (2011) and Myers (2010) simply compared theoretical luminosities with the observed protostellar luminosity distribution, whereas this study uses radiative transfer calculations to ``observe'' the underlying theory (in this case, the Vorobyov \\& Basu simulations) in a manner consistent with, and with direct comparison to, the observations.  The organization of this paper is as follows.  A brief description of the models is provided in \\S \\ref{sec_model}, focusing on the hydrodynamical simulations in \\S \\ref{sec_model_sim} and the radiative transfer models in \\S \\ref{sec_model_radtrans}.  \\S \\ref{sec_comparing} describes how the models are turned into observational signatures and the observational dataset to which the models are compared, while \\S \\ref{sec_results} describes the basic results.  A discussion of these results is presented in \\S \\ref{sec_discussion}, focusing on the degree to which the models resolve the luminosity problem in \\S \\ref{sec_discussion_lum}, the match between observed and model bolometric temperatures in \\S \\ref{sec_discussion_tbol}, the duration of the embedded phase in \\S \\ref{sec_discussion_durations}, and the number, duration, and importance of the mass accretion bursts in \\S \\ref{sec_discussion_bursts}. Finally, a basic summary of our results and conclusions is given in \\S \\ref{sec_summary}. ",
        "conclusions": "\\label{sec_discussion}  \\subsection{Resolving the Luminosity Problem}\\label{sec_discussion_lum}  As shown above in Figures \\ref{fig_blt} and \\ref{fig_lboltbolhist}, the models considered here reproduce the full range of observations in \\lbol\\ $-$ \\tbol\\ space and provide a reasonable match to the observed protostellar luminosity distribution.  Thus we conclude that the accretion process predicted by the Vorobyov \\& (2005b, 2006, 2010) simulations of collapsing cores resolves the luminosity problem, although we caution that these results must be revisited in the future once simulations that feature smaller sink cells closer to the stellar surface and fully capture both the physics in the inner disk and the actual protostellar accretion rates are possible (see \\S \\ref{sec_simulations_accretion} for further discussion). In Paper II we argued that episodic accretion is both necessary and sufficient to resolve the luminosity problem, but that conclusion is subject to uncertainty since the actual prescription for episodic accretion included in that paper was quite simple and did not fully capture the accretion process in the simulations. Here we have coupled the exact evolution of the collapsing cores with radiative transfer calculations and demonstrated that the Vorobyov \\& Basu simulations resolve the luminosity problem, although we defer the question of whether or not episodic accretion itself is necessary to Section \\ref{sec_discussion_bursts} below.  We also note here that these models predict a smooth distribution in \\lbol\\ $-$ \\tbol\\ space, whereas the models in Paper II featured white ``excluded zones'' where the models spent no time but sources were observed to exist. The fact that only 3 initial mass cores were considered in Paper II (0.3, 1, and 3 \\msun) was argued to artifically create these zones; the increased sampling of core masses in this paper confirms this argument.  Most of the remaining discrepancy between the model predictions and observations of protostellar luminosities actually appears in the form of a ``reverse luminosity problem'', evident in both Figures \\ref{fig_blt} and \\ref{fig_lboltbolhist} as an overabundance of time spent at \\lbol\\ $\\la$ 0.1 \\lsun\\ compared to observations.  Indeed, a K-S test on the observed and model luminosity distributions that only compares the distributions above 0.1 \\lsun\\ returns a value of 0.87, slightly higher than the value of 0.85 obtained for the full distributions.  What causes the disagreement at low luminosities? Only 7\\% of the observed sources have \\lbol\\ $\\leq 0.1$ \\lsun\\ whereas the models spent 21\\% of the total time at such luminosities.  Furthermore, the majority of this time (19\\% out of 21\\%) is spent at \\tbol\\ $\\leq 100$ K, where there are no observed sources.  At least some, and possibly all, of this discrepancy can be explained by observational completeness limits.  As noted above, the c2d survey is generally complete to protostars with \\lbol\\ $\\ga$ 0.05 \\lsun, although the exact completeness limit varies depending on distance and the detailed shape of the SED of each source.  Indeed, several extremely low luminosity sources undetected in the \\emph{Spitzer} c2d survey, with both internal and bolometric luminosities significantly below 0.1 \\lsun, have recently been discovered, most through detections of outflows driven by cores with no associated \\emph{Spitzer} c2d sources (Chen et al.~2010; Enoch et al.~2010; Dunham et al.~2011; Pineda et al.~2011).  With such low luminosities at least some of these sources may be first hydrostatic cores.  Sensitive interferometer outflow surveys and very deep \\emph{Herschel} and \\emph{James Webb Space Telescope} infrared surveys directed towards cores currently classified as starless are needed to fully identify and characterize the population of such extremely low luminosity protostars and/or first cores before an accurate comparison to the models can be made for luminosities below 0.1 \\lsun.  \\subsection{Bolometric Temperatures}\\label{sec_discussion_tbol}  As noted above in \\S \\ref{sec_results}, the models do not provide a good match to the observed \\tbol\\ distribution, with most of the discrepancy in a population of embedded objects at high \\tbol\\ ($\\ga 1000$ K) predicted by the models but lacking in the observations. In the models, most of this time spent at \\tbol\\ $\\ga 1000$ K arises when \\rdisk\\ is larger than a few hundred AU and the line-of-sight does not pass through the disk.  As described in detail in \\S \\ref{sec_model_radtrans}, we adopt analytic profiles for the disk and core density profiles since the hydro simulations do not provide the full vertical density structure.  In order to do this, we define the core inner radius to be equal to the disk outer radius so that there is no overlap between where the disk and core density profiles are defined.  However, as a consequence of this method, large cavities devoid of material exist above the surface of the disk but within the core inner radius, and these cavities increase as the disk sizes increase.  Lines of sight that pass through these cavities have reduced optical depths, allowing more short-wavelength emission to escape and thus increasing \\tbol.  In reality, such large cavities are unlikely to exist; instead, the disk and core density profiles should smoothly join together.  As we have mentioned in \\S \\ref{sec_model_radtrans}, an improved methodology that reconstructs the disk and core vertical structure and incorporates this exact structure into the radiative transfer calculations (rather than adopting simple analytic profiles) is currently under development and will be presented in a future paper.  This method will result in a more accurate distribution of material above the disk surface and will likely remove much of the discrepancy between observed and model values of \\tbol.  Since the distribution of luminosities is set mainly by the accretion rates and protostellar masses, we argue that including a more accurate physical structure should not significantly alter our results on the resolution of the luminosity problem, although this will be explicitly tested in a future paper.  We also note here the possibility that these models contain too much rotation and angular momentum, since this would cause both the rotational flattening of the cores and sizes of the disks to be overestimated, resulting in the models overpredicting \\tbol\\ compared to observations.  As discussed in some detail in \\S \\ref{sec_model_initial} above we do not consider this to be very likely, but we acknowledge that it is a possibility since $\\beta$, the ratio of rotational to gravitational energy, is not a directly observable quantity and observed ranges of $\\beta$ could be overestimated if infall and/or outflow motion is mistakenly attributed to rotation.  Even after including a more accurate physical structure, some discrepancy between observed and model values of \\tbol\\ may remain, particularly once the effects of outflow cavities are included.  Figure \\ref{fig_tbolsed}, which plots an example spectral energy distribution (SED) for a model with high \\tbol, shows that such models feature a double-peaked SED.  This figure also demonstrates that, from optical wavelengths to either the 24 or 70 \\um\\ bands probed by \\emph{Spitzer} observations, such models resemble transition disk Class II sources (sometimes also referred to as cold disk sources; see \\merin\\ et al.~[2010] and references therein).  Given that our models predict that 40\\% of embedded sources are classified as Class II by \\tbol, and that the observed fraction of Class 0+I to Class II sources is 0.19 when classifying via extinction-corrected values of \\tbol\\ (Evans et al.~2009), our models predict that 12\\% of Class II young stellar objects are actually embedded objects with SEDs like that shown in Figure \\ref{fig_tbolsed}.  While this is consistent with the upper range of the observed fraction of Class II sources with transition disk SEDs (3\\% -- 12\\%; \\merin\\ et al.~2010; Furlan et al.~2011), the models clearly overpredict this fraction, as described above.  Nevertheless, we caution that a small fraction of Class II sources with transition disk SEDs could in fact be embedded sources.  \\begin{figure} \\epsscale{1.0} \\plotone{f11.eps} \\caption{\\label{fig_tbolsed}Example of a Spectral Energy Distribution (SED) for a model with high \\tbol.  This particular SED is for the $i=$ 45$^{\\rm o}$ line-of-sight at 210,000 yr into the collapse of model 3, when the core mass is 0.9649 \\msun, the disk mass is 0.1656 \\msun, the protostellar mass is 0.5122 \\msun, and the disk radius is 744 AU.  This SED has \\tbol\\ $= 1037$ K. The solid line plots the model SED over all wavelengths, while the points show the model SED at standard optical and near-infrared wavelengths and the $3.6-70$ \\um\\ wavelengths provided by the \\emph{Spitzer Space Telescope}.  The points have been shifted down by a factor of 10 for clarity.} \\end{figure}  If the above statement is true, why are they missing from the observed sample of embedded sources?  As discussed in detail in Paper II, whether or not a population of embedded objects with high \\tbol\\ exists remains an open question.  The Evans et al.~(2009) sample is based primarily on the association of \\emph{Spitzer} sources featuring rising SEDs and red colors with millimeter continuum sources (Enoch et al.~2009a; Dunham et al.~2008) and is likely biased against such objects since they would often not meet these criteria and would instead be assumed to be chance alignments between millimeter continuum sources and background sources and/or later-stage YSOs.  Furthermore, the extinction corrections applied by Evans et al.~(2009) to the embedded sources are average extinctions over each individual cloud rather than true line-of-sight extinctions, and thus may underestimate the true extinction since current, active star formation (and thus the position of the youngest, embedded sources) is associated with the densest parts of molecular clouds (e.g., Heiderman et al.~2010; Lada et al.~2010).  As shown in Paper II, such underestimates could, in the worst cases, artifically lower \\tbol\\ from several thousand K to several hundred K.  Future work must revisit the observational samples and carefully evaluate whether or not a population of embedded sources with high enough \\tbol\\ to be classified as Class II or Class III exists.  \\subsection{Duration of the Embedded Phase}\\label{sec_discussion_durations}  \\begin{figure} \\epsscale{1.0} \\plotone{f12.eps} \\caption{\\label{fig_duration}\\emph{Left:} $t_{emb}$ vs.~\\mstar\\ for the 23 models listed in Table \\ref{tab_models} and considered in this paper. \\emph{Right:}  Average \\mdotstar\\ over the embedded phase duration vs.~\\mstar\\ for the same models.} \\end{figure}  For the 23 models listed in Table \\ref{tab_models} and considered in this paper, the duration of the embedded phase ($t_{emb}$) ranges from 0.052 (model 19) to 0.951 Myr (model 11).  For each model, Figure \\ref{fig_duration} plots both $t_{emb}$ (left panel) and the average accretion rate over the embedded phase duration (right panel, calculated as final stellar mass divided by $t_{emb}$) vs.~the final stellar mass produced by the model.  The duration of the embedded phase increases approximately linearly with the final stellar mass produced, and thus the average mass accretion rate (not to be confused with the \\emph{instantaneous} mass accretion rate, which is highly variable [see Figure \\ref{fig_mdot}]) is approximately constant and does not depend on final stellar mass.  Offner \\& McKee (2011) argued that models that tend toward constant accretion time rather than constant accretion rate are necessary to match the observed protostellar luminosity distribution, consistent with the conclusion by Myers (2010) that accretion rates that increase with mass can at least partially resolve the luminosity problem.  We disagree that such models are necessary.  Our models tend toward a constant average accretion rate of $\\sim 1-3 \\times 10^{-6}$ \\msun\\ yr$^{-1}$ rather than a constant accretion time yet still provide an excellent match to the observed protostellar luminosity distribution.  The average duration of the embedded phase, weighted by final stellar mass as described in \\S \\ref{sec_results}, is 0.12 Myr.  In contrast, Evans et al.~(2009) derived an embedded phase lifetime of $t_{emb} = 0.44$ Myr based on the ratio of Class 0+I sources to Class II sources in the c2d sample and the assumption of a Class II lifetime of 2 Myr.  Thus these models predict a significantly shorter $t_{emb}$ than suggested by recent observations.  However, a number of caveats apply to this comparison:  (1) We have taken the Class I/II boundary to be the point at which 10\\% of the initial core mass remains and terminate the models at this point.  In reality the exact point at which to set this Class boundary is uncertain and could easily shift the duration of the embedded phase by factors of $\\sim$ 2 in either direction; (2) The observationally determined $t_{emb}$ is pinned to a Class II lifetime of 2 Myr, but the uncertainty in this lifetime is about 1 Myr (50\\%; see discussion in Evans et al.~2009); and (3) The number of Class 0+I sources in the Evans et al.~sample may be overestimated.  The third point above is emphasized by three recent studies. First, van Kempen et al.~(2009) observed 22 Class I sources in Ophiuchus in \\hcop\\ $J = 4-3$, a tracer of warm and dense gas, and showed that 11 (50\\%) are not detected and thus show no evidence of being surrounded by a dense core. Second, McClure et al.~(2010) used their revised extinction law and classification method to show that greater than 50\\% of the Class I sources in Ophiuchus are highly extincted disk sources no longer embedded within cores.  Third, Heiderman et al.~(2010) observed 53 Class I sources in a variety of nearby star-forming regions in \\hcop\\ $J = 3-2$, a tracer of dense gas, and showed that 31 (58\\%) are not detected and thus not associated with a dense core.  All three studies likely represent upper limits to the true fraction of ``fake'' Class I sources.  In the cases of van Kempen et al.~(2009) and McClure et al.~(2010) this is because there is substantial evidence that Ophiuchus is located behind an extinction screen that, if not properly accounted for, will redden source SEDs and artifically increase the number of Class I sources (see, e.g., Figure 12 of Evans et al.~2009).  In the case of Heiderman et al.~(2010) this is because their study specifically targeted suspiscious Class I sources located in low extinction regions of clouds. We thus consider 50\\% as an upper limit to the fraction of Class I sources in the Evans et al.~(2009) sample that are actually misclassified Class II sources.  Shifting 50\\% of the Class I sources to Class II would decrease the Evans et al.~(2009) observationally determined value of $t_{emb}$ from 0.44 Myr to 0.2 Myr.  Combined with the other caveats mentioned above, the duration of the embedded phase predicted by our models is marginally consistent with observations but should be revisited in the future as uncertainties in the observations are improved.  Finally, the Stage 0 ([\\mstar+\\mdisk] / [\\mstar+\\mdisk+\\mcore] $\\leq 0.5$; \\andre\\ et al.~1993) duration for the 23 models considered in this paper ranges from 0.009 (model 19) to 0.256 (model 15) Myr, with an average (again weighted by final stellar mass as described in \\S \\ref{sec_discussion}) of 0.027 Myr.  Compared to the total embedded duration of 0.12 Myr, our models predict that the Stage 0 phase is only 23\\% of the total embedded duration.  This is a natural consequence of the fact that these models feature average mass accretion rates that decrease with time (Vorobyov \\& Basu 2010; see Figure \\ref{fig_mdot}), thus the first 50\\% of the mass will accrete from the core faster than the second 50\\% of the mass.  Our results are generally consistent with recent observationally determined estimates of the lifetime of Stage 0 relative to the total embedded phase. For example, Enoch et al.~(2009a) found 39 Class 0 and 89 Class I sources in Perseus, Ophiuchus, and Serpens, giving a relative Stage 0 lifetime of 30\\% of the total embedded phase duration. Additionally, Maury et al.~(2011) found that between $9-12$ of the 57 protostars they identified in the Serpens South cluster (Gutermuth et al.~2008) were Class 0 sources, giving a relative Stage 0 lifetime of 16\\%$-$21\\% of the total embedded phase duration.  While the Enoch et al.~and Maury et al.~results are difficult to compare quantitatively since they use different classification methods that trace the underlying physical Stage to different degrees of accuracy (\\tbol\\ in the case of Enoch et al.~and position in \\lbol$-$\\mcore\\ space in the case of Maury et al.), the general agreement between our models and these observational results is encouraging.  \\subsection{Number, Duration, and Importance of Bursts} \\label{sec_discussion_bursts}  The 23 models considered in this paper feature between 0 (model 16) and 97 (model 15) accretion bursts, where the exact criteria for defining bursts is given in Appendix \\ref{app_resamp}.  The percentage of total time spent in bursts ranges from 0\\% (model 16) to 11.8\\% (model 15), and the percentage of total mass accreted in bursts ranges from 0\\% (model 16) to 35.5\\% (model 4). On average (where the average is weighted by final stellar mass as described in \\S \\ref{sec_results}), 1.3\\% of the total time is spent in bursts and 5.3\\% of the total mass is accreted in these bursts.  These values represent the statistical average values of the fraction of total time spent and mass accreted in bursts assuming a standard Kroupa IMF (see Section \\ref{sec_results} for details).  We caution that the exact values depend on the criteria used for defining bursts and would increase if a lower accretion rate floor were used (see Appendix \\ref{app_resamp}).  The simple, parameterized models presented in Paper II spend between $1.5\\%-2$\\% of their total time in bursts, consistent with the 1.3\\% of total time featured by these models.  However, in the Paper II models between $50\\%-91\\%$ of the final stellar mass accretes in bursts, in stark contrast to the 5.3\\% of total mass that accretes in bursts in the models considered here.  The explanation for this large difference lies in the detailed implementation of accretion bursts in Paper II.  As described in \\S \\ref{sec_intro}, a burst was triggered each time the ratio of \\mdisk\\ to \\mstar\\ exceeded 0.2.  At this point the accretion rate onto the star was increased from 0 \\msun\\ yr$^{-1}$ to $10^{-4}$ \\msun\\ yr$^{-1}$ until the disk was fully drained of mass, allowing the cycle to begin anew.  However, in reality only the most extreme bursts reach accretion rates of $10^{-4}$ \\msun\\ yr$^{-1}$, as evident from Figure \\ref{fig_mdot}. Furthermore, only a very small amount of the mass in the disk accretes in a given burst (about $10^{-2}$ \\msun, on average), whereas in Paper II the assumption that the disk was fully drained in each burst led to situations where up to $0.1-0.2$ \\msun\\ were accreting in single bursts.  The models considered here, based on actual hydrodynamic simulations rather than simple parameterizations, are significantly more realistic.  Comparison to the results presented in Papers I and II demonstrate that the accretion process predicted by the Vorobyov \\& Basu (2005b, 2006, 2010) simulations essentially resolves the luminosity problem inherent in models with constant mass accretion.  As first noted by Kenyon et al.~(1990), there are two types of non-steady mass accretion that could potentially resolve the luminosity problem: (1) accretion rates that start high and then decrease with time, and (2) generally low accretion rates punctuated by short, episodic bursts of high accretion.  Figure \\ref{fig_mdot} clearly illustrates that the Vorobyov \\& Basu simulations feature both declining accretion rates with time \\emph{and} short, episodic accretion bursts.  Which of these is responsible for resolving the luminosity problem?  \\begin{figure} \\epsscale{1.0} \\plotone{f13.eps} \\caption{\\label{fig_lboltbolhist_avg}Histograms showing the fraction of total sources (observations; solid filled histogram) and fraction of total time spent by all models after time-averaging as described in \\S \\ref{sec_discussion_bursts} (dashed unfilled histogram; calculated from Equation \\ref{eq_bins}) at various \\lbol\\ (left) and \\tbol\\ (right).  The binsize is 1/3 dex in both quantities.  For the observations, only the 112 embedded sources (plotted as filled circles on the BLT diagrams) are included.} \\end{figure}  Given that, on average, only 5.3\\% of the total mass accretes in bursts, one might suspect that it is the declining accretion rates with time rather than the bursts that lower model luminosities and improve the match to observations.  However, this 5.3\\% excludes all of the mass that accretes in lower-amplitude accretion rate increases that do not meet the criteria for a burst as defined in \\S \\ref{app_resamp}.  Thus, to properly evaluate whether the bursts are required in order to resolve the luminosity problem, we have time-averaged the accretion rates to filter out the effects of the bursts and variability and re-run all 23 models. We averaged all models over 20,000 yr durations unless the total model duration was less than 0.2 Myr, in which case we decreased the averaging window to either 10,000 or 5,000 yr in order to preserve at least 10 timesteps between the formation of the protostar and the end of the embedded phase.  Figure \\ref{fig_lboltbolhist_avg} plots histograms comparing the fraction of total time the time-averaged models spend at various \\lbol\\ and \\tbol\\ to the observed distributions, similar to Figure \\ref{fig_lboltbolhist}.  While the time-averaged models provide a much better match to observations than models with constant mass accretion (see Papers I and II for such models), comparing Figures \\ref{fig_lboltbolhist} and \\ref{fig_lboltbolhist_avg} shows that the time-averaged models feature a small shift to higher luminosities and do not do quite as good of a job resolving the luminosity problem.  This is confirmed by a K-S test on the observed and time-averaged model luminosity distributions, which returns a value of 0.59, lower than the value of 0.85 returned for the original models.  \\begin{figure} \\epsscale{1.0} \\plotone{f14.eps} \\caption{\\label{fig_arate_avg}\\mdotstar\\ vs.~time for the original (thin black line) and time-averaged (thick gray line) timesteps for Model 11. A window of 20,000 yr was used for the time-averaging.  Note that the variability and bursts have not been fully filtered out.} \\end{figure}  These results indicate that the declining accretion rates are not solely responsible for resolving the luminosity problem, a finding that is only reinforced by the fact that, even with the large windows over which we have time-averaged, the variability and bursts are not fully filtered out\\footnote{In theory we could adopt even larger averaging windows, but to do so would decrease the total number of timesteps below 10 for many models and risk not fully sampling the decline of the accretion rate with time.} (see Figure \\ref{fig_arate_avg} for an example).  We thus conclude that it is a combination of both the accretion rates that decline with time and the variability and episodic bursts that resolve the luminosity problem. We consider this to be a plausible result given that the Vorobyov \\& Basu simulations self-consistently predict both, and we argue that the fact that the bursts are required is in general agreement with the other indirect evidence for accretion variability and bursts described in \\S \\ref{sec_intro}.    In this paper we have coupled radiative transfer models with hydrodynamical simulations of collapsing cores predicting accretion rates that both decline with time and feature episodic accretion bursts caused by fragments torqued onto the protostar.  We have calculated the time evolution of standard observational signatures (\\lbol, \\tbol, and \\lbolsmm) for cores collapsing following these simulations.  We have compared our results to a database of 1024 YSOs containing 112 embedded protostars recently compiled by Evans et al.~(2009).  We summarize our main conclusions as follows:  \\begin{enumerate} \\item The hydrodynamical simulations presented by Vorobyov \\& Basu (2005, 2006, 2010) reproduce the full spread of observed embedded protostars in a diagram of \\lbol\\ vs.~\\tbol.  The models resolve the luminosity problem and provide a reasonable match to the observed protostellar luminosity distribution (K-S value of 0.85). The models predict a large number of sources at very low ($\\la 0.1$ \\lsun) luminosities absent in the observations due to the observational sensitivity limit; removing such low luminosities from the comparison between models and observations increases the K-S value to 0.87.  The models predict that only 0.2\\% of the total time is spent at \\lbol\\ $\\geq$ 100 \\lsun; a larger dataset than the 112 protostars considered here is necessary to test this prediction. Time-averaged models that filter out the accretion variability and bursts do not provide as good of a match to the observed luminosity problem, suggesting that the bursts are required. \\item The models do not provide a good match to the distribution of observed \\tbol\\ for embedded protostars (K-S value of 0.42).  Instead, the models predict a subtantial population of embedded protostars at \\tbol\\ $\\ga$ 1000 K and thus classified as Class II or Class III sources.  Most of this discrepancy arises from the method by which we adopted analytic disk and core density profiles and will be alleviated with future models that incorporate the exact physical structure from the hydro simulations, but some of the discrepancy may also be due to a population of embedded protostars with high values of \\tbol\\ missing from the current database of protostars due to the selection criteria applied to construct this database.  The planned future models will not significantly alter the model luminosity distribution since this is primarily determined by the accretion rates and protostellar masses, not by the detailed physical structure adopted in the radiative transfer calculations. \\item The models reproduce the full spread of sources in a plot of \\lbol\\ vs.~\\mcore, including the existence of very low luminosity protostars with \\lbol\\ $\\sim 0.1-0.2$ \\lsun\\ but relatively high core masses of $1-2$ \\msun. \\item The duration of each model is approximately proportional to the final stellar mass produced, yet these models resolve the luminosity problem and provide an excellent match to the observed protostellar luminosity distribution.  This is in contrast to recent results in the literature claiming that models that tend toward constant accretion time rather than constant accretion rate are necessary to match the observed distribution. \\item The IMF-weighted average duration of the embedded phase in our models is 0.12 Myr, whereas Evans et al.~(2009) recently determined the embedded phase duration to be 0.44 Myr.  We have suggested a number of possible means by which these two estimates of the embedded phase duration may be reconciled.  The IMF-weighted average model Stage 0 duration is 0.027 Myr, or 23\\% of the total embedded phase duration.  Observationally determined values based on the ratio of Class 0 to Class I sources range from $16\\%-30\\%$ (Enoch et al.~2009a; Maury et al.~2011).  Our models are consistent with this range. \\item On average, these models spend 1.3\\% of their total time in accretion bursts, during which time 5.3\\% of the final stellar mass accretes.  In the most extreme models these values reach 11.8\\% and 35.5\\%, respectively. Thus accretion is not truly ``episodic'' since it actually occurs at all times rather than only in episodes.  A better description is that accretion is ``variable with episodic bursts.'' \\end{enumerate}  Future work must concentrate on improving the accuracy and self-consistency of the hydrodynamical simulations and radiative transfer models, and on compiling a more complete and more accurate database of protostars and their observational signatures.  Nevertheless, we expect our primary conclusion that the Vorobyov \\& Basu (2005, 2006, 2010) simulations resolve the luminosity problem and match the observed protostellar luminosity distribution to remain unchanged to such future improvements."
    },
    "1112/1112.4430_arXiv.txt": {
        "abstract": " ",
        "introduction": "With the advent of future wide-field surveys, weak gravitational lensing will become a premier probe of cosmology and an important tool to constrain dark energy, neutrinos and the initial conditions (see for instance \\cite{2010AnP...522..324D,Giannantonio:2011ya}). In order to fully exploit the potentiality of the  convergence and shear fields, it will be important to use their whole statistics. In particular, the lensing bispectrum represents a complementary probe to the power spectrum, as it will provide constraints  that are comparable to those obtained from the power spectrum alone \\cite{Berge:2009xj}.  One of the primary interests of wide-field surveys is to look for primordial non-Gaussianities generated in the early universe and the lensing bispectrum represents a natural tool to capture such features \\cite{Bernardeau:2002fc,Takada:2003sv,2011MNRAS.411..595P,Jeong:2011rh,2011arXiv1107.1656S}. Given the current and forecasted constraints on $\\fNL$, there is a plethora of second-order effects intervening between the initial conditions and the observations that could be potentially relevant. Some of these effects, such as the second-order growth of matter fluctuations in Eulerian perturbation theory, the Born correction, the lens-lens coupling and the non-linear conversion between the galaxy shape distortion and the (observable) reduced shear, have been thoroughly studied in the past years and are expected to dominate  the lensing 3-point statistics on small-angular scales \\cite{1997A&A...322....1B,Cooray:2000uu,2002A&A...389L..28B,Cooray:2002mj,2003MNRAS.344..857T,Dodelson:2005zj,2006JCAP...03..007S,2010A&A...523A..28K,Dodelson:2005rf,Schneider:1997ge}.  When studying weak lensing one usually restricts the analysis to small angular separations (or large multipole moments), which is justified by the fact that so far cosmic shear surveys have covered only a limited portion of the sky. However, nearly full-sky surveys are currently under preparation and on large angular scales general relativistic  second-order couplings will become relevant. These corrections are usually neglected because they are suppressed by the ratio between the scale probed and the Hubble scale, but are important on angular scales of the order of the angular diameter distance to the source. These are analogous to the second-order effects from general relativity affecting the CMB bispectrum when at least one of the scales probed is larger than the horizon at recombination \\cite{Bartolo:2004ty,Boubekeur:2008kn,Pitrou:2010sn}. In the CMB, these effects have been recently found to be negligible for the contamination of Planck searches for a primordial local signal in the squeezed limit \\cite{Creminelli:2011sq,Bartolo:2011wb}. In case of cosmic shear observations, these effects however cannot be a priori neglected. The aim here is to provide an exhaustive evaluation of the impact of these contributions to observations. Those results extend recent studies on the impact of relativistic corrections to the observations of large-scale galaxy clustering in \\cite{Yoo:2009au,Bonvin:2011bg,Challinor:2011bk,Baldauf:2011bh,Yoo:2011zc}.  The complete study of all second-order effects in the cosmic shear, including the general relativistic ones, has been recently undertaken in \\cite{Bernardeau:2009bm} by solving at second order the Sachs equation  \\cite{1961RSPSA.264..309S}, which describes the deformation of the cross-section of a light bundle, mapping galaxy shapes into their angular images. The main results of this paper are reviewed in \\sect~\\ref{sec:GReffects}.   The goal of the current paper is to compute the bispectrum of the electric part of the cosmic shear from all second-order effects and compare the importance of the general relativistic corrections with the standard couplings. We review the computation of the shear power spectrum in \\sect~\\ref{sec:PS} where we also present one of the numerical shortcuts we will use throughout the paper, the second-order Limber approximation. The formal expression of the bispectrum, computed in \\sect~\\ref{sec:BS}, is determined by two types of non-linear contributions: the dynamical couplings, which depend on the particular metric solution of  Einstein's equations at second order, and the geometrical couplings, so-called because they depend on the geometry of the solution of the Sachs equation at second order. Among the latter we also include the corrections coming from the inhomogeneities of a fixed-redshift source plane. The calculations of \\sect~\\ref{sec:BS} involve multiple integrals and complicated manipulations of spherical harmonics and Wigner symbols. (In appendix~\\ref{app:spart} one can find useful relations for these calculations.) We are afraid to say that those calculations are extremely lengthy and technical. Impatient readers can skip this section and go directly to \\sect~\\ref{sec:results}. In this section we present our results concentrating on the squeezed limit, where one of the multipole  moments $l$ involved in the bispectrum is small and we discuss the functional forms of the resulting terms as well as their relative importance. In \\sect~\\ref{sec:primordial} we compare our results with the bispectrum generated by  primordial non-Gaussianities of the local type and we compute their contamination to a primordial $f_{\\rm NL}^{\\rm loc}$. Finally,  in \\sect~\\ref{sec:conclusion} we conclude and discuss the results of the paper. ",
        "conclusions": "\\label{sec:conclusion}     In this paper we present an exhaustive calculation of the full-sky bispectrum of the shear field as it would be built by an observer measuring shapes of galaxies at a given redshift. The calculation is performed at tree order (i.e.~neglecting loop corrections), assuming that the initial curvature perturbation $\\zeta$ obeys Gaussian statistics, such as after single-field inflation. Furthermore, we study the relative importance of the different couplings and we compare them with primordial non-Gaussianities of the local type.  Our derivation fully exploits and completes the recent study of Ref.~\\cite{Bernardeau:2009bm},  where the full-sky second-order expression of the shear field has been computed. This article derives all second-order terms in the metric fluctuations that contribute to the shear field and it is thus the natural starting point of the computation of the tree order bispectrum presented here. In this paper we focus our calculations on the bispectrum of the scalar, i.e.~electric, part of the shear field alone. This is somewhat restrictive since the second-order shear field contains a pseudo-scalar, i.e.~magnetic, mode. We leave the study of the $B$-mode and of the correlators where it is relevant for the future.  The computation of the bispectrum from the formal second-order shear expression is an involved exercise, which includes two steps. The first consists in the formal derivation of the bispectrum in harmonic space. The electric part of the shear field itself is defined from its harmonic space expression, in a way similar to the mode decompositions of the CMB polarization. Its bispectrum is then computed with the help of the decomposition properties of the spherical harmonics and their spin weighted extensions.  These results are presented in \\sect~\\ref{sec:BS} of the paper where the contribution of each term is fully computed. Then, the second part consists in numerically  evaluating the terms that we have found, in the context of the $\\Lambda$CDM model.  Two approximations are commonly employed for weak lensing calculations, namely the Limber approximation \\cite{1953ApJ...117..134L}, which consists in neglecting longitudinal modes, as they tend to average out along the line of sight, and the flat sky approximation, which consists in replacing a small portion of the spherical sky with a  2-dimensional plane.\\footnote{Note that these two approximations are not equivalent and one does not necessarily imply the other.} In order to be consistent with our full-sky treatment, it is a priori not possible to employ neither of these two approximations. Abandoning the Limber approximation makes the computations  numerically challenging due to the large number of integrations involved and, as a consequence, only a limited number of configurations could be exactly computed. In order to circumvent this problem we have used the higher-order Limber approximation discussed in \\cite{1992ApJ...388..272K}. Apart from for a specific type of contributions, we have found that sufficiently high accuracy can be reached with the help of the second-order approximation. The procedure that we have used to implement this approximation to the angular power spectrum and the bispectrum is presented in appendix~\\ref{app:limber}.  Our results are described in Sect. 5. In this section we give a precise account of the relative importance of the various contributing terms, as well as of their dependence with the multipole configuration. We distinguish the different contributing terms as geometrical and dynamical, depending on whether they originate from pure general relativistic effects on the line-of-sights or due to the dynamical evolution of the metric field. Note however that this distinction is somewhat arbitrary, since it depends on the gauge choice. Furthermore, in each case we distinguish the standard terms as those which survive at small angular scale, from the non-standard ones that are mainly relevant at large angular scales.  We discuss the contribution to the bispectrum of the different nonlinear couplings focussing on the squeezed limit configuration, where the non-standard general relativistic corrections are more likely to play a significant role. In \\sect~\\ref{sec:results} we give a precise account of the relative importance of the various contributing  terms as well as of their dependence with the multipole configuration. We find that the general relativistic effects, i.e. the non-standard coupling terms,  are typically smaller than the standard non-linear couplings. However their relative importance increases at smaller source redshifts. Among the new couplings, the dynamical couplings are always much smaller than the others, while those due to the inhomogeneity in the redshift of the source dominate. The latter can even become of the same order of the standard couplings when the redshift of the sources is below $0.5$. These results are summarized in fig.~\\ref{fig:bz1} and \\ref{fig:bzp0}.   In \\sect~\\ref{sec:primordial} we finally compute the corresponding level of contamination induced by the standard and non-standard terms on the amplitude of primordial non-Gaussianities of the local type. Note that, as in the squeezed limit the standard non-linear couplings depend on the angle between the short and long mode, they differ significantly from a local signal. For $z_S \\gtrsim0.8$, the contamination to $f_{\\rm NL}^{\\rm loc}$ of the standard geometrical couplings amounts to $f_{\\rm NL}^{\\rm loc} \\gtrsim 100$. On the other hand, the relativistic corrections induce a non-Gaussianity which is mainly of the local type and can contaminate the search for a primordial signal by $f_{\\rm NL}^{\\rm loc} \\sim 1$ to $f_{\\rm NL}^{\\rm loc}\\gtrsim 10$ when the source redshifts vary from $1$ to $0.3$.   Our investigations did not go as far as computing the signal-to-noise ratio for the bispectrum induced by the general relativistic corrections. However, given the much larger amplitude of the standard contributions, we expect it to be rather small. This might not be the case when, besides the only $E$-mode three-point correlators we have investigated here, quantities involving $B$-modes are also included. Relativistic corrections indeed generically induce  $B$-modes with an amplitude a priori  comparable to the $E$-modes~\\cite{Bernardeau:2009bm}. This is at variance with the standard dynamical couplings that induce second-order $E$-modes only. It ensures that the correlations between $E$ and $B$ modes are only non vanishing in the presence of general relativistic corrections making such quantities more likely to be observable. We leave such calculations for a future study."
    },
    "1112/1112.2600_arXiv.txt": {
        "abstract": "The AGILE scientific instrument has been calibrated with a tagged $\\gamma$-ray beam at the Beam Test Facility (BTF) of the INFN Laboratori Nazionali di Frascati (LNF). The goal of the calibration was the measure of the Point Spread Function (PSF) as a function of the photon energy and incident angle and the validation of the Monte Carlo (MC) simulation of the silicon tracker operation. The calibration setup is described and some preliminary results are presented. ",
        "introduction": "AGILE (Astro-rivelatore Gamma a Immagini LEggero) is a Small Scientific Mission of the Italian Space Agency (ASI) launched on April 2007 and dedicated to high-energy astrophysics \\cite{agimis}. The AGILE satellite is designed to detect and image photons in the 18 - 60 keV, 30 MeV - 50 GeV and 350 keV - 100 MeV energy bands with excellent spatial resolution, timing capability, and large field of view.\\\\ AGILE is the most compact ($\\approx 0.25 m^3$), light (120 kg for the instrument, 350 kg for the whole satellite) and low power ($\\approx 60 W$) scientific instrument ever developed for high-energy astrophysics.\\\\ The AGILE scientific payload (shown in Fig.\\ref{agidraw}) consists of three detectors with independent detection capability. The Gamma-Ray Imaging Detector (GRID) consists of a Si-W converter-tracker \\cite{st1} sensitive in the $\\gamma$-ray energy range 30 MeV - 50 GeV, a shallow ($1.5\\, X_0$ on-axis) CsI Calorimeter \\cite{minical} and a segmented AntiCoincidence system based on plastic scintillators \\cite{ac}.\\\\ In addition to the GRID, a coded-mask hard X-ray imaging system (SuperAGILE), made of a Si detector plane and a W mask, ensures coverage in the range $18-60 \\mathrm{keV}$ \\cite{superagile}.\\\\ The AGILE main feature is the combination of two co-aligned imaging detectors (SuperAGILE and GRID) sensitive in the hard X-ray and in the $\\gamma$-ray ranges with large field of view ($\\approx 1.0\\mathrm{sr}$ and $\\approx 2.5 \\mathrm{sr}$ respectively).\\\\ Moreover the CsI MiniCalorimeter (MCAL) can operate in stand alone \"burst mode\" in the $350\\,\\mathrm{keV}- 100\\,\\mathrm{MeV}$ range to detect GRB.\\\\ On ground and subsequently on flight calibrations of a detector are essential to the interpretation of its results. The purpose of the calibration of a scientific instrument is to reproduce, under controlled condition, the detector response in operation.\\\\ This paper describes the on-ground calibration of the silicon tracker and some results on the instrument performances derived by it.  \\begin{figure}[!ht] \\begin{center} \\includegraphics[width=5cm]{AgileFigure.eps} \\end{center} \\caption{A schematic view of the AGILE scientific instrument.} \\label{agidraw} \\end{figure} ",
        "conclusions": "This paper presents some preliminary results of the calibration of the AGILE ST at the BTF of the LNF in 2005.\\\\ The setup is described in detail as well as the calibration requirements. We discussed the problems encountered in exploiting the PTS originally designed and a novel approach devised to circumvent those problems: the phase analysis.\\\\ We concentrated on the measurements of the PSF presenting two possible definitions: the Gaussian and the 68\\%\\ PSF.\\\\ The calibration results are compared with the MC simulations for a broad set of variables, showing good consistency with some poorer agreement for 68\\%\\ PSF mainly at low energies.\\\\ These results give confidence on the use of the MC simulation in the untested part of the $\\gamma$-ray parameters (e.g. higher $E_\\gamma$) especially in flight conditions, i.e. without low energy background, and in the measurement of detector parameters, like absolute efficiency and energy resolution, that are difficult to measure without exploiting the PTS information."
    },
    "1112/1112.5216_arXiv.txt": {
        "abstract": "In this paper, a decay vacuum model $\\bar{\\rho}_\\Lambda=3\\sigma M_p^2H_0 H$ is revisited by detailed analysis of background evolution and perturbation equations. We show the imprints on CMB temperature and matter power spectrum from the effective coupling terms between dark sectors by comparing to the standard cosmological constant model and observational data points (WMAP7 and SDSS DR7). We find that the decay vacuum model can describe the expansion rate at late times as well as the standard cosmological constant model but it fails to simultaneously reproduce the observed CMB and matter power spectrum. Its generalization $\\bar{\\rho}_\\Lambda=3M_p^2(\\xi_1 H_0 H+\\xi_2  H^2)$ is also discussed. Detailed analysis of the background evolution shows that the dimensionless parameter $\\xi_{2}$ would be zero to avoid the unnatural 'fine tuning' and to keep the positivity of energy density of dark matter and dark energy in the early epoch. ",
        "introduction": "Since the accelerating expansion of the universe has been found from the measures of the luminosity-redshift relation $d_L(z)$ of type Ia supernovae (SN Ia) \\cite{Riess}, a cosmic component called dark energy was often introduced to explain the acceleration within the framework of general relativity. Now more and more evidence, such as cosmic microwave background (CMB) \\cite{Bennett,ref:wmap7}, baryon acoustic oscillations (BAO) \\cite{Eisenstein}, weak gravitational lensing \\cite{Jarvis} and x-ray clusters \\cite{Allen,Allen2}, indicated that  the universe is spatially flat and dominated by dark energy at present. Apart from dark energy models, modified gravity \\cite{Carroll,Trodden} can also explain the acceleration of the present universe. However, we just focus on dark energy model in this paper. Among the various dark energy models \\cite{Copeland}, including scalar field \\cite{Ratra1}, vector field\\cite{Zhang94,Kiselev}, holographic dark energy \\cite{Holographic}, Chaplygin gas \\cite{Kamenshchik} and so on. The cosmological constant model ($\\Lambda$CDM) \\cite{Weinberg} is the simplest one. However, as is well known, the $\\Lambda$CDM suffers from a fine tuning problem: the observed vacuum energy density of order $\\sim10^{-47}\\text{GeV}^4$ is about $10^{121}$ orders of magnitude smaller than the value expected by quantum field theory for a cutoff scale that is the Plank scale, and is still about $10^{44}$ orders smaller even for a cutoff scale that is the QCD scale \\cite{Copeland}. As an extension to $\\Lambda$CDM, the decaying vacuum (DV) dark energy model was  proposed \\cite{Borges1,Carneiro}, based on the incomplete quantum field theory in the curved 4-dimension space-time. In this model, the vacuum serves as  dark energy, whose energy density decays with the expansion of the universe leading to an additional production of the matter component. In the late-time with a quasi-de Sitter background, the vacuum density is proportional to the Hubble rate, $\\rho_\\Lambda(t) \\propto H(t)$. However, the equation of state for the vacuum is a constant value $w=p_\\Lambda(t)/\\rho_\\Lambda(t)=-1$, the same as that in the $\\Lambda$CDM model. Moreover, as an interesting feature, the late-time dynamics of the DV model is similar to $\\Lambda$CDM \\cite{Borges1,Carneiro}.  The quasar APM 08279+5255 at $z=3.91$ was used to examine the DV model \\cite{mltong2}, and it was found that the DV model can greatly alleviate the high redshift cosmic age problem existing in the $\\Lambda$CDM model. In order to distinguish the DV model from other dark energy models at the late-time Universe, the state finder and $Om$ diagnostics of the DV model were also presented in \\cite{mltong2}. Moreover, the DV model has been tested by $\\chi^2$ analysis using the observational data of SN Ia \\cite{Carneiro2}, a joint data from SN Ia, BAO and CMB\\cite{Carneiro3,Pigozzo}, and the joint data that the Gamma-ray bursts, Hubble rate and x rays in galaxy clusters were added \\cite{tongnoh}. It was found that, the DV model favors a relatively larger value of the matter density contrast, $\\Omega_m=(0.34\\sim0.43)$.  In this paper, compared to previous work \\cite{Pigozzo,tongnoh}, we take the radiation component into account in a more reasonable way. We will demonstrate the temperature anisotropies of CMB induced by the matter perturbations in the DV model with various values of $\\Omega_m$. The matter power spectrum is investigated.  This paper is structured as follows. In section \\ref{sec:bg}, the background evolution equations are given in a spatially flat FRW universe. We give a brief review of cosmological perturbation theory in section \\ref{sec:cp}. The main results of this paper are presented in section \\ref{sec:results}. In section \\ref{sec:tvv}, we give a brief discussion on a generalized form  $\\bar{\\rho}_{\\Lambda}=3M_p^2(\\xi_1 H_0 H+\\xi_2  H^2)$ and point out that this model is not viable. Section \\ref{sec:con} is the conclusion. ",
        "conclusions": "\\label{sec:con}  In this paper, a decay vacuum model $\\bar{\\rho}_\\Lambda=3\\sigma M_p^2H_0 H$ and its generalization $\\bar{\\rho}_\\Lambda=3M_p^2(\\xi_1 H_0 H+\\xi_2  H^2)$, a time variable cosmological constant model, are revisited.  At first, the background evolution equation in a spatially flat FRW universe containing cold dark matter, radiation, baryon and time variable cosmological constant is given. The relative departure from $\\Lambda$CDM model is minor, please see the left panel of Fig. \\ref{fig:evolution}. So to discriminate the decay vacuum model from $\\Lambda$CDM model, high redshift observations are needed. In the decay vacuum model case, an effective interaction between cold dark matter and vacuum can be introduced. Then the evolution of cold dark matter will depart from the conventional power law $a^{-3}$. And the large scale structure formation will be strongly different from that of $\\Lambda$CDM model. Though the baryon component evolves in the scaling $a^{-3}$, the background evolves different from $\\Lambda$CDM model for the effective interaction between cold dark matter and decay vacuum. Then the dynamic evolution would be modified. So the cosmological perturbations are taken into account. As results, the angular power spectrum of CMB and matter power are presented with different parameter values of cold dark matter abundance $\\omega_c$, please see Fig. \\ref{fig:clspk}. From this figure, one can conclude that CMB observations and matter power spectrum can distinguish the decay vacuum model from $\\Lambda$CDM model markedly. When $\\Omega_{c} h^{2}=0.2158$, i.e. $\\Omega_{c0}=0.4404$, the purple dashed line in Fig. \\ref{fig:clspk} is close to observational data points. It means that increasing the abundance of cold dark matter will depress the acoustic peaks to cosmic observational data points in this model. However, in the right panel of Fig. \\ref{fig:clspk}, one sees that increasing the abundance of cold dark matter will enhance the matter power spectrum at small scale but depress that at large scale. That makes it difficult to match observational data points. With these observations, this model would be ruled out. But to know in what kind of levels to rule out this model, testing this model with current available cosmic observational data sets, for example type Ia supernovae, baryon acoustic oscillation, full CMB and SDSS DR7 etc, would be interesting. Furthermore, a generalized vacuum model $\\bar{\\rho}_{\\Lambda}=3M_p^2(\\xi_1 H_0 H+\\xi_2  H^2)$ was discussed. From a detailed analysis, one can find that the parameter space of $\\xi_{2}$ is a very small negative dimensionless parameter. To keep the positivity of energy density of dark matter and dark energy at early epoch, the parameter $\\xi_{2}$ suffers from the fine tuning problem. So to avoid this unnatural condition, the $\\xi_{2}$ would be set to zero. Then it reduces to the decay vacuum model. In this sense, it would not be a viable dark energy model."
    },
    "1112/1112.0928_arXiv.txt": {
        "abstract": "{Many physical parameters change with time in star forming regions. Here we attempt to correlate changes in infall and outflow motions in high mass star forming regions with evolutionary stage using JCMT observations. } {From a sample of 45 high mass star forming regions in three phases of evolution, we investigate the presence of established infall and outflow tracers to determine whether there are any trends attributable to the age of the source.} {We obtained JCMT observations of HCO$^+$/H$^{13}$CO$^+$ J=4-3 to trace large scale infall, and SiO J=8-7 to trace recent outflow activity.  We compare the infall and outflow detections to the evolutionary stage of the host source (high mass protostellar objects, hypercompact HII regions and ultracompact HII regions). We also note that the integrated intensity of SiO varies with the full width at half maximum of the H$^{13}$CO$^+$.} {We find a surprising lack of SiO detections in the middle stage (Hypercompact HII regions), which may be due to an observational bias.  When SiO is detected, we find that the integrated intensity of the line increases with evolutionary stage.  We also note that all of the sources with infall signatures onto Ultracompact HII regions have corresponding outflow signatures as well.} {} ",
        "introduction": "High mass stars are a vital player in the evolution of galaxies. During their formation they inject energy into their surroundings in the form of outflows and ionizing photons. As they evolve their stellar winds continue to stir turbulence in the surrounding medium, and when they explode as supernovae they additionally enhance the metallicity of their environs.  The processes involved in this earliest stage, the formation of high mass stars is not nearly as well constrained as are the mechanisms responsible for lower mass stars. The pre main sequence evolution of low mass stars \\citep[i.e.][]{shu87,andre93} has been well constrained and statistically based lifetimes of the stages have come about, in part, due to observing large numbers of individual sources \\citep[i.e][]{Enoch09,Spezzi08}.  Similar studies for the early evolution of high mass stars (M$\\gtrsim 8$ M$_\\odot$) have proven problematic due to limiting factors such as the large average distances to high mass star forming regions and the clustered nature of their formation.  However studies such as \\citet{Beltran11} suggest high mass stars are likely to form though disk mediated accretion in a similar but scaled up version of lower mass star formation. An evolutionary sequence is developing, however progress is hampered by a lack of large sample surveys of the molecular gas dynamics in these region \\citep[see, for example][]{ZY07,Beuther07}. Evolution in low mass systems is often characterized by changes in the spectral energy distributions (SEDs). This is not as easily done in high mass systems \\citep{Molinari08}, which is why studying the gas dynamics is very important.   In broad terms, the evolution of massive protostars begins in an infrared dark cloud \\citep[IRDC, i.e.][]{Egan98}. The central condensation then begins heating its environment, at which point the core starts emitting at IR wavelengths and it becomes a high mass protostellar object \\citep[HMPO, i.e.][]{Sridharan02} within a hot core.  The protostar continues to gain mass and heat its environment eventually forming a hypercompact (HC) HII region \\citep[i.e.][]{Keto07} which then grows into an ultracompact (UC) HII region \\citep[i.e.][]{WC89}.  This sequence has been shown graphically in \\citet{Zapata10}.  The goal of this paper is to follow the larger scale (that observable with a single dish) outflow and infall structures surrounding sources in the later three of these four evolutionary stages (HMPO, HCHII and UCHII region), and determine whether any evolutionary trends are distinguishable.   In the case of low mass star formation, \\citet{andre93} conclude that outflow motions can be seen quite early in the Class 0 phase, and \\citet{Bontemps96} state that the outflows from Class 0 sources are more powerful than those from Class I sources at the same luminosity.  They also suggest that all deeply embedded low mass protostars should have outflows. As for infall in the low mass regime, inverse P-Cygni profiles indicative of infall have been seen for a decade now \\citep{difrancesco01,Zapata08,Furuya11}, and some observations even suggest that infall is seen perpendicular to the outflow direction \\citep{Arce04}. That these trends can be applied to higher mass protostars is starting to be explored \\citep[i.e.][]{LS11}.  Here we present HCO$^+$ and H$^{13}$CO$^+$ (J=4-3) observations to study the bulk infall signatures in high mass star forming regions.  We also present SiO (J=8-7) observations to trace recent outflow activity from the same sources. We have collected a sample of 45 high mass star forming regions in three evolutionary stages: HMPOs, and sources with HC and UCHII regions.  In Section \\ref{sec:observations} we present our observations and in Section \\ref{sec:method} we present our methodology.  In Section \\ref{sec:results} we present our results, which we discuss in Section \\ref{sec:discussion} where we also conclude. ",
        "conclusions": ""
    },
    "1112/1112.0575_arXiv.txt": {
        "abstract": "We consider issues related to the conformal mapping between the Einstein and Jordan frames in $f(R)$ cosmology. We consider the impact of the conformal transformation on the gauge of a perturbed system and show that unless the system is written in a restricted set of gauges the mapping could produce an inconsistent result in the target frame. Newtonian gauge lies within the restricted group but synchronous gauge does not. If this is not treated carefully it could in principle contaminate numerical calculations. ",
        "introduction": "\\noindent Extended gravity theories, where the Einstein-Hilbert Lagrangian density $\\mathcal{L}_{\\rm EH}=R$ is replaced by a more general function including terms of higher-order in derivatives of the metric ($R^2$, $R_{\\mu\\nu}R^{\\mu\\nu}$, $R_{\\alpha\\beta\\mu\\nu}R^{\\alpha\\beta\\mu\\nu}\\ldots$) and couplings to new dynamical degrees of freedom, have long been of interest in relativity. The last few decades have seen increasing applications of these models to cosmology. See \\cite{Capozziello:2011et,Nojiri:2006ri} and their references for an overview and further details on these models and their applications to cosmology.\\footnote{Recently \\cite{Biswas:2011ar} presented the most general action guaranteed to be free of ghosts when quantising perturbations on a classical vacuum; a viable model of extended gravity will be constrained to conform with this action.} In the last decade attention has focused on exploiting extended gravity to model dark energy without the need to introduce exotic particle species. A model frequently employed in this context is $f(R)$ gravity (see for example \\cite{DeFelice:2010aj,Nojiri:2010wj} for recent reviews), where the Einstein-Hilbert Lagrangian density is replaced with an arbitrary function of the Ricci scalar, $\\mathcal{L}=f(R)$, while the matter couples minimally to the metric and follows its geodesics in free motion. This representation of an extended gravity model is known as the ``Jordan frame''.  The action can be transformed to a variety of different forms. One of the most % common transformations is into the so-called ``Einstein frame'' where the action is manipulated to isolate a Ricci scalar of a new metric. The residual terms can be interpreted as an effective scalar field to which matter is non-minimally coupled and deflected from the geodesics of the metric. The field equations are otherwise those of standard general relativity. We review the model in the Jordan and Einstein frames in \\S\\ref{Section:Model}.  Transforming from the Jordan frame to the Einstein frame is extremely useful in the study of $f(R)$ gravity. We employ the metric formulation, in which the action is varied with respect to the metric alone, and in which the field equations are fourth-order. In the Einstein frame, as in standard GR, the theory is second-order -- a significant simplification. Aspects of the transformation have been controversial for some time (see for example \\cite{Faraoni:1998qx,Deruelle:2010ht,Quiros:2011wb,Capozziello:2011et} and further references in \\S\\ref{Section:Equivalence}). The discussion has centred upon the nature of the equivalence, and authors can be separated \\cite{Faraoni:1998qx,Capozziello:2011et} into two camps: those who feel the equivalence is ``physical'' and observables can be calculated in either frame, and those who feel the equivalence is mathematical in nature and that observables should be calculated in a chosen ``physical frame''. We briefly discuss this issue in \\S\\ref{Section:Equivalence}.  Modern cosmology is the study of Friedmann-Lema\\^itre-Robertson-Walker (FLRW) spacetimes, perturbed to linear or higher orders, and this provides an illustration of the major motivation for our paper. Numerical studies targeting the cosmic microwave background (CMB) and the matter power spectrum (such as \\cite{Bean:2006up,MGCAMB11,Gu:2011kk,Li:2011vk,Dossett:2011tn}) often employ modifications to Boltzmann codes such as CAMB \\cite{CAMB}. While it is common to employ a general parameterisation \\cite{MGCAMB11,Gu:2011kk,Li:2011vk,Dossett:2011tn} in the Jordan frame, working in the Einstein frame (e.g. \\cite{Schimd:2004nq,Bean:2006up}) provides a flexible alternative. Exploiting the Einstein frame allows us to consider any modified gravity model which possesses an Einstein frame and for which background solutions in the Jordan frame might be extremely difficult to find, either analytically or numerically. A recent example is the study of the string gas cosmology (see e.g. \\cite{Brandenberger:2011et}), in which the primordial power spectrum is frequently calculated in the Einstein frame as opposed to the string frame, as in \\cite{Brandenberger:2006pr,Kaloper:2006xw}. The use of the Einstein frame allows us to import our intuition and understanding from standard gravity -- at least while we remain within the frame -- and, perhaps more importantly, it allows us to directly leverage well-tested codes developed for general relativity and which contain a wealth of physics and arbitrary collections of fluids. Codes of comparable breadth would be extremely lengthy to implement in the Jordan frame, and their fourth-order nature has consequences for speed and stability.  Considering the nature of the transformation for a perturbed model reveals an issue which to the best of our knowledge is as yet unappreciated. Perturbed systems in relativity exhibit the gauge issue -- the mapping between the fictitious background spacetime and the physical perturbed spacetime is non-unique. In cosmology, gauge freedoms allow one to eliminate four perturbative degrees of freedom (two scalar and two vector, where the classification is based on how the modes transform on a 3-surface), which specifies the slicing and threading we have chosen for our foliation. The transformation between the Einstein and the Jordan frames tangles this choice. We consider the impact of this and how one can resolve the resulting gauge ambiguities. We illustrate with a simple $f(R)$ model in a vacuum FLRW spacetime, but it should be emphasised that this issue will in principle occur in any perturbed spacetime and a wide range of extended theories of gravity.  We choose to work in the standard (``metric'') approach to cosmological perturbations (e.g. \\cite{MukhanovFeldmanBrandenberger,MaAndBertschinger}) and leave potential issues in the 3+1 gauge-invariant and covariant (GIC) approach \\cite{EllisBruni,Li:2007xn,Abebe:2011ry} for future study.\\footnote{The issue of perturbed models of modified gravity in the 3+1 GIC approach has been considered in \\cite{Carloni:2009gp}. The authors wish to thank Sergei Odintsov for bringing this work to our attention.} While many authors working on perturbed spacetimes in modified gravity choose to employ ``gauge-invariant'' variables (for example, \\cite{Hwang1990CQG,MukhanovFeldmanBrandenberger}) or work in the Jordan frame (such as \\cite{Hwang1996PRD,Song:2006ej,Carroll:2006jn,Tsujikawa:2007tg,delaCruzDombriz:2008cp,DeFelice:2010gb,Bertacca:2011wu}), the analysis of gauge problems is necessary for three reasons. Firstly, fixing a gauge (as in \\cite{Hwang:1996np,Bean:2006up,Zhang:2007nk,Tsujikawa:2007gd}), is a valid approach, and any problems when transforming to the Einstein frame must be understood. Secondly, gauge-invariant variables are themselves based on variables in a particular gauge \\cite{Malik:2008im}. The frequently-employed Bardeen potentials are, for example, gauge-invariant generalisations of the potentials in Newtonian gauge. If there is an issue with the choice of the ``natural'' gauge of a gauge-invariant perturbation, the perturbation itself cannot necessarily be trusted. Thirdly, our study is motivated not least by the use of Boltzmann codes in the Einstein frame. By far the widest used Boltzmann code is CAMB \\cite{CAMB}, which is written in a frame equivalent to synchronous gauge.\\footnote{Of the other modern Boltzmann codes, cmbeasy \\cite{CMBEasy} can be used in synchronous or Newtonian gauges, and CLASS \\cite{CLASS} is best tested in synchronous gauge although it was updated with support for Newtonian gauge after the release of the first version of this paper. Very recently, CosmoLib \\cite{CosmoLib} was released, programmed entirely in Newtonian gauge. COMICS \\cite{COSMICs} could be employed in both Newtonian and synchronous gauges, but was rapidly superseded by CMBFast \\cite{CMBFast} which could only be used in synchronous gauge. Both of these codes are now obsolete, with CMBFast in turn superseded by CAMB and cmbeasy.  Our comments concerning CAMB and synchronous gauge equally apply to other Boltzmann codes in other gauges.} If there is an issue with the frame transformation and synchronous gauge, then the initial conditions for modifications of CAMB could be suspect, as could the final results even after being transformed back into the Jordan frame. Worse, if there is an issue the error is compounded: each transformation between frames will induce errors. Given that the use of CAMB (or similar codes) is of interest, the study of transformations of systems in synchronous gauge is forced upon us, despite its known shortcomings: synchronous gauge is not fully fixed \\cite{Malik:2008im,Christopherson:2011ra} and ambiguities in the choice of threading remain.  The alternative formulation of $f(R)$ gravity employs the Palatini approach, where the metric and the connections are treated as independent objects. The Palatini approach is popular in cosmology (e.g. \\cite{Capozziello:2010ef,Capozziello:2011et,Olmo:2011uz} and their references) and often employed in studies of cosmological perturbations (e.g. \\cite{Koivisto:2005yc,Li:2006vi,Sotiriou:2008rp}). The nature of the transformation between the frames implies that the issues we highlight in this paper are also relevant in the Palatini approach. Since the field equations in the Palatini approach are intrinsically second-order the usefulness of the transformation is somewhat lessened, but if a system is more easily evaluated in the Einstein frame then the potential ambiguities should still be considered.  The gauge problem is discussed in \\S\\ref{Section:GaugeAmbiguity}A and illustrated in \\S\\ref{Section:GaugeAmbiguity}B. We close the paper in \\S\\ref{Section:Discussion} with a brief discussion.  We employ a signature $(-+++)$ and the Ricci tensor $R_{\\mu\\nu}=R^\\alpha_{\\phantom{\\alpha}\\mu\\alpha\\nu}=\\Gamma^\\alpha_{\\mu\\nu,\\alpha}+\\ldots$ (as in for example \\cite{Gravitation,Maeda,DeFelice:2010aj}) and generally follow the notation of \\cite{DeFelice:2010aj}. An overdot will represent differentiation with respect to conformal time $d\\eta=dt/a$, and the conformal Hubble rate is written $\\mathcal{H}=\\dot{a}/a$. ",
        "conclusions": "\\label{Section:Discussion} In this paper we have discussed a technical subtlety of the conformal transformation between the Jordan and Einstein frames that to the best of our knowledge has not been highlighted before: na\\\"ively transferring a perturbed system between the frames tangles your choice of gauge.  It should be emphasised that this effect is \\emph{calculational}, not physical, and that if sufficient care has been taken no previous results are changed by this. It is also extremely important to note that this issue is not restricted to $f(R)$ gravity or to the study of cosmological perturbations, and that it potentially arises whenever a perturbed system is conformally transformed regardless of the background metric or the model of gravity. $f(R)$ theory and vacuum cosmology provide a useful, straightforward illustration. While we have concretely demonstrated the issue only for this vacuum $f(R)$ system, in principle it arises generically, although demonstrating this explicitly is beyond the scope of this paper.  The properties of the conformal transformation restrict us to a particular set of gauges which we refer to as the ``restricted group'': gauges that possess a lapse, a spatial curvature, and a density perturbation. This restriction only applies when the transformation is performed. At all other times one is free to work in any convenient gauge. However, we additionally argued that there are alternatives to working in the restricted group -- one can instead choose to ``refix'' the gauge, or to evolve the system using the redundant Einstein field equations.  We have illustrated our arguments with a simple vacuum system, showing that there are no ambiguities introduced by equations (\\ref{BardeenTransformations}, \\ref{FieldTransformations}), so long as one takes such care. If this is not done and the results of a transformation from a gauge outside the restricted group are interpreted as if they themselves lie in that gauge, the system will be inaccurately specified.  Perhaps the most important place where this issue will arise is in the treatment of initial conditions. For instance, Boltzmann codes are frequently \\cite{COSMICs,CMBFast,CAMB,CLASS} written in synchronous gauge, which lies outside of the restricted group. Interpreting the Jordan frame as physical, the initial conditions must be set in the Jordan frame. However, regardless of the gauge chosen, the transformation into the Einstein frame induces a spurious lapse function which must be reabsorbed if the results of the calculation are to make any sense. This itself introduces an additional problem: the transformation to synchronous gauge is not fully specified, and introduces an arbitrary function of the 3-coordinates which must itself be removed with care. The alternative is to include the redundant degrees of freedom in the Boltzmann codes, which requires much additional effort. A careless study without attention to these issues would either leave the system with the spurious lapse, or risk introducing a gauge mode. Some Boltzmann codes \\cite{COSMICs,CMBEasy,CLASS} include modules written in Newtonian gauge, which lies within the restricted group, and studies of extended gravity that employ these codes remain consistent. However, COSMICS is very outdated, while cmbeasy and CLASS are programmed in C++ and C respectively; given the widespread use of Fortran in cosmology, its large codebase and the extensive testing it has undergone, CAMB remains the dominant Boltzmann code and the study of synchronous gauge is necessary.\\footnote{CosmoLib \\cite{CosmoLib} is also written in Fortran but in Newtonian gauge; however it has appeared too recently to yet have much of an impact. However, it could prove extremely useful for studies of modified gravity.} Similar arguments apply to the initial conditions employed in n-body simulations of modified gravity performed in the Einstein frame.  The other time at which the gauge ambiguity becomes important is at the end of a calculation. Taking the Jordan frame to be physical, the results of a simulation in the Einstein frame must be transformed back. Given that calculations are frequently performed in synchronous gauge then one must either transform into a safer gauge before performing the conformal transformation, or deal with the redundant system in the Jordan frame. Otherwise one runs a serious risk of contamination. At this stage it seems easier to refix the gauge in the Jordan frame to a fully-specified gauge (such as Newtonian or uniform curvature), or to leave the system unspecified. So long as one does so consistently, observables will not be affected, since gauge transformations cannot change an observation.  It is interesting to note that the lensing potential $\\Phi+\\Psi$ is uncontaminated by the transformation. As such, if one wishes to calculate the weak lensing signal on a constant-time hypersurface, this can be achieved transforming from the Jordan frame into the Einstein frame in any gauge and not worrying about refixing the system \\cite{Schimd:2004nq}. For a static (or quasi-static) system, this will be a good approximation. However, carelessly evolving the system without due care will in principle introduce errors, so that the frame invariance does not imply such a straightforward estimate of the integrated Sachs-Wolfe effect.  In summary, we recommend that authors exploiting the Einstein frame to study extended theories of gravity work in the restricted group. This removes the need for additional gauge transformations to reabsorb the extra apparent degrees of freedom, or the need to employ redundant Einstein equations. If one chooses the latter options, care should be taken to ensure that the result remains consistent."
    },
    "1112/1112.1745_arXiv.txt": {
        "abstract": "*{I discuss the effects of measurement error on regression and density estimation. I review the statistical methods that have been developed to correct for measurement error that are most popular in astronomical data analysis, discussing their advantages and disadvantages. I describe functional models for accounting for measurement error in regression, with emphasis on the methods of moments approach and the modified loss function approach. I then describe structural models for accounting for measurement error in regression and density estimation, with emphasis on maximum-likelihood and Bayesian methods. As an example of a Bayesian application, I analyze an astronomical data set subject to large measurement errors and a non-linear dependence between the response and covariate. I conclude with some directions for future research.}   \\vspace{-1.5 in} \\abstract{I discuss the effects of measurement error on regression and density estimation. I review the statistical methods that have been developed to correct for measurement error that are most popular in astronomical data analysis, discussing their advantages and disadvantages. I describe functional models for accounting for measurement error in regression, with emphasis on the methods of moments approach and the modified loss function approach. I then describe structural models for accounting for measurement error in regression and density estimation, with emphasis on maximum-likelihood and Bayesian methods. As an example of a Bayesian application, I analyze an astronomical data set subject to large measurement errors and a non-linear dependence between the response and covariate. I conclude with some directions for future research.} ",
        "introduction": "\\label{intro}  Measurement error is ubiquitous in astronomy. Astronomical data consists of passive observations of objects, whereby astronomers are able to directly measure the flux of an object as a function of wavelength, its location on the sky, and the time of the observation. Because the number of photons detected from an astronomical object follows a Poisson process, this makes the measurement of a source's intensity intrinsically subject to measurement error, even if none is introduced from the detector. Therefore, the very nature of astronomical data makes measurement error unavoidable. Moreoever, quantities that are derived from an object's observed emission, either by fitting a model to the spectral energy distribution (SED) or by employing scaling relationships, are also `measured' (derived) with error. Examples include mass, metallicity, and distance. Often the measurement error on the derived quantities is significant. This is unfortunate as inference on the derived quantities is often the goal of astronomical data analysis. Therefore, there has been considerable interest in how to perform statistical inference in the presence of measurement error.  Measurement error is a problem that affects, at various levels, all scientific research. Because of this, numerous methods for handling measurement errors have been developed (Fuller 1987, Cheng \\& Van Ness 1999, and Carroll et al. 2006 are good references). In this contribution, I will present a survey of methods for handling measurement error that have been developed and used in astronomical data analysis. Because astronomical measurement errors are, in general, heteroskedastic (having different variances), I will limit my discussion to methods developed for heteroskedasticity. I will focus on situations where a deterministic relationship is not assumed between the variables, but where all variables of interest are random and are measured with error. Because of this, I will ignore situations where the measurement error is the only source of randomness in one's data. An example of this type of situation is fitting a model to an observed spectrum, where the measurement error is the only source of randomness; i.e., in the absence of measurement error a deterministic relationship is assumed between, say, flux density and wavelength. Methods for handling measurement error in this case are relatively well-established, and typically one simply minimizes the usual $\\chi^2$ statistic (e.g., Bevington 2003). However, it is worth pointing out that many complications may still exist, and more sophisticated methods may be needed, especially when dealing with low-count X-ray and $\\gamma$-ray data (e.g., van Dyk et al. 2001) or when incorporating calibration unceratinties (Lee et al. 2011). Instead, I will focus on methods for analyzing data from astronomical samples, where the variables are a random sample from an underlying distribution. Within the context of regression, this implies that intrinsic scatter (referred to as equation error in the statistics literature) exists in the relationship among the variables, and thus a deterministic relationship is not assumed between the variables even without the presence of measurement error.  Most of the techniques I will discuss focus on accounting for measurement error in regression. The goal of regression is often to understand how one variable changes with another. For example, how does the mass of a black hole change as a function of the stellar velocity dispersion of the host galaxy's bulge? Typically one simply estimates how the average value and dispersion of one variable depends on another. Measurement error statistical models are typically divided into two types: `functional' and `structural' models. In functional modeling, one assumes that the unknown true values of the variables are fixed, whereas in structural modeling the unknown true values of the variables have their own intrinsic distribution. As a result, in structural modeling one must parameterically model the distribution of the true values of the variables, whereas in functional modeling one does not. Density estimation is another important technique in astronomical data analysis, being the foundation for luminosity and mass function estimation. The methods I will discuss for handling measurement error in structural models are also applicable to density estimation, as in this case regression and density estimation are based on the same formalism. When discussing regression methods, I will refer to the `dependent' variable as the response, and the `independent variables' as the covariates. ",
        "conclusions": ""
    },
    "1112/1112.2948_arXiv.txt": {
        "abstract": "M87 is the first detected non-blazar extragalactic Tera-Electron-Volt (TeV) source with rapid variation and very flat spectrum in the TeV band. To explain the two-peaks in the spectral energy distribution (SED) of the nucleus of M87 which is similar to those of blazars, the most commonly adopted models are the synchrotron self-Compton scattering (SSC) models and the external inverse Compton scattering (EIC) models. Considering that there is no correlated variation in the soft band (from radio to X-ray) matching the TeV variation, and the TeV sources should not suffer from the $\\g\\g$ absorption due to the flat TeV spectrum, the EIC models are advantageous in modeling the TeV emission from M87. In this paper, we propose a self-consistent EIC model to explain the flat TeV spectrum of M87 within the framework of fully general relativity, where the background soft photons are from the advection-dominated accretion flow (ADAF) around the central black hole, and the high energy electrons are from the mini-jets which are powered by the magnetic reconnection in the main jet \\citep{2010MNRAS.402.1649G}. In our model, both the TeV flares observed in the years of 2005 and 2008 could be well explained: the $\\g\\g$ absorption for TeV photons is very low, even inside the region very close to the black hole $20R_g\\sim50R_g$; at the same region, the average EIC cooling time ($\\sim 10^2\\sim10^3s$) is short, which is consistent with the observed time scale of TeV variation. Furthermore, we also discuss the possibility that the accompanying X-ray flare in 2008 is due to the direct synchrotron radiation of the mini-jets. ",
        "introduction": "Unlike blazars and the Galactic sources (e.g. SNRs), which composite the most population of TeV sources, M87 is the first discovered radio galaxy with TeV radiation. Recently three other radio galaxies Centaurus A, 3C66B and IC310  were also identified as TeV sources \\citep{ 2009ApJ...695L..40A,2009ApJ...692L..29A, 2010ApJ...723L.207A}. Comparing to blazars, M87 has much milder variations in both optical and X-ray band; most importantly, the prominent kpc-scale jet of M87 has a very large viewing angle $\\sim30^\\circ$ with respect to the line of sight \\citep{bick96,ahar06}. In 2005, rapid TeV variation ($1-2$days) was discovered with flat spectrum for the first time, but without correlated X-ray variation from the nuclear core \\citep{ahar06}. While in 2008, the radio, X-ray, TeV joint observation discovered that TeV flares lasts 2 weeks accompanied with an X-ray flare and a radio flare inside the unresolved nuclear core ($30\\times60R_s$) as well as with a radio blob moving out of the unresolved core \\citep{2009Sci...325..444A}. Here $R_s$ used in the VLBA observations are based on the black hole (BH) mass of M87 $M_{BH}=6\\times10^9M_{\\sun}$ \\citep{2009ApJ...700.1690G}. Recently, a TeV flare in 2010 which is very similar to the previous TeV flares in 2005 and 2008 was reported. Further observations at X-rays and radio find the correlated X-ray flare, but no enhanced radio flux from the core \\citep{Abramowski2011,Harris2011}. In this paper, we adopt $M_{BH}=3\\times10^9M_{\\sun}$ \\citep{1997ApJ...489..579M} to be consistent with the advection-dominated accretion flow (ADAF) models and the mini-jet models in the previous studies \\citep{2009ApJ...699..513L,2010MNRAS.402.1649G}.  Although M87 has a large viewing angle, it is very near to the Earth with a distance of $\\sim16.7$Mpc. Owing to its proximity, it is proposed that M87 might be a misaligned blazar \\citep{1998ApJ...493L..83T}. Whereas blazars are believed to have jets beaming towards us, TeV sources in M87 could be shocks with relativistic bulk velocity inside the jets. Recent observed minute-scale variation from galaxies like Mrk 421 (Fossati et al. 2008) and PKS 2155-304 \\citep{2007ApJ...664L..71A} indicates that TeV sources should be compact and very close to the BH. Comparing to the TeV flares of blazars, the slower variation ($1-2$ day) detected in M87 could be the results of too few observed TeV photons (each data point requires integration of photons for the whole night) or the much lower Doppler factor of the bulk velocity of the TeV sources due to the large inclination. Furthermore, unlike the steeper TeV spectrum of blazars, the much flatter TeV spectrum of M87 could be mainly due to the lack of $\\g\\g$ absorption. The large viewing angle may play an important role in avoiding $\\g\\g$ absorption within the jet; thanks to its proximity and the dimness of its host galaxy, the absorption to TeV photon from M87 caused by the galactic and the intergalactic background soft photons, as well as the cosmic background photons is rather weak \\citep{nero07}.  To explain the spectral energy distribution (SED) of the nucleus of M87, the one zone synchrotron self-Compton(SSC) models, which have been applied to the TeV flares in blazars \\citep[e.g.][]{1998ApJ...509..608T, 2001A&A...367..809K}, face difficulties in fitting the two peaks in the SED, however, in the multi-zone SSC models, the TeV source is separated from the soft ones, therefore the whole SED could be easily fitted as long as there are enough SSC blobs with specific locations, electron distributions and bulk velocity \\citep[e.g.][]{2008MNRAS.385L..98T,2007sf2a.conf..196L,2005ApJ...634L..33G}. The multi-zone SSC models could also provide better explanation for the rapid variation, and the orphan TeV flares (the variations in the soft band can not connect with the TeV variation). Furthermore, considering the very flat power-law SED of M87 in the TeV band with index $p\\simeq-2.2\\sim-2.6$, the $\\g\\g$ absorption provides crucial constrains on the SSC models. So far, most of the SSC models still have problems in explaining the very flat TeV spectrum due to the certain $\\g\\g$ absorption, therefore the external inverse Compton (EIC) process could be more likely responsible for the TeV flare \\citep{2008MNRAS.384L..19B,2009MNRAS.395L..29G,2010MNRAS.402.1649G}. In the EIC models, the TeV source is far away from the soft one, or has the relativistic bulk velocity with respect to the soft one, thus the $\\g\\g$ absorption can be reduced.  In those models mentioned above, how the very high energy (VHE) particles are produced is still an open question. It is generally believed that VHE particles are accelerated by the shocks in jets. There are some other possibilities include the mini-jets powered via the magnetic reconnection in the main jet \\citep[][and references therein]{2010MNRAS.402.1649G}; the magnetic centrifugal acceleration in the vicinity of BH \\citep{nero07,rieg08}, and so on. In the mini-jets model \\citep{2010MNRAS.402.1649G}, magnetic reconnection within the main jet produces two oppositely directed (in the rest frame of the main jet) mini-jets, in the laboratory frame, one of them always points within the angle of the main jet and is observable in blazars because their jets point at us. The other mini-jet (its counterpart) points outside the opening angle of the main jet and is potentially observable to off-axis observers in case of the misaligned jets, such as those of M87 and Centaurus A. In a word, the great advantage of the mini-jets model is that even at large inclination, mini-jets with a high bulk speed can still be detected, which is helpful for explanation of the fast TeV variation from M87. After all, in order to avoid the $\\g\\g$ absorption, the energy density of the soft synchrotron photons in the TeV source must be limited, therefore, the minimal Lorentz factor of the VHE particles is generally assumed to be $10^{3-4}$ and the strength of the magnetic field in the TeV source below several Gauss \\citep{2010MNRAS.402.1649G}. Beside the direct inverse Compton process of the VHE electrons, there are also some alternative models to produce the SED of the TeV flares, such as the hadronic models including, among which, the interaction between the VHE protons and soft photons \\citep{2004A&A...419...89R} and the proton-proton collision process when a red giant was passing the base of the jet \\citep{2010ApJ...724.1517B}.  In this paper, we develop a fully general relativity EIC model for explaining the TeV emission from M87, in which the soft photons emanate from the ADAF around the BH. In \\S2, we investigate the safe zone of TeV photons (the $\\g\\g$ optical depth is below unity). The technical details of our fully general relativity EIC model can be found in \\S3. The numerical results and the discussions are given in \\S4 and \\S5, respectively.  \\begin{figure}[t!] \\centering \\includegraphics[width=0.6\\textwidth, trim=0pt 80pt 0pt 0pt]{fig1.eps} \\caption{ The contour of $\\tau_{\\g\\g}=0.1$ or $1$ in the $\\alpha,\\beta$ plane. Where $r$ represents the distance to the BH in the Boyer-Lindquist frame. The viewing angle is taken to be $30^\\circ$.  The blank region in the lower panels is the region with $\\tau_{\\g\\g} \\leq 1$. } \\label{fig1} \\end{figure} ",
        "conclusions": "In this paper, by taking into account the fully general relativity effects, we propose a disk-dominating external Compton-scattering model for explaining the flat TeV radiation from M87. The external Compton-scattering model suffers much less self-$\\g\\g$ absorption, comparing to the one-zone self-synchrotron Compton model, in which the soft photons and the VHE electrons are from the same blobs. The advantage of the EIC model in which the soft radiation is from the accretion disk is that the soft radiation source could be more compact and further away from the TeV source than ones in the SSC models, in which the $\\g\\g$ absorption can be reduced significantly. Besides, our EIC model is also supported by the observed mild IR-UV-X flux variations of the nucleus and the large viewing angle of the jet of M87 (the nuclei flux with viewing angle $30^\\circ$ is unlikely to be contaminated by the jet, neither to be blocked by the disk itself). It should be noticed that the VHE electrons and TeV photons may be influenced by some outflows with the anisotropic radiation at the soft band which are not beaming towards us, such as the jet itself. Unfortunately, the base of the jet of M87 appears to be quite chaotic, thus it is hard to draw any conclusions about the detailed structure of jet/outflows \\citep{2011arXiv1109.6252P}. In our model, the TeV sources are the mini-jets which are easily able to beam towards us with high Doppler factor and give rise to the observed VHE radiation.  How to distinguish between SSC and EIC models by the future observations? The main difference between these two models is that the IR-UV variation and TeV flares are not necessarily correlated in the EIC models, but they are in the SSC models, especially the multi-zone SSC model. Therefore, we can distinguish between the SSC and EIC models by the simultaneous observations of the IR-UV emission during the TeV flare. Unfortunately, during the 2005 TeV flares, there is no accompanied HST observation of M87. With the better sensitivity of CTA (it is about $>$10 times better than those of the present Cerenkov telescopes), we will be able to detect the fainter flares and obtain the minimal variation time scale of the TeV flares, which increases the opportunity of finding the correlated variation between IR-UV and TeV emission.  In our models, with little $\\g\\g$ absorption, all the EIC spectra at the observational band have inherited the original power-law index of the VHE electrons. After fitting those spectra, we calculate the cooling time of the VHE electrons in the mini-jets (see section 4) to make sure that they would not travel longer than $ct_{\\rm var}\\delta_D$. It turns out when TeV source is located at $H_{\\rm TeV}<50R_{\\rm g}$, the IC cooling time is very much shorter than the upper-limits of the observed TeV variation of 1-2 days. Therefore, if possibly, the TeV source is located further away from the disk ($H_{\\rm TeV}>50R_{\\rm g}$), the cooling efficiency of the VHE electrons in the mini-jets is much lower, the mini-jets could move out of the main jet. As a result, the variation of the soft radiation background at the different locations above the disk should be considered.  We also discuss the probability that the direct synchrotron radiation from the mini-jets may cause the correlated X-ray flare in 2008. As expected, according to our model, if the magnetic field is about several Gauss, the direct synchrotron flux from these mini-jets which lies at the X-ray band can explain the X-ray flare very well but has no influence on the observed radiation at the GeV-TeV band. However, the magnetic field $B$ can not be constrained by any direct observation, as discussed above, and we only have a rough estimate of the strength of the magnetic field supposed to be consistent with the mini-jets model. Considering that the synchrotron flux is roughly about $F_{\\rm syn}\\sim N_0B^2$, it is obvious that the $N_0$ and $B$ are degenerated. If the TeV source is confined inside the region $20R_g<H_{\\rm TeV}<50R_{\\rm g}$, the corresponding strength of the magnetic field is required to be about $2G< B < 8G$, in order to produce the X-ray flare in 2008.  Dealing with the two weeks data of the 2008 flare from MAGIC, Albert et al. (2008) had separated the high state data ( $\\sim$ 1 day time scale) from the low state data and they found that during the high states, the spectra are harder: the spectrum index of the high state is about 2.2, while that of the low state is about 2.6. Because the IC cooling time and the synchrotron cooling time of VHE particles is short, the observed fastest variation should be dominated by the time of the acceleration, or even the shifting of the beaming angle of the mini-jets. If there is no continuous acceleration of the relativistic electrons, the spectra of the TeV flares will become steeper. In the year of 2008, VERITAS captured the TeV flare (Feb 9-13, 2008) which is close to the date of the X-ray flare (Feb 16, 2008). VERITAS observation shows that the spectrum index of the TeV flare is about 2.4, which is softer than the earlier flares captured by MAGIC \\citep{2010ApJ...716..819A}. We suggest that the TeV flare captured by VERITAS is happening during the rapid synchrotron cooling of the relativistic electrons in the mini-jet, and the X-ray flare could be a byproduct of that TeV flare, as the X-ray flares in 2008. Our suggestion could naturally explain the X-ray excess and the softer spectrum of TeV emission. There are two ways to check above suggestion: first, the flux of the X-ray via synchrotron radiation should be anti-correlated with the TeV flux via EIC, which could be tested by the future observations with the higher time resolution (less than 1 day); second, and the TeV flare should share the similar spectrum index with that of the X-ray flare. Unfortunately, due to the pile up effect of X-ray observation \\citep{2009ApJ...699..305H}, we can not obtain the correct spectrum index of the X-ray flare."
    },
    "1112/1112.0726_arXiv.txt": {
        "abstract": "In a series of our previous works, a model of radiatively inefficient accretion flows (RIAFs) in a global magnetic field (so called resistive RIAF model) has proved its ability to account for many physical processes taking place in such accretion flows as realized in the nuclei of the galaxies believed to be accreting at a very small fraction of each Eddinton accretion rate. Within the present status of this model, however, the model cannot describe the launch of a self-confined bipolar jet from the vicinity of disk's inner edge, although it allows the existence of a thermal wind widely distributed over the disk surfaces. This is because the electric field (and hence the Poynting flux) vanishes everywhere in the disk, whereas such a jet in a globally ordered magnetic field is most likely to be accelerated electrodynamically. We show in the present paper that this defect can be overcome naturally if we reformulate the problem so as to admit a quasi-stationary change of the magnetic field (and hence the appearance of a non-irrotational electric field), and also restore all the terms of order $\\epsilon\\equiv (v_r/v_{\\varphi})^2\\lesssim 1$ (where $v_r$ and $v_{\\varphi}$ denote radial and azimuthal components, respectively, of the fluid velocity) which have been neglected altogether in the previous treatments. The restored effects are the inertial and magnetic draggings on the infalling matter. As an illustrative example, a model solution which is correct up to ${\\cal O}(\\epsilon)$ is derived under a set of plausible restrictions. The new solution predicts the appearance of a localized Poynting flux in a region near the disk inner edge, suggesting strongly that a jet is launched from this region. Another interesting prediction is the appearance of a rapid change of the magnetic field also localized to this region. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.5335_arXiv.txt": {
        "abstract": "{ We review the effect of the formation of nonlinear structures on the expansion rate, spatial curvature and light propagation in the universe, focusing on the possibility that it could explain cosmological observations without the introduction of dark energy or modified gravity. We concentrate on explaining the relevant physics and highlighting open questions. }  \\preprint{HIP-2011-35/TH} ",
        "introduction": "\\label{sec:intro}  \\subsection{The role of structure formation in cosmology}  The formation of nonlinear structures from tiny almost Gaussian perturbations by gravitational instability is a central topic in cosmology. Most studies are done either in general relativity (GR) using linear perturbation theory around the locally homogeneous and isotropic Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) model or in Newtonian gravity, where extensive numerical computations have been done with N-body simulations, preceded and accompanied by analytical work, to access the nonlinear regime of structure formation. The formation of structures is sensitive to the expansion of the universe, so the observed large-scale structure is a valuable probe of cosmological evolution. There is also another side to the matter: perturbations have a significant effect on the expansion rate once they enter the nonlinear regime. In standard linear theory, the effect vanishes on average by construction. In Newtonian gravity, this turns out to be true also in the non-perturbative regime (see \\sec{sec:newton}). However, this result does not carry over to non-perturbative GR, where the geometry is dynamical. Therefore, there could be a large effect on the average expansion rate, the evolution of spatial curvature and light propagation.  The effect of deviations from exact homogeneity and isotropy on the average expansion is known as {\\it backreaction}. The effect was first studied in 1962 \\cite{Shirokov:1962}, and an in-depth discussion was given 20 years later by George Ellis \\cite{fitting}. Ellis introduced the term {\\it fitting problem} to describe the issue of finding a model that best describes the average of the real universe with its complex structures. The problems laid out by Ellis and collaborators did not become mainstream concerns in cosmology at the time, which may partly be due to the complexity of the issues. Another likely reason is that cosmological observations were not precise enough to indicate any tension with a homogeneous and isotropic cosmological model based on ordinary GR with ordinary (visible and dark) matter. As more precise data has revealed that such a model disagrees with observations, it has become topical to study backreaction in more detail to see whether it could account for this failure \\cite{Buchert:2000, Wetterich:2001, Schwarz:2002, Rasanen}. We review backreaction in classical GR due to the formation of nonlinear structures at late times, with a focus on the possibility that it could explain this discrepancy. We do not aim to give a historical overview of the literature, but rather try to explain the relevant physics and highlight open questions. There are a number of reviews with comprehensive lists of references \\cite{Rasanen:2006b,buchert:review,ellisbuchert,Ellis:2011a,buchertFOCUS, Rasanen:2011a, Clarkson:2011b}. We concentrate on structures that are known to exist and that are expected in the standard structure formation scenario based on usual models of inflation. In particular, we do not discuss models with Gpc-scale structures where, for instance, the observer is located near the centre of a large void \\cite{voidmodels}. We do not discuss the backreaction of classical or quantum perturbations during inflation or reheating \\cite{quantum}. We also do not cover schemes that introduce some extra structures to GR, which may be needed if the `fitting problem' is to be solved for tensorial objects such as the metric \\cite{nonscalar}.  In section 1 we briefly review some relevant cosmological observations and discuss assumptions underlying the FLRW model. In section 2 we put into perspective the equations that govern the average expansion, explain the meaning of average acceleration and discuss the difference between GR and Newtonian gravity. In section 3 we discuss estimates of backreaction in perturbation theory and in non-perturbative models. In section 4 we consider light propagation and the relation of spatially averaged quantities to observables, as well as observational signatures of backreaction. We conclude in section 5 with a summary.  \\subsection{Cosmological observations} \\label{sec:obs}  A linearly perturbed homogeneous and isotropic cosmological model with ordinary matter (i.e. matter with non-negative pressure) and ordinary gravity (based on the four-dimensional Einstein-Hilbert action) can explain all observations of the early universe from Big Bang Nucleosynthesis onwards. The simplest such model, known as the Standard Cold Dark Matter (SCDM) model, is matter-dominated at late times and has flat spatial sections. The SCDM model can also account for late-time observations at the accuracy that had been achieved in the 1980s. However, modern observations have shown that the distance to the last scattering surface at redshift $z=1090$ is a factor of $1.5$ to $2.0$ larger than predicted \\cite{Vonlanthen:2010, Hubble} (if the Hubble constant is held fixed and primordial perturbations have a power-law spectrum). Observations of type Ia supernovae and large-scale structure \\cite{SNIa, kinematic, BAO} show that this discrepancy occurs at redshifts of order unity, that is, at approximately 10 billion years. Most observations probe distances, and the expansion rate is usually inferred through the use of the FLRW relation between the distance and the expansion rate (see \\sec{sec:light}). In the FLRW model, the explanation for longer distances is that the expansion has accelerated, so objects have been pushed further away than expected.  There are few model-independent observations of the expansion rate as a function of redshift, but the Hubble parameter today is relatively well known \\cite{Hubble}. By combining with measurements of the matter density we obtain $\\Omega_\\mathrm{m0}\\equiv8\\pi\\GN\\rho_\\mathrm{m0}/(3H_0^2)\\approx0.2\\ldots0.4$ \\cite{Vonlanthen:2010, Peebles:2004}. Combining it with the age of the universe instead \\cite{Krauss:2003}, we get $H_0t_0\\approx0.8\\ldots1.1$. In the SCDM model, $\\Omega_\\mathrm{m0}=1$ and $H_0t_0=2/3$, so either way we find that the expansion is faster than expected by a factor ranging from $1.2$ to $2.2$. This finding supports the interpretation of the distance observations in terms of faster expansion. At the moment, model-independently it is possible to say only that the expansion has decelerated less, not that it would have accelerated \\cite{ages}. Analysis of large scale structure \\cite{BAOradial}, especially in combination with supernova data \\cite{Blake:2011} can provide strong constraints on the evolution of the Hubble parameter, though these constraints depend on the assumption that the universe is close to FLRW. Currently, we cannot conclude that the expansion would have accelerated without assuming the validity of the FLRW approximation.  Within the FLRW model, the conclusion that the expansion rate has accelerated follows directly from kinematical analysis of distance observations \\cite{kinematic}. To obtain acceleration in the FLRW model it is necessary to modify the dynamics either by introducing some exotic form of matter that has negative pressure or by changing the laws of gravitation on cosmological scales. In the simplest such model, the $\\Lambda$CDM model, constant vacuum energy or, equivalently, a cosmological constant is added. The only difference between the $\\Lambda$CDM model and the SCDM model is that in the former the expansion rate increases at late times, which is sufficient to bring predictions and observations into agreement (although problems appear to remain with structures on large scales \\cite{inhom, Einasto, Murphy:2010, Yaryura:2010, Nadathur:2011}). This point is worth emphasizing: anything that changes the expansion rate and the corresponding distances at late times by the right amount (and that does not modify other things too much) accounts for all of the observed discrepancies, whether of luminosities of supernovae, anisotropies of the cosmic microwave background (CMB), the growth rate of structures or other probes. Notably, all the relevant observations involve integrals over long distances or averages over large scales.  \\subsection{The hypotheses behind the FLRW model}  The common justification for using the FLRW model to describe the universe on average is that the universe appears to be homogeneous and isotropic on large scales. However, we have to distinguish between {\\it exact} and {\\it statistical} homogeneity and isotropy. Exact homogeneity and isotropy means that the spatial geometry has a local symmetry: all points and directions are equivalent. (The topological structure may break this symmetry globally.) Statistical homogeneity and isotropy means that if we consider a large domain anywhere in the universe, the mean quantities in the domain do not depend on its location, orientation or size, provided that the domain is larger than the {\\em homogeneity scale}\\footnote{Assuming that a homogeneity scale exists, that is, that the mean density and other average quantities converge to a scale-independent value as the volume increases. We make this assumption throughout. In statistical physics, this property is known as {\\it spatial homogeneity}  -- not to be confused with the GR property of the same name which is related to local symmetry \\cite{inhom}.}.  In the usual cosmological treatment, the early universe is assumed to be close to exact homogeneity and isotropy in the sense that the difference between the energy density, expansion rate and other quantities between any two spatial locations is small. Cosmological observations, especially the CMB, provide strong support for this assumption\\footnote{However, it has not been proven and in principle the effects we discuss could also be relevant in the early universe.}, and we retain it. Nevertheless, when density perturbations become nonlinear at late times, the universe is no longer locally close to homogeneity and isotropy. However, it remains statistically homogeneous and isotropic on large scales, assuming the initial distribution of perturbations had this property. Statistical homogeneity and isotropy (and a homogeneity scale) are predictions of the simplest inflationary models. We assume that these assumptions hold in the real universe, and that the homogeneity scale is of order 100 Mpc. Whether the existence of a homogeneity scale has been observationally established is a matter of contention \\cite{Hogg:2005, inhom}, and studies of the morphology of structures suggest that such a scale is no smaller than 300 Mpc \\cite{fluctuations}.  Two concepts that are sometimes used to motivate the use of the FLRW metric to describe the universe are the {\\em Copernican principle} and the {\\em cosmological principle}. According to the former, our location in the universe is not special, whereas according to the latter all locations in the universe are equivalent. However, the Copernican principle says nothing about possible symmetries of the geometry or the matter distribution. It is possible to have a space that has preferred directions or locations, but where the observer is not in a preferred location. On the other hand, the cosmological principle is a statement about spacetime symmetry, however strictly speaking it does not apply to the real universe because the universe contains structures. If the principle is weakened and interpreted to refer to the distribution of structures, then it is nothing but the statement of statistical homogeneity and isotropy. In modern cosmology, this statement is a prediction of simple models of inflation, rather than a principle, and it is subject to observational tests. Neither the Copernican principle nor the cosmological principle (interpreted to refer to large-scale statistical properties) show that the universe would be described by the FLRW model. In the case of observational tests \\cite{Hogg:2005, inhom, Clarkson:2007b, kSZ, Maartens:2011, Heavens:2011, Clifton:2011}, it is important to distinguish whether they probe statistical homogeneity and isotropy, the FLRW metric or the Copernican principle. (For example, the important check proposed in \\cite{Clarkson:2007b} tests the FLRW metric, not the Copernican principle.)  Put simply, the FLRW model describes universes that are locally homogeneous and isotropic on all scales, not universes that are only statistically homogeneous and isotropic. Because there are large local deviations, the average evolution may be far from the FLRW behavior even above the homogeneity scale. The possibility that the observed change in average quantities from those of the SCDM model at late times is due to the formation of structures may be termed the {\\it backreaction conjecture}. ",
        "conclusions": "We discuss the {\\it backreaction conjecture}, namely the possibility that structure formation changes the average expansion rate, spatial curvature and light propagation, thereby eliminating the need for dark energy or modified gravity. The change in the average properties due to structure formation is present in reality but is not taken into account either in the FLRW model and its linear perturbations or Newtonian non-linear models.  The basic mechanism of backreaction is simple and can be demonstrated in toy models: because the universe is inhomogeneous, different regions expand at different rates, so the fraction of volume in faster expanding regions can grow, and the average expansion rate can rise. Structure formation has a preferred time of 10 billion years, which agrees with the observed timescale for the change in the expansion rate, and the amplitude of the change can also be understood from simple considerations. However, the effect has not been quantified in a fully realistic model.  Various things remain to be done before robust quantitative conclusions can be drawn. If backreaction is significant, it cannot be understood simply as a change in the FLRW background model. To analyze signals such as baryon acoustic oscillations and other features in the distribution of large structure, we must develop perturbation theory around a non-FLRW background, where the mean is the average of deviations that have a large amplitude, but a small coherence length.  The treatment of light propagation also needs to be made more rigorous and extended both to cover phenomena such as weak lensing in detail and to include effects due to the discreteness of matter. The deviation of the relation between distance and the expansion rate from the FLRW case is an important prediction, which can be tested without a calculation of the average expansion rate. However, the goal should be to derive the change in the average expansion rate with quantified errors. In addition to statistical models, one way to address the question could be to generalize N-body simulations to include the relevant relativistic degrees of freedom.  The backreaction conjecture is conservative in the sense that it does not involve new fundamental physics, only neglected effects in non-linear GR. Before the effect of structure formation on the average expansion rate is reliably quantified, we will not know whether dark energy or modified gravity is needed. It is plausible that backreaction can explain all of the observations, and even if the large-scale average properties of the universe turn out to be close to the FLRW case, the corrections may nevertheless be quantitatively important."
    },
    "1112/1112.1982_arXiv.txt": {
        "abstract": "\\boldmath We present the design and testing of a 2--20~GHz continuum band analog lag correlator with 16 frequency channels for astronomical interferometry. The correlator has been designed for future use with a prototype single-baseline interferometer operating at 185--275~GHz. The design uses a broadband Wilkinson divider tree with integral thin-film resistors implemented on an alumina substrate, and custom-made broadband InGaP/GaAs Gilbert Cell multipliers. The prototype correlator has been fully bench-tested, together with the necessary readout electronics for acquisition of the output signals. The results of these measurements show that the response of the correlator is well behaved over the band. An investigation of the noise behaviour also shows that the signal-to-noise of the system is not limited by the correlator performance. ",
        "introduction": "\\IEEEPARstart{T}{he cross-correlator} is the heart of an interferometric radio telescope, where the signals from each antenna are combined to form the complex visibility measurements from which the image of the sky is derived.  Interferometry, with its inherently high stability and direct measurement of spatial spectra, is a particularly valuable technique for studying fine-scale structure in the Cosmic Microwave Background (CMB).  By measuring the power spectrum of the CMB, we can obtain a unique snapshot of the distribution and motions of matter in the Universe 400,000 years after the Big Bang, providing constraints on key cosmological parameters. Results from such observations have helped establish the overall geometry of the Universe, and therefore our understanding of its fate; new observations will test models of the earliest moments in the formation of the Universe.  Improvements in front-end receiving systems and in the back-end signal processing systems are needed to help progress the study of the CMB. Increasing the system bandwidth allows more precise radiometry and hence more sensitive measurements. Here we describe a cross-correlator that covers the very wide bandwidths needed to gather signals for measurements of tiny deviations from the CMB\u2019s 2.73~K blackbody radiation temperature. Wide bandwidths are needed to fully exploit the windows in the atmospheric transmission in the millimeter wavebands, which are typically many tens of GHz wide. This correlator operates over the frequency range 2--20~GHz, and is the back-end for a prototype wide-band interferometer, GUBBINS \\cite{GUBBINS}, currently under development. This telescope is a millimeter-wave (185--275~GHz) single-baseline interferometer employing SIS mixers and very wide bandwidth low-noise IF amplifiers. A schematic of the instrument is shown in Fig. \\ref{fig:schematic}.  Other broadband analog correlators which have been used in experiments devoted to the CMB have bandwidths up to 16~GHz (e.g.  \\cite{CBI}, \\cite{AMIBA}). The correlator we discuss here has a combination of broader bandwidth, lower complexity and lower power dissipation than previous correlators used in CMB interferometry, properties that enable the construction of larger interferometric arrays. ",
        "conclusions": "We have presented an ultra-broadband analog Fourier transform correlator. The instrument works well over 2--20~GHz with a possible maximum bandwidth of 1--23~GHz. The frequency band is split into 16 complex channels, each 1.25~GHz wide. The response of the correlator is well behaved over the band and the signal-to-noise of the system is not limited by the cross-correlator. We have also described the low frequency readout electronics necessary to digitise a fully phase-switched astronomical signal."
    },
    "1112/1112.3349_arXiv.txt": {
        "abstract": "{ The {\\it Fermi}, {\\it Swift} and {\\it INTEGRAL} satellites, together with ground based (especially Cherenkov) telescopes made possible a great progress in our understanding of relativistic jets. We can now start to attack the difficult questions of jet formation, collimation and content. We can also used them as probes to quantify the amount of IR and optical background radiation, and the amount of the cosmic magnetic field. Since they are the most powerful steady sources of the Universe, we can study them also at large redshifts, and this is a very fruitful field of research. To this aim, I will emphasize the importance of high energy X--rays, where very powerful blazars are predicted to emit most of their electromagnetic power. For them, the emission of the underlying accretion disk becomes unhidden by the non--thermal jet radiation, allowing to estimate the black hole mass and the accretion rate. In turn, this highlights the connection between the disk and the jet. Since the highest power blazars could have their emission peak in the $\\sim$MeV band, hard X--ray instruments could be more appropriate than the {\\it Fermi}/LAT to detect them. }  \\FullConference{The Extreme and Variable High Energy Sky - extremesky2011,\\\\ September 19-23, 2011\\\\ Chia Laguna (Cagliari), Italy}   \\begin{document} ",
        "introduction": "BL Lacs and Flat Spectrum Radio Quasars (FSRQs), collectively called blazars, have relativistic jets that point at us. Due to relativistic beaming, their flux is enhanced and they can therefore be visible up to large redshifts. This makes these sources good probes for studying the physics of jets and to explore some interesting properties of the far Universe. Their spectral energy distribution (SED) is always characterized by two broad humps (in $\\nu F_\\nu$), the first peaking at mm to UV frequencies, the second peaking in the MeV--GeV (and sometimes TeV) bands \\cite{fossati98}. While the origin of the first peak is certainly due to synchrotron, there is some debate about the origin of the high energy peak: the prevalent hypothesis is that it is due to the same electrons responsible for the synchrotron peak, scattering their own synchrotron photons in low power BL Lac sources (SSC, \\cite{maraschi92}), and scattering radiation produced externally to the jet (EC) in high power FSRQs \\cite{dermer93}, \\cite{sikora94}, \\cite{ghisellini98}. Through simultaneous data covering the IR to $\\gamma$--ray band, we can now derive several interesting parameters of the jet emitting region, and, in very powerful sources, we directly see the contribution of the accretion disk in the optical band. Thus, the black hole mass and accretion rate can be estimated, allowing to compare accretion and jet powers, both in absolute terms and when these quantities are measured in Eddington units.  The {\\it Fermi} satellite allowed a huge jump in strengthening the knowledge of the SED of blazars since the {\\it Compton Gamma Ray Observatory} era, detecting several hundreds of blazars of all kinds. But the added value of X--ray observations is also huge: hard X--rays above 10 keV in blazars are particularly important at the two extremes of the so--called ``blazar sequence\" \\cite{fossati98}: i) in low power BL Lacs they can be due to the tail of the synchrotron spectrum, making them good candidates as strong TeV emitters: ii) on the high power end of the sequence, namely in very powerful FSRQs, the hard X--ray flux is close to the emission peak, that in these sources is in the MeV energy range, and is dominating the bolometric output. Therefore in these sources hard X--rays carry a very significant fraction of the jet luminosity, making them visible and detectable at very high redshift.  This poses the question: to find out the most powerful blazars at high redshift, what is the best energy band and instrument? Hard X-rays (thus INTEGRAL and {\\it Swift}/BAT) or $\\gamma$--rays (i.e. {\\it Fermi}/LAT)? The answer of course depends on the average source flux in the two bands coupled with the corresponding sensitivity. I will here argue that hard X--rays are more promising.  \\begin{figure} \\vskip -0.4 cm \\center \\includegraphics[width=.65\\textwidth]{0227_369f.ps} \\vskip -0.7 cm \\caption{ The SED blazar 0227--369 from the radio to the $\\gamma$--rays. The lines refer to the leptonic, one--zone model used in \\cite{ghisellini09}. The different components are labelled. Note the synchrotron jet continuum that peaks in the mm band and is steep [$\\alpha>1$; $F(\\nu) \\propto \\nu^{-\\alpha}$] after the peak. This leaves the accretion disk contribution unhidden. By modelling the disk component we can derive the black hole mass and accretion rate. } \\label{0227} \\end{figure} ",
        "conclusions": ""
    },
    "1112/1112.1666_arXiv.txt": {
        "abstract": "I briefly review the flatness problem within the context of classical cosmology and examine some of the debate in the literature with regard to its definition and even the question whether it exists.  I then present some new calculations for cosmological models which will collapse in the future; together with previous work by others for models which will expand forever, this allows one to examine the flatness problem quantitatively for all cosmological models.  This leads to the conclusion that the flatness problem does not exist, not only for the cosmological models corresponding to the currently popular values of \\lnull\\ and \\onull\\ but indeed for all Friedmann\\ndash Lema\\^{\\i}tre models. ",
        "introduction": "The flatness problem has been called one of the outstanding puzzles in cosmology \\citep[\\eg][]{RDickePPeebles79a}.  This in itself is rather puzzling in view of the fact that the arguments in favour of it being a problem are rather vague and heuristic, while quantitative arguments have been presented against the claim that it is a problem, at least for some classes of cosmological models \\citep[\\eg][]{PColesGEllis97a,KLake05a}.  The flatness problem is one of the main motivations for the inflationary scenario \\citep{AGuth81a}. Of course, if there is no flatness problem (or, indeed, even if there were no motivation at all for inflation), this does not mean that inflation could not have occurred.  However, it does mean that inflation should not be taken as given based on the belief that it explains away the flatness problem and thus without it classical cosmology leads to absurd conclusions.  The plan of this paper is as follows.  In \\Sect~\\ref{cosmo} I present the basic equations needed in the rest of the paper, mainly to define my notation (unfortunately, there is not a uniform notation in the literature) and give an overview of the entire cosmological parameter space relevant to the discussion.  Section~\\ref{history} gives a brief historical overview of the flatness problem and some qualitative arguments against it.  In \\Sect~\\ref{collapse} I discuss a new quantitative argument regarding cosmological models which will collapse in the future.  Sections~\\ref{Lake} and~\\ref{anthro} discuss previous quantitative results by others for other classes of cosmological models. Section~\\ref{summary} summarizes the results for all cosmological models. ",
        "conclusions": "\\label{summary}   The qualitative flatness problem, \\ie~the puzzle why the universe was arbitrarily close to the Einstein\\ndash de~Sitter model\\footnote{Or, for an empty universe, the Milne or de~Sitter model.} at early times, does not exist.  It is merely a consequence of the way $\\lambda$ and $\\Omega$ are defined.  Neither does the quantitative flatness problem exist: although the cosmological parameters in general evolve with time, it is not puzzling that we don't observe extreme values for them today.  In the case of models which will collapse in the future this is because large (absolute) values of $\\lambda$ and $\\Omega$ occur only during a relatively short time in the lifetime of such a universe, namely near the time of maximum expansion.  $\\lambda$ and $\\Omega$ can become large only when $H$ becomes small, and this happens only during the time when the universe is at or near its maximum size.  [Arbitrarily small (absolute) values, if they occur at all, also occur for only a relatively short time].  For models which will expand forever, large values are possible only for $k=+1$.  However, this occurs only for $\\alpha \\approx 1$.  In this case, the fine-tuning argument is reversed; only in the case of fine-tuning do $\\lambda$ and $\\Omega$ become arbitrarily large.  Since all models which will expand forever asymptotically approach $\\Omega=0$, arbitrarily small values of $\\Omega$ can occur.  Those with $\\lambda=0$ (and hence $k=-1$) approach the Milne model with $\\Omega=0$; models with $\\lambda>0$, whatever the value of $k$, approach the de~Sitter model with $\\lambda=1$ (the Milne and de~Sitter models themselves are of course stationary points).  (If $\\lambda=0$ at any time then $\\lambda=0$ at all times.  Otherwise, arbitrarily small values of $\\lambda$, if they occur at all, occur only for a relatively short time.)  However, if \\hnull\\ has a value similar to or smaller than the observed value, small values of $\\Omega$ will occur only in the far future when anthropic arguments probably make the observation of such a low value of $\\Omega$ unlikely.  While (for $\\lambda>0$) a higher value of \\hnull\\ would allow a low value of $\\Omega$ even for an age near the observed age, such a universe would have spent only a very short time during which $\\Omega$ was not very small, so structure formation would have been strongly suppressed.  It is interesting to note that the three arguments presented here make it unlikely that we would observe extreme values of \\lnull\\ or \\onull. This automatically solves the so-called coincidence problem, which has been called deeply puzzling \\citep*[\\eg][]{MTegmarkAVilenkinLPogosian05a}.  Also interesting is that in the appendix to his seminal paper, \\citet{AGuth81a} anticipates much of the subsequent discussion.  He points out that even though essentially all cosmological models begin arbitrarily close to the Einstein\\ndash de~Sitter universe (the qualitative flatness problem), this still leaves the question as to why the universe is so close to the Einstein\\ndash de~Sitter universe today (the quantitative flatness problem).  He also rejects the idea that some basic principle must force the universe to conform exactly to the Einstein\\ndash de~Sitter model on the grounds that this is obviously only an approximation in the case of the real universe (see the discussion of this above).  He basically recasts the flatness problem as the longevity problem: fine tuning is required in order that the universe does not recollapse or thin out to extremely low density within a very short time.  However, this argument relies on an assumption for the value of \\hnull, while our arguments do not need this assumption except in the third category.  Also, as I have shown here, even in an extremely short-lived universe (which of course recollapses), extreme values of $\\lambda$ or $\\Omega$ are observed only during a relatively small fraction of the lifetime of the universe.  If there is no flatness problem, what does this mean for inflation?  If there is no flatness problem (or, indeed, even if there were no motivation at all for inflation), this does not mean that inflation could not have occurred.  However, it does mean that inflation should not be taken as a given based on the belief that it explains away the flatness problem and thus without it classical cosmology leads to absurd conclusions.  Inflation also solves the monopole and isotropy problems. However, the monopole problem seems to be more a problem with theories of particle physics than with cosmology \\citep{JNarlikarTPadmanabhan91a} while \\citet{JBarrow95a} claims that there is no isotropy problem.\\footnote{It is, however, debatable whether Barrow's conclusion is as general as he claims or depends too strongly on his assumptions.} If all of these claims are true, then this still does not prove that inflation didn't happen, but the necessity for inflation or something like it is weakened if not destroyed altogether."
    },
    "1112/1112.6270_arXiv.txt": {
        "abstract": "{} {We construct a  theoretical model to predict the number of orphan afterglows (OA) from  gamma-ray bursts (GRBs) triggered by  primordial  metal-free  (Pop III) stars  expected to be observed by the Gaia mission. In particular, we consider primordial metal-free stars  that were affected by radiation from other stars (Pop III.2) as a possible target. } {We use a  semi-analytical approach that includes  all relevant feedback effects   to construct  cosmic  star formation history and its connection with  the cumulative number of GRBs. The OA events are generated using the  Monte Carlo method, and realistic simulations of Gaia's scanning law are performed  to derive the  observation probability expectation.} {We show that Gaia can observe  up to 2.28 $\\pm$ 0.88  off-axis afterglows and  2.78 $\\pm$ 1.41  on-axis during the five-year nominal mission. This implies that a  nonnegligible percentage of afterglows   that may be  observed by  Gaia ($\\sim 10\\%$) could  have  Pop III stars as progenitors.  } {} ",
        "introduction": "The first stars (hereafter, Pop~III-primordial metal-free) in the Universe are thought to have played a crucial role in the early cosmic evolution by emitting the first light and producing the first heavy elements \\citep{bromm09}.  Understanding  such objects is very important since their detection would permit  the pristine regions of the Universe to be probed. However, there has been no direct observation of the so-called Pop~III  stars up to now.  Pop III stars may produce collapsar gamma-ray bursts (GRBs) whose total isotropic energy could be $\\approx 2$ orders of magnitude larger than average \\citep{barkov2010,komissarov2010,meszaros2010,suwa2011,toma2011}. Even if the Pop III star has a supergiant hydrogen envelope, the GRB jet can break out of it because of the long-lasting accretion of the envelope itself \\citep{nagakura2011,suwa2011}. It is of great importance to study the rate and detectability of Pop~III GRB prompt emissions and afterglows in current and future surveys.  We explore here the  possibility to observe these objects through  their  afterglows \\citep{toma2011}. Observations of GRB afterglows make it possible to derive physical properties of the explosion mechanism and the circumburst medium. It is intriguing to search for signatures of metal-poor stars in the GRB afterglows at low and high redshifts.   GRB optical afterglows are one of the possible transients to be detected by the Gaia\\footnote{http://www.rssd.esa.int/GAIA/} mission. Recently  \\citet{Japelj2011}  explored the detectability of such afterglows with Gaia using a Monte Carlo approach that inspired us.  As the GRB jet sweeps the interstellar medium, the Lorentz factor of the jet is decelerated and the  jet starts to expand sideways, eventually  becoming detectable by off-axis observers. These afterglows are not associated with the prompt GRB emission and are called orphan afterglows (OA) \\citep{nakar2002, rossi2008}.   De Souza et al. (2011) showed that, considering EXIST\\footnote{http://exist.gsfc.nasa.gov/design/} specifications, we can expect to observe a maximum of $\\approx 0.08$ GRBs with $z>10$  per year originating from primordial metal-free stars (Pop III.1) and $\\approx 20$ GRBs with $z>6$ per year coming from primordial metal-free stars that were affected by the radiation from other stars (Pop III.2). In the context of the current \\textit{Swift}\\footnote{http://swift.gsfc.nasa.gov/docs/swift/swiftsc.html} satellite, $\\approx 0.2$ GRBs with $z>6$ per year from Pop III.2 stars are expected.  These numbers reflect the fact that, compared to Pop~III.1 stars, Pop III.2 stars  are more abundant and can be observed in a lower redshift range, which makes them more suitable targets. In  the light of such results, the calculations presented here will focus on Pop III.2 stars alone.  Searches have been made of OAs  by both X-ray surveys \\citep{grindlay1999,greiner2000} and  optical searches \\citep{becker2004,rykoff2005,rau2006,malacrino2007}. The purpose of the present paper is to calculate the Pop~III.2 GRB OA  rate that might be detected by the Gaia mission \\citep[for  more details about Gaia,  see,  e.g.,][]{Perryman:2001p3838, Lindegren:2009p8828}.  The Gaia mission is one of the most ambitious projects of modern astronomy. It aims to  create  a very precise tridimensional, dynamical, and chemical  census of our Galaxy from astrometric, spectrophotometric,  and spectroscopic data. In order to do this, the Gaia satellite will perform observations of the entire sky in a continuous scanning created from the coupling of rotations and precession movements called  the scanning law. For point sources, these observations will be unbiased and the data of all the objects bellow a certain limiting magnitude (G=20)  will be transferred to the ground. Certainly, galactic and extragalactic sources will be among those objects.  Typically, Pop~III.2 stars are formed in an initially ionized gas \\citep{Johnson06,Yoshida07}. They are thought to be less massive  than Pop~III.1 stars but still massive enough to produce GRBs. Recent results from \\citet{greif2011} show that, instead of forming a single object, the gas in mini-halos fragments vigorously into a number of protostars with a range of different  masses. It is not clear up to now how this initial range of mass will be  mapped into the final mass function of Pop III stars.  The most likely conclusion is that Pop III stars are less likely to reach masses in excess of $\\sim 140 M_{\\odot}$, which consequently affect the estimated number of GRBs from Pop III.1.  \\citet{hosokawa2011}, performing state-of-the-art radiation-hydrodynamics simulations, showed that the typical mass of Pop III stars  could be $\\sim 43 M_{\\odot}$. Here we assume that this will not affect significantly the mass range  assumed for Pop III.2 ($\\sim 10-100 M_{\\odot}$).   The paper is organized as follows. In Sect. 2, we  calculate the formation rate of primordial GRBs. In Sect. 3, we calculate the OA light curves and their redshift distribution. In Sect. 4, we discuss the details of the   Gaia mission and  derive the probability of a given event to be observed by Gaia. In Sect. 5,  we discuss the results,  and  finally,  in Sect. 6,  we give our concluding remarks. Throughout the paper,  we adopt the standard $\\Lambda$ cold dark matter model with the best-fit cosmological parameters from   \\citet{jarosik2010} (WMAP-Yr7\\footnote{http://lambda.gsfc.nasa.gov/product/map/current/}), $\\Omega_{\\rm m} = 0.267, \\Omega_{\\Lambda} = 0.734$, and $H_0 = 71 {\\rm km}~{\\rm s}^{-1}{\\rm Mpc}^{-1}$. ",
        "conclusions": "It is important to emphasize that  our knowledge concerning first stars and their GRBs is still quite incomplete. Many of their properties (e.g.,  characteristic mass, SFR and efficiency to trigger GRBs) are still very uncertain, and more reliable information can only come once a detection is confirmed.  Recently, \\citet{hosokawa2011}, performing state-of-the-art radiation-hydrodynamics simulations, showed that the typical mass of primordial stars could be $\\sim 43 M_{\\odot}$, i.e.,  less massive than originally expected by theoretical models. Their results,  though,  are affected by assumptions on the initial conditions. This confirms that we are far away from understanding all characteristics of these objects and any observation would be of paramount importance to improve theoretical models.   In this work, we estimated the average number of OAs events originating from Pop III stars that the Gaia mission may observe to be up to   2.28 $\\pm$ 0.88  off-axis afterglows and  2.78 $\\pm$ 1.41   on-axis ones. In case such events are found among Gaia data, valuable physical properties associated with  the primordial stars of our Universe and their environment could  be constrained."
    },
    "1112/1112.5981_arXiv.txt": {
        "abstract": "This article reviews the present knowledge on oscillating main sequence A/F stars that belong to more than one class of pulsator. Due to recent results of asteroseismic space missions, we now know of many $\\delta$~Scuti/$\\gamma$~Doradus stars. However, $\\gamma$~Doradus variability was also detected in a rapidly oscillating Ap star, and solar-like oscillations were discovered in a $\\delta$~Scuti star. The astrophysical information that can be gained from these pulsators is discussed, and confronted with what is believed to be known about pulsational driving. ",
        "introduction": "Over the last 40 years, our knowledge about pulsating stars has greatly expanded. The ever increasing accuracy in stellar observations and advances in theory have revealed a large number of previously unknown classes of variable star. Figure 1 compares the knowledge on the positions of pulsating stars in the HR diagram over this time span.  \\begin{figure} \\plotfiddle{hakone1.ps}{8.3cm}{00}{72}{72}{-203}{-40} \\caption{Left: classes of pulsating star known in the late 1960's. Right: a selection of classes of pulsating star known to date. Areas hatched from lower right to upper left depict domains of g-mode pulsators, areas hatched from lower left to upper right delineate domains of p-mode pulsators; overlapping areas may contain \"hybrid\" pulsators. Parts of model evolutionary tracks for main sequence, horizontal branch and post-AGB stars are shown as dashed-dotted lines for orientation.} \\end{figure}  The most striking difference between the two panels in this figure is that the HR Diagram nowadays is well filled with pulsating variables. This has the important consequence that stellar oscillations can be utilized to sound the interiors of stars in many different evolutionary stages, and over wide ranges of mass and chemical composition. Besides better defined locations of the pulsating stars over the last 40 years, another new feature is obvious: several of the different instability strips overlap.  These instability strips contain stars with different pulsational behaviour, or in other words, with different types of mode excited. Consequently, it can be suspected that stars possessing more than a single set of mode spectra exist. Ground based observations have shown that they do, and the history of their study has been reviewed elsewhere (e.g., Handler 2009).  The most complicated region in the HR Diagram in terms of overlap of different classes of pulsating star is where the lower classical instability strip intersects the main sequence, at spectral types of A and F. There one finds: \\begin{itemize} \\item $\\delta$ Scuti stars, pulsating in radial and nonradial pressure and mixed modes of low radial order with periods between 20 minutes and 6 hours, driven by the classical $\\kappa$ mechanism in the He II ionization zone (Baker \\& Kippenhahn 1962) \\item rapidly oscillating Ap (roAp) stars, high-order pressure modes pulsators governed by a global magnetic field, with periods of 5 - 20 minutes, believed to be driven by the $\\kappa$ mechanism in the He I/H ionization zone (Balmforth et al.\\ 2001) \\item $\\gamma$ Doradus stars that oscillate in high-order g modes with periods between 0.3 - 3 days, and likely driven by blocking of flux at the base of the envelope convection zone (Guzik et al.\\ 2000) \\item solar-like oscillations with periods around 10 - 40 minutes are theoretically predicted to be present and driven by surface convection (Houdek et al.\\ 1999) \\end{itemize}  The present article reviews the present knowledge on hybrid pulsators in this domain in the light of recent results of asteroseismic space missions. There will be a strong focus on Kepler data, not only because they have the best quality, but for the practical reason that the author is involved in their analysis. ",
        "conclusions": "The $\\delta$ Scuti stars, the roAp stars, the $\\gamma$ Dor stars and solar like oscillators all overlap at the intersection of the classical instability strip and the main sequence. Each of these classes of pulsating star are driven in different regions of the star. To date, we know hybrid pulsators between: $\\delta$ Scuti and $\\gamma$ Dor stars, roAp stars and $\\gamma$ Dor stars as well as $\\delta$ Scuti stars and solar-like oscillators. The other three possible combinations of two classes of pulsator have not yet been observed.  This status makes sense from the point of view of pulsational driving (as we believe to understand it to date), because the mechanisms operating in the known hybrids, or their basic physical characteristics, do not affect each other negatively. On the other hand, the absence of roAp/solar-like hybrids is also not a surprise, as the strong magnetic fields of Ap stars are believed to suppress surface convection. Similar arguments may be invoked to explain the non-detection of roAp/$\\delta$ Scuti hybrids: the diffusion and settling of chemical elements is supposed to deplete the driving zone for $\\delta$ Scuti oscillations of the required Helium, and slow wave leakage due to strong magnetic fields is believed to provide additional damping.  What about $\\gamma$ Doradus/solar-like hybrids? Such stars have not been reported to date, but the author believes this is only a question of time. Along those lines, given the large fraction of Delgam Scudor stars, one may also expect the existence of $\\delta$ Scuti/$\\gamma$ Doradus/solar-like hybrids, which would be another goldmine for asteroseismology of main sequence stars."
    },
    "1112/1112.5723_arXiv.txt": {
        "abstract": "We calculate the angular power spectrum of galaxies selected from the Sloan Digital Sky Survey (SDSS) Data Release 7 (DR7) by using a quadratic estimation method with KL-compression. The primary data sample includes over 18 million galaxies covering more than 5,700 square degrees after masking areas with bright objects, reddening greater than 0.2 magnitudes, and seeing of more than 1.5 arcseconds.  We test for systematic effects by calculating the angular power spectrum by SDSS stripe and find that these measurements are minimally affected by seeing and reddening.  We calculate the angular power spectrum for $\\ell \\le 200$ multipoles by using 40 bandpowers for the full sample, and $\\ell \\le 1000$ multipoles using 50 bandpowers for individual stripes.  We also calculate the angular power spectrum for this sample separated into 3 magnitude bins with mean redshifts of $z = 0.171$, $z = 0.217$, and $z = 0.261$ to examine the evolution of the angular power spectrum.  We determine the theoretical linear angular power spectrum by projecting the 3D power spectrum to two dimensions for a basic comparison to our observational results.  By minimizing the $\\chi^2$ fit between these data and the theoretical linear angular power spectrum we measure a loosely-constrained fit of $\\Omega_m = 0.31^{+0.18}_{-0.11}$ with a linear bias of $b = 0.94 \\pm 0.04$. ",
        "introduction": "\\label{Introduction}  The angular power spectrum, $C_{\\ell}$, is a statistical measure that quantitatively characterizes the large scale angular distribution of matter \\citep{peebles73}. Therefore, calculating the angular power spectrum of galaxies is useful as both a method of data compression, reducing clustering information of an arbitrary number of galaxy positions down to a set of $C_{\\ell}$ and their corresponding window functions, and also since the $C_{\\ell}$ values derived from the observations can be easily compared to theoretical predictions.  Calculations of angular power spectra are well known to cosmologists for their usefulness in studying the Cosmic Microwave Background (CMB), as the CMB provides a detailed and precise measurement of the density variations in the early universe (e.g., \\citealt{smoot92,netterfield02,spergel07}).  However, to study large scale structure in other eras, it is necessary to analyze how mass clusters by using galaxies as a tracer of the underlying dark matter distribution.  Angular power spectra of galaxies have been calculated for galaxy surveys with various depth and survey area (e.g., \\citealt{huterer01,blake04,frith05}) including the SDSS (\\citealt{tegmark02}, hereafter T02; \\citealt{blake07,thomas10}). By using angular power spectra to calculate galaxy clustering, we study the Fourier modes of the galaxy distribution; this method is most sensitive to large scale effects. Recent galaxy surveys such as the APM Galaxy Survey \\citep{maddox90}, the Two Micron All Sky Survey \\citep{skrutskie06}, and the SDSS \\citep{abazajian09} have cataloged large areas of the sky, thereby providing enormous numbers of galaxies for which we can measure angular clustering.  However, to date the galaxy angular power spectrum has not been calculated for the full SDSS main galaxy sample.  In this paper, we address this deficiency.  The angular power spectrum is useful for large scale clustering, while it is complemented by the two-point angular correlation function on small scales.  The two-point angular correlation function (e.g., \\citealt{brunner00,myers07,ross10}), which is related to the angular power spectrum by the Legendre transform (T02), is more sensitive to smaller scale clustering because the calculation is done in configuration space where the distances between nearby pairs of galaxies can be calculated faster.  This makes the two-point angular correlation function advantageous to use on scales where non-linear evolution is important.  This regime is also where the angular power spectrum at large $\\ell$ is more difficult to measure and model, partly due to correlations introduced between the $C_{\\ell}$.  To calculate the angular power spectrum, we want to find the most probable parameters $C_{\\ell}$ that could produce the data we observe.  To do this, we need the likelihood function of the angular power spectrum, which is proportional to the probability of the data given the $C_{\\ell}$.  Though in theory we would like to know the entire likelihood function, calculating this $\\ell_{max}$-dimensional function is difficult \\citep{oh98}.  Fortunately, since we are only interested in the most probable $C_{\\ell}$, we only really need to know the maximum of this function.  To determine the $C_{\\ell}$ that maximize the likelihood function, we use the quadratic estimation method (\\citealt{tegmark97a,bond98}, hereafter BJK98).  This technique fits a quadratic function to the shape of the likelihood function for some initial angular power spectrum, finds the $C_{\\ell}$ that maximize this quadratic, and uses these $C_{\\ell}$ for a new quadratic fit to iteratively converge to the true maximum of the likelihood function. Once we have found the angular power spectrum of galaxies, we can use the results to infer what cosmological parameters are consistent with the measurement (e.g., \\citealt{jaffe99}).  In this paper, we discuss the SDSS DR7 data, our selected sample and subsamples, and our systematic tests and masks in Section \\ref{Data}. In Section \\ref{Method}, we discuss our pixelization scheme, KL-compression,  and the quadratic angular power spectrum estimation method of BJK98 in detail.  In Section \\ref{Results}, we apply this estimator to the complete SDSS DR7, selected subsamples, and individual SDSS stripes, and present the results.  We construct a theoretical linear angular power spectrum to compare with the observational results, and we extract cosmological matter density and linear bias from this computation in Section \\ref{Theory}.  Finally, we discuss our results in Section \\ref{Discussion}, and conclude the paper in Section \\ref{Conclusion}. ",
        "conclusions": "\\label{Conclusion}  We have used the quadratic estimation method with KL-compression to determine the SDSS DR7 angular power spectrum, first as a means of radical compression of the angular clustering information, and second to match these observed angular power spectra with theoretical angular power spectra to extract the linear bias and cosmological matter density.  We masked for observational effects and applied this method to over 18 million SDSS DR7 galaxies and three magnitude subsamples out to $\\ell \\le 200$.  We also measured the angular power spectrum for each individual stripe out to $\\ell \\le 1000$ for stripes 9--37.  We have used the photometric redshift distribution of these galaxies to project the 3D power spectrum to two dimensions to obtain theoretical linear angular power spectrum, and used $\\chi^2$ minimization to determine the best fit parameters given the observations.  As the linear angular power spectrum approximation is not valid for the entire range of our estimated angular power spectrum, these parameter constraints have a large allowed range of values.  We found that the linear bias of our samples was $b = 1.09 \\pm 0.05$ in the 18--19 magnitude range, $b = 1.03 \\pm 0.04$ for 19--20, and $b = 0.92 \\pm 0.04$ for 20--21, with an overall bias of $b = 0.94 \\pm 0.04$ for our combined 18--21 magnitude sample.  We have also calculated the cosmological density of matter as $\\Omega_m = 0.31^{+0.18}_{-0.11}$ from our entire sample."
    },
    "1112/1112.0330_arXiv.txt": {
        "abstract": "The dark energy dominated warm dark matter (WDM) model is a promising alternative cosmological scenario. We explore large-scale structure formation in this paradigm. We do this in two different ways: with the halo model approach and with the help of an ensemble of high resolution $N$-body simulations. Combining these quasi-independent approaches, leads to a physical understanding of the important processes which shape the formation of structures. We take a detailed look at the halo mass function, the concentrations and the linear halo bias of WDM. In all cases we find interesting deviations with respect to CDM. In particular, the concentration-mass relation displays a turnover for group scale dark matter haloes, for the case of WDM particles with masses of the order $m_{\\rm WDM}\\sim0.25 \\keV$. This may be interpreted as a hint for top-down structure formation on small scales. We implement our results into the halo model and find much better agreement with simulations.  On small scales the WDM halo model now performs as well as its CDM counterpart. ",
        "introduction": "\\label{sec:Intro}  Over the last decade the vacuum energy dominated cold dark matter (hereafter $\\Lambda$CDM) scenario, has emerged as a standard model for cosmology.  This owes largely to the combination of information from galaxy clustering surveys such as the 2dFGRS and SDSS with WMAP measurements of the temperature anisotropies in the microwave background \\citep{Coleetal2005short,Tegmarketal2006short,Komatsuetal2011short}. However, the nature of the two dark components in the $\\Lambda$CDM model are still completely unknown and it is therefore important to keep exploring alternative models and test their compatibility with observations.  In the $\\Lambda$CDM model the dark matter is assumed to be composed of heavy, cold thermal relic particles that decoupled from normal matter very early in the history of the Universe \\citep{Peebles1982,Blumenthaletal1984,KolbTurner1990,Jungmanetal1996}. Whilst there is a large body of indirect astrophysical evidence that strongly supports CDM, there are some hints that it has shortcomings. Firstly, CDM galaxy haloes contain a huge number of subhaloes \\citep{Mooreetal1999c,DiemandKuhlen2008,Springeletal2008,Stadeletal2009}, while observations indicate that only relatively few satellite galaxies exist around the Milky Way and M31 \\citep{Mooreetal1999c,Klypinetal1999}. Secondly, the highest resolution halo simulations show that the slope of the inner density profile decreases linearly at smaller radii \\citep{Navarroetal1997,Mooreetal1999c,Diemandetal2004,Springeletal2008,Stadeletal2009}, whereas the density profiles inferred from galaxy rotation curves are significantly shallower \\citep{Mooreetal1999c} \\citep[and for recent studies see][and references there in]{Swatersetal2003,Saluccietal2007,deBloketal2008,Gentileetal2009}. Thirdly, the observed number of dwarf galaxies in the voids appears to be far smaller than expected from CDM \\citep{Peebles2001,Tikhonovetal2009,PeeblesNusser2010}. Another example is the excess in the prediction of dwarf galaxy concentrations \\citep{Lovelletal2011}. Whilst, it has become clear that some of these discrepancies might be resolved through an improved understanding of galaxy formation, they have led some to consider changes to the $\\Lambda$CDM paradigm.  One possible solution might be warm dark matter (WDM) \\citep{BondSzalay1983,Bardeenetal1986,Bodeetal2001}. In this scenario, the dark particle is considered to be lighter than its CDM counterpart, and so remains relativistic longer and also retains a thermal velocity.  Since WDM particles are collisionless and decouple early, they may `free-stream' or diffuse out of perturbations whose size is smaller than the Jeans' length\\footnote{Although originally defined in the context of gas dynamics, the Jeans length can be generalized to collisionless systems by replacing the sound speed with the velocity dispersion. The reason for this tight analogy lies in the linearized equation of perturbations, which has the same structure for gas and collisionless fluids \\citep[see][for more details]{Peebles1982}.} in the radiation dominated Universe \\citep{KolbTurner1990}. This free-streaming of the WDM particles acts to damp structure formation on small scales. Two potential candidates are the sterile neutrino \\citep{DodelsonWidrow1994,ShaposhnikovTkachev2006}, and the gravitino \\citep{Ellisetal1984,Moroietal1993,Kawasakietal1997,Gorbunovetal2008}, both of which require extensions of the standard model of particle physics.  Recent observational constraints have suggested that sterile neutrinos can not be the dark matter: the Lyman alpha forest \\citep{Seljaketal2006b,Boyarskyetal2009a} and QSO lensing \\citep{MirandaMaccio2007} bounds are $m_{\\nu_s}>8\\keV$, whilst those from the X-ray background are $m_{\\nu_s}<4\\keV$ \\citep{Boyarskyetal2008}\\footnote{Lower bounds on the mass of a fully thermalized WDM particle can be obtained using \\Eqn{eq:massrescale} \\citep[see][]{Vieletal2005}.}.  However, a more recent assessment has suggested that a better motivated particle physics model based on resonant production of the sterile neutrino, may evade these constraints: the Lyman alpha forest bound is brought down to $m_{\\nu_s}\\gtrsim 2\\keV$ and the X-ray background is pushed to $m_{\\nu_s}<50\\keV$ (for very low mixing angles) \\citep{Boyarskyetal2009b}. It therefore seems that additional, independent methods for constraining the $\\Lambda$WDM scenario would be valuable.  In \\citet{Markovicetal2010} and \\citet{SmithMarkovic2011}, it was proposed that the $\\Lambda$WDM scenario could be tested through weak lensing by large-scale structure. The advantage of such a probe is that it is only sensitive to the total mass distribution projected along the line of sight. However, to obtain constraints on the WDM particle mass, an accurate model for the nonlinear matter clustering is required. In these papers, an approach based on the halo model was developed. Accurate predictions from this model require: detailed knowledge of the abundance of dark matter haloes, their spatial large-scale bias, and their density profiles. In these studies, it was assumed that the semi-analytic methods, which were developed for CDM, would also apply to WDM.  In this paper we perform a series of very high resolution CDM and WDM $N$-body simulations with the specific aim of exploring the halo model ingredients in the $\\Lambda$WDM scenario.  Over the past decade, there have been a limited number of numerical simulation studies of nonlinear structure formation in the WDM model \\citep{Colombietal1996,Mooreetal1999c,Colinetal2000,WhiteCroft2000,Avila-Reeseetal2001, Bodeetal2001,Bullocketal2002,ZentnerBullock2003,Colinetal2008,Zavalaetal2009,MaccioFontanot2010,Lovelletal2011,Vieletal2011,Dunstanetal2011}. In most of these previous studies, conclusions have been drawn from object-by-object comparison of a relatively small number of haloes simulated in boxes of typical size $L=25\\Mpc$. In this work we are more interested in the overall impact that the WDM hypothesis has on the statistical properties of large-scale structures. We therefore simulate boxes that are 10 times larger than have been typically studied before, hence having roughly $\\sim1000$ times larger sampling volume. This means, that our conclusions will have greater statistical weight, than those from previous studies. Furthermore, our results should be less susceptible to finite volume effects, which can lead to underestimates of the nonlinear growth.  The paper is structured as follows: In \\S\\ref{sec:WDMtheory} we provide a brief overview of the salient features of linear theory structure formation in the WDM model and we review the halo model approach. In \\S\\ref{sec:Simulations} we describe the $N$-body simulations. In \\S\\ref{sec:Ingredients} we explore the main ingredients of the halo model: the halo mass function, bias and density profiles. In \\S\\ref{sec:Comparison} we compare the halo model predictions for the matter power with our measurements from the simulations. Finally, in \\S\\ref{sec:Conclusion} we summarize our findings. ",
        "conclusions": "\\label{sec:Conclusion}  In this paper we have explored nonlinear structure formation in the WDM cosmological model, through a large suite of cosmological $N$-body simulations and through the halo model. The study was done for a set of fully thermalized WDM models with particle masses in the set $m=\\{0.25,0.5,0.75,1.0,1.25\\} \\keV$. These masses range from purely pedagogical models, towards more realistic scenarios for the dark particle.  For the simulations we chose a box size ${L=256\\Mpc}$, which was small enough to resolve both the small scales, where WDM effects play an important role, and the large scales, which are required for correct linear evolution of the box-modes. All models were simulated with $N=\\{256^3, 512^3, 1024^3\\}$ particles. This was done in order to disentangle physical effects from numerical ones.  In the original halo model calculation for WDM by \\citet{SmithMarkovic2011}, it was shown that in order to make robust predictions, one requires good understanding of dark matter halo profiles, the mass function and halo bias.  In this work we performed a detailed study of all of these ingredients. Our findings can be summarized as follows:  \\begin{enumerate}{\\leftmargin = 1.0em} \\item Mass function: Below a certain scale, the WDM mass function is suppressed with respect to CDM. This suppression is considerably stronger than that obtained by simply applying the Sheth-Tormen approach together with the linear power spectrum of WDM. In agreement with \\citet{SmithMarkovic2011}, we found that the mass functions for the different WDM models could be transformed into a single locus of points. This was achieved by taking the ratio of the WDM mass function with that for CDM, and then rescaling the masses by $M_{\\rm hm}$ (or equivalently $M_{\\rm fs}$). We used a fitting function similar to that proposed in \\citet{Dunstanetal2011} to link the Sheth-Tormen mass function to the measured one. The fitting function, which has only one free parameter, was able to reproduce all of the data with an accuracy of a few percent. We also found a strong boost in the mass function at very small mass scales. We showed that this was consistent with artificial halo formation around the initial particle lattice \\citep[cf.][]{WangWhite2007}.  \\item Halo bias: We measured the linear halo bias, using the four largest modes in our simulations. For smaller mass haloes, we found a small enhancement of the bias in WDM simulations, which was qualitatively consistent with the predictions of \\citet{SmithMarkovic2011}. However, owing to the simulation box being too small, we were unable to quantify this more robustly.  At very small masses we found a prominent boost in the bias. We found that this was again a sign of artificial halo formation.  \\item Density profiles: In the CDM model, the density profiles of dark matter haloes can be characterized by an NFW profile, with a monotonically decreasing concentration-mass relation. In the WDM scenario, we have shown that the NFW profile remains valid for the models and resolution limits of our simulations, and we saw no evidence for a central density core. A simple adaption of the CDM concentration-mass relation, would suggest a strong flattening towards small masses. Whilst, we found such a flattening, the measurements in fact revealed a turnover towards smaller masses. This somewhat surprising result may be interpreted as a sign of top-down structure formation.  We modelled the mean relation by adapting a fitting formula similar to that for the mass function. Our fit to the $c(M)$ data was good to an accuracy of $\\sim10\\%$. Interestingly, we found that the deviations from CDM in the WDM model, appear in the $c(M)$ relation for halo masses one order of magnitude larger than for the mass function.  \\end{enumerate}  After analyzing these ingredients in detail and developing new fitting functions for them, we were able to improve the small-scale performance of the WDM halo model. We found that for $k\\gtrsim3\\kMpc$, we could predict the absolute amplitude of the power spectrum to better than $\\sim10\\%$. However, we were able to predict the ratio of the WDM to CDM spectra, at better than $\\lesssim2\\%$. This was competitive with the latest fitting formulae \\citep{Vieletal2011}.  One of the many advantages of the halo model based approach, is that we may more confidently extrapolate our power spectra predictions to smaller scales than can be done from a fitting formula, since the model is built on physical quantities. Furthermore, we may also use the model to study the clustering of galaxies \\citep{Zehavietal2005}. It is hoped that this may lead to a method for constraining WDM models from galaxy clustering studies.  Lastly, one further issue for future study, is to establish a better theoretical understanding of what shapes the mass function and halo concentrations in WDM. In particular, in finding the turnover in the concentration mass relation, have we really seen the reversal of bottom-up structure formation. This promises to be an interesting future challenge."
    },
    "1112/1112.2103_arXiv.txt": {
        "abstract": "In the current-driven, kink-type Tayler instability (TI) a sufficiently strong azimuthal magnetic field becomes unstable against non-axisymmetric perturbations. The TI has been discussed as a possible ingredient of the solar dynamo mechanism and a source of the helical structures in cosmic jets. It is also considered as a size limiting factor for liquid metal batteries. We report on a liquid metal TI experiment using a cylindrical column of the eutectic alloy GaInSn to which electrical currents of up to 8 kA are applied. We present results of external magnetic field measurements that indicate the occurrence of the TI in good agreement with numerical predictions. The interference of TI with the competing large scale convection, resulting from Joule heating, is also discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.0106_arXiv.txt": {
        "abstract": "A molecular line survey has been carried out toward the carbon-rich asymptotic giant branch star CW~Leo employing the HIFI instrument on board of the {\\it Herschel} satellite. Numerous features from 480~GHz to beyond 1100~GHz could be assigned unambiguously to the fairly floppy SiC$_2$ molecule. However, predictions from laboratory data exhibited large deviations from the observed frequencies even after some lower frequency data from this survey were incorporated into a fit. Therefore, we present a combined fit of all available laboratory data together with data from radio-astronomical observations. ",
        "introduction": "\\label{introduction}  Silacyclopropynylidene, SiC$_2$, somewhat better known as silicon dicarbide, is a fascinating molecule for spectroscopists, structural and quantum chemists as well as astronomers. In 1926, uncataloged bands near 500~nm were discovered in the spectra of several carbon-rich asymptotic giant branch (AGB) stars~\\cite{SiC2_A-X_astro_1926a,SiC2_A-X_astro_1926b}. These are late-type stars which produce elements heavier than helium and which eject large quantities of gaseous material as well as dust which form a circumstellar envelope (CSE). Thirty years later, laboratory spectroscopy established that the molecule SiC$_2$ is the carrier of these bands~\\cite{SiC2_A-X_lab-ident_1956}. It was assumed that the molecule has a linear SiCC structure (silapropadienediylidene) in analogy with the isoelectronic propadienediylidene, C$_3$. Although some later studies cast doubt on the linear structure of the molecule, it took almost another 30 years until the analysis of the rotational structure of the electronic origin band unmistakably determined the structure as silacyclopropynylidene~\\cite{SiC2_A-X_lab-non-lin_1984}. In the course of their analysis, the authors instigated quantum chemical calculations which provided evidence that the cyclic isomer of SiC$_2$ may be lower in energy than the linear form~\\cite{SiC2_ai_1984}.  A plethora of quantum chemical calculations on various properties of SiC$_2$ have been published later, yielding energy differences between the linear and the cyclic form which depended strongly on the level of the calculation and the size of the basis set. A high level ab initio calculation concluded that the cyclic isomer of SiC$_2$ is the only minimum on the potential energy surface and that the linear transition state is 24.3~kJ/mol higher in energy~\\cite{SiC2_ai_1997}. However, the authors attached a caveat to this value as an anharmonic force field calculation provided a much too small value for the vibrational energy of $\\varv_3 = 1$ and much too large anharmonicity constants. Therefore, some of the authors revisited the problem of the energy difference between the two SiC$_2$ structures, the last time in 2003 when very high level calculations combined with very large basis sets, basis set extrapolation to infinite size as well as additional corrections yielded a value of 26.5~kJ/mol~\\cite{SiC2_ai_2003}.   The SiC$_2$ molecular parameters obtained in Ref.~\\cite{SiC2_A-X_lab-non-lin_1984} laid the foundation for progress in laboratory spectroscopy. The $J = 1 - 0$ rotational transition frequencies of the three isotopologs SiC$_2$\\footnote{Unlabeled atoms refer to $^{12}$C and $^{28}$Si.}, $^{29}$SiC$_2$, and $^{30}$SiC$_2$, as well as the permanent electric dipole moment were measured using Fourier transform microwave spectroscopy~\\cite{SiC2_1-0_dip_1989}. Subsequently, 34 additional transition frequencies were measured for the main isotopic species between 93 and 370~GHz~\\cite{SiC2_rot_1989}. Even though a comparatively large number of 15 spectroscopic parameters of a standard Watson-type Hamiltonian in the $A$-reduction, all parameters up to sixth order, were employed in the fit, the transition frequencies were reproduced on average to only four times the experimental uncertainties.  Similarly large bodies of laboratory transition frequencies were obtained for SiC$^{13}$C between 339 and 405~MHz~\\cite{SiCC-13_isos_astro_1991} and, only very recently, for $^{29}$SiC$_2$ and $^{30}$SiC$_2$ between 140 and 360~GHz~\\cite{Si-29_30-C2_rot_2011}.   Higher excited vibrational states have been studied for the main isotopolog. Rotational transitions in its low-lying $\\varv_3 = 1$ vibrational state between 186 and 399~GHz~\\cite{SiC2_001_rot_1991} and in $\\varv_3 = 1$ and 2 were obtained between 140 and 400~GHz~\\cite{SiC2_001_002_rot_1994}. The symmetry of $\\varv_3 = 1$ is $b_2$ and can be viewed as an asymmetric bending state which facilitates internal rotation of the C$_2$ unit with respect to the Si atom. Its vibrational energy has been determined as 196.37~cm$^{-1}$ from an investigation into the laser-induced and the dispersed fluorescence of a jet-cooled sample of SiC$_2$~\\cite{SiC2_LiF_DF_1991}.   Several attempts have been made to model rotational and sometimes also rovibrational data to a varying extent and accuracy. A semirigid bender (SRB) Hamiltonian was employed~\\cite{SiC2_modeling_1994} to reproduce $\\varv_3 = 0$ and 1 rotational transition frequencies~\\cite{SiC2_rot_1989,SiC2_001_rot_1991} as well as rovibrational data from their previous~\\cite{SiC2_LiF_DF_1991} and present work~\\cite{SiC2_modeling_1994}. The analysis suggested that the linear configuration is not a local minimum, and the cyclic form is 22.5~kJ/mol lower than the linear form; the estimated uncertainty was 2.4~kJ/mol. They found the energy difference to be in good agreement with results from their own ab initio calculations; their highest level value being 21.8~kJ/mol. The reproduction of the vibrational data was reasonable, that of the rotational data was only qualitative and thus of no use for radio-astronomical purposes.  The $\\varv_3 = 0$ and 1 state rotational data~\\cite{SiC2_rot_1989,SiC2_001_rot_1991} were also modeled with a dedicated internal rotation Hamiltonian~\\cite{SiC2_modeling_1993}. A greatly improved reproduction, within 1.5~times the experimental uncertainties, was achieved for the ground vibrational state, albeit at the expense of 16 spectroscopic parameters compared to 15 previously~\\cite{SiC2_rot_1989}. The barrier to linearity was derived as $\\sim$54~kJ/mol, more than twice the value from ab initio calculations of that time.  A reasonably successful result has been obtained by employing a conventional Watson-type Hamiltonian in the $S$-reduction~\\cite{IRC_10216_2008}. The reproduction of the ground state rotational data~\\cite{SiC2_1-0_dip_1989,SiC2_rot_1989} was converged after varying 17 spectroscopic and keeping one fixed. The data were reproduced on average as well as in the fit which employed an internal rotation Hamiltonian~\\cite{SiC2_modeling_1993}; the number of varied parameters, however, was larger by still another one. After scaling of the parameters, the transition frequencies of $^{29}$SiC$_2$ and $^{30}$SiC$_2$ available at that time, mostly from astronomical observation~\\cite{SiCC-13_isos_astro_1991,IRC_10216_2008}, could be reproduced well after releasing only 5 spectroscopic parameters~\\cite{IRC_10216_2008}; 13 parameter were released in the fit for SiC$^{13}$C because of the large amount of accurate laboratory data~\\cite{SiCC-13_isos_astro_1991}.   The unambiguous assignment of a spectroscopic feature observed in space relies on reliable predictions which are usually based on laboratory data which, for the most part, have been obtained in approximately the same frequency domain. The predictions may also be based in part on data from astronomical observations, recent examples include H$^{13}$CO$^+$~\\cite{HC-13-O+_2004}, DCO$^+$~\\cite{DCO+_2005,DCO+_etc_2009}, DNC and HN$^{13}$C~\\cite{DCO+_etc_2009}, C$_2$H~\\cite{C2H_2009}, and C$^{13}$CH~\\cite{CC-13-H_2010}. Sometimes, identifications are even possible in the absence of laboratory data, as demonstrated recently by the detection of C$_5$N$^-$ in the CSE of CW~Leo~\\cite{det_C5N-}. Frequencies have to be predicted with accuracies better than around one tenth of the line width to permit extraction of dynamical information. Some type of intensity information, such as that at a certain temperature, the line strength, or the Einstein $A$-value is needed as is additional auxiliary information such as quantum numbers, lower or upper state energies etc. Pickett's {\\scriptsize SPCAT} and {\\scriptsize SPFIT} programs \\cite{Herb} have been developed for that purpose and have evolved over time, see e.\\,g. Ref.~\\cite{editorial_Herb-Ed}. These programs are routinely used in the Cologne Database for Molecular Spectroscopy\\footnote{Internet address: http://www.astro.uni-koeln.de/cdms}, CDMS~\\cite{CDMS_1,CDMS_2}, to provide in its catalog section\\footnote{Internet address: http://www.astro.uni-koeln.de/cdms/catalog} predictions of (mostly) rotational spectra of molecules which may be found in various environments in space.   The spectroscopic parameters from the first rotational analysis of the electronic spectrum of SiC$_2$~\\cite{SiC2_A-X_lab-non-lin_1984} were accurate enough to identify nine unassigned emission features, previously observed between 93 and 171~GHz toward the carbon-rich AGB star CW~Leo, also known as IRC~+10216, as belonging to SiC$_2$ and improved the SiC$_2$ structural parameters~\\cite{SiC2_det-IRC_1984}. Using the structural parameters from that work, the three $J = 4 - 3$, $\\Delta K_a = 0$ transitions of $^{29}$SiC$_2$ and $^{30}$SiC$_2$ were detected toward the same source very soon thereafter~\\cite{Si-29_30-C2_det-IRC_1986}; two of these transitions were marginally detected for SiC$^{13}$C. Extensive sets of transition frequencies of $^{29}$SiC$_2$, $^{30}$SiC$_2$, and SiC$^{13}$C were obtained from radio-astronomical observation between 90 and 241~GHz together with laboratory rest frequencies for the latter isotopolog~\\cite{SiCC-13_isos_astro_1991}.   Radio lines of SiC$_2$ have also been detected in the CSEs of other C-rich AGB stars, e.\\,g. toward II~Lup, which is also known as IRAS 15194-5115~\\cite{SiC2_II_Lup_1993}. However, SiC$_2$ features are particularly strong toward CW~Leo, which, to a large extent, is due to its proximity to our Solar system. CW~Leo has been studied extensively, and many molecular species, such as CN$^-$~\\cite{det_CN-} and FeCN~\\cite{det_FeCN}, have been detected toward its CSE exclusively or for the first time, including many Si-containing ones, such as SiCN and SiNC~\\cite{det_SiNC}. The source is a subject of observations in several key projects carried out with the recently launched {\\it Herschel} satellite~\\cite{Herschel}. One of these projects is a molecular line survey carried out with the Heterodyne Instrument for the Far-Infrared (HIFI)~\\cite{HIFI}. This high-resolution instrument covers, in several bands, the 480$-$1250 and 1410$-$1910~GHz regions. A preliminary analysis of a lower frequency region (554$-$637~GHz) revealed that a rather small number of molecules account for a large fraction of the emission features~\\cite{SiC2_HIFI_2010}. These molecules are CS, SiO, SiS, HCN, and, in particular, SiC$_2$. Predictions of the rotational spectrum based on Ref.~\\cite{IRC_10216_2008} turned out to be very good for transitions with higher values of $K_a$, but showed increasing deviations of up to 10~MHz for transitions with decreasing $K_a$, contrary to common expectations. A combined fit of these transitions frequencies together with those from laboratory spectra~\\cite{SiC2_1-0_dip_1989,SiC2_rot_1989} required only one additional parameter~\\cite{SiC2_HIFI_2010} to reproduce the astronomical data within uncertainties, and the laboratory data about as well as before~\\cite{IRC_10216_2008}. Predictions of the SiC$_2$ rotational spectrum based on these results are currently available as version~2 in the CDMS\\footnote{http://www.astro.uni-koeln.de/cgi-bin/cdmssearch?file=c052527.cat; see http://www.astro.uni-koeln.de/cgi-bin/cdmsinfo?file=e052527.cat for the documentation}. Since they were still not appropriate to predict the observed emission features satisfactorily to high frequencies (beyond 1100~GHz), we present here a combined fit using these as well as laboratory data, supplemented with data from additional astronomical observations.      \\begin{figure*} \\begin{center} \\includegraphics[width=15.0cm]{sic2_fig_vs1.eps} \\end{center} \\caption{Sections of the molecular line survey of CW~Leo displaying selected SiC$_2$ transitions. Their quantum numbers $J'_{K'_a,K'_c} - J''_{K''_a,K''_c}$ are given in red (grey). The formulae and quantum numbers of other lines are given in blue (black).} \\label{spectrum} \\end{figure*} ",
        "conclusions": "\\label{Conclusions}  Astronomical observations with {\\it Herschel} have been used to improve spectroscopic parameters of SiC$_2$ in its ground vibrational state greatly. The transition frequencies obtained from laboratory data as well as from astronomical observations have been reproduced to within the estimated uncertainties employing a conventional Watson-type Hamiltonian. These data should be useful to test alternative models to describe the rotational or even rovibrational energy levels of this floppy molecule. Predictions of the rotational spectrum as well as line, parameter and other auxiliary files, both from present fits as well as previous ones, will be available in the CDMS~\\cite{CDMS_1,CDMS_2}."
    },
    "1112/1112.1046_arXiv.txt": {
        "abstract": "*{New instrumentation is providing new insights into intermediate mass pulsating Cepheids, particularly about their formation and history.  Three approaches are discussed, using space (Hubble and Chandra) and ground-based studies (radial velocities). First, we are conducting a survey of Cepheids with the Hubble Space Telescope Wide Field Camera 3 (WFC3) to identify possible resolved companions (for example Eta Aql)  and thus provide constraints on star formation. Followup  X-ray observations (Chandra and XMM-Newton) can confirm whether possible low mass companions are young enough to be physical companions of Cepheids. In a related study of intermediate mass stars, Chandra X-ray observations of late B stars in  Tr 16 have been used to determine the fraction which have X-ray active low mass companions. Finally, the Tennessee State Automatic Spectroscopic Telescope AST and the Moscow University group have obtained velocities of a number of Cepheids.  As an example, the orbit of V350 Sgr has been redetermined, providing a new level of accuracy to the orbital velocity amplitude, which is needed for mass determination.   }      \\abstract{New instrumentation is providing new insights into intermediate mass pulsating Cepheids, particularly about their formation and history.  Three approaches are discussed, using space (Hubble and Chandra) and ground-based studies (radial velocities). First, we are conducting a survey of Cepheids with the Hubble Space Telescope Wide Field Camera 3 (WFC3) to identify possible resolved companions (for example Eta Aql)  and thus provide constraints on star formation. Followup  X-ray observations (Chandra and XMM-Newton) can confirm whether possible low mass companions are young enough to be physical companions of Cepheids. In a related study of intermediate mass stars, Chandra X-ray observations of late B stars in  Tr 16 have been used to determine the fraction which have X-ray active low mass companions. Finally, the Tennessee State Automatic Spectroscopic Telescope AST and the Moscow University group have obtained velocities of a number of Cepheids.  As an example, the orbit of V350 Sgr has been redetermined, providing a new level of accuracy to the orbital velocity amplitude, which is needed for mass determination. }  \\vskip .1truein ",
        "introduction": "This contribution focuses on two aspects of binary Cepheids: information they provide about star formation and about stellar evolution (masses).  \\begin{figure} \\centering \\includegraphics[scale=.20]{evans_fig1a.eps} \\hspace{0.5cm} \\includegraphics[scale=.25]{evans_fig1b.eps}     \\caption{Left: The center of the HST WFC3 V image of $\\eta$ Aql. The image has a log scale and is approximately 10$''$ wide. Right: The difference image with the Cepheid T Mon subtracted from the $\\eta$ Aql image.  The companion is clearly visible 0.7$''$ from the Cepheid.  Only the vertical column bleeding remains uncorrected.   } \\label{fig:1}       % \\end{figure} ",
        "conclusions": ""
    },
    "1112/1112.3569_arXiv.txt": {
        "abstract": "Simple models of magnetic field generation by convection in rotating spherical shells exhibit properties resembling those observed on the sun. The {assumption of the Boussinesq approximation made in these models} prevents a  realistic description of the solar cycle, but through a physically motivated change in the boundary condition for the differential rotation the propagation of dynamo waves towards higher latitudes can be reversed at least at low latitudes. ",
        "introduction": "A well known difficulty in modeling the solar dynamo is the fact that the dynamo waves which describe the nearly time periodic dynamics of the magnetic field tend to propagate from lower to higher latitudes instead in the opposite sense as observed on the sun. This effect is well known in mean field models of the solar cycle, see for example Stix (1976), but it is also observed in numerical solutions {of convection-driven} dynamos in rotating spherical fluid shells which are supposed to model processes in the solar convection zone. This shortcoming and others are caused by the inadequate representation of the compressibility of the solar atmosphere. Even with the huge power of modern supercomputers it is not yet possible to resolve adequately the dynamics of convection in the presence of the large density variation between bottom and top of the solar convection zone and to model appropriately the {dependence of density on pressure.}  {In their early} direct numerical models of solar convection and magnetic field generation, Gilman and Miller (1981) assumed the Boussinesq approximation in which the fluid is regarded as incompressible except in connection with the gravity term where the temperature dependence of the density is taken into account. This approximation eliminates the need for a separate equation of state and leads to a system of equations describing long period processes while the short period acoustic modes no longer enter the analysis. The same effect is obtained in the anelastic approximation in which the horizontally averaged density variation is taken into account, but the fluctuating component of the density is still only represented in the gravity term. For applications of the anelastic approximation in models of the solar convection zone see Gilman and Glatzmaier (1981) and Elliot et al. (2000).  Convection in rotating spherical fluid shells heated from below is always associated with a differential rotation generated by the Reynolds stresses of convection. An analytical model demonstrating the preference of banana shaped convection cells girdling the equator and the associated solar like differential rotation was presented by Busse (1970, 1973). While the analytical solution for stress-free boundaries exhibits a depth independent differential rotation, a differential rotation decreasing with depth is always found in fully nonlinear numerical models. This property together with the fact that the differential rotation reaches  its maximum at the equator is responsible for the propagation of the dynamo waves towards higher latitudes (Yoshimura, 1975). \\begin{figure*} \\psfrag{Omg}{$\\Omega$} \\begin{center} \\epsfig{file=Fig01.eps,width=1.5\\columnwidth,clip=} \\end{center} \\caption[]{{(color online)} Differential rotation of the solar convection zone according to helioseismology (right, image courtesy NSF's National Solar Observatory) and differential rotation of the numerical model in the case $\\eta=0.65$, $P=1$, $\\tau=2$, $R=120000$, $P_m=4$, $\\beta=1.5$ and mixed velocity boundary conditions {(left)}. } \\label{fig010} \\end{figure*}  The solar differential rotation also decreases with depth throughout most of the convection zone with the exception of the tachocline region near its bottom and a region near its surface as indicated in figure \\ref{fig010}(b). There is no general agreement about the origin of {the} upper 30Mm deep layer in which the differential rotation increases with depth. Here we assume that it is caused by supergranular convection that is characterized by a strong asymmetry between rising hot plasma and descending cool plasma. This type of convection has been modeled by hexagonal convection cells in the presence of rotation (Busse, {2007}). As shown in this paper the asymmetry between rising and descending flow in hexagonal convection cells does indeed generate a differential rotation that increases with depth. We use this dynamical property of convection as motivation to modify the usually assumed stress-free boundary condition. Solely for the differential rotation we apply the condition given in expression (A11) of the Appendix. The resulting profiles shown in the example of figure \\ref{fig010}(a) are still not very solar like, but they show an increase with depth of the differential rotation in the upper layer and  tend to generate a more solar like behavior in the numerical simulations, {as discussed below.} ",
        "conclusions": "In this report an attempt is described to explore the extent to which the operation of the solar dynamo can be understood on the basis of a minimal, but physically consistent, {convection-driven} dynamo model. Although the Boussinesq approximation is a highly unrealistic assumption in the case of the Sun, the convection columns are similar to those found in numerical simulations based on {anelastic} models, see, for instance, Brun et al.~(2004) and Ghizaru et al.~(2010). Although there is little solar evidence for this type {of} convection, there are generally believed to exist as \"giant cells\" in the deeper region of the solar convection zone.  The structure of the magnetic field found in our simulations differs substantially from the commonly assumed structure of the solar magnetic field dominated by a strong axisymmetric toroidal component. When the evidence for this traditional view is examined, however, it is found that observations do not contradict fields dominated by $m=1$ or $m=2$ components such as those shown in figures \\ref{fig020} and \\ref{fig050}. On the contrary, the presence of active longitudes on the sun indicates at least that components with the azimuthal wavenumbers $m=1$ and $m=2$ play a significant role.  {A more detailed exploration of the parameter space of our minimal model is likely to reveal an even better correspondence with solar observations.}  \\appendix"
    },
    "1112/1112.4349_arXiv.txt": {
        "abstract": "{We present the analysis of the bright X-ray binary 4U~1820-30, based mainly on \\xmm-RGS data, but using complementary data from \\xmm-\\epi, \\inte, and \\ch-HETG, to investigate different aspects of the source. The broad band continuum is well fitted by a classical combination of black body and Comptonized emission. The continuum shape and the high flux of the source ($L/L_{\\rm Edd}\\sim0.16$) are consistent with a \"high state\" of the source. We do not find significant evidence of iron emission at energies $\\geq6.4$\\,keV. The soft X-ray spectrum contain a number of absorption features. Here we focus on the cold-mildly ionized gas. The neutral gas column density is $N_{\\rm H}\\sim 1.63\\times10^{21}$\\,cm$^{-2}$. The detailed study of the oxygen and iron edge reveals that those elements are depleted, defined here as the ratio between dust and the total ISM cold phase, by a factor $0.20\\pm0.02$ and $0.87\\pm0.14$, respectively. Using the available dust models, the best fit points to a major contribution of Mg-rich silicates, with metallic iron inclusion. Although we find that a large fraction of Fe is in dust form, the fit shows that Fe-rich silicates are disfavored. The measured Mg:Fe ratio is $2.0\\pm0.3$. Interestingly, this modeling may point to a well studied dust constituent (GEMS), sometimes proposed as a silicate constituent in our Galaxy. Oxygen and iron are found to be slightly over- and under-abundant, respectively (1.23 and 0.85 times the solar value) along this line of sight. We also report the detection of two absorption lines, tentatively identified as part of an outflow of mildly ionized gas ($\\xi\\sim-0.5$) at a velocity of $\\sim1200$\\,\\kms.} {} ",
        "introduction": "\\label{par:intro}  The interstellar medium (ISM) in the plane of our Galaxy is a dynamic and complex environment, composed of mainly neutral matter in both gas and dust form and by a warmer gas phase in the form of diffuse emission in, and above, the Galactic plane. The properties of the cold phase in the diffuse ISM have been extensively studied at long wavelengths, from the far-infrared to the far-UV \\citep[e.g.][ for a review]{draine03}. A sizeable fraction of the cold phase is locked up in dust grains \\citep[e.g.][ and references therein]{ss96,jenkins09}. Amorphous silicate materials together with graphite and policyclic aromatic carbon should account for the majority of the depleted elements measured in the ISM: C, O, Fe, Mg, Si \\citep[e.g.][]{wd01,wooden08}. One of the major spectroscopic signatures of the presence of Fe- and Mg- rich amorphous silicates is the 10\\,$\\mu$m emission feature. Extensive studies of this feature lead to the conclusion that the Fe:Mg ratio should be approximatively 1 \\citep[e.g.][]{lidraine01}. Main sources of both Mg and Fe silicates are O-rich asymptotic giant branch stars (AGB) and supernovae. However, the process of amorphization of the dust agglomerates, for instance by rapid cooling of the gas phase \\citep{wooden05} or cosmic rays bombardment \\citep{carrez02}, strongly favors the survival of Mg silicates. In fact, a recent analysis successfully models the 10\\,$\\mu$m feature in terms of Mg-rich silicates, if non-spherical grain shapes are used \\citep{min07}. Glassy material, consisting of Mg-rich silicates with metallic iron and sulfide inclusion \\citep[called GEMS,][]{bradley94} have also been commonly found during the {\\em Stardust} mission. Their origin is mostly from the interplanetary environment \\citep[e.g.][]{kel_mes04}, but a fraction have a composition compatible with an ISM origin \\citep{kel_mes08}. Iron is however a highly depleted element \\citep[70--99\\% of Fe is in dust,][]{wilms00,whi03} whose inclusion into solid grains is not completely understood \\citep[e.g.][]{whi03}. This is mainly due to the difficulty of modeling iron emission, which does not display any sharp feature in long-wavelength spectra.  The abundances of the most important metals in the Galactic disk smoothly decrease with the galactocentric distance. The average slope of the distribution is $\\sim0.06$\\,dex\\,kpc$^{-1}$ \\citep[][ and references therein]{chen03}. However, a large scatter in the abundance measurements as a function of the Galactic radius is reported. This is attributed to different factors which contaminate the smooth mixing due to the pure stellar evolution process. Indeed the medium can be locally influenced by e.g. supernovae ejecta, and in-falling metal-poor material into the disk \\citep{lugaro99,nittler05}.  In recent years it has become clear that the X-ray band could provide an excellent laboratory to study the silicate content of the diffuse ISM, as the absorption K edges of O ($E=0.538$\\,keV), Mg ($E=1.30$\\,keV) and Si ($E=1.84$\\,keV) together with the Fe LII and LIII edges ($0.71$ and $0.72$\\,keV, respectively) fall in the low-energy X-ray band. The method used is to study the absorbed spectrum of bright X-ray binaries, located in different regions in the disk, observed with high-energy resolution instruments. This allows to probe the interstellar dust (ID) content in a variety of environments, with different extinction and with possibly different dust formation history. Previous studies of X-ray spectra taken along different lines of sight led first to the recognition that not only gas but also dust plays a role in shaping the iron and oxygen edges \\citep{takei02,kaastra09} and later led to the quantitative modeling of those edges \\citep{lee09,devc09,ciro}. There is not yet a clear picture of the chemical composition of the ID as seen in X-rays. Silicates containing andratite (iron-rich silicates) were reported studying the oxygen edge of \\object{GS~1826-238} \\citep{ciro}, while iron oxides, rather than iron silicates were reported along the line of sight of \\object{Cyg~X-1} \\citep{lee09}. This may point to a chemically inhomogeneous distribution of ID.\\\\  \\object{4U~1820-30} is an extensively studied source, by virtue of its extraordinary intrinsic properties. It is an ultracompact \\citep[orbital period 11.4\\,minutes,][]{1987ApJ...312L..17S} X-ray binary consisting of a neutron star and a He white-dwarf \\citep{1987ApJ...322..842R}. The presence of X-ray bursts associated with the neutron star has been early recognized \\citep[first by][]{grindlay76}. \\uu\\ is classified as an atoll source \\citep{1989A&A...225...79H}, showing kilohertz quasi-periodic oscillation at different frequencies in its power spectrum \\citep[e.g.][]{smale97,zhang98}.\\\\ The broad band spectrum has been studied with several instruments. The source displays a Comptonized continuum and a soft black body component \\citep[e.g.][]{sidoli01}.  With the advent of high resolution spectroscopy, the low-energy X-ray spectrum also revealed interesting features. Absorption by ionized ions has been reported in several studies \\citep{futamoto04,yw05,juett06,cackett08}. The absence of blueshift in the lines of this gas and lack of variability generally points to an interstellar origin \\citep[but see][]{cackett08}. The hypothesis of a gas intrinsic to the source is intriguing as other X-ray binaries often display absorption by ionized gas either at the rest frame of the source \\citep{vanpeet} or outflowing \\citep[e.g.][]{jmiller06,neilsenlee09}. The source has also been used as a backlight to illuminate the interstellar dust along the line of sight. This results in a dust scattering halo, which in this source is moderate \\citep[$\\sim3.2$\\% at 1\\,keV,][]{peter95}, given the relatively low Galactic column density ($N_{\\rm H}\\sim1.63\\times10^{21}$\\,cm$^{-2}$, this study).  Absorption by cold interstellar dust has been never studied in detail in this source. In this paper we aim for a comprehensive view of the cold absorbing medium along the line of sight to this source. This study benefits from the combined information provided by the \\ch\\ and \\xmm\\ high resolution instruments, which allows us to meaningfully study both the Fe\\,L and O\\,K edges. In addition, we make use of dust models updated with the latest laboratory measurements \\citep[e.g.][]{lee08,lee09}. We also present the broad band continuum behavior underlying the absorption features.\\\\    The adopted protosolar abundances follow the prescription given by \\citet{lodders09} and discussed in \\citet{lodders10}. The broad band energy spectrum (\\epi\\ and \\inte) is fitted using the $\\chi^2$ minimization method, taking care that at least 20 counts per bin are present in each data set. For the high resolution spectra (RGS and \\ch-HETG) the Cash statistic has been used\\footnote{http://heasarc.gsfc.nasa.gov/docs/xanadu/xspec/manual/XSappendixCash.html}. The errors quoted are for 68\\% confidence level, corresponding to $\\Delta\\chi^2=1$ or $\\Delta C^2=1$. The spectral fitting package used in this paper is SPEX\\footnote{www.sron.nl/spex} \\citep{kaastra96}. The adopted distance is 7.6\\,kpc \\citep{kuulkers03}. The nominal hydrogen column density toward the source is $1.29\\times10^{21}$\\,cm$^{-2}$ \\citep{kal05}, to be compared with $1.52\\times10^{21}$\\,cm$^{-2}$ \\citep{dl90}.  This paper is organized as follows: in Sect.~\\ref{par:data}, we illustrate the data handling for the different instruments used. Sect.~\\ref{par:cont} is devoted to the continuum determination, using \\epi\\ and \\inte\\ data. In Sect.~\\ref{par:high} we describe in detail the modeling of the different absorption components using RGS and \\ch-MEG. The discussion can be found in Sect.~\\ref{par:discussion} and the conclusions in Sect.~\\ref{par:conclusions}.\\\\ ",
        "conclusions": "\\label{par:conclusions} In this paper we present the X-ray analysis of the continuum of \\uu\\ (using a quasi-simultaneous observation of \\xmm\\ and \\inte) and of the absorption features due to the cold matter in the line of sight (using \\xmm-RGS and \\ch-MEG data).\\\\  The continuum shape and the Eddington ratio show that the source was caught in a high-state. The continuum is well fitted by black body emission plus a Comptonization component which extends up to 40\\,keV. We do not find evidence of iron emission, either from neutral or ionized matter. This may be naturally explained by the metal-poor accretion stream expected from the white-dwarf companion.\\\\  The absorption spectrum shows the presence of many components with different ionization. We focused on the cold and mildly ionized phase only. Oxygen has been found slightly overabundant by a factor 1.23 times the solar value. Iron is on the contrary slightly underabundant ($\\sim$0.85 times solar). The abundance values are not dramatically deviating from the solar ones and do not allow us to assign a precise location of the absorbing gas. However, a location close to the observer seems likely.  Thanks to the simultaneous study of absorption by dust and gas we measured also the element depletion. Oxygen is mildly depleted by a factor about 0.20. The depletion of iron is more evident, as the depletion factor is 0.87. The depletion of Mg and Si are more difficult to determine. We find that they are depleted of a factor $>0.97$ and $>0.86$, respectively\\\\  We modeled the dust contribution with the currently availables absorption profiles of dust compounds. Our conclusions bear the uncertainty due to the still limited dust data-base and a lower sensitivity in selected spectral regions. However we clearly find that both the oxygen and iron edges cannot be fitted by iron-rich silicates. On the contrary, the oxygen edge is consistent to be mostly absorbed by enstatite (MgSiO$_3$, possibly in a glass-form). Metallic iron should be the main absorber in the iron edge. This leads to the interesting possibility that a GEMS-like form (Mg-rich silicates with metallic iron inclusion) of grain may be absorbing along this line of sight. A fraction of the studied GEMS, in particular the sulfur-poor grains, are believed to be of ISM origin and have also been proposed as constituents of ISM. For the first time an X-ray absorption analysis provides a tentative confirmation of this scenario.\\\\   Finally, we report the tentative detection of a mildly ionized outflow ($v_{\\rm out}\\sim1200$\\,\\kms), highlighted by the \\oiv\\ and \\ov\\ absorption lines. Both a photo- or collisional- ionizing process could fit the lines, leaving open the interpretation on the nature of this gas."
    },
    "1112/1112.4663_arXiv.txt": {
        "abstract": "The super-massive objects in galactic nuclei are thought to be the Kerr black holes predicted by General Relativity, although a definite proof of their actual nature is still lacking. The most massive objects in AGN ($M \\sim 10^9$~$M_\\odot$) seem to have a high radiative efficiency ($\\eta \\sim 0.4$) and a moderate mass accretion rate ($L_{\\rm bol}/L_{\\rm Edd} \\sim 0.3$). The high radiative efficiency could suggest they are very rapidly-rotating black holes. The moderate luminosity could indicate that their accretion disk is geometrically thin. If so, these objects could be excellent candidates to test the Kerr black hole hypothesis. An accurate measurement of the radiative efficiency of an individual AGN may probe the geometry of the space-time around the black hole candidate with a precision comparable to the one achievable with future space-based gravitational-wave detectors like LISA. A robust evidence of the existence of a black hole candidate with $\\eta > 0.32$ and accreting from a thin disk may be interpreted as an indication of new physics. For the time being, there are several issues to address before using AGN to test the Kerr paradigm, but the approach seems to be promising and capable of providing interesting results before the advent of gravitational wave astronomy. ",
        "introduction": "Gravity has been tested and verified for distances in the range $\\sim$1~mm to $\\sim$1~pc (mainly within its Newtonian limit) and for weak gravitational fields~\\cite{w,5th}. The research is now moving to check the validity of the theory at cosmological scales, sub-millimeter distances, and for strong gravitational fields. One of the most intriguing predictions of General Relativity (GR) is that the collapsing matter produces singularities in the space-time. According to the weak cosmic censorship conjecture, singularities of gravitational collapse must be hidden within black holes (BHs)~\\cite{p}. In 4-dimensional GR, uncharged BHs are described by the Kerr solution, which is completely specified by two parameters, the mass, $M$, and the spin angular momentum, $J$~\\citep{c}. The condition for the existence of the event horizon is $a_* \\le 1$, where $a_* = |J/M^2|$ is the spin parameter\\footnote{Throughout the paper I use units in which $G_N = c = 1$, unless stated otherwise.}. When $a_* > 1$, there is no horizon and the central singularity is naked, violating the weak cosmic censorship conjecture.   Astronomers have discovered at least two classes of BH candidates (for a review, see e.g. Ref.~\\cite{n}): stellar-mass objects in X-ray binary systems ($M \\sim 5 - 20$~$M_\\odot$) and super-massive objects in galactic nuclei ($M \\sim 10^5 - 10^9$~$M_\\odot$). The estimates of the masses of these objects are robust, because determined via dynamical measurements and without any assumption about the geometry of the space-time. The key-point is that the stellar-mass objects in X-ray binary systems are too heavy to be neutron or quark stars for any reasonable matter equation of state~\\cite{kb}, while the super-massive objects at the centers of galaxies are too heavy, compact, and old to be clusters of non-luminous bodies, as the cluster lifetime would be shorter than the age of these systems~\\cite{m}. All these objects are therefore thought to be the BHs predicted by GR, as they cannot be explained otherwise without introducing new physics. There are also some observations interpreted as an indirect evidence for the existence of the event horizon~\\cite{bln} (but see~\\cite{a}). On the contrary, there is no indication that the geometry around these objects is described by the Kerr metric.   Testing the Kerr BH hypothesis is thus the next step to progress in this research field and several authors have indeed suggested possible ways to do it using present and future data (for a review, see e.g. Ref.~\\cite{b4}). A very promising approach is the detection of extreme mass ratio inspirals (EMRIs, i.e. systems consisting of a stellar-mass compact object orbiting a super-massive BH candidate) with future space-based gravitational-wave antennas. Missions like LISA will be able to follow the stellar-mass compact object for millions of orbits around the central super-massive BH candidate, and therefore deviations from the Kerr geometry will lead to a phase difference in the gravitational waveforms that grows with the number of observed cycles~\\cite{r}. However, these data will not be available shortly, as the first mission will be at best in the early 2020s. The nature of BH candidates can also be tested by extending the methods currently used to estimate the spin of these objects, such as X-ray continuum~\\cite{bb1} and K$\\alpha$-iron measurements~\\cite{pj}, observations of quasi-periodic oscillations~\\cite{jpp}, and measurements of the cosmic X-ray background~\\cite{b2,b5}. These methods can in principle be applied even with present data, provided that the systematic errors are properly understood. Future observations of the shadow of nearby super-massive BH candidates are another exciting possibility to test the Kerr BH paradigm~\\cite{bf}.   Previous studies have clearly pointed out that ``rapidly-rotating'' objects are the best candidates to test the Kerr BH hypothesis: if the object rotates fast, even a small deviation from the Kerr background can cause significant differences in the properties of the electromagnetic radiation emitted by the gas of the accretion disk and peculiar features, otherwise absent in the Kerr geometry, may show up~\\cite{bb2,bm}. The aim of this paper is to investigate the potentialities of the most massive BH candidates in AGN, through the measurement of their radiative efficiency $\\eta$, defined by $L_{\\rm bol} = \\eta \\dot{M}$, where $L_{\\rm bol}$ is the bolometric luminosity of the source and $\\dot{M}$ is the mass accretion rate of the BH candidate. The estimate of the mean radiative efficiency of AGN through the Soltan's argument~\\cite{s} already suggests the presence of rapidly-rotating BHs~\\cite{erz,ho}. Recently, Davis and Laor have proposed a way to measure the radiative efficiency of individual AGN~\\cite{dl}. The authors found that the most massive BH candidates (with a mass $M \\sim 10^9$~$M_\\odot$) would have a high radiative efficiency, up to $\\eta \\sim 0.4$, and a moderate mass accretion rate, $L_{\\rm bol}/L_{\\rm Edd} \\sim 0.3$, where $L_{\\rm Edd}$ is the Eddington luminosity of the source. The standard accretion disk model in a Kerr background would predict a high value of the spin parameter $a_*$ for these objects, extremely close to 1. At the same time, the moderate luminosity (in Eddington units) may indicate a thin accretion disk and the applicability of the standard accretion disk model. If these estimates and these considerations are correct, the most massive compact objects in AGN would be excellent candidates to test the Kerr paradigm. The sole measurement of $\\eta$ can potentially constrain either $a_*$ and a deviation from the Kerr geometry.   The paper is organized as follows. In Section~\\ref{s-2}, I review the standard accretion disk model, its assumptions and properties, its effects on the evolution of the spin parameter of the central object, and its applicability. In Section~\\ref{s-k}, I consider an accretion disk in the Kerr background, summarizing well-known results that should be expected if the BH candidates are the BHs of GR. In Sections~\\ref{s-jp} and \\ref{s-mn}, I discuss accretion disks respectively in the Johannsen-Psaltis (JP)~\\cite{jp} and in the Manko-Novikov (MN)~\\cite{mn} space-times. These are two metrics that can be conveniently used to describe a background deviating from the Kerr geometry. The nature of the two metrics is definitively different: the JP metric describes non-Kerr BHs in a putative alternative theory of gravity, while the MN one is an exact solution of the Einstein's vacuum equations and can describe the exterior gravitational field of generic compact objects. In both cases, the body is characterized by a mass, a spin angular momentum, and an infinite number of  ``deformation parameters'', even if here, for the sake of simplicity, I will consider only a single deformation parameter at a time. I will show that the two metrics present common features. In particular, a high radiative efficiency necessarily requires a very rigid compact object, much stiffer than a self-gravitating fluid with ``normal'' equations of state. The confirmation of the existence of individual AGN with high radiative efficiency ($\\eta > 0.3$) can potentially either be used to put strong constraints on the Kerr nature of astrophysical BH candidates and to discover new physics, as in the framework of the standard accretion disk model $\\eta$ cannot exceed 0.32. In Section~\\ref{s-d}, I discuss the findings of this paper in relation with current estimates of the radiative efficiency of AGN. The conclusions and the issues that need to be addressed before using the most massive objects in AGN to really test the Kerr BH hypothesis are reported in Section~\\ref{s-c}. In Appendices~\\ref{a-1} and \\ref{a-2}, the reader can find the non-vanishing metric coefficients respectively of the JP and of the MN metric. ",
        "conclusions": "\\label{s-c}  There is some evidence that the most massive BH candidates in AGN have a high radiative efficiency and a thin accretion disk. In this case, they could be excellent candidates to test GR in the strong field regime and, in particular, the Kerr BH paradigm. A robust observation of a high radiative efficiency together with the confirmation of the validity of the standard accretion disk model for the accretion process onto these objects would constrain possible deviations from the Kerr geometry of the space-time around astrophysical BH candidates. If we restrict our attention only to objects more oblate than a Kerr BH with the same spin parameter, a radiative efficiency $\\eta > 0.30$ requires that the quadrupole moment of the compact object deviates not more than $\\sim 10^{-3}$ with respect to the one of a Kerr BH. If we allow for the existence of compact objects more prolate than a Kerr BH (whose existence, however, may be questionable, see the footnote~3), the constraints is much weaker, at the level of 20\\%. On the other hand, the observation of an accreting BH candidate with $\\eta > 0.32$ might indicate that the super-massive BH candidates in galactic nuclei are not the Kerr BH of GR, as there is no realistic astrophysical mechanism capable to spin up them to $a_* > 0.998$.   For the time being, there are at least two main issues to address before using the measurement of the radiative efficiency of individual AGN to test GR: \\begin{enumerate} \\item We do not really know if the NT model can describe thin accretion disks around astrophysical BH candidates. If it can or it is easy to estimate the necessary corrections, the measurement of $\\eta$ could provide an interesting way to probe the geometry of the space-time around BH candidates. If it cannot, it seems to be impossible to test GR (with this or other approaches) from the properties of the radiation emitted by the gas in the accretion disk and we should wait for the advent of gravitational wave astronomy. \\item We need robust and more precise measurements of $\\eta$. At present, there are several sources with $\\eta > 0.32$ at high masses, but we cannot really exclude they actually have $\\eta < 0.32$~\\cite{dl}. There are some sources of uncertainty in the final estimate of the radiative efficiency and the most important one is the determination of the mass of the BH candidate. Moreover, the results presented in~\\cite{dl} are based on a preliminary study and it is necessary to further investigate and verify the validity of the method proposed by these authors. \\end{enumerate}"
    },
    "1112/1112.4488_arXiv.txt": {
        "abstract": "{ We propose a general scenario for moduli stabilization where low-energy supersymmetry can be accommodated with a high scale of inflation. The key ingredient is that the stabilization of the modulus field \\emph{during} and \\emph{after} inflation is not associated with a single, common scale, but relies on two different mechanisms. We illustrate this general scenario in a simple example, where during inflation the modulus is stabilized with a large mass by a K\\\"ahler potential coupling to the field which provides the inflationary vacuum energy via its F-term. After inflation, the modulus is stabilized, for instance, by a KKLT superpotential. } ",
        "introduction": "\\label{Sec:Introduction}  Low-energy supersymmetry (SUSY) at the TeV-scale provides an attractive extension of the Standard Model (SM) for various reasons, in particular, since it can solve the gauge hierarchy problem. Furthermore, it is going to be tested in the ongoing LHC experiments. Regarding cosmology, the paradigm of cosmic inflation has emerged as the prime candidate for early universe physics to resolve the flatness and horizon problems associated with the hot big bang scenario. Typically, models of inflation require a comparatively high scale of inflation, close to the energies where the gauge interactions of the SM can be unified. The scale of inflation might soon be tested by an observation of tensor modes with the PLANCK satellite. It is therefore an interesting question whether a high scale of inflation can be realized together with TeV-scale SUSY.  A priori, the two topics of low-energy SUSY and high-scale inflation seem to be unrelated. However, they turn out to be connected due to the issue of ``moduli stabilization'', in particular in the context of string theory. String compactifications generically lead to many light scalar fields, aka moduli, which for example parametrize the size and shape of the compactified internal manifolds. Tracing the physics of extradimensions, the couplings in the 4D effective low-energy theory become functions of the moduli fields. To make sure that the internal dimensions do not decompactify, and also to avoid observational constraints on the space-time variability of the low-energy coupling constants, the moduli fields must be stabilized (i.e. must acquire a suitable mass) both during and after inflation.  In the context of type IIB string theory, the issue of moduli stabilization is most well understood in the well-known KKLT scenario \\cite{Kachru:2003aw}. It assumes that the dilaton and the complex structure moduli are stabilized via fluxes \\cite{Giddings:2001yu,Dasgupta:1999ss,Taylor:1999ii}.\\footnote{For reviews and an extensive list of references on moduli stabilization with fluxes and nonperturbative effects in type IIB string theory cf. e.g. \\cite{Grana:2005jc,Douglas:2006es,Blumenhagen:2006ci,Denef:2007pq}.} Therefore, only the dynamics of the volume moduli are important for the low-energy physics. The volume moduli are stabilized by the contributions from nonperturbative effects such as gaugino condensates. Moreover, moduli stabilization also nicely combines with dynamical supersymmetry breaking with the moduli comprising the hidden sector.  When the issue of moduli stabilization is discussed in conjunction with inflation, however, a severe problem emerges, namely adding the inflationary sector may destabilize the moduli, which was pointed out by Buchm\\\"uller, Hamaguchi, Lebedev and Ratz and by Kallosh and Linde in \\cite{Buchmuller:2004xr,Buchmuller:2004tz,Kallosh:2004yh}. Here, we are concerned with a particular version of this problem sometimes referred to as the Kallosh-Linde (KL) problem \\cite{Kallosh:2004yh}: One often finds an upper bound on the inflationary energy scale in terms of the present-day gravitino mass, \\begin{equation} \\Hinf \\leq m_{3/2}^{\\text{today}} \\, , \\end{equation} to avoid the destabilization of the volume modulus. For TeV-scale SUSY breaking, one has $m_{3/2} \\sim \\text{TeV}$ and thus the scale of inflation is bound to be very small, much below the scale required for many model building approaches (and also very much below observational sensitivities for gravitational waves). Basically, the problem appears since there is effectively only one scale in the problem, which sets both the gravitino mass today and the height of the barrier towards decompactification. A SUSY breaking uplifting term turns the AdS minimum into a nearly Minkowski minimum, and also sets the height of the barrier separating the metastable vacuum from the vacuum at infinity. Typically, inflation in such a setup can be viewed as an additional uplifting which induces a runaway potential for the modulus. Thus, if its contribution becomes too large, the barrier and hence the minimum disappear. In summary, the inflationary scale is constrained to be smaller than the height of the barrier, and therefore the gravitino mass. We will review the problem in more detail in the next section.  As a solution to the problem, KL \\cite{Kallosh:2004yh} suggested a form of the superpotential for which the gravitino mass in the present vacuum is unrelated to the height of the barrier. Therefore, choosing the gravitino mass at the TeV-scale, we can always independently increase the barrier height such that high-scale inflation models do not destabilize the modulus. This corresponds to a fine-tuning of the terms in the superpotential. Following the approach of KL, in the context of volume modulus inflation in a racetrack setup, the problem has been addressed when one of the exponents of the nonperturbative terms is positive \\cite{Abe:2005rx,Badziak:2008yg,Abe:2008xu,Badziak:2008gv,Badziak:2009eh}. However, this can be done only at the expense of introducing more parameters in the theory. In the large volume scenario, attempts have been made to accommodate a small gravitino mass with a high scale of inflation, when inflation happens exponentially far away from the present Minkowski vacua. Other than some inevitable fine tuning, the working models have several phenomenological difficulties \\cite{Conlon:2008cj}. Dynamical avenues  for the case of chaotic inflation \\cite{He:2010uk} and hybrid inflation \\cite{Kobayashi:2010rx} have also been explored, where the gravitino mass becomes inflaton-dependent in a suitable way. Difficulties related to the realizations of high-scale inflation together with low-energy SUSY breaking have been discussed in \\cite{Davis:2008sa,Davis:2008fv} and some resolutions have been proposed in \\cite{Mooij:2010cs,Davis:2008sa}. However, this has been successfully achieved only for the superpotential of \\cite{Kallosh:2004yh}. Combining chaotic inflation and supersymmetry breaking within the KL scheme has recently been discussed in \\cite{Kallosh:2011qk} for the general chaotic inflation models of \\cite{Kallosh:2010ug,Kallosh:2010xz,Demozzi:2010aj}. For other approaches to the KL problem see e.g. \\cite{Misra:2009ei}.  Here, we propose a different resolution of the problem, namely to use two different mechanisms to stabilize the moduli during and after inflation. During inflation, the modulus field receives a large mass proportional to the inflationary vacuum energy. This is achieved, for example, by a suitable moduli-dependence of the K\\\"ahler metric of the field driving supersymmetry breaking during inflation. The role of such a type of coupling for moduli stabilization was first noted in \\cite{Antusch:2008pn}. At the end of inflation, the vacuum energy goes away and we invoke a stabilization mechanism like KKLT relying as usual on nonperturbative terms in the superpotential. We illustrate our general idea in a simple example: chaotic inflation protected by a shift symmetry combined with a KKLT-type superpotential. However, we stress that our framework can be applied to more general setups.  We begin with a brief review of the Kallosh-Linde Problem in section~\\ref{Sec:ReviewKalloshLindeProblem}. Afterwards, we describe our general framework for resolving the problem in section~\\ref{Sec:ResolutionGeneralFramework} and illustrate it in a simple example with shift symmetric chaotic inflation and a KKLT superpotential in section~\\ref{Sec:AnExplicitExample}. Some results for a generic choice of $W_{\\mod}(T)$ are presented in appendix \\ref{Sec:AppendixGeneralWmod}. Finally, we give our conclusions in section~\\ref{Sec:Conclusion}. ",
        "conclusions": "\\label{Sec:Conclusion}  In this article, we have proposed a general scenario for moduli stabilization where low-energy supersymmetry breaking can be accommodated together with a high scale of inflation. In our proposal the KL problem (which is reviewed in section \\ref{Sec:ReviewKalloshLindeProblem}) is resolved because the stabilization of the modulus field \\emph{during} and \\emph{after} inflation is not associated with a single, common scale, but instead relies on two different mechanism to stabilize the modulus during and after inflation.  More explicitly (c.f.\\ section \\ref{Sec:ResolutionGeneralFramework}), we suggest to consider a K\\\"ahler potential which features a coupling between the modulus field and the field whose F-term drives inflation in such a way that the term $V_{\\inf} \\sim e^{K} \\, K^{X \\bar{X}} \\lvert D_{X} W \\rvert^{2}$ creates a minimum for the modulus which stablizes it with large mass during inflation. After inflation, when $D_{X} W$ vanishes, a ``standard'' mechanism involving nonperturbative terms in the superpotential can take over to stabilize the modulus as usual. The way we avoid the KL problem in this setup works essentially as follows: The gravitino mass $m_{3/2}$ now only sets the scale for moduli stabilization after inflation and therefore remains small all the time. During inflation, the scale for moduli stabilization is set by the inflationary energy scale $\\Hinf$ itself and no longer also by $m_{3/2}^{\\text{today}}$. This allows to consistently combine low scale SUSY and high scale inflation.  There is of course a price to pay: Since the two minima for the modulus during and after inflation generically do not coincide, we have to make sure that there is no overshoot problem after inflation. This requires, for instance, that the difference between the two minima for the modulus lie not too far apart (c.f.\\ section \\ref{Sec:ExampleOvershoot}). Without a dynamical mechanism to guarantee a smooth transition between the two minima, achieving $\\sigma_{0} \\approx \\sigma_{\\inf}$ may require some amount of tuning of the model parameters. However, notice that also the KL solution requires some tuning to disentangle the height of the barrier from the gravitino mass today. Also note that, for a KKLT-type stabilisation mechanism after inflation, the mass of the modulus amounts about $m_T \\sim 16\\pi^2 \\, m_{3/2}$, which means it can be haevy enough to decay before BBN, avoiding the standard cosmological moduli problem.  We have illustrated our general strategy in a simple model of chaotic inflation with a shift symmetry supplemented by a KKLT-type superpotential and uplifting term (c.f.\\ section \\ref{Sec:AnExplicitExample}). Moduli stabilization during inflation is achieved considering a rather general K\\\"ahler potential coupling between the field which provides the inflationary vacuum energy by its F-term and the modulus. We also showed that in the limit of high-scale inflation and low-energy supersymmetry breaking, i.e.~for $W, D_{T} W \\ll W_{X}$, the corrections to the inflationary observables from the modulus sector become negligible. The appendix contains results for a general moduli stabilizing superpotential $W_{\\mod}(T)$ combined with the simple chaotic inflation model. Results for hybrid (or tribrid) inflation models and models based on a Heisenberg symmetry instead of a shift symmetry will appear elsewhere \\cite{HeisenbergTribrid}.  Finally, emphasize that our general strategy may work for more general scenarios. Even though many ingredients of our scenario are motivated from a string theory point of view, we did not consider a particular embedding in a string theory compactification here and defer this discussion to the future.  \\subsection*{Acknowledgements} We thank P.~M.~Kostka for collaboration on the initial stage of this project. We would like to thank Alexander Westphal for several illuminating comments and discussions. S.H.~is supported by the German Science Foundation (DFG) Cluster of Excellence ``Origin and Structure of the Universe''. K.D.~is supported by the German Science Foundation (DFG) within the Collaborative Research Center 676 ``Particles, Strings and the Early Universe''. S.A.~acknowledges support by the Swiss National Science Foundation."
    },
    "1112/1112.3057_arXiv.txt": {
        "abstract": "We study the impact of the spectral dependence of the linear polarization rotation induced by an achromatic half-wave plate on measurements of cosmic microwave background polarization in the presence of astrophysical foregrounds. We focus on the systematic effects induced on the measurement of inflationary gravitational waves by uncertainties in the polarization and spectral index of Galactic dust. We find that for the experimental configuration and noise levels of the balloon-borne EBEX experiment, which has three frequency bands centered at 150, 250, and 410 GHz, a crude dust subtraction process mitigates systematic effects to below detectable levels for 10\\% polarized dust and tensor to scalar ratio of as low as $r=0.01$. We also study the impact of uncertainties in the spectral response of the instrument. With a top-hat model of the spectral response for each band, characterized by band-center and band-width, and with the same crude dust subtraction process, we find that these parameters need to be determined to within 1 and 0.8~GHz at 150 GHz; 9 and 2.0~GHz at 250 GHz; and 20 and 14~GHz at 410 GHz, respectively. The approach presented in this paper is applicable to other optical elements that exhibit polarization rotation as a function of frequency. ",
        "introduction": "The Cosmic Microwave Background (CMB) polarization field can be decomposed into two orthogonal E and B modes. On large angular scales the B-mode signal encodes information about inflation, a period of rapid expansion in the early universe \\citep{Kamionkowski97,seljak97}. The signal is characterized by the tensor-to-scalar ratio $r$ which quantifies the relative strength of inflationary gravitational waves (IGW) and density perturbations generated by inflation. The level of the IGW signal encodes information about the energy scale at which inflation occurred. The current upper limit is $r~<~0.2$ \\citep{WMAP7yr_Komatsu}. A number of experimental efforts are ongoing to search for the signal at levels as low as $r \\sim 0.01$ over the coming years. On small angular scales, the B-mode signal is dominated by the `lensing signal', which results from gravitational lensing of the CMB photons by the large scale structure of the universe. The lensing converts E-mode to B-mode polarization \\citep{Zaldarriaga1998}.  Galactic foregrounds are expected to be a source of confusion for measurements of the B-mode signal. Above 70~GHz the polarized emission from Galactic dust is predicted to dominate over much of the sky and be comparable to the IGW signal with an $r$ value of 0.1 or less even for the cleanest regions of the sky \\citep{Page2007_WMAP3yrpol, Gold2009, Fraisse2011}. Therefore many experimental efforts plan to employ multiple frequencies which will enable foreground identification and subtraction.  Some CMB polarimeters use a half-wave plate (HWP) to modulate the observed linear polarization, such as EBEX \\citep{BRKspie}, SPIDER \\citep{SpiderSPIE} and POLARBEAR \\citep{PolarbearSPIE}. When observing at multiple frequency bands simultaneously, an achromatic half-wave plate (AHWP) can be used. An AHWP is a stack of monochromatic HWPs with a particular set of orientation angles relative to each other \\citep{Pancharatnam1955}. While an AHWP has a higher modulation efficiency across a broad frequency range compared to a single HWP, it rotates the polarization angle of the incident light by an amount that depends on frequency \\citep{AHWP_original,Matsumura09}. The amount of rotation depends on the construction parameters of the AHWP and on the spectrum and polarized intensity of the constituent signals, which for this paper are CMB and Galactic dust. With knowledge of the spectrum and relative polarization intensity, the amount of rotation can be calculated and corrected. However, while the spectrum of the CMB component is well known, that of dust is not. The polarized intensities of dust and CMB are also not well known. These uncertainties may pose challenges in the extraction of the underlying IGW signal. Various authors studied the impact of HWP non-idealities on measurements of CMB polarization \\citep{Brown2009, Brian2010}. However, this particular frequency dependent rotation effect has not been studied in the context of B-mode measurements. The goal of this paper is to quantitatively assess this effect. For concreteness we adopt the AHWP model, frequency bands and approximate noise information that are applicable to the E and B experiment (EBEX) \\citep{BRKspie}, a balloon-borne CMB polarimeter targeting the IGW signal at the $r \\sim 0.04$ level.  In Section 2 we describe the basic components of the simulation. Section 3 focuses on quantifying the effect of rotation due to the AHWP in the 150 GHz band. In Section 4 we use multiple frequency information to account for rotation due to Galactic dust. In Section 5, we study the additional effects of uncertainties in the spectral response of the instrument, and in Section 6 we make concluding remarks. ",
        "conclusions": "The spectral response of an achromatic half-wave plate may induce bias in the estimation of polarized signals. We analyze the level of such bias in the context of measurements of the B-mode signal of the CMB in the presence of Galactic dust, the dominant source of foreground emission in cases of interest here. For concreteness, we use the specific experimental configuration corresponding to the EBEX balloon-borne experiment.  For the area of sky considered we find that with reasonable assumptions about the magnitude and spectral shape of dust, the effects of rotation induced by the AHWP are only appreciable when dust is polarized at a level of about 5\\% and above and the tensor-to-scalar ratio $r$ is less than $\\sim0.05$. In the regime when the effects of rotation are appreciable, even a crude process of dust estimation and subtraction mitigates the effects of AHWP rotation to below detectable levels. For example, using the crude dust subtraction process we find no bias in the estimation of the B-mode power spectrum for dust polarization fraction as large as 10\\% and $r$ as low as $0.01$. For 2\\% dust polarization fraction, $r$ of 0.02 or higher is recovered without bias. We also find that the dust subtraction causes the power spectrum error bars to increase by a modest 30\\% on average.   Employing the same dust estimation and subtraction process, but now assuming errors in knowledge of the experiment's detection band-center and band-width, we find the accuracy with which these need to be measured. For example, for the particular experimental configuration considered, we found that band-center and band-width of the 150 GHz band need to be determined to better than 1 and 0.8~GHz, respectively. It is possible that this requirement may not need to be as stringent if a more sophisticated foreground estimation and subtraction process is used. This research is ongoing.  We explore the sensitivity of the particular experimental configuration to high frequency spectral leaks. Using a rejection level that is readily achievable experimentally we show that spectral leaks are not expected to pose challenges for the operation with an AHWP.  The analysis and subtraction approach discussed in this paper are applicable to other optical elements for which polarization rotation is a function of frequency. For example, ~\\citet{OBrient_thesis} describes a broadband, mm-wave detection technique that is based on sinuous antenna. It is well documented that such antennas change the phase response of polarized signals, and that this effect is frequency dependent. Thus they exhibit fundamentally the same behavior as an AHWP. Our methods and approach apply to such cases."
    },
    "1112/1112.1578_arXiv.txt": {
        "abstract": "{The daily sunspot numbers of the whole disk as well as the northern and southern hemispheres from January 1, 1945  to December 31, 2010 are used to investigate the temporal variation of the rotational cycle length through the  continuous  wavelet transformation analysis method. The auto-correlation function analysis of daily hemispheric sunspot numbers shows that the southern hemisphere rotates faster than the northern hemisphere. The results obtained from the wavelet transformation analysis are: there exists no direct relationship between the variation trend of the rotational cycle length  and the variation trend of solar activity in the two hemispheres; the rotational cycle length of both hemispheres has no significant period appearing at the 11 years, but has significant period of about 7.6 years. Analysis concerning the solar cycle dependence of the rotational cycle length shows that in the whole disk and the northern hemisphere acceleration seems to appear before the minimum time of solar activity. Furthermore, the cross-correlation study indicates that the rotational cycle length of the two hemispheres has different phases, and the rotational cycle length of the whole disk as well as the northern and southern hemispheres also has phase shifts with the corresponding solar activity. What's more, the temporal variation of North-South (N-S) asymmetry of the rotational cycle length is also studied; it displays the same variation trend as the N-S asymmetry of solar activity  in a solar cycle as well as in the considered time interval, and it has two significant periods of 7.7 and 17.5 years. Moreover, the N-S asymmetry of the rotational cycle length and the N-S asymmetry of solar activity are highly correlated. It's inferred that the northern hemisphere should rotate faster at the beginning of solar cycle 24. ",
        "introduction": "% \\label{sect:intro}   There are two main  methods  used in investigating the solar rotation rate: the trace method and the spectroscopic method \\cite{Lucio}. And it is found that the Sun has a higher rotation rate in the equatorial region: 26 days at the equator while 30~days at $60^\\circ$ latitude \\cite{Lawrence2008, Le2007}. More details about  different measures of the Sun's rotation rate can be found in the review papers \\cite{ Howard1984a, Schroeter1985, Snodgrass1992, Beck2000, Lucio}. Hoping to reach a more synthetic view of solar rotation, Heristchi \\& Mouradian \\cite{Heristchi2009} suggested a method called global rotation applied to structures of solar activity. Using this method, they indicated that  individual structures, local proper motions, meridian drift or differential rotation could be analyzed together in the considered time.  How the solar differential rotation varies  in a solar cycle as well as in a long time is still  an  unsolved problem. Li et al. \\cite{li2011a, li2011b} used a continuous complex Morlet wavelet transformation  to investigate the temporal variations of the rotational cycle length of daily sunspot areas and daily sunspot numbers from a global point of view, and indicated that  the rotational cycle length of the Sun had a secular trend, and the rotational period had no relation with the Schwabe cycle. Li et al. \\cite{li2011b} pointed out that a lower than average rotation velocity should statistically appear around the maximum time of solar activity, while around the minimum time the rotation velocity was very close to the average. But Gilman \\& Howard \\cite{Gilman1984}, Zuccarello \\& Zappal\\'{a} \\cite{Zuccarello} and Braj\\v{s}a et al. \\cite{ Brajsa2006} claimed a higher than average rotation velocity appear in the minimum time of solar activity.   The North-South (N-S) asymmetry in solar activity is an important part of solar physics. A lot of research has been done based on various solar activity  indices on the solar surface. More details about the N-S asymmetry can be found in Vizoso \\& Ballester \\cite{Vizoso1990}, Verma \\cite{Verma1993}, Carbonell et al. \\cite{Carbonell1993, Carbonell2007}, Li et al. \\cite{li2001, li2002,li2010}, and S\\'{y}kora \\& Ryb\\'{a}k \\cite{Sykora2010}. Besides, the rotational periods are also subjected to a N-S asymmetry \\cite{Temmer2003}. Gilman \\& Howard \\cite{Gilman1984} found that in the northern hemisphere the rotation was more solid-body-like. Javaraiah \\& Ulrich \\cite{Jav2006} indicated that there existed difference in the hemispheric rotation rates. Howard et al. \\cite{Howard1984b} analyzed the large spots data and found that the rotation rate increased less in the northern hemisphere. Antonucci et al. \\cite{Antonucci1990} investigated the rotational period of the photospheric magnetic field during cycle 21 and their results showed that the  two hemispheres had different dominant periods--- 26.9 days for the northern hemisphere and 28.1 days for the southern hemisphere. Also, the result of Temmer et al. \\cite{Temmer2002a, Temmer2002b} concerning the rotational periods of $H\\alpha$ flare and sunspot numbers  accorded  with the periods found by Antonucci et al. \\cite{Antonucci1990}. However, the observational result of Balthasar et al. \\cite{Balthasar1986} indicated that the sunspots had a little higher rotation rate in the southern hemisphere by analyzing sunspot groups of all types in the period 1874-1976. Georgieva \\& kirov \\cite{Georgieva2003} indicated that the two hemispheres not only rotated differently but aslo had different periodicities  in the variations of the rotation parameters. The N-S asymmetry in hydrogen filament rotation has been studied by Gigolashvili \\cite{Gigolashvili2001} and Gigolashvili et al. \\cite{Gigolashvili2003}. They found that the sign of asymmetry changed with the hale period, and they suggested that the N-S asymmetry of the solar rotation might be connected with the N-S asymmetry of solar activity.  This work follows the previous study of  Li et al. \\cite{li2011a,li2011b}. We still use the continuous complex Morlet wavelet transformation to obtain the rotational signals reflected in the daily hemispheric sunspots' wavelet power spectrum from a global point of view, and then conduct further research on temporal variation of the solar  rotation separately into the northern and southern hemispheres and on their relationship with the hemispheric solar activity. In addition, we investigate the N-S asymmetry of the solar rotational cycle length, including its time-variation, its periodicity,  and also its relationship with the  N-S asymmetry of solar activity. ",
        "conclusions": "\\label{sect:discussion}  The long-time variations of the solar rotation rate are studied  in the northern and southern hemispheres respectively through a continuous wavelet transformation method from a global point of view, and the main results are listed as follows:  1. The autocorrelation  function indicates that the southern hemisphere rotates faster than the northern hemisphere in the considered time interval. Lusting \\cite{Lusting} studied the solar differential rotation by using positions of sunspots of the years from 1947 to 1981, and found that the southern hemisphere had a smaller gradient of the differential rotation than the northern hemisphere had, in  other words,  the southern hemisphere rotated faster than the northern hemisphere. And from the Table 1 of Javaraiah et al. \\cite{Jav2005}, we may find that the southern hemisphere indeed had a smaller gradient of differential rotation. To answer why the southern hemisphere rotates faster, one need to make further study of the reasons for the difference of the long-time variations of the gradient of differential rotation in two hemispheres.    2. The rotational cycle length of the northern and southern hemispheres has different variation trends, while solar activity in the two hemispheres has the same variation trend. It means that the long-time variation trend of hemispheric rotation rate has no direct relation  with the variation trend of solar activity in the considered time. Li et al. \\cite{li2011c} pointed out that secular trends of solar rotation on an average of latitudes or at a certain latitude should change with latitudes. Thus it may be one possible reason for the different trends of the two  hemispheres' rotational cycle length that the sunspots of the northern and southern hemispheres form in different average latitudes in the considered time. Further research is needed.    3. The rotational cycle length of both the hemispheres has no significant period (scale)  appearing at the 11-year Schwabe cycle, in accordance with   Li et al. \\cite{li2011a}, but both has significant period of about 7.6 years. And the period of about 7.6 years has been found in the periodicity of the surface equatorial rotation rate by Javaraiah et al. \\cite{Jav2009}.  4. In the whole disk and the northern hemisphere, a higher than average velocity appear before the minimum time of solar activity. This may be caused by the phase difference and periodic difference between the hemispheric rotational cycle length and the hemispheric solar activity, and  may be also influenced by spatial-temporal distribution of the sunspots. The solar-cycle dependence of the two hemispheres' rotational cycle length is also different, and this may be the result of  the phase shift between the northern and southern rotational cycle length, as well as the phase shift between the northern and southern solar activity.   5. The rotational cycle length of the northern and  southern hemispheres shows difference in their phases. Additionally, the phase shifts between the rotational cycle length and the sunspot numbers in the north, south and the whole disk are different from one another. Since the relation between the hemispheric rotation and the hemispheric sunspot activity is complex,  in-depth research is needed.  6. The N-S asymmetry of the rotational cycle length  has the same variation trend as the N-S asymmetry of solar activity in a solar cycle as well as in the considered time interval. The N-S asymmetry of the rotational cycle length  has two significant periods --- 7.7 and 17.5 years. Moreover, it has high correlation with the N-S asymmetry of sunspot activity. On the basis of the aforementioned characteristic  and the regularity advanced by Vizoso \\& Ballester \\cite{Vizoso1990} and Li et al. \\cite{li2002}, it's inferred that the northern hemisphere should rotate faster at the beginning  of  solar cycle 24.     \\normalem"
    },
    "1112/1112.3327_arXiv.txt": {
        "abstract": "Large H\\,\\textsc{ii} regions, with angular dimensions exceeding 10 pc, usually enclose numerous massive O-stars. Stellar winds from such stars are expected to play a sizeable role in the dynamical, morphological and chemical evolution of the targeted nebula. Kinematically, stellar winds remain hardly observable i.e., the typical expansion velocities of wind-blown bubbles being often confused with other dynamical processes also regularly found H\\,\\textsc{ii} regions. However, supersonic shock waves, developed by stellar winds, should favor shock excitation and leave a well-defined spectral signature in the ionized nebular content. In this work, the presence of stellar winds, observed through shock excitation, is investigated in the brightest portions of the Galactic IC\\,1805 nebula, a giant H\\,\\textsc{ii} region encompassing at least 10 O-stars from main-sequence O9 to giant and supergiant O4. The use of the imaging Fourier transform spectrometer SpIOMM enabled the simultaneous acquisition of the spectral information associated to the H{$\\alpha$}\\,$\\lambda$6563 $\\mbox{\\AA}$, $[$N\\,\\textsc{ii}$]$\\,$\\lambda$$\\lambda$6548, 6584 $\\mbox{\\AA}$, and $[$S\\,\\textsc{ii}$]$\\,$\\lambda$$\\lambda$6716, 6731 $\\mbox{\\AA}$ ionic lines. Diagnostic diagrams, first introduced by Sabbadin and collaborators, were used to circumscribe portions of the nebula likely subject to shock excitation from other areas dominated by photoionization. The gas compression, expected from supersonic shocks, is investigated by comparing the pre- and post-shocked material's densities computed from the $\\frac{[\\textnormal{S}\\,\\textsc{ii}]\\,\\lambda6716}{[\\textnormal{S}\\,\\textsc{ii}]\\,\\lambda6731}$ line ratio. The typical $\\frac{[\\textnormal{N}\\,\\textsc{ii}]\\,\\lambda6584}{[\\textnormal{N}\\,\\textsc{ii}]\\,\\lambda6548}$ line ratio slightly exceeds the theoretical value of 3 expected in low-density regimes. To explain such behavior, a scenario based on collisional de-excitations affecting the $[$N\\,\\textsc{ii}$]$\\,$\\lambda$6548 $\\mbox{\\AA}$ line is proposed. ",
        "introduction": "The presence of ionized material in the interstellar medium (hereafter, ISM) can be attributed to two distinctive mechanisms. First, photoionization of the surrounding neutral gas by the strong, energetic ultra-violet (UV) flux of nearby massive stars is largely responsible for the detection of the standard hydrogen Balmer series, typically used for the morphological and kinematical description of H\\,\\textsc{ii} regions. Secondly, stellar winds with high terminal velocities and violent supernova blasts are commonly associated with the propagation of transonic and supersonic shock waves in the surrounding medium. The important increase of the post-shocked gas' temperature favors its ionization through shock excitation.  Pioneering work by \\citet{Sab1977} has compared specific flux ratios for a variety of ionic transitions in the ISM optical gas. This allowed the authors to approximately separate standard H\\,\\textsc{ii} regions and planetary nebulae (hereafter, PNe) mostly governed by photoionization from shock-dominated supernova remnants (hereafter, SNRs). These diagnostic diagrams were often used, in literature, to classify large amount of ionized targets in large-scale objects, for example more-or-less distant galaxies (e.g., \\citealt{Mag2003,Rie2006}). This has led to emission-line ratio plots in which each object is usually statistically represented by a single point. Obviously, intrinsic variations, within a given object, can be investigated by targeting Galactic objects individually (e.g., \\citealt{Phi1998,Phi1999,Phi2010}). This allows to spatially resolve much smaller ($\\ll$\\,0.1 pc) structures and artifacts characterized by peculiar line ratios that are, otherwise, unrevealed or statistically negligible in observations using poorer angular resolutions. The investigation of close ionized objects has already revealed that photoionization and shock excitation can both be found in individual regions.  Kinematically speaking, standard H\\,\\textsc{ii} regions encompassing massive stars earlier than O6 have shown little indication of a strong impact on the surrounding gas attributed to stellar winds, referred earlier as a potential source for shock excitation in nebulae. This remains usually true even when using observations with high spectral resolutions that allow to measure velocity fluctuations down to a few km s$^{-1}$. The main reason for this resides in the double-shock model said to accurately describe the dynamical evolution of wind-blown bubbles \\citep{Wea1977}. The reverse shock quickly (i.e., often within less than a spatial element of resolution) converts high-velocity ($>$\\,1\\,000 km s$^{-1}$) terminal winds into low-velocity, high-temperature gas. The thermal energy of this hot, pressurized material then initiates and fuels the expansion of the dense shell of post-shock ISM material found at much greater distances with respect to the central star. The typical expansion velocity of the shocked shell is usually a few km s$^{-1}$ above 10 km s$^{-1}$, roughly the speed of sound for warm ($\\sim$10\\,000 K) H$^{+}$ gas. Unfortunately, in complex tridimensional geometries, this kind of velocities could be easily confused with standard accelerated outflows in H\\,\\textsc{ii} regions, such as Champagne \\citep{Ten1979,Bod1979} or photoevaporated \\citep{Art2006,Mel2006} flows, turbulent motions \\citep{Jon1984,Ars1988,Lag2011} or fluid instabilities.  This work mainly explores the possibility of detecting shock waves associated, in particular, to stellar winds using the emissive properties of the ionized gas rather than a more typical, and not always successful, approach based on the information retrieved from radial velocities and non-thermal line widths.  The IC\\,1805 nebula is located in the Perseus arm of our Galaxy. The most massive stars of the Melotte 15 star cluster are currently responsible for the energetic support of the large H\\,\\textsc{ii} region. The south-central portion of the nebula (see Figure 1), in the direct vicinity of the star cluster, is gas-rich and will be used in this work to fill the \\citet{Sab1977} diagrams. Our goals are to (1) obtain reliable, high signal-to-noise ratios (hereafter, S/N), spectra of the optical emission in the brightest portions of the IC\\,1805 star-forming complex, (2) provide the corresponding series of line-ratio diagnostic diagrams, and (3) investigate the impact of photoionization and shock excitation in the targeted gas volume.  We present, in $\\S$~2, the Galactic H\\,\\textsc{ii} region IC\\,1805 and its associated star cluster, Melotte 15. Information related to our spectrometric observations and methods used for data reduction are detailed in $\\S$~3. Results of our study and diagnostic diagrams are provided in $\\S$~4. Interpretation and discussion follow in $\\S$~5. Summarized results and general conclusions will finally be provided in $\\S$~6. ",
        "conclusions": "The use of the imaging Fourier transform spectrometer SpIOMM allowed us to obtain series of emission-line profiles of the optical gas in the brightest, central portions of the Galactic IC\\,1805 H\\,\\textsc{ii} region. The bandwidth used at data acquisition allowed the simultaneous observations of the H{$\\alpha$}\\,$\\lambda$6563 $\\mbox{\\AA}$, $[$N\\,\\textsc{ii}$]$\\,$\\lambda$$\\lambda$6548, 6584 $\\mbox{\\AA}$, and $[$S\\,\\textsc{ii}$]$\\,$\\lambda$$\\lambda$6716, 6731 $\\mbox{\\AA}$ ionic lines (see $\\S$~3).  The main goal of this work was to investigate on the presence of supersonic shock waves attributed to stellar winds in the vicinity of Melotte 15, the current star cluster actually fueling up the expansion of the IC\\,1805 nebula. Literature has long suggested that a kinematical detection of stellar winds, in H\\,\\textsc{ii} region, might represent a difficult task since the typical expansion velocity of shocked shells could be commonly confused with other dynamical processes revealing similar kinematical behaviors (see $\\S$~1). On the other hand, specific line ratios retrieved from the optical gas may indicate if the presence of ionized material is attributed to standard photoionization or to shock excitation. These line ratios therefore provide a non-kinematical tool for identifying shocks in the nebular volume. Our results are summarized as follow:  \\begin{enumerate} \\renewcommand{\\labelenumi}{\\Roman{enumi}.} \\item The $\\frac{[\\textnormal{N}\\,\\textsc{ii}]\\,\\lambda6584}{[\\textnormal{N}\\,\\textsc{ii}]\\,\\lambda6548}$ line ratio genuinely deviates from the theoretical value of 3 (see $\\S$~5.1). Values varying between 2.5 and 4 were commonly found. The distribution of the density measurements has not allowed to demonstrate that the $[$N\\,\\textsc{ii}$]$\\,$\\lambda$6548 $\\mbox{\\AA}$ line could be affected by collisional de-excitations, hence reducing its peak intensity. This scenario remains nonetheless plausible and could be verified if densities, in the N$^{+}$ volume, could be measured precisely (see $\\S$~5.1). Densities extracted from the $\\frac{[\\textnormal{S}\\,\\textsc{ii}]\\,\\lambda6716}{[\\textnormal{S}\\,\\textsc{ii}]\\,\\lambda6731}$ ratio may not perfectly reflect the conditions prevailing in the N$^{+}$ volume if both nitrogen and sulfur are not co-spatial in central IC\\,1805. \\item The diagnostic diagrams, introduced by \\citet{Sab1977}, indicate, in first approximation, that photoionization most likely dominate in IC\\,1805 (see $\\S$~4.2.5). This initially holds even for the densest, most emissive structures in our field-of-view, directly exposed to the radiation fields and stellar winds of the nearby, most massive stars (see $\\S$~5.2). \\item Evidences for shock excitation appear only following the subtraction of the diffuse foreground/background material. This component likely occupies a very large fraction of the nebular volume in IC\\,1805; a volume sufficiently large that, even though being tenuous in nature, the foreground/background material strongly dilutes the signal emanating from shocked condensations along the line-of-sight. Its mean density is estimated at 25 cm$^{-3}$ (see $\\S$~5.2.1). \\item Oriented on a south-north axis and surrounded by numerous O-stars, a bright, large ionized feature occupies the central area of our field-of-view. The last, tenuous fragments of an old molecular cloud can be found near its southern portion. Shocks may have contributed to ionize material found in the direct vicinity of the molecular clump while the northern parts of the ionized feature, deprived of molecular emission, appears to be largely dominated by photoionization (see $\\S$~5.2.1). \\item The shock-excited ionized gas has a mean density of 175 cm$^{-3}$ and the compression factor, between pre- and post-shocked gas, is typically between 2 and 10. For isothermal shocks, this suggests shock velocities between 15 and 30 km s$^{-1}$, in agreement with models describing the expansion of wind-blown bubbles (see $\\S$~5.2.1). Geometrically speaking, given the apparent proximity between the central, ionized structure and the most massive stars of Melotte 15, winds seem to have had sufficient time to reach the structure within a timescale corresponding to the age of the star cluster (see $\\S$~5.2.1). This gives credence to our assumption that shock excitation can be found in the south-central portions of our field-of-view. \\item Points identified with a high probability of shock excitation reveal compression factors typically greater than those points where shocks may be present although less certain (see $\\S$~5.2.1). \\item Shocks did not seem to have played a major role in the ionization of a molecular cloud's envelope located in the south-eastern portion of our field-of-view. This most likely results from the molecular fragment being located too far from the ionizing sources so that (1) shocks induced by stellar winds have not reached yet the cloud or (2) shocks have reached the cloud with velocities too low to initiate shock excitation (see $\\S$~5.2.2.1). \\item Shock development was clearly detected on the outskirts of an apparently weak, but well-defined elephant trunk located in the eastern portion of our field-of-view (see $\\S$~5.2.2.2). This is in agreement with theoretical works developed on such feature typically found in H\\,\\textsc{ii} regions. \\end{enumerate}"
    },
    "1112/1112.2926_arXiv.txt": {
        "abstract": "We present a mechanism for the dark matter stability  in the framework of a non-Abelian flavour symmetry renormalizable model. The same non-abelian discrete flavor symmetry which accounts for the observed pattern of neutrino oscillations, spontaneously breaks to a $Z_2$ subgroup which renders DM stable. The simplest scheme leads to a scalar doublet DM potentially detectable in nuclear recoil experiments, inverse neutrino mass hierarchy, hence a neutrinoless double beta decay rate accessible to upcoming searches. ",
        "introduction": "Non-baryonic Dark Matter (DM) is one of the most compelling problems of modern cosmology. Despite the fact that its existence is well established by cosmological and astrophysical probes, its nature remains elusive. Still, observations can constraint the properties of dark matter and give some hints about its identity. For instance, a fundamental requirement for a viable dark matter candidate is the stability over cosmological times. This suggests the existence of an exact or slightly-broken symmetry protecting or suppresing its decay.  It has been shown recently that such symmetry can also be related to the flavor structure of the Standard Model~\\cite{arXiv:1007.0871}. The model proposed in ~\\cite{arXiv:1007.0871} is based on a $A_4$ symmetry with four $SU(2)$ Higgs doublets. After the electroweak symmetry breaking, the $A_4$ (even permutation of four objects) group is spontaneously broken into a $Z_2$ subgroup which is responsible for the DM stability. The leptonic sector is also extended. It consist in four right handed neutrinos and the light neutrino masses are generated through the type-I seesaw mechanism and obey an inverted hierarchy mass spectrum with $m_{\\nu_3}=0$ and vanishing reactor angle $\\theta_{13}=0$.  We have also been consider models with a different matter content for the right handed neutrinos with the same DM stability mechanism but with different neutrino phenomenology~\\cite{arXiv:1011.1371} or a model based on the dihedral group $D_4$ where the some flavour changing neutral currents are present and constraint the DM sector~\\cite{arXiv:1104.0178}. Other models with flavor symmetries but with decaying DM have also been considered, see for instance~\\cite{arXiv:1010.4963,arXiv:1011.5753}. For a model with stable DM but with the same behavior of the DM, in the sense that instead of being inert totally, the DM couples to some right handed neutrinos in a similar way in our model\\footnote{This afther the $A_4$ is broken into the $Z_2$ symmetry.} see~\\cite{arXiv:1111.4938}. ",
        "conclusions": "We have studied a model where the stability of the dark matter particle arises from a flavor symmetry.  The $A_4$ non-abelian discrete group accounts for both the observed pattern of neutrino mixing and for DM stability. We have analyzed the constraints that follow from electroweak precision tests, collider searches and perturbativity. We have also analyzed the prospects for direct and indirect dark matter detection and found that, although the former already excludes a large region in parameter space, we cannot constrain the mass of the DM candidate. In contrast, indirect DM detection is not yet sensitive enough to probe our predictions. However, forecasted sensitivities indicate that Fermi-LAT should start probing them in the near future.   All of the above relies mainly on the properties of the scalar sector responsible for the breaking of the gauge and flavor symmetries. The motivation of our approach is to link the origin of dark matter  to the origin of neutrino mass and the understanding of the pattern of neutrino mixing, two of the most outstanding challenges in particle physics today. At this level one may ask what are the possible tests of this idea in the neutrino sector. We found an inverted neutrino mass hierarchy, hence a neutrinoless double beta decay rate accessible to upcoming searches, while $\\theta_{13}=0$ giving no CP violation in neutrino oscillations. Note however that the connection of dark matter to neutrino properties depends strongly on how the symmetry breaking sector couples to the leptons."
    },
    "1112/1112.0438_arXiv.txt": {
        "abstract": "We review basic properties of the population of dwarf galaxies in the Local Group focusing on dwarf spheroidal galaxies found in the immediate vicinity of the Milky Way. The evidence for dark matter in these objects is critically assessed. We describe the methods of dynamical modelling of such objects, using a few examples of the best-studied dwarfs and discuss the sources of uncertainties in mass estimates. We conclude with perspectives for dwarf galaxies as targets for dark matter detection experiments. ",
        "introduction": "The population of dwarf galaxies in the Local Group offers a unique opportunity to test our theories of structure formation in the Universe. Starting from the Magellanic Clouds which were known since antiquity, the census of the dwarf galaxies in our immediate vicinity still grows. The sample of dwarf galaxies in the Local Group can be divided using the morphological criteria into classes of dwarf irregulars (dIrr) that are flattened, rotating, bright and still forming stars and dwarf spheroidals (dSph) which are rounder, faint and contain mostly old stellar populations dominated by random motions.  Dwarf galaxies of the Local Group tend to cluster around the two main hosts, the Milky Way and Andromeda, exhibiting a pronounced morphology-density relation: while dSphs are typically found close to one of the big galaxies, dIrrs occupy more isolated regions. The origin of this relation is an interesting issue that the theories of structure formation must address. Another question in which dwarf galaxies played a role is the so-called problem of missing satellites: the number of observed satellites of the Milky Way is much smaller than predicted by theories based on cold dark matter. Dwarf galaxies, especially nearby dSphs, have also drawn attention once a significant number of stellar velocities could be measured and their masses could be determined. These masses turned out to be much larger than could be explained by the stellar content indicating large amounts of dark matter present.  \\begin{figure} \\epsfig{file=lokas_fig1.eps, height=10cm} \\caption{Observational parameters of dwarf galaxies in the Local Group: the central surface brightness $\\mu_V$ (upper left panel), the half-light radius $r_{1/2}$ (lower left panel), the ratio of the rotation velocity to the velocity dispersion $V/\\sigma$ (upper right panel) and the ellipticity $e=1-b/a$ (lower right panel) as a function of the total visual magnitude $M_V$. The open circles show the data for dIrr galaxies, filled black circles the data for the dSph and dSph/dE galaxies and gray circles for transitory dIrr/dSph dwarfs.} \\label{lokas_fig1} \\end{figure}  Observational properties of dwarf galaxies \\cite{mateo98} are usually characterized by the total visual magnitude $M_V$, the central surface brightness $\\mu_V$, the half-light radius (containing half the total luminosity) $r_{1/2}$, the ellipticity $e=1-b/a$ and the ratio of the rotation velocity to the velocity dispersion $V/\\sigma$. The latter four quantities are plotted in Figure~\\ref{lokas_fig1} as a function of magnitude for the dwarf galaxies in the Local Group with $-16$ mag $< M_V < -8$ mag using data compiled in \\cite{lokas11}. The sample includes the eight ``classical\" best-known dSph galaxies, satellites of the Milky Way: Carina, Draco, Fornax, Leo I, Leo II, Sculptor, Sextans and Ursa Minor with magnitudes in the range $-13$ mag $< M_V < -8$ mag which all have low central surface brightness $\\mu_V$ ($>22.4$ mag arcsec$^{-2}$) and small half-light radii ($< 0.7$ kpc). The properties of these eight dwarfs are listed in Table~\\ref{properties}. All classical dwarfs are believed to be characterized by large mass-to-light ratios between 10 and a few hundred solar units.  \\begin{table} \\begin{center} \\caption{Properties of the classical dSph galaxies, satellites of the Milky Way.} \\begin{tabular}{lrcccc} \\hline Dwarf\t     &  $M_V \\ $ & $r_{1/2}$ & $\\mu_V$             & $V/\\sigma$ & $e=1-b/a$   \\\\ galaxy       &  [mag] &  [kpc]    & [mag arcsec$^{-2}$] &            &             \\\\ \\hline Carina       &$\t-8.62 $ & 0.241 &  25.5\t& 0.43\t&\t0.33     \\\\ Draco        &$\t-8.74 $ & 0.196 &  25.3 & 0.21\t&\t0.29     \\\\ Fornax       &$\t-13.03$ & 0.668 &  23.4 & 0.18\t&\t0.31     \\\\ Leo I        &$\t-11.49$ & 0.246 &  22.4 & 0.33  &       0.21     \\\\ Leo II       &$\t-9.60 $ & 0.151 &  24.0\t& 0.28\t&\t0.13     \\\\ Sculptor     &$\t-10.53$ & 0.260 &  23.7 & 0.30\t&\t0.32     \\\\ Sextans      &$\t-9.20 $ & 0.682 &  26.2 & 0.48\t&\t0.35     \\\\ Ursa Minor   &$\t-8.42 $ & 0.280 &  25.5\t& 0.49\t&\t0.56     \\\\ \\hline \\label{properties} \\end{tabular} \\end{center} \\end{table}  In recent years we have witnessed discoveries of new dSph galaxies in the Local Group, mainly in the northern part of the Sloan Digital Sky Survey \\cite{belokurov07}. These new dwarfs are generally fainter and more irregular in shape than the classical ones but not more distant. If analogous discoveries follow in the southern hemisphere, they may significantly help to solve the problem of missing satellites. Spectroscopic studies of the stellar populations \\cite{simongeha} of the faint dwarfs revealed quite large velocity dispersions and thus large masses for these dwarfs were estimated suggesting even higher dark matter content than in well-known dSphs. It has been speculated \\cite{strigari} that the characteristic masses of all dwarfs are of the order of $10^7$ solar masses, independently of their luminosity. A more acceptable proposal \\cite{walker} suggests instead a universal relation between the half-light radius and the mass contained within it. ",
        "conclusions": "The case for a high dark matter content in dSph galaxies of the Local Group is quite strong. The objects are found to contain significant amounts of dark matter even if the modelling is performed with very conservative assumptions such as the one that mass follows light and on samples which were subject to restrictive algorithms of interloper rejection. It has been demonstrated that the classical dwarfs, even if strongly affected by tidal force from the Milky Way can be reliably modelled by standard methods assuming virial equilibrium.  The high mass-to-light ratios of the classical dwarfs may turn out to be even higher in the case of the newly discovered population of ultra-faint satellites of the Milky Way. If the trend is confirmed, the fainter dwarfs may become excellent targets for experiments aiming for the direct detection of dark matter in the Universe via self-annihilation of dark matter particles. Dwarf galaxies are good candidates for such targets not only because of the high dark matter content but also because they have low astrophysical gamma ray backgrounds and they are small and localized so may easily fit in the field of view of many instruments.  Recently, observations of dwarf galaxies with the Fermi Large Area Telescope began to yield interesting constraints on the parameters of a variety of supersymmetric models that provide candidates for weakly interacting massive particles as dark matter \\cite{abdo10}. Although no clear signal of annihilation has been detected from dwarf galaxies so far, upper limits on the annihilation cross-section can be obtained. The upper limits do not yet reach the interesting region of density parameter values compatible with the standard cosmological model, but are already able to exclude a significant number of supersymmetric models of particle physics."
    },
    "1112/1112.0324_arXiv.txt": {
        "abstract": "We present our \\emph{XMM-Newton} observation of the fastest rotating spiral galaxy UGC~12591.  We detect hot gas halo emission out to 110~kpc from the galaxy center, and constrain the halo gas mass to be smaller than $3.5\\times10^{11} \\msun$.  We also measure the temperature of the hot gas as $T=0.64\\pm0.03$~keV.  Combining our X-ray constraints and the near-infrared and radio measurements in the literature, we find a baryon mass fraction of 0.03--0.04 in UGC~12591, suggesting a missing baryon mass of 75\\% compared with the cosmological mean value.  Combined with another recent measurement in NGC~1961, the result strongly argues that the majority of missing baryons in spiral galaxies does not reside in their hot halos. We also find that UGC~12591 lies significantly below the baryonic Tully-Fisher relationship. Finally, we find that the baryon fractions of massive spiral galaxies are similar to those of galaxy groups with similar masses, indicating that the baryon loss is ultimately controlled by the gravitational potential well. The cooling radius of this gas halo is small, similar to NGC~1961, which argues that the majority of stellar mass of this galaxy is not assembled as a result of cooling of this gas halo. ",
        "introduction": "Observations show that nearby galaxies are missing most of their baryons (e.g., Hoekstra et al.\\ 2005; Heymans et al.\\ 2006; Mandelbaum et al.\\ 2006; Gavazzi et al.\\ 2007; Jiang \\& Kochanek 2007; Bregman 2007) when compared to the cosmological baryon to matter ratio (e.g., $f_b = 0.171\\pm0.009$ from WMAP, Dunkley et al.\\ 2009). For example, the Milky Way is missing two thirds of its baryon allotment (e.g., Sakamoto et al.\\ 2003) and less massive galaxies have retained less than 10\\% of their baryons (e.g., Corbelli 2003; Walker et al.\\ 2007). This situation is confirmed in other galaxies through a variety of methods (e.g., Hoekstra et al. 2005; McGaugh 2007). However, most of these studies do not include the baryon mass in the hot gas halo of galaxies, and it is possible that the majority of the missing baryons in galaxies actually resides in their hot has halos based on theoretical predictions (e.g., White \\& Frenk 1991; Sommer-Larsen 2006; Fukugita \\& Peebles 2006; also see the review by Benson 2010).  This component can be difficult to detect for spiral galaxies due to its faintness, especially if the gas density profile is flat. As another possibility, the missing baryons from galaxies may have escaped from the potential wells of the galaxies but reside in their parent groups or clusters (e.g., Humphrey et al.\\ 2011). Finally, the missing baryons can be in the form of warm-hot intergalactic medium (Cen \\& Okstiker 1999; 2006). Deep X-ray observations are needed to distinguish these different scenarios.  The situation is different for rich galaxy clusters, where the gas mass dominates the baryon content, and the baryon fractions in these massive systems are close to the cosmological value after combining the gas and stellar baryon contributions (e.g., Vikhlinin et al.\\ 2006; Allen et al.\\ 2008). The measurements of the baryon fraction in different systems suggest that the fraction depends on the dynamical mass of the systems: rich clusters retain their cosmological allotment of baryons, while galaxies are baryon-poor. We summarize the situation in Dai et al.\\ (2010) by combining the archival data points reported in the literature and our stacking analysis result using the ROSAT All-Sky Survey data of 4,000 nearby galaxy groups and clusters (Dai et al.\\ 2007). We find that the baryon fractions from dwarf galaxies to rich galaxy clusters can be fit by a broken power-law model with the break at the circular velocity of $V_c \\sim 440~\\kms$.  The scatter of the fractions about the mean relation is small considering the huge dynamic range of the systems.  Further examining the relation, we find that the baryon fractions are similar for different systems with similar total masses but different compositions. For example, the baryon fractions of poor galaxy groups, where the baryon mass is still dominated by the gas mass, are close to those of massive galaxies, where the baryon mass is dominated by the stellar mass.  Such a coincidence is puzzling considering the differences between their mass compositions and energy feedback mechanisms.  To test whether the missing baryons reside in galaxy halos and further constrain the baryon loss in different mass scales, we focus on the massive galaxy \\ugc. \\ugc\\ is a spiral galaxy with the largest measured rotational velocity to date (466-500 km/sec, Giovanelli et al. 1986; Paturel 2003). To appreciate this galaxy, its optical-IR luminosity is nine times that of M31. Nine big spirals were in a single group is richer than Hickson groups and is about half of the Fornax cluster, but much more compact. In this paper, we combine our \\xmm\\ observation of \\ugc\\ with 2MASS and other data to determine the baryon fraction and the composition of this massive spiral galaxy. Throughout the paper, we use the cosmological parameters from WMAP with $H_0 = 72~\\rm{km~s^{-1}~Mpc^{-1}}$, $\\Omega_{\\rm m} = 0.26$, and $\\Omega_{\\Lambda}= 0.74$ (Dunkley et al.\\ 2009). ",
        "conclusions": "\\begin{deluxetable}{lccccccccc} \\tabletypesize{\\scriptsize} \\tablecolumns{10} \\tablewidth{0pt} \\tablecaption{Gravitational and Baryon Mass Components in \\ugc \\label{tab:mass}} \\tablehead{ \\colhead{Component} & \\colhead{within 50~kpc radius} & \\colhead{within 500~kpc radius} } \\startdata Stellar           & $45\\pm10$     & $45\\pm10$ \\\\ Cold Gas          & 0.7           & 0.7       \\\\ Hot Gas           & $0.41\\pm0.03$ & 13        \\\\ Including Flattened Hot Gas & 0.55       & $\\ls 35$  \\\\ Gravitational     & 270           & 1900      \\\\ Stellar-to-gas ratio $r_{sg}$ & 39         & 3.3       \\\\ $r_{sg}$ including the flattened gas & 33 & 3.3--1.3 \\\\ Baryon Fraction $f_b$  & 0.17          & 0.03      \\\\ $f_b$ including the flattened gas & 0.17 & 0.03--0.04 \\\\ \\enddata \\tablecomments{The masses are in the unit of $10^{10} \\msun$.  The \\rtwo\\ is at $\\simeq 550$~kpc.} \\end{deluxetable}  \\subsection{Baryon Mass Components in \\ugc} We list the various baryon mass components in \\ugc\\ in Table~\\ref{tab:mass}. We have constrained the hot gas mass of \\ugc\\ using the \\xmm\\ observation as $(4.1\\pm0.3)\\times10^{9} \\msun$ within 50~kpc with an average temperature of $T=0.64\\pm0.03$~keV. We have also constrained the hot gas mass of $1.3\\times10^{11} \\msun$ within 500~kpc regions using our best fit $\\beta$ model, and the hot gas mass is below $3.5\\times10^{11} \\msun$ within 500~kpc even if we add another flatter $\\beta$ model component in our fits. Beside the hot gas mass, there are other baryon mass components in the galaxy including the stellar mass and cold gas mass components. For the cold gas mass component, \\citet{gi86} measure the \\HI\\ mass of $5.3\\times10^{9} h_{72}^{-2} \\msun$ from the radio data. Assuming the \\HI\\ mass is 75\\% of the total cold gas mass, we find that the total cold gas mass is $M_{cg} = 7.1\\times10^9 \\msun$, where we ignore the contribution from the molecular gas component. We estimate the stellar mass of \\ugc\\ as $(4.5\\pm1.0)\\times10^{11} \\msun$ within 29~kpc radius using its $K$ band total magnitude ($K = 8.89$~mag) from the 2MASS Extended Source Catalog and a range of mass-to-light ratio from 0.6 to 0.95 (e.g., Bell et al.\\ 2003). The 2MASS team calculates the total magnitude by integrating the surface brightness profile out to $\\sim 4$ disk scale lengths from the isophotal aperture well below the $1\\sigma$ noise level. For the mass-to-light ratio, Bell et al.\\ (2003) measure a value of 0.95 as the cosmic mean value.  However, since \\ugc\\ is a late-type galaxy and could have a lower mass-to-light ratio, we choose to use $0.78\\pm0.18$ in our calculation. We find the largest uncertainties in the baryon mass are from the systematical uncertainty in the stellar mass-to-light ratio and the flattened gas halo. Combining the two effects, we find an uncertainty of $2.4\\times10^{11} \\msun$ for the total baryon mass within 500~kpc.  In the central region within $\\sim50$~kpc, the baryon mass is clearly dominated by the stellar mass, and the stellar-to-gas mass ratio is $r_{sg} \\simeq 39$.  Using the rotational velocity of 466--500~\\kms, we measure a total mass of $m_{tot} = 2.7\\times10^{12} \\msun$ within 50~kpc and a baryon mass fraction of $f_b \\simeq 0.17$, consistent with the cosmological baryon fraction. Out to the 500~kpc region (\\rtwo $\\simeq 550$~kpc), we use the gravitational mass $m_{tot} = 1.9\\times10^{13} \\msun$ estimated from the X-ray data, because the rotational curve is only constrained within 28$h_{72}^{-1}$~kpc \\citep{gi86}, smaller than the total mass $m_{tot} = 2.7\\times10^{13} \\msun$ estimated using the rotational curve.  The baryon mass within 500~kpc is $m_b = 5.9\\times10^{11} \\msun$, and we measure a baryon fraction of $f_b \\simeq 0.03$.  Considering a second flattened gas component, the baryon fraction within 500~kpc can reach to $f_b \\ls 0.04$. Since we use the smaller total mass estimate in the calculation, the baryon fraction quoted should be treated as a conservative upper limit. The stellar mass component is still more important with $r_{sg} = 3.3$ within 500~kpc, or $r_{sg} \\gs 1.3$ with the additional flatter gas component.  To summarize, combining our \\xmm\\ observation and the 2MASS and radio data in the literature, we have constrained that \\ugc\\ has lost at least 75\\% of the baryons compared to the cosmological value.  The missing baryons do not reside in the hot halos for spiral galaxies.  Our result confirms the recent measurements in another giant spiral NGC~1961 using \\chandra\\ by Anderson \\& Bregman (2011), who find that NGC~1961 has also lost 75\\% of its baryon content.  \\subsection{Baryonic Tully-Fisher Relationship} \\begin{figure} \\epsscale{1} \\plotone{f9.eps} \\caption{The baryonic Tully-Fisher relationship of McGaugh (2005) adding NGC~1961 (Anderson \\& Bregman 2011) and \\ugc. \\label{fig:btf}} \\end{figure} Using the fastest rotating galaxy \\ugc, we are able to extend the baryonic Tully-Fisher relationship (BTF), a correlation between the baryon mass and rotational velocity of galaxies (McGaugh 2005; 2011) to the high rotational velocity regime of 500~\\kms. We plot \\ugc\\ in the BTF diagram together with the galaxies in McGaugh (2005) and the other massive galaxy NGC1961 (Anderson \\& Bregman 2011) with a rotational velocity of 402~\\kms\\ in Figure~\\ref{fig:btf}. Anderson \\& Bregman (2011) find that NGC1961 deviates slightly from the linear fits to BTF. However, the authors caution that the offset can be caused by systematic uncertainties. Indeed, assuming a $K$ band mass-to-light of 0.95, NGC~1961 would be on the BTF relation. Here, \\ugc\\ provides another challenge to the linear BTF relation, which predicts a baryon mass of $2.2\\times10^{12} \\msun$ for \\ugc, whereas we measure a baryon mass in the range of 5.9--8.1$\\times10^{11} \\msun$ with an uncertainty of $2.4\\times10^{11} \\msun$.  Thus, \\ugc\\ is $6\\sigma$ below the BTF relation.  If the offset from the BTF relation is caused by the uncertainties in the $K$ band mass-to-light ratio, a ratio of 3.1 is needed to put \\ugc\\ on the BTF relation, which is extremely unlikely (e.g., Bell et al.\\ 2003). Thus, it is possible that the BTF relation turns over for massive galaxies with $v_c \\gs 400 \\kms$. However, measurements from a larger sample of massive spiral galaxies are needed to confirm this result.  \\subsection{Overall Relationship of Baryon Fractions with Total Mass} \\begin{figure} \\epsscale{1} \\plotone{f10.eps} \\caption{The baryon fraction as a function of the gravitational potential well indicated by the circular velocity at \\rtwo.  We plot the new measurements from massive spiral galaxies \\ugc\\ and NGC~1961 (Anderson \\& Bregman 2011) over the archival data from Sakamoto et al.\\ (2003), McGaugh (2005), Flynn et al.\\ (2006), Vikhlinin et al.\\ (2006), Gavazzi et al.\\ (2007), Walker et al.\\ (2007), Stark et al.\\ (2009), Sun et al.\\ (2009), Dai et al.\\ (2010). The dotted line is the cosmological baryon fraction measured from CMB, and the dashed line is our best fit broken power law model for baryon losses.  For massive spiral galaxies like \\ugc\\ and NGC~1961, the baryon loss is at least 75\\%. \\label{fig:bf}} \\end{figure} We plot the baryon fractions of \\ugc\\ and NGC~1961 against their rotational velocities, a proxy for the depth of the gravitational potential well, in Figure~\\ref{fig:bf}. We include the data composed by Dai et al.\\ (2010) including the stacked results from optically-selected clusters (Dai et al.\\ 2007; 2010), individual galaxy clusters (Vikhlinin et al.\\ 2006; Sun et al.\\ 2008) corrected for the stellar component (Dai et al.\\ 2010), individual galaxies (McGaugh 2005; Walker et al.\\ 2007; Stark et al.\\ 2009), elliptical galaxies (Gavazzi et al.\\ 2007) and the Milky Way (Sakamoto et al.\\ 2003; Flynn et al.\\ 2006), and the recent addition of gas-rich late-type galaxies (Begun et al.\\ 2008; Trachternach et al.\\ 2009) composed by McGaugh (2011). This enables us to compare the baryon loss across a large range of systems from dwarf galaxies to rich galaxy clusters. For massive spiral galaxies such as \\ugc\\ and NGC~1961, their dark matter halo masses are close to that of a medium galaxy group, and we find that their baryon losses are comparable to the stacked results for galaxy groups in Figure~\\ref{fig:bf}. For different systems as spiral galaxies and galaxy groups, the consistency between their baryon fractions suggests that the baryon loss is ultimately controlled by the potential well of the dark matter halo. However, currently we still lack constraints from individual galaxy groups to confirm the stacking results of Dai et al.\\ (2010).  The overall baryon fractions for all systems can be fit by a broken power law model (Dai et al.\\ 2010). With the addition of new measurements, especially from those gas-rich late-type galaxies (Begun et al.\\ 2008; Trachternach et al.\\ 2009), we can better fit the power law slope for baryon fractions in less massive systems. Thus, we re-fit the data with a broken power law model to find that \\begin{equation} f_{b} = \\frac{0.16 (v_c/643~{\\rm km/s})^{a}}{{(1+(v_c/643~{\\rm km/s})^{c})}^{b/c}}, \\end{equation} where $a=1.26$, $b=1.24$, and $c=2$.  The baryon fraction, $f_{b}$, scales as $f_b\\propto v_c^{a-b=0.02}$ above the break and \u00a0$f_{b} \\propto v_c^{a=1.26}$ below the break, and the parameter $c$ in the equation is the smoothness of the broken power law model, which is fixed in our fit. Comparing with the fit in Dai et al.\\ (2010), we find the major difference lies in the shallower slope $f_b \\propto v_c^{1.26}$ for the baryon loss in galaxies. We also find a larger break location and a shallower slope for galaxy clusters.  \\subsection{Cooling of the Gas Halo} The gas halo is predicted to play an important role in galaxy formation and evolution. With our detection of the gas halo emission in \\ugc, we can estimate the cooling time of this hot halo and the implied accretion rate onto the galaxy, which can provide constraints on the gas available for new star formation. We define the cooling radius as the radius where the cooling time is 10~Gyr, using the expression of the cooling time (Fukugita \\& Peebles 2006), \\begin{equation} \\tau(r) = \\frac{1.5 nkT}{\\Lambda n_e \\left(n-n_e\\right)} \\approx \\frac{1.5kT\\times1.92}{\\Lambda n_e \\times 0.92}, \\end{equation} where the latter expression assumes a primeval He abundance resulting in a total particle density of $n = 1.92 n_e$. For $T = 10^{6.85}$~K, $Z/ Z_{\\odot} = 0.5$, and $\\Lambda = 10^{-22.85}$ erg cm$^{3}$ s$^{-1}$ (Sutherland \\& Dopita 1993), the cooling radius is at $n_e = 6.8\\times10^{-4}$ cm$^{-3}$. For the range of best-fit $\\beta$-model profiles constrained in this paper, this corresponds to a cooling radius between 15.6 and 18.0 kpc, and a hot halo mass of $6.2-9.2\\times10^8 M_{\\odot}$ within that radius. We can roughly estimate the cooling time and rate by dividing the thermal energy in the hot gas within the cooling radius by the luminosity within that radius, and this yields a wide range in cooling time of 2.8-6.3 Gyr for material within the cooling radius, but a fairly narrow range in the effective cooling rate of $0.15-0.21 M_{\\odot}$ year$^{-1}$. This halo accretion rate is two orders of magnitude too low to assemble the stellar mass of this galaxy within a Hubble time. Therefore, significant accretion must have occurred via some other mode, such as cold flows or mergers, to produce the stellar mass seen in this galaxy today, confirming the conclusion drew in NGC~1961 (Anderson \\& Bregman 2011)."
    },
    "1112/1112.5447_arXiv.txt": {
        "abstract": "We discuss the combined constraints on a Finite Scale Factor Singularity (FSF) universe evolution scenario, which come from the shift parameter $\\mathcal{R}$, baryon acoustic oscillations (BAO) $\\mathcal{A}$, and from the type Ia supernovae. We show that observations allow existence of such singularities in the $2\\times10^9$ years in future (at $1\\sigma$ CL) which is much farther than a Sudden Future Singularity (SFS), and that at the present moment of the cosmic evolution, one cannot differentiate between cosmological scenario which allow finite scale factor singularities and the standard {\\bf $\\Lambda$CDM} dark energy models. We also show that there is an allowed value of $m=2/3$ within $1\\sigma$ CL, which corresponds to a dust-filled Einstein-de-Sitter universe limit of the early time evolution and so it is pasted ito a standard early-time scenario. ",
        "introduction": "\\setcounter{equation}{0} Finite scale factor singularities (FSF) are one of types of new and exotic singularities which were found first in Ref. \\cite{BGT}. The inspiration to search these new types of singularities was due to the observations of high-redshift type Ia supernovae (SNIa) which provided strong evidence that the expansion of the universe is accelerating due to an unknown form of dark energy\\cite{supernovaeold} phenomenologically behaving as the cosmological constant. Further observational data \\cite{supernovaenew} made cosmologists think of more accelerating universe filled with phantom \\cite{phantom} which violated all energy conditions: the null ($\\varrho c^2 + p \\geq 0$), weak ($\\varrho c^2 \\geq 0$ and $\\varrho c^2 + p \\geq 0$), strong ($\\varrho c^2 + p \\geq 0$ and $\\varrho c^2 + 3p \\geq 0$), and dominant energy ($\\varrho c^2 \\geq 0$, $-\\varrho c^2 \\leq p \\leq \\varrho c^2$) ($c$ is the speed of light, $\\varrho$ is the mass density in $kg m^{-3}$ and $p$ is the pressure). Phantom-driven dark energy leads to a big-rip singularity (BR or type I according to \\cite{nojiri}) in which the infinite values of the energy density and pressure ($\\rho$, $p\\to\\infty$) are accompanied by the infinite value of the scale factor ($a\\to\\infty$) \\cite{caldwellPRL}.  The list of new types of singularities contains: a big-rip (BR), a sudden future singularity (SFS) \\cite{barrow04}, which can appear in inhomogeneous and anisotropic models too \\cite{barrow042, sfs1}, a generalized sudden future singularity (GSFS), a finite scale factor singularity (FSF) \\cite{aps}, a big-separation singularity (BS). They are characterized by violation of all or some of the energy conditions which results in a blow-up of all or some of the physical quantities: the scale factor, the energy density and the pressure. The finite scale factor singularity, which is the subject of this paper, is a weak singularity according to Tipler's definition, but strong according to Kr\\'{o}lak's definition \\cite{lazkoz}. Apart from mentioned above there are also $w$-singularities \\cite{wsin}, which are not physical singularities, but are singularities of a barotropic index $w$ which are present in different cosmological models of $f(R)$ gravity \\cite{star1980}, in scalar field models \\cite{Setare}, and in brane cosmologies \\cite{Sahni}.  For a finite scale factor singularity (FSF) $\\rho \\to \\infty$ and $p \\to \\infty $ diverge, while the scale factor $a$ remains constant. This means that it is similar to a big-bang singularity with only one exception - the scale factor is a constant $a = a_s=$ const instead of zero $a_{BB} = 0$. Besides, FSF singularities are stronger than a sudden future singularities (SFS) [10], so they placed themselves in between the big-bang (which is strong and geodesically incomplete) and the SFS. It has been shown that SFS appear in physical theories such as $f(R)$ gravity, scalar field cosmologies, brane cosmologies \\cite{yuri} and, in particular, they plague loop quantum cosmology \\cite{LQC}. It seems then that after an appropriate choice of the scalar field, $f(R)$ function or brane, FSF will be easily deduced to appear in such physical theories as well.  \\indent The paper is organized as follows. In section \\ref{s2} we present FSF scenario. In section \\ref{oc} we derive the expressions for the observables: baryon acoustic oscillations, distance to the last scattering surface, type Ia supernovae, used to test an FSF scenario. In section \\ref{rac} we give the results and discussion. ",
        "conclusions": "\\label{rac} \\setcounter{equation}{0}  Overall $\\chi^2$ that has been used is \\be \\chi^2=\\chi^2_{SN}+\\chi^2_{R}+\\chi^2_{A} \\ee   In Fig. \\ref{fig1} we present contour plots showing the joint marginal posterior distribution for each pair of FSF model parameters. Each sub-panel shows three contours, denoting roughly 68\\%, 95\\% and 99\\% (from light gray to dark grey respectively) credible regions.  In Fig. \\ref{fig1} we see that there are two disjoined regions in the parameter space for which FSF singularities are allowed in the universe. $\\delta$ is always positive which is not surprising since only for positive $\\delta$ there can be an accelerated expansion in the investigated model. Characteristic feature of the picture is that there are two qualitatively different regions divided by the value of the parameter $m$. There are two branches, the first for $m>1$, and the second for $m<1$. In the following tables we present ranges of the values of the parameters for three confidence levels. For the case when $m>1$: \\begin{center} \\begin{tabular}{cccc }  {\\bf $1\\sigma$ CL}     &    {\\bf $2\\sigma$ CL}    &  {\\bf $3\\sigma$ CL}        \\\\\\hline $m\\in(1.00,1.14)$      &  $m\\in(1.00,1.21)$       & $m\\in(1.00,1.28)$          \\\\\\hline $\\delta > 1.00$        &  $\\delta > 1.00$         & $\\delta > 1.00     $     \\\\\\hline $n\\in(0.0,0.97)$       &  $n\\in(0.0,0.98)$        & $n\\in(0.0,0.99)$           \\\\\\hline $y_0\\in(0.0,0.50)$    &  $y_0\\in(0.00,0.66)$     & $y_0\\in(0.00,0.76)$        \\\\  \\end{tabular} \\end{center}  Below the case of $m<1$:   \\begin{center} \\begin{tabular}{ccc}  {\\bf $1\\sigma$ CL}     &    {\\bf $2\\sigma$ CL}    &  {\\bf $3\\sigma$ CL}        \\\\\\hline $m\\in(0.41,0.66)$     &  $m\\in(0.39,0.67)$       & $m\\in(0.39,0.72)$          \\\\\\hline $\\delta\\in(0.52,0.99)$ &  $\\delta\\in(0.36,0.99)$  & $\\delta\\in(0.26,0.99)$     \\\\\\hline $n\\in(0.0,0.85)$       &  $n\\in(0.0,0.91)$        & $n\\in(0.0,0.94)$           \\\\\\hline $y_0\\in(0.45,0.88)$    &  $y_0\\in(0.43,0.94)$     & $y_0\\in(0.41,0.97)$        \\\\  \\end{tabular} \\end{center} \\newpage For the case of the branch $m>1$ the maximum value at $3\\sigma \\ y_0$ parameter is $0.76$, which means that such a singularity can happen later than for the second branch $m<1$. While $y_0$ approaches $0.76$, $\\delta$ grows stronger and as the value of $\\delta$ more grows, $m$ becomes better constrained.  What seems to be important is that there is an allowed value of $m=2/3$ within $1\\sigma$ CL, which could correspond to the dust filled Einstein-de-Sitter universe  in a close to big-bang limiting case. What is also interesting for this branch is that while the allowed values of the nonstandarcity parameter $\\delta$ are small, i.e. $\\delta<1$, the parameter $y_0$ approaches unity which means that the singularity may happen in the nearest future for this case. For $1\\sigma$ CL, $y_0=0.885$ corresponds to $\\sim 2\\times 10^9$ years to the time of the singularity. For $3\\sigma$ CL, $y_0=0.97$ corresponds to the present time at $\\sim 0.37 \\times 10^9$ years before the time of the singularity.\\\\ \\indent In conclusion, we have shown that for a finite scale factor singularity there is an allowed value of $m = 2/3$ within $1\\sigma$ CL, which corresponds to the dust-filled Einstein-de-Sitter universe for the close to big-bang limiting case. The finite scale factor singularity may happen within $2\\times 10^9$ years in future for $1\\sigma$ CL and its prediction at the present moment of cosmic evolution cannot be distinguished, with current observational data, from the prediction given by the standard quintessence scenario of future evolution in the Concordance Model \\cite{chevallier00, linder02, linder05, koivisto05, caldwell07, zhang07, amendola07, diporto07, hu07b, linder09}."
    },
    "1112/1112.2784_arXiv.txt": {
        "abstract": "VO-KOREL is a  web service exploiting the technology of Virtual Observatory for providing the astronomers with the intuitive graphical front-end and distributed computing back-end running the most recent version of Fourier disentangling code KOREL.  The system integrates the ideas of the  e-shop basket, conserving the privacy of every user by transfer encryption and access authentication, with  features of laboratory notebook, allowing the easy housekeeping of both input parameters and final results,  as well as it explores a newly emerging technology of cloud computing.  While the web-based front-end allows the user to submit data and parameter files, edit parameters, manage a job list, resubmit or cancel running jobs and mainly watching the text and graphical results of a disentangling process, the main part of the back-end  is a simple job queue submission system executing in parallel multiple instances of  FORTRAN code KOREL. This  may be easily extended for GRID-based deployment on massively parallel computing clusters.  The short introduction into underlying technologies is given, briefly mentioning advantages as well as  bottlenecks of the  design used. ",
        "introduction": "The Virtual observatory (hereafter VO) is a global infrastructure of distributed astronomical archives and data processing services enabling the standardised discovery and access to the astronomical data worldwide as well as a large set of powerful tools for scientific  analysis and visualisation (\\cite[Solano 2006]{solano}). VO supports mainly the  multi-wavelength research or discovery of rare objects by cross-matching huge catalogues.  For interaction with the user as well as other computers the VO uses the modern technology of  Web Services (WS).  The WS is typically complex processing application using the web technology  to transfer input data  to the main processing back-end and the results (after intensive number crunching) back to user.  All this can be done using only an ordinary web browser (and in principle the science may be done on the fast palmtop or advanced mobile phone).  An example of WS is the e-shop or ticket reservation system. ",
        "conclusions": "The VO-KOREL service is not only giving comfortable environment for Fourier disentangling of spectra, but it is a test-bed of general cloud infrastructure for execution of most scientific computationally intensive codes, like models of stellar atmospheres or special processing of complex data sets."
    },
    "1112/1112.1672_arXiv.txt": {
        "abstract": "{} {A thorough multiwavelength, multiheight study of the vector magnetic field in a compact active region filament (NOAA 10781) on 2005 July 3rd and 5th is presented. We suggest an evolutionary scenario for this filament. } {Two different inversion codes were used to analyze the full Stokes vectors acquired with the Tenerife Infrared Polarimeter (TIP-II) in a spectral range which comprises the chromospheric \\ion{He}{i} 10830\\,\\AA\\ multiplet and the photospheric \\ion{Si}{i} 10827\\,\\AA\\ line. In addition, we used \\textit{SOHO}/MDI magnetograms as well as BBSO and \\textit{TRACE} images to study the evolution of the filament and its active region (AR). High resolution images of the Dutch Open Telescope were also used. } {An active region filament (that was formed before our observing run) was detected in the chromospheric Helium absorption images on July 3rd.  The chromospheric vector magnetic field in this portion of the filament was strongly sheared (parallel to the filament axis) whereas the photospheric field lines underneath had an inverse polarity configuration. From July 3rd to July 5th, an opening and closing of the polarities at either side of the polarity inversion line (PIL) was recorded, resembling the recently discovered process of the sliding door effect seen by Hinode. This is confirmed both with TIP-II and \\textit{SOHO}/MDI data. During this time, a newly created region that contained pores and orphan penumbrae at the PIL was observed. On July 5th, a normal polarity configuration was inferred from the chromospheric spectra, while strongly sheared field lines aligned with the PIL were found in the photosphere. In this same data set, the spine of the filament is also observed in a different portion of the FOV and is clearly mapped by the Silicon line core. } {The inferred vector magnetic fields of the filament suggest a flux rope topology. Furthermore, the observations indicate that the filament is divided in two parts, one which lies in the chromosphere and another one that stays trapped in the photosphere. Therefore, only the top of the helical structure is seen by the Helium lines. The pores and orphan penumbrae at the PIL appear to be the photospheric counterpart of the extremely low-lying filament. We suggest that orphan penumbrae are formed in very narrow PILs of compact ARs and are the photospheric manifestation of flux ropes in the photosphere.} ",
        "introduction": "Solar filaments, also called prominences when observed in emission outside the disk, are elongated structures formed of dense plasma which lies above polarity inversion lines (PILs) of the photospheric magnetic field \\citep{babcock55} and has a lower temperature than its surroundings. A functional definition of the PIL would be an imaginary line that separates two close areas of opposite polarity. On disk filaments are readily identifiable in the quiet Sun (quiescent filaments) and in active regions (active region filaments) when observed using common chromospheric wavelengths, e.g., the \\ion{He}{i} 10830\\,\\AA\\ multiplet and the H$\\alpha$ 6563\\,\\AA\\ line. The magnetic field plays a major role in their formation, stability and evolution. Hence, spectropolarimetric observations at different heights of the atmosphere combined with high angular resolution images are needed to obtain an overall picture of the physical processes which take place in filaments. The magnetic field strength and its orientation can be inferred using appropiate diagnostic techniques that are able to interpret the Zeeman effect as well as scattering polarization and its modification through the Hanle effect \\citep[][and references therein]{tandberg95}. Although many line-of-sight observations have been carried out in the past decades, it is important to emphasize the need for full-Stokes measurements in order to have complete information of the vector magnetic field.  According to observational studies, the magnetic field topology under active region (AR) filaments has sheared or twisted field lines along the PIL which can support dense plasma in magnetic dips \\citep{lites05, lopez06}. Previous findings about prominence magnetic structure have shown models with dipped field lines in a normal \\citep{kippen57} or inverse\\footnote{Throughout this paper, by {\\em inverse configuration} we mean a magnetic field vector, perpendicular to the filament's axis, that points from the negative to the positive polarity. This is the opposite to what one would expect in a potential field solution - where the field points from the positive to the negative side (and is referred to as normal configuration).} \\citep{kuperus74, pneuman83} polarity configuration, the latter having a helical structure. From a large sample of quiescent prominence observations, \\citet{leroy84} found both types of configurations and a strong correlation between the magnetic field topology and the height of the prominence. On the other hand, \\citet{bommier94} found predominantly inverse polarity configurations. However, more recent photospheric spectropolarimetric observations of AR filaments have revealed that both configurations can coexist in the same filament \\citep{guo10} or even evolve with time from one type to the other, as presented by \\citet{okamoto08} from the analysis of vector magnetogram sequences.  The formation process of filaments is still a controversial issue in solar physics and has been widely debated. On the one hand, there is the \\emph{sheared arcade model} which, by large scale photospheric footpoint motions (such as shear at, and convergence towards the PIL, and subsequent reconnection processes) is able to form dips, and even helical structures, where plasma can be gravitationally confined \\citep[e.g.,][]{pneuman83,balle89, antiochos94,aulanier98, devore00,martens01,welsch05,karpen07}. This model is capable of reproducing both inverse and normal polarity configurations in the same filament \\citep{aulanier02}. However, recent observational works in a quiescent filament \\citep{hindman06} and along an AR filament channel \\citep{lites10} did not find evidence for these systematic photospheric flows that converged at the PIL and triggered reconnection processes. In our opinion, more observations are needed to support this important result.  On the other hand, the \\emph{flux rope emergence model} assumes that the twist in the field lines is already present in the convection zone before emerging into the atmosphere \\citep[eg.,][]{kurokawa87, low94, low95, lites05}. The rise of twisted magnetic fields has been studied by various authors through observations \\citep{lites95, leka96, lites10} as well as three-dimensional (3D) numerical simulations \\citep[e.g.,][]{fan01,archontis04,magara04,sykora08,fan09,yelles09}. Such an emerging helical flux rope scenario, combined with the presence of granular flows, is suggested to be the main process of formation and maintenance of the AR filament studied by \\citet{okamoto08, okamoto09}. Recent nonlinear force-free field (NLFFF) extrapolations of photospheric magnetic fields underneath AR filaments \\citep{guo10, canou10, jing10} agreed that the main structure has to be a flux rope, although \\citet{guo10} also found dipped arcade field lines in the same filament. Flux rope emergence from below the photosphere encounters severe difficulties to reach chromospheric layers and above, as it is loaded with mass \\citep[e.g.,][]{archontis08}. Indeed, the previous work suggests that a second flux rope formed from instabilities and atmospheric reconnection is what lifts up the mass and forms the observed active region filaments.  The main difference between the sheared arcade and the emerging flux rope models is the existence of a flux concentration stuck at photospheric levels in the latter scenario. The reader is referred to the paper of \\citet{mackay10} for a recent review on the magnetic structure of filaments.  In the past, only a few measurements of the magnetic field strength in AR filaments have been done. For instance, \\citet{kuckein09}, using the same data set that is described in this paper, found a predominance of Zeeman-like signatures in the Stokes profiles. Using three different methods, they inferred very strong magnetic fields in the filament (600--700\\,Gauss), with a dominant transverse component (with respect to the line-of-sight).  The aim of this paper is to study the strength and topology of the magnetic field in an active region filament at photospheric and chromospheric heights simultaneously. Recently, several studies have presented analyses of the vector magnetic field in filaments or prominences from observations either in the photosphere \\citep[e.g.,][]{lites05,lopez06,okamoto08,guo10,lites10} or the chromosphere \\citep[e.g.,][]{lin98,casini03,merenda06,kuckein09}, but non of them have inferred the field at both heights at the same time. The 10830\\,\\AA\\ spectral region, which includes a chromospheric \\ion{He}{i} triplet and a photospheric \\ion{Si}{i} line, offers a unique spectral window to understand the physical processes that take place in AR filaments as already shown by \\citet{sasso11}. In this work, we focus on the overall magnetic configuration observed simultaneously in the photosphere and the chromosphere. ",
        "conclusions": "The main conclusions of the present work are:  \\begin{enumerate} \\item The ``sliding doors'' effect described by Okamoto et al. (2008) was seen in our observations during the period between July 3rd and July 5th. This has been identified in \\textit{SOHO}/MDI data as well as in the TIP-II \\ion{Si}{i} magnetograms for those days.  \\item The observed AR filament can be separated into two areas. The first one is well observed on July 3rd and shows the filament axis, or spine, in the Helium absorption images. The \\ion{He}{i} data show the magnetic field (with horizontal field strengths of 400--500\\,G) aligned with the filament, indicating a strong sheared configuration. This region is also observed in the FOV of July 5th, when portions of the spine are even seen in the Silicon line core images. This would indicate that the part of the spine observed this day lies in lower layers than the portion observed two days earlier.  \\item The photospheric fields for this area of the filament show an inverse configuration of the magnetic field lines that we interpret to be the bottom of a flux rope structure. Similarly, we interpret the spine seen in Helium as the signature of the flux rope axis.  \\item The second area corresponds to the orphan penumbrae that appear during the sliding door timeframe and clearly shows a different magnetic configuration. Helium observations exhibit a normal polarity configuration whereas the Silicon data show a strongly sheared region with very intense horizontal fields. This magnetic topology is interpreted as a scenario in which the chromospheric magnetic fields trace the top part of the flux rope, while the spine -or flux rope axis- is still lying in the photosphere.  \\item The observed orphan penumbrae is thus the photospheric counterpart of this active region filament. Since the spine is also seen in the Silicon line, we emphasize that active region filaments can have a clear photospheric signature. In particular, it would be interesting to find out if all orphan penumbral regions observed inside active regions are related to their filaments.  \\item The whole emergence process of these orphan penumbrae does not leave a clear imprint in the flux history of the active region. This is interpreted as due to the flux system being mostly transverse, with a low enough twist to not generate identifiable signals in \\textit{SOHO}/MDI longitudinal magnetograms.  \\item We do not observe flux loops with a normal polarity configuration in the photosphere, as suggested by some filament generation models based on footpoint motions and reconnection \\citep[see][]{balle89, vanballe07}. While the preexisting filament observed on July 1st in the BBSO data could have been created in the way described in these works, the configuration and the evolution described here for the period from July 3rd to July 5th suggests otherwise.  \\end{enumerate}  The current study was limited, not only by the small FOV of TIP-II, which nowadays has a slit twice as long as during our observing campaign in 2005, but also by the observational gap on July 4th. Magnetic field extrapolations could undoubtedly shed more light on clarifying the magnetic structure of this filament. It is now crucial to carry out more multiwavelength measurements, as the one presented in here, with higher cadence and bigger FOVs to fit the pieces of the puzzle together, in order to fully understand the origin, evolution and magnetic topology of AR filaments. In particular, continuous vector magnetograms of active regions together with simultaneous imaging of the corona should be able to prove/disprove the proposed scenario for the last stages of their evolution. The instrument suite on board the NASA/SDO satellite is the best candidate for such an study."
    },
    "1112/1112.6052_arXiv.txt": {
        "abstract": "We investigate the constraints on the scalar, vector and spin-3/2 dark matter interaction with the standard model particles, from the observations of dark matter relic density, the direct detection experiments of CDMS and XENON, and the indirect detection of the $\\bar p/p$ ratio by PAMELA. A model independent way is adopted by constructing general 4-particle operators up to dimension 6 for the effective interaction between dark matter and standard model particles. We find that the constraints from different experiments are complementary with each other. Comparison among these constraints may exclude some effective models of dark matter and limit some parameters of others. The spin-independent direct detection gives strong constraints for some operators, while the indirect detection of $\\bar p/p$ data can be more sensitive than direct detection or relic density for light dark matter (whose mass $\\lesssim 70$~GeV) in some cases. The constraints on some operators for spin-3/2 dark matter are shown to be similar to those on their analogous operators for Dirac fermionic dark matter. There are still some operators not sensitive to the current dark matter direct and indirect search experiments. ",
        "introduction": "It is by now established~\\cite{Tegmark:2006az,Komatsu:2008hk,Komatsu:2010fb} that about 23\\% of the constituents of the Universe are composed of dark matter (DM)~\\cite{kolb-turner,Jungman:1995df,Bertone:2004pz, Murayama:2007ek,Feng:2010gw}. However, its nature remains unclear. A well-motivated candidate for DM is the weakly interacting massive particle (WIMP), which must be stable, nonrelativistic, electrically neutral and colorless. If the WIMP mass is from a few GeV to TeV while their interaction strength is of the weak scale, they can naturally yield the observed relic density of DM~\\cite{Feng:2010gw}. Although there is no stable WIMP in the Standard Model (SM), WIMP candidates exist in various theoretical models trying to solve the SM problems at the weak scale, such as supersymmetric models \\cite{Jungman:1995df,Goldberg:1983nd,Ellis:1983ew,Arnowitt:1992aq,Nath:1992ty, Kane:1993td}, extra dimensional models~\\cite{Kolb:1983fm,Cheng:2002ej, Hooper:2007qk,Servant:2002aq,Servant:2002hb,Agashe:2004ci, Agashe:2004bm,Agashe:2007jb}, little Higgs models \\cite{Cheng:2004yc,Low:2004xc,Birkedal:2006fz,Freitas:2009jq,Kim:2009dr}, left-right symmetric models~\\cite{Dolle:2007ce,Guo:2008hy,Guo:2008si, Guo:2010vy}, and some other models (e.g.~\\cite{Khlopov:2008ty,Li:2010rz, Cui:2011wk,Kanemura:2011mw}).  The specific models mentioned above are very attractive, but still lack experimental support. The well-running LHC experiment may find some important signals of these models in the near future. However, if other new particle species are all so heavy that the DM particle is the only new particle within the reach of LHC, it will be very difficult to know which model the DM particle belongs to. In addition, it is possible that the DM particle may be first observed by direct or indirect detection experiments. These early observations may only provide information about some general properties of the DM particle, and may not be able to distinguish the underlying theories. Therefore, the model-independent studies of the DM phenomenology can play an important role as they may avoid theoretical bias \\cite{Birkedal:2004xn,Giuliani:2004uk,Kurylov:2003ra,Beltran:2008xg, Cirelli:2008pk,Shepherd:2009sa}. Recently there have been quite a few papers to study various phenomenologies related with DM in the model-independent way~\\cite{Cao:2009uv,Cao:2009uw, Beltran:2010ww,Fitzpatrick:2010em,Goodman:2010yf,Bai:2010hh, Goodman:2010ku,Goodman:2010qn,Bell:2010ei,Zheng:2010js,Cheung:2010ua, Cheung:2011nt,Mambrini:2011pw,Rajaraman:2011wf,Fox:2011pm,Goodman:2011jq, Kamenik:2011vy,Shoemaker:2011vi}. Especially the relic density measured by WMAP~\\cite{Komatsu:2010fb}, direct detection from CDMS~\\cite{Ahmed:2009zw,Akerib:2005za} and XENON \\cite{Aprile:2011hi,Angle:2008we}, and possible collider signals from Tevatron (e.g.~\\cite{Aaltonen:2008hh,CDF:monojet}) and from LHC (e.g. \\cite{Chatrchyan:2011nd,ATLAS:monojet}) are considered in these studies.  In our previous work~\\cite{Zheng:2010js}, we investigated a general set of 4-fermion operators for the effective interaction between the spin-1/2 fermionic DM and the SM particles, and gave the phenomenological constraints from the observed DM relic density, the direct detection experiments by CDMS and XENON, and the indirect detection of the $\\bar{p}/p$ ratio by PAMELA~\\cite{Adriani:2010rc}. It was found that the constraints from different observations are quite complementary. Besides the possibility of spin-1/2 WIMPs, it is also possible that DM is composed of scalar, vector, or spin-3/2 WIMPs, which belong to different representations of Lorentz group. In these cases, the forms of possible effective operators are different. These differences may lead to distinguishable phenomenological results. In this work we extend our previous analysis to the cases of scalar, vector and spin-3/2 DM. We will consider a general set of 4-particle operators up to dimension 6 for the effective interaction between the WIMPs and the SM fermions and compute their phenomenological constraints. We will use the updated limit of XENON100 SI direct detection \\cite{Aprile:2011hi}, which is stronger than that adopted in our previous work~\\cite{Zheng:2010js}.  This paper is organized as follows. In Sections \\ref{sec-scal} and \\ref{sec-vect}, the effective models of scalar and vector DM are discussed, respectively. In the subsections of these two sections, we explore the constraints on these models from the DM relic density, direct and indirect detection searches and the validity of effective theory, and then present the combined constraints on the effective coupling constants of these models. In Section \\ref{sec-3/2}, the study on the effective models of spin-3/2 DM are carried out briefly. The conclusions are given in Section \\ref{sec-con}. ",
        "conclusions": "}  In this work, we give a general analysis of the 4-particle interaction between SM particles and DM which consists of scalar, vector or spin-3/2 WIMPs. The most general forms of the 4-particle operators up to dimension 6 have been considered. We find that for scalar, vector and spin-3/2 DM, the constraints from DM relic density, DM direct and indirect detection are complementary to each other. Thus the comparison among different kinds of experimental results gives us a complete picture about the current DM searches. In general, the constraints from SI direct detection are the most stringent, while those from SD direct detection are quite weak. On the other hand, for light DM (whose mass $\\lesssim 70$~GeV) the cosmic-ray $\\bar p/p$ data can be more sensitive than direct detection or relic density in some cases.  Assuming one operator dominates the effective interaction between DM and SM fermions, we find that in some cases the constraints are so strong that the Universe will be overclosed by DM thermal production. If the standard cosmology is still retained, the DM models in such cases should be excluded, which are indicated in Tables \\ref{tab:scalar_sum}, \\ref{tab:vector_sum} and \\ref{tab:3/2_sum}. In the case of scalar DM, recent direct detection experiments exclude some $M_\\phi$ regions for S and V interactions, while the PAMELA $\\bar p/p$ ratio excludes some small $M_\\phi$ regions ($\\lesssim 70$~GeV) for S and SP interactions. In the case of vector DM, recent direct detection experiments exclude some $M_X$ regions for S and V interactions, and for $\\widetilde{\\mathrm{VA}}$ interaction only with universal couplings. The PAMELA $\\bar p/p$ ratio, however, excludes some small $M_X$ regions ($\\lesssim 70$~GeV) for S, T, SP and $\\widetilde{\\mathrm{T}}$ interactions, and most of these interactions cannot be excluded by direct detection. In the case of spin-3/2 DM, the constraints on most interactions (S, P, V, A, T1, T4, SP, PS, VA and AV) are much similar to those on their analogous interactions of Dirac fermionic DM in our previous work \\cite{Zheng:2010js}. Besides, the constraints on T2 and T3 interactions are so weak that only direct detection can exclude a little region of $M_\\chi$ in the case of T2 interaction with universal couplings. For T5 and T6 interactions, some small $M_\\chi$ regions ($\\lesssim 70$~GeV) are excluded by the PAMELA $\\bar p/p$ ratio. Among the scalar, vector and spin-3/2 DM effective models, there are still some effective interactions to which the recent DM direct and indirect search experiments are not sensitive at all."
    },
    "1112/1112.4866_arXiv.txt": {
        "abstract": "We have recently proposed a special class of scalar tensor theories known as {\\it the Fab Four}.  These arose from attempts to  analyse the  cosmological constant problem  within the context of Horndeski's most general scalar tensor theory.  {\\it The Fab Four} together give rise to a model of self-tuning,  with the relevant solutions evading Weinberg's no-go theorem by relaxing the condition of Poincar\\'e invariance in the scalar sector.  {\\it The Fab Four}  are made up of four geometric terms in the action with each term containing a free potential function of the scalar field.  In this paper we rigorously derive this model from the general model of Horndeski, proving that {\\it the Fab Four} represents  the {\\it only} classical scalar tensor theory of this type that has any hope of tackling the cosmological constant problem. We present the full equations of motion for this theory, and give an heuristic argument to suggest that one might be able to keep radiative corrections under control. We also give  {\\it the  Fab Four} in terms of the potentials presented in  Deffayet {\\it et al}'s version of Horndeski. ",
        "introduction": "The cosmological constant problem has been % described as the most embarrassing fine-tuning problem in Physics today. According to our current understanding of particle physics and effective {quantum} field theory, the vacuum receives zero point energy contributions from each particle species right up to the {UV} cut-off, which may be as high as the Planck scale. The trouble is that in General Relativity, any matter, including vacuum energy, gravitates and  the only way to make it compatible with observation is to demand  considerable fine-tuning between the vacuum energy and the bare cosmological constant.  The situation is exacerbated by phase transitions in the early universe that can give rise to constant shifts in the vacuum energy contribution.  To date, particle physicists have failed to come up with a satisfactory solution to this problem, so some recent attempts have instead focussed on gravitational physics. This  alternative approach requires a non-trivial  modification  of Einstein's theory at large distances (see \\cite{review} for a detailed review of modified gravity).    One particularly interesting direction involves scalar-tensor theories of gravity. It seems sensible to require that any theory maintains second order field equations in order to avoid an Ostrogradski instability \\cite{ostro}, and the most general scalar-tensor theory satisfying that criteria  in four dimensions was written down back in 1974 by Horndeski  \\cite{horndeski:1974} (it has recently been rediscovered independently in \\cite{general}). Such theories of modified gravity cover a wide range of models, ranging from Brans-Dicke gravity \\cite{bdgravity}  to the recent models \\cite{covgal,galmodels} inspired by galileon theory \\cite{galileon}. Galileon models are examples of  higher order scalar tensor Lagrangians with second order field equations, and, as a result, they are closely related to Kaluza-Klein compactifications of higher dimensional Lovelock theories \\cite{kkl, VanAcoleyen:2011mj}.  Of course all of these scalar-tensor models  can be considered as special cases of Horndeski's original action.    In \\cite{Charmousis:2011bf} we obtained a new class of solutions arising out of Horndeski's theory on FLRW backgrounds. The new solutions gave a viable self-tuning mechanism for solving the (old) cosmological constant problem, at least at the classical level, by completely screening the spacetime curvature from the net cosmological constant. This would seem to be in violation of Weinberg's famous no-go theorem \\cite{nogo} that forbids precisely this kind of self-adjustment mechanism. However, Weinberg assumes Poincar\\'e invariance to hold universally across all fields whereas we allow it to be broken in the scalar field sector.  In other words, we continue to require  Poincar\\'e invariance at the level of   spacetime curvature, but not  at the level of the self-adjusting scalar field.   A similar approach was adopted in the context of bigalileon theory  \\cite{bigalileon} where only a small vacuum energy could be successfully screened away. In \\cite{Charmousis:2011bf}, we provided a brief sketch of how the system works for  scalar tensor theories where matter is only minimally coupled to the metric (required to ensure compatibility with Einstein's Equivalence Principle (EEP)). By demanding the presence of a viable  self-tuning mechanism we were able to place powerful restrictions on the allowed form of Horndeski's original Lagrangian. Whereas the  original model is complicated, with many arbitrary functions of both the scalar and its derivatives, we showed that  once the model is passed through our self-tuning filter (to be defined shortly), it reduces in form to just four base Lagrangians each depending on an arbitrary function of the scalar only, coupled to a curvature term. We called these base Lagrangians {\\it the Fab Four}: John, Paul, George and Ringo.  Together, {\\it the Fab Four} make up the most general scalar-tensor theory capable of self-tuning. Individually they are given by the following \\begin{eqnarray} \\label{eq:john} {\\cal L}_{john} &=& \\sqrt{-g} V_{john}(\\phi)G^{\\mu\\nu} \\nabla_\\mu\\phi \\nabla_\\nu \\phi \\\\ \\label{eq:paul} {\\cal L}_{paul} &=&\\sqrt{-g}V_{paul}(\\phi)   P^{\\mu\\nu\\alpha \\beta} \\nabla_\\mu \\phi \\nabla_\\alpha \\phi \\nabla_\\nu \\nabla_\\beta \\phi \\\\ \\label{eq:george} {\\cal L}_{george} &=&\\sqrt{-g}V_{george}(\\phi) R \\\\ \\label{eq:ringo} {\\cal L}_{ringo} &=& \\sqrt{-g}V_{ringo}(\\phi) \\GB \\end{eqnarray} where $R$ is the Ricci scalar, $G_{\\mu\\nu}$ is the Einstein tensor, $P_{\\mu\\nu\\alpha \\beta}$ is the double dual of the Riemann tensor \\cite{mtw}, $\\GB=R^{\\mu\\nu \\alpha \\beta} R_{\\mu\\nu \\alpha \\beta}-4R^{\\mu\\nu}R_{\\mu\\nu}+R^2$ is the Gauss-Bonnet combination, and in what follows the Greek indices $\\mu,\\nu =0..3$. The purpose of this paper is  to rigorously derive the conditions that lead to these four base Lagrangians, showing how they naturally lead to self-tuning solutions, provided that $\\{V_{john}, V_{paul}, V_{george}\\} \\neq \\{0,0, constant\\}$. Note that this constraint means that GR is {\\it not} a {\\it Fab Four} theory, consistent with the fact that it does not have self-tuning solutions.  To be clear as to what is meant by ``self-tuning\", let us define our self-tuning filter. We require that \\begin{itemize} \\item the theory should admit a Minkowski vacuum\\footnote{For simplicity throughout the introductory part of the text we have simply written, ``Minkowski vacuum\" to stand for``a patch of Minkowski vacuum\". This technical issue will be made clear later on in section \\ref{self-tune}} for any value of the net cosmological constant \\item this should remain true before and after any phase transition where the cosmological constant jumps instantaneously by a finite amount. \\item the theory should permit a non-trivial cosmology \\end{itemize} The last condition ensures that Minkowski space is not the only cosmological solution available, something that is certainly required by observation. The idea is that the cosmological field equations should be dynamical, with the Minkowski solution corresponding to some sort of fixed point. In other words, once we are on a Minkowski solution, we stay there -- otherwise we evolve to it dynamically. This last statement would indicate that the self-tuning vacuum is an attractive fixed point. We do not prove this here, but in our companion paper on cosmology \\cite{fab4cos} we will see plenty of examples where it is indeed the case.  The first two conditions are the basic requirements of any successful self-tuning mechanism. There are many examples in the literature which pass the first condition, but fall down at the second. This includes the much explored co-dimension two braneworld models in which the compact extra dimensions are shaped like a rugby ball \\cite{Carroll}. The  brane tension controls  the deficit angle, while the brane geometry  is completely determined by the bulk cosmological constant and the magnetic flux. Therefore, this passes our first condition.  However, when the brane tension changes after a phase transition it affects the brane curvature via the backdoor, by altering magnetic flux and the  theory falls foul of our second condition \\cite{cline}.  It is interesting to note that any diffeomorphism invariant theory that passes both the first and second condition will admit a Minkowski solution in the presence of {\\it any} cosmological fluid, not just a cosmological constant.  The point is that our vacuum energy density corresponds to a piecewise constant function, with discontinuities at the phase transitions. In principle these transitions can occur at any given time, so a Minkowski solution can be returned for all piecewise constant energy densities. The energy density of  an arbitrary cosmological fluid can be well approximated by a piecewise constant function, and so it follows that  it must also admit a  Minkowski solution. Like we said, this property must hold for {\\it any} diffeomorphism invariant  theory passing our first two conditions, and not just {\\it the Fab Four}.  We might worry that this prevents any hope of a sensible matter dominated cosmology. However, this is where the third condition comes into play, and we once again refer the reader to our companion paper \\cite{fab4cos} for evidence that sensible cosmologies are indeed possible within this theory.  Even so, the main aim of this paper is not to extoll the virtues of  {\\it the Fab Four} but to push a very general class of modified gravity theories through our self-tuning filter and to see what happens. In a sense we are testing the scope of Weinberg's theorem, relaxing one of his assumptions and seeing how far we can go.  It turns  out that our filter is very efficient -- it removes most of Horndeski's original theory-- but it is not $100\\%$ efficient. We are left with  {\\it the Fab Four}.     The layout of the paper is as follows: in section \\ref{horndeski-action} we present the original action of Horndeski  \\cite{horndeski:1974}, minimally coupled to matter, and derive the Hamiltonian and scalar field equations of motion for the system. In section \\ref{self-tune} we demonstrate how a self-tuning solution can in principle be obtained by relaxing Weinberg's no-go theorem to allow the scalar field to evolve in time. This is followed in section \\ref{self-tune-horndeski} with a derivation of the self-tuned Horndeski action, where we show how the initial complicated Lagrangian reduces to four simple terms each one being an arbitrary function of the scalar field alone coupled to a curvature term. Of particular note is that any dependence on the kinetic energy of the scalar field drops out.  In section \\ref{conc} we bring everything together and discuss further demands we may wish to make on our theory, over and above our original filter, ranging from cosmological and solar system tests, to issues of stability. We also elucidate the elegant geometrical structure possessed by {\\it the Fab Four} and present their equations of motion in full.  We have a number of appendices, most of which are technical additions to the main text. The exceptions are appendices \\ref{app:dgsz} and \\ref{radiative}. In  appendix \\ref{app:dgsz} we present {\\it the Fab Four} in the language of the  potentials of  Deffayet {\\it et al}'s version of Horndeski \\cite{general}.  In appendix \\ref{radiative} we discuss the issue of radiative corrections to {\\it the Fab Four}. This is an important question, because radiative corrections are at the heart of the cosmological constant problem.    We do not attempt a detailed analysis -- that is certainly beyond the scope of the current paper --- but we do perform some heuristic calculations. It seems that radiative corrections can be kept under control given some not too restrictive conditions. ",
        "conclusions": ""
    },
    "1112/1112.2501_arXiv.txt": {
        "abstract": "Over the past decade advancements in the understanding of several astrophysical phenomena have allowed us to infer a concordance cosmological model that successfully accounts for most of the observations of our universe. This has opened up the way to studies that aim to better determine the constants of the model and confront its predictions with those of competing scenarios. Here, we use strong gravitational lenses as cosmological probes. Strong lensing, as opposed to weak lensing, produces multiple images of a single source. Extracting cosmologically relevant information requires accurate modeling of the lens mass distribution, the latter being a galaxy or a cluster. To this purpose a variety of models are available, but it is hard to distinguish between them, as the choice is mostly guided by the quality of the fit to the data without accounting for the number of additional parameters introduced. However, this is a model selection problem rather than one of parameter fitting that we address in the Bayesian framework. Using simple test cases, we show that the assumption of more complicate lens models may not be justified given the level of accuracy of the data. ",
        "introduction": "Over the past years, our understanding of the universe has greatly improved. The concordance model explains most of the cosmological observations. We have now entered a phase where finding new observational ways of measuring the constants of this model as well as confronting predictions with those of competing scenarios is crucial to further advance the research in cosmology. Strongly lensed quasars constitute such a cosmological probe.  In a strong gravitational lens, each image is the result of a different light-path. As a result, if the source behind the lens has a variable luminosity, as quasars do, it will manifest with a time delay between the two images.  This time delay $\\Delta t$ depends on the gravitational potential of the lens, and the underlying cosmological model.  The time delay between two images A and B is: \\begin{equation} \\Delta t_{A,B} = (1 + z_l) \\frac{d_l d_s}{d_{ls}} \\left( \\frac{1}{2}((\\theta_A - \\beta)^2 -(\\theta_B -\\beta)^2) + \\psi(\\theta_A) - \\psi(\\theta_B) \\right) \\end{equation} where $\\Delta t_{A,B}$, $z_l$, $\\theta_A$ and $\\theta_B$ are observables, $\\beta$, $\\psi(\\theta_A)$ and $\\psi(\\theta_B)$ depend on the lens model and $d_l$, $d_s$ and $d_{ls}$ depend on cosmology.  Using the above relation, we can derive constraints on cosmological parameters, provided we assume a lens model. Time delays are particularly sensitive to the value of the Hubble constant $H_0$.  Unfortunately, a change in the lens model can shift the inferred value of $H_0$ by a factor of two. Hence, the modeling of the lens, as well as a robust discrimination between lens models, is  critical to the study of time delays.  Here, we first discuss the lens models used in our study, then describe our methodology based on Bayesian statistical analysis, and finally present our results. ",
        "conclusions": "Bayesian techniques are a good way to discriminate between lens models and allow us to decide which double lenses can be accurately modeled by a simple power-law model. With the result from this analysis, we now have a good sample of double lenses, together with lens models, to determine cosmological parameters more accurately. %"
    },
    "1112/1112.2737_arXiv.txt": {
        "abstract": "We present absorption line indices measured in the integrated spectra of globular clusters both from the Galaxy and from M~31.  Our samples include 41 Galactic globular clusters, and more than 300 clusters in M~31.  The conversion of instrumental equivalent widths into the Lick system is described, and zero-point uncertainties are provided. Comparison of line indices of old M~31 clusters and Galactic globular clusters suggests an absence of important differences in chemical composition between the two cluster systems.  In particular, CN indices in the spectra of M~31 and Galactic clusters are essentially consistent with each other, in disagreement with several previous works.  We reanalyze some of the previous data, and conclude that reported CN differences between M~31 and Galactic clusters were mostly due to data calibration uncertainties.  Our data support the conclusion that the chemical compositions of Milky Way and M~31 globular clusters are not substantially different, and that there is no need to resort to enhanced nitrogen abundances to account for the optical spectra of M~31 globular clusters. ",
        "introduction": "The history of star formation and chemical enrichment of galaxies is encoded in the ages and chemical compositions of their stellar populations.  In particular, powerful insights on the processes leading to the assembly of the Galactic halo are gained by studies of the chemical abundances of their constituent populations of field stars and globular clusters (GCs).  It is only natural to extend such studies to the nearest giant spiral galaxy, M~31.  While detailed abundances of individual stars in the Galactic halo, disk, and bulge, have been obtained and extensively analyzed in the past several decades, similar studies of comparable samples of M~31 halo stars await the advent of 20-30~m class telescopes, equipped with efficient high-resolution spectrographs.  In the meantime, integrated light studies of M~31 GCs have historically provided valuable qualitative information about the chemical composition of the halo of M~31, leading to important clues on its formation history.  As stellar population synthesis models and instruments both grow in sophistication, quantitative information on elemental abundances from integrated-light studies of M~31 GCs is also becoming available \\citep[e.g.,][]{co09,ca10}.  The history of this field has been punctuated by heroic observational efforts, based on optical spectra painstakingly collected with, according to today's standards, relatively small telescopes, rather inefficient spectrographs, and low-quantum efficiency (and often difficult to calibrate) detectors \\citep[e.g.,][]{bu84,bh90}.  The main results emerging from these efforts are: 1) M~31 GCs are on average slightly more metal rich than their Galactic counterparts, while spanning roughly the same range of metallicities \\citep{bh91}; 2) M~31 GCs are $\\alpha$-enhanced, just like those in the MW \\citep[e.g.,][]{pu05}; 3) M~31 GCs are CN-enhanced compared to MW GCs \\citep{bu84,bh90,be04,pu05}.  In particular, \\cite{bu84} found that the CN band at $\\lambda$ 4170 ${\\rm\\AA}$ is stronger by $\\sim$ 0.05 mag in the spectra of moderately metal-poor M31 GCs than in those of their MW counterparts of same metallicity; and 4) M~31 GCs are possibly younger/older than those in the MW.  \\cite{bu84} found that $H\\beta$ is stronger in the spectra of M~31 GCs than in those of their MW counterparts by about 0.5 ${\\rm\\AA}$, indicating a younger age or differences in horizontal branch morphology.  The latter two results are contradicted by the findings presented in this paper.  While \\cite{ca10} discuss the Balmer line strengths in the spectra of the two GC systems, comparing their ages, CN is the main focus of this paper.  Those significant conclusions may teach us about important aspects of the formation of the two galaxies.  For instance, the comparable mean metallicities may be indicative of similar overall chemical enrichment, suggesting that the early star formation efficiency has been similar in the two systems---or perhaps in the sub-components that eventually assembled to form them.  The similar $\\alpha$-enhancement, assessed by measurements of Mg and Fe-sensitive absorption lines, suggests that either the time scale for star formation, the IMF, or a combination thereof, were similar in both galaxies.  The difference in CN strength, which has been ascribed to a difference in nitrogen abundance, is difficult to interpret, owing mostly to uncertainties in the models for the nucleosynthesis of that element. The issue is further complicated by the fact that CNO elements are seen to present strong star-to-star variations in Galactic GCs \\citep[e.g.,][and references therein]{gr04}, which may be associated with the presence of multiple stellar populations in those GCs \\citep[][and references therein]{cs10,pi09,ms09,can98}, and by inference, in their M~31 counterparts.  Regardless, any scenario for the formation of the MW and M~31 haloes, and their GC systems, will be challenged by the large nitrogen abundance differences between systems that look otherwise very similar.    This is the fourth of a series of papers dedicated to analyzing the kinematics and chemistry of a large sample of M~31 GCs, based on high-quality integrated spectra for several hundred M~31 clusters, obtained with MMT/Hectospec.  In Paper I \\citep{ca09a} we characterized the population of young ($\\simless$ 2 Gyr) M~31 clusters in terms of their ages, metallicities, masses, and kinematics.  In Paper II \\citep{ca10}, the ages and metallicities of M~31 old GCs were studied.  In particular, we found no differences between the ages of the two old GC systems, in disagreement with previous claims. Moreover, we found that the M~31 GC system does not have a bimodal metallicity distribution, in agreement with recent findings \\citep{yo06}.  In Paper III \\citep{mo10}, we suggest that the old bulge GCs in M~31 are characterized by a bar-like kinematics.  In this paper, we present absorption line indices measured in integrated spectra of a large sample of M~31 and MW GCs. The data for M~31 GCs come from spectra described by \\cite{ca09a}, and those for MW GCs are based on the integrated GC spectra from \\cite{s05}.  While the focus of this paper is on the comparison of CN strengths between the two GC systems, \\cite{s11} will present an analysis of abundance ratios, based on application of stellar population synthesis models to these index measurements. This paper is organized as follows: in Section~\\ref{data}, the measurement of Lick absorption line indices is described, and a comparison between CN indices in M~31 and MW GCs is presented in Section~\\ref{analysis}.  Our conclusions are summarized in Section~\\ref{epilogue}. ",
        "conclusions": "\\label{epilogue}  We present absorption line index measurements taken in integrated spectra of a large number of M~31 and MW GCs, from \\cite{ca09a} and \\cite{s05}, respectively.  We discuss the conversion of instrumental measurements to a common equivalent system \\citep[defined in][]{s07}, as well as the uncertainties in the zero points of these systems.  By comparing measurements of CN indices in old GCs belonging to both data sets, we conclude that, in disagreement with previous work by \\cite{bu84}, \\cite{bh90}, \\cite{da90}, \\cite{bu04}, \\cite{be04}, and \\cite{pu05}, among others, M~31 and MW GCs of the same [Fe/H] {\\it generally} have similar CN-band strengths.  Because this result disagrees with conclusions by many different groups, we have reanalyzed the data by some of these different authors and suggest that in most cases their conclusions were a result of calibration problems.  In particular, the blue Lick CN indices are relatively weak and are very sensitive to flux calibration uncertainties.  We also find that the two MW GCs with [Fe/H] $\\simgreater -0.4$, NGC~6528 and 6553, have weaker CN bands than their M~31 counterparts at the same [Fe/H].  It is not entirely clear whether this latter difference is real or a by-product of sky-subtraction errors, or of issues with the sampling of the GCs' brightest stars.  Other data sets \\citep[e.g.,][]{pu05} do not show similar CN discrepancies between these GCs and their M~31 metal-rich counterparts.  Even if our results are correct, and these two GCs indeed have weaker CN indices than metal-rich M~31 GCs, the fact that they seem to depart significantly from a $\\langle$Fe$\\rangle$--$CN_1$ trend that is present in {\\it both} samples suggests that NGC~6528 and 6553 are the exception rather than the rule.   The result by \\cite{bu84} that M~31 GCs have stronger CN bands than their MW counterparts of same metallicity, even though confirmed by several later studies, has always been difficult to understand because it is very hard to explain how nitrogen abundances could differ substantially between otherwise very similar systems.  \\cite{lb03} proposed a scenario where GCs were formed from zero-metallicity material, pre-enriched by hypernovae explosions in the center of $\\sim 10^6 M_\\odot$ gas clouds.  Nitrogen abundance differences between the two GC systems was the one piece of evidence that could not be explained in that scenario.  We hope that the task of devising a mechanism for the formation of the haloes of the Milky Way and Andromeda galaxies is made easier by the findings presented in this paper.  In a forthcoming paper, we contrast measurements of the abundance patterns of M~31 and Galactic GCs, and discuss their implications to our understanding of the formation of the two galaxy haloes."
    },
    "1112/1112.2995_arXiv.txt": {
        "abstract": "{We address the issue of cosmological backreaction from non-linear structure formation by constructing an approximation for the time evolved metric of a dust dominated universe based on a gradient expansion. Our metric begins as a perturbation of a flat Friedmann-Robertson-Walker state described by a nearly scale invariant, Gaussian, power-law distribution, and evolves in time until non-linear structures have formed. After describing and attempting to control for certain complications in the implementation of this approach, this metric then forms a working model of the universe. We numerically calculate the evolution of the average scale factor in this model and hence the backreaction. We argue that, despite its limitations, this model is more realistic than previous models that have confronted the issue of backreaction. We find that the \\emph{instantaneous} effects of backreaction in this model could be as large as $\\sim10\\%$ of the background. This suggests that a proper understanding of the \\emph{cumulative} effects of backreaction could be crucial for precision cosmology and any future exploration of the dark sector.}  \\begin{document} ",
        "introduction": "The issue of cosmological backreaction \\cite{Buchert:1999er} has fueled a lively debate in the literature. The question is: can inhomogeneities in our universe backreact and affect the average dynamics within our causal horizon such that the observed acceleration can be attributed to their influence? If the answer is positive then the mystery of the cosmological constant will have found a solution requiring no new physics but which will nevertheless demonstrate the subtlety of gravitational phenomena. Even if the answer is negative, backreaction will be operative at some level due to the non-linear nature of gravity and its effects may be visible in the new generation of cosmological observations. Either way, calculating its magnitude is an important cosmological question.  The problem with assessing the magnitude of this backreaction lies with the complexities of the non-linear structures forming under gravity in the late universe. For example, using second order perturbation theory falls short of providing an answer since the effect is expected to emerge in the non-linear regime where perturbation theory breaks down. Attempts to model the non-linear structures involving toy models of voids and overdensities may be useful for developing intuition but are nevertheless simplistic. N-body simulations imply that the effect is small but they are Newtonian with zero global backreaction by construction; the backreaction is a purely relativistic effect \\cite{Buchert:1995fz}. A strong argument for the smallness of the effect uses the fact that even with non-linear overdensities, the local metric perturbations and peculiar velocities are still much smaller than unity (away from black holes) and that a perturbative FRW framework should therefore hold. However, counterarguments have been put forward involving subtleties in the choice of background employed in these treatments. For references finding a a small backreaction see e.g. \\cite{Ishibashi:2005sj, Paranjape:2008jc, Baumann:2010tm, Alonso:2010zv, Green:2010qy, Mattsson:2010ky} while arguments for significant backreaction can be found in e.g. \\cite{Rasanen:2006kp, Rasanen:2008it, Clarkson:2009hr, estimate, Collins:2010ea, Bochner:2011dr}. For recent articles and reviews on the subject with more extensive references  see \\cite{Buchert:2011sx, arXiv:1103.2335, arXiv:1102.1449,arXiv:1103.5974, arXiv:1102.1015, arXiv:1105.0909, arXiv:1106.1693, arXiv:1103.2016, arXiv:1102.0408, Kolb, arXiv:1105.1886, arXiv:1005.0788}.  In this paper we discuss cosmological backreaction in a novel manner, with a fully relativistic framework for a universe that begins as a perturbed Friedmann-Robertson-Walker, $\\Omega=1$, CDM universe. To be precise, we employ a gradient expansion to express the metric as a series of terms with an increasing number of spatial gradients and coefficients which are functions of proper time. The initial conditions are the standard adiabatic and Gaussian post-inflationary primordial perturbations. We use the synchronous gauge: our coordinate lines comove with CDM particles and our time hypersurfaces are labeled by their proper time. Of course the gradient series has to be truncated and thus does not capture the developing non-linearities entirely realistically. However, we argue that even if truncated this series can still provide a well motivated \\emph{model} for the true geometry which can be made increasingly accurate in principle. Furthermore, it extends into the non-linear regime, describing the collapse of initial over-densities and the rarefaction of initial under-densities which go on to form the voids dominating the cosmological volume.\\footnote{The synchronous gauge and a gradient expansion to study backreaction was first used in \\cite{Kolb:2005da} to argue that a non-perturbative approach is really needed to settle the issue. In the present work we use a truncated gradient series but also consider it in the non-perturbative regime, beyond its apparent radius of convergence. As explained in the text we argue that this can be used as a model for the metric of the universe which captures the evolution of over- and under-dense regions.}  Here is the outline of the rest of the paper: In the following section we obtain the series solution for the metric in a gradient expansion using a Hamilton-Jacobi formulation. This approach, first developed in e.g. \\cite{Croudace:1993yt} and \\cite{Stewart:1994wq}, simplifies the calculation significantly compared to a more straightforward expansion of the standard Einstein equations. Our treatment follows a slightly different logical development. Then, in section 3, we apply these results to the backreaction problem by numerically evaluating the evolution of the average scale factor and the backreaction parameter $Q$. We find that, up to certain qualifications which we explain, backreaction leads to non-negligible deviations from the unperturbed background model. Our results indicate that the effect might be relevant, indeed crucial, for precision cosmology and the exploration of the dark sector. We summarize and discuss our findings in section 4 where future directions are also laid out. ",
        "conclusions": "\\subsection{Potential improvements to the model}\\label{sec:extensions}  While our model is interesting and internally consistent and perhaps even an improvement in sophistication compared to earlier models, what we want is to know is the magnitude and nature of the backreaction in the real universe. The most obvious next step in improving our model is to add a cosmological constant term, $\\Lambda$. With that achieved, it would be possible to do a full Markov Chain Monte Carlo over the parameters of the model to see which set of parameters fits the observed data best. If our model, with backreaction, fitted the data better than the concordance model, without backreaction, then this would indicate that the assumptions we made to solve the various complications were well-motivated.  However, another complication that is less easily solved would need to be addressed before comparisons to the real universe can be meaningfully made. In this paper we have calculated the backreaction that arises as structures form over a small range of scales. We can vary where this range lies, but we cannot make it arbitrarily wide. To compare the model to observation we need a reliable calculation of the \\emph{total} backreaction that arises over \\emph{all} scales. In principle this problem can be solved by considering a larger box; however, in practical terms, this can never be fully achieved.  Given the hierarchical nature of structure formation, an ideal solution to this problem is a time-dependent grid spacing and box size. The time-dependence could be global, with the grid spacing $R$ set as a function of time $R(t)$. $R(t)$ would be chosen based on which scales start forming structures at which times and at which times structure growth freezes out. Alternatively, the time-dependence could be set locally. There could be a progressive ``zooming out'' as individual gridpoints expand into neighbouring gridpoints. Whenever two gridpoints collide it would be possible to stop considering them separately and to recalculate the metric from that time onwards as if they were one gridpoint. Precisely what metric to use for this combined gridpoint would need to be determined. Either a global or a local time-dependence to the grid spacing would face this difficulty. That is, the question of precisely what initial conditions, or background model, to use for the zoomed out metric. Should $t_0$ (which was set as a function of $h$) become a time-dependent quantity? If so, should the $h$ used in the transfer function also be dependent on the grid spacing and thus time?  If the zooming out problem could be solved it would remove any problems that relate to how to fix the over expansion of gridpoints. However, if the zooming out problem is not solved, there are some improvements that can be made to our fixing methods. Firstly, it should be possible to pick $g_{\\rm rat}$ from observations, rather than to treat it as a free parameter, as we have here. Secondly, we chose to fix our expanding regions to the initial background model. This is not ideal. When any region reaches the point where we wish to fix its expansion, the \\emph{average} expansion rate, and more importantly, the average volume and density, are not the same as those of the initial background, extrapolated to this time. If the universe is to reach a new asymptotic FRW state it should be related to the new average density, not the one extrapolated from the initial background model. However, doing this is not easy. The deviation of the average density from the initial background model changes with time. Therefore, the correct background model to fix the expanding regions to also changes with time. When the total backreaction is $<10\\%$ these issues will be less important than the many other features of the model. However, when comparing the model to precise measurements, perhaps $10\\%$ of $10\\%$ is no longer an ignorably small quantity.  \\subsection{Conclusions}  Let us close by summarizing our methodology and the conclusions we can draw from it. We have approximated the evolving inhomogeneous synchronous gauge metric of a CDM $\\Omega=1$ universe with realistic initial conditions using a gradient expansion, and studied the backreaction on the average dynamics. This gradient expansion is an approximation to the true metric of such a universe which goes beyond the description of standard perturbation theory. Instead of being limited by the smallness of the inhomogeneities, it has a temporal range of validity set by the initial local 3-curvature. For timescales $t>t_{\\rm con}$ the gradient expansion does not converge anymore. This does not pose a real problem for studying regions with positive initial curvature which collapse at around that timescale and eventually virialize. However, it does present limitations in following the late time evolution of under-dense regions with negative initial curvature. Nevertheless, if the late time evolution of such regions is fixed in some prescribed manner, this approach offers a way to construct a realistic \\emph{model} even for the late-time metric of such a universe.  Our methodology improves over previous approaches in that it is fully relativistic, goes beyond perturbation theory and does not treat the universe as a collection of unconnected under-dense and over-dense structures with high symmetry. Of course it does require modeling of the late time behaviour of the under-dense regions which is not naturally captured by the gradient expansion but such modeling can be simply parameterized and possibly informed by observation. We have attempted to achieve this in a simple manner in this paper but it can still be made more realistic in a number of ways.  Given the approximations mentioned, we computed the backreaction of inhomogeneities in this model metric. Our qualitative results indicate that backreaction may constitute more than a percent effect and is thus highly relevant in future considerations of precision cosmology. Of course more realism is needed before definite statements can be made regarding the backreaction in our universe.  There is another aspect in which such a model would aid in understanding the effects of inhomogeneities for cosmology. We have focused on the average dynamics but what we actually see is light propagating through the inhomogeneous universe and not the average scale factor. One could question whether the latter is at all relevant for our observations. To really assess the impact of inhomogeneities on observations one should trace light rays through the inhomogeneous spacetime and determine what observers would see \\cite{arXiv:1002.1232,arXiv:1109.2484}. This is easy to do in a box described by our metric. In particular we can compute what happens to the redshift of photons or the travel time though over-dense and under-dense regions. The great advantage of our approach is that this is easily calculable within our model. Therefore, even irrespective of whether our model describes the universe entirely accurately, we can answer the question of whether a large backreaction in the synchronous gauge corresponds to a large shift in the time it takes a photon to traverse the universe and/or the redshift it experiences. We will return to this issue and further improvements to our model in future work."
    },
    "1112/1112.0281.txt": {
        "abstract": "We present the initial-final mass relation derived from 10 white dwarfs in wide binaries that consist of a main sequence star and a white dwarf. The temperature and gravity of each white dwarf was measured by fitting theoretical model atmospheres to the observed spectrum using a $\\chi^{2}$ fitting algorithm. The cooling time and mass was obtained using theoretical cooling tracks. The total age of each binary was estimated from the chromospheric activity of its main sequence component to an uncertainty of about 0.17 dex in log \\textit{t} The difference between the total age and white dwarf cooling time is taken as the main sequence lifetime of each white dwarf. The initial mass of each white dwarf was then determined using stellar evolution tracks with a corresponding metallicity derived from spectra of their main sequence companions, thus yielding the initial-final mass relation. Most of the initial masses of the white dwarf components are between 1 - 2 M$_{\\odot}$. Our results suggest a correlation between the metallicity of a white dwarf's progenitor and the amount of  post-main-sequence  mass loss it experiences - at least among progenitors with masses in the range of 1 - 2 M$_{\\odot}$.  A comparison of our observations to theoretical models suggests that low mass stars preferentially lose mass on the red giant branch. %Wide binaries tend to be much older than nearby clusters. Thus, they contribute valuable limits to the low mass end of the initial-final mass relation where clusters provide sparse information. %The lack of a similar correlation among cluster stars suggests that the mechanism for mass loss in wide binaries differs from that in open clusters or that not enough low mass white dwarfs are observed in clusters to see such a relation. %Willson's model when the age-metallicity correction for progenitors of white dwarfs is taken into account.  %% what keyword punctuation is appropriate. ",
        "introduction": "Over 90 percent of all stars shed at least half their mass as they evolve towards their final state - a white dwarf (WD). The initial-final mass relation (IFMR) represents a mapping between the mass of a WD remnant and the mass of its hydrogen-burning main-sequence (MS) progenitor. It also characterizes the amount of material stars with primordial masses M $\\lesssim$ 8 M$_{\\odot}$ recycle to the interstellar medium. Thus, it is of paramount importance to understanding the chemical enrichment and the efficiency of star formation in galaxies. This relation is also a key constraint on stellar evolution theory. %Unfortunately, currently it is not well constrained either by theory or by observation.  One of the first attempts to empirically determine the IFMR was undertaken by Weidemann (1977). Reimers $\\&$ Koester (1982) and Koester $\\&$ Reimers (1996) presented observations of WDs in the open cluster NGC 2516 and obtained an IFMR using these WDs.  Weidemann (2000; W00 hereafter) updated the IFMR by incorporating new theoretical and observational data. Claver et al. (2001) observed six WDs in the Praesepe open cluster and determined a monotonic IFMR. Williams et al. (2004; 2009) presented an empirical determination of the IFMR based on a spectroscopic analysis of massive white dwarfs in NGC 2168 (M35). They showed that the resultant white dwarf mass increases monotonically with progenitor mass for masses greater than 4 M$_{\\odot}$. Ferrario et al. (2005) re-evaluated the ensemble of data that has been used to determine the IFMR and characterized a mean IFMR about which there is an intrinsic scatter. They showed that a linear IFMR predicts a mass distribution in reasonable agreement with the Palomar-Green survey.  Kalirai et al. (2005) determined the IFMR from observations of very faint WDs in the rich open cluster NGC2099 (M37). They found stars with initial masses between 2.8 and 3.4 M$_{\\odot}$ lose 70 - 75\\% of their mass during post-MS evolution. Dobbie et al. (2006) also constructed a new IFMR based on 11 WDs in the Praesepe open cluster.  Rubin et al. (2008) constructed an IFMR  based on 19 spectroscopically identified WDs in NGC 1039 (M34).  Catal\\'{a}n et al. (2008a) studied the IFMR using six WDs in wide binaries and suggested the IFMR may not be a single-valued function. Kalirai et al. (2008) presented constraints on the low-mass end of the IFMR using older open clusters NGC7789 (\\textit{t} = 1.4 Gyr ), NGC 6819 (\\textit{t} = 2.5 Gyr) and NGC6791 (\\textit{t} = 8.6 Gyr). Later, Kalirai et al. (2009) extended the IFMR to lower masses using the globular cluster M4. Since most nearby clusters are relatively young, the initial masses of those WDs tend to be high. %So far, WDs with low initial masses of only a few clusters have been included in a determination of the IFMR (Kalirai et al. 2008; 2009).  In this paper we investigate the IFMR using wide ``fragile\" binary systems containing a WD with a MS companion. The systems have relatively large orbital separations ($<a>$ $\\sim$ 10$^{3}$ AU; Oswalt et al. 1993; Silvestri et al. 2001, 2002, 2005). Thus, one can safely assume that each component has evolved independently, unaffected by mass exchange or tidal coupling that complicate the evolution of closer pairs. Components of a given binary are coeval (Greenstein 1986). Essentially, each may be regarded as an open cluster with only two components.  Although it is difficult to obtain ages for wide binaries as accurate as ages for open clusters, they tend to be nearer, brighter and are far more numerous than nearby clusters. Moreover, they span a much more continuous range in age. Catal\\'{a}n et al. (2008a; C08a hereafter) used six wide binaries to investigate the IFMR; three of the WDs  had low initial mass ($<$ 2M$_{\\odot}$). Our sample contains additional WDs at the low initial mass limit.    Previous research has indicated a large scatter in the empirical IFMR. What is the source of this scatter? Kalirai et al. (2005) found some weak evidence of a metallicity dependence in the IFMR. Kalirai et al. (2007) found evidence for enhanced mass loss at extremely high metallicities by studying the white-dwarf mass distribution in the supersolar-metallicity star cluster NGC 6791 ([Fe/H] = +0.4). Kalirai et al. (2009) found a relatively flat relation between mass loss and metallicity ([Fe/H] between -1.1 to solar metallicity) by extending these studies to WDs in the globular cluster M4.  However, a clear relation between the metallicity and scatter in the IFMR has not been demonstrated (cf. Williams 2007; Catal\\'{a}n et al. 2008b, C08b hereafter). We investigated whether there is a metallicity dependence on the IFMR  in our wide binary sample using the spectra of their MS companions.    Section 2 provides an overview of the observations and reductions for our sample. The astrophysical properties of the WDs are discussed in section 3. The MS companions are discussed in section 4. In section 5, we present and discuss our IFMR. Section 6 compares our observations to theoretical models of post-MS mass loss. A discussion of the implications of our findings is given in section 7. ",
        "conclusions": "%Eight are DA type and two are DB type. The $T$$\\rm_{eff}$ and log $g$ of DA WDs were derived from the H line fits between the observed spectra and %theoretical spectra, while those of two DB WDs are obtained from the literature. The mass and cooling times of the WDs %were determined from model cooling tracks. The total ages of the binaries were estimated from the CA of their MS components. First, S$\\rm_{HK}$ is measured. Then we construct the calibration between our $S\\rm_{HK}$ and $S\\rm_{MW}$. After $S\\rm_{HK}$ transformed into $S\\rm_{MW}$, $R\\arcmin\\rm_{HK}$ could be calculated by Equation Noyes et al. (1984). Finally, the ages of each binary were determined by the relation between $R\\arcmin\\rm_{HK}$ and age (Soderblom et al. 1991). The metallicities of wide binaries were derived by template matches between the MS spectra and theoretical spectra. The difference between the total age and the cooling time of each WD was taken as the lifetime of its progenitor. The initial masses of 10 WDs were then determined by the evolution track. Our IFMR contains six WDs whose M$\\rm_{i}$ are lower than 2 M$_{\\odot}$. They contribute to the low initial mass limit not well sampled by clusters.  Our WDs in wide binaries suggest a linear IFMR over the initial mass range 1.1 M$_{\\odot}$ to 4.1 M$_{\\odot}$: In this study, we constructed an empirical IFMR using 10 WDs in wide binaries. Our IFMR contains six WDs whose M$\\rm_{i}$ are lower than 2 M$_{\\odot}$. They contribute to the low initial mass limit that is not well-sampled by clusters.  Our WDs in wide binaries suggest a linear IFMR over the initial mass range 1.1 M$_{\\odot}$ to 4.1 M$_{\\odot}$ (Equation 5).  We compared our mass loss vs. metallicity relation to theoretical models for evolving lower MS stars ($<$ 2 M$_{\\odot}$). In general, the models predict a net mass loss and IFMR that agree with the values found from our observation within the current uncertainty of measurement.  Kalirai et al. (2007) and Kalirai et al. (2009) tentatively found a metallicity dependence on the IFMR. We find that at least part of the scatter seen in the IFMR is correlated with metalllcity. Stars with lower metallicity apparently shed less mass when they become WDs. %{\\bf Except NGC 6791 (Kalirai et al. 2007) and M4 (Kalirai et al. 2009), IFMRs from most of open clusters have shown no such dependance on metallicity, }it is curious that evolution for stars in wide binaries seems to be different from that in open clusters. Perhaps this is because the Z dependance of mass loss is most pronounced in low mass progenitors, which are poorly sampled by observation of cluster.  %Although only two DBs are in our current sample, we suspect the metallicity effect on the IFMR is different for DA and DB stars. More observational data will be needed to prove this. %  %We made a comparison between our mass loss vs. [Fe/H] relation and theoretical model for lower mass WDs ($<$ 2 M$_{\\odot}$). %Overall, the mass loss models predict net mass loss and M$\\rm_{f}$ vs. M$\\rm_{i}$ that are very close to the values found from the white dwarf pairs.  This will be true for other mass loss formulae with similar death-lines, as discussed by Willson (2007, 2008a, 2008b, 2009).  In addition, the Z-dependence of our models is consistent with the data from these white dwarf stars.  The bulk of the Z-dependence in the mass loss rates comes from the dependence of R on Z, with a smaller effect from the efficiency of forming dust (the gas/dust ratio) in the models, it is important to note.    %% The displaymath environment will produce the same sort of equation as %% the equation environment, except that the equation will not be numbered %% by LaTeX.   %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the"
    },
    "1112/1112.3595_arXiv.txt": {
        "abstract": "We have conducted a series of numerical experiments with the spherically symmetric, general relativistic, neutrino radiation hydrodynamics code \\aboltz\\ to examine the effects of several approximations used in multidimensional core-collapse supernova simulations. Our code permits us to examine the effects of these approximations quantitatively by removing, or substituting for, the pieces of supernova physics of interest. These approximations include: (1) using Newtonian versus general relativistic gravity, hydrodynamics, and transport; (2) using  a reduced set of weak interactions, including the omission of non-isoenergetic neutrino scattering, versus the current state-of-the-art; and (3) omitting  the velocity-dependent terms, or observer corrections,  from the neutrino Boltzmann kinetic equation. We demonstrate that each of these changes has noticeable effects on the outcomes of our simulations. Of these, we find that the omission of observer corrections is particularly detrimental to the potential for neutrino-driven explosions and exhibits a failure to conserve lepton number. Finally, we discuss the impact of these results on our understanding of current, and the requirements for future, multidimensional models. ",
        "introduction": "Colgate and White \\citeyearpar{CoWh66}  were the first to propose that core-collapse supernovae may be neutrino-driven and performed the first numerical simulations of such events, launching more than four decades of research that continues to this day. A significant milestone occurred nearly two decades later with Wilson's discovery that delayed neutrino-driven explosions could be obtained. Based on his models, Wilson concluded \\citep{Wils85,BeWi85} that the stalled supernova shock wave could be revived via neutrino absorption on a time scale of several hundred milliseconds given the intense flux of neutrinos emerging from the proto-neutron star liberating the star's gravitational binding energy. Observations of the neutrinos from SN1987A, the first such observations of supernova neutrinos \\citep{BiBlBr87,HiKaKo87}, provided support for the central role of neutrinos in the explosion mechanism. State-of-the-art simulations today continue to explore Wilson's neutrino-driven explosion mechanism in the context of two- and three-dimensional models \\citep[e.g., see][]{BuLiDe07,MaJa09,BrMeHi09b,SuKoTa10}.  Neutrinos are weakly interacting particles whose cross sections are energy dependent. Thus, unlike all other components in a supernova model, they are not well described as a fluid, except in the deepest layers, and their transition in space to non-fluid-like behavior  depends on their energy. Instead, the evolution of the neutrino radiation field, particularly in the semi-transparent regime, is far better characterized by classical kinetics---specifically, the general relativistic Boltzmann kinetic equation  \\citep[e.g., see][]{CaMe03}, \\notetoeditor{Superscripts should be placed to the left of subscripts in the following equation:} \\begin{equation} p^{\\hat \\mu} \\left(  {\\Lambda^{\\bar\\mu}}_{\\hat\\mu} {e^\\mu}_{\\bar\\mu} \\frac{\\partial f}{\\partial x^\\mu} - {\\Gamma^{\\hat\\nu}}_{\\hat\\rho \\hat\\mu} p^{\\hat\\rho} \\frac{\\partial f}{\\partial p^{\\hat\\nu}}\\right) ={C}[f], \\label{eq:boltz} \\end{equation} where, for spherically symmetry, \\begin{eqnarray} \\frac{1}{E}{C}[f] &=& \\label{eq:collision} (1-f)j-\\chi f \\\\ \\nonumber &+&\\frac{1}{c}\\frac{1}{h^{3}c^{3}}E^{2}\\int d\\mu' R_{{\\rm IS}}(\\mu, \\mu',E)f \\\\ \\nonumber &-&\\frac{1}{c}\\frac{1}{h^{3}c^{3}}E^{2}f\\int d\\mu'   R_{{\\rm IS}} (\\mu, \\mu',E)\\\\ \\nonumber &+&\\frac{1}{h^{3}c^{4}} (1-f) \\int dE'E'^{2}d\\mu' R_{{\\rm NIS}}^{{\\rm in}}(\\mu, \\mu',E,E') f \\\\ \\nonumber &-&\\frac{1}{h^{3}c^{4}} f \\int dE'E'^{2}d\\mu' R_{{\\rm NIS}}^{{\\rm out}}(\\mu, \\mu',E,E') (1-f) \\\\ \\nonumber &+&\\frac{1}{h^{3}c^{4}} (1-f) \\int dE'E'^{2}d\\mu' R_{{\\rm PR}}^{{\\rm in}}(\\mu, \\mu',E,E') (1-\\bar{f}) \\\\ \\nonumber &-&\\frac{1}{h^{3}c^{4}} f \\int dE'E'^{2}d\\mu' R_{{\\rm PR}}^{{\\rm out}}(\\mu, \\mu',E,E') \\bar{f}. \\end{eqnarray} Equation (\\ref{eq:boltz}) describes the evolution of the neutrino distribution function $f(t,x_1,x_2,x_3,\\mu_1,\\mu_2,E)$, which at time $t$ and spatial location $(x_1,x_2,x_3)$ supplies the distribution of neutrinos in direction cosines $(\\mu_1,\\mu_2)$ and energy $E$--i.e., the angular and spectral distribution of neutrinos. One such Boltzmann equation is solved for each flavor of neutrino---electron, muon, and tau (\\nue, \\numu, and \\nutau, respectively)---and for their antineutrino partners (\\nuebar, \\numubar, and \\nutaubar). The invariant collision term, ${C}[f]$, in equation~(\\ref{eq:collision}) is written using emission, $j$,  absorption, $\\chi$, and scattering and pair kernels, $R$, following the forms often used for neutrino transport \\citep[e.g., see][]{MeMe99}, where $\\bar{f}$ is the distribution function for the partner antineutrino and $\\mu$  is the neutrino direction cosine. In equations (\\ref{eq:boltz}) and (\\ref{eq:collision}), $f$ is a function of $(\\mu,E)$, as well as position and time. The $(\\mu',E')$ dependence of $f$ and $\\bar{f}$ inside the integrals illustrates the physical coupling of all energies and angles for each neutrino species and of neutrino and antineutrino partners.  The first term on the left-hand side of equation (\\ref{eq:boltz}) describes the time evolution of the local neutrino distribution owing to spatial transport through the volume of interest. The second, far more complex term on the left-hand side (\\momspder) describes the evolution of the local neutrino distribution in angle and energy as the result of (A) the coordinate system chosen, (B) special relativistic effects, and (C) general relativistic effects. In what follows, we will refer to (B) as ``observer corrections.''  Terms describing (A) depend on the choice of coordinate system. For example, in spherical-polar coordinates, the neutrino direction cosine relative to the outwardly pointing radial vector changes as the neutrino propagates through a local volume. This coordinate-system effect is included in \\momspder\\ and is present even in the absence of fluid motion or general relativity. For Cartesian coordinates, the neutrino direction cosines do not change as a result of the coordinate system choice alone and, consequently, such a term is absent.  Terms describing (B) depend on the frame of reference chosen to measure the neutrino direction cosines and energies. The comoving frame, with neutrino direction cosines and energies  measured in an inertial frame of reference instantaneously comoving with the stellar core fluid with which the neutrinos interact, is often used. Neutrino--matter interactions are naturally expressed in this frame. Given this choice, the  terms in \\momspder\\  present a significant numerical challenge. Finding discrete representations that guarantee conservation of lepton number and energy is one of the most difficult aspects of modeling neutrino transport in stellar cores. This has been achieved for general relativistic, spherically symmetric flows \\citep{LiMeMe04}, providing the conceptual and implementation groundwork for achieving the same  in axisymmetric (2D) and non-symmetric (3D) flows. Further theoretical foundations have been laid \\citep{CaMe03,CaLeMe05}; steps toward the development of a 2D Boltzmann solver have been taken \\citep{OtBuDe08}; and the challenge now is to fully implement lepton energy and number conserving discretizations in 2D and 3D models.  For another choice of reference frame---measuring neutrino angles and energies relative to the inertial, ``lab\" frame of a distant observer---the terms encapsulating the special relativistic effects in \\momspder\\ are absent, simplifying the left-hand side of the Boltzmann equation. In such a frame of reference, the neutrino direction cosines and energies do not change from observer to observer in the frame. However, this simplification comes at a price because the neutrino--matter interactions are naturally described in the comoving frame. In the lab frame,  a Lorentz transformation is required in order to express the comoving-frame neutrino--matter interactions in terms of the lab-frame direction cosines and energies, which introduces non-trivial velocity dependencies into the lab-frame collision term.  One approach to the complexity of the lab-frame collision term is the ``mixed frame'' approach, which uses the lab-frame 4-momenta and an $\\mathcal{O}(v/c)$ Taylor-series expansion in energy of the comoving-frame emissivities and opacities \\citep{MiKl82}. \\citet{HuBu07} have proposed to use the mixed-frame approach for core-collapse simulations with extensions for non-isotropic and non-isoenergetic scattering. The mixed-frame approach may be difficult to extend to arbitrarily relativistic flows, and has not yet been used in the context of a full-physics core-collapse supernova simulation.  In a general relativistic setting, such as core-collapse supernovae, we must contend with (B) and/or (C) regardless of the frame of reference chosen to describe the neutrino direction cosines and energies. Even for static general relativistic environments, angular aberration, gravitational red shift, and other effects occur, and the resulting terms in (C) are always present.  Regardless of  approach, comoving- or lab-frame, it is problematic to adapt the simplicity of both approaches, simultaneously  simplifying the left- and right-hand sides of the Boltzmann equation, as has been done in \\citet{BuLiDe06,BuLiDe07}, \\citet{OtBuDe08}, and other models using the \\vulcan\\ code, although one can view the implementation in these works as steps toward a more complete description. They deploy a lab-frame approach for terms describing angular aberration and energy shift on the left-hand side (or assume such terms are unimportant in a comoving-frame approach), while simultaneously deploying a comoving-frame approach for the collision term describing the neutrino--matter interactions on the right-hand side. This is not a mixed-frame approach in the sense described above. It is an approach not based in any reference frame, and it is physical only for static cases in which there is no distinction between lab and comoving frames. One of the goals of this study is to investigate the importance of the terms in \\momspder\\ in a comoving-frame approach, and whether they can be ignored while using a comoving-frame approach for the collision term.  Modeling general relativistic Boltzmann kinetics is also challenging because of the complexity of the collision term on the right-hand side of the Boltzmann equation, even in a comoving-frame formulation. Looking at equation (\\ref{eq:collision}), we see that the collision term describes the full, direct coupling of all neutrino angles and energies for each neutrino species, owing to neutrino isoenergetic (IS) scattering on nuclei and non-isoenergetic (NIS) scattering on electrons and nucleons. The pair creation and annihilation processes (PR) such as electron--positron annihilation and nucleon--nucleon bremsstrahlung also couple the angles and energies of the neutrino and antineutrino species of  each  flavor together. The coupling of all neutrino angles and energies through the relevant set of weak interactions dominates the computation associated with the solution of the neutrino Boltzmann equations. It has been argued \\citep{BuLiDe06,BuLiDe07,NoBuAl10} these couplings are subdominant and can be ignored,  greatly simplifying the neutrino Boltzmann equations and significantly reducing the computational cost associated with their solution. A second goal of this study is to investigate whether or not such  approximations to the collision term are realistic.  The complete general relativistic Boltzmann equation was solved in spherically symmetric models of core-collapse supernovae by the Oak Ridge-Basel collaboration \\citep{LiMeTh01,LiMeMe04} and  by Sumiyoshi and collaborators \\citep{SuYaSu05}. Achieving this  in three-dimensional models of core-collapse supernovae presents a major challenge, one that will likely require sustained exascale resources to meet.   The overarching goal of this study is to use  general relativistic, spherically symmetric Boltzmann simulations to guide and, more importantly, set minimum requirements for accurate 2D and 3D simulations. We use the Oak Ridge-Basel code \\aboltz\\  in these studies to compare general relativistic--full weak interaction physics (\\grfull), Newtonian--full weak interaction physics (\\nfull), Newtonian--reduced weak interaction physics (\\nreduc), and Newtonian--reduced weak interaction physics--no observer correction (\\noc) models. These models will demonstrate the importance of general relativity, a complete weak interaction set and treatment, and the terms in \\momspder\\ to stellar core collapse and the post-core-bounce evolution. Current multidimensional models suggest that spherical symmetry is a reasonable approximation for the first 100--150 ms after bounce \\citep{MaJa09,BrMeHi09b,SuKoTa10}. Thus, the simulations presented here are relevant for discussing the initial conditions present for all multidimensional phenomena that might ensue; e.g., neutrino-driven convection and the standing accretion shock instability (SASI). ",
        "conclusions": "We have examined the consequences of removing (1) GR effects, (2) non-isoenergetic scattering and detailed nuclear EC opacities, and (3)  observer corrections from spherically symmetric models of core-collapse supernovae. We have found that all of these changes, individually and especially when taken together, affect the progress of stellar collapse and the post-bounce evolution of the shock and core thermodynamic properties in significant ways, in constrast to the assessments made by \\citet{BuLiDe06,BuLiDe07} and \\citet{NoBuAl10}. We have computed variations in the shock radius, neutrino luminosities, and neutrino RMS energies as large as 60~km, 35~\\Bethes, and 10~\\mev, respectively, across the four models considered here.  Omission of GR results in a less compact core and an unrealistically more favorable shock progression after bounce. Eliminating non-isoenergetic scatterings and simplifying electron capture on nuclei drastically reduces the core deleptonization and expands the homologous core at bounce. Omission of the observer corrections dramatically alters  core deleptionization, the shock position, and neutrino luminosities after bounce, in part resulting from a complete breakdown of lepton number conservation.   The lepton non-conservation and non-promotion of neutrino energy resulting from omitting observer corrections in our \\noc\\ model results in a compact, low-$Y_L$ core and a shock trajectory that is the least favorable of our models. The artificial loss of lepton number, lower neutrino luminosities, and the consequent lower neutrino heating rate and shallower shock trajectory may explain the lack of neutrino-driven explosions in models computed with \\vulcan\\ \\citep[see][]{BuLiDe07}, in contrast to the results reported by others \\citep{MaJa09,BrMeHi09b,SuKoTa10,TaKoSu11}.  Moreover, the changes in $Y_{e}$ and $Y_{L}$, their gradients, and the entropy gradients that we see as we traverse the models shown here will change the location and strength of convectively unstable regions in the proto-NS and between the proto-NS and the shock. The lepton and entropy gradients in the proto-NS drive prompt convection, the entropy gradients between the proto-NS and the shock drive neutrino-driven convection, and these in turn seed and are seeded by the SASI. That is, the changes we have documented in this transport study have implications for all of the multidimensional phenomena we know to be important in multidimensional supernova models once spherical symmetry is broken.  All of the ingredients (1)--(3) above must be included in multidimensional simulations of core-collapse supernovae to ensure  physical fidelity. Their omission is not the only approximation used in current multidimensional simulations, some of which (like the ray-by-ray approximation) are inadequately understood and need to be better understood or phased out. Certainly, further examination of these approximations is required within the context of multidimensional simulations."
    },
    "1112/1112.3240_arXiv.txt": {
        "abstract": "{} % {This work aims to study the unexplained sulfur depletion observed toward dense clouds and protostars.} {We made simulation experiments of the UV-photoprocessing and sublimation of H$_2$S and H$_2$S:H$_2$O ice in dense clouds and circumstellar regions, using the Interstellar Astrochemistry Chamber (ISAC), a state-of-the-art ultra-high-vacuum setup. The ice was monitored in situ by mid-infrared spectroscopy in transmittance. Temperature-programmed desorption (TPD) of the ice was performed using a quadrupole mass spectrometer (QMS) to detect the volatiles desorbing from the ice.} {Comparing our laboratory data to infrared observations of protostars we obtained a more accurate upper limit of the abundance of H$_2$S ice toward these objects. We determined the desorption temperature of H$_2$S ice, which depends on the initial H$_2$S:H$_2$O ratio. UV-photoprocessing of H$_2$S:H$_2$O ice led to the formation of several species. Among them, H$_2$S$_2$ was found to photodissociate forming S$_2$ and, by elongation, other species up to S$_8$, which are refractory at room temperature. A large fraction of the missing sulfur in dense clouds and circumstellar regions could thus be polymeric sulfur residing in dust grains.} {} ",
        "introduction": "\\label{intro}  Sulfur is depleted in molecular clouds by a factor of 1000 compared to its estimated cosmic abundance (Tieftrunk et al. 1994), while in the diffuse interstellar medium the abundance of sulfur in the gas phase is comparable to the cosmic abundance. This suggests that there is a form of sulfur in the gas phase that was not observed in molecular clouds, which could be atomic sulfur, or alternatively that sulfur chemistry on icy grain mantles, present in dense clouds and regions around YSOs but not in the diffuse interstellar medium, plays an important role.  There is compelling evidence that supports the role of dust grains on sulfur chemistry. Several gas-phase S-containing molecules were observed in hot cores, such as OCS, H$_2$S, H$_2$CS, SO, SO$_2$, HCS$^+$, and NS (van der Tak et al. 2003). Current gas-phase chemical models are unable to explain the abundances of S-species like HCS$^+$ and OCS measured toward protostars (Doty et al. 2004). The depletion of sulfur is observed not only in dense clouds, but also toward Class 0 and Class I sources (Buckle \\& Fuller 2003) and toward hot cores (Wakelam et al. 2004). The abundances of S-bearing species, including H$_2$S, suggest that these molecules are formed on grain surfaces and subsequently released to the gas phase.  Several S-bearing molecules have been detected in comets. Among them H$_2$S has the largest abundance, from 0.2 to 1.5\\% relative to H$_2$O (Irvine et al. 2000). There seems to be a general agreement between the molecular abundances observed in circumstellar ices and in comets (Bockel\\'ee-Morvan et al. 2000). That suggests that H$_2$S is expected to be present in circumstellar ice mantles. The H$_2$S ice will be strongly processed by UV and ion irradiation (Garozzo et al. 2010; Grim \\& Greenberg 1987; Moore et al. 2007).  We explore here the possibility that the missing sulfur atoms may be present in icy grain mantles. The cosmic abundance of sulfur is 1.23 $\\times$ 10$^{-5}$ N$_{\\rm H}$, or 37 times less abundant than oxygen (Snow \\& Witt 1996), and like oxygen, sulfur belongs to group 16 of the periodic table. The electronegativity values of S and O are 2.58 and 3.44 on the Pauling scale, and the electron affinities are 200 and 141 kJ mol$^{-1}$. The bond energy of an S-H bond is 363 kJ mol$^{-1}$, while that of an O-H bond is 458.9 kJ mol$^{-1}$, and therefore S-H bond formation is favored over O-H bond formation. Given that H is the most abundant element, S atoms will tend to form H$_2$S molecules because they impinge on icy grain mantles. The feature due to the  \\ stretching mode of H$_2$S at 3.925 $\\mu$m (2548 cm$^{-1}$) has a band strength of $A$ $\\approx$ 2.9 $\\times$ 10$^{-17}$ cm molecule$^{-1}$ (Smith 1991), and $A$(H$_2$O) = 2.0 $\\times$ 10$^{-16}$ cm molecule$^{-1}$ (Hagen et al. 1981). Therefore, if we make the crude assumption that the H$_2$S abundance in the ice is roughly about 1/37 that of H$_2$O ice, based on the S/O = 1/37 cosmic abundance ratio, the absorbance area of H$_2$S relative to that of H$_2$O might just be 0.4\\%, which is close to the detection limit of most observations. With the exception of a weak CH$_3$OH absorption, the feature falls on a relatively {\\it clean} part of the mid-infrared spectrum. If present, it might be observable in the spectra of circumstellar or dense interstellar icy grains.  Solid H$_2$S has not been detected in the interstellar medium. The presence of H$_2$S ice was inferred in W33A, a high-mass protostar (Geballe 1985; 1991), and a band at 4.9 $\\mu$m (2040 cm$^{-1}$) was attributed to OCS (Geballe 1985, Palumbo et al. 1995). The detection of OCS is possible since the 4.9 $\\mu$m band has a considerable band strength, $A$ = 1.5 $\\times$10$^{-16}$ cm molecule$^{-1}$ (Hudgins et al. 1993). In addition to OCS, SO$_2$ was detected in ice mantles (Boogert et al. 1997). The H$_2$S detection by Geballe is not fully supported in the literature. Van der Tak et al. (2003) argue that infrared observations do not support the assumption that H$_2$S is the main S reservoir in grain mantles, and they provide the ISO-SWS observations of W33A as an example. One of the problems for identifying the 3.925 $\\mu$m (2548 cm$^{-1}$) band of H$_2$S in H$_2$O-rich ice mantles was the lack of laboratory spectra of H$_2$S embedded in an H$_2$O matrix, which is expected to affect this band significantly.  Ultraviolet emission spectra of the coma of comet IRAS-Araki-Alcock showed the presence of S$_2$, and the spatial profiles indicate a release of this species directly from, or very close to, the nucleus (A'Hearn et al. 1983). Diatomic sulfur was later found in the comet Hyakutake (Laffont et al. 1996). Based on the formation of S$_2$ at 12 K from irradiation of dirty ice containing H$_2$S, Grim \\& Greenberg (1987) suggested that S$_2$ was formed in interstellar ice mantles that ultimately aggregate into comets. Subsequently, fast reactions between OCS and metastable S, produced during the dissociation of CS$_2$, were proposed. This supports the formation of S$_2$ in the coma (e.g., A'Hearn et al. 2000). The formation of polymeric sulfur in interstellar environments involves the dissociation of H$_2$S and might, therefore, serve as a probe of energetic processing of the precometary ice.  Ion irradiation experiments of ice analogs containing CO, CH$_3$OH, and S-bearing species lead to formation of OCS and CS$_2$ molecules (Ferrante et al. 2008; Garozzo et al. 2010).  Previously, we performed an experiment consisting of the photoprocessing of H$_2$O:CO:NH$_3$:H$_2$S ice followed by warm-up to room temperature. This residue was analyzed by means of gas chromatography coupled to mass spectroscopy. Among the residue products several N-heterocycles and a number of S-bearing molecules were detected. S-polymers, S$_6$ through S$_8$, were formed, resulting from the S atoms released after photodissociation of H$_2$S ice. But also pentathian (S$_5$CH$_2$), hexathiepan (S$_6$CH$_2$), and c-(S-CH$_2$-NH-CH$_2$-NH-CH$_2$) were detected (Mu\\~noz Caro 2002). The presence of S-containing refractory molecules in icy grain mantles, like sulfur polymers, might be the reservoir of the missing sulfur in dense clouds and circumstellar environments (Wakelam et al. 2005).  To approach the S-depletion dilemma in dense clouds and YSOs, we carried out a series of experiments on the deposition of pure H$_2$S or H$_2$S in an H$_2$O-matrix under UHV conditions at 7 K, to mimic interstellar/circumstellar conditions followed by warm-up. We measured the mid-infrared spectra of both pure H$_2$S ice and H$_2$S in an H$_2$O ice matrix at different temperatures from 7 K to sublimation. Similar spectra were reported by Moore et al. (2007) at higher temperatures. Using a quadrupole mass spectrometer (QMS), we obtained the temperature programmed desorption (TPD) plots showing the abundances of the molecules released to the gas phase as a function of temperature during warm-up. These experiments were repeated, including UV irradiation of the ice. The infrared spectra of H$_2$S ice measured in the laboratory were compared with spectroscopic observations performed by ISO, providing new upper limits on the H$_2$S abundance toward protostars.  The layout of this paper is as follows. In Sect.~2 we describe the experimental protocol. The experimental results are presented in Sect.~3. The astrophysical implications are discussed in Sect.~4, and the main conclusions summarized in Sect.~5. ",
        "conclusions": "\\label{conc} We determined the desorption temperature of H$_2$S ice, which depends on the initial H$_2$S/H$_2$O ratio. Pure H$_2$S ice desorbs around 82 K. When H$_2$S is present in an H$_2$O ice matrix, a fraction of the H$_2$S co-desorbs with H$_2$O in the 130--170 K temperature range, showing two maxima around 145 and 163 K. These results agree with Collings et al. (2004). Comparison of the laboratory infrared spectra of H$_2$S, pure or in an H$_2$O matrix, with ISO observations of protostars W33A and IRAS18316-0602 provided upper limits of 0.7 and 0.13\\% on the solid H$_2$S abundance relative to H$_2$O. These values are too low to explain the S depletion observed toward dense clouds and circumstellar regions. Another reservoir of S in these regions could be the products of H$_2$S ice photoprocessing. It was found that solid H$_2$S, which expected in icy grain mantles, photolyzes very readily, leading to the formation of several species, including H$_2$S$_2$, HS$^{\\cdot}$, HS$_2^{\\cdot}$, S$_2$, and H$_2$SS. The set of reactions leading to these species is known, and it was possible to determine the values of some of the rate constants involved. If H$_2$S was present in an H$_2$O ice matrix the SO$_2$, SO$_4^=$, HSO$_3^-$, HSO$_4^-$, H$_2$SO$_2$, H$_2$SO$_4$, and H$_2$S$_2$ photoprocessing products were formed in our experiments. Proton bombardment of H$_2$S:H$_2$O ice led to formation of H$_2$S$_2$ and SO$_2$, while a similar processing of SO$_2$:H$_2$O ice formed H$_2$O$_2$, H$_3$O$^+$, HSO$_3^-$, HSO$_4^-$, and SO$_4^=$ (Moore et al. 2007). With the exception of H$_2$SO$_4$, all the above products of H$_2$S:H$_2$O ice photoprocessing desorbed between 100-200 K. S-polymers from S$_3$ to S$_8$ are also products of H$_2$S photoprocessing; these species can stick on grains even after sublimation of the ice mantle (Mu\\~noz Caro 2002, this work). In general, given the expected relative low abundances of S-polymers in ice mantles, they would be difficult to detect by infrared spectroscopy, but some of them will be observable with ALMA in the gas phase (Wakelam et al. 2005). The UV-irradiation of H$_2$O:CO:CH$_3$OH:NH$_3$:H$_2$S ices is left for future work. It was found that S$_8$ was by far the most abundant refractory product in these experiments, even for low initial H$_2$S ice abundances on the order of 2\\% relative to H$_2$O ice (Mu\\~noz Caro 2002). Including C in H$_2$S-containing ices, in the form of abundant ice components such as CH$_3$OH or CO, may lead to species like OCS, which is observed in circumstellar ice mantles."
    },
    "1112/1112.0902_arXiv.txt": {
        "abstract": "{ We review our knowledge of the origin and phenomenology of low states (also known as VY Scl states) in cataclysmic variables, making special emphasis on the cataclysmic variable population found just above the period gap. This orbital period range between approximately 3 and 4 hours is the preferred land for the elusive SW Sextantis stars, which are believed to be about to enter the gap when their donor stars become fully convective. Despite their main role in our understanding of the whole picture of cataclysmic variable evolution, the study of their component stars is almost impossible due to the extreme veiling brightness of the accretion disc during the high state. Here we present the first steps toward the characterisation of the white dwarfs and the donor stars in these VY Scl systems in the low state, and the discovery of sporadic accretion events and satellite emission lines like those observed in the AM Her systems in the low state, which are likely related to magnetic activity of the fast rotating secondary stars. ",
        "introduction": "Cataclysmic variables (CVs) are occasionally caught fading toward states of greatly diminished brightness, or `low states' (also known as VY Scl states). During these unusual quiescent stages CVs shine $\\sim$3-5 mag fainter than they do in the high state, and can stay at that level for days, months or even years before returning to the high state. Low states seem to occur independently of the magnetic field of the white dwarf (WD): they are observed in practically all AM Her stars (polars), some intermediate polars, and many weakly-magnetic CVs, including a large fraction of nova-likes and a number of dwarf novae and Z Cam stars. The causes and effects of low states have been previously reviewed by \\cite{king+cannizzo98}, \\cite{warner99}, and \\cite{hessman00}.   The main driver of the low-state phenomenology came from the study of the discless AM Her stars. Without an accretion disc any large brightness drop must be caused by a significant decrease in the magnetic accretion luminosity, which is proportional to the mass transfer rate from the donor star. Therefore, the typical photometric changes seen during low states are directly related to variations in the rate at which the donor star provides matter.  The picture, however, is not as simple as a brightness drop followed by a relatively long quiescence phase and a return to the high state. As Fig.~\\ref{fig1} shows CVs in the low state are always {\\em eager} to recover, displaying quick brightness changes between the high and the low state, and even getting stuck at intermediate states for a while.  As mentioned, there is broad agreement that mass transfer from the donor star through the L1 point temporarily stops or greatly reduces during a low state, but the exact mechanism still remains unknown. \\cite{livio+pringle94} proposed the accumulation of starspots close to the L1 point as a way of inhibiting Roche-lobe overflow. Only gas provided by the magnetic activity or stellar wind of the donor star would then be available for accretion \\citep[e.g.][]{hessmanetal00}. In fact, entanglement of the magnetic fields of both stars has been proposed to explain the line emission patterns observed in AM Her during the low state \\citep[][and references therein]{kafkaetal08}.  The observation of low states is therefore a unique opportunity to study the WDs and donor stars in CVs, especially the magnetic activity of the fast rotating donors, their solar-like activity cycles and the interplay with the strong WD magnetic field if present. On the other hand, the large mass accretion rate of nova-like CVs in the high state make their discs the dominating source of light, thus veiling the two stars and preventing dynamical measurements of their masses and other fundamental parameters such as the donor spectral type or the mean magnetic field of the WD from being done. It is only during deep low states that the WDs and the donor stars become visible for study. ",
        "conclusions": "Although we still don't exactly know what causes the low states in CVs, the observation with 8--10m class telescopes is providing new insights to the problem. Obtaining accurate stellar masses and radii in CVs thought to be close to enter the period gap is fundamental to test the current CV evolution theories. In addition, time-resolved, multifrequency observations during low states will help address open questions like the possible presence of a magnetic white dwarf or a remnant accretion disc in the CVs populating the 3--4 h orbital period region.   \\begin{figure} \\includegraphics[width=7cm,angle=-90]{PabloRodriguezGil_CVs_2011_01_fig08.ps} \\caption{ \\footnotesize H$\\alpha$ trailed spectra diagram showing the satellite lines and their orbital variation. Black corresponds to maximum emission.} \\label{fig5} \\end{figure}"
    },
    "1112/1112.2136_arXiv.txt": {
        "abstract": "A new three-body method is used to compute the rate of the triple-alpha capture reaction which is the primary source of $^{12}$C in stars. In this work, we combine the Faddeev hyperspherical harmonics and the R-matrix method to obtain a full solution to the three-body $\\alpha+\\alpha+\\alpha$ continuum. Particular attention is paid to the long range effects caused by the pairwise Coulomb interactions. The new rate agrees with the NACRE rate for temperatures greater than $0.07$ GK, but a large enhancement  at lower temperature is found ($\\approx 10^{12}$ at $0.02$ GK). Our results are compared to previous calculations where additional approximations were made. We show that the new rate does not significantly change the evolution of stars around one solar mass. In particular, such stars still undergo a red-giant phase consistent with observations, and no significant differences are found in the final white dwarfs. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.5185_arXiv.txt": {
        "abstract": "We investigate neutrino processes for conditions reached in simulations of core-collapse supernovae. Where neutrino-matter interactions play an important role, matter is partially degenerate, and we extend earlier work that addressed the degenerate regime. We derive expressions for the spin structure factor in neutron matter, which is a key quantity required for evaluating rates of neutrino processes. We show that, for essentially all conditions encountered in the post-bounce phase of core-collapse supernovae, it is a very good approximation to calculate the spin relaxation rates in the nondegenerate limit. We calculate spin relaxation rates based on chiral effective field theory interactions and find that they are typically a factor of two smaller than those obtained using the standard one-pion-exchange interaction alone. ",
        "introduction": "Understanding the mechanisms responsible for supernovae and neutron star formation requires a knowledge of the equation of state and transport mechanisms in matter at densities of the order of that in atomic nuclei and at temperatures ranging up to $10^{11}$K.  Despite almost half a century of work on the subject, the question of how a fraction of the large thermal energy from the collapse of the core is transferred to the outer stellar layers, thereby causing a supernova explosion, is one that has yet to find a convincing answer.  Because of the high matter density in the new-born protoneutron star, most macroscopic transport processes are ineffective, and a variety of other mechanisms have been considered. These include energy transfer by neutrinos that interact with matter via weak interactions \\citep{Colgate.White:1966,Bethe.Wilson:1985}, magnetic fields in combination with rotation \\citep{Ardeljan.Bisnovatyi.ea:2004,Kotake.Sato.Takahashi:2006}, convection \\citep{Herant.Benz.ea:1994}, or pressure waves and Alfv\\'en waves that may be emitted by the protoneutron star \\citep{Burrows.Livne.ea:2007,Sagert.Fischer.ea:2009}.  Some of these mechanisms lead to very characteristic features in the gravitational wave \\citep{Ott:2009} and neutrino signatures \\citep{Dasgupta.Fischer.ea:2010}, and these will be constrained by observations of the next galactic supernova. Other mechanisms have been tested in simulations, which demonstrate that, with current microphysical input, neutrino transport alone is insufficient to generate explosions \\citep{Liebendoerfer.Mezzacappa.ea:2001,% Rampp.Janka:2002,Thompson.Burrows.Pinto:2003}. However, it is possible to obtain explosions when neutrino heating and fluid instabilities are combined in axisymmetric models of stellar core collapse \\citep{Buras.Rampp.ea:2006,Marek.Janka:2009,Bruenn2009,% Suwa:2009py,Brandt:2010xa}.  In all the above mechanisms, transport by neutrinos is crucial: it is the dominant process for energy emission, it contributes to the transport of energy and lepton number, it influences the radial entropy and lepton fraction gradients that determine the stellar structure and fluid instabilities, and it gives rise to the neutrino luminosities and spectra, which will become the main observables for probing the properties of matter at high density in the next nearby supernova event.  Calculations of rates of neutrino processes in dense matter shortly after the discovery of weak neutral currents were reviewed by \\citet{Freedman}.  In the early work, the particles participating in the reactions were taken to be free, but, subsequently, effects of strong interactions were taken into account \\citep{sawyer, Iwamoto}. More recently, detailed calculations of rates of neutrino processes have been performed within a mean-field approach (the random phase approximation) by \\citet{Reddy,BurrowsSawyer,Reddy2,% BurrowsReddyThompson}. One effect not included in these calculations is that excitations can decay due to interactions in the dense medium. As stressed in \\citet{RaffeltSeckel,RaffeltSeckelSigl,HR}, this can lead to an energy transfer in neutrino processes considerably greater than that predicted on the basis of excitations with infinitely long lifetimes (i.e., from recoil effects alone). \\citet{LOP} showed how to include these damping effects in a mean-field approach, and a unified approach to structure factors was described by \\citet{LPS}. Detailed calculations were performed based on chiral effective field theory (EFT) interactions in \\citet{Bacca} for degenerate neutrons. The prime purpose of the present paper is to extend these studies to partially degenerate and nondegenerate conditions.  The plan of the paper is as follows. In Section~\\ref{conditions}, we analyze the results of simulations of core-collapse supernovae and determine the conditions for which it is important to know rates of neutrino processes.  These results point clearly to the need for a better understanding of neutrino properties in regions where nucleons are partially degenerate or nondegenerate. In Section~\\ref{formalism}, we develop a general formalism based on Landau's theory of normal Fermi liquids for calculating structure factors of strongly interacting matter. The spin relaxation rate for partially degenerate conditions is derived in Section~\\ref{time} and the nondegenerate limit is studied in Section~\\ref{nondegen}. A key quantity is the spin relaxation rate in partially degenerate neutron matter, and we calculate this in Section~\\ref{results} based on the one-pion-exchange approximation for nucleon-nucleon interactions, which is the standard one used in simulations \\citep{HR}, and from chiral EFT interactions. Particular attention is paid to conditions of importance in supernova simulations. Finally, we summarize and give future perspectives in Section~\\ref{summary}. ",
        "conclusions": "\\label{summary}  We have analyzed simulations of core-collapse supernovae and have shown that in the post-bounce phase the relevant conditions for neutrino processes are partially degenerate or nondegenerate. We then developed a formalism for calculating neutrino rates in strongly interacting matter by generalizing to the partially degenerate regime the approach used for degenerate matter in \\citet{LPS}. The resulting spin dynamical structure factor takes into account both mean-field effects and collisions between excitations in neutron matter. We then calculated the spin relaxation rate, a key ingredient in the structure factor which enters expressions for the rates of neutrino processes. This was done at two levels of NN interactions, the OPE approximation, which is commonly used in simulations \\citep{HR}, and chiral EFT interactions. We have found that chiral NN interactions lead to a reduction of the spin relaxation rate typically by a factor of two for a broad range of conditions. This reduction is similar to what previously has been found in the degenerate regime \\citep{Bacca}. We have also found that our OPE rate, where the width is obtained consistently by solving the Boltzmann equation, differs from the OPE results of \\citet{HR}. This may be a consequence of the imposition of a normalization condition in \\citet{HR}. Moreover, our results demonstrate that the nondegenerate limit is an excellent approximation for the conditions encountered in the post-bounce phase of matter at subnuclear densities in supernova simulations.  Future directions include the study of many-body contributions and of many-body forces and electroweak currents in chiral EFT, and the extension of the calculations to mixtures of neutrons and protons, where the central parts of nuclear interactions can cause relaxation of the axial charge, because of the different axial charges of the neutron and proton. These extensions will be greatly simplified by the finding that the nondegenerate limit provides an excellent approximation for the relevant supernova conditions.  Throughout the paper we have assumed that the basic coupling of the weak neutral field (that of the $Z$ boson) to nucleons is via a one-body operator. However, the spin relaxation effect that we have considered amounts to coupling a single quasiparticle-quasihole pair to two quasiparticle-quasihole pairs. In other words, strong interactions have generated a two-body contribution to the weak charge operator. However, not all two-body contributions to the operator are included in calculating relaxation effects by this procedure. The approximation we have made includes only contributions with a single quasiparticle in an intermediate state, which are enhanced by $1/\\omega$. Consideration of other two-body contributions to the weak charge operator is left for future studies."
    },
    "1112/1112.1009_arXiv.txt": {
        "abstract": "{} {A high density portion of the Orion Molecular Cloud 1 (OMC-1) contains the prominent, warm Kleinmann-Low (KL) nebula  % plus a farther region in which intermediate to high mass stars are forming. Its outside is affected by ultraviolet radiation from the neighboring Orion Nebula Cluster and forms the archetypical photon-dominated region (PDR) with the prominent bar feature. Its nearness makes the OMC-1 core region a touchstone for research on the dense molecular interstellar medium and PDRs.} {Using the Atacama Pathfinder Experiment telescope (APEX), we have imaged the line emission from the multiple transitions of several carbon monoxide (CO) isotopologues over the OMC-1 core region. Our observations employed the $2\\times7$ pixel submillimeter \\champp\\ array to produce maps ($\\sim300\\arcsec\\times350\\arcsec$) of \\TWCO, \\THCO, and \\CXVIIIO\\ from mid-$J$ transitions ($J=6-5$ to $8-7$). We also obtained the \\THCO\\ and \\CXVIIIO\\ $J=3-2$ images toward this region.} {The \\TWCO\\ line emission shows a well-defined structure which is shaped and excited by a variety of phenomena, including the energetic photons from hot, massive stars in the nearby Orion Nebula's central Trapezium cluster, active high- and intermediate-mass star formation, and a past energetic event that excites the KL nebula. Radiative transfer modeling of the various isotopologic CO lines implies typical \\HH\\ densities in the OMC-1 core region of $\\sim 10^{4}-10^{6}$ cm$^{-3}$ and generally elevated temperatures ($\\sim 50-250$ K). We estimate a warm gas mass in the OMC-1 core region of 86--285 \\Msol.} {} ",
        "introduction": "\\begin{table*}% \\centering \\caption{\\label{table1}Observational parameters} \\begin{tabular}{lcrrcccr} \\hline \\hline  Molecule/Line & Frequency\\tablefootmark{a} & $E_{\\rm up}/k$ & $\\theta_{\\rm MB}$ & $\\eta_{\\rm MB}$ & Receiver    & PWV        & $T_{\\rm{sys}}$\\\\ &(MHz)    & (K) & (\\arcsec)            & & & (mm) & (K)\\\\ \\hline \\CXVIIIO\\ $J=3-2$&   329330.553  & 32  &18.0   &  0.75  &  APEX-2a   & $\\sim1.0$      & 220  \\\\ \\THCO\\ $J=3-2$   &   330587.965  & 32  &18.0   &  0.75  &  APEX-2a   & $\\sim1.0$      & 210  \\\\ \\CXVIIIO\\ $J=6-5$&   658553.278  & 111 & 9.0   &  0.47  &  \\champp\\  & $\\lesssim0.7$  & 1600  \\\\ \\TWCO\\ $J=6-5$   &   691473.076  & 116 & 8.6   &  0.47  &  \\champp\\  & $\\lesssim0.5$  & 1800  \\\\ \\TWCO\\ $J=7-6$   &   806651.806  & 155 & 7.4   &  0.45  &  \\champp\\  & $\\lesssim0.5$  & 3900  \\\\ \\THCO\\ $J=8-7$   &   881272.808  & 190 & 6.7   &  0.45  &  \\champp\\  & $\\lesssim0.7$  & 3700  \\\\ \\hline  \\end{tabular}  \\tablefoot{Columns are, from left to right, CO isotopologue and transition, frequency, energy of the transition's upper energy level above the ground-state, the HPBW beam size, the main beam efficiency, the receiver used, the average precipitable water vapor column and the average system temperature during the observations. \\tablefoottext{a}{Taken from the Cologne Database for Molecular Spectroscopy \\citep[CDMS,][]{Mueller2005}\\footnote{http://www.astro.uni-koeln.de/cdms/}.} }    \\end{table*}   The Orion Molecular Cloud 1 (OMC-1) is a complex region of the interstellar medium (ISM) stretching over more than 2.4 pc ($20\\arcmin$, roughly north-south) on the sky \\citep{Kutner1976}. Its densest part, toward which the Great Orion Nebula (M42), a classical compact ''blister'' \\HII\\ region, and its associated Orion Nebular Cluster (ONC) appear in projection, is one of the best-studied regions in astronomy. It is a test bed for studies of (proto)stars and clusters and the formation of low-, intermediate-, and high-mass stars \\citep[for a review of this region, see][]{O'Dell2001,O'Dell2008}. Much of this region's prominence is due to its distance of just $414\\pm7$ pc \\citep{Menten2007}, which makes the ONC and OMC-1 the closest regions of recent (few million years old) and ongoing high-mass star formation. In the following, we shall use the term ``OMC-1 core'' or even just OMC-1 for the roughly $8\\arcmin\\times8\\arcmin$- or 1 pc$^{2}$-sized dense molecular cloud region, which is located closely (0.1--0.2 pc) behind M42 and most of the stars in the ONC \\citep[see][the latter give a comprehensive overview of this region and its phenomena]{Zuckerman1973,Genzel1989}.  The OMC-1 core region may be divided into three main zones, all of which show bright (sub)millimeter wavelength emission from warm dust and molecular gas: the Becklin-Neugebauer/Kleinmann-Low (BN/KL) region, Orion South (OMC-1S or Orion-S) and the Orion Bar.  The Orion BN/KL and Orion South regions are considered to be ``twin'' high-mass star forming regions because they have similar masses ($\\lesssim$ 100 \\Msol) and bolometric luminosities \\citep[$10^4-10^5$ \\Lsol,][]{Mezger1990,Drapatz1983} and show comparable levels of star-forming activity \\citep{O'Dell2008}. The BN/KL region harbors the eponymous ``hot core'' \\citep{Masson1984}, which was (and commonly still is) taken to be the prototype of the hot dense regions observed around many newly formed stars \\citep[see, e.g.,][]{Kurtz2000}. An interesting alternative explanation for this region's energetics (other than being powered by an embedded central heating source) is a protostellar merger event that released a few times $10^{47}$ erg of energy about 500 years ago \\citep{Bally2005,Zapata2011,Bally2011}.   The Orion Bar is a well-described photon-dominated region (PDR) located at the side of the OMC-1 core, facing M42, which is heated and partially ionized by far-ultraviolet (FUV) photons from the young massive stars (most of them from $\\theta^1$ C, a spectral type O5--O7 star) that form the ``Trapezium'' at the center of the ONC \\citep[see, e.g.,][]{Hollenbach1997,Walmsley2000}. In addition, the Orion Bar appears to be located at the edge of the \\HII\\ blister tangential to the line of sight.  On giant molecular cloud (GMC) scales, low-rotational level ($J$) \\TWCO\\ emission (commonly from the $J=1-0$ line) of the ambient, low density gas is usually used to trace the mass of the molecular ISM under a range of assumptions \\citep{Bloemen1984,Dame1985}. However, in massive star forming regions with much higher densities and temperatures, observations of the submillimeter and far-infrared (FIR) wavelength mid- or high-$J$ transitions of \\TWCO\\ and its $^{13}$C and $^{18}$O isotopologues are required for determinations of the gas temperature and density, usually in conjunction with Large Velocity Gradient (LVG) radiative transfer and PDR modelings.  Much of the \\TWCO\\ emission from the OMC-1 core has been proposed to arise from the neutral and partially ionized back side of the \\HII\\ blister, i.e., the PDR \\citep{Genzel1989}. The pioneering \\textit{submillimeter} observations of the \\TWCO\\ $J=7-6$, $6-5$ and \\THCO\\ $J=7-6$ transitions by \\citet{Schmid-Burgk1989} and \\citet{Graf1990} revealed high density gas ($n\\geq$ $10^{4}$ cm$^{-3}$) and elevated temperatures ($T\\geq$ 50 K) in all parts of the OMC-1 core region. Observations of the even more highly excited \\TWCO\\ $J=9-8$ line by \\citet{Marrone2004} with the Receiver Lab Telescope (84\\arcsec\\ resolution) showed that the hot and broad velocity emission in the line profiles arises mainly from the BN/KL region, while much of the narrower ($\\sim3$--6 \\kms\\ wide) emission arises from the PDR. This was also observed in the spectra of the \\TWCO\\ $J=9-8$ transition taken by \\citet{Kawamura2002}, which indicated warm molecular gas from the extended ``quiescent ridge'' region, north of Orion BN/KL. Furthermore, the observations of \\TWCO\\ $J=7-6$ and $J=4-3$ made by \\citet{Wilson2001} together with LVG modeling suggest that the broad line widths of these lines from the BN/KL region are probably due to shock heating, while most of the narrow line extended emission can be explained by PDR models. Recently, \\citet{Furuya2009} presented the BN/KL region images in the \\TWCO\\ $J=7-6$ line.    Clearly, multi-transition, high angular resolution \\textit{imaging} studies are needed to disentangle the relative contributions of shock and radiative excitation to the CO emission, and to identify possible energy sources. All of the spectral line observations described above report either single pointing toward selected positions in OMC-1 or limited mapping of special areas such as the Bar \\citep{Lis1998} or the BN/KL region \\citep{Furuya2009}. The only exceptions are the \\TWCO\\ $J=7-6$ map by \\citet{Schmid-Burgk1989} and the \\TWCO\\ $J=4-3$ and $J=7-6$ maps by \\citet{Wilson2001}. However, the first of these, which covers the whole OMC-1 core, was taken with the poor resolution of the Kuiper Airborne Observatory.  Here we present the first high (better than $10\\arcsec$) resolution large scale maps of the OMC-1 core using the $2\\times7$ pixel submillimeter \\champp\\ array receiver on the APEX telescope. Seven of the receiver units cover the 620 to 720 GHz frequency range and seven the 780 to 950 GHz range, allowing imaging of two lines simultaneously. As we shall see in \\S\\ref{obs}, we imaged the OMC-1 core in \\TWCO\\ $J=6-5$ plus $J=7-6$ in one run and \\THCO\\ $J=8-7$ plus \\CXVIIIO\\ $J=6-5$ in another. The fast mapping afforded by the 14 unit instrument resulted in a superior image consistency. The mid-$J$ CO data were complemented by the maps of the \\THCO\\ and \\CXVIIIO\\ $J=3-2$ lines. ",
        "conclusions": "\\subsection{Large scale maps -- overall morphology}    Figure \\ref{map1} shows $\\approx6\\arcmin\\times8\\arcmin$-size maps of the OMC-1 core in the intensity of the \\TWCO\\ $J=6-5$ and $J=7-6$ lines integrated over the local standard of rest (LSR) velocity interval from $-$25 to $+$30~\\kms. The emission distributions of both lines, clearly resolved in our observations, resemble each other closely and are consistent with the JCMT/SCUBA 850 \\micron\\ image obtained by \\citet{Johnstone1999}, where several regions are discernible. The most prominent of which are: (a) the emission maximum centered on the BN/KL region, (b) a secondary maximum coinciding with Orion South (OMC-1S), (c) a region $\\approx3\\arcmin$ east of the Trapezium, which is known as Orion East (OMC-1 E; see below), and (d), the prominent straight Orion Bar PDR. The bar actually appears connected to the Orion South region, forming a mirrored letter L shape. Furthermore, there are two minor maxima at $1\\farcm8$ N and $3\\farcm3$ NNE of the BN/KL peak position.  Apart from their spatial distributions, the regions above can be distinguished by their centroid LSR velocities, $V_{\\rm LSR}$, and their line widths, $\\Delta V$. The Orion BN/KL region displays a spectacularly broad line up to $200$~\\kms\\ wide, centered around 6~\\kms\\ with a narrow ($\\Delta V \\approx 5$~\\kms) feature at $V_{\\rm LSR}\\approx8$~\\kms\\ superposed. Orion South is characterized by $\\sim6$~\\kms\\ wide lines, with higher values at positions with outflow activities. The \\TWCO\\ $J=6-5$ and $J=7-6$ lines observed toward the Orion Bar have FWHM widths of $\\sim4$~\\kms\\ and are centered around $V_{\\rm LSR}=10$~\\kms. The Orion East region has previously been mapped by \\citet{Houde2004} and \\citet{Herrmann1997}, but otherwise has received relatively little attention. The observed line widths and LSR velocities of Orion East, as well the high intensity of its CO emission ($T_{\\rm{MB}}\\gtrsim 160$ K) clearly demonstrate its PDR nature. All our measurements are in good agreement with the results of \\citet{Schmid-Burgk1989}, \\citet{Graf1990}, \\citet{Wilson2001}, \\citet{Kawamura2002}, and \\citet{Marrone2004}. The measurements of the CO isotopologues in the OMC-1 core are summarized in Table \\ref{table2}.     Figure \\ref{map2} shows the integrated intensity of \\THCO\\ $J=8-7$ and \\CXVIIIO\\ $J=6-5$ maps in a similar area as \\TWCO\\ but in a smaller range of radial velocities (from $+$5 to $+$15 km s$^{-1}$ for \\THCO\\ and from $+$4 to $+$12 km s$^{-1}$ for \\CXVIIIO). Both maps also show a good correspondence but with the integrated emission from \\CXVIIIO\\ $J=6-5$ fainter than \\THCO\\ $J=8-7$. The two isotopologues have fainter and less broader lines compared with the \\TWCO\\ line emission which is expected for optically thick lines (see the spectra in Figs. \\ref{map1} and \\ref{map2}). The Orion Bar is well resolved in both lines and also shows a sharp edge. In the mid-$J$ images (Figs. \\ref{map1} and \\ref{map2}), the strong peak in the center of the Bar is also seen in the recent {\\it Herschel}/SPIRE high-$J$ \\TWCO\\ and \\THCO\\ data obtained by \\citet{Habart2010}. In addition, the line widths of \\THCO\\ $J=8-7$ and \\CXVIIIO\\ $J=6-5$ toward all four main regions are narrow, even in the outflow zones, i.e., Orion BN/KL and South.  The \\THCO\\ $J=3-2$ and \\CXVIIIO\\ $J=3-2$ integrated intensity maps are shown in Figure \\ref{map3}. These images were made in a similar area as higher-$J$ \\TWCO\\ and in a range of radial velocities equal to \\THCO\\ $J=8-7$ and \\CXVIIIO\\ $J=6-5$. The \\CXVIIIO\\ $J=3-2$ emission is only present where the \\THCO\\ $J=3-2$ emission is strong. The \\THCO\\ and \\CXVIIIO\\ line profiles at this rotational transition toward the Orion BN/KL and South regions are very broad and show well-defined wings. The southwestern part of the Bar is especially strong in this lower-$J$ transition for both isotopologues, which is also shown in the {\\it Herschel}/SPIRE high-$J$ \\TWCO\\ and \\THCO\\ images reported by \\citet{Habart2010}.     \\subsection{Velocity structure}     The \\TWCO\\ $J=6-5$ and $J=7-6$ velocity channel maps are shown in Figure \\ref{chmap}. These maps reveal a well-resolved and complex structure of the OMC-1 core at ambient velocities as well as higher radial velocities. The outflows that are located mainly at Orion BN/KL, Orion South, and toward the north of Orion BN/KL are clearly seen in these images \\citep[see][]{Zapata2006,Zapata2009}, e.g., \\vlsr=0--6 \\kms\\ for the BN/KL region and \\vlsr=12--14 \\kms\\ for the Orion South region. Besides, no clear north-south velocity gradients across the OMC-1 core region are observed.     The compact and warm structures toward the BN/KL region traced by the \\TWCO\\ $J=6-5$ and $J=7-6$ lines, which display a very broad range of velocities (up to about $\\pm100$ \\kms), are part of the enigmatic molecular outflow that seems to be produced by a violent explosion during the disruption of a massive young stellar system \\citep{Bally2005,Zapata2009,Bally2011}. Our observations reveal that some faint filamentary structures are very likely associated with the high-velocity \\TWCO\\ bullets (Peng et al, in prep.) reported by \\citet{Zapata2009}. %       \\subsection{Line ratios}  In Figures \\ref{ratio3}--\\ref{ratio2}, we show the line intensity ratio maps between the CO isotopologues and the different rotational transitions.  From the image of the \\TWCO\\ to \\CXVIIIO\\ $J=6-5$ ratio (Fig. \\ref{ratio3} upper panel), we can see the spatial distribution of high optical depth regions at different velocities. The north-south dense ridge from Orion BN/KL to Orion South is clearly seen at 9 \\kms. This high optical depth ridge is similar to the filamentary structures of the NH$_{3}$ emission \\citep{Wiseman1996,Wiseman1998}, but has some morphological differences in the north of the Orion BN/KL region. Besides, the straight shape of the Orion Bar is pronounced at \\vlsr=10--11 \\kms, where three high optical depth regions are seen at the two ends and center of the Bar at \\vlsr=10 and 11 \\kms, respectively. In the lower panel of Figure \\ref{ratio3}, the \\CXVIIIO\\ $J=6-5$ to $J=3-2$ ratios around 10--11 \\kms\\ show a gradient in the Orion Bar which goes in the direction from the Trapezium stars. It indicates that \\CXVIIIO\\ $J=6-5$ is strongly excited at the edge of the Bar by UV photons from the Trapezium stars.   \\begin{figure} \\centering \\includegraphics[angle=270,width=0.48\\textwidth]{Orion-A-12CO_7-6_6-5-ratio-chmap-LowRes.eps} \\caption{Velocity channel maps (5, 10, and 15 km s$^{-1}$) of the \\TWCO\\ $J=7-6$ to $J=6-5$ ratio ($\\theta_{\\rm HPBW}=10\\arcsec$) in the OMC-1 core. The 10 \\kms\\ image is overlaid with the \\TWCO\\ $J=6-5$ contours as shown in Fig. \\ref{map1}. The empty area are due to data blanking (3 $\\sigma$ level of the \\TWCO\\ $J=7-6$ temperature). The dashed-lines indicate that these ratio gradients are likely caused by outflows or filaments. Black triangles mark the Orion Hot Core position. \\label{ratio1}} \\end{figure}   \\begin{figure} \\centering \\includegraphics[angle=270,width=0.45\\textwidth]{all-inte-ratio-multi-maps.eps} \\caption{(a) The OMC-1 core \\CXVIIIO\\ $J=3-2$ to \\THCO\\ $J=3-2$ integrated intensity ratio map ($\\theta_{\\rm HPBW}=20\\arcsec$) overlaid with the \\TWCO\\ $J=6-5$ contours (grey) as shown in Fig. \\ref{map1}. The black contour represents the \\CXVIIIO/\\THCO\\ abundance ratio of 0.122 (60/490). (b) The OMC-1 core \\CXVIIIO\\ $J=3-2$ to $J=6-5$ integrated intensity ratio map ($\\theta_{\\rm HPBW}=20\\arcsec$) overlaid with the \\TWCO\\ $J=6-5$ contours as shown in Fig. \\ref{map1}. Black triangles mark the Orion Hot Core position. \\label{ratio2}} \\end{figure}    Figure \\ref{ratio1} shows the ratio between \\TWCO\\ $J=7-6$ and $J=6-5$ at three different radial velocities of 5, 10, and 15 \\kms. The three panels show clear variations across the integrated line emission. It is interesting to note that at the cloud velocity of 10 km s$^{-1}$, there are high 7--6/6--5 ratios located very close to the position of the Trapezium stars toward Orion BN/KL. These gradients are likely produced by those massive stars that heat the molecular cloud, causing the stronger \\TWCO\\ $J=7-6$ line emission compared with the $J=6-5$ line. A gradient is also seen in the Orion Bar which goes in the perpendicular direction to the Trapezium stars (Fig. \\ref{ratio1} at $V_{\\rm LSR}=10$ \\kms), and some patches with higher \\TWCO\\ $J=7-6$ brightness temperatures are seen inside the Bar or behind the ionization front. Some horizontal and vertical strips, artifacts from the OTF mapping, are also seen in the Orion Bar and South regions. These artifacts can affect the intensity ratio by $\\sim 14 \\%$ given the calibration errors of $10 \\%$ for both \\TWCO\\ $J=6-5$ and $J=7-6$. However, some structures seen in the \\TWCO\\ $J=7-6$ to $J=6-5$ ratio maps, e.g., the dashed-lines in Figure \\ref{ratio1}, are unlikely due to the calibration error or scanning strips during the observation, and seem to be footprints of filaments or outflows. Some similar filament structures have been noticed in the dust continuum emission observed by \\citet{Johnstone1999}.        In the ratio map of \\CXVIIIO\\ $J=3-2$ to \\THCO\\ $J=3-2$ shown in Figure \\ref{ratio2}, a clear elongated and dense ridge is seen in the north-south direction, where the \\THCO\\ $J=3-2$ intensity is weaker in the Orion BN/KL region than at Orion South and the north of BN/KL. Therefore, the lower CO column density derived from the lower \\CXVIIIO\\ $J=3-2$ to \\THCO\\ $J=3-2$ ratio toward Orion BN/KL may be misleading. Even though the \\THCO\\ $J=3-2$ line is optically thick and may be self-absorbed, this will result in a higher \\CXVIIIO\\ $J=3-2$ to \\THCO\\ $J=3-2$ ratio instead of a lower ratio in Orion BN/KL. As \\citet{Goldsmith1997} pointed out, the column density calculated using low-$J$ \\CXVIIIO\\ lines is accurate except for the temperature $\\geq150$ K, which is the case in the OMC-1 core. Hence, higher-$J$ CO observations are critical to determine the CO column density in the OMC-1 core region. In addition, already shown in Figure \\ref{ratio1}, the low \\CXVIIIO\\ $J=3-2$ to $J=6-5$ ratios indicate that the \\CXVIIIO\\ $J=6-5$ intensity is strongly enhanced toward the Orion BN/KL, South, and Bar regions, and is likely due to their PDR nature. Nevertheless, shocks/outflows may also play an important role in the Orion BN/KL and South regions, leading to an extra heating in the cloud.    \\subsection{Excitation temperature and density estimates}             The \\TWCO\\ $J=6-5$ and \\CXVIIIO\\ $J=6-5$ data are used here to derive the gas excitation temperature and density of the OMC-1 core. In the submillimeter regime where the Rayleigh-Jeans approximation is often not valid, the observed radiation temperature in local thermodynamic equilibrium (LTE) can be expressed as \\begin{equation}\\label{cal-orion-1} {T^{*}_{\\rm R}}=\\frac{h\\nu}{k}\\biggl[\\frac{1}{e^{h\\nu/kT_{\\rm ex}}-1}-\\frac{1}{e^{h\\nu/kT_{\\rm bg}}-1}\\biggr](1-e^{-\\tau_{\\nu}}), \\end{equation} where $T_{\\rm ex}$ is the excitation temperature and $\\tau_{\\nu}$ is the optical depth at a specific molecular line transition. The background brightness temperature $T_{\\rm bg}$ includes the cosmic background radiation of 2.73 K and the radiation from warm dust, ranging from about 5 K for less dense gas to about 26 K for the Orion Hor Core \\citep{Goldsmith1997}. The second term inside the brackets can be neglected since it only contributes $\\lesssim2\\%$ of $T^{*}_{\\rm R}$. By making the assumption of the same excitation temperature for \\TWCO\\ and \\CXVIIIO\\ $J=6-5$, the optical depths of \\TWCO\\ and \\CXVIIIO\\ $J=6-5$ can be determined from the relation \\begin{equation}\\label{cal-orion-2} \\frac{T^{*}_{\\rm R}(\\rm{^{12}CO)}}{T^{*}_{\\rm R}(\\rm{C^{18}O)}}\\approx\\frac{1-e^{-\\tau(\\rm{^{12}CO})}}{1-e^{-\\tau(\\rm{C^{18}O})}}. \\end{equation} Since the optical depth of \\TWCO\\ $J=6-5 \\gg 1$, the optical depth of \\CXVIIIO\\ $J=6-5$ can be directly obtained ($\\sim 0.08$) by assuming that the optical depth ratio is approximated to the isotopologic abundance ratio, and here we adopted [\\TWCO]/[\\CXVIIIO] of 490 \\citep[][]{Boreiko1996,Wilson1992} and a beam filling factor of unity ($T^{*}_{\\rm R}=T_{\\rm MB}$). The excitation temperatures of \\TWCO\\ and \\CXVIIIO\\ $J=6-5$ can be estimated via \\begin{equation}\\label{cal-orion-3} T_{\\rm ex}=\\frac{h\\nu}{k}\\biggl[{\\rm ln}\\biggl(1+\\frac{h\\nu}{kT_{\\rm MB}\\rm(^{12}CO)}\\biggr)\\biggr]^{-1}. \\end{equation} Then the total \\CXVIIIO\\ column density can be derived from \\begin{equation} N({\\rm C^{18}O})=\\frac{3kQ_{\\rm rot}}{8\\pi^{3}\\nu S\\mu^{2}}e^{\\frac{E_{\\rm up}}{kT_{\\rm ex}}}\\int T_{\\rm MB}dV \\end{equation} \\begin{equation} \\simeq1.33\\times10^{12}\\ (T_{\\rm ex}+0.88)\\ e^{\\frac{E_{\\rm up}}{kT_{\\rm ex}}}\\int T_{\\rm MB}dV\\ \\rm{cm^{-2}}, \\end{equation} where $E_{\\rm up}$ is 110.6 K for the \\CXVIIIO\\ $J=6-5$ transition and $Q_{\\rm rot}$ is the rotational partition function. In the end, the \\HH\\ column density can be estimated by adopting a [\\TWCO]/[\\CXVIIIO] abundance ration of 490 and a \\TWCO\\ abundance of $8\\times10^{-5}$ \\citep{Wilson1992}.            \\begin{figure*} \\centering \\includegraphics[angle=270,width=0.9\\textwidth]{Orion-A-12CO-C18O_6-5-den-ext-map.eps} \\caption{(a) The \\CXVIIIO\\ column density image of the OMC-1 core overlaid with black contours of the JCMT 850 \\micron\\ dust continuum image. The 850 \\micron\\ image is taken from the JCMT data archive. The contours represent 0.2$\\%$, 0.5$\\%$, 1$\\%$, 2$\\%$, 4$\\%$, 7$\\%$, 20$\\%$, 40$\\%$, 60$\\%$, and 80$\\%$ of the peak intensity (87 Jy beam$^{-1}$). (b) The excitation temperature image for \\TWCO\\ and \\CXVIIIO\\ overlaid with the 35 mm VLA-GBT continuum emission \\citep{Dicker2009} in black contours running from 15$\\%$ to 95$\\%$ in steps of 8$\\%$ of the peak intensity (0.94 Jy beam$^{-1}$). The stars mark the positions of the five Trapezium stars, and black triangles represent the position of the Orion Hot Core.\\label{den-temp}} \\end{figure*}   The results of the \\HH\\ column density and excitation temperature distributions are shown in Figure \\ref{den-temp}, where the north-south dense ridge near the Trapezium cluster and the Orion Bar in the southeast are clearly seen. The average $\\tau(\\rm C^{18}O)$ at the $J=6-5$ transition over the whole OMC-1 core region is about 0.08 with an average excitation temperature of about 115 K, which are consistent with the \\TWCO\\ $J=9-8$ observations by \\citet{Kawamura2002}. Besides, the minimum excitation temperature in the OMC-1 core is $\\sim 30$ K, indicating a generally warm environment. The average \\CXVIIIO\\ column density in the OMC-1 core is about $8.7\\times10^{15}$ cm$^{-2}$, corresponding to an \\HH\\ column density of $5.5\\times10^{22}$ cm$^{-2}$ assuming a \\TWCO\\ abundance of $8\\times10^{-5}$ \\citep{Wilson1992}. Therefore, the dense gas mass is $\\sim 140$ \\Msol\\ in an area of 0.13 pc$^2$ in the OMC-1 core. The \\HH\\ column density agrees with the density of $5\\times10^{21}-7\\times10^{22}$ cm$^{-2}$ derived from the mid-$J$ CO isotopologue observations by \\citet{Wirstrom2006}. Our results agree with the recent measurement by \\citet{Wilson2011} using the \\THCO\\ $J=6-5$ line with a lower [\\TWCO]/[\\HH] ratio of $2\\times10^{-5}$. Additionally, \\citet{Habart2010} derived an \\HH\\ column density of about $9\\times10^{22}$ cm$^{-2}$ toward the Orion Bar from the {\\it Herschel}/SPIRE high-$J$ CO observations, assuming a lower excitation temperature of 85 K and a higher \\TWCO\\ to \\CXVIIIO\\ ratio of 560. Their result is consistent with ours, where we obtained  an \\HH\\ column density of $\\approx10^{23}$ cm$^{-2}$ toward the Orion Bar with a higher excitation temperature ($\\gtrsim150$ K).  \\begin{figure*} \\centering \\includegraphics[angle=270,width=0.9\\textwidth]{radex-co-13co-c18o-grid-dens-peak.eps} \\caption{The RADEX modeling results shown in the \\TWCO\\ $J=6-5$ peak temperatures and the \\THCO/\\CXVIIIO\\ ratios for three different \\TWCO\\ column densities. The column densities of \\THCO\\ and \\CXVIIIO\\ were fixed at the isotopologic abundance ratio of $1/60$ and $1/490$ of the \\TWCO\\ column density. The line widths of \\TWCO, \\THCO, and \\CXVIIIO\\ were fixed at 5, 5, and 4 \\kms, respectively. The red and blue contours denote the temperature and \\HH\\ number density, respectively, and the color points represent the data in the different regions of the OMC-1 core. Each point represents a peak \\TWCO\\ $J=6-5$ temperature and a temperature ratio between \\THCO\\ $J=8-7$ and \\CXVIIIO\\ $J=6-5$ from a single pixel with a resolution of 20$\\arcsec$. (a) The modeling input column density of \\TWCO\\ is $1\\times10^{18}$ cm$^{-2}$, and the data are plotted in a \\TWCO\\ column density density of $7.5\\times10^{17}-2.5\\times10^{18}$ cm$^{-2}$. (b) The modeling input column density of \\TWCO\\ is $5\\times10^{18}$ cm$^{-2}$, and the data are plotted in a density range of $2.5\\times10^{18}-7.5\\times10^{18}$ cm$^{-2}$. (c) The modeling input column density of \\TWCO\\ is $1\\times10^{19}$ cm$^{-2}$, and the data are plotted in the density of $\\geq7.5\\times 10^{18}$ cm$^{-2}$. Upper right panel shows the \\TWCO\\ column density map toward Orion BN/KL in contours of 5, 10, 20, 30, and $40\\times 10^{18}$ cm$^{-2}$. The stars mark the positions of the five Trapezium stars, and the Orion Hot Core (HC) position is also marked. \\label{radex}} \\end{figure*}       The JCMT 870 $\\mu$m dust continuum emission image \\citep{Johnstone1999} shown in Figure \\ref{den-temp} agrees with our density map well, where the dust emission traces mostly the dense ridge (Orion BN/KL and South) and Bar regions. It is interesting to note that the dense ridge has an offset from the peak temperature positions, especially in the Orion South region ($\\sim30\\arcsec$) and the north of the Orion BN/KL region ($\\sim20\\arcsec$). This dense ridge seems to extend farther to the north of the OMC-1 core, and is probably related to the NH$_3$ filamentary structure seen by \\citet{Wiseman1998,Wiseman1996}. Additionally, as Figure \\ref{den-temp} (b) demonstrates, the interface between the ionized and the molecular warm dense gas is very pronounced, where Orion BN/KL and Orion South are located at the north-south dense ridge, and the Orion Bar and Orion East are part of the high temperature enclosed structure well illustrated by the free-free continuum emission. Furthermore, our large-scale images are consistent with the [\\CII] and [\\OI] emission maps obtained by \\citet{Herrmann1997}, where the rather uniform [\\CII] emission indicates that PDRs are present over the whole OMC-1 core region.        \\subsection{RADEX modeling}       The non-LTE radiative transfer program RADEX \\citep{vanderTak2007} was used to investigate the accuracy of the temperature and density calculation in LTE shown above. The optical depth effects are taken into account by RADEX using the escape probability approximation, where a uniform sphere geometry was chosen. The different transitions of \\TWCO, \\THCO, and \\CXVIIIO\\ were used in the modeling, and their line widths were fixed at 5, 5, and 4 \\kms, respectively, which are the average values from our spectra. The molecular data using in the modeling are taken from the Leiden Atomic and Molecular Database\\footnote{http://www.strw.leidenuniv.nl/$\\sim$moldata/} \\citep[LAMDA;][]{Schoeier2005}. In addition, only \\HH\\ was chosen as a collision partner, and 2.73 K was adopted as the background temperature. The input parameters are kinetic temperature and \\HH\\ number density. Each model was iterated with a kinetic temperature ranging from 50 to 450 K and an \\HH\\ number density ranging from $10^4$ to $10^8$ cm$^{-3}$. The input column densities of these three molecules were chosen, so that their ratios were fixed at their isotopologic abundance ratios, i.e., 490 for [\\TWCO]/[\\CXVIIIO] and 60 for [\\TWCO]/[\\THCO]. Three models were iterated with three different \\TWCO\\ column densities of $1\\times10^{18}$, $5\\times10^{18}$, and $1\\times10^{19}$ cm$^{-2}$ which were obtained from the LTE calculation. Since these models did not take outflows into account, the output molecular line radiation temperatures ($T_{\\rm R}$) were directly compared with the peak temperatures instead of the integrated intensities to avoid the strong line wing emission. The modeling results are shown in Figure \\ref{radex}.    Figure \\ref{radex} reveals that most regions in the OMC-1 core have an \\HH\\ number density of $\\sim 10^{4}-10^{6}$ cm$^{-3}$. The Orion Bar and Orion East emission peaks (Fig. \\ref{radex} a), however, show higher densities. The distribution of the \\HH\\ number density in the Bar is consistent with $10^{4}-10^{5}$ cm$^{-3}$ for a homogeneous medium \\citep{Wyrowski1997}. The modeling results in the Orion BN/KL, South, and Bar regions all show similar trends of a diverse kinetic temperature (about 100--200 K), which are consistent with the LTE calculation results. In addition, the non-LTE modeling results indicate that the central region of Orion BN/KL close to the Hot Core (in 20\\arcsec\\ size) has a number density of about $10^5$ cm$^{-3}$ and a temperature of about 200 K, which is also close to the excitation temperature calculated in the LTE case. Noteworthily, the Orion BN/KL region and part of the Orion South region (Fig. \\ref{radex} a) show very high kinetic temperatures ($>350$ K) which may indicate an extra heating mechanism, e.g., outflow/shock heating. The suggestion of outflow/shock heating is supported by the larger \\TWCO\\ line widths (Fig. \\ref{map1}) and a great amount of known outflows in these regions \\citep[see][]{Henney2007}. However, the high kinetic temperatures ($\\geq250$ K) in the Orion Bar and Orion East regions (Fig. \\ref{radex} a) are hardly explained by the outflow/shock heating for the lack of outflow activities or broad \\TWCO\\ line widths. Instead, as PDRs, the Orion Bar and Orion East are expected to be heated by FUV photons from the Trapezium OB stars. Moreover, a recent study of the Orion Bar \\citep{Pellegrini2009} suggests an extra heating by the excess density of cosmic rays, which are trapped in the compressed magnetic field.  \\subsection{Mass of the OMC-1 core}  By adopting a radius of 40\\arcsec--60\\arcsec\\ and an \\HH\\ density of $5\\times10^{4}-2\\times10^{5}$ cm$^{-3}$ for Orion BN/KL, we can estimate a gas mass of 7--85 \\Msol\\ with the assumption of a spherical volume at a distance of 414 pc \\citep{Menten2007}. In a similar way, the gas mass in Orion South is 3--49 \\Msol\\ with a 30\\arcsec--50\\arcsec\\ radius and an \\HH\\ density of $5\\times10^{4}-2\\times10^{5}$ cm$^{-3}$. For Orion East, the \\HH\\ density is $\\sim2\\times10^5$ cm$^{-3}$ with a 15\\arcsec--25\\arcsec\\ radius, and the gas mass is 2--7 \\Msol. In contrast, it is very difficult to estimate the gas mass in the Orion Bar because of its geometry. We adopted a cylindrical geometry and an \\HH\\ density of $5\\times10^{4}-2\\times10^{5}$ cm$^{-3}$ as the lower limit since an edge-on plane may contain much more gas. With a length of 300\\arcsec\\ and a radius of 20\\arcsec--30\\arcsec, the gas mass of the Orion Bar is 9--79 \\Msol. The estimated total dense \\HH\\ gas mass in the OMC-1 core is 21--220 \\Msol\\ within a radius of 0.3 pc. This is also consistent with the mass estimated ($\\sim140$ \\Msol) from the LTE calculation above. If we include the warm diffuse gas ($10^{4}$ cm$^{-3}$) in the same volume, the total warm gas mass is 86--285 \\Msol\\ in the OMC-1 core. This  mass estimate for the OMC-1 core  is close to that derived from the Odin \\TWCO\\ and \\THCO\\ $J=5-4$ observations \\citep{Wirstrom2006}, i.e., 320 \\Msol\\ in the molecular ridge for a farther distance of 500 pc. The result of \\citet{Wilson2001} also suggests a similar estimate of a warm gas mass of 310--430 \\Msol\\ from the observations of \\TWCO\\ $J=7-6$ and $J=4-3$, again for a distance of 500 pc. Scaling these two literature values % to the distance of 414 pc, results in 213--295 \\Msol\\ for the total mass of the OMC-1 core. This is comparable to the value derived by us. Hence, we conclude that the higher-$J$ CO lines do trace most of the gas in the OMC-1 core which is heated mainly by the Trapezium stars."
    },
    "1112/1112.3710_arXiv.txt": {
        "abstract": "A growing  number of recent observations have revealed that the Galactic globular cluster (GC) $\\omega$ Cen is not the only GC that shows abundance spread in heavy elements (e.g., Fe). In order to understand the origin of the Galactic GCs with heavy element abundance spread (``HEAS''), we investigate the formation processes of massive GCs (MGCs) with masses larger than $10^{6} {\\rm M}_{\\odot}$ in gas-rich dwarf galaxies interacting and merging with the very young Galaxy. We find that massive and compact stellar clumps with masses larger than $10^{6} {\\rm M}_{\\odot}$, which can be regarded as progenitors of MGCs, can form from massive gas clumps  that are developed through merging of gaseous regions initially at different radii  and thus with different  metallicities. Therefore it is inevitable that MGCs formed in dwarfs have HEAS. The abundance spread in each individual MGC depends on the radial metallicity gradient of the host  dwarf such that it can be larger for the steeper metallicity gradient. For example,  MGCs formed in a dwarf with a central metallicity of [Fe/H]$=-1.1$ and the radial gradient of $\\sim -0.2$ dex kpc$^{-1}$ can have the abundance spread of $\\Delta {\\rm [Fe/H]} \\sim 0.2$. The simulated MGCs appear to be significantly flattened owing to their dissipative formation from gas disks of their host dwarfs. Based on these results, we discuss possibly diverse formation mechanisms for the Galactic GCs such as M22, M54, NGC 2419, $\\omega$ Cen, and Terzan 5. ",
        "introduction": "One of remarkable  recent developments in observational studies of the Galactic GCs is that most of the investigated GCs show varying degrees of chemical abundance spread: helium abundance spread in  $\\omega$ Cen and NGC 2808 (e.g., Bedin et al. 2004; Piotto et al. 2007), a larger dispersion in $s$-process elements abundances in  NGC 1851 (e.g., Yong \\& Grundahl 2008; Milone et al. 2011), abundance spread in light element  in ``normal'' GCs (e.g., Carretta et al. 2010a), and HEAS in M22, Terzan 5, and NGC 2419 (e.g., Da Costa et al. 2009; Ferraro et al. 2009; Marino et al. 2009; Cohen et al. 2010). Lee et al. (2009) investigated color-magnitude diagrams of stars in $hk$-bands for the Galactic GCs and suggested that a significant fraction of the Galactic GCs have two different populations with HEAS (but see also Marino et al. 2009; Carretta et al. 2010b  for similar and different claims).        The origin of the HEAS in $\\omega$ Cen has been discussed by theoretical models in the context of formation and evolution of stellar galactic nuclei in dwarfs (``the stripped nucleus scenario'';  e.g., Bekki \\& Freeman 2003; Romano et al. 2010). Since stellar populations of stellar galactic nuclei can have different metallicities and ages owing to their possible formation by merging of different stellar and gaseous clumps (e.g., Bekki 2007), it would  not be so  surprising that $\\omega$ Cen is observed to  have HEAS owing to its possible origin from the stellar nucleus of its host dwarf. Also, it would be possible that the formation process of other GCs with HEAS (e.g., Terzan 5) can be different from that of $\\omega$ Cen (e.g., GC mergers;   Br\\\"uns \\& Kroupa 2011; Bekki \\& Yong 2011, BY11). Ideally speaking,  chemodynamical simulations  of GC formation based on a theoretical model are essential for discussing both the formation processes of GCs from gas and the resultant internal chemical abundance spread within GCs in a fully self-consistent manner. However no chemodynamical simulations on the origin of GCs with HEAS (other than $\\omega$ Cen) have been done so far.    The purpose of this Letter is to clearly show whether and how GCs with  HEAS can be formed by using chemodynamical simulations of GC formation  processes based on a  GC formation scenario. We here consider that most of the Galactic GCs were formed in galactic building blocks (e.g.,  Searle \\& Zinn 1978), in particular, when they were interacting and merging with the very young Galaxy ($>10$ Gyr ago). We focus exclusively  on [Fe/H] spread ($\\Delta {\\rm [Fe/H]}$) in MGCs with masses as large as or larger than $10^6 {\\rm M}_{\\odot}$ formed  within their host dwarfs (i.e., galactic building blocks).    \\begin{figure*} \\psfig{file=f1.eps,width=18.0cm} \\caption{ The time evolution of the mass distribution of the dwarf with respect to the Galactic center (left four) and to the dwarf's center (right four) projected onto the $x$-$y$ plane for the standard model (M1). Gas and new stars are shown by cyan and magenta, respectively, and $T$, which is the time that has elapsed since the simulation started, is shown in the upper left corner in units of Gyr for each panel. The green lines in the left four panels represent the orbit of the dwarf. The locations of MGC1, 2, and 3 are indicated by blue, red, and green circles at $T=1.7$ Gyr, respectively. } \\label{Figure. 1} \\end{figure*} ",
        "conclusions": "The present chemodynamical study has first demonstrated that MGCs with HEAS can be formed from massive gas clumps developed from merging of different gaseous regions with different metallicities. Therefore, MGCs with HEAS can be formed even without chemical enrichment by supernovae during their formation. The present study suggests that more massive GCs are more likely to have a larger degree of HEAS owing to their formation from gas clouds developed from merging between a larger number of different gaseous regions with different metallicities. The low-mass GCs can be formed from single gas clouds so that they are likely to have no/little HEAS. The present study thus suggests that the origin of HEAS in some Galactic GCs can be closely associated with HEAS of ISM in their host dwarfs.  The large metallicity spread of various elements and possible age variation in $\\omega$ Cen (e.g., Sollima et al. 2005) can be consistent with the stripped nucleus scenario.  However, this does not necessarily mean that all of the Galactic GCs with HEAS (e.g., NGC 2419, M22, and Terzan 5) were formed from stripped nuclei of dwarfs. A GC with distinct two peaks in the [Fe/H] distributions of the stars might well be consistent with GC merging with different [Fe/H] (e.g., Br\\\"uns \\& Kroupa 2011; BY11). A GC with a smaller  degree of HEAS yet no clear bimodal [Fe/H] distribution might well form from massive gas clumps of the host dwarf and then be stripped from the dwarf without sinking into the center (i.e., without  becoming the stellar nucleus) to finally become the Galactic halo  GC. We lastly suggest that  nucleation,  GC merging,  and merging of gas clouds with different [Fe/H] can  be all promising mechanisms  for the formation of the Galactic GCs with different degrees of HEAD."
    },
    "1112/1112.2520_arXiv.txt": {
        "abstract": "{High-mass microquasars consist of a massive star and a compact object, the latter producing jets that will interact with the stellar wind. The evolution of the jets, and ultimately their radiative outcome, could depend strongly on the inhomogeneity of the wind, which calls for a detailed study.} {The hydrodynamics of the interaction between a jet and a clumpy wind is studied, focusing on the global wind- and single clump-jet interplay.} {We have performed, using the code \\textit{Ratpenat}, three-dimensional numerical simulations of a clumpy wind interacting with a mildly relativistic jet, and of individual clumps penetrating into a jet.} {For typical wind and jet velocities, filling factors of about $\\ga 0.1$ are already enough for the wind to be considered as clumpy. An inhomogeneous wind makes the jet more unstable when crossing the system. Kinetic luminosities $\\sim 10^{37}$~erg/s allow the jet to reach the borders of a compact binary with an O star, as in the smooth wind case, although with a substantially higher degree of disruption. When able to enter into the jet, clumps are compressed and heated during a time of about their size divided by the sound speed in the shocked clump. Then, clumps quickly disrupt, mass-loading and slowing down the jet.} {We conclude that moderate wind clumpiness makes already a strong difference with the homogeneous wind case, enhancing jet disruption, mass-loading, bending, and likely energy dissipation in the form of emission. All this can have observational consequences at high-energies and also in the large scale radio jets.} ",
        "introduction": "\\label{intro}  Microquasars are binary systems hosting a star and an accreting black hole or neutron star. Matter from the star  is transferred to the compact object, part of it being launched through magnetocentrifugal forces \\citep[e.g.,][]{bla77,bla82,bar11}.  This triggers the formation of bi-polar jets, which generate non-thermal radio emission \\citep[e.g.,][]{mir99,rib05}, and are thought to be the location from where the gamma rays observed in some  sources are emitted \\citep[e.g.,][]{alb07,tav09,abd09,sab10}. Jets could be magnetically dominated at their base, but magnetohydrodynamical processes occurring at higher jet height would accelerate the flow,  efficiently converting magnetic energy into kinetic one \\citep[e.g.][]{kom07}.  At the scales of the binary system ($\\sim 10^6\\,R_{\\rm Sch}$, where $R_{\\rm Sch}$ is the Schwarzschild radius), the jet is likely to be already a hydrodynamical (HD) flow. We focus here on the persistent jets thought to be present during the low-hard state of microquasars, although some considerations for transient ejecta, associated to low-hard to high-soft state transitions, are done below \\citep[see, e.g.,][for reviews on microquasar states]{fen04,fen09}.  Part of the energy carried by the jet can be dissipated in the form of magnetic reconnection, recollimation and internal shocks, shear layers in the jet walls, and turbulence. Part of the dissipated energy can go to non-thermal particles, generating low- and high-energy emission via different mechanisms, synchrotron from radio to X-rays, and inverse Compton (IC) and hadron-related processes up to gamma rays \\citep[see, e.g.,][and references therein]{bos09}.  As shown by \\cite{per08} (PB08 hereafter) and \\cite{per10} (PBK10 hereafter), in microquasars hosting an OB star (high-mass microquasars; hereafter HMMQ) the jet may be strongly influenced by the stellar wind. The one-side impact of the wind on the (presumably) already HD jet leads to strong and asymmetric recollimation shocks, bending and different types of instabilities: the recollimation shocks seem suitable candidates for particle acceleration and non-thermal emission; bending may be noticeable in radio at milliarcsecond scales; instabilities may destroy the jet flow even within the binary system. For typical wind and jet velocities, say $v_{\\rm w}\\sim 2\\times 10^8$ and $v_{\\rm j}\\sim 10^{10}$~cm/s, respectively, studies show that for compact binaries and jet-to-wind momentum flux ratios $\\la 0.1$ the jet can be already disrupted (PB08, PBK10). This number is rather constraining, since only a few HMMQ might be above this threshold. This could be the reason for the low number of HMMQ detected, as suggested in PBK10. In any case, even if not destroyed within the binary system, jets can suffer strong perturbations with both dynamical and radiative consequences.  Previous work in HMMQ wind-jet interactions was done under the assumption that the wind is homogeneous, but in fact stellar winds are thought to be clumpy \\citep[e.g.,][]{owo06,mof08}. For this reason, it has been proposed that wind clumpiness should be taken into account when studying HMMQ (\\citealt{owo09,ara09} -ABR09 hereafter-; \\citealt{rom10,ara11}).  An important parameter that determines the inhomogeneity of the wind is the wind filling factor $f$, which determines the wind volume fraction with higher density. For a significant departure from homogeneity, the intraclump medium mass and momentum fluxes will be negligible and only clumps will have a dynamical impact on the jet. Since the interaction between a HD jet and a clumpy wind has not been studied in detail, we have carried out 3-dimensional (3D) simulations of this scenario. Simulations have been done for two different jet powers and jet-to-wind momentum ratio, $L_{\\rm j}=3\\times 10^{36}-10^{37}$~erg/s and $\\approx 0.03-0.08$, respectively, to explore what could be the transition between jet destruction and long-term collimation. Another simulation has focused on the evolution of individual clumps injected in the jet at different heights. Unlike in PBK10, in which the simulation started with the jet being injected at its base, here the jet is conical and crossing the whole grid, and the clumpy wind is injected from one of the jet sides.  The first goal of this work is the study of the hydrodynamical evolution of a jet when the wind interacting with it is clumpy. The second goal is to quantify for which values of the clump and jet parameters, clumpiness becomes a relevant factor. It is also interesting to study the evolution of a clump under the impact of a microquasar jet. The results can also be used to refine radiation models or interpret radio observations, although this will be treated qualitatively. The paper is organized as follows: in Section~\\ref{phys}, the scenario studied here is briefly introduced; in Sect.~\\ref{sim}, the simulations are described (Sect.~\\ref{sim1}); results are shown in Sect.~\\ref{sim2}; finally, in Sect.~\\ref{disc}, the results are discussed in the context of jet propagation in HMMQ (Sect.~\\ref{disc1}), individual clump-jet interactions (Sect.~\\ref{disc2}), and their implications for the non-thermal emission (Sect.~\\ref{disc3}). Throughout the paper, we will use cgs units. ",
        "conclusions": "\\label{disc}   \\subsection{Jet propagation}\\label{disc1}  The simulated jets are strongly deviated from their original direction and show disrupted morphologies at the end of the calculations. Initially, the first interactions at small values of $z$ trigger helical patterns that, in the absence of further perturbations, could couple to Kelvin-Helmholtz unstable modes. However, the reconfinement shock triggered by the ram pressure of the clumpy wind on the jet plus individual interactions with multiple clouds downstream of this shock, generate non-linear structures that are most important for the jet evolution. The position of the reconfinement shock is located within the binary region, as predicted by PB08. There, the jet is decelerated and heated, thus becoming even weaker relative to the wind thrust, as also shown in PB08 and PBK10. In the present case, the clumps from the up-wind region can penetrate in the jet already at $z\\sim 1-2\\times 10^{12}$~cm, generating additional bow-like shocks and decelerating even more the jet flow. The deviation of the jet favors the presence of a region, on the down-wind side of the jet, to which, by the end of simulations, the clump bow shocks and jet mixing have still not reached. In this region, the flow keeps a relatively large axial velocity in this direction. However, both shocks and clump-mixed jet material should eventually fill the whole jet at larger $z$, so the description just presented (see Figs~\\ref{fig:maps8}, \\ref{fig:maps9} and \\ref{fig:maps10}) can be extended to the whole jet cross section. As a  result, farther downstream the jet becomes mass-loaded, slow and transonic. Although a similar result was already obtained in the homogeneous wind case (PB08, PBK10), wind clumping significantly potentiates jet destruction.   When comparing jet~B in this work with jet~2 in PBK10, it is remarkable that the presence of a forward shock and  the absence of clumps in the latter have a strong influence on the long-term evolution of the jet. The jet-driven forward shock in the wind and the corresponding cocoon (see PBK10) keep the wind at some distance from the jet. If a homogeneous wind were already in contact with jet~2 in PBK10, a smooth shear layer would form. Otherwise, the presence of clumps in jet~B enhances  the thrust locally in the wind, increasing mass, momentum and energy exchange. Jet~2 in PBK10 thus shows a larger  degree of collimation and is just slightly deviated from its original direction of propagation, contrary to what  we observe in the case of jet~B in this paper. From this work we can conclude that jet luminosities  $L_{\\rm j}\\gtrsim 10^{37}\\,\\dot{M}_{-6}\\mathrm{erg/s}$ ($\\dot{M}_{-6}=\\dot{M}/10^{-6}\\,M_\\odot$/yr) are needed if the jet is not to be destroyed when crossing the binary system. We notice that these dynamical arguments  favor $L_{\\rm j}\\gtrsim 10^{37}$ and $10^{38}$~erg/s in the HMMQ Cygnus~X-1 and Cygnus~X-3, respectively. The cocoon/clumpy wind case deserves few words. This situation takes place when the forward shock is well within the binary system, and the cocoon pressure is high, preventing clumps from entering the cocoon. When the forward shock has reached the outskirts of the binary, the cocoon pressure drops quickly (PBK10), allowing wind clumps to penetrate into the cocoon and reach the jet, and eventually dissipate the cocoon away.  \\subsection{Clump evolution}\\label{disc2}  In jets~A and B, those clumps reaching the jet relatively close to its base are destroyed just by jet expansion or erosion. These interactions trigger a shock wave that propagates inside the jet, and when interactions are frequent such a wave forces the strongly disruptive asymmetric recollimation shock. This phenomenon is illustrated in Fig.~\\ref{fig:maps12}), in which the first clump is completely destroyed by jet expansion, whereas the second one, even when it does not fully penetrate into the jet, triggers a shock strong that propagates all through the latter lasting for the whole simulation (several $t_{\\rm d}$; see Sect.~\\ref{phys}). The clump at the highest $z$ in jet~A' ($1.4\\times 10^{12}$~cm) is shocked but still not significantly disrupted by instabilities after few $t_{\\rm d}$. Later, this clump could be destroyed or may eventually escape the jet, although jet bending in the down-wind direction makes the latter unlikely. When clumps are disrupted inside the jet, all the clump mass is entrained by the flow. The level of mass loading can be easily estimated from the amount of clumps entering into the jet per time unit: $\\dot{N}_{\\rm cj}\\sim (\\eta/4\\pi)(3\\dot{M}/4\\pi R_{\\rm c}^3\\rho_{\\rm c})\\approx 0.02$~clump/s or $\\approx 5\\times 10^{17}$~g/s. This is $\\sim 3$ times more mass flux than in the jet. Therefore, jet deceleration due to mass loading can be very efficient, as is most clearly seen in the simulation results for jet~B. The implications also apply to faster and lighter, but equally powerful jets. Even if a strongly relativistic flow is injected at the jet base, for similar jet (relativistic) ram pressures clump penetration into, and mixing with, the jet will also occur. This will mass-load and brake the jet very quickly.  The qualitative dynamical scenario presented in ABR09 is validated by our numerical work in a semi-quantitative way. Remarkably, as hinted in ABR09, the clump interaction with many clumps will reduce the jet kinetic luminosity and also its ram pressure. In this way, clumps can actually keep their integrity longer, and penetrate farther inside the jet, even for moderate $f$-values.  \\subsection{Radiation considerations}\\label{disc3}  In ABR09, the radiation produced by a clump inside a HMMQ jet was calculated. Given the compactness of the considered region, i.e. the bow shock formed around the clump, radio emission was negligible. The luminosities of synchrotron X-ray, and (stellar photon) IC gamma-rays, dominated the non-thermal output, with their values anticorrelating depending on the magnetic field. Based on that work, and extending the study to the multi-clump/jet interaction case simulated here, we qualitatively discuss in what follows the expected emission.  The typical number of clumps inside the jet, at the binary scales, can be estimated as $N_{\\rm cj}\\sim$ few times $t_{\\rm d}\\,N_{\\rm cj}$, i.e. $\\gtrsim 10$ for the wind and jet properties adopted in this work (recall that the real clump lifetime is longer than $t_{\\rm d}$). Since $t_{\\rm d}\\sim 100\\,{\\rm s}<R_{\\rm j}/v_{\\rm w}\\sim 1000$~s, clump destruction will likely limit $N_{\\rm j}$ instead of jet-crossing escape. Although the jet can be quite disrupted and mass-loaded at the borders of the system, one can derive an estimate for the non-thermal emission, which depending on $B^2/8\\pi u_*$ ($u_*=L_*/4\\pi d^2 c$) will be released either through synchrotron (keV-MeV) or IC (GeV-TeV). We note that, for $B$-values well below equipartition with the non-thermal electron energy density, synchrotron self-Compton may also contribute to the IC component, although we neglect this component here for analysis simplicity. We also do not consider non-radiative losses or hadronic processes, but refer to ABR09.  Assuming a 10\\% efficiency for the kinetic-to-non-thermal energy transfer in clump bow shocks, and accounting for the clump-jet covering fraction $\\xi_{\\rm cj}\\sim (R_{\\rm c}/\\eta d)^2$, the high-energy emission can reach luminosities of $\\sim 0.1\\,N_{\\rm cj}\\,\\xi_{\\rm cj}\\,L_{\\rm j}\\sim 0.01\\,L_{\\rm j}$. This value may lead to a source that may be detectable in GeV ($B^2/4\\pi u_*<1$) and even in TeV ($B^2/4\\pi u_*\\ll 1$), showing persistent plus fluctuating variability components \\citep{owo09}. Strong TeV photon absorption may be avoided since most of the emission would come from the borders of the binary \\citep[e.g.,][]{bos09}. For $B^2/4\\pi u_*\\gtrsim 1$, synchrotron soft gamma-rays would dominate the non-thermal output. Thermal radiation, peaking at several keV, from the shocked clumps may be also present (ABR09). but it is minor for the adopted clump properties.  It is noteworthy that for a steep clump-size distribution (see Sect.~\\ref{phys}), i.e. with most of the clump-jet interaction occurring in a quasi-homogeneous wind regime, still several big clumps could enter the jet. This would likely lead to a lower level of persistent emission, but time to time chance would bring few big and dense clumps together inside the jet, rendering sudden and bright events with durations of a few $1000\\,(R_{\\rm c}/10^{11}\\,{\\rm cm})\\,(c_{\\rm c}/10^8\\,{\\rm cm})^{-1}$~s (ABR09). These statistically expected violent events could be related to a temporary complete destruction of the jet for large enough clumps.  Although the jets of microquasars present a complex phenomenology of their own \\citep{fen04,fen09}, on top of that  it is expected additional phenomena linked to the jet-wind interaction in HMMQ, as this and previous works (PB08, PBK10) show. Since the jet power is strongly linked to the mass-loss rate,  HMMQ jets detectable in X-rays or gamma-rays will be probably found in systems with moderate-to-strong winds. Therefore, unless particle acceleration is negligible in clump bow shocks, HMMQ  phenomenology at high-energies must be strongly affected by them, showing wind-related strong variability at high energies unless the wind is quite homogeneous or clumps with  $R_{\\rm c}>10^{10}$~cm are completely missing. Concerning (powerful) transient ejections, usually associated to X-ray  state transitions, we note that such an ejections will require some time to form. If this takes hours, the wind will have  time to surround the transient jet. Then, the clump impact will be as described here unless the jet is too powerful. If powerful blobs would appear as discrete even at the scales of the binary, they may have too much inertia to be significantly affected by the wind.  The dynamical impact of the stellar wind on the jet, enhanced by clumpiness, should not be neglected when interpreting radio emission from HMMQ. Even if jets escape from disruption, the enhanced jet entropy and bending, even by small angles ($\\sim 10^\\circ$), could have observational consequences. The reason is that bending pushes jet material farther from the orbital plane axis, increasing the strength of the Coriolis force exerted by the stellar wind and the orbital motion. The farther away from this axis the jet reaches, the stronger this force gets, enhancing jet heating, turbulence and bending. Although quantitative predictions call for a detailed study, this region may be observationally probed using VLBI techniques. Far enough from the binary, once the jet has become too wide to be affected by the orbital motion, a collimated supersonic flow may form again. If the energy and momentum fluxes kept enough anisotropy after crossing the system and suffer from orbital motion, given a negligible external pressure a jet-like structure could form again and propagate unstopped up to pc scales, terminating in the ISM \\citep[e.g.][]{bor09,bos11,yoo11}. Despite the different outflow geometry, similar phenomena are also expected in high-mass binaries hosting non-accreting pulsars \\citep[see][]{bos11b}."
    },
    "1112/1112.2158_arXiv.txt": {
        "abstract": "Pulsar timing now has a rich history in placing limits on the stochastic background of gravitational waves, and we plan soon to reach the sensitivity where we can detect, not just place limits on, the stochastic background.  However, the capability of pulsar timing goes beyond the detection of a background.  Herein I review efforts that include single source detection, localization, waveform recovery, a clever use of a ``time-machine\" effect, alternate theories of gravity, and finally studies of the noise in our ``detector\" that will allow us to tune and optimize the experiment.  Pulsar timing arrays are no longer ``blunt\" instruments for gravitational-wave detection limited to only detecting an amplitude of the background.  Rather they are shrewd and tunable detectors, capable of a rich and dynamic variety of astrophysical measurements. ",
        "introduction": "Pulsars are basically celestial clocks, and as such, can be used to construct a Galactic-scale gravitational wave detector using the same concept as ground-based interferometric detectors, i.e. one looks for phase changes in the arrival of the signal at the vertex station, in this case, earth.  The length scales of our detector `arms' (1000 light years) as compared to the length of ground-based arms (4km) allow us to probe a different gravitational-wave frequency regime (nHz), a complement to the ground-based kHz regime \\citep{Yardley10}.  For almost 30 years pulsar timers have been putting limits on the energy density of the stochastic background using pulsar timing \\citep{Romani83, Stinebring90, Kaspi94, Lommenthesis, Jenet06, vanHaasteren11, Yardley11}. They point out that at some moment in the future, we will detect rather than limit the stochastic background.  This moment is predicted to be sometime within this decade \\citep{Demorest09, Verbiest09}.  In the last 10 years the field of gravitational-wave detection using pulsars has matured, and we are now considering much more than just the background of gravitational waves.  We are demonstrating that very precise work on specific sources can be done, and that we need to `tune' this detector in order to maximize our sensitivity to these sources.  This manuscript briefly reviews these efforts, and is organized as follows. In \\S \\ref{sec:overview} I give some more details about the concept and current thought behind using pulsar timing arrays (PTAs) to detect gravitational waves.  In the subsequent sections I review the work that shows that PTAs can be (\\S \\ref{sec:directional}) directional detectors,  (\\S \\ref{sec:waveform}) used to recover the gravitational waveform, (\\S \\ref{sec:time}) used to recover information about the source at some time in past, (\\S \\ref{sec:distance}) used to measure luminosity distance to gravitational-wave sources, (\\S \\ref{sec:altgravity}) used to test alternate theories of gravity, and (\\S \\ref{sec:noise}) characterized as a formal `detector' using measurements of their noise. Finally, in \\S \\ref{sec:summary} I summarize the ways in which a PTA is no longer a `blunt' instrument for gravitational-wave detection, but rather a tunable, pointable, and adjustable detector that can be used to gain very specific astrophysical information about the gravitational-wave source being detected. ",
        "conclusions": ""
    },
    "1112/1112.2534_arXiv.txt": {
        "abstract": "We report the observation in the GeV band of the blazar 1ES~0229+200, which over recent years has become one of the primary sources used to put constraints on the Extragalactic Background Light (EBL) and Extragalactic Magnetic Field (EGMF). We derive constraints on both the EBL and EGMF from the combined Fermi-HESS data set taking into account the direct and cascade components of the source spectrum. We show that the limit on the EBL depends on the EGMF strength and vice versa. In particular, an EBL density twice as high as that derived by Franceschini et al. (2008) is allowed if the EGMF is strong enough. On the other hand, an EGMF strength as low as $6 \\times 10^{-18}$~G is allowed if the EBL density is at the level of the lower bound from the direct source counts. We present the combined EBL and EGMF limits as an exclusion plot in two dimensional parameter space: EGMF strength vs. EBL density. ",
        "introduction": "Very-high-energy (VHE) \\gr\\ flux from distant blazars is absorbed on the way from the source to the Earth through its interaction with the Extragalactic Background Light (EBL) photons (\\cite{gould}). The measurement of the induced distortions of the VHE \\gr\\ flux from distant hard-spectrum blazars by the effect of absorption on the EBL was used to derive constraints on the EBL density (\\cite{AhaEBL,Aha1ES0229,Orr_EBL_1ES0229}).  The conventional derivation of the upper bound on the EBL from \\gr\\ observations adopts the assumption that the intrinsic powerlaw-type spectrum of the primary source (a distant blazar) is characterized by the slope $dN_\\gamma/dE\\sim E^{-\\Gamma}$ with $\\Gamma \\ge 1.5$. This assumption appears reasonable in the framework of the most simple synchrotron-self-Compton (SSC) models for the broad band spectra of blazars. However, particular blazars considered for the derivations of EBL limits (the blazars with hardest intrinsic \\gr\\ spectra) may well not fit into this simplest SSC model framework, so that the assumption of $\\Gamma\\ge 1.5$ might not be applicable (\\cite{aharonian08,bottcher08,katarzinski06,Lefa_hard1,Lefa_hard2,neronov11}). If the intrinsic spectra of the blazars used for the derivation of the upper bound on the EBL density are harder than $\\Gamma=1.5$, this upper limit is relaxed (\\cite{mazin}).  Constraints on the intrinsic slope of the spectra of blazars  can be obtained from the observations by the Fermi Large Area Telescope (LAT) (\\cite{atwood09}) in the energy band below $\\sim 100$~GeV, where the effect of absorption on the EBL becomes negligible.  However, the blazars used for the derivation of constraints on the EBL are characterized by hard spectra, which makes it difficult to observe their flux below 100~GeV. In fact, the blazar 1ES~0229+200, which provides the tightest constraints on the EBL (\\cite{Aha1ES0229}) is not  listed in the catalogue of sources detected by LAT in two-year exposure (\\cite{fermi_catalog}), with only upper limits on the source flux derived from the LAT data (\\cite{NeronovEGMF,Tavecchio:2010mk,TaylorEGMF}) and a weak detection reported by \\cite{Orr_EBL_1ES0229}.  An additional difficulty for such constraints is that the spectrum of hard blazars might be composed of two contributions. Apart from the direct \\gr\\ emission from the primary source, an additional contribution is expected from the \\gr\\ cascade initiated in the intergalactic medium (IGM) by the absorbed VHE $\\gamma$ rays (\\cite{aharonian94,plaga95}). The overall flux and the spectral shape of the cascade contribution are determined by the strength of the extragalactic magnetic field (EGMF) (\\cite{Aharonian_cascade,plaga95,neronov07,NerSem_prediction}). Uncertainty of the EGMF strength introduces an uncertainty of the importance of the cascade contribution and prevents the measurement of the slope of the intrinsic \\gr\\ spectrum of the source. In fact, the limits on the EBL derived up to now are based on an underlying assumption about the EGMF strength (the EGMF should be strong enough to suppress the cascade contribution up to the $\\sim$TeV energy band), which is not justified a-priori.  If the assumption  on the EGMF strength is relaxed, the \\gr\\ data can be used to measure the EGMF strength. The electron-positron pairs, created as a result of the absorption of multi-TeV photons, up-scatter the Cosmic microwave background (CMB) as they cool, creating secondary emission in the GeV domain (\\cite{NerSem_prediction}). With the mean-free-path of TeV $\\gamma$ rays being $D_\\gamma \\simeq 100~(E_\\gamma/10~ {\\rm TeV})^{-1}$~Mpc, the production of these multi-TeV energy $e^+$-$e^-$ pairs occurs predominantly in extragalactic space. Furthermore, if non-negligible magnetic fields are present in the region these pairs are born into, they deflect significantly from their initial directions, resulting in the secondary GeV photons having a reduced probability of reaching the observer. Thus, the apparent flux is suppressed at low energies. Observational evidence for the presence of such a suppression places constraints on the EGMF strength to be $\\gtrsim 10^{-17}$~G (\\cite{Dermer,Dolag:2010ni,NeronovEGMF,Tavecchio:2010mk,TaylorEGMF}).  This limit on the EGMF is derived assuming a certain level of EBL density (a low EBL density derived by \\cite{Franceschini_EBL} is used in most of the publications on the EGMF in order to derive conservative bounds). However, a variety of EBL models do exist (\\cite{gilmore09, kneiske04, stecker_ebl}), which spans the full range of the present uncertainty in the EBL density. This uncertainty therefore introduces an uncertainty into the bounds on the EGMF derived from the \\gr\\ data. Higher(lower) EBL density leads to stronger(weaker) absorption of the primary \\gr\\ flux along with a stronger(weaker) cascade contribution to the observed flux in the GeV range. The suppression of such stronger(weaker) cascade contribution would require a stronger(weaker) EGMF.  Thus, the limits on the EBL derived from the \\gr\\ data depend on the assumptions made about the EGMF strength and vice versa. This implies that the correct procedure for the derivation of  limits on the EBL from the \\gr\\ data should include marginalization over the possible EGMF values. Conversely, the correct procedure for the derivation of the limits on the EGMF should include marginalization over the possible EBL densities and spectra.  Practical implementation of this correct procedure for the derivation of the EBL and EGMF bounds, however, present a challenging task since the quality of the \\gr\\ data is usually insufficient for the exploration of the entire EBL-EGMF parameter space. Exploration of this parameter space requires the measurement of the source spectra both in the TeV and GeV energy ranges.  In this Letter we report the observations of the blazar 1ES~0229+200 in the 1-300~GeV energy range using LAT's data with three years of exposure. The detection of the source below 100~GeV provides the information necessary for the correct analysis of the EGMF-dependent upper bound on the EBL density and of the EBL-dependent lower bound on the EGMF strength.  We present such bounds in the form of a two-dimensional exclusion plot in the ``EBL density'' vs. ``EGMF strength'' parameter space. ",
        "conclusions": "Following our investigations into the newly detected blazar 1ES~0229+200 with Fermi/LAT we find that, if the \\cite{Franceschini_EBL} EBL level is adopted, an EGMF strength of at least $10^{-17}$~G is required in order for the inevitable GeV cascade spectral component to be consistent with observations. However, we find that, more generally, the EGMF lower bound is dependent on the EBL level adopted, as demonstrated in Fig.~\\ref{EBL_exclusion}.  One should keep in mind that this bound is derived for large ($\\lambda_B > 1$~Mpc) correlation lengths of the EGMF, and scales approximately as $\\lambda_B^{-1/2}$ for $\\lambda_B \\lesssim 1$~Mpc. Our result lies in agreement with previous findings (\\cite{TaylorEGMF,Huan:2011kp}). Under the assumption that the source is stable on a time-scale much longer than 3 years, a larger bound on the EGMF is obtained (\\cite{NeronovEGMF,TaylorEGMF,Dolag:2010ni,Tavecchio:2010mk}).  We also find that an intrinsic blazar spectral index of $\\Gamma=1.5$ is able to sit in agreement with observations of 1ES~0229+200 by Fermi/LAT ($\\Gamma=1.36\\pm 0.25$). Thus our results can find agreement with the underlying assumptions behind previous EBL constraint calculations by \\cite{AhaEBL}. However, the remaining uncertainty in the origin of the blazar's spectral shape in the GeV domain leaves open the possibility for a harder intrinsic spectrum ($\\Gamma<1.5$).  We note, that in previous works of \\cite{AhaEBL} and \\cite{Orr_EBL_1ES0229} the authors made the implicit assumption for the absence of the cascade contribution in the GeV-TeV domain, which is equivalent to the presence of the strong EGMF with $B \\gtrsim 10^{-15}$~G. As can be seen from Fig. \\ref{EBL_exclusion}, for such a strong field our result is also compatible with the level of EBL found by \\cite{Orr_EBL_1ES0229}. For weaker fields, $B \\sim 10^{-17}$~G, we are also compatible with the findings of \\cite{dominguez11}. However, for such weak fields a proper account should be given of the GeV-TeV cascade component contribution.  Future observations with ground-base Cherenkov telescopes, such as MAGIC Stereo, HESS 2 and CTA, will dramatically improve the measurements of blazar spectra below $\\sim 100$~GeV, allowing much better constraints on the EGMF and EBL. In the particular case of 1ES~0229+200, if the intrinsic spectrum is hard, the presented EGMF lower bound may be transformed to a \\textit{measurement}, helping to better understand the nature of the extragalactic magnetic field."
    },
    "1112/1112.2228_arXiv.txt": {
        "abstract": "We report on the \\xray\\ and multiwavelength properties of 11 radio-quiet quasars with weak or no emission lines identified by the Sloan Digital Sky Survey (SDSS) with redshift $z=0.4$--$2.5$. Our sample was selected from the Plotkin et~al. catalog of radio-quiet, weak-featured AGNs. The distribution of relative \\xray\\ brightness for our \\hbox{low-redshift} weak-line quasar (WLQ) candidates is significantly different from that of typical radio-quiet quasars, having an excess of \\xray\\ weak sources, but it is consistent with that of high-redshift WLQs. Over half of the \\hbox{low-redshift} WLQ candidates are \\xray\\ weak by a factor of $\\gtrsim 5$, compared to a typical SDSS quasar with similar UV/optical luminosity. These \\xray\\ weak sources generally show similar UV emission-line properties to those of the \\xray\\ weak quasar \\phl\\ (weak and blueshifted high-ionization lines, weak semi-forbidden lines, and strong UV Fe emission); they may belong to the notable class of \\phl\\ analogs. The average \\xray\\ spectrum of these sources is somewhat harder than that of typical radio-quiet quasars. Several other \\hbox{low-redshift} WLQ candidates have normal ratios of \\xray-to-optical/UV flux, and their average \\xray\\ spectral properties are also similar to those of typical radio-quiet quasars. The \\xray\\ weak and \\xray\\ normal WLQ candidates may belong to the same subset of quasars having high-ionization ``shielding gas'' covering most of the wind-dominated broad emission-line region, but be viewed at different inclinations. The mid-infrared-to-X-ray spectral energy distributions (SEDs) of these sources are generally consistent with those of typical SDSS quasars, showing that they are not likely to be BL~Lac objects with relativistically boosted continua and diluted emission lines. The mid-infrared-to-UV SEDs of most radio-quiet weak-featured AGNs without sensitive X-ray coverage (34 objects) are also consistent with those of typical SDSS quasars. However, one source in our \\xray\\ observed sample is remarkably strong in \\hbox{X-rays}, indicating that a small fraction of \\hbox{low-redshift} WLQ candidates may actually be BL~Lacs residing in the radio-faint tail of the BL~Lac population. We also investigate universal selection criteria for WLQs over a wide range of redshift, finding that it is not possible to select WLQ candidates in a fully consistent way using different prominent emission lines (e.g., Ly$\\alpha$, \\civ, \\mgii, and H$\\beta$) as a function of redshift. ",
        "introduction": "Strong and broad line emission is a common feature of quasar spectra in the optical and UV bands. However, since multi-color quasar selection at high redshift in the Sloan Digital Sky Survey (SDSS; York et~al. 2000) is mostly based upon the presence of the Ly$\\alpha$ forest and Lyman break (e.g., Richards et~al.\\ 2002), the SDSS can also effectively select high-redshift quasars with weak or no emission lines. About 90 such weak-line quasars (WLQs) at high redshift have been found with \\lyanv\\ rest-frame equivalent widths of \\hbox{REW~$<15$~\\AA} (e.g., Fan et~al.\\ 1999, 2006; Anderson et~al.\\ 2001; Collinge et~al.\\ 2005; Diamond-Stanic et~al.\\ 2009, hereafter DS09). Some of these objects show a hint of weak Ly$\\alpha$ emission but no other lines; others are completely bereft of detectable emission lines even in high-quality spectra. High-redshift SDSS quasars show an approximately log-normal distribution of \\lyanv\\ REW with a mean of $\\approx 62$~\\AA\\ (DS09). The WLQs constitute $\\simgt 3\\sigma$ negative deviations from the mean, and there is no corresponding population with $\\simgt 3\\sigma$ positive deviations. The majority of these high-redshift WLQs are radio quiet ($\\alpha_{\\rm ro}>-0.21$; \\aro\\ is the slope of a nominal power law between 5~GHz and 2500~\\AA\\ in the rest frame; see \\S\\ref{xray} for a full definition).  WLQs have mainly been studied at high redshifts due to the fact that the Ly$\\alpha$ forest enters into the SDSS spectroscopic coverage for quasars at $z>2.2$. However, there is no apparent reason to believe that these objects should not also exist at lower redshifts. Indeed, a few apparent analogs of WLQs at lower redshifts have been found serendipitously over the past $\\approx 15$ years; e.g., PG~1407+265 (McDowell et~al.\\ 1995; \\hbox{$z=0.94$}), 2QZJ2154--3056 (Londish et~al.\\ 2004; \\hbox{$z=0.49$}), and PHL~1811 (Leighly et~al.\\ 2007ab; \\hbox{$z=0.19$}). As a byproduct of a systematic survey for optically selected BL~Lacertae objects (hereafter BL Lacs) in SDSS Data Release~7 (DR7; Abazajian et~al. 2009), Plotkin et~al. (2010a) discovered about 60 additional radio-quiet WLQ candidates at $z<2.2$ for which all emission features have \\hbox{REW~$<5$~\\AA}. These objects are perhaps the first \\hbox{low-redshift} SDSS counterparts of the previously identified high-redshift SDSS WLQs. Following the nomenclature that has been established by previous work on WLQs (e.g., Shemmer et~al. 2009), we define ``high redshift'' as $z>2.2$ and ``low-redshift'' as $z\\leqslant 2.2$, because WLQs are selected with different approaches for these redshift ranges (see above). Although WLQs are rare, their exceptional characteristics constitute a challenge to our overall understanding of quasar geometry and physics, especially the quasar broad emission-line region (BELR). Analogously, physical insights have been gained by investigating other minority populations with exceptional emission-line or absorption-line properties, such as Narrow-Line Seyfert~1 (NLS1) galaxies and Broad Absorption Line (BAL) quasars. Therefore, extensive studies of the multi-band properties of WLQs should have scientific value.  There are several candidate explanations for the physical nature of WLQs. Their UV emission lines may be weak due to an ``anemic'' BELR with a significant deficit of line-emitting gas (e.g., Shemmer et al. 2010). It has also been speculated that WLQs may represent an early stage of quasar evolution in which an accretion disk has formed and emits a typical continuum, but BELR formation is still in progress (e.g., Hryniewicz et~al. 2010; Liu \\& Zhang 2011).  The weak UV emission lines may also be a consequence of a spectral energy distribution (SED) which lacks high-energy ionizing photons. This soft SED may be a result of unusual accretion rate. For example, an extremely high accretion rate might produce a UV-peaked SED (e.g., Leighly et~al. 2007). In this scenario, high-ionization lines, like \\civ, should be suppressed relative to low-ionization lines like H$\\beta$. However, Shemmer et~al. (2010) estimated the normalized accretion rates, $L/L_{\\rm Edd}$, of two high-redshift WLQs via near-infrared spectroscopy and found their accretion rates were within the range for typical quasars with similar luminosities and redshifts. Alternatively, a combination of low accretion rate and large black hole mass may lead to a relatively cold accretion disk that emits few ionizing photons. Laor \\& Davis (2011) predicted a steeply falling SED at $\\lambda\\;<\\;1000$~\\AA\\ for quasars with cold accretion disks, and such an SED was observed in the WLQ SDSS~J0945+1009 by Hryniewicz et~al. (2010).  High-energy ionizing photons (including \\hbox{X-rays}) may be heavily absorbed before they reach the BELR. Wu et~al. (2011) studied a population of \\xray\\ weak quasars with unusual UV emission-line properties like those of \\phl\\ (weak and highly blueshifted high-ionization lines, weak semi-forbidden lines, and strong UV Fe emission). All of their radio-quiet \\phl\\ analogs were found to be \\xray\\ weak by a factor of $\\approx13$ on average. These objects also show a harder average \\xray\\ spectrum than those for typical quasars which suggests the presence of \\xray\\ absorption. \\phl\\ analogs appear observationally to be a significant subset ($\\approx 30\\%$) of WLQs. The existence of a class of quasars with high-ionization ``shielding gas'' covering most of the BELR, but little more than the BELR, could potentially unify the \\phl\\ analogs and WLQs via orientation effects (see \\S4.6 of Wu et~al. 2011). The shielding gas would absorb high-energy ionizing photons before they reach the BELR, resulting in weak high-ionization emission lines. When such a quasar is observed through the BELR and the shielding gas, a \\phl\\ analog would be seen; when it is observed along other directions, an \\xray\\ normal WLQ would be observed.  Another possibility is that instead of being intrinsically weak, the UV emission lines of WLQs could in principle be diluted by a relativistically boosted UV/optical continuum as for BL~Lac objects. However, this scenario is not likely for most WLQs. Shemmer et~al. (2009) found that the X-ray properties of high-redshift WLQs are inconsistent with those of BL~Lac objects. Furthermore, there is no evidence of strong optical variability or polarization for these WLQs (see DS09; Meusinger et~al. 2011). The UV-to-infrared SEDs of high-redshift WLQs are also similar to those of typical quasars, while the SEDs of BL~Lac objects are much different (DS09; Lane et~al. 2011). Nevertheless, it is possible that the population of BL~Lac objects has a small radio-quiet tail (e.g., Plotkin et~al. 2010b) and that a small fraction ($\\lesssim5\\%$; see Lane et~al. 2011) of the general WLQ population may be BL~Lac objects.  Most previous studies of WLQs were based on high-redshift objects. To investigate the nature of the overall WLQ population, we obtained new \\xray\\ observations of \\hbox{low-redshift} WLQs selected mainly from the catalog of radio-quiet BL~Lac candidates in Plotkin et~al. (2010a). We also utilized sensitive archival \\xray\\ coverage of the sources in their catalog. Our closely related science goals are the following: (1) enable comparison of the broad-band SEDs of \\hbox{low-redshift} WLQs to those of high-redshift WLQs, typical radio-quiet quasars, and BL~Lac objects; (2) provide basic constraints upon \\hbox{X-ray} spectral properties via band-ratio analysis and joint spectral fitting; (3) clarify if there is broad-band SED diversity among \\hbox{low-redshift} WLQs; and (4) allow reliable planning of future long, spectroscopic \\hbox{X-ray} observations.  In \\S\\ref{sample} we describe the selection of our sample of \\hbox{low-redshift}, radio-quiet WLQ candidates. In \\S\\ref{uvo} we detail their UV/optical observations and the measurement of their rest-frame UV spectral properties. In \\S\\ref{xray} we describe the relevant \\hbox{X-ray} data analyses. Overall results and associated discussion are presented in \\S\\ref{discuss}. Throughout this paper, we adopt a cosmology with $H_0=70.5$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm M}=0.274$, and $\\Omega_{\\Lambda}=0.726$ (e.g., Komatsu et~al. 2009). ",
        "conclusions": "\\label{discuss}  \\subsection{Relative X-ray Brightness}\\label{discuss:daox}  The \\daox\\ parameter (see Eqn.~\\ref{daoxeqn} for definition) is utilized to assess the \\xray\\ brightness of a quasar relative to typical radio-quiet quasars with similar UV luminosity. We compare the \\daox\\ distribution of our \\hbox{low-redshift}, radio-quiet WLQ candidates (see Fig.~\\ref{daox_fig}) to that of the 132 radio-quiet, non-BAL quasars in Sample B of Gibson et~al. (2008),\\footnote{We used an improved version of the Sample~B quasars in Gibson et~al. (2008) from which we further removed seven BAL quasars  (see Footnote~16 in Wu et~al. 2011).} which represent typical radio-quiet SDSS quasars. All of the 132 Sample~B quasars are \\xray\\ detected. The Peto-Prentice test (e.g., Latta 1981), implemented in the Astronomy Survival Analysis package (ASURV; e.g., Lavalley et~al. 1992), is used to assess whether our \\hbox{low-redshift} WLQ candidates follow the same \\daox\\ distribution as that for typical quasars (see results in Table~\\ref{twost_table}). We prefer the Peto-Prentice test to other possible similar tests because it is the least affected by the factors of different censoring patterns or unequal sizes of the two samples which exist in our case. We also compare the \\daox\\ distribution of high-redshift, radio-quiet WLQs in Shemmer et~al. (2006, 2009) to that of our \\hbox{low-redshift}, radio-quiet WLQ candidates and that of typical SDSS quasars (also see Fig.~\\ref{daox_fig} and Table~\\ref{twost_table}).  The \\daox\\ distribution of our \\hbox{low-redshift}, radio-quiet WLQ candidates is significantly different from that of typical SDSS quasars. The probability of null-hypothesis (two samples following the same distribution) is only $6.3\\times 10^{-7}$. This result is mainly due to the presence of a skew tail of \\xray\\ weak WLQs (see Fig.~\\ref{daox_fig}). Seven out of the 11 objects in our sample of \\hbox{low-redshift}, radio-quiet WLQ candidates have \\daox$\\;<-0.2$, giving a fraction of \\xray\\ weak objects of ($64^{+34}_{-24}$)\\% (68\\% confidence level). The mean \\daox\\ value for the \\hbox{low-redshift}, radio-quiet WLQ candidates is $-0.214\\pm0.078$, calculated with the Kaplan-Meier estimator also implemented in the ASURV package, while that for the Sample B quasars is $-0.001\\pm0.011$. The \\daox\\ distribution of the nine high-redshift, radio-quiet WLQs in Shemmer et~al. (2006, 2009) is also different from that of typical radio-quiet SDSS quasars, but less significantly (the probability of null-hypothesis is $7.1\\times 10^{-3}$).\\footnote{This result is somewhat inconsistent with the finding by Shemmer et~al. (2009) that their high-redshift WLQs have a similar \\daox\\ distribution to that of typical SDSS quasars. Shemmer et~al. (2009) used the Kolmogorov-Smirnov test and ignored the two high-redshift, radio-quiet WLQs with \\daox\\ upper limits. We include all nine high-redshift, radio-quiet WLQs since the Peto-Prentice test can properly treat censored data. The utilization of an improved Sample~B (see Footnote~17) does not substantially contribute to the inconsistency here. The Peto-Prentice test using all nine high-redshift, radio-quiet WLQs and the original Sample~B (as used by Shemmer et~al. 2009) provides a null-hypothesis probability of $1.53\\times 10^{-2}$.} Five of the nine objects in the \\hbox{high-redshift}, radio-quiet WLQ sample have \\daox$\\;<\\;-0.2$, giving a fraction of \\xray\\ weak objects of ($56^{+37}_{-24}$)\\% (68\\% confidence level); we note that the fraction could be somewhat higher (6/9) owing to the weak \\xray\\ upper limit for J1237+6301. The mean \\daox\\ value for the \\hbox{high-redshift}, radio-quiet WLQs is $-0.144\\pm0.075$. As expected the combined \\hbox{low-redshift} and high-redshift, radio-quiet WLQ sample (mean \\daox\\ value of $-0.187\\pm0.056$) also follows a different \\daox\\ distribution from that of typical SDSS quasars (null-hypothesis probability of $4.5\\times10^{-6}$). The \\daox\\ distribution of \\hbox{low-redshift}, radio-quiet WLQ candidates is consistent with that of high-redshift, radio-quiet WLQs (null-hypothesis probability of 0.35), though the sample sizes being compared are limited.  \\subsection{Classifying Radio-Quiet WLQs}\\label{discuss:class}  \\begin{figure*}[t] \\centering \\includegraphics[width=6.3in]{fig07color.ps}  \\vspace{0.1cm} \\caption{\\footnotesize{$\\alpha_{\\rm ro}$--$\\alpha_{\\rm ox}$ diagram for WLQ candidates (red filled circles for \\hbox{low-redshift} objects in our sample; blue filled squares for high-redshift objects in Shemmer et al. 2006, 2009), BL~Lac objects (black open circles; Shemmer et al. 2009), and typical radio-quiet SDSS quasars (small green dots; Sample~B quasars of Gibson et~al. 2008). The red asterisk represents PHL~1811. Rightward (downward) pointing arrows represent \\aox\\ (\\aro) upper limits. The two dashed lines mark the criteria for radio-quiet (\\aro $> -0.21$), radio-intermediate ($-0.39 <$\\aro$<-0.21$), and radio-loud (\\aro$<-0.39$) objects. The boxes bordered by dotted lines show the three suggested groups of WLQ candidates based on their multi-band properties. Note that the WLQ sample has an excess of objects with large negative \\aox\\ values, compared to both typical radio-quiet quasars and BL~Lac objects.} \\label{aroaox_fig}} \\end{figure*}  \\begin{figure*}[t] \\centering \\includegraphics[width=6.0in]{fig08color.ps} \\caption{\\footnotesize{REW(\\civ) plotted against the \\civ\\ blueshift for our radio-quiet WLQ candidates (filled squares for \\xray\\ detected sources, filled upside-down triangles for \\xray\\ undetected sources), \\phl\\ (asterisk), radio-quiet \\phl\\ analogs in Wu et~al. (2011) (stars), and \\hbox{radio-quiet}, \\hbox{non-BAL} quasars in Sample~B of Gibson et~al. (2008a) (circles). These sources are color-coded according to their \\daox\\ values (three color bins are used, corresponding to the \\xray\\ weak, \\xray\\ normal and \\xray\\ strong sources described in \\S\\ref{discuss:class}, respectively). The color bar shows the \\daox\\ range for each color. Source names for WLQs are labeled in the format of 'J$hhmm$' for brevity. The grey dots show the 13,582 \\hbox{radio-quiet} quasars in Sample A of Richards et~al. (2011; see their Fig.~7).} \\label{c4bew_fig}} \\end{figure*}  To investigate the multi-band properties of \\hbox{low-redshift}, radio-quiet WLQ candidates, we plotted the sources of our sample in an \\aro-\\aox\\ diagram (Fig.~\\ref{aroaox_fig}) along with the high-redshift, radio-quiet WLQs in Shemmer et~al. (2006, 2009), the BL~Lac sample in Shemmer et~al. (2009), and the Sample~B quasar in Gibson et~al. (2008). The \\hbox{low-redshift}, radio-quiet WLQ candidates have similar multi-band properties to those of high-redshift, radio-quiet WLQs. They are generally much fainter in radio and \\hbox{X-rays} than most of the BL~Lac objects. The weak emission lines of \\hbox{low-redshift}, radio-quiet WLQ candidates are therefore not likely due to the dilution by relativistically boosted continua as for BL~Lac objects (see discussion in \\S4.1 of Shemmer et~al. 2009). However, it is possible that a small percentage of the WLQ candidates actually belong to the radio-faint tail of the BL~Lac population (see below).  The population of radio-quiet WLQ candidates (both at low redshift and high redshift) has a wide dispersion of relative \\xray\\ brightness and UV emission-line properties. Motivated by their \\daox\\ distribution, their emission-line properties discussed below, and observations of related objects (e.g., Wu et~al. 2011), we will discuss them in three groups.   The majority of WLQ candidates are not only \\xray\\ weaker than BL~Lac objects, but also weaker than typical radio-quiet SDSS quasars. These WLQ candidates may belong to the notable class of \\xray\\ weak quasars termed ``\\phl\\ analogs'' which were recently studied in detail by Wu et~al. (2011). The \\phl\\ analogs generally have weak and highly blueshifted high-ionization lines (e.g., \\civ, Si~{\\sc iv}), weak semi-forbidden lines (e.g., C~{\\sc iii}]), and strong UV Fe~{\\sc ii} and/or Fe~{\\sc iii} emission. Some of our \\hbox{low-redshift}, radio-quiet WLQs have similar UV emission-line properties to those of \\phl, as listed below: \\begin{enumerate} \\item J0812+5225 (\\daox$=-0.42$) has weak C~{\\sc iii}] and strong Fe~{\\sc ii} emission. \\item J0945+1009 (\\daox$\\;<\\;-0.34$) has weak \\civ\\ and C~{\\sc iii}] emission lines. Its \\civ\\ line is highly blueshifted ($\\approx -7000$~km~s$^{-1}$). \\item J1252+2640 (\\daox$=-0.39$) has weak C~{\\sc iii}] and strong Fe~{\\sc ii} emission. \\item J1139$-$0201 (\\daox$=-0.38$) has weak and highly blueshifted ($\\approx -2950$~km~s$^{-1}$) \\civ\\ emission, weak C~{\\sc iii}] emission, and strong Fe~{\\sc iii} emission. \\end{enumerate} A high-redshift, radio-quiet WLQ J1302+0030 (\\daox$=-0.38$) also has a weak and highly blueshifted \\civ\\ emission line (DS09; Wu et~al. 2011). All of the above mentioned sources are \\xray\\ weak by a factor of $>\\;7$ (see Table~\\ref{aox_table}). Fig.~\\ref{c4bew_fig} shows the distribution of our WLQ candidates in the \\daox$-$\\ion{C}{4} blueshift$-$ REW(\\ion{C}{4}) parameter space. J0945+1009, J1139$-$0201 and J1302+0030 are similar to \\phl\\ analogs in this diagram. Based on the model in \\S4.6 of Wu et~al. (2011), these \\phl\\ analogs may have high-ionization shielding gas with large column density and a large covering factor of the BELR which blocks most of the ionizing photons, resulting in weak high-ionization emission lines. If a quasar of this kind is viewed through the BELR and shielding gas, it would be an \\xray\\ weak WLQ with weak and highly blueshifted high-ionization lines (e.g., \\civ). Based on the estimate in \\S4.6 of Wu et~al. (2011), \\phl\\ analogs should make up $\\approx 30\\%$ of the total WLQ population. However, our sample of low-redshift, radio-quiet WLQ candidates appears to have a higher fraction ($\\gtrsim 50\\%$) of \\phl\\ analogs, which may indicate our sample has some selection bias toward \\xray\\ weak WLQ candidates. This bias could perhaps be the result of a more strict criterion upon the strengths of emission lines for most sources in our sample (REW~$\\lesssim$~5~\\AA). Quasars with weaker emission lines (e.g., \\civ) are perhaps more likely to be weak in \\hbox{X-rays} (e.g., see \\S4.5 of Wu et~al. 2011). The apparently higher fraction of \\phl\\ analogs in our sample than that in Wu et~al. (2011) may perhaps also be caused simply by small-sample statistics. A Fisher's exact test (Fisher 1922) gives an $11.1\\%$ probability for the different fractions of \\phl\\ analogs among these two samples under the null hypothesis (i.e., the two samples have the same fraction of \\phl\\ analogs).  Some of our WLQ candidates have similar \\xray\\ brightness to that of typical radio-quiet quasars ($-0.2\\lesssim$~\\daox~$\\lesssim0.2$). Their high-ionization lines are also weak, but perhaps not highly blueshifted (e.g., see Fig.~\\ref{c4bew_fig} for J1408+0205). Some of them (e.g., J1612+5118) have very weak UV Fe~{\\sc ii} and/or Fe~{\\sc iii} emission. J1612+5118 does seem to have a highly blueshifted \\civ\\ line, for which the reason is unclear. However, the \\civ\\ line of this source is close to the blue border of its SDSS spectral coverage. Further UV spectroscopy with better \\civ\\ coverage is needed to confirm its \\civ\\ blueshift. Based on the model in Wu et~al. (2011), these sources are similar to \\phl\\ analogs physically, but they are viewed at different orientations. These sources are observed along lines of sight that avoid the shielding gas and the BELR. Therefore they appear normal in \\hbox{X-rays}. Their high-ionization lines are generally not highly blueshifted.  In our WLQ candidate sample, one source (J1109+3736) is remarkably strong in \\hbox{X-rays}. It also shows similar UV/optical spectral properties to those of BL~Lac objects. This source may belong to the radio-faint tail of the BL~Lac population; we will discuss it further in \\S\\ref{discuss:j1109}. It is worth noting that the division of our radio-quiet WLQ candidates into the three groups discussed above (as shown in Fig.~\\ref{aroaox_fig}) is somewhat arbitrary. We do have some ``border-line'' sources with \\daox$\\;\\approx\\;\\pm0.2$ (e.g., J1212+5341). It is difficult to classify these sources clearly based on current information.  \\subsection{The Infrared-to-X-ray SEDs of the Radio-Quiet WLQ Candidates}\\label{discuss:seds}  \\begin{figure*}[t] \\centering \\includegraphics[width=6.2in]{fig09.ps}  \\vspace{0.2cm} \\caption{\\footnotesize{Rest-frame spectral energy distributions (SEDs) of our \\hbox{low-redshift} WLQs with the best multiwavelength coverage, and of the radio-quiet BL~Lac candidate J1109+3736, ordered by \\daox. The photometric data points are from {\\it WISE} (open diamonds), 2MASS (open triangles), SDSS (filled circles), {\\it GALEX} (open circles), and X-ray observations (asterisks). The average SED of all SDSS quasars from the sample of Richards et al. (2006) is also shown (solid curve), scaled to the flux at rest-frame $10^{15}$~Hz. A parabolic SED for typical BL~Lac objects is shown by the dotted line in the top-left panel. The `+' signs in the bottom right-panel show the POSS photometry for J1109+3736.} \\label{sed1_fig}} \\end{figure*}  For the purpose of investigating further the multi-wavelength SEDs of our \\hbox{low-redshift}, radio-quiet WLQs, we gathered photometry for our sample from the following bands: (1) near- and mid-infrared from {\\it WISE} (The Wide-field Infrared Survey Explorer; Wright et~al. 2010); (2) near-infrared from 2MASS (The Two Micron All Sky Survey; Skrutskie et~al. 2006); (3) optical from the SDSS; (4) UV from {\\it GALEX} (The Galaxy Evolution Explorer; Martin et~al. 2005); and (5) \\hbox{X-rays} from this work. Fig.~\\ref{sed1_fig} shows the SEDs of the five \\hbox{low-redshift} WLQs in our sample which have the best multi-band coverage. We also include the SED of the radio-quiet BL~Lac candidate J1109+3736 (see more discussion in \\S\\ref{discuss:j1109}). A key point to keep in mind is that these multi-band observations are non-simultaneous. The SEDs are therefore subject to potential distortions due to variability.  We examined the {\\it WISE} and {\\it GALEX} image tiles by eye to identify potential cases of source blending, confusion, or incorrect matching caused by the low angular resolution of {\\it WISE} and {\\it GALEX}. None of the sources in our sample is subject to these kinds of problems. The 2MASS magnitudes in the SDSS DR7 quasar catalog were utilized; this catalog provided aperture photometry for additional sources detected down to $2\\sigma$ (see \\S5 of Schneider et~al. 2010). For the sources without detections at $\\geq2\\sigma$, we adopted flux upper limits obtained following the same photometry procedure (C.~M.~Krawczyk \\& G.~T.~Richards 2011, private communication). The first five sources in Fig.~\\ref{sed1_fig} are from the groups of \\xray\\ weak and \\xray\\ normal WLQ candidates discussed above. Their mid-infrared-to-UV SEDs are generally consistent with the composite SEDs of typical SDSS quasars in Richards et~al. (2006), and they are significantly different from the SEDs of BL~Lac objects (see the dotted parabolic line in the top-left panel of Fig.~\\ref{sed1_fig}; e.g., Nieppola et~al. 2006). We also investigated the SEDs for the {\\it WISE}-covered radio-quiet objects cataloged in Plotkin et~al. (2010a) which do not have sensitive \\xray\\ coverage (see the Appendix). The majority of them also have mid-infrared-to-UV SEDs consistent with those of typical radio-quiet quasars. Lane et~al. (2011) obtained similar results for their high-redshift WLQs; the composite SED of their high-redshift WLQs is inconsistent with SEDs of BL~Lac objects. For one source in our sample, J0945+1009, the flux in the UV band is lower than for typical SDSS quasars (see the {\\it GALEX} data points in the top-right panel of Fig.~\\ref{sed1_fig}). The UV deficiency of this source may be caused by Ly$\\alpha$-forest intervening absorption. However, Laor \\& Davis (2011) argued that such intervening absorption is not significant ($\\sim11\\%$ at most) for this source with $z=1.66$. The near-infrared-to-UV SED of J0945+1009 can be well fitted with their local black-body model for a cold accretion disk.  \\subsection{X-ray Spectral Properties of \\hbox{Low-Redshift} WLQ Candidates}\\label{discuss:stack}  Most of the \\hbox{low-redshift}, radio-quiet WLQ candidates do not have sufficient \\xray\\ counts for an individual \\xray\\ spectral analysis. We therefore investigate the average \\xray\\ spectral properties for these sources via stacking analyses and joint fitting.  A stacked spectral analysis was performed for the six \\hbox{low-redshift}, radio-quiet WLQ candidates with \\daox$\\;<\\;-0.3$ (J0812+5225, J0945+1009, J1139$-$0201, J1252+2640, J2115+0001, and J2324+1443). These sources are the weakest in \\hbox{X-rays} among the full sample of \\hbox{low-redshift} WLQ candidates. The detected sources have similar numbers of \\xray\\ counts, so that any one of them will not dominate the stacking analysis. These six sources span a relatively wide range of redshift ($z=1.15$--$2.50$), and thus the observed-frame bands of each source correspond to different energy ranges in the rest frame. However, under the assumption of a simple power-law spectral model, one can stack the \\xray\\ counts to obtain the average effective power-law photon index. We added the \\xray\\ counts of these sources in the observed-frame soft band and hard band, respectively. The numbers of total net counts are $14.4^{+4.9}_{-3.7}$ in the soft band and $5.2^{+3.4}_{-2.2}$ in the hard band ($68\\%$ confidence level), and the resulting band ratio is $0.36^{+0.27}_{-0.18}$. With the average Galactic neutral hydrogen column density of these sources ($N_{\\rm H}=3.50\\times10^{20}$~cm$^{-2}$), the band ratio was converted to an effective power-law photon index $\\Gamma=1.66^{+0.63}_{-0.51}$. The average \\xray\\ spectrum of the \\xray\\ weak \\hbox{low-redshift} WLQs is perhaps somewhat harder than that for typical radio-quiet quasars ($\\Gamma\\approx2$), but it is consistent within the error bars. This average \\xray\\ spectrum is likely softer than that of the \\phl\\ analogs at $z=2.19$--$2.38$ ($\\Gamma=1.10^{+0.45}_{-0.40}$) in Wu et~al. (2011) which was also obtained via a stacking analysis, but consistent within 2$\\sigma$. Both stacking analyses suffer from large uncertainty due to limited \\xray\\ counts. For a sample combining the six \\xray\\ weak WLQs analyzed here (which are likely to be \\phl\\ analogs) and the radio-quiet \\phl\\ analogs in Wu et~al. (2011), the average \\xray\\ spectrum has a flat effective power-law photon index of $\\Gamma=1.35^{+0.33}_{-0.31}$. Deeper \\xray\\ observations are necessary to give tighter constraints on the \\xray\\ spectral properties of these \\xray\\ weak WLQ candidates.  Two sources (J0945+1009 and J2115+0001) are undetected by \\chandra. Adding the \\xray\\ counts of these two sources cannot generate a stacked source that would be detected by \\chandra, because both of these two sources have zero \\xray\\ counts. However, we are able to obtain a tighter average constraint on their \\xray\\ brightness via stacking analysis. The upper limit upon the soft-band count rate of the stacked source is $2.51\\times10^{-4}$~s$^{-1}$. Average values of redshift, Galactic $N_{\\rm H}$, and $f_{2500\\mbox{\\rm~\\scriptsize\\AA}}$ are adopted in the following calculation. The upper limit upon the average flux density at rest-frame 2~keV is $5.16\\times10^{-33}$~erg~cm$^{-2}$~s$^{-1}$~Hz$^{-1}$ under the assumption of the Galactic-absorbed power law with $\\Gamma=2$. The upper limits upon \\aox\\ and \\daox\\ are calculated to be \\aox~$<\\;-2.21$ and \\daox~$<\\;-0.50$. Therefore, these two sources are \\xray\\ weak by a factor of $>\\;20$ on average.  For the two \\xray\\ normal, \\hbox{low-redshift} WLQs (J1604+4326 and J1612+5118), we performed joint fitting to study the average \\xray\\ spectral properties of these sources. The \\xray\\ spectra were extracted from apertures of $3''$ radius centered on the \\xray\\ positions of these sources via the standard CIAO routine {\\sc psextract}. The background spectra were extracted from annular regions with inner radii of $6''$ and outer radii of $9''$, which are free of \\xray\\ sources. Two spectra were extracted individually for J1604+4326 from its two \\chandra\\ observations. Spectral fitting was performed with XSPEC v12.6.0 (Arnaud 1996). The $C$-statistic (Cash 1979) was used in the spectral fitting instead of the standard $\\chi^2$ statistic because the $C$-statistic is well suited to the limited \\xray\\ counts in our analysis (e.g., Nousek \\& Shue 1989). We fit the spectra jointly using a power-law model with a Galactic absorption component represented by the \\verb+wabs+ model (Morrison \\& McCammon 1983). We also used another model similar to the first, but adding an intrinsic (redshifted) neutral absorption component, represented by the \\verb+zwabs+ model. Both sources were assigned their own values of redshift and Galactic neutral hydrogen column density; the Galactic column density was fixed to the values calculated with COLDEN (Column~4 of Table~\\ref{aox_table}). The joint fitting results are shown in Table~\\ref{xspec_table}.  The quoted errors or upper limits are at the 90\\% confidence level for one parameter of interest ($\\Delta\\;C=2.71$; Avni 1976; Cash 1979). The average \\xray\\ spectral properties of these two sources are similar to those of typical radio-quiet quasars. They have an average photon index ($\\Gamma=2.07^{+0.31}_{-0.30}$), consistent with that of typical radio-quiet quasars ($\\Gamma\\approx2$). The average photon index is also consistent with those from their band-ratio analyses. We did not find evidence of strong intrinsic neutral absorption for J1604+4326 and J1612+5118 ($N_{\\rm H}\\;\\lesssim\\;1.58\\times10^{22}$~cm$^{-2}$); the spectral fitting quality was not improved after adding the intrinsic neutral absorption component.  \\begin{figure}[t] \\centering \\includegraphics[width=3.0in]{fig10a.eps}  \\vspace{0.10cm}  \\includegraphics[width=3.0in]{fig10b.eps} \\caption{\\footnotesize{\\hbox{X-ray} spectra of SDSS~J1530+2310 (top panel; binned to a minimum of 10 counts per bin) and SDSS~J1109+3736 (bottom panel; binned to a minimum of 20 counts per bin) fitted with a power-law model with Galactic absorption. The residuals are shown in units of $\\sigma$. The inset of each panel shows contours of the photon index vs. intrinsic neutral hydrogen column density parameter space, at confidence levels of 68\\%, 90\\%, and 99\\%, respectively.} \\label{xspec_fig}} \\end{figure}  Two sources (J1109+3736 and J1530+2310) have sufficient \\xray\\ counts for individual spectral analysis. We will discuss the \\xray\\ spectral properties of J1530+2310 here and leave the radio-quiet BL~Lac candidate J1109+3736 for the next subsection. The \\xray\\ spectrum of J1530+2310 was extracted following the same procedure as for J1604+4326 and J1612+5118 (see above). The spectrum of J1530+2310 was first grouped to have at least 10 counts per bin (see Fig.~\\ref{xspec_fig}). We used the same spectral models as those for the joint fitting above. The standard $\\chi^2$ statistic was utilized. The fitting results are also shown in Table~\\ref{xspec_table}. Fig.~\\ref{xspec_fig} shows the \\xray\\ spectrum, the best-fit power-law model with Galactic absorption, and the contour plot of the $\\Gamma-N_{\\rm H}$ parameter space for the spectral model with intrinsic neutral absorption for J1530+2310. The \\xray\\ spectral properties of J1530+2310 are similar to those of J1604+4326 and J1612+5118. The photon index of J1530+2310 from the spectral fitting ($\\Gamma=2.11^{+0.37}_{-0.34}$) is consistent with that from the band-ratio analysis (see Table~\\ref{cts_table}). J1530+2310 also shows no evidence of strong intrinsic neutral absorption ($N_{\\rm H}\\;\\lesssim\\;2.67\\times10^{22}$~cm$^{-2}$). In summary, the \\xray\\ spectral properties of \\xray\\ normal \\hbox{low-redshift} WLQs are consistent with those of typical radio-quiet quasars. Shemmer et~al. (2009) found similar results for their high-redshift, radio-quiet WLQs. Their sources have a somewhat harder average \\xray\\ spectrum ($\\Gamma=1.86^{+0.72}_{-0.48}$), perhaps because they fit both \\xray\\ weak and \\xray\\ normal WLQs jointly. Shemmer et~al. (2009, 2010) also performed individual \\xray\\ spectral analysis on two high-redshift, radio-intermediate WLQs (J1141+0219 and J1231+0138).   \\subsection{J1109+3736: The Radio-quiet BL Lac Candidate}\\label{discuss:j1109}  J1109+3736 is a radio-quiet BL~Lac candidate based on its multi-band properties. It is strong in \\hbox{X-rays} by a factor of $\\approx6.3$ (\\aox$\\;=\\;-1.10$, \\daox$\\;=\\;0.31$). Its \\aox\\ value is similar to that of the majority of the BL~Lac population (see Fig.~\\ref{aroaox_fig}). \\xray\\ spectral analysis was carried out for J1109+3736 following similar procedures as those in \\S\\ref{discuss:stack}. The \\chandra\\ spectrum of J1109+3736 was grouped to have at least 20 counts per bin (see Fig.~\\ref{xspec_fig}). % The best-fit parameters are listed in Table~\\ref{xspec_table}. The \\xray\\ spectrum, the best-fit power-law model with Galactic absorption, and $\\Gamma-N_{\\rm H}$ contours for J1109+3736 are also shown in Fig.~\\ref{xspec_fig}. The best-fit photon index of J1109+3736 ($\\Gamma=1.77\\pm0.14$) is consistent with that from the band-ratio analysis. This photon index value indicates a harder \\xray\\ spectrum than those for the majority of the high-energy peaked BL~Lac objects (HBL),\\footnote{J1109+3736 would be classified as an HBL if it were indeed a BL~Lac object, based on the criterion $\\alpha_{\\rm rx}\\;<\\;0.75$ of Padovani \\& Giommi (1995), where $\\alpha_{\\rm rx}$ is the radio-to-X-ray spectral index, $\\alpha_{\\rm rx}=-0.130\\;\\log(f_{\\rm 1~keV}/f_{\\rm 5~GHz})$. The $\\alpha_{\\rm rx}$ value of J1109+3736 is estimated to be 0.46.} which have a mean photon index $\\Gamma\\approx2.2$, but it is still consistent with the broad distribution of HBL photon-index values (e.g., see the bottom panel of Fig.~1 in Donato et~al. 2005). Although this source has a high \\xray\\ count rate, we do not expect strong photon pile-up effects because a $1/2$ subarray mode was used for its \\chandra\\ observation. There is no evidence for significant intrinsic absorption in the J1109+3736 spectrum ($N_{\\rm H}\\;\\lesssim\\;1.4\\times10^{21}$~cm$^{-2}$).  The SDSS spectrum of J1109+3736 (see the bottom-right panel of Fig.~\\ref{spec_fig}) shows a strong power-law continuum without any detectable emission lines, in particular no Balmer lines. Note that some of the \\xray\\ weak, \\hbox{low-redshift} WLQ candidates (e.g., \\phl; see Leighly et~al. 2007b) have fairly strong Balmer lines. The strength of its Ca~{\\sc ii} H/K break\\footnote{The strength of the Ca~{\\sc ii} H/K break is defined as the fractional change of the average flux densities in continuum regions blueward (3750--3950~\\AA) and redward (4050--4250~\\AA) of the H/K break (see Equation 1 of Plotkin et~al. 2010a).} ($C=0.185$) indicates a relatively small contribution to the SDSS spectrum from the host galaxy. J1109+3736 shows moderate \\xray\\ variability between its \\chandra\\ observation and the epoch of the \\rosat\\ All Sky Survey in 1990 November (RASS; Voges et~al. 1999). It was not detected by RASS. The upper limit for \\aox\\ is estimated to be \\aox$\\;<\\;-1.29$, showing a factor of $>3$ variation compared to its \\chandra\\ observation. The quasar should have been detected in the RASS if it had the same \\xray\\ brightness as that in the \\chandra\\ epoch (expecting $\\approx20$ counts in a 250~s RASS observation). There is also evidence for optical variability of J1109+3736. This source was optically identified during the Cambridge APM (Automated Plate Measuring Machine) scans of POSS (Palomar Observatory Sky Survey) plates (e.g., McMahon et~al. 2002). We converted the APM magnitudes in the $O$ and $E$ bands ($O=20.76,\\;E=19.12$) into a $B$-band magnitude ($B=20.57$) using Equations \\hbox{(1--3)} in McMahon et~al. (2002). We also converted the SDSS magnitudes to a $B$-band magnitude ($B=18.88$) using the equations in Table~1 of Jester et~al. (2005). The $B$-band magnitude difference is 1.69, indicating a factor of $\\sim 5$ variability over a time span of $\\sim50$ years. This source shows greater variability than most quasars during this time span (e.g., see Fig.~24 of Sesar et~al. 2006). Therefore, J1109+3736 is a BL~Lac candidate on the radio-faint tail of the full BL~Lac population.  The broad-band SED of J1109+3736 is shown in Fig.~\\ref{sed1_fig}. It is relatively weak in the near infrared compared to its SDSS brightness. However, the SED profile of this source could be strongly affected by its variability (see the comparison between its SDSS and POSS fluxes in Fig.~\\ref{sed1_fig}) given its possible BL~Lac nature. Infrared photometry from {\\it WISE}, which will be available in the {\\it WISE} full data release,\\footnote{J1109+3736 is not covered by the currently available {\\it WISE} preliminary data release. See coverage map at http://wise2.ipac.caltech.edu/docs/release/prelim/.} will be helpful because the infrared-to-UV SEDs of BL~Lac objects are very different from those of typical quasars (e.g., Lane et~al. 2011). Future polarization measurements of J1109+3736 will also be useful in understanding its nature, since BL~Lac objects are usually highly polarized in the optical/UV band. However, polarimetry surveys of optically selected BL~Lac samples (e.g., Smith et~al. 2007; Heidt \\& Nilsson 2011) did not find any highly polarized radio-quiet BL~Lac candidates, indicating this kind of source, if it indeed exists, should make up only a small fraction of the total population of radio-quiet WLQ candidates."
    },
    "1112/1112.3704_arXiv.txt": {
        "abstract": "We present near-infrared spectroscopic observations of star-forming galaxies at $z\\sim1.4$ with FMOS on the Subaru Telescope. We observed K-band selected galaxies in the SXDS/UDS fields with $K\\leq23.9$ mag, $1.2\\leq z_{ph} \\leq 1.6$, $M_{*}\\geq 10^{9.5} M_{\\odot}$, and expected F(H$\\alpha$) $\\geq$ 10$^{-16}$ erg s$^{-1}$ cm$^{-2}$. 71 objects in the sample have significant detections of H$\\alpha$. For these objects, excluding  possible AGNs identified from the BPT diagram, gas-phase metallicities are obtained from [N\\emissiontype{II}]/H$\\alpha$ line ratio. The sample is split into three stellar mass bins, and the spectra are stacked in each stellar mass bin. The mass-metallicity relation obtained at $z\\sim1.4$ is located between those at $z\\sim0.8$ and $z\\sim2.2$. We constrain an intrinsic scatter to be $\\sim0.1$ dex or larger in the mass-metallicity relation at $z\\sim1.4$; the scatter may be larger at higher redshifts. We found trends that the deviation from the mass-metallicity relation depends on the SFR and the half light radius: Galaxies with higher SFR and larger half light radii show lower metallicities at a given stellar mass. One possible scenario for the trends is the infall of pristine gas accreted from IGM or through merger events. Our data points show larger scatter than the fundamental metallicity relation (FMR) at $z\\sim0.1$ and the average metallicities slightly deviate from the FMR. The compilation of the mass-metallicity relations at $z\\sim3$ to $z\\sim0.1$ shows that they evolve smoothly from $z\\sim3$ to $z\\sim0$ without changing the shape so much except for the massive part at $z\\sim0$. ",
        "introduction": "The gas-phase metallicity (hereafter, metallicity) is a fundamental parameter for the understanding of formation and evolution of galaxies, because it traces past star formation activity; metals in gas are produced in stars formed and returned into the inter-stellar medium (ISM) of galaxies. The metallicity is also affected by inflow and outflow of gas. Thus, by investigating the metallicity and its cosmological evolution, the star formation history of galaxies together with the gas inflow and outflow can be constrained.  The presence of a correlation between stellar mass and metallicity (hereafter, mass-metallicity relation) in the local universe is well known. \\citet{Tremonti:2004p4119} established the mass-metallicity relation at $z\\sim0.1$ with a large sample ($\\sim53,000$) of  Sloan Digital Sky Survey (SDSS) galaxies. Tracing the cosmological evolution of the mass-metallicity relation is indispensable to reveal how galaxies have been evolving. At $z\\sim1$, the mass-metallicity relations and the downsizing-like evolution to $z\\sim0.1$ were found; the less evolution can be seen in the massive part (e.g., \\cite{Savaglio:2005p3325,Zahid:2011p11939}). However, the anti-downsizing-like evolution of the mass-metallicity relations were also presented at similar redshifts \\citep{Lamareille:2009p5295, PerezMontero:2009p5308}. The mass-metallicity relations were obtained at $z\\sim2$ by using $\\sim9 0$ galaxies \\citep{Erb:2006p4143} and at $z\\sim3$ by $\\sim20$ galaxies \\citep{Maiolino:2008p5212, Mannucci:2009p8028} and the smooth evolution of the mass-metallicity relation from $z\\sim3$ to $z\\sim0$ is suggested. However, at $z\\sim2$, higher metallicities at a given stellar mass have been reported \\citep{Hayashi:2009p4235,Yoshikawa:2010p4286,Onodera:2010p4273} than those found by \\citet{Erb:2006p4143}. The discrepancy at $z\\sim2$ may be partly due to the small size of their samples ($10-20$). Since the redshift of $z=1-2$ is close to the peak epoch in the cosmic star-formation history, establishing the mass-metallicity relation at this redshift is very important.  Although the metallicity correlates well with stellar mass,  there is a scatter in the relation. Its origin may provide a clue to understand the process of the chemical enrichment and may be related to how the mass-metallicity relation evolves. At $z\\sim0.1$, \\citet{Tremonti:2004p4119} found that the mass-metallicity relation has an intrinsic scatter of $\\sim0.1$ and the scatter is larger at the lower stellar mass. At a given stellar mass, galaxies with higher surface stellar mass density tends to show higher metallicity, suggesting that they transformed more gas into stars raising the metallicity. \\citet{Ellison:2008p7997} showed that galaxies with larger specific star-formation rate (SFR) and size show lower metallicity at a given stellar mass at $z\\sim0.1$. More recently, \\citet{Mannucci:2010p8026} and \\citet{Yates:2011p16030} found that galaxies with larger SFRs tend to show lower metallicities at $z\\sim0.1$ and the scatter around the mass-metallicity relation is reduced very much (i.e., the fundamental metallicity relation) by introducing SFR as the second parameter affecting the metallicity. The scatter of the mass-metallicity relation, however, is still unclear at higher redshifts because of the limited sample size. In order to investigate these trends further, including the dependency of other parameters, at higher redshift, a large near-infrared spectroscopic survey is necessary, which is very time consuming and hard to achieve.  The Fiber Multi Object Spectrograph (FMOS) on the Subaru Telescope enables us to survey a large spectroscopic sample at high redshifts.  FMOS is a near infrared (NIR) fiber multi-spectrograph \\citep{Kimura:2010p11396}.   The fiber positioner ``Echidna'' at the prime focus feeds 400 fibers in a FoV of \\timeform{30'} diameter. Two NIR spectrographs (IRS1 and IRS2) with an OH airglow suppression system are capable of obtaining both low resolution ($R\\sim650$) and high resolution ($R\\sim3000$) spectra in the wavelength range of $0.9-1.8$$\\mu$m. Taking advantage of the multiplicity available with FMOS, we are conducting a large spectroscopic survey for star-forming galaxies at $z=1-2$. This redshift range is of much interest, because it is near or shortly after the peak epoch in the cosmic star formation history (e.g., \\cite{Hopkins:2006p4539}) and thus important for the history of the chemical evolution of galaxies. It is also suitable for FMOS, because in the redshift range [N\\emissiontype{II}]$\\lambda\\lambda$6548,6584, H$\\alpha$, [O\\emissiontype{III}]$\\lambda\\lambda$4959,5007, and H$\\beta$ lines are located in the wavelength range observable with FMOS. In this paper we present the first results from this survey as to metallicity measurements of galaxies at $z\\sim1.4$.  Throughout this paper, we adopt the concordance cosmology, ($\\Omega_{M}$ , $\\Omega_{\\Lambda}$ , $h$) = (0.3, 0.7, 0.7). All magnitudes are in the AB system \\citep{Oke:1983p15127}. ",
        "conclusions": "We present the first results obtained from near-infrared spectroscopic observations of star-forming galaxies at $z\\sim1.4$ with FMOS on the Subaru Telescope. We observed K-band selected galaxies in the SXDS/UDS fields with $K\\leq23.9$ mag, $1.2\\leq z_{ph} \\leq 1.6$, $M_{*}\\geq 10^{9.5} M_{\\odot}$, and expected F(H$\\alpha$)$\\geq10^{-16}$ erg s$^{-1}$ cm$^{-2}$. H$\\alpha$ emission lines of 71 objects are detected significantly. For these objects, excluding possible AGNs identified from the BPT diagram, gas-phase metallicities and upper limits are obtained from [N\\emissiontype{II}]/H$\\alpha$ line ratio by carefully considering the effects of the OH-masks. We separate the sample into three stellar mass bins and stack the spectra. We obtain a mass-metallicity relation at $z\\sim1.4$ from the stacking analysis. The mass-metallicity relation is located between those at $z\\sim0.8$ and $z\\sim2.2$. We tried to constrain an intrinsic scatter of the mass-metallicity relation and found that the scatter is $\\sim0.1$ dex and is comparable to that at $z\\sim0.1$ \\citep{Tremonti:2004p4119}. The scatter increases as stellar mass decreases if we take the intrinsic scatters at face-value. The scatters of the mass-metallicity relation at $z\\sim1.4$, however, should be lower limits; this implies that the scatter may be larger at higher redshifts. We found that the deviation from the mass-metallicity relation depends on the SFR and the half light radius: Galaxies with higher SFR and larger half light radii show lower metallicities at a given stellar mass. One possible scenario for the trends is the infall of pristine gas accreted from IGM or through merger events. Our results do not show clear FMR as proposed by \\citet{Mannucci:2010p8026}. Trends of the dependence of the SFR and the size on the mass-metallicity relation are, however, still not so clear. A larger sample from the further survey with FMOS in the future may be able to reveal clearer trends, not only the dependence of the SFR and the size but also that of other parameters. The compilation of the mass-metallicity relations at $z\\sim3$ to $z\\sim0.1$ shows that they evolve smoothly from $z\\sim3$ to $z\\sim0$ without changing its shape so much (except for the massive part at $z\\sim0.1$). The metallicity at $M_{*}=10^{10}M_{\\odot}$ on the mass-metallicity relation increases with decreasing redshift and can be described as $12+\\textrm{log(O/H)}=8.69 - 0.086 (1+z)^{1.3}$. The metallicity evolution rate was the highest at the cosmic age of $\\lesssim$ 3 Gyr which was before redshift of 2. However, the result may be influenced by sample selection and further studies with large samples at high redshifts are desirable.  \\bigskip We would like to thank an anonymous referee for useful comments. We are grateful to the FMOS support astronomer Kentaro Aoki for his support during the observations. We also appreciate Soh Ikarashi, Kotaro Kohno, Kenta Matsuoka, and Tohru Nagao sharing fibers in their FMOS observations. KY is financially supported by a Research Fellowship of the Japan Society for the Promotion of Science for Young Scientists. KO's activity is supported by the grant-in-aid for Scientific Research on Priority Areas (19047003).  We acknowledge support for the FMOS instrument development from the UK Science and Technology Facilities Council (STFC). DB and ECL acknowledge support from STFC studentships. We would like to express our acknowledgment to the indigenous Hawaiian people for their understanding of the significant role of the summit of Mauna Kea in astronomical research."
    },
    "1112/1112.1701_arXiv.txt": {
        "abstract": "We present three new eclipsing white-dwarf / M-dwarf binary systems discovered during a search for transiting planets around M-dwarfs. Unlike most known eclipsing systems of this type, the optical and infrared emission is dominated by the M-dwarf components, and the systems have optical colors and discovery light curves consistent with being Jupiter-radius transiting planets around early M-dwarfs. We detail the PTF/M-dwarf transiting planet survey, part of the Palomar Transient Factory (PTF). We present a Graphics Processing Unit (GPU)-based box-least-squares search for transits that runs approximately 8$\\times$ faster than similar algorithms implemented on general purpose systems. For the discovered systems, we decompose low-resolution spectra of the systems into white-dwarf and M-dwarf components, and use radial velocity measurements and cooling models to estimate masses and radii for the white dwarfs. The systems are compact, with periods between 0.35 and 0.45 days and semimajor axes of approximately $\\rm2 R_{\\odot}$ (0.01 AU). The M-dwarfs have masses of approximately 0.35$\\rm{M_{\\odot}}$, and the white dwarfs are all hydrogen-atmosphere with temperatures of around 8000K, and have masses of approximately 0.5$\\rm{M_{\\odot}}$. We use the Robo-AO laser guide star adaptive optics system to tentatively identify one of the objects as a triple system. We also use high-cadence photometry to put an upper limit on the white dwarf radius of 0.025$\\rm{R_{\\odot}}$ (95\\% confidence) in one of the systems. Accounting for our detection efficiency and geometric factors, we estimate that $\\rm 0.08\\%^{+0.10\\%}_{-0.05\\%}$ (90\\% confidence) of M-dwarfs are in these short-period, post-common-envelope white-dwarf / M-dwarf binaries where the optical light is dominated by the M-dwarf. Similar eclipsing binary systems can have arbitrarily small eclipse depths in red bands and generate plausible small-planet-transit light curves. As such, these systems are a source of false positives for M-dwarf transiting planet searches. We present several ways to rapidly distinguish these binaries from transiting planet systems. ",
        "introduction": "Large numbers of non-eclipsing white-dwarf / main-sequence binaries have been discovered in the Sloan Digital Sky Survey and other surveys (e.g. \\citealt{Rebassa2011}; \\citealt{Bianchi2007} and references therein). For low-mass stars in particular there is a bridge in color between white dwarfs and M-dwarfs. The bridge is interpreted as being due to rare white-dwarf / M-dwarf binaries, at a ratio with respect to single stars of $\\sim$1:2300 \\citep{Smolcic2004}.  White-dwarf / M-dwarf \\textit{eclipsing} systems are much rarer, and almost all have been discovered by searching for white dwarfs displaying very deep eclipses of up to several magnitudes (e.g. \\citealt{Drake2010}). These searches find systems containing relatively hot ($>12000$K) white dwarfs and mid-to-late M-dwarfs. The discovery rate of these systems (e.g. \\citealt{Drake2010, Parsons2011, Parsons2011b}) is increasing with the advent of large sky surveys. These binaries survived the common-envelope phase of their evolution and many will become cataclysmic variables (e.g. \\citealt{Nebot2011}), and so the properties and number statistics of these systems can provide windows into two important areas of stellar evolution. Precision measurements of the systems allow the determination of the masses and radii of two types of stars for which there are relatively few measurements \\citep{Nebot2009, Kraus2011, Pyrzas2011, Pyrzas2009}.  In this paper we present three eclipsing white-dwarf/M-dwarf systems discovered during the PTF/M-dwarfs search for transiting planets around M-dwarfs. In contrast to most known eclipsing systems of this type, the systems detected in this survey have optical and infrared emission dominated by the M-dwarf component and contain relatively low-temperature (8000K) white dwarfs and relatively early M-dwarfs. The shape of the light curves of the detected systems is similar to that expected for transiting giant planets around M-dwarfs, in particular in having a flat-bottomed eclipse with a depth of 1-20\\% in red optical bands.  The PTF/M-dwarfs survey \\citep{Law2011} is a search for transiting planets around 100,000 M-dwarfs. The survey is performed with the Palomar Transient Factory (PTF) camera \\citep{Rahmer2008, Law2009, Law2010} on the 48-inch Samuel Oschin telescope at Palomar Observatory, and is a Key Project of the Palomar Transient Factory \\citep{Law2009, Rau2009}. The PTF/M-dwarfs survey is designed to complement other M-dwarf transiting planet surveys such as MEarth (e.g. \\citealt{Charbonneau2009, Irwin2010}), the WFCam transit survey \\citep{Sipocz2011} and the \\mbox{M-dwarfs} in the Kepler mission target list \\citep{Borucki2011}, by covering a much larger number of M-dwarfs at somewhat lower sensitivity. The survey achieves photometric precisions of a few percent for $\\approx$100,000 targets, and few-millimag precision around a subset of $\\approx$10,000 M-dwarfs. These systems offer much larger transit depths compared to solar-type stars, while their very red colors compared to most other stars in the field greatly reduce the probability of a blended eclipse producing a difficult-to-detect transit false positive.   \\begin{figure*} \\centering \\resizebox{1.0\\textwidth}{!} { \\includegraphics{field.eps} } \\caption{A PTF Camera image of a survey field centered at 17:28, +57:22 and covering $3.50^{\\circ}\\times2.31^{\\circ}$. The highlighted points show the 2,851 stars with photometrically-estimated spectral types later than K4 and photometric stability better than 5\\%. The colors of the points correspond to the stellar temperatures, with K4 as yellow and late-M-dwarfs as red. The pipeline has removed stars with possible photometric precision problems such as proximity to a bright star or a bright ghost image. North is up and East is to the right.}  \\label{fig:field} \\end{figure*}  The three eclipsing systems presented here were originally detected as Jupiter-sized planet candidates during the first year of operations of the survey. Follow up of the candidates showed large color changes during eclipse and very large radial velocity signals, suggesting a hidden hot companion. In this paper we detail the properties of these eclipsing M-dwarf - white-dwarf systems and explore ways to distinguish them from true planetary transits.  The paper is organized as follows: in section \\ref{sec:survey} we describe the PTF/M-dwarfs survey, its precision photometry methods, and its target detection strategies, including a new method of performing a box-least-squares transit search in parallel on GPU hardware. In section \\ref{sec:disc_and_followup} we detail the three new eclipsing white-dwarf - M-dwarf systems and describe follow up photometric, low-resolution spectroscopic, and radial velocity observations, which are used to determine masses and radii for the system components in section \\ref{sec:models}. Section \\ref{sec:discussion} determines the frequency of eclipsing binaries such as these and discusses ways to distinguish them from transiting planets. ",
        "conclusions": "\\label{sec:discussion}  The discovery of only three white-dwarf / M-dwarf binaries in the PTF/M-dwarfs survey suggests a low incidence for such systems. This phase of the survey has covered $\\approx$45,000 M-dwarfs with sufficient precision to detect short-period systems with similar eclipse depths to those shown here (figure \\ref{fig:n_targets}). The survey's detection efficiency for systems with fractional-day periods is near 100\\% (figure \\ref{fig:detec_effic}). The geometric probability of eclipse is 5-15\\% for an early-M-dwarf primary in a roughly equal-mass half-day-period binary with a white dwarf. Taken together, the detection of three systems implies that $\\rm 0.08\\%^{+0.10\\%}_{-0.05\\%}$ (90\\% confidence) of M-dwarfs are in short-period post-common-envelope white-dwarf / M-dwarf binaries where the optical light is dominated by the M-dwarf.   \\begin{figure} \\centering \\resizebox{1.0\\columnwidth}{!} { \\includegraphics{n_targets.eps} } \\caption{The cumulative number of M-dwarf targets searched by the PTF/M-dwarfs survey in the first year of operations, as a function of achieved photometric stability. M-dwarfs in the PTF/M-dwarf fields are selected and confirmed on the basis of 2MASS, USNO-B1 and (where available) SDSS colors, along with proper motions determined from 2MASS and USNO-B1 positions. We require a high-confidence photometric color fit, along with a 2$\\sigma$ proper motion detection to rule out giant interlopers.} \\label{fig:n_targets} \\end{figure}   \\begin{figure} \\centering \\resizebox{1.0\\columnwidth}{!} { \\includegraphics{period_detect.eps} } \\caption{The eclipse detection efficiency for a typical PTF/M-dwarfs observing scenario, based on simulations including the actual observing cadences of the survey. The detection efficiency is the probability that at least two eclipses will be observed from an eclipsing target with a specified period, after observing window and weather effects have been taken into account. The low efficiencies around integer numbers of days at short periods are due to the likelihood of an eclipse always occurring in daytime; the same effect allows some long period  binaries to be picked up if their eclipses always occur at night.} \\label{fig:detec_effic} \\end{figure}  The periods of the detected systems are all longer than most known eclipsing M-dwarf / white-dwarf systems (figure \\ref{fig:period_dist}). Our survey is designed to detect longer-period transiting exoplanets, so this is not in itself surprising. However, our lack of detections at shorter periods, despite our near-100\\% detection efficiency for such systems, suggests that binaries including these relatively low-temperature white dwarfs are preferentially found at relatively large orbital radii.  \\begin{figure} \\centering \\resizebox{1.0\\columnwidth}{!} { \\includegraphics{period_hist.eps} } \\caption{The periods of our three detected systems compared to the distribution of periods of the white-dwarf dominated eclipsing systems found by \\citet{Drake2010}, which were all found in a single survey and so offer a useful comparison sample.} \\label{fig:period_dist} \\end{figure}  These eclipsing binaries appear as M-dwarfs when selected by optical and infrared colors, and they display transit-like light curves which can have arbitrarily small depths. They appear to be at least three times more common than transiting giant planets in the PTF/M-dwarfs survey. These properties make these systems an important false-positive for current and future transiting planet searches around M-dwarfs. Fortunately, they can be distinguished from true transiting planets with small amounts of extra data, using the methods detailed below. The methods are ordered by difficulty, starting with those which require only the discovery light curves.  \\begin{enumerate} \\item{Eclipse duration. The increased mass of the system compared to a M-dwarf/planet system leads to a shorter eclipse; high-SNR and high-cadence light curves are however required to distinguish this case from a lower-inclination system.} \\item{Eclipse shape. The eclipse of the white dwarf by the M-dwarf provides an exactly-flat-bottomed eclipse outside the short ingress and egress periods. With high-precision photometry this can be distinguished from a true transiting planet which will show the effects of limb darkening as it passes in front of the M-dwarf.} \\item{Out-of-eclipse variability. All our systems show out-of-eclipse variability at levels which are orders of magnitude greater than that expected for planetary systems (see \\citet{Drake2003} for details of a similar selection method for planetary transits around solar-type stars).} \\item{UV emission. Two of our new systems showed UV emission and the third has possible emission. Where data is available, the presence of UV (or even u-band) emission suggests the existence of something other than an M-dwarf in the system.} \\item{Low-resolution spectroscopy. Like the three targets presented in this paper, a low-resolution spectrum could be decomposed into white dwarf and M-dwarf components. Faint, low-temperature white dwarfs may be, however, hard to detect with this technique.} \\item{Multi-color photometry. These systems show a strong variation in eclipse depth with wavelength (depending on the relative temperatures of the white dwarf and M-dwarf and their ratio of radii).} \\item{Radial velocities. These systems have radial velocity amplitudes at least 500$\\times$ larger than expected for a planet with the same period. Just two RV few-km/sec-precision datapoints are sufficient to discriminate these systems from transiting planet systems.} \\end{enumerate}  Of these methods, mulitcolor photometry through eclipse is probably the most time-efficient method of detecting systems like this. As large scale transiting planet surveys of M-dwarfs start up and continue, many more systems in this interesting parameter space are likely to be found.   Follow-up precision photometry and radial velocities will allow direct measurements of the masses and radii of all components of these systems. As the white dwarf transits across the disk of the M-dwarf the transit depth is expected to be around 1 millimag, although lensing by the white dwarf (e.g., \\citealt{Marsh2001}) will make the transit shallower than expected from only geometric considerations (e.g., \\citealt{Steinfadt2010}). High-cadence and high-precision photometry of the white dwarf eclipse (occultation) ingress and egress may be the best approach to directly measure the white dwarf radii. Furthermore, small asymmetries in the ingress and egress light curve due to the photometric Rossiter-McLaughlin effect \\citep{Shporer2012, Groot2011} can allow a measurement of the white dwarf spin-orbit alignment and rotation velocity. With these methods, this new group of systems will fill a poorly-covered range of the white-dwarf and M-dwarf mass/radius relations."
    },
    "1112/1112.6121_arXiv.txt": {
        "abstract": "The application of high end computing to astrophysical problems, mainly in the galactic environment, is developing for many years at the Dep. of Physics of Sapienza Univ. of Roma. The main scientific topic is the physics of self gravitating systems, whose specific subtopics are: i) celestial mechanics and interplanetary probe transfers in the solar system; ii) dynamics of globular clusters and of globular cluster systems in their parent galaxies; iii) nuclear clusters formation and evolution; iv) massive black hole formation and evolution; v) young star cluster early evolution. In this poster we describe the software and hardware computational resources available in our group and how we are developing both software and hardware to reach the scientific aims above itemized. ",
        "introduction": "Celestial mechanics is one of the most classic examples of chaos in physics: the mutual gravitating systems show a chaotic behaviour, being extremely sensitive to differences in initial conditions. This problem can be only partially controlled using high-order integration algorithms. The intrinsic difficulty of the problem is summarized by the so called double divergence of the pair interaction potential $U_{ij} \\propto 1/r_{ij}$, where $r_{ij}$ is the distance between particle $i$ and particle $j$.  The ``ultraviolet'' divergence corresponds to gravitational encounters at vanishing distance, while the ``infrared'' divergence means that the force never vanishes. The computational problems arising from these divergences make the classic gravitational \\emph{N}-body problem unique. ",
        "conclusions": ""
    },
    "1112/1112.2064_arXiv.txt": {
        "abstract": "Current observations suggest that the universe was reionized sometime  before $z\\sim 6$. One way to observe  this epoch of the universe is through the Near Infrared Background (NIRB),  which contains information about galaxies which may be too faint to be observed individually. We calculate the angular power spectrum ($C_l$) of the NIRB fluctuations caused by the distribution of these galaxies.  Assuming a complete subtraction of any post-reionization component, $C_l$ will be dominated by galaxies responsible for completing reionization (e.g., $z\\sim 6$).  The shape of $C_l$ at high $l$ is sensitive to the  amount of {\\it{non-linear}} bias of dark matter halos hosting  galaxies.  As the non-linear bias  depends  on the mass of these halos,  we can use the shape of $C_l$ to infer typical masses of  dark matter halos responsible for completing reionization. We extend our previous study by using a higher-resolution $N$-body simulation, which can resolve halos down to  $10^8~M_\\sun$.  We also include improved radiative transfer, which allows for the  suppression of star formation in small-mass halos due to  photo-ionization heating.  As the non-linear bias enhances the dark-matter-halo power  spectrum on small scales, we find that $C_l$ is steeper for  the case with a complete suppression of small sources or  partial suppression of star formation in small halos (the minimum galaxy mass is $M_{\\rm min}=10^9~M_\\sun$ in ionized regions and  $M_{\\rm min}=10^8~M_\\sun$ in neutral regions) than the case in which these small halos were unsuppressed.  In all cases, we do not see a turn-over toward high $l$ in  the shape of $l^2 C_l$. ",
        "introduction": "\\label{sec:introduction}  Probing the beginnings of star and galaxy formation in the early universe is one of the goals of modern cosmology.  These high redshift stars provided a wealth of ionizing photons, which caused all of the hydrogen in the universe to be ionized, a process called reionization. Therefore, understanding high-redshift star formation is closely coupled with our understanding of reionization history.  Recent measurements of the polarization of the Cosmic Microwave Background have shown that reionization started early and was extended in time, with an equivalent instantaneous reionization redshift of $z \\sim 11$  \\citep{kogut/etal:2003,spergel/etal:2003,spergel/etal:2007, page/etal:2007,dunkley/etal:2008,komatsu/etal:2008}.  Therefore, we expect there to be significant star formation before this.  With improving observations, there are many ways to be able to observe these first few generations of stars and galaxies.  We can now observe high-redshift ($z>6$) star formation via galactic surveys, which provide statistical properties of high-redshift sources.  These surveys provide a wealth of information on early star formation; however, they only probe the ``tip of the iceberg,'' e.g., a small fraction of the whole population, mainly those galaxies that are bright enough to be identified in surveys. {\\footnote{Those exceptionally bright galaxies also might be missed in surveys, because they are rare enough that they will not be found within the survey's limited area \\citep{trenti/stiavelli:2008}}}.  Another way to observe these high-redshift galaxies is to look for the {\\it diffuse} background light originating from them. Ultra-violet photons produced by these early galaxies at $z\\gtrsim 6$ would be redshifted into the near infrared band ($\\lambda\\gtrsim 1~\\mu$m). Therefore, any background above and beyond that from the low-redshift galaxies could be attributable to these early galaxies. Observing the Near Infrared Background (NIRB) has the benefit of probing a population of galaxies {\\it{as a whole}}, not just the unusually bright objects which can be detected individually. For this purpose, one can observe the mean NIRB \\citep{santos/bromm/kamionkowski:2002, magliocchetti/salvaterra/ferrara:2003,salvaterra/ferrara:2003, cooray/yoshida:2004,madau/silk:2005, fernandez/komatsu:2006}  as well as fluctuations  \\citep{kashlinsky/odenwald:2000,kashlinsky/etal:2002, kashlinsky/etal:2004,kashlinsky/etal:2005,kash/etal:2007,kashlinsky/etal:2007, kashlinskyb/etal:2007c,kashlinsky:2005,magliocchetti/salvaterra/ferrara:2003, odenwald/etal:2003,cooray/etal:2004,matsumoto/etal:2005,thompson/etal:2007a, thompson/etal:2007b, fernandez/etal:2010}.  In our previous work \\citep{fernandez/etal:2010}, we calculated the angular power spectrum of the NIRB fluctuations using an $N$-body simulation to trace the large-scale structure and the formation of galactic halos, % coupled with radiative transfer simulations of reionization. Combining this numerical data with the latest analytic modeling of the internal, unresolved properties of radiation processes such as stellar emission, Lyman-$\\alpha$ line, two-photon emission, free-free and free-bound emission from the formed galaxies \\citep{fernandez/komatsu:2006}, we analyzed the effects of changing the star formation efficiency $f_*$, the escape fraction $f_{\\rm esc}$\\footnote{Here, by the ``escape fraction,'' we refer to a fraction of hydrogen-ionizing photons escaping from the galaxies into the intergalactic medium. We ignore dust extinction.}, and the mass and metallicity of the stars on the NIRB angular power spectrum.  Our previous simulation was able to resolve dark matter halos down to the minimum mass of $M_{\\rm min}=2.2\\times 10^9~M_\\sun$. In the current paper, we extend this analysis by improving the mass resolution by more than an order of magnitude to $M_{\\rm min}=10^8~M_\\sun$. We also include a new physical effect: the suppression of star formation in small-mass halos due to photo-ionization heating (the so-called ``Jeans-mass filtering'', which prevents baryons from collapsing into small dark matter halos \\citep[e.g.,][]{shapiro/etal:1994,2007MNRAS.376..534I}).  The structure of this paper is as follows. The simulations used for this analysis are described in \\S~\\ref{sec:simulation}. The fiducial stellar model used is discussed in \\S~\\ref{sec:stellarpop}, and our results are presented in \\S~\\ref{sec:results}. This model is then expanded upon in \\S~\\ref{sec:pops}.  We compare our results to observations in \\S~\\ref{sec:obs}, and conclude in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  The NIRB mean and fluctuation signals can be a very powerful probe of the high-redshift star formation ($z\\gtrsim 6$). Building upon our previous work on this subject \\citep{fernandez/komatsu:2006,fernandez/etal:2010}, we have extended the calculation of the angular power spectrum of the NIRB fluctuation by improving the mass resolution of $N$-body simulations by an order of magnitude. With the simulation now resolving the dark matter halo mass down to $M=10^8~M_\\sun$, we have confirmed our previous finding: due to the non-linear halo biasing, the shape of $l^2C_l$ of the NIRB fluctuation does not exhibit a turn over. However, we do observe a flattening of the shape as we lower the minimum halo mass (because the non-linear bias decreases as masses go down), and thus the shape of $l^2C_l$ can be used to infer typical masses of dark matter halos hosting galaxies in a high redshift universe.  We have gone beyond simply increasing the mass resolution of the simulation. For the first time, we consider cases with a radiative feedback suppressing the star formation in low-mass galaxies with $M<10^9~M_\\sun$, due to the Jeans-mass filtering in the ionized and heated IGM. We find that, when low-mass sources are partially suppressed by photo-ionization heating, the predicted angular power spectrum becomes similar to the one with a complete suppression of halos below $M_{\\rm min}=10^9~M_\\sun$, yielding a steep power spectrum at high $l$. Therefore, the shape of the angular power spectrum at high $l$ can directly provide information on the typical mass of sources responsible for completing reionization.  The amplitude of the angular power spectrum is less robust, as it depends on a number of parameters such as the stellar initial mass spectrum, metallicity of stars, star formation efficiency, and escape fraction. However, one robust feature is that it is largely determined by the properties of the halo populations at late stages of reionization.  Therefore, the angular power spectrum is higher if the stars produce more ionizing photons that do not escape from the halo.  (Therefore, maximizing $f_*$ and $N_i$ while minimizing $f_{\\rm esc}$ within the limitations of reionization.)  Finally, we find that our predictions, all of which are tuned to satisfy the reionization constraints, are below the current measurements. Given that these measurements should probably be taken as upper limits on the contributions from galaxies in $z\\gtrsim 6$, we conclude that our calculations are consistent with the current measurements."
    },
    "1112/1112.5191_arXiv.txt": {
        "abstract": "We present \\textsc{NBursts+phot}, a novel technique for the parametric inversion of spectrophotometric data for unresolved stellar populations where high-resolution spectra and broadband SEDs are fitted simultaneously helping to break the degeneracies between parameters of multi-component stellar population models. ",
        "introduction": "Optical spectra of galaxies contain important information about their internal kinematics and stellar content that can be extracted by fitting them against stellar population models. There is a number of approaches for the full spectrum fitting that were proved to be efficient in studying different classes of galaxies and star clusters.  On the other hand, the broad-band photometry in different spectral domains that became available for large samples of galaxies thanks to modern wide-field surveys allows one to study certain properties that cannot be derived from the spectra (e.g. internal extinction) and to brake degeneracies between parameters in case of complex star formation histories (SFH). Usually, photometric measurements and spectra are used independently. Here we propose a new approach \\textsc{NBursts+phot} that fits in a single minimization loop spectral and photometric information and recovers both, parametric SFH, and internal kinematics of galaxies. ",
        "conclusions": ""
    },
    "1112/1112.5158_arXiv.txt": {
        "abstract": "We consider the possibility that the cosmological dark matter consists of particles very close in mass to new colored particles below the TeV scale. While such a scenario is inherently difficult to directly confirm at colliders, we find that indirect dark matter searches may be a powerful alternative. In particular, we show that in this case dark matter annihilation to $\\bar q q g$ final states can give rise to significant antiproton (but also gamma-ray) fluxes, and compare the resulting constraints to bounds from direct searches at LEP, the Tevatron and the LHC. For supersymmetric neutralinos degenerate with squarks, \\fex, antiprotons can give rise to more stringent constraints for masses below around 45--80~GeV. ",
        "introduction": "Collider experiments have pushed the scale for possible new physics beyond the standard model (BSM) to ever higher energies in recent years. The CERN LHC, in particular, is now running at a center of mass energy of $7\\,$TeV and the non-observation of any clear BSM signal, so far, has allowed to place impressive limits~\\cite{Aad:2011qa,Chatrchyan:2011zy}; after an integrated luminosity of $1.04$ fb$^{-1}$, scalar quarks (squarks) in  supersymmetric extensions of the standard model, \\fex, have already been excluded for masses below $\\sim 900$ GeV~\\cite{Aad:2011ib,cmssimp}.  One has to keep in mind, however, that these results rest on the assumption that the new colored states quickly decay into the lightest neutralino, which is assumed to be massless for the sake of the analysis, and emit high-energy QCD jets in this process. These limits thus do not apply if the only accessible lighter states are very close in mass~\\cite{ATLAS:limit,Chatrchyan:2011ek,LeCompte:2011cn,LeCompte:2011fh}, in which case the resulting jets and missing transverse energy would be too soft to pass the signal selection criteria.  Scenarios with degenerate particle spectra are thus generically very difficult to probe at hadron colliders; the situation is considerably better for electron-positron colliders like LEP, but even in this case the limits from direct searches for colored states do not apply to highly degenerate spectra~\\cite{Heister:2002hp, Abbiendi:2002mp, Achard:2003ge, Abdallah:2003xe, LEP_SUSYWG}.  Here, we investigate whether such a situation could be probed by indirect dark matter (DM) searches, assuming that the lightest BSM state is a weakly interacting massive particle (WIMP) that makes up the cosmological DM.  The self-annihilation of WIMP pairs  could then leave an imprint in the spectrum of cosmic rays (see, \\fex, Ref.~\\cite{Bertone:2004pz} for a review on particle DM and indirect searches).  As we will demonstrate, WIMP annihilation into $\\bar qq g$ final states, a channel so far hardly explored in indirect DM searches, may give rise to the most stringent constraints in degenerate scenarios (mostly through antiproton, but also through gamma-ray production). In fact, this channel may help to fill remaining loop-holes for the existence of new colored states below masses of around 100 GeV that are left from direct collider searches for such particles.  The rest of this Letter is organized as follows: In Section 2 we discuss the importance and main characteristics of DM annihilation into $\\bar{q}qg$ final states, focussing on the case of neutralino DM; in Section 3 we present the energy spectra of antiprotons that result from $\\bar{q}qg$ hadronization and discuss antiproton propagation through our Galaxy; Section 4 contains the limits on the annihilation cross-section that can be derived from cosmic-ray observations and Section 5 confronts these limits with direct collider searches for charged colored particles. Finally, in Section 6 and 7, respectively, we present a discussion of our results and conclude. ",
        "conclusions": "In this Letter, we have described a DM annihilation channel, and discussed its phenomenological implications in some detail, that so far has hardly been explored in the context of indirect dark matter searches.  If DM is a Majorana fermion $\\chi$, or a scalar, direct annihilation via $\\chi\\chi\\to\\bar{q}q$ is strongly helicity suppressed. In this case, the overall annihilation rate \\emph{today} can be dominated by internal gluon Bremsstrahlung, $\\chi\\chi\\to\\bar{q}qg$, and can be boosted by several orders of magnitude with respect to the tree-level result;  while this enhancement is much smaller in the early universe, during \\emph{freeze-out}, it can also have a non-negligible effect on the relic density~\\cite{Flores:1989ru,hep-ph/9306325,hep-ph/0608215}. The precise branching ratios depend on the difference between the DM particle's mass $m_\\chi$ and the squark mass $m_{\\tilde{q}}$; the process $\\chi\\chi\\to\\bar{q}qg$ becomes strongest in the mass degenerate case, i.e.~for $m_{\\tilde{q}}/m_\\chi\\simeq1$.  We studied the spectra of cosmic rays generated by this  channel, and found that the energy spectrum of antiprotons coming from $\\bar{q}qg$ is generically enhanced  with respect to the spectrum from $\\bar{q}q$ final states (Fig.~\\ref{fig:dNdx}). The corresponding limits on the annihilation cross-section for $\\chi\\chi\\to\\bar{q}qg$ are hence somewhat stronger than limits on the traditional $\\bar{b}b$ or $W^+W^-$ final states (Fig.~\\ref{fig:svlim}). From this figure we can see, as noted before~\\cite{Evoli:2011id,Bottino:2005xy,Bringmann:2009ca,Lavalle:2010yw}, that antiprotons are most constraining for light DM masses, which in particular seems to disfavor the simplest DM interpretations of  recent results from direct DM detection experiments~\\cite{Bernabei:2010mq,Aalseth:2011wp,Angloher:2011uu}.  Even more interestingly, the region of small mass splittings between squarks and the LSP (where the annihilation fluxes are largest) is precisely the region that is generically difficult to probe directly with collider searches for squarks; this is because the energy stored in hadronic jets coming from the decay $\\tilde{q}\\to\\chi q$ becomes too small to pass the trigger selection in this limit. In this sense, indirect DM searches are complementary to collider searches and constrain a part of the parameter space that is not directly accessible by either LEP, Tevatron or  LHC (Fig.~\\ref{fig:muconstraints}).\\footnote{ Note that also in the situation where the scalar tops are heavy, which generally helps to lift the lightest neutral Higgs boson mass to the high value of $\\sim125\\,$GeV consistent with recent indications found by both ATLAS and CMS, our results from indirect DM searches would not change and constitute important independent limits on light first or second generation squarks.} Depending on the cosmic-ray propagation model, these limits can exclude Bino masses up to 45--80 GeV in the degenerate case.  Ongoing AMS-02 measurements are expected to strengthen the limits on the annihilation cross-section by up to an order of magnitude (see \\fex~Ref.~\\cite{Evoli:2011id}), which will make it possible to probe Bino masses of 100 GeV and beyond.   \\medskip \\paragraph{Note added} During the final stages of this work, we became aware of another project analyzing the annihilation of DM into $\\bar q q g$ final states~\\cite{TAS}.  \\paragraph"
    },
    "1112/1112.1955.txt": {
        "abstract": "We present new radial velocity measurements from the Bulge Radial Velocity Assay ({\\sl BRAVA}), a large scale spectroscopic survey of M-type giants in the Galactic bulge/bar region.   The sample of $\\sim$4500 new radial velocities, mostly in the region $-10^\\circ < l < +10^\\circ$ and $b\\approx -6^{\\circ}$ more than doubles the existent published data set.   Our new data extend our rotation curve and velocity dispersion profile to $+20^\\circ$, which is $\\sim$2.8 kpc from the Galactic Center.  The new data confirm the cylindrical rotation observed at $-6^\\circ$ and $-8^\\circ$, and are an excellent fit to the Shen et al. (2010) N-body bar model.  We measure the strength of the TiO$\\varepsilon$ molecular band as a first step towards a metallicity ranking of the stellar sample, from which we confirm the presence of a vertical abundance gradient.  Our survey finds no strong evidence of previously unknown kinematic streams.  We also publish our complete catalog of radial velocities, photometry, TiO band strengths, and spectra, which is available at the IRSA archive: {\\tt http://irsa.ipac.caltech.edu/} as well as at UCLA: {\\tt http://brava.astro.ucla.edu/}. ",
        "introduction": "Only for stars in the Milky Way is it currently possible to examine both the three dimensional kinematics and composition of a central bulge/bar population, offering a unique laboratory for the study of galaxy formation and evolution.   Up until now, there have been very few optical spectra and radial velocities of bulge stars, published in a catalog form, with the spectra and measurements made publicly available.   Here we present a catalog of the {\\it Bulge Radial Velocity Assay}  (BRAVA) that has a sample of $\\sim 10,000$ M giant stars selected from the red giant branch (RGB) of the Two micron All-sky Survey \\citep[2MASS;][]{skrutskie06}.  This catalog gives the 2MASS magnitudes and positions, our measured radial velocities, and our spectra.  This publication presents our complete sample of low resolution spectra for the {\\sl BRAVA} survey, which covers approximately $-10^\\circ < l <+10^\\circ $ and $-4^\\circ <b < -8^\\circ $.  We also present a rich new data set at $b=-6^\\circ$ that complements the first results presented in Howard et~al. (2008, hereafter Paper I) and the properties of which conform to our other latitudinal studies at $-4^\\circ$ and $-8^\\circ$.  We expect that the database will ultimately be useful in constraining the dynamical model of the Galactic  bulge, and in placing limits on the detection of cold streams from fossil infall events.  The 2MASS astrometry is of sufficient quality that proper motions for our sample will be measurable in the next few years.  Although evidence for a barred potential has long been known \\citep[e.g.][]{liszt80}, the modeling of the 2 micron surface brightness distribution in the bulge as a bar \\citep{blitz91} established that the stellar distribution of the bulge is best modeled as a bar with the position angle of its major axis as $20-45^\\circ$.  The self-consistent dynamical model of a rapidly rotating bar \\citep{zhao96} was a significant advance, and a number of models have followed, however, additional progress requires constraining new models with a much larger kinematic data set.  The question of exactly how the bar structure formed -- presumably via some kind of secular evolution \\citep{kormendy04}, but including rapid star formation, winds, and dissipation -- can be better addressed with these new large samples.  Additionally, the relationship between the bulge, inner disk, thick disk and halo, may be illuminated further.  As the coverage and sample sizes analyzed in the bulge increase, it becomes possible to place meaningful limits on the influence of infalling satellite and other types of merger events.  The ultimate goal is to move beyond the characterization of the bar's morphology and orientation to the characterization of the bulge/bar as a stellar system with a unique and complex history.  Early surveys of the bulge \\citep{nassau58} immediately revealed a distinguishing characteristic, namely large numbers of M giants compared with the giant population in globular clusters and the halo.  The breakthrough in achieving a physical understanding of the M giant population arose from the surveys undertaken by Victor and Betty Blanco using the newly commissioned grating/prism on the 4m telescope at CTIO (Blanco, McCarthy, \\& Blanco 1984) -- an effort that yielded thousands of photographic low resolution classification spectra.  These studies provided the raw material that fed the first radial velocity study \\citep{mould86} and infrared surveys \\citep[e.g.][]{frogel87}.  The cool bulge M giants could be easily identified and exploited because of their brightness in the near IR (where extinction is low) and by their strong molecular (chiefly titanium oxide $\\rm[TiO]$) bands.  Sharples, Walker \\& Cropper (1990) used an early fiber spectrograph on the Anglo-Australian Telescope to obtain $\\sim 300$ spectra of bulge M giants in Baade's Window.  However, no study specifically exploited the  M giants as kinematic probes over the entire bulge, even though they could be easily identified from their extremely red $V-I$ colors.  Following the self consistent, rapidly rotating bar model of \\citep{zhao96}, two efforts modeled newly available, substantial data sets.  \\citet{sevenster99} employed the OH/IR star population to constrain a bar model, while the \\citet{beaulieu00} study used planetary nebulae (PNe) as kinematic probes of the bulge and fit a range of bar models to the data.   Both probes offered an advantage in that they are ubiquitous and easy to identify over the whole of the bulge.   In the case of the PNe, the numbers are relatively few (373).  Unfortunately, PNe are not specific to a given stellar population so disk contamination is a possibility, and the total number in the PNe sample is modest.  However, \\citet{beaulieu00} were able to test both self-consistent and N-body bar models against the PNe kinematic data.  \\citet{sevenster99} were able to draw significant constraints from the OH/IR star kinematics, finding a bar-shaped bulge at a roughly 45$^\\circ$ bar angle.  The OH/IR stars are also rare, with numbers similar to those of the PNe.  With the release of the 2MASS database, it became possible to easily select M giants over the entirety of the Galactic bulge.  The 2MASS survey offered an essentially unlimited supply of kinematic probes.  The {\\sl Bulge Radial Velocity Assay} or {\\sl BRAVA}  (see also Rich et al. 2007b and Paper I) was conceived to exploit the high quality, uniform, photometry and astrometry for this data set.    We select red giants from the K, J$-$K red giant branch that are approximately at the distance to the Galactic center, R$_{\\rm 0}$=8 kpc, and are highly likely to be bulge/bar members.  These stars can thus be selected from highly obscured regions of the bulge.  The first results of the {\\sl BRAVA} project  are given additionally in Rich et al. (2007b), Paper I, Howard et al. (2009), and Shen et al. (2010). The principle result, based on $\\sim$4500 stars, was the confirmation of cylindrical rotation, and that the simple bar/boxy bulge model matches the BRAVA kinematics strikingly well with no need for a merger-made classical bulge \\citep{shen10}.  The sample of red giant kinematic probes is now 8585.   We also have in progress a new version of the self-consistent model (Wang et al. 2011).  As described in \\citet{rich07b}, we used the Hydra spectrograph on the Blanco 4m telescope, set at a central wavelength of $\\sim$7900\\AA \\footnote{The wavelength range for each {\\sl BRAVA} field is given in our data release, as it varies slightly depending on what year the observations were taken.}.  Many details of the observation and analysis are given in Paper I and will not be repeated here.  After some experimentation, we ultimately determined that even with the prominent, partially overlapping, molecular band absorption, the near-infrared Ca triplet (CaT) offered the best opportunity to obtain excellent radial velocity cross correlation.  While originally conceived as a purely kinematic study, the data quality of {\\sl BRAVA} allows us, in principle, to investigate the metallicity distribution of the bulge.  Only with the addition of chemical abundance information will we ultimately be able to pin down the formation mechanisms that laid the kinematic traces found in our survey.  In particular, the presence of a radial or vertical metallicity gradients in the bulge may be indicative formation mechanisms \\citep{melendez08, zoccali08, babusiaux10, johnson11} and warrants a well calibrated, large-number sample, such as {\\sl BRAVA}. % Here we report the coordinates, photometry,  TiO molecular band strengths, and radial velocities of individual red giants in the {\\sl BRAVA} fields, as well as the mean velocities and velocity dispersions for each field.  The data release also includes all of our actual reduced spectra.  It is our intention to include proper motion data to the compilation, as well. % The present paper describes the final data set of {\\sl BRAVA} and presents the {\\sl BRAVA} data release and website, which includes all available data and spectra from our observing campaign. We also refine the analysis presented in \\cite{shen10} and confirm that the Galactic bulge appears to consist of a massive bar undergoing pure cylindrical rotation and yield strong limits on the presence of a classical bulge.  The structure of the paper is as follows. We present the observations and spectroscopy, especially insofar as these were changed from the setups used by Paper I and \\citet{howard09} in \\textsection1. Velocity calibrations and the main results are discussed in the next section, \\textsection2, where we also introduce our public data release. Our molecular line index measurements are then introduced in \\textsection3.  \\textsection4 shows the new data fit to the \\citet{shen10} N-body bar model, before concluding the data release in \\textsection5. ",
        "conclusions": "% % We have presented the final data set of the {\\sl BRAVA} survey and provided a website for data access, which is available at the IRSA archive: {\\tt http://irsa.ipac.caltech.edu/} as well as at UCLA: {\\tt http://brava.astro.ucla.edu/}. The observations at the $b$=$-$6$^{\\circ }$ latitude fields as well as more observations at the $b$=$-$8$^{\\circ }$ latitude fields confirms previous suggestions that the Galactic bulge rotates cylindrically, as do boxy bulges of other galaxies \\citep[e.g.,][]{kormendy82, jarvis90, shaw93, kormendy04, howard09}. The data in this paper double the number of stars shown in \\cite{howard08,howard09} and used in \\cite{shen10}.  Analysis of these data in the same fashion shows that the Galactic bulge consists of a single massive bar formed by secular evolution. The existence of such a massive bar and no classical bulge implies that our Galaxy has not suffered any large mergers (enough to form a classical bulge) since the epoch at which the disk formed. % % Based on TiO$\\varepsilon$ index measurements we qualitatively see the signatures of a vertical metallicity gradient -- an attribute that is in contrast to our kinematic characterization of the bulge as a pure pseudobulge. A detailed analysis of the {\\sl BRAVA} metallicity scale and its spatial distribution is clearly warranted and underway.  % %"
    },
    "1112/1112.5969_arXiv.txt": {
        "abstract": "{} {Estimating the marginal likelihoods is an essential feature of model selection in the Bayesian context. It is especially crucial to have good estimates when assessing the number of planets orbiting stars and the different models explain the noisy data with different numbers of Keplerian signals. We introduce a simple method for approximating the marginal likelihoods in practice when a statistically representative sample from the parameter posterior density is available.} {We use our truncated posterior mixture estimate to receive accurate model probabilities for models with differing number of Keplerian signals in radial velocity data. We test this estimate in simple scenarios to assess its accuracy and rate of convergence in practice when the corresponding estimates calculated using deviance information criterion can be applied to receive trustworthy results for reliable comparison. As a test case, we determine the posterior probability of a planet orbiting HD 3651 given Lick and Keck radial velocity data.} {The posterior mixture estimate appears to be a simple and an accurate way of calculating marginal integrals from posterior samples. We show, that it can be used to estimate the marginal integrals reliably in practice, given a suitable selection of parameter $\\lambda$, that controls its accuracy and convergence rate. It is also more accurate than the one block Metropolis-Hastings estimate and can be used in any application because it is not based on assumptions on the nature of the posterior density nor the amount of data or parameters in the statistical model.} {} ",
        "introduction": "The selection between a collection of candidate models is of significant in all fields of astronomy but especially so, when the purpose is to extract weak planetary signals from noisy data. The ability to tell whether a signal is present in data as reliably as possible is essential in several searches for low-mass exoplanets orbiting nearby stars, whether made using the Doppler spectroscopy method; e.g. the Anglo-Australian Planet Search \\citep[e.g.][and references therein]{tinney2001,jones2002}, High-Accuracy Radial Velocity Planet Searcher \\citep[e.g.][and references therein]{mayor2003,lovis2011}, Hich Resolution Echelle Spectrometer \\citep[e.g.][and references therein]{vogt1994,vogt2010}; by searching photometric transits; e.g. Convection Rotation and Planetary Transits \\citep[e.g.][and references therein]{barge2007,hebrard2011}, WASP \\citep[e.g.][and references threrein]{colliercameron2007,faedi2011}; or other possible techniques, such as astrometry \\citep[e.g.][]{benedict2002,pravdo2009} and transit timing \\citep[e.g.][]{holman2005} or other current or future methods.  Using Bayesian tools, it is possible to determine the relative probabilities for each statistical model in some selected collection of models to assess their relative performance, or relative ability to explain the data in a probabilistic manner. This is also important in the context of being able to assess their inability to explain several data sets in terms of the model inadequacy of \\citet{tuomi2011b}. Especially, when different statistical models contain different numbers of planets orbiting the target star, assessing their relative posterior probabilities given the measurements is extremely important to detect all the signals in the data \\citep[e.g.][]{gregory2005,gregory2007a,gregory2007b,tuomi2009,tuomi2012} and to avoid the detection of false positives \\citep[e.g.][]{bean2010,tuomi2011}. However, determining the posterior probabilities require the ability to calculate marginal integrals that are complicated multidimensional integrals of likelihood functions and priors over the whole parameter space. While there are several methods of estimating the values of these integrals, those that are computationally simple and easy to implement are more often than not the poorest ones with respect to their accuracy and convergence properties \\citep[e.g.][]{kass1995,clyde2007,ford2007}. There are also more complicated methods for estimating multidimensional integrals but they may provide more difficult computational problems themselves than typical data analyses are, which makes it difficult to use them in practice.  Because of these difficulties and the need to be able to assess the marginal integrals reliably, we introduce a simple method for estimating the marginal integrals in practice if a statistically representative sample from the parameter posterior density exists. As such a sample is usually calculated when assessing the posterior densities of model parameters using posterior sampling algorithms \\citep[e.g.][]{metropolis1953,hastings1970,haario2001}, the ability to use the very same sample in determining the marginal integral is extremely useful in practice. There are methods for taking advantage of the posterior sample in this manner \\citep[e.g.][]{newton1994,kass1995,chib2001,clyde2007} but their performance, despite some studies \\citep[e.g.][]{kass1995,ford2007}, is not generally well known, especially so in astronomical problems, and some of them may also require samplings from other densities simultaneously, such as the prior density or the proposal density of the Metropolis-Hastings (M-H) output, making their application difficult.  In this article, we introduce a simple method that can be used to receive accurate estimates for the marginal integral. We test our estimate, called the truncated posterior mixture (TPM) estimate in scenarios where the marginal integral can be calculated accurately using simple existing methods. The deviance information criterion \\citep[DIC,][]{spiegelhalter2002} is asymptotically an accurate estimate when the sample size, i.e. the sample drawn from the posterior density, increases and can be used if the posterior is a multivariate Gaussian. Therefore, we compare our estimate with the DIC estimate in such cases to test its accuracy in practice. If accurate, our estimate is applicable whenever a statistically representative sample from the posterior is available because we do not make any assumptions regarding the shape of the posterior density when deriving the TPM estimate. The only assumptions are, that such a sample exists and it is statistically representative. We also calculate the marginal likelihoods using the simple Akaike information criterion (AIC) for small sample size \\citep{akaike1973,burnham2002}, the harmonic mean (HM) estimate that is a special case of the TPM with poor convergence properties, and the One Block Metropolis-Hastings (OBMH) method of \\citet{chib2001} that requires the simultaneous sampling of posterior and proposal densities. \\citet{kass1995} and \\citet{clyde2007} give detailed summaries of different methods in the context of model selection problems.  Finally, we also test the performance of the TPM estimate and the effects of prior choice in simple cases where it is possible to calculate the marginal integral from a sample from the prior (with the common mean estimate) and/or using direct numerical integration. Especially, we show the undesirable effects of Bartlett's paradox on the marginal integrals and demonstrate that the TPM estimate actually circumvents these effects in practice. ",
        "conclusions": "Calculating the marginal integral for model selection purposes is generally a challenging computational problem. While there are several good estimates for these integrals, they are usually only applicable under certain limiting assumptions about the nature of the posterior density, the amount of parameters in the statistical model, or the number of measurements available. Therefore, we have introduced a new method for estimating these integrals in practice. Given the availability of a sample from the posterior density of model parameters, our truncated posterior mixture estimate is a reasonably accurate one and very easily calculated in practice. We have only assumed that a statistically representative sample drawn from the posterior density exists when deriving our posterior mixture estimate (see Appendix). Therefore, it is applicable to any model comparison problems in astronomy and other fields of scientific inquiry and is not restricted to problems where the posterior has a certain shape and dimension.  The comparisons of different estimates given the radial velocities of HD 3651 revealed that the TPM yields estimates very close to the DIC estimate, which is known to be a reliable one in case of Gaussian posterior density. In fact, we chose the HD 3651 as an example star because of the planet orbiting it is known to have an eccentric orbit that enables the Gaussianity of the probability distributions of eccentricity and the two angular parameters of the Keplerian model, namely, longitude of pericentre and mean anomaly. However, the simple small-sample version of the AIC proved reasonably accurate as well when the number of measurements well exceeded the number of free parameters of the model (e.g. Table \\ref{combined_bf}). We also note that the OBMH estimate of \\citet{chib2001}, while converging rapidly, tends to yield somewhat biased results that exaggerate the posterior probability of the more complicated model, making it possibly -- at least in the test cases of the current work -- prone to detections of false positives.  In practice, the TMP can be used by calculating its value directly from the sample drawn from the posterior density of the model parameters. Selecting a suitable value for parameter $\\lambda$ is then of essence when calculating its value in practice. In all the three different test cases studied in this article, a choice of $\\lambda = 10^{-4}$ yielded estimates that converged rapidly for all the models in all the test cases and resulted in posterior probabilities that differed little from those calculated using the DIC estimate. When the difference between the two models was the smallest (case 3.), there was practically no bias in the TMP estimate with respect to the DIC. Also, when the posterior odds of the one-Keplerian model was the greatest (case 2.), the TPM, with $\\lambda = 10^{-4}$, overestimated the posterior probability of the one-Keplerian model by a factor of 10, though, in that case, the Bayes factor used in model selection was already so heavily in favour of the one-Keplerian model that this overestimation is not significant in practice in terms of being able to select the best model.  Because of the possible biases caused by too large $\\lambda$, it would then be convenient in practice to calculate the TPM estimate using few different values of parameter $\\lambda$. With the sample from the posterior density available, this could be done with little computational cost. Then, it would be possible to use the lowest value for $\\lambda$ that still converges to receive a trustworthy TPM estimate and correspondingly trustworthy model selection results in any model selection problem.  Finally, because any constant coefficients in the prior probability densities have an effect on the marginal integrals by corresponding to different prior weights for different models, we have shown how the TPM estimate deals with this problem. Effectively, it corresponds to setting the constant coefficients in the prior equal to unity, which makes the TPM estimate independent of the unit choice of the parameters."
    },
    "1112/1112.2247.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory \\abstract {The nova outburst of V407 Cyg in 2010 Mar. 10 was the first observed for this star but its close resemblance to the well known symbiotic-like recurrent nova RS Oph suggests that it is also a member of this rare type of Galactic novae.  The nova was the first detected at $\\gamma$-ray energies and is the first known nova explosion for this system.   The extensive multiwavelength coverage of this outburst makes it an ideal comparison with the few other outbursts known for similar systems. We extend our previous analysis  of the Mira and the expanding shock from the explosion to detail the time development of the photoionized Mira wind, circumstellar medium, and shocked circumstellar environment to derive their physical parameters and how they relate to large scale structure of the environment, extending the previous coverage to more than 500 days after outburst.  We use  optical spectra obtained at high resolution with the Nordic Optical Telescope (NOT) (R$\\approx$45000 to 65000) and medium resolution Ond\\v{r}ejov  Observatory (R$\\approx$12000) data and compare the line variations  with publicly available archival measurements at 30 GHz OVNR and at X-rays with Swift during the first four months of the outburst, through the end of the epoch of strong XR emission.   We use nebular diagnostics and high resolution profile variations to derive the densities and locations of the extended emission.  We find that the higher the ionization and/or the higher the excitation energy, the more closely the profiles resemble the He II/Ca V-type high velocity shock profile discussed in Paper I.  This also  accounts for the comparative development of the [N II] and [O III] isoelectronic transitions: the [O III] 4363\\AA\\ profile does not show the low velocity peaks while the excited [N II] 5754\\AA\\ does.  If nitrogen is mainly N$^{+3}$ or higher in the shock, the upper state of the [N II] nebular lines will contribute but if the oxygen is O$^{+2}$ then this line is formed by recombination, masking the nebular contributor, and the lower states are collisionally quenched but emit from the low density surroundings.    Absorption lines of Fe-peak ions formed in the Mira wind were visible as P Cyg profiles at low velocity before Day 69, around the time of the X-ray peak and we identified many absorption transitions without accompanying emission for metal lines.   The H Balmer lines showed strong P Cyg absorption troughs that weakened during the 2010 observing period, through Day 128.   The Fe-peak  line profiles and flux variations were different for permitted and forbidden transitions: the E1 transitions were not visible after Day 128 but had shown a narrow peak superimposed on an extended (200 km s$^{-1}$) blue wing,  while the M1 and E2 transitions persisted to Day 529, the last observation, and showed extended redshifted wings up of the same velocity.    We distinguish the components from the shock, the photoionized environment, and the chromosphere and inner Mira wind using spectra taken more than one year after outburst.  The multiple shells and radiative excitation phenomenology are similar to those recently cited for GRBs and SNIa.} % } % conclusions heading (optional), leave it empty if necessary ",
        "introduction": "The 2010 outburst of V407 Cyg, a previously classified D-type symbiotic star (Munari, Margoni, \\& Stagni 1990) whose historical variability had been classified as ``nova-like'' by Hoffmeister (1949) and Meinunger (1966), developed photometrically and spectroscopically almost identically to the most recent eruptions of the symbiotic-like recurrent nova (SyRN) RS Oph (1985, 2006).  This event moved the system into that rare group of explosive cataclysmic variables.   In Shore et al. (2011, hereafter  paper I)  we presented the first results of the outburst based on multiwavelength photometry and spectroscopy from the Nordic Optical Telescope, Ond\\v{r}ejov Observatory, {\\it Swift}, and {\\it Fermi/LAT} concentrating on the Mira properties and the time development of the shock propagating through the red giant (RG) wind.  In this second paper, we focus on the evolution of the circumstellar environment  that was photoionized and illuminated by the emission from the white dwarf (WD) and the shock.  We extend the previous results to include additional proxies for the high energy excitation mechanisms.   %\\justification %__________________________________________________________________ ",
        "conclusions": "Among the known symbiotic stars, V407 Cyg is unique.  Even compared with the limited sample of the symbiotic-like recurrent nova systems, the V407 Cyg system is unusual (Anupama 2008, Mikolajewska 2008, Bode 2010).  Among classical symbiotic systems, its Mira is the longest period pulsator known (Munari et al. 1990, Belczy\\'nski et al. 2002).   The progenitor of the Mira must have been massive, $>$4 M$_\\odot$, based on the lithium excess (Tatarnikova  et al., 2003a,b; Paper I;  Munari et al. 2011).  Only T CrB and RS Oph, both SyRNe, show significantly enhanced $^7$Li (Wallerstein et al. 2008), reinforcing the identification of V407 Cyg as a member of this class notwithstanding that the 2010 outburst is the first {\\it known} such event in this system.  It is the only SyRN system to show maser emission (Deguchi et al. 2011), the individual components of which varied during the outburst.  But perhaps most unusual was the event that led to its being labeled ``nova-like'' in the literature, a symbiotic nova-like event in 1936 that lasted for about five years during which time the visual brightness of the system increased by more than a factor of 10, a behavior that is {\\it not} seen historically among the other SyRNe (in addition to T CrB and RS Oph, neither V745 Sco nor V3890 Sgr appear to have such variability).  There were no contemporaneous spectroscopic observations during the event of the 1930s but, if it were like other systems, there would have been continuous mass loss through a wind from the white dwarf induced by the onset of a thermonuclear envelope process following a low level of accretion.\\footnote{The historical record is quite poorly sampled before 1990 and we know very little of what V407 Cyg was doing between the outburst of the 1930s and the beginning of photometric and spectroscopic monitoring in the 1990s.  The unanticipated explosion of V407 Cyg means it was not included in the photographic survey by Schaefer et al. (2009) of Galactic recurrent novae.}  The nova explosion produced a rich variety of phenomena, as we have detailed, that distinguish this outburst from other SyRNe.  The first indication of [O III] 4959, 5007\\AA\\ was on JD 55272 (Day 6) ; [N II]  6548, 6583\\AA\\  only appeared after JD  55289 (Day 24).  The narrow [O I] 6300, 6363\\AA\\ and Mg I 4671\\AA\\ emission lines were, in contrast, clearly visible from Day 4.   In the Day 4 spectrum, narrow emission lines from permitted Fe-peak ions were seen without the underlying broad wings; those appeared by Day 6 and had become asymmetric in opposite senses (as we discussed above) by Day 13.   The first appearance of the double peak emission on [O III] 4959\\AA\\ can be estimated from the comparison of the lower resolution Ond\\v{r}ejov spectrum on JD 55292 and the first NOT spectrum on JD 55288.  This gives a limit to the radius of the ionized region of $\\leq$2$\\times$10$^{16}$ cm.  Assuming that this is matter ejected during the outburst in the 1930s, the width is $\\Delta R/R \\leq 0.15$ based on the duration of the event (about 5 yr)  and using the electron density from the [N II] analysis, $n_e \\approx 3\\times$ 10$^4$ cm$^{-3}$,  the ionized mass sampled by the nova line profile is $\\leq$10$^{-4}$ M$_\\odot$.  If this is due to a wind during a symbiotic nova-like event, the inferred mass loss rate was  $\\leq 2\\times$ 10$^{-5}$ M$_\\odot$yr$^{-1}$ independent of the distance.   The density is lower than, but roughly compatible with, the electron density inferred from the recombination analysis of the Na I D lines at this velocity discussed in Paper I, about 10$^5$ cm$^{-3}$.  The disappearance of the blue component can be interpreted as recombination in a zone located on the side of the giant toward the observer and shadowed from the WD.  The redward extension would be, in this picture, the low density periphery that is still illuminated by the WD and also collisionally excited by the shock.  The similar maximum positive radial velocity on all the forbidden transitions and the absence of the permitted lines, point to the emission arising in a very low density region.  The distinct taxonomic separation of the line profile asymmetry between permitted and forbidden transitions supports the picture that the blue wing of the E2 and M1 lines is suppressed by higher densities being encountered in the approaching part of the shock. and the disappearance of this component on almost all of the strong lines in the last spectrum being due to recombination.    The gap between the 2010 and 2011 observing sessions is, however, too long to constrain the densities.  \\begin{figure} \\centering \\includegraphics[width=8cm]{p2f34.ps} %%%\\includegraphics{empty.eps} %%%\\includegraphics{empty.eps} \\caption{An example of the separation of the various contributors to the observed line profiles in the 2011 Aug. 21 NOT spectrum.  Top: [O I] 6300\\AA.  Bottom: [O III] 5007\\AA\\ (solid), Mg I 4571\\AA\\ (dot), [S II] 6716\\AA\\ (dash), [N II] 6548\\AA\\ (dot-dash).  The Mira continuum has been removed in each but the scaling is for display (arbitrary flux units).  The spike at +60 km s$^{-1}$ is a cosmic ray hit near Mg I.   See text for description.} \\end{figure}  The late-epoch spectrum from JD 55794 further clarifies the various contributions to the profiles.  As an example, we used the [O I] 6300\\AA\\ line (which was identical in profile to the related 6363\\AA\\ line) for the decomposition, shown in the upper panel of Fig. 33.  In the lower panel of Fig. 33 we show an empirical decomposition using a set of representative profiles from the data itself with the continuum of the Mira subtracted.  The relative intensities are chosen for display, they are not results from a fitting procedure.   The Mg II line was narrow with a FWHM of 25 km s$^{-1}$ centered at -61$\\pm$1 km s$^{-1}$, the same radial velocity previously derived for the Mira wind absorption lines.   The [N II] 6548\\AA\\ showed an extended red wing to +200 km s$^{-1}$ and the persistent red peak.  This latter emission peak coincided with one on the [S II] 6716 and 6730\\AA\\ lines that also showed the reduced blue peak, like [Ar III] 7135 (Fig. 32).   The [O III] 4959\\AA\\ and 5007\\AA\\ lines no longer displayed the low velocity emission peaks, instead having a pure shock profile similar to [Ca V] 5309\\AA.  The comparison indicates,however,  that a part of the [O I] profile is still being contributed by the part of the wind that has not been completely ionized and, as we proposed in Paper 1, the wing may be due to charge exchange behind the shock.\u00caThe [S II] and [N II] lines still showed  contributions from the extended environment.  Deguchi et al. (2011) present the variations of the SiO maser during the first month of outburst.  Taking $V_{Mira,LSR} = -26$ km s$^{-1}$, the main SiO  (J=1-0, v=2) component is at -30 km s$^{-1}$ (LSR), corresponding in heliocentric $v_{rad}$ to -58 km s$^{-1}$,  at slightly higher velocity ($\\approx$ 4 km s$^{-1}$) than the peak of the [O I] 6300\\AA\\ line on JD 55286 and the mean absorption line velocity but compatible with the wind line velocities. \\footnote{There is an ambiguity here.  Deguchi et al. give the  system heliocentric radial velocity of -41 km s$^{-1}$, and discuss orbital motion with a 43 year period  citing the radial velocity from Tatarnikova et al. (2003b) using the Ca I 6573\\AA\\ line.  We detected this line in the last NOT spectra from 2010 and in the 2011 spectrum at $v_{\\rm rad}$=-43 km s$^{-1}$.  There was, as we mentioned above, a weak absorption feature on Na I D at a similar velocity and a weak absorption component on Mg I 5167\\AA.  The velocity disagreed, however, with all other absorption line values and we have no explanation for this other than either a blend or emission filling in the blue wing.  We adopt the velocity determined from the Li I 6707\\AA\\ and [Ca II] 7291, 7323\\AA\\ lines in our comparison.}    The (J=1-0, v=1) transition also shows a weaker line at -47 km s$^{-1}$ that does not corresponding to any absorption on the Na I D lines.   The other SiO maser line (J=1-0, v=2) shows only the -58 km s$^{-1}$ line.  The velocity is consistent with that of the Mira wind as determined from the absorption features discussed in the previous sections.  The rapid turn-off of the maser after the explosion, if due to the irradiation by the initial shock and the advance of the ionized region as a precursor, indicates that the cold material is located at a distance of about 1 to 10 AU from the Mira, consistent with the distance in Deguchi et al (2011). %We should note that the timescale for the disappearance of the wind absorption lines, $\\leq$45 days, is longer than the timescale of about 20 days on which the maser lines turned on again.  The orientation proposed by Deguchi et al. (2011) from the maser emission places the WD near quadrature at the time of the nova explosion.  The chromospheric lines, and the ``emission hole'' in the forbidden lines, are more consistent with a phase closer to superior conjunction for the WD -- that is, being seen through a substantial portion of the Mira wind.  Two results point to this.  The absorption line spectrum during the early outburst was strong, both on P Cyg components and in isolated absorption lines, and are more consistent with a long pathlength.  The other is the asymmetry of the high velocity component of the lines.  Were the WD seen nearer quadrature, it would not be possible to produce so asymmetric a shock line profile.   The persistence of the neutral lines also support the picture that the inner portion of the wind was screened from the UV of the explosion and shock while the bulk of the wind was ionized within the first 60 days.  If the fluorescent emission was excited by the UV of the WD, the disappearance of those transitions could indicate either that the ionization had removed the ions (mainly singly ionized Fe-peak transitions) or that the WD had turned off.  Monitoring observations of V407 Cyg at centimeter wavelengths with the EVLA have also shown the appearance, after Day 160, of a resolved nonthermal component at a few times $10^{15}$ cm, depending on the distance adopted for the Mira.\\footnote{The EVLA spectra are available at the URL:\\\\   https://safe.nrao.edu/evla/nova/\\#v407cyg.}   This is compatible with the location of the gas that we suggest was ejected during the 1930s event.   The shock, if expanding as inferred in Paper I with a simple power law $v \\sim t^{-1/3}$, should have reached a distance of a few times $10^{15}$ cm in about one year.  The 2011 Apr. 23 Ond\\v{r}ejov spectrum shows that the [N II] lines had developed extended redward emission with a maximum velocity of about +300 km s$^{-1}$, which is the velocity we expect from the shock at that time.\\footnote{As an aside we note that from the [O III] and [N II] diagnostics, the density at this velocity is compatible with a compression ratio of about an order of magnitude, indicating that the shock is likely now isothermal, and for this compression ratio we would expect a steep spectrum with an exponent of about -3.3 (Blandford \\& Ostriker 1978) and a strong shock of order $10^3$ in pressure ratio. }    %  correction for LSR to Hel = 12.5 km/s"
    },
    "1112/1112.4532_arXiv.txt": {
        "abstract": "We extend the two-dimensional Cartesian shapelet formalism to $d$-dimensions. Concentrating on the three-dimensional case, we derive shapelet-based equations for the mass, centroid, root-mean-square radius, and components of the quadrupole moment and moment of inertia tensors. Using cosmological $N$-body simulations as an application domain, we show that three-dimensional shapelets can be used to replicate the complex sub-structure of dark matter halos and demonstrate the basis of an automated classification scheme for halo shapes. We investigate the shapelet decomposition process from an algorithmic viewpoint, and consider opportunities for accelerating the computation of shapelet-based representations using graphics processing units (GPUs). ",
        "introduction": "Complex, three-dimensional structures abound in astronomy on all scales from ``fluffy'' dust aggregrates in molecular clouds (Ossenkopf 1993; Stepnik et al. 2003), to cosmological large-scale structure that has been described as ``sponge-like'' (Gott, Dickinson \\& Melott 1986), or a ``skeleton'' (Sousbie et al. 2008) of clusters, filaments and voids (Barrow, Bhavsar \\& Sonoda 1985; White et al. 1987).  While aspects of these structures can be expressed in terms of simple, geometrically-motivated properties such as their triaxiality or quadrupole moment, these quantities are not able to capture the higher order complexity of the true shape.  The challenge, therefore, is to provide an accurate description of an arbitrary three-dimensional (3-d) shape, possibly over many physical length scales, in the hope that this can lead to improved theoretical or analytical insight into the structure in question.   The human visual system is more than capable of identifying structures and sub-structures for an individual 3-d object, but such qualitative interpretations only have limited use -- it is not practical to attempt a classification of shapes by eye when there are many thousands of objects to inspect.\\footnote{Although, if there are enough individual eyes available to assist, then this approach is feasible, as the Galaxy Zoo project ({\\tt http://www.galaxyzoo.org}) has demonstrated.}  The preferred alternative is an automated approach including: \\begin{itemize} \\item decomposition via an appropriate basis set (e.g. Fourier analysis, wavelet transformations); \\item partitioning [e.g. Voronoi tesselation - see Icke \\& van de Weygaert (1987) for an early cosmological application]; \\item the use of minimal spanning trees to identify connected structures (Barrow et al. 1985; Pearson \\& Coles 1995); \\item Minkowski functionals [which return global geometric properties such as volume, surface area and edge density -- Mecke, Buchert \\& Wagner (1994); Sahni, Sathyaprakash \\& Shandarin  (1998)]; and \\item segementation [e.g. ``dendrograms'' used by Goodman et al. (2009) to identify self-gravitating structures in molecular clouds]. \\end{itemize}  The approach we present in this paper is the extension of the two-dimensional (2-d) shapelet method (Refregier 2003) to three dimensions. Shapelets are sets of orthonormal basis functions based on the Hermite polynomial solutions of the quantum harmonic oscillator (QHO). Simple analytic forms can be derived for the physical properties of 3-d structures (e.g. centre of mass, root-mean-square radius and the components of the quadrupole moment and moment of intertia tensors), which can be efficiently calculated in shapelet space.   In astronomy, 2-d shapelets have been applied to image simulation (Massey et al. 2004; Ferry et al. 2008), the morphological classification of galaxies (Kelly \\& McKay 2004; Andrae, Jahnke \\& Melchior 2011) and sunspots (Young et al.  2005), and weak gravitational lensing (Refregier \\& Bacon 2003). The latter includes the measurement of shear (Kuijken 2006), flexion (Goldberg \\& Bacon 2005), point-spread function modelling and deconvolution (Melchior et al. 2009; Paulin-Henriksson, Referegier \\& Amara 2009), and weak lensing by large-scale structure from the FIRST radio survey (Chang, Refregier \\& Helfand 2004).  Massey et al. (2007) investigated weak lensing with polar shapelets (Massey \\& Refregier 2005), a form more suitable for images with rotational symmetry.  Further properties of shapelets, including integral relations and convolution sums are presented in Coffey (2006).  The importance of the shapelet approach lies not so much in the basis functions, but in the simplifed computation of quantities relating to shape and structure that can be determined once a shapelet decomposition has been obtained.  These analytic quantities are expressed as linear sums of weighted shapelet states, greatly reducing the calculation complexity compared to (numerically) solving the related integral formulations.  Shapelet decomposition is not without its problems [see Berry, Hobson \\& Withington (2004) for an extensive discussion].  Melchior, Meneghetti \\& Bartelmann (2007) examined the limitations of shapelet image analysis in cases where the orthonormality condition [see equation (\\ref{eqn:ortho}) below] fails, and proposed a  decomposition procedure that preserves physical properties of images.  Melchior et al. (2010) and Bosch (2010) considered problems with using circular Gaussian basis functions to model galaxies with high ellipticity or a large S\\'{e}rsic index. Ngan et al. (2009) proposed an alternative orthonormal basis based on the S\\'{e}rsic profile (hence S\\'{e}rsiclets) for use in weak lensing analysis.  While helping to avoid issues with poor shape recovery from overfitting low signal-to-noise galaxies, and fitting with too many degrees of freedom, S\\'{e}rsiclets do not possess the analytic properties of shapelets, and the basis functions must be generated numerically.  Indeed, it is the existence of analytic functions that has motivated our choice of 3-d Cartesian shapelets as an appropriate tool for quantifying properties of three-dimensional structures.  The remainder of this paper is set out as follows. In Section \\ref{sec:cartesian}, we present the mathematics of 3- and $d$-dimensional Cartesian shapelets. New analytic expressions are presented for several important physical properties of 3-d structures in Section \\ref{sct:analytic}.  In Section \\ref{sct:implement}, we describe issues relating to implementing an efficient 3-d shapelet decomposition code.  We highlight the inherent high-degree of paralellism in the shapelet decomposition algorithm, which makes it a promising target for graphics processing units. In Section \\ref{sct:cosmoapp}, we present first applications of 3-d shapelets to problems in cosmological simulations, with an emphasis on studying sub-structure in dark matter halos, demonstrating how an automated shape classifier can work in shapelet space. We end with a summary and outlook for 3-d shapelets in astronomy in Section \\ref{sct:conc}.   \\begin{figure*} \\begin{center} \\includegraphics[width=7in]{figure1.pdf} \\end{center}  \\caption{Examples of three-dimensional Cartesian shapelets ($\\beta = 1$). Top row: (left) ${\\bmath n} = (0,0,0)$; (right) ${\\bmath n} = (0,1,0)$. Bottom row: (left) ${\\bmath n} = (2,0,2)$; (right) ${\\bmath n} = (1,2,4)$. For each panel, we calculate the maximum data value, $f_{\\rm max}$, and generate 10 equally spaced iso-surfaces over the range $(-f_{\\rm max}, f_{\\rm max})$. Individual isosurfaces are coloured with a two-ended intensity colour map: blue $\\rightarrow$ black $\\rightarrow$ orange.  \\label{fig:B000}} \\end{figure*} ",
        "conclusions": "\\label{sct:conc} We have extended the two-dimensional Cartesian shapelet formalism of Refregier (2003) to three dimensions, deriving analytic expressions for the zeroth moment, object centroid, root-mean-square radius, and the components of the quadrupole moment and moment of inertia tensors. We also presented generalisations to $d$-dimensions.  Further work is necessary to develop a robust and systematic optimisation strategy for the decomposition parameters, and the development of specfic applications for the three-dimensional shapelet technique requires such a strategy. There are also opportunities to develop the formalism further, specifically extending it to include spherical shapelet functions [c.f. the alternative presentation of two-dimensional Cartesian shapelets as polar shapelets by Massey \\& Refregier (2005)].  The shapelet decomposition algorithm exhibits attributes that make it an ideal target for implementation on modern, massively-parallel GPUs. Our algorithm analysis demonstrates that the computation is entirely (or embarassingly) parallel; has minimal or no branching; maintains a high ratio of arithmetic operations to memory-access operations; and has a memory access pattern that will result in aligned or contiguous access to memory, required for achieving a high memory throughput. With our proposed scheme of precomputing shapelet voxel-integral terms, the computation reduces to a parallel series of multiply-add operations, which are almost ideal for GPUs -- we anticipate achieving close to peak processing performance.  Significantly reducing the computation time for the shapelet  decomposition, compared to CPU, means that more processing time is then available for optimisation.  As an example application, we have demonstrated how three-dimensional shapelets can be used to study the complex sub-structures of dark matter haloes from cosmological $N$-body simulations, including providing an alternative approach to classifying the properties of haloes.  Our preliminary investigation suggests that halo triaxiality measured purely from the moment of inertia tensor may be incorrect due to limitations of group finders that are not able to separate out what may be truly distinct sub-clumps.  Improvements to our current `by eye' approach to classification could include development of a decision tree or neural network classifier, or the use of principle component analysis to significantly reduce the number of shapelet terms required for classification (Kelly \\& McKay 2004).  The shapelet formalism is virtually unexplored in the three-dimensional domain,  offering opportunities for the further development of a methodology that can be used to quantify and analyse complex three-dimensional structures.  Future applications of the three-dimensional shapelet techinique may include classification and parameterisation of sources identified in H{\\sc i} spectral line data cubes; studying the shapes of voids in cosmological simulations (by considering an inverted density field); and the possibility to generate mock dark matter haloes through an extensive study of the distribution of shapelet amplitudes as a function of mass and triaxiality."
    },
    "1112/1112.1910_arXiv.txt": {
        "abstract": "{It has been proposed that the envelopes of luminous stars {{\\changedA may}} be subject to substantial radius inflation. The peculiar structure of such inflated envelopes, with an almost void, radiatively dominated region beneath a thin, dense shell could mean that many in reality compact stars are hidden below inflated envelopes, displaying much lower effective temperatures.  {{\\changedA The inflation effect has been discussed in relation to the radius problem of Wolf-Rayet (WR) stars, but has yet failed to explain the large observed radii of Galactic WR stars.}}  } { We wish to obtain a physical {{\\changedA perspective of the inflation effect}}, and study the consequences for the radii of Wolf-Rayet (WR) stars, and luminous blue variables (LBVs).  For WR stars the observed radii are up to an order of magnitude larger than predicted by theory, whilst S Doradus-type LBVs are subject to humongous radius variations, which remain as yet ill-explained.}  {{{\\changedA We use a dual approach to investigate the envelope inflation, based on numerical models for stars near the Eddington limit, and a new analytic formalism to describe the effect.  An additional new aspect is that we take the effect of density inhomogeneities (clumping) within the outer stellar envelopes into account.}}} {{{\\changedA Due to the effect of clumping we are able to bring the observed WR radii in agreement with theory.  Based on our new formalism, we find that the radial inflation is a function of a dimensionless parameter $W$, which largely depends on the topology of the Fe-opacity peak, i.e., on material properties.}} For $W>1$, we {{\\changedA discover}} an instability limit, for which the stellar envelope becomes gravitationally unbound, i.e.\\ there no longer exists a static solution.  Within this framework we {{\\changedA are also able to explain the S\\,Doradus-type instabilities}} for LBVs like AG\\,Car, with a possible triggering due to changes in stellar rotation.  } {The stellar effective temperatures in the upper Hertzsprung-Russell (HR) diagram are potentially strongly affected by the inflation effect.  This may have particularly strong effects on the evolved massive LBV and WR stars just prior to their final collapse, as the progenitors of supernovae (SNe)\\,Ibc, SNe\\,II, and long-duration gamma-ray bursts (long GRBs).} ",
        "introduction": "\\label{sec:intro}  In the standard picture for the evolution of the most massive stars (with $M \\ga 30\\,\\msun$), O stars have been proposed to evolve through successive luminous blue variable (LBV) and Wolf-Rayet (WR) phases before exploding as hydrogen (H) free supernovae (SNe) Ibc \\citep[e.g.,][]{lan1:94,mey1:03,heg1:03}.  This evolutionary path is however still open as LBVs have more recently been suggested to be the direct progenitors of some H-rich type II SNe \\citep{kot1:06,smi2:07,gal1:09}. Additional renewed interest in WR stars arises from their connection to long-duration gamma ray bursts \\citep[long GRBs, cf.][]{yoo1:05,woo1:06}.  What LBVs and WR stars have in common is their proximity to the Eddington limit.  Not only is this thought to be instrumental for their mass-loss behavior \\citep{vin1:02,gra1:08,vin1:11,gra1:11}, but it may also have key consequences for their stellar structures.  Using modern OPAL opacities, \\citet{ish1:99}, \\citet{pet1:06}, and \\citet{yun1:08} studied the interior structure of massive stars, revealing a ``core-halo'' configuration that involves a relatively small convective core and an extended radiative envelope. This is referred to as the ``inflation'' of the outer envelope.  It has been known for many decades that the observed radii of WR stars are almost an order of magnitude larger than canonical WR stellar structure models indicate. This is oftentimes attributed to their so-called pseudo-photospheres, which concern an ``effective'' photosphere several times larger than the hydrostatic radius, as a result of a dense stellar outflow \\citep[e.g.,][]{cro1:07}. Pseudo-photospheres have also been discussed in the context of LBVs, although the issue is still under debate \\citep[e.g.,][]{smi1:04}.  An alternative explanation for the large radii of WR stars may involve envelope inflation. The issue is particularly relevant in the context of WR stars as GRB progenitors, as there have been several suggestions that the required radii of GRB progenitors are up to an order of magnitude ($\\sim$\\,$1\\,R_\\odot$ vs. $\\sim$\\,$10\\,R_\\odot$) smaller than the radii of observed WR stars \\citep{mod1:09,cui1:10}.  The WR radius issue is thus highly relevant for GRB modeling.  For LBVs the issue of their radii is equally interesting. The defining property of LBVs concerns their ``S Doradus'' cycles. On timescales of years to decades LBVs are seen to vary between effective temperatures of $\\sim$30\\,kK (early B spectral type) to $\\sim$8\\,kK (early F). Due to the lack of convincing counter-evidence, these S Dor excursions in the upper Hertzsprung-Russell diagram have generally been assumed to occur at constant bolometric luminosity \\citep{hum1:94}, but recent studies have challenged this behavior for a couple of objects \\citep{gro1:09,cla1:09}. Either way, the issue of the increased LBV radii during S Dor cycles remains un-questioned, but a satisfying explanation for it has yet to be provided \\citep{vin1:09}.  Outer envelope inflation may turn out to be an interesting explanation for S Dor variations, as will be detailed in the following.  The paper is organized as follows.  In Sect.\\,\\ref{sec:modelcomp} we present numerical stellar structure models that show the envelope inflation effect. In Sect.\\,\\ref{sec:anal} we describe the effect analytically, {\\changedA including a new instability limit. We also provide a recipe that can be used to estimate the inflation effect for arbitrary stars with given (observed) parameters}. The implications of our results are discussed in Sect.\\,\\ref{sec:implications}, with focus on WR\\,stars and LBVs. The conclusions are summarized in Sect.\\,\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In the present work, we have discussed the possible formation of radially inflated stellar envelopes for stars approaching the Eddington limit, as originally proposed by \\citet{ish1:99}, and \\citet{pet1:06}. {\\changedA In addition to numerical models, we could provide an analytic description of this process, leading to the discovery of a new instability limit, and a clumping dependence. Both effects turn out to be of profound importance. While the former may be connected to S Doradus-type instabilities in LBVs, the latter can account for the large observed radii of Galactic WR stars, and thus resolve the long standing WR radius problem.}  {\\changedA Within our analytical approach we could describe the inflation effect in terms of a dimensionless parameter $W$} (Eqs.\\,\\ref{eq:Qrecipe}, \\ref{eq:W}) that characterizes the ratio between internal and gravitational energy at the base of the envelope. $W$ can be estimated on the basis of opacity tables and a combination of stellar parameters. The envelope inflation occurs when $W$ approaches one. For $W\\ge 1$, we find that the envelope becomes gravitationally unbound, i.e., there exists no static solution.  A prerequisite for envelope inflation is the proximity of the star to the Eddington limit, with Eddington factors of $\\Gamma_{\\rm e}>0.3$--0.35, at solar metallicity.  For a given chemical composition and $\\Gamma_{\\rm e}$, the resulting stellar radius depends on the ratio $M/R$, and on the degree of (in)homogeneity of the material within the envelope (characterized by a clumping factor $D$).  Envelope inflation has a strong impact on the effective temperatures of stars in the upper HR diagram. For Wolf-Rayet stars it can account for the 'radius problem', i.e., that the observed stellar temperatures of H-free WR stars are much lower than predicted by canonical stellar structure models. To reproduce the observed HRD positions of these objects, it is necessary to assume that the material within the envelope is clumped. The resulting clumping factors lie in the range of $D=1$--16, in notable agreement with typical clumping factors detected in the winds of WR stars \\citep[e.g.,][]{ham1:98}.  As the envelope inflation is expected to depend on metallicity \\citep{ish1:99,pet1:06}, it may also account for the observed $Z$-dependence of the spectral subtype distribution of WR stars, i.e., that WR stars in high-$Z$ environments show lower effective temperatures \\citep[e.g.,][]{cro1:07}.  For the luminous blue variable AG\\,Car we suggest that the observed HRD position of a $\\sim$\\,$70\\,M_\\odot$ star, with a relatively compact core, might be subject to substantial envelope inflation. During its S\\,Dor cycle, AG\\,Car is found to increase its radius by a factor of two, while its luminosity decreases by a factor of 1.5. The luminosity decrease might be explained as a result of the formation of a dense outer shell with a mass of $\\sim 2$$\\cdot$$10^{-3}\\,M_\\odot$ as predicted by our models. The energy needed to lift this material from the core to the outer radius is of the same order of magnitude as the 'missing energy' due to the decrease in luminosity. Similar to the mechanism proposed by \\citet{gro2:09,gro1:11}, the S\\,Dor variability of AG\\,Car may be due to the effect of stellar rotation, which reduces the effective stellar mass, and thus leads to an instability with $W>1$, that forces the star to expand.  We conclude that the stellar effective temperatures in the upper HR diagram are potentially strongly affected by the inflation effect. The peculiar structure of the inflated envelopes, with an almost void, radiatively dominated region, beneath a thin and dense shell could mean that many, rather very compact stars, are hidden below inflated envelopes, and thus display much lower effective temperatures to the observer. Also compact, fast rotating stars may be hidden this way. This may particularly affect massive stars just before the final collapse, and the question of whether they form a SN\\,Ibc, or a GRB \\citep[for the latter smaller core radii are inferred than the observed WR radii, cf.][]{cui1:10}.  The complex envelope structure may also affect the early X-ray afterglow of GRBs \\citep[e.g.][]{li1:07}."
    },
    "1112/1112.0322.txt": {
        "abstract": "%(Note: I will write abstract and conclusions once we have decided the very final contents of the paper)  We discuss chemical enrichments of $\\sim$4000 SDSS early-type galaxies using as tracers a large variety of element abundance ratios, namely [C/Fe], [N/Fe], [O/Fe], [Mg/Fe], [Ca/Fe] and [Ti/Fe]. We utilise the stellar population models of absorption line indices from \\citet*{TMJ10} which are based on the MILES stellar library.  We confirm previous results of increasing age, [Z/H] and [O/Fe] ratios (most often represented by [$\\alpha$/Fe] in the literature) with velocity dispersion. We further derive identical correlations with velocity dispersion for the abundance ratios [O/Fe], [Mg/Fe] and [C/Fe], implying that C/Mg and C/O are close to solar values. This sets a lower limit on the formation time-scales and star-burst components of early-type galaxies to $\\sim$0.4 Gyr, which is the lifetime of a 3M$\\odot$ star, since the full C enrichment must be reached. [N/Fe] correlates with velocity dispersion, but offset to lower values and with a steeper slope compared to the other element ratios. We do not find any environmental dependencies for the abundances of C and N, contrary to previous reports in the literature. [Fe/H] does not correlate with velocity dispersion over the entire parameter range covered, but for fixed age we find a steep trend for the [Fe/H]-$\\sigma$ relation. This trend is weaker than the analogous for total metallicity (which also shows steeper trends at fixed age) owing to the lower Fe contribution from SN Ia for more massive early-type galaxies. We find [Ca/Fe] ratios that are close to solar values over the entire velocity dispersion range covered. Tentative, due to large scatter, the results for [Ti/Fe] indicate that Ti follows the trends of Ca. This implies a significant contribution from SN Ia to the enrichment of heavy $\\alpha$-elements and puts strong constraints on supernova nucleosynthesis and models of galactic chemical evolution. %Interestingly, we find no correlations with environment for any of the element abundance ratios. This is in disagreement with previous studies that found indications for over-abundances of N and C in low density environments and poses tight constrains to the formation histories of massive elliptical galaxies. %We present a method for simultaneously deriving the stellar population parameters log(age), [Z/H] and the element abundance ratios [C/Fe], [N/Fe], [O/Fe], [Mg/Fe], [Ca/Fe] and (for the first time) [Ti/Fe] for unresolved stellar populations. We utilise the new flux-calibrated stellar population models of absorption line indices from \\citet{TMJ10}. %The new method is used for deriving full chemical enrichment histories for a sample of 3802 early-type galaxies from the Sloan Digital Sky Survey \\citep[SDSS,][]{york00}. We confirm previous results of increasing age, [Z/H], [O/Fe] (most often represented by [$\\alpha$/Fe] in the literature), [Mg/Fe], [C/Fe] and [N/Fe] with stellar %velocity dispersion. As expected we find very similar trends for [O/Fe] and [Mg/Fe], but also for [C/Fe] which means that C/Mg and C/O are close to solar values. This sets a lower limit on the formation time-scales for early-type galaxies to $\\sim$0.3 Gyr, which is the age of a 4M$\\odot$ star, since the full C enrichment must be reached. %the elements [C/Fe], [N/Fe] and [O/Fe] confirming the production of C and N in massive stars. %On the contrary we find an under-abundance of Ca compared to Mg and O. Although tentative due to large errors, a similar pattern is found for Ti as for Ca. %scale with Fe, hence there are no such correlations with galaxy mass for either [Ca/Fe] %or [Ti/Fe]. %These results indicate that the contribution from SNIa to the enrichment of $\\alpha$-elements is dependent on atomic number. %more to the enrichment of heavier alpha-elements than previously %thought. %This puts strong constraints on supernova nucleosynthesis and models of galactic chemical evolution. %Interestingly, we find no correlations with environment for any of the element abundance ratios. This is in %disagreement with previous studies that found indications for over-abundances of N and C in low density %environments and poses tight constrains to the formation histories of massive elliptical galaxies. ",
        "introduction": "\\label{intro}  %- Search for different formation histories for galaxies, bensby, feltzing used for Milky-Way %- Search for different formation mechanisms for different elements %- unresolved SP, integrated light important to reveal parameters %- introduce models %- systematic differences in models, how they affect the result. in this case differences in %  stellar libraries and in stellar evolution. %- introduce the production of elements, Ca thought to be similar to Mg but evidence of not being %  (cenarro et al  2003, thomas et al. 2003, smith et al), %  C and N produced in low mass star according to smith et al, but in massive stars according to %  sanchez-blazquez et al. %- smith et al. find a slope for Ca/Fe with Fe/H and a dependence of C/Fe with age that when taken %  into account give a similar behavior for Ca with sigma as Mg, indicating that %- C and N thought to be produced in low- and intermediate mass stars, however evidence %  for production of C in massive stars (intro sanchez-blazquez). Could differences found %  for N and C both here and for GC be evidence for different production mechanisms for %  C and N. %- Kelson find differences between N and C %- sanchez-blazquez show Evidence for environmental differences for C and N. %- environmental study in clemenz et al 2006, only influenced, not alpha, C or metallicity %- arheological downsisizing get support from xxx downsizing      %- Stellar light useful for formation histories and chemical evolution %- Search for different formation histories for galaxies, %- downsizing found - results. Support from xxx downsizing %- more ratios to reveal different histories - results for galaxies and milky-way %- for first time, full range of elments up to Fe, detailed differences %- Search for different formation mechanisms for different elements %- introduce the production of elements, Ca thought to be similar to Mg but evidence of not being %  (cenarro et al  2003, thomas et al. 2003, smith et al), %  C and N produced in low mass star according to smith et al, but in massive stars according to %  sanchez-blazquez et al. %- C and N thought to be produced in low- and intermediate mass stars, however evidence %  for production of C in massive stars (intro sanchez-blazquez). Could differences found %  for N and C both here and for GC be evidence for different production mechanisms for %  C and N. %- Ca underabundance possible evidence for SNIa contributing to alpha-elements close  to Fe-peak elements %  the inclusion of Ti strengthen this, Ti included here for first time. %- unresolved SP, integrated light important to reveal parameters %- Lick indices and models %- systematic differences in models, how they affect the result. in this case differences in %  stellar libraries and in stellar evolution.  %Stars carry useful information on the conditions present at the time of their formation. The chemical compositions of stellar atmospheres are tracers of the element abundances of the parent gas clouds forming the stars throughout the formation history of a galaxy. Some elements are also affected by dredge-up during stellar evolution \\citep*[e.g.][]{sweigart89}. %up to several billion of years ago. Stellar populations are therefore a powerful tool to extract information on chemical evolution in the Universe. Element abundances can be directly determined for individual stars of resolved stellar populations in the Milky Way or in nearby dwarf galaxies, using absorption lines measured in high resolution stellar spectra \\citep*[e.g.][]{edvardsson93,fuhrmann98,bensby04,feltzing09,bensby10}. Light averaged spectra must instead be used for determining element abundances of distant unresolved stellar populations. The absorption features of such spectra are sensitive to multiple elements due to velocity dispersion broadening. The Lick system of absorption line indices \\citep*[e.g.][]{worthey94,trager98} have been frequently used for measuring 25 prominent absorption features in galaxy spectra.  Elements are produced in stellar nucleosynthesis besides the primordial nucleosynthesis of H and He. The chemical enrichment of stellar populations depends on the star formation history, initial mass function, fraction of exploding supernovae etc. The chemical pattern of the parent gas clouds will be carried on to new stellar generations. Thus the chemical enrichment of stellar populations is also affected by mechanisms affecting the interstellar medium (ISM) such as the efficiency of stellar winds to mix newly synthesised elements with the ISM, efficiency of outflow from galactic winds to remove enriched gas, inflow of less enriched gas from gas reservoirs etc. \\citep[e.g.][]{matteucci89,matteucci94}. Chemical enrichment sets stringent constraints on galaxy formation and evolution. %, which can be done by studying element abundance ratios of stellar populations and the ISM.  %The chemical pattern of a stellar population will be affected by several different mechanisms such as the time-scale of star formation, the efficiency to mix newly synthesised elements with the interstellar medium (ISM), the initial mass function, efficiency of galactic winds to remove the ISM, fraction of Supernovaes etc. Tracing the chemical enrichment of stellar populations is therefore useful for constraining galaxy formation which can be done by studying element abundance ratios of the stellar populations and the ISM.  Studies beginning in the late 1970's have revealed non-solar abundance ratios for the stellar populations of early-type galaxies \\citep*[e.g.][]{oconnell76,peterson76,burstein84,worthey92,davies93,surma95}, indicating different chemical enrichment histories. %This information can be used for setting constraints on star formation histories and the evolution %of the Universe. %on the production of various elements. This triggered more detailed investigations showing that the ratio between $\\alpha$-elements and Fe-peak elements increases with increasing galaxy mass for early-type galaxies \\citep[e.g.][]{trager00b,thomas05,bernardi06,clemens06,thomas10}. %Containing most of the stellar mass in the Universe, early-type galaxies are %good tracers of mass assembly and have been used extensively to constrain galaxy formation. Most interestingly, the [$\\alpha$/Fe] ratio participates in the E-E dichotomy, i.e. elliptical galaxies with low [$\\alpha$/Fe] ratios have core-less central profiles, while [$\\alpha$/Fe]-enhanced galaxies with short formation time-scales have cores \\citep{Kormendy09}.  %While the $\\alpha$-elements are believed to mainly be ejected into the interstellar %medium through Type II supernovae (SNII), Fe-peak elements are also contributed by Type Ia supernovae (SNIa). %The delay of SNIa relative star formation results in weaker $\\alpha$/Fe for more extended periods of star formation. %Hence the observed $\\alpha$/Fe-mass trends have been interpreted as shorter formation time-scales for more massive early-type galaxies%(mass most often representented %by velocity dispersion) %, referred to as %astro-archeological downsizing. %a contribution of $\\alpha$-elements to the interstellar medium (ISM) %from Type II supernovea (SNII) only, while SNIa contribute Fe-peak elements as well under delayed time-scales relative to star formation. %The consequence is that a higher $\\alpha$/Fe ratio indicate shorter formation time-scales for more massive %early-type galaxies, referred to as astro-archeological downsizing.  %Other explanations to the observed $\\alpha$/Fe-$\\sigma$ relationship have been proposed, in addition to astro-archeological downsizing, %such as mass dependent initial mass function (IMF) and more effective SNII yields at higher velocity dispersion \\citep{trager00b}. %In the debate, astro-archeological downsizing gets support %from the cosmological downsizing found by \\citet{cimatti06} and \\citet{pozzetti09}, where the populations of more massive galaxies %are in place at higher redshifts. The light-averaged ages of early-type galaxies have also been found to increase with increasing velocity dispersion \\citep[e.g.][]{thomas10}, giving %further support to the downsizing scenario.  %The downsizing scenario is in contradiction to todays models of hierarchical galaxy formation that instead predict %the most massive systems to have the longest formation time-scales \\citep[e.g][]{lucia07}. Hierarchical models are unable to %reproduce the observed $\\alpha$/Fe-$\\sigma$ relationship \\citep{thomas99}, but are on the other hand %able to very well reproduce the observed large scale clustering in the universe \\citep[e.g.][]{springel05}. %The picture of galaxy evolution is clearly far from fully understood.  Individual element abundance ratios, in addition to $\\alpha$/Fe, can further disentangle the formation of different stellar populations. Since the individual elements are produced in different stellar evolutionary phases they trace varying formation histories. \\citet{edvardsson93} and \\cite{bensby10} derive as many as 12 different element abundance ratios to distinguish between the different formation histories of the stellar populations in the Milky~Way.  %Several studies have derived individual element abundances for early-type galaxies, in most cases for fairly small samples, %with varying results. \\citet{sanchez03} predict over-abundances of carbon and nitrogen %in environments of lower density for 98 early-type galaxies, while \\citet{clemens06} study a large sample of SDSS %\\citep[Sloan Digital Sky Survey,][]{york00} early-type galaxies and find environment not to influence the enhancement of carbon. %\\citet{sanchez06} study 98 early-type galaxies and find evidence for $\\alpha$/Fe and C/Fe to increase with increasing $\\sigma$ with a similar slope, while %N/Fe show a steeper slope and Mg/Fe an even steeper slope. %\\citet{kelson06} on the other hand find no significant trends for neither $\\alpha$/Fe nor C/Fe with $\\sigma$, but a strong anticorrelation of %$\\alpha$/N with $\\sigma$. %\\citet{graves08} introduced a method of simultaneously derviing individual element abundance ratios for unresolved stellar populations. %This method was adopted in \\citet{graves07} and \\citet{smith09} for a large of SDSS galaxies and 147 cluster galaxies, respectively. \\citet{smith09} %find similar trends of increasing abundancs for Mg, C and N with increasing $\\sigma$ for a sample of 147 red-sequence galaxies. %The $\\alpha$-element Ca is a puzzle as it is has been found to trace Fe instead of other $\\alpha$-elements for early-type galaxies %\\citep{cenarro03,thomas03b,smith09}.  A number of studies have derived individual element abundance ratios for early-type galaxies \\citep{sanchez03,sanchez06,clemens06,kelson06,graves07,graves08,smith09,price11}, in several cases for fairly small samples. %Results have been \\citet{sanchez03}, \\citet{sanchez06}, \\citet{clemens06}, \\citet{kelson06}, \\citet{graves07}, \\citet{graves08} and \\citet{smith09}, leaving the picture of chemical evolution in early-type galaxies incomplete. Different methods are applied in these studies, but they are all based on absorption line indices. The results are dependent on the method applied and the sample used. The aim of this work is to simultaneously derive all element abundance ratios allowed by the sensitivity of the adopted absorption line indices. The maximum amount of information is extracted from the indices to reliably derive the abundance ratios and state of the art models of stellar populations of absorption indices are utilised to obtain as accurate results as possible.  To fully interpret observed element abundance ratio trends, % derived for unresolved stellar populations stellar nucleosynthesis needs to be understood. \\citet{pipino10} find up-to-date models of chemical evolution to struggle in simultaneously reproducing observed abundance ratios for Carbon and Nitrogen from the unresolved stellar populations of early-type galaxies. The $\\alpha$-element Ca is a puzzle as it is has been found to trace Fe instead of other $\\alpha$-elements for early-type galaxies \\citep{cenarro03,thomas03b,smith09,Saglia02}. This has been interpreted as Ca being contributed by SNIa as well as SNII \\citep*{TJM10}. Thus element abundance ratios of unresolved stellar populations are also useful for constraining stellar nucleosynthesis.  % \\citet{sanchez03},  predict over-abundances of carbon and nitrogen %in environments of lower density for 98 early-type galaxies, while \\citet{clemens06} study a large sample of SDSS %\\citep[Sloan Digital Sky Survey,][]{york00} early-type galaxies and find environment not to influence the enhancement of carbon. %\\citet{sanchez06} study 98 early-type galaxies and find evidence for $\\alpha$/Fe and C/Fe to increase with increasing $\\sigma$ with a similar slope, while %N/Fe show a steeper slope and Mg/Fe an even steeper slope. %\\citet{kelson06} on the other hand find no significant trends for neither $\\alpha$/Fe nor C/Fe with $\\sigma$, but a strong anticorrelation of %$\\alpha$/N with $\\sigma$. %\\citet{graves08} introduced a method of simultaneously derving individual element abundance ratios for unresolved stellar populations. %This method was adopted in \\citet{graves07} and \\citet{smith09} for a large sample of SDSS galaxies and 147 cluster galaxies, respectively. \\citet{smith09} %find similar trends of increasing abundancs for Mg, C and N with increasing $\\sigma$ for a sample of 147 red-sequence galaxies. %The $\\alpha$-element Ca is a puzzle as it is has been found to trace Fe instead of other $\\alpha$-elements for early-type galaxies %\\citep{cenarro03,thomas03b,smith09}.  %The picture is again far from fully understood and studies of the full range of element abundances for large samples of galaxies %are needed. %More precisely, this work presents We present a technique for deriving a wide range of element abundance ratios %the full chemical enrichment histories for unresolved stellar populations, including %the abundance ratios [O/Fe] (representing [$\\alpha$/Fe]), [C/Fe], [N/Fe], [Mg/Fe], [Ca/Fe] and [Ti/Fe]. The method is based on new flux-calibrated stellar population models of absorption line indices presented in \\citet{TMJ10}. %The sensitivities to individual %element abundances of the absorption line indices are used for simultaneously derive all element abundance ratios. In \\citet{TJM10} we applied the method to galactic globular clusters and found it to be well calibrated. We analyse a sample of 3802 SDSS early-type galaxies for which we investigate element ratio scaling relations with velocity dispersion.% and environmental density.  The paper is organised as follows. The SDSS early-type galaxy sample used is presented in Section~\\ref{data} and the technique for deriving the element abundance ratios is described in Section~\\ref{spp}. The results of derived element abundance ratios for the data sample are presented in Section~\\ref{results} and further discussed and compared with the literature in Section~\\ref{disc}. Conclusions are given in Section~\\ref{conc}.  %The use of individual element abundances can not only be used for %The varying results of individual element abundances questions the  %Individual element abundance ratios are not only useful for ratios obviously can be used for SFH as well as constraining processes producing elements, IMF variations etc.  %For the first time include all elements.  %The picture of galaxy formation and evolution is clearly far from fully understood %and the star, also chemical evolution not well understood. Studying several ratios can disentangle galaxy evolution and %chemical evolution. For the first time include all elements  %To make the picture even more complex, several studies have found that not all $\\alpha$-like elements are equally enhanced %in early-type galaxies \\citep*[e.g.][]{vazdekis97,worthey98,trager98,henry99,TMB03,thomas03b,smith09}. %This implies that   %The picture of galaxy evolution is clearly not fully understood and further analyses are needed to break down %the details of galaxy formation. To do so   %All these studies indicated that the formation and evolution of galaxies are %more complex than previously believed. Stellar population models of absorption %indices are a key tool for unveiling these complexities surrounding galaxy %evolution, since absorption indices are sensitive %to the fundamental stellar population parameters age, metallicity and %element abundance ratios.  %Integrated light for unresolved stellar populations.  %It is believed that in the explosion of SNII the full range of elements up to Fe is produced, %including the $\\alpha$-elements. The explosion of SNIa is instead believed mainly produce %the F-epeak elements. The ratio  %Stellar population models of absorption line indices are a key tool for %the analysis of galaxy absorption spectra and the derivation of fundamental %galaxy properties such as formation age and element abundance %ratios (e.g. [$\\alpha$/Fe]. The absorption  %Several studies have later found that not all $\\alpha$-like elements are equally enhanced %in early-type galaxies \\citep*[e.g.][]{vazdekis97,worthey98,trager98,henry99,TMB03,thomas03b,smith09}. %Important cosmological results have also been made from studying element abundance ratio effects using stellar %population models of absorption indices, with the indication of %down-sizing %effects for early-type galaxies (early-type galaxies), where the most massive early-type galaxies seem to %have formed under the shortest time-scales \\citep{trager00b,thomas05}. %For such studies the %[$\\alpha$/Fe]-ratio is useful as it traces formation time-scales, with %the $\\alpha$-elements being produced during short time-scales when massive %stars end their life as SNII and the Fe-elements being produced under more %extensive %time-scales when binaries end up as SNIa. %The indications of down-sizing effects %were in contradiction to hierarchical models of galaxy formation %that could not reproduce the enhancement of $\\alpha$-elements in luminous %ellipticals \\citep{thomas99} %uitable for tracing mass assembly  %- element abundance ratios  %- Enivronment studies for early-type galaxies  %- Clemens et al. (2006): no gradient of Ca4227 with velocity dispersion %                         Velocity dispersion correlate with environment, higher sigma in denser areas  %- Introduce the use of absorption line indices  %- Motivate the study of early-type galaxies  %----------------------------- %    CHAPTER 2 %----------------------------- ",
        "conclusions": "\\label{conc}  We present light-averaged ages, metallicities and element abundance ratios for 3802 SDSS early-type galaxies drawn from the MOSES catalogue \\citep{schawinski07} with visual morphology classifications. %The 25 Lick indices measured on the observed galaxies are compared to the flux-calibrated models of stellar populations based on absorption line indices from \\citet*{TMJ10}. The sensitivity of the Lick indices to the variation of different element abundances allow us to determine element abundance ratios for C, N, O, Mg, Ca, Ti and~Fe. %In this work we use updated stellar population models and a new method that on top of age and metallicity computes [C/Fe], [N/Fe], [O/Fe], [Mg/Fe], [Ca/Fe] and [Ti/Fe]. This method is of an iterative nature due to the sensitivity to multiple elements for several Lick indices. Using the flux-calibrated TMJ models of absorption line indices, which are based on the MILES stellar library, we have developed a method for simultaneously deriving the element abundance ratios [C/Fe], [O/Fe] (inferred from [$\\alpha$/Fe]), [N/Fe], [Mg/Fe], [Ca/Fe] and [Ti/Fe]. The models are well calibrated with galactic globular clusters with independent measurements of stellar population parameters and element ratios \\citep[TMJ,][]{TJM10}.  %Light-average ages, total metallicities and [$\\alpha$/Fe] ratios were derived for the same data sample in \\citet{thomas10}. In this work we use updated stellar population models and a new method that on top of age and metallicity computes [C/Fe], [N/Fe], [O/Fe], [Mg/Fe], [Ca/Fe] and [Ti/Fe]. This method is of an iterative nature due to the sensitivity to multiple elements for several Lick indices. %The element abundance ratios [C/Fe], [N/Fe], [O/Fe], [Mg/Fe], [Ca/Fe] and [Ti/Fe] are discussed along with light-average age, total metallicity and iron abundance. %Despite new models and method the results of this work agree with the results of \\citet{thomas10}, i.e. strong correlations for age, metallicity, [O/Fe] and [Mg/Fe] with velocity dispersion. Hence the the conclusions from \\citet{thomas10} remain valid.  We study the relationships between the stellar population parameters and galaxy stellar velocity dispersion. In agreement with the literature stellar population age and total metallicity correlate with velocity dispersion. [Fe/H] instead does not show such a correlation over the entire parameter range covered, but for a fixed age a steep trend is found for the [Fe/H]-$\\sigma$ relation. This trend is shallower than the analogous for [Z/H] due to suppressed Fe enrichment in more massive galaxies because of time-scale dependent contribution from SN Ia  Similar trends are found for [O/Fe], [Mg/Fe] and [C/Fe], i.e. strong correlations with velocity dispersion in agreement with the literature. The first two are expected to be similar, since both O and Mg belong to the group of $\\alpha$-elements produced in massive stars through type II Supernovae. This is also in favour of the down-sizing scenario of early-type galaxies that set an upper limit on the star formation time-scales and where more massive systems experience shorter time-scales \\citep[e.g.][]{thomas10}.  The C/Mg ratios are close to solar values, which instead sets a lower limit for the formation time-scales of early-type galaxies. Stars with masses down to $\\sim$3 M$_{\\odot}$ contribute significantly to the production of C. To reach solar C/Mg ratios formation time-scales need to be long enough for such stars to eject C into the ISM. The inferred lower formation time-scale limit is then $\\sim$0.4 Gyr, which is the life-time of a 3 M$_{\\odot}$ star. %contribute significant amounts of C and have a life-times of $\\sim$0.4 Gyr, which then sets the lower formation time-scale limit. %A slightly shallower and steeper slope compared to the [Mg/Fe]-$\\sigma$ relationship is found for [O/Fe]-$\\sigma$ and [C/Fe]-$\\sigma$, respectively. This can be explained by metallicity dependent mass-loss rates in massive stars, ejecting more C before it can be converted to O.  The [N/Fe] ratios are overall lower by $\\sim$0.2 dex compared to [O/Fe] and [Mg/Fe] and the trend with velocity dispersion is very steep, i.e. more massive galaxies have significantly higher [N/Fe] ratios. The observed [N/Fe]-$\\sigma$ trends are difficult to interpret due to uncertainties in the origin of N. %Systematic uncertainties in the modelling of N abundances arise from the choice of stellar evolutionary tracks and further complicates the interpretation. These systematic uncertainties affect the zero-point of the [N/Fe]-$\\sigma$ relationship, while the slope remains steep. The zero-point and slope of this relationship can not be simultaneously matched by up-to-date models of chemical evolution \\citep{pipino10}. Either the theoretical stellar yields have to be increased by a significant factor or other prescriptions have to be incorporated into the models that affect the N yields. Such prescriptions could be: 1. N yields with a stronger dependence on metallicity, since more massive early-type galaxies are more metal-rich. 2. A dependence on galaxy mass for the ratio between the time-scale of star formation and the time-scale of primordial gas inflow, which affects the N/O, N/Mg and N/Fe ratios due to the secondary nature of N. %3. Star formation fraction in globular clusters, which show high N abundances, that is dependent on galaxy mass. %At least three different scenaThis slope is stronger than for [Mg/Fe]-$\\sigma$ and [O/Fe]-$\\sigma$, indicatingThe N abundances are under-abundant compared to O and Mg with $\\sim$0.2 dex.  We do not find any dependence on environmental density for the element ratios studied. This is in contradiction to previous studies that have reported environmental dependencies for C and N abundances. Hence difference formation scenarios for field and cluster early-type galaxies can not be inferred from the element ratios studied in this work.  The [Ca/Fe] ratios do not correlate significantly with velocity dispersion and are close to solar values over the entire velocity dispersion range covered. Although tentative, due to large errors, Ti shows a behaviour similar to Ca. %abundances are also overall under-abundant compared to O and Mg. This indicates an atomic number dependent contribution from type Ia Supernovae to the production of $\\alpha$-elements, i.e. the yields from type Ia Supernovae are higher for heavier $\\alpha$-elements. This is now universally found since similar patterns have been found in the stellar populations of the Milky Way \\citep[][and references therein]{TJM10} and puts strong constraints on supernova nucleosynthesis.  %The ? elements follow a pattern such that the elements with higher atomic number, namely Ca and Ti, are less enhanced. More specifically, [Ca/Fe] ratios are lower than [O/Fe] and [Mg/Fe] by about 0.2 dex. Ti continues this trend. We compare this result with recent determinations of element abundances in globular cluster and field stars of the Milky Way. We come to the conclusion that this pattern is now universally found. It suggests that Type Ia su- pernovae contribute significantly to the enrichment of the heavier ? elements as predicted in supernova explosion calculations and galactic chemical evolution models. This explains the presence of a Ca under-abundance (close to solar [Ca/Fe] ratios) in massive early-type galaxies and predicts similarly low [Ti/Fe] ratios in pop- ulations with short formation time-scales.  %[Ti/Fe] is for the first time derived for extragalactic unresolved stellar populations in this work."
    },
    "1112/1112.1067_arXiv.txt": {
        "abstract": "We present a resolved-star spectroscopic survey of  \\ndwarfs{} dwarf spheroidal (dSph) satellites of the Andromeda Galaxy (M31). We filter foreground contamination from Milky Way (MW) stars, noting that MW substructure is evident in this contaminant sample. We also filter M31 halo field giant stars, and identify the remainder as probable dSph members. We then use these members to determine the kinematical properties of the dSphs. For the first time, we confirm that And XVIII, XXI, and XXII show kinematics consistent with bound, dark matter-dominated galaxies. From the velocity dispersions for the full sample of dSphs we determine masses, which we combine with the size and luminosity of the galaxies to produce mass-size-luminosity scaling relations. With these scalings we determine that the M31 dSphs are fully consistent with the MW dSphs, suggesting that the well-studied MW satellite population provides a fair sample for broader conclusions. We also estimate dark matter halo masses of the satellites, and find that there is no sign that the luminosity of these galaxies depends on  their dark halo mass, a result consistent with what is seen for MW dwarfs.  Two of the M31 dSphs (And XV, XVI) have estimated maximum circular velocities smaller than $12$ km/s (to 1$\\sigma$), which likely places them within the lowest mass dark matter halos known to host stars (along with Bo\\\"{o}tes I of the MW).  Finally, we use the systemic velocities of the M31 satellites to estimate the mass of the M31 halo, obtaining a virial mass consistent with previous results. ",
        "introduction": "\\label{sec:intro}     Dwarf spheroidal (dSph) galaxies are among the most extreme objects in the pantheon of galaxies.  Their low luminosities ($10^3 < L/L_\\odot <  10^8$), lack of significant gas \\citep{grcevich09}, and low numbers compared to \\LCDM{} expectations  \\citep{kl99ms,moo99ms} are all puzzles that remain to be solved.  The difficulty in understanding the count of dSph galaxies around the Milky Way (MW) and M31 is known as the missing satellites problem, an issue that has prompted a flurry of activity modeling these galaxies \\citep[recently,][and references therein]{krav10satrev,bullock10msp,kaz11stirr,font11}.  Most models rely heavily on feedback scenarios that are tied directly to the masses of the dark matter halos that presumably host dSph galaxies. In this sense, mass determinations for dwarfs are among the most important diagnostic measurements for testing theoretical predictions at the frontier of galaxy formation.   { \\let\\thefootnote\\relax\\footnotetext{* The data presented herein were obtained at the W.M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California and the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W.M. Keck Foundation} } \\addtocounter{footnote}{-1}   While the brightest dSphs can be detected at the distance of nearby clusters \\citep[e.g.,][]{hilker03,durrell07},  faint, diffuse  dSphs can only be detected via resolved star counts, which limits detection to the Local Group (LG).  Kinematics of these galaxies thus requires resolved star spectroscopy at extragalactic distances. Thus, despite the motivations outlined above, detailed study of a large population of dSphs has been limited to the MW satellites \\citep[e.g.,][]{mateo98,sandg,walker09,simon11seg1}.  Studies of the MW dSph population have resulted in puzzles that have provided interesting challenges to \\LCDM{} and galaxy formation models. The  missing satellite problem noted above is the classic example, for which a variety of solutions have surfaced \\citep[e.g,][] {bullock00,stri07msp,tollerud08,bovill09,kop09,krav10satrev}.   However, there remain other questions such as the cause of their low gas fractions \\citep {grcevich09,nichols11}, their morphologies \\citep{kaz11stirr}, or the curiously small number with the high densities expected by \\LCDM{} \\citep{BKBK11}. These studies are based entirely on the MW dSph population, as this has been the only available data set.   Yet there is evidence that the MW has had an unusual merger history relative to similarly bright galaxies such as M31 \\citep{guhathakurta06,hammer07}. Furthermore, there are hints that the dSph populations of M31 and the MW exhibit different scaling relations \\citep{mcc05andbig,kalirai10}. Hence, expanding the sample of satellite systems is crucial to generalizing the information dSphs provide about  galaxy formation.  Fortunately, the past few years have seen much growth in the known satellite population of M31.  This is primarily due to the advent of deep surveys of the region surrounding M31 specifically designed to search for substructure like dSphs or their remnants \\citep[e.g.,][]{ibata07,pandas09nat}.   While the distance to M31 means that detection limits do not reach those of the MW's ultra-faint dSph system,  smaller angular coverage is needed to survey M31's environs. Thus, the full population of known M31 satellites (shown in Figure \\ref{fig:overview}) now includes 27 dSphs \\citep{zucker04,martin06,zucker07,maj07and14,irwin08,mc08,martin09,richardson11,Slater11And28,Bell11And29}.  These are further supplemented by 4 dwarf Elliptical (dE) satellites, which are similar in morphology to the dSph, but have somewhat higher luminosities.   The imaging surveys that detected these dSphs provide data sets that allow characterization of the photometric properties of M31 satellites, but do not provide the kinematics of these dSphs necessary for characterizing their masses and dark matter content.  These kinematical data sets are also crucial for confirming the candidates' status as self-bound galaxies.  This is illustrated by two  M31 dSphs (And IV and VIII) that were originally identified as dSph satellites, but were later shown by spectroscopic follow-up to be non-satellites \\citep{ferg00no4,ibata04and8,merrett06and8}.  Beyond this, understanding the kinematics of M31 satellite system as a whole also provides insight into M31, its dark matter halo, and its accretion history \\citep[e.g.,][]{evans00,vdm08}. While kinematics exist for some of the M31 dSphs \\citep{cote99and2,guhathakurta00,chapman05and9,Chapman07,chapman08,Letarte09,collins10,kalirai10}, the large number of recent discoveries leave many yet to be spectroscopically confirmed, and a homogeneously observed and reduced sample of a significant fraction of these satellites is necessary to properly determine characteristics of the satellites system \\emph{as a whole}.   With these ends in mind, we report here on kinematics of \\ndwarfs{} dSphs from the Spectroscopic and Photometric Landscape of Andromeda's Stellar Halo (SPLASH) Survey.  This ongoing survey of the environs of M31 aims to characterize the stellar halo of M31 and its satellite population via resolved star spectroscopy.  A companion paper on structural parameters and photometric properties is forthcoming \\comppapp.  This paper is organized as follows. In \\S \\ref{sec:obs}, we describe the observations performed for this data set. In \\S \\ref{sec:analysis}, we describe the reduction and membership analysis performed homogeneously across the spectroscopic data set, as well as our method for  estimating total velocity dispersions for each satellite in the sample.  In \\S \\ref{sec:data}, we present the results of our full spectroscopic sample, and consider each satellite in turn, describing the results of our analysis and unique aspects of each galaxy. In \\S \\ref{sec:scaling}, we consider the scaling relations of M31 dSphs and compare them to the MW.  In \\S \\ref{sec:M31mass}, we use the galaxies' $\\vsys$ to estimate the mass of M31.  Finally, we present our conclusions in \\S \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  In this paper, we have described spectroscopy of M31 dSph satellites as part of the SPLASH Survey.  We  filter out MW foreground and M31 halo field contamination to identify M31 dSph member stars, and use these data to determine $\\vsys$ and $\\slos$ for the satellites.  Based on these kinematics, we determine for each dSph the implied mass within the half light radius, and (under the assumption that these objects are dark matter-dominated) we estimate their dark halo properties.  This paper can be summarized as follows:  \\begin{enumerate}  \\item We provide a homogenous spectroscopic survey of \\ndwarfs{} M31 dSphs and provide radial velocities of resolved stars in these galaxies.  \\item We confirm that And XVIII, XXI, and XXII are kinematically cold and hence likely true satellite galaxies.  \\item We find that And XXII has a $\\vsys$ close to with M33,  suggesting it is associated with M33 rather than M31. If so, this is likely the first large mass ratio sub-subhalo (or satellite of a satellite) known.  \\item We find that the M31 dSphs obey very similar mass-size-luminosity scalings to those of MW satellites.  This suggests that the MW satellite population is not particularly unique and may be typical of a starforming $L_*$ galaxy.  \\item We use the scalings of the M31 dSphs to infer properties of their dark matter halos.  The masses of these halos show no sign of scaling with luminosity, similar to the MW dSphs \\citep{stri08commonmass,walker09}.  \\item The density of And XIV, as well as perhaps And XV and And XVI, is consistent with dark matter halos with $\\vmax<10$ \\kps{} (although consistent with higher masses at $\\sim 1\\sigma$).  If the most-likely masses for these systems are correct, these (along with the MW satellite Bo\\\"{o}tes I) are the lowest-mass dark matter halos hosting stars, with potential well depths indicative of field halos that are below the atomic hydrogen cooling limit.  \\item Using the systemic velocities of M31 dSphs as tracer particles and adopting an empirical mass estimator suggested by n-body simulations, we estimate the mass of M31 within \\satsrad{} kpc to be $\\andmass$. This corresponds to a virial mass for M31's dark matter halo of $\\andmvir$.  \\end{enumerate}   This analysis of M31 dSphs thus represent a major step forward in understanding the faintest known class of galaxies. The M31 satellites present an opportunity to understand these galaxies in a new way, as a system, along with their host halo, providing a rich  set of opportunities for examining  galaxy formation and \\LCDM.  Their similarity to MW dSphs also confirm of the Copernican principle, affirming that the MW may be a typical galaxy with typical satellites, albeit in an extraordinary universe."
    },
    "1112/1112.4752_arXiv.txt": {
        "abstract": "We consider the effect of galaxy intrinsic alignments (IAs) on dark energy constraints from weak gravitational lensing. We summarise the latest version of the linear alignment model of IAs, following the brief note of Hirata \\& Seljak (2010) and further interpretation in Laszlo et al. (2011). We show the cosmological bias on the dark energy equation of state parameters $w_0$ and $w_a$ that would occur if IAs were ignored. We find that $w_0$ and $w_a$ are both catastrophically biased, by an absolute value of just greater than unity under the Fisher matrix approximation. This contrasts with a bias several times larger for the earlier IA implementation. % Therefore there is no doubt that IAs must be taken into account for future Stage III experiments and beyond. We use a flexible grid of IA and galaxy bias parameters as used in previous work, and investigate what would happen if the universe used the latest IA model, but we assumed the earlier version. We find that despite the large difference between the two IA models, the grid flexibility is sufficient to remove cosmological bias and recover the correct dark energy equation of state. In an appendix, we compare observed shear power spectra to those from a popular previous implementation and explain the differences. ",
        "introduction": "The gravitational lensing of distant galaxy images has the potential to be a powerful cosmological tool. The lensing effect directly probes the matter distribution as a function of redshift, and thus tells us about the expansion history and growth of structure in the Universe. In this way it allows us to constrain the dark energy equation of state, or whatever is causing the apparent accelerated expansion.  Weak gravitational lensing poses a number of tough technical challenges if its true potential is to be exploited. The typical cosmic shear induced on galaxies of interest % is of order 1\\%, which is significantly smaller than the intrinsic ellipticity of the galaxies themselves. Of course we, as observers, have no access to the \\textit{unlensed} galaxy images so we must treat a population of galaxies statistically to recover the cosmological information contained in the cosmic shear signal. The correlation function of galaxy shapes as a measure of gravitational lensing was first proposed by~\\citet{kaiser92} and first observed by \\citet{Bacon:2000yp,Kaiser:2000if,Wittman:2000tc} and \\citet{van_Waerbeke:2000rm}. For reviews see \\citet{Bartelmann:1999yn,Munshi:2006fn,Refregier:2003ct}, and \\citet{hoekstra_jain_2008}.  A naive approach to cosmic shear assumes that the intrinsic distribution of galaxy ellipticities is random across the sky. If this was the case observed ellipticities on a certain patch of sky could be averaged to recover the cosmic shear, because the intrinsic ellipticity would average to zero. However it was soon pointed out that this assumption of random intrinsic ellipticity distribution is unjustified \\citep{HRH,catelankb01,crittendennpt01,croftm00}.  In fact galaxies may be expected to align with the large scale gravitational potentials in which they form so we expect physically close galaxies to be preferentially aligned with each other (known as the Intrinsic-Intrinsic (II) correlation). \\citet{hiratas04} noted an additional negative correlation between foreground galaxies shaped by a particular gravitational potential and background galaxies which are lensed by the same potential (known as the Gravitational-Intrinsic (GI) correlation) which can be of greater magnitude than the II term.  After the alignment of galaxy ellipticities from linear response to the gravitational potential % was proposed as an effect % by \\citet{catelankb01} the study was put on a firm analytic footing by the introduction of the Linear Alignment (LA) model by \\citet{hiratas04}, hereafter HS04. This approach, in which the orientation of galaxies responds linearly to the large-scale gravitational potential in which they form, has become the standard model when including the effects of Instrinsic Alignments (IA).  Correlated shears of galaxies have been observed at low redshift, where they can be attributed to intrinsic alignments, by Brown et al. 2002, and have been measured in galaxies selected to be physically close in \\citet{mandelbaumea06,hirataea07,Okumura:2008bm,brainerdea_2009} and \\citet{mandelbaumea06}. % Number density - shear correlations are easier to observe since the random galaxy ellipticity only enters the calculation once, and this has been constrained in \\citet{mandelbaumea06,hirataea07,Okumura:2008,mandelbaumea09} and \\citet{joachimiea_megazlrg}.  \\cite{hirataea07} noted that the scale dependence of the signal they measured was better matched to theory if the non-linear matter power spectrum was used in the LA instead of the linear power spectrum implied in HS04. \\cite{bridleandking} then used the non-linear matter power spectrum in their cosmological forecast calculations. This approach has been called the Non-Linear Alignment (NLA) ansatz. \\citet{schneiderb09} attempted a more motivated solution for assigning power to the IA model at small scales through the use of the halo model of galaxy clustering. By design their halo model reproduced the LA model at large scales.  The II contribution to intrinsic alignments can be taken into account in cosmological analyses by removing galaxies with small physical separation from the analysis e.g. \\citet{kings02,kings03}; \\citet{heymansh03}; \\citet{takadaw04}. This was attempted in a real analysis of COSMOS data in \\citet{schrabbackea09} by removing the autocorrelation tomographic bin. An extension of this method for GI was suggested in  \\citet{king05} and developed to a sophisticated level in \\citet{joachimi_schneider_2008,joachimi_schneider_2009}. Alternatively a model may be assumed for all the intrinsic alignment contributions, and the free parameters can be marginalised over, as demonstrated at the Fisher matrix level in \\citet{bridleandking,bernstein_2008,joachimi_bridle_2009,MGpaper1,MGpaper2} and demonstrated on real data in \\citet{kirk_bs_2010}.  The redshift evolution of the IA contributions in the LA model was found to be incorrect due to a mistake in HS04 in the conversion between the primordial potential and the matter power spectrum. This was corrected in a new version of HS04, issued as \\citet{hiratas10_posterratum}, hereafter HS10. In addition there is an ambiguity in HS04 as to which cosmological epoch is responsible for the ``imprinting'' of galaxy intrinsic alignments. Most starkly the question is, are IAs frozen in at some redshift of formation or do they evolve with the growth of structure, particularly nonlinear clustering on small scales? \\citet{blazek2011} note this as an issue in the HS04 approach and estimate its effect on the GI amplitude as of order 20\\%. \\citet{MGpaper1}, \\citet{MGpaper2} implement one physically motivated solution to this question (as well as including the redshift evolution correction of \\citealt{hiratas10_posterratum}) which we adopt in this work.  The result of this evolution in the treatment of IAs is that much useful theoretical and observational work has been conducted using an incorrect implementation of the LA model, using often unjustified treatments of the small scale seeding of galaxy IAs. In this work our aim is to present basic results for the most up to date implementation of the LA model for IAs. We present angular power spectra for components of the shear-shear ($\\epsilon\\epsilon$), position-position ($nn$) and position-shear ($n\\epsilon$) observables as well as the reduction in constraining power in measuring dark energy caused by a robust treatment of IAs and the biasing of cosmological parameters that results from ignoring IAs or treating them using an old model. Many of these results reproduce previous work in the literature which was conducted using an old implementation of the LA model. Specifically we reproduce Fig. 1 from \\citet{joachimi_bridle_2009} for the new implementation of the LA model which we hope will act as a reference for those wishing to apply it in the future. Furthermore we investigate the ability of a flexible parameterisation of intrinsic alignments to compensate for using the old LA model compared to the latest version we summarise here.  This paper is organised as follows: In section \\ref{sec:LA_now} we describe the most up to date implementation of the LA model of IAs, detailing our formalism for shear-shear ($\\epsilon\\epsilon$), position-position ($nn$) and position-shear ($n\\epsilon$) correlations. Section \\ref{sec:bias} presents the bias on cosmological parameter estimation caused by mistreatment of IAs for the latest model, the old model and an intermediate model; we then conclude in section \\ref{sec:conclusions}. In an appendix we summarise the history of the LA model and present the major differences between the latest implementation and the previous widely used version. ",
        "conclusions": "\\label{sec:conclusions}  Galaxy IAs are the most important physical systematic in the study of cosmic shear. As increasing volumes of data become available from WGL surveys more interest is being paid to their correct treatment. Much of this work has been dominated by the LA model, originally introduced in HS04 (and the NLA extension which we refer to as HS04NL). This original implementation was subsequently corrected in \\citet{hiratas10_posterratum} and detailed attention paid to the treatment of non-linear clustering in \\citet{MGpaper1,MGpaper2}.  However, much of the existing literature has been produced using the uncorrected HS04NL model. In this paper we have provided a brief explanation of the evolution of, and context surrounding, the LA model for IAs, highlighting the most important differences between HS04NL and the latest implementation.  We have calculated the angular power spectra of the cosmic shear observables, correcting the implicit mistake of \\citet{joachimi_bridle_2009}, which was based on HS04, which we hope will act as a new reference for those interested in applying IAs to their cosmic shear analysis.  The main motivation for the study of IAs is the measurement of unbiased constraints of key cosmological parameters. We show that the new LA implementation significantly reduces the impact of IAs, and hence the bias, but that the effect is still very significant, producing a bias at the tens of $\\sigma$ level when we assume perfect knowledge of IAs.  If we did know the IA signal perfectly then we could produce unbiased measurements by simply subtracting the IA signal from our measured cosmic shear. In practice it is useful to parameterise our ignorance of the true IA signal through a set of nuisance parameters which are marginalised over to produce weaker but hopefully unbiased cosmological estimates. We show that a robust grid of 130 nuisance parameters for IAs and magnification uncertainties, % allowed to vary in scale and redshift, effectively removes the bias due to assuming an incorrect IA model.  The same effect holds when our observables are extended to include shear-shear, position-position and shear-position correlations. The extra observables increase constraining power so that we are able to produce unbiased constraints on $\\epsilon\\epsilon+n\\epsilon+nn$ which recover 40\\% of the constraining power of the highly biased constraint from $\\epsilon\\epsilon$ alone when the wrong IA model is assumed.   The LA model has allowed a firm foothold on the study of IAs to develop over the last decade. We have detailed an updated implementation of the most used IA model and its physical motivation, showing a reduction in biasing of cosmological constraints. In addition we reiterate that a robust nuisance parameter model can control the biasing due to IAs for either the old or new implementations. Together, these results should give us confidence that the IA effect is under control as we begin to analyse the first data from large cosmic shear surveys.  However our somewhat pessimistic approach carries the cost of reduced constraining power. Future work, focused on accurate simulation and measurement of IAs is sure to provide more detail on the physical mechanisms responsible for the initial intrinsic ellipticity distribution and its evolution. This better knowledge of IAs will improve our ability to model them and reduce our dependence on brute-force marginalisation over nuisance parameters. As such the marginalised constraints which we present in Fig.~\\ref{fig:bias_oldnew} may be a worst-case scenario. With better knowledge of IAs, better constraints on cosmology will be possible."
    },
    "1112/1112.3943_arXiv.txt": {
        "abstract": "We report the petrology, O isotopic composition, and Al-Mg isotope systematics of a chondrule fragment from the Jupiter-family comet Wild 2, returned to Earth by NASA's Stardust mission. This object shows characteristics of a type II chondrule that formed from an evolved oxygen isotopic reservoir. No evidence for extinct $^{26}$Al was found, with ($^{26}$Al/ $^{27}$Al)$_0$ $<$ 3.0$\\times$10$^{-6}$. Assuming homogenous distribution of $^{26}$Al in the solar nebula, this particle crystallized at least 3 Myr after the earliest solar system objects---relatively late compared to most chondrules in meteorites. We interpret the presence of this object in a Kuiper Belt body as evidence of late, large-scale transport of small objects between the inner and outer solar nebula. Our observations constrain the formation of Jupiter (a barrier to outward transport if it formed further from the Sun than this cometary chondrule) to be more than 3 Myr after calcium-aluminum-rich inclusions. ",
        "introduction": "High-temperature objects---fragments of chondrules and calcium-aluminum-rich inclusions (CAIs)---have been discovered in the collection of samples returned by NASA's Stardust mission from comet Wild 2 (e.g. \\citet{nak08}, \\citet{mat10}). The diversity of the Wild 2 samples shows that this comet does not consist entirely of unaltered nebular material rich in presolar grains and interstellar amorphous silicates \\citep{bro06}, nor is comet Wild 2 close in composition to a known class of meteoritic material or chondritic-porous interplanetary dust particles (e.g. \\citet{ogl10}). The presence of chondrules and CAIs in a Jupiter-family comet requires a transport mechanism between the hot inner solar system, where these objects likely crystallized, and the scattered disk, where they were eventually accreted into the comet \\citep{dun04}.  Large-scale transport in the disk must have occurred before the formation of Jupiter because the growing embryo of Jupiter strongly suppressed outward transport across its orbit \\citep{cie09}. A nascent Jupiter efficiently accreted material that diffused into its orbital gap, effectively creating a barrier to outward transport of high-temperature objects if its orbit was farther from the Sun than where these objects formed.  Measurements of the decay products of short-lived radionuclides in high-temperature objects in the Stardust collection will yield an earliest time when this large-scale transport occurred. It is then possible to constrain the timing of the formation of Jupiter, which because of its large mass controlled much of the dynamics of planetesimals forming in the young solar system (e.g \\citet{wal11}). ",
        "conclusions": "Since CAIs are believed to have formed when the Sun was a class 0 or class I protostar, Iris formed late in the evolution of the solar nebula, at a time when $\\sim$95\\% of the scattered disk is thought to have cleared \\citep{her01}. The residual disk material may have coated Jupiter-family comets with a ``late veneer'' enriched in inner solar system material \\citep{ogl10}, and this may explain significant differences between the Stardust sample, which sampled Wild 2 coma material ejected by jets that entrained near-surface dust \\citep{bel10}, and chondritic porous  interplanetary dust particles, which probably sample the bulk of Kuiper Belt comets \\citep{nes10}.  These analyses are for a single object in the Stardust collection. However, Iris is not unique or even particularly unusual in the suite of particles returned from comet Wild 2. \\citet{nak08} identified four chondrule-like objects in the Stardust samples which are more $^{16}$O-enriched than Iris (Figure \\ref{oiso}) and olivines that are less Fe-rich (Fo$_{79\\text{--}80}$, Fo$_{95}$, Fo$_{91}$).  At least one chondrule fragment was also identified during the Stardust preliminary examination \\citep{zol06}. All of these objects show igneous textures similar to Iris, though Iris appears to have formed in a more oxidizing environment. Recent measurements by \\citet{jos11} show two fragments of Fe-rich olivine (Fo$_{62\\text{--}67}$, Fo$_{58\\text{--}61}$), that also have O isotopic composition close to Iris. Chondrule-like objects from Stardust show a broad range of isotopic and mineralogical compositions; a subset of these objects could be genetically related to Iris.  The formation time of Iris is most consistent with the late-forming chondrules in CR chondrites. Chondrules from CB chondrites also formed relatively late (as determined by $^{207}$Pb--$^{206}$Pb measurements), likely by a giant impact in the early solar system \\citep{kro05}. The parent bodies of both CR and CB chondrites were scarcely heated \\citep{sco06}, similar to Wild 2 \\citep{bro06}. However, type II chondrules like Iris (and similar Stardust fragments) are very rare in CR and CB meteorites: $<$1\\% of all chondrules \\citep{wei93,kro02}. The Iris mesostasis is Fe-poor and Na-rich compared to mesostasis in type II chondrules in unequilibrated CR chondrules, and the Iris olivine is Ca- and Al-rich relative to olivine in CR type II chondrules \\citep{ber11}. Although Iris is similar to chondrules from CR and CB meteorites in that it formed relatively late, Iris and similar Stardust fragments are unlikely to have originated from the CR- or CB-chondrite-forming region.  The Fe-rich olivines in Iris must have formed in an environment with oxygen fugacity higher than the typical redox conditions of the steady-state solar nebula \\citep{kro00}. Shocks in the outer solar nebula, beyond the water-snow line at $\\sim$5 AU \\citep{cyr98}, were rich in water vapor \\citep{cie03} which could have formed a chondrule like Iris near the current orbit of Jupiter. However, the predicted composition of chondrule olivine in such a shock is strongly peaked between Fo$_{76}$ and Fo$_{89}$, Fo$_{70}$ olivine and lower is predicted to be $\\sim$3\\% \\citep{fed08} of all chondrule olivine generated in the shock. Therefore we conclude that Iris was probably not created in an outer-nebula shock, but likely formed in the inner solar nebula, from material with relatively high Fe/Mg or in a region of high oxygen fugacity \\citep{jon90}, more than 3 Myr after CAIs.  Recent measurements of oxygen isotope variations in the rim of a CAI from the CV3 chondrite Allende \\citep{sim11} indicate that these early objects experienced circulation in the solar nebula. A Type C CAI found in the Stardust samples was constrained to have crystallized at least 1.7 Myr after the onset of CAI formation (assuming a homogeneous nebular reservoir of canonical ($^{26}$Al/$^{27}$Al)$_0$), though it likely experienced a complex history \\citep{mat10}. Additionally, high-temperature components (CAIs and chondrules) do not appear to be scarce in the Stardust samples. These observations provide strong evidence of a dynamic early solar system transporting material between the inner and outer nebula.  The formation of Iris in the inner nebula requires it to be transported to the scattered disk at $\\sim$35 AU \\citep{dun04,tir09} where it was incorporated into comet Wild 2.  This transport could have occurred outside the plane of the disk (e.g. \\citet{cie07}), by diffusion (e.g. \\citet{cuz03}), aerodynamic lofting and radial drift (e.g.  \\citet{cie08}), or outward advective flows (e.g. \\citet{hug11}). Jupiter's growing embryo would open a gap in the disk \\citep{bat03} which would prohibit outward transport \\citep{cie09} if Iris was created in an inner-nebula shock (unless Jupiter had migrated inside of 2 AU and was closer to the Sun than Iris when Iris formed \\citep{wal11}). Iris likely formed in an event prior to Jupiter's formation, such as in an inner-nebula spiral shock that existed before Jupiter opened a gap in the disk \\citep{bol05}. Therefore, our measurements set a constraint on the formation time of Jupiter of at least 3 Myr after CAI formation making it unlikely that Jupiter formed early (e.g., by disk instability \\citep{bos01}).  Our constraint on the formation time of Jupiter is consistent with arguments based on the requirement to accrete asteroids (with constituent radiometric-dated chondrules and CAIs) before Jupiter grows large enough to inhibit accretion ($>$3--5 Myr, \\citet{sco06}). Jupiter was estimated to form $\\sim$3.3 Myr after the onset of planetesimal fragmentation in the main belt by \\citet{bot05}, also consistent with the outward transport of Iris. The formation of chondrules in planetesimal bow shocks caused by Jovian resonances \\citep{wei98} requires Jupiter to form $\\sim$1 Myr after CAIs, a scenario which is disallowed by our measurements."
    },
    "1112/1112.4614_arXiv.txt": {
        "abstract": "The Fermi Large Area Telescope (Fermi LAT) provides long term systematic monitoring observations of the gamma-ray emission from blazars. The variability properties and the correlation with other wavelength bands are important clues for the evaluation of blazar models. We present results from timing and multiwavelength correlation analysis and discuss differences between blazar classes. ",
        "introduction": "More than 60 AGNs, almost all blazars, were detected by EGRET on Compton GRO, which established these sources as a powerful class of gamma-ray emitters [1]. In its first year of operation Fermi LAT has increased the number of known gamma-ray blazars by a factor of 10. Even more importantly the instrument is mapping the full sky every three hours, which allow regular monitoring of these sources on time scales from hours to years. These monitoring observations now form an important part of many ongoing efforts to study the variability and multiwavelength properties of blazars. Such observations contains information about the relative origin of different spectral components and about dynamical and radiation processes in blazar jets. ",
        "conclusions": ""
    },
    "1112/1112.1031_arXiv.txt": {
        "abstract": "The $\\lya$ emission has been observed from galaxies over a redshift span $z \\sim 0 - 8.6$. However, the evolution of high-redshift $\\lya$ emitters (LAEs), and the link between these populations and local galaxies, remain poorly understood. Here, we investigate the $\\lya$ properties of progenitors of a local $L^{*}$ galaxy by combining cosmological hydrodynamic simulations with three-dimensional radiative transfer calculations using the new $\\art$ code. We find that the main progenitor (the most massive one) of a Milky Way-like galaxy has a number of $\\lya$ properties close to those of observed LAEs at $z \\sim 2 - 6$, but most of the fainter ones appear to fall below the detection limits of current surveys. The $\\lya$ photon escape fraction depends sensitively on a number of physical properties of the galaxy, such as mass, star formation rate, and metallicity, as well as galaxy morphology and orientation.  Moreover, we find that high-redshift LAEs show blue-shifted $\\lya$ line profiles characteristic of gas inflow, and that the $\\lya$ emission by excitation cooling increases with redshift, and becomes dominant at $z \\gtrsim 6$. Our results suggest that some observed LAEs at $z \\sim 2-6$ with luminosity of $\\La \\sim 10^{42-43}~\\ergs$ may be similar to the main progenitor of the Milky Way at high redshift, and that they may evolve into present-day $L^{*}$ galaxies. ",
        "introduction": "The $\\lya$ emission from young galaxies can be a powerful probe of the early universe \\citep{Partridge67, Charlot93}. Recent narrow-band deep imaging surveys using large-aperture telescopes have detected a large number of $\\lya$ emitting galaxies, or  $\\lya$ emitters (LAEs), at redshifts $z \\gtrsim 3$ \\citep[e.g.,][]{Hu96, Cowie98, Steidel00, Malhotra04, Taniguchi05, Kashikawa06, Shimasaku06, Iye06, Hu2006, Gronwall07, Ouchi08, Hu2010, Ouchi10, Lehnert10}. By combining $\\lya$ emission with broad-band continuum, multi-wavelength observations are beginning to address the physical properties of these high-redshift LAEs \\citep[e.g.,][]{Gawiser06, Gronwall07, Lai07, Nilsson07, Pirzkal07, Lai08, Ouchi08, Pentericci09, Ono10A, Ono10B, Hayes10, Finkelstein11, Nilsson11, Acquaviva11}. It has been suggested that these objects are mostly compact, young galaxies with low metallicity. In addition, \\citet{Ouchi08} studied the evolution of equivalent widths (EWs) and the characteristic $L^{*}_{\\lya}$ with redshift from $z \\sim 3$ to $z\\sim 6$, and found that the mean EW increased with redshift, while the $L^{*}_{\\lya}$ did not change significantly. More recently, \\citet{Ciardullo11} studied the evolution of luminosity function (LF) and EW from $z=2.1$ to $3.1$, and found that $L_{*}$ increases from $10^{42.3}~\\rm \\ergs$ at $z=2.1$ to $10^{43}~\\rm \\ergs$ at $z=3.1$. \\citet{Blanc11} studied the $\\lya$ properties of LAEs in the redshift range $z=1.9-3.8$ from the Hobby Eberly Telescope Dark Energy Experiment (HETDEX) Pilot Survey, and showed that the median $\\lya$ escape fraction ($\\fesc$) was $\\sim 29 \\%$, and it does not evolve significantly with redshift. On the other hand, \\citet{Hayes11} suggested that $\\fesc$ monotonically increases between redshift 0 and 6, which implies that high-z galaxies tend to be LAEs.  While high-redshift LAEs have been studied with large samples in the redshift range of $z \\sim 2.2-6.6$, there is only a limited number of observations on LAEs at $z \\lesssim 1$. Some local star-forming galaxies have been studied by various wavelengths and show a complex structure of $\\lya$ and UV continuum \\citep{Hayes07, Ostlin09}. \\citet{Atek09} showed that the $\\fesc$ of local LAEs have a large dispersion, ranging from $\\sim 3$ to $100~\\%$. In addition, \\citet{Deharveng08} studied a sample of 96 local LAEs at $z=0.2-0.35$ from UV space telescope {\\it GALEX}, and found that these LAEs have similar EW distribution as those at $z=3.1$. Recently, \\citet{Cowie10} have studied $z \\sim 0.3$ LAEs with a larger sample, and showed that these LAEs are more compact, and have lower metallicity than UV-continuum selected galaxies at the same redshift. In addition, \\citet{Finkelstein09} suggested, from fitting of spectral energy distributions (SEDs), that low-z LAEs are significantly more massive and older galaxies than their high-z counterparts.  One of the important issues in galaxy evolution is how high-redshift LAEs evolve into galaxies in the local universe. \\citet{Gawiser07} suggested, from clustering analysis, that most $z=3.1$ LAEs evolve to present-day galaxies of $\\lesssim 2.5~ L^{*}$, unlike other populations which typically evolve into more massive galaxies. Moreover, \\citet{Guaita10} indicated that LAEs at $z=2.1$ were building blocks of present-day $L^{*}$ galaxies such as the Milky Way (MW).  However, the link between high-redshift LAEs and local galaxies, and the probability of these LAEs evolving into present-day $L^{*}$ galaxies are not well constrained from observations. In order to address these questions, one may use the Milky Way as a local laboratory. Moreover, since $\\lya$ emission has been detected from the most distant galaxies, understanding of the $\\lya$ properties of the Milky Way progenitors will provide an important clue to the formation of early galaxies. To date, there are only a limited number of theoretical studies on this important issue \\citep[e.g.,][]{Salvadori10, Dayal11b}. Both \\citet{Salvadori10} and \\citet{Dayal11b} focused on MW progenitors at $z\\sim 6$ constructed from semi-analytical merger trees and a cosmological smoothed particle hydrodynamics (SPH) simulation, respectively. They both used the same analytical prescription of $\\lya$ emission in which the intrinsic $\\lya$ luminosity scales linearly with the star formation rate \\citep{Dayal08}. However, because the $\\lya$ properties depend sensitively on a number of factors, such as the scattering and propagation of the photons in the inhomogeneous medium, the dust content of the gas, the ionization structure, the UV continuum, and the photon escape fraction. Such a complicated process can only be probed by comprehensive $\\lya$ radiative transfer calculations combined with realistic simulation of galaxy formation. As we will show in this work, our detailed $\\lya$ modeling on a high-resolution cosmological simulation produce a number of $\\lya$ properties such as the luminosity functions at different redshifts in good agreement with observations. Moreover, in order to investigate the evolution of LAEs, we need to study the progenitors of the MW at different redshifts systematically, not just at a specific time.  In this paper, we investigate the $\\lya$ properties of MW progenitors over a wide redshift range of $z \\sim 0 - 10$, by combing cosmological SPH simulation of a MW-like galaxy from Zhu et al. (in preparation) with 3D RT calculations using the newly improved $\\art$ code by \\cite{Yajima11A}. Our hydrodynamic simulation includes important physics of both dark and baryonic matter, and has high resolutions to track the formation history of the MW. Our RT calculations include both $\\lya$ resonant scattering and continuum emission, and are done on an adaptive-mesh refinement grid, which covers a large dynamical range and resolves the small-scale structures in high-density region. Interstellar dust is also taken into account to accurately estimate the $\\fesc$ of $\\lya$ photons and UV continuum, and the EWs.  The paper is organized as follows. We describe our cosmological simulation in \\S2, and the RT calculations in \\S3. In \\S4, we present results of the $\\lya$ properties of MW progenitors from redshift 10 to 0, which include the $\\lya$ surface brightness, $\\lya$ luminosity, $\\fesc$, EW and line profile. In \\S5, we discuss the dependence of $\\fesc$ on physical properties, LAE fraction in our galaxy sample, $\\lya$ escaping angle and the contribution from excitation cooling to $\\lya$ emissivity, and we summarize in \\S6. ",
        "conclusions": "\\subsection{Dependence of $\\lya$ Properties on Galaxy Properties} \\label{sec:fesc}  \\begin{figure*} \\begin{center} \\includegraphics[scale=0.9]{fesc_all.eps} \\caption{Dependence of $\\lya$ properties on various physical properties of a galaxy. {Upper~:} The relation between $\\lya$ escape fraction and halo mass, SFR and metallicity. The different color represents the different redshift, as indicated in the color bar. The dash lines are least-absolute-deviations fittings. {Lower~:} The relation between $\\lya$ luminosity and halo mass, SFR and metallicity. } \\label{fig:fesc_all} \\end{center} \\end{figure*}  \\begin{figure} \\begin{center} \\includegraphics[scale=0.4]{frac.eps} \\caption{ Number fraction of LAEs in our galaxy sample. The green, blue and red filled circles represent the fraction at $\\lya$ luminosity threshold, $10^{40}, ~10^{41}$, and $10^{42}~\\ergs$, respectively. } \\label{fig:frac} \\end{center} \\end{figure}  As shown in previous Sections, the $\\lya$ properties vary significantly in different galaxies. Here we explore the dependence on a number of physical properties of a galaxy. Figure~\\ref{fig:fesc_all} shows the dependence of escape fraction $\\fesc$ (top panels) and $\\lya$ luminosity $\\La$ (bottom panels) on the galaxy mass, SFR, and metallicity Z. We apply a least-absolute-deviations fitting to the data using a power-law function, ${\\rm log} Y = \\alpha {\\rm log}X + \\beta$.  The mass dependence of $\\fesc$ has a large dispersion, but from our fitting, $\\alpha \\sim - 0.02$ and $\\beta \\sim -0.17$, which suggests that $\\fesc$ roughly decreases with the total mass, consistent with the results of \\citet{Laursen09b}. At $M \\sim 10^{10-11}~\\Msun$, the $\\fesc$ is mostly constant at $\\sim 10 - 30~\\%$. In contrast, $\\fesc$ is more tightly correlated with the SFR, with $\\alpha \\sim -0.08$ and $\\beta \\sim -0.41$. At high SFR, dust  can be enriched quickly by type II supernovae, and can effectively absorb the $\\lya$ photons. In addition, galaxies with high SFR have more hydrogen gas. The gas decreases the mean free path of $\\lya$ photons, resulting in the increase of the dust optical depth which reduces the escape fraction. In addition, the $\\fesc$ decreases with metallicity, $\\alpha \\sim -0.35$ and $\\beta \\sim -0.65$. Since the dust content linearly increases with metallicity in our model, the $\\lya$ photons can be absorbed effectively by gas with high metallicity. This trend is consistent with observational indication by \\cite{Atek09} and \\cite{Hayes10}.  On the other hand, the luminosity $\\La$ has different relationships with these properties from the $\\fesc$. The $\\La$ is also roughly correlated with the mass, $\\La \\simeq 10^{37.7}\\times\\Mtot^{0.38}$, with a large dispersion. Only massive galaxies with $\\Mtot \\gtrsim 10^{11}~\\Msun$ have the $\\lya$ luminosity of $\\La \\ge 10^{42}~\\ergs$. This is consistent with suggestions from clustering analysis of observed LAEs at $z=2-3$ \\citep[e.g.,][]{Gawiser07, Guaita10}. In our model, the massive galaxies at $z=2-3$ evolve into $L^{*}$ galaxies at $z=0$. Hence, our results support the suggestion by \\cite{Gawiser07} that the observed LAEs with $\\La \\gtrsim 10^{42}~\\ergs$ at $z=2-3$ are likely progenitors of local $L^{*}$ galaxies.  The $\\La$ has the tightest correlation with SFR among the properties investigated here: $\\La \\simeq 10^{41.7}\\times \\rm SFR^{0.53}$ In the literature, a simple linear relation is commonly used, with $\\La \\;(\\ergs) = 1.1 \\times 10^{42}\\times \\rm{SFR} \\;(\\Msunyr)$, assuming that $\\La / L_{\\rm H\\alpha} = 8.7$ (case B). However, our result suggests that the relation between $\\La$ and SFR becomes somewhat shallower due to the dependence of $\\fesc$ on SFR. Finally, the emergent $\\La$ does not show a strong dependence on metallicity, $\\La \\simeq 10^{41.1}\\times (Z/Z_{\\odot})^{-0.27}$. This is due to the fact that, although the intrinsic $\\La$ increases with halo mass (so does SFR and metallicity), the $\\fesc$ decreases with metallicity, so $\\La$  of higher-metallicity galaxies is suppressed by dust absorption.  We should point out that the large scatter in the correlations in Figure~\\ref{fig:fesc_all} may be due to the small volume of our simulation and the small number of our galaxy sample. In addition, as we discuss in Section~\\ref{sec:limit}, a number of limitations of our model, such as the simplified ISM model and insufficient resolutions, may contribute to uncertainty in these relations. Moreover, the luminosity scaling relations may change under some specific detection limits. We will study these relations in detail with improved model and simulations in future work.   \\subsection{Redshift Dependence of LAE Fraction}  The number fraction of LAEs ($\\fa$) in our sample is shown in figure~\\ref{fig:frac}. The detection limit of $\\lya$ varies in different surveys. At high redshifts ($z \\gtrsim 3$), the LAE detection in most of observations has been confined to $\\La \\gtrsim 10^{42}~\\ergs$. Here, we derive the $\\fa$ with three $\\La$ thresholds, $\\La \\gtrsim 10^{40}, ~10^{41}, ~10^{42}~\\ergs$ with EW of $\\gtrsim 20~\\A$. The $\\fa$ with $\\La \\gtrsim 10^{40}, 10^{41}~\\ergs$ rapidly increases from $z=0$ to $\\sim 5$, and then remains nearly constant with higher values $\\gtrsim 0.8$ at $z \\gtrsim 5$. The trend is roughly similar to the SFR history (Figure~\\ref{fig:sfr}). Since the $\\La$ is tightly correlated with the SFR (Figure~\\ref{fig:fesc_all}), the number of galaxies with $\\La \\gtrsim 10^{40}, ~10^{41}$ increases at $z \\sim 0 - 4$. On the other hand, although the SFR decreases at $z \\gtrsim 4$, the $\\fesc$ increases due to low metallicity. Hence, the $\\fa$ does not decrease at $z \\gtrsim 4$.  Meanwhile, the $\\fa$ with $\\La \\gtrsim 10^{42}~\\ergs$ is nearly constant, and shows $\\sim 2 - 10~\\%$. Since the SFR tightly correlates with $\\La$, and it roughly increases with the galaxy mass, some massive galaxies can be LAEs with $\\La \\gtrsim 10^{42}~\\ergs$. In addition, the $\\fa$ of LAEs having intrinsic $\\La \\gtrsim 10^{42}~\\ergs$ change with cosmic star formation history. However, the $\\fesc$ decreases around the phase of SFR peak, and therefore suppresses the $\\fa$.  On the other hand, at lower redshift, the observations indicate that number density of LAEs decreases by some factors \\citep[e.g.,][]{Cowie10}. The discrepancy may come from the difference in density field and the small box in our simulation. Our initial condition is a somewhat special one which is focused on a MW-size galaxy, and the zoom-in simulation region is $\\sim 5^3~ h^{-3} \\Mpc^{3}$. Therefore, our simulation cannot reproduce the global statistics in observations. Moreover, the LAEs fraction having $EW > 25~\\A$ in this work shows $\\sim 44\\; \\%$ at z=4 and $\\sim 100\\; \\%$ at $z = 6$, which is somewhat higher than the LAE fraction in LBG sample \\citep{Stark10, Stark11, Pentericci11, Schenker12, Ono12}. However, in observation, the LAE fraction increases with decreasing UV brightness. Most of our model galaxies at $z \\gtrsim 3$ are fainter than the detection threshold in the LBG observation. Since the number of galaxies brighter than the threshold of LBG observation is quite small (less than ten), we need a larger sample covering a wide mass range to verify the model of LAEs. In addition, although some LAEs have been observed with UV continuum, and hence categorized as LBGs, it is inadequate to study LAEs from LBG-only sample, because a large fraction of LAEs may have UV continuum under the detection limit of current observations. We will address the general properties such as luminosity function, EW distribution and clustering systematically by using a set of uniform simulations with mean density field in larger volumes in future work.  \\subsection{The Viewing-angle Scatter of Escaping $\\lya$ Photons} \\label{sec:angle}  \\begin{figure} \\begin{center} \\includegraphics[scale=0.45]{fescmap.eps} \\caption{The ``viewing-angle scatter'' of escaping $\\lya$ photons depends strongly on galaxy morphology and orientation. Shown here is the $\\lya$ escape probability in a irregular galaxy progenitor at z=3.1 (top panel) and the spiral MW galaxy at z=0 (bottom panel). The color bar indicates the probability per unit  $\\rm d\\phi, dcos\\theta$. There is no clear direction in the irregular galaxy, but in the disk galaxy, the $\\lya$ photons escape in a preferred direction normal to the disk. } \\label{fig:fescmap} \\end{center} \\end{figure}  Despite their high metallicity, a fraction of galaxies at low redshift $z \\lesssim 1$ show high escape fraction $\\fesc$ of $\\lya$ photons (Figure~\\ref{fig:fesc_all}). We find that the escaping angle of the $\\lya$ photons depends strongly on the galaxy morphology and orientation, a phenomenon we dub as the ``viewing-angle scatter''. % Disky objects seen edge on can be hundred times fainter than  the same objects seen face on. In a galaxy which has a gas disk, the $\\lya$ photons escape in a preferred direction normal to the disk, but there is no clear escaping direction in compact or irregular galaxies without a gas disk. We demonstrate this effect in Figure~\\ref{fig:fescmap}. We first estimate the normal direction to the gas disk according to the total angular momentum of the gas, and set $\\theta = 0^{\\circ}$ along this direction. In a galaxy with irregular morphology such as the main progenitor at $z = 3.1$, there is no clear preferred escaping angle, as illustrated in the top panel of Figure~\\ref{fig:fescmap}. However, in a spiral galaxy with rotationally supported gas disk such as the MW galaxy at z=0 in our simulation, the escaping angle is strongly confined to $\\rm{cos\\theta \\simeq \\pm 1}$, corresponding to $\\theta \\simeq 0^{\\circ}$ or $180^{\\circ}$, as shown in the bottom panel of Figure~\\ref{fig:fescmap}. This is due to the fact that the $\\lya$ photons have the minimum optical depth along the normal direction to the gas disk. More than $60~\\%$ $\\lya$ photons escapes to the direction of $\\rm |cos~\\theta| \\lesssim 0.5$. Generally the $\\lya$ flux from our model galaxies can scatter around the mean value typically by a factor of ten just from different orientations.  As illustrated in Figure~\\ref{fig:img}, most galaxies in our simulation have highly irregular shapes at high redshift due to accretion and gravitational interaction. At z=0, a number of them evolve into spiral disks. The ``viewing-angle scatter'' explains why we see high $\\lya$ escape fractions in a number of low-z galaxies, and the fact that $\\lya$ is detected in a large number of face-on spiral galaxies in the nearby universe \\citep[e.g.,][]{Cowie10}.   \\subsection{Contribution of Excitation $\\lya$ Cooling} \\label{sec:exc}  \\begin{figure} \\begin{center} \\includegraphics[scale=0.4]{exc.eps} \\caption{ The fraction of excitation cooling $\\lya$ to the total intrinsic $\\lya$ luminosity as a function of redshift. The red filled circles indicate the values of the main progenitor, while the blue filled circles represent the median value of the galaxy sample indicated with black open circles. The error bars show the quartiles. } \\label{fig:exc} \\end{center} \\end{figure}  There are two major mechanisms to generate $\\lya$ emission, the recombination of ionizing photons and the collisional excitation of hydrogen gas. However, the relative contribution between the two mechanisms is not well understood. From our calculations, we find that the contributing fraction of excitation $\\lya$ emission to the total intrinsic $\\lya$ luminosity increases with redshift, as shown in figure~\\ref{fig:exc}.  In our cosmological simulation, galaxy evolution is accompanied by cold, filamentary gas streams with temperature $T\\sim 10^{4-5}~\\rm K$, which penetrate deep inside dark matter halos (Zhu et al. in preparation, Yajima et al. in preparation), a phenomenon also reported by other groups \\citep{Katz03, Keres05, Keres09, Birnboim03, Dekel06, Ocvirk08, Brooks09, Dekel09}. Such cold gas can efficiently produce the excitation $\\lya$ cooling photons \\citep{Dijkstra09, Faucher09, Goerdt10}. At higher redshift, galaxies form through more efficient gas accretion and more frequent merging event. As a result, the contributing fraction increases with redshift, and becomes dominant at $z \\gtrsim 6$. This excitation mechanism does not depend on the stellar radiation, and can therefore produce high $\\lya$ EWs. We find that the EWs of LAEs increases significantly at $z \\gtrsim 6$, reaching $\\gtrsim 10^3~\\A$ at $z \\sim 10$. This is larger than the upper-limit of EW, $240~\\A$, which considers only stellar sources assuming a Salpeter IMF with solar abundance of metallicity \\citep{Charlot93}. Although the upper-limit increases with decreasing metallicity, it was suggested that top-heavy IMF like Pop III stars are needed for making $\\rm EW > 400~\\A$ \\citep{Schaerer03, Raiter10}. However, even though Salpeter-IMF is used in this work and the stellar metallicity is mostly $Z / Z_{\\rm \\odot} \\gtrsim 10^{-3}$, we find that the EW can be higher than the upper-limit by the efficient excitation $\\lya$ emission. On the other hand, the $\\lya$ line is strongly damped by IGM correction at $z \\gtrsim 6$ \\citep{Haiman02, Laursen11, Dayal11b}, which can result in a lower EW. The suppression by IGM highly depends on the inhomogeneous ionization structure around LAEs \\citep[e.g.,][]{McQuinn07, Mesinger08, Iliev08}. We will address the detectability of high-redshift LAEs and EW after IGM correction by running large-scale $\\lya$ RT in IGM in future work.  \\subsection{$\\lya$ Luminosity Functions} \\label{sec:LF}  The simulation box in this work is too small to study global statistics directly. As a rough estimate, we use  the luminosity -- halo mass correlation we find above may be used to construct $\\lya$ luminosity functions (LFs) at different redshift when combined with halo mass functions from large-box, general cosmological simulations. For example, at $z = 3.1$, we divide all galaxies in the snap shot by the halo mass with 0.25 dex, and fit to the median value of each bin, this gives a correlation of $\\La ~({\\ergs})= 10^{32.94} \\times (M_{\\rm halo}^{0.79} / \\Msun)$. We then use this to convert the halo mass function of \\citet{Sheth99} to the $\\lya$ LF.  Figure~\\ref{fig:LF} shows the resulting $\\lya$ LFs in comparison with observations at redshift $z = 3.1, 5.7$, respectively. The red solid curves are LFs above a  detection threshold of $S_{\\rm Ly\\alpha} = 10^{-18}~\\rm ergs \\; s^{-1} \\; cm^{-2} \\; arcsec^{-2}$ \\citep{Ouchi08, Ouchi10}, while the red dashed lines represent LFs from total luminosity (counting all escaped photons without a flux cut).  While the un-filtered LFs seem to agree with observations of \\cite{Gronwall07} and \\cite{Ouchi08}, the filtered ones are significantly off. The difference comes from the reduction of $\\La$ and $\\fesc$ due to the flux cut. Moreover, the dispersion in the luminosity -- halo mass relation at different redshift may cause a large scatter in the LFs.  This plot suggests that the current simulation in this work is not suitable to study a large galaxy population and its statistical properties, because there are too few observable LAEs. Moreover, as discussed earlier, the predicted $\\lya$ properties may be affected by a number of numerical and physical limitations of our model. For example, the one-phase model currently used in the present work may  underestimates the density of cold hydrogen gas, and hence underestimates the $\\lya$ flux. We will study the $\\lya$ LFs at different redshift in a forthcoming paper with the improved  $\\art$ which incorporates a two-phase ISM model, and a general simulation with mean overdensity in a larger volume (Yajima et al, in preparation).  \\begin{figure} \\begin{center} \\includegraphics[scale=0.4]{LF_wSB.eps} \\caption{ Derived $\\lya$ luminosity functions from our model in comparison with observations. The red, solid lines are filtered LFs above a  detection threshold of $S_{\\rm Ly\\alpha} = 10^{-18}~\\rm ergs \\; s^{-1} \\; cm^{-2} \\; arcsec^{-2}$ \\citep{Ouchi08}, while the red, dashed lines represent unfiltered ones. The black dashed lines are the Schechter function at $z=3.1$ and $5.7$ derived from LAEs observation by \\citet{Ouchi08}, and the black dotted line is from observation at $z=3.1$ by \\citet{Gronwall07}. } \\label{fig:LF} \\end{center} \\end{figure}  \\subsection{Limitations of Our Model} \\label{sec:limit}  As demonstrated above, our model is able to explain a number of observed properties of LAEs at different redshift. However, we should point out that our current simulations suffer from a number of major limitations which may affect the predicted $\\lya$ properties.  \\begin{itemize}  \\item In the current work, we use a one-phase ISM model, which considers the average density and temperature of the gas. Such a model likely underestimates the density of cold hydrogen gas, which may lead to significant underestimate of the $\\lya$ emission coming from cold ($\\sim 10^4$~K), dense gas. On the other hand, such a model also underestimates the amount of dust associated with cold molecular gas, which likely results in underestimate of absorption of $\\lya$ photons by gas and dust. We will investigate the $\\lya$  RT and ionization structures in a two-phase ISM model in a forthcoming paper.  \\item The absorption and transmission of IGM are not taken into account in the RT calculations. As discussed in the previous section, these two effects may suppress the $\\lya$ flux and change the line profiles.  \\item The simulations do not have sufficient resolutions to resolve dense regions and outflow, which requires a high spatial resolution of  $\\sim \\rm pc$ \\citep[e.g.,][]{Fujita09}. It is a challenge for cosmological simulations to resolve both the inflow gas from large scales of $\\sim \\rm Mpc$ and the outflow from pc-scale star forming regions. For a simulation with a box of 100 Mpc like the one we have, this requires a large dynamical range over eight orders of magnitude, which is beyond the scope of our current work.  \\item The simulation box in this work is too small to study a large galaxy population, as well as effects of environment on galaxy properties and their evolution. One needs uniform simulations in large volumes in order to systematically investigate the formation and evolution of L* galaxies.  \\end{itemize}  Finally, we stress once again that caution should be taken when comparing directly the results from our calculations  to data from a given survey, because, as discussed above, the observed $\\lya$ properties depend sensitively on a number of factors, including galaxy properties, viewing angle, model parameters, and observational threshold."
    },
    "1112/1112.0037_arXiv.txt": {
        "abstract": "We present the results of numerical simulations of wave propagation and jet formation in solar atmosphere models with different magnetic field configurations. The presence in the chromosphere of waves with periods longer than the acoustic cutoff period has been ascribed to either strong inclined magnetic fields, or changes in the radiative relaxation time. Our simulations include a sophisticated treatment of radiative losses, as well as fields with different strengths and inclinations. Using Fourier and wavelet analysis techniques, we investigate the periodicity of the waves that travel through the chromosphere. We find that the velocity signal is dominated by waves with periods around 5~minutes in regions of strong, inclined field, including at the edges of strong flux tubes where the field expands, whereas 3-minute waves dominate in regions of weak or vertically oriented fields. Our results show that the field inclination is very important for long-period wave propagation, whereas variations in the radiative relaxation time have little effect. Furthermore, we find that atmospheric conditions can vary significantly on timescales of a few minutes, meaning that a Fourier analysis of wave propagation can be misleading. Wavelet techniques take variations with time into account and are more suitable analysis tools. Finally, we investigate the properties of jets formed by the propagating waves once they reach the transition region, and find systematic differences between the jets in inclined field regions and those in vertical field regions, in agreement with observations of dynamic fibrils. ",
        "introduction": "Chromospheric wave propagation has been an extensively studied, but poorly understood, subject in solar physics for some time. One puzzle has been the presence of propagating waves with periods on the order of 5 minutes or more. Such waves were observed by, e.g., \\citet{Orrall1966} and \\citet{Giovanelli+etal1978}; this was considered surprising since the acoustic cutoff period in the chromosphere, above which waves should not be able to propagate, is on the order of 200~s. Such long-period propagation was later found to be widespread, generally occurring wherever the local magnetic field is strong. \\citet{Lites+etal1993} show an example of observations of neighboring internetwork and network regions that have a marked difference in their Fourier spectra.  It was realized \\citep{Michalitsanos1973,Bel+Leroy1977,Suematsu1990} that, since magnetoacoustic waves in a strongly magnetized medium are restricted to propagating along field lines, the effective gravity (i.e., the component of gravity along the magnetic field) would be reduced in magnetic regions. Since the cutoff frequency (in an isothermal atmosphere) is given by \\begin{equation} \\nu_{ac}=\\frac{\\gamma g_{eff}}{4\\pi c_s}, \\label{cutoffeq} \\end{equation} where $c_s$ is the sound speed and $\\gamma$ is the ratio of specific heats, the cutoff frequency would also be lower in regions of strong inclined field, potentially allowing 5-minute ($\\sim$3~mHz) waves to propagate. This hypothesis has since been tested in a number of increasingly advanced numerical simulations, e.g. \\citet{DePontieu+etal2004}, \\citet{Hansteen+etal2006}, and \\citet{Heggland+etal2007}. All have found that it is an effective mechanism for transmitting 5-minute power through the chromosphere; some models have suggested that the leakage of 5-minute waves into the chromosphere can also explain the presence of 5-minute oscillatory signal in the corona \\citep{DePontieu+etal2005}.  One criticism of this explanation has been that not all observations of 5-minute propagation are in regions of obviously inclined field. An alternative explanation has therefore been suggested, originally by \\citet{Roberts1983} and later developed and tested by \\citet{Centeno+etal2006,Centeno+etal2009} and \\citet{Khomenko+etal2008}, in which changes in the radiative relaxation time associated with small scale magnetic structures are responsible for increasing the cutoff period. This mechanism has also been demonstrated to enable propagation of 5-minute oscillations, even in vertical magnetic structures. However, the underlying basis of this theory and the related simulations is a highly simplified energy equation in which radiative losses are approximated by Newton's law of cooling (\\citeauthor{Khomenko+etal2008} \\citeyear{Khomenko+etal2008} and references therein).  A related subject that has been studied extensively is the generation and propagation of chromospheric jets such as spicules (type I and II), macrospicules, surges, fibrils, and mottles. A variety of models have been proposed over the years (see \\citeauthor{Sterling2000} \\citeyear{Sterling2000} for a review), and although it remains unclear whether the many types of jets are different manifestations of the same underlying physical phenomenon or not, there are at least strong indications that the jets known as dynamic fibrils are driven by shock waves traveling through the chromosphere \\citep{Hansteen+etal2006,DePontieu+etal2007,Heggland+etal2007,Martinez+etal2009a}. Many of these jets have lifetimes around 5~minutes, and so the waves driving them need to have been channeled into the upper chromosphere via one of the processes outlined above.  In this paper, we present the results of two-dimensional simulations of wave propagation from the convection zone to the transition region and corona. The simulations include a sophisticated treatment of radiative losses and heat conduction, and study the effects of different magnetic field geometries and strengths on the propagation of waves through the chromosphere. We also look at the jets that these waves produce once they reach the transition region, and perform a statistical comparison of the jets produced in a model with vertical field and in one with inclined field.  In \\textsection{}2 we describe the different simulations and the code. \\textsection{}3 contains an analysis of wave propagation and periodicities in the various models, and the results are discussed in \\textsection{}4. \\textsection{}5 looks at jet formation and properties, and a summary and conclusions follow in \\textsection{}6. ",
        "conclusions": "\\begin{figure*} \\plottwo{f24a.ps}{f24b.ps} \\caption{Fourier spectra of the vertical velocity taken from our simulations (left panels) and the Doppler velocity derived from synthesized spectral lines (right panels). The synthesized spectra have been smoothed over 2 arcseconds. In all panels, we show the spectrum of case B in the center (between $x=16.5$~Mm and $x=33$~Mm), flanked on both sides by the spectrum of the non-magnetic case A. This figure can be compared with Figure 4 in \\citet{Lites+etal1993}.} \\label{litesrad} \\end{figure*}  The main result from our simulations is that 5-minute waves are able to propagate through the chromosphere in regions where the magnetic field is inclined and sufficiently strong. In regions where the field is vertical or weak, the velocity field is dominated by waves with periods of around 3 minutes.  As mentioned in the introduction, a model which uses the radiative relaxation time, rather than the inclination of the magnetic field, as the mechanism for increasing the cutoff period, has been proposed, and has been demonstrated to work in numerical simulations \\citep{Khomenko+etal2008}. However, it relies on a simple Newtonian cooling model for approximating radiative losses.  Our simulations use an advanced and realistic method for computing the radiative losses \\citep{Gudiksen+etal2011} that includes all the important mechanisms involved in radiative losses in the photosphere and chromosphere, and as such it is considerably more realistic than a simple Newtonian cooling model. While the radiative relaxation time model expects 5-minute propagation above all strong small-scale magnetic structures, regardless of the field inclination, our simulations show 3-minute propagation above the central region of flux concentrations, and 5-minute propagation in inclined-field regions to the sides. This is exactly the result predicted by the field inclination model, but is in conflict with the predictions of the radiative relaxation model. Our conclusion, based on the simulations presented here, is therefore that the radiative relaxation model is not effective when a more realistic energy equation is considered.  While large-scale regions of homogeneous inclined field, like our cases D and E, represent idealizations of the conditions at the edge of plages, our case C represents realistic conditions for an isolated magnetic element or pore. In this model, the magnetic field is largely vertical in the photosphere and above the transition region, while the field expansion leads to significant inclination where it is needed the most, in the chromosphere. This can then account for the 5-minute propagation, but we expect 3-minute propagation in the center of the flux concentration. Is this supported by observations?  The answer is in fact not very clear. \\citet{Lites+etal1993} show an example of a network region which has very little power at frequencies above 4~mHz over most of its area. There is, however, a pronounced spike of higher-frequency power located in the center of the network ($x=30$~Mm in their Figure 4). The authors attribute this to noise from seeing, which it may very well be; it is visible at frequencies up to 20~mHz, and a study of the coherence spectra also lends credence to this explanation. However, the signal at frequencies of 5-6~mHz in this spike is much stronger than that at higher frequencies, and indeed is of the same magnitude as the (real) signal at those frequencies in the neighboring quiet regions. It is notable that such a spike, with a spatial extent of 1-2~Mm, is in fact exactly what we find in our case C. The central spike observed by \\citeauthor{Lites+etal1993} could thus be real.  In order to make more direct comparisons to their results, which were based on Doppler shifts in observations of the Ca~{\\sc ii}~H line at $3968.49$~\\AA{} and the Fe~{\\sc i} line at $3966.82$~\\AA, we have computed synthetic spectra in these lines from our simulations using the non-LTE radiative transfer code MULTI \\citep{Carlsson1986}. The Ca line is treated in full non-LTE, while the Fe line, which is formed in the wing of the Ca line, is included in the same computation using additional localized opacity and source function terms based on LTE. The calculations have been performed column by column, i.e., neglecting any radiative interaction in the $x$-direction. The computed spectra have been smoothed over $2''$, which appears to be close to the effective resolution of the observations in \\citet{Lites+etal1993}, and we have then calculated Doppler shifts from these smoothed spectra.  It should be noted that Ca~{\\sc ii}~H has a very complex line profile, usually with several emission peaks within the deeper absorption line. In a dynamic atmosphere, particularly in the presence of strong shocks, these emission peaks can become very large and are often asymmetric (so called bright grains), and the line can undergo central reversal as well. Defining and identifying the line center of such a line is a non-trivial task. We have used a method that finds the center of the region where the intensity is below a certain threshold above the minimum intensity. This method generally gives acceptable results, but can not be expected to correspond directly to the velocity at any one given height in the simulations.  The power spectra of the calculated Doppler velocities are shown in the two right-hand panels of Figure~\\ref{litesrad}; the two left-hand panels show the power spectra of the vertical velocity taken directly from our simulation data, at heights corresponding to the approximate formation heights of the lines. Since the observations of \\citet{Lites+etal1993} covered a larger area than our simulation boxes, we show a combination of the spectra from two of our simulations in the figure. In the center of all panels, between $x=16.5$~Mm and $x=33$~Mm, we show the spectrum of our case B, representing network conditions. On both sides, from $x=0$~Mm to $x=16.5$~Mm and from $x=33$~Mm to $x=49.5$~Mm, we show the spectrum of our case A, representing conditions in the weakly magnetized internetwork. This figure should be compared with Figure 4 in \\citet{Lites+etal1993}.  While our synthesized Ca power spectrum (upper right) is not a perfect match to the observations of \\citeauthor{Lites+etal1993}, there are many similarities. Of particular note is the difference in the dominant frequencies between the non-magnetic regions on both sides and the network region in the center. The network (case B) is dominated by lower frequencies, particularly in the area between $x=22$~Mm and $x=29$~Mm, where the dominant frequencies are 3-$3.5$~mHz. The internetwork (case A) has a more scattered spectrum, but the dominant frequencies are mostly between 5 and 7~mHz. Furthermore, the network has significant power at very low frequencies, below 3~mHz. These effects are also found by \\citeauthor{Lites+etal1993}.  This power at low frequencies is not found in the simulation velocity at $z=1$~Mm (upper left), and although the velocity spectrum and the Ca spectrum are similar in many ways, there are also several differences. These differences are partly due to the smoothing applied to the Ca data, partly due to the difficulty of defining a meaningful Doppler shift of the highly complex Ca line profile, and partly due to the fact that the line is formed over a range of heights rather than at one given height, and this height range can also vary with both horizontal position and time. Using the Ca Doppler velocity as a proxy for atmospheric velocity on the real Sun can therefore be misleading.  The velocity spectrum of the Fe line (lower right) shows a very good correspondence with the velocity at $z=0.25$~Mm in the models (lower left). This line has a much simpler profile and is formed in a region with few strong shocks affecting local conditions. The dominant frequency is 3~mHz in both the network and the internetwork, though there is also some power at 5~mHz in most locations. \\citeauthor{Lites+etal1993} also find that the power in this line is in the 3-5~mHz band, although most of it is between 3 and 4~mHz. Like us, they find no significant difference between the spectra in the internetwork and in the network in this line.  More recent observations related to the question of wave periodicity have been performed by \\citet{Centeno+etal2009}, \\citet{deWijn+etal2009}, and \\citet{Stangalini+etal2011} \\citep[see also][]{Jefferies+etal2006}. \\citeauthor{Centeno+etal2009} used the Tenerife Infrared Polarimeter of the German Vacuum Tower Telescope at the Observatorio del Teide, with a seeing-limited spatial resolution of $1''-1\\farcs{}5$. They then found propagating 5-minute waves in the chromosphere above a facular region. The photospheric magnetic field as determined from Stokes inversions was within 20$^{\\circ}$ of the vertical. \\citeauthor{deWijn+etal2009}, using the Solar Optical Telescope on {\\it Hinode} with a resolution of $0\\farcs{}16$, found 3-minute signal in the center of a plage region, but more 5-minute propagation towards the sides in the direction of the expanding field. \\citeauthor{Stangalini+etal2011} used a combination of data from {\\it Hinode} and the IBIS instrument at the Dunn Solar Telescope, with an estimated average resolution of $0\\farcs{}36$. They found propagating 5-minute waves along the inclined field on the edges of a pore, and some power in 3-minute oscillations at the center. They also found both 5-minute and 3-minute propagation, though with more power in the 5-minute band, in a nearby region with small magnetic elements where they estimate that the chromospheric magnetic field is close to vertical, based on a force-free field extrapolation.  The results of our simulations are in agreement with \\citet{deWijn+etal2009}, and with the propagation patterns observed by \\citet{Stangalini+etal2011} around their pore. In the more vertical magnetic structures observed by \\citet{Stangalini+etal2011} and \\citet{Centeno+etal2009}, we would expect more power in the 3-minute band than in the 5-minute band based on the results of our simulations, if indeed the field is mainly vertical and the flux tubes do not move around very much. There are, however, several possible mechanisms that could explain the 5-minute dominance and resolve this apparent difference.  For one, all flux tubes naturally expand with height, and this expansion creates a region of inclined field (as illustrated, on a large scale, by our case C). Thus, even if the field is close to vertical in the photosphere, there will be regions between the photosphere and the chromosphere where the field at the edges of the flux tubes is inclined, and the long-period waves can propagate there. \\citet{Centeno+etal2009} do study coherence spectra to look for signs of a possible horizontal shift in the signal as a result of propagation along inclined fields, and find good coherence between the photospheric and chromospheric signal, but our results show that the field does not need to be inclined throughout the photosphere and chromosphere in order to enable 5-minute wave propagation. A few hundred km along the edges of an expanding flux tube may be enough, and any horizontal shift could then be less than one resolution element ($1''-1\\farcs{}5$ in their data). In the observations of \\citet{Stangalini+etal2011}, the photospheric field is not uniformly vertical, and a force-free extrapolation is not a very good approximation in the chromosphere. Furthermore, although there is more power in the 5-minute signal, they also find significant signal at periods around 3 minutes in the region with smaller magnetic elements. We believe that our model, where field inclination is the dominant mechanism for allowing long-period wave propagation, is compatible with these findings.  Flux tube movement and limited resolution may also be partially responsible for the relative dominance of 5-minute power in the observations of \\citet{Centeno+etal2009}. In case B, we found that the strongest 3-minute power appeared above flux tubes that undergo little horizontal motion, while 5-minute power was found in inclined field regions at the edges of flux tubes. If the flux tubes move around, both the 5-minute and 3-minute power will be spread out and one would not see a clear distinction between the (average) flux tube center and the sides in a Fourier analysis. The flux tube center also covers a rather small area at any given time, and this makes the related 3-minute waves difficult to observe in low-resolution data. In the higher-resolution data of \\citet{Stangalini+etal2011}, regions of 3-minute propagation are found, and this could possibly be because the flux tubes at these locations move around less. We would encourage observers to look for differences in the periodicity of oscillations at the center and edges of flux tubes in future high-resolution datasets.  A third possibility is that heating may play a role. The temperature structure of the chromosphere is in general not well known. In particular, the real Sun may have more magnetic heating of the upper chromosphere than our 2D models. Higher temperature would reduce the cutoff frequency and allow 5-minute waves to propagate more easily.  Yet another suggestion, as mentioned by \\citet{deWijn+etal2009}, is that the field may be twisted. The waves could then travel along field lines that are everywhere inclined with respect to the local vertical, but without significant horizontal displacement. Such field twist is a 3D effect and can not be tested in our 2D simulations, but should be considered in later work.  A different point, that we have already mentioned in our analysis, is also worth making: although a Fourier analysis can be a powerful tool, it has some important limitations. In order to achieve sufficient spectral resolution, one usually needs time series on the order of one hour. The solar atmosphere, however, is dynamic on timescales of minutes. Atmospheric conditions can and do change, and the Fourier transform is not well suited for picking up such changes. At least in areas where the general signal is weak, non-recurring events can end up dominating the power spectrum (see Figure~\\ref{bampang456} for an example; there are several other examples in the dataset). In such cases, the Fourier analysis does not say all that much about the general conditions at that location. This can be particularly dangerous because the eye is naturally drawn to peaks in the power spectrum. Also, there can be times when local conditions are notably different from the time average, and these can be correlated with changes in the signal (e.g., Figure~\\ref{vertpwrudang185}). A wavelet analysis, which takes time variations into account, or at least a careful examination of the time series, is essential for identifying the actual processes going on in the atmosphere.    We have presented several simulations investigating the periodicity of waves propagating through the chromosphere. We find that waves with periods of around 5 minutes can propagate in regions where the magnetic field is strong and inclined, including at the edges of flux tubes. In regions where the magnetic field is weak or vertical, we find primarily 3-minute waves; this also applies above vertical flux tubes and in the center of strong expanding flux tubes. These results indicate that field inclination is critical to the propagation of long-period waves. Where the flux tubes undergo significant horizontal motion, both the 5-minute and the 3-minute power is spread out and the distinction is not as clearly visible. Since we have included an advanced treatment of radiative losses in our simulations and find 3-minute propagation above vertical field regions, we conclude that variation in the radiative relaxation time is not an effective mechanism for increasing the cutoff period. Our simulations are in agreement with the results of recent high-resolution {\\it Hinode} observations.  We have also studied the jets produced by these waves once they reach the transition region. We find systematic differences between the jets produced in a model with mostly vertical field, and in a model with mostly inclined field. The results are in agreement with observations of dynamic fibrils.  It is also important to point out, for the purpose of future analyses of wave propagation and periodicity, whether from simulations or observations, that the Fourier analysis can be misleading and hide important information about the state of the medium. This is because the solar atmosphere is dynamic and changes on timescales much shorter than one hour, which is often the minimum timescale needed to achieve sufficient spectral resolution in a Fourier analysis. A wavelet analysis takes variations in time into account and can be an invaluable tool for figuring out what processes are important for the dynamics."
    },
    "1112/1112.2497_arXiv.txt": {
        "abstract": "{} {We investigate the surface nitrogen content of the six magnetic O stars known to date as well as of the early B--type star $\\tau$ Sco. We compare these abundances to predictions of evolutionary models to isolate the effects of magnetic field on the transport of elements in stellar interiors.} {We conduct a quantitative spectroscopic analysis of the sample stars with state-of-the-art atmosphere models. We rely on high signal-to-noise ratio, high resolution optical spectra obtained with ESPADONS at CFHT and NARVAL at TBL. Atmosphere models and synthetic spectra are computed with the code CMFGEN. Values of $N/H$ together with their uncertainties are determined and compared to predictions of evolutionary models.} {We find that the magnetic stars can be divided into two groups: one with stars displaying no N enrichment (one object); and one with stars most likely showing extra N enrichment (5 objects). For one star ($\\Theta^{1}$ Ori C) no robust conclusion can be drawn due to its young age. The star with no N enrichment is the one with the weakest magnetic field, possibly of dynamo origin. It might be a star having experienced strong magnetic braking under the condition of solid body rotation, but its rotational velocity is still relatively large. The five stars with high N content were probably slow rotators on the zero age main sequence, but they have surface $N/H$ typical of normal O stars, indicating that the presence of a (probably fossil) magnetic field leads to extra enrichment. These stars may have a strong differential rotation inducing shear mixing. Our results should be viewed as a basis on which new theoretical simulations can rely to better understand the effect of magnetism on the evolution of massive stars.} {} ",
        "introduction": "The evolution of massive stars is governed by two main physical processes: their mass loss \\citep{cm86} and their rotation \\citep{mm00}.  All additional processes directly affecting these properties are thus major contributors to the fate of massive stars. This is the case of magnetism. Since the pioneering work of \\citet{bm97}, it is known that the presence of a large scale surface magnetic field will modify massive stars winds. By deflecting material along the field lines, magnetism affects the wind geometry and structure. Simulations by \\citet{ud02,ud08,ud09} have shown that equatorial overdensities can be created. Material can either fall back onto the stellar surface or be ejected in the magnetic equatorial plane depending on the interplay with rotation. In the case of very strong magnetic fields, observations and simulations show that a rigid magnetosphere strongly affects the wind structure, creating caps of stellar material in the minima of the magneto-rotational potential \\citep{townsend07}. Since the wind flow is deflected by the presence of a magnetic field, the way the star loses its angular momentum is affected. In addition, the internal structure of the star is also affected, forcing the star to follow new evolutionary paths. Given the relevance of massive star evolution to many fields of astrophysics (galactic chemical evolution, interstellar dynamics, supernovae and gamma--ray bursts, nucleosynthesis...), it is crucial to understand the magnetic properties of these objects in order to constrain their influence on the evolution and fate of OB stars.  The recent development of powerful spectropolarimeters has revolutionized our view of massive stars magnetism. The first detection of a magnetic field on the Trapezium O--type star $\\Theta^{1}$~Ori~C \\citep{donati02} dates back to less than a decade. This spectroscopically varying and strong X--ray emitting star has revealed a surface kilogauss field organized mostly as a dipole. Subsequent detections have been made on HD~191612 \\citep{donati06}, HD~148937 \\citep{hubrig08,wade11}, HD~57682 \\citep{grunhut09} and HD~108 \\citep{martins10}. Field strengths are of several hundred Gauss to a few kilogauss and the geometry, when constrained, is consistent with a dipole. The only exception is the O supergiant $\\zeta$~Ori~A \\citep{bouret08} for which a more complex field topology was tentatively derived, together with a remarkably weak field of 50--100 G. This raises the question of the origin of magnetism in massive stars. The debate is ongoing. A fossil origin, stable over Myrs and with a rather simple geometry, is supported by some simulations \\citep{brait04,duez10}. In that case, the star retains the original magnetic field of its parental molecular cloud or dynamo field of its convective progenitor. Alternatively, dynamo processes may be at work. In analogy with low mass stars where a convective envelope (or even a fully convective structure) exists, magnetic field might be generated in the convective core of high mass stars and subsequently transported to the surface \\citep{charb01,brun05}. A dynamo may also operate in the radiative zones of massive stars \\citep{spruit02,mcdo04}. The field would be produced by the shear in this region. This idea has been tested in numerical simulations, with contradictory results \\citep{brait06,zahn07}.  However, assuming that this dynamo process was able to produce and maintain a magnetic field, evolutionary calculations have been conducted to test the effects of this type of magnetism on the evolution of massive stars \\citep{mm04}. They find that the internal structure is strongly affected, with a completely different rotation profile in the star compared to non-magnetic calculations including rotation. A consequence of this re-organization of the stellar interior is a modification of the transport of angular momentum and of chemical elements. In particular, \\citet{mm05} studied the effects on the surface chemical appearance of the star. They show that the presence of a dynamo generated magnetic field significantly increases the efficiency of mixing, leading to strong nitrogen and even helium surface enhancement. Further analysis of the feedback effects of magnetic braking balanced these results though, highlighting the importance of the type of rotation profile (differential versus solid-body) in the final chemical appearance of the star \\citep{meynet11}.  Although still limited, the sample of magnetic O stars is now populous enough to test these predictions. An important side effect of the spectropolarimetric observations is the production of very high S/N ratio optical spectra at high resolution. These are perfectly suited for abundance studies. In this paper, we thus analyze the surface nitrogen content of the magnetic O stars known to date as well as of the B0.2V star $\\tau$ Sco. Our aim is to see if peculiar chemical patterns are observed and if so, to what extent they can be explained by the current generation of evolutionary models. The goal is to identify specific patterns that can only be attributed to the presence of a magnetic field. This is particularly important in the context of explaining the population of N--rich slowly rotating stars in the Magellanic Cloud reported by \\citet{hunter08,hunter09}. Such stars are not explained by rotating evolutionary tracks and have been suggested to be magnetic. The observations are presented in Sect.\\ \\ref{s_obs}. We describe the analysis with atmosphere models in Sect.\\ \\ref{s_mod}. The results are discussed in Sect.\\ \\ref{s_mag}. We show how the chemical properties of the magnetic O stars can be interpreted in terms of evolutionary models with rotation and with or without magnetic field. We finally summarize our conclusions in Sect.\\ \\ref{s_conc}. ",
        "conclusions": "\\label{s_conc}  We have conducted a quantitative spectroscopic analysis of the six magnetic O stars known to date together with the early B star $\\tau$ Sco. Spectra collected with ESPADONS at CFHT and NARVAL at TBL have been used. Atmosphere models computed with the code CMFGEN have been computed. We have determined the stellar parameters and the surface nitrogen abundance. A careful investigation of the uncertainties associated with $N/H$ has been pursued.  Our results can be divided into two parts:  \\begin{itemize}  \\item[$\\bullet$] The surface $N/H$ of $\\zeta$ Ori A -- the star with the weakest magnetic field -- is lower than the predictions of evolutionary models. For that star the presence of the magnetic field probably quenches the transport of chemical elements. The conditions under which such a process can happen are not clear. Evolutionary models accounting for magnetic braking predict no chemical enrichment in solid body rotating stars. But their rotational velocities are also rather low, whereas  $\\zeta$~Ori~A appears to have a rather standard \\vsini.  \\item[$\\bullet$] The other magnetic O stars, especially the three Of?p stars, display surface nitrogen content consistent with those of non-magnetic O stars with initial \\vsini\\ of 300 \\kms. However, these magnetic O stars were most likely slow rotator on the zero age main sequence (with values of \\vsini\\ of a few tens of \\kms\\ at most). Consequently, they feature larger surface $N/H$ than they should if they did not host magnetic fields as also indicated by their nitrogen excess compared to rotating evolutionary models.  \\end{itemize}  The presence of a magnetic field in O stars thus appears to produce two types of chemical peculiarities: either an extra enrichment, or the absence of enrichment. Confrontations with evolutionary predictions in the case of dynamo fields indicate that solid body rotation might explain the properties of the non--enriched star, but its relatively normal rotational velocity is a puzzle in that context. For the N--rich magnetic O stars, a fossil field is usually assumed. Our results should be a basis for comparison with future evolutionary calculations including fossil fields since such evolutionary models do not exist at present."
    },
    "1112/1112.6388_arXiv.txt": {
        "abstract": "Earliest AUGER (the Pierre Auger Observatory) UHECR (Ultra High Energy Cosmic Rays) anisotropy correlated with AGN (active galactic nuclei) within a  GZK (Greisen-Zatsepin-Kuz'min  cut off)  Universe almost fade away. Recent UHECR mass compositions did show a  negligible nucleon composition and an UHECR nuclei (light or heavy) signature. Indeed  last map miss the Super-Galactic Plane. The absence of  UHECR  events toward the Virgo cluster, an unique spread clustering of events around  Cen-A, our nearest AGN, suggested a  He-like  nuclei as the main extragalactic UHECR component from Cen A, coexisting with Auger and HIRES  (High-Resolution Fly's Eye) composition. Because the light nuclei fragility  such He UHECR cannot arrive from Virgo (being too far). Multiplet at twenty EeV along Cen A recently discovered by Auger confirm this interpretation as being foreseen to be indebt to fragments ($D,He^{3}, p$) that had to reach us along the same UHECR. However  remaining majority of UHECR clustering are partially correlated  with a gamma noise at (1-3 MeV) in Comptel sky, linked to $Al^{26}$ galactic radioactive map as well as to a few TeV gamma (ICECUBE-ARGO) anisotropy maps; rare UHECR triplet are overlapping on Vela  TeV anisotropy and other nearest galactic gamma sources (as partially Crab and a Galactic core corona). Therefore UHECR   might be also (or mostly) heavy radioactive galactic nuclei  as $Ni^{55}$, $Ni^{56}$,  $Ni^{57}$  and $Co^{57}$ bent from the sources whose $\\beta$ and $\\gamma$ radioactivity and decay in flight is boosted (by Lorentz factor $\\Gamma_{Ni}\\simeq 10^{9}- 10^{8}$),  leading to  TeV correlated sky anisotropy. Galactic UHECR signals inside the inner center maybe suppressed by the largest spreading repulsive Lorentz bending forces. More clustering around external galactic plane is nevertheless expected in present and future data. Magellanic Cloud and Magellanic  Stream may also rise more and more in UHECR maps (as well as in multiplet signals). Future UHECR clustering might be observed  around Cas A and Cygnus by T.A. (Telescope Array). The UHECR spectra cut off may be not an extragalactic GZK feature but just the imprint of a galactic confinement and-or  a spectroscopic heavy composition decrease step. ",
        "introduction": "Ultra High Energy Cosmic Rays, UHECR, Astronomy and nuclear composition are more and more in severe conflict. This contradiction was already inscribed in  early apparent discover of a Super Galactic UHECR correlation \\cite{Auger-Nov07} with a first UHECR nuclei composition signature. It was note realistic for UHECR to arrive exactly  from Super Galactic plane, as apparent map were suggesting, while being heavy nuclei (and not nucleon). A very heavy UHECR composition being greatly bent by galactic magnetic fields cannot open to any sharp astronomy nor  correlate to Super Galactic Plane, nor it can explain easily the unique observed mild UHECR Cen A event clustering.     Heavy nuclei as Fe, Ni, Co   may  however induce large scale anisotropy or spread inhomogeneity (as possibly from a nearby source, Vela). Moreover let us remind that  UHECR (if nucleon and if  extragalactic) are making photo-pion interacting on cosmic CMBR (Cosmic Microwave Background Radiation) and a consequent EeV gamma and neutrinos secondaries. These expected signals at EeV for gamma and tau neutrinos are more and more constrained being  at the edge of AUGER detection, see \\cite{Auger08}.  If UHECR are light nuclei they may produce  fragment secondaries nuclei or nucleon, as well as gamma-ray and neutrino tails at tens of TeV-PeV because nuclei photo-dissociation; also heavy nuclei, if extragalactic, may lead by photodisintegration on extragalactic flight leading to such TeVs-PeV gamma and neutrino tail. If UHECR are very heavy  they maybe bent and bound and confined within their local birth source group. For instance UHECR heavy nuclei maybe bent and confined within Virgo group ; in analogy heaviest UHECR may be bent, crowded and nearly contained within our galaxy, but they cannot produce much photodisintegration secondaries. However \\emph{UHECR radioactive nuclei} as  $Ni^{56}$, $Ni^{57}$ and $Co^{57}$ $Co^{60}$, while flying and decaying, may trace their presence by boosted gamma (and positron) tails and beta decay neutrinos. If such UHECR are produced by Super-novae remanent, SNRs (SuperNova Remanents) , or better by their twin micro-quasars jets (in GRB-SN asymmetric jet explosions and late precessing jet model, see \\cite{Fargion1998}) these tails might be observable. The relic gamma radiation (in rest frame of the nuclei around hundred keV) will be boosted by UHECR $E \\geq 6 \\cdot 10^{19} eV$ energy, heavy nuclei huge Lorentz factor ($\\Gamma_{Ni^{55}} \\simeq 10^{9}$), shining at our laboratory system around tens-hundred TeV gamma region, observable at Milagro-ARGO-ICECUBE  TeVs CR maps as their correlated anisotropy, see Fig \\ref{fig4-5}. A comparable trace somehow  correlated, but at more non relativistic stages, is offered by heavy nuclei (as  the $Al^{26}$ at MeV Comptel gamma maps see Fig. \\ref{1r}) in galactic plane  \\cite{Fargion2011}. Inspired by this  first radioactive gamma-UHECR connection possibility   now we extend such understanding  for the partial gamma TeV-UHECR connection anisotropy assumed born by  boosted nuclei decay, see Fig.\\ref{fig4-5}. As a natural consequence present UHECR cut off may be related not to any real nucleon extragalactic GZK cut-off but  to a  heavy nuclei spectroscopy step, a confinement in our own galaxy, an  nearby anisotropic distances dilution. The same heavy nuclei ejection by nearby Supernova maybe the source of local nuclei anisotropy made by asymmetric SN-jet explosions (see:\\cite{Hansen:SN-Uranium}).  \\begin{figure}[htb] \\begin{center}  \\includegraphics[width=3.1 in]{Disk-Ellipse1.eps} \\includegraphics[width=3.1 in]{Comptel-1-2-3.eps} \\caption{ Left: The AUGER UHECR event map and two of the three AUGER multiplet clustering toward Cen A; a thin narrow elliptical area and a small disk mark the place.  A third Multiplet clustering points toward Large and Small Magellanic clouds and it overlaps on Magellanic stream. Right: last UHECR event map by AUGER where the clustering near Cen A overlap the MeV Comptel gamma (1-3 MeV)  map; note connected map beside nearby AGN Cen A, along the nearest pulsar Vela, the Magellanic stream and the galactic core. }\\label{1r} \\end{center} \\end{figure}  \\begin{figure}[!t] \\vspace{5mm} \\centering \\includegraphics[width=4.8 in]{ALL-TEV-ARGO-AUGER-HIRES.eps} \\caption{ The AUGER 2010 UHECR (red) and Hires (blue) event map in celestial coordinate on recent (2010) TeV diffused CR map (ARGO-Milagro-ICECUBE) and labels. See \\cite{ARGO},\\cite{Desiati}, and references. The triplet UHECR event clustering toward Vela and the TeV spread activity around  is remarkable. The Cen A clustering at the fig. center is the main feature in the UHECR map. ARGO TeV anisotropy born around Crab connect and overlap the UHECR events in nearby Orion TeV region. Note doublet along the galactic plane and the  TeV near Cygnus and Cas A regions, where Hires did and we foresee TA may detect UHECR events.} \\label{fig4-5} \\end{figure}  Extragalactic UHECR around CenA  formed (mostly) by  lightest nuclei may explain a partial clustering of events, and the UHECR absence  around Virgo. Light nuclei are fragile and fly few Mpc before being  halted by photo-disruption \\cite{Fargion2008},\\cite{Fargion2009},\\cite{Fargion2010}. The fragments $He + \\gamma \\rightarrow D+D,D+\\gamma \\rightarrow p+n+ \\gamma  , He + \\gamma \\rightarrow He^{3}+n, He + \\gamma \\rightarrow T +p $  may nevertheless trace the same UHECR maps by a secondary  clustering at half or even fourth of the UHECR primary  energy \\cite{Fargion09a}, \\cite{Fargion2011}. At lower energy (at ten EeV or below) the huge smeared cosmic ray isotropy and homogeneity may hide these tiny inhomogeneity traces.  Nevertheless the very recent AUGER EeV map did not show the expected solar Compton Getting anisotropy but an anisotropy in a different region. This result partially disagree with a (widely accepted) extragalactic EeV UHECR component but it favors a more local (galactic) EeV  UHECR presence whose smearing is related to the large Lorentz angle bending and whose non-Compton-Getting inhomogeneity is related to nearby sources, as we are also suggesting in present article. ",
        "conclusions": "$ , $Co^{60}$ radioactive nuclei? } The correlation of UHECR with Cen A, the absence of Virgo, the hint of correlation with  Vela and  with galactic TeV anisotropy, might be in part solved by an extragalactic lightest nuclei, mainly He, for Cen A , (see Fig.\\ref{1r}) and mostly by heavy radioactive nuclei for other sources mostly in our own galaxy (see Fig.\\ref{fig4-5}).  A partial confirm is the predicted \\cite{Fargion09b} ,\\cite{Fargion2011} and observed \\cite{Auger11} multiplet clustering (as Deuterium or proton fragments) at half UHECR edge energy aligned  (see Fig.\\ref{1r}): He like UHECR  maybe bent by a characteristic angle as large as  $\\delta_{rm-He}  \\simeq 16^\\circ$; expected  proton or D fragments multiplet along tails spread  at $\\delta_{rm-p}  \\simeq 32^\\circ$ \\cite{Fargion2011},\\cite{Auger11}. Moreover as shown  here, UHECR Ni,Co  and heavier nuclei too may be deflected by $\\delta_{rm-Ni}  \\simeq 18,7^{o}$ for nearby Vela,  $\\delta_{rm-Ni}  \\simeq  128^{o}$ (or less) for Crab, Gum Nebulae at few kpc away, forcing UHECR  toward nearby TeV inhomogeneities. Such TeV gamma anisotropy is made by boosted UHECR decay at hundred keV gamma and beta positrons shining to us at tens TeV.  In conclusion therefore we foresee  analogous UHECR traces around Cygnus and Cas A (as in TeV map) in growing future TA map. We foresee also crowding and clustering around ($\\mp 18^{o}$) the galactic plane for similar arguments, but not much in galactic center where the stronger magnetic field may spray away the UHECR signals around an external corona ($\\mp 20^{o}$). The UHECR radioactive beta decay in flight  may spread as  gamma TeV traces as well as $\\nu$  TeVs-PeVs anisotropy whose muon signal is  sink into dominant atmospheric muon $\\nu_{\\mu}$ noise. A new spread $\\nu_{\\tau}$ neutrino astronomy, noise free,  related to astronomical (parasite oscillated) tau neutrino and its boosted tau (\\emph{mini-double bang}  within a 5-50 meter size) in Deep Core  or Antares may reveal hundred TeV tau decay (seeing  similar PeVs ones in ICECUBE \\cite{Learned}). Also Tau airshowers may rise in Cherenkov beamed air-showers. \\cite{Fargion1999}, \\cite{FarTau} or fluorescence telescopes at higher energies \\cite{FarTau},\\cite{Feng02},\\cite{Auger08}.  The discover of such expected  Neutrino astronomy  may shed additional light on the UHECR nature, origination  and mass composition, while  opening our eyes to mysterious  UHECR sources. Also future gamma TeV maps and UHECR multiplet correlation as foreseen may  lead to a more conclusive fit of this unsolved, century old, cosmic ray puzzle. \\\\"
    },
    "1112/1112.6341_arXiv.txt": {
        "abstract": "{The formation of spectroscopic binaries (SB) may be a natural byproduct of star formation. The early dynamical evolution of multiple stellar systems after the initial fragmentation of molecular clouds leaves characteristic imprints on the properties of young, multiple stars. The discovery and the characterization of the youngest SB will allow us to infer the mechanisms and timescales involved  in their formation. Our work aims to find spectroscopic companions around young stellar objects (YSO). We present a near-IR high-resolution (R~$\\sim$~60000) multi-epoch radial velocity survey of 7 YSO in the star forming region (SFR) $\\rho$ Ophiuchus. The radial velocities of each source were derived using a two-dimensional cross-correlation function, using the zero-point established by the Earth's atmosphere as reference. More than 14 spectral lines in the CO $\\Delta\\nu$ = (0-2) bandhead window were used in the cross-correlation against LTE atmospheric models to compute the final results. We found that the spectra of the protostars in our sample agree well with the predicted stellar photospheric profiles, indicating that the radial velocities derived are indeed of stellar nature. Three of the targets analyzed exhibit large radial velocity variations during the three observation epochs. These objects - pending further confirmation and orbital characteristics - may become the first evidence for proto-spectroscopic binaries, and will provide important constraints on their formation. Our preliminary binary fraction (BF) of $\\sim$71\\% (when merging our results with those of previous studies) is in line with the notion that multiplicity is very high at young ages and therefore a byproduct of star formation.} ",
        "introduction": "\\indent Since the early twenties that there is growing evidence that multiple stars may be the rule and not the exception. This conception is more recentlt supported by observations of \\cite{duq91} who found that the stellar multiplicity among main sequence (MS) stars in the solar neighborhood may be as high as $\\sim$ 60\\%. Since then, more refined surveys probed the dependencies of the properties of stellar systems on physical parameters like age, mass, separation, and environment. These observations are highly relevant, because there is still no comprehensive theory that explains the properties of multiple stars and their relation to star formation. Recently, \\cite{raga10} confirmed that stellar multiplicity decreases with stellar mass (or later spectral type) in the vicinity of the Sun, a finding that was already suggested earlier \\cite[\\textit{e.g.}][]{sieg05,lada06}. Because the stellar population in the solar neighborhood is dominated by late-type M-stars, this implies that about half of the solar-like main sequence stars in the solar neighborhood occur in multiple systems. \\\\ \\indent It is likely that the multiplicity rate of stars is established in the early phases of star formation or during the dynamical evolution that occurs afterwards \\citep{good04, dedo04, bate09}. And it is plausible that both initial and environmental conditions constrain multiple star formation, and play a vital role in explaining their evolution over time, and their relation to the observed MS stellar multiplicity \\citep{mat00,ster03,cone08,ko08,kacz11}. It is therefore essential to probe the stellar multiplicity as early as possible, to gain a full understanding of the formation and the dynamical evolution of multiple systems. \\\\ \\indent \\cite{gh93,lein93,simon95,pat02,bec03} and \\cite{duc04,duc07} all noted differences in the BF of pre-main-sequence (PMS) stars in different SFRs. In particular, the multiplicity of T Tauri stars with ages $\\lesssim$ 10$^6$ yrs, appears to be higher by a factor of 2 in less dense regions such as Taurus compared to denser SFR such as Ophiuchus \\citep[\\textit{e.g.}][]{simon95,bec03,ratz05} or the field \\citep{pat02, duc04}. Toward even younger ages, radio continuum \\citep[\\textit{e.g.}][]{rei04} and near- and mid- Infrared (IR) \\citep{bars05, hai02, hai04, hai06} investigations found additional evidence for an even higher multiplicity among Class I and Flat-spectrum ($\\sim$ 2 - 5 x 10$^5$ yrs) protostars. \\\\ On the other hand, \\cite{mau10} failed to detect companions in a small sample of 5 protostars at separations 50 \\textless $a$ \\textless 5000 AU ($a$ being the semimajor axis) in a millimeter study of self-embedded, young (Class 0) objects ($\\sim$ 10$^4$ - 10$^5$ yrs). Taken together with the observations of \\cite{loon00}, the authors argue that multiplicity in the separation ranges from 100 to 600 AU would only be defined in a later stage of star formation (namely after the Class 0 phase) and that early multiplicity in pristine systems may not be as ubiquitous and primeval. \\\\ \\indent However, all multiplicity studies in the earliest phases of stellar evolution are plagued with relatively small number statistics. In particular, very little is known about companions of embedded protostars at the sub-AU separation scales, a regime impossible to address with conventional imaging techniques. Spectroscopic Binaries (SB)  may provide additional important constraints on the star-formation mechanism itself. \\cite{toko02} and \\cite{toko06} estimated the SB frequency in MS stars and inferred that 65\\% of their sample of 165 spectroscopic binaries were members of higher order multiple systems, often found in a hierarchical configuration. Remarkably, the frequency of SB in multiple systems strongly depends on the orbital period of the SB: 96 of the SBs with orbital periods shorther than 3 days are in multiple systems, while this rate drops to 30\\% for SBs with orbital periods longer than 13 days. The higher frequency of spectroscopic system within triple or higher order systems suggests an imprint of the formation mechanism itself. Spectroscopic pairs may have lost their orbital angular momentum owing to the presence of a third body that interacts with the inner binary via Kozai cycles and tidal friction \\citep{kos72, ste05, fabr07}. This interaction has the effect of tightening the spectroscopic binary orbit, while the tertiary would eventually be evacuated to an outer region of the system. The SB may therefore hold an important fossil record of (prevailing) initial conditions during their formation. However, this scenario is not conclusive by itself, and it is not clear at which evolutionary stage this process operates efficiently. Observations are required to better constrain the detailed mechanisms of how tight binaries are formed at very young ages. \\\\ \\indent The first successful PMS SB detections were performed by \\cite{herbig77,hart86,mat89} and \\cite{mat94} in the Taurus and Ophiuchus SFRs with a radial velocity precision limited to $\\sim$1kms$^{-1}$ by optical high-resolution spectroscopy. Since then, improving instrumental capabilities and better radial velocity precision allowed to address PMS sources and the lower mass regimes \\citep[\\textit{e.g.}][]{mel03, cov06, ku06, joe06, joe08, prat07}. Only recently, infrared (IR) high-resolution spectroscopy allowed the search for SB among the youngest, embedded PMS and the very low mass stars \\citep[\\textit{e.g.}][]{prat07,bla07,bla10,fi10b,rice10,cro11}. In particular \\citep[][]{fi10b,bla10,cro11} reached a precision of less than 50 ms$^{-1}$ in PMS stars and very low mass stars. Even the planetary mass regime may be in reach using high-stability, high-precision IR spectroscopy. \\\\ \\indent Our own work aims to find SBs in a sample of 38 Class I/ II protostars in the SFR $\\rho$ Ophiuchus. We performed a high-resolution IR spectroscopic survey to derive precision radial velocities (RV) using the NIR spectrograph CRIRES at the VLT \\citep{Kaeufl04}. We describe our sample, the data analysis methodology and results in Sect. 2. Sect. 3 discusses the RV derived in the context of the $\\rho$ Ophiuchus SFR and its implications on the formation of SB at very young ages in general. In Sect. 4. we draw our conclusions. \\begin{figure}[ht] \\centering \\includegraphics[height=8.0cm]{tel_star.eps} \\caption{Process of telluric signature removal in DET4. The extracted nodded spectra of GSS26 (A) is divided by the spectrum of the telluric standard stars with the same exposure time and observed closest in time (same airmass) (B). Lower panel exhibits 'cleaned' nodded spectrum (C).} \\label{allspec8} \\end{figure} ",
        "conclusions": ""
    },
    "1112/1112.4491_arXiv.txt": {
        "abstract": "We explore a class of dark matter models with two dark matter candidates, only one interacts with the standard model sector. One of the dark matter is thermalized with the assistance of the other stable particle. While both stable particles contribute to the total relic density only one can elastically scatter with nuclei, thus effectively reducing the direct detection rate. ",
        "introduction": "The most natural explanation for the astrophysical and cosmological indications of a large component of dark matter in the Universe is a new weakly interacting massive particle that behaves as cold dark matter. Such dark matter candidates are found in different  extensions of the standard model (SM)~\\cite{Bertone:2004pz}. However, dark matter (DM) could be composed of more than one particle, and several multi component DM models have been suggested recently~\\cite{Zurek:2008qg, Profumo:2009tb, Feldman:2010wy, Winslow:2010nk, Batell:2010bp, Liu:2011aa, Adulpravitchai:2011ei}. Interest for multi component DM arose, in particular, from hints for signals in cosmic rays as well as in direct detection experiments corresponding to two completely different DM mass scales. On the one hand, the PAMELA experiment~\\cite{PAMELA} has reported an cosmic ray excess of positrons which would be consistent with a TeV scale DM candidate annihilating mainly into leptons. On the other hand, DAMA~\\cite{DAMA}, CoGeNT~\\cite{Cogent} and CRESST~\\cite{CRESST} have all reported excesses in the direct detection rate that would be compatible with DM around 10 GeV. Even though there is no conclusive evidence that these signals are due to DM, it is interesting to investigate multi component DM models to expand the range of DM models at a time when DM searches in indirect, direct and collider experiments are increasing their sensitivities. Furthermore, replacing the usual R-parity symmetry that guarantees the stability of the lightest R-parity odd particle by an enlarged symmetry group allows not only multiple DM candidates but also new freeze-out mechanisms. In particular, the semi-annihilation mechanism where two DM particles annihilate into another DM particle and a SM particle was proposed in Ref.~\\cite{semiannihilation}.   In this paper, we propose a new type of freeze-out mechanism, assisted freeze-out, within the framework of multi component DM models. In this new freeze-out mechanism, one DM candidate can be thermalized only through the assistance of the other stable particle. Thus, the decoupling of one DM particle from the thermal bath is influenced by the other DM particle. Consequently, the relic density of DM is solved by using two coupled Boltzmann equations for two stable fields. In the analysis of the right-handed sneutrino DM~\\cite{arXiv:1105.1652}, a similar situation has been already considered even though this model has only one stable particle. After setting the Boltzmann equations, we construct a simple model with two hidden DM sectors corresponding to two new $U(1)$ gauge symmetries, only one of which interacts with the SM sector. This is achieved through kinetic mixing of the new and standard gauge bosons. The DM particles are assumed to be Dirac fermions. We then show how both particles can contribute to the relic density of DM while only one can scatter elastically on nucleons. Although the direct detection rate tends to be rather high in this model, we show examples, where all constraints can be satisfied, including a case with a DM candidate around 10 GeV.  This paper is organized as follows. The basic set-up is presented in section 2. An explicit model is constructed in section 3 and the implications for the relic density of DM as well as for the direct detection rates on nucleons are studied. Section 4 contains our conclusions. ",
        "conclusions": "We have illustrated with a simple toy model containing two stable dark matter particles, how the assisted freeze-out mechanism worked and could reproduce the measured value for the relic density of dark matter. The main feature of this type of model is that only one of the DM particles is involved in direct detection searches while both contribute to the relic density. In particular, when the DM particle that interacts with SM particles is the subdominant DM component, it is possible to reconcile models with large elastic scattering rates on nuclei with the exclusion bounds of XENON100. Moreover, the lighter particle can be a light DM candidate around 10 GeV as indicated  by the DAMA, CoGeNT and CRESST results."
    },
    "1112/1112.3564_arXiv.txt": {
        "abstract": "Global frequentist fits to the CMSSM and NUHM1 using the {\\tt MasterCode} framework predicted $\\Mh \\simeq 119 \\gev$ in fits incorporating the \\gmt\\ constraint and $\\simeq 126 \\gev$ without it. Recent results by ATLAS and CMS could be compatible with a Standard Model-like Higgs boson around $\\Mh \\simeq 125 \\gev$. We use the previous {\\tt MasterCode} analysis to calculate the likelihood for a measurement of any nominal Higgs mass within the range of 115 to 130~GeV. Assuming a Higgs mass measurement at $\\Mh \\simeq 125 \\gev$, we display updated global likelihood contours in the $(m_0, m_{1/2})$ and other parameter planes of the CMSSM and NUHM1, and present updated likelihood functions for $\\mgl, \\msqr$, \\bmm\\ and the spin-independent dark matter cross section \\ssi. The implications of dropping \\gmt\\ from the fits are also discussed. We furthermore comment on a hypothetical measurement of $\\Mh \\simeq 119 \\gev$.   \\bigskip \\begin{center} {\\tt KCL-PH-TH/2011-40, LCTS/2011-21, CERN-PH-TH/2011-305, \\\\ DCPT/11/168, DESY 11-242, IPPP/11/84, FTPI-MINN-11/31, UMN-TH-3023/11, {arXiv:1112.3564 [hep-ph]}} \\end{center} \\vspace{2.0cm} ",
        "introduction": "\\label{sec:intro}  Taking into account the relevant experimental constraints, the CMSSM and NUHM1 predict that the lightest Higgs boson should have couplings similar to those of the Standard Model (SM) Higgs boson~\\cite{Ellis:2001qv,Ambrosanio:2001xb,mc7}, and that it should weigh no more than $\\sim 130 \\gev$~\\cite{mh,FeynHiggs,asbstev}. We recently reported the results of global frequentist fits within the CMSSM and NUHM1 to the first $\\sim 1$/fb of LHC data, also including precision electroweak and flavour measurements and the XENON100 upper limit on elastic spin-independent dark matter scattering~\\cite{mc7}, updating the results of previous global fits by ourselves~\\cite{mc1,mc2,mc3,mc35,mc4,mc5,mc6,mcweb} and others~\\cite{pre-LHC,post-LHC} (see also \\cite{LR}). The results reported in~\\cite{mc7} included likelihood contours in the $(m_0, m_{1/2})$, $(\\tan \\beta, m_{1/2})$ and $(\\MA, \\tan \\beta)$ planes of the CMSSM and NUHM1, as well as $\\Delta \\chi^2$ functions for $\\mgl$, \\bmm, $\\Mh, \\MA$ and sparticle production thresholds in $e^+ e^-$ annihilation.  Notable predictions of these global fits included $\\Mh = 119.1^{+3.4}_{-2.9} \\gev$ in the CMSSM and $\\Mh = 118.8^{+2.7}_{-1.1} \\gev$ in the NUHM1 (which should be combined with an estimated theory error $\\Delta \\Mh = \\pm 1.5 \\gev$). These two fits are based solely on the Higgs-{\\it independent} searches including the \\gmt\\ constraint, i.e., they do not rely on the existing limits from LEP~\\cite{Barate:2003sz,Schael:2006cr}, the Tevatron~\\cite{TevHiggs}, or the LHC~\\cite{ATLASHA,CMSHA}. These predictions increase to $\\Mh = 124.8^{+3.4}_{-10.5} \\gev$ in the CMSSM and $126.6^{+0.7}_{-1.9} \\gev$ in the NUHM1 if the \\gmt\\ constraint is dropped.  Subsequently, the ATLAS and CMS Collaborations have released their official combination of the searches for a SM Higgs boson with the first $\\sim 1$/fb of LHC luminosity at $E_{\\rm cm} = 7$~TeV~\\cite{ATLAS+CMS}. Impressively, the combination excludes a SM Higgs boson with a mass between 141 and 476~GeV. Most recently, the ATLAS and CMS Collaborations have presented preliminary updates of their results with $\\sim 5$/fb of data~\\cite{Dec13}. These results may be compatible with a SM-like Higgs boson around $\\Mh \\simeq 125 \\gev$, though CMS also report an excess at $\\Mh \\simeq 119 \\gev$ in the $ZZ^*$ channel. We recall that, for low values of $\\Mh$, the SM electroweak vacuum would be unstable~\\cite{unstable}, decaying into a state with Higgs vev $> 10^8 (10^{10}) \\gev$ if $\\Mh = 119 (125) \\gev$, and that a very plausible mechanism for stabilizing the vacuum is supersymmetry (SUSY)~\\cite{ER}.   In this paper, we first report the likelihood function for an LHC measurement of $\\Mh$ with a nominal value $\\in (115, 130) \\gev$, incorporating the theoretical error $\\pm 1.5 \\gev$ and an estimate $\\pm 1 \\gev$ of the possible experimental error. In both the CMSSM and NUHM1, this likelihood function is minimized for $\\Mh \\simeq 119 \\gev$ if \\gmt\\ is included, and is contained within the theoretical uncertainty range shown previously as a `red band'~\\cite{mc7}.  We then discuss the consequences of combining a measurement of $\\Mh \\simeq 125 \\gev$ (assuming that the current excess will be confirmed with more integrated luminosity) with our previous analysis~\\cite{mc7} of constraints on the CMSSM and NUHM1 including \\gmt.  We find that the best-fit values of $m_0$ and $m_{1/2}$ in the CMSSM and NUHM1 are moved to substantially higher values, especially in the case of $m_{1/2}$. We also update our results on the best-fit regions in the $(m_{1/2}, \\tb)$ and $(\\MA, \\tb)$ planes, where we find again the substantial increase in $m_{1/2}$, as compared with our pre-LHC $\\Mh$ results. We present the corresponding one-dimensional likelihood functions for the gluino mass $\\mgl$, an average right-handed squark mass $\\msqr$, the lighter scalar tau mass, $\\mstaue$, as well as in the $(\\mneu{1}, \\ssi)$ plane, where $\\mneu{1}$ is the mass of the lightest neutralino and \\ssi\\ is the spin-independent dark matter scattering cross section. As could be expected, we find larger values of $\\mgl, \\msqr, \\mneu{1}$ and $\\mstaue$ than in our pre-LHC $\\Mh$ fit, and smaller values of \\ssi, though \\bmm\\ is little affected.  Since $\\Mh \\simeq 125 \\gev$ is the value that was favoured in the CMSSM/NUHM1 fits omitting the \\gmt\\ constraint~\\cite{mc7}, we also show some results for fits where \\gmt\\ is dropped. In this case, we find that preferred regions of the $(m_0, m_{1/2})$ planes are localized at relatively high values, corresponding to relatively large sparticle masses. Correspondingly, the spin-independent dark matter scattering cross section \\ssi\\ would be relatively small in this case, though again there would be relatively little effect on \\bmm.  Finally, we show selected results for a hypothetical measurement of $\\Mh \\simeq 119 \\gev$. ",
        "conclusions": "The ATLAS and CMS searches for the Higgs boson have already excluded a very large range of masses, with the only remaining windows for a SM-like Higgs boson being in the ranges $\\Mh \\in (115.5, 127) \\gev$ or $> 600 \\gev$~\\cite{ATLAS+CMS,Dec13}. The latter range is disfavoured by precision electroweak data, so attention naturally focuses on the low-mass range. It may or not be a coincidence that this range includes the range $\\Mh \\lsim 130 \\gev$ accessible in simple supersymmetric models such as the CMSSM and NUHM1. Within this range, our previous global fits of these models including \\gmt\\ predicted $\\Mh \\sim 119 \\gev$ if the \\gmt\\ constraint was included in the fit, and $\\Mh \\sim 126 \\gev$ if \\gmt\\ was omitted~\\cite{mc7}. The latest ATLAS and CMS results display an interesting fluctuation at $\\Mh \\sim 125 \\gev$ (i.e.\\ close to the latter result from \\cite{mc7}) and we have combined a hypothetical measurement of $\\Mh = 125 \\gev$ with the global likelihood functions obtained in our previous fits~\\cite{mc7}.  As we have shown in this paper, this combination refines our previous predictions for the CMSSM and NUHM1 model parameters within global fits incorporating \\gmt. In particular, the combination prefers a range of larger values of $m_{1/2}$, resulting in larger values of $\\mgl$ and other sparticle masses being preferred, restricting the prospects for discovering supersymmetry at the LHC within these models. The predictions for \\ssi\\ are pushed to higher masses and lower cross sections, particularly in the CMSSM. There are also smaller changes in the predictions for other observables such as \\bmm\\ .  We have also shown the analogous CMSSM and NUHM1 fit results for a hypothetical measurement of $\\Mh \\simeq 125 \\gev$ if the \\gmt\\ constraint is omitted. In this case we find a stronger preference for larger values of $(m_0, m_{1/2})$, and correspondingly larger values of $\\tb$ and $\\MA$, as well as larger values of $\\mgl, \\msqr$, potentially lying beyond the reach of the LHC. We have also commented on the potential implications of a hypothetical Higgs discovery at $\\Mh \\simeq 119 \\gev$.  Time will soon tell where the LHC experiments are indeed discovering the Higgs boson. However, we have shown that $\\Mh = 125 \\gev$ is a possibility within the CMSSM and NUHM1, although it lies at the upper range of what is possible within the CMSSM or NUHM1, and might suggest reduced prospects for discovering these particular models of supersymmetry at the LHC. Alternatively, it could well be that one should look beyond the frameworks of the models discussed here.   \\subsubsection*{Note Added}  After acceptance of this paper for publication, we became aware of issues in the implementation of the {\\tt FeynHiggs} code and in the cold dark matter density calculation, which required extra sampling and reprocessing of the NUHM1 parameter space. We are grateful to Nazila~Mahmoudi and Azar Mustafayev for discussions on dark matter density calculations.  \\subsubsection*{Acknowledgements}  The work of O.B., K.J.D., J.E., J.M. and K.A.O. is supported in part by the London Centre for Terauniverse Studies (LCTS), using funding from the European Research Council via the Advanced Investigator Grant 267352. The work of S.H. is supported in part by CICYT (grant FPA 2010--22163-C02-01) and by the Spanish MICINN's Consolider-Ingenio 2010 Program under grant MultiDark CSD2009-00064. The work of K.A.O. is supported in part by DOE grant DE-FG02-94ER-40823 at the University of Minnesota."
    },
    "1112/1112.1561_arXiv.txt": {
        "abstract": " ",
        "introduction": "The Wide-field ASKAP Legacy L-band All-sky Blind surveY (WALLABY)\\footnote{Principal Investigators: Baerbel Koribalski and Lister Staveley-Smith. See www.atnf.csiro.au/research/WALLABY for more details about the survey.} (\\citet{Koribalski_2009}; Koribalski, B., Staveley-Smith, L. et~al., in preparation) is an ambitious project that aims to detect neutral hydrogen to a redshift of $z \\sim 0.26$, across $\\sim 70\\%$ of the sky. It is one of the two top ranked projects that will be carried out using the Australian SKA Pathfinder (ASKAP). WALLABY is possible because of ASKAP's unprecedented $\\sim 30$ sq. degree field-of-view, which is achieved using Phased Array Feeds (PAFs). WALLABY will use all 36 of ASKAP's antennae, but due to limitations on computing resources will only process the inner 30 antennae (with a maximum baseline of 2km) to image the sky with a 30$^{\\prime\\prime}$ synthesised beam and produce datacubes with voxels\\footnote{Voxels are often referred to as pixels when discussing a single channel of a datacube. Technically though these `pixels' are still voxels. For this reason the term voxel is used throughout instead of pixel to aid consistency.} that project to $\\sim$10$^{\\prime\\prime}$ on the sky. The high spatial resolution is complemented by an anticipated spectral resolution of 3.86 km s$^{-1}$. ASKAP spectral datacubes will therefore cover a large area of the sky to high resolution, which results in very large datacubes containing at least 2048 x 2048 x 16 384 voxels. WALLABY will consist of $\\sim$ 1200 of these large datacubes. The size and number of these datacubes renders manual source finding unfeasible. The performance of the automatic source finder used by WALLABY will determine how many (\\HI) galaxies are found by WALLABY.  The majority of source finders in existence use intensity thresholding to find sources. {\\sc SExtractor} \\citep{1996A&AS..117..393B}, {\\sc SFind} \\citep{2002AJ....123.1086H} and {\\sc Duchamp} \\citep{2008glv..book..343W,Duchamp2} are good examples of source finders based on intensity thresholding. Conceptually, intensity threshold source finders check every pixel (voxel) in an image (datacube) to see if the pixel (voxel) value is sufficiently extreme that it's unlikely to be noise. Once all of the source pixels (voxels) have been identified, they are combined into objects. The various intensity threshold source finders differ in how they estimate the noise, set a threshold for identifying source pixels (voxels), pre-process the image (datacube) to improve the source finder results and the manner in which they create objects from source pixels (voxels). All intensity threshold source finders share an inherent limitation though.  Consider an arbitrary source in a spectral datacube. Improved spatial and spectral resolutions result in the source occupying more voxels in the datacube. Dispersing the source's signal over more voxels means that it contributes less to the flux value of each voxel that it occupies. This makes it harder for an intensity threshold method to detect the source. Using a simple model this effect is illustrated in Figure \\ref{fig:VoxelSN}, where the maximum voxel $S/N$ of an object with an integrated $S/N$ of 5 is plotted for various asymmetries. The maximum voxel $S/N$ is calculated to be $S/N_{integrated} \\times \\beta / \\sqrt{n}$, where $\\beta$ describes the asymmetry of the object's flux distribution and $n$ is the number of voxels.  By overlaying the minimum expected size (in voxels) of WALLABY sources on Figure \\ref{fig:VoxelSN}, we can assess the impact of this inherent limitation on WALLABY. The neutral hydrogen detected in emission is warm, which gives it an intrinsic amount of dispersion. We will assume that any real WALLABY source extends over at least 3 channels. We will also assume that in every channel, a source occupies at least 3x3 voxels for a 30$^{\\prime\\prime}$ synthesised beam and 10$^{\\prime\\prime}$ voxels). Galaxies rarely lie at the middle of a voxel though, so a more realistic minimum is a grid of 4x4 or 5x5 voxels in every channel. It is expected that most galaxies detected by WALLABY will be unresolved or at most marginally resolved (Duffy, A., Meyer, M. \\& Staveley-Smith, L. 2011, in preparation), so we also consider a 7x7 grid of voxels in every channel to account for marginally resolved, off-centre galaxies. After multiplying by the minimum number of channels to obtain the expected minimum size of WALLABY galaxies in voxels, the minimum size for these different grids is overlaid in Figure \\ref{fig:VoxelSN}. This demonstrates that off-centre and/or marginally resolved galaxies will be difficult to detect with a basic implementation of an intensity thresholding source finder, unless the flux is asymmetrically distributed. Figure \\ref{fig:VoxelSN} also illustrates that this effect is amplified in 3-D datasets such as future WALLABY datacubes. For example, in a 2-D image the 7x7 and 5x5 vertical lines would approximately lie at the position of the 4x4 and 3x3 vertical lines.  \\begin{figure}[ht] \\begin{center} \\includegraphics[width=0.98\\linewidth]{VoxelSN_5.ps} \\caption{The maximum voxel $S/N$ of an object (with an integrated $S/N = 5$) plotted against the number of voxels comprising the object, n. The various lines correspond to different asymmetries in the distribution of the voxel's flux over the n voxels. The vertical lines (labelled) denote the minimum size of a point source extending over three channels and occupying 3x3, 4x4, 5x5 or 7x7 voxels in every channel.} \\label{fig:VoxelSN} \\end{center} \\end{figure}  This inherent limitation is compounded further by using a size-based rejection criterion to weed out false detections. If you only detect a few, unconnected voxels the source will be flagged as a false positive by a size-based rejection criterion. Off-centre and/or marginally resolved galaxies that are detected because of asymmetric flux distributions are most likely to be detected in the form of a few, unconnected voxels. This effect will also show up as an enhanced fracture rate of extended, well resolved sources.  The inherent limitation of intensity thresholding based source finders can be offset by using more aggressive intensity thresholds (i.e., lower intensity thresholds), but it often results in many false (source) detections. The solution to this problem is to run the source finder on the datacube multiple times with the datacube smoothed to a different scale each time. This is however a very inefficient solution to the problem. {\\sc Duchamp} provides multiple options for dealing with this inherent limitation: a secondary `growth' threshold, smoothing and a 3-dimensional wavelet reconstruction of the dataset.  As part of its design study, WALLABY has investigated novel methods of source detection as an alternative to multiple passes of an intensity threshold source finder. The goal of this investigation was to develop source detection methods that are optimised for large datacubes with high spatial and spectral resolution. The Characterised Noise \\HI ({\\sc CNHI}) source finder that I present here is one of the novel source detection methods that have been developed.  The rest of the paper is structured as follows. The conceptual framework for the {\\sc CNHI} source finder is presented in Section \\ref{section:concept}. The inherent limitations of the {\\sc CNHI} source finder are then discussed in Section \\ref{section:limits}. Next, the current implementation of the {\\sc CNHI} source finder is presented in Section \\ref{section:implement}. Finally, some example results are discussed in Section \\ref{section:examples} before finishing with a summary. ",
        "conclusions": "\\label{section:summary} WALLABY and other projects that will be carried out on next-generation radio telescopes ASKAP and MeerKAT herald the start of the data deluge era in radio astronomy. The sheer size of WALLABY datacubes necessitates automation of many tasks in the data reduction pipeline, that previously would have been carried out with some level of manual input by an astronomer. Complete automation of finding \\HI galaxies in spectral datacubes is one of the challenges that is actively being investigated by WALLABY.  The resolution and size of WALLABY observations poses a challenge for many existing automated source finders. This challenge arises from the underlying conceptual framework and algorithm that these source finders are based on, and not a flaw in the implementation. The {\\sc CNHI} source finder has been developed using a conceptual framework that can handle the large size of WALLABY datacubes and takes advantage of the resolution. Treating a datacube as a set of spectra (akin to an IFU observation), it attempts to find sources by looking for regions in each spectrum that do not look like noise. This is achieved using a novel implementation of matched filtering. Instead of using multiple filters that describe various types of sources, a single filter describing the noise is used. Sources are detected using this noise filter by identifying regions that do not look like noise.  The performance of the {\\sc CNHI} source finder was tested using the PS and ES datasets in \\citet{WP_2011}. Analysis of the {\\sc CNHI} source finder output demonstrated that a reasonable combination of completeness and refined reliability can be achieved. A refined completeness of $\\sim 80\\%$ and $\\sim 50\\%$ was achieved for the PS and ES datasets, respectively, with a refined reliability of $\\sim 95\\%$. The PS dataset is better than the $80\\%$ completeness would suggest though, because the {\\sc CNHI} source finder found $\\sim 95\\%$ of all PS objects with a maximum voxel flux $\\geq 5\\sigma$, with a refined reliability of $\\sim 95\\%$. This analysis also demonstrated that the {\\sc CNHI} source finder recovers a significant fraction of the source flux. The recovery fraction asymptotes towards $100\\%$ as the total flux increases. More aggressive thresholds (larger) result in a recovery fraction that asymptotes faster, and starts higher. This suggests an alternative use of the {\\sc CNHI} source finder as a source parametrisation tool, that is used in tandem with another source finder. Finally, the performance analysis demonstrated that {\\sc CNHI} detections of a source are an unbiassed representation of the source and its flux distribution. These results are very promising, and warrant further testing and refinement of the {\\sc CNHI} source finder.  There are three development goals for the {\\sc CNHI} source finder. Further development of the {\\sc CNHI} source finder will initially focus on incorporating multi-scale bundling. This will effectively achieve independently-scaled matched filtering in both the frequency and spatial dimensions. Additionally, the {\\sc CNHI} source finder will have a simple intensity thresholding test added to it. Incorporating an intensity thresholding test will make the {\\sc CNHI} source finder sensitive to sources occupying 3 or fewer channels. The final development goal is to incorporate fourier analysis, polynomial fitting and existing baseline structure removal techniques. Upon completing this next development cycle, the {\\sc CNHI} source finder will be tested and tweaked using the next round of ASKAP simulations and the \\HI Parkes All Sky Survey (HIPASS) \\citep{2000ASPC..217...50S} datacubes. The HIPASS datacubes have been selected because they have a well defined source catalogue, contain a mixture of resolved and unresolved sources, have known artifacts and calibration issues and there is a potential for the {\\sc CNHI} source finder to detect new sources."
    },
    "1112/1112.4805_arXiv.txt": {
        "abstract": "DAV stars, also  called ZZ Ceti variables, are  pulsating white dwarfs with atmospheres rich in H.  Asteroseismology of DAV stars can provide valuable clues about  the origin, structure and evolution  of DA white dwarfs. Recently,  a new DAV star, WD  J191643.83$+$393849.7, has been discovered  in  the field  of  the  \\emph{Kepler}  spacecraft.  It  is expected that further  monitoring of this star in  the next years will enable astronomers  to obtain  the best lightcurve  of a  pulsating DA white dwarf ever recorded,  and thus to know with unprecedented precision the hidden  details of the internal structure  of this star.  In this paper,  we perform a first asteroseismological  analysis of WD J191643.83$+$393849.7   on  the   basis  of   fully   evolutionary  DA white-dwarf  models.   Specifically,  we  employ  a  complete  set  of evolutionary  DA  white-dwarf  structures  covering a  wide  range  of effective  temperatures, stellar masses,  and H  envelope thicknesses. These models have  been obtained on the basis  of a complete treatment of the evolutionary history of progenitors stars.  We compute $g$-mode adiabatic pulsation  periods for this  set of models and  compare them with  the  pulsation periods  exhibited  by WD  J191643.83$+$393849.7. Based  on a tentative  estimation of  the mean  period spacing  of the star, we  find that  the stellar mass  should be  substantially large ($\\gtrsim  0.80  M_{\\odot}$), in  agreement with  the  spectroscopically derived stellar  mass.  Also, from  period-to-period fits we  find an asteroseismological   model   characterised   by   a   low effective temperature, rather high stellar mass and a thin H envelope. The possibility that this rather massive pulsating white dwarf can be further monitored with \\emph{Kepler} with a high degree of detail turns the star WD  J191643.83$+$393849.7 into a promising and unique object to study the physics of crystallization and carbon/oxygen phase diagrams at high densities. ",
        "introduction": "Pulsating DA (H-rich atmospheres) white  dwarfs, also known as ZZ Ceti or DAV variable stars, comprise the most numerous and studied class of degenerate pulsators.   Their instability strip, probably  a pure one, is centered at an effective  temperature of around $12\\,000$ K (Winget \\& Kepler  2008; Fontaine  \\& Brassard 2008;  Althaus et  al.  2010a). The  luminosity  variations  (of  up  to $0.30$  mag)  are  caused  by nonradial  $g$-mode  pulsations of  low  degree  ($\\ell  \\leq 2$)  and periods between $\\approx  70$ and $\\approx 1500$ s,  excited through a combination of  the $\\kappa-\\gamma$  mechanism acting in  the hydrogen partial ionization zone (Dolez \\&  Vauclair 1981; Winget et al.  1982) and the ``convective driving'' mechanism --- once a sufficiently thick outer convective zone has been developed (Brickhill 1991; Goldreich \\& Wu  1999).  ZZ Ceti  asteroseismology ---  the comparison  between the observed   periods  and   the  periods   computed  for   a   suite  of representative stellar models --- has the potential of disentangle the internal structure  of DA white dwarfs, allowing  to place constraints on the  stellar mass, the thickness  of the outer  envelopes, the core chemical  composition, weak  magnetic fields  and slow  rotation rates from the observed period patterns  (Winget \\& Kepler 2008; Fontaine \\& Brassard 2008; Althaus et al. 2010a).  A total of  148 ZZ Ceti stars are currently  known (Castanheira et al. 2010).  To this  list we must add the recently  discovered DAV star WD J191643.83$+$393849.7   (\\emph{Kepler}  ID   4552982   hereinafter  WD J1916$+$3938), a  ZZ Ceti located  in the \\emph{Kepler}  mission field and identified through ground-based times series photometry (Hermes et al.  2011; hereinafter HEA11).  This  star ($T_{\\rm eff}= 11\\, 129 \\pm 115$ K  ; $\\log g=  8.34 \\pm 0.06$) exhibits  low-amplitude luminosity variations  with periods between  $\\approx 800$  and $\\approx  1450$ s (see column 2 of Table \\ref{table1}).  It is expected that an extended monitoring  of this  object with  \\emph{Kepler} could  bring  the best lightcurve of a pulsating white dwarf ever recorded (HEA11), even with a better quality  than the lightcurves of the  brighter stars observed through  uninterrupted  ground-based  campaigns  of  the  Whole  Earth Telescope (WET; Nather et al. 1990).  In  order  for  white  dwarf  asteroseismology  to  provide  realistic constraints on the  internal structure of white dwarfs,  it is crucial the  use of stellar  models characterised  by consistent  and detailed chemical  profiles  to   accurately  assess  the  adiabatic  pulsation periods. This requirement has been tackled successfully in the case of the hot DOVs or GW Vir stars (see C\\'orsico et al. 2007a, 2007b, 2008, 2009). Very  recently, Romero et  al. (2011) have performed  the first asteroseismological analysis of a set of 44 ZZ Ceti stars based on the new generation  of fully evolutionary DA white  dwarf models presented in  Althaus  et  al.   (2010b).   These models  are  characterised  by realistic  chemical profiles  throughout the  star and  cover  a wide range of stellar  masses, thicknesses of the H  envelope and effective temperatures. The analysis of a large  number of ZZ Ceti stars has the potential  to characterise  the  global properties  of  the class,  in particular the thicknesses  of the H envelope and  the stellar masses. The study  of Romero et al.   (2011) revealed that DA  white dwarfs in the  solar  neighbourhood  could  harbor  a  broad  range  of  H-layer thickness.   On  the  other  hand,  this analysis  was  successful  in finding, for the first time, an univocal asteroseismological model for the archetypal  ZZ Ceti star G117$-$B15A,  one of the  targets of that analysis.  Motivated by  the exciting discovery  of the first pulsating  DA white dwarf in the field of the \\emph{Kepler} mission, and encouraged by the availability  of the most  complete set  of detailed  evolutionary and pulsation models for DAV stars up to date (Romero et al 2011), we present in this paper a   first  asteroseismological  analysis   of  WD J1916$+$3938.   We   are  well  aware   that  the  asteroseismological potential of this pulsating star will probably be fully exploited with further  observations  by  \\emph{Kepler},  which will  allow  to  have available a period spectrum substantially  richer than that we have at hand at present just  from ground-based photometry.  However, we think that a  first asteroseismological analysis  of this star based  on the seven  periods ---  a representative  number  of periods  for ZZ  Ceti standards  ---  currently  available  is worth  being  attempted.   By considering the  mean period spacing  exhibited by the star,  we found that  WD J1916$+$3938  should be  substantially massive,  in excellent agreement  with  the  spectroscopic  inference.  We  also  found  that period-to-period    fits   favour    an    asteroseismological   model characterised by a high stellar mass, a low effective temperature, and a rather thin H envelope.  The paper  is organized  as follows.  In  the next Section  we briefly describe the  computer codes and  stellar models employed.   The short Section \\ref{smd}  is devoted to  the derivation of  the spectroscopic mass of  WD J1916$+$3938.  We  make inferences about the  stellar mass from the  period spacing data in  Section \\ref{period-spacing}, taking into account  its dependence not  only with the  effective temperature and mass, but  instead also with the thickness of  the H envelope.  In Section \\ref{period-fit}  we extract  information of the  star through asteroseismological period fits.  Section \\ref{rotation} is devoted to infer  a rotation  period  of  the star.   Finally,  we enumerate  our results in the Conclusions (Section \\ref{conclu}).  \\begin{table} \\centering \\caption{The observed frequencies, periods, and amplitudes of WD J1916$+$3938 (columns 1 to 3), according to HEA11, and the periods of a phenomenological model (see Section \\ref{period-spacing}) with a fixed ($\\ell= 1$) period spacing of $38.54$ s (columns 4 to 6).} \\begin{tabular}{cccccc} \\hline \\hline $\\nu$ & $\\Pi_i^{\\rm obs}$  &  $A_i$   & $\\Delta k$ & $\\Pi_{\\rm mod}$ &  $\\Pi_i^{\\rm obs}-\\Pi_{\\rm mod}$ \\\\ $[\\mu$Hz$]$ & $[$ s $]$ & $[\\%]$    &             &   $[$ s $]$    &    $[$ s $]$            \\\\ \\hline 1213.7 &      823.9 &  0.38   & $2$  & $819.53$ & $ 4.37$ \\\\ 1198.9 &      834.1 &  0.32   & $5$  & $935.60$ & $-1.01$ \\\\ 1070.1 &      934.5 &  0.36   & ---   &   ---   &  ---   \\\\ 1032.1 &      968.9 &  0.44  & $6$  & $974.14$ & $-5.24$ \\\\ 918.3 &     1089.0 &  0.25  & $9$  & $1089.75$& $-0.75$ \\\\ 854.8 &     1169.9 &  0.23  & $11$ & $1166.82$& $ 3.08$ \\\\ 696.0 &     1436.7 &  0.24 & $18$ & $1436.58$& $ 0.12$ \\\\ \\hline \\hline \\end{tabular} \\label{table1} \\end{table} ",
        "conclusions": "\\label{conclu}  Tempted by  the exciting discovery  of the first pulsating  DA white dwarf in  the \\emph{Kepler} mission  field by HEA11, and encouraged by  the availability of  the most complete set  of detailed evolutionary and pulsation models for  DAV stars up to date (Romero et al.    2011),   in   this    paper   we   have   presented   a   first asteroseismological analysis of WD  J1916$+$3938.  This, even with the certainty that the full asteroseismological potential of this DAV star will be  exploited with further  observations by \\emph{Kepler}  in the next years.  We summarize our main findings below:  \\begin{itemize}  \\item[-]    We    derive    a    spectroscopic    mass    of    $M_*= 0.797-0.805\\ M_{\\odot}$ for WD J1916$+$3938  on the basis of the new DA white dwarf  evolutionary tracks of Romero et  al.  (2011).  This value is smaller than the estimate of HEA11, of $M_*= 0.82\\  M_{\\odot}$, based on  the older models of  Wood (1990). In our analysis,  we have taken into account the dependence of the tracks in the $T_{\\rm eff}-\\log g$ diagram with the thickness of the H envelope  (see Fig.  \\ref{gteff}).  \\item[-] By assuming that all the periods exhibited by WD J1916$+$3938 are normal  modes of  the star, and  supposing that all  the periods except one (at 834.1 s) are associated to $\\ell= 1$ modes, we derive an average period spacing of $\\overline{\\Delta \\Pi}_{\\rm obs}= 38.54 \\pm  0.29$ s.   On the  basis of  the asymptotic  theory  of stellar pulsations,  we estimate  a stellar  mass of  $M_*= 0.851  \\pm 0.031 M_{\\odot}$,  corresponding  to  a  thick  H  envelope. If  we take into  account the strong dependence of  the asymptotic  period spacing with  the thickness of the H envelope, we found an evident degeneracy of $\\Delta \\Pi_{\\ell}^{\\rm a}$ with $M_*$ and $M_{\\rm H}$. In virtue of this, we can conclude that $M_* \\gtrsim  0.837 M_{\\odot}$, but we cannot say anything about the thickness of the H envelope.  Finally, if  we ignore  the constraint  imposed  by the  effective temperature   as  derived   by  spectroscopy,   we  found   that  WD J1916$+$3938 should have a stellar mass $M_* \\gtrsim 0.77 M_{\\odot}$ in order for the star to be within the ZZ Ceti instability strip.  \\item[-]  By using the  individual pulsation  periods exhibited  by WD J1916$+$3938, and  taking into  account several assumptions  for their $\\ell$- and $k$-identification, we found an asteroseismological model  with $T_{\\rm  eff}= 11\\,380$ K,  $\\log g= 8.39$,  $M_*= 0.837  M_{\\odot}$, $\\log(M_{\\rm  H}/M_*)=  -7.36$, and $\\log(M_{\\rm   He}/M_*)=   -2.50$.    The   stellar  mass   of   the asteroseismological   model   is   in   good  agreement   with   the spectroscopic  estimate ($M_*\\approx  0.8 M_{\\odot}$)  and  with the derivation   from   the   period   spacing   ($M_*\\gtrsim   0.837 M_{\\odot}$). This is quite  encouraging, because the three estimates have been derived from very different approaches and assumptions. On the other hand, the effective temperature of the asteroseismological model is somewhat higher than  the spectroscopic one  ($T_{\\rm eff}= 11\\,129 \\pm  115$ K).  \\item[-] In passing,  we have estimated a rotation  period of 18.77 hs for WD  J1916$+$3938.  The reliability  of this result rests  on the assumption that the modes with periods  at 823.9 s and 834.1 are the $m=  +1$ and  $m=  -1$ components,  respectively, of  a $\\ell= 1$ frequency triplet due to rotation.  \\end{itemize}  In closing, we  think that the conclusion that  WD J1916$+$3938 should be a  rather massive ZZ Ceti  star, with $M_*  \\gtrsim 0.8 M_{\\odot}$, appears to  be robust. The possibility that this rather massive pulsating white dwarf can be further monitored with \\emph{Kepler} with a high degree of detail turns this star into a promising and unique target to study the physics of crystallization and carbon/oxygen phase diagrams at high densities. In this regard, we  note that, for  a stellar mass of  $M_* \\sim 0.88  M_{\\odot}$, crystallization of  matter should occur at  the centre of  a white dwarf star  at $T_{\\rm eff}  \\approx  10\\,500$  K according to the  phase  diagram  of Segretain \\& Chabrier (1993), and at $T_{\\rm eff} \\approx 9\\,600$ K if the new phase  diagram recently proposed by   Horowitz et  al.  (2010) is considered  (Althaus et al. 2011). Indeed,  based on direct molecular dynamics simulations for dense carbon-oxygen mixtures, Horowitz et al. (2010) find substantially lower crystallization temperatures and that the shape of the carbon-oxygen phase diagram is of the azeotropic form --- and not of the spindle type --- as previously believed. Our analysis  from the period  spacing of  WD J1916$+$3938 suggests that  this star  could well have  a stellar mass  much higher than $0.8  M_{\\odot}$ and  necessarily a rather  thin H  envelope.  If this  were  the case,  this  star  could have  a  substantial fraction  of its  interior  crystallized depending on the adopted carbon-oxygen phase  diagram, an issue that could be explored with future asteroseismological analysis  that exploit the new high-quality photometric   data  coming   from  the   ongoing  monitoring   by  the \\emph{Kepler} mission."
    },
    "1112/1112.3108_arXiv.txt": {
        "abstract": "We have catalogued and analysed cosmological parameter determinations and their error bars published between the years 1990 and 2010. Our study focuses on the popularity of measurements, their precision and their accuracy. The accuracy of past measurements is gauged by comparison with the most recent WMAP results of Komatsu et al. (2011). The 637 measurements in our study are of 12 different parameters and we place the techniques used to carry them out into 12 different categories. We find that the popularity of parameter measurements (published measurements per year) in all 12 cases except for the dark energy equation of  state parameter $w_0$ peaked between 1995 and 2004. Of the individual techniques, only Baryon Oscillation measurements were still rising in popularity at the end of the studied time period. The quoted precision (fractional error) of most measurements has been declining relatively slowly, with several parameters, such as the amplitude of mass fluctutations $\\sigma_{8}$ and the Hubble constant $H_{0}$ remaining close to the $10\\%$ precision level  for a 10-15 year period. The accuracy of recent parameter measurements is generally what would be expected given the quoted error bars, although before the year 2000, the accuracy was significantly worse, consistent with an average underestimate of the error bars by a factor of $\\sim 2$. When used as complement to traditional forecasting techniques, our results suggest that future measurements of parameters such as fNL, and $w_{a}$ will have been informed by the gradual improvment in understanding and treatment of systematic errors and are likely to be accurate. However, care must be taken to avoid the effects of confirmation bias, which may be affecting recent measurements of dark energy parameters. For example, of the 28 measurements of $\\Omega_{\\Lambda}$ in our sample published since 2003, only 2 are more than 1 $\\sigma$ from the WMAP results. Wider use of blind analyses in cosmology could help to avoid this. ",
        "introduction": "Modern cosmological parameters have been measured since Hubble's (1929) discovery of the expansion of the Universe. The number of model parameters increased during the late 1980s with the introduction of what is often referred to as the ``Standard cosmological model'' (e.g. Dodelson 2005). The idea of ``Precision cosmology'' emerged more recently, and by the present time,  many of the parameters in this model are well known (see e.g., Komatsu \\etal 2011, hereafter WMAP7). This presents us with an interesting opportunity: by comparing the past measurements of parameters and their error bars with the currently known values, we can evaluate how well the measurements were carried out in the past, how realistic the quoted uncertainties were, and which methods gave the most statistically reliable results. We can also study how both their precision and accuracy has varied with time. Such research will help us in our quest to make critical evaluations of what will be possible in the future, and by working with past data serves as a complement to more conventional future extrapolations of technology and techniques (e.g., the report of the Dark Energy Task Force, hereafter DETF, Albrecht \\etal, 2006). In the present paper we make a first attempt at such a study, by compiling published  parameter values taken from the NASA Astrophysics Data System over the years 1990-2010.  Previous studies of cosmological parameter determinations have tended to focus on the Hubble Constant, $H_{0}$, for which there is a longer than 80 year baseline for analysis. Several papers have used  the comprehensive database compiled by John Huchra\\footnote{https://www.cfa.harvard.edu/~dfabricant/huchra/} to generate their dataset, such as the study of the non-Gaussian error distribution in those measurements (Chen \\etal 2003). Gott \\etal (2001)  used median statistics in a metanalysis of these $H_{0}$  measurements to find the most probable value (and also analysed early measurements of $\\Omega_{\\Lambda}$). This median statistics approach has also been used to combine individual estimates of $\\Omega_{m0}$, the present mean mass density in non-relativistic matter by Chen and Ratra (2003). In the present paper we do not seek to combine the measurements from various works into best determinations of parameters. Instead we start from the assumption that the parameters we look at have been well measured (and their correct values are close to the WMAP7 values) and see what this implies about past measurements.  We therefore will be starting with the assumption that $\\Lambda$CDM is the correct cosmological model. This should be borne in mind when interpreting our results. Even if the true cosmology turns out in the future to be something else, we expect that the effective values of the $\\Lambda$CDM parameters are not likely to be very different (given the good fits to current data), so that our approach will have some value even then. Parameters which at the moment are unknown, or very poorly constrained, such as the non-Gaussianity parameter fNL (e.g., Slosar \\etal 2008), or the time derivatives of the dark energy equation of state parameter $w$ (e.g., Chevallier \\& Polarski 2001) can obviously not be studied at present with our approach. Instead we hope that the general lessons from the past about the reliability of error bars, methods and achievable precision and accuracy can usefully to inform future efforts to measure those parameters.  The DETF report explains how four different techniques are being used and will be used in the future to constrain dark energy parameters. These techniques, gravitational lensing, baryon oscillations, galaxy cluster surveys and supernova surveys all have a history and have been involved in a large number of previous measurements of different parameters. It is interesting to see how they have performed in the past, and evaluate  them based on this data. By looking over the published record, we can also show how measurement precision has changed, in terms of the quoted fractional error bars, and see how this compares with predicted future trends. One can ask whether for example the earlier error bars were unrealistically small, so that the quoted precision of measurements has not changed much. This should have consquences for the accuracy of measurements, which we will define and measure. In general, our motivation for this study can be summarized by the idea that once cosmological parameter measurements are published, for the most part they are ignored when future work arrives. The dataset left behind can instead become a valuable resource to inform future work.  Our plan for the paper is as follows. In Section 2., we detail the source for the cosmogical parameter estimates and how the data was collated. We explain the different categorizations of measurements and methods and standardization that was carried out. In Section 3 we outline the steps involved in our analysis of the data, and present results including historical trends in some individual parameters and the precision and accuracy of measurements. In Section 4 we summarize our findings and discuss our results. ",
        "conclusions": "\\subsection{Summary} We have compiled cosmological parameter measurements published between 1990 and 2010 and the techniques used to measure them. Using this data we have carried out an analysis of historical trends in popularity, precision and accuracy. The accuracy of past measurements has been estimated by assuming that WMAP7 parameter values of  Komatsu \\etal (2011) (combined with $\\Lambda$CDM standard values for e.g. $w_{0}$) are the correct ones. Our findings can be summarised as follows:  (1) The number of published measurements for different parameters peaks between 1995 and 2004 for all cases, except for $w_{0}$ for which the number was still rising in 2010.  (2) Of all techniques used to measure the parameters, only baryon oscillation and ``combined'' measurements were still rising in terms of publications per year by 2010.  (3) The quoted precision of measurements has been declining relatively slowly for most parameters, with several (e.g. $\\sigma_{8}$, $H_{0}$ remaining flat for 10-15 years.  (4) The accuracy of recent parameter measurements is generally what should be expected based on the quoted error bars i.e. the error bars overall are neither understimated nor overestimated (an accuracy, $N_{\\sigma}=1.0$, within the Poisson uncertainty on the measurement). Before 2000, the accuracy $N_{\\sigma}$ as closer to 2, indicating underestimation of the error bars by a factor of 2. Overall, there is a small non-Gaussian tail to the error distributions (we find that $20\\%$ of measurements are more that $2\\sigma$ away from the true values.  (5) The accuracy of most methods has become consistent with $N_{\\sigma}=1.0$, with the historically most innaccurate parameter measurement technique being the use of galaxy peculiar velocities. Measurements of $\\Omega_{M}$ and particularly $\\Omega_{\\Lambda}$ made since 2000 tend to have accuracy $N_{\\sigma}$ significantly less than 1.0, indicating ``confirmation bias'' and/or an overestimation of error bar sizes.    \\subsection{Discussion} Over the 20 year period covered in this study, it is apparent that many of the parameters in what is now the concordance CDM cosmological model went from the status of no information or only limits to being known at the $10\\%$ level or better. It is also apparent from Figure 2 that there was a ``golden age'' of parameter measurements between $\\sim 1995$ and $\\sim 2005$ during which the number of published measurements peaked sharply and then declined. This seems to indicate that for many purposes (such as the use of a background cosmology in galaxy formation models), the precision to which the $\\Lambda$CDM parameters were known by the time of the first WMAP results is sufficient, and many of the reasons for pinning down the model better had diminished after that.  This said, however, the exception to this rule, measurements of $w_{0}$ (which are still rising in terms of number per year at the end of our study) seems to point to a coming new era in parameter measurement. Certainly, the motivation for the large number of ongoing and future large-scale structure, lensing and other surveys is to hunt for the signatures of dynamical dark energy and modified gravity, and given the number of researchers carrying out these studies it is likely that measurements will continue to rise. Many parameters which we have not catalogued are now within reach of quantitiative study. These include the modified gravity parameter, $E_{G}$ (Reyes \\etal 2010) and the time derivative of the equation of state parameter, $w_{a}$. Measurements of such parameters involve searching for deviations from the concordance CDM model and fall into a different category from most of the parameters we have studied in this paper.  Inflationary parameters such as the non-Gaussianity fNL, or tensor to scalar ratio $r$ will pinned down with higher precision in the future, and these should also represent a growth area. The motivation for most future measurements being largely framed in terms of a quest for fundamental physics, it would be logical to assume that they will continue until the cause for the Universe's acceleration are better understood. Likewise, parameters describing the dark matter particle should be added to this category.  Possible behaviours for the precision of future parameter measurements can be predicted by looking at the past results (Figure \\ref{prepa}). There is a very wide range, but most parameters improve slowly, with a factor of 10 improvement in precision over the 20 years representing the extreme (2 out of 12 parameters). The precision of some parameters has remained relatively flat for the whole period, so this is a possibility for future so far unconstrained parameters. An argument against this slow progress however is the fact that many new surveys (such as of Baryon Oscillations) are targeted primarily at measures of specific parameters, and this aggressive approach (for example including specific precisions to be obtained at a given time in survey proposal documents) could lead to faster progress.  Our investigation of the accuracy of results could potentially lead to some of the most interesting findings. We have seen that in the earlier half of our studied time period there is evidence that the error bars were significantly underestimated, but that this has changed over time.  When discussing the accuracy, we are should be aware  that it was not possible in our analysis to take into account several factors which have the potential to affect our conclusions. For example, we do not keep track of the priors that people have assumed in their measurements, and in many of the later cases, this may include the WMAP results as priors. That this is happening is likely to be responsible for much of the post WMAP1 tightening of constraints seen in Figures \\ref{h0}- \\ref{w0}. When computing the error bars on the mean accuracies of measurements (Figs \\ref{accp} and \\ref{accm}) we have used Poisson errors based on the number of measurements in each bin. This will tend to underestimate the uncertainty on the accuracy because some of those measurements could be using the same underlying data, or be using similar priors, or a combination of the two. There will therefore be correlations between the error bars so that our estimates of the accuracy will be affected. Equivalently the chi-square of the fiducial result compared to the data points will be incorrectly determined to be low because of the correlations are not included.   Bearing the above points in mind, we return to the panels in Figures \\ref{accp} and \\ref{accm} where the accuracy seems to be significantly below $N_{\\sigma}=1$. This is most obvious in the second panel ($\\Omega_{\\Lambda}$) of Figure \\ref{accp}. Such as result could be a sign that either the error bars have been significantly overestimated, or else that researchers have been influenced by prior results (``confirmation bias''), or a combination of the two. If we return to the data points which led to the last two bins of panel two of Figure \\ref{accp}, we find something especially striking. Of those 28 measurements, only 2 are more than $1 \\sigma$ from the fiducial results of Table 1. These 28 measurements were carried out by approximately 11 separate groups (as determined by authorship lists) using several different techniques.  This closeness of published results to the ``correct'' ones is somewhat worrying for future measurements. One can interpret this as coming partly from  error bars being overerestimated by cautious cosmologists, for example by including possible systematic errors in the error bars which are not actually present to such a large degree, or in a related point  authors marginalizing over parameters which are actually better known than was assumed. We note that including or excluding the actually quoted systematic error bars (Section 3.4) has little effect on this result. An additional question is why some parameters have $N_{\\sigma} <1$ and others do not (e.g., $\\sigma_{8}$). The relatively low number statistics of our whole dataset preclude us from making any strong statements about this issue. If it does partly result from confirmation bias, one can also wonder how observers knew which value of $\\Omega_{\\Lambda}$ (for example) would be the ``correct'' one, given that our fiducial (mostly WMAP7) results from Table 1 were published in 2011. If this bias is present, it is probably related to the mean level for  $\\Omega_{\\Lambda}$ resulting from several prior measurements. For example in Figure \\ref{lambda} and others, the value of the parameter seems to be pretty well determined at least by 2003.  If we look at the techniques which are often associated with dark energy measurements, SN and BAO, we can see in Figure \\ref{accm} that these two have low $N_{\\sigma}$ for recent measurements. Of the 23 measurements which where included in the last bins of the SN panel of Figure \\ref{accm}, only 2 are more than 1 $\\sigma$ from the fiducial result. We note  that this fiducial result from Table 1 does include BAO measurements, but not SN estimates of dark energy. In the case of SN, however, only 4 measurements of $\\Omega_{\\Lambda}$ are included in these bins, and only 2 separate groups of researchers, so that for that subset of data, statistical fluctuations may well be responsible for the low $N_{\\sigma}$ seen. If confirmation bias is present, on the other hand, one could argue about who is confirming who- certainly the first SN results on dark energy predate those from BAO and from most other techniques. These sorts of questions might be addressed by a more detailed look at the published measurements, including details of priors, jointly used datasets and analysis techniques. Then again small number statistics probably would not allow firm conclusions to be drawn. These hints should instead serve as a warning that care and perhaps concrete steps be taken to avoid any confirmation bias in the future.  In conclusion, we have seen that huge progress has been made in the 20 year period covered by our study. Important questions have been resolved (e.g., is the Universe open?, do massive neutrinos contribute substantially to the dark matter density?), a model has been found which agrees with essentially all observational data  so far ($\\Lambda$CDM), and the parameters of that model have been pinned down at the $1-10\\%$ level. The first WMAP results (e.g. as presented in Spergel \\etal (2003)) form a watershed which is easy to pick up in most plots of parameters with time, and serves as a reminder that statistically measurable progress is not always gradual. Perhaps the most interesting aspect of our study, of the accuracy of past results compared to our most recent knowledge has found that understanding of systematic errors and uncertainties in cosmological measurements has demonstrably improved since the early 1990s. On average, results in the last 10 years are consistent with expectations, given their error bars, something which should instill confidence in future measurements. There are some signs that recent measurements of dark energy parameters are closer to the ``expected'' values for $\\Lambda$CDM than statistically likely. These may be explainable by correlations between measurements which we have not included. On the other hand this may serve as a sign that as cosmology collaboration sizes increase carrying out more blind analyses (as in particle physics) may be a good idea."
    },
    "1112/1112.2612_arXiv.txt": {
        "abstract": "We have used the Robert C.\\ Byrd Green Bank Telescope to time nine previously known pulsars without published timing solutions in the globular clusters M62, NGC~6544, and NGC~6624.  We have full timing solutions that measure the spin, astrometric, and (where applicable) binary parameters for six of these pulsars.  The remaining three pulsars (reported here for the first time) were not detected enough to establish solutions.  We also report our timing solutions for five pulsars with previously published solutions, and find good agreement with past authors, except for PSR J1701$-$3006B in M62. Gas in this system is probably responsible for the discrepancy in orbital parameters, and we have been able to measure a change in the orbital period over the course of our observations.  Among the pulsars with new solutions we find several binary pulsars with very low mass companions (members of the so-called ``black widow'' class) and we are able to place constraints on the mass-to-light ratio in two clusters.  We confirm that one of the pulsars in NGC~6624 is indeed a member of the rare class of non-recycled pulsars found in globular clusters.  We also have measured the orbital precession and Shapiro delay for a relativistic binary in NGC~6544.  If we assume that the orbital precession can be described entirely by general relativity, which is likely, we are able to measure the total system mass ($2.57190(73)\\; \\Msun$) and companion mass ($1.2064(20)\\; \\Msun$), from which we derive the orbital inclination ($\\sin{i} = 0.9956(14)$) and the pulsar mass ($1.3655(21)\\; \\Msun$), the most precise such measurement ever obtained for a millisecond pulsar. The companion is the most massive known around a fully recycled pulsar. ",
        "introduction": "}  Millisecond pulsars (MSPs) form by ``recycling'' a dormant neutron star through the accretion of mass and angular momentum from a binary companion \\citep{acr+82}.  This leads to a very rapidly rotating and stable pulsar with a relatively low magnetic field ($\\sim 10^9\\; \\gauss$) and long lifetime ($\\gtrsim 10^9\\; \\yr$).  MSPs form naturally in the dense environments of globular clusters (GCs), thanks to frequent exchange interactions that may lead to the formation of mass transferring binaries \\citep{cr05}.  Sensitive searches of GCs have uncovered 144 pulsars in 28 clusters\\footnote{See \\url{http://www.naic.edu/~pfreire/GCpsr.html} for an up-to-date list.}, and the vast majority of these are recycled MSPs.  Indeed, nearly half of all MSPs have been discovered in GCs\\footnote{There are 160 field pulsars with period $P < 20\\; \\ms$ in the ATNF catalog, and 131 GC pulsars that meet the same criteria (see \\url{http://www.atnf.csiro.au/people/pulsar/psrcat/}).}.  The same interactions that form MSPs so efficiently in clusters also lead to many exotic systems that are rarely seen in the disk of the Galaxy.  These include the fastest spinning MSP \\citep{hrs+06}, highly eccentric binaries \\citep{rhs+05,frg07}, massive neutron stars \\citep{frb+08}, pulsar-main sequence binaries \\citep{dpm+01}, and many ``black widow'' systems \\citep{kbr+05}.  These discoveries demonstrate the huge scientific payoff that can come from the discovery of unique pulsars thanks to two factors---the extreme nature of neutron stars, which opens windows on otherwise inaccessible realms of physics, and the extraordinary clock-like nature of MSPs.  The bedrock of pulsar astronomy is \\emph{timing}, the process of creating a model that unambiguously accounts for every rotation of the pulsar, thus probing the pulsar and its environment.  For MSPs, the arrival time of a pulse can typically be measured to within a few microseconds or better \\citep[e.g.][]{vbv+08}, which enables very precise timing models. Timing GC pulsars leads to unique challenges, though.  Because GCs are usually at a distance of several kiloparsecs, the flux density of their constituent MSPs is usually very low, necessitating very long integration times.  Another consequence of cluster distances is very high dispersion measures, which necessitate moving to higher observing frequencies (e.g. $2\\; \\GHz$) where pulsars tend to be weaker. Luckily, since many clusters contain several MSPs that can each be observed during a single observation, the required observing time is well spent.  In addition to recycled MSPs, a small population of slow, non-recycled pulsars have also been observed in a handful of clusters \\citep{lmd96,blt+11}.  Their presence is something of a mystery, because they resemble in every way the ``normal'' pulsars that are so numerous in the Galactic disk, but which have lifetimes $\\sim 10^7$--$10^8\\; \\yr$ and are thought to form through core collapse supernovae of massive stars.  However, GCs are old stellar systems with typical ages $10^{10}\\; \\yr$ \\citep{cgcf00}, so all stars massive enough to form pulsars should have died some $10^{10}\\; \\yr$ ago.  As such, the core collapse of massive stars cannot be the avenue through which non-recycled pulsars in GCs have formed.  The most popular alternative formation scenario involves the collapse of a massive white dwarf via electron capture \\citep{nom84,nom87}, though the details of these so-called electron capture supernovae are not well understood.  In this paper, we present new timing solutions for nine pulsars spread across three GCs---M62, NGC~6544, and NGC~6624.  Each of these clusters contains multiple MSPs and NGC~6624 contains two non-recycled pulsars.  All of these pulsars were discovered elsewhere, but full, phase-connected timing solutions have only been published for five of them (see Table \\ref{table:psrs} for a summary and list of references).  We also present our solutions for these five.  We began a timing campaign using the Robert C.\\ Byrd Green Bank Telescope (GBT) for the nine pulsars without full solutions, starting in 2009 February.  We have managed to obtain full solutions that include measurement of the first period derivative for six of these pulsars. A combination of low flux densities and an irreversible data processing error conspired to prevent us from obtaining full solutions for the two remaining MSPs, but we are in a position to comment on them in further detail.  Our most exciting result is the discovery of a massive companion orbiting a fully recycled MSP.  In \\S\\ref{sec:obs} we describe our observations and in \\S\\ref{sec:timing} outline our method for timing the pulsars and further data analysis.  A discussion of individual systems can be found in \\S\\ref{sec:psrs}, and we provide a summary in \\S\\ref{sec:conc}. ",
        "conclusions": "}  We set out to obtain full timing solutions for nine previously known GC pulsars in M62, NGC~6544, and NGC~6624 (three of which have never been published).  Our efforts were successful for six of these pulsars, and a seventh now has a partial solution.  We have confirmed the nature of three black widow pulsars (M62E and F, and NGC~6544A). NGC~6544A has $\\dot{P} < 0$, providing unambiguous evidence that $\\dot{P}$ is dominated by acceleration in the cluster potential, and providing us with a probe of $\\Upsilon\\rmsub{V}$ in the cluster core. We find $\\Upsilon\\rmsub{V} \\gtrsim 0.072$ which, while not very constraining, is consistent with small amounts of low-luminosity matter such as massive black holes or other stellar remnants.  We are also able to confirm that NGC~6624C belongs to a rare class of non-recycled GC pulsars similar to the slow pulsars found in the Galactic disk.  The highlight of our work is unquestionably NGC~6544B.  This binary MSP is in a highly eccentric orbit and exhibits clear orbital precession and Shapiro delay.  Under the well justified assumption that this is due almost entirely to GR, we are able to obtain a most-probable total system mass of $2.57190(73)\\; \\Msun$, companion mass of $1.2064(20)\\; \\Msun$, and pulsar mass of $1.3655(21)\\; \\Msun$. This is the highest mass companion known to orbit a fully recycled MSP, and raises the possibility that NGC~6544B is part of either double neutron star system, which could form via exchange interactions in the GC.  We are deeply grateful to Andrea Possenti and Alessandro Corongiu for providing us with data collected by the Parkes Telescope that was used to check for phase connection in our solutions for M62D, E, and F.  We would also like to thank an anonymous referee for helpful comments. R.\\ Lynch acknowledges support from the GBT Student Support program and the National Science Foundation grant AST-0907967 during the course of this work.  The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc."
    },
    "1112/1112.0567_arXiv.txt": {
        "abstract": "Weak-gravitational-lensing distortions to the intensity pattern of 21-cm radiation from the dark ages can be decomposed geometrically into curl and curl-free components.  Lensing by primordial gravitational waves induces a curl component, while the contribution from lensing by density fluctuations is strongly suppressed. Angular fluctuations in the 21-cm background extend to very small angular scales, and measurements at different frequencies probe different shells in redshift space.  There is thus a huge trove of information with which to reconstruct the curl component of the lensing field, allowing tensor-to-scalar ratios conceivably as small as $r\\sim 10^{-9}$---far smaller than those currently accessible---to be probed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.0084_arXiv.txt": {
        "abstract": "We discuss the effects of cosmic phase transition on the spectrum of primordial gravitational waves generated during inflation.  The energy density of the scalar condensation responsible for the phase transition may become sizable at the epoch of phase transition, which significantly affects the evolution of the universe.  As a result, the amplitudes of the gravitational waves at high frequency modes are suppressed.  Thus the gravitational wave spectrum can be a probe of phase transition in the early universe. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.0745_arXiv.txt": {
        "abstract": "The cross-correlation of the Ly-$\\alpha$ forest and redshifted 21-cm emission has recently been proposed as an observational tool for mapping out the large-scale structures in the post-reionization era $z \\le 6$.  This has a significant advantage as the problems of continuum subtraction and foreground removal are expected to be considerably less severe in comparison to the respective auto-correlation signals.  Further, the effect of discrete quasar sampling is considerably less severe for the cross-correlation in comparison to the Ly-$\\alpha$ forest auto-correlation signal.  In this paper we explore the possibility of using the cross-correlation signal to detect the baryon acoustic oscillation (BAO). To this end, we have developed a theoretical formalism to calculate the expected cross-correlation signal and its variance.  We have used this to predict the expected signal, and estimate the range of observational parameters where a detection is possible.  For the Ly-$\\alpha$ forest, we have considered BOSS and BIGBOSS which are expected have a quasar density of $16 \\ {\\rm deg}^{-2}$ and $64 \\ {\\rm deg}^{-2}$ respectively. A  radio interferometric array that covers the redshift range $z=2$ to $3$ using antennas of size $2 \\, {\\rm m} \\times 2 \\, {\\rm m}$, which have a $20^{\\circ} \\times 20^{\\circ}$ field of view, is well suited for the redshifted 21-cm observations. It is  required  to observe $25$ independent fields of view, which corresponds to the entire angular extent of BOSS. We find that  it is necessary to  achieve a noise level of $1.1 \\times 10^{-5} \\ {\\rm mK}^2$ and $6.25 \\times 10^{-6} \\ {\\rm mK}^2$ per field of view  in the redshifted 21-cm observations to detect the angular and radial BAO respectively with BOSS. The corresponding figures are $3.3 \\times 10^{-5} \\ {\\rm mK}^2$ and $1.7 \\times 10^{-5} \\ {\\rm mK}^2$ for BIGBOSS. We also discuss possible observational strategies for detecting the BAO signal.   Four to five independent radio interferometric arrays, each containing $400$ antennas uniformly sampling all the baselines within $50 \\, {\\rm m}$ will be able to carry out these observations in the span of a few years. ",
        "introduction": "Neutral hydrogen (HI) in the post-reionization epoch ($z < 6$) is known to be an important cosmological probe seen in both emission and absorption.  Here, the redshifted 21-cm emission \\citep{hirev1, hirev3} and the transmitted QSO flux through the Ly-$\\alpha$ forest \\citep{wien, Mandel} are both of utmost observational interest. In a recent paper \\citet{tgs5} have proposed the cross-correlation of the 21-cm signal with the Lyman-$\\alpha$ forest as a new probe of the post-reionization era. While it is true that the emission and the absorption signals both originate from neutral hydrogen (HI) at the same redshift (or epoch), these two signals, however, do not originate from the same set of astrophysical sources. The 21-cm emission originates from the HI housed in the Damped Lyman-$\\alpha$ systems (DLAs) which are known to contain the bulk of the HI at low redshifts \\citep{proch05}. The collective emission from the individual clouds appears as a diffuse background in low frequency radio observations \\citep{poreion2}. On the contrary, the Ly-$\\alpha$ forest consists of a large number of Ly-$\\alpha$ absorption lines seen in the spectra of distant background quasars. These absorption features arises due to small density fluctuations in the predominantly ionized diffuse IGM. On large scales, however, the fluctuations in the 21-cm signal and the the Ly-$\\alpha$ forest transmitted flux are both believed to be excellent tracers of the underlying dark matter distribution. \\citet{tgs5} have proposed that the cross-correlation between these two signals can be used to probe the power spectrum during the post-reionization era.  The HI power spectrum can be determined separately from observations  of the Ly-$\\alpha$ forest \\citep{pspec} and the redshifted 21-cm emission \\citep{poreion1}. The detection of the individual signals, however, face severe observational challenges. The Ly-$\\alpha$ auto-correlation power spectrum is expected to be dominated by the Poisson noise arising from the discrete sampling of QSO lines of sight. The cross-correlation signal has the advantage that it is not affected by the Poisson noise which affects only its variance. Uncertainties in fitting the QSO continuum \\citep{pspec4, kim04, bolton} pose another challenge for using Ly-$\\alpha$ observations to determine the power spectrum.  Astronomical sources like the galactic synchrotron emission, extra-galactic point sources, etc. appear as foregrounds for the redshifted 21-cm signal. These foregrounds, which are several orders of magnitude larger than the signal, pose a great challenge for detecting the post-reionization 21-cm signal \\citep{fg10,fg11}. The foregrounds and systematics in the Ly-$\\alpha$ forest and the redshifted 21-cm emission  are, however expected to be uncorrelated and we therefore expect the problem to be much less severe for the cross-correlation signal. A detection of a cross-correlation signal shall, hence, conclusively ascertain its cosmological origin.  Apart from being an independent probe of the large scale matter distribution, the cross-correlation signal can potentially unveil the same astrophysical and cosmological information as the individual auto-correlations.  Cosmological density perturbations drive acoustic waves in the primordial baryon-photon plasma which are frozen once recombination takes place at $ z \\sim 1000$, leaving a distinct oscillatory signature on the CMBR anisotropy power spectrum \\citep{peeb70}. The sound horizon at recombination  sets a standard ruler that maybe used to calibrate cosmological distances.  Baryons contribute to $15 \\% $ of the total matter density,  and the baryon  acoustic oscillations  are imprinted in the late time clustering of non-relativistic matter. The signal, here, is however suppressed by a factor $ \\sim \\Omega_b/\\Omega_m \\sim 0.1$,  unlike the  CMBR temperature anisotropies where it is an order unity effect \\citep{komatsu}. The baryon acoustic oscillation (BAO) is a powerful  probe of cosmological parameters \\citep{seoeisen,white}. This is particularly useful since the effect occurs on large scales ($\\sim 150 \\, \\rm Mpc$), where the fluctuations are still in the linear regime. It is possible to measure the angular diameter distance and the Hubble parameter as functions of redshift using the the transverse and the longitudinal oscillations respectively. These provide   means for estimating cosmological parameters and placing  stringent constraints on dark energy models. Nonlinear effect  of gravitational clustering  tend to wipe out the BAO signal, and it is preferable to avoid very low redshifts where this is a potential problem. However, very high redshifts too are not very useful for constraining dark energy models. Several authors have reported a $2$ to $3 \\sigma$ detection of the BAO in  low redshift galaxy surveys \\citep{baoeisen05, percival07, baosdss}. The possibility of detecting the BAO signal in the Ly-$\\alpha$ forest has  been extensively studied by \\citet{cosparam1}. Several groups have also considered the   possibility of detecting the BAO signal  using  redshift 21-cm emission \\citep{pen2008, maowu, masui10}.  Several QSO surveys are now being considered  with the intent of measuring the BAO using the Ly-$\\alpha$ forest (eg. BOSS \\citep{mcd1}  and BIGBOSS \\citep{bigboss}). The possibility of a  wide field redshifted 21-cm survey to detect the BAO is also under serious consideration. In this paper, we consider  the possibility of studying the BAO using the cross-correlation signal.  In Section 2. of this paper we quantify the cross-correlation between the Ly-$\\alpha$ forest and the 21-cm emission, and  present theoretical  predictions of  the expected signal. We have used  the multi-frequency  angular power spectrum $C_{\\ell}(\\Delta z)$ (MAPS,  \\citealt{datta1})   in preference to the more commonly used three dimensional power spectrum $P(k)$ to quantify the cross-correlation signal. This has several advantages which we briefly discuss here. MAPS  refers to the  angular multipole $\\ell$ (or  equivalently angle) and redshift interval $\\Delta z$ which are the directly relevant observational quantities for both  Ly-$\\alpha$ surveys and redshifted 21-cm observations. This is particularly important if we wish  to determine the  angular  scales and redshift intervals  which  need to be covered in order to detect a particular feature in the signal.  The foregrounds in the redshifted 21-cm observations and the continuum in the Ly-$\\alpha$ forest are both expected to have a smooth, slow variation along the frequency  axis, and this plays a crucial role in removing these from the respective data. It is therefore advantageous to use MAPS which retains the distinction between the frequency (redshift)  and angular information, unlike $P(k)$ which mixes these up.  \\citet{fg3} have analyzed GMRT observations using  MAPS to jointly characterize the angular and frequency dependence of the foregrounds at $150 \\, {\\rm MHz}$.  \\citet{fg10} and \\citet{fg11} have applied  MAPS to Analyze $610 \\, {\\rm MHZ}$  GMRT observations, and have used this to characterize the foregrounds for the  post-reionization 21-cm signal. In fact, they show that it is possible  to completely remove the foregrounds from the  measured $C_{\\ell}(\\Delta \\nu)$ by subtracting out polynomials in $\\Delta \\nu$. Finally, the signal could have a significant contribution from the  light cone effect, particularly if the observations span a large redshift interval. It is, in principle, possible to account for this in the MAPS, though we have not done this here. It is, however, not possible to account for this effect in the three dimensional power spectrum which mixes up the information from different epochs through a Fourier transform along  the radial direction.  In Section 3 of this paper we quantify the imprint of the BAO feature on the cross-correlation signal. Here MAPS has the  advantage that it allows us to separately study the radial and the transverse oscillations through the $\\Delta \\nu$ and the $\\ell$ dependence respectively. In Section 4 we introduce an estimator for the MAPS of the cross-correlation signal and derive its statistical properties. In particular, we present a detailed analysis of the noise for the cross-correlation estimator. In Section 5 we present several observational considerations which are relevant for the cross-correlation signal, some   pertaining to the Ly-$\\alpha$ forest and others to the redshifted 21-cm signal or both. Finally, in  Section 6  we discuss the detectability of the cross-correlation signal and the BAO features. We have used the cosmological parameters $(\\Omega_m h^2, \\Omega_b h^2, \\Omega_{\\Lambda}, h, n_s, \\sigma_{8})= (0.136, 0.023, 0.726, 0.71, 0.97, 0.83)$  from \\citet{komatsu} throughout this paper. ",
        "conclusions": "In this paper we have developed a theoretical formalism for  estimating the cross-correlation signal between the fluctuations in the Ly-$\\alpha$ forest and the fluctuations in the redshifted 21-cm emission from neutral hydrogen. Both of these quantities are measured as functions of  frequency (redshift) and  angular scale. Consequently, we have  used the Multi-frequency Angular Power Spectrum (MAPS) to quantify the statistical properties of the cross-correlation signal. This deals directly with the observed quantities, and  retains the distinction between the angular and the frequency information. This, as we shall elaborate shortly, is very important in the light of foreground removal and continuum subtraction.  Continuum fitting and subtraction is a very critical step in calculating $\\delta_{\\F}$ for the Ly-$\\alpha$  forest  and several different methods have been proposed  for handling this  \\citep{pspec4, mcd06}.  Errors in continuum subtraction can be a serious problem for  the Ly-$\\alpha$  forest auto-correlation signal \\citep{kim04}. The redshifted 21-cm signal is buried deep under astrophysical foregrounds which are several orders of magnitude larger \\citep{shaver99,fg4,fg1,fg6, fg3, pen09, fg10}.  Several different techniques have been proposed for separating the cosmological 21-cm signal from the foregrounds  \\citep{ fg3, jelic, fg7, fg8, fg10}. Foreground removal is a rather severe  problem for observations of the redshifted 21-cm auto-correlation signal.  In addition to the 21-cm signal, the foregrounds make a very large contribution to the expectation value of the auto-correlation estimator. However, the foregrounds are believed to have a slowly varying, smooth frequency  (or $\\Delta z$) dependence which is  quite distinct from the signal which decorrelates rapidly with increasing $\\Delta z$ (Figure \\ref{fig:kappa1}).  It is therefore possible to remove  the foregrounds from the measured MAPS ($P_{\\F T o}(\\ell,\\Delta z)$ ) by subtracting  out any component that varies slowly with $\\Delta z$. In fact, \\citet{fg11} have recently  used MAPS to analyze 610-MHZ GMRT  observations and show that it is possible to remove the foregrounds from  the m auto-correlation by fitting and subtracting out slowly varying polynomials in $\\Delta \\nu$.  We do not expect the continuum  in the Ly-$\\alpha$ forest to have any correlation with the foregrounds in the redshifted 21-cm observations,  and  consequently they will  not contribute to the expected cross-correlation signal. The continuum and the foregrounds will, however, appear as extra contributions to the variance of the cross-correlation estimator. Since these contributions appear in the variance, it is possible to reduce these by  combining different independent estimates of the cross-correlation.  We can  reduce the continuum and foreground contributions  by combining  estimates of the cross-correlation at different baselines $\\U$ and different fields of view. The problem, therefore, is much less severe in comparison  to the auto-correlation. Further, the continuum and the foregrounds are both expected  to have a slowly varying frequency  (or $\\Delta z$) dependence and it should be possible to remove these from the measured MAPS by  subtracting  out the  component that varies slowly with $\\Delta z$.  We plan to perform a detailed analysis of these issues in future.  Our study shows that it is possible to have a  $5-\\sigma$ detection of the imprint of the first BAO peak in the cross-correlation signal   using BOSS, an upcoming QSO survey. For this, we have considered a radio interferometric array that covers the $z$ range $z=2$ to $3$ using antennas of size $2 \\, {\\rm m} \\times 2 \\, {\\rm m}$ which have a $20^{\\circ} \\times  20^{\\circ}$ field of view. We find that  we need to observe $25$ fields of view, approximately the full angular coverage of BOSS. with a noise level of $N_T=1.1 \\times 10^{-5} \\, {\\rm mK}^2$  and $N_T=  6.25 \\times 10^{-6} \\, {\\rm mK}^2$ in order to achieve  a $5-\\sigma$ detection of the angular  and radial oscillations respectively. The corresponding noise levels are $N_T=3.3 \\times 10^{-5} \\, {\\rm mK}^2$  and $N_T=  1.7 \\times 10^{-5} \\, {\\rm mK}^2$ for BIGBOSS whose quasar density is expected to be four times larger than that of BOSS.  We now briefly discuss how it may be possible to carry out such observations. In our analysis we have made the simplifying assumption that the antennas are distributed such that all the baselines within $50 \\, {\\rm m}$ are uniformly sampled. We consider an interferometric array with $N_{ann}=400$ antennas which  roughly corresponds to the maximum number of $2 \\, {\\rm m} \\times 2 \\, {\\rm m}$ antennas that can fit in a $50 \\, {\\rm m} \\times 50 \\, {\\rm m}$ region.  We see that we need $\\Delta t = 5700 \\, {\\rm hr}$ and $10,000 \\, {\\rm hr}$ of observation per field of view (\\ref{eq:oc1})  to reach  the  noise levels  required to detect the angular and radial oscillations respectively with BOSS. The corresponding figures are  $1900 \\, {\\rm hr}$ and $3700 \\, {\\rm hr}$ for the BIGBOSS.  Note that we it is required to observe a single field of view for $8 \\, {\\rm hr}$ a day for a whole year in order to achieve $3000 \\, {hrs}$ of  observing time.  It is quite evident that we require to observe $25$ fields of view, with $2$ to $3$ years of dedicated observations for each field, in  in order to detect the BAO. Operating sequentially, considering one field after the next, the required observations would possibly run over a period of $50 \\, {\\rm yr}$ to a century, which raises the need to consider alternative observational strategies. It is a viable  possibility  to  have antennas that  can simultaneously observe several independent  fields of view.  However, it is unlikely (if not impossible) to have antennas that can simultaneously observe $16$ or $25$  such $20^{\\circ} \\times 20^{\\circ}$  fields of view.   For the purpose of this discussion, we assume that we have antennas that can simultaneously observe $4$ fields of view. We then see that it would approximately require observations over a decade (or more) in order to detect the BAO. Another possibility is to have several radio interferometric arrays, each located at a different  location and observing a different parts of the sky. Four to five separate arrays, each with $400$  antennas, would be required to carry out these observations in the span of a few years. It is important to note that it may be possible to reduce the observational requirements to some extent by optimally distributing the baselines instead of considering them to be uniformly distributed. We propose to investigate  these issues   in a future study.    Observations of the BAO can be used to constrain the values of various cosmological parameters. The equation of state of the Dark Energy is particularly important in this context. In this paper we have mainly  estimated the range of observational parameters for which  it  will be possible to detect the  BAO using  the cross-correlation signal. We plan to study a variety of issues including the optimal array configuration  and cosmological parameter estimation in future."
    },
    "1112/1112.2789_arXiv.txt": {
        "abstract": "{ The majority of VO-compatible spectra handling applications operates only with a few spectra entirely downloaded from single or several SSAP servers. We try to identify the scientific cases which could immediately benefit from future SSAP applications designed for GRID deployment. Their key feature is the sophisticated spectra pre-selection and preprocessing done on distributed servers using the intelligent agent summarising the results and performing final high-level processing or analysis. ",
        "introduction": "There is a wealth of calibrated astronomical spectra accessible in current data archives, however most of them are not suitable for direct physical analysis. Being pre-reduced by some automatic pipeline or individually by manual reduction, they  are mostly stored in archives in form of  FITS file or ASCII table as the relation of intensity in arbitrary numbers (instrumental counts, digital numbers etc.) and wavelength (or frequency).  The  scientific analysis of such  spectra requires further processing by the variety of different methods. In certain studies a huge number of spectra has to be collected from different servers (e.g. in different spectral regions ) and transformed into  common units.  The clear definition of metadata description and easy unit conversion as well as transparent  lookup and download of spectra in Virtual Observatory has a great potential to become a new advanced way of astronomical research. The connection of powerful distributed computing available by GRID, Web Services and VO protocols can establish an innovative  research environment allowing the data mining  of huge datasets. ",
        "conclusions": "Astronomical spectroscopy uses a wide range of techniques with different level of complexity to achieve its final goal --- to  estimate the most precise and reliable  information about celestial objects. The  large part of spectroscopic analysis today  has been accomplished by several independent non VO-compatible legacy packages, where  each works  with different local files in its own data format. Analysis of large number of spectra is thus very tedious work requiring good data bookkeeping.  Accomplishing the analysis in VO infrastructure may benefit from automatic aggregation of distributed archive resources (e.g. the multispectral research), seamless on-the-fly data conversion, common interoperability of all tools  (using PLASTIC protocol) and powerful graphical visualisation of measured and derived quantities.  The currently available VO clients supporting the SSA protocol can provide with only basic functions in a interactive environment.  The more advanced work with spectra (e.g. automatic preprocessing or advanced analysis) is not supported in VO at all.  The most of the legacy applications should be turned into the VO server side services and the conversion of legacy scripts and recipes into Workflows will allow their easy deployment on GRIDs.  Modern techniques of astronomical analysis require the considerable amount of computing power and very offten the iterative comparison with theoretical models is indispensable.  The properly designed VO service getting both observed and synthetic spectra transparently from VO resources and processing the number of objects in parallel on GRID might become a killer application turning the attention of wide astronomical community to the Virtual observatory as a viable and innovative way of modern astronomical research.  By introduction of  modern  VO-aware tools into the  astronomical spectral analysis  a remarkable increase of effectiveness of astronomical research can be achieved."
    },
    "1112/1112.2740_arXiv.txt": {
        "abstract": "{ Studies have suggested that there is a strong correlation between the masses of nuclear star clusters (NSCs) and their host galaxies, a correlation which said to be an extension of the well-known correlations between supermassive black holes (SMBHs) and their host galaxies. But careful analysis of disk galaxies -- including 2D bulge/disk/bar decompositions -- shows that while SMBHs correlate with the stellar mass of the \\textit{bulge} component of galaxies, the masses of NSCs correlate much better with the \\textit{total} galaxy stellar mass. In addition, the mass ratio $\\mnsc/\\mtot$ for NSCs in spirals (at least those with Hubble types Sc and later) is typically an order of magnitude smaller than the mass ratio $\\mbh/\\mbulge$ of SMBHs.  The absence of a universal ``central massive object'' correlation argues against common formation and growth mechanisms for both SMBHs and NSCs. We also discuss evidence for a break in the NSC--host galaxy correlation: galaxies with Hubble types earlier than Sbc appear to host systematically more massive NSCs than do types Sc and later. ",
        "introduction": "As far as we can tell, all massive galaxies in the local universe harbor supermassive black holes (SMBHs, with masses $\\mbh \\sim 10^{6}$--$10^{9} \\Msun$). The masses of these SMBHs correlate strongly with several global properties of the host galaxies, particularly with the central velocity dispersion \\sigzero{} \\citep{ferrarese00,gebhardt00} and with the bulge luminosity or mass \\citep[e.g.,][]{marconi03,haring04}. These correlations imply that the processes which drive galaxy growth and the processes which drive black hole growth are intimately linked -- perhaps even the \\textit{same} processes.  It is now also clear that many galaxies, particularly later-type spirals, host luminous nuclear star clusters \\citep[NSCs; e.g.,][]{carollo97,boker02}, with masses in the range $10^{5}$--$10^{8} \\Msun$; see the review by \\citet{boker08} for more details. Recently, several authors have argued that NSCs and central SMBHs have the \\textit{same} host-galaxy correlations: in particular, that SMBHs and NSCs have the same correlation with bulge luminosity and mass \\citep{wehner06,ferrarese06,cote06,rossa06} (but see \\citet{balcells07}).  The suggestion, then, is that NSCs and SMBHs are in a sense members of the same family of ``Central Massive Objects'' (CMOs), and thus that they may have grown via the same mechanisms \\citep[e.g.,][]{mclaughlin06,li07,nayakshin09,devecchi10}.  We argue, however, that one should be cautious about assuming that NSCs and SMBHs are really part of the same family, with the same host-galaxy relationships. To begin with, the samples of \\citet{wehner06} and \\citet{ferrarese06}, which were used to make the CMO argument, were almost entirely early-type galaxies -- mostly ellipticals and dwarf ellipticals. These are galaxies which are, in essence, ``pure bulge'' systems, so one could just as easily argue for a correlation with \\textit{total} galaxy mass.  But we know that SMBHs in \\textit{disk} galaxies correlate better with just the bulge, and \\textit{not} with the total galaxy mass or light \\citep[e.g.,][]{kormendy-gebhardt01,kormendy11}. Given that there have been previous claims that NSCs in spiral galaxies correlate with the \\textit{total} galaxy light \\citep[e.g.,][]{carollo98}, we are prompted the ask the question: do nuclear clusters \\textit{in disk galaxies} correlate with the bulge (like SMBHs), or with the whole galaxy? ",
        "conclusions": ""
    },
    "1112/1112.5356_arXiv.txt": {
        "abstract": "The most extended, closed magnetic loops inferred on T Tauri stars confine hot, X-ray emitting plasma at distances from the stellar surface beyond the the X-ray bright corona and closed large-scale field, distances comparable to the corotation radius. Mechanical equilibrium models have shown that dense condensations, or ``slingshot prominences'', can rise to great heights due to their density and temperatures cooler than their environs. On T Tauri stars, however, we detect plasma at temperatures hotter than the ambient coronal temperature. By previous model results, these loops should not reach the inferred heights of tens of stellar radii where they likely no longer have the support of the external field against magnetic tension. In this work, we consider the effects of a stellar wind and show that indeed, hot loops that are negatively buoyant can attain a mechanical equilibrium at heights above the typical extent of the closed corona and the corotation radius. ",
        "introduction": "Observations of T Tauri Stars (TTS) have revealed the presence of magnetic structures much larger than those seen on our Sun. Indeed, analysis of these structures has invoked solar flare-based models: to infer the length scales of confining loops, \\citet{Reale:1997} developed the uniform cooling loop (UCL) model, linking solar flare X-ray decay slopes to spatially resolved lengths of the loops confining the emitting gas. The UCL model was applied to data from the Chandra Orion Ultradeep Project \\citep[COUP,][]{Getman:2005}, and it was discovered that evidently the post-flare loops of TTS can reach up to multiple stellar radii in extent \\citep[][results summarised in Figure \\ref{F-f05loops}]{Favata:2005}. This result is surprising: T Tauri coronae are typically believed to be compact, as evidenced by the common detection of rotationally modulated coronal X-ray emission \\citep{Flaccomio:2005}. X-ray bright regions are generally thought to be confined within the corona, with loops generally not reaching more than a stellar radius in height.  In the mechanical equilibrium model of \\citet{Jardine:2005}, prominences form after reconnection occurs in the open field, and it is a straightforward matter to show that cool, dense loops can reach heights well beyond corotation (the radius at which the orbital velocity equals the stellar rotational velocity) and the closed stellar coronal field. Hot loops such as those inferred by the COUP, however, appear to need the support of the stellar field, and thus should be limited to heights within the closed coronal magnetic field: a fraction of the heights inferred by UCL models. It was initially suggested that perhaps these hot, large-scale loops are anchored to circumstellar disc material for stability. Testing this idea, followup studies of the COUP objects that searched for discs were unable to find conclusive evidence for dusty disc material within reach of the loops \\citep[e.g.,][]{Getman:2008,Aarnio:2010}. Furthermore, the apparent anti-correlation of large-scale loops and close-in disc material suggests it is the absence of inner disc material that allows the loops to reach greater extents than they would were close-in disc material present \\citep[as suggested by][close-in circumstellar discs could strip away these stars' outer coronae]{Jardine:2006}. Potentially, this indicates some missing physics in earlier models, and it raises the question of how such large loops remain stable for the duration of the X-ray flare decay phase (in some cases, multiple rotation periods).  \\begin{figure} \\begin{center} \\includegraphics[width=85mm,clip=true,trim=30.0 0 10. 20.]{f05loops.eps} \\end{center} \\caption{A summary of the results of \\citet{Favata:2005}: inferred loop height as a function of measured loop temperature. Vertical lines connect the loop height to that star's corotation radius. We have scaled the figure for ease of viewing; the tallest loop (its vertical corotation line shown) is $\\sim$63 R$_{*}$ in height. Filled points denote cases in which any parameters in calculation of the corotation radius were assumed (i.e., any values of R$_{*}$, M$_{*}$, or P$_{rot}$ not drawn from surveys of the ONC; fiducial values were used as discussed in section \\ref{SS-fiducial}). The grey horizontal line at 3 R$_{*}$ indicates the location of the source surface (the extent of the closed coronal magnetic field) in our models. There is a pileup of points at 270 MK because, to be conservative, the authors adopted a maximum best-fit temperature and applied it to all cases above a given threshold (beyond this temperature, the fits to their X-ray data are statistically equivalent, and a higher temperature simply yields a shorter loop).} \\vspace{-15pt} \\label{F-f05loops} \\end{figure}  The potential ramifications of highly energetic events on stellar and circumstellar evolution are far-reaching: large-scale magnetic structures, in their formation and eventual disruption, could regulate stellar angular momentum evolution as well as circumstellar disc evolution, and thus planet formation. The combination of the high frequency with which large-scale magnetic structures have been inferred and the great heights at which they form above the stellar surface indicates circumstances are ideal for the generation of substantial torques opposing stellar rotation. Large-scale magnetic reconnection, in addition to shedding mass and angular momentum, also bombards circumstellar material with high energy photons and particles. This could potentially play a pivotal role in planet formation by stimulating grain growth \\citep[cf. flash-heated chondrule formation,][]{Miura:2007}. High energy radiation and particle fluxes can influence disc chemistry \\citep[cf., solar nebula isotopic abundances,][]{Feigelson:2002} and structure: \\citet{Glassgold:1997} showed that hard stellar X-rays could potentially ionise the disc sufficiently for accretion to occur via magnetohydrodynamic instability. The large heights inferred for these loops would put X-ray emitting plasma closer to the disc, reducing the column depth through which the hard X-rays would have to travel before ionising disc material.  Given these potential impacts of large scale magnetic loops (and their destabilization) on early star and disc evolution, we aim to ascertain whether the existence of hot, large, magnetic loops, as inferred from the analysis of the brightest, most energetic events in X-ray surveys of TTS, is physically plausible. To do this, we address an evident buoyancy defecit in the previous models by including a stellar wind in the model physics. We do not distinguish between the single loop UCL model or the possibility of multiple loop events; we only seek to demonstrate that it is possible to find hot loops beyond the typical ``quiescent'' X-ray emitting corona, and in some cases, beyond corotation. ",
        "conclusions": "\\label{S-discus}  We have addressed the issue of mechanical equilibrium for large, hot loops. Since cases have been found in which hot loops several stellar radii in height seem to not be disc-supported, we propose that these loops reach mechanical equilibria beyond the closed coronal magnetic field due to the stellar wind. We have shown that, for a generic T Tauri star, our equilibrium loop heights are consistent with the inferred loop lengths for the most powerful flares observed by the COUP. We also recover solutions for hot loops within the range of heights below the source surface anticipated by previous models, suggesting that large and small stable post-flare loops can occur under similar conditions. In the physics of previous models, which were hydrostatic and did not consider the role of stellar winds, loops hotter than the quiescent corona were confined to heights within the closed coronal field, needing its support. Our present work indicates that, given a stellar wind, hot loops can be found at heights well beyond the source surface, and even beyond corotation.  Young stars are known to be very magnetically active in the T Tauri phase, with frequent flaring and high X-ray luminosities. The apparent abundance of large scale loops, favourable conditions for their formation in TTS coronae and winds, and potential for consequence on pre-main sequence stellar (and circumstellar) evolution makes understanding these large-scale loops very important. Having verified that such large-scale magnetic structures can indeed exist, interesting subsequent work will be to determine how the loops' disruption can impact circumstellar disc material as well as the star itself. Large-scale disruptions of this nature could effect stellar rotational evolution and activity as well as circumstellar disc evolution and planet formation."
    },
    "1112/1112.5483_arXiv.txt": {
        "abstract": "We investigate non-spherical behavior of gas accreting onto a central supermassive black hole. Assuming optically thin conditions, we include radiative cooling and radiative heating by the central X-ray source. Our simulations are performed using the 3D SPH code GADGET-3 and are compared to theoretical predictions as well as to 1D simulations performed using the grid code ZEUS. As found in earlier 1D studies, our 3D simulations show that the accretion mode depends on the X-ray luminosity ($L_X$) for a fixed density at infinity and accretion efficiency. In the low $L_X$ limit, gas accretes in a stable, spherically symmetric fashion. In the high $L_X$ limit, the inner gas is significantly heated up and expands, reducing the central mass inflow rate. The expanding gas can turn into a strong enough outflow capable of expelling most of the gas at larger radii. For some intermediate $L_X$, the accretion flow becomes unstable developing prominent non-spherical features. Our detailed analysis and tests show that the key reason for this unstable non-spherical nature of the flow is thermal instability (TI). Small perturbations of the initially spherically symmetric accretion flow that is heated by the intermediate $L_X$ quickly grow to form cold and dense clumps surrounded by overheated low density regions. The cold clumps continue their inward motion forming filamentary structures; while the hot infalling gas slows down because of buoyancy and can even start outflowing through the channels in between the filaments. We measured various local and global properties of our solutions. In particular, we found that the ratio between the mass inflow rates of the cold and hot gas is a dynamical quantity depending on several factors: time, spatial location, and $L_X$; % and ranges between $0$ and $4$. We briefly discuss astrophysical implications of such TI-driven fragmentation of accreting gas on the formation of clouds in narrow and broad line regions of AGN,  the formation of stars, and the observed variability of the AGN luminiosity. ",
        "introduction": "\\label{sec-intro}  Active galaxies are believed to be powered by accretion of matter onto the central supermassive black holes (SMBHs) \\citep[e.g.,][]{Salpeter64, LyndenBell69, Blandford76, Ozernoi78, Balick82, Rees84, Fabian90, Kormendy95, Kauffmann00, Ferrarese05}. Active Galactic Nuclei (AGN) are argued to influence cosmological galaxy formation in the form of feedback, whereby the overall properties of a galaxy can be regulated by its central BH \\citep[e.g.,][]{Silk98, King03, Wyithe03, Granato04, Murray05, Begelman05, Best07, Ciotti07, Pipino09, Ostriker10, DiMatteo11}. Observationally AGN are characterized as a compact continuum source around a central SMBH, surrounded by a much larger emission line region consisting of multiphase gas clouds. The continuum emission shows a    % relatively flat spectral energy distribution over a broad wave-band from radio to gamma rays, and is modeled as coming from the central engine composed of a gas accretion disk of size $10^{-4} - 10^{-3}$ pc. The broad-line region (BLR) extends up to $0.01 - 1$ pc and is composed of higher-density, higher-velocity gas; while the outer narrow-line regions (NLR) are considered to be located at distances up to $10 - 1000$ pc from the central SMBH.  Modeling gas accretion onto SMBHs and resulting AGN feedback in a cosmological context is computationally challenging, because a large dynamic range of length scales is involved: BH accretion on sub-pc scale, to galaxy physics on hundreds of kpc scale. Therefore SMBH accretion and feedback is usually treated by a sub-grid model in cosmological simulations \\citep[e.g.,][]{DiMatteo08, Booth09}. At the same time, gas accretion flow within the Bondi radius of a SMBH has just started to be resolved in observations \\citep{Wong11}. Hence numerical studies resolving the Bondi radius are required for a meaningful comparison, and it would be       % useful to improve a sub-grid model of AGN feedback for cosmological simulations based on the results of small-scale simulations. In \\citet[][hereafter Paper~I]{Barai11}, we studied spherically-symmetric Bondi accretion of gas onto a SMBH,     % and started to explore the effects of radiative heating and cooling.   A number of recent studies resolving the Bondi radius have focused on accretion onto intermediate-mass BHs regulated by feedback from X-ray and UV radiation emitted isotropically near the BH. Analytical work \\citep{Milosavljevic09b} as well as 1D and 2D simulations \\citep{Milosavljevic09, Park11, Park12} find that photoionization heating and radiation pressure significantly reduce the steady-state accretion rate to a fraction of the Eddington rate, or over 2 orders of magnitude below the Bondi rate. However, in 1D simulations by \\citet{Li11} the self-gravity of the gas eventually overcomes the radiative feedback effects, and boosts the accretion to the Eddington rate after one free-fall timescale. \\citet{Park11, Park12} also find the development of Rayleigh-Taylor instabilities in their 2D simulations. In these work the BH luminosity is made to be proportional to the rate at which gas accretes across the inner boundary, whereas in our study the central luminosity has a constant value in each simulation because we want to have a simplified setup with a non-variable source to enable us cleanly investigate the physical processes influencing the accretion flow.  In Paper~I we discovered indications of thermal instability (TI) in some runs, which we investigate in detail in the current work. Studying TI in a dynamical flow in 3D is challenging, because the effects of a gravitational field significantly complicate the development of TI \\citep[e.g.,][]{Balbus89}.  The foundations for understanding astrophysical thermal instability were laid in the classical paper by \\citet{Field65}. The general criterion for TI to occur is: \\begin{equation} \\label{eq-TI-criterion} \\left [ \\frac{\\partial {\\cal{L}}} {\\partial S} \\right ]_A < 0, \\end{equation} where ${\\cal{L}}$ is the net cooling-heating rate, $S$ is the specific entropy, and $A$ is a thermodynamic variable that is kept constant for a given perturbation. Subsequently the theory of TI has been studied over decades, including analytical developments: local TI in static and dynamical systems \\citep[e.g.,][]{Balbus86, Balbus89, Balbus95}, global TI by linear perturbation analysis \\citep[e.g.,][]{Yamada01}, and numerical hydrodynamical calculations to investigate the perturbations in the non-linear regime \\citep[e.g.,][]{Hattori90}.   TI has been invoked to explain observed features in various astrophysical domains: solar prominences \\citep[e.g.,][]{Parker53, Karpen88}; formation of cold gas clouds and clumpy substructure in the interstellar medium \\citep[e.g.,][]{Schwarz72, Parravano87, Hennebelle99, Burkert00, Koyama02, Sanchez-Salcedo02, Inutsuka05}; origin of globular clusters during the collapse of a protogalaxy \\citep[e.g.,][]{Fall85}; formation of galaxies \\citep[e.g.,][]{Gold59, Sofue69}; general cooling flows in galaxies and galaxy cluster cores \\citep[e.g.,][]{Mathews78, Nulsen86, Malagoli87, Meiksin88, Yoshida91, Guo08}; the multi-phase structure of cold filaments condensing out of the hot intracluster medium in galaxy clusters and groups \\citep[e.g.,][]{Sharma10, McCourt11}.  More relevant for our work are the applications of TI in black hole accretion systems. Spherically symmetric accretion flow onto a compact object is preheated by radiation from the central X-ray source. \\citet{Ostriker76} showed the preheating suppresses accretion, such that above a X-ray luminosity $0.005$ of the Eddington limit steady flow is impossible. \\citet{Cowie78} found time-dependent behavior of the preheated accretion for a much wider range of source parameters, and the presence of instabilities. However, \\citet{BisnovatyiKogan80} found that numerical solutions exist for any luminosity below the Eddington limit, but they did not investigate the stability of the solutions in detail. Performing stability analyses \\citet{Stellingwerf82} found globally stable solutions, but those become subject to a drift instability at luminosities over the Eddington limit exceeding $\\sim 0.03$. Hydrodynamical simulations by \\citet{Proga07} show that with an accretion luminosity of $0.6$ of the Eddington value, the gas flow settles into a steady state with two components: an equatorial inflow, and a bipolar inflow/outflow with the outflow leaving the system along the disk rotational axis. There the temperature at the outer radius was fixed at a relatively high value ($2 \\times 10^{7}$ K), so the gas did not go through a thermally unstable phase. Also \\citet{Park12} find that a quasi-steady BH accretion is possible at the Eddington luminosity when the gas density is above a critical value inversely proportional to the BH mass.  Investigating spherical accretion of gas irradiated by a quasar-like continuum, \\citet{Krolik83} found specific criteria for when TI may disrupt single-phase steady flow, in the luminosity -- efficiency (of mass to energy conversion) plane. \\citet{Krolik88} argued that broad emission line clouds in AGN originates in cooler condensations forming within a hot medium via finite-amplitude TI. However, \\citet{Mathews90} concluded that broad line-emitting clouds are unlikely to arise from thermally unstable condensations in a hotter medium, but are likely composed of higher-density gas ejected from nearby stars or from a dense accretion disk. \\citet{Wang12} investigated the dynamics of clumps, considered to exist by thermal instability, embedded in and confined by advection-dominated accretion flows onto black holes.   We find that a spherically accreting gas distribution acted upon by a central gravitational potential, undergoing radiative heating and cooling, can be thermally unstable, cool non-spherically and become clumpy. Hence results from our simulations show that fragmentation can occur in a simple system with minimal physical processes. Our study can give insight into the formation and evolution of clouds      % near AGN. TI operating in such a way can be responsible for the origin of two phase medium in the NLR, support star-formation near AGN, contribute to the observed variability of the luminosity.     %  This paper is organized as follows. We describe our simulations in \\S\\ref{sec-numerical}. We present and explain the results in \\S\\ref{sec-results}, and discuss in \\S\\ref{sec-discussion}. We summarize and conclude in \\S\\ref{sec-conclusion}. ",
        "conclusions": "\\label{sec-conclusion}  Gas accretion onto a central SMBH was studied in Paper~I by performing simulations of spherical Bondi accretion using the 3D SPH code GADGET-3. The simulations included radiative heating by a central X-ray corona and radiative cooling. All the results presented in Paper~I were of spherically-symmetric gas properties. However, for a certain parameter range of X-ray luminosity $L_{X}/L_{\\rm Edd} = 0.01, 0.02, 0.05$ with a fixed density at the outer boundary $\\rho_{\\infty} = 10^{-23}$\\,g/cm$^3$ (i.e., within the unstable $\\xi$-range, \\S\\ref{sec-TI}), we observe non-spherical, multi-phase gas motion. In the current work we present detailed analyses of the inhomogeneous gas behavior. The main results are summarized below:  (1) The gas goes through various modes of accretion for different $L_{X}$. There is a stable, spherically-symmetric inflow with $L_{X} = 0.005$ (Run 25). The gas is heated up and expands as $L_X$ is increased, which decreases the central mass inflow at a varying rate. For some intermediate $L_X$ the accretion flow undergoes unstable motion developing non-spherical features. With high enough $L_X$ a strong outflow is produced, ejecting most of the gas outside.   (2) With $L_X = 0.01$ and $0.02$, the gas cools and becomes denser in a non-spherical manner, fragments, and forms multiple clumps. The clumping occurs over $r < 30$\\,pc for $L_{X} = 0.01$, and $r \\sim 40 - 100$\\,pc for $L_{X} = 0.02$. It starts to develop from $\\sim 1.6 - 1.7$ Myr, and continues up to the simulation end. The clumps are filamentary in structure, supersonic, and have a lower photoionization. Hotter, lighter, subsonic gas exist in-between the filaments, and start to flow outward at later times. The clumps undergo a dynamic evolution: become stretched, fragment, merge with each other, finally accreting into $r_{\\rm in}$.  The gas particle properties versus radius show large scatters (few $100$s to $1000$) at the radial ranges where fragmentation occurs, corresponding to the multi-phase structure of the gas. The mass inflow rate shows large fluctuations in the run with $L_X = 0.01$ after $1.6$\\,Myr, and in the run with $L_X = 0.02$ after $2.3$\\,Myr, because of the accretion of multi-phase gas, with the spikes corresponding to the infall of cold dense clumps.   (3) As they move inward the cooled clumps remain denser, however heat up to above $10^5$\\,K. The central heating at $r < 1$\\,pc is mostly by adiabatic compression reaching $T > 10^7$\\,K at $r_{\\rm in}$, while radiative heating would dominate at a few pc for some particles. The photoionization parameter increases while moving to smaller radii, and the clumps remain supersonic. Artificial viscosity heating increases in magnitude toward $r_{\\rm in}$, however it always remains sub-dominant compared to adiabatic heating. This is in contrast to Paper~I due to the different conditions here because of the higher $L_{X}$.   (4) The cooling and clumping is caused by the interplay of thermal instability in the inflowing gas acted upon by radiative processes, as well as being adiabatically compressed and heated. Gas particles in our simulations undergo TI within the unstable $\\xi$-range, where the $\\Trad - \\xi$ equilibrium curve has the steepest slope (between $\\xi \\sim 100 - 500$), and most physically relevant processes (those with constant-pressure or free falling) can give rise to perturbations with a shallower slope.  We detect a distinct mode of cooling/heating in the $[T - \\xi]$ plane. A particle while being dominated by adiabatic process moves up to higher-$T$ along the unstable part of $\\Trad - \\xi$. This happens because the adiabatic free-fall relation of $T$ versus $\\xi$ ($T_{ff, a} \\propto \\xi_{ff, a}^2$) has the same slope as the radiative equilibrium curve in the unstable regions. At a later time the particle has a transition to being dominated by radiative process. Then classical TI operates, and it undergoes catastrophic cooling or heating (whichever dominates), and moves away from the unstable $\\Trad - \\xi$ part. It finally settles on one of the stable branches where $\\Trad - \\xi$ has a shallower slope.   (5) Typically the amplitude of temperature and density perturbations versus time shows that the perturbations have an exponential growth at first, and then come to a saturation for the remaining simulation run. The minimum temperature has an exponential decrease initially, indicating the growth of cooling instability. It remains almost constant after that at the temperature of the lower stable $\\Trad - \\xi$ branch, where $T$ is a weak function of $\\xi$. In terms of density the maximum value rises indicating the growth of cooled, dense clumps caused by TI.   (6) We have multi-phase medium (cold gas with $2 \\times 10^4$\\,K and hot gas with $10^6 - 10^7$\\,K) at the same radius, forming self-consistently in the simulation volume. Furthermore we resolve the Bondi radius of the accretion flow. This can be used to constrain the $\\alpha$-parameter  that is multiplied to the Bondi-Hoyle-Lyttleton mass accretion rate in the sub-grid model of SMBH feedback in large-scale galaxy simulations.  Considering a limiting temperature of $10^5$\\,K between cold and hot phases, we compute the ratio ($\\chi$) of mass inflow rates of cold gas over hot. We find that $\\chi$ is a dynamical quantity depending on several factors: time, spatial location, and X-ray luminosity. The maximum value of $\\chi$ varies between $1 - 5$. At a fixed radius, $\\chi$ increases with time initially, and later remains almost-constant at a value $\\chi \\sim 3 - 4$. Thus we do not find a factor as large as $\\alpha \\sim 100$, since our $\\chi \\leq 1 - 5$. In other words, in our simulations we do not find such a large fraction of gas hidden in the cold phase that is unresolved in cosmological simulations.   (7) With $L_X = 0.02$ (Run 27) we see both fragmented cold clumps and buoyantly rising hot bubbles. Initially there is a central spherical outflow at $r < 40$\\,pc, which decreases the mass inflow rate at $r_{\\rm in}$ by $\\sim 1000$ times between $1.4 - 2.2$\\,Myr. The surrounding inflowing gas between $[40 - 70]$\\,pc cools and fragments by TI. The dense clumps grow with time, and continue to move inward, disrupting the spherical symmetry of the central hotter outflow.  The clumps start accreting into $r_{\\rm in}$ at $t \\geq 2.2$\\,Myr, causing the central mass inflow rate to rise again. The gas motion becomes inhomogeneous with cooler, denser clumps falling in to the center but heating up as they move in; and hotter, lighter gas moving out. There are also few hot bubbles buoyantly rising from the center.   (8) With a high enough value of the X-ray luminosity ($L_X = 0.05$ in Run 28), the gas is heated up strongly, creating a spherically symmetric outflow. This expels a large fraction of gas out of $r_{\\rm out}$, reduces the central mass inflow rate drastically, and produces a gas density profile that decreases towards center. There is also the formation of an anisotropic, hot bubble buoyantly rising from the center. It has a narrow elongated structure, while its head has a bow shock-like morphology, where the interior gas collides with the relatively cold ambient medium and produces a shock region larger than the jet. The bubble has a greater outward velocity, lower density, higher temperature and higher entropy compared to the gas at the same radius."
    },
    "1112/1112.2430_arXiv.txt": {
        "abstract": "A spectrum of cosmic rays within energy range $10^{15} - 3 \\times 10^{17}$~eV was derived from the data of the small Cherenkov setup, which is a part of the Yakutsk complex EAS array. In this, work a new series of observation is covered. These observations lasted from 2000 till 2010 and resulted in increased number of registered events within interval $10^{16} - 10^{18}$~eV, which in turn made it possible to reproduce cosmic ray spectrum in this energy domain with better precision. A sign of a thin structure is observed in the shape of the spectrum. It could be related to the escape of heavy nuclei from our Galaxy. Cosmic ray mass composition was obtained for the energy region $10^{16} - 10^{18}$~eV. A joint analysis of spectrum and mass composition of cosmic rays was performed. Obtained results are considered in the context of theoretical computations that were performed with the use of hypothesis of galactic and meta-galactic origin of cosmic rays. ",
        "introduction": "Energy spectrum of cosmic rays (CR) in energy range $3 \\times (10^{15} - 10^{18})$~eV could not be studied in detail with compact arrays due to their small acceptance at energy above $10^{17}$~eV. At the same time this area of the spectrum is of a great interest, since local irregularities are manifested there: production of kinks (thin structure at $3 \\times 10^{15} - 10^{17}$~eV) arising from non-uniform distribution of heavier CR components in our Galaxy. On the other hand, this effect is smoothed by addition of a new component (of meta-galactic or other origin) to the cosmic ray flux near Earth. As a result, presence/absence of significant irregularities in spectra measured by various compact arrays allows one to speculate on the CR origin and propagation in our Galaxy~\\cite{Knurenko2011, Kalmykov2008}.  The Yakutsk array in this sense appears as a unique scientific tool. It is related to medium-sized arrays, capable of effective measuring of cosmic rays flux in a wide energy range ($10^{15} - 10^{19}$~eV). Other important traits of the array are its model-independent technique of energy estimation of extensive air showers (EAS) and the ability to track longitudinal EAS development by detecting the Cherenkov light emission. Factors mentioned above enable adopting the unique method, combining the studies of CR spectrum and mass composition aimed at exploration of astrophysical aspect of cosmic rays~\\cite{Knurenko2006, Knurenko2007}. ",
        "conclusions": ""
    },
    "1112/1112.6025_arXiv.txt": {
        "abstract": "In the past decade or so observations of supernovae, Large Scale Structures (LSS), and the Cosmic Microwave Background (CMB) have confirmed the presence of what is called dark energy - the cause of accelerating expansion of the Universe. They have also measured its density as well as the value of other cosmological parameters according to the concordance $\\Lambda$CDM model with few percent uncertainties. Next generation of surveys should allow to constrain this model with better precisions, or distinguish between a $\\Lambda$CDM and alternative models such as modified gravity and (interacting)-quintessence models. In this work we parametrize both homogeneous and anisotropic components of matter density in the context of interacting dark energy models with the goal of discriminating between $f(R)$ modified gravity and its generalizations, and interacting dark energy models, for which we also propose a phenomenological description of energy-momentum conservation equations inspired by particle physics. It is based on the fact that the simplest interactions between particles/fields are elastic scattering and decay. The parametrization of growth rate proposed here is nonetheless general and can be used to constrain other interactions. As an example of applications, we present an order of magnitude estimation of the accuracy of the measurement of these parameters using Euclid and Planck surveys data, and leave a better estimation to a dedicated work. ",
        "introduction": "\\label{sec:frieman} Apriori the measurement of the equation of state of dark energy is simple. It is enough to measure the expansion rate of the Universe $H(z) \\equiv \\dot {a}(z)/a$, or a quantity related to it such as the luminosity distance $D(z)$ at different redshifts. Then, by modeling known constituents of the Universe as non-interacting perfect fluids, one can fit the data and measure the effective equation of state of dark energy $w_{eff}(z)$, defined as $P_{eff}(z)/\\rho_{eff}(z)$. The suffix {\\it ``eff''} is used to remind that pressures and densities obtained in this way can be effective quantities rather than physical pressure and density of constituents, because we have neglected any interaction between components. Therefore, from now on {\\it effective quantities} mean quantities determined from data by considering a null hypothesis.  In practice however the life is not so simple. The density of a perfect fluid changes with redshift as $(1+z)^{3\\gamma}$ ($\\gamma$ is defined in (\\ref{gammade})). Therefore, at low redshifts when $z \\rightarrow 0$, the total density is not very sensitive to the value of $\\gamma$ or equivalently $w(z)$ and their variation with $z$, see Appendix \\ref{app:a} for more details. This statement is independent of the type of data or proxy used for determining $H(z)$ or $D(z)$. On the other hand, at high redshifts where $H(z)$ is more sensitive to the equation of state, dark energy is subdominant. Moreover, it is more difficult to measure $H(z)$ and $D(z)$ at higher redshifts and measurement uncertainties can make the estimation of $w(z)$ and its evolution unusable for discrimination between models.  If constituents of the Universe do not interact with each other, Friedman equation which determines the evolution of expansion function $a(t)$ can be written as: \\bea &&\\frac{H^2}{H_0^2} = \\frac {\\rho (z)}{\\rho_0} = \\Omega_m (1+z)^3 + \\Omega_h (1+z)^4 + \\Omega_K (1+z)^2 + \\Omega_{de} (1+z)^{3\\gamma(z)}, \\quad {\\rho_c (z)} \\equiv \\frac{3H^2}{8\\pi G} \\label{friedmannoint} \\\\ &&\\gamma (z) = \\frac{1}{\\ln (1+z)} \\int_0^z dz'\\frac{1 + w(z')}{1+z'}, \\quad P_{de} (z) \\equiv w(z) \\rho_{de} \\label{gammade} \\\\ && \\text{\\it m = cold dark matter, b = baryons, h = hot matter, k = curvature, and de = dark energy} \\nonumber \\eea In this class of models matter and radiation densities evolve only due to the expansion. This is a good approximation for all redshifts $z < z_{cmb} \\sim 1100$.  In interacting dark energy models matter and radiation terms in the right hand side of the Friedman equation (\\ref{friedmannoint}) can contain an additional redshift-dependent factor: \\be \\frac{H^2}{H_0^2} = \\frac {\\rho_c (z)}{\\rho_{c0}} = \\sum_i \\Omega_i {\\mathcal F}_{i}(z) (1+z)^{3\\gamma_i} \\quad \\quad i=\\text{{\\it m, b, h,k,} and {\\it de}} \\label{friedmanintde} \\ee Without lack of generality we assume that ${\\mathcal F}_{de} = 1$ and all redshift dependent terms are included in $\\gamma (z)$. In quintessence models the coefficient of the curvature term also is constant because it is assumed to be related to geometry/gravity and independent of the behaviour of other components. At present observations are consistent with only gravitational interaction between various components in (\\ref{friedmanintde}), thus additional interactions must be very weak. By definition and without lack of generality we consider ${\\mathcal F}_i(z=0) = 1$. Observations also show that $\\Omega_k \\approx 0$, therefore throughout this work we assume $\\Omega_k = 0$ unless it is explicitly mentioned. Note that in the case of modified gravity models, a parametrization similar to (\\ref {friedmanintde}) can be defined both in Einstein and Jordan frames.  A simple example for which an approximate expression for ${\\mathcal F}_i(z)$ coefficients can be found is a model with a cosmological constant as dark energy and a slowly decaying dark matter. The decay remnants are assumed to be visible relativistic particles~\\cite{houridecay}. In this case: \\bea && \\frac{H^2}{H_0^2} \\approx {\\Omega}_m (1 + z)^3 \\exp (\\frac {{\\tau}_0 - t}{\\tau}) + {\\Omega}_b (1 + z)^3 + {\\Omega}_h (1 + z)^4 + {\\Omega}_m (1 + z)^4 \\biggl (1 - \\exp (\\frac {{\\tau}_0 - t}{\\tau}) \\biggr ) + {\\Omega}_{\\Lambda} \\nonumber \\\\ && \\label {totdens} \\\\ && {\\mathcal F}_m(t) \\approx \\exp (\\frac {\\tau_0 - t}{\\tau}) + (1+z) \\biggl (1 - \\exp (\\frac {t_0 - t}{\\tau}) \\biggr ), \\quad \\tau \\gg \\tau_0, \\quad {\\mathcal F}_b = {\\mathcal F}_h = 1 \\quad \\gamma (z) = 0 \\label{dmdecaycoeff} \\eea where $\\tau$ is the lifetime of dark matter and $\\tau_0$ is the age of the Universe. It is demonstrated that in this example, if the decay/interaction of dark matter is not considered in the data analysis, a $w_{eff} < -1$ can be obtained for dark energy, see~\\cite{houriustat,das05} for more details about the set up and the proof.  Note that in (\\ref{dmdecaycoeff}), we have included the contribution of relativistic remnants in ${\\mathcal F}_m$. However, as this component has a redshift dependence similar to hot matter, it also makes sense to consider it as hot matter and add it to hot component. It is even possible to add this term to dark energy contribution, as long as it is small and induces only a slight deviation from a cosmological constant. In this case, one can show that the {\\it effective dark energy} will have $w_{eff} < -1$~\\cite{houriustat}. The reason for such freedom is that we do not measure or take into account the decay remnants. This example clearly shows that parametrization (\\ref{friedmanintde}) is not unique when all the components and their interactions are not know. Therefore, one has to be very careful about degeneracies when data are analyzed and interpreted. In particular, prior assumptions such as stability of matter and radiation components can affect measurements and conclusions. This example also show that for ruling out $\\Lambda$CDM model, it is enough to prove that at least one of ${\\mathcal F}_i \\neq 1$, or $\\gamma_{de} \\neq 0$\\footnote{This statement is true if baryon pressure is negligible. Future surveys can be sensitive to small baryon pressure. In this case it must be taken into account before any conclusion about $\\Lambda$CDM model is made.}.  Extension of this example to quintessence models without coupling to matter is straightforward and one simply needs to consider $\\gamma (z) \\neq 0$. A more interesting extension is to assume that the quintessence scalar field is one of the remnants of the decay of dark matter, which during cosmological time condensates and makes a classical quintessence field. In this case, it has been shown~\\cite{houriquin,houriquin1,houricondens} that coefficients ${\\mathcal F}_m$, ${\\mathcal F}_h$, and equation of state of dark energy $w(z)$ (or equivalently $\\gamma (z)$) are not independent. However, their relations are too sophisticated and cannot be described in an analytical form and numerical techniques should be employed~\\cite {houriquin}.  According to (\\ref{dmdecaycoeff}): \\be {\\mathcal F}_m(z) > {\\mathcal F}_m(z=0) \\label{quinfdmtode} \\ee and because $\\tau \\gg \\tau_0$, ${\\mathcal F}_i(z)$ coefficients are close to 1 at all redshifts. In general, for an interaction which transfers energy from dark matter to other components, the inequality (\\ref{quinfdmtode}) is applied because at high redshifts one expects a larger contribution of dark matter in the total density than in a non-interacting model. Inversely, if energy is transferred from other components, for instance from dark energy, to dark matter: \\be {\\mathcal F}_m(z) < {\\mathcal F}_m(z=0) \\label{quinfdetodm} \\ee An example of such models is {\\it scaling dark energy}~\\cite{scalingde,scalingde1} in which at early times dark energy has a much larger contribution in the total energy density, but it gradually decays to dark matter and only recently its equation of state approaches $w \\sim -1$. Another example is the class of models called {\\it early dark energy}. Although the original model~\\cite{earlyde,earlyde1,earlyde2} is a pure quintessence/k-essence, there are variants of this model in which, there is an interaction in the dark sector~\\cite{earlydeint} or between dark energy and visible sector~\\cite{earlydeintvarpalpha}.  In models with elastic interaction between two sectors, no energy is transferred between them, and ${\\mathcal F}_m(z) = {\\mathcal F}_h(z) = 1$. Nonetheless, the phase space of matter and dark energy in these models can change and thereby $w_{de}$ can depend on $z$.  For $\\mathrm {f}(R)$ modified gravity models homogeneous Einstein equations and energy conservation equation in Jordan frame are~\\cite{frbean}:\\footnote{When equations apply to both dark matter and baryons, we indicate them collectively with subscript $m$.} \\bea && (1+\\mathrm {f}_R) H^2 + \\frac{1}{6} \\mathrm {f} - \\frac{a''}{a^3}\\mathrm {f}_R + \\frac{\\mathrm {f'}_R}{a} H = \\frac {8\\pi G}{3} \\sum_i \\rho_i \\label{grmg0} \\\\ && \\frac {a''}{a^3} = -\\frac{4\\pi G}{3} \\sum_i (\\rho_i + 3P_i) + (1+\\mathrm {f}_R) H^2 + \\frac{\\mathrm {f}}{6} - \\frac {H\\mathrm {f}'_R}{a} - \\frac{\\mathrm {f}''_R}{2a^2} \\label{grmg1} \\\\ && \\dot{\\rho}_i + 3H\\rho_i = -\\frac{\\dot{f}_R}{2 (1+\\mathrm {f}_R)}(\\rho_i -3 P_i), \\quad i=\\text{\\it m, h, k} \\label{grgmener} \\eea where $a' = a\\dot{a}$.\\footnote{Here we have written Einstein and conservation equations in Jordan frame because they lead to expressions for ${\\mathcal F}_i$ coefficients which are explicitly very different from quintessence case.} Dot and prime mean derivation with respect to comoving and conformal time, respectively. Subscript $R$ means derivation with respect to scalar curvature $R \\equiv R_{\\mu\\nu}g^{\\mu\\nu}$. We remind that at linear order the effect of matter perturbations on $R$ is zero, thus $R$ only depends on $z$ and the effect of $\\mathrm {f}(R) \\neq 0$ on the evolution of perturbations manifests itself by changing the background cosmology.  After solving density conservation equation (\\ref{grgmener}), Friedman equation (\\ref{grmg0}) can be written as the following: \\bea \\rho_i (z) &=& \\rho_i (z=0) (1+z)^{3\\gamma_i} \\biggl (\\frac{{1+\\mathrm {f}_R (z=0)}} {1+\\mathrm {f}_R (z)} \\biggr)^{-\\frac{1-3w_i}{2}} \\label{rhosolmg} \\\\ \\frac{H^2}{H_0^2} &=& \\frac {\\rho_c (z)}{\\rho_{c0}} = \\sum_i \\Omega_i {\\mathcal F}_{i}(z) (1+z)^{3\\gamma_i}  \\label{friedmanmg} \\\\ \\rho_{de} &=& \\frac{3}{8\\pi G}~\\frac{1}{1+\\mathrm {f}_R}(-\\frac{\\mathrm {f}(R)}{6} + \\frac{a''}{a^3}\\mathrm {f}_R - H\\dot{\\mathrm {f}}_R) \\label{mgdedens} \\\\ {\\mathcal F}_i(z) &=& \\biggl (\\frac{1+\\mathrm {f}_R(R(z=0))}{1+\\mathrm {f}_R(R(z))}\\biggr)^ {-\\frac{1-3w_i}{2}}, \\quad w_m = 0, \\quad w_h = \\frac{1}{3}, \\quad w_k = -\\frac{1}{3} \\label{fmgdef} \\eea Equation (\\ref{mgdedens}) is the energy density of {\\it effective dark energy} in $\\mathrm {f}(R)$ gravity models. Similar to quintessence models we can assume ${\\mathcal F}_{de} = 1$. The only explicit difference between (\\ref{friedmanmg}) and the same equation for an interacting quintessence model is the presence of a nontrivial coefficient for the curvature term if $\\Omega_k \\neq 0$. Nonetheless, the evolution of coefficients ${\\mathcal F}_i(z)$ with redshift is different from their counterparts in interacting quintessence models, in particular from models in which energy is transferred to dark energy at low redshifts, see equation (\\ref{quinfdmtode}). In fact, the function $\\mathrm {f}(R)$ is not completely arbitrary and must satisfy a number of constraints. Notably, $\\mathrm {f}(R)|_{|R| \\gg 0} \\rightarrow 0$ to make the model consistent with Einstein theory of gravity in mild or strong gravity fields, and $\\mathrm {f}_R > 0$ due to stability constraint~\\cite{modgrstabil}. Under these conditions: \\be {\\mathcal F}_i(z) > {\\mathcal F}_i(z=0)  \\label{mgfz} \\ee Comparing (\\ref{mgfz}), (\\ref{quinfdmtode}), and (\\ref{quinfdetodm}) one can immediately conclude that the measurement of ${\\mathcal F}_m(z)$ and its evolution with redshift can discriminate between dark energy models in which energy is transferred from dark energy to dark matter such as scaling models, and $\\mathrm {f}_R$ modified gravity models. But it cannot discriminate modified gravity from models in which energy is transferred from dark matter to dark energy such as the model discussed in~\\cite{houriquin,houriquin1,houricondens}. To discriminate the latter and other models of this category from $\\mathrm {f}_R$ modified gravity, the coefficient of relativistic (hot) component ${\\mathcal F}_h(z)$ and its evolution must be measured. Evidently, such measurements are very difficult. For instance, one has to measure very precisely the temperature of CMB at high redshifts or $H(z)$ at a large number of redshift bins and fit the data with ${\\mathcal F}_h \\neq 1$. In Einstein frame the evolution of matter density is the same as in equation (\\ref{rhosolmg})~\\cite{frbean}, but evolution equation of hot matter is similar to $\\Lambda$CDM. In what concerns the discrimination from interacting quintessence what is discussed from Jourdan is applicable.  \\subsection {Model-independent discrimination of interacting dark energy models} \\label{sec:discrim} In this section we show that if $\\Lambda$CDM or a simple quintessence are considered as null hypothesis, measurements of effective dark energy density and effective equation of state from $H (z)$ and the function $A(z)$ defined in Appendix \\ref {app:a} separately, give different values for these quantities if dark energy interacts with matter. Similarity of ${\\mathcal F}_m(z)$, specially if the curvature of the Universe is zero, means that we cannot distinguish between interacting quintessence and modified gravity models in a model-independent manner - except for the cases explained above. For this reason in this section we only study the discrimination between interacting dark energy models parametrized as in equation (\\ref{friedmanmg}) and a cosmological constant and/or non-interacting quintessence.  For analyzing cosmological data, $\\Lambda$CDM with a stable and non-interacting dark matter is usually used as null hypothesis. Nonetheless, the methodology explained below is not sensitive to redshift dependence of $\\gamma_{de}$, and we can consider the more general case of non-interacting quintessence as the null hypothesis. The expansion of the Universe for such cosmologies is ruled by equation (\\ref{friedmannoint}). Therefore, we rearrange terms in equation (\\ref{friedmanmg}) such that it looks similar to equation (\\ref{friedmannoint}). Then, we determine effective quantities which are measured by fitting a $\\Lambda$CDM or a non-interacting quintessence model to data: \\be \\frac{H^2}{H_0^2} = \\sum_i \\Omega_i (1+z)^{3\\gamma_i} + \\sum_i \\Omega_i ({\\mathcal F}_{i}(z) - 1) (1+z)^{3\\gamma_i} + \\Omega_{de} (1+z)^{3\\gamma_{de} (z)} \\label{friedmanmgequiv} \\ee In null hypothesis model only $\\gamma_{de}$ is redshift dependent and $\\gamma_i,~i=m,~h,~k$ are constant. By comparing (\\ref{friedmanmgequiv}) with (\\ref{friedmannoint}) the {\\it effective} contribution of dark energy is expressed as: \\be \\Omega_{eff}^{(H)} (1+z)^{3\\gamma_{eff}^{(H)}(z)} = \\sum_i \\Omega_i ({\\mathcal F}_i(z) - 1) (1+z)^{3\\gamma_i} + \\Omega_{de} (1+z)^{3\\gamma_{de} (z)} \\label{deeffdef} \\ee In both interacting quintessence and modified gravity models coefficients $F_i$'s are defined such that $F_i (z=0) = 1$, therefore at $z = 0$ the first term in (\\ref{deeffdef}) is null and we can separate $\\Omega_{eff}$ and $\\gamma_{de} (z)$: \\bea && \\Omega_{eff}^{(H)} = \\Omega_{de}, \\quad \\quad \\gamma_{eff}^{(H)}(z=0) = \\gamma_{de}(z=0) \\label{omegaffh} \\\\ && \\gamma_{eff}^{(H)}(z) = \\frac{\\log \\biggl (\\sum_i \\frac{\\Omega_i}{\\Omega_{de}} ({\\mathcal F}_i (z) -1)(1+z)^{3\\gamma_i} + (1+z)^{3\\gamma_{de} (z)}\\biggr )}{3 \\log(1+z)} \\label{deeffexp} \\eea where superscript $(H)$ means measured from Hubble constant $H$.  Suppose we can also measure $A(z)$ defined in (\\ref{azdef}), and use it to determine the effective density and equation of state of dark energy. For an interacting dark energy model parametrized according to (\\ref{friedmanmgequiv}) quantities $B(z)$ and $A(z)$ are: \\bea B(z) &\\equiv& \\frac{1}{3 (1+z)^2 \\rho_0} \\frac {d\\rho}{dz} = \\nonumber \\\\ && \\sum_{i=m,h,k} \\Omega_i \\biggl (\\gamma_i {\\mathcal F}_i (z) + (1+z) \\frac{d{\\mathcal F}_i}{dz} \\biggr )(1+z)^{3 (\\gamma_i - 1)} + \\Omega_{de}(w(z)+1) (1+z)^{3(\\gamma_{de}(z) - 1)} \\label{bzintdmde} \\\\ A(z) &\\equiv& B(z) - \\sum_{i=m,h,k} \\Omega_i \\gamma_i (1+z)^{3(\\gamma_i - 1)} = \\nonumber \\\\ && \\sum_{i=m,h,k} \\Omega_i \\biggl (\\gamma_i ({\\mathcal F}_i (z) - 1) + (1+z) \\frac{d{\\mathcal F}_i}{dz} \\biggr ) (1+z)^{3 (\\gamma_i - 1)} + \\Omega_{de}(w(z)+1) (1+z)^{3(\\gamma_{de}(z) - 1)} \\label{azintdmde} \\eea Using (\\ref{azdef}) in Appendix \\ref{app:a} as the definition of $A(z)$, we find the following expression for its parameters: \\bea \\Omega_{eff}^{(A)} (w_{eff}^{(A)}(z) + 1) (1+z)^{3\\gamma_{eff}^{(A)}(z)} &=& \\sum_i \\Omega_i \\biggl (\\gamma_i ({\\mathcal F}_i (z) - 1) + (1+z) \\frac{d{\\mathcal F}_i}{dz} \\biggr ) (1+z)^{3\\gamma_i} + \\nonumber \\\\ && \\Omega_{de}(w(z)+1)(1+z)^{3\\gamma_{de}(z)} = (1+z) A(z) \\label{deeffdefaz} \\eea where superscript $(A)$ means measured from $A(z)$. Equations (\\ref{deeffdef}) and (\\ref{deeffdefaz}) are fundamentally different. In particular: \\be \\Omega_{eff}^{(A)} = \\frac{\\sum_i \\Omega_i \\frac{d{\\mathcal F}_i (z=0)}{dz} + \\Omega_{de} (w(z=0)+1)}{w_{eff}^{(A)}(z=0) + 1} \\label{omegaedeffaz} \\ee which in contrast to $\\Omega_{eff}^{(H)}$, in general is not equal to $\\Omega_{de}$. Equality arises only when ${\\mathcal F}_i$ do not vary with redshift i.e. ${\\mathcal F}_i = 1$ at all redshifts. This condition is satisfied by the null hypothesis $\\Lambda$CDM and by non-interacting quintessence models. Therefore, assuming that $\\Omega_m$ and $\\Omega_k$ are known (e.g. from CMB), simultaneous measurements of $H(z)$ and $A(z)$ at even one $z > 0$ is apriori enough for testing the presence of an interaction between dark matter and dark energy independent of the underlying model. Evidently, in practice the measurements must be performed at many redshift bins to improve statistics and to compensate for measurement errors.  Apriori one can use other quantities such as angular diameter distance $D_A$ or luminosity distance $D_L$ which are easier to measure rather than $A(z)$. However, both these quantities are functional of $H(z)$ - through integration of $1/H^{1/2}(z)$. Thus, in general they do not have an analytical expression. Besides, their derivatives depend on ${\\mathcal F}_i$'s only, in contrast to (\\ref{deeffdefaz}) that depends on both ${\\mathcal F}_i$'s and their derivatives. Therefore, $\\Omega_{eff}$ and $\\gamma_{eff}$ obtained from $dD_A/dz$ or $dD_L/dz$ will be equal to ones determined from $H(z)$ irrespective of the underlying cosmology. This shows that the function $A(z)$ (or equivalently $B(z)$) introduced in~\\cite{houriaz} has special properties and is well suited for discriminating between dark energy models. It can be measured from supernovae data, see~\\cite{houriaz} for the methodology. As for LSS data, one needs to determine both $H(z)$ and its evolution $dH(z)/dz$ to be able to calculate $A(z)$, for instance from the BAO and the power spectrum of matter fluctuations~\\cite{psobs}. This is not an easy task. As an example consider supernovae observations that measure the luminosity distance $D_L$ to a supernova from its standardized apparent magnitude. The angular luminosity distance $D_A$ is related to the luminosity distance, see (\\ref{dah}). To determine $dD_A/dz$ apriori one can use the measured $D_A$, and determine its derivative (slope). However, due to scattering and discreteness of data, such a measurement will have large uncertainties. The same problem arises for $dH(z)/dz$ or $A(z)$ because they depend on derivatives of $D_L$, see equations (\\ref{rhoderdl}) and (\\ref{rhoderda}). Nonetheless, there are various methods such as binning of data, using a fit in place of discrete data, etc. that allow to improve the estimation. Near future large area surveys such as Euclid~\\cite{euclid}, BigBOSS~\\cite{bigboss}, LSST~\\cite{lsst} will be able to determine these quantities with relatively good precision, see also Sec. \\ref{sec:forecast} for measurement methodology. In particular, large surface spectroscopic and lensing surveys such as Euclid are able to determine the variation of total density with redshift $d\\rho/dz \\propto B(z)$ with good precision. In Appendix \\ref{app:b} we obtain the Fisher matrix for dark energy parameters without considering a specific parametrization for the equation of state $w(z)$.  \\subsection{Discrimination precision} \\label{sec:errordiscrim} Measurements of cosmological parameters show that $w^{obs}_{de} \\sim -1$ irrespective of which proxy or measurement method - supernovae, CMB, or LSS has been used. This means that $|{\\mathcal F}_i(z) - 1| \\approx 0$ and $d{\\mathcal F}_i (z)/dz \\approx 0$. Moreover, addition of ${\\mathcal F}_i(z)$ to the model increases the number of parameters. Giving the fact that we have essentially two observables: $H(z)$ and one of $D_A(z)$, $D_L(z)$ or $B(z)$, greater number of parameters means also greater degeneracy, thus more uncertainty for discrimination between $\\Lambda$CDM, a non-interacting quintessence, and interacting dark energy models.  One way of measuring the presence of interaction without having to fit data to the large number of parameters in equations (\\ref{friedmanmgequiv}) and (\\ref{azintdmde}), is to measure how different $\\Omega_{eff}^{(H)},~\\Omega_{eff}^{(A)},~\\gamma_{eff}^{(H)},$ and $\\gamma_{eff}^{(A)}(z)$ are, because as we discussed in the previous section, when ${\\mathcal F}_i \\neq 1$ these effective quantities are not the same. To this end, a natural criteria is: \\be \\Theta (z) \\equiv \\frac{\\Omega_{eff}^{(A)} (w_{eff}^{(A)}(z) + 1) (1+z)^{3\\gamma_{eff}^{(A)}(z)} - \\Omega_{eff}^{(H)} (w_{eff}^{(H)}(z) + 1) (1+z)^{3\\gamma_{eff}^{(H)}(z)}}{\\Omega_{eff}^{(H)} (w_{eff}^{(H)}(z) + 1) (1+z)^{3\\gamma_{eff}^{(H)}(z)}} \\label{deftheta} \\ee This quantity can be explained explicitly as a function of $\\Omega_i,~{\\mathcal F}_i,~\\gamma_i$, and is zero when $F_i= 1,~dF_i/dz = 0$. Note that we have chosen expression (\\ref{deeffdefaz}) for comparison rather than (\\ref{deeffdef}) because it is not possible to determine $\\Omega_{eff}^{(A)}$ in a model independent manner, see equation (\\ref{omegaedeffaz}). By contrast $\\Omega_{eff}^{(H)} = \\Omega_{de}$, thus $\\gamma_{eff}^{(H)}$ and thereby $w_{eff}^{(H)}$ can be determined without any reference to ${\\mathcal F}_i$ coefficients. In~\\cite {houriaz} we suggested to use the sign and evolution of $A(z)$ to discriminate between dark energy with $\\gamma (z) \\neq 0$ and a cosmological constant. Here $\\Theta (z)$ plays a similar role for discriminating between interacting or non-interacting dark energy.  Assuming that $\\Omega_m$ and $\\Omega_h$ are determined independently and with very good precision, for instance from CMB anisotropies with marginalization over $\\gamma_{de}$, $\\Theta$ can be determined from the measurement of $H(z)$ and $B(z)$. The latter can be measured from whole sky or wide area spectroscopic surveys data such as Euclid, or multi-band photometric surveys such as DES. Evidently determination of $B(z)$ that depends on $dH/dz$ is very difficult. However, it is easy to see that there is no other quantity which can be measured more easily and discriminates between $\\Lambda$CDM and dynamical dark energy models with a better precision. For instance, the BAO method determines $H(z)$ and $D_A(z)$ directly. But, $D_A(z)$ depends on $w(z)$ or equivalently $\\gamma (z)$ through an integral, see equation (\\ref{dah}). Therefore, it is less sensitive to the variation of $\\gamma (z)$ with redshift. This is analogue to binning a data. Evidently, a binned data is less noisy and has a smaller uncertainty. But, if the goal is to measure the variation of data, the binning can completely smear out small variations. Therefore, irrespective of methods and measured proxies, we are limited by inherent properties of the physical system. In this respect, the precision with which $\\Theta (z)$ can be measured gives the ultimate sensitivity of an observation/data set to deviation from $\\Lambda$CDM. ",
        "conclusions": ""
    },
    "1112/1112.1096_arXiv.txt": {
        "abstract": "We use SDSS spectra and optical to far-infrared photometry for a sample of 31 FeLoBAL QSOs to study the relationship between the AGN-driven outflows, and obscured star formation in their host galaxies. We find that FeLoBAL QSOs invariably have IR luminosities exceeding $10^{12}$L$_{\\odot}$. The AGN supplies 75\\% of the total IR emission, on average, but with a range from 20\\% to 100$\\%$. We find a clear anticorrelation between the strength of the AGN-driven outflows and the contribution from star formation to the total IR luminosity, with a much higher chance of seeing a starburst contribution in excess of 25\\% in systems with weak outflows than in systems with strong outflows. Moreover, we find no evidence that this effect is driven by the IR luminosity of the AGN. We conclude that radiatively driven outflows from AGN act to curtail obscured star formation in the host galaxies of reddened QSOs to less than $\\sim 25\\%$ of the total IR luminosity. This is the most direct evidence yet obtained for `quasar mode' AGN feedback. ",
        "introduction": "Over the last decade, several problems have arisen in modelling the cosmological evolution of galaxies. These include (1) the difficulties that models faced in explaining the observed galaxy luminosity function at low and high redshifts simultaneously (e.g. \\citealt{benson03}), (2) the prediction that rich galaxy clusters should harbour cooling flows when few are observed (e.g. \\citealt{peterson01}), and (3) the difficulties that models face in reproducing the number of IR-luminous galaxies observed at high redshift (e.g. \\citealt{baugh05}).  One of the most promising solutions to these issues is AGN feedback. AGN feedback is the exertion of influence of an SMBH on $\\gtrsim$kpc scales to curtail star formation in the host galaxy, and/or accretion onto the SMBH itself. Galaxy evolution models usually assume this feedback to occur in one or both of two simplified modes; `quasar' mode and `radio' mode. Quasar mode feedback occurs via radiation from an accretion disk, while radio mode feedback occurs via a relativistic jet that transfers momentum to the ISM. Both these feedback paradigms have led to improvements in the ability of models to reproduce observations (e.g. \\citealt{bower08,somer08}), but observational evidence for either paradigm remains sparse.  Our group has been looking for evidence of quasar mode feedback by examining systems in which such feedback may be ongoing. To do so, we have been studying the `FeLoBAL' QSOs \\citep{hazard87,hall02}. We selected this population for two reasons. First, their UV absorption troughs are unambiguous signatures of radiatively driven outflows powered by an AGN. Second, FeLoBAL QSOs are invariably reddened and have high IR luminosities \\citep{farrah07,farrah10}. We here present results for 31 FeLoBAL QSOs, in which we compare the strength of their outflows as estimated from their UV spectral properties, to the luminosity of obscured star formation in their host galaxies as measured from optical through far-IR photometry. We define \"IR luminosity\" as the luminosity integrated over 1-1000$\\mu$m in the rest-frame. ",
        "conclusions": "We find that FeLoBAL QSOs are luminous in the IR. All 31 of our sample have total IR luminosities (L$_{Tot}$) in excess of $10^{12}$L$_{\\odot}$, with nine objects exceeding $10^{13}$L$_{\\odot}$. These luminosities are comparable to those of the wider population of BAL QSOs \\citep{gallagher07}, red QSOs \\citep{geo09}, and ULIRGs \\citep{farrah03,farrah09}. The dominant power source is usually an AGN. A pure AGN is either the most likely power source, or consistent within the 90\\% confidence interval, for $35\\%$ of the sample. A starburst component is required for the remaining objects, but in only twelve objects is the starburst more luminous than $10^{12}$L$_{\\odot}$, and in only three objects is the starburst more luminous than the AGN. The mean starburst contribution to the total IR luminosity (f$_{SB}$) is $\\sim24\\%$, comparable to that of local ULIRGs with `warm' IR colours, but lower than those of PG QSOs \\citep{veill09}.  \\begin{figure} \\includegraphics[width=52mm,angle=90]{farrah_fig1a.ps} \\includegraphics[width=52mm,angle=90]{farrah_fig1b.ps} \\caption{Absorption strength vs (left) starburst (stars) and AGN (squares) IR luminosities, and (right) starburst contribution to the total IR luminosity.}\\label{luminosities} \\end{figure}  We now examine whether or not there is a relationship between the AGN-driven outflows and the obscured star formation. To quantify the strength of the outflows we use the same, single species (MgII$\\lambda$2799\\AA) across the whole sample, and adopt the BI$_{0}$ parameter of \\citet{gibson09}. We measured the MgII absorption strengths using the methods described in \\citet{f2ms}.  We first examine if absorption strength depends on AGN or starburst IR luminosity (Fig \\ref{luminosities}, left). There is no correlation between absorption strength and L$_{SB}$ (a Spearman rank test gives $\\rho = -0.10$, $P = 0.58$) and at best a weak correlation between absorption strength and L$_{AGN}$ ($\\rho = 0.39$, $P = 0.04$). Conversely, we see a weak but clear anticorrelation between absorption strength and f$_{SB}$ ($\\rho = -0.49$, $P = 0.005$, Fig \\ref{luminosities} right). Moreover, all the systems with BI$_{0}\\gtrsim 3500$km s$^{-1}$ have $f_{SB}<25\\%$, while the systems with BI$_{0}\\lesssim 3500$km s$^{-1}$ have a wide dispersion in starburst contributions, from $0\\%$ to $\\sim80\\%$. A two-sided Kolmogorov-Smirnov test for the objects above and below BI$_{0} = 3500$km s$^{-1}$ yields a difference in distributions at 99.84\\% confidence, though the number of objects in the two subsamples is too low for this test to be robust.  We examine this result in more detail by constructing probability distribution functions (using all the fit solutions) for f$_{SB}$ for two subsamples divided at BI$_{0}=3500$km s$^{-1}$ (Figure \\ref{pdfs}, left). We see a clear difference. The low absorption strength subsample shows a much higher chance of a higher starburst contribution than the high absorption strength subsample. We quantify this by extracting the probabilities of obtaining $f_{SB}>25\\%$. For the whole sample we find P($f_{SB}>25\\%$)$=50.3^{+5.3}_{-5.4}$\\%, for the low absorption strength sample we find P($f_{SB}>25\\%$)$=67.3^{+4.5}_{-4.1}$\\%, while for the high absorption strength sample we find P($f_{SB}>25\\%$)$=17.8^{+3.7}_{-6.5}$\\%.  This anticorrelation between absorption strength and f$_{SB}$ is straightforwardly interpreted as the outflow from the AGN curtailing star formation in the host galaxy. In this scenario, the systems with $f_{SB}>25\\%$ are those in which an outflow has yet to curtail star formation, so the outflows always have BI$_{0}\\lesssim 3500$km s$^{-1}$. Conversely, the systems with $f_{SB}<25\\%$ are those in which an outflow has curtailed star formation, {\\itshape and} those in which such an outflow has subsequently waned, making the dispersion in absorption strengths wide.  There are however four other ways that this anticorrelation could arise. The first is that the starburst `suppresses' AGN outflows, so when the starburst wanes (via a cause unrelated to the AGN) an AGN driven outflow can appear, possibly because the ISM density has been reduced by the starburst. We cannot test this scenario, but it would likely require a serendipitous conjunction of ISM and SMBH parameters, so we do not consider it further. The second is if a high starburst contribution meant that the Mg II troughs were {\\it observed} to be weaker than they really are. Again, we cannot test this, but the rest-frame UV continua of starbursts in ULIRGs are at least an order of magnitude too weak to provide this effect, and can show absorption in the same species \\citep{farrah05}. The third is if BAL QSOs with strong starbursts preferentially drop out of the SDSS QSO selection compared to BAL QSOs with weak starbursts, and were thus not included in \\citet{trump06}. This possibility is also not testable by us, but the SDSS is now turning up FeLoBAL QSOs in large numbers, and the SDSS QSO followup colour selections are fairly relaxed, so we do not consider this possibility likely either.  The fourth possibility is that stronger outflows reflect an increase in the IR emission from the AGN, but have no effect on the starburst; if this is the case we would see a decline in the contribution from star formation to the total IR luminosity with increasing absorption strength, but with no direct relationship behind the decline. This is a possibility we can test, as follows. If outflow strength is just a proxy for AGN luminosity, then we should see a bigger difference between the starburst contribution PDFs for subsamples divided by AGN luminosity than between subsamples divided by absorption strength. In the right panel of Figure \\ref{pdfs} we show starburst contribution PDFs for two subsamples divided by AGN luminosity at L$_{AGN}=10^{12.5}$L$_{\\odot}$. The difference between the PDFs divided by AGN luminosity is {\\it weaker} than the difference between the PDFs divided by absorption strength.  Furthermore, we are more likely (albeit only at $\\sim$2$\\sigma$) to obtain a smaller starburst contribution by selecting high absorption strength systems than we are by selecting high AGN luminosity systems; for P($f_{SB}>25\\%$): the BI$_{0}>3500$km s$^{-1}$ subsample is $17.8^{+3.7}_{-6.5}$\\% while the L$_{AGN}>10^{12.5}$L$_{\\odot}$ subsample is $38.7^{+8.5}_{-10.0}$\\%). In other words, we are more successful in finding systems with a large starburst contribution to the total IR emission by selecting on weak outflows than we are by selecting on low AGN luminosity. Also, we are more likely to find star formation with a lower absolute luminosity in the {\\itshape lower} luminosity AGN subsample (e.g. for P(L$_{Sb}<10^{12}$L$_{\\odot})$: the L$_{AGN}<10^{12.5}$L$_{\\odot}$ subsample is $69.4^{+5.2}_{-4.6}$\\% while the L$_{AGN}>10^{12.5}$L$_{\\odot}$ subsample is $43.2^{+6.9}_{-10.4}$\\%)  Overall therefore, we find that radiatively driven outflows from an AGN with absorption strengths $\\gtrsim3500$ km s$^{-1}$ act to curtail star formation in their host galaxies. We also find that this effect is (at least largely) relative; such outflows reduce the contribution from star formation to the total IR luminosity to less than $\\sim25\\%$. We also propose that the {\\itshape infrared} luminosity of the AGN is not a good proxy for the degree of AGN feedback that is taking place.  \\begin{figure} \\includegraphics[width=52mm,angle=90]{farrah_fig2a.ps} \\includegraphics[width=52mm,angle=90]{farrah_fig2b.ps} \\caption{Probability Distribution Functions for the starburst contribution to the total IR luminosity. In both panels, the solid line is the PDF for the whole sample. The other lines are for subsamples divided according to two criteria. {\\itshape Left panel}: by absorption strength. {\\itshape Right panel}: by AGN IR luminosity.}\\label{pdfs} \\end{figure}"
    },
    "1112/1112.3093_arXiv.txt": {
        "abstract": " ",
        "introduction": "With the advent of the \\textit{Square Kilometre Array} (SKA; \\citealt{dewdney2009}) and its precursors and pathfinders, including the \\textit{Australian SKA Pathfinder} (ASKAP; \\citealt{deboer2009}), the \\textit{Karoo Array Telescope} (Meer\\-KAT; \\citealt{jonas2009}), and the \\textit{Aperture Tile In Focus} (APERTIF; \\citealt{oosterloo2009}), the prospect of deep radio continuum and \\ion{H}{i} surveys of large areas on the sky demands for new strategies in the areas of data reduction and analysis, given the sheer volume of the expected data streams, in particular for spectroscopic surveys.  Of particular importance is the automatic and accurate identification and parametrisation of sources with high completeness and reliability. Due to the large data volumes to be searched, source finding algorithms must be fully automated, and the once common practice of source finding `by eye' will no longer be feasible. Moreover, accurate source para\\-metrisation algorithms need to be developed to generate reliable source catalogues free of systematic errors, as otherwise the integrity of scientific results based on the survey data could be compromised.  In this paper we will take a closer look at the \\textsc{Duchamp} source finder\\footnote{\\textsc{Duchamp} website: http://www.atnf.csiro.au\\slash{}people\\slash{}Matthew.Whiting\\slash{}Duchamp/} \\citep{whiting2011a,whiting2012}. \\textsc{Du\\-champ} has been developed by Matthew Whiting at CSIRO as a general-purpose source finder for three-dimensional data cubes as well as two-di\\-mensional images and will serve as the default source finder in the processing of data from the ASKAP survey science projects. The software identifies sources by searching for regions of emission above a specified flux threshold. To improve its performance, \\textsc{Duchamp} offers several methods of preconditioning and filtering of the input data, including spatial and spectral smoothing as well as reconstruction of the entire image or data cube with the help of wavelets. In addition to source finding, \\textsc{Duchamp} provides the user with basic source parametrisation, including the measurement of position, size, radial velocity, line width, and integrated flux of a source. More information about the capabilities of the software is available from the \\textsc{Duchamp} User Guide \\citep{whiting2011b}. A brief overview of \\textsc{Du\\-champ}'s basic functionality is provided in Section~\\ref{sect_duchamp}.  \\begin{table}[ht] \\begin{center} \\caption{Summary of the parameters used to generate the visibility data set and noise image for the point source models.} \\label{tab_model} \\begin{tabular}{lrl} \\hline Parameter (visibility) &         Value & Unit           \\\\ \\hline Number of antennas     &           $36$ &                \\\\ System temperature     &           $50$ & $\\mathrm{K}$   \\\\ Declination            &  $-45^{\\circ}$ &                \\\\ Total integration time &            $8$ & $\\mathrm{h}$   \\\\ Hour angle range       &        $\\pm 4$ & $\\mathrm{h}$   \\\\ Cycle time             &            $5$ & $\\mathrm{s}$   \\\\ Stokes parameters      &              I &                \\\\ Number of channels     &           $31$ &                \\\\ Frequency              &         $1.42$ & $\\mathrm{GHz}$ \\\\ Channel width          &        $18.31$ & $\\mathrm{kHz}$ \\\\ &         $3.86$ & $\\mathrm{km \\, s}^{-1}$ \\\\ \\hline Parameter (image)      &          Value & Unit                 \\\\ \\hline Final image size       & $31 \\times 31$ & $\\mathrm{px}$        \\\\ Field diameter         &            $5$ & $\\mathrm{arcmin}$    \\\\ Pixel size             &           $10$ & $\\mathrm{arcsec}$    \\\\ Robustness             &            $0$ &                      \\\\ Gaussian $uv$ taper    &         $7.28$ & $\\mathrm{k} \\lambda$ \\\\ &         $1.54$ & $\\mathrm{km}$        \\\\ \\textsc{rms} noise     &         $1.95$ & $\\mathrm{mJy}$       \\\\ Synthesised beam       &                &                      \\\\ \\quad major axis       &         $27.1$ & $\\mathrm{arcsec}$    \\\\ \\quad minor axis       &         $26.7$ & $\\mathrm{arcsec}$    \\\\ \\quad position angle   &   $87\\fdeg{}9$ &                      \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  So far, \\textsc{Duchamp}'s source finding and parametrisation capabilities have never been systematically tested on a large set of artificial sources with well-defined parameters. The aim of this paper is to bridge this gap by thoroughly testing the performance of \\textsc{Duchamp} on sets of artificial point sources and galaxy models as well as a data set containing real galaxies and telescope noise. The tests were originally motivated by the need to identify suitable source finding algorithms for the \\textit{Widefield ASKAP L-band Legacy All-sky Blind Survey} (WALLABY; \\citealt{koribalski2009}),\\footnote{Principal investigators: B\\\"{a}rbel Koribalski and Lister Staveley-Smith; public website: http://www.atnf.csiro.au\\slash{}research\\slash{}WALLABY/} one of the large, extragalactic ASKAP survey science projects currently in pre\\-paration \\citep{westmeier2010}. Hence, the tests presented here will focus on the detection of compact and extended \\ion{H}{i}~sources, in particular galaxies, in three-dimensional data cubes with ASKAP characteristics.  However, we believe that the results and conclusions presented in this paper will be of interest not only to those involved in SKA precursor science, but to a larger community of astronomers interested in the automatic detection and parametrisation of sources in their data sets, regardless of the wavelength range involved. For a comparison of \\textsc{Duchamp}'s performance with that of other source finding algorithms we refer to the paper by \\citet{popping2011} in this issue.  This paper is organised as follows: In Section~\\ref{sect_surveys} we summarise the source finding strategies of other large \\ion{H}{i}~surveys in the past, followed by a brief overview of the \\textsc{Duchamp} source finder in Section~\\ref{sect_duchamp}.  In Section~\\ref{sect_pointsources} we present the outcome of our test of \\textsc{Duchamp} on point source models with simple Gaussian line profiles. Section~\\ref{sect_galaxies} describes our testing of \\textsc{Duchamp} on models of disc galaxies with varying physical parameters. In Section~\\ref{sect_whisp} we apply \\textsc{Duchamp} to a data cube containing real galaxies and genuine noise extracted from radio observations. A discussion of our results is presented in Section~\\ref{sect_discussion}. Finally, Section~\\ref{sect_summary} summarises our main results and conclusions. ",
        "conclusions": "\\label{sect_discussion}  In general, \\textsc{Duchamp} does what it promises to do. It is able to reliably detect sources down to low signal-to-noise ratios and accurately determine their position and radial velocity. These are the most fundamental requirements for any source finder. Our tests also demonstrated that by using and fine-tuning the options of `\\`{a} trous' wavelet reconstruction and growing of sources to lower flux levels the performance of \\textsc{Duchamp} can be greatly enhanced.  \\subsection{Improving Reliability} \\label{sect_reliability}  The reliability figures reported throughout this paper have all been ``raw'' reliabilities, i.e.~reliabilities as achieved by \\textsc{Duchamp} prior to any filtering of the output source catalogue. The user would normally wish to substantially improve these through appropriate filtering of the source catalogue based on the source parameters as measured by \\textsc{Duchamp}.  The left-hand panel of Figure~\\ref{fig_false} shows the measured integrated flux plotted against measured line width for all genuine (black data points) and false (red data points) detections found by \\textsc{Duchamp} in the point source models discussed in Section~\\ref{sect_pointsources}. It is obvious that genuine and false detections occupy largely disjunct regions of $F_{\\rm int}$--$w_{50}$ parameter space, with false detections generally occurring near the low end of the integrated flux spectrum. Similar plots can be generated for other combinations of source parameters, but $F_{\\rm int}$ and $w_{50}$ usually provide the best distinction between genuine and false detections.  The easiest way to improve the reliability of \\textsc{Du\\-champ}'s source finding results is to simply apply a cut in $F_{\\rm int}$ to exclude most false detections while retaining most of the genuine sources. In our example, applying a cut at $F_{\\rm int} = 40~\\mathrm{mJy \\, km \\, s}^{-1}$ will discard $97.2\\%$ of all false detections while at the same time retaining $96.9\\%$ of all genuine sources, thereby increasing the overall reliability from $77.1\\%$ to $99.2\\%$ while only moderately decreasing the overall completeness from $83.0\\%$ to $80.6\\%$.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=\\linewidth]{fig_whisp.eps} \\caption{Completeness (filled circles, solid lines) and reliability (open circles, dashed lines) of \\textsc{Duchamp} on the WSRT model cube with WHISP galaxies for different flux thresholds and input parameters. The colours, as shown in the legend, distinguish the different wavelet reconstruction modes (one-dimensional versus three-dimensional) and wavelet reconstruction thresholds ($3 \\sigma$ versus $4 \\sigma$) used in the tests (see Table~\\ref{tab_input3} for details). The numbers alongside the data points refer to the corresponding runs as listed in Table~\\ref{tab_input3}.} \\label{fig_whisp} \\end{center} \\end{figure}  \\begin{figure}[t] \\begin{center} \\includegraphics[width=\\linewidth]{fig_completeness_3.eps} \\caption{Completeness as a function of integrated flux for selected runs (see legend) of \\textsc{Duchamp} on the WSRT model cube with WHISP galaxies. The choice of colours is the same as in Figure~\\ref{fig_whisp}.} \\label{fig_completeness_3} \\end{center} \\end{figure}  A similar cut can be applied to the results from run~7 on the test data cube containing artificially redshifted WHISP galaxies, as plotted in the right-hand panel of Figure~\\ref{fig_false}. Again, applying a simple flux threshold of $0.5~\\mathrm{Jy \\, km \\, s}^{-1}$ will improve reliability from $10\\%$ to $84\\%$, while only moderately decreasing completeness from $59\\%$ to $50\\%$. The method is not quite as successful as for the point sources, as we are now dealing with real galaxies and real noise with interference and artefacts, but nevertheless a significant improvement in reliability can be achieved without any severe impact on the number of genuine detections.  This simple example illustrates that the ``raw'' reliability figures quoted throughout this paper should not be considered as the final numbers. Reliability can be greatly improved through very basic filtering in parameter space of the \\textsc{Duchamp} output catalogue. In principle, this applies to the output of almost any source finder. Alternatively, instead of removing sources from the output catalogue, it may be desirable to calculate a reliability number for each catalogue entry based on the source's location in parameter space and leave it to the catalogue's users to decide as part of their scientific analysis at which reliability level they wish to make the cut.  \\begin{figure*}[t] \\begin{center} \\includegraphics[width=\\linewidth]{fig_false.eps} \\caption{Measured integrated flux, $F_{\\rm int}$, versus measured line width, $w_{50}$, of all genuine (black) and false (red) detections made by \\textsc{Duchamp} in the point source models with Gaussian line profiles (left) and the test cube with artificially redshifted WHISP galaxies (right). The dashed, black lines indicate the flux levels of $0.04$ and $0.5~\\mathrm{Jy \\, km \\, s}^{-1}$ used to filter false detections.} \\label{fig_false} \\end{center} \\end{figure*}  \\subsection{Source Parametrisation Issues}  When it comes to source parametrisation, the measurements provided by \\textsc{Duchamp} are affected by several systematic errors. These systematic errors are not due to errors in the software itself, but a consequence of the the presence of noise in the data as well as the methods and algorithms used for measuring source parameters.  Spectral line widths determined by \\textsc{Duchamp} are generally very accurate and not much affected by noise-induced, systematic errors as far as the $w_{50}$ parameter is concerned. The two other line width parameters calculated by \\textsc{Duchamp}, $w_{20}$ and $w_{\\rm vel}$, appear to be systematically too large over a wide range of signal-to-noise ratios and should not be used unless explicitly required in special, well-defined circumstances.  Peak fluxes, as reported by \\textsc{Duchamp}, are in general slightly too large for bright sources and significantly too large (on a relative scale) for faint sources. This is due to the fact that \\textsc{Duchamp} determines the peak flux by simply selecting the value of the brightest pixel encountered. This method introduces a bias towards positive noise peaks sitting on top of the brightest region of a source, and hence, in the presence of noise, peak fluxes measured by \\textsc{Duchamp} will be systematically too high.  Integrated fluxes determined by \\textsc{Duchamp} are significantly and systematically too small, in particular for faint sources. This is likely caused by the fact that \\textsc{Duchamp} simply sums over the flux of discrete elements above a given threshold to determine the integrated flux, thereby missing some of the flux from elements below the flux threshold. Hence, the raw integrated flux measurements currently provided by \\textsc{Duchamp} are not useful and need to be corrected to compensate for the systematic offset. This issue is particularly sensitive as many scientific projects, including the ASKAP survey science projects WALLABY and DINGO\\footnote{\\textit{Deep Investigation of Neutral Gas Origins}; principal investigator: Martin Meyer; public website: http:/\\slash{}internal.physics.uwa.edu.au\\slash{}$\\sim$mmeyer\\slash{}dingo/} \\citep{meyer2009}, rely on accurate flux measurements, for example for determining the \\ion{H}{i}~mass function of galaxies.  Finally, a particular problem in the case of galaxies is that under certain circumstances galaxies either get broken up into multiple detections or only one half of a galaxy is detected. This problem mainly affects faint, edge-on galaxies with broad spectral profiles that are partly hidden in the noise and results in systematic errors in the measurements of essentially all source parameters, including basic parameters such as position and radial velocity. Such cases of multiple or partial detections must be identified and treated separately to prevent biases in any scientific analysis based on the source finding results.    In this paper we present and discuss the results of basic, three-dimensional source finding tests with \\textsc{Du\\-champ}, the standard source finder for the Australian SKA Pathfinder, using different sets of unresolved and extended \\ion{H}{i}~model sources as well as a data set of real galaxies and noise obtained from \\ion{H}{i}~observations with the WSRT.  Overall, \\textsc{Duchamp} appears to be a successful, gen\\-eral-purpose source finder capable of reliably detecting sources down to low signal-to-noise ratios and accurately determining their position and velocity. In the case of point sources with simple Gaussian spectral lines we achieve a completeness of about 50\\% at a peak signal-to-noise ratio of~3 and an integrated flux level of about $0.1~\\mathrm{Jy \\, km \\, s}^{-1}$. The latter corresponds to an \\ion{H}{i} mass sensitivity of about $2 \\times 10^{8}~\\mathrm{M}_{\\odot}$ at a distance of $100~\\mathrm{Mpc}$ which is slightly better than what the WALLABY project is expected to achieve for real galaxies \\citep{koribalski2009}. The situation is less ideal for extended sources with double-horn profiles. In this case we achieve 50\\% completeness at an integrated flux level of about $2.5~\\mathrm{Jy \\, km \\, s}^{-1}$ for the model galaxies and $0.7~\\mathrm{Jy \\, km \\, s}^{-1}$ for the WHISP galaxies. The latter is equivalent to an \\ion{H}{i} mass sensitivity of about $1.7 \\times 10^{9}~\\mathrm{M}_{\\odot}$ at a distance of $100~\\mathrm{Mpc}$, illustrating that the performance of \\textsc{Duchamp}, as well as any other source finder, will strongly depend on source morphology. However, these figures may well be improved by carefully optimising the various input parameters offered by \\textsc{Duchamp}.  In its current state \\textsc{Duchamp} is not particularly successful in parametrising sources in the presence of noise in the data cube, and other, external algorithms for source parametrisation should be considered instead. It appears, however, that most, if not all, para\\-metrisation issues are due to intrinsic limitations in the implemented algorithms themselves and not due to errors in their implementation, suggesting that most of the problems can in principle be solved by implementing more sophisticated parametrisation algorithms in \\textsc{Duchamp}. Alternatively, corrections would have to be applied to all parameters derived by \\textsc{Duchamp} to compensate for systematic errors. Such corrections, however, would have to be highly specialised and tailored to the particular survey and source type concerned."
    },
    "1112/1112.0679_arXiv.txt": {
        "abstract": "We provide a quantitative description and statistical interpretation of the optical continuum variability of quasars. The Sloan Digital Sky Survey (SDSS) has obtained repeated imaging in five UV-to-IR photometric bands for 33,881 spectroscopically confirmed quasars. About 10,000 quasars have an average of 60 observations in each band obtained over a decade along Stripe 82 (S82), whereas the remaining \\about25,000 have 2--3 observations due to scan overlaps. The observed time lags span the range from a day to almost 10 years, and constrain quasar variability at rest-frame time lags of up to 4 years, and at rest-frame wavelengths from 1000\\AA\\ to 6000\\AA. We publicly release a user-friendly catalog of quasars from the SDSS Data Release 7 that have been observed at least twice in SDSS or once in both SDSS and the Palomar Observatory Sky Survey, and we use it to analyze the ensemble properties of quasar variability. Based on a damped random walk (DRW) model defined by a characteristic time scale and an asymptotic variability amplitude that scale with the luminosity, black hole mass, and rest wavelength for individual quasars calibrated in S82, we can fully explain the ensemble variability statistics of the non-S82 quasars such as the exponential distribution of large magnitude changes. All available data are consistent with the DRW model as a viable description of the optical continuum variability of quasars on time scales of \\about 5--2000~days in the rest frame. We use these models to predict the incidence of quasar contamination in transient surveys such as those from PTF and LSST. ",
        "introduction": "The optical continuum variability of quasars has been recognized since their first optical identification (Matthews \\& Sandage 1963), and it has been proposed and utilized as an efficient method for their discovery (van den Bergh, Herbst, Prit-\\\\ chet 1973; Hawkins 1983; Hawkins \\& Veron 1995; Ivezi\\'{c} et al.\\ 2004a; Rengstorf et al.\\ 2006).  The observed characteristics of the variability can then be used to constrain the origin of their emission (e.g., Kawaguchi et al.\\ 1998; Trevese, Kron \\& Bunone 2001, and references therein). The variability of quasars has typically been quantified using a structure function (SF) analysis (e.g., Hughes et al.\\ 1992; Collier \\& Peterson 2001; Bauer et al.\\ 2009; Koz{\\l}owski et al.\\ 2010b; Welsh, Wheatley, \\& Neil 2011), where the SF is the root-mean-square (rms) magnitude change (\\dm) as a function of the time lag ($\\Delta t$) between measurements (similar to an auto-correlation function). It is fairly well established that quasar variability properties depend on physical properties such as the quasar luminosity, wavelength, time scale, and the presence of radio emission. However, despite the considerable observational effort invested over last few decades, many conflicting claims about various correlations exist in the literature (see Giveon et al.\\ 1999 for a detailed discussion).  The traditional method for studying variability has been to monitor a small, select sample of quasars over a long time baseline (e.g., Hawkins 2002; Giveon et al.\\ 1999; Rengstorf et al.\\ 2004). In this case, it is possible to compute the SF for each quasar, which can later be sample-averaged or studied individually. An alternative, utilized in more recent studies based on Sloan Digital Sky Survey (SDSS; York et al.\\ 2000) and Palomar Observatory Sky Survey (POSS; Minkowski \\& Abell 1963) data, is to compute a single SF for all quasars in a particular wavelength or luminosity range.  This approach, mandated by the fact that typically only a few epochs were available per object, only measures {\\it ensemble} properties and {\\it assumes} that all quasars selected from a narrow range of physical properties vary in the same way. Nevertheless, this approach enabled studies of quasar optical variability based on tens of thousands of objects and several hundred thousand photometric observations, as well as explorations of the long-term variability (Vanden Berk et al.\\ 2004, hereafter VB04; Ivezi\\'{c} et al.\\ 2004c, hereafter I04; Wilhite et al.\\ 2005; Mahabal et al.\\ 2005; De~Vries, Becker, White, \\& Loomis 2005, hereafter dV05; Sesar et al.\\ 2006, hereafter Ses06). For example, the size and quality of the sample analyzed by VB04 (two-epoch photometry for 25,000  spectroscopically confirmed quasars) allowed them to constrain how quasar variability  in the rest frame optical/UV regime depends upon rest frame time lag (up to  $\\sim$2 years), luminosity, rest wavelength, redshift, the detection of radio or X-ray emission, and the presence of broad absorption line systems. By comparing SDSS and POSS measurements for \\about 20,000 quasars spectroscopically confirmed by the SDSS, Ses06 constrained the optical quasar variability on time scales from 10 to 50 years (in the observer's frame). They report that there is a characteristic time scale of order 1 year in the quasar rest frame beyond which the SF flattens to a constant value.  The SDSS has also facilitated both individual- and ensemble-based approaches by providing a large multi-epoch sample of quasars over the Northern Galactic Cap and well-sampled light curves in the Southern Stripe 82 (S82) survey.  It is reassuring that the two approaches lead to similar SFs, as discussed by dV05. A test of this assumption is also described in MacLeod et al.\\ (2008), who show that indeed the {\\it mean} behavior is the same. With such large samples, the ensemble SF(\\dt) slopes are well-constrained, and the values suggest that accretion disk instabilities are the most likely mechanism causing the observed optical variability (VB04; Kawaguchi et al.\\ 1998; see also Lyubarskii 1997). However, attempts to constrain physical models using the ensemble SF are invalid as soon as one realizes that the ensemble SF(\\dt) is a weighted sum of individual quasars with different structure functions (MacLeod et al.\\ 2008).  While studies have traditionally examined  ``non-parametric'' statistical measures of variability such as the SF, a major challenge has been to describe the variability of individual quasars in a compact way.  Recently, the introduction of a damped random walk (DRW) model has provided a way to mathematically characterize quasar light curves in terms of a characteristic time scale ($\\tau$) and an amplitude ($SF_{\\infty}$) which are then correlated with the physical properties such as luminosity and black hole mass. Kelly et al.\\ (2009, hereafter KBS09) modeled a sample of 100 quasar light curves as a DRW and suggested that thermal fluctuations driven by an underlying stochastic process such as a turbulent magnetic field may be the dominant cause for the optical flux fluctuations. Koz{\\l}owski et al.\\ (2010a; hereafter Koz{\\l}10) applied the DRW model to the well-sampled Optical Gravitational Lensing Experiment (OGLE) light curves (Udalski et al. 1997; Udalski et al. 2008) of mid-infrared-selected quasars behind the Magellanic Clouds from Koz{\\l}owski \\& Kochanek (2009). Their analysis shows that the DRW model is robust enough to efficiently select quasars from other variable sources, despite the large surface density of foreground Magellanic Cloud stars (see also Butler \\& Bloom 2011; MacLeod et al.\\ 2011; Koz{\\l}owski et al.\\ 2011a).  In MacLeod et al.\\ (2010, hereafter Mac10), we applied the DRW model to the light curves of \\about 10,000 quasars in S82 and found a correlation between $SF_{\\infty}$ and black hole mass which is independent of the anti-correlations with luminosity and wavelength (see also Ai et al.\\ 2010; Meusinger et al.\\ 2011). We also found that $\\tau$ increases with increasing wavelength, remains nearly constant with redshift and luminosity, and increases with increasing black hole mass (see also KBS09; Koz{\\l}10).  In Kelly et al.\\ (2011), it was shown that a similar stochastic model but with multiple time scales for a single object can accurately reproduce the X-ray variability of active galactic nuclei (AGN) and microquasars. An inhomogeneous accretion disk model, where the temperature fluctuations throughout the disk are driven by a DRW process, can explain the disk sizes derived from microlensing light curves (see Morgan et al.\\ 2010) while matching the observed level of optical variability, and predicts SEDs which are in better agreement with observations than standard thin disk models (Dexter \\& Agol 2011).  One defining feature of the DRW model for a {\\em single} quasar is that it predicts a Gaussian distribution of magnitude differences \\dm\\ for a given \\dt. On the other hand, the observed \\dm\\ distribution in the optical for an {\\em ensemble} of quasars observed at two times separated by \\dt\\ deviates strongly from a Gaussian but is well fit by an exponential distribution (I04).  This conflict represents an important puzzle for understanding the DRW model and its applicability to quasar light curves.  Also, the high likelihood of extreme values of \\dm\\ has important implications for the interpretation of observations of transients. For example, Vanden Berk et al.\\ (2002) reported the detection of an orphan gamma-ray burst afterglow based on the 2.5 magnitude decrease in optical flux. Such a large flux change was inconsistent with a quasar based on a Gaussian model for their variability, but the source was nevertheless confirmed to be a highly variable quasar (Gal-Yam et al.\\ 2002). An accurate statistical description of the two-epoch photometry for ensembles of quasars will be important for transient detection in large surveys, where quasars represent a major contaminant.  Our goal here is to produce a unified view of ensemble and individual optical variability in the context of the DRW model.  In this study, we show that the differences in shape between the ensemble SF and the DRW SF are well-explained by averaging over the properties of individual quasars. We also show that the exponential distributions of magnitude changes for ensembles of quasars at fixed time lag are naturally constructed by summing the intrinsically Gaussian distributions of magnitude changes produced by individual quasars. There are several residual issues which we discuss as part of the comparison. An overview of the SDSS and POSS data used in this study is presented in Section~\\ref{data}. In Section~\\ref{OV}, we describe the observed ensemble quasar variability in terms of the SF as a function of wavelength and time lag in the observer's frame.  We then convert to rest-frame quantities and compare the data to a model ensemble SF based on our previous DRW analysis of S82 light curves. In Section~\\ref{LT}, we combine the constraints on short-term quasar variability based on SDSS data with the constraints on long-term variability derived from matching the SDSS and POSS catalogs. In Section~\\ref{FS}, we discuss the implications our results have on transient identification, with a focus on future time-domain surveys such as the Palomar Transient Factory and the Large Synoptic Survey Telescope. Our results are discussed and summarized in Section~\\ref{Disc}. ",
        "conclusions": "\\label{Disc}  We have assembled, organized and publicly released a dataset including \\about 3.5 million photometric measurements for 80,000 spectroscopically confirmed quasars.  The available time lags span 0.8~days to almost 20~years in the observer's frame. We have analyzed and quantified the observed variability in the observer's and rest frames. By assuming a DRW model for each quasar in our sample, we reconcile the observed variability of individual quasars in S82 with their ensemble statistics. Our principal results are as follows: \\begin{enumerate} \\item  Long-term quasar variability measurements, constrained using SDSS and POSS data for time lags up to 50~years (in the observer's frame), conclusively show that a simple power-law dependence for the structure function cannot be extrapolated beyond a decade, and suggests an average characteristic time scale for quasar variability in the rest frame of $\\sim$2~years and an average long-term dispersion of $\\sim$0.26~mag (for rest wavelengths 2000--3000\\AA).  This behavior extrapolates well to the UV results of Welsh, Wheatley \\& Neil (2011), who find that the SF for GALEX NUV data reaches about 0.4~mag and flattens at $\\dt_{RF}>300$~days. This SF limit corresponds to a limiting SF value of 0.33~mag when using our definition of the SF (see Section~\\ref{SF}), and 0.27~mag when also scaling to the $u$-band using the wavelength dependence from Eq.~\\ref{eq:form}. This result is in close agreement with the SF at the shortest wavelengths in our data set. Voevodkin (2011) found that a broken power-law provides a good fit to the S82 ensemble SF with a slope of 0.33 at long time scales steepening to 0.79 below 42~days.  Our 2-epoch SDSS data are consistent with the shallower slope of 0.33, but our data do not support the conclusion of a much steeper SF(\\dt) for small \\dt\\ found by Voevodkin (2011). While we cannot rule out a broken power-law dependence with the available data, the observed SF is fully consistent with the form expected for a DRW (Eq.~\\ref{eq:sfdt}). \\item We tested the DRW model results based on SDSS Stripe 82 data on an independent dataset, and confirm that the variability parameters $\\tau$ and SF$_{\\infty}$ correlate with physical parameters as found for individual quasars (e.g., Mac10 and references therein).  This is evident from the agreement of our model with the observed ensemble variability of SDSS quasars.  However, the results indicate that the measured $\\tau$ and SF$_{\\infty}$ distributions are biased low for the S82 sample by a factor of about 2 and $\\sqrt{2}$, respectively. This bias most likely results from the 10-year length limit of the S82 light curves, although it could also be due to uncertain behavior for long time scales.  The best-fit $A$ coefficients from Eq.~\\ref{eq:form} (Mac10) need to be shifted upwards by 0.15 dex in the case of SF$_{\\infty}$, and by 0.30 dex in the case of $\\tau$, in order to explain the long time scale constraints provided by the SDSS--POSS dataset. These shifts leave the shorter time scale variability statistics unchanged. \\item For a given time lag and wavelength, the magnitude difference (\\dm) distribution is exponential rather than Gaussian for large magnitude changes. This is well explained as a cumulative effect of averaging over quasars with a range of different $\\tau$ and ${\\rm SF}_{\\infty}$. This is a remarkable result given that the \\dm\\ distribution of every individual quasar is Gaussian. \\item We made predictions for the incidence of quasar contamination in transient surveys using detailed simulations of quasar light curves from a mock LSST catalog. Due to the exponential nature of the \\dm\\ distributions for quasars, the probability of observing $\\dm>1$~mag reaches 0.02 in the $u$ band (where variability is strongest) for time lags of 300~days, and $6\\times 10^{-4}$ for $\\dm>2$~mag. Assuming Gaussian \\dm\\ distributions will result in erroneous likelihood estimates that are about 10 and 1000 times smaller, respectively. \\end{enumerate}  It is clear that a major limitation for the S82 quasars is the quality of light curves in both sampling density and time span. It is also clear that our variability model needs to be better tested given the evidence for a likely bias in the S82 time scale estimates. The best current sample for these improvements is that from the OGLE microlensing survey, since the light curves are more densely-sampled and longer than for S82. Here, the probem is the lack of spectroscopic identification of quasar candidates, although the follow-up confirmation of quasars is rapidly improving (Koz{\\l}owski et al.\\ 2011b). The next-generation surveys will also greatly improve the constraints on the long-term SF both individually and for ensembles of quasars. The best short-term prospects are PTF (Law et al.\\ 2009), Pan-STARRS (Kaiser et al.~2002), and the Dark Energy Survey (DES; Honscheid et al.~2008). In particular, the DES supernova program will greatly expand many of the S82 quasar light curves with $griz$ sampling once per week for $\\sim$3~months per year over 5~years. The combination of SDSS, Pan-STARRS, DES, and LSST will yield well-sampled light curves covering over 25 years for 10,000 quasars in Stripe 82.  The success of the model presented here suggests that a range of characteristic time scales exists among an ensemble of quasars, which can be related to physical time scales in the accretion disk. While we assumed a single $\\tau$ per quasar, there is evidence that multiple time scales can exist for a given quasar (Collier \\& Peterson 2001; Kelly et al.\\ 2011). Therefore, the study presented here can be extended to adopt the model in Kelly et al.\\ (2011), which fits more than one $\\tau$ for a given object. This model, also called a mixed Ornstein-Uhlenbeck (OU) process, reproduces a PSD of the form exhibited by the X-ray light curves of galactic black holes and AGN, which is flat below a low-frequency break, decays as $1/f$ above the low-frequency break, and steepens to $1/f^2$ above a high-frequency break.  In this case, with two characteristic time scales for each quasar, the long-term ensemble SF can be revisited and possibly explained in the context of a mixed OU process.  Note that recent optical data from the Kepler mission (Mushotzky et al.\\ 2011), which have a sampling of 1 data point roughly every 30 minutes and 0.1\\% errors, suggest an additional break to a steeper slope ($\\sim 1/f^3$), but this dependence is seen on time scales shorter than can be resolved in SDSS data.  \\vskip 0.4in \\leftline"
    },
    "1112/1112.0509_arXiv.txt": {
        "abstract": "\\emph{WMAP} data when combined with ancillary data on free-free, synchrotron and dust allow an improved understanding of the spectrum of emission from each of these components. Here we examine the sky variation at intermediate and high latitudes using a cross-correlation technique. In particular, we compare the observed emission in several large partitions of the sky plus 33 selected sky regions to three ``standard'' templates. The regions are selected using a criterion based on the morphology of these template maps.  The synchrotron emission shows evidence of steepening between GHz frequencies and the \\emph{WMAP} bands. There are indications of spectral index variations across the sky but the current data are not precise enough to accurately quantify this from region-to-region.  The emission correlated with the \\halpha template shows clear evidence of deviation from a free-free spectrum. The emission can be decomposed into a contribution from both free-free and spinning dust in the warm ionised medium of the Galaxy. The derived free-free emissivity corresponds to a  mean electron temperature of $\\sim 6000$~K, although the value depends critically on the impact of dust absorption on the  \\halpha intensity. The WIM spinning dust emission has a peak emission in intensity in the range 40--50~GHz.  The anomalous microwave emission associated with dust is detected at high significance in most of the 33 fields studied. The anomalous emission correlates well with the Finkbeiner et al. (1999) model 8 predictions (FDS8) at 94 GHz, and is well described globally by a power-law emission model with an effective spectral index between 20 and 60~GHz of $\\beta \\approx -2.7$.  It is clear that attempts to explain the emission by spinning dust models require multiple components, which presumably relates to a complex mix of emission regions along a given line-of-sight. An enhancement of the thermal dust contribution over the FDS8 predictions by a factor $\\sim 1.2$ is required with such models. Furthermore, the emissivity varies by a factor of $\\sim 50\\%$ from cloud to cloud relative to the mean.  The significance of these results for the correction of CMB data for Galactic foreground emission is discussed. ",
        "introduction": "A major goal of observational cosmology is to determine those parameters that describe the Universe. Observations of the Cosmic Microwave Background (CMB) at frequencies in the range 20 -- 200 GHz provide unique data to achieve this by establishing the statistical properties of temperature (and polarisation) measurements. However, an impediment to such studies arises due to foreground emission in our own Galaxy from at least three sources -- synchrotron, free-free and thermal dust emission. As CMB studies move to ever higher precision it is essential to determine the properties of these components to similarly high accuracy. Indeed, although the combined foreground level reaches a minimum in this frequency range ($\\nu \\approx 70$\\,GHz), it remains the dominant signal over large fractions of the sky. Of particular relevance to this discussion is the fact that each of the foreground components has a spectral index that varies from one line of sight to another so using a single spectral index can lead to significant uncertainties in the corrections required. It is therefore essential to study the nature of the Galactic signal at microwave wavelengths in its own right.  All-sky observations by the Wilkinson Microwave Anisotropy Probe \\citep[\\emph{WMAP},][]{Bennett_WMAP1:2003a} at the 5 frequencies of 23, 33, 41, 61 and 94 GHz can provide the basis for improving our understanding of local foregrounds.  By comparing these maps with templates for synchrotron, free-free and dust emission, made at frequencies where the specific emission mechanisms dominate, it is possible to clarify important properties of the emission. Indeed, new insights into the nature of the Galactic diffuse emission have arisen from studies of the \\emph{WMAP} data, including the detection of several unexpected new contributions.  Anomalous dust-correlated emission \\citep{Leitch:1997} was originally observed in the \\emph{COBE}-DMR data \\citep{Kogut_DMR:1996a} but was thought to be due to free-free emission. Draine \\& Lazarian (1998a,b) shifted attention to the dust itself as the source of emission through a mechanism now referred to as ``spinning dust'', or dipole emission from very rapidly spinning grains. A reanalysis of the intermediate and high Galactic latitude data taken by \\emph{COBE}-DMR and supplemented by 19~GHz observations \\citep{Banday_DMR:2003} led to evidence for dust at intermediate Galactic latitudes emitting a spectrum consistent with the the form expected for spinning dust, specifically a hint of a turnover at frequencies below $\\sim$ 20~GHz. However, study of the \\emph{WMAP} data has allowed further refinement of our understanding of the emission. Cross-correlation of the K-band data with observations at 15~GHz \\citep{dOC:2004} again indicated a plateau or downturn in foreground emission inconsistent with a free-free or synchrotron origin. \\citet{Lagache_WMAP:2003} compared the \\emph{WMAP} data to HI column density measurements,revealing an increase in emission with decreasing density, suggesting that the anomalous emission is connected to small, transient heated grains. In addition, a series of papers \\citep{Finkbeiner_WMAP1:2004, Dobler_WMAP3:2008a, Dobler_WMAP3:2008b} have strongly confirmed the presence of anomalous dust emission in the \\emph{WMAP} data, and claim to have found evidence of such a component from the diffuse warm ionised medium (WIM) of the Galaxy.  Specifically, the correlation with a H$_{\\alpha}$ template, commonly utilised as a proxy for the free-free emission, is not consistent with the spectrum expected for ionised gas, and a broad bump is seen peaking towards $\\sim$ 40~GHz. Subsequently, \\citet{Dobler_WMAP5:2009} have attempted to constrain specific spinning dust parameters such as the density and typical electric dipole moment of the grains. Recently, \\cite{Peel:2011} have shown that the K-band dust-correlated component is not strongly affected by the inclusion of a 2.3\\,GHz synchrotron template, reducing the possibility of a flat-spectrum synchrotron component.   Finally, the so-called \\emph{WMAP}-haze was identified by \\citet{Finkbeiner_WMAP1:2004} although it was already clearly apparent in the foreground residuals of \\citet{Bennett_WMAP1:2003b} and subsequently in the SMICA analysis of \\citet{Patanchon:2005}. The initial physical interpretation of the haze was that it was associated with free-free emission from hot gas in excess of $10^5$\\,K, but this was refuted on the basis of the lack of associated X-ray emission. The detection of the haze relies upon the use of standard templates to remove known foreground emission utilising spatially independent spectra over the entire (high-latitude) sky. It has been argued that such a crude approximation to the true behaviour of the foreground emission at microwave frequencies may well lead to unphysical results. Indeed, \\citet{Gold_WMAP7:2011} find that a spatial variation of spectral index of order 0.25 between 408 MHz and K-band is sufficient to reproduce the haze amplitude. Furthermore, a corresponding polarised signal was found to be absent from the \\emph{WMAP} data. However, the lack of polarised emission can be explained by the entanglement of the Galactic magnetic field towards the Galactic centre, leading to a lower polarisation fraction there as compared to the outer Galaxy. Nevertheless, the microwave haze remains an active area of research, particularly given its possible association with a gamma-ray counterpart observed in the \\emph{Fermi} data \\citep{Dobler_Fermi:2010}.  In this paper, we characterise the spatial variation in the foreground emission in terms of the various emission mechanisms noted above by introducing a new partition of the sky into morphologically selected regions.  In previous work \\citep[hereafter D06]{Davies_WMAP1:2006}, our approach was to identify regions away from the Galactic plane which were expected to be dominant in one of the three foreground components, free-free, synchrotron or dust and to derive the spectrum for each component. Five regions covering angular scales of $3\\deg$ to $20\\deg$ were chosen for each component, based on foreground template maps, making 15 in all. Here, we generalise this approach and introduce an algorithm to define a set of 35 regions for further study. The selection is intended to minimise the potential cross-talk between the various physical components and to select regions over which the spectral behaviour is uniform thus supporting the use of a template-based comparison. We use a well-known and understood cross-correlation technique to undertake the analysis.  The paper is organised as follows. Section~\\ref{sec:data} describes the \\emph{WMAP} data and foreground templates used in this analysis while Section~\\ref{sec:masks} defines the regions of interest for investigation. The methodology of the cross-correlation analysis is outlined in Section~\\ref{sec:methods} and the corresponding results are presented in Section~\\ref{sec:results}.  Model-dependent spectral fits for each component are considered in Section~\\ref{sec:modelfits} and overall conclusions given in Section~\\ref{sec:conclusions}. Appendix~\\ref{app:halpha} discusses in detail issues related to the \\halpha template used in the analysis, Appendix~\\ref{app:regions} defines the detailed method for partitioning the sky, and finally Appendix~\\ref{app:complete_results} tabulates all of the template-fit coefficients for all the regions analysed. ",
        "conclusions": "\\end{figure*}   \\subsection{Synchrotron}\\label{sec:results_sync}  \\begin{table} \\scriptsize \\begin{center} \\begin{tabular}{lccc} \\hline & \\multicolumn{3}{c}{\\bf Synchrotron spectral indices}\\\\ \\hline Region & K/408 &  Ka/408 & Q/408\\\\ \\hline EBV      &$\\bf   -2.90         ^{+    0.01         }_{   -0.01     }$&$\\bf   -2.91         ^{+    0.02         }_{   -0.02     }$&$\\bf   -2.92         ^{+    0.03         }_{   -0.03        }$\\\\ KQ85       &$\\bf   -2.91         ^{+    0.01         }_{   -0.01     }$&$\\bf   -2.93         ^{+    0.02         }_{   -0.02     }$&$\\bf   -2.94         ^{+    0.03         }_{   -0.04        }$\\\\ NPS      &$\\bf -3.03     ^{+ 0.03}_{-0.03}$&$                      -3.06 ^{+0.08}_{-0.12} $&$ < -2.90 $\\\\ GN       &$\\bf   -2.90         ^{+    0.01         }_{   -0.01     }$&$\\bf   -2.90         ^{+    0.02         }_{   -0.02     }$&$\\bf   -2.91         ^{+    0.04         }_{   -0.04        }$\\\\ GN$_{\\rm reduced}$       &$\\bf   -2.88         ^{+    0.01         }_{   -0.01     }$&$\\bf   -2.88         ^{+    0.02         }_{   -0.02     }$&$\\bf   -2.88         ^{+    0.03         }_{   -0.04        }$\\\\ EN      &$\\bf   -2.95         ^{+    0.01         }_{   -0.01     }$&$\\bf   -2.99         ^{+    0.03         }_{   -0.03        }$&$   -3.04         ^{+    0.06         }_{   -0.09        }$\\\\ EN$_{\\rm  reduced }$       &$\\bf   -2.94         ^{+    0.01         }_{   -0.01     }$&$\\bf   -2.97         ^{+    0.03         }_{   -0.04     }$&$  -3.02         ^{+    0.06         }_{   -0.09        }$\\\\ GS       &$\\bf   -2.91         ^{+    0.01         }_{   -0.01     }$&$\\bf   -2.91         ^{+    0.02         }_{   -0.02     }$&$\\bf   -2.92         ^{+    0.04         }_{   -0.05        }$\\\\ ES       &$\\bf   -2.90         ^{+    0.01         }_{   -0.01     }$&$\\bf   -2.90         ^{+    0.02         }_{   -0.02     }$&$\\bf   -2.89         ^{+    0.03         }_{   -0.03        }$\\\\ \\hline 1       &$ <    -2.96      $&$ <    -2.88      $&$ <    -2.80         $\\\\ 2          &$ \\bf  -2.94         ^{+    0.04         }_{   -0.05     }$&$ <    -2.79      $&$ <    -2.74         $\\\\ 3          &$ \\bf   -3.03         ^{+    0.03         }_{   -0.03        }$&$   -3.04         ^{+    0.07         }_{   -0.10     }$&$ <    -2.89         $\\\\ 4       &$ <    -2.77      $&$ <    -2.60      $&$ <    -2.50         $\\\\ 5       &$ <    -2.75      $&$ <    -2.62      $&$ <    -2.50         $\\\\ 6          &$   -3.11         ^{+    0.10         }_{   -0.16     }$&$ <    -2.87      $&$ <    -2.79         $\\\\ 7          &$   -3.06         ^{+    0.09         }_{   -0.15     }$&$ <    -2.84      $&$ <    -2.78         $\\\\ 8          &$  \\bf  -2.86         ^{+    0.01         }_{   -0.01        }$&$ \\bf   -2.88         ^{+    0.04         }_{   -0.04        }$&$ \\bf   -2.89         ^{+    0.06         }_{   -0.09        }$\\\\ 9          &$   -3.15         ^{+    0.09         }_{   -0.13     }$&$ <    -2.91      $&$ <    -2.84         $\\\\ 10          &$ \\bf   -2.90         ^{+    0.06         }_{   -0.09        }$&$   -2.76         ^{+    0.08         }_{   -0.14        }$&$   -2.64         ^{+    0.08         }_{   -0.14        }$\\\\ 11          &$ \\bf  -2.78         ^{+    0.02         }_{   -0.03        }$&$   -2.88         ^{+    0.08         }_{   -0.14     }$&$ <    -2.76         $\\\\ 12          &$  \\bf  -2.92         ^{+    0.06         }_{   -0.08     }$&$ <    -2.72      $&$ <    -2.63         $\\\\ 13       &$ <    -3.00      $&$ <    -2.93      $&$ <    -2.85         $\\\\ 15          &$ \\bf   -2.84         ^{+    0.04         }_{   -0.04     }$&$ <    -2.78      $&$ <    -2.80         $\\\\ 16          &$ \\bf   -2.98         ^{+    0.03         }_{   -0.03        }$&$   -2.97         ^{+    0.07         }_{   -0.09     }$&$ <    -2.82         $\\\\ 17          &$ \\bf   -2.94         ^{+    0.04         }_{   -0.05     }$&$ <    -2.85      $&$ <    -2.82         $\\\\ 18          &$ \\bf   -2.90         ^{+    0.06         }_{   -0.07     }$&$ <    -2.68      $&$ <    -2.59         $\\\\ 19       &$ <    -2.85      $&$ <    -2.67      $&$ <    -2.58         $\\\\ 20       &$ <    -2.82      $&$ <    -2.65      $&$ <    -2.53         $\\\\ 21          &$ \\bf   -2.98         ^{+    0.07         }_{   -0.09     }$&$ <    -2.78      $&$ <    -2.70         $\\\\ 22          &$   -2.98         ^{+    0.09         }_{   -0.15     }$&$ <    -2.69      $&$ <    -2.59         $\\\\ 23       &$ <    -2.94      $&$ <    -2.87      $&$ <    -2.81         $\\\\ 24          &$ \\bf   -2.80         ^{+    0.04         }_{   -0.05        }$&$   -2.71         ^{+    0.07         }_{   -0.10        }$&$   -2.65         ^{+    0.08         }_{   -0.14        }$\\\\ 25       &$ <    -2.92      $&$ <    -2.76      $&$ <    -2.65         $\\\\ 26          &$   -2.82         ^{+    0.10         }_{   -0.16     }$&$ <    -2.53      $&$ <    -2.45         $\\\\ 27          &$   -2.77         ^{+    0.08         }_{   -0.11     }$&$ <    -2.51      $&$ <    -2.42         $\\\\ 28       &$ <    -2.79      $&$ <    -2.66      $&$ <    -2.57         $\\\\ 29          &$   -2.80         ^{+    0.10         }_{   -0.17     }$&$ <    -2.50      $&$ <    -2.40         $\\\\ 30       &$ <    -2.82      $&$ <    -2.67      $&$ <    -2.57         $\\\\ 31          &$   -2.87         ^{+    0.08         }_{   -0.13     }$&$ <    -2.70      $&$ <    -2.67         $\\\\ 32          &$ \\bf   -2.95         ^{+    0.04         }_{   -0.04        }$&$   -2.95         ^{+    0.08         }_{   -0.14     }$&$ <    -2.76         $\\\\ 33          &$   -2.89         ^{+    0.08         }_{   -0.11     }$&$ <    -2.62      $&$ <    -2.52         $\\\\ \\hline Average  &$ \\bf  -2.95         ^{+    0.01         }_{   -0.01 }$&$ \\bf   -2.97   ^{+    0.03         }_{   -0.03     }   $&$ \\bf   -3.00 ^{+    0.05         }_{   -0.06     }         $\\\\ \\hline \\end{tabular} \\end{center} \\caption{Synchrotron fits between 408 MHz and \\emph{WMAP }K-, Ka- and Q-bands for various regions of the sky. For those regions that are detected at 2-$\\sigma$ significance, the fit values from Table~\\ref{tab:results_synch_33regions} are converted into a spectral index using the usual power-law relation $\\beta$ ($T_b \\propto \\nu^{\\beta}$). Associated errors are determined by using the fit values plus and minus the 1$\\sigma$ error bars.  Otherwise, one-sided 95\\% confidence level upper limits on the index are quoted based on the fit value plus  1.64$\\sigma$. Bold text denotes detections at 3-$\\sigma$ significance or more.  For the global fits, NPS -- North Polar Spur, GN -- Galactic North, GN$_{\\rm reduced}$ -- Galactic North with the NPS removed, EN -- Ecliptic North, EN$_{\\rm reduced}$ -- Ecliptic North with the NPS removed, GS -- Galactic South, ES -- Ecliptic South.} \\label{tab:simple_synch_spectra} \\end{table}  As expected, the template fit coefficients between the Haslam 408 MHz data and the \\emph{WMAP} sky maps all fall with frequency in a manner consistent with power-law emission.  In Table~\\ref{tab:simple_synch_spectra} we define simple pairwise spectral indices between 408 MHz and the K-, Ka- and Q-bands. The global fits exhibit a typical index of order $-2.90$ at K-band, with the results from all masks slightly flatter than the value of $-3.01$ found in the \\citetalias{Davies_WMAP1:2006} analysis of the \\emph{WMAP} first-year data for the then-preferred Kp2 sky-coverage. Analyses of the lower-frequency surveys at 408, 1420 and 2326\\,MHz by \\citet{Giardino_synch_pol:2002} and \\citet{Platania_indices:2003} suggest a spectral index over this lower frequency range of approximately $-2.7$, thus our results support the idea of spectral steepening, continuing beyond K-band. We will consider this further in Section~\\ref{sec:modelfits}.  However, there are differences in the coefficients depending on the exact sky coverage that must reflect genuine spectral variations on the sky. The NPS is recognised to be an arc of steep spectrum emission at lower frequencies, thus it is not surprising that it is notably steeper than the rest of the high latitude sky.  The spectral index value of $-3.03$ at K-band is quite consistent with the value adopted in \\citet{Finkbeiner_WMAP1:2004} to remove the prominent emission from the \\emph{WMAP} data. The presence of the NPS  also impacts measures of the spectral index in both the Galactic and Ecliptic northern sky, resulting in a modest steepening of the index. Interestingly, the northern Ecliptic hemisphere is notably steeper than the other hemispheres, and exhibits an increasingly steep index with frequency. Conversely, the corresponding southern hemisphere hints at spectral flattening, whereas both the north and south Galactic hemispheres are consistent with simple  power-law behaviour.  From the 33 regions of interest, there are 13 regions where the synchrotron fit coefficients are detected at greater than $3 \\sigma$ confidence at K-band. Most of these regions include contributions close to the Galactic plane, although regions 2, 3 and 32 are mostly at high Galactic latitude. The inferred spectral indices span the range $-2.78$ to $-3.03$, inconsistent with statistical variation alone, and likely representing genuine spectral variations on the sky. We note that region 3  contains the NPS and its behaviour seems to be dominated by that component. Significant emission is detected for region 8 at K-, Ka- and Q-band. Regions 13 and 14 correspond to an area of the sky containing the Gum nebula, and show no evidence for detection of synchrotron emission, with upper limits consistent with a steep spectrum, particularly in the southern region. Interestingly, region 9 has one of the steepest spectral indices on the sky, despite the putative presence of the \\emph{WMAP} haze. However, the overlap between the brightest regions of the haze emission as seen in \\citet{Dobler_WMAP3:2008a} and regions 9, 16 and 17 is very small, and unlikely to affect any studies here.  It is interesting that the spectral indices inferred from the mean of the regional fit coefficients are steeper than the typical global fit values. This trend was also seen in our previous work \\citepalias{Davies_WMAP1:2006}, and may be due to a selection effect in that the regional subdivisions partly favour stronger synchrotron emission regions which may exhibit steeper spectra than normal due to synchrotron losses. In addition, evidence of spectral steepening with frequency is again generally seen although regions 10 and 24 show inconsistent behaviour with the other regions in that a flattening of the spectrum is indicated.   \\subsection{Free-free}\\label{sec:results_freefree}  In this paper, we use an H$_{\\alpha}$ template as a proxy for the free-free emission. We see significant correlation between the \\emph{WMAP} data and the template for the global fits at all frequencies (Table~\\ref{tab:results_freefree_33regions}). As with the synchrotron results, there are interesting variations depending on the exact sky coverage, with the northern Ecliptic hemisphere showing significantly enhanced amplitude, whilst the south indicates a lower emissivity. 14 individual regions are detected at 3$\\sigma$ significance at K-band. Most of these regions lie close to the Galactic plane ($|b|\\lesssim 20^{\\circ}$).   \\citetalias{DDD:2003} detail the relationship between the expected free-free brightness temperature and the related H$_{\\alpha}$ intensity, and its dependence on both frequency and electron temperature ($T_e$) in the ionised medium. As can be seen from their Fig.~5, the spectral dependence of the emission shows weak curvature, but over the range of frequencies covered by \\emph{WMAP} a reasonable approximation is a power-law with index $-2.15$.  Inspection of the coefficients in Table~\\ref{tab:results_freefree_33regions} indicates that there are departures from this behaviour. Following \\citet{Dobler_WMAP3:2008b} we plot these results in intensity units in Fig.~\\ref{fig:results_summary}. It should be apparent that a bump in emission is seen around 40--50~GHz for the global fits, a feature that \\citet{Dobler_WMAP3:2008b} argue is indicative of contributions from both classical free-free emission plus a spinning dust component in the WIM. The mean spectrum of the regional fits perhaps indicates a slightly broader bump in the spectrum. However, it is also the case that there is a range of behaviour seen amongst the individual regions, some of which are consistent with emission from a single physical component only - either free-free emission or a more steeply falling spectrum as expected from spinning dust. In Section~\\ref{sec:modelfits}, we will undertake a more detailed modelling of the emission in terms of these components.  Such a feature in the spectrum of course has implications for the determination of physical parameters such as $T_e$. Nevertheless, we can make some general inferences, particularly by examining the K- and W-band amplitudes that are least affected by a putative WIM spinning dust component. The global fits seem to be consistent with values in the range 6000 -- 8000\\,K, with some dependency on the exact mask used. The average free-free electron temperature inferred from the 33 regions is also in this range, although for individual regions there is a spread of values between 4000 and more than 10000\\,K. These values are somewhat higher than seen previously in \\citet{Davies_WMAP1:2006}, and this is due to the use of 3$^{\\circ}$ smoothed data here, rather than the 1$^{\\circ}$ resolution data used earlier for reasons provided in Appendix~\\ref{app:halpha}. It is also interesting to note that the 7000\\,K temperature inferred from the 31.5 GHz $COBE$-DMR data at  7$^{\\circ}$ resolution in \\citet{Banday_DMR:2003} is quite consistent with the Ka-band values determined here. Moreover, the higher values in this paper are in better agreement with the electron temperatures derived from radio recombination line studies of extended HII regions \\citep{Shaver:1983,Paladini_HII:2004,Alves:2011} which derive an average $T_e$ value of $\\approx 7000$\\,K in the vicinity of the solar neighbourhood. There is therefore no need to invoke a large fraction of scattered \\halpha light to account for this discrepancy (e.g., \\citealt{Witt_Scattered_Halpha:2010}).  \\subsection{Impact of Dust Extinction}\\label{sec:dust_extinct}  \\begin{figure*} \\begin{tabular}{ccc} \\epsfig{file=figures/fig_41.ps,width=0.3\\linewidth,angle=0,clip=}& \\epsfig{file=figures/fig_42.ps,width=0.3\\linewidth,angle=0,clip=}& \\epsfig{file=figures/fig_43.ps,width=0.3\\linewidth,angle=0,clip=}\\\\ \\epsfig{file=figures/fig_44.ps,width=0.3\\linewidth,angle=0,clip=}& \\epsfig{file=figures/fig_45.ps,width=0.3\\linewidth,angle=0,clip=}& \\epsfig{file=figures/fig_46.ps,width=0.3\\linewidth,angle=0,clip=}\\\\ \\end{tabular} \\caption{Dependence of the correlation coefficients between the \\emph{WMAP} data and the \\halpha template as a function of the absorption correction $f_{d}$. The union mask refers to the combination of all 33 small regions defined in this paper. Thus the results for ``Union + Galactic North'' refers to the inverse-noise-weighted average coefficients for the subset of regions in the northern Galactic hemisphere. The observed dependencies for regions 13, 14 and 23 are representative of that seen for all regions with a significant detection amplitude at K-band.} \\label{fig:coefficients_fd} \\end{figure*}  \\citet{Bennett_WMAP1:2003b} summarise various uncertainties in the use of the H$_{\\alpha}$ template to trace free-free emission in the Galaxy. These include uncertainties in the value of the electron temperature $T_e$, and in the value of the dust absorption correction, specified here by the $f_d$ parameter. In principle, there will be variations in both throughout the Galaxy. As noted previously, \\citetalias{DDD:2003} have determined that for local regions such as Barnard's Arc and the Gum nebula, there is little absorption by dust, and this is the default situation that we have assumed in our analyses. However, for mid-Galactic latitudes a value $f_{d} \\sim$ 0.3 is preferred, and \\citet{Finkbeiner:2003} adopted the assumption that the H$_{\\alpha}$ emission is co-extensive with dust emission along the line-of-sight, ie. $f_{d} \\sim$ 0.5.  In this section we discuss the impact of varying $f_d$ on the template fit coefficients. In Fig.~\\ref{fig:coefficients_fd} we show this variation as a function of $f_d$, and compare against the behaviour expected for a range of values of $T_e$. In general, we expect that as the amplitude of the H$_{\\alpha}$ template is corrected upwards by increasing $f_d$, then the fit coefficients will decrease, implying a lower value for the electron temperature. This is indeed what is seen, but the extent of the correction depends on the region of sky under consideration.  For the largest sky area that we consider, the EBV mask, the use of a template corrected absorption specified by $f_{d} = 0.5$ results in coefficients consistent with values of $T_e$ of order 3000\\,K. However, assuming no dust absorption correction yields values closer to 6000\\,K.  A similar result is seen for the southern extension of the Gum nebula (region 14). If it is indeed the case that there is no evidence for dust absorption, then a higher temperature of 6000\\,K is determined. Conversely, regions 15, 18 and 21 yield coefficients at K-band consistent with temperatures closer to 10000\\,K if no absorption correction is applied, whereas values of order 8000\\,K are found for $f_d = 0.5$. Given that these regions exhibit rising spectra that may require the presence of significant emission from a spinning dust component, then the latter value would be more consistent. However, there are also regions of the sky typified by region 23, where apparently acceptable values of the electron temperature are inferred over a range of values of $f_d$. Indeed, it may be that the large spread in coefficients seen for the different regions reflects changes in the fraction of dust mixed with the WIM as much as variations in $T_e$.  It does appear, therefore, that there are a range of values for $T_e$ and $f_d$ throughout the Galaxy, and reaching conclusions about their values solely from studies of H$_{\\alpha}$ correlations suffers from degeneracies between the parameters. The adoption of a value for the dust absorption correction of 0.5 in, for example,  \\citet{Dobler_WMAP3:2008b}, thus seems to be associated, at least in part, with the low electron temperatures inferred. This uncertainty then has implications for modelling the emission of the diffuse component.   A  more serious complication would arise if the validity of using  H$_{\\alpha}$ as a tracer of free-free emission were questioned. \\citet{Mattila_Scattered_Halpha:2007} have argued that the H$_{\\alpha}$ excess towards the high-latitude interstellar cloud LDN~1780 is the result of scattering of  H$_{\\alpha}$ photons produced elsewhere in the Galaxy by dust in the cloud. Earlier computations by \\citet{Wood_Scattered_Halpha:1999} suggested that this contribution would typically be 5--20\\% at high-latitudes, although their model has been criticised due to the assumption of a smooth distribution of material in the ISM. \\citet{Witt_Scattered_Halpha:2010} have used an empirical relation to relate the scattered H$_{\\alpha}$ intensity in the translucent cloud LDN1780 to the \\emph{IRAS} 100~$\\mu$m diffuse background intensity and conclude that this estimate is reasonable for 50\\% of the high-latitude sky, but that the scattered contribution can be highly structured and result in contributions of between 25 and 50\\% of the observed intensity for a further 40\\% of this region of the sky. Such a result would clearly have implications for using an H$_{\\alpha}$ template to trace free-free emission, and a complex relationship between the template and dust would arise both due to the scattering contribution and due to any dust absorption correction applied to the data. Applying a correction for the former would effectively result in an increase of the  H$_{\\alpha}$-\\emph{WMAP} correlation coefficients and therefore of the inferred electron temperature, whereas application of the latter results in the opposite behaviour. More recently, \\citet{Brandt_Scattered_Halpha:2011} have estimated the fraction of high-latitude H$_{\\alpha}$ that is scattered to be $19\\pm 4\\%$, a value consistent with that proposed by \\citet{Dong_WIM_Halpha:2011} to reconcile the low ratio of radio free-free to  H$_{\\alpha}$.  Ultimately, unravelling the degeneracy between the electron temperature, H$_{\\alpha}$ scattering and dust absorption requires additional observations. Detailed RRL surveys in the Galactic plane together with radio-continuum surveys at frequencies of $\\sim$5~GHz, as expected from the C-BASS \\citep{King_CBASS:2010} will be important in this respect.  Finally, we would like to make some remarks about region 11, the coefficients of which suggest an exceptionally high temperature of more than 25000\\,K. If this result were considered unphysical, then naively a value of $f_d = 1$ would be required in order to lower the inferred temperature to the 6000\\,K seen in the EBV fit. In fact, a more realistic assessment of the situation is that the EBV mask is not large enough to eliminate some parts of region 11 close to the Galactic plane where the simple dust absorption correction is untrustworthy. If instead we apply the KQ85 mask before analysing the region, then for a more plausible value of $f_{d} = 0.5$ the inferred electron temperature is again consistent with 6000\\,K.   \\subsection{Dust}\\label{sec:results_dust}  As can be seen in the right-hand panels of Fig.~\\ref{fig:results_summary}, the template fit coefficients determined between the \\emph{WMAP} data and the FDS8 dust template prediction at W-band are consistent with a rising thermal dust contribution at frequencies higher than V-band, and a rising spectrum to lower frequencies below it. The latter corresponds to the now widely identified anomalous microwave emission (AME).  It appears that the FDS8 template underpredicts the W-band amplitude by approximately 30\\% for the global fits, consistent with the results of \\citetalias{Davies_WMAP1:2006}. A broad range of values for the individual regions is seen, but only three detect emission at a statistically significant level. The mean emissivity at W-band of all regions shows a more modest enhancement, but is nevertheless consistent with the FDS8 predictions, as indeed are the three significant regions.  At K-band,  there are variations in the global fit amplitudes depending on  sky coverage. The KQ85 mask indicates a lower emissivity compared to the EBV as might be expected. The Southern Galactic and Ecliptic hemispheres have the highest amplitudes. All coefficients are higher than those determined for the Kp2 sky coverage in \\citetalias{Davies_WMAP1:2006}. 26 of the individual regions detect emission at 3$\\sigma$ significance or higher. The mean emission amplitude lies between that for EBV and KQ85,  and again somewhat higher than in \\citetalias{Davies_WMAP1:2006}.  The regions indicate a variation around the mean of approximately 50\\% of its amplitude, inconsistent with statistical errors alone and indicating genuine spatial variations in the AME emissivity.  The overall spectrum would appear to be well-described by a superposition of two power-law emissivities. The thermal dust emission described by the FDS8 model is adequately represented by an emissivity index $\\beta_d = 1.55$ over the \\emph{WMAP} range of frequencies. We fit the power-law AME spectrum with the thermal dust index fixed to this value, and find AME spectral indices of order $-2.7$ for both the global fits and regional mean.  \\citet{Draine_spinning:1998} first proposed that the AME could be explained by electric dipole radiation from rotationally excited small interstellar grains, or spinning dust. Our results have ramifications for such models of the emission. In particular, given that the spinning dust spectra typically fall off steeply with frequency beyond their peak, then it is unlikely that a single such spectrum could account adequately for the effective power-law emission. Indeed, the observed spectrum is presumably formed from a superposition of components with varying spectra as a consequence of their differing physical environments along a given line-of-sight. We will discuss detailed fits of the observed emission in Section~\\ref{sec:modelfits}, and their implications for more refined models of the AME.    The following analytic forms are used to fit the synchrotron coefficients. \\begin{itemize}  \\item Model SI : Given the power law energy distribution of cosmic ray electrons, we assumed a power-law emissivity in terms of brightness temperature over the WMAP frequencies as \\begin{equation} T_A (\\nu) = A_{s} \\times \\left( \\frac{\\nu}{23} \\right )^{\\beta_{s}}_{\\text{GHz}} \\end{equation} where $\\beta_s$ is the spectral index and $A_s$ is the normalised amplitude expressed in $\\mu \\text{K}$ with respect to the frequency $\\nu_0 = 23$\\,GHz.   \\item Model SII : A power-law emissivity is assumed to extend from 408~MHz up to and through the WMAP frequencies. Since we use the 408~MHz survey as a template for the synchrotron emission, the amplitude at the low frequency must be reproduced perfectly. This results in an effective constraint to be applied to the analytical form above, and we then fit for the spectral index $\\beta_{s}$ only.  \\begin{equation} T_A (\\nu) = 10^6 \\times \\left( \\frac{\\nu}{0.408} \\right )^{\\beta_{s}}_{\\text{GHz}} \\end{equation}  \\item Model SIII : The cosmic ray electron energy spectrum is expected to steepen with time due to radiation energy loss. A review of cosmic-ray propagation including electrons can be found in \\citet{Strong_etal:2007}, whilst \\citet{Strong_etal:2011} directly test propagation models based on cosmic-ray and gamma-ray data against synchrotron data from 22~MHz to 94~GHz as averaged over mid-latitude regions ($10^{\\circ} < \\mid b \\mid < 45^{\\circ}$). The latter analysis confirms the need for a low-energy break in the cosmic-ray electron injection spectrum to account for the steepening synchrotron spectrum. Since we do not include synchrotron information at frequencies intermediate to 408~MHz and the \\emph{WMAP} data, we follow the treatment of \\citet{Gold_WMAP5:2009}. Specifically, the emissivity is assumed to follow a power-law from 408~MHz until K-band and then to exhibit spectral curvature as follows, \\begin{align*} T_A (\\nu) &= A_{s} \\times \\left( \\frac{\\nu}{23} \\right )_{\\text{GHz}}^{\\beta_{s}}  &\\nu < \\nu_\\text{K} \\\\ &= A_{s} \\times \\left( \\frac{\\nu}{23} \\right )_{\\text{GHz}}^{\\beta_{s} + c_s ln(\\frac{\\nu}{\\nu_\\text{K}})}  & \\nu > \\nu_\\text{K} \\end{align*} As above, the 408~MHz point is fixed, thus $A_s$ can be written in terms of $\\beta_s$, and we are left to fit this spectral index and the curvature $c_s$. For a WMAP frequency point, the above equation reduces to the form: \\begin{equation} T_A (\\nu) = 10^6 \\times \\left( \\frac{\\nu}{0.408} \\right )^{\\beta_s}_{\\text{GHz}} \\times \\left( \\frac{\\nu}{23} \\right )^{ c_s ln(\\frac{\\nu}{\\nu_K})}_{\\text{GHz}} \\end{equation}  \\end{itemize}  \\begin{table} \\scriptsize \\begin{center} \\begin{tabular}{lcccccccc} \\hline &\\multicolumn{3}{c}{SI} &\\multicolumn{2}{c}{SII} &\\multicolumn{3}{c}{SIII}\\\\ \\hline Region & $A_{s}$ & $\\beta_{s}$ & $\\chi^2$   & $\\beta_{s}$  & $\\chi^2$ & $\\beta_{s}$ & $c_s$ & $\\chi^2$  \\\\ \\hline EBV    &    7.96 $\\pm$    0.18    &   -2.99 $\\pm$    0.14    &   0.047    &   -2.91 $\\pm$    0.01    &   0.106    &   -2.91 $\\pm$    0.01    &   -0.17 $\\pm$    0.28    &   0.011   \\\\ KQ85    &    7.64 $\\pm$    0.20    &   -3.11 $\\pm$    0.19    &   0.246    &   -2.92 $\\pm$    0.01    &   0.470    &   -2.92 $\\pm$    0.01    &   -0.44 $\\pm$    0.39    &   0.071   \\\\ NPS     &    4.68 $\\pm$    0.55    &   -3.43 $\\pm$    0.92    &   0.021   &   -3.05 $\\pm$    0.03    &   0.069  &   -3.04 $\\pm$    0.03    &   -0.91 $\\pm$    2.18    &   0.003   \\\\                                                                 GN    &    8.16 $\\pm$    0.25    &   -2.94 $\\pm$    0.19    &   0.007    &   -2.91 $\\pm$    0.01    &   0.013    &   -2.91 $\\pm$    0.01    &   -0.06 $\\pm$    0.35    &   0.006   \\\\ EN    &    6.41 $\\pm$    0.26    &   -3.55 $\\pm$    0.33    &   1.077    &   -2.97 $\\pm$    0.01    &   1.914    &   -2.96 $\\pm$    0.01    &   -1.38 $\\pm$    0.81    &   0.628   \\\\ GS    &    7.84 $\\pm$    0.26    &   -2.99 $\\pm$    0.21    &   0.082    &   -2.92 $\\pm$    0.01    &   0.093    &   -2.92 $\\pm$    0.01    &   -0.20 $\\pm$    0.42    &   0.031   \\\\ ES    &    7.95 $\\pm$    0.22    &   -2.74 $\\pm$    0.16    &   0.337    &   -2.91 $\\pm$    0.01    &   0.504    &   -2.91 $\\pm$    0.01    &    0.37 $\\pm$    0.23    &   0.075   \\\\ \\hline 2    &    6.82 $\\pm$    1.25    &   -3.05 $\\pm$    1.23    &   0.018    &   -2.95 $\\pm$    0.04    &   0.015    &   -2.95 $\\pm$    0.05    &   -0.34 $\\pm$    2.59    &   0.012   \\\\ 3    &    4.66 $\\pm$    0.53    &   -3.03 $\\pm$    0.76    &   0.006    &   -3.04 $\\pm$    0.03    &   0.005    &   -3.05 $\\pm$    0.03    &    0.08 $\\pm$    1.30    &   0.005   \\\\ 8    &    9.23 $\\pm$    0.53    &   -3.01 $\\pm$    0.38    &   0.002    &   -2.88 $\\pm$    0.01    &   0.037    &   -2.87 $\\pm$    0.01    &   -0.25 $\\pm$    0.72    &   0.005   \\\\ 10    &    7.67 $\\pm$    2.04    &   -0.56 $\\pm$    0.42    &   0.076    &   -2.86 $\\pm$    0.05    &   2.305    &   -2.87 $\\pm$    0.05    &    1.69 $\\pm$    0.28    &   0.700   \\\\ 11    &   12.61 $\\pm$    1.24    &   -4.33 $\\pm$    1.16    &   0.560    &   -2.81 $\\pm$    0.02    &   1.403    &   -2.79 $\\pm$    0.03    &   -3.67 $\\pm$    3.07    &   0.441   \\\\ 12    &    7.33 $\\pm$    1.88    &   -2.09 $\\pm$    1.03    &   0.028    &   -2.92 $\\pm$    0.06    &   0.121    &   -2.93 $\\pm$    0.06    &    1.01 $\\pm$    0.88    &   0.010   \\\\ 15    &    9.74 $\\pm$    1.62    &   -5.38 $\\pm$    2.78    &   1.284    &   -2.88 $\\pm$    0.04    &   1.801    &   -2.85 $\\pm$    0.04    &   -5.99 $\\pm$    7.08    &   1.194   \\\\ 16    &    5.88 $\\pm$    0.62    &   -2.96 $\\pm$    0.69    &   0.002    &   -2.99 $\\pm$    0.03    &   0.002    &   -2.99 $\\pm$    0.03    &    0.01 $\\pm$    1.32    &   0.002   \\\\ 17    &    6.42 $\\pm$    1.26    &   -5.08 $\\pm$    3.08    &   0.574    &   -2.98 $\\pm$    0.05    &   0.822    &   -2.95 $\\pm$    0.05    &   -4.99 $\\pm$    7.95    &   0.527   \\\\ 18    &    7.95 $\\pm$    1.96    &   -1.62 $\\pm$    0.77    &   0.070    &   -2.89 $\\pm$    0.05    &   0.373    &   -2.90 $\\pm$    0.06    &    1.26 $\\pm$    0.55    &   0.055   \\\\ 21    &    5.76 $\\pm$    1.69    &   -3.18 $\\pm$    1.99    &   0.001    &   -2.99 $\\pm$    0.07    &   0.003    &   -2.99 $\\pm$    0.07    &   -0.40 $\\pm$    4.20    &   0.001   \\\\ 22    &    5.65 $\\pm$    2.39    &   -1.44 $\\pm$    1.17    &   0.031    &   -2.97 $\\pm$    0.09    &   0.195    &   -2.98 $\\pm$    0.09    &    1.40 $\\pm$    0.83    &   0.038   \\\\ 24    &   11.98 $\\pm$    2.20    &   -1.36 $\\pm$    0.48    &   0.104    &   -2.78 $\\pm$    0.04    &   1.056    &   -2.79 $\\pm$    0.04    &    1.27 $\\pm$    0.35    &   0.231   \\\\ 32    &    6.64 $\\pm$    0.98    &   -2.91 $\\pm$    0.93    &   0.007    &   -2.96 $\\pm$    0.04    &   0.006    &   -2.96 $\\pm$    0.04    &    0.05 $\\pm$    1.63    &   0.007   \\\\ \\hline \\end{tabular} \\end{center} \\caption{Model fits to the synchrotron coefficients determined between the 5 WMAP frequencies and the Haslam 408 MHz template for large sky areas (upper part of table) and for those regions that indicate a 3$\\sigma$ significant amplitude at K-band (lower part). The models SI, SII and SIII are described in Section~\\ref{sec:discussion_synch}. $A_{s}$ represents the normalisation amplitude at K-band, $\\beta_{s}$ the synchrotron spectral index and  $c_{s}$ the spectral curvature. The key for the global fits is as for Table~\\ref{tab:simple_synch_spectra}. \\label{tab:synch_model_results}} \\end{table}  \\begin{center} \\begin{figure} \\begin{tabular}{ccc} \\epsfig{file=figures/fig_51.ps,width=0.3\\linewidth,angle=0,clip=} & \\epsfig{file=figures/fig_52.ps,width=0.3\\linewidth,angle=0,clip=} & \\epsfig{file=figures/fig_53.ps,width=0.3\\linewidth,angle=0,clip=} \\\\ \\epsfig{file=figures/fig_54.ps,width=0.3\\linewidth,angle=0,clip=}  & \\epsfig{file=figures/fig_55.ps,width=0.3\\linewidth,angle=0,clip=}  & \\epsfig{file=figures/fig_56.ps,width=0.3\\linewidth,angle=0,clip=} \\\\ \\end{tabular} \\caption{Synchrotron spectral fits for various regions of the sky. A comparison is made of the observed template fits amplitudes derived from the 5 \\emph{WMAP} frequency bands and the 408~MHz survey with the 3 models SI, SII and SIII as defined in Section~\\ref{sec:discussion_synch}. The observed spectral dependencies of regions 3, 10 and 15 are representative of those seen for all regions with a significant detection at K-band.\\label{fig:synch_model_plots}} \\end{figure}  \\end{center}  The results are summarised in Table~\\ref{tab:synch_model_results}. Fig.~\\ref{fig:synch_model_plots} presents a comparison of the model fits for three of the global masks, plus three of the regions that can be considered representative of the general results seen. The $\\chi^2$ results generally should be considered indicative rather than definitive, as would be expected given the small number of data points and the comparable number of model parameters.  The synchrotron fit coefficients obtained for the global masks that include the 408~MHz datum as a reference point (model SII) indicate a typical power-law spectral index of $\\beta_s=-2.9\\pm0.1$. Model SI fits to the 5 \\emph{WMAP} frequencies are generally steeper, however, only the Ecliptic North region shows evidence of such behaviour at a significant level (a 2$\\sigma$ shift in the spectral index). Indeed, a model with negative spectral curvature is the best fit to the data, indicating steepening due to cosmic ray energy loss mechanisms. Curiously, the Ecliptic South indicates spectral flattening, although at much lower significance. The SII fit to the NPS region finds a spectral index $\\beta_s=-3.05\\pm0.03$ which is steeper than the average spectral index variation over the remaining sky, as expected \\citep{Lawson:1987}.  For the individual regions, the results are generally similar to the global fits. The spectral behaviour is consistent with emission from power-law cosmic ray electron spectra with only hints of a steepening between 408~MHz and K-band. Region 11 might be considered to show weak evidence for negative spectral curvature. More interestingly, three regions (10, 12 and 24) indicate positive curvature that might be understood as due to the presence of multiple emission regions with varying spectral behaviour along the line-of-sight.  However, region 10 has a very flat spectrum between K- and W-band that is strongly favoured over a steeper spectrum extending to 408~MHz, and perhaps even mildly inconsistent with the curvature model adopted here. Whether this indicates some problem with the analysis in this region, due to problems with the templates or cross-talk between components, is difficult to determine. It is apparent that the region lies close to the North Celestial Pole where the 408~MHz template still retains considerable striations from the original survey, whereas we certainly consider that the template fitting methodology has by now been extensively tested. However, the spectral flattening seen in other analyses is less dramatic. For example, \\citet{Gold_WMAP7:2011} find a range of values of $c_{s}$ between $\\pm 0.7$ with a variance of $\\sim 0.4$ for their MCMC analysis of the \\emph{WMAP} data combined with the 408~MHz data.   \\subsection{Free-free}\\label{sec:discussion_freefree}  The free-free spectral index is almost constant over the \\emph{WMAP} frequency range and changes only slightly with the the electron temperature. Over this frequency range, the spectral dependence of the corresponding \\halpha  emission can then be approximated with a power-law model of fixed spectral index $-0.15$. However, \\citet{Dobler_WMAP3:2008a} observed deviations from such power-law behaviour, which they attributed to the presence of spinning dust emission in the WIM also traced by \\halpha.  We investigate our template fit coefficients in terms of the following three models.  \\begin{itemize} \\item Model FI : We consider that the emission is due entirely to the free-free mechanism, ie. a power-law model emissivity with a fixed spectral index of $-0.15$ is assumed, \\begin{equation} I(\\nu) = A_{f}\\times \\left( \\frac{\\nu}{23} \\right )^{-0.15}_{\\text{GHz}} \\end{equation} Thus only the free-free amplitude ($A_{f}$)  needs to be estimated from the template fit coefficients. \\\\  \\item Model FII : The coefficients are fitted with an empirically motivated power-law model emissivity. This is particularly illustrative as to the extent that the standard free-free emissivity index is a poor fit to the data. \\begin{equation} I(\\nu) = A_{PL} \\times \\left( \\frac{\\nu}{23} \\right )^{\\beta_{PL}}_{\\text{GHz}} \\end{equation} In this case, both the foreground amplitude ($A_{PL}$) and spectral index  ($\\beta_{PL}$) are to be estimated.\\\\   \\item Model FIII : The coefficients are fitted with a combination of free-free emission and a spinning dust model computed by the {\\tt{SPDUST2}} code \\citep{SpDust2:2011} for typical WIM conditions. This model peaks around 30~GHz, but we allow a simple shift $\\Delta \\nu_{\\text{WIM}}$ to be applied to the spectrum, as an approximation that represents the effect of varying the WIM physical parameters to match the model spectrum to the data. \\begin{equation} I(\\nu) =   A_{f} \\times\\left( \\frac{\\nu}{23} \\right)^{-0.15}_{\\text{GHz}}  + A_{\\text{WIM}}  \\times D_{\\text{WIM}}(\\nu - \\Delta \\nu_{\\text{WIM}}) \\end{equation} Here, $D_{\\text{WIM}}(\\nu)$ represents the normalised spinning dust spectral model for the WIM at a given frequency. Clearly, we must now fit 3 parameters -- the free-free amplitude ($A_{f}$),  WIM amplitude ($A_{\\text{WIM}}$ ) and WIM frequency shift ($\\Delta \\nu_{\\text{WIM}}$). We have also considered fits of this model to the coefficients derived using the \\halpha after correction for dust absorption ($f_{d}=0.5$).\\\\  \\end{itemize}   \\begin{table} \\scriptsize \\begin{center} \\begin{tabular}{lccccccccc} \\hline & \\multicolumn{2}{c}{FI} & \\multicolumn{3}{c}{FII} & \\multicolumn{4}{c}{FIII} \\\\ \\hline Region & $A_{f}$ & $\\chi^2$ & $A_{PL}$ & $\\beta_{PL}$ & $\\chi^2$ &  $A_{f}$ & $A_{WIM}$ & $\\Delta \\nu_{\\text{WIM}}$ & $\\chi^2$ \\\\ \\hline EBV    &   0.156 $\\pm$   0.001    &   6.374    &   0.153 $\\pm$   0.002    &   0.010 $\\pm$   0.036    &   1.633    &   0.150 $\\pm$   0.004    &   2.094 $\\pm$   0.429    &  13.575 $\\pm$   4.727    &   0.340   \\\\ KQ85    &   0.166 $\\pm$   0.003    &   2.449    &   0.163 $\\pm$   0.003    &   0.061 $\\pm$   0.067    &   0.205    &   0.161 $\\pm$   0.004    &   2.822 $\\pm$   1.204    &  15.930 $\\pm$   6.463    &   0.439   \\\\ GN    &   0.163 $\\pm$   0.002    &   2.057    &   0.161 $\\pm$   0.003    &  -0.073 $\\pm$   0.065    &   2.264    &   0.141 $\\pm$   0.017    &   3.483 $\\pm$   1.979    &   5.343 $\\pm$   3.317    &   0.704   \\\\ EN    &   0.210 $\\pm$   0.004    &   2.177    &   0.207 $\\pm$   0.004    &   0.031 $\\pm$   0.070    &   0.730    &   0.203 $\\pm$   0.008    &   3.442 $\\pm$   1.245    &  14.290 $\\pm$   7.196    &   0.062   \\\\ GS    &   0.156 $\\pm$   0.002    &   4.622    &   0.153 $\\pm$   0.002    &   0.039 $\\pm$   0.044    &   0.296    &   0.152 $\\pm$   0.002    &   2.386 $\\pm$   0.806    &  16.354 $\\pm$   4.503    &   1.163   \\\\ ES    &   0.146 $\\pm$   0.001    &   3.499    &   0.144 $\\pm$   0.002    &  -0.003 $\\pm$   0.041    &   1.031    &   0.141 $\\pm$   0.005    &   1.775 $\\pm$   0.491    &  12.932 $\\pm$   6.538    &   0.202   \\\\ \\hline 7    &   0.126 $\\pm$   0.029    &   1.003    &   0.141 $\\pm$   0.030    &  -2.043 $\\pm$   2.257    &   0.601    &   0.000 $\\pm$   0.000    &  15.936 $\\pm$   3.512    &   0.000 $\\pm$   0.000    &   1.381   \\\\ 8    &   0.111 $\\pm$   0.021    &   0.086    &   0.110 $\\pm$   0.024    &  -0.043 $\\pm$   0.831    &   0.109    &   0.053 $\\pm$   0.152    &   8.168 $\\pm$  18.511    &   3.214 $\\pm$  10.241    &   0.024   \\\\ 9    &   0.162 $\\pm$   0.004    &   3.619    &   0.156 $\\pm$   0.004    &   0.205 $\\pm$   0.101    &   1.046    &   0.152 $\\pm$   0.006    &   5.470 $\\pm$   1.890    &  15.818 $\\pm$   5.131    &   0.115   \\\\ 11    &   0.388 $\\pm$   0.011    &   0.298    &   0.386 $\\pm$   0.012    &  -0.105 $\\pm$   0.128    &   0.358    &   0.388 $\\pm$   0.011    &   0.000 $\\pm$   0.000    &   0.000 $\\pm$   0.000    &   0.596   \\\\ 12    &   0.135 $\\pm$   0.034    &   0.731    &   0.150 $\\pm$   0.036    &  -1.693 $\\pm$   2.087    &   0.481    &   0.000 $\\pm$   0.000    &  17.099 $\\pm$   4.139    &   0.000 $\\pm$   0.000    &   0.909   \\\\ 13    &   0.145 $\\pm$   0.005    &   0.037    &   0.145 $\\pm$   0.006    &  -0.119 $\\pm$   0.169    &   0.039    &   0.137 $\\pm$   0.038    &   1.238 $\\pm$   4.593    &   4.601 $\\pm$  20.905    &   0.003   \\\\ 14    &   0.156 $\\pm$   0.003    &   2.336    &   0.153 $\\pm$   0.003    &   0.060 $\\pm$   0.066    &   0.096    &   0.152 $\\pm$   0.003    &   2.858 $\\pm$   1.513    &  17.556 $\\pm$   5.663    &   1.060   \\\\ 15    &   0.200 $\\pm$   0.011    &   0.836    &   0.191 $\\pm$   0.012    &   0.227 $\\pm$   0.192    &   0.051    &   0.192 $\\pm$   0.012    &   6.819 $\\pm$   6.145    &  20.334 $\\pm$   7.279    &   0.627   \\\\ 18    &   0.235 $\\pm$   0.027    &   0.772    &   0.216 $\\pm$   0.029    &   0.620 $\\pm$   0.358    &   0.018    &   0.219 $\\pm$   0.031    &  19.115 $\\pm$  17.644    &  21.162 $\\pm$   7.240    &   0.612   \\\\ 20    &   0.135 $\\pm$   0.012    &   1.495    &   0.123 $\\pm$   0.013    &   0.685 $\\pm$   0.276    &   0.116    &   0.127 $\\pm$   0.014    &  14.890 $\\pm$   9.172    &  23.658 $\\pm$   4.864    &   1.101   \\\\ 21    &   0.250 $\\pm$   0.043    &   1.159    &   0.212 $\\pm$   0.044    &   1.062 $\\pm$   0.411    &   0.014    &   0.222 $\\pm$   0.048    &  42.774 $\\pm$  29.924    &  23.582 $\\pm$   5.653    &   0.893   \\\\ 23    &   0.164 $\\pm$   0.023    &   0.157    &   0.168 $\\pm$   0.025    &  -0.447 $\\pm$   0.708    &   0.134    &   0.066 $\\pm$   0.170    &  12.174 $\\pm$  20.912    &   0.000 $\\pm$   0.000    &   0.144   \\\\ 24    &   0.118 $\\pm$   0.005    &   0.340    &   0.116 $\\pm$   0.006    &   0.074 $\\pm$   0.195    &   0.048    &   0.114 $\\pm$   0.015    &   1.977 $\\pm$   1.888    &  13.646 $\\pm$  21.049    &   0.093   \\\\ 32    &   0.164 $\\pm$   0.041    &   0.374    &   0.144 $\\pm$   0.042    &   0.914 $\\pm$   0.657    &   0.019    &   0.150 $\\pm$   0.046    &  19.687 $\\pm$  27.751    &  22.391 $\\pm$  10.970    &   0.368   \\\\ \\hline \\end{tabular} \\end{center} \\caption{Model fits to the free-free coefficients determined between the 5 WMAP frequencies and the DDD \\halpha template for large sky areas (upper part of table) and for those regions that indicate a 3$\\sigma$ significant amplitude at K-band (lower part). The models FI, FII and FIII are fully defined in Section~\\ref{sec:discussion_freefree}. $A_{f}$ represents the normalisation amplitude of the free-free emission at K-band, $A_{PL}$ represents the normalisation amplitude for the power-law model emission at K-band, $\\beta_{PL}$ the corresponding power-law spectral index, $A_{WIM}$ is the amplitude of the WIM spinning dust model in units of 10$^{20}$ R cm$^{-2}$, and $\\Delta \\nu_{\\text{WIM}}$ the shift in frequency of the peak of the dust model to better the fit the data. The key for the global fits as for Table~\\ref{tab:simple_synch_spectra}. \\label{tab:freefree_model_hemispheres}} \\end{table}  \\begin{center} \\begin{figure} \\begin{tabular}{ccc} \\epsfig{file=figures/fig_61.ps,width=0.3\\linewidth,angle=0,clip=} & \\epsfig{file=figures/fig_62.ps,width=0.3\\linewidth,angle=0,clip=} & \\epsfig{file=figures/fig_63.ps,width=0.3\\linewidth,angle=0,clip=} \\\\ \\epsfig{file=figures/fig_64.ps,width=0.3\\linewidth,angle=0,clip=} & \\epsfig{file=figures/fig_65.ps,width=0.3\\linewidth,angle=0,clip=} & \\epsfig{file=figures/fig_66.ps,width=0.3\\linewidth,angle=0,clip=} \\\\ \\end{tabular} \\caption{Free-free spectral fits for various regions of the sky. A comparison is made of the observed template fits amplitudes derived from the 5 \\emph{WMAP} frequency bands and the  DDD \\halpha template with the 3 models FI, FII and FIII as defined in Section~\\ref{sec:discussion_freefree}. Also shown in the plot are the separate free-free (FF) and WIM spinning dust components that constitute model FIII. The observed spectral dependencies of regions 7, 13 and 14 are representative of those seen for all regions with a significant detection at K-band. \\label{fig:freefree_model_hemispheres}} \\end{figure} \\end{center}  The results are summarised in Table~\\ref{tab:freefree_model_hemispheres}. Fig.~\\ref{fig:freefree_model_hemispheres} presents a comparison of the model fits for three of the global masks, plus three of the regions that can be considered representative of the general results seen.  The global fits all indicate a deviation from the free-free emission law (FI) at high significance. Unconstrained power-law fits (FII) are generally flatter, and models containing both free-free and spinning dust emission (FII) are typically preferred. For the EBV mask, the ratio of free-free to spinning dust emission (at the shifted peak frequency) is of order 7,  but this drops to 4 for the Galactic North. All regional fits are consistent with a 15\\,GHz shift to higher frequencies of the peak intensity for the WIM spinning dust component with the exception of the Galactic North region which favours a smaller value of 5\\,GHz.  The fits to individual regions are generally consistent with the global mask results, and only region 9 shows a clear rejection of model FI. The typical frequency shift is again of order 15~GHz, with a larger dispersion with several regions preferring no shift at all. Of those, regions 7 and 12 can be explained by a single component - either free-free, power-law emission or spinning dust alone with no frequency shift. However, it should be noted that the \\halpha fit coefficients for these regions are only significant at K-band. Region 11 is consistent with free-free only emission, and yields no evidence for spinning dust, but the region itself has been flagged as anomalous as discussed previously in section \\ref{sec:dust_extinct}.  Region 13 corresponds to the Northern Gum nebula, and is dominated by free-free emission with a small contribution from spinning dust at a lower peak frequency, $\\sim 35$~GHz, than is typical. The Southern part of the nebula is contained in region 14. This is the most clearly detected structure on the sky, significant at all frequencies, as traced by the \\halpha template, and shows significant evidence for a spinning dust contribution with a free-free to spinning dust ratio of $\\sim$ 5. Otherwise, this ratio varies considerably from region to region.  Some regions that have a rising spectrum in terms of the \\halpha coefficients and naturally favour models with a spinning dust contribution over pure free-free are better fitted still by a flatter/rising power-law emission model. This might be alleviated with more detailed spinning dust modelling (beyond the scope of our paper), or by including physical effects that increase the spinning dust amplitude at frequencies higher than the peak. Such modifications have been investigated by \\citet{Hoang:2010}.  The $A_{f}$ fit coefficients can, of course, be converted to estimates of the thermal electron temperature in the ionised medium. Since the global mask fits require the presence of a spinning dust component, we only consider the $T_e$ results from model FIII, as presented in Table~\\ref{tab:freefree_electron_temperatures}. The table also includes results derived from fits to an \\halpha template corrected for dust absorption assuming $f_d=0.5$. We do not include the detailed coefficient results here since the interpretation presented above remains essentially unchanged, and only conclusions about $T_{e}$ are affected. The global masks are consistent with values of the electron temperature of order 6000\\,K without any dust absorption correction, falling to 3000\\,K when $f_d=0.5$ is assumed. These values are for guidance only -- since the regions are not independent, an average is meaningless. It is interesting to note that the Ecliptic North shows an enhanced temperature some 50\\% higher than these typical values. Whether this reflects some property of the \\halpha template is difficult to say, but the coverage is dominated by measurements from \\emph{WHAM} data. Conversely, the Ecliptic South global fit gives a lower temperature that may also suggest issues with the template, given that it is largely comprised of the \\emph{SHASSA} fields. New observations from the southern extension of the \\emph{WHAM} survey \\citep{Haffner_WHAM-S:2011} should help resolve this issue in the near future. Nevertheless, problems may still remain near the ecliptic poles given difficulties with removing the geocoronal \\halpha contribution. For the individual regions (ignoring region 11 which is considered to be anomalous), we find a weighted average of $6300 \\pm  200$\\,K (dropping to 5900\\,K if region 11 is excluded) without a correction for dust absorption, and $2900\\pm 100$\\,K otherwise. These values are in good agreement with the global averages as might be expected. However, note that the dispersion of values is considerably larger than the quoted uncertainty, implying true variations in temperature on the sky. Moreover, as discussed previously, it is likely that some of the dispersion seen reflects the existence of a range of values for both $T_{e}$ and $f_{d}$ throughout the Galaxy.   \\begin{table} \\scriptsize \\begin{center} \\begin{tabular}{lcc} \\hline &  \\multicolumn{2}{c}{Inferred $T_{e}$  (K) for model FIII} \\\\ Region &  $f_{d} = 0.0$ & $f_{d} = 0.5$ \\\\ \\hline EBV  &   5900 $\\pm$    300   &   2600 $\\pm$    200   \\\\ KQ85  &   6600 $\\pm$    300   &   3200 $\\pm$    200   \\\\ GN  &   5300 $\\pm$   1000   &   2500 $\\pm$    500   \\\\ EN  &   9500 $\\pm$    600   &   5300 $\\pm$    600   \\\\ GS  &   6000 $\\pm$    100   &   2600 $\\pm$    100   \\\\ ES  &   5300 $\\pm$    300   &   2200 $\\pm$    200   \\\\ \\hline 7$^{(a)}$  &   4400 $\\pm$   1700   &   5600 $\\pm$   1100   \\\\ 8$^{(a)}$  &   3600 $\\pm$   1200   &   2300 $\\pm$    900   \\\\ 9  &   6000 $\\pm$    400   &   3200 $\\pm$    300   \\\\ 11  &  25500 $\\pm$   1100   &  11800 $\\pm$    500   \\\\ 12$^{(a)}$  &   5000 $\\pm$   2100   &   3300 $\\pm$   1300   \\\\ 13  &   5100 $\\pm$   2200   &   3500 $\\pm$   1500   \\\\ 14  &   6000 $\\pm$    200   &   2700 $\\pm$    100   \\\\ 15  &   8600 $\\pm$    800   &   5900 $\\pm$    600   \\\\ 18  &  10700 $\\pm$   2300   &   7700 $\\pm$   1700   \\\\ 20  &   4500 $\\pm$    800   &   2400 $\\pm$    500   \\\\ 21  &  10900 $\\pm$   3600   &   4000 $\\pm$   1800   \\\\ 23$^{(a)}$  &   6800 $\\pm$   1400   &   5400 $\\pm$   1000   \\\\ 24  &   3700 $\\pm$    600   &   2300 $\\pm$    700   \\\\ 32  &   5900 $\\pm$   2900   &   5600 $\\pm$   2500   \\\\ \\hline \\end{tabular} \\end{center} \\caption{Inferred free-free electron temperature $T_{e}$ in Kelvins corresponding to model fit FIII with dust corrections $f_d=0.0$ and $0.5$. $^(a)$ These regions show an effective degeneracy between free-free only or WIM spinning dust only solutions, thus FIII solutions have no contribution from free-free emission. We have therefore used the FI results to compute $T_{e}$ in these cases.} \\label{tab:freefree_electron_temperatures} \\end{table}   \\subsection{Dust}\\label{sec:discussion_dust}  The total dust emission is modelled as a combination of the relatively well-understood thermal dust emission and the AME. The former is assumed to have a fixed spectral index relative to the FDS8 94~GHz template over the \\emph{WMAP} frequency range as determined directly from the FDS8 model. We consider the following two models in order to fit the dust coefficients.  \\begin{itemize} \\item Model DI : The dust coefficients are fitted with a combination of thermal (vibrational) dust and a power law dust-correlated AME. \\\\ \\begin{equation} T_A (\\nu)  =  A_{PLD} \\times \\left( \\frac{\\nu}{23.} \\right)^{\\beta_{PLD}}_{\\text{GHz}} + A_{TD}\\times \\left( \\frac{\\nu}{94} \\right)^{1.55}_{\\text{GHz}} \\end{equation} Three parameters -- the thermal dust amplitude ($A_{TD}$), the power law dust amplitude ($A_{PLD}$) and power law dust spectral index ($\\beta_{PLD}$) -- are fitted to the coefficients.  \\item Model DII : The dust coefficients are fitted with a combination of thermal dust and two spinning dust components (CNM and WNM). The two spectra are generated using the {\\tt{SPDUST2}} code assuming typical CNM and WNM conditions \\citep{Draine_spinning:1998}. Given that both spectra peak at approximately the same frequency ($\\approx 30$\\,GHz), and the limited number of degrees of freedom available in the fit, it is only possible to apply a frequency shift to one component in order to match observations. \\citet{Hoang:2011} found that modifying the CNM properties to increase its peak frequency yields a closer match to the \\emph{WMAP} observations, and therefore we elect to allow a frequency shift of this component. \\\\ \\begin{equation} T_A (\\nu)  =  A_{\\text{WNM}}  \\times D_{\\text{WNM}}(\\nu) + A_{\\text{CNM}}  \\times D_{\\text{CNM}}(\\nu - \\Delta \\nu_{\\text{CNM}}) + A_{TD} \\times  \\left( \\frac{\\nu}{94} \\right )^{1.55}_{\\text{GHz}} \\end{equation} We fit four parameters to the template fit coefficients: the thermal dust amplitude ($A_{TD}$), the WNM amplitude ($A_{\\text{WNM}}$) normalised at 23~GHz, the CNM amplitude ($A_{\\text{CNM}}$) normalised at 41~GHz, and the CNM peak frequency shift ($\\Delta \\nu_{\\text{CNM}}$).  \\end{itemize}  \\begin{table} \\scriptsize \\begin{center} \\begin{tabular}{lccccccccc} \\hline Region &\\multicolumn{4}{c}{Model DI} &\\multicolumn{5}{c}{Model DII}\\\\ \\hline & $A_{PLD}$ &$\\beta_{PLD}$ & $A_{TD}$ & $\\chi^2$ & $A_{WNM}$ & $A_{CNM}$ & $\\Delta {\\nu_{\\text{CNM}}}$ & $A_{TD}$ & $\\chi^2$  \\\\ \\hline EBV    &    9.15 $\\pm$    0.09    &   -2.74 $\\pm$    0.06    &    1.01 $\\pm$    0.08    &   1.710    &    9.35 $\\pm$    0.10    &    0.67 $\\pm$    0.09    &   22.42 $\\pm$    1.22    &    1.22 $\\pm$    0.08    &   0.825   \\\\ KQ85    &    7.61 $\\pm$    0.08    &   -2.78 $\\pm$    0.07    &    1.14 $\\pm$    0.08    &   0.667    &    7.78 $\\pm$    0.09    &    0.53 $\\pm$    0.08    &   23.34 $\\pm$    1.33    &    1.29 $\\pm$    0.08    &   1.208   \\\\ GN   &    8.91 $\\pm$    0.14    &   -2.86 $\\pm$    0.11    &    1.00 $\\pm$    0.13    &   0.769    &    9.12 $\\pm$    0.15    &    0.50 $\\pm$    0.14    &   23.35 $\\pm$    2.51    &    1.18 $\\pm$    0.13    &   0.829   \\\\ EN    &    8.32 $\\pm$    0.13    &   -2.83 $\\pm$    0.11    &    0.99 $\\pm$    0.13    &   0.617    &    8.51 $\\pm$    0.14    &    0.49 $\\pm$    0.14    &   23.35 $\\pm$    2.42    &    1.16 $\\pm$    0.13    &   0.679   \\\\ GS  &    9.73 $\\pm$    0.13    &   -2.64 $\\pm$    0.08    &    1.02 $\\pm$    0.13    &   0.699    &    9.93 $\\pm$    0.13    &    0.85 $\\pm$    0.13    &   21.99 $\\pm$    1.41    &    1.28 $\\pm$    0.12    &   0.114   \\\\ ES   &    9.82 $\\pm$    0.13    &   -2.69 $\\pm$    0.08    &    0.99 $\\pm$    0.12    &   1.162    &   10.01 $\\pm$    0.13    &    0.78 $\\pm$    0.12    &   22.07 $\\pm$    1.47    &    1.24 $\\pm$    0.11    &   0.285   \\\\ \\hline 2    &    8.03 $\\pm$    1.60    &   -4.14 $\\pm$    2.16    &    0.00 $\\pm$    0.00    &   0.093    &    8.17 $\\pm$    1.63    &    0.00 $\\pm$    0.00    &    0.00 $\\pm$    0.00    &    0.00 $\\pm$    0.00    &   0.378   \\\\ 3    &    8.73 $\\pm$    0.99    &   -2.60 $\\pm$    0.66    &    2.03 $\\pm$    0.95    &   0.016    &    8.94 $\\pm$    1.04    &    0.88 $\\pm$    0.97    &   23.97 $\\pm$    8.94    &    2.20 $\\pm$    0.91    &   0.001   \\\\ 6    &    7.49 $\\pm$    2.33    &   -3.06 $\\pm$    2.31    &    0.84 $\\pm$    2.20    &   0.007    &    7.68 $\\pm$    2.44    &    0.17 $\\pm$    2.50    &   25.27 $\\pm$  100.72    &    0.99 $\\pm$    2.23    &   0.002   \\\\ 7    &    9.79 $\\pm$    0.63    &   -2.56 $\\pm$    0.40    &    1.19 $\\pm$    0.64    &   0.034    &    9.96 $\\pm$    0.67    &    0.98 $\\pm$    0.65    &   21.14 $\\pm$    6.79    &    1.49 $\\pm$    0.60    &   0.008   \\\\ 8    &    6.70 $\\pm$    0.63    &   -2.73 $\\pm$    0.62    &    1.41 $\\pm$    0.62    &   0.019    &    6.86 $\\pm$    0.66    &    0.56 $\\pm$    0.68    &   24.81 $\\pm$    8.69    &    1.48 $\\pm$    0.61    &   0.018   \\\\ 9    &   13.09 $\\pm$    0.45    &   -2.61 $\\pm$    0.22    &    1.09 $\\pm$    0.52    &   0.121    &   13.34 $\\pm$    0.48    &    1.20 $\\pm$    0.48    &   21.58 $\\pm$    3.83    &    1.49 $\\pm$    0.50    &   0.003   \\\\ 10    &    6.55 $\\pm$    0.71    &   -3.16 $\\pm$    0.76    &    0.67 $\\pm$    0.60    &   0.051    &    6.69 $\\pm$    0.74    &    0.11 $\\pm$    0.70    &   25.45 $\\pm$   44.79    &    0.78 $\\pm$    0.61    &   0.110   \\\\ 11    &    4.20 $\\pm$    0.63    &   -2.45 $\\pm$    0.84    &    0.81 $\\pm$    0.64    &   0.061    &    4.21 $\\pm$    0.86    &    0.44 $\\pm$    0.59    &   16.51 $\\pm$   25.45    &    1.02 $\\pm$    0.59    &   0.006   \\\\ 12    &    5.29 $\\pm$    0.36    &   -3.79 $\\pm$    0.68    &    0.35 $\\pm$    0.38    &   0.304    &    5.40 $\\pm$    0.37    &    0.00 $\\pm$    0.00    &    0.00 $\\pm$    0.00    &    0.34 $\\pm$    0.35    &   1.691   \\\\ 13    &    5.61 $\\pm$    0.71    &   -3.16 $\\pm$    0.88    &    0.78 $\\pm$    0.60    &   0.048    &    5.73 $\\pm$    0.75    &    0.09 $\\pm$    0.69    &   24.66 $\\pm$   62.57    &    0.88 $\\pm$    0.61    &   0.091   \\\\ 14    &    9.70 $\\pm$    0.45    &   -2.53 $\\pm$    0.28    &    1.20 $\\pm$    0.43    &   0.033    &    9.87 $\\pm$    0.48    &    1.03 $\\pm$    0.46    &   21.51 $\\pm$    4.49    &    1.50 $\\pm$    0.40    &   0.000   \\\\ 15    &    8.25 $\\pm$    0.54    &   -2.41 $\\pm$    0.38    &    1.64 $\\pm$    0.55    &   0.015    &    8.40 $\\pm$    0.58    &    1.06 $\\pm$    0.56    &   21.79 $\\pm$    5.14    &    1.92 $\\pm$    0.50    &   0.037   \\\\ 16    &    6.29 $\\pm$    0.62    &   -4.72 $\\pm$    1.36    &    0.00 $\\pm$    0.00    &   0.327    &    6.38 $\\pm$    0.63    &    0.00 $\\pm$    0.00    &    0.00 $\\pm$    0.00    &    0.00 $\\pm$    0.00    &   3.088   \\\\ 17    &    6.59 $\\pm$    1.79    &   -3.90 $\\pm$    2.82    &    0.17 $\\pm$    1.62    &   0.055    &    6.73 $\\pm$    1.83    &    0.00 $\\pm$    0.00    &    0.00 $\\pm$    0.00    &    0.15 $\\pm$    1.52    &   0.148   \\\\ 18    &    6.99 $\\pm$    0.90    &   -3.41 $\\pm$    1.10    &    0.57 $\\pm$    0.78    &   0.053    &    7.14 $\\pm$    0.91    &    0.00 $\\pm$    0.00    &    0.00 $\\pm$    0.00    &    0.67 $\\pm$    0.72    &   0.127   \\\\ 19    &    9.29 $\\pm$    1.88    &   -1.87 $\\pm$    0.99    &    2.81 $\\pm$    2.23    &   0.050    &    9.48 $\\pm$    2.01    &    2.19 $\\pm$    1.97    &   22.96 $\\pm$    7.81    &    3.36 $\\pm$    1.81    &   0.062   \\\\ 20    &    9.85 $\\pm$    0.72    &   -2.90 $\\pm$    0.52    &    0.87 $\\pm$    0.71    &   0.070    &   10.09 $\\pm$    0.76    &    0.45 $\\pm$    0.76    &   23.50 $\\pm$   14.84    &    1.08 $\\pm$    0.71    &   0.024   \\\\ 21    &    6.92 $\\pm$    0.45    &   -2.98 $\\pm$    0.47    &    0.96 $\\pm$    0.40    &   0.049    &    7.08 $\\pm$    0.47    &    0.28 $\\pm$    0.47    &   24.34 $\\pm$   12.08    &    1.08 $\\pm$    0.41    &   0.124   \\\\ 22    &    8.73 $\\pm$    0.81    &   -2.40 $\\pm$    0.54    &    1.69 $\\pm$    0.87    &   0.019    &    8.87 $\\pm$    0.87    &    1.08 $\\pm$    0.85    &   20.69 $\\pm$    8.30    &    2.03 $\\pm$    0.80    &   0.023   \\\\ 23    &    8.96 $\\pm$    0.36    &   -2.54 $\\pm$    0.22    &    1.12 $\\pm$    0.32    &   0.079    &    9.12 $\\pm$    0.38    &    0.95 $\\pm$    0.30    &   21.38 $\\pm$    3.44    &    1.40 $\\pm$    0.30    &   0.005   \\\\ 24    &    7.75 $\\pm$    0.54    &   -3.56 $\\pm$    0.63    &    0.04 $\\pm$    0.49    &   0.207    &    7.92 $\\pm$    0.55    &    0.00 $\\pm$    0.00    &    0.00 $\\pm$    0.00    &    0.10 $\\pm$    0.45    &   0.867   \\\\ 28    &    8.46 $\\pm$    1.80    &   -2.03 $\\pm$    1.09    &    2.47 $\\pm$    2.01    &   0.038    &    8.64 $\\pm$    1.91    &    1.73 $\\pm$    1.89    &   23.34 $\\pm$    9.27    &    2.83 $\\pm$    1.71    &   0.031   \\\\ 30    &    9.73 $\\pm$    3.14    &   -2.00 $\\pm$    1.62    &    2.30 $\\pm$    3.47    &   0.006    &    9.85 $\\pm$    3.36    &    2.09 $\\pm$    3.18    &   21.92 $\\pm$   14.67    &    2.81 $\\pm$    2.91    &   0.002   \\\\ 31    &   12.00 $\\pm$    3.23    &   -1.93 $\\pm$    1.32    &    2.70 $\\pm$    3.66    &   0.005    &   12.13 $\\pm$    3.46    &    2.67 $\\pm$    3.27    &   21.47 $\\pm$   12.17    &    3.45 $\\pm$    3.00    &   0.011   \\\\ 32    &    8.15 $\\pm$    0.81    &   -2.87 $\\pm$    0.70    &    0.85 $\\pm$    0.81    &   0.003    &    8.33 $\\pm$    0.85    &    0.53 $\\pm$    0.87    &   24.26 $\\pm$   12.86    &    0.95 $\\pm$    0.81    &   0.098   \\\\ 33    &    8.53 $\\pm$    1.88    &   -3.08 $\\pm$    1.60    &    0.78 $\\pm$    1.70    &   0.025    &    8.74 $\\pm$    1.99    &    0.12 $\\pm$    1.91    &   23.77 $\\pm$  135.62    &    0.99 $\\pm$    1.72    &   0.004   \\\\ \\hline \\end{tabular} \\end{center} \\caption{Model fits to the dust coefficients determined between the 5 WMAP frequencies and the FDS8  template  for large sky areas (upper part of table) and for those regions that indicate a 3$\\sigma$ significant amplitude at K-band (lower part). The models DI and DII are fully defined in Section~\\ref{sec:discussion_dust}. $A_{PLD}$ and $\\beta_{PLD}$ represent the normalisation amplitude  at K-band and spectral index respectively of a power-law anomalous dust emission component. $A_{WNM}$ is the amplitude of the WNM spinning dust model at K-band, $A_{CNM}$ is the amplitude of the CNM spinning dust model normalised at 41~GHz, and $\\Delta \\nu_{CNM}$ is the shift in frequency of the peak of the CNM dust model to better the fit the data. $A_{TD}$ is the amplitude of the thermal dust emission with an assumed spectral index  $\\beta_{TD} = 1.55$ as determined directly from the FDS8 dust model. The key for the global fits as for Table~\\ref{tab:simple_synch_spectra}.} \\label{tab:dust_model_hemispheres} \\end{table}   \\begin{center} \\begin{figure} \\begin{tabular}{ccc} \\epsfig{file=figures/fig_71.ps,width=0.3\\linewidth,angle=0,clip=} & \\epsfig{file=figures/fig_72.ps,width=0.3\\linewidth,angle=0,clip=} & \\epsfig{file=figures/fig_73.ps,width=0.3\\linewidth,angle=0,clip=} \\\\ \\epsfig{file=figures/fig_74.ps,width=0.3\\linewidth,angle=0,clip=} & \\epsfig{file=figures/fig_75.ps,width=0.3\\linewidth,angle=0,clip=} & \\epsfig{file=figures/fig_76.ps,width=0.3\\linewidth,angle=0,clip=} \\\\ \\end{tabular} \\caption{Dust spectral fits for various regions of the sky. A comparison is made of the observed template fit amplitudes derived from the 5 \\emph{WMAP} frequency bands and the FDS8 template with the 2 models DI and DII as defined in Section~\\ref{sec:discussion_dust}. Also shown in the plot are the separate power-law anomalous dust (PLD) and vibrational (thermal) dust (VD) that constitute model DI, and the cold neutral medium (CNM) and warm neutral medium (WNM) components that, together with the vibrational dust, constitute model DII. The observed spectral dependencies of regions 12, 14 and 21 are representative of those seen for all regions with a significant detection at K-band.\\label{fig:dust_model_hemispheres}} \\end{figure} \\end{center}  The results are summarised in Table~\\ref{tab:dust_model_hemispheres}. Fig.~\\ref{fig:dust_model_hemispheres} presents a comparison of the model fits for three of the global masks, plus three of the regions that can be considered representative of the general results seen.  It should be apparent from the table that both models DI and DII provide adequate fits to the global masks. However, in order to explain the emission at V-band in terms of spinning dust models, two contributions are required in combination with an enhanced thermal dust normalisation some 20\\% larger than that predicted at 94\\,GHz by FDS8. Such models also require a shift in the peak CNM emission of $\\sim$22\\,GHz and a ratio of WNM to CNM spinning dust emission of order 4:1.  In fact, if the CNM emission were to be omitted, or the thermal dust amplitude kept at the canonical FDS8 value, then the V- and/or W-band amplitudes would be substantially underpredicted. In practise, given the low number of data points sensitive to the thermal dust emission, there is an effective degeneracy between the CNM and vibrational dust components, that may be alleviated somewhat if the spinning dust models can be enhanced in amplitude at frequencies higher than their peak. Indeed, spinning dust can contribute significantly to emission in the $\\sim$60 to 90 GHz range depending on the local physical conditions. However, one might then expect to see a much flatter spectrum around the Q-, V-, and W-bands than is actually observed.  It may also be that some modification of the assumed thermal dust spectral index is required.  The fits to individual regions are less informative, as expected given the lack of clear detections for the template fit coefficients at V- and W-band. Given this, the analysis might be criticised for overfitting the data points. Nevertheless, both a power-law and a WNM spinning dust component are good candidates to explain the lower frequency points. In almost all cases, a significant contribution from CNM or thermal dust emission is not required.  This discrepancy between the regional fits and global analysis may simply reflect a signal-to-noise issue at higher frequencies than the former, or the fact that the latter are essentially averages over many regions which exhibit some variation in dust properties. Regions 12 and 16 are interesting in that the power-law model is clearly preferred over model DII, although the analysis is really only constrained by the data points at K- and Ka-band.  For those regions in which a significant vibrational dust contribution is required, the amplitude is systematically higher than the canonical FDS8 value, although still consistent with it. However, the required DI enhancement is lower, thus the DII results more likely again point to a deficiency in the specific spinning dust models used in this analysis.  The FDS8 model value for the dust spectral index used here is quite flat, whereas early \\emph{Planck} papers \\citep{Planck_ERXXIV_ISM:2011,Planck_ERXXV_Molecular:2011} studying emission at frequencies higher than $\\sim$100~GHz favour values closer to 1.8. However, \\citet{Planck_ERXIX_DarkGas:2011} suggests that the dust SED flattens in the millimetre wavelength range. At lower frequencies, there is considerable uncertainty, and detailed modelling for candidate AME regions (including Perseus and $\\rho$-Ophiucus) of the foreground spectra for all components and across the \\emph{Planck} frequencies suggests values in the range 1.5--2 \\citep{Planck_ERXX_AME:2011}.  We have repeated the calculations above but imposing a thermal dust spectral index of 1.7 on the fits.  For model DI, there is a general increase in amplitude of the power-law dust component, with an associated flattening of the spectral index at levels below the 1-$\\sigma$ error bar, and no impact on the thermal dust amplitude. Model DII shows similarly modest increases in the amplitudes of the two spinning dust components, with little change in the CNM peak frequency shift, and again no change in the thermal dust amplitude.  However, adopting the FDS8 thermal dust model at W-band and then imposing a steeper index to lower frequencies is not self-consistent.  Therefore, a more detailed treatment is required including specific modelling of all dust components combined with higher frequency measurements. An obvious potential error in our analysis is the selective use of the CNM and WNM spectra for typical conditions in those phases of the ISM. What is clear is that the frequency range of 60--100\\,GHz is a remarkably interesting regime for dust astrophysics."
    },
    "1112/1112.2817_arXiv.txt": {
        "abstract": "Detection and characterization of exoplanets by direct imaging requires a coronagraph designed to deliver high contrast at small angular separation. To achieve this, an accurate control of low order aberrations, such as pointing and focus errors, is essential to optimize coronagraphic rejection and avoid the possible confusion between exoplanet light and coronagraphic leaks in the science image. Simulations and laboratory prototyping have shown that a Coronagraphic Low Order Wave-Front Sensor (CLOWFS), using a single defocused image of a reflective focal plane ring, can be used to control tip-tilt to an accuracy of $10^{-3}$ $\\lambda{/D}$. This paper demonstrates that the data acquired by CLOWFS can also be used in post-processing to calibrate residual coronagraphic leaks from the science image. Using both the CLOWFS camera and the science camera in the system, we quantify the accuracy of the method and its ability to successfully remove light due to low order errors from the science image. We also report the implementation and performance of the CLOWFS on the Subaru Coronagraphic Extreme AO (SCExAO) system and its expected on-sky performance. In the laboratory, with a level of disturbance similar to what is encountered in a post Adaptive Optics beam, CLOWFS post-processing has achieved speckle calibration to 1/300 of the raw speckle level. This is about 40 times better than could be done with an idealized PSF subtraction that does not rely on CLOWFS. ",
        "introduction": "Since the discovery of 51 Pegasi by \\citet{1995Natur.378..355M}, several hundreds extrasolar planets have been uncovered, for the most part via indirect detection methods such as radial velocity measurements of the reflex motion of their host star and photometric transit. Direct imaging was recently able to produce the first high contrast images of  \\emph{solid} planetary candidates orbiting nearby main sequence stars: Fomalhaut \\citep{2008Sci...322.1345K}, Beta Pictoris \\citep{2009A&A...493L..21L} and HR 8799 \\citep{2008Sci...322.1348M, 2009ApJ...694L.148L}. Following up on these direct detections, \\citet{2010ApJ...710L..35J} obtained the first direct extrasolar planet spectrum, opening the way to better characterization of planetary atmospheres. These direct detections are currently limited to planets at large orbital separations of several dozens of AU, and only probe a small fraction of the distribution of known extrasolar planets, for which  the median orbital separation is likely closer to 1 AU.  At near infrared wavelength, an 8-meter class telescope provides sufficient angular resolution (40 milli-arcseconds at $\\lambda$=$\\sim$1.6 ${\\mu}$m) to be able to detect companions in the Habitable Zone of nearby stars (d $<$ 30 pc). High contrast imaging near the diffraction limit however requires very good control and calibration of the wavefront aberrations in the optical system. From the ground, this task is complicated by the rapidly changing atmospheric wavefront creating speckled images. While Adaptive Optics (AO) offers a major improvement, performance is still limited and direct imaging of planets often requires post-processing techniques such as Angular Differential Imaging (ADI) \\citep{2006ApJ...641..556M}, which is most efficient at large ($\\gtrapprox$ 10 $\\lambda/D$) angular separations.  In the near future, high contrast imaging systems employing new coronagraphic and wavefront control techniques will greatly improve our ability to image exoplanets. These Extreme-AO systems include the Gemini Planet Imager (GPI) \\citep{2006SPIE.6272E..18M}, ESO's SPHERE \\citep{2008SPIE.7014E..41B} and Subaru's SCExAO \\citep{2009SPIE.7440E..20M}. They will use efficient coronagraphs and high-speed high-order wavefront corrections to reduce speckles in the coronagraphic image. The images of the HR 8799 planetary system obtained by \\citet{2010AAS...21537706S} on a well-corrected 1.5 meter aperture using a high-performance coronagraph demonstrate the relevance of this approach.  Imaging companions close to the edge of the occulting mask in a coronagraph (that is $\\sim$40 mas on SCExAO) however remains an unachieved feat. The current state of the art for PSF calibration is an optimized version of ADI called LOCI introduced by \\citet{2007ApJ...660..770L}. LOCI can calibrate out high-order aberrations to a very high level of contrast (12 magnitudes and higher), but like any technique relying on ADI, only for angular separations greater than 10 $\\lambda/D$. Detection at the edge of a $\\sim$ 1 $\\lambda/D$ occulter is extremely sensitive to low-order aberrations such as pointing and focus.  Of these aberrations, pointing is especially critical, since near the occulter, a tip-tilt excursion along a given direction mimics the signal of a true companion in a coronagraphic image. This issue, first identified for high contrast space borne coronagraphs, has been addressed by \\cite{2009ApJ...693...75G}, with the Coronagraphic Low-Order Wavefront Sensor (CLOWFS), a scheme using the light occulted by a modified focal plane mask as an accurate pointing tracker.  The idea of using the light otherwise lost in the coronagraphic focal plane for tracking is not a novelty. It was for instance successfully implemented on the LYOT project \\citep{2006ApJ...650..484D}. The calibration unit of GPI also uses the light occulted by the focal plane mask to measure low order aberrations, after re-imaging the pupil in a Shack Hartmann type wavefront sensor \\citep{2010SPIE.7736E.179W}: pointing performance with this scheme reaches 2 mas for typical expected observing conditions.  Maximum sensitivity to pointing errors is reached when the light from opposite edges of the pupil is allowed to interfere, which naturally happens in the focal plane. In this respect, while robust to a wide range of errors, a Shack-Hartman doesn't appear optimal to measure pointing: because it splits the pupil into sub-pupils, this capability is lost, resulting in lesser performance than a focal-plane based wavefront sensing scheme \\citep{2005ApJ...629..592G}.  The originality of the CLOWFS design resides in its dual-zone focal plane mask, designed to suppress a strong offset to the signal, carrying most of the power but no information, in a manner reminiscent of strioscopy. The suppression of this offset turns the otherwise imperceptible changes due to small pointing errors into a macroscopic change of the CLOWFS image. Using this scheme, \\citet{2009ApJ...693...75G} were able to stabilize tip-tilt at the level of $10^{-3} \\lambda/D$ in a closed-loop system for $\\lambda=633$ nm, in a laboratory experiment. The current implementation of CLOWFS on SCExAO exhibits pointing residuals $<$0.2 mas, with a 50 Hz frame rate.  While this level of performance is quite remarkable, we demonstrate in this work that additional calibration can be achieved in post-processing, and lead to an improved subtraction of coronagraphic leaks due to low-order aberrations in a long exposure. Using the SCExAO system as a testbed, we experimentally demonstrate a 40 times improvement of the detection limit over a classical calibration procedure at angular separation of a few $\\lambda/D$.  This paper is organized as follows : in Section~\\ref{sec:method}, we introduce how to use CLOWFS for post processing of coronagraphic images. In Section~\\ref{sec:exp}, we describe the implementation of the concept on the SCExAO experiment of the Subaru telescope, and we present our results in Section~\\ref{Sec:results}. In Section~\\ref{sec:discussion}, we summarise our results and discuss possible updates to the presented CLOWFS configuration.   \\begin{figure*}[htb!] \\centerline{\\includegraphics[scale=0.4]{f1.eps}} \\caption{Optical layout of SCExAO used in Lyot-coronagraph mode. A mask lit by an IR ($\\lambda=1.55$ ${\\mu}m$) laser diode emulates the Subaru Telescope pupil. The beam is focused on the dual-zone focal-plane occulting mask described in the text. The light reflected by the focal-plane mask feeds the (slightly defocused) CLOWFS detector, used to characterize pointing errors. The light that is not intercepted by the mask is then re-imaged on the \"science\" detector. Examples of images obtained on both cameras are presented in Fig.~\\ref{fig:clowfs_img} and \\ref{fig:sci_img}.} \\label{fig:bench} \\end{figure*} ",
        "conclusions": "\\label{sec:discussion}  The results presented in Sec~\\ref{Sec:results} demonstrate that, used in post-processing, CLOWFS improves the sensitivity of a coronagraph by a factor 40, in comparison to more conventional -yet idealized- type of calibration. The method described here for this proof of concept is yet still quite rudimentary, and could benefit from several improvements.  Instead of relying on an exhaustive search by a MMA-like algorithm on a massive database, necessarily containing redundant information, the dictionary could be simplified using approaches like Principal Component Analysis, in which CLOWFS images would be projected onto a smaller sub-set of CLOWFS modes. In addition, an extended knowledge of several experimental parameters, such as the telescope elevation, star color or magnitude, could be implemented in the dictionary so as to improve on the actual calibration.  The approach is nevertheless powerful as it requires neither modelling of the coronagraph, nor assumptions such as linearity of the response in order to work. The absence of such requirements makes it extremely robust, and able to handle multiple situations and types of coronagraphic leaks.  The Subaru Coronagraphic Extreme AO Project implements a CLOWFS, used in close-loop. The system architecture however allows for a simple implementation of the CLOWFS post-processing, with no impact on the hardware, to calibrate residual low-order modes induced coronagraphic leaks. CLOWFS on SCExAO currently generates a steady 4 Mbits/s data stream. While not prohibitively large for an experiment, the storage of entire nights of CLOWFS data will quickly prove impractical as SCExAO enters a more aggressive observing phase. Eventually, the calibration of coronagraphic leaks using CLOWFS will require the matching to be performed on-the-fly, so that only the final coronagraphic leak image is actually saved.  Because it makes no assumption on targets and the current state of the optical bench, CLOWFS used both for live and post-processing is likely to play a significant part to the success of SCExAO, focused on the detection of planetary companions at small angular separation."
    },
    "1112/1112.5430_arXiv.txt": {
        "abstract": "For a narrow band of values of the top quark and Higgs boson masses, the Standard Model Higgs potential develops a false minimum at energies of about $10^{16}$ GeV, where primordial Inflation could have started in a cold metastable state. A graceful exit to a radiation-dominated era is provided, {\\it e.g.}, by scalar-tensor gravity models. We pointed out that if Inflation happened in this false minimum, the Higgs boson mass has to be in the range $126.0 \\pm 3.5$ GeV, where ATLAS and CMS subsequently reported excesses of events. Here we show that for these values of the Higgs boson mass, the inflationary gravitational wave background has be discovered with a tensor-to-scalar ratio at hand of future experiments. We suggest that combining cosmological observations with measurements of the top quark and Higgs boson masses represents a further test of the hypothesis that the Standard Model false minimum was the source of Inflation in the Universe. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.1025_arXiv.txt": {
        "abstract": "We present $K$-band integral field observations of the circumstellar envelope of the evolved star \\ohtto. Spatial and spectral information were simultaneously acquired using the {\\sc sinfoni} integral field unit, with adaptive optics, on the Very Large Telescope. The observations reveal the discovery of H$_2$ emission (1) around the centre of the nebula and (2) located in clumps along the Western side of the Northern lobe, presumably associated with the strong shocks that stimulate the previously reported \\halpha~emission at the same location. An observed H$_2$ \\rto~line ratio of ~8.3$\\pm$1.9 was calculated for the central field, a value consistent with shock excitation. ",
        "introduction": "\\label{sec:Intro}  OH~231.8+4.2 (hereafter OH231) is an O-rich late spectral type (M) central star (Mira variable, QX Pup) with bipolar high-velocity dust and gas outflows~\\citep{2001A&A...373..932A},  filamentary structures observed in scattered and molecular line emission, and large angular size (10$^{\\prime\\prime}$$\\times$60$^{\\prime\\prime}$). Often labeled as a post-AGB object or pre-planetary nebula, the presence of both a Mira central star and a main-sequence companion of spectral type A \\citep{2004ApJ...616..519S} seems to contradict this classification. OH231 is more likely a D-type bipolar symbiotic system \\citep[][]{2010PASA...27..129F}. However in some cases, morphological similarities do exist between post-AGB and symbiotic objects, most strikingly the presence of highly collimated bipolar nebulae. It is via fast-collimated outflows that these stars shape their surrounding nebula. Understanding the development and origin of these fast outflows is critical for advancing hydrodynamical models of wind interaction. Recent work by \\citet{2009ApJ...696.1630L} attempting to reproduce the high velocity molecular emission in AFGL~618 using collimated fast wind models, emphasises the need for further observations and model development in this area.   OH231 has been the subject of many studies spanning multiple wavelength ranges, for example: \\citet{1985ApJ...297..702C} were first to propose the existence of a binary companion; \\citet{2002A&A...389..271B} imaged the shape of the shocks using H${\\alpha}$ (reproduced in Fig.~\\ref{fig:oh231_naco_sin_wfpc2} [a]) detected with the {\\em Hubble Space Telescope (HST)}; \\citet{2003ApJ...585..482M} report {\\em HST}/NICMOS NIR images of the dust distribution and hence a high resolution map of the extinction through the nebula. \\citet{2006ApJ...646L.123M} using the MIDI and NACO instruments on the Very Large Telescope (VLT) detected a compact circumstellar disc. The envelope of OH231 is also known to be rich in molecular species (e.g. H$_2$O, OH, and SiO) however previous studies in the NIR have all returned null detections of H$_2$ \\citep[e.g.][]{1998ApJ...509..728W,2006ApJ...646L.123M}.  In this Letter, we present the results of preliminary observations of OH231 at NIR ($K$-band) wavelengths showing for the first time the presence of shock-excited H$_2$. Throughout this work we assume OH231 is a member of the open cluster M46 at a distance of 1.3~kpc~\\citep{1985ApJ...292..487J}. The origin of the coordinate system used in Figures~1,3, and 4 is given by the location of the SiO maser emission at RA=07$^h$42$^m$16$^s$.93, Dec=-14$^\\circ$42$^{\\prime}$50$^{\\prime\\prime}$.2 (J2000)~\\citep{2002A&A...385L...1S}, and the inclination angle of the bipolar axis is 36$^\\circ$~to the plane of the sky \\citep{1992ApJ...398..552K}. ",
        "conclusions": "\\label{sec:discuss}   We have presented VLT/{\\sc sinfoni} integral field observations of OH231, revealing the presence of several ro-vibrational H$_2$ lines. The main conclusions are: \\begin{list}{\\makelabel{-}}{\\itemsep=-0.2em \\leftmargin=1em} \\item The discovery of H$_2$ emission near the centre of OH231, possibly located at the edge of an axisymmetric shell or an incomplete disc. \\item A \\rto~value of 8.3$\\pm$1.9 is found for the equatorial H$_2$, suggesting a collisional excitation mechanism. \\item Our observations of the central shell/disc of H$_2$ show no velocity structure. However, higher S/N and/or velocity resolution data are needed to accurately probe the kinematics in this region. \\item We detect fast-moving H$_2$ (\\simm 110 \\kms, along the bipolar axis) via the \\ozso transition along the North-Western tip of the nebula, a region where a strong \\halpha~bow shock exists. Due to the small FOV of our observations, the full extent of this H$_2$ is unknown. \\end{list}"
    },
    "1112/1112.3020_arXiv.txt": {
        "abstract": "Recent measurements of cosmic-ray spectra of several individual nuclear species by the CREAM, TRACER, and ATIC experiments indicate a change in the spectral index of the power laws at TeV energies. Possible explanations among others include non linear diffusive shock acceleration of cosmic-rays, different cosmic-ray propagation properties at higher and lower energies in the Galaxy and the presence of nearby sources. In this paper, we show that if supernova remnants are the main sources of cosmic rays in our Galaxy, the effect of the nearby remnants can be responsible for the observed spectral changes. Using a rigidity dependent escape of cosmic-rays from the supernova remnants, we explain the apparent observed property that the hardening of the helium spectrum occurs at relatively lower energies as compared to the protons and also that the spectral hardening does not persist beyond $\\sim (20-30)$ TeV energies. ",
        "introduction": "Recently, cosmic-ray (CR) measurements by the new-generation balloon-borne experiments such as the ATIC (Panov et al. 2007), CREAM (Yoon et al. 2011), and TRACER (Ave et al. 2008) seem to indicate that the CR spectrum deviates from a single power law. The spectra of all individual elements seem to be harder at TeV energies than at lower energies. Such a hardening is not easy to explain under the standard models of CR acceleration and their propagation in the Galaxy. Under the standard theory,  CRs below the knee ($\\sim 3$ PeV) are considered to be produced by supernova remnant (SNR) shock waves by diffusive shock acceleration  mechanism (Bell 1978, Blandford $\\&$ Eichler 1987). Such a mechanism naturally predicts a power law spectrum of $E^{-\\gamma}$ with the index $\\gamma=2$ for strong shocks. On the other hand, CR propagation in the Galaxy is considered to be of diffusive nature which is due to scattering by magnetic field irregularities and the CR self excited Alfven and hydromagnetic waves present in the Galaxy. Measurements of CR secondary-to-primary ratios indicate that the diffusion is energy dependent with the diffusion coefficient $D(E)\\propto E^a$ with $a\\approx (0.3-0.7)$. Under these considerations, the CR spectrum in the Galaxy is expected to follow a single power law with index $(\\gamma+a)$ up to the knee, which do not seem to agree quite easily with the observed hardening at TeV energies.  The observed data can be explained if either the source spectrum or the diffusion index flattens at higher energies. Non linear diffusive shock acceleration theories where CRs modify the shock structure predict concave spectra (flatter at higher energies) at the shocks. But, the total spectrum injected into the interstellar medium which is the sum of the instantaneous spectra over the SNR life time is very close to a pure power-law (Caprioli et al. 2010). The concave signature can be even more diluted when summed over an ensemble of SNRs (Ptuskin $\\&$ Zirakashvili 2005). From the propagation point of view, there are models which assume a harder or constant CR diffusion coefficient at higher energies in the Galaxy (Ave et al. 2009). Such models are motivated not only by the apparent flattening of the observed boron to carbon ratio above $\\sim 100$ GeV energies, but also by the observed CR anisotropy which is almost independent of energy. Recently, it has also been proposed that dispersion in the spectral indices of CR source spectrum from many sources can also be responsible for the observed spectral hardening (Yuan et al. 2011).  Another possible explanation, as also pointed out in Ahn et al. 2010, is the presence of nearby sources. Erlykin $\\&$ Wolfendale 2011 suggested that an extra component of CRs with a steep spectrum could be contributing below $\\sim 200$ GeV/n while above that, the spectrum is entirely determined by a harder CR background. They proposed that the sources of the extra component could be in OB associations in the Local Bubble. Recently, Ohira $\\&$ Ioka 2011 proposed that the hardening could be due to decreasing Mach number in hot superbubbles with multiple supernovae. In another recent work, Vladimirov et al. 2011 investigated several possible interpretations (including local source effect) for the observed  spectral features at low and high energies using the GALPROP propagation code. They also presented the possible effects on other observed properties such as CR anisotropy, isotopic ratios and the Galactic diffuse $\\gamma$-ray emissions.  In our present study, we investigate whether the spectral hardening observed at TeV energies could be an effect of the nearby SNRs. Although there has not been any direct detection of CRs from any sources, SNRs remain the most favorable candidates both theoretically as well as observationally. At least the presence of high energy particles up to few TeVs inside SNRs have been confirmed by the detections of non-thermal X-rays (Parizot et al. 2006) and TeV $\\gamma$-rays from several SNRs (Aharonian et al. 2006, 2008a). Moreover, the detection of TeV electrons by the HESS experiment (Aharonian et al. 2008b) indicates the presence of one or more CR sources within a distance of $\\sim 1$ kpc from us. If these sources produce both, electrons as well as nuclei, we expect to see some effects on the spectra of CR nuclei observed at the Earth.  \\section {Model} The diffusive propagation of CRs in the Galaxy neglecting the effects due to nuclear spallation can be described by the following equation, \\begin{equation} \\nabla\\cdot(D\\nabla N)+Q=\\frac{\\partial N}{\\partial t} \\end{equation} where $N(\\textbf{r},E,t)$ is the differential number density at a distance $\\textbf{r}$ from the source at time $t$, $E$ is the kinetic energy/nucleon and $Q(\\textbf{r},E,t)$ is the source term. The diffusion coefficient is taken as $D(\\Re)=D_0(\\Re/\\Re_0)^a$ for $\\Re>\\Re_0$, where $\\Re$ denotes the particle rigidity which is given by $\\Re=AE/Z$ for charge $Z$ and mass number A. For our study, we consider two sets of values for $(D_0, \\Re_0, a)$: one based on purely diffusion model (hereafter Model A) and the other based on models including CR reacceleration due to interstellar turbulence (hereafter Model B). We choose $(D_0, \\Re_0, a)=(2.9, 3, 0.6)$ for Model A (Thoudam 2008) and $(5.8, 4, 0.33)$ for Model B (Strong et al. 2010), where $D_0$ is in units of $10^{28}$ cm$^2$ s$^{-1}$ and $\\Re_0$ is in GV.  Under diffusive shock acceleration theory, CRs are confined within the remnant due to the magnetic turbulence generated by the CRs themselves. They can escape when their upstream diffusion length defined as $l_{diff} = D_s(E)/u_s$ is greater than the escape length from the shock front which is usually taken as $l_{esc}\\approx 0.1 R_s$, where $u_s$ and $R_s$ denote the shock velocity and the shock radius respectively. In the Bohm diffusion limit, the upstream diffusion coefficient scales linearly with energy as $D_s(E)\\propto E$ which implies that higher energy particles can escape the remnant at early times followed later by the lower energy ones. But, the exact energy dependence of $D_s$ is still not well understood and depends on some poorly known quantities which include the spectral distribution of the CR self-excited turbulence waves, the level of magnetic field amplification by the CRs and the dynamical reaction of CRs on the shock structure. Therefore, we follow a simple but reasonable parameterization for the CR escape time similar to that adopted by Gabici et al. 2009 as given below, \\begin{equation} t_{esc}(\\Re)=t_{sed}\\left(\\frac{\\Re}{\\Re_{max}}\\right)^{-1/\\alpha} \\end{equation} where $t_{sed}$ denotes the start of the Sedov phase, $\\Re_{max}$ denotes the maximum CR rigidity and $\\alpha$ is a positive constant. We assume that the maximum CR energy accelerate by an SNR scales with the charge number $Z$ as $ZU_{max}$, where $U_{max}$ denotes the maximum kinetic energy of the protons which is taken as $1$ PeV for our study (Berezhko 1996). This scaling gives $\\Re_{max}=1$ PV. In units of energy/nucleon, the maximum energy for helium is $E_{max}=0.5$ PeV/n.  Eq. (2) assumes that the highest energy CRs of all the species start escaping at the onset of the Sedov phase itself. Writing Eq. (2) in terms of total kinetic energy, it is easy to check that for the same kinetic energy, the escape time of CRs scales with the charge number as $Z^{1/\\alpha}$, i.e, higher charged particles escape at relatively later stages of the SNR evolution. Thus, our escape model takes into account the general understanding of diffusive shock acceleration theory that higher charged particles can be confined for relatively longer duration within  the remnant. In terms of energy/nucleon, we can write Eq. (2) as \\begin{equation} t_{esc}(E)=t_{sed}\\left(\\frac{AE}{Z\\Re_{max}}\\right)^{-1/\\alpha} \\end{equation} Eq. (3) shows that for the same energy/nucleon, all nuclei with charge $Z>1$ escape earlier than the protons by a factor of $(A/Z)^{-1/\\alpha}$. We further assume that no particles remain confined after the shock completely dies out which we assume to occur when the SNR age $10^5$ yr. Taking this into account, the CR escape time for our study is taken as $T_{esc}(E)= \\mathrm{min}\\left[t_{esc}(E), 10^5 \\mathrm{yr}\\right]$. For detailed studies on particle escape from SNRs, see e.g., Ptuskin $\\&$ Zirakashvili 2005, Caprioli et al. 2009 and Ohira et al. 2010.  The corresponding escape radius of CRs is calculated using the age-radius Sedov relation for SNRs as given below, \\begin{equation} R_{esc}(E)=2.5u_0\\;t_{sed}\\left[\\left(\\frac{T_{esc}}{t_{sed}}\\right)^{0.4}-0.6\\right] \\end{equation} where $u_0$ represents the initial shock velocity, i.e the velocity at $t=t_{sed}$.  The source term in Eq. (1) is taken as, \\begin{equation} Q(\\textbf{r},E,t)=\\frac{q(E)}{A_{esc}}\\delta(t-T_{esc})\\delta(r-R_{esc}) \\end{equation} where $A_{esc}=4\\pi R_{esc}^2$ denote the surface area of the SNR at the time when CRs of energy $E$ escape the remnant. It should be noted that our consideration of the rigidity dependent escape time and the finite source size are different from the commonly adopted burst-like point source approximation where CRs of all rigidities are assumed to escape at the same time from a point source. For CR study near the sources, the point source approximation can break down and it looks more realistic to take their sizes into account (Thoudam $\\&$ H\\\"orandel 2011). Recently, such importance has also been highlighted in Ohira et al. 2011 in the study of $\\gamma$-ray emission from SNRs interacting with molecular clouds.  The source spectrum in Eq. (5) is taken as $q(E)=Aq(U)$ with $q(U)$ given by, \\begin{equation} q(U)=k(U^2+2Um)^{-(\\gamma+1)/2}(U+m) \\end{equation} where $U=AE$ represents the particle total kinetic energy, $m$ is the rest mass energy and $k$ is the normalization constant which is related to the CR injection efficiency.  Solving Eq. (1), the spectrum at a distance $r_s$ from the SNR follows, \\begin{eqnarray} N(r_s,E,t)=\\frac{q(E)\\,R_{esc}}{r_sA_{esc}\\sqrt{\\pi D(t-T_{esc})}}\\mathrm{exp}\\left[-\\frac{\\left(R_{esc}^2+r_s^2\\right)}{4D(t-T_{esc})}\\right]\\nonumber\\\\ \\times\\;\\mathrm{sinh}\\left(\\frac{r_sR_{esc}}{2D(t-T_{esc})}\\right)\\; \\end{eqnarray} For high energy particles for which the diffusion radius defined as $r_{diff}=\\sqrt{D(t-T_{esc})}$ is much larger than $(r_s, R_{esc})$, Eq. (7) follows a power-law of the form $N(E)\\propto E^{-\\left(\\Gamma+\\frac{3}{2}a\\right)}$.  Eq. (7) can be used to calculate the CR spectra from the nearby SNRs. We choose proton and helium for our study and consider only those SNRs with distances $<1$ kpc from the Earth and ages $< 2\\times 10^5$ yr. From the available literature, we found 10 SNRs listed as follows with their distances (kpc) and ages (yr)  respectively given in parentheses: Cygnus Loop $(0.54, 10^4)$, HB21 $(0.8, 1.9\\times 10^4)$, HB9 $(0.8, 6.6\\times 10^3)$, S147 $(0.8, 4.6\\times 10^3)$, Vela $(0.3, 1.1\\times 10^4)$, G299.2-2.9 $(0.5, 5\\times 10^3)$, SN185 $(0.95, 1.8\\times 10^3)$, Monogem $(0.3, 1.1\\times 10^5)$, G114.3+0.3 $(0.7, 4.1\\times 10^4)$ and Vela Junior $(0.75, 3.5\\times 10^3)$.  In addition to the contributions from the nearby SNRs, we assume that there exists a steady CR background in the Galaxy which dominates the overall CR spectrum. For the CRs observed at the Earth, we assume that this background component consists of contributions from distant SNRs plus any other possible sources in the Galaxy. For our study, we obtain the background by fitting the observed CR spectrum between $(20-200)$ GeV/n. This is the energy region where the contamination from the nearby sources is expected to be less and at the same time, not much affected by the Solar modulation. In fact, it has been shown in Thoudam 2008 that the presence of nearby sources can produce stronger density fluctuations at higher energies than at lower energies because of the energy dependent nature of CR diffusion. Therefore, we believe that it is reasonable to assume that the low energy CRs that we observe at the Earth are not much affected by the presence of nearby SNRs and they largely represent the averaged background spectrum in the Galaxy. We will show in the following that this is indeed the most likely case. ",
        "conclusions": ""
    },
    "1112/1112.3957_arXiv.txt": {
        "abstract": "Approximately $14$\\% of known Galactic open clusters possess absolute errors $\\le 20$\\% as evaluated from $n\\ge3$ independent distance estimates, and the statistics for age estimates are markedly worse.  That impedes such diverse efforts as calibrating standard candles and constraining masses for substellar companions. New data from the VVV survey may be employed to establish precise cluster distances with comparatively reduced uncertainties ($\\le10$\\%).  This is illustrated by deriving parameters for Pismis~19 and NGC~4349, two pertinent open clusters which hitherto feature sizable uncertainties ($60$\\%). Fundamental parameters determined for Pismis~19 from new VVV $JHK_s$ photometry are $d=2.40\\pm0.15$ kpc, $<{\\it E_{J-H}}>=0.34\\pm0.04$, and $\\log{\\tau}=9.05\\pm0.10$, whereas for NGC~4349 the analysis yielded $d=1.63\\pm0.13$ kpc, ${\\it E_{J-H}}=0.09\\pm0.02$, $\\log{\\tau}=8.55\\pm0.10$.  The results exhibit a significant ($\\ge 5 \\times$) reduction in uncertainties, and indicate that: $i$) existing parameters for the substellar object NGC~4349 127b require revision, in part because the new cluster parameters imply that the host is $20$\\% less-massive (${\\cal M_{*}/M_{\\sun}} \\sim 3.1$);  $ii$) R Cru is not a member of NGC~4349 and should be excluded from period-Wesenheit calibrations that anchor the distance scale; $iii$) and results for Pismis~19 underscore the advantages gleaned from employing deep VVV $JHK_s$ data to examine obscured ($A_V \\sim 4$) and differentially reddened intermediate-age clusters. ",
        "introduction": "} Approximately $30$\\% of the 395 open clusters featuring $n\\ge3$ independent distance estimates exhibit absolute errors $\\ge 20$\\%  \\citep[][their Fig.~2]{pa06}.  There are $\\ge 2 \\times 10^3$ cataloged Galactic open clusters \\citep{di02}, implying that merely $\\sim 14$\\% of the known sample possess errors $\\le 20$\\% as evaluated from three distance estimates.  The uncertainties permeate into analyses which rely on the cluster zero-point, such as the calibration of any constituent standard candles or substellar companions \\citep{lm07,ma11b}.   Consider that published parameters for NGC~4349 span $d=900-2200$ pc and $\\tau=0.1-0.6$ Gyr (\\S \\ref{s-ngc4349}).  Yet physical parameters for the substellar companion to TYC 8975-2601-1 \\citep{lm07,ka08} rely on those inferred for the host from cluster membership (NGC~4349).  Furthermore, the nearer distance and younger age for NGC 4349 potentially imply cluster membership for the classical Cepheid R Cru, which lies within the cluster's corona. Establishing cluster membership would enable the subsequent calibration of Cepheid period-luminosity and period-Wesenheit relations \\citep{tu10}.  Such functions bolster efforts to establish extragalactic distances and zero-point the SNe Ia scale \\citep[e.g.,][]{pg04}.  The aforementioned examples underscore the broad ramifications of an uncertain cluster scale.  Admittedly, age estimates for open clusters are less reliable since a third exhibit absolute errors $>50$\\% \\citep[][$n\\ge3$]{pa06}, and presumably the statistics worsen for obscured clusters near the Galatic plane.  In this study, new VVV (VISTA Variables in the V\\'{i}a L\\'actea) $JHK_s$ photometry is employed to illustrate the marked improvement that can be achieved \\textit{vis \\`a vis} open cluster distances.  Two important clusters featuring particularly discrepant published parameters are examined, namely Pismis~19 and NGC~4349.  Distances for the clusters display a $\\sim 60$\\% spread and individual uncertainties of $\\sim 30$\\%.  Efforts to secure precise parameters for Pismis~19 via optical photometry have been complicated by differential reddening and $A_V\\sim4$ magnitudes of obscuring dust. Parameters for Pismis~19 and NGC~4349 derived here exhibit a marked ($>5 \\times$) reduction in uncertainties (\\S \\ref{s-analysis}), and highlight the advantages of using VVV data to determine reliable cluster distances and compliment existing efforts. ",
        "conclusions": "} VVV $JHK_s$ observations may be employed to help establish precise cluster distances that feature comparatively reduced uncertainties ($\\le 10$\\%).  That is illustrated by deriving fundamental parameters for Pismis~19 and NGC~4349, two important clusters which hitherto exhibit sizable uncertainties ($60$\\%, \\S \\ref{s-pismis19} and \\ref{s-ngc4349}).  A precise distance determination for Pismis~19 from optical photometry was hampered in part by significant reddening (Fig.~\\ref{fig-dr}, $A_V \\sim 4$).  The existing ambiguity surrounding the distance to NGC~4349 ensured that the pertinence of invaluable putative constituents were mitigated (i.e., the classical Cepheid R Cru and a substellar companion for the  member TYC 8975-2601-1).  Parameters derived for Pismis~19 are: $d=2.40\\pm0.15$ kpc, $<{\\it E(J-H)}>=0.34\\pm0.04$, $\\log{\\tau}=9.05\\pm0.10$ (Fig.~\\ref{fig-dr}), whereas NGC~4349 exhibits $d=1.63\\pm0.10$ kpc, ${\\it E(J-H)}=0.09\\pm0.02$, $\\log{\\tau}=8.55\\pm0.10$ (Fig.~\\ref{fig-cm}).  The nature of the VVV survey ensured that the revised results, which have pertinent ramifications, compliment existing estimates and display a marked improvement ($\\ge 5 \\times$) in precision. New VVV $JHK_s$ for stars in NGC~4349 and Pismis~19 imply that: existing physical parameters derived for NGC~4349 127b need to be redetermined in part since the mass for the host star was revised downward to ${\\cal M_{*}/M_{\\sun}}\\sim 3.1$ (\\S \\ref{s-ngc4349}); the classical Cepheid R Cru is not a member of NGC~4349 (\\S \\ref{s-ngc4349});  and VVV $JHK_s$ photometry is particularly suited for constraining parameters of obscured and differentially reddened intermediate-age clusters (e.g., Pismis~19, $A_V \\sim 4$, Fig.~\\ref{fig-dr}).  The VVV and UKIDSS surveys \\citep{lu08,mi10} may be employed to achieve significant gains toward strengthening the open cluster distance scale.  Yet considerable work remains, and unknown systematic errors may be discovered given the nascent nature of the aforementioned surveys.  Consequently, obtaining independent multiband observations are desirable to corroborate derived cluster parameters \\citep[e.g.,][]{ca11}.  \\subsection*{{\\rm \\scriptsize ACKNOWLEDGEMENTS}} \\scriptsize{DM is grateful to the following individuals and consortia whose efforts lie at the foundation of the research: 2MASS, P. Stetson (DAOPHOT), W. Lohmann, WEBDA (E. Paunzen), DAML (W. Dias), C. Lovis, M. Mayor, CDS, arXiv, and NASA ADS.  We gratefully acknowledge use of data from the ESO Public Survey programme ID 179.B-2002 taken with the VISTA telescope, the Cambridge Astronomical Survey Unit, and funding from the FONDAP Center for Astrophysics 15010003, the BASAL CATA Center for Astrophysics and Associated Technologies PFB-06, the MILENIO Milky Way Millennium Nucleus from the Ministry of Economics ICM grant P07-021-F, and Proyecto FONDECYT Regular 1090213. WG, CMB, and DG are grateful for support from the Chilean Center for Astrophysics FONDAP 15010003 and the BASAL Centro de Astrofisica y Tecnologias Afines (CATA) PFB-06/2007. RK acknowledges support from Proyecto DIUV23/2009, Universidad de Valparaiso.  RS acknowledges financial support from CONICYT through GEMINI Project Nr. 32080016.}"
    },
    "1112/1112.4746_arXiv.txt": {
        "abstract": "The Ly-${\\alpha}$ line of He {\\sc ii} at 304 \\AA\\ is one of the spectral lines of choice for EUV channels of narrow-band imagers on board space telescopes, which provide spectacular intensity images of the outer solar atmosphere. Since the magnetic field information is encoded in the polarization of the spectral line radiation, it is important to investigate whether the He {\\sc ii} line radiation from the solar disk can be polarized, along with its magnetic sensitivity. Here we report some theoretical predictions concerning the linear polarization signals produced by scattering processes in this strong emission line of the solar transition region, taking into account radiative transfer and the Hanle effect caused by the presence of organized and random magnetic fields. We find that the fractional polarization amplitudes are significant (${\\sim}1\\%$), even when considering the wavelength-integrated signals. Interestingly, the scattering polarization of the Ly-${\\alpha}$ line of He {\\sc ii} starts to be sensitive to the Hanle effect for magnetic strengths $B{\\gtrsim}100$ G (i.e., for magnetic strengths of the order of and larger than the Hanle saturation field of the hydrogen Ly-${\\alpha}$ line at 1216 \\AA). We therefore propose simultaneous observations of the scattering polarization in both Ly-${\\alpha}$ lines to facilitate magnetic field measurements in the upper solar chromosphere. Even the development of a narrow-band imaging polarimeter for the He {\\sc ii} 304 \\AA\\ line alone would be already of great diagnostic value for probing the solar transition region. ",
        "introduction": "The understanding of the complex interface region between the photosphere and corona of the Sun is a very important challenge in astrophysics. In this highly structured and dynamic region of the outer solar atmosphere, where magnetic and hydrodynamic forces compete for dominance, most of the non-thermal energy that creates the corona and solar wind is released. Novel measurements of key physical quantities, like the magnetic field, would improve our understanding of this boundary region. Spectroscopic observations are needed for determining temperatures, flows and waves, while the magnetic field information is encoded in the spectral line polarization (e.g., Stenflo 1994; Landi Degl'Innocenti \\& Landolfi 2004). The aim of this paper is to propose a new technique, based on spectro-polarimetric measurements, which may be particularly useful for determining the magnetic field vector in the solar transition region.  The spectral lines that originate in the outer solar atmosphere (chromosphere, transition region and corona) are mainly located in the FUV (912-3000 \\AA) and EUV (100-912 \\AA) spectral ranges, such as the Ly-${\\alpha}$ lines of H {\\sc i} and He {\\sc ii} at 1216 \\AA\\ and 304 \\AA, respectively. Their intensity $I(\\lambda)$ profiles are practically insensitive to the magnetic fields one may expect in the plasma of the outer solar atmosphere, so they cannot be used to obtain quantitative information on the strength ($B$), inclination ($\\theta_B$) and azimuth ($\\chi_B$) of the magnetic field vector. A similar drawback applies to the circular polarization (quantified by the Stokes $V(\\lambda)$ profile) produced by the longitudinal Zeeman effect, because the Stokes-$V$ amplitude scales with the ratio, ${\\cal R}$, between the Zeeman splitting and the Doppler line width. For such solar spectral lines ${\\cal R}{\\ll}1$, especially outside sunspots (${\\cal R}={{1.4{\\times}10^{-7}{\\lambda}B}{/}{\\sqrt{1.663{\\times}10^{-2}T/\\alpha+{\\xi}^2}}}$, where $\\lambda$ is the spectral line wavelength in \\AA, $T$ the kinetic temperature in K, $\\xi$ the microturbulent velocity in ${\\rm km}\\,{\\rm s}^{-1}$, $\\alpha$ the atomic weight of the atom under consideration, and $B$ is in gauss; see Landi Degl'Innocenti \\& Landolfi 2004). The situation is even worse for the linear polarization (quantified by the Stokes $Q(\\lambda)$ and $U(\\lambda)$ profiles) produced by the transverse Zeeman effect because their amplitude scales with ${\\cal R}^2$.  In order to obtain quantitative information on the magnetic field of the outer solar atmosphere we need to measure the polarization caused by scattering processes in the spectral lines that form in such regions, ideally using two or more spectral lines with different sensitivities to the Hanle effect (e.g., the review by Trujillo Bueno 2010; see also Stenflo, Keller \\& Gandorfer 1998; Manso Sainz, Landi Degl'Innocenti \\& Trujillo Bueno 2004). We recall that the Hanle effect is the modification of the linear polarization produced by scattering processes in a spectral line, caused by the presence of a magnetic field inclined with respect to the symmetry axis of the incident radiation field. Approximately, the emergent linear polarization is sensitive to magnetic strengths between 0.2$B_H$ and 5$B_H$, where $B_H{=}\\,{1.137{\\times}10^{-7}}/{t_{\\rm life}g}$ is the critical Hanle field for which the Zeeman splitting of the line's level under consideration is equal to its natural width ($t_{\\rm life}$ is the level's radiative lifetime, in seconds, and $g$ its Land\\'e factor).  In a recent paper we showed that the hydrogen Ly-${\\alpha}$ line is expected to show measurable scattering polarization when observing the solar disk and that via the Hanle effect the line-center polarization amplitude is  sensitive to the presence of magnetic fields in the solar transition region, with good sensitivity to magnetic strengths between 10~G and 100~G (see Trujillo Bueno, \\v{S}t\\v{e}p\\'an \\& Casini 2011). The observational signatures of the Hanle effect might however be confused with the symmetry breaking effects caused by the presence of horizontal atmospheric inhomogeneities (e.g., Manso Sainz \\& Trujillo Bueno 2011). Although for the hydrogen Ly-${\\alpha}$ line such symmetry breaking effects can often be distinguished from the Hanle effect (e.g., \\v{S}t\\v{e}p\\'an \\& Trujillo Bueno 2012), it is of great interest to find another transition region line with measurable scattering polarization but such that is practically immune to the weak magnetic fields expected for the quiet regions of the upper solar chromosphere ($B{<}100$ G).  The aim of this Letter is to show some theoretical predictions concerning the linear polarization produced by scattering processes in the Ly-${\\alpha}$ line of He {\\sc ii}, whose significant emission originates in the solar transition region. As we shall see below, the line-center fractional polarization signals are significant (${\\sim}1\\%$) and have a very interesting magnetic sensitivity, which lead us to argue that the development of a space-based instrument capable of obtaining high-resolution $I$, $Q/I$ and $U/I$ images in the He {\\sc ii} Ly-${\\alpha}$ line would be very useful to determine the three-dimensional magnetic structure of the solar transition region. ",
        "conclusions": "The Stokes $I(\\lambda)$ profile of the Ly-${\\alpha}$ line of He {\\sc ii} at 304 \\AA\\ is about 10 times narrower than that of the Ly-${\\alpha}$ line of H {\\sc i} at 1216 \\AA. While only the core of the hydrogen Ly-${\\alpha}$ line originates in the solar transition region, the Ly-${\\alpha}$ line of He {\\sc ii} is emitted mostly in the solar transition region. Noteworthy is also that the observed line-center intensities in the Ly-${\\alpha}$ lines of H {\\sc i} and He {\\sc ii} are of the same order of magnitude (e.g., Fontenla et al. 1993).  The radiative transfer investigation reported in this paper indicates that the line-center fractional linear polarization amplitude of the Ly-${\\alpha}$ line of He {\\sc ii} should be significantly larger than that expected for the Ly-${\\alpha}$ line of H {\\sc i} (e.g., a factor three larger at $\\mu\\,{\\approx}\\,0.3$ in the unmagnetized case). Moreover, the fractional linear polarization amplitude that results from the wavelength-integrated Stokes profiles of the He {\\sc ii} line turns out to be similar to the line-center signal\\footnote{In a forthcoming publication we will show that the fractional linear polarization amplitude obtained from the wavelength-integrated Stokes profiles is even larger when partial frequency redistribution effects are taken into account.}. These results partially compensate the fact that the total number of photons emerging from any solar disk position in the Ly-${\\alpha}$ line of He {\\sc ii} is significantly smaller than within a small (0.1 \\AA) wavelength interval around the hydrogen Ly-${\\alpha}$ line center.  We have shown also that for magnetic strengths $B{<}100$ G the Ly-${\\alpha}$ line of He {\\sc ii} is nearly immune to magnetic fields\\footnote{The fact that with a narrow-band instrument, designed to measure the intensity and polarization of the Ly-${\\alpha}$ line of He {\\sc ii}, we may have a small contribution from the nearby Si {\\sc XI} emission at 303.3 \\AA\\ should not be a problem because the critical Hanle field of this line is also very large (i.e., $B_H{\\approx}730$ G).}. This is particularly interesting because for magnetic strengths between 10 and 100 G the Ly-${\\alpha}$ line of H {\\sc i} is indeed sensitive to the Hanle effect. Therefore, outside active regions the Ly-${\\alpha}$ line of He {\\sc ii} can be used as a reasonable reference line for facilitating magnetic field ``measurements\" via the Hanle effect in the Ly-${\\alpha}$ line of H {\\sc i}.  All these results encourage the development of the following instruments for a space telescope:  \\begin{itemize} \\item A spectropolarimeter for measuring the line-core polarization of the Ly-${\\alpha}$ line of H {\\sc i} with a spectral resolution of at least 0.1 \\AA. \\item A narrow-band polarimeter for obtaining intensity and linear polarization images of the solar transition region in the Ly-${\\alpha}$ light of He {\\sc ii}. \\end{itemize}  Although the combined use of the two Ly-$\\alpha$ lines opens up a new diagnostic window for magnetic-field measurements in the upper solar chromosphere, the interpretation of such spectro-polarimetric observations will still require radiative transfer calculations in realistic atmospheric models because the two spectral lines are not formed in exactly the same way (e.g., Jordan 1975; Fontenla, Avrett \\& Loeser 2002; Pietarila \\& Judge 2004). In spite of such a complication the comparison between spectro-polarimetric observations in the two Ly-$\\alpha$ lines can provide unique insights into the physics and geometry of the transition region.  Finally, it is of interest to note that off-limb observations of resonant scattering polarization in the He {\\sc ii} 304 \\AA\\ line may be also useful for exploring the geometry and magnetic field structure of spicules, prominences and of the solar corona.  {\\bf Acknowledgments} Financial support by the Spanish Ministry of Science and Innovation through projects AYA2010-18029 (Solar Magnetism and Astrophysical Spectropolarimetry) and CONSOLIDER INGENIO CSD 2009-00038 (Molecular Astrophysics: The Herschel and Alma Era) is gratefully acknowledged."
    },
    "1112/1112.0414_arXiv.txt": {
        "abstract": "We present $UBV$ photometry of the highly reddened and poorly studied open cluster Berkeley~55, revealing an important population of B-type stars and several evolved stars of high luminosity. Intermediate resolution far-red spectra of several candidate members confirm the presence of one F-type supergiant and six late supergiants or bright giants. The brightest blue stars are mid-B giants. Spectroscopic and photometric analyses indicate an age $50\\pm10$~Myr. The cluster is located at a distance $d\\approx4$~kpc, consistent with other tracers of the Perseus Arm in this direction. Berkeley~55 is thus a moderately young open cluster with a sizable population of candidate red (super)giant members, which can provide valuable information about the evolution of intermediate-mass stars. ",
        "introduction": "\\label{introduction}  After exhaustion of H in their cores, stars evolve toward lower $T_{{\\rm eff}}$, becoming, according to their masses, red giants or supergiants (RSGs). Both high- and intermediate-mass stars are subject to complex physical processes in their later evolution, which result in important changes in their observable characteristics ($T_{{\\rm eff}}$, $L_{{\\rm bol}}$). As a consequence, their evolutionary tracks trace loops in the HR diagrams \\citep[e.g.,][]{chiosi}. The shape of these loops depends on the physics of the stellar interior, generally modeled via poorly understood parameters \\citep[e.g.,][]{chiosi,mf94,salas99,mm00}. The extent of semi-convection and overshooting has very important consequences on many aspects of stellar evolution, most notably the ratio of initial mass to white dwarf mass \\citep[e.g.,][]{jeff97,weide00} and the boundary between stars that leave white dwarfs as remnants and those that explode as supernovae \\citep[SNe; e.g.,][]{poel08}.  M-type supergiants represent the final evolutionary stage of moderately massive stars, with typical initial masses in the 8--25$\\,M_{\\odot}$ range \\citep{levesque05}. Such objects are the immediate progenitors of Type II-P SNe, the most frequent type of supernova explosion in the Local Universe \\citep{smartt09,smith11}. Most of the explosions come from low-mass RSGs, stars with initial masses $M_{*}\\la12\\,M_{\\sun}$. The lower mass limit for a star to produce a supernova has been estimated at $\\approx 8.5^{+1}_{-1.5}\\,M_{\\sun}$ \\citep{smarttal}. Nevertheless, stars with lower masses are also classified as supergiants, generally K\\,Ib objects, showing that the morphological separation between red supergiants and red giants does not coincide exactly with the boundary between high- and intermediate-mass stars.  Since stars become more luminous as they experience blue loops and later enter the AGB branch, a given star may appear first as a bright giant and later in its evolution as a supergiant. Open clusters with large populations of evolved stars can help constrain the inputs of models and therefore improve our understanding of such basic questions. Identification and study of such clusters is thus an important astrophysical issue.  To this aim, we have analyzed a sample of poorly-studied, cataloged, optically-visible clusters in the Northern Hemisphere, using 2MASS data \\citep{skru06}. We made use of the $Q_{{\\rm IR}}$ index, a reddening-free parameter, defined as $(J-H)-1.70(H-K_{{\\rm S}})$, which has been proved to be a very useful tool for identification of intrinsically blue stars \\citep[see, e.g.,][]{cp05,ns07} and separation of luminous red stars from red dwarfs \\citep{neg11}. Berkeley~55 (Be~55) was readily identified as containing an obvious clump of red luminous stars associated with a sequence of intrinsically blue stars.  Be~55 is a faint, compact open cluster in the constellation Cygnus. The WEBDA database\\footnote{At {\\tt http://www.univie.ac.at/webda/}} provides coordinates RA:~21h 16m 58s, Dec:~$+51^{\\circ}\\:45\\arcmin\\:32\\arcsec$ ($\\ell=93\\fdg03$, $b=+1\\fdg80$). \\citet{macie07} presented $BV$ photometry of an extended field around the obvious cluster core and found it to be extremely compact, with $r_{{\\rm core}}=0\\farcm7$. They derived a distance $d=1.2$~kpc and an age $\\log t=8.5$ (315~Myr), which was later used by \\citet{tadross} to calibrate a fit to the 2MASS data in the region. However, the cluster looks too compact for this age and distance, while the isochrone fit of \\citet{tadross} would suggest that most of the evolved cluster members are AGB stars, an extremely unusual situation.  In this paper, we present $UBV$ photometry of the cluster area and a spectroscopic survey of likely members, showing that its population is much younger than implied by the single-color fit of \\citet{macie07}. From our new data, we show that Be~55 is a young ($\\log t\\approx7.7$) open cluster with a rich population of evolved stars. ",
        "conclusions": "\\label{discussion}  We have presented spectroscopy of a sizable sample of stars in the open cluster Berkeley~55, revealing a population of 7 low-luminosity supergiants or bright giants (one of them, not a certain member), 4 mid-B giants and a main sequence starting around B4\\,V.  \\subsection{Cluster parameters}  Only two previous studies of the cluster have been published. \\citet{macie07} derive an age of $\\sim315$~Myr. \\citet{tadross} used their result to calibrate his fitting procedure and therefore derives a similar age of 300~Myr, from a fit to the 2MASS data alone.  This age is completely incompatible with the observed population, as stars earlier than B5 are only seen in clusters younger than 100~Myr. Clusters with $\\sim100$~Myr have stars of spectral type $\\sim$B6 around the main-sequence turn-off (e.g., NGC~7790, \\citealt{matthews95}; Berkeley~58, \\citealt{turner08}).  The difference between our age determination and that of \\citet{macie07} is almost certainly due to the identification of the evolved population. The decontaminated $V/(B-V)$ in \\citet{macie07} does not show a clump of red stars, as in our Fig.~\\ref{rawcmd}. Some of the (super)giants seem to be blended in their images and do not appear in their original photometry. The others are not conspicuous when a wide field is considered, because of contamination by field interlopers with similar $(B-V)$ color. The stars forming this clump, however, cannot be overlooked when the infrared data are considered. S1 is the brightest star in $K_{{\\rm S}}$ within $10\\arcmin$ of the cluster center, and there are only 5 field stars with $K_{{\\rm S}}$ magnitudes comparable to the 7 red (super)giants. The 7 candidate members form a clearly separated clump in the $(J-K_{{\\rm S}})$/$K_{{\\rm S}}$, $H-K_{{\\rm S}}$/$K_{{\\rm S}}$ and $Q$/$K_{{\\rm S}}$ diagrams, and have similar (and unusual) spectral types. Though accurate radial velocities will be needed to confirm individual memberships, there can be no doubt that they represent the population of evolved stars in the cluster. The small systematic offset in $(B-V)$ between our photometry and that of \\citet{macie07} is unlikely to have biased their age determination. It is simply very difficult to obtain an accurate fit when photometry in only two bands is available.  At Galactic longitude $\\ell=93\\fdg0$, a distance $\\sim4$~kpc (corresponding to our $DM=13.0$) seems compatible  with tracers of the Perseus arm, though some Galactic models suggest a higher distance for this feature. Spiral tracers in this direction are scarce. NGC~7067 ($\\ell=91\\fdg2$) has distance estimates of 3.6~kpc \\citep{ys04} and 4.4~kpc \\citep{hassan73}. NGC~7128 ($\\ell=97\\fdg3$), with $DM=13.0\\pm0.2$ (4~kpc), is another likely tracer at some distance from Be~55 \\citep{balog}.  As the age of Berkeley~55 is $\\sim50$~Myr, the red stars in the cluster should have initial masses $M_{*}\\ga7.0\\,M_{\\odot}$ \\citep{marigo}, or slightly higher if they started their lives rotating quickly \\citep[cf.][]{mm00}. Such stars define the boundary between intermediate-mass and massive stars, and therefore can provide important constraints for stellar evolution models.  Milky Way clusters with similar ages and rich populations of RSGs (or bright giants) are relatively scarce. NGC~6649 ($\\sim50$~Myr; \\citealt{turner81}) with one Cepheid and four red (super)giants is one of the few examples. NGC~6664 (also $\\sim$50~Myr; \\citealt{schmidt82}) has a similar population, but is less well studied. As a consequence, the sample of Galactic RSGs studied by \\citet{levesque05} contains almost no stars in the $6\\,M_{\\odot}<M_{*}<10\\,M_{\\odot}$ range.  The presence of 7 candidate evolved members makes Berkeley~55 a very suitable target to study the physical effects operating during the He-core burning phase of intermediate-mass stars.  \\subsection{The lowest mass for a supergiant}  As mentioned in the introduction, the separation between red supergiants and red bright giants is morphological, and does not coincide with the separation between high-mass and low-mass stars. Of course, making such a distinction necessarily implies dividing a continuum into artificial boxes\\footnote{It should be noted that many authors classically count luminosity class II objects as supergiants}. In spite of that, defining an age limit after which clusters do not contain RSGs seems interesting, as new clusters are being discovered where high extinction renders only the bright red stars accessible for spectroscopy \\citep[e.g.,][]{neg10,davies11}.  Looking at the lists of yellow evolved stars in clusters from \\citet{harris76} or \\citet{sowella,sowellb}, it is difficult to set an age limit after which Ib supergiants are not seen. A few objects have been classified as K\\,Ib supergiants in clusters $\\ga 100$~Myr old, but in all cases there are alternative lower-luminosity classifications based on more recent data (for example, star 164 in NGC~129, classified K2.5\\,II-III by \\citealt{keenan}\\footnote{This star seems not to be a member, based on its radial velocity \\citep{mermi08}, though it was considered a member for a long time \\citep[see references in][]{turner92}.}; or star 110 in NGC~2516, classified M1.5\\,IIa by \\citealt{kp85}). Most classical Cepheids appear as F/G Ib supergiants.  Many clusters containing Cepheids are sparsely populated and have the Cepheid as only supergiant \\citep[e.g.,][]{turner96}. In clusters with larger populations, the ratio of F/G supergiants to later spectral types is very variable. For example, NGC~7790 contains three F-type Cepheids, but no red (super)giants, while NGC~6067 (with two Cepheid members) contains a large number of red giants and supergiants \\citep{mer87}.  If we consider a number of clusters for which accurate ages in the range of interest have been derived through careful photometric analysis or equivalent methods, we find by decreasing age, that NGC~2516 ($\\sim140$~Myr), which has no Cepheids, contains four bright red giants, with luminosity classes II or III \\citep{kp85,gl00}. IC~4725, which is $\\sim90$~Myr old, contains one Cepheid supergiant and two class II red giants \\citep{kp85,mer87}. Berkeley~82, with an age $\\sim70$~Myr contains two yellow bright stars with spectral types G2\\,II/Ib and K2\\,II \\citep{forbes86}, while Harvard~20 contains a G5\\,Ib supergiant at $\\sim60$~Myr \\citep{turner80}. On the other hand, NGC~6664, with an age $\\sim50$~Myr, contains several objects classified as bright giants by \\citet{harris76}. Finally, NGC~6067 contains red stars with luminosity classes II and Ib, leading \\citet{mer87} to doubt its estimated age $\\ga100$~Myr, which is also in contradiction with the observed period of its fundamental-mode Cepheid member V340~Nor.  In view of this, it seems that red stars close to the clump position are classified as K supergiants in clusters up to 50--60~Myr, while some stars are classified as such in somewhat older clusters. No cluster older than $\\sim100$~Myr seems to contain a bona-fide K supergiant. Stars with F and G spectral types, such as Cepheids, are classified as supergiants at older ages, e.g., SZ Tau (F5-9\\,Ib) in NGC~1647 \\citep[$\\sim150$~Myr;][]{turner92}.  The classifications derived for stars in Be~55 seem to confirm these trends, with stars on both sides of the divide at an age $\\approx50$~Myr."
    },
    "1112/1112.5007_arXiv.txt": {
        "abstract": "Hu\\,1-2 is a planetary nebula that contains an isolated knot located northwestern of the main nebula, which could be related to a collimated outflow. We present a subsarcsecond H$\\alpha$+[N\\,{\\sc ii}] image and a high-resolution, long-slit spectrum of Hu\\,1-2 that allow us to identify the southeastern counterpart of the northwestern knot and to establish their high velocity ($>$ 340 km\\,s$^{-1}$), collimated bipolar outflow nature. The detection of the northwestern knot in POSS red plates allows us to carry out a proper motion analysis by combining three POSS red plates and two narrow-band H$\\alpha$+[N\\,{\\sc ii}] CCD images, with a time baseline of $\\simeq$ 57 yr. A proper motion of 20 $\\pm$ 6\\,mas\\,yr$^{-1}$ along position angle 312$^{\\circ}$ $\\pm$ 15$^{\\circ}$, and a dynamical age of 1375$^{+590}$\\llap{$_{-320}$}\\,yr are obtained for the bipolar outflow. The measured proper motion and the spatio-kinematical properties of the bipolar outflow yield a lower limit of 2.7\\,kpc for the distance to Hu\\,1-2. ",
        "introduction": "Planetary nebulae (PNe) represent the final evolutionary stages of low-- and intermediate--mass stars. They are formed when the envelope ejected during the previous Asymptotic Giant Branch phase is ionized by the hot central star. In the last years the presence of collimated bipolar outflows in PNe has been widely recognized. These outflows probably play a crucial role in the shaping and dynamical evolution of PNe \\citep{b25}. Therefore, the knowledge of their dynamical properties is critical to understand how PNe form. The kinematics of the outflows has been studied by means of high-resolution spectroscopy that permits to obtain the radial component of the velocity (e.g., Miranda et al. 1999, 2001; Guerrero et al. 2000). The measurement of the tangential component, i.e., the proper motion of collimated outflows, is restricted to a very few cases and it has been obtained mostly from {\\it HST} observations (NGC\\,7009: Fern\\'andez et al. 2004; He\\,2-90: Sahai et al. 2002; Hen\\,3-1475: Borkowski \\& Harrington 2001; Riera et al. 2003) or {\\it VLA} radio continuum data (NGC\\,7009: Rodr\\'{\\i}guez \\& G\\'omez 2007) with the noticeable exceptions of the works by Liller (1965) and Meaburn (1997) who used photographic material to measure proper motions of the collimated outflows in NGC\\,7009 and KjPn\\,8, respectively. The measurement of proper motions is important because it allows us to obtain the distances to the PN, if the expansion velocity of the outflows is known (e.g., Meaburn 1997), or the expansion velocity vector, if the distance is known (e.g., Fern\\'andez et al.2007). Therefore, further measurements of proper motions of collimated outflows in PNe are highly desirable.  Collimated outflows may exist in the PN Hu\\,1-2 (PN\\,G086.5$-$08.8) that has been classified as elliptical with ansae \\citep{b4}. The kinematics of the inner nebular regions ($\\simeq$ 10$''$) of Hu\\,1-2 was analyzed by Sabbadin et al. (1983, 1987). These authors identified a toroid and bipolar lobes but a reconstruction of the spatio-kinematical structure was not possible because of the peculiar velocity field. In addition, one single knot can be hinted in the images by Manchado et al. (1996) northwestwards and well outside the main nebula, while a possible southeastern counterpart is superimposed by a field star. Remarkably, the northwestern knot can be identified in POSS red plates of the Digitized Sky Surveys. In this respect, Hu\\,1-2 constitutes a rare case (together with NGC\\,7009 and KjPn\\,8) because knots/collimated outflows in PNe usually are weak and/or located close to the much brighter main nebula, which do not favor their detection in POSS plates. Moreover, the H$\\alpha$ and [N\\,{\\sc ii}] emission lines are by far the dominant emissions from knots/collimated outflows in PNe in the optical (e.g., Balick 1993, 1994), as it is the case of the northwestern knot of Hu\\,1-2. In consequence, the emission from the northwestern knot of Hu\\,1-2 detected in the POSS red plates can be attributed to these two emission lines. Therefore, Hu\\,1-2 offers an excellent opportunity to attempt a proper motion analysis of knots/collimated outflows in PNe by combining POSS red plates with modern H$\\alpha$+[N\\,{\\sc ii}] imagery.  In this paper we present a new H$\\alpha$+[N\\,{\\sc ii}] image of Hu\\,1-2 obtained under subarsecond conditions, and a high resolution, long-slit spectrum that allow us to identify the southeastern counterpart of the northwestern knot and to establish that these two knots constitute a high velocity, collimated bipolar outflow. In addition, we carry out an analysis of five images obtained at different epochs, including three images from the POSS, to measure the proper motion of the northwestern knot and to constrain the distance to Hu\\,1-2. ",
        "conclusions": "We have presented a subarsecond H$\\alpha$+[N\\,{\\sc ii}] image and a high resolution long-slit spectrum of Hu\\,1-2 covering its bipolar knots. In addition, we have analyzed five images of Hu\\,1-2 with a time baseline of $\\simeq$ 57 yr, including three POSS red plates in which one of the bipolar knots can be identified, to measure its proper motion and to constrain the distance to the nebula. The main conclusions of this paper are:  \\begin{itemize}  \\item The subarcsecond image allows us to identify the two bipolar knots of the pair, while in previous images only one of them is clearly detected. The knots present bow-shock-like morphology, are located at 27$\\farcs$5 from the central star and oriented at PA 320$^{\\circ}$ that coincides with the orientation of the main nebular axis.  \\item The radial velocity of the knots is $\\pm$ 60 km\\,s$^{-1}$. If the knots moved along the main nebular axis (tilted $<$ 10$^{\\circ}$ with respect to the plane of the sky), their expansion velocity would be $>$ 340\\,km\\,s$^{-1}$. Therefore, morphology and kinematics demonstrate that these knots constitute a true high velocity, collimated bipolar outflow and, most probably, represent bow-shocks associated to high velocity bullets.  \\item A proper motion of 20 $\\pm$ 6 mas\\,yr$^{-1}$ along PA 312$^{\\circ}$ $\\pm$ 15$^{\\circ}$ is obtained for the bipolar knots. The corresponding dynamical age is 1375$^{+590}$\\llap{$_{-320}$}\\,yr.  \\item A lower limit of 2.7\\,kpc for the distance to Hu\\,1-2 is required to make compatible the measured proper motion with the spatio-kinematical properties of the bipolar outflow. This lower limit rules out a large fraction of the statistical distances and the extinction distance previously determined for Hu\\,1-2.  \\end{itemize}"
    },
    "1112/1112.0622_arXiv.txt": {
        "abstract": "We investigate the feasibility of implementing a system that will coordinate ground-based optical telescopes to cover the Fermi GBM Error Circle (EC). The aim of the system is to localize GBM detected GRBs and facilitate multi-wavelength follow-up from space and ground. This system will optimize the observing locations in the GBM EC based on individual telescope location, Field of View (FoV) and sensitivity. The proposed system will coordinate GBM EC scanning by professional as well as amateur astronomers around the world. The results of a Monte Carlo simulation to investigate the feasibility of the project are presented. ",
        "introduction": "Gamma-ray bursts (GRB) are bursts of gamma-rays that arguably signal the birth of a black hole somewhere in the universe. Based on the duration and spectrum, two classes of bursts have been observed~\\citep{Kouveliotou1993}: those that last less than two seconds and have on the average hard spectra (short GRBs), and those that last longer than two seconds and are spectrally softer (long GRBs). The exact nature of the GRB progenitors is unknown, although it is possible that long GRBs come from the collapse of massive, rapidly rotating stars~\\citep{Woosley2006, WoosleyBloom2006} and short GRBs result from the merger of compact objects~\\citep{Eichler1989, Narayan1992}. Regardless of the progenitor system, accretion onto the resulting compact object is thought to create a highly relativistic jet. The prompt gamma-ray emission from the GRBs may arises from the internal shocks due to collisions of faster shells with slower ones ejected earlier by the central engine. The subsequent softer multi-wavelength emission, referred to as the afterglow, may be due to the collision of the fireball with the extra-stellar material~\\citep{Rees1994,Piran1999}.  Our understanding of GRBs progressed very rapidly after the detection of multi-wavelength afterglows. Well localized, favorably positioned GRBs get fairly good multi-wavelength afterglow coverage. Currently, the leading GRB afterglow detection mission is $Swift$ which detects 90--100 GRBs annually. Most of the $Swift$ GRBs get observed by various instruments around the world because of its rapid arc-minute localization capability. Compared to $Swift$, $Fermi$ Gamma-ray Burst Monitor (GBM) detects about 250 burst per year but with poor localization. The Error Circle (EC) of $Fermi$ GBM detected bursts is too large for a single telescope to observe effectively. The typical statistical uncertainty of the GBM burst location is about 3.3 degrees. However, when combined with the systematic uncertainty of 3.8 degrees~\\citep{Briggs2009}, the total burst location uncertainty is $\\sim$ 5.0 degrees (i.e., 5.0 degree error radius). Naturally, a brighter burst will have a smaller GBM EC than a weaker burst.  Even though the localization is poor, GBM detected bursts have very good timing and spectral information including crucial $E_{\\rm peak}$ measurements ($E_{\\rm peak}$ is the peak energy of the GRB $\\nu F_{\\nu}$ spectrum). If there is a method to localize GBM detected GRBs to a few arc-seconds uncertainty, then large telescopes can do deeper follow--up observations to determine the redshift of the burst and also potentially identify any emerging supernova. In addition, $Swift$ can also slew quickly to the GBM burst in order to observe the X-ray afterglow and obtain its light curve in X-ray wavelengths.  Based on $Swift$ observations about $\\sim$ 60\\% of GRBs have optical counterparts~\\citep{Gehrels2009}. These optical counterparts are detected by various observatories with R magnitudes ranging from 14 to 22 within few hours after the burst~\\citep{Fiore2007}. Thus, it is reasonable to assume about 60\\% of the GBM detected GRBs also have optical counterparts with similar brightness distribution. If we were able to cover the entire GBM EC within about 24 hours after the burst it is conceivable that we would be able to find optical afterglows of $\\sim$150 GRBs per year, which is more than the total number of burst $Swift$ detects per year.  Due to the small energy range (15-150 keV) of the $Swift$ Burst Alert Telescope (BAT), $Swift$ measurements alone cannot constrain the $E_{\\rm peak}$ of all BAT detected bursts~\\citep{Sakamoto2009}. In contrast, due to the wide energy range (8 keV - 40 MeV) of GBM, all GRBs detected by GBM have fairly good $E_{\\rm peak}$ measurements. Hence, addition of possibly another $\\sim$100 bursts per year with good $E_{\\rm peak}$ and redshift measurements may allow us to explore the validity of various GRB luminosity relations and to conduct detailed GRB Hubble diagram studies.  We have investigated the feasibility of using a system to do coordinated monitoring of the BAT field-of-view (FoV) for prompt optical emission from GRBs~\\citep{Ukwatta2011}. The study showed that with the current instrumentation, performing such a coordinated monitoring is not practical mainly due to the BAT's very large FoV. However, a similar coordinated observing campaign can be used to find the optical afterglow of GBM detected bursts. The GBM EC is much smaller than the BAT FoV and observers do not need to continuously monitor the field to detect the optical afterglow. This enables a given observatory to perform multiple observations inside the GBM EC and thereby increase the chance of a afterglow detection. ",
        "conclusions": "We investigate the feasibility of implementing a system that will coordinate ground based telescopes (both amateur and professional) to scan the GBM EC in order to localize GBM bursts. Unlike the GCN system, proposed system will send individual customized messages to telescopes to observe certain patches in the GBM EC. The scientific objective of the system is by localizing GBM detected burst, we will be able to increase the number of GBM bursts with mutil-wavelength followups potentially with redshifts measurements. These measurements are scientifically very important because there are hints that GBM bursts may represent significantly different burst population. Based on our simulation, we can detect about 25 GRB afterglows per year using just 10 telescopes with $1.0^{0} \\times 1.0^{0}$ field-of-view. With more telescopes participating in the program, we should be able to detect many more afterglows and study a potentially interesting burst population that is currently inaccessible to the GRB community."
    },
    "1112/1112.0308_arXiv.txt": {
        "abstract": "We demonstrate a new technique for detecting radio transients based on interferometric closure quantities. The technique uses the bispectrum, the product of visibilities around a closed-loop of baselines of an interferometer. The bispectrum is calibration independent, resistant to interference, and computationally efficient, so it can be built into correlators for real-time transient detection. Our technique could find celestial transients anywhere in the field of view and localize them to arcsecond precision. At the Karl G. Jansky Very Large Array (VLA), such a system would have a high survey speed and a 5$\\sigma$\\ sensitivity of 38 mJy on 10 ms timescales with 1 GHz of bandwidth. The ability to localize dispersed millisecond pulses to arcsecond precision in large volumes of interferometer data has several unique science applications. Localizing individual pulses from Galactic pulsars will help find X-ray counterparts that define their physical properties, while finding host galaxies of extragalactic transients will measure the electron density of the intergalactic medium with a single dispersed pulse. Exoplanets and active stars have distinct millisecond variability that can be used to identify them and probe their magnetospheres. We use millisecond time scale visibilities from the Allen Telescope Array (ATA) and VLA to show that the bispectrum can detect dispersed pulses and reject local interference. The computational and data efficiency of the bispectrum will help find transients on a range of time scales with next-generation radio interferometers. ",
        "introduction": "Since the discovery of the first pulsar, radio observations have steadily revealed a wide range of fast ($\\lesssim1$s) transient phenomena. Pulsars \\citep{1968Natur.217..709H,1982Natur.300..615B,2006Natur.442..892C}, planets and active stars \\citep{2007ApJ...663L..25H,2007AA...475..359G}, and peculiar new transients \\citep{2006Natur.439..817M,2007Sci...318..777L} all have distinct radio phenomenology related to the physics of the host and the medium through which the signal propagates.  While large, single-dish telescopes have pioneered the study of fast radio transients, the nature of their design has several limits. Single-dish telescopes can only localize an individual pulse to somewhere in its field of view \\citep[typically arcminutes;][]{2001MNRAS.328...17M}. Timing of periodic emission can localize very precisely, but individual pulses can't be localized well enough to easily identify optical or X-ray counterparts. Since interferometers simultaneously have a large collecting area and a large field of view, they survey more efficiently than single-dish telescopes. Survey speed scales roughly the square of the product of the number of dishes and their diameter; this predicts that the VLA surveys almost 40 times faster than the Green Bank Telescope, assuming all other factors are equal. Despite this potential, next-generation radio interferometers (VLA, PAPER, LOFAR, ASKAP, and MeerKAT) have fast transients observing modes that are essentially identical to the single-dish concept: point one or more beams on the sky and search for pulses \\citep{2011ApJ...734...20M,2011arXiv1109.2659S}.  In principle, interferometers have much more information available to them because they detect the electromagnetic wave over a distributed array of receivers. Instead of forming a beam, interferometers can measure visibilities by correlating signals across the array. Visibilities contain spatial information that can reject interference and localize a source to \\emph{arcsecond} precision over a large field of view. Traditionally, visibilities haven't been used in this way because it was difficult to run correlators this fast and the data volume was too large to handle or efficiently search for pulses \\citep{2011ApJ...742...12L}. As a result, a standard radio interferometer averages data to second time scales, losing nearly all information about millisecond variability in the sky. Even on these longer time scales, imaging algorithms are complex and subject to subtle statistical biases \\citep{1998AJ....115.1693C}. The challenge of data rates and volumes is expected to grow with next-generation interferometers, which will have more antennas, wider bandwidths, and larger fields of view \\citep{2011ApJ...739L...1P}.  Here, we present a technique that reduces the challenge of detecting transients in visibilities to the much simpler single-dish problem, while maintaining the utility of visibilities. The technique is based on the ``bispectrum'', a quantity formed from visibilities that is sensitive to a point source anywhere in the field of view \\citep{1987AA...180..269C,1989AJ.....98.1112K,1995AJ....109.1391R}. Differencing visibilities makes it possible to interpret the bispectrum as a statistic for transients anywhere in the field of view on time scales from milliseconds to seconds. We describe an algorithm for real-time, arcsecond localization of transients using raw visibilities from a radio interferometer.  The theory of closure quantities and their application to transients is presented in \\S \\ref{theory}. In \\S \\ref{demo} we test the theoretical predictions with millisecond time scale visibilities from the VLA and the ``Pocket Correlator'' (PoCo) instrument at the ATA \\citep{2011ApJ...742...12L}. As shown in \\S \\ref{demand}, this technique has a moderate computational demand and can reduce the flow of data to a manageable size. We conclude with thoughts about new science accessible with this technique and how it can address the growing challenge of big data in the study of transients with radio interferometers. ",
        "conclusions": "We demonstrate how to use the bispectrum to detect and localize transients with radio interferometers. The technique produces a single statistic proportional to the brightness of celestial transients, thus reducing the problem of interferometric pulse detection to that of single-dish pulse detection. For a small cost in sensitivity, this technique makes an interferometer into a massive multibeam receiver with a proportionately large survey speed. Using coherent beamforming to survey a similar area with next-generation radio interferometers would require $10^2$\\ to $10^5$\\ times more computational power.  The bispectrum technique is efficient enough to build into a correlator for real-time, millisecond transient detection. This system could save visibility data associated with a candidate and apply standard calibration and imaging techniques to localize it to arcsecond precision. The growing capacity of correlators for interferometers will make it possible to run such a system in real time to commensally search all data for transients. The technique can also be used to find transients offline and on longer time scales. In that case, the bispectrum is useful because it is computationally simple, resistant to interference, and requires no calibration.  A real-time, commensal system would probe an unprecedented volume of data, giving it sensitivity to rare transients. Each pulse could be localized in real time to arcsecond precision, opening a range of new science. Pulses from rotating radio transients \\citep{2006Natur.439..817M} could be associated with X-ray counterparts to help address their relation to the normal pulsar population. The dispersion of extragalactic transients would measure the electron density of the intergalactic medium for the first time \\citep{2003ApJ...596..982M,2007Sci...318..777L}; localization will help find optical counterparts and redshifts to host galaxies for each pulse. A fast transients survey of the Galactic center could detect individual pulses from pulsars that would probe the dispersion toward the region and potentially find a pulsar in orbit around a supermassive black hole \\citep{2004ApJ...615..253P,2010ApJ...715..939M,2011arXiv1111.4216W}. Magnetospheres of exoplanets and active stars can emit bright, complex bursts that could uniquely identify them \\citep{2004ApJ...612..511L,2006ApJ...653..690H}.  The strength of the bispectrum technique is its real-time ability to find transients throughout the field of view and reduce the flow of data to select times. This shows how algorithms will be key to efficiently extracting the best science in the era of big data. The ability to find pulses with interferometers may have application to communications problems, such as developing smart antennas for the efficient use of spectrum. If so, this technique would follow in the footsteps of previous transients searches that contributed to the development of Wi-Fi \\citep{1978Natur.276..590O}."
    },
    "1112/1112.3598_arXiv.txt": {
        "abstract": "The rapid, irreversible change of the photospheric magnetic field has been recognized as an important element of the solar flare process. This Letter reports such a rapid change of magnetic fields during the 2011 February 13 M6.6 flare in NOAA AR 11158 that we found from the vector magnetograms of the Helioseismic and Magnetic Imager with 12-min cadence. High-resolution magnetograms of Hinode that are available at \\sm$-$5.5, $-$1.5, 1.5, and 4 hrs relative to the flare maximum are used to reconstruct three-dimensional coronal magnetic field under the nonlinear force-free field (NLFFF) assumption. UV and hard X-ray images are also used to illuminate the magnetic field evolution and energy release. The rapid change is mainly detected by HMI in a compact region lying in the center of the magnetic sigmoid, where the mean horizontal field strength exhibited a significant increase by 28\\%. The region lies between the initial strong UV and hard X-ray sources in the chromosphere, which are cospatial with the central feet of the sigmoid according to the NLFFF model. The NLFFF model further shows that strong coronal currents are concentrated immediately above the region, and that more intriguingly, the coronal current system underwent an apparent downward collapse after the sigmoid eruption. These results are discussed in favor of both the tether-cutting reconnection producing the flare and the ensuing implosion of the coronal field resulting from the energy release. ",
        "introduction": "It has been known that the long-term evolution of photospheric magnetic field (PMF) driven by new flux emergence and surface flows plays important roles in building up free energy in the corona, and that this free magnetic energy powers flares and coronal mass ejections \\citep[CMEs;][]{priest02}. On the other hand, a short-term variation of the PMF associated with flares has not been considered because PMFs are strongly line-tied to the dense high-$\\beta$ photosphere and thus are thought unlikely to be altered by any flare-related  disturbances created in the tenuous low-$\\beta$ corona. Only recently, a back reaction of the coronal magnetic field (CMF) on the PMF during the reconfiguration of the CMF has been seriously considered from the theoretical point of view. The idea is that the CMF should contract inward, as the magnetic energy of the CMF decreases after flares and/or CMEs \\citep[called ``implosion'';][]{hudson00}. This may make the PMF be oriented more horizontally, resulting in a Lorentz force acting downwardly on the solar surface \\citep{fisher10,hudson08}. In fact, a few theoretical flare models predict the appearance of more horizontal PMFs after flares/CMEs in the form of newly formed, low-lying field lines close to the surface \\citep[e.g.,][]{melrose97,moore01}.  The phenomenon of coronal implosion has been reported in several recent observational studies \\citep{liur+implosion09,liur+wang09,liur+wang10}. The related phenomenon of the rapid change of the PMF to incline more horizontally after flares had also been reported almost two decades ago \\citep{wang92,wang94}. Namely, the transverse field near the magnetic polarity inversion line (PIL) at the core flaring region is enhanced rapidly and irreversibly, which is often accompanied by an increase of magnetic shear. Such a rapid change of vector PMF change closely associated with flares/CMEs has been consistently found later on using ground-based \\citep{schmieder94,wang02b,wang04,wang+liu05,liu05,wang07a,wang10} and \\hinode\\ observations \\citep{jing08,li09}. Indirect evidence includes the unbalanced flux evolution of the line-of-sight magnetic field \\citep{wang10}, the variation of the sunspot white-light structure in flaring regions \\citep{wang04a,deng05,liu05,chen07b}, and the change in the pattern of the penumbral Evershed flow \\citep{deng11}. Additional support also comes from the three-dimensional (3D) MHD simulations of an emerging twisted flux tube \\citep{fan10}, the eruption of which induces similar impact on the low-altitude magnetic fields \\citep{li10}.  \\begin{figure*}[t] \\epsscale{1} \\plotone{f1.eps} \\caption{$B_v$ maps are superimposed with arrows (for clarity, those in negative and positive fields are coded in different colors) representing horizontal plasma flows (a) and magnetic field vectors (b--d). The NLFFF lines in (e) stem from the four flare kernels FP1--FP4 in 1700~\\AA\\ (f) at the flare onset. The white contours in (e) outline the double J-shaped flare ribbons in Ca~{\\sc ii}~H. Contours (with levels of 38\\%, 48\\%, 58\\%, 68\\%, 78\\%, 88\\% of the maximum flux) in (f) denote \\hsi\\ CLEAN images reconstructed with detectors 2--8 (except that 17:32:08~UT image is made without the detector 2 due to the relatively low counts at this time). The white box overplotted on all panels denotes the region R under study with enhanced $B_h$ after the flare. \\label{f1}} \\end{figure*}  While the evolution of vector PMFs closely associated with flares were mostly studied using ground-based observations subject to seeing variations, \\citet{wang+shuo11} and \\citet{sun11} used data from the \\hmi~\\cite[HMI;][]{schou11} on board the recently launched \\Sdo\\ (\\sdo) to find a rapid enhancement of transverse field associated with the 2011 February 15 X2.2 flare. Since solar activity is now increasing, this is an opportune moment to advance the study of the flare-related evolution of the PMF using the state-of-the-art, seeing-free photospheric vector magnetograms acquired by space missions.  \\begin{figure*}[t] \\epsscale{.8} \\plotone{f2.eps} \\caption{Temporal evolution of $\\langle B_h \\rangle$ of the region R in Figure~\\ref{f1}, in comparison with the flare light curves of soft and hard X-rays. The mean error of $B_h$ of the region R at each time instance measured by HMI is plotted as error bars. The error estimates for \\hinode\\ data are currently not usable$^2$. \\label{f2}} \\end{figure*}   In this Letter, we investigate the PMF evolution associated with an M6.6 flare on 2011 February 13 that was well covered by both HMI and \\hinode\\ (see Section~\\ref{data}), in which the transverse field enhancement at the flaring PIL is unequivocally detected. We will discuss the rapid change of the 3D CMF associated with that of the PMF using nonlinear force-free field (NLFFF) extrapolations from the active region. ",
        "conclusions": "We have used the vector magnetograms from HMI and \\hinode\\ to analyze the change of the PMF associated with the 2011 February 13 M6.6 flare, and interpreted the results with the aid of NLFFF modeling based on the \\hinode\\ data under the context of 3D magnetic reconnection. Major results are summarized as follows.  \\begin{enumerate} \\item A compact region R at the central flaring PIL shows a rapid and permanent enhancement of $\\langle B_h \\rangle$ by \\sm320~G (\\sm28\\% of the preflare magnitude as observed by HMI) closely associated with the flare. A similar trend of $\\langle B_h \\rangle$ evolution is also detected by \\hinode. Overall, the PMF at the region R becomes more inclined to the surface after the flare.  \\item The magnetic fields lying immediately above R, which correspond to the central part of the sigmoidal region, are highly sheared and have strong horizontal electric currents. The most intense flare kernels are initially located at the two ends of this channel-like current system, implying that the start of electron acceleration could be closely associated with the strong horizontal currents \\citep{leka93}. The configuration of the NLFFF lines is favorable for a subsequent tether-cutting reconnection \\citep{moore01}, when disturbed by converging surface flows. This ends up with the short and low-lying loops, which explains the enhanced $B_h$ and $\\tilde{S}$ at the region R.  \\item Across the flare energy-release time, the nonpotentiality as represented by $\\tilde{S}$ increases in the height range from the surface up to \\sm1~Mm and then decreases at higher altitudes, which is closely correlated with the evolution of $J_h$. This is in accordance with the reconnection of current-carrying loops modeled by \\citet{melrose97}. Remarkably, the vertical profiles of $J_h$ as well as $\\varphi$ before and after the flare clearly demonstrate a downward collapse of the CMF by \\sm1.3~Mm toward the surface, which is probably related to coronal implosion \\citep{hudson00}. \\end{enumerate}  In summary, the change of the PMF and the corresponding change of the CMF observed in this study are not caused by the long-term development of the active region but occurred as rapid as the flare. The magnetic field change may involve two closely related physical processes: the tether-cut reconnection producing the flare and the ensuing collapse of the CMF resulting from the energy release."
    },
    "1112/1112.4976_arXiv.txt": {
        "abstract": "{Very-high-energy (VHE, E $>$ 100 GeV) \\g\\ radiation has already been detected from several supernova remnants (SNRs). These objects, which are well-studied in radio, optical and X-ray wavelengths, constitute one of the most intriguing source classes in VHE astronomy. H.E.S.S., an array of four imaging atmospheric Cherenkov telescopes in Namibia, has recorded an extensive dataset of VHE \\gr\\ observations covering the central region of the Milky Way, both from pointed observations as well as from the Galactic Plane Survey conducted in the inner region of the Galaxy. From radio observations, several hundred SNRs are known in the Milky Way, but until now only few of them have been identified as VHE \\gr\\ emitters. Using the H.E.S.S. dataset and a large ensemble of radio SNRs localized in the inner region of the Galaxy, the standard framework that links the origin of cosmic rays to the \\gr\\ visibility of SNRs can now be tested. Here we present the ensemble of investigated SNRs and discuss constraints on the parameter space used within a theoretical model of hadronic VHE \\gr\\ production.} ",
        "introduction": "SNRs are widely believed to be the main sources of Galactic Cosmic Rays (CRs) with energies of up to \\unit[$10^{15}$]{eV}. The amount of energy released during a supernova (SN) explosion, the rate of Galactic SNe and the energy density of CRs together with diffusive shock acceleration as a possible acceleration mechanism make SNRs the prime candidates for the acceleration of CRs. In radio, optical and X-ray wavelengths deep observations on SNRs have been performed and several SNRs have been discovered in the relatively young field of VHE \\gr\\ astronomy . With H.E.S.S. the shell-type morphology of individual SNRs such as RX~J1713.7-3946~\\cite{HESS:1713I}, SN~1006~\\cite{HESS:SN1006}, Vela Junior~\\cite{HESS:VelaJunior} and HESS~J1731-347~\\cite{HESS:1731} has been resolved. These discoveries support the hypothesis of particle acceleration in the shell of SNRs. Using the large dataset of H.E.S.S. accumulated during the Galactic plane survey (GPS)~\\cite{ICRC2011:Survey} in recent years we study SNRs in the inner Galactic region, ranging from 275$^\\circ$ to 60$^\\circ$ in Galactic longitude and from -3$^\\circ$ to 3$^\\circ$ in Galactic latitude.  Here, the VHE \\gr\\ signal of known radio SNRs listed in Green's radio SNR catalog~\\cite{GREEN:RadioSNRs} is investigated. For each SNR within the GPS region, an upper limit on the VHE \\gr\\ flux has been calculated and can be compared to the model predictions. In this work we test the commonly used model of the hadronic \\gr\\ emission expected from SNRs presented by Drury~\\etal~\\cite{Drury:SNRVisibility}. They calculate the integral \\gr\\ flux from hadronic interactions based on intrinsic SN parameters. VHE \\g\\ rays result from the decay of neutral pions generated in proton--proton interactions in the expanding shell of the SNR. For the first time, the high sensitivity of H.E.S.S. and the large dataset of the GPS allow us to derive upper limits that are low enough to test model parameters. Using the upper limit on the VHE \\gr\\ flux represents a conservative test since an additional leptonic component might be present in the flux. ",
        "conclusions": "\\label{sec:paramUL} Supernova remnants are the most plausible sites of cosmic-ray acceleration. As a consequence, they are predicted to be emitters of VHE $\\gamma$ radiation by $\\pi^0$ decay during the early stage of their evolution. The good sensitivity of H.E.S.S.~and the large dataset accumulated in its Galactic Plane Survey now allow us to compare quantitative estimates for this emission to observations, and to study the SNR population in the Inner Galaxy as a whole. From Green's catalog of radio SNRs, candidates were selected whose estimated age makes them predicted sources of VHE \\g\\ radiation and the existance of a distance measurement allows for the expected flux to be calculated.  Fig.~\\ref{log_p_dist} shows the distribution of the upper limits on the product $p\\equiv \\theta{}E_\\mathrm{SN}n/\\unit[10^{51}]{erg}\\,\\mathrm{cm}^{-3}$ as calculated from the H.E.S.S.~flux upper limit in the framework of the model of Drury \\etal. The vertical line marks a value of $p=0.1$ which would be expected for a supernova explosion energy of $E_\\mathrm{SN}=10^{51}\\,\\mathrm{erg}$, an ambient number density of $n=1\\,\\mathrm{cm}^{-3}$, and an acceleration efficiency of $\\theta=0.1$. The H.E.S.S.~upper limits fall within the same order of magnitude of this value.  SNRs with an upper limit on the $p$ parameter above 0.1 can be divided into two groups: those with low exposure and those with detection in VHE \\g\\ rays. In the first case the H.E.S.S. coverage of the region is not deep enough to put a constraining upper limit on the SNR, whereas in the latter case VHE \\gr\\ emission has been detected, which can originate from the shell itself (e.g. G347.3-0.5, HESS~J1713-397), from a possible bright pulsar wind nebulae (PWN, e.g. G0.9+0.1, HESS~J1747-281) or can result from an overlapping bright source (e.g. G27.4+0.0, HESS~J1841-055).  There are a few SNRs that appear to be underluminous, and they are interesting because, in principle, they have the potential of falsifying the model of hadronic \\gr\\ production. But the uncertainties in the input parameters to the model are large. In particular, the ambient density is usually uncertain by an order of magnitude or more and the acceleration efficiency is also not well known.  In summary, the present study constitutes a first falsification test for a commonly used model of the \\gr\\ visibility of supernova remnants. Further observations and future instruments like the Cherenkov Telescope Array (CTA) with its vastly increased sensitivity and good survey capabilities will contribute to a better understanding of acceleration mechanisms of cosmic rays.   {\\small \\subsubsection*{Acknowledgements} The support of the Namibian authorities and of the University of Namibia in facilitating the construction and operation of H.E.S.S. is gratefully acknowledged, as is the support by the German Ministry for Education and Research (BMBF), the Max Planck Society, the French Ministry for Research, the CNRS-IN2P3 and the Astroparticle Interdisciplinary Programme of the CNRS, the U.K. Science and Technology Facilities Council (STFC), the IPNP of the Charles University, the Polish Ministry of Science and Higher Education, the South African Department of Science and Technology and National Research Foundation, and by the University of Namibia. We appreciate the excellent work of the technical support staff in Berlin, Durham, Hamburg, Heidelberg, Palaiseau, Paris, Saclay, and in Namibia in the construction and operation of the equipment.}"
    },
    "1112/1112.0002_arXiv.txt": {
        "abstract": "We present the results of a multiplicity survey of $212$ T~Tauri stars in the Chamaeleon~I and Taurus-Auriga star-forming regions, based on high-resolution spectra from the Magellan Clay $6.5$\\,m telescope. From these data, we achieved a typical radial velocity precision of $\\sim\\!80$\\,m\\,s$^{-1}$ with slower rotators yielding better precision, in general. For $174$ of these stars, we obtained multi-epoch data with sufficient time baselines to identify binaries based on radial velocity variations. We identified eight close binaries and four close triples, of which three and two, respectively, are new discoveries. The spectroscopic multiplicity fractions we find for Cha~I ($7\\%$) and Tau-Aur ($6\\%$) are similar to each other, and to the results of field star surveys in the same mass and period regime. However, unlike the results from imaging surveys, the frequency of systems with close companions in our sample is not seen to depend on primary mass. Additionally, we do not find a strong correlation between accretion and close multiplicity. This implies that close companions are not likely the main source of the accretion shut down observed in weak-lined T~Tauri stars. Our results also suggest that sufficient radial velocity precision can be achieved for at least a subset of slowly rotating young stars to search for hot Jupiter planets. ",
        "introduction": "\\label{sec:Introduction}  Most stars, both in the solar neighborhood and in young clusters are members of binary or multiple systems. Yet, the formation and early evolution of binary and multiple stars is poorly constrained observationally, and not well understood theoretically. For instance, the fraction of wide binaries in dense star-forming regions such as the Orion Nebula Cluster and IC~348 are comparable to that of field stars \\citep{1999AA...343..831D,2006AA...458..461K} whereas the frequency of binaries is much higher among young stars in dispersed T~associations like Taurus-Auriga \\citep[e.g.,][]{1995ApJ...443..625S,1997ApJ...481..378G,2003AJ....126.2009B}; for reviews see \\citet{2000prpl.conf..703M} and \\citet{2007prpl.conf..379D}. Furthermore, high-order multiples are more common in nearby star-forming regions than for solar-type main-sequence stars in the solar neighborhood \\citep{2006AA...459..909C}.  Some simulations suggest that stars usually form in triples and higher-order multiple systems only to be dispersed later, with the fraction of stars in multiple systems decreasing from $80\\%$ down to $40\\%$ by about $10$\\,Myrs \\citep{2004MNRAS.351..617D}. Predictions are that the binary fraction is higher among higher mass stars, and that brown dwarfs are never close companions to stars \\citep{2004AA...414..633G}.  While past multiplicity surveys, using speckle imaging and adaptive optics on $4$-m class telescopes, have drawn attention to the ubiquity of binaries in star-forming regions \\citep[e.g.,][]{1993AJ....106.2005G,1993AA...278..129L}, their limited contrast and angular resolution have left many key questions unanswered or only partially answered. For instance, the frequency of higher-order multiples is uncertain, and so is the frequency of very low-mass stellar and sub-stellar companions \\citep{2007ApJ...671.2074A}. Multiple systems are probably common.  With adaptive optics on $8$-m class telescopes, it has become straightforward to detect all stellar and even all brown dwarf companions down to the deuterium-burning limit with separations of tens of AU for nearby young stars \\citep[e.g.,][]{2008ApJ...683..844L}. Like \\citet{2008MNRAS.385.2210M}, our radial velocity study is complementary and covers the close separations.  We present the results of a high-resolution spectroscopic survey of $212$ stars spanning $\\sim\\!0.2$--$3\\,M_{\\sun}$ in the nearby $\\sim\\!2$\\,Myr old star-forming regions Chamaeleon~I (hereafter Cha~I) and Taurus-Auriga (hereafter Tau-Aur). Cha~I and Tau-Aur are at distances of $\\sim\\!160$\\,pc and $\\sim\\!140$\\,pc, respectively \\citep{1997AA...327.1194W, 1994AJ....108.1872K}. Previously, we presented a study of rotation, disk, and accretion signatures for a subsample of $144$ stars showing no evidence of spectroscopic companions and broadening functions showing only a single source \\citep{2009ApJ...695.1648N}; we use the derived projected rotational velocities ($v\\,\\sin\\,i$) and accretion signatures (H$\\alpha\\,10\\%$\\,width) for this work. Furthermore, we also analyzed the variability in accretion-related emission lines for a subsample of $40$ classical T~Tauri stars \\citep{2009ApJ...694L.153N}. \\citet{2008ApJ...683..844L} presented a census of wide binaries in Cha~I that encompasses our sample while \\citet{2007ApJ...670.1337D} investigated circumstellar disks, including the effect of companions on disks, for a subsample of Cha~I targets with available near-IR data.  Among the issues we address in this work is the dependence of multiplicity on primary mass, i.e., whether higher mass stars are more likely to be in binaries and multiples than their lower mass counterparts. An increase in wide binaries with increasing mass has been observed for both young stars and field dwarfs \\citep[e.g.,][]{2007ApJ...662..413K,2008ApJ...683..844L}. Furthermore, we look at how close-in multiplicity varies between different star-forming regions and the field; the multiplicity of $\\sim\\!0.1$--$2\\,M_{\\sun}$ field dwarfs in the solar neighborhood has been studied extensively \\citep[e.g.,][]{1991AA...248..485D,1992ApJ...396..178F,2010ApJS..190....1R}. We also explore whether close companions contribute to the observed difference between classical T~Tauri stars (CTTS), which are accreting, and weak-lined T~Tauri star (WTTS), which are not accreting based on weak H$\\alpha$ emission. Although the source of dichotomy between CTTS and WTTS is currently unknown, the presence of non-accreting $2$\\,Myr old stars is surprising. It has been suggested that the inner disks around weak-lined objects may have been truncated by close binary companions \\citep[e.g.,][and references therein]{2000prpl.conf..703M}. ",
        "conclusions": "\\label{sec:SummaryAndConcludingRemarks}  Binary and higher-order multiple systems are prevalent in young clusters, and studying the statistical properties of these systems offers indirect constraints on star formation. Here, we present a spectroscopic study of mulitiplicity for young stars in the Chamaeleon~I and Taurus-Auriga star-forming regions. Our main results are as follows.  \\begin{enumerate}  \\item The spectroscopic multiplicity fractions for Cha~I and Tau-Aur are similar to each other, and to those of field stars in the same mass and period regime (for details, see \\S\\ref{sec:BinaryPopulations}). This finding implies that the overall fraction of short period stellar companions could stablize after initial formation.  \\item Close multiplicity is not seen to depend on primary mass. The mass of the host star can become important in the separation regime of imaging surveys. For discussion, see \\S\\ref{sec:MassDependence}.  \\item There is no strong correlation between accretion and the frequency of systems with close companions, and thus, close stellar companions are unlikely the principal source of the accretion cutoff observed in WTTS (see \\S\\ref{sec:BinaryPopulations}). Moreover, if close companions are responsible solely for the accretion difference between CTTSs and WTTSs then there should be a disparity in the proportion of spectroscopic binaries between these two populations, but this is not seen in our survey of Cha~I and Tau-Aur.  \\end{enumerate}  By undertaking this extensive spectroscopic survey of T~Tauri stars, we have gained some insight that could be beneficial to future efforts. First and foremost, we now know the main limiting factor for a radial velocity study of young stars is the strong intrinsic noise present in some objects. This noise would need to be reduced in some way in order both to detect lower mass companions more effectively, and to extend the measurement to longer orbital periods. Our radial velocity precision is sufficient to detect, at the time of this writing, a few dozen of the known hot Jupiter planets, which have radial velocity semi-amplitudes of a few $100{\\rm\\,m\\,s}^{-1}$. If such planets were to be detected around pre-main sequence stars, it would provide direct constraints on their formation and migration timescales. For young stars, one way to reduce the intrinsic noise is to observe at infrared wavelengths \\citep{2006ApJ...644L..75M,2008AA...489L...9H}. Another option is to average out the noise by observing over its characteristic period, i.e., the rotation period of the noisy star. To some extent, we have already done this by taking multiple spectra of the same targets during an observing run. Ideally, one would multiplex observations, e.g., by using multiple fibers as in VLT/FLAMES. This solution may not be suitable for star-forming regions like Taurus-Auriga, which span a large area of the sky, but it would probably be quite useful for compact regions like Chamaeleon.  We have also learned that our high resolution and S/N were extremely beneficial. It was probably a major reason we could find so many SB2 candidates; being able to resolve the stellar rotation helped us to determine the stability of the broadening functions. Less clear is whether one would really need the large spectral range, or whether a few well-chosen echelle orders would suffice. This consideration is relevant since it would determine whether or not multiplexing is feasible."
    },
    "1112/1112.0234_arXiv.txt": {
        "abstract": "With the discovery of massive neutron stars such as PSR J1614-2230, the question has arisen whether exotic matter such as hyperons can exist in the neutron star core. We examine the conditions under which hyperons can exist in massive neutron stars. We consistently investigate the vector meson-hyperon coupling, going from SU(6) quark model to a broader SU(3) symmetry. We propose that the maximum neutron star mass decreases linearly with the strangeness content $f_s$ of the neutron star core as $M_{max}(f_s) = M_{max}(0) - 0.6 M_{\\odot} (f_s/0.1)$, which seems to be independent of the underlying nuclear equation of state and the vector baryon-meson coupling scheme. Thus, pulsar mass measurements can be used to constrain the hyperon fraction in neutron stars. ",
        "introduction": "The discovery of the massive neutron star PSR J1614-2230 has raised new challenges for theories of dense matter beyond nuclear saturation density. Shapiro delay measurements from radio timing observations of the binary millisecond pulsar indicate a large mass of 1.97$\\pm$0.04$M_{\\odot}$ of the neutron star \\cite{Demorest10}. The core of a neutron star harbors a dense matter environment, which could be the site for strangeness containing matter, such as hyperons. Though nuclear interactions in the saturation regime are well understood, one has to utilize neutron star observations to find clues about the physics of cold and dense matter beyond several times saturation density. Any theory of ultra dense matter has to explain the recently observed large neutron star mass. According to existing models of dense matter, the presence of hyperons leads  to a considerable softening of the equation of state (EoS), resulting in a corresponding reduction of the maximum mass of the neutron star \\cite{Glendenning92derivativecoupling,Baldo00,Vidana00,Djapo}. Then the existing theories involving hyperons are in conflict with the large pulsar masses \\cite{Lattimer10}. On including hyperons, most relativistic models obtain maximum neutron star masses in the range $1.4-1.8 M_{\\odot}$ \\cite{Glen85,GM1,Knorren,BalbergGal,PalHanauske,Hanauske00,Zschiesche,Long12}. However, in exceptional cases, neutron stars with maximum masses larger than 2$M_{\\odot}$ have been obtained, either by pushing the threshold for appearance of hyperons to higher densities, or due to strong hyperon vector repulsion \\cite{Long12,Huber,Hofmann,Rikovska,Dhiman07,Dexheimer08,Bombaci08,Cavagnioli11}. Taurines et al. \\cite{Taurines} achieved large neutron star masses including hyperons by considering a model with density dependent coupling constants, which were varied nonlinearly with the scalar field. Recently, Bednarek et al. \\cite{Bednarek2011} also achieved a stiffening of the EoS by using a non-linear relativistic mean field (RMF) model with quartic terms involving the hidden strangeness vector meson. In addition to the inclusion of such a meson into a density dependent RMF model, Lastowiecki et al. \\cite{Lastowiecki12} assumed a quark matter core in order to obtain massive stars. Bonanno and Sedrakian \\cite{Bonanno12} also modeled massive neutron stars with a hyperon and quark core using a fairly stiff EoS and vector repulsion among quarks. In several studies, the maximum neutron star masses obtained when including hyperons were not very different from those containing nucleons only \\cite{Hofmann,Rikovska,Dexheimer08}. In more sophisticated models such as the Brueckner-Hartree-Fock model, the maximum neutron star masses were generally found to be lower than $1.6M_{\\odot}$ which is in contradiction with observed pulsar masses \\cite{Baldo00,Vidana00,Djapo,Baldo98,Nishizaki,SchulzeVidana,SchulzeRijken,Logoteta12}.\\\\ From the studies cited in the previous paragraph, it seems that the possible presence of hyperons in massive neutron stars is in many cases reconciled by incorporating large vector repulsion in an ad hoc way. In contrast, we investigate the role of vector repulsion starting from symmetry arguments. Assuming SU(3) symmetry, we perform a controlled parameter study and constrain the parameters using the observed mass of PSR J1614-2230. This procedure is in line with modern microscopic models for realistic baryon-baryon potentials such as the Bonn potentials \\cite{Machleidt87} and the Nijmegen potentials \\cite{Rijken99} which adopt SU(3) symmetry to describe the baryon interactions for the baryon octet. For our investigations, we employ a RMF model, in which the parameters are calibrated around nuclear saturation density \\cite{Glen85,Knorren,Schaffner96}. However, the extrapolation of such properties to supranuclear densities presents uncertainties. In a previous paper \\cite{WsbChtt12}, we investigated how the uncertainty in nuclear saturation properties, such as effective nucleon mass or nuclear compressibility, or hypernuclear properties, such as potential depths of hyperons in nuclear matter, could influence our conclusions about the presence of hyperons in the core of massive neutron stars. In this work, we question the fundamental assumption of SU(6) symmetry, which relates the hyperon couplings to the nuclear couplings.\\\\ \\indent This paper is organized in the following way. In Sec. \\ref{model}, we describe the model to calculate the EoS. The parameters of the model are listed in Sec. \\ref{parameters}. The results of our calculations are discussed in Sec. \\ref{results}, and a summary of our conclusions is given in Sec. \\ref{summary}. \\\\ ",
        "conclusions": "\\label{results} \\subsection{Varying the $g_8/g_1$ Ratio $z$} We probe the effects of the $g_8/g_1$ ratio $z$ on the stiffness of the hadronic EoS. We plot the coupling constants as functions of $z$ in Fig. \\ref{figurezfree}, for the explicit formulae given in equation (\\ref{zfreecouplings}). For $z=0$, all coupling constants $g_{\\it{B\\omega}}$ are the same and similarly are all coupling constants $g_{\\it{B\\phi}}$ equal. This is due to the fact that $z=0$ corresponds to $g_8=0$, which results in the equality of the corresponding baryon-meson couplings as the baryons couple only to the flavor singlet state. With increasing $z$, i.e. with increasing contribution from the coupling to the octet $g_8$, the resulting couplings all become smaller except for $g_{\\it{\\Xi\\phi}}$. At $z=1/\\sqrt{6}\\approx0.4082$ the SU(6) case is reached where the $\\phi$ does not couple to the nucleon. Thereafter, for $z>1/\\sqrt{6}$ the coupling constant $g_{\\it{N\\phi}}$ changes its sign so that it is not a repulsive but now an attractive interaction. Note, that the $\\omega$ coupling constants for $\\Lambda$ and $\\Sigma$ hyperons are equal for all values of $z$. As anticipated in Sec. \\ref{beyondsu6}, we restrict $z$ to the interval $z\\in[0:2/\\sqrt{6}]$. \\\\ \\begin{figure} \\includegraphics[height=8.6cm,angle=270]{fig1_color} \\caption{Relative vector meson coupling constants as functions of the $g_8/g_1$ ratio $z$ for fixed $\\alpha_V=1$. The value $z=1/\\sqrt{6}$ corresponds to the SU(6) case.} \\label{figurezfree} \\end{figure} \\begin{figure} \\includegraphics[width=8.6cm]{fig2_color} \\caption{EoS for different $g_8/g_1$ ratios $z$ within a nonlinear $\\sigma-\\omega$ model with additional $\\phi$ meson and the full baryon octet for GM1 parameter set. The EoS get stiffer with decreasing $z$.} \\label{zfreeeosfigure} \\end{figure} \\indent The EoS for $z=0.1,0.2,...,0.8$ are plotted in Fig. \\ref{zfreeeosfigure}. At first glance it becomes clear that the EoS stiffens with decreasing $z$. This can be explained with the help of Fig. \\ref{figurezfree} where we had plotted the $z$ dependence of the vector meson coupling constants. We noticed that with increasing $z$, all couplings except $g_{\\it{\\Xi\\phi}}$ decrease. Since the $\\Xi$ hyperons only play a subordinate role compared to the neutrons, the increase of $g_{\\it{\\Xi\\phi}}$ does not prevent that part of the overall interaction between the baryons, which is mediated by the vector mesons, to become less repulsive. Therefore, the EoS must soften with increasing $z$. Together with further decreasing coupling strengths between the vector mesons and the other baryons (except again $g_{\\it{\\Xi\\phi}}$) the EoS becomes even softer until $z=2/\\sqrt{6}$ is reached. \\\\ \\begin{figure} \\includegraphics[width=08.6cm]{fig3_color} \\caption{Mass-radius relations for the EoS displayed in Fig. \\ref{zfreeeosfigure}. The maximum mass is obtained for the case $z=0$, where all baryons couple to the vector mesons with equal strengths.} \\label{zfreemrfigure} \\end{figure} \\indent In Fig. \\ref{zfreemrfigure} we plot the mass-radius relations for the various EoS we just discussed in the context of Fig. \\ref{zfreeeosfigure}. As expected from the influence of $z$ on the stiffness of the EoS, the lowest maximum mass is obtained for $z=2/\\sqrt{6}\\approx0.8165$, or in the case of Fig. \\ref{zfreemrfigure} at $z=0.8$, namely $M=1.49M_{\\odot}$. The maximum mass grows up to the value $M=2.36 M_{\\odot}$ for $z=0$. We notice that the maximum mass of a neutron star in our RMF model reacts rather strongly to the variation of $z$: over the whole $z$ range the change in the maximum mass is $\\Delta_M$=0.87$M_{\\odot}$.\\\\ \\indent After varying $z$ in rather big steps, we now plot in Fig. \\ref{zfreedifferentmodelsfigure} the maximum neutron star mass as a continuous function of $z$ for the models GM1, NL3 and TM1 and for nuclear matter as well as for baryonic matter. \\begin{figure} \\includegraphics[width=8.6cm]{fig4_color} \\caption{Maximum masses as functions of the $g_8/g_1$ ratio $z$ for NL3, GM1 and TM1 parameter sets. For each parameterization the case of pure nucleonic matter is also displayed.} \\label{zfreedifferentmodelsfigure} \\end{figure} As already analyzed in the discussion of Fig. \\ref{figurezfree}, we see that the branch for nucleonic matter is insensitive to the changes in $z$ for the NL3 and GM1 models. For TM1, the quartic self interaction term in the Lagrangian spoils this property since we had to keep the corresponding coupling constant $d$ from the SU(6) value of $z$ fixed for the whole $z$ range: the maximum masses for pure nucleonic stars therefore depend on the actual $z$ value, showing a minimum for the SU(6) case $z\\approx0.4082$, namely $M=2.18M_{\\odot}$. \\\\ \\begin{figure} \\includegraphics[width=8.6cm]{fig5_color} \\caption{Particle fractions for the GM1 parameter set for 3 different values of $z$: $z=0.8$ (a), $z=0.408$ (b) corresponding to the SU(6) case, and $z=0$ (c). The threshold for the appearance of the hyperons $\\Lambda$, $\\Xi^-$ and $\\Xi^0$ is pushed to higher densities with decreasing $z$.} \\label{GM1frac} \\end{figure} Considering the neutron stars containing hyperons, we see in Fig. \\ref{zfreedifferentmodelsfigure} that the maximum masses depend on $z$ as already observed in Fig. \\ref{zfreemrfigure}: for the largest $z$ values the maximum masses are the smallest and continually grow with decreasing $z$. It is interesting to see that towards $z=0$ the maximum masses of these stars seem to approach the maximum masses of the corresponding pure nucleonic stars for all parameter sets studied. A look at the particle number fractions for the GM1 parameter set at $z=0.8$ (which we plot in Fig. \\ref{GM1frac}(a)) shows that the first hyperons to appear in the hadronic matter are the $\\Xi^-$ and the $\\Lambda$ at total baryon number densities of $n_b\\approx$ 0.28 ${fm}^{-3}$ and $n_b\\approx$ 0.29 ${fm}^{-3}$ respectively, while the $\\Xi^0$ appears much later at $n_b\\approx$ 0.76 ${fm}^{-3}$. On increasing $z$ to its SU(6) value (Fig. \\ref{GM1frac}(b)), the $\\Lambda$ hyperon appears first at $n_b\\approx$ 0.36 ${fm}^{-3}$, followed by $\\Xi^-$ at $n_b\\approx$ 0.4 ${fm}^{-3}$ and $\\Xi^0$ at $n_b\\approx$ 0.89 ${fm}^{-3}$. At $z=0$ (Fig. \\ref{GM1frac}(c)), the threshold of appearance of hyperons is pushed to even higher densities: $n_b\\approx$ 0.73 ${fm}^{-3}$ for $\\Lambda$, $n_b\\approx$ 0.74 ${fm}^{-3}$ for  $\\Xi^-$ and  $n_b\\approx$ 1.38 ${fm}^{-3}$ for $\\Xi^0$. Thus, for $z=0$ the neutron stars consist mainly of nuclear matter which is why the maximum masses are so close to those of the pure nucleonic stars. For the case of the parameter sets NL3 and TM1, the particle fractions are qualitatively the same as in the GM1 case. In the case of NL3 parameterization, a well-known instability occurs at high densities when the effective nucleon mass becomes zero \\cite{Schaffner96}. The critical density for the appearance of the instability depends on the value of the hyperon coupling constants. However for the present investigation, this instability plays no role as it appears beyond the maximum densities reached in the neutron star interior. \\subsubsection{Combining $m_N^*$ and $z$ variations} The impact of z on the maximum mass of neutron stars is as comparably large as the influence of the effective nucleon mass at saturation $m_N^*$ as investigated in our previous study \\cite{WsbChtt12}. We therefore combine both parameters in a single plot, Fig. \\ref{figuremeffz}, where we show the maximum neutron star mass as a function of $m_N^*/m_N$ for different $z$ values. The incompressibility in this case is fixed to $K=240$ MeV, but the exact value is irrelevant as shown in our previous paper \\cite{WsbChtt12}. We see in Fig. \\ref{figuremeffz} that the effective mass has basically the same effect for all $z$ values and $z$ the same effect for all effective masses: for fixed $z$, the maximum masses decrease drastically for increasing effective mass. For low $z$ values, where the EoS is stiffer than for higher $z$ values, the dependence of the maximum mass on the effective mass is slightly larger: the difference along the whole range of $m_N^*/m_N$ is $\\approx$ 0.6${M}_\\odot$ for $z=0.8$ while for $z=0$ it is $\\approx$ 0.9${M}_\\odot$. The influence of $z$ on the maximum masses is slightly more pronounced than that of the effective masses: the difference in the maximum mass between $z=0$ and $z=0.8$ is $\\approx$ 0.65${M}_\\odot$ for $m_N^*/m_N=0.8$ and about 1${M}_{\\odot}$ for $m_N^*/m_N=0.55$. For comparison, we also plot the maximum masses of purely nucleonic neutron stars, and also mark in the figure the points corresponding to the SU(6) case, for several other RMF sets fitted to properties of nuclei e.g., TM1, NL3 or NL-SH \\cite{NL3,Sugahara94,NL-Z,PL-Z,NL-SH}. In Fig. \\ref{zfreedifferentmodelsfigure} the masses of baryonic stars were found to approach the limit of purely nucleonic stars for decreasing values of z. With Fig. \\ref{figuremeffz} we can now visualize that the maximum masses in these two cases differ slightly ($\\Delta_M<0.01M_\\odot$).\\\\ \\indent We close this section by concluding that a maximum mass of at least 1.97 $\\pm$ 0.04${M}_{\\odot}$ requires very small values of $z$ for large effective masses ($z\\leq0.1$ at $m_N^*/m_N=0.8$), $z$ values around SU(6) for effective masses close to $m_N^*/m_N \\approx 0.7$ and very low effective masses for large $z$ values ($m_N^*/m_N < 0.58$ for $z=0.7$). We note, that $z$ values close to the maximum of $z=2/\\sqrt{6}$ are now allowed configurations within the plot range: the maximally allowed $z$ value for the investigated model is $z < 0.77$ at $m_N^*/m_N=0.55$. \\begin{figure} \\includegraphics[width=8.6cm]{fig6_color} \\caption{Maximum masses of hyperonic neutron stars as functions of effective nucleon mass $m_N^*$ for different values of the $g_8/g_1$ ratio $z$. For comparison, a line for nucleonic stars and points to mark RMF sets (e.g. TM1, NL3) corresponding to the SU(6) case are also given.} \\label{figuremeffz} \\end{figure} \\subsection{Varying the F/(F+D) ratio $\\alpha_V$} After systematically investigating a wider $z$ range we repeat the formalism for the $\\alpha_V$-ratio in the present section. The $F/(F+D)$ ratio is by definition restricted to the interval $\\alpha_V\\in[0;1]$, where the lower bound corresponds to a pure D-type coupling and the upper limit (i.e. the SU(6) value) corresponds to a pure F-type coupling. Analogous to the case studied above, we adopt ideal mixing as well as a $g_8/g_1$ ratio fixed to its SU(6) value $z=1/\\sqrt{6}$ and we allow for the $\\phi$ meson to couple to the nucleon. We plot the coupling strengths in Fig. \\ref{figurealphafreeB} where we vary $\\alpha_V$ between 0 and 1. Note that the coupling constants $g_{\\omega \\Lambda}$ for $\\Lambda$ hyperons do not change considerably and that all vector couplings remain repulsive (do not change their sign).\\\\ \\begin{figure} \\includegraphics[height=8.6cm,angle=270]{fig7_color} \\caption{Vector meson coupling constants as functions of the $F/(F+D)$ ratio $\\alpha_V$. The $g_8/g_1$ ratio is fixed to its SU(6) value $z=1/\\sqrt{6}$ and ideal mixing is assumed. The SU(6) case is given by $\\alpha_V=1$.} \\label{figurealphafreeB} \\end{figure} \\indent The ratio of the F- and D-type couplings can be continuously varied between the two extremes of a pure F- and a pure D-type coupling. We see that from $\\alpha_V=1$ down to $\\alpha_V=0$ all couplings become stronger except for $g_{\\it{\\Sigma\\phi}}/g_{\\it{N\\omega}}$. Since the $\\Sigma$ hyperons have but very little influence on the EoS up to neutron star densities, we can expect that the EoS become stiffer with decreasing $\\alpha_V$. This is exactly what we find in Fig. \\ref{figurealphaeosB} where we plot the EoS for several values of $\\alpha_V$ using \"model $\\sigma \\omega \\rho \\phi$\" and GM1 parameter set.\\\\ \\begin{figure} \\includegraphics[width=8.6cm]{fig8_color} \\caption{EoS for ``model $\\sigma \\omega \\rho \\phi$`` in GM1 parameterization for different values of $\\alpha_V$. $z$ is fixed to its SU(6) value $z=1/\\sqrt{6}$ and ideal mixing is assumed. The EoS become stiffer with decreasing $\\alpha_V$.} \\label{figurealphaeosB} \\end{figure} \\indent The stiffness depends monotonously on $\\alpha_V$ and we get the softest EoS for the SU(6) case $\\alpha_V=1$, i.e. a pure F-type coupling, while the stiffest EoS is obtained for $\\alpha_V=0$ which corresponds to the pure D-type coupling of the baryon and meson multiplets. The EoS for $\\alpha_V\\leq0.2$ appear to be indistinguishable at neutron star densities. This is evident in Fig. \\ref{figurealphamrB}, where we plot the mass-radius relations corresponding to the EoS from Fig. \\ref{figurealphaeosB}: for the values $\\alpha_V=0.0-0.2$ the maximum masses but also the radii of the corresponding stars coincide (M$_{\\it{max}}=2.36$M$_\\odot$, R$=11.8$ km). We note, that this value of the maximum mass is also obtained for the purely nucleonic case (compare e.g. Fig. \\ref{zfreedifferentmodelsfigure}). Thus, for the \"model $\\sigma \\omega \\rho \\phi$\" the nuclear matter limit is reached below $\\alpha_V < 0.2$ in the case of GM1, GM3 and NL3 parameter sets, while for the very stiff PL-Z EoS (having effective mass $m^*_N/m_N \\simeq 0.55$) \\cite{PL-Z} pure nucleonic stars are already obtained for $\\alpha_V < 0.3$. \\\\ \\begin{figure} \\includegraphics[width=8.6cm]{fig9_color} \\caption{Mass-radius relations obtained from the EoS in Fig. \\ref{figurealphaeosB}.} \\label{figurealphamrB} \\end{figure} \\indent In this way, for hyperonic stars the limit of nucleonic stars is continuously approached for decreasing values of the $g_8/g_1$ ratio $z$, or for decreasing values of the $F/(F+D)$ ratio $\\alpha_V$ away from the SU(6) value, respectively. \\\\ \\\\ \\begin{figure} \\includegraphics[width=8.6cm]{fig10_color} \\caption{Maximum masses of neutron stars as functions of strangeness fraction $f_s$ in the neutron star core for four different EoS. The maximum mass decreases linearly with the strangeness fraction approximately as $M_{max}(f_s) = M_{max}(0) - 0.6 M_{\\odot} (f_s/0.1)$.} \\label{mmax_fs} \\end{figure} \\indent To generalize our findings, we plot in Fig. \\ref{mmax_fs} the maximum masses of neutron stars as a function of strangeness fraction $f_s$ (the number of strange quarks divided by the total number of quarks) for four different EoS: for $m_N^*/m_N$ = 0.55, 0.75 and for the NL3 ($m^*_N/m_N = 0.6$) and GM1 ($m^*_N/m_N = 0.7$) parameter sets, by varying $\\alpha_V$ (solid lines) and $z$ (dashed lines). It is evident from the figure that on decreasing $\\alpha_V$ or $z$, the strangeness fraction in the core decreases, and there is a corresponding increase in the maximum mass of the star. At zero strangeness fraction the nucleonic limit is reached, and this corresponds to the highest value of the maximum mass. For $m^*_N/m_N \\lessapprox 0.7$, the relation between the maximum mass of the star and its strangeness fraction can be fitted linearly according to the formula: \\begin{equation} \\frac{M_{max}}{M_{\\odot}} = \\frac{M_{max}  (f_s = 0)}{M_{\\odot}} - c \\left(\\frac{f_s}{0.1}\\right)~, \\end{equation} where $c\\approx0.6M_{\\odot}$. This has interesting consequences when we use these results to predict the maximally allowed strangeness fraction in maximum mass neutron stars. In an associated study \\cite{TolosSagert}, we applied the results from measurements of sub-threshold kaon production in heavy-ion collisions (HIC) to study the implications on neutron star properties. It was found that the heavy-ion data and causality imply a firm upper limit on maximum mass of compact stars of 3 solar masses. Substituting this value into the derived formula above, we can show that the strangeness fraction in a 2 $M_{\\odot}$ star cannot be more than: \\begin{eqnarray} f_s^{max} &=& \\frac{M_{max}^{HIC} - M_{max}^{obs}}{6 M_{\\odot}} \\nonumber\\\\ &=& \\frac{3 M_{\\odot} - 2 M_{\\odot}}{6 M_{\\odot}} = 0.17 ~. \\end{eqnarray} \\indent We also point out that for the EoS compatible with the observed mass limit, the hyperon fraction in a canonical $1.4M_{\\odot}$ star is zero. Only for small effective nucleon masses and values of $z$ above the SU(6) value can a hyperon fraction of less than $0.5\\%$ be reached.\\\\ \\indent We have now found that we can reach the nuclear matter limit starting from a RMF model including hyperons and continuously changing the model parameters, instead of making a discrete ``on/off'' decision about whether to include or exclude hyperons. Instead, we should take the position of saying that one can only exclude a RMF parameter set as soon as the corresponding maximum mass for nucleonic stars is incompatible with observations. As long as the observational mass limit is below the nucleonic mass limit of the model, it is possible to have hyperons in the core of maximum mass neutron stars."
    },
    "1112/1112.1697_arXiv.txt": {
        "abstract": "We present continued radio observations of the tidal disruption event Swift\\,J164449.3+573451 extending to $\\delta t\\approx 216$ days after discovery.  The data were obtained with the EVLA, AMI Large Array, CARMA, the SMA, and the VLBA+Effelsberg as part of a long-term program to monitor the expansion and energy scale of the relativistic outflow, and to trace the parsec-scale environment around a previously-dormant supermassive black hole (SMBH).  The new observations reveal a significant change in the radio evolution starting at $\\delta t\\approx 1$ month, with a brightening at all frequencies that requires an increase in the energy by about an order of magnitude, and an overall density profile around the SMBH of $\\rho\\propto r^{-3/2}$ ($0.1-1.2$ pc) with a significant flattening at $r\\approx 0.4-0.6$ pc.  The increase in energy cannot be explained with continuous injection from an $L\\propto t^{-5/3}$ tail, which is observed in the X-rays.  Instead, we conclude that the relativistic jet was launched with a wide range of Lorentz factors, obeying $E(>\\Gamma_j)\\propto \\Gamma_j^{-2.5}$.  The similar ratio of duration to dynamical timescale for \\sw\\ and GRBs suggests that this result may be applicable to GRB jets as well.  The radial density profile may be indicative of Bondi accretion, with the inferred flattening at $r\\sim 0.5$ pc in good agreement with the Bondi radius for a $\\sim {\\rm few}\\times 10^6$ M$_\\odot$ black hole.  The density at $\\sim 0.5$ pc is about a factor of 30 times lower than inferred for the Milky Way galactic center, potentially due to a smaller number of mass-shedding massive stars.  From our latest observations ($\\delta t\\approx 216$ d) we find that the jet energy is $E_{\\rm j,iso}\\approx 5\\times 10^{53}$ erg ($E_j\\approx 2.4\\times 10^{51}$ erg for $\\theta_j=0.1$), the radius is $r\\approx 1.2$ pc, the Lorentz factor is $\\Gamma_j\\approx 2.2$, the ambient density is $n\\approx 0.2$ cm$^{-3}$, and the projected angular size is $r_{\\rm proj}\\approx 25$ $\\mu$as, below the resolution of the VLBA+Effelsberg.  Assuming no future changes in the observed evolution and a final integrated total energy of $E_j\\approx 10^{52}$ erg, we predict that the radio emission from \\sw\\ should be detectable with the EVLA for several decades, and will be resolvable with VLBI in a few years. ",
        "introduction": "\\label{sec:into}  The discovery of the unusual $\\gamma$-ray/X-ray transient Swift\\,J164449.3+573451 (hereafter, \\sw), which coincided with the nucleus of an inactive galaxy at $z=0.354$, has opened a new window into high-energy transient phenomena, with potential implications to our understanding of relativistic outflows in systems such as gamma-ray bursts (GRBs) and active galactic nuclei (AGN).  The prevailing interpretation for this event is the tidal disruption of a star by a dormant supermassive black hole (SMBH) with a mass of $M_{\\rm BH}\\sim 10^6-10^7$ M$_\\odot$ (\\citealt{bgm+11,bkg+11,ltc+11,zbs+11}; but see \\citealt{kp11,osj11,qk11} for alternative explanations).  The argument for a tidal disruption origin is based on: (i) a positional coincidence ($\\lesssim 0.2$ kpc) with the host galaxy nucleus; (ii) rapid time variability in $\\gamma$-rays and X-rays ($\\lesssim 10^2$ s), which requires a compact source of $\\lesssim 0.15$ AU, a few times the Schwarzschild radius of a $\\sim 10^6$ M$_\\odot$ black hole; (iii) high $\\gamma$-ray and X-ray luminosity of $\\sim 10^{47}$ erg s$^{-1}$, which exceeds the Eddington limit of a $\\sim 10^6$ M$_\\odot$ black hole by $2-3$ orders of magnitude; (iv) a lack of previous radio to $\\gamma$-ray activity from this source to much deeper limits than the observed outburst, pointing to a rapid onset; and (v) long-term X-ray luminosity evolution following $L_X\\propto t^{-5/3}$, as expected from the fallback of tidally disrupted material (e.g., \\citealt{ree88,sq09}).  Equally important, \\sw\\ was accompanied by bright radio synchrotron emission, with an initial peak in the millimeter band ($F_\\nu\\approx 35$ mJy) and a steep spectral slope at lower frequencies indicative of self-absorption (\\citealt{zbs+11}; hereafter, ZBS11).  The properties of the radio emission established the existence of a relativistic outflow with a Lorentz factor of $\\Gamma\\sim {\\rm few}$ (ZBS11, \\citealt{bgm+11}).  The spectral energy distribution also demonstrated that the lack of detected optical variability required significant rest-frame extinction ($A_V\\gtrsim 5$ mag; ZBS11, \\citealt{ltc+11}), and that the X-rays were produced by a distinct emission component, rather than inverse Compton scattering by the radio-emitting relativistic electrons (ZBS11).  Finally, the evolution of the radio emission on a timescale of $\\delta t\\sim 5-22$ d pointed to an ambient density with a radial profile of roughly $\\rho\\propto r^{-2}$, as well as a mild increase in the energy of the outflow (ZBS11).  The formation of a relativistic jet with dominant X-ray and radio emission were not predicted in standard tidal disruption models (e.g., \\citealt{ree88,sq09}), which instead focused on the thermal optical/UV emission from the long-term accretion of the stellar debris.  A signature of the latter process is a mass accretion rate that evolves as $\\dot{M}\\propto t^{-5/3}$, presumably leading to emission with the same temporal dependence (e.g., \\citealt{kg99,gbm+08,vfg+11}). Shortly before the discovery of \\sw, \\citet{gm11} investigated the potential signature of a putative relativistic outflow, and concluded that the interaction of the outflow with the ambient medium will lead to radio emission on a timescale of $\\delta t\\sim 1$ yr (for typical off-axis observers).  While the mechanism for the radio emission from \\sw\\ is interaction with an external medium, the actual light curves differ from the off-axis prediction.  To address this issue, in a follow-up paper \\citet{mgm11} (hereafter, MGM11) reconsidered the model for a relativistic jet interacting with an ambient medium.  They draw on the inferences from the early radio emission described in ZBS11 to infer the properties of the environment and the jet kinetic energy, and use this information to predict the future evolution of the radio emission.  This long-term radio evolution is of great interest because it can provide several critical insights: \\begin{itemize} \\item The integrated energy release in the relativistic outflow, including the anticipated injection from on-going accretion. \\item The density profile around a previously-dormant SMBH on $\\sim 0.1-10$ pc scales, which cannot be otherwise probed in AGN. \\item The potential to spatially resolve the outflow with very long baseline interferometry (VLBI), and hence to measure the dynamical evolution (expansion and potentially spreading) of a relativistic jet. \\item Predictions for the radio emission from tidal disruption jets as viewed by off-axis observers on timescales of months to years to decades. \\end{itemize} The energy scale and jet dynamics are of particular importance since the total energy input and the structure of the jet may also have implications for relativistic jets in GRBs and AGN.  The ability to trace the environment on parsec scales provides a unique probe of gas inflow or outflow around an inactive SMBH on scales that cannot be probed outside of the Milky Way.  Finally, the long-term radio emission from \\sw\\ will inform future radio searches for tidal disruption events (TDEs) that can overcome the low detection rate in $\\gamma$-rays/X-rays (due to beaming), and obscuration due to extinction in the optical/UV (as in the case of \\sw).  To extract these critical properties we are undertaking long-term monitoring of the radio emission from \\sw\\ using a wide range of centimeter- and millimeter-band facilities.  Here we present radio observations of \\sw\\ that extend to $\\delta t\\approx 216$ d, and use these observations to determine the evolution of the total energy and ambient density.  We find that the evolution of both quantities deviates from the behavior at $\\delta t\\lesssim 1$ month (presented in ZBS11), thereby providing crucial insight into the structure of the relativistic outflow and the ambient medium.  This paper is the first in a series that will investigate the long-term radio evolution of \\sw\\ and the implications for relativistic jets and parsec-scale environments around supermassive black holes, including efforts to resolve the source with VLBI and to measure polarization.  The current paper is organized as follows.  We describe the radio observations in \\S\\ref{sec:obs}, and summarize the radio evolution at $\\delta t\\approx 5-216$ d in \\S\\ref{sec:prem}.  In \\S\\ref{sec:model} we present our modeling of the radio emission, which utilizes the formulation of MGM11.  The implications for the energy scale and ambient density are discussed in \\S\\ref{sec:energy} and \\S\\ref{sec:density}, respectively, and we finally consider the implications for relativistic jets and the parsec-scale environments of SMBHs in \\S\\ref{sec:implic}. ",
        "conclusions": "\\label{sec:conc}  We presented radio observations of \\sw\\ extending to $\\delta t\\approx 216$ d and spanning a wide range of frequencies.  The evolution of the radio emission changes dramatically at $\\delta t \\gtrsim 1$ month, requiring an increase in the total energy by about an order of magnitude, a density profile of $\\rho\\propto r^{-3/2}$ ($0.1-1.2$ pc), and a flattening at $r\\approx 0.4-0.6$ pc.  A comparison of the model to optical limits and near-IR detections indicates a cooling break at $\\nu_c\\sim 10^{13}$ Hz and host galaxy extinction of $A_{\\rm V,host} \\gtrsim 3.5$ mag.  The increase in energy cannot be explained by injection from an $L\\propto t^{-5/3}$ tail that is expected in tidal disruption events and which matches the evolution of the X-ray emission.  We conclude that a natural explanation is a structured outflow with $E(>\\Gamma_j)\\propto \\Gamma_j^{-2.5}$.  The inferred density profile and the radial scale of the density enhancement are in rough agreement with the expectation for Bondi accretion from a circumnuclear medium. The jet energetics and structure, as well as the detailed density profile on $\\sim 0.1-1$ pc scale are a testament to the important insight that can be gained from continued radio observations of \\sw. In particular, the radial density profile is traced in greater detail than even the inner parsec of the Milky Way.  Continued radio observations will probe the environment to a scale of $\\sim 10$ pc in the coming decade.  Using the results of our analysis we can predict the future evolution of the radio emission (modulo any future unpredictable changes in energy and/or density as we have found here).  We use the evolution of $L_{\\rm j,iso}$ and $n_{\\rm CNM}$ as inferred from the data at $\\delta t\\lesssim 216$ d, and assume that the density will continue to evolve as $\\rho\\propto r^{-1.5}$ and that the energy will increase to a maximum beaming-corrected value of $E_j=L_{\\rm j,iso}t_j[1-{\\rm cos} (\\theta_j)]$ with $E_j=10^{52},\\,3\\times 10^{52},\\,10^{53}$ erg.  The resulting light curves at 6 and 22 GHz are shown in Figure~\\ref{fig:predlc}.  The long-term evolution is marked by a break when $E_j$ achieves its maximum value, corresponding to about $5$, $28$, and $180$ yr for our three choices of maximum energy.  Using the $5\\sigma$ sensitivity of the EVLA in an observation of a few hours\\footnotemark\\footnotetext{At 22 GHz we use the sensitivity for the full 8 GHz bandwidth that will become available some time in 2012.}, we find that the emission at 22 and 6 GHz should be detectable for at least $\\sim 40$ yr and $\\sim 80$ yr, respectively.  Indeed, any significant upgrades to the EVLA or the construction of more sensitive radio facilities in the coming decades may extend the range of detectability to centuries\\footnotemark\\footnotetext{Significant budget cuts to radio facilities in the future may lead to the opposite effect.}.  The same is true if the total energy scale is $\\sim 10^{53}$ erg.  An equally important question is whether the jet will be resolvable with VLBI in the future.  The projected radius is $r_{\\rm proj}\\approx r\\theta_j$, as long as the jet maintains its collimation.  In Figure~\\ref{fig:predrad} we plot the predicted future evolution of $r$ using the prescription described above.  We find that for $\\theta_j=0.1$ and a best-case VLBI angular resolution of $\\approx 0.2$ mas (FWHM), the source should become resolvable at $\\delta t\\approx 6$ yr.  On this timescale the 22 GHz flux density is expected to be only $\\approx 2$ mJy (Figure~\\ref{fig:predlc}), still accessible with VLBI.  While the flux density at 6 GHz is expected to be larger by about a factor of $2.6$, the angular resolution at this frequency is poorer by about a factor of 3.7, making it less competitive than 22 GHz.  Thus, we conclude that the radio emission from \\sw\\ may be marginally resolved in a few years.  On the other hand, if the jet undergoes significant spreading on the timescale at which it becomes non-relativistic (as expected for GRB jets: e.g., \\citealt{lw00}) it is possible that it will become resolvable at $\\delta t\\sim 1-2$ yr when the expected 22 GHz flux density is still $\\sim 10$ mJy.  We are undertaking continued multi-frequency radio monitoring of \\sw\\ to follow the long-term evolution of the relativistic outflow and the radial profile of the ambient medium.  Even in the absence of any future dramatic changes relative to the current evolution, we expect that in the next few years we may be able to determine the total energy of the relativistic outflow, measure the spreading of the jet, and study the radial density profile to a scale of $\\sim 10$ pc. Future papers in this series will detail these results."
    },
    "1112/1112.2668.txt": {
        "abstract": "V\\,393 Scorpii is a bright Galactic Double Periodic Variable showing a long photometric cycle of $\\approx$ 253 days.  We present new $VIJK$ photometric time series for V\\,393 Scorpii along with the analysis of ASAS $V$-band photometry. We disentangled all light curves into the orbital and long cycle components. The  ASAS $V$-band {\\it orbital} light curve was modeled with two stellar components  plus a circumprimary optically thick  disc assuming a semidetached configuration. We present the results of this calculation, giving physical parameters for the stars and the disc, along with general system dimensions. Our results are in close agreement with those previously found by Mennickent et al. (2010) from IR spectroscopy and the modeling of the spectral energy distribution. The stability of the orbital light curve suggests that the stellar $+$ disc configuration remains stable during the long cycle. Therefore, the long cycle should be produced by an additional variable and not-eclipsed emitting structure. We discuss the evolutionary stage of the system finding the best match with one of the  evolutionary models of van Rensbergen et al. (2008).  According to these models, the system is found  to be after an episode of fast mass exchange that transferred 4 M$_{\\odot}$ from the donor to the gainer  in a period of 400.000 years. We argue that a significant fraction of this mass has not been accreted by the gainer but remains in an optically thick {\\it massive} ($\\sim$ 2 M$_{\\odot}$) disc-like surrounding pseudo-photosphere  whose luminosity is not driven by viscosity but probably by reprocessed stellar radiation. Finally, we provide the result of our search for Galactic Double Periodic Variables and briefly discuss the outliers $\\beta$ Lyr and $RX$ Cas. ",
        "introduction": "V\\,393 Scorpii is one of the Galactic  Double Periodic Variables (DPVs),  a group of interacting binaries showing a long photometric cycle lasting roughly 33 times  the orbital period (Mennickent et al. 2003, Mennickent \\& Ko{\\l}aczkowski 2009, Michalska et al. 2009, Poleski et al. 2010). DPVs  have been interpreted as semi-detached interacting binaries   with ongoing cyclic episodes of mass loss into the interstellar medium  (Mennickent et al. 2008, Mennickent \\& Ko{\\l}aczkowski 2010). The 253-d long photometric cycle of V\\,393\\,Scorpii was discovered by Pilecki \\&  Szczygiel (2007)  after inspection of the ASAS catalogue for eclipsing binaries with additional variability. The IUE-UV properties of V\\,393 Scorpii were studied by Peters (2001) who found evidence for a hot temperature region produced by the tangential impact of the gas stream into the gainer photosphere. This region should be the origin of the superionized lines observed in the UV,  like N\\,V, C\\,IV and Si\\,IV, that are likely produced by resonance scattering in a plasma of  temperature T $\\sim$ 10$^{5}$ K and electron density $N_{e} \\sim$  10$^{9}$ cm$^{-3}$ (Peters \\& Polidan 1984). The star was also studied by means of  multi-epoch high-resolution  IR spectroscopy by Mennickent et al. (2010,  hereafter M10), who also studied broad-band photometry and IUE ultraviolet spectra. After summarizing the available literature of this object, these authors argued for a semidetached B3 + F0  binary with masses 8  M$_{\\odot}$ and  2 M$_{\\odot}$  for the gainer and donor  (hereafter also called  primary and secondary, respectively) and orbital separation of 35 $R_{\\odot}$. Most remarkably, M10 found evidence  for large mass loss through the Lagrangian L3 point during epochs of long cycle minimum and claim that their observations  suggest that the mass loss producing the long cycle is probably concentrated in equatorial regions. %None detailed spectroscopic study of the long cycle of V\\,393\\,Sco has been still published.  In this paper we refine stellar and  system parameters of V\\,393 Scorpii by fitting the light curve with a multicomponent model including a stationary circumprimary disc. %and present  new insights on the nature of the long cycle based on a multi-site and multi-epoch high-resolution optical spectroscopic study. A detailed report of the observations used in this paper is given in Section 2, our results are presented in Section 3, a detailed discussion of these results is given in Section 4 and our conclusions are presented in Section 5. %Preliminary reports of this research were recently published in conference proceedings (Mennickent et al. 2010x, 2010x). %showing larger H$\\alpha$ emission during the long cycle maximum and a modulation of the radial velocity and full width of half maximum (FWHM) of the He\\,I\\,5875 line during the long cycle. ",
        "conclusions": "In this paper we have modeled the {\\it orbital} light curve of the intermediate-mass interacting binary \\var  to obtain stellar and system parameters. We also disentangled the long-term light curve at optical and infrared photometric bands. We have found insights on the system evolutionary stage and  long cycle nature. The main results of our research are:\\\\  \\begin{itemize} \\item The orbital period change, if present, is shorter than 0.5 sec per year. % $\\dot{P_{o}} < $ 0.5 sec yr$^{-1}$. \\item The long-term light curves  are characterized by a smooth oscillation in a time scale of 253 days with larger amplitudes in redder bandpasses. \\item The best fit to the orbital light curve requires a non-stellar component that was modeled with an optically thick disc model. The disc radius is about half of the Roche lobe radius of the gainer and two bright spots are required to fit the observations. \\item We found the stellar and system parameters that best match the observations,  which are given in Table 2 along with parameters for the  disc and the bright spots. \\item The stability of the orbital light curve suggests that the stellar $+$ disc configuration remains stable during the long cycle. Variability of the optically thick disc is not the main source for long cycle. \\item Therefore and in order to fit the redder color at long maximum, we argue that the long cycle is produced by free-free emission in a variable structure, probably visible perpendicular to the orbital plane, something reminiscent of the jets found in $\\beta$ Lyr (Harmanec et al. 1996). The suggestion of equatorial mass loss as the cause of the long cycle by M10 is probably  biased by detection of equatorial outflows not necessarily related to the long cycle. \\item A comparison with published evolutionary tracks provides an estimate for the  age of the system, namely log t =  7 $\\times$ 10$^{7}$  yr. We find the system after a mass exchange episode, where 4 M$_{\\odot}$ were transferred from the donor to the gainer in a period of 400.000 years. \\item The evolutionary  model with initial stellar masses of 6 and 3.6 M$_{\\odot}$ reproduces relatively well the present donor parameters and orbital period, but overestimates the gainer temperature and luminosity, a fact that could be ascribed to the optically thick disc not considered in the evolutionary tracks. In order to explain these features, we argue that part of the mass, maybe up to 2  M$_{\\odot}$ has not been accreted by the gainer, but remains in the {\\it massive} optically-thick disc. \\item We find a discrepancy between  values of $\\dot{M}$ derived from theoretical model fitting of observationally-derived parameters and those derived from accretion-theory analysis of observationally-derived parameters,  the former being smaller by 2 orders of magnitude. If the optically thick disc is a representation of a  massive  disc-like pseudo-photosphere, then its luminosity could not be accretion-driven. The constancy of the orbital period  supports the view that the disk luminosity is not driven by viscosity, but probably by reprocessed stellar radiation. %Alternatively, V\\,393\\,Sco could be at the end of the first mass transfer burst or starting the second one (i.e. 3\\% younger and 7\\% older, respectively). These models fit $\\dot{M}$ obtained from accretion-disc theory, but their $\\chi^{2}$ are larger by a factor of {\\bf 10} than the best fit model.    For both models, our result of a massive disc-like pseudo-photosphere remains unaltered. \\item We present the results of our search for Galactic DPVs and make a comparison  with hot emission-line binaries with cyclic long-term brightness changes.  13 Galactic DPVs show a similar correlation between  $P_{o}$  and $P_{long}$   to  that observed in LMC and SMC DPVs. The systems $\\beta$ Lyr and RX Cas  deviate from this tendency. These systems could be in an earlier evolutionary stage compared with DPVs.  \\end{itemize}"
    },
    "1112/1112.4598_arXiv.txt": {
        "abstract": "The existence of 10$^9$ M$_{\\odot}$ black holes (BH) in massive galaxies by $z \\ sim$ 7 is one of the great unsolved mysteries in cosmological structure formation. One theory argues that they originate from the black holes of Pop III stars at $z \\sim$ 20 and then accrete at the Eddington limit down to the epoch of reionization, which requires that they have constant access to rich supplies of fuel.  Because early numerical simulations suggested that Pop III stars were $\\gtrsim$ 100 M$_{\\odot}$, the supermassive black hole seeds considered up to now were 100 - 300 M$_{\\odot}$. However, there is a growing numerical and observational consensus that some Pop III stars were tens of solar masses, not hundreds, and that 20 - 40 M$_{\\odot}$ black holes may have been much more plentiful at high redshift.  However, we find that natal kicks imparted to 20 - 40 M$_{\\odot}$ Pop III BHs during formation eject them from their halos and hence their fuel supply, precluding them from Eddington-limit growth.  Consequently, supermassive black holes are far less likely to form from low-mass Pop III stars than from very massive ones. ",
        "introduction": "The existence of 10$^9$ M$_{\\odot}$ black holes (BH) in massive galaxies by $z \\sim 7$, only a billion years after the Big Bang \\citep[e.g.][]{mort11}, poses one of the great unsolved problems in cosmological structure formation.   In the $\\Lambda$CDM paradigm, early structure formation is hierarchical, with small objects at high redshifts evolving into ever more massive ones by accretion and mergers through cosmic time. For this reason it is generally supposed that the supermassive black holes (SMBH) that power the $z \\sim 7$ \\textit{Sloan Digital Sky Survey} (\\textit{SDSS})  quasars grow from much smaller seeds at earlier epochs.  The origin of SMBH and how they reach such large masses in such short times is a subject of ongoing debate.  Three modes of formation have been proposed for SMBH seeds: the collapse of Pop III stars into 100 - 300 M$_{\\odot}$ black holes at $z \\sim 20$ \\citep{awa09}, baryon collapse in 10$^8$ M$_{\\odot}$ dark matter halos that have somehow bypassed previous star formation into 10$^4$ - 10$^6$ M$_{\\odot}$ BH at $z \\sim 15$ \\citep{wta08,rh09,sbh10}, and more exotic pathways like the relativistic collapse of dense primeval star clusters into 10$^4$ - 10$^6$ M$_{\\odot}$ BH \\citep[see section 3.3 of][for a recent review]{brmvol08}.  Stellar-mass SMBH seeds form at $z \\sim 20$ when Pop III stars die in either core-collapse supernovae (SNe, 15 - 45 M$_{\\odot}$) or by direct collapse to a BH (45 - 100 M$_{\\odot}$, $\\gtrsim 260$ M$_{\\odot}$) \\citep{hw02}.  This formation channel is favored by some because most dark matter halos will form a Pop III star at this epoch if they reach masses of $\\sim$ 10$^5$ M$_ {\\odot}$ \\citep{abn02, bcl02}.  However, these BH have such low initial masses that they must continuously accrete at the Eddington limit to reach 10$^9$ M$_{\\odot}$ by $z \\sim 7$.  This is problematic for several reasons. First, numerical simulations have shown that Pop III stars usually evaporate the halos that give birth to them, so the BH are 'born starving' \\citep[e.g.][]{ wan04,ket04,wet08a}.  Filamentary inflows and mergers later restore baryons to the halo but only after 50 - 100 Myr \\citep{yet07}, during which crucial e-foldings in mass are lost.  Second, preliminary studies indicate that once accretion commences, the BH itself emits ionizing radiation that disperses its own fuel supply, limiting its growth rate to a fraction of the Eddington limit \\citep{ milos09,pm11,pm12} \\citep[but see][]{li11}. Furthermore, if the seed BH is not confined to the halo, its duty cycle as it meanders through cosmological density fields is intermittent, which also curtails its growth \\citep{awa09}.  Until now, 20 - 40 M$_{\\odot}$ Pop III BH \\citep{zwh08,wet08b} have been overlooked as candidates for SMBH seeds because previous studies assume that primordial stars are $\\gtrsim$ 100 M$_{\\odot}$.  However, there is a growing numerical and observational consensus that some Pop III stars are tens of solar masses, not hundreds.  More recent, much larger ensembles of numerical simulations found many halos with central collapse rates consistent with 20 - 60 M$_{\\odot}$ for the final mass of the star \\citep{on07} and that a fraction of the halos form binaries in this mass range \\citep{turk09}.  Furthermore, new simulations of the formation of Pop III protostellar accretion disks suggest that they were prone to fragmentation into as many as a dozen smaller stars \\citep{ stacy10,clark11,sm11,get11}.  Very preliminary calculations of I-front breakout from these disks indicate that ionizing UV radiation may terminate accretion onto the nascent star at $\\sim$ 40 M$_{\\odot}$ \\citep{hos11,stacy12}.  On the observational side, recent attempts to reconcile the nucleosynthetic yields of Pop III supernovae with the chemical abundances found in ancient, dim extremely metal-poor stars in the Galactic halo suggest that 15 - 40 M$_{\\odot}$ primordial stars may have been responsible for most of the heavy elements expelled into the primeval IGM \\citep{jet09b}.  The failure to detect the distinctive 'odd-even' nucleosynthetic signature of 140 - 260 M$_{\\odot}$ pair-instability SNe in metal-poor stars to date reinforces the fact that some Pop III stars might not be very massive, but this pattern may have been masked by selection effects in the observations \\citep{karl08}.  Low-mass Pop III BH are crucially different from more massive BH because they are born in supernova explosions rather than by direct collapse.  Asymmetries in the core-collapse engine can impart kicks of 200 - 1000 km/s to 20 - 40 M$_{ \\odot}$ BH, ejecting them from the halos that gave birth to them.  In this Letter we examine the implications of natal kicks for low-mass Pop III black holes as candidates for SMBH seeds.  In $\\S \\, 2$ we review the formation pathways for low-mass Pop III BH.  In $\\S \\, 3$ we calculate their post-supernova kinematics and retention fractions in halos.  In $\\S \\, 4$ we conclude. ",
        "conclusions": "The number of stars that 10$^5$ - 10$^7$ M$_{\\odot}$ halos typically form is not well constrained.  The studies of Pop III protostellar disk fragmentation performed thus far do not follow the evolution of the disk for enough dynamical times to determine the ultimate fate of the fragments, which may later merge with the central object or be destroyed by gravitational torques before becoming distinct stars.  Ionizing UV radiation from one star-forming fragment or even from a nearby halo can also prematurely halt the collapse of other fragments in the disk, lowering the number of stars that eventually form in the halo \\citep[e.g.][]{su06,wet08b,suh09,wet10}.  We also note that while the evolution of the fragments in the disk is expected to be roughly coeval, their 5 - 10 Myr quasistatic collapse times raise the possibility that the first star to form in the halo may explode and pre-empt the collapse of other fragments \\citep[e.g.][]{sak09}.  Consequently, the number of low-mass Pop III stars that occupy the halo likely ranges from one to at most ten.  Ejecta-driven natal kicks will evict most 20 - 32 M$_{\\odot}$ BH from their host halos, neutrino-driven kicks can drive more than 90\\% of 32 - 40 M$_{\\odot}$ BH from their halos, as we show in the right panel of Figure \\ref{fig:BH_MR}.  This guarantees that on average all the BH will vacate the halo even if ten stars originally formed in it. Post-supernova kinematics thus strongly discourages 20 - 40 M$_{\\odot}$ Pop III BH from becoming supermassive because they are ejected from their fuel supply and deprived of crucial early e-foldings in mass.  This process greatly reduces the parameter space in stellar mass from which SMBH can originate \\citep[e.g.][]{th09, lfh09}, especially if Pop III stars were mostly less than 50 M$_{\\odot}$. Also, if a given halo is capable of supporting early continuous Eddington rate accretion, a 20 - 40 M$ _{\\odot}$ BH is much less likely to become supermassive than a 100 M$_{\\odot}$ BH, either because it is ejected from the halo at birth or because it must undergo additional e-folding times to reach large masses.  If most low-mass Pop III black holes were ejected from their halos at $z \\sim$ 20, where are they today?  If on average they depart their host halos at $\\sim$ 500 km/s, they are unlikely to encounter another halo capable of capturing them in less than a Hubble time, and so many of these BH were exiled to the voids between galaxies. Over time, they may have gradually gained mass as they encountered high-density regions.  In contrast, Pop III BH above 40 M$_{\\odot}$ are unlikely to be born with kicks and remain in the halo, intermittently accreting and growing over cosmic time. These black holes are much more likely to reside in the galaxies into which their host halos were taken, a few of which may have become the supermassive black holes found in the SDSS quasars today."
    },
    "1112/1112.3147_arXiv.txt": {
        "abstract": "We analyze a high-resolution spectrum of the A3m star HD\\,27411.   We compare abundances derived from ATLAS9 model atmospheres with those using the more computationally-intensive ATLAS12 code.  We found very little differences in the abundances, suggesting that ATLAS9 can be used for moderate chemical peculiarity.  Our abundances agree well with the predictions of diffusion theory, though for some elements it was necessary to calculate line profiles in non-thermodynamic equilibrium to obtain agreement. We investigate the effective temperatures and luminosities of Am/Fm stars using synthetic Str\\\"{o}mgren indices derived from calculated spectra with the atmospheric abundances of HD\\,27411.  We find that the effective temperatures of Am/Fm stars derived from Str\\\"{o}mgren photometry are reliable, but the luminosities are probably too low.  Caution is required when deriving the reddening of these stars owing to line blanketing effects. A comparison of the relative proportions of pulsating and non-pulsating Am stars with $\\delta$~Scuti stars shows quite clearly that there is no significant decrease of helium in the driving zone, contrary to current models of diffusion. ",
        "introduction": "The ``metallic-lined'' or Am stars are A-type stars which have strong absorption lines of some metals such as Zn, Sr, Zr and Ba and weaker lines of other metals such as Ca and/or Sc relative to their spectral type as determined by the strength of the hydrogen lines \\citep{preston74}.  The strong metallic lines are more typical of an F star rather than an A star.  The work of \\citet{michaud70} established radiative diffusion in a strong magnetic field as the likely cause of the chemical peculiarities in Ap stars.  When the magnetic field is absent, diffusion leads to the Am/Fm stars \\citep{watson71}. The presence of magnetic fields in Am stars has been investigated, but with negative results, (e.g. \\citet{fossati07}). A peculiarity of Am stars is that their projected rotational velocities are generally much smaller than normal A stars and they are nearly always members of close binary systems.  Rotational braking by tidal friction in a binary system is regarded as a possible explanation for the low rotational velocities in Am stars. Slow rotation further assists the segregation of elements by diffusion.  The abundance anomalies predicted by the diffusion hypothesis are usually much larger than observed.  \\citet{richer00} developed detailed models of the structure and evolution of Am/Fm stars using OPAL opacities, taking into account atomic diffusion and the effect of radiative acceleration.  These models develop a convective zone due to ionization of iron-group elements at a temperature of approximately 200,000~K.  In addition to this convective zone, these stars also have a thin superficial convective zone in which H and He{\\sc i} are partially ionized.  By assuming sufficient overshoot due to turbulence, these separate convective zones become one large convective zone.  The resulting mixing dilutes the large abundance anomalies predicted by previous model, leading to abundances which closely resemble those observed in Am/Fm stars.  A detailed abundance analysis of eight Am stars belonging to the Praesepe cluster \\citep{fossati07} show good agreement with the predictions of \\citet{richer00} for almost all the common elements except for Na and possibly S.  The models of \\citet{richer00} assume a certain ad-hoc parametrization of turbulent transport coefficients which are adjusted to reproduce observations.  Other parameterizations of turbulence have been proposed for other types of stars.  \\cite{talon06} have investigated to what extent these are consistent with the anomalies observed on Am/Fm stars.  They find that the precision of current abundances is insufficient to distinguish between models.  More recently, \\citet{michaud11} have studied the abundance anomalies of the mild Am star Sirius A.  They find that except for B, N and Na, there is good agreement with the predicted anomalies but turbulent mixing or mass loss is required.  It is not clear whether it is turbulence or mass loss which competes with diffusion to lower the abundance anomalies.  For example, \\citet{vick11} find that diffusion in the presence of weak mass loss can explain the observed abundance anomalies of pre-main-sequence stars.  This is in contrast to turbulence models which do not allow for abundance anomalies to develop on the pre-main-sequence.   Most of the pulsational driving in $\\delta$~Scuti stars is caused by the $\\kappa$~mechanism operating in the He{\\sc ii} ionization zone.  Diffusion tends to drain He from this zone and therefore pulsational driving may be expected to be weaker or absent in Am/Fm stars \\citep{baglin72}. In fact, for many years it was thought that classical Am/Fm stars did not pulsate, though claims were made for some stars \\citep{kurtz89}.  Recently, intensive ground-based observations by SUPER-WASP \\citep{smalley11}, and also from  the {\\it Kepler} mission \\citep{balona11} have shown that many Am/Fm stars do pulsate.  \\citet{smalley11}, for example, found that about 200 Am/Fm stars out of a total of 1600 (12.5 percent) show $\\delta$~Sct pulsations, but with generally lower amplitudes. They found that the pulsating Am/Fm stars are confined between the red and blue radial fundamental edges, in agreement with \\citet{balona11}.  While there are many $\\delta$~Sct stars hotter than the fundamental blue edge, this does not seem to be the case for pulsating Am/Fm stars.  The significance of this result remains to be evaluated.  The effect of draining of He from the He{\\sc ii} ionization zone is to reduce the width of the instability strip, the blue edge moving towards the red edge, eventually leading to the disappearance of the instability strip when He is sufficiently depleted \\citep{cox79}.  \\citet{turcotte00} has discussed the effect of diffusion on pulsations in Am/Fm stars using the models by \\citet{richer00}.  One significant difference with earlier models is that a substantial amount of He remains in the He{\\sc ii} ionization zone. The blue edge of the instability strip for Am/Fm stars is sensitive to the magnitude of the abundance variations and is thus indicative of the depth of mixing by turbulence.  \\citet{turcotte00} predict that pulsating Am/Fm stars should lie in a confined region of the HR diagram close to the red edge of the $\\delta$~Sct instability strip.  However, \\citet{balona11} show that there is no relationship between the predicted Am/Fm instability strip and the actual location of these stars in the HR diagram.  A particularly interesting result of the pulsation analysis of \\citet{turcotte00} is the prediction of long-period g modes in A-type stars.  As the star evolves, the driving regions shift deeper into the star and the g modes become gradually more and more excited.  Whereas p modes are stabilized through diffusion, g modes tend to be excited as a result of that process. It appears that diffusion may act to enhance driving of long-period g modes due to a significant increase in opacity due to iron-group elements. This may have a bearing on the fact that nearly all A-type stars observed by {\\it Kepler} have unexplained low-frequencies \\citep{balona11a}.   When dealing with objects with non-standard chemical composition, such as Am stars, it is crucial that the opacities are correctly calculated.  This question has been investigated by several authors in recent years.  These studies show that a non-standard chemical composition of the stellar atmosphere alters the flux distribution of the star or modifies the profiles of the Balmer lines (\\citet{leo97}, \\citet{cat04}). Therefore a determination of T$_{\\rm eff}$ and $\\log g$ based on a comparison between observed and computed Balmer-line profiles will not be correct unless one takes into account the metallicity of the star. Thus, even estimates based on standard analysis of the spectra may be in error when applied to Am/Fm stars.  In this paper, we investigate the determination of effective temperature and surface gravity of the Am star HD\\,27411 (HR~1353, A3m) using spectra in the ESO archives.  The purpose is to determine whether the stellar parameters of this star agree with those obtained from Str\\\"{o}mgren photometry and hence to test the reliability of the effective temperature calibration applied to Am/Fm stars.  The star was used by \\citet{ryab08} as a comparison in their study on the calcium stratification in Ap stars.  HD\\,27411 is not known to pulsate. However, as we know from {\\it Kepler} observations, pulsations in A and F stars with amplitudes too low to be detected from the ground are common.  Atmospheric models obtained with ATLAS9 \\citep{kur93} use precomputed line opacities in the form of opacity distribution functions (ODFs).  These are tabulated for multiples of the solar metallicity and for various microturbulent velocities.  This approach allows very fast computation of model atmospheres, but with very little flexibility in choice of chemical profile and microturbulent velocity. While this is satisfactory for most applications, it fails for chemically peculiar stars where a non-standard chemical composition profile is required.  This can be done with ATLAS12 \\citep{kur97}, which is essentially identical to ATLAS9, but uses the opacity sampling (OS) method to evaluate line opacities.  In this study we compare the abundances of HD\\,27411 obtained with both codes to determine if the use of ATLAS12 is essential.  \\begin{figure*} \\centering \\includegraphics[width=8.8cm]{balmer1.ps} \\includegraphics[width=8.8cm]{balmer2.ps} \\caption{Comparison between the observed (crosses) and computed (solid red line) hydrogen line profiles. From top to bottom: the Balmer line profiles from H$_{\\delta}$ to H$_{\\alpha}$. The synthetic profiles were computed with SYNTHE using an ATLAS12 model atmosphere with T$_{\\rm eff}$\\,=\\,7400~$\\pm$~150~K, $\\log g$\\,=\\,4.0~$\\pm$~0.1, $\\xi$\\,=\\,4.2~$\\pm$~0.3~km~s$^{-1}$, $v_e \\sin i$\\,=\\,20.5\\,$\\pm$\\,0.5~km~s$^{-1}$ and individual abundances shown in Table~\\ref{abund}.} \\label{balmer} \\end{figure*}  The result that Am stars are not confined to particular region of the $\\delta$~Sct instability strip depends, to a large extent, on effective temperatures and luminosities estimated from  Str\\\"{o}mgren photometry \\citep{smalley11}.  It is not clear whether the calibration, derived from normal AF stars, can be applied to Am/Fm stars.  In this paper we use synthetic Str\\\"{o}mgren photometry applied to models of Am/Fm stars to investigate the reliability of fundamental parameters estimated from the photometry.  Finally, we discuss the relative numbers of pulsating and non-pulsating Am stars and compare these to the relative numbers of $\\delta$~Scuti and constant stars in the instability strip.  From this comparison, one can deduce the effectiveness of pulsational driving in the He{\\sc ii} ionization zone and compare the He abundance to that expected from diffusion calculations.  \\begin{table} \\centering \\caption{Ions used to determine the microturbulent velocity in HD\\,27411.  The number of spectral lines used, the microturbulent velocity, $\\xi$, the derived abundance and the radial velocity, RV, are listed.} \\begin{tabular}{lcccc} \\hline \\hline Elem & N & $\\xi$ & Abundance & RV \\\\ &   & km s$^{-1}$ & $\\log N_{\\rm el}/N_{\\rm Tot}$ & km s$^{-1}$ \\\\ \\hline Fe{\\sc i}  & 71 & 3.9\\,$\\pm$\\,0.3 & $-$3.80\\,$\\pm$\\,0.03 & 40.4\\,$\\pm$\\,0.6\\\\ Fe{\\sc ii} & 15 & 4.4\\,$\\pm$\\,0.2 & $-$3.84\\,$\\pm$\\,0.04 & 40.9\\,$\\pm$\\,0.6\\\\ Ni{\\sc i}  & 24 & 4.4\\,$\\pm$\\,0.6 & $-$4.63\\,$\\pm$\\,0.03 & 40.5\\,$\\pm$\\,1.3\\\\ \\hline Adopted   & --  & 4.2\\,$\\pm$\\,0.3 & -- & 40.6\\,$\\pm$\\,0.3\\\\ \\hline \\end{tabular} \\label{xi_vr} \\end{table} ",
        "conclusions": "We analyzed the spectrum of the Am star HD\\,27411 from the UVES POP archive. Our aim was to investigate if the ATLAS12 model atmosphere code provides more reliable results than the ATLAS9 code for chemically peculiar stars. We found that there is very little difference in the abundances derived from ATLAS9 and from ATLAS12.   Since ATLAS12 demands considerably greater resources, it seems safe to use ATLAS9, at least for moderate metallic enhancement. We find that the derived abundances in HD\\,27411 are in good agreement with the predictions of diffusion models by \\citet{richer00}.  There were discrepancies for Na, Al, Si, S, and Cr, but these are resolved by using NLTE model atmospheres.  We investigated the reliability of effective temperatures and luminosities of Am/Fm stars determined by Str\\\"{o}mgren photometry by synthesizing spectra having the abundances of HD\\,27411 for a range of effective temperatures.  The resulting synthetic colours indicate that effective temperatures can be reliably determined from photometry, but owing to line blanketing in the $c_1$ passband, the resulting surface gravities are systematically to high, leading to lower luminosities.  This result appears to be verified by comparing luminosities of Am/Fm stars obtained from their parallaxes and from photometry.  {\\bf Determination of reliable luminosities for Am stars remains a difficult problem.  At this stage, parallaxes offer the best results, but this can only be done for very few stars.  As we have seen, luminosities obtained from Str\\\"{o}mgren photometry are subject to a systematic bias which depends on the overabundances of metals.  The error in the surface gravity from high-resolution spectroscopy is typically 0.1 in $\\log g$.  For A--F main sequence and giants, this translates into an error of about 0.12 in $\\log L/L_\\odot$ when using the calibration of \\citet{torres10}.  For HD\\,27411, for example, we derive $\\log g = 4.0 \\pm 0.1$ from spectroscopy, whereas the value derived from the parallax is $\\log g = 3.8 \\pm 0.1$ (Table\\,2), Although the two values only differ by two standard deviations, this is enough to cause a difference of 0.36 in $\\log L/L_\\odot$.  Although spectroscopic determinations of luminosities may lead to quite large errors in the luminosity, they should at least not be biased.}  By far the most serious problem confronting the diffusion model is that there seems to be no appreciable settling of He in the He{\\sc ii} ionization zone, as predicted by the models. This is demonstrated by the fact that pulsating Am/Fm stars occur throughout the $\\delta$~Scuti instability strip, though they tend to be cooler than normal $\\delta$~Sct stars.  In fact, the relative proportions of pulsating Am stars to non-pulsating Am stars is no different from the proportion of $\\delta$~Sct stars to constant stars in the $\\delta$~Sct instability strip.   There is clearly a need to revise current ideas of diffusion to explain the Am phenomenon."
    },
    "1112/1112.1967_arXiv.txt": {
        "abstract": "\\baselineskip=16pt   We present the holographic predictions for cosmological 3-point correlators, involving both scalar and tensor modes, for a universe which started in a non-geometric holographic phase. Holographic formulae relate the cosmological 3-point functions to stress tensor correlation functions of a holographically dual three-dimensional non-gravitational QFT. We  compute these correlators at 1-loop order for a theory containing massless scalars, fermions and gauge fields, and present an extensive analysis of the constraints due to Ward identities showing that they uniquely determine the correlators up to a few constants. We define shapes for all cosmological bispectra and compare the holographic shapes to the slow-roll ones, finding that some are distinguishable while others, perhaps surprisingly, are not. ",
        "introduction": "In a recent series of papers \\cite{McFadden:2009fg,McFadden:2010na, McFadden:2010jw,McFadden:2010vh,Easther:2011wh,McFadden:2011kk} we put forward a holographic framework for inflationary cosmology and discussed a novel class of models describing a universe that started in a non-geometric phase, which is described holographically via a large-$N$ three-dimensional QFT. The power spectra and the scalar bispectrum were computed in \\cite{McFadden:2009fg,McFadden:2010na} and \\cite{McFadden:2010jw}, respectively, and in this paper we complete this program by computing the non-Gaussianities that involve tensors. Non-Gaussianities involving tensors are not expected to be measurable in near-future experiments. Nevertheless, they are still interesting theoretically and their structure has been the topic of several recent papers \\cite{Maldacena:2011nz,Soda:2011am,Shiraishi:2011st,Gao:2011vs}.  The holographic model is specified by providing the dual QFT and the holographic dictionary that relates QFT correlation functions to cosmological observables. We worked out the holographic dictionary for non-Gaussianities involving tensors in \\cite{McFadden:2011kk}, and in this paper we compute the relevant QFT correlation functions. The models we discuss are based on perturbative three-dimensional QFTs that admit a large-$N$ limit and have a generalised conformal structure \\cite{Jevicki:1998ub,Kanitscheider:2008kd}. An example of such a theory is $SU(N)$ Yang-Mills theory coupled to massless  scalars and fermions, with all fields transforming in the adjoint of $SU(N)$. The non-Gaussianities are extracted from the 3-point function of the stress tensor of this theory. The leading order 1-loop computation of this 3-point function is independent of the interactions of the QFT, and thus our main task is to compute this 3-point function for free QFTs.  Since all leading order results depend only on the free theory, let us briefly discuss the case in which the holographic model is a free QFT. In such a model, the spectrum is the exactly scale-invariant Harrison-Zel'dovich spectrum and the bispectrum is given exactly by the results reported here, {\\it i.e.}, the leading order results are the exact answer in the free theory. The shapes associated with the bispectrum may thus be considered as the analogue of the exact scale-invariant spectrum for higher point functions. We have seen in \\cite{McFadden:2010vh} that the scalar bispectrum shape for this model is indeed special: it is exactly equal to the factorisable equilateral shape\\footnote{ The holographic model is also the only model that yields exactly this shape.} originally introduced in \\cite{Creminelli:2005hu}. One may thus anticipate that shapes associated with the other 3-point functions will also have special properties. Possible shapes for the bispectrum involving only tensors have been discussed recently in \\cite{Maldacena:2011nz} and here we will define and discuss shapes for the bispectrum involving both tensors and scalars.  The computation of the 3-point function of the stress tensor at 1-loop is a non-trivial task even for free QFTs. We discuss and develop several methods for evaluating the relevant Feynman diagrams. The 1-loop result is constrained by Ward identities and these provide a very non-trivial check of the expression we obtained by a direct computation.  As mentioned above, to 1-loop order only the free part of the QFT enters. The QFT consists of gauge fields, fermions, minimal and conformal scalars. Conformal scalars and fermions are conformal field theories and their 3-point functions are constrained by conformal Ward identities. As is well known (from a position space analysis) \\cite{Osborn:1993cr}, the 3-point function of the stress tensor in $d=3$ is uniquely fixed by conformal invariance to be a linear combination of two conformal invariants, and is thus parametrised by two constants. (We assume parity is preserved). Our computation is done in momentum space and we thus provide the most general such 3-point functions in momentum space, where the two parameters are the number of conformal scalars and the number of fermions. The same computation was also recently reported in \\cite{Maldacena:2011nz}\\footnote{Our results agree with the ones in v2 of \\cite{Maldacena:2011nz}. Relative to \\cite{Maldacena:2011nz}, we also computed semi-local terms.}. We have explicitly verified that our results satisfy the conformal Ward identities.  Let us now turn to minimal scalars and gauge fields. In three dimensions vectors are dual to scalars, so one may expect that gauge fields  contribute the same as minimal scalars at 1-loop order. We will indeed verify that this is the case. Note that beyond 1-loop the two are expected to contribute differently. Minimal scalars differ from conformal scalars in the way they couple to gravity, which in flat spacetime is reflected in their having a different stress tensor. More precisely, the stress tensor $T^\\phi_{ij}$ for a minimal scalar may be decomposed into a part $\\wT^\\phi_{ij}$ corresponding to the stress tensor for a conformal scalar plus an ``improvement term'': \\[ \\label{minconfdecomp} T_{ij}^\\phi = \\wT_{ij}^\\phi - \\frac{1}{8} \\left( \\delta_{ij} \\partial^2 - \\partial_i \\partial_j \\right) \\phi^2. \\] It follows that the 3-point function of $T_{ij}^\\phi$ may be computed from the 3-point functions involving $\\wT_{ij}^\\phi$ and the dimension one operator ${\\cal O}_1= \\phi^2$. In turn, these 3-point functions are uniquely determined by conformal invariance, up to constants \\cite{Osborn:1993cr}.  Thus, effectively all 3-point functions are determined by conformal 3-point functions at this order, even though the underlying theory is not conformal.  In \\cite{Maldacena:2011nz}, the 3-point functions for tensors were computed in a de Sitter background.  The de Sitter isometries act as the conformal group at late times, and the 3-point functions are then constrained by conformal invariance to be specific linear combinations of the 3-point functions of conformal scalars and fermions. Note that the de Sitter result is the leading order approximation for slow-roll inflation. In general one expects that (broken) conformal invariance would constrain cosmological correlators in asymptotically de-Sitter slow-roll inflation, see also \\cite{Antoniadis:2011ib, Creminelli:2011mw}.  Our holographic results are for a very different universe, but we have seen that all relevant 3-point functions are essentially determined by conformal 3-point functions. One may then wonder how our results compare with those of slow-roll inflation. The 3-point functions involving only tensors are determined by the 3-point functions of conformal scalars and fermions, and, as in the discussion of \\cite{Maldacena:2011nz}, they agree exactly with slow-roll inflation if the field content of the dual QFT is appropriately chosen. The other 3-point functions (involving both scalar and tensor perturbations) are different but, perhaps surprisingly, they are rather similar. To quantify the differences we define (and plot) shape functions for all correlators, generalising the notion of shape functions for 3-point functions of only scalars or only tensors.  This paper is organised as follows. In Section \\ref{sec:dual} we discuss the dual QFT and in Section \\ref{sec:dict} we present the holographic dictionary. In Section \\ref{sec:eva} we compute all relevant QFT correlation function by direct evaluation of the relevant Feynman integrals, and in Section \\ref{sec:Wid} we explain their structure using Ward identities. The holographic predictions for the cosmological observables are presented in Section \\ref{sec:HolPredictions} and these results are compared with the slow-roll ones in Section \\ref{sec:slow-roll}. We discuss our results in Section \\ref{sec:disc}. Several technical results are presented in four appendices: in Appendix \\ref{sec:hel} we summarise our notation and conventions for the helicity tensors, in Appendix \\ref{methods_app} we present three different methods for evaluating the relevant diagrams, in Appendix \\ref{App_ghosts} we show that ghosts and gauge fixing terms do not contribute in correlators of the stress tensor and in Appendix \\ref{app_WI} we present the conformal and diffeomorphicm Ward identities. ",
        "conclusions": "\\label{sec:disc}  In this paper we computed the complete set of bispectra (and defined and extracted the corresponding shapes\\footnote{ To our knowledge, shapes other than those relevant for purely scalar or purely tensor bispectra have not been discussed before.}) for a class of holographic models of the very early universe % based on perturbative QFT. The leading 1-loop result actually depends only on the free part of the QFT, so in particular our results are also the complete answer when the dual QFT is free. The field content of the dual theory includes gauge fields, massless fermions, massless minimal and conformal scalars and thus the parameters that can appear in the results are the number of species for each type of field.  The bispectra could, a priori, depend on these in a complicated way, but it turns out that we get instead (nearly) universal results that are independent of all details of the dual QFT, within the class of the theories we consider. Thus, % these models make clean and precise predictions.  One can trace this universality to the specific form of the holographic map, the fact that to leading order the QFT is free, symmetry considerations and properties of $d=3$ theories. Let us explain this. Firstly, in three dimensions, vectors are dual to scalars so one may anticipate that the contribution due to gauge fields (at 1-loop order) is equal to that of the contribution due to mininal scalars, and we indeed find this to be the case. Taking this into account, the answer could then depend on three parameters, the number of conformal scalars, ${\\cal N}_\\chi$, the number of fermions, ${\\cal N}_\\psi$ and the total number of gauge fields plus minimal scalars, ${\\cal N}_{(B)}$.  The trace Ward identity of the dual QFT and the specific form of the holographic formulae then imply that, in all correlators involving at least one factor of $\\zeta$, the field content appears only as a multiplicative factor and is such that the corresponding shape functions are completely independent of the field content.  Let us now turn to correlators involving only tensors: these are effectively determined by the 3-point function of the stress tensor of a CFT. In three dimensions, this 3-point function is parametrised by two constants, which in our case are related to the field content. Indeed, the shape corresponding to three positive helicity gravitons does depend on the field content, but surprisingly the shape for two positive and one negative helicity graviton is independent of the field content. This can be explained in part by the fact that the $\\<T^{(+)} T^{(+)} T^{(-)}\\>$ correlator (at separated points) is actually uniquely fixed by conformal invariance up to a single constant. We emphasize however that this by itself is not sufficient to explain the independence of the corresponding shape function from the field content, as the specific form of the semi-local terms (both in the holographic map and in $\\<T^{(+)} T^{(+)} T^{(-)}\\>$) is crucial for this to happen.  Our calculations carefully include all such semi-local contributions. In the holographic formulae for the bispectra, these contributions appear as terms in the numerator that are non-analytic in only one of the three momenta. Since the denominator of the holographic formulae is however non-analytic in all three momenta, the net contribution  of these semi-local terms to the bispectra is in fact non-analytic in {\\it two} of the three momenta. Semi-local terms in the holographic formulae may thus contribute, for example, to `local'-type non-Gaussianity behaving as $1/q_1^3 q_2^{3} + \\mathrm{perms}$. Contributions of this nature therefore play a crucial role in allowing different cosmological shapes to be distinguished.  To get a feeling for our results we also computed the corresponding slow-roll results and compared them with the holographic results. Firstly, comparing the power spectra one obtains a relation between the parameters  $N^2$, ${\\cal N}_{(A)}$ and ${\\cal N}_{(B)}$ of the QFT and the parameters  $\\kappa^2$, $H_*^2$ and $\\epsilon_*$ of the slow-roll model. Comparing the 3-point functions, we find that the $\\bg^{(+)}\\bg^{(+)}\\bg^{(-)}$ correlators agree exactly, while the $\\bg^{(+)}\\bg^{(+)}\\bg^{(+)}$ correlators can be made to agree if one imposes that that the field content satisfies the relation $2\\mathcal{N}_\\psi = \\mathcal{N}_\\phi+\\mathcal{N}_A+\\mathcal{N}_\\chi$. As explained in \\cite{Maldacena:2011nz}, these slow-roll correlators are constrained by the late-time de Sitter isometries to satisfy conformal Ward identities, and thus at separated points they should be expressible in terms of the 3-point  functions of conformal scalars and free fermions. Indeed, the linear combination found in v2 of \\cite{Maldacena:2011nz} is the same as the one we find (setting $\\mathcal{N}_\\phi=\\mathcal{N}_A=0$ in our relation). By taking into account the contribution from semi-local terms, however, we are further able to correctly recover every individual term appearing in the graviton bispectra.    There is no apparent reason for the remaining slow-roll and holographic correlators to agree. Nevertheless we find rather similar results. To quantify the difference we used the cosine orthogonality measure of \\cite{Fergusson:2008ra} to obtain a first indication of the distinguishability of the corresponding shapes. We find that the shapes for $\\z \\z \\bg^{(s)}$ are nearly indistinguishable, while for $\\z \\bg^{(+)} \\bg^{(+)}$, the two shapes may be distinguished (as a consequence of differing behaviour in the squeezed limit where the momentum associated with the $\\z$ goes to zero), with the case of $\\z \\bg^{(+)} \\bg^{(-)}$ lying in between.  All in all, we have a rather complete understanding of this class of models and their phenomenology. There are still a few things to be understood better: what constrains the semi-local contributions to the tensor correlators, and why are the holographic results apparently close to slow-roll ones? One can presumably also understand the squeezed limit of the correlators using Ward identities. On a whole, however, the structure of these models is reasonably firmly understood. It would be interesting to arrive at a similar level of understanding for the class holographic models that are based on deformations of conformal field theories."
    },
    "1112/1112.4290_arXiv.txt": {
        "abstract": "{The physical origin of the strong magnetic activity in T Tauri stars and its relation to stellar rotation is not yet well-understood. } {We investigate the relation between the X-ray activity, rotation, and Rossby number for a sample of young stars in the $\\approx 3$~Myr old cluster IC~348. } {We use the data of four \\textit{Chandra} observations of IC~348 to derive the X-ray luminosities of the young stars. Basic stellar parameters and rotation rates are collected from the literature. This results in a sample of 82 X-ray detected stars with known rotation periods. We determine the Rossby numbers (i.e.~the ratio of rotation period to convective turnover time) of 76 of these stars from stellar structure- and evolution-models for pre-main sequence stars. } {The young stars in IC~348 show no correlation between X-ray activity and rotation period. For the Rossby numbers, nearly all IC~348 stars are in the saturated regime of the activity--rotation relation defined by main-sequence stars. Searching for possible super-saturation effects, we find a marginal (but statistically in-significant) trend that the stars with the smallest Rossby numbers have slightly lower X-ray activity levels. There are no significant differences in the X-ray activity level for stars of different spectral types and no relation between spectral type and Rossby number is seen. In addition, for stars belonging to different IR-classes, no significant differences  are present for the X-ray activity level as well as for their Rossby numbers. We compare the dispersion in the fractional X-ray luminosities of the stars in the saturated rotation regime in IC~348 to that seen in younger and older stellar populations. The scatter seen in the  $\\approx 3$~Myr old IC~348 [$\\sigma\\left(\\log\\,(L_{\\rm X}/L_{\\rm bol}) \\right) = 0.43$] is considerably smaller than for the $\\approx 1$~Myr old Orion Nebula Cluster [\\,$\\sigma\\left(\\log (L_{\\rm X}/L_{\\rm bol}) \\right) = 0.63$], but, at the same time, considerably larger than the dispersion seen in the $\\approx 30$~Myr old cluster NGC~2547 [$\\sigma\\left(\\log (L_{\\rm X}/L_{\\rm bol}) \\right) = 0.24$] and in main-sequence stars. } { The results of our X-ray analysis of IC~348 show that neither the rotation rates nor the presence/absence of circumstellar disks are of fundamental importance for determining the level of X-ray activity in TTS. Our results suggest that the scatter in the X-ray activity levels for the rapidly rotating members of young clusters decreases with the age of the  stellar population. We interpret this as a signature of the changing interior structure of pre-main sequence stars and the consequent changes in the dynamo mechanisms that are responsible for the magnetic field generation.} ",
        "introduction": "Young stellar objects (YSOs) in all evolutionary stages, from class~I protostars, to T Tauri stars (TTS) to zero-age main-sequence stars, have highly ($\\sim 10^3 - 10^4$ times) elevated levels of X-ray activity (for reviews of the X-ray properties of YSOs and stellar coronal astronomy in general, see Feigelson \\& Montmerle 1999 and Favata \\& Micela 2003). Despite many years of research, the physical origin of this X-ray activity remains poorly understood. Although there is strong evidence that in most TTS the X-ray emission originates from magnetically confined coronal plasma \\citep[e.g.][]{Preibisch05}, it is unclear what kind of dynamo processes create the required magnetic fields.    For main-sequence stars, the level of the magnetic activity (and thus the strength of the X-ray emission) is mainly determined by their rotation rate. Observations have revealed a clear rotation--activity relation of the form  $L_{\\rm X}/L_{\\rm bol} \\propto P_{\\rm rot}^{-2.6}$ \\citep[e.g.][]{Pallavicini81,Pizzolato03}, which is in good agreement with the expectations of solar-like $\\alpha\\!-\\!\\Omega$ dynamo models \\citep[e.g.][]{Maggio87,Ossendrijver03}. The solar dynamo is thought to be anchored in the ``tachocline'', a thin zone between the inner radiative core and the outer convection zone. However, the connection between the observed surface magnetic activity and the properties of the dynamo generating the magnetic flux is not yet fully understood \\citep[e.g.,][]{Isik11}.  For main-sequence stars, the increase in magnetic activity towards shorter rotation periods stops for periods shorter than $\\sim 2-3$ days, where the activity saturates around $\\log\\left(L_{\\rm X}/L_{\\rm bol}\\right) \\approx -3$. The physical reasons for this saturation effect are poorly understood \\citep[see, e.g.,][]{Jardine99}. For extremely rapidly rotating stars, the activity level seems to decrease slightly with increasing rotation rate \\citep{Prosser96,Randich00,James00,Jeffries11}, a phenomenon that is denoted as ``super-saturation'' and also not well-understood.      Studies of stellar populations with different ages show that that there is a continuous evolution from the very high X-ray activity levels in the youngest stages to the much lower activity seen in older (more than a few hundred Myr old) stars \\citep[e.g.,][]{Guedel97,PF05,Telleschi05}. This evolution can be explained by the temporal decrease in the stellar angular momentum \\citep{Bouvier97,Herbst07}. Furthermore, the presence of strong magnetic fields on the surface of T Tauri stars has been clearly established \\citep[e.g.][]{Johns-Krull07}. These pieces of evidence suggest that the X-ray activity of YSOs originates in dynamo processes similar to those present in our Sun.  However, the relation between rotation and X-ray activity in TTS remained unclear until recently, since in most studies of star forming regions the number of X-ray detected TTS with known rotation periods was too small to draw sound conclusions.  A few years ago, the $Chandra$ Orion Ultradeep Project (COUP), a ten-day long observation of the Orion Nebula Cluster (ONC) with $Chandra$/ACIS \\citep[for details of the observation and data analysis, see][]{Getman05} and the XMM-Newton Extended Survey of the Taurus Molecular Cloud \\citep[XEST, see][]{Guedel07} provided very sensitive X-ray data sets for large samples of TTS. The COUP and the XEST data have both shown that the TTS in Orion and Taurus do \\textit{not} follow the relation between rotation period and X-ray luminosity for main-sequence stars \\citep{Preibisch05,Briggs07}. In addition the TTS spin up during the first $\\sim$10-30 Myr \\citep{Herbst05} does not lead to an increase in the X-ray luminosity. This places doubt on the solar-like dynamo activity scenario for TTS. Another argument against solar-like dynamos in young TTS comes from theoretical considerations: at ages of  $\\le 2$~Myr, most TTS are expected to be fully convective, and thus should not possess a tachocline. The conventional $\\alpha\\!-\\!\\Omega$ dynamo cannot work in such a situation. Several alternative dynamo concepts have been developed for fully convective stars \\citep[e.g.][]{Giampapa96,Kueker99,Dobler06,Chabrier06,Browning07,Voegler07,Graham10}. Although the reliability of these theoretical models is not entirely clear, there is good evidence for the simultaneous presence of an $\\alpha\\!-\\!\\Omega$ dynamo \\textit{and} some kind of a small-scale turbulent dynamo in the convection zone of our Sun \\citep[e.g.,][]{Durney93,Bueno04}.    Direct observations of the relation between rotation and magnetic activity for stellar samples spanning a wide range of ages can provide fundamental constraints on the theoretical models. Numerous large datasets are available for samples of main-sequence stars as well as for young stellar clusters with ages as young as $30$~Myr \\citep[e.g.][]{Prosser96,Stauffer97a,Stauffer97b,Randich00,Jeffries11}. For younger ages, however, there is still a clear lack of reasonably large stellar samples for which good activity \\textit{and} rotation data are available. The data obtained in the COUP and XEST observations do both indicate that there are very young stellar populations of only $\\la 1$~Myr old \\citep[see, e.g.,][]{Weights09,Dib10,Luhman10}. Since the stellar rotation, the magnetic activity levels, and other basic stellar parameters evolve strongly in the age range between 1~Myr and 30~Myr, a sample of stars with an age of a few Myr can provide valuable information. With an age of $\\approx 3$~Myr \\citep{Luhman03,Mayne07}, the young cluster IC~348 is very well-suited in this respect. This age is particularly interesting because it corresponds to the point in time when the structure of solar-mass stars changes from a fully convective interior to a radiative core plus convective envelope structure, and this should affect the dynamo processes that are the ultimate source of the magnetic activity. ",
        "conclusions": "Our analysis of the relation between X-ray activity (as traced by deep \\textit{Chandra} observations) and the rotation properties of the TTS in the young cluster IC~348 has yielded the following results. First, we have shown that there is no correlation between the fractional X-ray luminosity and the rotation periods of the TTS; even the rather slowly rotating stars (periods $\\ga 10$~days) display very strong X-ray activity. Second, according to the Rossby numbers (that we have determined based on detailed stellar models for pre-main sequence stars), all TTS are in the saturated regime of the rotation--activity relation defined by main-sequence stars. Third, we have found no significant evidence of super-saturation among the most rapid rotators, although our data suggest that stars with extremely low Rossby numbers ($Ro \\leq 0.006$) have slightly lower activity levels. This seems to agree with the results of \\citet{Jeffries11}, who claim that in their NGC~2547 sample supersaturation occurs only for Rossby numbers lower than 0.005. In all three of these aspects, the TTS in IC~348 behave similarly to the TTS in the Orion Nebula Cluster.  A remarkable difference is however seen in the scatter in the X-ray activity levels for the TTS in the saturated rotation regime. The scatter seen in the $\\approx 3$~Myr old IC~348 sample is considerably smaller than that in the $\\approx 1$~Myr old Orion Nebula Cluster sample, but, at the same time, considerably larger than seen in the $\\approx 30$~Myr old stars in NGC~2547 as well as for main-sequence stars. This suggests that some process reduces the wide distribution of activity levels seen in the youngest stars towards the much narrower distribution in older pre-main sequence and (young) main sequence stars during the first $\\sim 30$~Myr period, whereas the absolute level of the X-ray activity remains roughly constant during that time.  A possible explanation of this effect may be related to the stellar interior structure and the corresponding dynamo mechanisms that are the basis of the magnetic activity that produces the observed X-ray emission. At an age of $\\leq 1$~Myr (e.g., in the Orion Nebula Cluster), almost all low- and intermediate mass stars ($M \\leq 2\\,M_\\odot$) are fully convective. As already mentioned, the solar-like $\\alpha\\!-\\!\\Omega$ dynamo cannot operate in such a situation and some kind of small-scale convective dynamo must be the source of the magnetic activity. The more or less chaotic nature of such a dynamo may be responsible for the very wide scatter in activity levels seen for stars with similar rotation rates and stellar parameters. According to the pre-main sequence stellar evolution models of \\citet{Siess00}, a $1.5\\,M_\\odot$ star develops a radiative core at an age of about 1~Myr. As time proceeds, stars with increasingly lower masses develop radiative cores (at $\\approx 1.7$~Myr for $1.2\\,M_\\odot$ stars, at $\\approx 2.3$~Myr for $1.0\\,M_\\odot$ stars, and at $\\approx 4.2$~Myr for $0.8\\,M_\\odot$ stars). As soon as a star has a radiative core, the conditions for the operation of a solar-like $\\alpha\\!-\\!\\Omega$ dynamo are met, which then gradually replaces the convective dynamo. When the total magnetic activity is dominated by the $\\alpha\\!-\\!\\Omega$ dynamo, stars in the saturated rotation regime have a rather homogeneous level of X-ray activity, i.e.~a small scatter in their $L_{\\rm X}/L_{\\rm bol}$ values.  At the 3~Myr age of IC~348, all stars with $M \\ge 0.9\\,M_\\odot$ (corresponding to a spectral type of K6) should have a radiative core; this concerns about 25\\% of our stellar sample in IC~348 and can explain why the dispersion in the X-ray activity levels in IC~348 is smaller than in the Orion Nebula Cluster. As time proceeds, an increasing fraction of the lower-mass stars make the transition from a fully convective to a core/envelope structure, where the operation of an $\\alpha\\!-\\!\\Omega$ dynamo becomes possible, and this should continuously decrease the scatter in their $L_{\\rm X}/L_{\\rm bol}$ values. By the age of $\\approx 30$~Myr, nearly all stars have attained their final stellar structure, and thus the scatter in the activity levels has settled to the rather small dispersion as typical of (rapidly rotating) main-sequence stars. This scenario could be a qualitative explanation of the observed decrease in the scatter in the X-ray activity level of stars in the saturated rotation regime."
    },
    "1112/1112.6406_arXiv.txt": {
        "abstract": "One key goal of the {\\it Hubble Space Telescope} Cosmic Assembly Near-Infrared Deep Extragalactic Legacy Survey is to track galaxy evolution back to $z\\approx 8$. Its two-tiered ``wide and deep'' strategy bridges significant gaps in existing near-infrared surveys. Here we report on $z\\approx 8$ galaxy candidates selected as F105W-band dropouts in one of its deep fields, which covers 50.1~arcmin$^2$ to 4~ks depth in each of three near-infrared bands in the Great Observatories Origins Deep Survey southern field. Two of our candidates have $J<26.2$~mag, and are $> 1$~mag brighter than any previously known F105W-dropouts. We derive constraints on the bright-end of the rest-frame ultraviolet luminosity function of galaxies at $z\\approx 8$, and show that the number density of such very bright objects is higher than expected from the previous Schechter luminosity function estimates at this redshift. Another two candidates are securely detected in {\\it Spitzer} Infrared Array Camera images, which are the first such individual detections at $z\\approx 8$. Their derived stellar masses are on the order of a few $\\times 10^9$~$M_\\odot$, from which we obtain the first measurement of the high-mass end of the galaxy stellar mass function at $z\\approx 8$. The high number density of very luminous and very massive galaxies at $z\\approx 8$, if real, could imply a large stellar-to-halo mass ratio and an efficient conversion of baryons to stars at such an early time. ",
        "introduction": "In the recent years, deep near-infrared imaging surveys have begun to yield a significant number of candidate galaxies at very high redshifts. The wide-field surveys from the ground have produced a handful of bright candidates at $z\\approx 7$ (e.g., Ouchi et al. 2009; Hickey et al. 2010; Castellano et al. 2010; Capak et al. 2011; Hsieh et al. 2012; Hathi et al. 2012), while the pencil-beam survey by the {\\it Hubble Space Telescope} (HST) Wide Field Camera 3 (WFC3) within the historical Advanced Camera for Surveys (ACS) Hubble Ultra Deep Field (hereafter ACS-HUDF; Beckwith et al. 2006) and its two parallel fields (HUDF09; PI. Illingworth) have probed the fainter populations at $z\\approx 7$--8 (Oesch et al. 2010; Bouwens et al. 2010, B10; Bunker et al. 2010; McLure et al. 2010, M10; Yan et al. 2010, Y10; Finkelstein et al. 2010) and possibly out to $z\\approx 10$ (Yan et al. 2010; Wyithe et al. 2011; Bouwens et al. 2011a). The WFC3 Early Release Science program (ERS, PI. O'Connell; Windhorst et al. 2011) has played an important role in connecting the ``wide'' and the ``deep'' ends of the exploration, allowing construction of $z\\approx 7$--8 samples at intermediate brightness levels (Wilkins et al. 2011; B11; Lorenzoni et al. 2011, L11). Meanwhile, the Hubble Infrared Pure Parallel Imaging Extragalactic Survey (HIPPIES; Yan et al. 2011) and the Brightest of Reionizing Galaxies Survey (Trenti et al. 2011a) have been exploring the bright end of the population through a large number of random, discrete WFC3 pointings obtained during the {\\it HST} parallel orbits. Surveys in foreground cluster fields utilizing gravitational lensing magnification have also resulted in a handful of promising candidates at $z\\gtrsim 7$--9 (e.g., Richard et al. 2006; Bradley et al. 2008, 2012; Laporte et al. 2011; Hall et al. 2012), which complement the surveys in blank-sky fields.  The Cosmic Assembly Near-Infrared Deep Extragalactic Legacy Survey (CANDELS; PIs: Faber \\& Ferguson; see Grogin et al. 2011 and Koekemoer et al. 2011) employs a two-tiered, ``wide and deep'' strategy, which makes it uniquely positioned in bridging the significant gaps among existing surveys. In particular, its ``Deep'' component will cover $\\sim 125$~arcmin$^2$ in two fields upon completion, and its data will be ideally suited for studying galaxies at very high redshifts. In this paper, we report our preliminary results from the Deep observations in the Great Observatories Origins Deep Survey (GOODS; Giavalisco et al. 2004) southern field, where we use the nearly complete data set to study the galaxy population at $z\\approx 8$. We use the following cosmological parameters throughout: $\\Omega_M=0.27$, $\\Omega_\\Lambda=0.73$ and $H_0=71$~km~s$^{-1}$~Mpc$^{-1}$. The quoted magnitudes are all in the AB system. ",
        "conclusions": "Our samples are unique in two aspects. (1) We have two very bright candidates at $J_{125}\\leq 26.0$~mag, which are at least $\\sim 1$~mag brighter than the brightest $Y_{105}$-dropouts previously found. They are comparable in brightness to the brightest $Y_{098}$-dropouts reported by Yan et al. (2011) and Trenti et al. (2011a,b) in their WFC3 parallel surveys, but could be at higher redshifts because the redshift window of this current $Y_{105}$-dropout selection is higher than that of the existing $Y_{098}$-dropouts. (2) Two other $Y_{105}$-dropouts, which are not among the brightest in the WFC3 bands,  are securely detected in IRAC, which is in sharp contrast to previous results where only non-detection in IRAC have been reported for $Y_{105}$-dropouts.  Here we discuss the implications of both.  \\subsection{Stellar Population and Stellar Mass}  We investigate the stellar populations of the IRAC-detected $Y_{105}$-dropouts, {\\tt AUTO\\_035} and {\\tt AUTO\\_293}, by analyzing their spectral energy distributions (SED) through template fitting. The approach here is similar to that of Finkelstein et al. (2011). We first use the EAZY code (Brammer et al. 2008) and the photometry of the GOODS ACS \\& CANDELS WFC3 data to estimate photometric redshifts ($z_{phot}$). The full SED (including the IRAC photometry) is then fitted to a suite of templates based on the updated models of BC03 (the so-called ``CB07'' models). A Salpeter initial mass function (Salpeter 1955) and the metallicities of 0.02--1$Z_\\odot$ are adopted. The templates have a range of exponentially decreasing and increasing SFHs (see also Papovich et al. 2011), and can include nebular lines based on the number of ionizing photons and metallicity of a given model (B. Salmon et al. in preparation). We assume the extinction law of Calzetti (2001) with $E(B-V)$$=$0--0.5~mag and the H I absorption as formulated in Madau (1995). The results are summarized in Figure 4. We note that fitting redshift and other properties simultaneously does not signficantly change these results, and the differences are captured in the errors that we quote here.  The best-fit photometric redshifts are $8.5^{+0.2}_{-4.5}$ and $8.9^{+0.2}_{-0.5}$ for {\\tt AUTO\\_035} and {\\tt 293}, respectively, which are consistent with the redshift window of our color selection. Adding the formal $I_{814}$ limit to the fitting process does not change the results.  We fit the SEDs with and without the contribution from nebular lines, and both results suggest high stellar mass ($\\mathcal M_*$) in the range of $10^{9}$ to $10^{10} M_\\odot$ for both objects. Including nebular emission lines, we obtain $2.5^{+4.7}_{-1.4}\\times 10^{9}$ and $0.9^{+3.0}_{-0.1}\\times 10^{9}$~$M_\\odot$ for {\\tt AUTO\\_035} and {\\tt 293} respectively, while without nebular emission lines these are $10.5^{+1.8}_{-8.1}\\times 10^{9}$ and $2.6^{+1.8}_{-1.4}\\times 10^{9}$~$M_\\odot$, respectively. The error bars reflect the 68\\% confidence level of the fit. Labb\\'e et al. (2010a) derived the stellar masses of galaxies at lower redshifts of $z\\approx 7$ using models without nebular emission, and the average value is $\\sim 10^9$~$M_\\odot$. The two objects in our sample have comparable or even larger values using similar models, which are surprisingly high at such an early epoch.  Using the $\\mathcal M_*$ estimates, we obtain the first measurement of the mass function (MF) of galaxies at $z\\approx 8$ at the high-mass end. As we only have two objects whose $\\mathcal M_*$ values are reasonably similar, we opt to count them within one mass bin of a $\\pm 0.4$~dex bin size. The two different sets of models result in significantly different mass estimates, and therefore the bin center is different for each case. Assuming a top-hat selection function of our survey within $7.7\\leq z\\leq 8.7$, we get $\\phi(log\\frac{M_*}{M_\\odot})|_{9.2\\pm 0.4} = (1.9^{+2.4}_{-0.6})\\times 10^{-5}$~Mpc$^{-3}$~dex$^{-1}$ if we adopt the values derived using the models with nebular emission. If we apply the effective volume ($V_{eff}$) correction of the survey instead (see \\S 3.4), we obtain $(3.1^{+4.0}_{-1.0})\\times 10^{-5}$~Mpc$^{-3}$~dex$^{-1}$. Here we use $V_{eff}=\\int V_{eff}(m_J)dm_J=\\int\\int dm_JdzP(m_J,z)dV/dz$, where $dzdV/dz$ is the unit co-moving volume at redshift $z$, and $P(m_J,z)$ is the redshift selection function at different magnitudes $m_J$ as derived through simulations (see Figure 3). The error bars here reflect the 68\\% interval of the uncertainties caused by the Poisson noise in the sample. For the case of using the $\\mathcal M_*$ estimates based on the models without nebular lines, all the above values are applicable at $\\phi(log\\frac{M_*}{M_\\odot})|_{9.7\\pm 0.4}$.  The contribution of these objects to the global stellar mass density, $\\rho_{*}$, at $z\\approx 8$ within our survey volume is $(3.3^{+7.4}_{-1.4}) \\times 10^{4}$ and $(12.6^{+3.4}_{-9.0}) \\times 10^{4}$~$M_\\odot$~Mpc$^{-3}$ based on the models with and without nebular emission lines, respectively, if assuming a top-hat selection function at $7.7\\leq z \\leq 8.7$. The errors reflect the 68\\% confidence level of the fitting results. If we apply the same correction for $V_{eff}$ as above, we obtain $\\rho_{*}=(5.5^{+12.6}_{-2.4}) \\times 10^{4}$ and $(21.2^{+5.8}_{-15.3}) \\times 10^{4}$~$M_\\odot$~Mpc$^{-3}$.  We derive the mass-to-light ratio of these two objects using their 4.5~$\\mu$m flux as a proxy to $L_V$. Using the fit results with contribution from nebular lines, we obtain $\\mathcal M_*/L_V=0.07^{+0.14}_{-0.04}$ and $0.04^{+0.13}_{-0.00}$ for {\\tt AUTO\\_035} and {\\tt 293}, respectively. Without the contribution from nebular lines, we obtain $\\mathcal M_*/L_V=0.31^{+0.06}_{-0.24}$ and $0.12^{+0.08}_{-0.06}$. The latter values are in general agreement with those obtained by Labb\\'e et al. (2010a,b) through stacking analysis of the IRAC data of the $Y_{098}$-dropouts and the $Y_{105}$-dropouts in the WFC3 ERS field (see also Gonz\\'alez et al. 2011) and the HUDF using similar models without nebular lines. On the other hand, Labb\\'e et al. (2010b) estimated that nebular lines have a small impact on $\\mathcal M_*/L_V$, reducing it by $\\sim 0.2$~dex. Our results suggest a stronger effect, reducing $\\mathcal M_*/L_V$ by $\\sim 0.5$--0.6 dex.  \\begin{figure}[tbp] \\centering \\includegraphics[width=\\linewidth]{fig4.eps} \\caption{Summary of SED fitting results for {\\tt AUTO\\_035} (left) and {\\tt 293} (right). The results obtained from the fit with and without the contributions from nebular lines are coded in blue and red, respectively. The top panels show the observed SEDs (grey circles) and the best-fit templates, and the insets display their 3.6 and 4.5~$\\mu$m image ($18\\farcs 6 \\times 18\\farcs 6$). The bottom panels show the likelihood functions of $z_{phot}$ and $M_*$, and the best-fit values are indicated by the vertical lines. } \\end{figure}  \\subsection{Constraint on the LF}  Here we use our sample to constrain the LF at $z\\approx 8$. We adopt the $J_{125}$ {\\tt MAG\\_AUTO} values of the candidates as their total magnitudes.  After applying the correction for $V_{eff}$, we obtain the stepwise LF in the five 0.5-mag bins (25.45, 25.95, 26.45, 26.95, 27.45)~mag as $\\phi = (2.5^{+5.8}_{-0.6}, 2.5^{+5.8}_{-0.6}, 2.8^{+6.3}_{-0.6}, 3.6^{+8.3}_{-0.8}, 24.2^{+22.6}_{-5.6})\\times 10^{-5}$~Mpc$^{-3}$~mag$^{-1}$. These number densities are shown in Figure 5, and compared to the predictions from a number of Schechter LF estimates at $z\\approx 8$ (M10; Y10; B11; L11). The black squares are our observed densities, while the red squares are the densities after correcting for $V_{eff}$. The error bars represent a Bayesian 68\\% credible interval, indicating the central 68 percentile range for the posterior distribution of the true number density assuming Poisson statistics. These uncertainties account for the incompleteness, but do not account for any systematic uncertainties from the contamination due to low-z interlopers. For ease of direct comparison to observations, surface density and apparent magnitude scales are also provided on the same figures to present these results in terms of differential number densities versus apparent $J_{125}$ magnitudes. For comparision, we also plot the {\\it observed} densities (i.e., before corrections for their corresponding $V_{eff}$ values) extracted from the sample of Y10 in the HUDF09 proper and those of B11 in the HUDF09, HUDF09P1 and HUDF09P2 {\\it after} applying the additional criteria of $Y_{105}-J_{125}\\geq 0.8$~mag and $J_{125}-H_{160}\\leq 0.3$~mag.  \\begin{figure}[tbp] \\centering \\includegraphics[width=\\linewidth]{fig5.eps} \\caption{Constraints on the very bright-end of the LF at $z\\approx 8$ based on the sample of $Y_{105}$-dropouts in CANDELS/Deep from this current work. The black squares are the differential surface densities inferred from the number counts in our sample, while the red squares are those after correction for $V_{eff}$ with respect to the volume within $7.7\\leq z\\leq 8.7$. The count predictions over $7.7\\leq z\\leq 8.7$ are from the Schechter LF estimates at $z\\approx 8$ from B11, M10, Y10, and L11 are shown as various black curves, together with a non-evolution one at $z\\approx 7$ from B11 as the blue curve. For comparison, different symbols at the faint-end show the raw (observed) surface densities of $Y_{105}$-dropouts based on the following samples: Y10 (HUDF -- stars), B11 (HUDF -- crosses; HUDFP1 -- pentagons; HUDFP2 -- triangles). } \\end{figure}  An intriguing feature that our samples reveal is that there could be an excess of bright objects at $J_{125}\\lesssim 26.2$~mag with respect to any of the existing Schechter LF estimates at $z\\approx 8$ in the literature would predict. The inferred number density at this bright end is even higher than the prediction from a non-evolution $z\\approx 7$ LF (e.g., the one from B11). This excess is still evident with the observed counts {\\it before} applying the incompleteness correction. The observed counts at $J_{125}\\lesssim 26.2$~mag in Figure 5 comes from objects {\\tt AUTO\\_048} and {\\tt AUTO\\_100}. Based on their photometry in $J_{125}$ and assuming $z=8$, they have $M_{UV}$ of $-21.54$ and $-21.21$, respectively. While such high luminosities have been observed in the spectroscopically confirmed $z\\approx 7$ galaxies of Ono et al. (2012), they are 5--9$\\times$ more luminous than any previously known $Y_{105}$-dropouts. Such an excess is unlikely due to cosmic variance, whose effect is expected to be small compared to the Poisson errors for our observed number density if we use the formalism of Trenti \\& Stiavelli (2008).  We note that, however, that O12 do not see such an excess, even though their sample also includes our brightest candidate {\\tt AUTO\\_048}. The main reason for this discrepancy could be that O12 use a larger volume in the calculation. They have included part of the CANDELS Wide field in the GOODS-S region, which is a major contributor to their larger volume. However, this area does not have sufficiently deep $Y_{105}$ data to apply our selection criterion of $Y_{105}-J_{125}\\geq 0.80$~mag down to $J_{125}\\approx 26.0$~mag. O12 have adopted a less stringent criterion of $Y_{105}-J_{125}\\geq 0.45$~mag, which allows them to construct a $Y_{105}$-dropout sample in this area at the cost of including candidates at lower redshifts ($z\\approx 7.2$). The advantage of their approach is that they can incorporate the WFC3 ERS field in the GOODS-S region as well, which further increases their survey volume. While the WFC3 ERS field has $Y_{098}$ but not $Y_{105}$ data, the redshift range probed by their $Y_{098}$-dropouts in this area overlaps that of their $Y_{105}$-dropouts elsewhere, and therefore O12 has combined the $Y_{098}$- and the $Y_{105}$-dropouts to address the LF averaged over a wide redshift range ($z\\approx 7$--8.7).  In contrast, our color criteria select galaxies in a higher range of redshifts, $7.7\\lesssim z\\lesssim 8.7$, and the bright-end excess (with respect to a smooth Schechter function) that we find suggests that our understanding of the LF at $z\\approx 8$ is still far from being complete. The significance of this excess is still subject to small number statistics, and we cannot rule out the possibility that one or both of our brightest objects could be interlopers at low redshifts. However, we should point out that similar bright-end excess at $z\\sim 8$ has also been suggested in other studies, for example in the bright $Y_{098}$-dropout study in HIPPIES (Yan et al. 2011) at similar depths, in the ground-based survey of Laporte et al. (2011) at a brighter level, and in the much wider survey in Capak et al. (2011) at an even brighter limit \\footnote{While Bowler et al. (2012) derived different $z_{ph}$ ($\\sim 2$) for the candidates in Capak et al. (2011), the true nature of these objects still remains uncertain.}.  \\subsection{Implications for Star Formation in the Early Universe }  Our $z\\approx 8$ candidate galaxy samples have two very luminous objects and two high stellar mass objects. We have presented tentative evidence that the bright end of the galaxy LF at $z\\approx 8$ might not follow the exponential cut-off of the Schechter function. While the stellar mass function at this redshift has not yet been determined, we have provided the first, albeit still tentative measurement at the high-mass end and shown that some galaxies with stellar masses as high as a few $\\times 10^9$~$M_\\odot$ might already in place at $z\\approx 8$.  The physical origin of the exponential cut-off at the bright end of the LF or the high mass end of the MF seen at lower redshifts remains a matter of debate. At $z\\sim 0$, star formation is ``quenched'' in halos with masses greater than $\\sim 10^{12}$~$M_\\odot$ (``quenching mass''), perhaps by AGN feedback (e.g., Croton et al. 2006; Somerville et al. 2008; Gabor et al. 2011). However, at higher redshift ($z\\lesssim 2$), the quenching mass may be higher (Dekel et al. 2009; Behroozi et al. 2010), and the existence of massive, rapidly star forming galaxies at $z>2$ is well established. At $z\\sim 2$, the observed exponential cut-off in the rest-frame UV LF appears to be due to dust (Reddy et al. 2010) rather than quenching. At very high redshifts, it has been noted before that the very blue rest-frame UV colors of $z\\approx 7$ candidate galaxies (some of them confirmed) suggest that they may contain little dust (e.g., Oesch et al. 2010; M10; Y10; Bunker et al. 2010). Thus, at $z\\approx 8$, when (1) the quenching mass is much higher than the typical halo mass, and (2) dust has little effect, perhaps we should not expect to see an exponential cutoff in the luminosity or mass functions of UV-bright galaxies.  It is interesting that the estimated stellar masses and number densities could imply a rather high efficiency of conversion of baryons into stars. For the lower stellar mass estimates, which are obtained using the models with nebular emission (the average is $\\sim 2\\times 10^9$~$M_\\odot$), and assuming that $\\sim 20$\\% of the available baryons have been converted to stars, the implied halo masses are on the order of $\\sim 10^{11}$~$M_\\odot$, for which the expected number density of dark matter halos in the currently favored LCDM cosmology is comparable to the observed number density of objects. This is consistent with the predicted star formation efficiencies and host halo masses at $z\\approx 8$ from cosmological hydrodynamic simulations. However, if the true stellar masses are higher, or the star formation efficiency is lower, a rapidly growing tension arises between the number density of dark matter halos and the observed number density of galaxies (above halo mass $\\sim 10^{11}$~$M_\\odot$, the halo number density declines by about two orders of magnitude for a factor of $\\sim 3$ increase in mass). While converting 20\\% of the available baryons into stars may not sound excessive, this is in fact the \\emph{maximum} value that has been inferred at any epoch. Due to the presumably very low metallicity of the gas in these early objects, we might have expected much lower star formation efficiencies than are seen locally (e.g. Krumholz \\& Dekel 2010)."
    },
    "1112/1112.3824_arXiv.txt": {
        "abstract": "{Recently, for the first time the abundance of P has been measured in disk stars. This provides the opportunity of comparing the observed abundances with predictions from theoretical models.} {We aim at predicting the chemical evolution of P in the Milky Way and compare our results with the observed P abundances in disk stars in order to put constraints on the P nucleosynthesis.}  {To do that we adopt the two-infall model of galactic chemical evolution, which is a good model for the Milky Way, and compute the evolution of the abundances of P and Fe. We adopt stellar yields for these elements from different sources. The element P should have been formed mainly in Type II supernovae. Finally, Fe is mainly produced by Type Ia supernovae.}  {Our results confirm that to reproduce the observed trend of [P/Fe] vs. [Fe/H] in disk stars, P is formed mainly in  massive stars. However, none of the available yields for P can reproduce the solar abundance of this element.  In other words, to reproduce the data one should assume that massive stars produce more P than predicted by a factor of $\\sim$ 3. } {We conclude that all the available yields of P from massive stars are largely underestimated and that nucleosynthesis calculations should be revised. We also predict the [P/Fe] expected in halo stars. } ",
        "introduction": "Recently, Caffau et al. (2011a) have measured the P abundance in a sample of 20 cool stars in the Galactic disk. They found that the [P/Fe] ratio behaves like the [S/Fe] ratio, namely increasing towards lower metallicity stars. This was the first time that P was observed in Galactic stars in spite of the fact that P is among the top 20 most abundant elements in the Universe.  There is only one single stable isotope of P, $^{31}$P, which is thought to be formed by neutron capture on $^{29}$Si and $^{30}$Si in massive stars. By means of a detailed model for the chemical evolution of the Milky Way one can compute the P evolution and compare it to the observations. This can allow us to understand the origin of this element and impose constraints on its formation inside stars.  In this paper, we adopt a good model for the chemical evolution on the Galaxy, already tested on a large number of chemical species (Fran\\c cois et al. 2004). In particular, the model takes into account detailed stellar nucleosynthesis and supernova progenitors (SNe Ia, II, Ib/c) as well as the stellar lifetimes.  The stellar yields of P have been computed by several authors such as Woosley \\& Weaver (1995), Kobayashi et al. (2006) as functions of the initial stellar metallicity.  In this paper we like to compare the predictions of our chemical evolution model for P, obtained by including different stellar yields, with the recent data.  In Section 2 we briefly describe the chemical evolution model adopted, in Section 3 we describe the observational data. In Section 4 the results are compared to the data and finally, in Section 5, some conclusions are drawn. ",
        "conclusions": "In this paper we have compared model predictions adopting different sets of yields with recent data on the abundance of P in the Galactic disk. Our conclusions can be summarized as follows:  \\begin{itemize}   \\item The observed fall of the abundance of [P/Fe] in Galactic disk stars suggests that P is mainly produced by core-collapse SNe with a small contribution from SNe Ia. Otherwise, if SNe Ia were important the  [P/Fe] ratio would remain fairly constant.   \\item The metallicity dependent yields of P from massive stars of K06 together with the P yields from SNe Ia from Iwamoto et al. (1999) well reproduce the data if the yields from massive stars are increased by a factor of $\\sim$ 3. This suggests that the yields of P available in the literature are underestimated. Both the neutron rich isotopes of $^{29,30}$Si and $^{31}$P, which derives from neutron capture on the two Si isotopes, are produced in the oxygen and neon burning shells in massive stars (WW95). Therefore, in order to have a larger P production, the O and Ne shell- burnings as well as the neutron capture on the $^{29,30}$Si isotopes and the destruction of $^{31}$P by (p, $\\alpha$) reactions should be revised.   \\item We predict also the behaviour of [P/Fe] in halo stars and we suggest that it should show a plateau between [Fe/H] =-1.0 and -3.0 dex corresponding to [P/Fe] $\\sim$ +0.5 dex if yields (corrected) from normal SNe II are adopted, and to $\\sim$+0.2dex if hypernova yields (corrected) are adopted. In order to distinguish between these two cases we await for observations of the P abundance in halo stars.  \\end{itemize}"
    },
    "1112/1112.5319_arXiv.txt": {
        "abstract": "{We used the new IRAM~30-m FTS backend to perform an unbiased $\\sim$15~GHz wide survey at 3~mm toward the \\pipe\\ young diffuse starless cores. We found an unexpectedly rich chemistry. We propose a new observational classification based on the 3~mm molecular line emission normalized by the core visual extinction (\\Av). Based on this classification, we report a clear differentiation in terms of chemical composition and of line emission properties, which served to define three molecular core groups. The ``diffuse'' cores, \\Av$\\la$15, show poor chemistry with mainly simple species (e.g. \\cs\\ and \\cdh).  The ``oxo-sulfurated'' cores, \\Av$\\simeq$15--22, appear to be abundant in species like \\so\\ and \\sod, but also in  HCO, which seem to disappear at higher densities. Finally,  the ``deuterated'' cores, \\Av$\\gsim$22, show typical evolved chemistry prior to the onset of the star formation process, with nitrogenated and deuterated species, as well as carbon chain molecules. Based on these categories, one of the ``diffuse'' cores (Core~47) has the spectral line properties of the ``oxo-sulfurated'' ones, which suggests that it is a possible failed core.} ",
        "introduction": "A new generation of sensitive receivers and wideband backends allows to study in detail the chemistry of faint starless cores. Several surveys have been performed toward them reporting a rich but relatively simple chemistry: essentially carbon chemistry with significant sulfur and nitrogen bearing molecules, followed by later deuteration which can be used as a chemical clock (e.g., \\citealp{turner94,turner00,hirota06,tafalla06,bergin07}). Recently, from the theoretical side, several papers have tried to model the starless core chemistry self-consistently \\citep{aikawa01,garrod05,keto08}.  The \\pipe\\ is a nearby (145~pc: \\citealp{alves07}) cloud  that harbors more than one hundred of low mass ($\\sim$1~M$_{\\odot}$) starless cores, most of them gravitational unbound but confined by the thermal/magnetic pressure  of the whole cloud \\citep{alves08,lada08,franco10}. The \\pipe\\ differs from the other nearest dark cloud complexes such as Taurus  or $\\rho$~Ophiuchus because it has a very small star formation efficiency \\citep{onishi99,forbrich09,roman10,roman11}.   Thus, the \\pipe\\ is an ideal target to study the physical and chemical conditions in a pristine environment prior to the onset of the star formation process, as the recent numerous studies have shown (e.g., \\citealp{brooke07,muench07,rathborne08}). \\citet{frau10} present the first results of an extensive continuum and molecular line study on a subset of a selected sample of cores distributed in the different regions of the \\pipe: {\\it bowl}, {\\it stem}, and B59.  The cores are in general less dense and less chemically evolved than starless cores in other star forming regions studied (e.g. \\citealp{crapsi05}).  We find very different morphologies and densities, and  no clear correlation of the chemical evolutionary stage of the cores with their  location in the cloud.  The \\pipe\\ starless cores have shown to be more heterogeneous than expected.  In this work, we present a wide ($\\sim$15~GHz) unbiased chemical survey at 3~mm toward a larger sample of \\pipe\\ starless cores, spanning a factor of 6 in their visual extinction (\\Av) peaks. This is a first step to characterize their varied chemistry in order to proceed to future modeling. ",
        "conclusions": "}  The chemistry detected toward the sample of fourteen starless cores is unexpectedly rich taking into account their low temperatures (10--15 K: \\citealp{rathborne08}) and visual extinctions. The apparent correlation within the sample of the 3~mm molecular transition normalized intensities to visual extinction allow us to propose an observational classification (see Fig.~\\ref{fig_spec}). We define three groups of starless cores, which are probably related with their dynamical age: {\\it ``diffuse''},  {\\it ``oxo-sulfurated''}, and {\\it ``deuterated''} cores. This classification can be useful in future wide band 3~mm observations of molecular clouds.  {\\it -- ``Diffuse'' cores}: a set of cores with small column densities (\\Av$\\lsim$15~mag $\\sim$ $N_{\\rm H_2}$$\\lsim$1.2$\\times$10$^{22}$~cm$^{-2}$) lies above the blue dot-dashed horizontal line in Fig.~\\ref{fig_spec}.  Their spectra is rather poor, showing only significant normalized intensity in the  transitions of the main isotopologues of abundant species like \\cdh,  HCN (and likely HNC), \\hcom\\ and SO.  Such a simple observational chemistry suggests that these are very young starless cores, or even transient clumps on which essentially the cloud chemistry is better detected due to density enhancements. Core~47 is a clear exception and it is discussed later in the text.  {\\it -- ``Oxo-sulfurated'' cores}: a group of denser cores  (\\Av$\\simeq$15--22~mag  $\\sim N_{\\rm H_2} \\simeq$1.2$\\times$10$^{22}$--1.7$\\times$10$^{22}$~cm$^{-2}$) that show richer chemistry but not yet significant deuteration to be observed.  In Fig.~\\ref{fig_spec} this group  lie between the  blue dot-dashed and red dashed horizontal lines. All the transitions detected in the {\\it ``diffuse''} cores are also present.  The \\so~3$_2$--2$_1$ transition is the main signpost as it is very  bright. Many other oxo-sulfurated molecules (\\tso, \\sod, and \\ocs), as well as \\hco, exhibit a same trend, but they are not detected at higher densities. This suggests an increase of these chemically related species in the gas-phase in this \\Av\\ range, followed by a later depletion/destruction as density increases. These cores might be in-the-making cores, which have developed a richer chemistry  and piled up more material, probably in a stage near to the onset of collapse \\citep{ruffle99}.   Core~102 is an exception in this group and is discussed later in the text.  {\\it -- ``Deuterated'' cores}: the densest cores of the sample (\\Av$\\gsim$22~mag $\\sim$ $N_{\\rm H_2}$$\\gsim$1.7$\\times$10$^{22}$~cm$^{-2}$), shown below the red dashed horizontal line in Fig.~\\ref{fig_spec}. Core~12, the densest one, sets the upper limit at \\Av=67.2~mag ($N_{\\rm H_2}$$\\simeq$5.3$\\times$10$^{22}$~cm$^{-2}$). These cores are generally bright in the transitions typical of the other two groups. The oxo-sulfurated molecules are the exception, hardly present, probably depleted/destructed at the densities reached. The main signpost are the emission, only present in this group, in rare isotopologues of the nitrogenated ubiquitous lines (\\htcn, \\hcqn, \\hntc, and \\hqnc), deuterated forms of abundant species (\\cthdt, \\nhdd, and \\dctn), and carbon-chain molecules (\\cqh\\ and \\chtcdh).  These cores might be stable starless cores with a life-time long enough to achieve the densities needed to synthesize efficiently carbon chains and deuterated species \\citep{roberts00,gwenlan00}.  As already said, Core~47 does not match the chemical properties of the diffuse cores. It shows a chemistry matching the oxo-sulfurated group, which proved to be very sensitive to density. This suggests that it might be a failed core which developed a rich chemistry and is now merging back into the cloud. This scenario can increase the oxo-sulfurated chemistry detected \\citep{garrod05}.  Core~47 is located close to Core~48 in the only \\pipe\\ region with superalfv\\'enic turbulence, as shown by optical polarization observations \\citep{franco10}. Therefore, it is possible that an external source of turbulence is disrupting the medium in this area and dispersing the cores.  On the other hand, Core~102, in the oxo-sulfurated group, shows a chemistry similar to that of the diffuse cores. Similarly, Core~87, among the deuterated cores, shows features of the oxo-sulfurated group. This suggests that these cores might have piled up material so quickly that a more complex chemistry had no time to be synthesized. Both cores lie in the same N-S oriented high-density structure \\citep{roman10} where \\citet{franco10} report a N-S magnetic field. The fast evolution might have been driven by magnetic fields with the surrounding mass collapsing in this direction.  The FTS chemical survey toward the starless cores of the \\pipe\\ showed a chemistry much more rich than expected for a cloud giving birth to low-mass stars at very low efficiency. A good interpretation of the results demands chemical modeling to investigate a possible evolutive track, which will be the purpose of a forthcoming study."
    },
    "1112/1112.0016_arXiv.txt": {
        "abstract": "VFTS 682, a very massive and very hot Wolf-Rayet (WR) star recently discovered in the Large Magellanic Cloud near the famous star cluster R136, might be providing us with a glimpse of a missing link in our understanding of Long Gamma-Ray Bursts (LGRBs), including dark GRBs. It is likely its properties result from chemically homogeneous evolution (CHE), believed to be a key process for a massive star to become a GRB. It is also heavily obscured by dust extinction, which could make it a dark GRB upon explosion. Using {\\em Spitzer}\\/ data we investigate the properties of interstellar dust in the vicinity of R136, and argue that its high obscuration is not unusual for its environment and that it could indeed be a slow runaway (``walkaway'') from R136. Unfortunately, based on its current mass loss rate, VFTS 682 is unlikely to become a GRB, because it will lose too much angular momentum at its death. If it were to become a GRB, it probably would also not be dark, either escaping or destroying its surrounding dusty region. Nevertheless, it is a very interesting star, deserving further studies, and being one of only three presently identified WR stars (two others in the Small Magellanic Cloud) that seems to be undergoing CHE. ",
        "introduction": "Ideally, in order to fully understand the mapping between massive star progenitors and when and how they explode (or not, \\citealt{Kochanek08}), we would have extensive multi-wavelength data obtained for many such explosions both BEFORE and AFTER the event. The AFTER part has certainly undergone an explosive growth in the last decade or so, due to many successful supernova (SN) searches such as the galaxy-targeted {\\em Lick Observatory Supernova Search} (LOSS), the {\\em Catalina Real-time Transient Survey} (CRTS), the {\\em Robotic Optical Transient Search Experiment} (ROTSE), the {\\em Palomar Transient Factory} (PTF); and gamma-ray burst searches such as the {\\em High Energy Transient Explorer-2 (HETE-2)} and {\\em Swift}.  The BEFORE part is naturally limited by the fact that massive stars, while relatively bright, are only observable in the local Universe ($d\\lesssim10\\;$Mpc). Here we have had to rely on proximity (SN 1987A, \\citealt{SN1987A}), luck based on archival data (e.g., \\citealt{Smartt09}), or the first systematic campaign to monitor future SN progenitors (\\citealt{Szczygiel11}), where many nearby galaxies are observed to sufficient depth so eventually they will provide progenitor data for a significant number of future SNe. We have by now compiled a fair amount of information on SN progenitors, such as the blue supergiant progenitor for SN 1987A (\\citealt{White87}) or dusty progenitors of SN 2008S and 2008 NGC 300 transient (\\citealt{Prieto08}; \\citealt{Prieto08b}).  However, for very rare explosive events, such as long-duration gamma-ray bursts (LGRBs), the systematic approach that works for the normal core-collapse SNe is not yet possible because the events are so rare. Therefore the prospect of future systematic observational efforts that could identify nearby LGRB progenitors is very dim. Here we have to rely on a more extended chain of reasoning, with data being supplemented with reasonable, theoretical guesses. Discovery of the connection between LGRBs and broad-lined Type Ic SNe (e.g., \\citealt{Stanek03}; \\citealt{Hjorth03}), thought to result from core-collapse of hydrogen-free massive stars, favors two possible progenitor models: single massive Wolf-Rayet (WR) stars with rapidly rotating cores (\\citealt{FS85}), or lower mass helium stars stripped by a close binary companion (\\citealt{Podsiadlowski04}). The prompt emission of an initial LGRB, lasting seconds or minutes, is usually followed by a multi-wavelength afterglow which lasts days to even years. Although X-ray afterglows of LGRBs are nearly always detected by {\\em Swift} and {\\em Fermi}, detection of optical and infrared afterglows is less common. A dark GRB is defined by either an absent or faint optical afterglow relative to its X-ray emission. The rapid detection of X-ray afterglows with {\\em Swift} revealed that the dark fraction of LGRBs is about $30\\%$ of GRBs (\\citealt{AS07}; \\citealt{Perley09}). In general, besides the intrinsically optically faint GRBs, the optical attenuation of dark GRBs can be caused by dust extinction in GRB host galaxies, foreground extinction, or Lyman-$\\alpha$ absorption by neutral hydrogen at high redshifts (\\citealt{Perley09}).  Study of massive stars in very nearby galaxies can supply the other missing links in our understanding of LGRBs.  The recently uncovered very massive stars up to $300M_{\\odot}$ in the Large Magellanic Cloud (LMC) young cluster R136 extend our knowledge of massive star formation and evolution (\\citealt{Crowther10}).  In particular, the newly discovered very massive WR star VFTS 682, located 30 pc away from R136, drew our attention, mainly because of its very high foreground dust extinction and possibility of being a GRB progenitor (\\citealt{VFTS682}). Its unusually high effective temperature can be understood as the consequence of chemically homogeneous evolution (CHE), which was proposed as the crucial part of the process of producing LGRBs (\\citealt{Yoon05}). If VFTS 682 will indeed make a GRB at the end of its evolution, the high foreground dust extinction could mean that it will become a dark GRB. However, its high mass and strong mass-loss seems to prevent it from being a GRB progenitor, and the potential afterglow might destroy the surrounding dust even if VFTS 682 eventually produces a GRB. Therefore, it is worthwhile to look into the details of VFTS 682's evolution and the fate of its circumstellar environment. Section \\ref{Data} describes the dust properties in the VFTS 682 vicinity region. Section \\ref{Progenitor} investigates the possibility of VFTS 682 producing a LGRB and Section \\ref{Dark} discusses the fate of dusty clouds around VFTS 682. Conclusions are presented in Section \\ref{Con}. ",
        "conclusions": "\\label{Con}  We give a first order estimate whether VFTS 682 can be a LGRB progenitor in the sense of CHE, which is still a question mark in previous work. VFTS 682 will most likely lose more than 85\\% of its current angular momentum and cease its CHE evolution before it leaves MS with a remaining mass $\\sim50-70M_{\\odot}$. In general VFTS 682 will fail to produce a GRB due to its strong mass-loss and the resulting angular momentum loss, unless it has a currently extreme rapid rotation ($v_{0}>700$ km s$^{-1}$) which helps it maintain CHE and eventually produce a GRB associated with a hypernova. Wind anisotropies and shorter lifetime could leave a more massive faster rotating star, but it is doubtful whether a relativistic jet could travel through such a thick stellar envelope and break out of the stellar surface, unless it is a rare long-lived ($\\gtrsim100$ s) LGRBs. Similarly, it is unlikely that VFTS 682 will be heavily obscured at death and produce a dark GRB. Its proper motion will probably cause it out of the dusty region, and a GRB of an early optical afterglow $L_{\\rm opt}>10^{50}$ erg s$^{-1}$ would destroy the dust clouds within 30 pc, otherwise the dusty clouds can be heated and ionized up to a region of $\\sim$ 100 pc by X-ray and optical radiation given by the death of VFTS 682.  CHE is believed to be the crucial process in the evolution path towards LGRBs. However, the observed sample of WRs likely undergoing CHE is quite small. The only observation of WR stars besides VFTS 682 which might be undergoing are two WNh stars in the SMC (i.e., SMC-WR1 and WR2 in \\citealt{Martins09}), which makes these three WR stars extremely important to understanding the evolution of WR stars undergoing CHE and the related problem of LGRB formation. The angular momentum losses of the two WR stars in the SMC will be much less significant than VFTS 682 in the LMC. Taking the stellar parameters in \\cite{Martins09}, we estimate that the nuclear timescale of SMC-WR1 (WR2) $\\tau_{\\rm nuc}\\sim4$ Myr (5 Myr) is shorter than the wind timescale $\\tau_{\\rm wind}\\sim10-20$ Myr, and the final stellar rotation velocity should be 90\\% of the current velocity. Therefore, the two WR stars in the SMC are more likely GRB progenitors in the scenario of CHE, although the metallicity threshold is still an issue (\\citealt{Martins09}). In any case, finding CHE WR stars in the LMC and SMC promotes a future work on theory models.  VFTS 682 is a very interesting star. As mentioned in Section 2, because there is no observation evidence to show that VFTS 682 is unusual compared to other stars in the 30 Dor region, the interesting possibility of existence of others massive stars like VFTS 682 in the vicinity of R136 is not excluded. Note that R136 is sufficient young and massive ($\\leq5.5\\times10^{4}M_{\\odot}$) to generate runaway stars beyond $150M_{\\odot}$. Recently \\cite{Banerjee11} gives a theoretical model study on the dynamical ejection of runaway massive stars from R136. We suggest that there might be other massive stars that ``walked away'' from R136, but are currently hidden behind even more dust than VFTS 682. Since mid-IR date cannot be used to flag such stars, spectroscopic observations, which are beyond the scope of this Letter, should be further investigated to show the possibility of their existence."
    },
    "1112/1112.1589_arXiv.txt": {
        "abstract": "{Generation and dissipation of magnetic fields is a fundamental physical process on the Sun. In comparison to flux emergence and the initial stages of sunspot formation, the demise of sunspots still lacks a comprehensive description. } {The evolution of sunspots is most commonly discussed in terms of their intensity and magnetic field. Here, we present additional information regarding the three-dimensional flow field in the vicinity of sunspots towards the end of their existence.} {We present a subset of multi-wavelengths observations obtained with the Japanese \\textit{Hinode} mission, the \\textit{Solar Dynamics Observatory} (SDO), and the \\textit{Vacuum Tower Telescope} (VTT) at \\textit{Observatorio del Teide}, Tenerife, Spain during the time period from 2010 November 18--23. Horizontal proper motions were derived from G-band and Ca\\,\\textsc{ii}\\,H images, whereas line-of-sight velocities were extracted from VTT Echelle H$\\alpha$~$\\lambda656.28$~nm spectra and Fe\\,\\textsc{i} $\\lambda630.25$~nm spectral data of the \\textit{Hinode/Spectro-Polarimeter}, which also provided three-dimensional magnetic field information. The \\textit{Helioseismic and Magnetic Imager} on board SDO provided continuum images and line-of-sight magnetograms as context for the high-resolution observations for the entire disk passage of the active region.} {We have performed a quantitative study of photospheric and chromospheric flow fields in and around decaying sunspots. In one of the trailing sunspots of active region NOAA~11126, we observed moat flow and moving magnetic features (MMFs), even after its penumbra had decayed. We also noticed a superpenumbral structure around this pore. MMFs follow well-defined, radial paths from the spot all the way to the border of a supergranular cell surrounding the spot. In contrast, flux emergence near the other sunspot prevented it from establishing such well ordered flow patterns, which could even be observed around a tiny pore of just 2~Mm diameter. After the disappearance of the sunspots/pores a coherent patch of abnormal granulation remained at their location, which was characterized by more uniform horizontal proper motions, low divergence values, and diminished photospheric Doppler velocities. This region, thus, differs significantly from granulation and other areas covered by G-band bright points. We conclude that this peculiar flow pattern is a signature of sunspot decay and the dispersal of magnetic flux.} {} ",
        "introduction": "Sunspots are a thought-provoking aspect of solar activity because of the close interaction between plasma motions and magnetic fields. Recent progress in MHD simulations \\citep[e.g.,][]{Rempel2011} provide a comprehensive framework for the interpretation of high-resolution sunspot observations. The formation of a penumbra around a sunspot is a rapid phenomenon, i.e., within a few hours a sunspot can develop a penumbra \\citep{Leka1998, Yang2003a}, which is intimately linked to more inclined magnetic field lines and the onset of the Evershed flow. \\citet{Schlichenmaier2010b} observed the growth of a penumbra where only the newly formed penumbra contributed to the increase in spot size while the umbra remained stable. The formation of a penumbra, which would surround the entire spot, was hindered by continuous flux emergence between the spots of the bipolar region.  Quite the opposite, the decay of a sunspot is a slow process. Decay rates for stable leading sunspots and irregular follower spots are different \\citep{MartinezPillet2002}. A number of decay laws were proposed such as a linear decay law described by \\citet{Bumba1963} and a parabolic decay law proposed by \\citet{Petrovay1997}. \\citet{MartinezPillet2002} critically reviews various diffusion models, concludes that they explain well how flux is spread over larger areas while the spot is decaying, but they fail to satisfactorily describe the flux removal process. The initial stages of sunspot decay, i.e., when the spot looses its penumbra, are exemplary described in \\citet{BellotRubio2008}, who discovered finger-like structures, which are neither related to penumbral filaments nor the Evershed flow. These features might be penumbral field lines rising to the chromosphere, thus contributing to the decay of the sunspot penumbra. When a sunspot looses its penumbra, its decay reaches a critical point. Magnetic field lines become more vertical and convective motions in its vicinity begin to change. These ideas of a critical inclination angle and convective motions were put together by \\citet{Rucklidge1995}, who explains in a simple model why small sunspots can have a penumbra while larger pores do not possess one.  The moat flow is a large-scale flow pattern commonly observed around sunspots \\citep{Meyer1974}. However, flux removal and dispersal can only be understood in the context of the moat flow's fine structure. Moving magnetic features (MMFs) play a major role in the flux dispersal process and they are only associated with decaying sunspots \\citep{Harvey1973}. The total flux carried by MMFs is several times larger than the flux contained within the sunspot itself. Thus, the polarity of MMFs has to be considered for a balance of the net flux. MMFs move radially outward with a velocity of 1~km~s$^{-1}$ before they reach and dissolve within the network, i.e., at the boundaries of the supergranular cell containing the sunspot. \\citet{Zuccarello2009} showed evidence that MMFs and moat flow  are present even in the vicinity of pores, i.e., in the absence of penumbral filaments and Evershed flow \\citep[cf.,][]{CabreraSolana2006}. \\citet{Deng2007} also detected a persistent moat flow after the penumbra around spot disappeared leaving only a pore. Even though the moat flow might not be closely tied to the Evershed flow, the sub-photospheric interaction of magnetic field lines and flows can still produce the observed flow patterns.  \\citet{Verma2011} described a local correlation tracking (LCT) method to measure horizontal flows based on \\textit{Hinode} G-band images. In this study, we perform a case study, where we put such horizontal flow fields in the context of other photospheric and chromospheric data. In particular, we are interested in the final stages of sunspot decay. In Sect.~\\ref{SEC02}, we present a subset of multi-wavelengths observations, which were obtained within the scope of \\textit{Hinode Operation Plan} (HOP) 0176. The temporal evolution of active region NOAA~11126 in terms of intensity, morphology, and flow as well as magnetic fields is described in Sect.~\\ref{SEC03} and discussed in Sect.~\\ref{SEC04}.    \\begin{table*}[t] \\caption{Observing characteristics and physical parameters}\\label{TAB01} \\tiny \\begin{tabular}{llcccccp{65mm}} \\cline{1-7}\\vspace*{-2.5mm}\\\\\\cline{1-7} \\multicolumn{2}{l}{November} & 18 & 19 & 20 & 21 & 22 & \\rule[-3pt]{0pt}{12pt}\\\\ \\cline{1-7} $B$                  & & S32.6$^{\\circ}$ & S32.6$^{\\circ}$ & S32.6$^{\\circ}$ & S32.6$^{\\circ}$ & S32.6$^{\\circ}$ & \\rule{0pt}{9pt}\\\\ $L$                  & & E5.5$^{\\circ}$  & W6.3$^{\\circ}$  & W20.5$^{\\circ}$ & W31.5$^{\\circ}$ & W44.1$^{\\circ}$ & \\\\ $\\mu$                &  & $0.81$  & $0.81$ &  $0.77$ & $0.70$ & $0.58$ & \\\\ $t_{0,\\mathrm{GB}}$  & UT & 10:23 & 09:05 & 09:07 & 09:05 & 09:05 &  \\\\ $t_{0,\\mathrm{SP}}$  & UT & 11:46 & 09:05 & 09:07 & 09:04 & 09:04 &  \\\\ $t_{0,\\mathrm{ES}}$  & UT & 10:22 & 08:59 & 14:16 & 09:31 & 08:47 &  \\rule[-4pt]{0pt}{6pt}\\\\ \\cline{1-7} $A_\\mathrm{Spot\\ A}$ & Mm$^{2}$          & 33.8               &  21.5           &     &     &     & \\rule{0pt}{9pt}\\\\ $v_\\mathrm{Spot\\ A}$ & km s$^{-1}$       & $0.22 \\pm 0.19$    & $0.15 \\pm 0.15$ &     &     &     & \\\\ &                   & $0.20 \\pm 0.25$    & $0.15 \\pm 0.14$ &     &     &     & \\\\ $v_\\mathrm{Ring\\ A}$ & km s$^{-1}$       & $0.44 \\pm 0.23$    & $0.37 \\pm 0.19$ &     &     &     & \\\\ & km s$^{-1}$       & $0.44 \\pm 0.23$    & $0.38 \\pm 0.19$ &     &     &     & \\\\ $\\left|\\nabla v\\right|_\\mathrm{Ring\\ A}$ & $10^{-3}$ s$^{-1}$ & $0.83 \\pm 0.78$ & $0.65 \\pm 0.54$ &      &     &     & \\\\ &                    & $1.28 \\pm 1.33$ & $1.13 \\pm 1.06$ &      &     &     & \\rule[-4pt]{0pt}{6pt}\\\\ \\cline{1-7} $A_\\mathrm{Spot\\ B}$ & Mm$^{2}$          & 54.7               & 21.6            & 54.2            & 17.5            & 2.7 & \\rule{0pt}{9pt}\\\\ $v_\\mathrm{Spot\\ B}$ & km s$^{-1}$       & $0.22 \\pm 0.15$    & $0.15 \\pm 0.15$ & $0.16 \\pm 0.07$ & $0.18 \\pm 0.05$ & $0.28 \\pm 0.06$ & \\\\ &                   & $0.20 \\pm 0.14$    & $0.11 \\pm 0.05$ & $0.15 \\pm 0.08$ & $0.25 \\pm 0.06$ & $0.37 \\pm 0.06$ & \\\\ $v_\\mathrm{Ring\\ B}$ & km s$^{-1}$       & $0.31 \\pm 0.19$    & $0.20 \\pm 0.12$ & $0.28 \\pm 0.14$ & $0.25 \\pm 0.14$ & $0.30 \\pm 0.09$ & \\\\ &                   & $0.32 \\pm 0.20$    & $0.20 \\pm 0.12$ & $0.29 \\pm 0.15$ & $0.28 \\pm 0.13$ & $0.40 \\pm 0.09$ & \\\\ $\\left|\\nabla v\\right|_\\mathrm{Ring\\ B}$ & $10^{-3}$ s$^{-1}$ & $0.63 \\pm 0.68$ & $0.37 \\pm 0.36$ & $0.46 \\pm 0.49$ & $0.47 \\pm 0.51$ & $0.28 \\pm 0.33$ & \\\\ &                    & $0.91 \\pm 1.20$ & $0.47 \\pm 0.59$ & $0.76 \\pm 1.04$ & $0.63 \\pm 0.81$ & $0.44 \\pm 0.44$ &  \\rule[-4pt]{0pt}{6pt}\\\\ \\cline{1-7} $A_\\mathrm{mag}$  &  Mm$^{2}$   & $136.9$ & $83.5$& $98.9$ & $44.7$ & $22.2$ & \\rule{0pt}{9pt}\\\\ $N_\\mathrm{mag}$  &             & 10      & 5     & 9      & 6      & 2 & \\\\ $v_\\mathrm{mag}$  & km s$^{-1}$ & $0.22 \\pm 0.12$ & $0.15 \\pm 0.09$ & $0.16 \\pm 0.07$ & $0.18 \\pm 0.05$ & $0.28 \\pm 0.06$ & \\\\ $v_\\mathrm{bp}$   & km s$^{-1}$ & $0.27 \\pm 0.19$ & $0.27 \\pm 0.16$ & $0.28 \\pm 0.16$ & $0.23 \\pm 0.12$ & $0.26 \\pm 0.10$ & \\\\ $v_\\mathrm{gran}$ & km s$^{-1}$ & $0.40 \\pm 0.24$ & $0.35 \\pm 0.20$ & $0.37 \\pm 0.20$ & $0.34 \\pm 0.19$ & $0.34 \\pm 0.18$ & \\\\ $\\left|\\nabla v\\right|_\\mathrm{gran}$ & $10^{-3}$ s$^{-1}$ & $0.88 \\pm 0.89$ & $0.72 \\pm 0.67$ & $0.85 \\pm 0.79 $ & $0.95 \\pm 0.88$ & $1.04\\pm 0.98$ & \\\\ $\\left|v\\right|_\\mathrm{mag, LOS}$    & km s$^{-1}$        & $0.03 \\pm 0.17$ & $0.02 \\pm 0.12$ & $0.04 \\pm 0.20$  & $0.02 \\pm 0.12$ & $0.00 \\pm 0.02$ &\\\\ $\\left|v\\right|_\\mathrm{bp, LOS}$     & km s$^{-1}$        & $0.43 \\pm 0.37$ & $0.44 \\pm 0.37$ & $0.44 \\pm 0.38$  & $0.41 \\pm 0.38$ & $0.34 \\pm 0.30$ &\\\\ $\\left|v\\right|_\\mathrm{gran, LOS}$   & km s$^{-1}$        & $0.51 \\pm 0.41$ & $0.50 \\pm 0.41$ & $0.55 \\pm 0.45$  & $0.57 \\pm 0.47$ & $0.55 \\pm 0.45$ &\\rule[-4pt]{0pt}{6pt} \\raisebox{23mm}[-23mm]{\\parbox{65mm}{ The parameters in the first column of the table refer to heliographic latitude $B$, heliographic longitude $L$, cosine of the heliocentric angle $\\mu$, start of the G-band (GB), spectro-polarimeter (SP), and Echelle spectrograph (ES) observing sequences $t_{0}$, spot area $A$, horizontal velocity $v$, mean divergence $\\left|\\nabla v\\right|$, number of magnetic elements $N_\\mathrm{mag}$. The indices refer to G-band bright points (bp), granulation (gran), magnetic elements (mag), the two sunspots (Spot~A and Spot~B), and the four-megameter wide annuli around both spots (Ring~A and Ring~B). If two rows are given for a physical parameter, then the top row refers to G-band data, whereas the bottom row was derived from Ca\\,\\textsc{ii}\\,H data. If present, the standard deviation refers to the variation of the physical parameters  within the specified regions rather than to any formal error.}}\\\\ \\cline{1-7} \\end{tabular} \\end{table*} ",
        "conclusions": "}  We have presented a detailed account of the final stages in the decay of the active region NOAA~11126,  which did not obey the Hale-Nicholson polarity law \\citep[e.g.,][]{Zirin1988a}.  Since only one out of ten active regions shows such a behavior \\citep{Howard1990} and we only present a case study, our results might not be representative for sunspot decay in general. However, space missions such as \\textit{Hinode} and SDO provide nowadays data of sufficient coverage, resolution, and cadence that statistical properties of sunspot decay become accessible. Furthermore, previous studies of non-Hale regions \\citep[e.g.,][and reference therein]{LopezFuentes2000} were centered on flux emergence, $\\delta$-spots, and strong solar flares. The present study can consequently be considered as an extension of these studies with a focus on a much quieter magnetic field topology, which might be representative for the lower solar activity during cycle No. 24 \\citep{Petrovay2010, Nielsen2011}.  The major findings of our study can be summarized as follows: (1) MMFs were observed in the vicinity of spot~\\textsf{A} until it decayed. Mostly type~II and a few interspersed \\textsf{U}-shaped type~I MMFs contributed to the observed moat flow, which also left a clear signature in the time-averaged Ca\\,\\textsc{ii}\\,H images \\citep{MartinezPillet2002}. (2) Even though penumbral filaments had almost completely disappeared in photospheric G-band images of spot~\\textsf{A} on November~18, H$\\alpha$ line core images clearly exhibited a structure reminiscent of a superpenumbra. Thus, filamentary structures including the inverse Evershed flow \\citep{Maltby1975, Georgakilas2003} might be more persistent in the chromosphere. (3) We have also observed MMFs in the vicinity of a tiny pore with a diameter of about 2~Mm, which did not show any indication of penumbral filaments. Such an observation argues strongly against a close tie between Evershed flows and MMFs \\citep[cf., ][]{VargasDominguez2008, VargasDominguez2010, CabreraSolana2006}. \\citet{Rempel2011} argues based on MHD simulations of sunspots that penumbral flows can be separated in two components, where the shallow one corresponds to the Evershed flow and the deeper one is related to moat flow. (4) The strong rotation and twist seen in spot~\\textsf{B} might explain, why this trailing spot never advected sufficient magnetic flux to establish more than a rudimentary penumbra and remained highly fragmented during its entire life cycle. (5) The photospheric and chromospheric maps of horizontal flows show a peculiar pattern, once the last dark feature of the active region has disappeared. In general, the flow field in this region is less structured than regions covered by granulation, i.e., the dispersed magnetic field still significantly affects the convective pattern. Chromospheric flows have increased notably compared to times when spots and pores were still present. Most prominently, a contiguous area of low divergence appears towards the end of sunspot decay."
    },
    "1112/1112.4479_arXiv.txt": {
        "abstract": "We study the buildup of magnetic fields during the formation of Population III star-forming regions, by conducting cosmological simulations from realistic initial conditions and varying the Jeans resolution.  To investigate this in detail, we start simulations from identical initial conditions, mandating 16, 32 and 64 zones per Jeans length, and studied the variation in their magnetic field amplification.   We find that, while compression results in some amplification, turbulent velocity fluctuations driven by the collapse can further amplify an initially weak seed field via dynamo action, provided there is sufficient numerical resolution to capture vortical motions (we find this requirement to be 64 zones per Jeans length, slightly larger than, but consistent with previous work run with more idealized collapse scenarios).  We explore saturation of amplification of the magnetic field, which could potentially become dynamically important in subsequent, fully-resolved calculations.  We have also identified a relatively surprising phenomena that is purely hydrodynamic: the higher-resolved simulations possess substantially different characteristics, including higher infall-velocity, increased temperatures inside 1000 AU, and decreased molecular hydrogen content in the innermost region.  Furthermore, we find that disk formation is suppressed in higher-resolution calculations, at least at the times that we can follow the calculation.  We discuss the effect this may have on the buildup of disks over the accretion history of the first clump to form as well as the potential for gravitational instabilities to develop and induce fragmentation. ",
        "introduction": "The formation of the first stars in the Universe is a process well-suited to numerical computation.  While direct observations of these stars are still an optimistic, yet-unrealized prospect \\citep{2005ApJ...629..615W, 2009ApJ...698L..68F, 2005Natur.434..871F, 2011Natur.477...67C}, simulations are able to begin with well-posed initial conditions and directly simulate their formation.  These calculations, typically spanning many orders of magnitude in both spatial and density scales \\citep{2008AIPC..990...16T, 2008Sci...321..669Y}, include the effects of dark matter, hydrodynamics, radiative cooling, chemical heating and cooling, and the non-equilibrium, self-consistent evolution of multiple ionization and molecular states of the gas \\citep{abel97, anninos97, RA04, 2008AIPC..990...25G, 2008MNRAS.388.1627G}.  For nearly a decade, the consensus viewpoint had been that the first stars formed in isolation, between $30-300~\\Msun$ in mass, and only one per $\\sim 10^{6}~\\Msun$ halo.  This had been confirmed by divergent methods: both Smoothed Particle Hydrodynamics (SPH) simulations and Adaptive Mesh Refinement (AMR) calculations showed similar results \\citep{ABN02, 2003ApJ...592..645Y, oshea07a}.  However, recently aspects of this picture are being challenged. Advancements in the current generation of simulations have taken place in three primary areas.  The first is that the effect of heating from the formation of molecular hydrogen via three-body reactions \\citep{2011ApJ...726...55T} is now being taken into account; this will heat the gas during later stages of collapse and through thermal regulation of the accretion rate produce higher infall velocities onto the central core.  The second aspect is related to the so-called ``Courant myopia'' of calculations proceeding to high densities. Because the Courant timescale becomes extremely short at high densities, out of necessity previous simulations stopped after the formation of the first high-density core.  Recent simulations have followed accretion for up to several thousand years, using the numerical technique of sink particles, accreting Lagrangian particles representing protostars below the resolution limit of the simulation \\citep{2011ApJ...737...75G, 2011Sci...331.1040C, 2010MNRAS.403...45S}.  While this approach allows to follow the calculations further in time it does lack the mathematical rigor that may allow one to prove the result to be correct.  The final effect now being added is simply that of sampling: in the published literature, very few calculations had been performed; compared to the wealth of calculations of galaxy mergers, cluster formation, and so on, the sampling of Population III star formation was dramatically underserved, with only a handful of papers discussing even multiple simulations \\citep[see, e.g.,][]{oshea07a}.  These recent shifts have suggested that these halos have fragmented into either a small-number multiplicity of protostars \\citep{2009Sci...325..601T, 2010MNRAS.403...45S} or many small-mass pre-stellar objects \\citep{2011Sci...331.1040C, 2011ApJ...737...75G}. However, none of these simulations have included the effects of magnetic fields in their calculations, nor have they self-consistently followed the growth of seed magnetic fields over the collapse and virialization of these first minihalos. Similarly, the fragmentation has been only studied for times of less than 1\\% of the accretion time scale of the massive stars being studied. It has been difficult to show whether possible early fragments would survive and not grow substantially in mass and simply merge with the central proto-star.  The influence of magnetic fields on the collapse of the first stars has received a renewed interest in recent years \\citep{2008ApJ...685..690M}. Early numerical simulations \\citep{2008ApJ...688L..57X} including the Biermann battery effect suggested that the dynamical effect of magnetic fields on primordial gas is very small.  Though their resolution was rather low, these simulations found that fields were primarily built up by collapse, leading to $B \\propto \\rho^{2/3}$.  The thermodynamics of primordial gas are quite sensitive to the $\\mathbf{B} - \\rho$ relationship via heating from ambipolar diffusion \\citep{2010ApJ...721..615S}.  Although magnetic fields are believed to be unlikely to affect the characteristic fragmentation scale for Population III stars, they may increase the temperature in very high density gas, leading to higher accretion rates onto the protostars \\citep{2009ApJ...703.1096S}. The ability of turbulent fluid motions to drive dynamo action during primordial protostellar collapse has been predicted analytically by \\citet{2010A&A...522A.115S} and confirmed numerically by \\citet{2010ApJ...721L.134S} and \\citet{2011ApJ...731...62F} for idealized collapse calculations. The latter work also suggested that a minimum resolution of 32 elements per Jeans length is required to capture dynamo action, giving a very strong motivation for higher resolution cosmological MHD calculations. Properly resolved, dynamo amplification leads to exponential growth of the field over a turbulent eddy turnover time, and can radically change the strength of the magnetic field at any given density during collapse.  The \\citet{2011ApJ...731...62F} simulations considered the collapse of a nearly-isothermal Bonnor-Ebert sphere with turbulence seeded by velocity perturbations. While these conditions were idealized from earlier hydrodynamic cosmological computations, we seek a more complete picture by following the full magnetohydrodynamic (MHD) evolution of primordial gas from cosmological initial conditions. This allows us to self-consistently study the interaction of primordial gas dynamics, turbulence, and dynamo action.  In this paper, we present the first highly-resolved calculations of the formation of the first stars in the Universe from cosmological initial conditions, taking into account a full suite of chemical reaction rates, chemical heating due to three-body reactions, magnetic fields from primordial seed fields, and a resolution of 64 zones per Jeans length.  Furthermore, it has been conducted with an open source, community-built simulation code available for inspection and contribution.  In this paper, we examine both the large-scale and small-scale amplification of the magnetic field, as well as the manner in which resolution affects the chemo-kinetic state of the inner molecular cloud.  A forthcoming paper will study the growth of turbulence in more detail. ",
        "conclusions": "We have presented on the first fully-cosmological calculations of Population III star formation that include all relevant chemical processes, as well as magnetic fields.  We see, in agreement with \\cite{2011ApJ...731...62F}, that a critical resolution exists above which we are able to resolve small-scale dynamo action resulting in increased magnetic-field field growth.  While we adequately resolve the \\textit{action} of this small-scale dynamo, we have not yet resolved the amplification caused by that dynamo, as demonstrated by the nascent J128 simulation.  In fact, while these calculations have yet to show a dynamically-important magnetic field, estimates of the possible saturation level of the magnetic field indicate that with increased resolution or stronger seed fields, magnetic fields may in fact become dynamically important.  A side effect of conducting these resolution studies has been that we see substantial variation in the chemical, kinetic, and velocity structure of collapsing metal-free, star-forming clouds.  In particular, the substantial difference in the J64 run suggests that under-resolving the hydrodynamics can result in incorrect solutions: the fraction of hydrogen gas that is in molecules, the speed of sound, the infall velocity and the turbulent support all depend strongly on hydrodynamic resolution.   Our calculations, while falling short of following potential fragmentation in detail, suggest that resolution effects may be an additional complication in determining the initial mass function of the first stars; in fact, they suggest that previous calculations in the literature have not yet converged on a solution.  Future calculations, of much higher resolution and attaining a greater saturation level of magnetic fields, will help to resolve these outstanding questions."
    },
    "1112/1112.1956.txt": {
        "abstract": "We present single-S\u00e9rsic two-dimensional model fits to $167,600$ galaxies modelled independently in the $ugrizYJHK$ bandpasses using reprocessed Sloan Digital Sky Survey Data Release Seven (SDSS DR7) and UKIRT Infrared Deep Sky Survey Large Area Survey (UKIDSS-LAS) imaging data available from the GAMA database. In order to facilitate this study we developed SIGMA, an R wrapper around several contemporary astronomy software packages including Source Extractor, PSF Extractor and GALFIT 3. SIGMA produces realistic 2D model fits to galaxies, employing automatic adaptive background subtraction and empirical PSF measurements on the fly for each galaxy in GAMA. Using these results, we define a common coverage area across the three GAMA regions containing $138,269$ galaxies. We provide S\u00e9rsic magnitudes truncated at $10$ $r_{e}$ which show good agreement with SDSS Petrosian and GAMA photometry for low S\u00e9rsic index systems ($n<4$), and much improved photometry for high S\u00e9rsic index systems ($n>4$), recovering as much as $\\Delta m=0.5$ magnitudes in the $r$ band. We employ a $K$ band S\u00e9rsic index/$u-r$ colour relation to delineate the massive ($n>\\sim2$) early-type galaxies (ETGs) from the late-type galaxies (LTGs). The mean S\u00e9rsic index of these ETGs shows a smooth variation with wavelength, increasing by $30\\%$ from $g$ through $K$. LTGs exhibit a more extreme change in S\u00e9rsic index, increasing by $52\\%$ across the same range. In addition, ETGs and LTGs exhibit a $38\\%$ and $25\\%$ decrease respectively in half-light radius from $g$ through $K$. These trends are shown to arise due to the effects of dust attenuation and stellar population/metallicity gradients within galaxy populations. ",
        "introduction": "\\label{sec:intro}The shapes and sizes of galaxies are not random but are defined by the orbital motions of their constituent stellar populations, arranging themselves into elliptical, bulge, disk, and bar like structures. Exactly why and how these structures come about is somewhat a mystery which no doubt relates to a complex formation history involving collapse, merging, infall, secular evolution, and feedback processes as well as the precise nature of the coupling between the dark matter, gas, dust and stars and the influence of the larger halo in which the galaxy might reside (group, cluster, etc.), and the broader environment (filament, void, nexus, etc.). The combination of variations in, for example, galaxy structure, formation history, evolution and relative environment lead to distinct measurable effects on global galaxy parameters such as colour, concentration and size. The ultimate goal of structural analysis is to inform this discussion by robustly isolating and quantifying these parameters and exploring correlations between these properties and those obtained by other means, such as dynamical information.  Once the underlying structure of a galaxy is understood the overarching morphological class may be determined, and from this we can explore correlations with, for example, environment through the well known morphology-density relation \\citep{Dressler1980}, i.e., the apparent preference for red, passive galaxies in the dense cores of galaxy groups and clusters. Several mechanisms have been suggested to explain this feature, most notably the combined effects of strangulation \\citep{Larson1980,Kauffmann1993,Diaferio2001}, ram pressure stripping \\citep{Gunn1972}, harassment \\citep{Moore1996} and tidal interactions and merging \\citep{Park2008}. Recent studies by, e.g., \\citet{VanderWel2008,Welikala2008,Welikala2009,Bamford2009} and others confirm this morphology-environment connection; however, they suggest that the relation between structure and morphology is less apparent. Indeed, it appears that the mass of a galaxy rather than the environment in which it resides is more influential in determining its structure, highlighting the importance of stellar mass estimates.  As an example of the connection between galaxy structure and the physical processes of galaxy formation \\citet*{Dalcanton1997} and independently \\citet*{Mo1998}, both following on from \\citet*{Fall1980}, relate the scale length of the disk to the angular momentum of a galaxy's dark matter halo. In addition, numerous properties of the bulge component are now known to relate to the mass of the super-massive black hole \\citep*[e.g.,][]{Haring2004,Novak2006,Graham2007}. Variations in structural properties as a function of wavelength \\citep[e.g.,][]{LaBarbera2010a} enable the extraction of colour gradients, potentially implying the direction of disk growth (e.g., inside out, \\citealt{Barden2005,Bakos2008,Trujillo2009,Wang2010}), or arguing for the redistribution of populations from the inner to outer regions \\citep{Roskar2008}, possibly coupled with bar formation \\citep{Debattista2006}.  The physics underpinning galaxy structure is relatively immature, despite the very long history dating back to Knox-Shaw, Reynolds and Hubble and essentially consists of spot-check simulations which focus on a particular phenomena in a non-cosmological context (for recent developments see \\citealt{Roskar2010} and \\citealt{Agertz2010}). For example, numerical models can readily produce bar, pseudo-bulge, spiral patterns, and spheroidal structures through coupled rotation, secular evolution, shock-wave propagation, and merging history. Until recently the very thin nature of the spiral disks has presented a particular challenge for numerical models, with numerical simulations in particular forming small thick disks, mainly because of the high level of merging \\citep{Navarro1997}. This is in stark contrast to a number of independent empirical studies \\citep[e.g.,][]{Driver2007,Gadotti2009,Tasca2011} which estimate that approximately $60\\%$ of the stellar mass in the universe today lies within disk systems, suggesting a more quiescent merger history (however, see \\citealp{Hopkins2010} on the stability of gas rich disks). In addition, studies by \\citet{MenendezDelmestre2007} suggest that up to $67\\%$ of spiral galaxies contain a barred structure, further complicating simulation efforts. However, numerical simulations are now starting to produce realistic disk systems \\citep{Governato2007,Governato2009,Agertz2009,Agertz2010} albeit with heavily controlled initial conditions, more quiescent merger histories and a greater degree of gas infall.  Beyond a distance of $\\sim100$ Mpc detailed structural studies have been relatively rare and mostly confined to the deep yet very narrow pencil beam surveys from the Hubble Space Telescope (HST). It was only following the refurbishment of HST that structural analysis once again became a mainstream study \\citep{Driver1995a,Driver1995b,Driver1998}. HST provides kpc resolution across the full path length of the Universe, which is now also becoming possible with AO ground-based systems \\citep{Huertas-Company2007}. The conjunction of development in numerical models and this new ability to resolve the shapes and sizes of galaxies at any distance has led to a dramatic renewed interest in structural analysis. One interesting claim is the apparent remarkable growth of galaxy sizes since intermediate redshifts (e.g., \\citealp{Trujillo2005,Barden2005,McIntosh2005,Trujillo2006,Trujillo2007,Weinzirl2011}), potentially supporting the notion of recent growth in disk systems following an earlier aggressive merger phase at $z\\sim2$ \\citep[see][]{Driver1996,Driver2005}. An alternative suggestion which does not require galaxy growth through mergers is the transformation of some of these so-called {}``red nuggets'' into the bulges of disk galaxies via the accretion of a cold gas disk \\citep{Graham2011}.  However, structural analysis is not trivial to implement and interpret correctly, and is plagued by a number of key issues. In particular: \\begin{enumerate} \\item Wavelength bias. At different wavelengths, light traces varying stellar populations (\\citealp{Block1999}). Typically this is a young stellar population at shorter wavelengths and an older stellar population at longer wavelengths. For this reason, it is vital when comparing structural properties to compare properties measured at the same rest wavelength. \\item Dust attenuation. Dust is predicted to modify not only the recovered total flux as a function of wavelength (e.g., \\citealp{Tuffs2004,Pierini2004}) but also galaxy sizes, shapes and concentrations (see for example \\citealp{Mollenhoff2006,Graham2008b}). Dust can vary enormously from system to system with significant environmental dependencies and strong evolution with redshift. Each individual galaxy ultimately requires either a dust correction or analysis at rest-NIR wavelengths where dust will have a smaller impact (see photon escape fraction curve in \\citealt{Driver2008}). It is obvious that any attempt to model the dust in galaxies raises the larger problem of degeneracies appearing between additive and subtractive flux components. \\item Local minima during the minimisation process. For a single profile fit there are often 7 free parameters, with that number rising for multi-component fits. The surface within this parameter space is known to be complex, containing multiple local minima representing potentially non-physical results, e.g., a bulge which contributes significantly more flux to the outer regions of a galaxy than the disk (\\citealt{Graham2001}). Other than manual checks of the output, various methods may be employed to reduce the risk of divergence on an incorrect result including; constraints applied during the minimisation routine and employing an automated logical filter (e.g., \\citealt{Allen2006}). \\item What lies below the limiting isophote. Whilst the surface brightness profile of some galaxies behaves as expected out to very faint magnitudes (e.g., NGC 300:\\citealp{Bland-Hawthorn2005,Vlajic2009}, NGC 7793: \\citealp{Vlajic2011}), the potential myriad of phenomena present in the outer wings of many systems may cause deviations away from a typical light profile. These include truncated and anti-truncated disks \\citep{Erwin2005,Pohlen2006}, UV excesses \\citep{Bush2010}, tidal debris, halos \\citep{Barker2009,McConnachie2009} and minor merger fossil records \\citep{Martinez-Delgado2010}. In fact, the outer regions of galaxies may defy any systematic profile fitting into a restricted number of structures. The accuracy of any estimation of the background sky and gradients therein will also no doubt affect analyses of these outer structures. \\item The number of components required. When considering the structure of very nearby galaxies, the deeper one looks the more one finds. Some galaxies, even in the dust free $3.8$ micron bands, require up to six components \\citep{Buta2010} before a satisfactory fit can be obtained. In many cases there is uncertainty as to how many components are required, how to quantitatively decide this in an automated and repeatable fashion, and which components are fundamental and which perhaps secondary. For example, should bar and pseudo-bulge flux be incorporated into a single disk model or kept distinct. \\item Sky estimation. Understanding the background sky level at the position of your primary object of interest is crucial in producing meaningful measurements of that galaxy. Considerations must be made in regards accuracy and speed of estimating the background. \\item Systematic selection bias. Sample bias will be introduced due to size, resolution, orientation, profile shape and smoothing scale limitations \\citep{Phillipps1986,Driver1999}. Samples of galaxies are usually selected based on global criteria, such as magnitude. However, it becomes non-trivial to transcribe these global limits into appropriate limits for galaxy sub-components, e.g., a certain type of disk may only have been detected because it also contains a prominent bulge. \\end{enumerate} There are several publicly-available galaxy modelling codes in common usage including GIM2D \\citep{Simard2002}, BUDDA \\citep{DeSouza2004}, GASPHOT \\citep{Pignatelli2006} and GALFIT 3 \\citep{Peng2010a}. In addition, there are a number of software pipelines, wrappers around contemporary astronomy software, that aim to automate the process of galaxy modelling including GALAPAGOS (Barden et al., 2011, submitted) and \\noun{PyMorph} \\citep{Vikram2010}. These packages all have their advantages and disadvantages and have been compared in a number of external studies (e.g. \\citealp{Haussler2007,Hoyos2011}) in addition to their own internal comparisons, and so we refer the reader to these papers for discussions of the pros and cons between 1D v 2D fitting and the actual minimisation algorithms employed. For this body of work GALFIT was chosen for its ease of use and high-quality realistic model outputs, plus the ability to perform simultaneous modelling of nearby neighbours to the primary galaxy.  In this series of papers we introduce and utilise SIGMA, an automated code designed to produce single-S\u00e9rsic and multi-component profile fits for galaxies in the GAMA dataset. Using SIGMA, this paper presents one of the largest catalogues of multi-wavelength single-S\u00e9rsic model fits; $167,600$ galaxies modelled independently across $9$ bandpasses. This catalogue is currently in use to aid in measurement of the evolution in the size-(stellar mass) distribution of galaxies (Baldry et al., 2011); explore star formation trends as a function of morphology (Bauer et al., in prep.); to further understand the cosmic SED from 0.1 micron to 1 mm (Driver et al., 2011); to apply dust corrections to galaxies observed at multiple inclinations (Grootes et al., 2012); to explore the dust properties and star-formation histories of local submillimetre selected galaxies \\citet{Rowlands2011}; better constrain stellar mass measurements by providing total flux corrections \\citep{Taylor2011}; comment on the quenching of star formation in the local universe (Taylor et al., in prep.); explore the relation between galaxy environments and their star formation rate variations (Wijesinghe et al., in prep.); provide a new method for automatic morphological classification (Kelvin et al., in prep.); and further explore the relation between environment (i.e., halo mass; \\citealp{Robotham2011}), morphology and structure (Kelvin et al., in prep.). Future studies building on SIGMA will incorporate advanced logical filtering and profile management to produce multi-component fits for a low redshift sample, allowing full structural decomposition into bulge-disk-bar, etc. (Kelvin et al., in prep).  This paper is organised as follows. We first outline the GAMA data in Section \\ref{sec:data}. We describe the SIGMA (Structural Investigation of Galaxies via Model Analysis) pipeline developed to process this data and produce robust 2D galaxy models in Section \\ref{sec:sigma} and present SIGMA single-S\u00e9rsic results for $167,600$ objects modelled independently in $ugrizYJHK$ from the GAMA Phase I study in Section \\ref{sec:output}. From this large catalogue we establish a common coverage sample of $138,269$ galaxies. Finally, we further explore the wavelength dependence on recovered structural parameters in Section \\ref{sec:wavelength}. A standard cosmology of $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}=0.3$, $\\Omega_{\\Lambda}=0.7$ is assumed throughout. ",
        "conclusions": "\\label{sec:conclusion}We have produced high-fidelity automated two-dimensional single-S\u00e9rsic model fits to $167,600$ galaxies selected from the GAMA input catalogue. These have been modelled independently across $ugrizYJHK$ using reprocessed SDSS and UKIDSS-LAS imaging data. These data have subsequently been delivered to the GAMA database in the form of the catalogue \\emph{SersicCatv07}. In order to facilitate the construction of this dataset, SIGMA, an extensive multi-processor enabled galaxy modelling pipeline, was developed. SIGMA is a wrapper and handler of several contemporary astronomy software packages, employing adaptive background subtraction routines and empirical PSF generation on a per-galaxy per-band basis to tailor input data into the galaxy modelling software GALFIT 3. Output results from GALFIT are analysed for pre-determined modelling errors such as positional migration, extreme model shape and/or size parameters and adverse nearby neighbour flux. Nearby object masking is employed as a last resort, with secondary neighbours being preferentially modelled simultaneously with the primary galaxy in the first instance.  Using this dataset, we have defined a common coverage area across the three GAMA regions that encompasses $138,269$ galaxies, $82.5\\%$ of the full sample. This common area contains only those galaxies which have been observed in all nine bands, providing a useful basis upon which to further explore wavelength trends. We define a S\u00e9rsic magnitude system that truncates S\u00e9rsic magnitudes at $10$ $r_{e}$. This ensures that flux is not extrapolated below the typical limiting isophote into regions where data quality and quantity is not sufficient to constrain the form of the galaxy light profile. Truncated S\u00e9rsic magnitudes appear to be a good descriptor of global galaxy colours and total galaxy flux. For well-resolved disk-like galaxies ($n<2$), traditional aperture-based methods are in good agreement with truncated S\u00e9rsic magnitudes. For high centrally-concentrated systems however ($n>4$), it appears that traditional aperture-based, such as Petrosian magnitudes, may miss as much as $\\Delta m_{r}=0.5$ magnitudes from the total flux budget which is only recovered through S\u00e9rsic modelling.  When considering the dataset in n--colour space we find galaxies appear to exist in two distinct groups. For the most massive systems, we associate these two groups with the spheroid-dominated early-type galaxy (ETG) and disk-dominated late-type galaxy (LTG) populations. Owing to the nature of our input sample selection, these definitions do not extend down to the fainter dwarf population, and so subsequent trends will not represent those systems. We use the longest wavelength $K$ band S\u00e9rsic index measurements in conjunction with rest-frame $u-r$ colour to define these two populations. Using these definitions, we are able to further probe the variations in recovered structural parameters with wavelength for each population.  We find that the S\u00e9rsic indices of ETGs remain reasonably stable at all wavelengths, increasing by $0.11$ dex ($+30\\%$) from $g$ to $K$ and becoming very stable beyond the $z$/$Y$ interface. In contrast to this, we find that LTGs exhibit larger variations in S\u00e9rsic index with wavelength, increasing by $0.18$ dex ($+52\\%$) across the same wavelength range. Recovered sizes for both the spheroid and disk systems show a significant variation with wavelength, showing a reduction in half-light radii of $0.20$ dex ($-38\\%$) in ETGs and $0.13$ dex ($-25\\%$) in LTGs from $g$ to $K$. Size variation of this scale for disk systems has been well predicted by dust models, highlighting the important role dust attenuation plays when considering structural variations across a broad wavelength range.  We note that spheroidal systems exhibit a larger size variation with wavelength than that found in disk systems. Possible physical explanations for this behaviour include low levels of unresolved dust or the effects of AGN feedback in the core of the galaxy, both of which would affect S\u00e9rsic profiling. Significant amounts of dust, such as an increased dust attenuation optical depth parameter $\\tau_{\\mathrm{B}}^{f}$, may allow current dust models to accurately describe the variation in half-light radii we find. It is unlikely however that a significant fraction of our spheroid-dominated population contain sufficient amounts of dust for this to be the case. Large stellar population/metallicity gradients present within individual structures of the galaxy would cause galaxies to look markedly different in different wavelengths, contributing to any concentration-wavelength/size-wavelength variation. In addition to these factors, uncertainties on the measured PSF and background sky must be considered.  However, when considering variations in half-light radius and S\u00e9rsic index together with wavelength we find that the large fluctuations in spheroidal parameters amount to a relatively modest impact on the recovered light profile. A comparatively larger effect is noted for the disk systems, particularly in the core region, supporting the presence and effect of dust attenuation in addition to stellar population/metallicity gradients. At a distance of $1$ pixel from the central region, spheroid systems display a variation in surface brightness of $0.49$ magnitudes from $u$ through to $K$. In disk systems, the comparative figure is $0.86$ magnitudes, an increase of $75\\%$. This highlights the importance of not considering recovered parameters in isolation, as the interplay between them has the possibility of masking underlying trends.  The effects of dust attenuation appear to be the dominant factor constraining the variations in structural parameters with wavelength, notably so for the disk-dominated population. In contrast with this, apparent large structural variations in the spheroid-dominated population appear to have a relatively minor effect on the underlying surface-brightness profile than might have been expected. Future studies in Kelvin et al. (2011; in prep.), focussing on a limited sub-sample of this dataset, will provide a deeper understanding of these structural variations with wavelength, enabling us to comment further on the key mechanisms involved in varying structural parameters with wavelength for a host of different morphologies."
    },
    "1112/1112.4223_arXiv.txt": {
        "abstract": "Any symmetry reduces a second-order differential equation to a first integral: variational symmetries of the action (exemplified by central field dynamics) lead to conservation laws, but symmetries of only the equations of motion (exemplified by scale-invariant hydrostatics) yield first-order {\\em non-conservation laws} between invariants. We obtain these non-conservation laws by extending Noether's Theorem to non-variational symmetries and present an innovative variational formulation of spherical adiabatic hydrostatics.  For the scale-invariant case, this novel synthesis of group theory, hydrostatics, and astrophysics allows us to recover all the known properties of polytropes and define a {\\em core radius}, inside which polytropes of index $n$ share a common core mass density structure, and outside of which their envelopes differ. The Emden solutions (regular solutions of the Lane-Emden equation) are obtained, along with useful approximations. An appendix discusses the $n=3$ polytrope in order to emphasize how the same mechanical structure allows different thermal structures in relativistic degenerate white dwarfs and zero age main sequence stars. ",
        "introduction": "Noether's Theorem relates every {\\em variational symmetry,} a symmetry of an action or similar integral, to a conservation law, a first integral of the equations of motion~\\cite{Bluman}. By an extension of Noether's Theorem, {\\em non-variational symmetries,} --- symmetries of the equations of motion which are not in general variational symmetries --- also lead to first integrals, which are not conservation laws of the usual divergence form, as discussed in a previous article~\\cite{BludKenI}. There it was shown that a Lagrangian $\\mathcal{L}(t,q_i,\\dot{q_i})$ and action $S=\\int{\\mathcal{L}(t,q_i,\\dot{q_i})} dt$, with degrees of freedom $q_i$, can be transformed under an infinitesimal point transformation $\\delta (t,q_i), \\delta q_j (t,q_i)$: \\bea \\delta \\mathcal{L}=\\dot{\\mathcal{L}}\\delta t+(\\partial\\mathcal{L}/\\partial q_i)\\delta q_i+(\\partial\\mathcal{L}/\\partial\\dot{q_i}) \\Bigl[\\frac{d \\delta q_i}{d t}-\\dot{q_i}\\frac{d \\delta t}{d t}\\Bigr]=\\Bigl [\\frac{dG}{dt}- \\mathcal{L}\\cdot\\frac{d(\\delta t)}{dt}+\\mathcal{D}_i\\cdot(\\delta q_i-\\dot{q_i}\\delta t)\\Bigr] \\quad, \\eea in terms of the total derivative of the {\\em Noether charge,} $G:=\\mathcal{L}\\cdot\\delta t+p_i\\cdot(\\delta q_i-\\dot{q_i}\\delta t),$ and the variational derivative $\\mathcal{D}_i:= \\partial\\mathcal{L}/\\partial q_i-d(\\partial\\mathcal{L}/\\partial\\dot{q_i})/dt$. For transformations that leave initial and final states unchanged, the variation in action is \\be \\delta S_{if}= G(f)-G(i)+\\int_i^f dt\\ \\Bigl[\\delta q_i\\cdot\\mathcal{D}_i+\\delta t\\cdot \\Bigl(\\frac{d \\mathcal{H}}{dt}+ \\frac{\\partial\\mathcal{L}}{\\partial t}\\Bigr) \\Bigr]\\quad , \\ee if the term in $d(\\delta t)/dt$ is integrated by parts. If the system evolution obeys an action principle, that this variation vanish for independent variations $\\delta q_i, \\delta t$ that vanish at initial and final times, the system obeys the Euler-Lagrange equations $\\mathcal{D}_i=0$ and $d \\mathcal{H}/dt=-\\partial\\mathcal{L}/\\partial t$, the rate of change of the Hamiltonian in non-conservative systems. On-shell, where $\\mathcal{D}_i=0$, \\bea \\delta S_{if}=\\int_i^f{ \\bar{\\delta}\\mathcal{L}}\\ dt=G(f)-G(i) \\\\ \\frac{dG}{dt} =\\bar{\\delta}\\mathcal{L}:=\\delta\\mathcal{L}+\\mathcal{L}\\cdot (d\\delta t/dt) . \\eea This is {\\em Noether's equation}, giving the evolution of a symmetry generator or Noether charge, in terms of the Lagrangian transformation that it generates. It expresses the Euler-Lagrange equations of motion as the divergence of the Noether charge. This divergence vanishes for a variational symmetry, but not for any other symmetry transformation.  Noether's equation (4) could have been derived directly from the definition of the Noether charge. But using the action principle make manifest the connection between Noether's equation and the Euler-Lagrange equations. We use the action principle and this connection to reformulate the theory of hydrostatic barotropic spheres, which is integrable if they are scale symmetric, even where this scale symmetry is not a symmetry of the action (Section II). The first integrals implied by any symmetry of the equations of motion, while generally not vanishing-divergence conservation laws, are still useful dynamical or structural first-order relationships.  Because it neglects all other structural features, scaling symmetry is the most general simplification that one can make for any dynamical system.  For the radial scaling transformations we consider, $\\delta r=r$, the Lagrangian scales as some scalar density $\\delta \\mathcal{L}=-2\\tilde{\\omega}\\mathcal{L};$ and the action scales as $\\delta S=(1-2\\tilde{\\omega}) S$. The Noether charge generating the scale transformation evolves according to a {\\em non-conservation law} $dG/dt =(1-2\\tilde{\\omega})\\mathcal{L}$, a first-order equation encapsulating all of the consequences of scaling symmetry~\\cite{BludKenI}. From this first-order equation follow directly all the properties of index-$n$ polytropes, as established in classical works~\\cite{Chandra,Schwarzschild}, modern textbooks~\\cite{Kippen,Hansen}, and the recent, excellent treatments of Horedt and Liu~\\cite{Horedt,Liu}.  Our secondary purpose is to present an original variational formulation of spherical hydrostatics and to extend Noether's Theorem to non-variational scaling symmetry, which yields a {\\em scaling non-conservation law} (Section II). For spherical hydrostasis, we define a {\\em core radius}, inside which all stars exhibit a common mass density structure. Outside this core, polytropes of different index $n$ show different density structures as the outer boundary is felt (Section III). Section IV completes the integration of the Lane-Emden equation by quadratures and obtains useful approximations to the Emden function $\\theta_n (\\xi)$.  An appendix reviews the thermodynamic properties of the physically important polytropes of index $n=3$~\\cite{Hansen,Kippen,BludKenI}. What is original here is the explanation of the the differences between relativistic degenerate white dwarf stars and ideal gas stars on the zero-age main sequence (ZAMS), following from their different entropy structures. Our original approximations to $\\theta_3 (\\xi)$ should prove useful in such stars. ",
        "conclusions": "%  We have explored how a symmetry of the equations of motion, but not of the action, reduces a second-order differential equation to first-order, which can be integrated by quadrature. In scale-invariant hydrostatics, the symmetry of the equations yields a first integral, which is a first-order equation between scale invariants, and yields {\\em directly} all the familiar properties of polytropes.  We observe that, like all stars, polytropes of index $n$ share a common core density profile and defined a {\\em core radius} outside of which their envelopes differ. The Emden functions $\\theta_n(\\xi)$, solutions of the Lane-Emden equation that are regular at the origin, are finally obtained, along with useful approximations.  The Appendix reviews the astrophysically most important $n=3$ polytrope, describing relativistic white dwarf stars and zero age main sequence stars. While reviewing these well-known applications~\\cite{Hansen,Kippen}, we stress how these {\\em same} mechanical structures differ {\\em thermodynamically} and the usefulness of our original (Section IV) approximations to these Emden functions.  \\appendix*"
    },
    "1112/1112.6226_arXiv.txt": {
        "abstract": "The ejection of matter in the close vicinity of a young stellar object is investigated, treating the accretion disk as a gravitationally bound reservoir of matter. By solving the resistive MHD equations in 2D axisymmetry using our version of the Zeus-3D code with newly implemented resistivity, we study the effect of magnetic diffusivity in the magnetospheric accretion-ejection mechanism. Physical resistivity was included in the whole computational domain so that reconnection is enabled by the physical as well as the numerical resistivity. We show, for the first time, that quasi-stationary fast ejecta of matter, which we call {\\em micro-ejections}, of small mass and angular momentum fluxes, can be launched from a purely resistive magnetosphere. They are produced by a combination of pressure gradient and magnetic forces, in presence of ongoing magnetic reconnection along the boundary layer between the star and the disk, where a current sheet is formed. Mass flux of micro-ejection increases with increasing magnetic field strength and stellar rotation rate, and is not dependent on the disk to corona density ratio and amount of resistivity. ",
        "introduction": "Highly collimated outflows have been observed from AGNs to Young Stellar Objects (YSOs) and young brown dwarfs \\citep{wh1,wh2}. Accreting compact stars, like accreting white dwarfs in symbiotic binaries \\citep{sk03} and neutron stars like Cir X-1 \\citep{hei07}, also show similar outflowing phenomena. An outflow is characterized as a jet if it is super-magnetosonic, collimated into an apparent narrow opening, and reaches a stationary or quasi-stationary state. Such high-velocity outflowing fluxes of matter are an integral part of stellar evolution. Observations in multiple wavelengths are reaching closer and closer to the objects that drive them.  Among all the systems, models of launching outflows in YSOs are closest to scrutiny by observations due to available data from star--forming regions. An accretion disk, through which matter accretes onto the young star with velocities close to a free--fall, is often associated with a jet--driving YSO \\citep{e94,e06}. A correlation between the accretion rate and the high-velocity jet power was found in many Classical T-Tauri Stars (CTTSs) \\citep{Cab90}. The ratio of mass loss in the outflow to disk accretion rate, $\\dot{M}_{\\rm w}/\\dot{M}_{\\rm a}$, extracted from observations is hard to constrain. It is best estimated to be approximately 0.1 through measurements of optical forbidden lines and veiling --- see e.g.\\ \\citet{hart95} and \\citet{edw08}. Recently, He~I~$\\lambda$10830 line has been used as a probe into the high-velocity winds originating from the inner region where the star interacts with the disk \\citep{e03,kw07}. Despite the potential diagnostic power of such emission lines, the actual structure and physical conditions of outflows can be more complex. Star-disk interaction is also investigated by means of X-rays, which offer deepest probe into the launching region. During flare events from the protostellar objects X-rays are emitted and have been observed (\\citet{fav05}, \\citet{gia06}, \\citet{aar10}, \\citet{mclw11}). Production of such X-rays is possible by either high temperature or high flow velocities, which can be related to the stellar surface, shocks in the stellar magnetosphere or, as we will suggest, to the reconnection events in the star-disk magnetosphere. Such events will leave imprint in the chemical and physical properties of the object \\citep{shuetal07}. To further interpret the observed line profiles, and the origins of outflows from the close vicinity of CTTSs, predictions from both theoretical and numerical models are required. \\begin{table*} \\caption{In many resistive-MHD simulations performed up to date, assumptions for initial conditions vary much and duration of simulations is very different. We list some of important works to position our work in the context. In the second column, ``fast'' stellar rotation means that corotation radius is smaller than the disk truncation radius, $R_{cor}<R_t$ and ``slow'' rotation means that $R_{cor}>R_t$. Disk is rotating in all the cases. Columns about physical resistivity and viscosity specify only if the option is at all included in the code; for where it is really effective, reader should check the related publication. When present, viscosity is important only in the disk. } \\begin{tabular}{|c||c|c|c|c|c|c|} \\hline & & & & & & \\\\ \\ \\  & $\\Omega_\\ast$ &  corona & $\\kappa=\\rho_d/\\rho_c$ & $T_{\\mathrm max}/T_0$ (days) & resistive & viscous\\\\ & & & & & & \\\\ \\hline & & & & & & \\\\ \\ \\  \\citet{hay96} & non-rotating & non-rotating & 10$^3$ & 5 & yes & no\\\\ \\ \\  \\citet{hir} & non-rotating & rotating & 10$^4$ & 16.5 & yes & no\\\\ & & rotating different than disk & & & & \\\\ \\ \\  \\citet{ms97} & slow & solid body rotation & 10$^2$ & 0.3 & yes & no \\\\ & & corotating with star at $R_{\\rm cor}$ & & & & \\\\ \\ \\ \\citet{gods99a} & fast & rotating & $10^4$ & 100 & yes & no \\\\ \\ \\  \\citet{rukl02} & slow & corotating with star & 10$^2$ & 100 & no & yes \\\\ & & for R$\\leq R_{\\rm cor}$, else with disk & & & & \\\\ \\ \\ \\citet{kuk03} & slow & not in hydrostatic balance, & 10$^3$ & 1000 & yes & yes \\\\ & & non-rotating & & & &  \\\\ \\ \\  \\citet{u06} & fast & corotating with star & 10$^3$ & 2000 & yes & yes \\\\ & & for R$\\leq R_{\\rm cor}$, else with disk & & & & \\\\ \\ \\  \\citet{R09} & fast & corotating with star  & 10$^4$ & 2000 & yes & yes\\\\ & and slow & for R$\\leq R_{\\rm cor}$, else with disk & & & & \\\\ \\hline & & & & & & \\\\ \\ \\  M\\v{C} et al. (present paper) & slow & corotating with star  & 10$^4$ & 1500 & yes & no \\\\ & & for R$\\leq R_{\\rm cor}$, else with disk & & & &  \\\\ \\hline \\end{tabular} \\label{tabla1} \\end{table*}  Outflows driven by energy derived from accretion are particularly appealing in the scenarios of jet launching. Many models have been proposed based on the concept of magnetocentrifugal wind mechanisms \\citep{bp82}, differing in the origins of the underlying magnetic fields and locations of matter launching. An outflow could be a disk wind driven by magnetic fields dragged in from the envelope or generated by the disk dynamo, or an inner disk wind anchored to the narrow innermost region as in the X-wind model powered by an enhanced dynamo from the star-disk interaction \\citep{shuetal94, shuetal97}, simultaneously with an accretion funnel \\citep{osshu95}. It might also be a stellar wind driven along the open field lines from the stellar surface by thermal or magnetic pressure \\citep{mattpu05,mattpu08}, or some combination of the different possibilities. Related to the launch of winds, magnetospheric accretion has been described in works by \\citet{koe91}, \\citet{osshu95} and \\citet{kold02} in the context of a magnetosphere interacting with the surrounding disk, sharing some similarities with the compact objects like neutron stars \\citep{ghla79a,ghla79b}. Except for the pure disk wind models, a magnetically connected star-disk system plays an important role in the making of the young stellar system and the evolution of angular momentum through the generation of outflows during the main phase of accretion.  Numerical investigations have been followed up on the time-dependent evolution of a system where the central star is magnetically connected to its accretion disk and their connection to jet formation and accretion. In one of the earliest attempts by \\citet{hay96}, where simulations of only a few rotation periods were obtained, a dipole magnetosphere corotating with the central star threaded the accretion disk that was in Keplerian rotation. Magnetic field lines connecting both the disk and the star inflate outwards due to shear, and reconnection blows out the matter along with the field, partially opening up the originally closed dipole loops. Gas can outflow from those opened field lines and might form part of the X-emission that is often associated with flares. Reconnection as a possible origin of X-rays from such systems has also been indicated in \\citet{EdP10}. \\citet{hir} investigated a magnetized star interacting with a truncated disk that was threaded with an initially uniform field dragged in from the outer core, in the same direction of the magnetosphere, but separated by a neutral current sheet  in the equatorial plane as a result of interaction between the fields brought together. For simplicity, the star was not rotating, but the differentially rotating disk could anyway provide enough shear to make the field inflate outwards, followed by a reconnection event and mass transfer onto the magnetosphere. The transferred mass diverted into two directions: one that falls onto the star and the other that flows out along the opened stellar field lines.  Longer simulations by \\citet{gods97} with an aligned dipole and a conducting accretion disk showed that differential rotation of the disk can drive episodes of loop expansion. Such expansion can drive two outflow components of gas: one hot convergent flow along the rotation axis, and another, slower cold flow on the disk side of the expanding loop. \\citet{ms97}, on the other hand, investigated interactions of magnetospheres with accretion disks under three different magnetic configurations and their respective dynamical evolution.  \\citet{kuk03} solved the disk in 1D with a radiative hydrodynamic code by \\citet{kley89}, and then extrapolated the solution to 2D as their initial condition. For the full 2D axisymmetric MHD problem, the induction equation, Lorentz force and Ohmic dissipation were now included into Kley's code, with the assumption of equal viscous and resistive dissipations. With $Pr\\sim 1$ magnetic field lines would not bend towards the axis. However, because of non-equilibrium initial conditions, they bunched close to the star. The main result was that, with the assumed mass accretion rate of $\\dot{M}_0=10^{-7}M_\\odot$~yr$^{-1}$, for a smaller magnetic field than 1~kG the disk was not disrupted; but for a larger field of the order of 1-10 kG, an outflow could be launched from the disk. In \\citet{rukl02} and \\citet{lrl05}, much better initial equilibrium has been set than in any previous simulations, with matter continuing to inflow through the disk because of viscosity. It was an improvement, as it was proceeding with the viscous time-scale, because of a slow accretion of matter. Without such initial equilibrium, the non-stationary initial conditions determine the flow in the disk, and influence the simulation. A star and part of the magnetosphere corotated up to the corotation radius, and the magnetosphere corotated with the disk farther out. The simulations were performed in the ideal MHD regime, with effective numerical resistivity diffusing magnetic field in the radial direction. They found funnel flows onto the central object, spinning up or down the star, depending on the ratio of rotation rate of the star to the rotation rate of the disk inner rim.  In \\citet{u06} and \\citet{R09} the effects of physical viscosity and resistivity on the outflow were investigated. In \\citet{R09} one of studied cases is with the magnetic Prandtl number, the ratio between the viscosity and the resistivity $\\mathrm{Pr}=\\nu/\\eta$, greater than one. Magnetic field lines are bent towards the magnetosphere in the gap and magnetic energy increases, enabling the outflow. It has been found that, in addition to the fast and light jet, there is another, new conical wind flowing up to 30 percent of the matter from the innermost portion of the disk. Two different cases were considered in their simulations, one for a fast (with the setup as in \\citet{u06} and a slightly different parameters), and the other for a slowly rotating star. In both cases, to enable the smooth start of the simulation, initially slow rotation of the star was gradually speeded up to its maximum value, with matter slowly inflowing from the outer boundary, to obtain stellar magnetic field compressed towards the magnetosphere in the gap. In the latter case, disk was not initially present in the computational box, but was formed from the matter inflowing from the outer boundary. Such setup was different from most other simulations in the literature.  With non-stationary initial conditions in the disk, accretion tends to be too fast. Most of the mentioned works involve resistive MHD models of accretion disks. Introduction of viscosity in the disk helps to obtain slow accretion, with a viscous time-scale. We do not include physical viscosity in our simulations, but we set the resistivity in the whole computational box, not only in the disk. Resistivity, which controls the onset of magnetic reconnection, triggers the necessary change in the magnetic field geometry needed for any launching. In our previous work, in \\citet{cf04}, we reported on the result (with the same code) when only a disk is present, without the magnetosphere of the central object taken into account. Propagation of the outflow through the resistive corona (with the disk set as a boundary condition) we investigated in \\citet{FC02}.  One important parameter to distinguish the investigated MHD regime is the already mentioned magnetic Prandtl number. As there is no physical viscosity included, our resistive simulations are in the regime of $\\mathrm{Pr}\\lesssim 1$. Violent initial conditions now helped to bring magnetic flux closer towards the star, helping the launching of matter outwards. Another parameter whose effect we study is the density ratio between the disk and the corona. It is usually included as a free parameter of the order of $10^2$ or $10^3$, at best $10^5$. We investigate the influence of this ratio on the mass and angular momentum flux in the launching of outflows. There are other possibilities in the setup, which we did not investigate here, e.g.\\ inclusion of the stellar wind, which would probably affect the open stellar field.  Table \\ref{tabla1} lists the kinematic and thermodynamic assumptions adopted in earlier works, each of them being usually repeated with a variety of parameters or methods, with or without physical resistivity and viscosity included in the code. We put our work in context of setups and assumptions of those works, as our results differ in the mass fluxes carried in the ejected gas.  The organization of the paper is as follows. We first describe our implementation of the boundary and initial conditions. In \\S\\ref{props} we report regimes we found under a broad range of parameters. We investigated the influence of corona to disk density ratio, strength of magnetic field and the physical resistivity. In \\S\\ref{recon} we address the role of reconnection in the launching, in \\S\\ref{elsase} we check a criterion for the site of launching, and in \\S\\ref{trunc} we compare position of the disk truncation radius in our simulations with some theoretical predictions. Then we discuss investigated parameters and the resulting ejections. ",
        "conclusions": "In this work, we for the first time demonstrate launching fast, light quasi-stationary flows of matter, which we call micro-ejections, from the magnetosphere above the gap in the star-disk system. In our simulations we included physical resistivity not only in the disk, but in the whole computational box. We investigate effects of anomalously large resistivity (when compared to microscopic resistivity), modeled as a function of matter density, to the launching of such micro-ejections in the case of slowly rotating star.  The physical resistivity, when included in previous simulations in the literature, was limited to the disk, decreasing effectively to zero out of it. The reconnection above the disk, which is necessarily happening during such simulations, was at the mercy of numerical resistivity, whose effects are different from physical resistivity In order to better understand differences between a variety of setups, we compare resistive simulations performed to date --- some of them are shown in Table \\ref{tabla1}. Resistivity can only moderately modify the flow shape, on a slower time scale than the formation of a magnetic wall \\citep{lyb96,lyb03}, so that if the resistivity is not large enough for reconnection to occur, the geometry of the flow will be different. As sketched in Figure \\ref{geommag}, common phenomena have been identified in our simulations, which occur, in those simulations of interacting star-disk systems from Table \\ref{tabla1} which have violent initial relaxation and/or weak magnetic field. After the relaxation of the system from unrealistic initial condition to a more evolved configuration, magnetic reconnection adjusts the topology of compressed field lines. Originally closed loops partially open. In this picture, the resistivity plays an important role, as it enables the reconnection.  We model physical resistivity as an anomalous turbulent resistivity dependent on the matter density $\\rho$, following \\cite{FC02}, with $\\eta=\\eta_0 \\rho^{1/3}$. Previously, resistivity has been included in the whole computational domain in \\citep{fen09}, but with the disk only as a boundary condition. Those simulations followed the propagation of outflow during few thousands of rotations of the inner disk radius. With the disk included in the computation, relaxation from initial conditions is much more complicated, and it is not easy to follow the evolution of the system for hundreds of rotations. We set up our resistive version of Zeus-3D code without employing any special procedure to prevent the violent relaxation from non-equilibrium initial conditions\\footnote{Such procedure could consist of the preparation of the particularly suitable disk model, as in \\citet{kuk03} or, as in the seminal paper of \\citet{R09}, slow introduction of the disk matter into the computational box, with gradual increase of the stellar rotation rate.}. Because of numerical difficulties with resistive MHD simulations in our setup with violent relaxation, our simulations tend to cease during relaxation, or last for too short after it for concluding about the results with a stable disk gap at a longer timescale. To perform the parameter study, we modified boundary conditions in the disk equatorial plane, to enforce a stable disk gap.  For a sufficiently large stellar magnetic field, a quasi-stationary micro-ejection forms. The magnetic field lines are opened above the star, with the more or less episodic, very light axial flow nearby the axis, which is artificial, of numerical origin. The magnetosphere between the star and the disk, along the boundary layer between the stellar and the disk field, is the site of launching of slower and more massive micro-ejection under a wider angle. The mechanism is similar to the launching of solar micro-flares, only that in our case magnetosphere of the star-disk system is the site of launching. Such events would leave trace in the chemical and physical properties of the object, as indicated in \\citet{shuetal07}. The role of reconnection in accretion disks with jets has been discussed also in \\citet{EdP10}. We will address this question in future work, together with the question of reconnection with different models of resistivity.  In our parameter study we investigated influence of the density contrast between the disk and corona, $\\kappa$. There is no difference in results for $\\kappa$ from $10^{3}$ to $10^{6}$. For $\\kappa=10^{2}$, the mass flux is always larger for a factor of few, and we conclude that results of simulations with $\\kappa=10^{2}$ could be unrealistic. Another parameter we checked was dependence of mass flux on stellar rotation rate. We found that mass flux increases with stellar rotation rate. Both fluxes also increase with increasing magnetic field--we vary the strength of the stellar dipole in the range of (0.1--200)\\,G for the disk accretion rate $10^{-6}M_\\odot/$yr. Both fluxes increase with increasing magnetic field. There is a limit to increase of magnetic field in simulation. A very strong initial field may cause strong shear motion at the beginning, and can largely disturb the relaxation process, especially in the case of a non-resistive corona, because pinching and reconnection of the magnetic field depend critically on the conditions in the magnetosphere \\citep{lov95,gods97}. Because resistivity enables the reconnection and reshaping of magnetic field, we also investigated if fluxes in micro-ejection depend on resistivity. We found that there is no such dependence. Without resistivity, numerical or physical, there is no reconnection, but once it occurs, the level of resistivity does not change the outcome of our simulations. Further study of this dependence is needed, with different resolutions. Our simulations do not scale well, and in the present setup we could not perform such study. In ideal MHD simulations performed with quasi-equilibrium initial conditions and slightly higher resolution than in our simulations --- see e.g.\\ \\citet{rukl02} --- numerical resistivity can mimic resistive effects, and such results could be more realistic than those in the high resolution ideal-MHD simulations \\citep{Y86}. We note that in high resolution simulations, physical resistivity should be included, or results could depend on numerical resistivity, which is not related to physical quantities of the setup. Any study of results is then necessarily flawed by purely numerical effects. Different models for resistivity should be investigated, too. A model for physical resistivity itself can affect the results, as was pointed out in \\cite{gods97}.  To identify locations in which the launching is possible in our computational box, we compare different indicators for a successful launching. The best such indicator in general is the line where the magnetic and matter pressure are equal ($\\beta=1$). We also find that in cases with weak magnetic field, the Elsasser number $\\Lambda$, could serve as an indicator for successful launching.  There is no disk component of flow in our results, and this is why they are very light. Matter in micro-ejections comes from the magnetosphere above the disk gap, which is just a small part of the matter from the disk inner radius, which slipped through the magnetic field lines."
    },
    "1112/1112.5539_arXiv.txt": {
        "abstract": " ",
        "introduction": "}   The Stellar Observations Network Group (SONG) is an initiative to design and build a global network of small telescopes. The goal of SONG is to become a key facility in both asteroseismology and planet search research programmes, a facility that provides state-of-the-art data of a quality that could not be achieved by use of any other space-based or ground-based facility \\cite{jcd06,gru06,gru08a,gru08b,gru09}.  SONG is proposed to have a total of eight nodes, four located in the northern hemisphere and four in the south. By placing the telescopes at roughly equally-spaced longitudes, long-term nearly continuous observations can be obtained. Each telescope has an aperture of 1m and will be equipped with a high resolution spectrograph for measuring very precise doppler velocities and dual--color lucky-imaging cameras for photometry of faint stars in crowded fields.  The prototype of SONG has been fully financed and a site on the Observatorio del Teide is being prepared for its installation. ",
        "conclusions": ""
    },
    "1112/1112.5363_arXiv.txt": {
        "abstract": "We study the stability of the (87) Sylvia system and of the neighborhood of its two satellites. We use numerical integrations considering the non-sphericity of Sylvia, as well as the mutual perturbation of the satellites and the solar perturbation. Two numerical models have been used, which describe respectively the short and long-term evolution of the system. We show that the actual system is in a deeply stable zone, but surrounded by both fast and secular chaotic regions due to resonances. We then investigate how tidal and BYORP effects modify the location of the system over time with respect to the instability zones. Finally, we briefly generalize this study to other known triple systems and to satellites of asteroids in general, and discuss about their distance from mean-motion and evection resonances. ",
        "introduction": "A large number of satellites of asteroids have been discovered since the discovery of the satellite Dactyl, thanks to the Galileo flyby of (243) Ida (\\citealt{belton1996}). As of today, there is 206 known systems (binary, triple and quintuple), following the Johnston's archive online database\\footnote{\\url{http://www.johnstonsarchive.net/astro/asteroidmoons.html}} (see also the online database\\footnote{\\url{http://www.asu.cas.cz/~asteroid/binastdata.htm}} described by \\cite{Pravec2007} and \\cite{Pravec2011}) and it is believed that small binaries could represent a fraction of 15\\% of the NEA population (\\citealt{Margot2002}; \\citealt{Pravec2006}).  Triple systems are rare and only nine known systems have been reported up to now in the entire Solar System.  The dynamical evolution and formation mechanisms of these systems are highly dependent on the size ratio between the secondaries and the primary. If this ratio is very small, as in the case of the Main-Belt asteroid (243) Ida, the systems are similar to the classical dynamical problem of a massless satellite orbiting a planet (see for example \\citealt{Kozai1959,Kozai1962}), this one being replaced by a possibly highly elongated ellipsoid (\\citealt{Chauvineau1993}; \\citealt{Scheeres1994}; \\citealt{Scheeres1996}; \\citealt{Compere2012a}). On the other hand, systems with similar size components, as the Near-Earth asteroid (66391) 1999 KW4, have to be described taking into account both their shapes and their rotations. A lot of studies have been realized on the expression of the full two-body problem and the study of its characteristics (\\citealt{Maciejewski1995}; \\citealt{Scheeres2002}; \\citealt{Fahnestock2008}; \\citealt{Boue2009}). Similarly, emphasis has been given during the past decade on the description of dissipative effects on binary systems, like tidal effects (\\citealt{Mathis2009}; \\citealt{Goldreich2009}; \\citealt{Taylor2010,Taylor2011}) or BYORP (\\citealt{Cuk2005}; \\citealt{Cuk2010}; \\citealt{Mcmahon2010}; \\citealt{Steinberg2011}).   We studied in this paper the dynamics and stability of the system (87) Sylvia, which was the first triple asteroid system discovered (\\citealt{Marchis2005a}). The specificities of this system place it in the first class described above. Sylvia, discovered in 1866, is a low-eccentric and midly-inclined asteroid located in the outer Main Belt. Its long-term evolution has been investigated through the AstDys project (\\citealt{Milani1998}; \\citealt{Knezevic2003}) giving its proper orbital elements ($\\overline{a}$ = 3.486 AU, $\\overline{e}$ = 0.0537, $\\overline{i}$ = $9.85^{o}$) and its secular fondamental frequencies ($n=55.297^{o}$/yr, $g = 134.798$\"/yr, $s = -130.782$\"/yr). Its orbit has been found to be slightly chaotic, exhibiting a Lyapunov time of $\\sim$ 1.4 Myr.  The two satellites of Sylvia present near-circular and near-equatorial orbits, and have a mass ratio of about $10^{-4}$ and $10^{-5}$ with Sylvia. The outermost satellite, Romulus, is approximately ten times more massive than the innermost one, Remus, for a semi-major axis twice as important. \\cite{Winter2009} studied the system and found that the satellites could be highly unstable when the oblateness of Sylvia (even a small fraction) is not taken into account. Indeed, the oblateness of the asteroid, as well as the short distance of the satellites from its surface ($\\sim$ 5 and 10 radius of Sylvia), critically increase the precession frequencies and prevent them from commensurabilities with frequencies arising from other gravitational perturbations.  Our aim is the understanding of the dynamical mechanisms present in the system and in its neighborhood. We then generalize some of the results to the other triple systems, and, in a general way, to the systems similar to (87) Sylvia, e.g. with a small size ratio and a primary diameter of the order of $\\sim$ 100 km. ",
        "conclusions": "We studied the dynamics and stability of the (87) Sylvia triple system by using numerical integrations of the complete and averaged equations of motion. We used a shape of Sylvia derived from light-curves observations and up to $C_{4,4}, S_{4,4}$ for the complete integrations. The position of the actual system lies in a very stable zone. We showed the possible evolutions of the system trough tidal and BYORP effects, and showed that the system, currently lying between the mean-motion resonances 2:1 and 3:1, will likely evolve through the evection resonance before the MMR 2:1 in the future. The other known triple system considered here, except (216) Kleopatra, also lie between the MMR 2:1 and 3:1. Finally, we show that the evection resonance could limit the outward evolution of the satellites.    \\paragraph*{Aknowledgments} The authors warmly thank A. Lemaitre for her constructive comments and her careful reading of the manuscript. Numerical simulations were made thanks to the local computing resources (Clusters ISCF and URBM-SYSDYN) at the University of Namur (FUNDP, Belgium). This work was supported by FAPESP (process n\\degre 2010/52715-5)."
    },
    "1112/1112.2775_arXiv.txt": {
        "abstract": "Using data mining techniques applied on emission line characteristics  of Be stars spectra  we attempted  to  find new Be stars candidates in SDSS SEGUE survey. The mid-resolution spectra of confirmed Be stars obtained from VO-compatible archive of Ond\\v{r}ejov observatory 2m telescope  were transformed to the spectral resolution of SDSS  and important characteristics of emission line profiles were estimated, to be used as a training base of supervised learning methods. The obtained knowledge base of the characteristic shapes and sizes of Be emission lines was finally used to identify new potential candidates in SDSS spectral survey. The several newly found Be stars candidates justify our  approach and approve Astroinformatics as a viable research methodology. ",
        "introduction": "Current data deluge in astronomy requires applying data mining techniques to extract new information about the physical nature of celestial objects. The possibility of cross-matching several surveys via Virtual Observatory protocols may play a key role in future discoveries. Data mining of large collections of spectra seems to be one of promising as well as challenging topics.  We have focused on obtaining new candidates of H$\\alpha$ emission stars using supervised data mining method of Decision Trees on almost 200,000 spectra in SDSS SEGUE \\citep{yanny2009segue} spectral survey. ",
        "conclusions": "\\begin{figure}[!htb] \\plotfiddle{P141_f4.eps}{6.2cm}{0}{70}{50}{-200}{20} \\caption{The example of candidate Be star found in SDSS SEGUE survey} \\end{figure}  \\noindent The classifier has identified 1110 Be stars candidates in SEGUE, however most of them are probably of different nature (e.g. AGNs, young stellar objects or reduction artifacts). Nevertheless, there are as well several highly probable  Be stars like the one on Fig.~3."
    },
    "1112/1112.0066_arXiv.txt": {
        "abstract": "In this work we propose a new orbital architecture for the two proposed circumbinary planets around the polar eclipsing binary HU Aquarii. We base the new two-planet, light-travel time model on the result of a Monte Carlo simulation driving a least-squares Levenberg-Marquardt minimisation algorithm on the observed eclipse egress times. Our best-fitting model with $\\chi_{r}^2=1.43$ resulted in high final eccentricities for the two companions leading to an unstable orbital configuration. From a large ensemble of initial guesses we examined the distribution of final eccentricities and semi-major axes for different $\\chi_{r}^2$ parameter intervals and encountered qualitatively a second population of best-fitting parameters. The main characteristic of this population is described by low-eccentric orbits favouring long-term orbital stability of the system. We present our best-fitting model candidate for the proposed two-planet system and demonstrate orbital stability over one million years using numerical integrations. ",
        "introduction": "In the past two decades, astrophysical timing measurements has been used to infer the existence of multiple low-mass planetary objects.  \\cite{WolszczanFrail1992} announced the first detection of a planetary system around the pulsar PSR1257+12. A two-planet system orbiting the short-period subdwarf B of the eclipsing binary HW Virginis was first presented in the works by \\cite{Lee2009}, and recently \\cite{Beuermann2010}, \\cite{Potter2011}, and \\cite{Doyle2011} announced the existence of two circumbinary planets in possible mean motion resonances around NN Serpentis, two giant planets orbiting the eclipsing binary UZ Fornacis, and a single circumbinary planet around the stars of the binary system Kepler 16, respectively.  The studies by \\cite{Beuermann2010} and \\cite{Potter2011} inferred the presence of additional massive objects by explaining the observed timing anomalies with the light-travel time (LTT, hereafter) effect. In the ideal case, the stellar components of an eclipsing binary system are orbiting their common centre of mass with a constant period. Timing irregularities can occur if the binary system is accompanied by an additional massive object. In this case, the binary centre of mass is orbiting the system's centre of mass. At some times the binary will be closer to the observer while at other times, it will be farther away, giving rise to the LTT effect on measured eclipse egress times.  In a recent study, the measurements of eclipse egress times of the eclipsing polar binary HU Aquarii (HU Aqr, hereafter) was used to infer the presence of a circumbinary planet around this system \\citep{Qian2011}.  These authors modeled the complete timing data set by adding the LTT effects from two circumbinary planets and found that the pericentre of the outer planetary companion is inside the orbit of the inner planet. This orbital architecture implies a crossing orbit configuration which points to strong mutual perturbations between the two planets. \\cite{Horner2011} subsequently carried out a detailed stability analysis of these bodies and found that almost all their initial conditions lead to unstable orbits on short-time scales. These authors concluded that the HU Aqr planetary system has most likely a different orbital architecture than proposed by \\cite{Qian2011}.  Three possibilities exist to explain the instability of the proposed planetary system. Either i) the LTT parameter space was not explored thoroughly while orbital stability constrains were imposed [omitted in \\cite{Qian2011}], or, ii) the applied LTT model was not complete and was unable to explain the timing data set properly, or iii) the data set is not yet large enough to draw reliable conclusions on the existence of additional low-mass companions.  In this work we aim to find a more plausible orbital architecture that best describes the timing data while also being conform with orbital stability requirements. We regard the stability condition as an observable that places additional constraints on the fitting process \\citep{Gozdziewski2005}. We carried out a large-scale Monte Carlo, least-squared parameter survey by fitting various LTT models to the complete data set. In parallel to the fitting process, we performed a stability analysis of various LTT model orbits which best describe the data set. By parameterising planet's Hill radii, we imposed orbital stability constraints to our parameter survey in order to explore orbital architectures that result in stable orbits. We extended the models by \\citet{Qian2011} to systems that also allow the inner planet to attain an eccentric orbit. Our results suggest that two-planet LTT models in general tend to produce orbits with higher eccentricities which are likely unstable.  The outline of this paper is as follows. In section 2 we present details of the adopted LTT model and provide a brief description of the least-squared minimisation algorithm. In section 3 we introduce our orbital stability constrains and in section 4 we describe results of numerical experiments. Finally, we give a summary and discussion in section 5. ",
        "conclusions": "In a recent study \\cite{Qian2011} announced the discovery of a two-planet system around the eclipsing polar binary HU Aquarii. In a subsequent paper, the stability of their best-fitting orbital parameters was examined by \\citet{Horner2011} who showed that the system is highly unstable. The most likely reason for this instability is the high eccentricity of the outer companion. In an attempt to find a system with stable orbits, we calculated a large number of two-planet LTT models using a Monte Carlo approach. This approach provides statistics on the distribution of the final goodness-of-fit $\\chi_{r}^2$ parameter in combination with constrains on requiring initial orbital stability.  In this work, we analyzed the complete available timing data of HU Aquarii. We made the following assumptions in our analysis. First we allowed the inner planet to attain an eccentric orbit in response to possible eccentricity excitation by an outside perturber. Compare to the study by \\cite{Qian2011}, this assumption resulted in introducing three extra free parameters. We also did not include a secular term which would account for timing variations due to mass-/angular momentum transfer and/or magnetic activity of the binary component. \\citet{Qian2011} argue that the secular term in their work is most likely caused by the presence of a third circumbinary massive object. However, the existence of this body in their data is far from obvious as the observational time span is much smaller than the orbital period of the proposed third circumbinary (planetary) companion. The confirmation of the third planet would require a few decades of future photometric follow-up monitoring. Finally, we did not consider mutual gravitational interactions between the two planets or gravitational perturbations on the binary stars.  Our best fit model had a goodness-of-fit parameter of $\\chi_{r}^2=1.43$ occurring fourteen times out of 111.844 initial guesses. From studying the histogram distribution of final $\\chi_{r}^2$ parameters, we were faced with a dichotomy. On the one hand, our best-fitting model of $\\chi_{r}^2=1.43$ had a low occurrence frequency (Fig.~\\ref{chi2histogram}) and resulted in high eccentricities (Fig.~\\ref{e-distribution}a). On the other hand, we encountered a significant higher occurrence frequency of larger $\\chi_{r}^2$ values. At the moment we have no explanation for this trend. One possibility is the inability of the LM minimisation algorithm to escape a local minimum in the non-linear parameter space. To test this possibility, one can use a least-squared minimisation procedure based on a genetic algorithm (GA) or a Bayesian Markov Chain Monte Carlo method.  The results of our study suggest that we should be able to trust the nominal $\\chi_{r}^2=1.43$ best-fitting parameters as this fit represents, most likely, the global minimum in $\\chi_{r}^2$ parameter-space, and provides the best possible description of the observed timing data set. However, as we have shown, the resulting best-fitting orbit with $\\chi_{r}^2=1.43$ was unstable. Models in the neighbourhood of this system were also unstable with a collision or ejection taking place on very small time scales. This instability can be attributed to high values of the final fitted eccentricities for one or both planetary companions. It is worth mentioning that models with goodness-of-fit close to the nominal best-fitting model did not allow low-eccentricity/circular orbits which would have rendered them most likely to be stable. We would like to mention that other authors have encountered similar (in)stability problem for the two planets in the systems of NN Serpentis and UZ Fornacis as well \\citep{Potter2011, Beuermann2010}. The two-planet system HW Virginis \\citep{Lee2009} also appears to be unstable (B. Funk, private communication). In particular, it is interesting to note that the best-fitting models of the proposed UZ Fornax circumbinary planetary system also attained high final eccentricities \\citep{Potter2011}.  Examining the $\\chi_{r}^2$ frequency distribution in detail resulted in the emergence of a second population of models with fitted parameters that fulfill our orbital stability constraints. Orbital stability was examined for models with $\\chi_{r}^{2}=1.89$. Although this population had a larger $\\chi_{r}^2$ goodness-of-fit than the nominal case, we were able to find stable orbits due to significantly lower orbital eccentricities. Combining/pairing a least-squared minimisation algorithm with a stability analysis enabled us to determine a new set of orbital elements for the two planets which are different from those given by \\citet{Qian2011}. The latter is another implication of our study which points to the existence of a variety of stable configurations when considering models with increasingly larger $\\chi_{r}^2$.  Assuming that no other physical effects are capable of explaining the observed timing variations in the light curve of HU Aquarii, we are left with the two-planet model as the most probable cause for the observed timing variation. Our study indicates that single one-planet LTT sinusoidal timing variations are inadequate to describe the observational data. However, it is necessary to mention that timing anomalies could also be caused by direct perturbations of a circumbinary planet on the binary orbit, though these perturbations would introduce a high-frequency (short-period) component to the eclipse timing variations due to the short orbital period of the binary system. A second possibility worthy to consider in future work would be to include mutual perturbations between the two proposed circumbinary planets. Explaining eclipse timing variations caused by the latter effect has not been explored for circumbinary planetary systems.  Finally, a similar study by \\cite{Wittenmyer2011} also addresses the instability of the two proposed companions around HU Aqr from revised LTT models. In their work the authors point out the need to obtain a better understanding of the mutual interaction between the two binary components. Binary orbital period modulation caused by magnetic interaction and/or mass-transport between the two binary components could also possibly explain the observed timing anomaly for HU Aqr. However, a combination of the above effects might also provide a satisfactory description of the observed timing data. At this point, we recommend future discoveries of multi-planet circumbinary systems also include a preliminary orbital stability study as a necessary condition to increase the likelihood of the existence of such systems. At the moment it seems that best-fitting single LTT models in superposition ($\\tau_1 + \\tau_2$) are inadequate to result in stable multi-planetary circumbinary systems. Though such a model provide a good description to the observational timing data. Furthermore, we also recommend future studies to obtain and publish timing data of the secondary eclipse (if available). Any timing variations introduced from gravitational interactions and/or LTT should also provide an adequate description of timing measurements obtained from the secondary eclipse.  Further constraining of our model requires obtaining additional timing data of the HU Aqr system. Using the results of our study involving stable planetary orbits, we predict that future mid-egress timing measurements would result in $O-C$ timing differences of -10.0 to 0.0 seconds (predicted at epoch $E = 40000$ in  Fig.~\\ref{BestFitOmC_2}). We emphasise that high-accuracy timing measurements are crucial to unveil the true nature of the observed timing anomalies as well as the development of improved models describing eclipse (egress) timing variations of a binary star system."
    },
    "1112/1112.2580_arXiv.txt": {
        "abstract": "LOFAR is a groundbreaking low-frequency radio telescope currently nearing completion across northern europe. As a software telescope with no moving parts, enormous fields of view and multi-beaming, it has fantastic potential for the exploration of the time-variable universe. In this brief paper I outline LOFAR's capabilities, as well as our plans to use it for a range of transient searches and some crude estimated rates of transient detections. ",
        "introduction": "LOFAR, the Low Frequency Array, is a large, low-frequency radio telescope in northern Europe, led by ASTRON. Construction of the array, which has its core collecting area in The Netherlands, with international stations in France, Germany, Sweden and The UK, is nearly complete, and astronomically interesting data are now being taken. LOFAR operates in the 30--80 and 120--240 MHz frequency ranges. The 80--120 MHz frequency gap corresponds to the FM radio bands at which frequencies astronomical observations would be impossible\\footnote{Unless northern europe could be persuaded to stop night-time FM radio broadcasts for a few weeks in the interests of finding the {\\em Epoch of Reionisation} signal..}. Construction of the array is almost complete -- See Fig \\ref{europe} for the distribution of operating LOFAR stations across Europe.  In addition, observations are occasionally possible to frequencies as low as 15 MHz. For a full reference paper on LOFAR see van Haarlem et al. (2012).  \\begin{figure}[t] \\begin{center} \\includegraphics[width=11.0cm]{europe.eps} \\caption{The distribution of complete LOFAR stations across Europe.} \\label{europe} \\end{center} \\end{figure}  LOFAR has six {\\em key science projects} (KSPs), one of which is {\\em Transients} (principal investigators Fender, Stappers, Wijers).  The remit of the TKSP covers essentially all transient and variable astrophysics, including commensal searches of all data (ultimately in near-real-time, although this functionality is not yet implemented). The TKSP covers both time-series and image-plane searches for transients and variables, including pulsars (Stappers et al. 2011). The adoption of transients and variables as key science drivers for the project is a theme for most of the large SKA pathfinders and precursors, in general unlike older radio facilities. However, time-series and image-plane transients have been separated for both ASKAP (which has {\\em CRAFT} and {\\em VAST} respectively) and MeerKAT ({\\em TRAPUM} and {\\em ThunderKAT}). This makes some sense from a techniques point of view, although there is some overlap in the science. ",
        "conclusions": "LOFAR has a wide range of diverse capabilities (multi-beaming, simultaneous timing and imaging modes, splitting of array in large numbers of individual stations, re-imaging in the past with the TBBs) which are ideally suited for exploring transient parameter space. Just as importantly, there is the will to support this science, with a Key Science Project dedicated to precisely this exploration. Based on current estimates, it should find many 1000s of interesting radio transients per year, providing a huge target list for multiwavelength follow-up, and providing new tests of our ability to automatically detect, classify and report such events efficiently."
    },
    "1112/1112.2549_arXiv.txt": {
        "abstract": "One of the main debated astrophysical problems is the role of the AGN feedback in galaxy formation. It is known that massive black holes have a profound effect  on the formation and evolution of galaxies, but how black holes and galaxies communicate is still an unsolved problem. For Radio Galaxies, feedback studies have mainly focused on jet/cavity systems in the most massive and X--ray luminous galaxy clusters. The recent high--resolution detection of warm absorbers in some Broad Line Radio Galaxies allow us to investigate the interplay between the nuclear engine and the surrounding medium from a different perspective. We report on the detection of warm absorbers in two Broad Line Radio Galaxies, 3C~382 and 3C~390.3, and discuss the physical and energetic properties of the absorbing gas.  Finally, we attempt a comparison between radio--loud and radio--quiet outflows. ",
        "introduction": "In the last few years, high--resolution X--ray spectroscopy has made progress in the exploration of the circumnuclear environment of radio--loud (RL) AGNs. While the presence of X--ray emitting and absorbing gas is well established in Seyfert galaxies, for RL sources  the investigation of the nuclear environment through this technique is very recent. Specifically, in Broad Line Radio Galaxies (BLRG), the RL counterpart of Seyfert 1s, the detection of warm absorbers (WA) \\footnote{With the term ``warm absorber'' we intend ionized outflowing gas in our line--of--sight that produces narrow absorption lines in the soft X--ray spectrum.}  was expected to be more difficult because of their small number in the local Universe and  because the Doppler amplification of the jet emission could mask the absorption features. However, steps forward have been recently made thanks to high--resolution X--ray spectroscopy. Here we summarize the results concerning the discovery of WAs in two BLRGs, 3C~382 and 3C~390.3. ",
        "conclusions": "We report on the detection of WAs in two BLRGs, 3C~382 and 3C~390.3, and discuss the physical and energetic properties of the absorbing gas: {\\bf (i)} the outflows are highly ionized (log$\\xi >$2~erg~cm~s$^{-1}$) and slow, with velocities ranging between 10$^{2}$--10$^{3}$~km~s$^{-1}$; {\\bf (ii)} the mass outflow rates are higher than the mass accretion rates if a volume filling factor (C$_{\\rm v}$) equal to 1 is assumed. Therefore a gas clumpy configuration (C$_{\\rm v} <$1) is expected; {\\bf (iii)} the kinetic luminosity associated to these slow outflows is always lower than the accretion luminosity and the jet kinetic power; {\\bf (iv)} although RL and RQ WA physical properties appear very similar, at least at zeroth order, the mass outflow rate and the radio--loudness parameter (R) seem to be correlated. This correlation could indicate a different gas distribution or alternatively, if the gas distribution is the same, powerful jets could favor the escape of more massive winds.  \\begin{figure}[pbh] \\begin{center} \\psfig{file=Torresi_1_f2.eps,width=3.7 cm, angle=0} \\vspace*{8pt} \\caption{Mass outflow rate plotted against the radio--loudness parameter (R). {\\it Red circles}: BLRGs considered in this work; {\\it black squares}: type 1 RQ AGNs belonging to our reference sample. \\label{f2}} \\end{center} \\end{figure}"
    },
    "1112/1112.0299_arXiv.txt": {
        "abstract": "Reports of the death of the precursor of Supernova (SN) 1961V in NGC 1058 are exaggerated.  Consideration of the best astrometric data shows that the star, known as ``Object 7,'' lies at the greatest proximity to SN 1961V and is the likely survivor of the ``SN impostor'' super-outburst. SN 1961V does not coincide with a neighboring radio source and is therefore not a radio SN. Additionally, the current properties of Object 7, based on data obtained with the {\\sl Hubble Space Telescope}, are consistent with it being a quiescent Luminous Blue Variable (LBV). Furthermore, post-explosion non-detections by the {\\sl Spitzer Space Telescope\\/} do not necessarily and sufficiently rule out a surviving LBV. We therefore consider, based on the available evidence, that it is yet a bit premature to reclassify SN 1961V as a {\\it bona fide\\/} SN. The inevitable demise of this star, though, may not be too far off. ",
        "introduction": "The evolution of the most massive stars is not well known. It is thought that main sequence stars with $M_{\\rm ZAMS} \\gtrsim 30\\ M_{\\sun}$ proceed through a blue supergiant phase, possibly to a red supergiant phase, or straight to a luminous blue variable (LBV) phase, and to the Wolf-Rayet (WR) phase prior to explosion as supernovae (SNe). That WR stars typically possess $M\\lesssim 20\\ M_{\\sun}$ \\citep{crowther07} requires their precursor stars to shed most of their mass, presumably through eruptive mass ejections as a LBV \\citep{humphreys94,smith06}. The term ``supernova impostor'' has been coined \\citep{vandyk00} to describe these eruptive events for massive stars, such as $\\eta$~Carinae \\citep[e.g.,][]{davidson97}. This is because various observational aspects of the eruptions can mimic the properties of true SNe. During $\\eta$~Car's Great Eruption in the 1800s, the star greatly exceeded the Eddington limit, with the bolometric luminosity increasing by $\\sim$ 2 mag. The total luminous output of such an eruption ($\\sim 10^{49}$--$10^{50}$ erg) can rival that of a SN.  However, a SN, by strict definition, is an explosive event that ends the life of a star (although a compact object may form in the process). If an optical transient is observed with energetics comparable to that of a true SN, and, after a sufficient period of time has elapsed, a star still exists at the exact position of the transient, then that transient is a SN impostor.  The most studied example of a SN impostor, SN~1961V in NGC 1058, remains controversial to this day. A debate has continued as to whether or not this event was a true SN.  In fact, it is truly impressive how many journal pages have been spent on this one object. This includes two recent studies, both of which present arguments that claim SN 1961V as a genuine SN. We argue here that SN 1961V is an impostor and that its precursor has survived what was a powerful eruption. Many of our points below were first presented in our review of SN impostors \\citep{vandyk11}. ",
        "conclusions": "Based on our analysis of the available data, we consider SN 1961V (still) to be a SN impostor and that Object 7 is the survivor of this event. All of the arguments we have made above, as part of this analysis, support this conclusion. We have also attempted to dispel several erroneous suppositions about the nature of this event. The survivor is clearly observed in very recent years to exist. The star quite plausibly has the properties of a quiescent, massive LBV. The star has not yet exploded and is not a radio SN.  The star is surrounded by a circumstellar shell or nebula, however, this shell is not nearly as dusty as required by \\citet{kochanek11} and previous investigators. We find that the visual extinction to Object 7 is in the range $A_V=1.8$--2.3 mag, and we suggest that most of this extinction is arising from the interstellar medium along our line-of-sight within the host galaxy, rather than from the shell itself. As a result, therefore, the shell need not be nearly as luminous an IR emitter as $\\eta$ Car. This conclusion is further supported by the mid-IR observations of other SN impostors.  The positional proximity of Object 7 to SN 1961V, and the positional offset between SN 1961V and the radio source centroid, are both inescapable facts. The only conceivable way to counter the former fact is to conclude that Object 7 is either a physical or an optical double to the progenitor of an actual SN. This, of course, is possible. The uncertainty alone in the absolute position, $0{\\farcs}1$, corresponds to $\\sim 5$ pc, certainly allowing for the possibility, in the relatively crowded cluster environment of SN 1961V, that Object 7 is merely a neighbor to the progenitor. Also, the probability of core-collapse SNe for high-mass stars in binaries is relatively high \\citep[e.g.,][]{kochanek09}. Although, if Object 7 {\\em were\\/} the binary companion to the star that exploded, we might expect Object 7 to have been stripped or otherwise affected by what would have been a very powerful explosion, potentially leading to an unusual brightness or color. This is not supported by the available photometry or inferred luminosity of Object 7. Instead, it appears to have the properties of a ``run-of-the-mill,'' evolved, high-mass star. Nonetheless, the positional offset between SN 1961V and the radio centroid is insurmountable --- these two are not one and the same, which is essentially at the heart of the case that has been made for SN 1961V being a true SN.  No need exists for the supposed core-collapse explosion ``hybrid'' of a SN II-P and SN IIn. Furthermore, we disagree with the statement made by \\citet{smith11} that  SN 1961V is somehow unique among the impostors. It has direct analogs in both SN 2000ch \\citep{wagner04,pastorello10} and SN 2009ip in NGC 7259 \\citep{smith10,foley11}, specifically in terms of the inferred expansion velocities at eruption and of the light curve behavior. The recent, more powerful revival of SN 2000ch displays ``unusually'' high velocities (FWHM) of 1500--2800 km s$^{-1}$ \\citep{pastorello10}. SN 2009ip shows blueshifted He\\,{\\sc i} absorption at 3000--5000 km s$^{-1}$, which \\citet{smith10} speculate for SN 2009ip could be due to a fast blast wave, which occurred quasi-contemporaneously with the origin of the slower ejecta. SN 2000ch exhibits multiple P Cygni absorption components, with profile edges up to 3000--3500 km s$^{-1}$. A sustained high-luminosity pre-eruption state was seen for both SNe 2000ch and 2009ip; SN 2000ch was at $M_R\\simeq -10.7$ mag prior to the 2000 eruption \\citep{wagner04}, and SN 2009ip had a pre-eruption luminosity of $M_V\\simeq -10$ mag \\citep{smith10}. Although not nearly the $M_{\\rm pg}\\approx -12$ mag brightness of SN 1961V, both SNe 2000ch and 2009ip, therefore, both may also have been in a super-Eddington phase prior to their giant eruptions. Even though the line widths were different, it is interesting to note that the H$\\alpha$ line profiles, in particular, of SN 1961V \\citep[][see his Figure 2]{zwicky64} and SN 2000ch \\citep[][see their Figure 7]{wagner04}, near maximum of the 2000 eruption, were very similar --- a sharp drop-off to the blue wing and extended wing to the red. The H$\\alpha$ emission line in 2002 for Object 7/SN 1961V also showed a similar asymmetric profile (and, we emphasize that this profile had a very similar shape as the much broader line profile for SN 1961V in 1962, as shown by \\citeauthor{zwicky64}). A possible explanation could be the effects of dust extinction in the expanding ejecta, although we have shown that the extinction from the ejecta is likely relatively modest for SN 1961V. A detailed analysis of this effect should be explored, although we consider this to be beyond the scope of this paper.  We do agree with \\citet{smith11} and \\citet{kochanek11}, however, that the ``undulations'' in SN 1961V's post-maximum light curve may well have arisen from the blast wave overtaking previously-ejected shells of matter ahead of the shock. We speculate here that the star was already in a sustained, eruptive outburst prior to 1960 and that the onset of the super-outburst, which peaked in luminosity in 1961 December, could possibly have been due to the interaction of the fast-moving ($\\sim$2000 km s$^{-1}$), dense, massive shell with the pre-existing, slower, less dense mass loss. A potential analog is the behavior of the SN IIn 1994W in NGC 4041 --- \\citet{dessart09} modeled the high luminosity, sustained plateau, and sudden drop-off of this SN as the interaction of a fast ($\\gtrsim 1500$ km s$^{-1}$), dense shell interacting with a slower, less dense shell. Furthermore, both the brief and more sustained plateaus in the SN 1961V light curve could have arisen from the interaction of the fast blast wave with various regimes of previously ejected matter, or to periods of partial recombination in the expanding ejecta, analogous to the recombination wave in SN II-P ejecta. Alternatively, these plateaus could also have been due to subsequent lesser eruptions by the star (Humphreys et al.~\\citeyear{humphreys99}). The fast-moving massive shell radiated for many years after the super-outburst, albeit more faintly. \\citet{goodrich89} detected the broad emission-line component at $F_{{\\rm H}\\alpha} \\simeq 2.2 \\times 10^{-16}$ erg cm$^{-2}$ s$^{-1}$ in 1986. We can set a limit on detection of this component in the {\\sl HST}/STIS spectrum from 2002 at $F_{{\\rm H}\\alpha} < 3 \\times 10^{-16}$ erg cm$^{-2}$ s$^{-1}$ (see Figure~\\ref{figspec}).  \\bibpunct[; ]{(}{)}{,}{a}{}{;} It is correct to point out the similarities in the variety of SN impostor and SN IIn properties. The two are very likely intimately intertwined, as we now know that at least one SN IIn arose from a SN impostor \\citep[SN 2005gl;][]{galyam07,galyam09}. The SN IIn are also a diverse class of objects (if the term ``class'' is truly applicable in the case of such heterogeneity), and it is still not clear how many SNe IIn have, in fact, really been impostors. For instance, SN 1994W itself may be an impostor \\citep{dessart09}.  Obviously, we have extrapolated fairly extravagantly from two photometric measurements of Object 7 from recent {\\sl HST\\/} data. We have had to assume that the colors for the star from 1991 are applicable to 2008 as well. What is ultimately required is a set of deep, multiband optical and near-IR observations of the SN 1961V environment using the modern {\\sl HST} with the Wide Field Camera 3, which would vastly improve upon the pre-furbishment WF/PC-1 images and the motley assortment of WFPC2 and unfiltered STIS image data analyzed to date. Such future observations would only require a rather modest investment of {\\sl HST\\/} time --- to image in $BVRIJH$ at $S/N\\gtrsim 10$ would require only 3 orbits. Looking beyond {\\sl HST}, observations with the {\\sl James Webb Space Telescope\\/}, particularly with the Mid-IR Instrument (MIRI), would reveal the actual dust emission from the surviving star. This object has clearly garnered much attention and speculation over the decades and continues to provide us with invaluable insights into the evolution of the most massive stars. With these space-based data, the definitive nature of the SN 1961V precursor may well be established once and for all.  \\bibpunct[; ]{(}{)}{;}{a}{}{;} Finally, NGC 1058 should continue to be regularly monitored by SN searches in nearby galaxies. For Object 7 as the SN 1961V survivor, we might expect, much as was the case for SN 2006jc \\citep{foley07,pastorello07}, that the star will {\\em truly\\/} explode as the SN that we believe other authors have erroneously concluded has already occurred. The super-outburst from the 1960s could well be the forerunner of a core-collapse event. It is possible that the star will explode as a high-luminosity, hydrogen-rich SN IIn, such as SN 2006gy \\citep{ofek07,smith07b}, or as a more helium-rich SN Ibn, such as SN 2006jc \\citep[as suggested by][]{kochanek11}."
    },
    "1112/1112.2255_arXiv.txt": {
        "abstract": "{}{We analyze coronagraph observations of a polar jet observed by the Sun Earth Connection Coronal and Heliospheric Investigation (SECCHI) instrument suite onboard the Solar TErrestrial RElations Observatory (STEREO) spacecraft.} {In our analysis we compare the brightness distribution of the jet in white-light coronagraph images with a dedicated kinetic particle model. We obtain a consistent estimate of the time that the jet was launched from the solar surface and an approximate initial velocity distribution in the jet source. The method also allows us to check the consistency of the kinetic model. In this first application, we consider only gravity as the dominant force on the jet particles along the magnetic field.} {We find that the kinetic model explains the observed brightness evolution well. The derived initiation time is consistent with the jet observations by the EUVI telescope at various wavelengths. The initial particle velocity distribution is fitted by Maxwellian distributions and we find deviations of the high energy tail from the Maxwellian distributions. We estimate the jet's total electron content to have a mass between $3.2\\times10^{14}$ and $1.8\\times10^{15}$ g. Mapping the integrated particle number along the jet trajectory to its source region and assuming a typical source region size, we obtain an initial electron density between $8\\times10^9$ and $5\\times10^{10}$ cm$^{-3}$ that is characteristic for the lower corona or the upper chromosphere. The total kinetic energy of all particles in the jet source region amounts from $2.1\\times10^{28}$ to $2.4\\times10^{29}$ erg.}{} ",
        "introduction": "Polar coronal jets were originally observed by the Extreme-Ultraviolet Imaging Telescope (EIT) onboard the Solar and Heliospheric Observatory (SOHO)~ \\citep{Domingo:etal:1995}. In white light, the polar jets discovered by LASCO/SOHO \\citep{StCyr:etal:1997} appear narrow and collimated, and expand rapidly as they travel through polar regions. They are often associated with an Extreme-Ultraviolet (EUV) jet seen near the solar surface \\citep{Wang:etal:1998}. These jets are often rooted in bright, low-lying loop features and are similar in appearance to Soft X-ray (SXR) jets \\citep{Shibata:etal:1992,Moses:etal:1997}. The essential acceleration mechanism for all these jets is very likely provided by magnetic reconnection. The difference between the above and other jet-like features, e.g. chromospheric jets, is the altitude where the magnetic reconnection is assumed to occur. The higher energy jets tend to be accelerated at a higher altitude than the lower energy jets \\citep{Shibata:etal:2007}. The EUV and SXR jets are often caused by the reconnection in the upper chromosphere or the lower corona.  \\citet{Wang:etal:1998} analyzed 27 jets in EIT and LASCO data and characterized their motion by three different velocities: the leading-edge velocity $v_{lead}$, the centroid velocity $v_{cen}$ and the initial velocity $v_{init}$ of the centroid. In all cases $v_{cen}$ was much less than $v_{lead}$, indicating that the jets stretched out rapidly as they propagated through the corona. The authors also found that the bulk of the jet material decelerated as it propagated from the limb to the C2 field of view (FOV). This deceleration was attributed to solar gravity. However, the combined results of $v_{init} < v_{escape}$ and $v_{lead} > v_{escape}$, and the lack of evidence for downflow in EIT and C2 led Wang et al. to propose some in situ acceleration that prevents the bulk of the jet material from falling back onto the Sun after it was ejected.  \\citet{Wood:etal:1999} improved the jet velocity estimates by \\citet{Wang:etal:1998} using on height-time plots. The authors determined the trace of the jet centroid from the EIT to the C2 FOV and found that the observed kinematic trajectories could be fitted with some success by ballistic orbits. They concluded, however, that gravity alone was not the only force controlling the jet propagation. Because of the similar behavior of the jets studied, both \\citet{Wood:etal:1999} and \\citet{Wang:etal:1998} suggested that by the time jets reached the C2 FOV, they were incorporated into the ambient solar wind.  More recently, \\citet{Ko:etal:2005} studied a jet observed jointly by several instruments above the limb. These authors found that a ballistic model could explain most of the dynamical properties of this jet. In their model it was assumed that the gas was ejected upward from the surface with a range of initial speeds. The smooth change of the upflow-to-downflow speed at a certain altitude derived from the ballistic model was found to be consistent with the change of the line intensities from Doppler dimming observed by UVCS/SOHO. Owing to a lack of high cadence coronagraph observations Ko et al.'s study was essentially confined to heights below 1.64 $r_\\odot$.  On June 7, 2007 a big eruptive jet was observed by EUVI, COR1, and COR2 on board the STEREO mission with higher spatial and temporal resolution compared to the data from EIT and LASCO C2. It extended from the solar surface to 5~$r_{\\odot}$. The event was also studied by \\citet{Patsourakos:etal:2008} from the stereoscopic viewpoint. They estimated the jet positions and the speed of the leading front at different times in the EUVI FOV.  In this paper we will attempt to analyze the same jet based on white-light coronagraph observations at heights beyond about 1.5 $r_\\odot$. We extend the ballistic approach by \\citet{Ko:etal:2005} by quantitatively comparing the density variation from Thomson scattered white-light brightness at different heights to the variation expected from a ballistic model of the jet particles. This method has the advantage that it avoids the estimates of jet centroids and fronts. These fronts are not well-defined for a diffusively spreading plasma cloud, the jet centroids are often difficult to determine because a substantial part of jet material is hidden behind the occulter. Another advantage of our particle kinetic analysis lies is that it also provides a test of the validity of the ballistic model. We therefore receive much more information than the conventional leading edge, centroid velocity measurements. After the description of the observations in \\S2, we will introduce the ballistic model in \\S3. In \\S4 we present the results and try to extrapolate our findings to the jet source region and also discuss the limits of our model. Finally, we summarize our conclusions in the last section. ",
        "conclusions": "We have followed the evolution of a big eruptive jet event observed by SECCHI in both EUVI images and the COR1 and COR2 white-light coronagraphs. Based on the distance-time brightness analysis, we found that a ballistic model for the jet particles in general can explain quite well the brightness variation beyond 1.5 $r_\\odot$ in the COR1 and COR2 fields of view with gravity as the dominant acceleration of the jet particles.  Additional parameters were derived or at least estimated, such as the initiation time, the initial velocity distribution, and the number of the jet particles. The derived initiation time is consistent with the EUVI observations at lower altitudes. The initial velocity distribution was fitted by two Maxwellian distributions with different mean kinetic energies. The good agreement with a Maxwellian for lower initial velocities may be due to collisions at heights below 1.4 $r_\\odot$. At high initial velocities, the distribution deviates from the Maxwellian toward a power law tail that may be a result of the jet acceleration process. The total jet particle number and kinetic energy sum up to about 1.6 to 8.9$\\times10^{38}$ and 2.1 to 24$\\times10^{28}$ erg, respectively.  We neglected the effect of the magnetic mirror force and of Coulomb collision. As discussed, they might have some effect on the kinetics of the jet particles at lower altitudes. Note that these two forces counteract each other: while the collisions with the background plasma will decelerate the particles, the mirror force accelerates them away from the Sun. Especially the correct modeling of the Coulomb collisions below 1.4 $r_\\odot$ requires additional assumptions, e.g., about the coronal background density and its velocity.  We outlined the basic idea of a new kinetic jet analysis. In the future, a more sophisticated kinetic model of the jet may be compared to white-light observations. We have shown that this comparison allows one to constrain details of the jet which could not be derived in previous studies. More work needs to be done in these directions. Moreover, more jet samples are required to find out to which extent a jet is embedded in the ambient solar wind and how the jet interacts with it.  PROBA-3, which will be launched in a few years, will have a coronagraph with a FOV of 1.04 to 3 $r_\\odot$. It will provide us with a broader initial velocity coverage because the lower velocity limit depends on the occulter's size . Therefore we will have less uncertainty in the initial velocity distribution, the electron density in the jet source region, etc. The higher temporal observations with more wavelength coverage from AIA/SDO will help determine the jet initiation time more precisely.  Recently, \\citet{Raouafi:etal:2008} found that a jet was very often succeeded by a plume above the jet launch site. Interestingly, in our jet study a plume was visible in both EUVI and COR1 before the jet. This phenomenon was also observed by \\citet{Lites:etal:1999} and other white-light observations in LASCO C2 \\citep{Llebaria:etal:2002}. However, no definite conclusion is given concerning the relation between plume and jet. A time series of 3D reconstruction of both plume and jet needs to be made to find the answer to this question."
    },
    "1112/1112.4229_arXiv.txt": {
        "abstract": "Individual dark matter halos in cosmological simulations vary widely in their detailed structural properties such as shape, rotation, substructure and degree of internal relaxation. Recent non-parametric (principal component) analyses suggest that a few principal components explain a large fraction of the scatter in halo properties. The main principal component is closely linked with concentration, which in turn is known to be related to the mass accretion history of the halo. Here we examine more generally the connection between mass accretion history and structural parameters. The space of mass accretion histories has principal components of its own. We find that the strongest two can be interpreted as the overall age of the halo and the acceleration or deceleration of growth at late times. These two components only account for $\\sim70$\\%\\ of the scatter in mass accretions histories however, due to the stochastic effect of major mergers. Relating structural parameters to formation history, we find that concentration correlates strongly with the early history of the halo, while relaxation correlates with the late history. We examine the inferences about formation history that can be drawn by splitting haloes into subsamples, based on observable properties such as concentration and shape at some final time. This approach suggests interesting possibilities, such as the possibility of defining young and old samples of galaxy clusters in a rigorous, quantitative way, or testing the dynamical assumptions of galaxy formation models empirically. ",
        "introduction": "\\label{sec:introduction} Dark matter halos provide the framework for visible structure in the universe over a span of eight decades in mass, from the scale of rich galaxy clusters down to the scale of individual dwarf galaxies.  Mass is normally assumed to be the main determinant of a halo's baryonic contents, and analytic models such as the halo occupation distribution \\citep[HOD --][]{Peacock, Seljak, Ma, White, Berlind, Cooray2002-halomodel} make this assumption explicitly. Yet the halos that form in cosmological simulations vary greatly in shape, concentration, spin, substructure and other structural properties. As simulations of increasing size and resolution provide a more and more detailed picture of halo properties, and as observational techniques including weak and/or strong gravitational lensing, X-ray and Sunyaev--Zel'dovich measurements \\citep[e.g.][]{Corless2009-WeaklensingShape,Oguri2010-WeaklensingShape25,Sereno2011-StronglensingShape,Morandi2011-xraylensing} reach a precision where they can determine structural properties such as shape and concentration reliably, it is important to understand how the structural features of a halo are interrelated, and what they can tell us specifically about its formation and evolution.  Recently two groups, \\citet{Skibba2011-PCA} and \\citet{JeesonDaniel2011-PCAcorrelations} (S11 and J11 hereafter), have taken the important step of performing non-parametric principal component analyses of halo properties. Principal component analysis (PCA) searches for simplifying trends in a complex data set by finding the axes in a multi-dimensional parameter space that account for the largest fraction of the scatter. In the simplest case, it can uncover linear correlations in the data (e.g. fundamental lines or planes) and reveal hidden patterns or simplifications. The results of S11  and J11 agree on some basic aspects of halo structural properties. Overall, the scatter in halo properties spans a fairly high-dimensional space, with 4 principal components required to explain about 70\\% of the scatter. Nonetheless, a few strong components emerge.  In terms of structural parameters, the first (strongest) principal component is best traced by concentration. The strong correlation with concentration suggests that this first principal component is linked to the overall age of the halo, and the strength of the correlation with $z_{0.5}$ confirms this. The origin of halo concentration has been studied extensively since it is a crucially important factor in many calculations, including strong lensing \\citep[e.g.][]{Broadhurst2005-stronglensing,Broadhurst2008-stronglensing2} and dark matter annihilation \\cite[e.g.][]{Taylor2003-darkmatterannihilation}. Several analytic models have been developed over the years to explain concentration in terms of formation history \\citep[e.g.][]{Bullock2001-concentrationMAH, Wechsler2002-MAH, Zhao09}, so the link between formation history and this particular structural property is fairly well understood. Older systems, that is systems that had already assembled most of their mass into one or a few progenitors at early times, are more concentrated on average, although the exact details of the connection between age and concentration vary from one analytic model to another.  If the first principal component of halo structural properties is thus linked to age, we might naturally ask what the others correspond to. Are shape, spin or relaxation also related to the formation history, and if so how? To put this question in a quantitative framework, we first have to decide how to describe the ``formation history\" itself -- what should we take this to mean exactly, given the complex set of merger and accretion events through which halos form?  We can take the mass accretion history (MAH) as a starting point, defining this as the function ${\\mathcal{M}}(z) \\equiv M(z)/M(0)$ which describes the mass of the main progenitor of a halo as a function of redshift, normalized to the value at $z = 0$ \\citep{vandenBosch2002-MAH}. Since ${\\mathcal{M}}(z)$ is a continuous function of a real variable, it contains an arbitrarily large amount of information about the history of a halo; equivalently, describing the MAH fully means specifying values of ${\\mathcal{M}}(z)$ at an infinite set of redshifts. To characterize the MAH in simpler terms, we can turn once again to principal component analysis, approximating each MAH as a vector of values ${\\mathcal{M}}(z_i)$ corresponding to the MAH evaluated at a finite, fixed set of redshifts $z_i$. PCA of these vectors can then tell us whether mass accretion histories are well described by a set of basis functions characterzed by a single parameter, as suggested by \\citet{vandenBosch2002-MAH} and \\citet{Wechsler2002-MAH}, or whether they require two or more variables to explain their diversity, as suggested by \\citet{Tasitsiomi} or \\citet{McBride2009-MAH}.  Assuming a few principal components capture the main features of a halo's MAH, this will allow us to study correlations between structure and history in a well-defined and quantitative way. We note that this is only a first step towards understanding structure in terms of growth history; the MAH does not contain all the information about a halo's past by any means, and an alternate approach is to study the origin of a particular property such as shape or relaxation in detail, considering the physical processes involved \\citep[e.g.][]{VeraCiro2011-ShapeAquarius,Power2011-CorrelationsMAH}.  In this paper, we generate halos in a set of cold dark matter (CDM) simulations covering three different mass scales. We record the MAH of each well-resolved halo and decompose the space of MAHs, taken as vectors of values ${\\mathcal{M}}(z_i)$, into principal components. These components provide a non-parametric description of the MAH and clarify its basic properties. We also record the structural properties of each halo at the present-day. We study the correlations between these properties themselves (reproducing the trends found by S11 and J11), and between structural properties and the principal components of the merger history. Finally, we discuss an application of our results, showing how samples of halos selected by concentration or shape will have systematically different formation histories. In this way, the observable properties of groups or clusters of galaxies can be used to infer their (unobservable) formation history.  The outline of the paper is as follows. Section \\ref{sec:measuring_halo_properties} describes our n-body simulations and group finding, and how the structural properties of halos are defined and measured. Section \\ref{sec:MAHs} analyses the vector space of MAHs, decomposing them into principal components. In Section \\ref{sec:structural_cor}, we then analyze the correlations between structural components, and derive principal components in this second vector space. In Section \\ref{sec:application} we discuss applications of this work. We summarize our results in Section \\ref{sec:summary}.  Throughout the paper, we consider a WMAP7 cosmology with parameters $\\Omega_M = 0.27$, $\\Omega_\\Lambda=0.73$, $\\Omega_b=0$ (in our CDM-only simulations), $H_{0} = 72$ $\\t{km} \\t{s}^{-1} \\t{Mpc}^{-1}$, $n_s = 0.96$, and $\\sigma_8 = 0.80$. ",
        "conclusions": ""
    },
    "1112/1112.3123_arXiv.txt": {
        "abstract": "{  Recent LHC data showed excesses of Higgs-like signals at the Higgs mass of around 125\\,GeV.  This may indicate supersymmetric models with relatively heavy scalar fermions to enhance the Higgs mass. The desired mass spectrum is realized in the anomaly-mediated supersymmetry breaking model, in which the Wino can naturally be the lightest superparticle (LSP).  We discuss possibilities for confirming such a scenario, particularly detecting signals from Wino LSP at direct detection experiments, indirect searches at neutrino telescopes and at the LHC.  }  \\end{center} \\end{titlepage}  \\setcounter{page}{2}    Higgs mass contains very important information about low-energy supersymmetry (SUSY) models, which is well motivated because it provides a viable candidate of dark matter (DM) and also because it realizes the gauge coupling unification.  In particular, in the minimal SUSY standard model (MSSM), the lightest Higgs boson cannot be heavier than the $Z$-boson at the tree level, while a sizable radiative correction may enhance the Higgs mass ~\\cite{Okada:1990vk}. The size of the radiative correction depends on the masses (and other parameters) of superparticles.  The lightest Higgs mass becomes larger as superparticles (in particular, stops) become heavier.  Thus, once the lightest Higgs mass is known, mass scale of superparticles is constrained.  Recently, the ATLAS collaboration reported $3.6\\sigma$ local excess of the standard model (SM) Higgs-like event at $m_h\\simeq 126\\ {\\rm GeV}$ \\cite{ATLAS-CONF-2011-163}.  In addition, the CMS collaboration also showed more than $2\\sigma$ local excess at $m_h\\simeq 124\\ {\\rm GeV}$ \\cite{HIG-11-032}.\\footnote{The excesses based on global probabilities, which take account of the look-elsewhere effect, are $2.3\\sigma$ (ATLAS) and $1.9\\sigma$ (CMS).}  In order to achieve such a value of the lightest Higgs mass in the MSSM, relatively large values of the superparticle masses are required; the typical scale of the sfermion masses to realize $m_h\\simeq 125$\\,GeV is $10$\\,TeV--$10^3$\\,TeV~\\cite{hep-ph/0408240,arXiv:1108.6077}.  Then, if the masses of all the superparticles are of the same order, it is difficult to find experimental signals of low-energy SUSY and the existence of SUSY is hardly confirmed.  Although the sfermion masses are much larger than the electroweak scale, gauginos may be much lighter than sfermions and within the reach of collier and other experiments.  One interesting possibility is the model in which the SUSY breaking scalar masses are from direct coupling to the SUSY breaking field while the gaugino masses are generated by the anomaly-mediation mechanism~\\cite{hep-th/9810155,hep-ph/9810442}; in this letter, we call such a model as anomaly-mediated SUSY breaking (AMSB) model. Even in the AMSB model, however, if the pure anomaly-mediation relation holds among the gaugino masses, gluino mass is about 8 times larger than the mass of Wino.  Thus, if the Wino mass is a few hundred GeV, which is the lower bound on it from astrophysical and cosmological considerations as will be reviewed later, the gluino mass becomes multi-TeV; with such a heavy gluino, the discovery of the SUSY signal at the LHC becomes challenging because we consider the case that all the squarks are extremely heavy.  Even so, there still exist possibilities of discovering signals of the AMSB scenario.  In particular, in the present framework, the neutral Wino is the lightest superparticle (LSP) and may be DM.  In such a case, pair annihilation cross section of the LSP and the scattering cross section of the LSP off the nuclei are both enhanced compared to the Bino LSP case, which has significant implications to direct and indirect detection of DM.  Because the search of the superparticles at the LHC may be difficult, it is important to pursue these possibilities and explore how well we can study the AMSB scenario with these procedures.\\footnote{ The heavy SUSY particle spectrum and their detectability were discussed in a different context in Ref.~\\cite{Hall:2011jd}. }  In this letter, motivated by the recent Higgs searches at the LHC, we discuss the detectability of the signals of AMSB scenario.  We pay particular attention to the case of the Wino LSP.  We focus on direct/indirect detection of the Wino DM at underground laboratories and neutrino telescopes.  We also comment on the LHC reach for the direct Wino production.  Since superparticles except for gauginos are heavy, standard methods for SUSY searches may not work.  Even in this case, we will show that there are some windows for the confirmation of the SUSY.  Let us first briefly discuss important properties of the AMSB scenario.  We assume that the soft SUSY breaking scalar masses are generated by the direct coupling between the scalars and the SUSY breaking hidden sector field, while the gaugino masses are generated by the anomaly mediation mechanism.  Adopting the pure AMSB relation, the gaugino masses are given by \\cite{hep-th/9810155, hep-ph/9810442} \\begin{eqnarray} M_a^{\\rm (AMSB)} = \\frac{b_a}{16\\pi^2} g_a^2 m_{3/2}, \\end{eqnarray} where $g_a$ ($a=1$--$3$) are gauge coupling constants of the SM gauge groups, $m_{3/2}$ is the gravitino mass, and $(b_1,b_2,b_3)=(11,1,-3)$.  Then, the Wino becomes the lightest among the gauginos, and gaugino masses largely separate: $m_{\\tilde{B}}: m_{\\tilde{W}}: m_{\\tilde{g}}\\simeq 3:1:8$.  Although the AMSB relation may be affected by Higgs and Higgsino loop diagrams \\cite{hep-ph/9810442, Gherghetta:1999sw}, we adopt the pure AMSB mass relation.  With the gaugino masses being of $O(100)\\ {\\rm GeV}$--$O(1)\\ {\\rm TeV}$, the gravitino mass becomes of $O(10)\\ {\\rm TeV}$--$O(100)\\ {\\rm TeV}$.  The sfermion masses are expected to be of the same order of the gravitino mass, which is preferred from the point of view of realizing $m_h\\simeq 125\\ {\\rm GeV}$.  In particular, if the scalar masses are (almost) equal to the gravitino mass, $m_h\\simeq 125\\ {\\rm GeV}$ requires relatively small value of $\\tan\\beta\\sim$ a few (where $\\tan\\beta$ is the ratio of the vacuum expectation values of up- and down-type Higgs bosons) \\cite{arXiv:1108.6077}.  Before discussing the detectability of the signals of AMSB model, we comment on the supersymmetric Higgs mass parameter (so-called $\\mu$-parameter).  In the present setup, the soft SUSY breaking scalar mass parameters of up- and down-type Higgs bosons are expected to be of $O(10)\\ {\\rm TeV}$--$O(100)\\ {\\rm TeV}$.  In order to have viable electroweak symmetry breaking, the $\\mu$-parameter (as well as heavy Higgs boson masses) is also expected to be of the same order; then, the Higgsinos become extremely heavy and the Wino becomes the LSP. Thus, we pay particular attention to the case of Wino LSP in the following.  In some of our following analysis, however, we consider the case with $\\mu\\sim O(100)\\ {\\rm GeV}$--$O(1)\\ {\\rm TeV}$ taking account of the possibility of an accidental tuning of the parameters. This is because detection rates of some of signals (in particular, the direct detection rates) strongly depend on the value of $\\mu$.  Taking account of the radiative correction due to the gauge boson loops, the neutral Wino becomes lighter than the charged one. Therefore, we focus on the case of neutral Wino LSP.  In addition, we assume that the LSP (i.e., the neutral Wino) is the dominant component of DM.  The Wino LSP accounts for the present DM density for $m_{\\tilde W} \\simeq 3$\\,TeV if it is produced only from thermal bath~\\cite{hep-ph/0610249}.  In the AMSB scenario, however, the Wino LSP can be non-thermally produced from the gravitino or moduli decay \\cite{hep-ph/9810442, hep-ph/9906527}.  If the reheating temperature takes an appropriate value, for example, the decay of gravitino produces the Wino LSP with correct relic density \\cite{hep-ph/0012052}, while thermal leptogenesis~\\cite{Fukugita:1986hr} works successfully~\\cite{Ibe:2011aa}.  Thus the Wino is a good DM candidate in the present setup.  Hereafter, we assume that the right amount of Wino is somehow produced in the early universe to be DM.   We start with discussing direct detection experiments of DM.  The scattering cross section of the Wino LSP off the nucleon significantly depends on $\\mu$.  Since all scalars except for the lightest Higgs boson are expected to be heavy enough, it is only the lightest Higgs boson that mediates the spin-independent (SI) scattering.  The DM-proton scattering cross section is given by~\\cite{Jungman:1995df} \\begin{equation} \\sigma = \\frac{4}{\\pi} \\left( \\frac{m_{\\tilde \\chi^0} m_N}{m_{\\tilde \\chi^0} + m_N} \\right)^2 \\left[ \\left( n_p f_p + n_n f_n \\right)^2 + 4 \\frac{J+1}{J} \\left( a_p \\langle s_p \\rangle + a_n \\langle s_n \\rangle \\right)^2 \\right], \\end{equation} where the first and the second term in the bracket are the contributions of SI and spin-dependent (SD) interaction, respectively. Here $m_{\\tilde \\chi_0}$ is the LSP mass, $m_N$ is the mass of the target nucleus, $n_p (n_n)$ is the number of proton (neutron) in the target nucleus, $J$ is the total nuclear spin, $a_p$ and $a_n$ are the effective DM-nucleon SD couplings, and $\\langle s_{p(n)} \\rangle$ are the expectation values of the spin content of the proton and neutron groups within the nucleus.  The effective DM-proton coupling, $f_p$, is given by \\begin{equation} f_{p} = \\sum_{q=u, d, s} \\frac{f_q^{H}}{m_q} m_p f_{T_q}^{(p)} + \\frac{2}{27}f_{T_G} \\sum_{q=c, b, t} \\frac{f_q^{H}}{m_q} m_p, \\end{equation} where $f_{T_G}=1-\\sum_{u,d,s}f_{T_q}^{(p)}$, $m_p$ and $m_q$ denote the proton and quark masses, respectively, and $f_q^H$ is the effective DM-quark coupling obtained by the exchange of the Higgs boson.  Since the DM-Higgs coupling is proportional to the magnitude of Wino-Higgsino mixing,\\footnote{ In the limit of pure Wino DM, the Wino-nucleon scattering cross section is too small to be detected~\\cite{arXiv:1004.4090}.  } the cross section is enhanced if the Wino-Higgsino mixing is large.  In Fig.~\\ref{fig:SI} we plot the Wino-proton SI and SD scattering cross section.  In this plot we have used following values for the quark contents in the proton~\\cite{Ohki:2008ff} : $f_{T_u}^{(p)}=0.023, f_{T_d}^{(p)}=0.034, f_{T_s}^{(p)}=0.025$ and taken $\\tan \\beta = 3$ and $\\tan\\beta=20$. The XENON100 experiment~\\cite{arXiv:1104.2549} most severely constrains the SI cross section.  The sensitivity is improved by a few orders of magnitude for the next generation 1 ton scale detectors, and then broad parameter regions up to $m_{\\tilde W}\\sim \\mu \\sim 1$\\,TeV will be explored.  The IceCube searches for neutrino events arising from the DM annihilation in the Sun.  Since the efficiency for the DM trapping into the Sun depends on the DM-proton scattering cross section, the high-energy neutrino observations give limits on it.  For the SD cross section, the IceCube gives the most stringent limit, and it will be further improved by about one order of magnitude with the DeepCore instrument~\\cite{arXiv:1111.2738}.  \\begin{figure}[tbp] \\begin{center} \\includegraphics[scale=0.5]{SI.eps} \\vskip 1cm \\includegraphics[scale=0.5]{SD.eps} \\caption{ Contours of spin-independent (SI) and spin-dependent (SD) Wino-proton scattering cross sections are plotted on the plane of $m_{\\rm wino}$ and $\\mu$. Shaded regions are excluded by the XENON100 experiment for SI, and IceCube experiment for SD. } \\label{fig:SI} \\end{center} \\end{figure}    We have also calculated the detection rate at the IceCube DeepCore, arising from high-energy neutrinos produced by the Wino annihilation at the Galactic Center (GC).  We distinguish two event classes following Refs.~\\cite{arXiv:0911.5188,arXiv:0912.3521,arXiv:1009.2068} : contained muon events and shower events.  The contained muons correspond to those emerge inside the instrumental volume through the high-energy neutrino interactions with nucleons.  The shower events are caused by charged current interactions of electron and tau neutrinos, and neutral current interactions of all neutrino species. They leave electromagnetic/hadronic shower inside the instrumental volume.  The event rate of the contained muons is given by \\begin{equation} N_{\\mu^+\\mu^-} = \\int dE_{\\nu_\\mu}\\int_{E_{\\rm th}}^{E_{\\nu_\\mu}} dE_\\mu \\left[ \\frac{d\\Phi_{\\nu_\\mu}}{dE_{\\nu_\\mu}} \\left( \\frac{d\\sigma_{\\nu_\\mu p}^{\\rm (CC)}}{dE_\\mu}n_p +\\frac{d\\sigma_{\\nu_\\mu n}^{\\rm (CC)}}{dE_\\mu}n_n \\right) +(\\nu_\\mu \\leftrightarrow \\bar \\nu_\\mu) \\right] V_{\\rm eff}(E_\\mu),   \\label{cont} \\end{equation} where $E_{\\nu_\\mu}$ is the incident neutrino energy, $E_\\mu$ is the muon energy resulting from the neutrino-proton (neutron) interactions, $E_{\\rm th}$ is the threshold energy above which the muon can be detected, $d\\Phi_{\\nu_\\mu}/dE_{\\nu_\\mu}$ is the neutrino flux at the Earth, $d\\sigma_{\\nu_\\mu p(n)}^{\\rm (CC)}/dE_\\mu$ denotes the neutrino-proton (neutron) charged current cross section for producing the muon energy with $E_\\mu$, $n_p (n_n)$ is the proton (neutron) number density in the detector material, and $V_{\\rm eff}$ is the effective volume for the muon detection.  The incident neutrino flux generated by DM annihilation from the GC within cone half angle of $\\theta$ is given by \\begin{equation} \\frac{d\\Phi_{\\nu_i}}{dE_{\\nu_i}} = \\frac{R_{\\odot}\\rho_{\\odot}^2}{8\\pi m_{\\tilde W}^2} \\left( \\sum_{j=e,\\mu,\\tau} \\langle \\sigma v\\rangle \\frac{dN_{\\nu_j}}{dE_{\\nu_j}}P_{j\\to i} \\right) \\langle J_2 \\rangle_\\Omega \\Delta \\Omega. \\end{equation} Here $R_\\odot = 8.5$\\,kpc and $\\rho_\\odot=0.3\\,{\\rm GeV}\\,{\\rm cm^{-3}}$, $\\langle \\sigma v\\rangle$ is the Wino self-annihilation cross section including the non-perturbative effect~\\cite{hep-ph/0307216}, and $dN_{\\nu_j}/dE_{\\nu_j}$ is the energy spectrum of the neutrino produced by DM annihilation, which is calculated by the PYTHIA package for the $WW$ final state~\\cite{hep-ph/0603175}, $P_{j\\to i}$ is the probability that the $\\nu_j$ at the production is converted to $\\nu_i$ because of the neutrino oscillation effect, $\\Delta\\Omega=2\\pi(1-\\cos\\theta)$, and $\\langle J_2 \\rangle_\\Omega$ includes the information about the DM density profile in the Galaxy~\\cite{arXiv:0812.0219}.  The shower event is evaluated in a similar way to the contained muon events (\\ref{cont}), except that the charged current interactions from $\\nu_e$ and $\\nu_\\tau$ as well as the neutral current interactions for all neutrino flavors are included.  The background event is evaluated by inserting the atmospheric neutrino flux into the expression (\\ref{cont}).  Fig.~\\ref{fig:ICDC} shows the signal-to-noise ratio at the IceCube DeepCore as a function of the Wino mass.  Sensitivities for contained muon events (upper panel) and shower events (lower panel) with 1 year and 10 year observations are shown.  We have adopted the NFW density profile and considered the neutrino flux from the cone half angle $\\theta = 10^{\\circ}$ and $\\theta = 25^{\\circ}$ around the GC.  As noted in Ref.~\\cite{arXiv:1009.2068}, the sensitivity is maximized for $\\theta \\simeq 10^{\\circ}$.  For this cone half angle, the flux dependence on the DM density profile is not large~\\cite{arXiv:0812.0219}.  The effective volume for the contained and shower events are set to be $0.04\\,{\\rm km^3}$ and $0.02\\,{\\rm km^3}$, respectively~\\cite{arXiv:0911.5188}.  The atmospheric background is taken from Ref.~\\cite{astro-ph/0611418}.  It is seen that the signal-to-noise ratio is at most order one for the Wino mass of a few hundred GeV.  We have also checked that the upward muon events expected at the KM3NeT detector~\\cite{848444}, assuming the effective area of $1\\,{\\rm km^2}$ and taking account of the energy loss of muons~\\cite{hep-ph/0012350}, provide similar sensitivities to the DeepCore.  \\begin{figure}[tbp] \\begin{center} \\includegraphics[scale=1.7]{ICDC_cont.eps} \\vskip 1cm \\includegraphics[scale=1.7]{ICDC_show.eps} \\caption{ Signal-to-noise ratio at the IceCube DeepCore as a function of the Wino mass. Sensitivities for contained muon events (upper panel) and shower events (lower panel) with 1 year and 10 year observations are shown. We have considered the neutrino flux from the cone half angle $\\theta = 10^{\\circ}$ and $\\theta = 25^{\\circ}$ around the GC. } \\label{fig:ICDC} \\end{center} \\end{figure}   The Wino DM annihilation may leave characteristic signatures on astrophysical observations.  Gamma-ray observations by Fermi-LAT and HESS severely restrict the DM annihilation cross section (see, e.g., Refs.~\\cite{arXiv:1108.3546,arXiv:1110.6151} for recent works).  The non-observations of DM-induced gamma-rays from dwarf galaxies excludes the Wino mass below $\\sim 400$\\,GeV~\\cite{arXiv:1108.3546}, although there are astrophysical uncertainties.  On the other hand, the cosmic-ray positron excess observed by PAMELA satellite~\\cite{arXiv:0810.4995} may be explained by the Wino DM annihilation with mass of 200\\,GeV~\\cite{arXiv:0810.1892,arXiv:0812.4555} although it may confront the constraints from gamma-rays and anti-protons.  The observations of light element abundances also give stringent bound on the DM annihilation cross section so as not to destroy light elements during Big-Bang Nucleosynthesis (BBN).  It gives a lower bound on the Wino mass as $m_{\\tilde W} \\gtrsim 200$\\,GeV~\\cite{astro-ph/0405583,arXiv:0810.1892}.  It may be encouraging that the cosmic lithium problem may be solved for the Wino mass of around this bound, which simultaneously may explain the PAMELA anomaly.  DM annihilation also affects the recombination history of the Universe, which results in the modification on the cosmic microwave background (CMB) anisotropy~\\cite{arXiv:0905.0003,arXiv:0906.1197,arXiv:0907.3985}. The constraint is comparable to that from BBN.  Taking these constraints into account, we conservatively consider that the Wino must be heavier than $\\sim 200$\\,GeV if it is the dominant component of DM.   Finally, we comment on a possibility of discovering a signal of AMSB model at the LHC.  If we adopt the AMSB mass relation among gauginos, gluino becomes relatively heavy.  Then, colored superparticles are hardly produced at the LHC.  Thus, we focus on the detection of a Wino signal.  If the neutral Wino $\\tilde{W}^0$ is the LSP, we have a chance to observe the track of charged Wino $\\tilde{W}^\\pm$ \\cite{Feng:1999fu, hep-ph/0610277}.  This is because the mass difference between charged and neutral Winos is so small ($\\sim 160\\ {\\rm MeV}$) that the decay length of $\\tilde{W}^\\pm$ becomes macroscopic ($c\\tau_{\\tilde{W}^\\pm}\\simeq 5\\ {\\rm cm}$).  Some of the produced charged Winos may travel through several layers of inner trackers and their track may be reconstructed.  In the ATLAS experiment, for example, the charged Wino track can be reconstructed with almost 100 \\% efficiency if $\\tilde{W}^\\pm$ hits the 3rd layer of the semiconducter tracker (SCT) before it decays \\cite{arXiv:0807.4987}. Then, because of the smallness of $c\\tau_{\\tilde{W}^\\pm}$ compared to the detector size, $\\tilde{W}^\\pm$ decays before going through the whole detector.  Such a charged Wino is identified as a high $p_T$ track which disappears in the middle of the detector.  Such a signal does not exist in the SM, and hence is a smoking gun evidence of the production of $\\tilde{W}^\\pm$.  The Wino pair can be produced by the Drell-Yan process at the LHC. However, there is no high $p_T$ jet nor track in the final state in such an event, and hence the event cannot be recorded.  In order to trigger on the Wino production events, one can use the event with high $p_T$ jet; such a jet can be from the initial state radiation.  Then, at the parton level, the Wino production processes relevant for the present study are the following: \\begin{eqnarray*} && q\\bar{q} \\rightarrow \\tilde{W}^+ \\tilde{W}^- g,~~~ gq \\rightarrow \\tilde{W}^+ \\tilde{W}^- q,~~~ g\\bar{q} \\rightarrow \\tilde{W}^+ \\tilde{W}^- \\bar{q}, \\\\ && q\\bar{q}' \\rightarrow \\tilde{W}^\\pm \\tilde{W}^0 g,~~~ gq \\rightarrow \\tilde{W}^\\pm \\tilde{W}^0 q',~~~ g\\bar{q} \\rightarrow \\tilde{W}^\\pm \\tilde{W}^0 \\bar{q}'. \\end{eqnarray*} We calculate the cross section of the process $pp\\rightarrow\\tilde{W}\\tilde{W}j$; we perform the parton level calculation, and we approximate the $p_T$ of jet by that of final-state quark or gluon.  In the calculation of the cross section, the helicity amplitude package HELAS \\cite{Murayama:1992gi} and the CT10 parton distribution functions \\cite{Lai:2010vv} are used.  For the phase space integration, we use the BASES package \\cite{BASES}. In the calculation, we require that the transverse momentum of the jet be larger than 170, 270, and 370 GeV, and that at least one charged Wino travels more than 44.3 cm which is the distance to the 3rd layer of SCT from the beam pipe in the ATLAS detector \\cite{Aad:2008zzm}.  In Fig.\\ \\ref{fig:cs}, the cross section is plotted as a function of the Wino mass.  In the high luminosity run with ${\\cal L}=2\\times 10^{33}\\ {\\rm cm^{-2}s^{-1}}$, for example, the so-called j370 trigger is planned to be available, which requires a jet with $p_T>370\\ {\\rm GeV}$ \\cite{Aad:2008zzm}.  Then, requiring 10 events with $p_T>370\\ {\\rm GeV}$ for the discovery, for example, Wino mass smaller than $270\\ {\\rm GeV}$ ($330\\ {\\rm GeV}$) is covered by the LHC with the luminosity of ${\\cal L}=100\\ {\\rm fb}^{-1}$ ($300\\ {\\rm fb}^{-1}$), where we have assumed that the background is negligible.  Thus, in particular when $\\mu$ is large, the LHC experiment still have a chance to cover the parameter region which has not been excluded yet by the current direct and indirect DM searches.  If the $p_T$ of the jet for the trigger can be reduced, the LHC can cover the region with larger Wino mass.  So far, we have assumed the pure AMSB relation among gaugino masses. However, as we have mentioned, such a relation may be largely affected by the Higgs and Higgsino loop diagrams.  With such an effect, the gluino mass may become $\\sim 1\\ {\\rm TeV}$ even when the Wino mass is a few GeV.  In such a case, the conventional procedures of the SUSY search using the missing energy distribution may work.  \\begin{figure}[tbp] \\begin{center} \\includegraphics[scale=1.7]{csCTEQ10.eps} \\caption{ Cross section for the process $pp\\rightarrow\\tilde{W}\\tilde{W}j$ (with $j=q$ or $g$), for $\\sqrt{s}=14\\ {\\rm TeV}$.  The transverse momentum of $j$ is required to be larger than 170, 270, and 370 GeV from above. } \\label{fig:cs} \\end{center} \\end{figure}  In summary, motivated by the recent report on the Higgs searches at the LHC, which indicated excesses of Higgs-like events at around $m_h\\simeq 125\\ {\\rm GeV}$, we have investigated prospects for confirmation of the AMSB scenario, particularly the detection of Wino LSP.  We have considered the situation that the scalars except for gauginos and Higgsinos are heavy enough so that they cannot be produced at colliders.  Even in this unfortunate case, the Wino DM may be detected through direct/indirect detection experiments.  Direct detection efficiency crucially depends on the Higgsino mass, and if the Wino and Higgsino masses happen to be close, future experiments may find their signals.  The neutrino telescopes such as IceCube DeepCore and KM3NeT also have a potential to discover the Wino LSP through the observation muon and/or shower events induced by high-energy neutrinos from DM annihilation at GC. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.3409_arXiv.txt": {
        "abstract": "We consider the acceleration of very small dust grains including Polycyclic Aromatic Hydrocarbons (PAHs) arising from the electrostatic interactions of dust grains that have charge fluctuations in time due to charging events. We simulate the charge fluctuations of very small grains due to their sticking collisions with electrons and ions in plasma, and the emission of photoelectrons by UV photons using Monte Carlo method. We identify the acceleration induced by the charge fluctuations as the dominant acceleration mechanism of very small grains in the diffuse interstellar medium (ISM). We show that this acceleration mechanism is efficient for environments with low degree of ionization (i.e., large Debye length), where the charge fluctuations are slow but have a large amplitude. We also discuss the implications of the present mechanism for grain coagulation and shattering in the diffuse ISM, molecular clouds, and protoplanetary disks. ",
        "introduction": "Dust is an important constituent of the interstellar medium (ISM), molecular clouds, and accretion disks (see Whittet 2003; Draine 2009, 2011). It gets involved in many key processes, for instance, it controls heating and cooling of the ISM (see Draine 2003; Tielens 2005), reveals magnetic fields through grain alignment (see Lazarian 2007 for a review), and interferes with attempts to measure the temperature anisotropy and polarization of the cosmic microwave background (CMB) radiation (see Lazarian \\& Finkbeiner 2003; Fraisse et al. 2009; Dunkley et al. 2009).  Very small dust grains of size $a\\le 100\\Angstrom$ with a notable fraction of polycyclic aromatic hydrocarbons (PAHs, hereafter very small grains and PAHs are used interchangeably) radiate electric dipole emission, which is an important component of CMB foregrounds (Draine \\& Lazarian 1998; Hoang et al. 2010, 2011). PAHs in protoplanetary disks (PPDs) around young intermediate mass (Herbig Ae/Be) stars (Acke \\& van den Ancker 2004) and in outer layers of low mass (T-Tauris) stars (Geers et al. 2006; Oliveira et al. 2010) may affect the magnetorotational instability (MRI) due to their influence on the environment ionization (Bai 2011; Perez-Becker \\& Chiang 2011).  Most properties of dust, including light extinction, electron photoemission, and chemical activity depend not only on grain chemical composition, but also on their sizes. In astrophysical environments, grain size is affected by grain-grain collisions, which depends on their relative velocities. The minimal velocity of grains corresponds to their thermal velocity at the ambient gas temperature arising from Brownian motions. These velocities are usually assumed in the models of dust coagulation (Ossenkopf 1993; Dominik \\& Dullemond 2008). It has been long known that large scale hydrodynamic motions associated with turbulence can make grains move faster (see Draine 1985), but detailed calculations applicable to astrophysical environments started to appear only a few years ago.  Since most astrophysical media are magnetized and grains are charged, the hydrodynamic treatment of acceleration is not adequate. A proper treatment of the grain acceleration through the resonant interactions of charged grains with magnetohydrodynamic (MHD) turbulence has been developed relatively recently (Lazarian \\& Yan 2002; Yan \\& Lazarian 2003; Yan et al. 2004; Yan 2009; Hoang et al. 2011b). This treatment makes an extensive use of the advances in the theory/numerical studies of compressible MHD turbulence (see Cho \\& Lazarian 2003; Kowal \\& Lazarian 2010 and ref. therein) and provides the mathematical formalism of the second-order Fermi acceleration of charged grains interacting with MHD turbulence.  The acceleration due to the resonant interactions of MHD turbulence with charged grains decreases with decreasing grain size. Such decrease arises from the fact that the Larmor radius of charged grains becomes smaller as the grain mass decreases, and correspondingly, grains have to interact with smaller and less powerful turbulent fluctuations. In addition, compressible fluctuations (i.e., fast modes), which were identified in Yan \\& Lazarian (2003) as the most efficient source of acceleration, get suppressed at small scales due to plasma viscous damping, while the Alfv\\'{e}nic mode gets inefficient for acceleration at the small scales due to anisotropy (see Yan \\& Lazarian 2003 for more discussion). The resonant acceleration for grains with sizes $\\leq 10^{-5}$~cm becomes rather inefficient in most media discussed in Yan \\& Lazarian (2003) and Yan et al. (2004, hereafter YLD04). These conclusions were also confirmed by Hoang et al. (2011b) who considered the non-linear effects of resonant acceleration (Yan \\& Lazarian 2008) and the transit time damping acceleration (TTD) of charged grains.  A new mechanism of astrophysical grain acceleration was sketched in Ivlev et al. (2010), where rough estimates of grain velocities arising from charge fluctuations in the ISM were provided. The study employed Fokker-Planck equation and assumed that (1) charge fluctuations are fast, i.e., the relaxation timescale for charge fluctuations is much shorter than the gaseous damping time and the Larmor precession period, and (2) charge fluctuations have small amplitude (i.e., $\\delta Q \\ll Q_{0}$) such that these fluctuations can be approximated as being continuous and modeled by a Gaussian distribution. These assumptions are usually not realistic for very small grains for which the charge fluctuations are substantial. Therefore, the unrealistically high velocities of very small grains were obtained. This motivates our present study which is intended to determining the actual velocities of PAHs in typical astrophysical conditions.  In this paper, we aim to quantify the acceleration of very small dust grains induced by grain charge fluctuations using numerical simulations. The structure of the paper is as follows. We first discuss major processes responsible for grain charging and charge fluctuations in \\S 2. In \\S 3 we present an algorithm to simulate the charge fluctuations using Monte Carlo method. In \\S 4 we present a numerical simulation method to investigate grain acceleration, including dynamical equations and Monte-Carlo simulations. Velocities of very small grains in the ISM from numerical simulations are presented in \\S 5. Discussions and summary are presented in \\S 6 and 7, respectively. ",
        "conclusions": "\\subsection{Grain Acceleration from Various Mechanisms} Grain acceleration by hydrodrag (see Draine 1985) was discussed as the dominant mechanism driving grain motions in the ISM as well as in protoplanetary disks. Recently, thanks to significant progresses in understanding of MHD turbulence, Lazarian \\& Yan (2003), YLD04, and Yan (2009) identified the gyro-resonant interactions of grains with fast modes in MHD turbulence as the dominant mechanisms to accelerate large grains ($a>10^{-5}$ cm) to super-Alfv'{e}nic velocities. For very small grains (e.g., PAHs and nanoparticles), the gyro-resonance acceleration is inefficient for most phases of the ISM, because PAHs and nanoparticles with smaller gyro radii only resonantly interact with small scale and weak fluctuations. Hoang et al. (2011b) revisited the treatment of gyroresonance acceleration by MHD turbulence by accounting for the fluctuations of the grain guiding center with respect to the mean magnetic field. They improved estimates of the previous authors, and proposed a new way of acceleration through transit time damping (TTD) of fast modes, which is efficient for super-Alfv\\'{e}nic grains (see solid lines in Fig. \\ref{fig_TTD}).  \\subsection{Acceleration induced by Charge Fluctuations} In the present paper, we have numerically investigated the grain acceleration mechanism induced by charge fluctuations for very small grains. The combination of MC simulations for charge fluctuations with direct simulations of electrostatic interactions between grains with fluctuating charge allows us to follow the evolution of grain velocities in time and estimate their terminal velocities. We showed that the acceleration induced by charge fluctuations is indeed an important mechanism to accelerate PAHs. Resulting grain velocities can be several times higher than their thermal velocities, which is obviously important.  \\begin{figure} [h] \\includegraphics[width=0.49\\textwidth]{f9.ps} \\caption{Grain velocity as a function of grain size $a$ due to charge fluctuations (shaded area) and due to resonance acceleration by MHD turbulence (Hoang et al. 2011b) for the CNM, WNM, and WIM. MHD turbulence is the efficient mechanism for large grains, and grain charge fluctuations work efficiently at the small end of grain size. Dotted and solid lines correspond to the acceleration from gyroresonance and gyroresonance plus transit time damping.} \\label{fig_TTD} \\end{figure}  We also found that grains velocities obtained from MC simulations are between two and three orders of magnitude lower than the prediction by the FP equation. Such overestimate of the FP approach arises because the charge fluctuations of very small grains in the ISM are slow, in contrast to the assumption of fast charge fluctuations assumed in the FP treatment, which makes electric fields fluctuating rather slowly. In addition, the FP treatment disregarded the screening effect of plasma. In the ISM, the mean separation between two grains is comparable to the Debye length. Thus, the Coulomb interaction is suppressed considerably, and the grain acceleration is decreased accordingly.  Figure \\ref{fig_TTD} compares grain velocities induced by charge fluctuations (shaded area) with those due the resonance interactions in MHD turbulence. As shown, MHD turbulence is efficient for large grains, while charge fluctuations are important for very small grains.  \\subsection{Grain coagulation and shattering}  The grain coagulation is a process in which small dust particles collide with each other to form aggregates of submicron size (Spitzer 1978). The grain coagulation can occur in dense clouds (e.g., Ossenkopf 1993; Ormel et al. 2009) as well as in protoplanetary disks (see van Boekel et al. 2003).  Grain coagulation and shattering result from grain-grain collisions, which depends on grain relative velocity, denoted by $v_{dd}$. The threshold velocity for the grain shattering is a function of the grain size: \\bea v_{\\rm shat}=2.7\\left(\\frac{a}{10^{-7}~\\cm}\\right)^{-5/6} {\\rm km}~\\s^{-1},\\label{v_cri} \\ena (Chokshi et al. 1993). If $v_{dd}< v_{\\rm shat}$, the grains collide and stick together. When $v_{dd} > v_{\\rm shat}$, the collisions with high velocity produce shock waves inside the grain, and shatter them in smaller fragments. For $v_{dd}\\rightarrow 20$ km$\\s^{-1}$, the evaporation of dust grains occurs and grains are destroyed.  We note that due to resonant acceleration by MHD turbulence, large grains $a>10^{-5}$ cm move frequently with velocity larger than $v_{\\rm shat}$ (see Fig. \\ref{fig_TTD}). Thus, the shattering of large grains can be efficient to replenish very small grains into the ISM. The effect of grain acceleration by MHD turbulence on grain coagulation and shattering was studied in Hirashita \\& Yan (2009). In some phases, they found that grain acceleration can modify grain size distribution. At the same time, very small grains move faster than thermal speed due to the acceleration discussed in the present paper and enhance grain coagulation. Therefore, future models of grain evolution should take into account the acceleration of very small grains.  \\subsection{Effects of PAH acceleration in protoplanetary disks}  Grain coagulation is the first step to planetesimal formation in protoplanetary disks (PPDs). It is also believed that the coagulation from very small grains to micron size occurs rapidly (see e.g., van Boekel et al. 2003), so that PAHs and nanoparticles are rapidly depleted in PPDs. However, recent observations by Infrared Space Observatory (ISO) and {\\it Spitzer} satellite reveal PAH emission features, which indicates the existence of PAHs in young intermediate mass (Herbig Ae/Be) stars (Acke \\& van den Ancker 2004) and in the surface layers exposed to UV radiation in young low mass (T-Tauri) stars (Geers et al. 2006; Oliveira et al. 2010; Bern\\'{e} et al. 2009).  PPDs have highly stratified structure, that means, there exists some gradient of grain size from the midplane to the surface. When PAH emission is seen, it is probably only from PAHs in the surface layers that \"see\" direct stellar radiation. The abundance of PAHs and other nanoparticles below the surface is rather uncertain.  {\\it (a) Effects of acceleration on dust coagulation and planet formation}  Current models of grain growth in PPDs disregarded grain charges and their fluctuations. Recently, Okuzumi (2009) took into account the effects of grain charge and showed that Coulomb potential barrier can suppress grain coagulation. Okuzumi (2009) appealed to turbulence as an energy source to overcome the potential barrier. Below, we show that charge fluctuations can induce very small grains to overcome the Coulomb barrier assuming that PPDs have a sufficiently high degree of ionization for which the charge fluctuations are important.  Assuming that the charge fluctuations are Gaussian, then, the mean charge and charge dispersion are respectively given by \\bea \\langle Z\\rangle=z\\frac{ak_{\\B}T_{\\gas}}{e^{2}}\\simeq-0.1\\left(\\frac{a}{10^{-6}\\cm}\\right) \\left(\\frac{T_{\\gas}}{100~\\K}\\right),~~~\\label{zmean} \\ena where $z=2.5$ is adopted (see Morfill \\& Ivlev 2009).  The dust mass density in PPDs is related to the gas density as \\bea \\rho_{d}=f_{dg}\\rho_{g},\\label{rhod} \\ena where $f_{dg}=0.014$ for the PPD (see Tanaka et al. 2005). The dust density of the grain size $a$ is then \\bea n_{d}=10^{1}\\left(\\frac{a}{10^{-6}~\\cm}\\right)^{-3} \\left(\\frac{n_{\\gas}}{10^{10}~\\cm^{-3}}\\right).\\label{nd_ppd} \\ena \\\\ Using above equations for Equation (\\ref{Tterm}), we obtain \\bea \\frac{T_{d}^{\\infty}}{T_{\\gas}}=9.31\\times 10^{17}\\epsilon^{2} \\left(\\frac{T_{\\gas}}{100~\\K}\\right)^{2} \\left(\\frac{a}{10^{-6}~\\cm}\\right)^{-4}\\nonumber\\\\ \\times x_{i} \\left(\\frac{0.1~\\cm^{-3}}{n_{i}}\\right)^{2}\\left(\\frac{n_{\\gas}}{10^{10}~\\cm^{-3}}\\right),~~~~~\\label{vd_ppd} \\ena where $\\epsilon<1$ is introduced to account for the overestimate of grain velocities obtained using the FP approach, $n_{i}$ is the ion density, $x_{i}=n_{i}/n_{\\gas}$ is the ionization fraction of gas, and PPDs are assumed to be isothermal plasma with $T_{\\gas}=T_{i}$.  Now, let us consider whether the charge fluctuations can help grain grow against the Coulomb potential barrier. The ratio of the grain kinetic energy arising from charge fluctuations to electrostatic potential at a distance equal to $2a$ with $a$ being the grain size is \\bea \\frac{E_{\\kin}}{E_{\\rm Coul}}=\\frac{3a k_{\\B}T_{d}^{\\infty}}{\\langle Z\\rangle^{2}e^{2}}.\\label{ratio} \\ena  From Equations (\\ref{zmean}), (\\ref{vd_ppd}), and (\\ref{ratio}) we obtain \\bea \\frac{E_{\\kin}}{E_{\\rm Coul}}\\simeq 2.3\\times10^{1}\\epsilon^{2}\\left(\\frac{a}{10^{-6}~\\cm}\\right)^{-5} \\left(\\frac{T_{\\gas}}{100~ \\K}\\right)\\nonumber\\\\ \\times\\left(\\frac{10^{10}~\\cm^{-3}}{n_{\\gas}}\\right)\\left(\\frac{10^{-4}}{x_{i}}\\right),~~~\\label{ekin_coul} \\ena where the ionization fraction $x_{i}\\ne 0$ was assumed to be sufficiently high so that the charge fluctuations are important.  Grain growth occurs only if $E_{\\kin}/E_{\\rm Coul}\\ge1$.  Thus, from Equation (\\ref{ekin_coul}) we can derive the range of grain size with the grain growth. Using parameters for PPDs conditions, $T_{\\gas}=300~\\K$, $n_{\\gas}=10^{10}~\\cm^{-3}$ and $x_{i}=10^{-4}$ and taking $\\epsilon=10^{-1}$, we found that very small grains with size $a\\leq 6\\times 10^{-7}~\\cm$. Therefore, the growth of larger grains can be induced by other processes (e.g., turbulence as suggested in Okuzumi 2009).  The presence of PAHs and nanoparticles in PPDs can arise from the fragmentation and shattering of micron graphite and silicate grains (see e.g., Dullemond \\& Dominik 2008). The level of MHD turbulence in PPDs is very uncertain, but it is believed that the layers near the surface are highly turbulent. As a result, large grains can be accelerated by the resonant interactions of fast modes to super-Alfv\\'{e}nic speed (Yan et al. 2004; Hoang et al. 2011b). Thus, the grain shattering may be efficient that return very small grains to PPDs.  {\\it (b) Effects of PAH acceleration on magnetorotational instability}  PAHs play an important role in the ionization equilibrium in PPDs. Bai (2011) found that PAHs can favor the accretion due to magneto-rotational instability (MRI) by decreasing ambipolar diffusions in outer layers of PPDs. In fact, if charged PAHs are more abundant than electron ($n_{\\rm PAH}> n_{e}$), then the ambipolar diffusion is less important. However, since the electrical conductivity by charged PAHs is more efficient than electron, the MRI becomes more efficient. Suprathermal velocities of PAHs due to charge fluctuations in PPDs can enhance the accretion rate of electron and ion onto grains, which certainly affect the ambipolar diffusion\\footnote{The ambipolar diffusion is usually invoked to describe the removal of magnetic field during star formation. Another powerful mechanism is reconnection diffusion based on the ability of turbulent magnetic field to reconnect and redistribute the entrained matter (Lazarian 2005; see also Lazarian 2011).} and MRI. Thus, understanding grain motions in PPDs is important for understanding the dynamics of PPDs and planetesimal formation."
    },
    "1112/1112.1898_arXiv.txt": {
        "abstract": "Compact binary coalescences, such as binary neutron stars or black holes, are among the most promising candidate sources for the current and future terrestrial gravitational-wave detectors. While such sources are best searched using matched template techniques and chirp template banks, integrating chirp signals from binaries over the entire Universe also leads to a gravitational-wave background (GWB). In this paper we systematically scan the parameter space for the binary coalescence GWB models, taking into account uncertainties in the star formation rate and in the delay time between the formation and coalescence of the binary, and we compare the computed GWB to the sensitivities of the second and third generation gravitational-wave detector networks. We find that second generation detectors are likely to detect the binary coalescence GWB, while the third generation detectors will probe most of the available parameter space. The binary coalescence GWB will, in fact, be a foreground for the third-generation detectors, potentially masking the GWB background due to cosmological sources. Accessing the cosmological GWB with third generation detectors will therefore require identification and subtraction of all inspiral signals from all binaries in the detectors' frequency band. ",
        "introduction": "The ground-based gravitational-wave detectors are rapidly increasing their sensitivities. The first generation detectors LIGO \\cite{LIGOinstr,S5detector} and Virgo \\cite{Virgo1,Virgo2} have reached their design sensitivities and collected excellent data over several years of exposure. The second generation detectors, Advanced LIGO \\cite{aLIGO, aLIGO2}, Advanced Virgo \\cite{aVirgo}, GEO-HF \\cite{GEOHF}, and LCGT \\cite{CLIO,LCGT} are currently being built and commissioned. With 10 times better strain sensitivity, these detectors are expected to yield first direct detections of gravitational-wave signals, and their first data is expected as early as 2014. Furthermore, there are already efforts under way to design the third-generation gravitational wave detectors, with another factor of 10 improvement in sensitivity. This includes the Einstein Telescope project \\cite{ET,ET2}, for which the design study was recently completed. These detectors are expected to open a new era in astronomy and astrophysics, providing new observations of various events and objects in the Universe, complementary to the standard electromagnetic observations.  Among the many sources of gravitational waves, the coalescences of binary systems, such as binary neutron stars (BNS), binary black holes (BBH), or a black hole and a neutron star (BHNS) stand out as the most likely candidates for first detections. These systems generate well understood \"chirp\" gravitational-wave signals, which have been computed using post-Newtonian approximation \\cite{phenom_waveform} or numerical relativity simulations \\cite{numrel}. One can then search for the chirp signals using matched template techniques - indeed a number of such searches have been performed using LIGO and Virgo data \\cite{CBC1,CBC2,CBC3}.  It has also been argued that adding the gravitational-wave signals from all binaries in the Universe will produce a gravitational-wave background (GWB) - for example, see \\cite{regman,regfrei,regchauv,regreview,ros11,mar11,BBHstoch} for the most recent studies in the context of terrestrial detectors. The LIGO and Virgo collaborations have developed techniques for searching for GWB by cross-correlating data from pairs of gravitational wave detectors \\cite{allenromano}. Such searches have also been performed using LIGO and Virgo data \\cite{S3stoch,S4stoch,S5stoch}, and have produced competitive upper limits on the energy density carried by gravitational waves.  The goal of this paper is to perform a detailed study of the accessibility of the GWB produced by the binary coalescences to the second and third generation gravitational-wave detectors. Our study follows the work of Regimbau and Mandic \\cite{regman}, and includes detailed scans of the parameter space in these models, as well as possible effects due to the uncertainty in the star formation rate and in time-delays associated with the formation of the binaries. We will show that this background is likely to be observed by the network of second generation detectors, and that the third-generation detectors will likely be able to explore most of the parameter space for these models. In Section 2 we summarize the calculation of the energy density for these models. In Section 3 we present results of our systematic study, and we include concluding remarks in Section 4. ",
        "conclusions": "In this paper we computed the gravitational wave background due to coalescences of binary neutron stars, binary black holes, and black hole - neutron star binaries, following the approach of \\cite{regman}. While such computations have been done in the past, in this study we performed a systematic scan of the parameter space, taking into account the possible variations in the result due to the choice of the star formation rate, and due to the choice of the distribution $P(t_d)$ of the delay time between the formation and coalescence of the binary. For each point in the parameter space, we compare the model prediction to the expected sensitivities of the second and third generation gravitational-wave detector networks. We find that models corresponding to the optimistic and realistic (in the case of BNS and BHNS models) local coalescence rates will be accessible to the second generation detector network. We also find that models corresponding to the pessimistic local coalescence rates will be accessible to the third-generation detector network. The binary coalescence GWB will, in fact, be a foreground for the third-generation detectors, and it will mask the GWB background due to early-Universe sources. Accessing the cosmological GWB with third generation detectors will therefore require identification and subtraction of all inspiral signals from all binaries in the relevant frequency band.  {\\it Acknowledgments:} The authors thank C. Belczinski, M. Dominik, and T. Bulik for informative discussions about the probability distribution for delay time, $P(t_d)$. CW and VM were supported in part by NSF grant PHY0758036."
    },
    "1112/1112.4473_arXiv.txt": {
        "abstract": "We present the discovery of a long-period, rapidly oscillating Ap star, \\hd. Using high-resolution time-series observations obtained with UVES at the ESO VLT telescope, we found radial velocity variations with amplitudes 7--150~\\ms\\ and a period of 23.6~min, exceeding that of any previously known roAp star. The largest pulsation amplitudes are observed for Eu\\iii, Ce\\iii\\ and for the narrow core of H$\\alpha$. We derived the atmospheric parameters and chemical composition of HD\\,177765, showing this star to be similar to other long-period roAp stars. Comparison with theoretical pulsational models \\mybf{indicates} an advanced evolutionary state for HD\\,177765. \\mybf{Abundance analyses of this and other roAp stars suggest} a systematic variation with age of the rare-earth line anomalies seen in cool Ap stars. ",
        "introduction": "\\label{intro}  The rapidly oscillating (roAp) stars are magnetic, chemically peculiar stars which pulsate in high-overtone acoustic  modes with typical periods of $\\approx$\\,10~min. These stars are located close to the instability strip crossing the main sequence between the early F and late A spectral types. First roAp pulsators were discovered by \\citet{1982MNRAS.200..807K}. Currently, about 40 of such stars are known.   Several excitation mechanisms were suggested in the past to drive pulsations in roAp stars  \\citep{1988IAUS..123..291D,1984trss.conf..346D,1983ApJ...275L...5S,1988MNRAS.235P...7M}.  Currently, it is widely accepted that the high frequency oscillations observed in these stars are excited by the opacity mechanism working on the hydrogen ionization region.  However, full non-adiabatic calculations show that the excitation of  high frequency acoustic oscillations by this mechanism can only be achieved in non-standard models, such as models with a modified T-tau relation \\citep{1998MNRAS.301...31G}, or models with envelope convection partially  or fully suppressed \\citep{2001MNRAS.323..362B,2005MNRAS.360.1022S}.Among these, models with convection suppressed seem to reproduce better the observed instability strip, however, the predicted red edge remains significantly hotter than the observed one.   The majority of roAp stars have been identified and analysed using high-speed photometric methods \\citep{2000BaltA...9..253K}. However, it has been realised that the ground-based photometry is not particularly suitable for discovering roAp stars. Time-resolved spectroscopy has a major advantage for detection of smaller-amplitude and longer-period pulsations. This has been demonstrated by, for example, the discovery of very low-amplitude pulsations in HD\\,75445 \\citep{2009A&A...493L..45K} and by the detection of long-period oscillations in HD\\,137909 \\citep[\\bCrB,][]{2004MNRAS.351..663H} and HD\\,116114  \\citep{2005MNRAS.358..665E}, which have been repeatedly classified as non-pulsating by photometric observations \\citep{1994MNRAS.271..129M,2005IBVS.5651....1L}.  The two latter objects, along with a faint roAp star KIC\\,10195926 found by the Kepler satellite \\citep{2011MNRAS.414.2550K}, form an unusual sub-group with pulsation periods of 16--21~min and spectroscopic properties different from the ``classical'', shorter-period roAp stars. Here we present the spectroscopic discovery of another member of this class, HD\\,177765, whose pulsation period of 24~min is the longest known for any roAp star. ",
        "conclusions": "\\label{concl}  We have analysed time-resolved spectra of the cool Ap star \\hd\\ obtained with the UVES instrument at VLT. The radial velocity and frequency analysis of these data reveals this object to be a roAp star with the pulsation period of 23.6~min. These oscillations, clearly present with amplitudes 40--150~\\ms\\ in the combined radial velocity curves of Ce\\ii, Eu\\ii\\ and Gd\\ii\\ lines, as well in the \\ha\\ core, occur with the longest known pulsational period for any roAp star. This discovery makes \\hd\\ a key object for testing predictions of pulsation theories because the frequency limits help in distinguishing alternative driving mechanisms and can provide useful asteroseismic constraints on the atmospheric and interior stellar structure.  No trigonometric parallax measurement is availble for \\hd. Therefore, we can only approximately place it in the HR-diagram for a comparison with pulsational models. Using the effective temperature, derived from photometry and spectroscopy, the pulsation frequency $\\nu=0.7$ mHz, and the evolutionary tracks from \\citet{2002MNRAS.333...47C}, which provide the frequency of the most unstable mode for models with given effective temperature and luminosity, we obtain a mass $M$\\,$\\approx$\\,2.2$M_\\odot$ and  a luminosity $\\log L/L_\\odot\\approx$\\,1.5. These parameters imply that the star is significantly evolved from the zero age main sequence, making it similar to the three other long-period, evolved roAp stars: \\bCrB, HD\\,116114, and KIC\\,10195926.  We have carried out an abundance analysis of \\hd\\ using the equivalent width and spectrum synthesis methods. A slow variation of the mean field modulus suggests a long rotational period. The diversity of Zeeman splitting patterns and the dependence of the abundance on the line strength and excitation potential (most clearly seen for Fe) indicates strong horizontal and vertical inhomogeneities.  We found that \\hd\\ is chemically similar to \\bCrB. Both stars are distinguished by a high Ce abundance and discordant abundances inferred from different ions of \\mybf{Eu and Ce}. But both stars lack a strong PrNd ionisation anomaly which is characteristic of higher frequency roAp stars. This anomaly is also absent in HD\\,116114. Thus, instead of being a spectroscopic signature of roAp stars as suggested by \\citet{2004A&A...423..705R}, the PrNd anomaly is probably related to the evolutionary state of roAp stars, distinguishing evolved, long-period pulsators from the shorter-period stars closer to the zero age main sequence. This finding represents one of the most compelling evidence for a systematic variation of surface chemical composition of Ap stars with age.  Longer time-series observations of \\hd\\ are highly desirable to determine additional pulsation frequencies which are necessary for an astroseismic analysis. Observations during a single night might be sufficient to resolve the expected large frequency separation of $\\Delta\\nu_0\\sim$\\,50~$\\mu$Hz."
    },
    "1112/1112.6247.txt": {
        "abstract": "We construct a simple model of the star-formation- (and resultant supernova-) driven mass and energy flows through the inner $\\sim$200 pc (in diameter) of the Galaxy. % Our modelling is constrained, in particular, by the non-thermal radio continuum and $\\gamma$-ray signals detected from the region. % The modelling points to a current star-formation rate of $0.04-0.12 \\msun$/year at $2\\sigma$ confidence within the region with best-fit value in the range $0.08-0.12 \\msun$/year which -- if sustained over 10 Gyr -- would fill out the $\\sim 10^9 \\msun$ stellar population of the nuclear bulge. % Mass is being accreted on to the Galactic centre (GC) region  at a rate $\\dot{M}_{IN} \\sim 0.3 \\msun/$year. % The region's star-formation activity drives an outflow of plasma, cosmic rays, and entrained, cooler gas. % Neither the plasma nor the entrained gas reaches the gravitational escape speed, however, and all this material fountains back on to the inner Galaxy. % The system we model  can naturally account for the recently-observed $\\gtrsim 10^6 \\msun$ `halo' of molecular gas surrounding the Central Molecular Zone out to 100-200 pc heights. % The injection of cooler, high-metallicity material into the Galactic halo above the GC may catalyse the subsequent cooling and condensation of hot plasma out of this region and explain the presence of relatively pristine, nuclear-unprocessed gas in the GC. % This process may also be an important ingredient in understanding the long-term stability of the GC star-formation rate. % The plasma outflow from the GC reaches a height of a few kpc and is compellingly related to the recently-discovered Fermi Bubbles by a number of pieces of evidence. % These include that the outflow advects precisely i) the power in cosmic rays required to sustain the Bubbles' $\\gamma$-ray luminosity in saturation; ii) the hot gas required to compensate for gas cooling and drop-out from the Bubbles; and iii) the magnetic field required to stabilise the walls of these structures. % Our modelling demonstrates that $\\sim 10^9 \\msun$ of hot gas is processed through the GC over 10 Gyr. % We speculate that the continual star-formation in the GC over the age of the Milky Way has kept the SMBH in a quiescent state thus preventing it from significantly heating the coronal gas, allowing for the continual accretion of  gas on to the disk and the sustenance of star formation on much wider scales in the Galaxy. % In general, our investigations  explicitly reveal  the GC's important role in the Milky Way's wider  stellar ecology. ",
        "introduction": "The inner 200 pc (in diameter) of the Milky Way features a spectacular confluence of unusual and energetic astrophysical phenomena. % Within this region of the Galaxy -- circumscribed by the Inner Lindblad Resonance associated with the non-axisymmetric gravitational potential of the Galactic bar -- the distribution of stars cusps sharply into the distinct population of the so-called nuclear bulge \\citep{Serabyn1996}. % Correspondingly, over the same region the current, inferred areal density of star formation, $\\dot{\\Sigma}_*$, sharply peaks to $\\sim 200 \\msun/$kpc$^2$/yr. % This is approximately three orders of magnitude higher than the mean value in the Galactic disk. % With such a high $\\dot{\\Sigma}_*$,  observations of the nuclei of external, star-forming galaxies tell us to expect a star-formation-driven outflow; there is much empirical and theoretical evidence that such an outflow exists in the GC as we have explored in a number of recent papers \\citep{Crocker2010b,Crocker2011a,Crocker2011b}. % This paper adds significantly to that evidence.   The very high star formation rate (SFR) density likewise sustains a very high energy density in all phases of the GC ISM. % Most directly, the optical and UV output of the many young, hot stars in the region is reprocessed by thick, ambient dust into a dominantly infrared photon background of $\\sim$100 eV cm$^{-3}$. % Radio continuum and $\\gamma$-ray observations allow one to place a robust lower limit of $\\sim$50 $\\mu$G on the typical magnetic field throughout the entire inner $\\sim$800 pc (in diameter) of the Galaxy \\citep{Crocker2010a}; modelling \\citep{Crocker2010b,Crocker2011b} points to a magnetic field in the inner 200 pc that is at least 100 $\\mu$G. % In association with and, in fact, as a necessary precondition to,  the high SFR, observations reveal an enormous agglomeration of hot, dense, and highly-turbulent molecular gas of mass $3 \\times 10^7 \\msun$ \\citep{Dahmen1998,Molinari2011}, 5-10\\% of the Milky Way's entire $H_2$ allocation. % This gas forms an asymmetric distribution extended along the plane referred to as the Central Molecular Zone \\citep[CMZ:][]{Serabyn1996}. % Recent infrared observations of the CMZ by $Herschel$ \\citep{Molinari2011} place much of this $H_2$ on a $\\sim 100 $ pc-radius ring, apparently akin to the star-forming rings observed in the nuclei of many face-on galaxies.  The CMZ $H_2$ is constantly bombarded by an extended, hard-spectrum cosmic ray ion population which results in a diffuse glow of hard-spectrum, $\\sim$TeV $\\gamma$-rays coextensive with the gas \\citep{Aharonian2006}. % GC X-ray observations \\citep{Koyama1989} {\\it apparently} reveal the existence of a very hot ($\\sim$ 7 keV), extended thermal plasma which would have an energy density similar to the lightfield, turbulent molecular gas, and magnetic field \\citep{Spergel1992}. %   It must also be remarked that the GC hosts  the Milky Way's resident supermassive ($\\sim4.3 \\times 10^6 \\msun$: \\citealt{Gillessen2009}) black hole (SMBH). % Though currently in a state of apparently unusual quiescence, this must certainly have been much more active at various times in the past \\citep[e.g.][]{Ponti2010}. % On the other hand, we have found from our recent work that the mechanical power delivered by supernovae -- occurring at a rate consistent with that pointed to by the region's current star-formation as traced by FIR emission -- is completely sufficient to sustain the currently-observed non-thermal emission ($\\sim$GHz radio continuum and $\\sim$TeV $\\gamma$-ray) from the $\\sim$200 pc scales of interest here. % Thus, from the point of view proffered by the non-thermal data, it is not {\\it necessary} that the SMBH have any significant role beyond the inner few pcs; our investigations below confirm this in general.     Finally, one of the most interesting recent discoveries in high energy astrophysics is of the `Fermi Bubbles', so-called because these structures were revealed \\citep{Su2010} in $\\sim$ GeV $\\gamma$-ray data collected by the $Fermi$-LAT. % The Bubbles are north-south symmetric about the Galactic plane and centred on the Galactic nucleus. % Given this morphology they are compellingly associated with some sort of activity in the GC. % Given, then, the Bubbles' large angular extent (they rise to $\\pm 50^\\circ$ in $b$) they are enormous structures extending  10 kpc from the plane. % The Bubbles' $\\gamma$-ray emission might be due to inverse Compton emission from a rather young (given short energy loss time) population of cosmic ray electrons. % Alternatively, the emission might be due to hadronic collisions experienced by a hard-spectrum cosmic ray proton and heavier ion population \\citep{Crocker2011a,Crocker2011c,Zubovas2011}. % We set out the evidence connecting the Fermi Bubbles with multi-Gyr-scale GC-star-formation-driven injection of cosmic ray {\\it protons} into the Galactic halo below.    \\subsection{Motivating Questions}  The preceding tour of GC and inner Galaxy phenomenology helps motivate a number of questions which we seek to address in this paper: \\begin{enumerate}  \\item{\\bf Gas accretion:} What is the rate at which the GC region typically accretes gas through the Galactic plane? % How do we understand the presence \\citep{Lubowich2000,Riquelme2010} of relatively pristine gas (i.e., that has undergone relatively little nuclear processing) in the GC?  \\item {\\bf Star formation:} What is the efficiency with which the GC converts gas into stars? % Given the unusual conditions in the GC environment, is GC star-formation biased towards the production of massive stars? %  \\item {\\bf GC ISM conditions:} Is the very hot plasma putatively revealed by X-ray observations real or not? % What contribution do the non-thermal ISM phases, in particular the cosmic rays, make to the overall energy density in the region? % What is the dynamical importance, if any, of the cosmic rays?  \\item {\\bf Outflows:}  There is multi-wavelength evidence (reviewed below) for outflows from the GC over size scales from pcs to 10 kpcs. % Are these outflows all different aspects of the same overarching phenomenon and how are they driven? % What is the wider importance of the GC outflow(s) to the Galactic ecology? % Is the material expelled from the nucleus lost to extra-galactic space or does it fountain back on to the Galactic disk? % How do the recently-discovered Fermi Bubbles \\citep{Su2010} relate to activity in the GC?  \\end{enumerate} More broadly, we aim in this paper to produce a first draft of a coherent explanation of all the disparate phenomena listed above that is itself physically plausible and motivated. % Overall, we shall see that it is  star-formation (driven by secular accretion processes over long timescales) -- rather than processes associated directly with the SMBH -- that seems to control the overall dynamics of the GC. %  \\subsection{Conventions and Assumptions}  We assume a distance to the Galactic centre of 8 kpc in this work. % We  use {\\tiny MATHEMATICA} notation: $f[x]$ denotes $f$ in a function of parameter $x$. % Formally,  the region we are investigating and attempting to model is that centred on $(l,b) = (0,0)$ and extending to $\\pm 0.8^\\circ$ in Galactic longitude and $\\pm 0.3^\\circ$ in Galactic latitude; this is the region for which the HESS telescope reported \\citep{Aharonian2006} a diffuse flux of $\\sim$TeV $\\gamma$-rays. ",
        "conclusions": "\\label{sctn_Conclusions}  We have modelled the mass and energy flows through the Galactic centre in a one-zone model and shown that  star-formation -- and resultant supernovae -- in the GC are well-sufficient to drive the gross dynamics of the system and to explain its non-thermal phenomenology. % None of this is to deny the importance of, e.g., stellar-radiation driving of gas dynamics in particular regions of the GC, but  radiation pressure is not required to explain the large-scale mass movements we infer here.  Current evidence does not seem strong enough for us to promote the empirically-inspired scaling of supernova mechanical energy with zero-age, main sequence progenitor mass that we have tested above; if anything, the `standard' $E_{SN} = 10^{51}$ erg assumption seems to work better in a number of cases (e.g., in suggesting a mass loading factor closer to expectation for star-burst-like environments; a SFR closer to other, independent estimates; supplying the total stellar mass of the nuclear stellar bulge for less extreme values of $\\dot{M}_{IN}$). % In any case, as emphasised above, the $E_{SN}$ scaling we adopted  certainly constitutes an upper limit to the true, population-averaged evolution of $E_{SN}$ with stellar mass. % Likely, further modelling, probably involving other constraints, will be needed before a better handle on  $E_{SN}[M_{ZAMS}]$ can be arrived at. % Given these considerations, we now particularise our discussion to the case of  $E_{SN} = 10^{51}$ erg.  A number of indicators come together to suggest that our control parameter, $\\dot{M}_{IN}$ -- the total mass being fed into the system -- has a lower limit at around 0.4 $\\msun$/year.% We find a 2$\\sigma$ upper limit on $\\dot{M}_{IN}$  (for the $E_{SN} = 10^{51}$ case) at 1.8 $\\msun$/year. % We emphasise, however, that  $\\dot{M}_{IN}$ is an overestimate of the total mass flux accreting out of the plane of the Galaxy on to the GC. % Our modelling suggests that there is an  outflow of plasma and cosmic rays from the system that entrains cool gas. % This entrained mass constitutes most of the mass flux but will fountain back on the GC. % This naturally accounts for the recently-observed halo of warm molecular hydrogen found to be surrounding the CMZ and may self-catalyse the accretion of relatively pristine corona gas into the system.    We find that our modelling robustly predicts an almost invariant $10^{39}$ erg/s for the power going into the freshly accelerated cosmic ray proton population in the GC region. % We nowhere constrain our model to produce this result: it emerges from the  numerical minimisation procedure. % As we have previously emphasised \\citep{Crocker2011a,Crocker2011b} this power is {\\it precisely} enough to sustain the $\\gamma$-ray emission from the Fermi Bubbles in a hadronic saturation scenario and allow for the inflation of the Bubbles against the pressure of the external medium in a few Gyr (probably assisted by the injected magnetic field and injected plasma). % Equally, we find that the modelling robustly predicts that the GC system injects $\\sim 10^{38}$ erg/s into hard-spectrum cosmic ray electrons; this is sufficient \\citep{Crocker2010a,Crocker2011b} to explain the non-thermal synchrotron radiation detected from the GC lobe and wider diffuse, non-thermal source region detected around the GC \\citep{LaRosa2005,Crocker2010a}.   Together with the evidence that the GC SFR has been quasi-stationary for Gyr timescales and that the outflow from the GC advects most non-thermal particles out of the acceleration region before they lose much energy, we consider our finding that the GC accelerates $10^{39}$ erg/s in cosmic rays as a very strong indication that GC star-formation essentially explains the Fermi Bubbles. % This argument is strengthened by the facts that the same star-formation processes can inject the plasma mass and thermal power required to fill-out the Bubbles and sustain their X-ray emission {\\it and} to inject the magnetic fields that can stabilise the Bubble surfaces against fluid instabilities, trap their cosmic ray and plasma contents for long timescales and explain their microwave synchrotron emission.  \\vspace{1 cm}  The Galactic centre is not particularly distinguished in the night sky -- but this belies its true activity: the many orders of magnitude of visual extinction arising from the column of dust we view it through is the reason for this $extrinsic$ dimness. % Equally, the GC is not particularly impressive as a non-thermal radiation source;  its $\\sim$ GeV $\\gamma$-ray luminosity, at few $\\times 10^{36}$ erg/s, is an order of magnitude short of the 5-10\\% of the Galaxy's output one might guess on the basis of the amount of massive star-formation happening in the system \\citep[the Galactic $\\sim$ GeV $\\gamma$-ray luminosity is about $3 \\times 10^{38}$ erg/s as inferred from fig.~1 of][]{Strong2010}. % As we have emphasised previously \\citep{Crocker2010b,Crocker2011a,Crocker2011b} {\\it this is because the cosmic rays accelerated in the region are mostly leaving before they can radiate}. % The radiation that these particles finally do emit is writ large in the Fermi Bubbles: it is from these structures that we detect the $\\sim$ 10\\% of Galactic $\\gamma$-ray luminosity (i.e., $4 \\times 10^{37}$ erg/s) that we expect on the basis of the GC's share of Galactic star-formation.  The fact that we can detect the Fermi Bubbles at all is testament to the long-term stability of the GC as  star-forming system. % We have seen hints above as to how this stability can be established, in particular, how GC star-formation activity can be insulated from the vicissitudes of conditions in the Galactic disk: it seems that a minimal level of accretion in the GC system is self-catalysed. % This is consistent with the presence of relatively pristine gas accreted out of the halo via a mechanism or mechanisms directly related to the star-formation-driven outflow (i.e., injection of cool, high metallicity gas and/or dust and/or shocks into the halo plasma).   We find complete consistency between the long-timescale-averaged power required to drive the GC system and the power  injected by the star-formation we can  infer is currently taking place in the system. % We need not invoke periods of very bright AGN-type activity of the supermassive black hole to explain the dynamics of the GC. % In fact, as a final speculation: these studies hint that the importance of sustained, GC star-formation is that it effectively erects a curtain wall around the SMBH, either using-up gas directly by creating new stars or blowing it away before much of it can reach Sgr A$^*$. % This prevents mass accretion at rates that would allow the SMBH to undergo phases of  activity sufficient to heat the Galaxy's coronal gas to such high temperatures that further accretion on the Galactic disk would become impossible\\footnote{We emphasise that this certainly does not imply there is {\\it no} accretion on to the SMBH or that it must be in a state of absolute quiescence.}. % This, in turn, enables  the long-term sustenance of {\\it disk} star formation   \\citep[cf.][]{Binney2011}. % Such a mechanism would explain the emerging finding \\citep{Erwin2011} -- to which the Milky Way adheres -- that the mass of a nuclear star cluster correlates with the {\\it total} stellar mass of its host galaxy rather than the galaxy's bulge."
    },
    "1112/1112.5955.txt": {
        "abstract": "We report on a Suzaku measurement of the shock feature associated with the western radio relic in the merging cluster A~3376\\@.  The temperature profile is characterized by an almost flat radial shape with $kT \\sim 4$ keV within $0.5 r_{200}$ and a rise by about 1 keV inside the radio relic. Across the relic region ($0.6-0.8 r_{200}$), the temperature shows a remarkable drop from about 4.7 keV to 1.3 keV\\@. This is a clear evidence that the radio relic really corresponds to a shock front possibly caused by a past major merger.  The observed sharp changes of the temperature and electron density indicate the Mach number ${\\cal M} \\sim 3$. The radial entropy profile is flatter than the prediction ($r^{1.1}$) of numerical simulations within $0.5 r_{200}$, and becomes steeper around the relic region. These observed features and time-scale estimation consistently imply that the ICM around the radio relic has experienced a merger shock and is in the middle of the process of dynamical and thermal relaxation. ",
        "introduction": "\\label{sec:intro} %%\\textcolor{red}{xx\u0082\u00cd\u0081AA~3667 paper\u0082\u00c9\u0082\u00c8\u0082\u00e8\u0082\u00dc\u0082\u00b7. \u0089\u00fc\u0092\u00f9\u0081A\u0092\u00c7\u0089\u00c1\u0082\u00b5\u0082\u00bd\u0095\u00b6\u008f\u00cd\u0082\u00cd\u0081A\u0090\u00d4 %%  \u008e\u009a\u0082\u00c5\u008e\u00a6\u0082\u00b5\u0082\u00c4\u0082\u00a2\u0082\u00dc\u0082\u00b7} \\textcolor{blue}{ \u008d\u00c4\u0089\u00fc\u0092\u00f9\u0082\u00b5\u0082\u00bd\u0095\u00aa\u0082\u00cd\u0081A\u0090\u00c2\u008e\u009a\u0082\u00c9\u0082\u00c8\u0082\u00e8\u0082\u00dc\u0082\u00b7\u0081B} %%   \\textcolor{green}{ \u008d\u00c4\u008d\u00c4\u0089\u00fc\u0092\u00f9\u0082\u00b5\u0082\u00bd\u0095\u00aa\u0082\u00cd\u0081A\u0090\u00c2\u008e\u009a\u0082\u00c9\u0082\u00c8\u0082\u00e8\u0082\u00dc\u0082\u00b7\u0081B} Clusters of galaxies are believed to grow through gas accretion from large-scale filaments and mergers of subclusters.  Cluster mergers are the most energetic events in the Universe after the Big Bang with the total kinetic energy of the colliding subclusters reaching $10^{65}$ ergs.  The kinetic energy is converted to thermal energy by driving shocks and turbulence. Existence of mega-parsec scale radio emission (halos and relics) in merging clusters indicates that those shocks and turbulence operate as main mechanisms of particle acceleration (see \\cite{ferrari08} and reference therein). Radio relics are considered to be the synchrotron emission generated through the interaction between relativistic electrons, accelerated by a shock, and amplified magnetic field in the intracluster medium (ICM)\\@. Although these radio features provide direct evidence of particle acceleration, the detection of non-thermal (hard) X-ray emission caused by relativistic particles is still controversial \\citep{nevalainen04,ajello09}. Currently, hard X-ray observations set lower limits of magnetic fields in ICM~\\citep{wik09,sugawara09, clarke06}. Based on recent studies, it is believed that magnetic field of 0.1--10 $\\mu$G exists in ICM\\@.   %To understand cluster formation, gas heatings and particle accelerations, %cluster mergers are very important fundamental process.  Theoretical studies of cluster merger shocks provide rich information about the magnetic field, turbulence, particle acceleration and their time evolution (see e.g.\\ \\cite{brueggen11} for a review). Merger shocks are characterized by a Mach number ${\\cal M} \\lesssim 3$, causing amplification of magnetic fields to $B \\sim 10\\ \\mu$G and production of relativistic particles \\citep{takizawa00,takizawa08}. Also, turbulence is considered to give a significant pressure support in the outer regions, with around 10\\% of the thermal pressure at least for a few Gyr after a major merger \\citep{parrish11}.  Recent numerical simulations include interaction of merger shocks with the filamentary cosmic web structures outside of clusters, and the results indicate that shock fronts are more enhanced in the filament direction \\citep{paul11}.   Though past X-ray observations have shown images and temperature structures of cluster mergers, there are only few clusters for which clear evidence of the shocks has been obtained (1E0657-56, A520, and A2146: \\cite{markevitch05, clowe06,russell10}). The low X-ray surface brightness in the outer regions has hampered precise measurements of gas temperature and density in the pre-shock region in particular. Recently, existence of shock fronts with a Mach number of about 2 was confirmed in the radio relic region of A~3667 \\citep{finoguenov10,akamatsu11b}.  Additionally, a remarkable radio relic indicated a shock front with a Mach number of 4.6 based on the measurement of radio spectral index (CIZA2242: \\cite{weeren10}). Thus, radio relics are the good probe to identify merger shocks and combination with X-ray data will yield important information about the gas dynamics associated with cluster mergers.  Since the radio relics are mainly found in the outskirts of clusters where the gas density is low, sensitive X-ray observations are needed, especially in the upstream side of the shocks.  The X-ray data will allow us to estimate the parameters of the gas and the magnetic field, and enable us to look into the actual process occurring during the cluster evolution.  In this paper, we present a new X-ray evidence of the shock associated with the radio relic in Abell 3376, a well-known merging cluster with irregular morphology at $z=0.046$. Previous studies showed the global mean temperature to be 4.0 keV and the existence of a pair of Mpc-scale radio relics revealed by 1.4 GHz VLA NVSS observations~\\citep{bagchi06}.  Another feature of A~3376 is the report of the hard X-ray signal with BeppoSAX/PDS~\\citep{nevalainen04}. %However, because of the Suzaku/HXD could't detect those signal\\citep{kawano09}, the detection of hard X-ray signal were still controversial. However, Suzaku HXD observation gave an upper limit which did not exclude the BeppoSAX flux \\citep{kawano09}.  We need to await further sensitive observations in the hard X-ray band.  We use $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M}=0.27$ and $\\Omega_\\Lambda = 0.73$, respectively, which give 54 kpc per arcminute at $z = 0.046$.  The virial radius is approximated by $r_{200} = 2.77 h_{70}^{-1} (\\langle T\\rangle /10 {\\rm keV})^{1/2}/E(z)\\ {\\rm Mpc} $, where $E(z)=(\\Omega_{\\rm M}(1+z)^{3}+1-\\Omega_{\\rm M})^{1/2}$ \\citep{henry09}.  For our cosmology and redshift, $r_{200}$ is 1.86 Mpc ($= \\timeform{34.6'}$) with $kT = 4.7$ keV\\@.  We employ solar abundance defined by \\citet{anders89} and Galactic absorption with $N_{\\rm H}=5.8 \\times10^{20}$cm$^{-2}$~\\citep{dickey90}.  Unless otherwise stated, the errors correspond to 90\\% confidence for a single parameter.  \\begin{table*}[t] \\label{tab:obslog} \\begin{center} \\caption{Suzaku observations of A~3376 and the background region.} \\begin{tabular}{cccccccccc} \\hline Target (Obs.\\ ID) & Position &Start time &Exposure$^\\ast$ &Exposure$^\\dagger$ \\\\ & (R.A., Decl.)   & (UT) & (ks) & (ks) \\\\ \\hline A~3376c (100034010) & $(90.56, -39.94)$ & 2005/10/06 14:46:08 & 110.4 & 78.0 \\\\ A~3376w (800011010) & $(90.05, -39.98)$ & 2005/11/07 14:15:05 & 120.0 & 99.9 \\\\ Q~0551-3637 (703036020) & $(88.19, -36.63)$ & 2008/05/14 13:25:55 & 18.8 & 15.3\\\\ \\hline $\\ast$: no selection with COR2 &\\multicolumn{2}{l}{ ${\\dagger}$: COR2 $> 8$ GV}\\\\\\ \\end{tabular} \\end{center} \\end{table*}  \\begin{figure}[t] \\begin{center} \\includegraphics[scale=0.45]{./a3376_0.5_8.0.ps} \\end{center} \\caption{ X-ray image of A~3376 in the energy band 0.5-8.0 keV, after subtraction of the NXB with no vignetting correction and after smoothing by a 2-dimensional Gaussian with $\\sigma =16$ pixel =$\\timeform{17''}$.  The large blue circle shows the virial radius of A~3376 with $r=34'.6$, and white dotted circles show the annular regions used for the spectral analysis.  Small magenta circles show point sources which are detected by {\\it wavdetect} (see text).  The VLA 1.4 GHz radio image is shown with green contours.  } \\label{fig:suzaku_image} \\end{figure} ",
        "conclusions": "\\label{sec:discussion} Suzaku performed 2 pointing observations of Abell 3376 along its merger axis.  The ICM temperature was found to be fairly flat at about 4 keV with a slow rise with radius in the region from the center to $0.6 r_{200}$~(1 Mpc).  This is followed by a sharp drop in the radio relic region, which implies there is a shock front.  We will attempt to evaluate the gas properties related with the shock (temperature, Mach number, magnetic field, and entropy) and discuss their implications. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %Temperature profiles %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsection{Temperature Profiles}\\label{sec:kt_comp} %\\textcolor{red}{ %In most of the relaxed clusters, the ICM temperature % structures  reflect the gravitational % potential and can be used to constrain the dark matter distribution. %} The ICM temperature is an important observational parameter which reflects the depth of the gravitational potential of the cluster.  Their profiles also enable us to look into the history of the ICM heating and the growth of clusters.  Previous Chandra, XMM-Newton results for other clusters indicated a significant temperature decline toward the outer regions.  Recent studies with Suzaku showed temperature profiles of ICM to the virial radius ($r_{200}$) for several clusters \\citep{george08, reiprich09, bautz09, hoshino10, kawaharada10, simionescu11}.  These profiles for relaxed clusters are approximated by a ``universal'' profile which declines by a factor of $\\sim 3$ at $r_{200}$ from the center.  On the other hand, merging clusters show complex temperature distribution depending on the mass ratio and collision geometry.  In the outer regions toward $r_{200}$, radial temperature profiles show some excesses and sharp drops which are not seen in relaxed systems (\\cite{akamatsu11a,akamatsu11b}). However, very little systematic studies have been made about the ICM properties in the outer regions of merging clusters.  We normalized the radial temperature profile by the average temperature determined within $0.4 r_{200}$.  Since A~3376 does not show a strong cool core, this should give a reasonable average temperature. %As mentioned above section, the $r_{200}$ derived from \\citet{henry09}. The resulting scaled temperature profile is shown in~Fig.\\ \\ref{fig:comp} along with the previous Suzaku results (black crosses and diamonds: taken from \\cite{akamatsu11a}).  Clearly, there is a marked difference between the profiles of relaxed and merging clusters.  The relaxed clusters (black crosses) generally show a smooth decline from the center to the outer regions, consistent with the result of hydrodynamic simulations~\\citep{burns10}.  Compared with these profiles, the A~3376 profile shows a large excess within $0.7 r_{200}$ before the sharp drop at $0.7-0.8 r_{200}$.  Numerical simulation of merging clusters by \\citet{paul11}shows a shock feature very similar to the one in A~3376. This simulation after about 3 Gyr after a major merger shows an arc-like high temperature region whose width is about 200 kpc, very similar to the observed temperature enhancement in Fig.\\ \\ref{fig:comp}.  This suggests that A~3376 is still evolving after several Gyrs elapsed from the last major merger.   The $\\Lambda$CDM model predicts that clusters grow through subcluster mergers and matter accretion from large-scale structures.  Those dynamical events should give strong impacts on the ICM temperature structure, such as seen in A~3376.  Since the ICM density is very low in the cluster outskirts ($n_e = 10^{-5}-10^{-3}\\rm\\ cm ^{-3}$), the radiative cooling and conduction times are expected to be very long with the cooling time largely exceeding the age of the Universe.  We can naturally expect that the temperature structure carries important information about the history of the cluster growth.  However, as mentioned above, there are no clear signs of dynamical evolution in relaxed clusters.  This indicates that the observed peculiar temperature structure in the merging clusters will be settled to the universal profile in a fairly short time scale, much shorter than the Hubble time.  Since cooling time is too long, we have to consider other possibilities such as (i) ICM cooling by adiabatic expansion, and (ii) ICM diffusion caused by pressure gradient. We will discuss these time scales in section~\\ref{sec:timescale}.   A fraction of the merger energy would be channeled into the acceleration of ultra-relativistic particles and amplification of magnetic fields. However, it is unknown that how these processes actually take place and distribute the merger energy to other forms in an efficient way. The turbulent motion can be detected through line broadening by X-ray microcalorimeters such as with SXS instrument on ASTRO-H \\citep{mituda10} and Athena, since X-ray calorimeters provide superior energy resolution by a factor of 20--30 better than that with the CCD instruments. The wealth of information will open a new window in our understanding of heating and particle acceleration caused by cluster merger shocks.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%Mach number!! %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsection{Mach Number} \\label{sec:shock} Recent studies with XMM-Newton and Suzaku showed sharp temperature drops at the region of the radio relic in A~3667~\\citep{finoguenov10, akamatsu11b}.  \\citet{finoguenov10} reported a sharp edge in the surface brightness at the outer boundary of the NW radio relic in A~3667\\@.  \\citet{akamatsu11b} reported a significant jump in the ICM parameters (temperature, surface brightness, electron density, and pressure) across the same relic.  The derived Mach number is ${\\cal M} \\sim 2$.  We estimate the Mach number based on the Suzaku data in the same way as \\citet{akamatsu11b}.  The Mach number can be obtained by applying the Rankine-Hugoniot jump condition,  \\begin{equation} \\frac{T_2}{T_1}=\\frac{5{\\cal M}^4+14{\\cal M}^2-3}{16{\\cal M}^2}, ~~ \\frac{1}{C}=\\frac{3}{4{\\cal M}^2}+\\frac{1}{4}, \\end{equation}  where $C=\\frac{n_{e1}}{n_{e2}}$ is the shock compression and subscripts 1 and 2 denote pre-shock and post-shock values, respectively, assuming the ratio of specific heats as $\\gamma=5/3$.  Here, we assume the regions \\timeform{21'}-\\timeform{24'} and \\timeform{27'}-\\timeform{31'} to be the post and pre shock regions, respectively. Because the compression factor, $C$, exceeds the value for the strong shock limit, we only derived mach number from temperature jump. Table~\\ref{tab:mach} shows the resultant Mach number, whose values are ${\\cal M}=2.94\\pm0.77$ based on the jumps of temperature. Compared with other clusters which also show shock fronts ~(A520, 1E0657-558, A2246, A~3667: ~\\cite{markevitch05,clowe06, russell10, akamatsu11b}), the Mach number in A~3376 is slightly larger than the other cases.  We derive the compression factor, $C = 7.8 ~\\pm~4.7$, which almost exceeds the value for the strong shock limit, even though with a large error.  The electron density estimated here is based on the assumption of spherical symmetry, which is not strictly correct when the shock propagates along the merger axis.  Since the spherical assumption relates the observed flux to a shell-like volume with a larger line-of-sight depth than the case of a relic volume, the density in particular for the pre-shock region is underestimated by a factor of roughly 1.5.  In addition to this uncertainty, we checked the dependence on the centroid position of the spherical shock.  We set the center of the sphere at the secondary X-ray peak which is closer to the radio relic.  The resultant compression parameter is $C=5.01$.  This indicates that the determination of the centroid have a significant impact on the shock parameter, which is mainly resulted from the estimated line of sight depth of the shocked region.   The shock front in A~3376 is located far from the cluster center~(1.2 Mpc $\\sim~0.6~r_{200}$) just as in A~3667\\@.  The temperature in the upstream region of the shock is $\\sim 1\\rm ~keV$, indicating the pre-shock sound speed to be $v_{\\rm ss}\\sim 520$ km~s$^{-1}$. Combining this with the shock compression $C$, we can evaluate the shock speed by $v_{\\rm shock}=C\\cdot v_{\\rm ss}$ to be $>2080$ km s$^{-1}$ assuming $C=4.0$ from the above discussion. The estimated shock speed well agrees with those in other clusters (1E0657-558: 4500 km s$^{-1}$, A520: 2300 km s$^{-1}$), but twice higher than the A~3667 case ($1360\\pm 120$ km s$^{-1}$). The shock speed is large enough to account for the radio relic by a merger shock model \\citep{takizawa00,ricker01, mathis05}. On the other hand, this $v_{\\rm shock}$ seems too high considering the ICM temperature of 4 keV\\@. We may need to consider a possibility that the true gas temperature is much higher than the measured electron temperature.  \\begin{figure*}[t] \\begin{tabular}{ccc} \\begin{minipage}{0.33333\\hsize} (a) Normalized temperature profiles \\begin{center} \\includegraphics[scale=0.32,angle=-90]{./comp-latest.ps} \\end{center} \\end{minipage} \\begin{minipage}{0.33333\\hsize} (b)  Entropy profiles \\begin{center} \\includegraphics[scale=0.32,angle=-90]{./entropy.ps} \\end{center} \\end{minipage} \\begin{minipage}{0.33333\\hsize} (c) Equilibrium time scale \\begin{center} \\includegraphics[scale=0.32,angle=-90]{./eqtime.ps} \\end{center} \\end{minipage} \\end{tabular} \\caption{\\label{fig:comp} Radial profiles of  normalized temperature, entropy and equilibration time scale. (a) Scaled projected temperature profiles compared with relaxed clusters~\\citep{akamatsu11a}. The profiles have been normalized to the mean (within 0.4 $r_{200}$) temperature. The $r_{200}$ value is derived from \\citet{henry09}. Dotted line shows simulation result~\\citep{burns10}. Two gray dashed lines show standard deviation. Black crosses are for the relaxed clusters. Diamonds show merger clusters (black:~A2142, blue:~A3667, red:~A3376). In entropy profile (b),  black solid line shows $K \\propto r^{1.1}$ given by~\\cite{voit03}. (c) Ion-electron  equilibration time scale. Horizontal scale is normalized by the radius of the radio relic. Solid curve shows the time after a shock heating assuming a constant shock speed $v=2000$ km s$^{-1}$. Red diamonds show the electron-ion equilibration time $t_{ie}$. Gray dashed diamond show the ionization time for $n_{e}t=3\\times10^{12} \\rm cm^{-3}s$. }  \\end{figure*}   \\begin{table}[t] \\begin{center} \\caption{Mach number estimation for A~3376 assuming $\\gamma=5/3$.} \\label{tab:mach} \\begin{tabular}{cccc}\\hline \\hline Parameter & Pre shock & Post shock & Mach number\\\\ \\hline $kT$ [keV] & $1.34\\pm 0.69$ & $4.68 \\pm 0.48$ & $2.97\\pm 0.77$ \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %Entropy profiles %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsection{Entropy Profile} \\label{sec:entropy}  Theoretical studies of the ICM~\\citep{tozzi01,voit03} predicts the entropy profile approximated by $r^{1.1}$, assuming that the gravitational energy of the accreting gas is efficiently converted into the thermal energy. However, Chandra observations revealed that the entropy in the central regions of relaxed clusters often showed a flat profile \\citep{cavagnolo09}.  The origin of this behavior is considered to be the feedback of star formation and central AGN activities, which generate extra entropy \\citep{peterson06}. Entropy profile is thus sensitive to the non-gravitational energy injection and can be used to explore such a process even in the outer regions of clusters.  We will define the ``entropy'' of the ICM by $K = kT\\cdot n_e^{-2/3}$.  Fig.\\ \\ref{fig:comp} shows the resultant entropy profile of A~3376.  The slope shows a large deviation from the standard value of 1.1 and close to 0.4 in the radius range \\timeform{0.5'}-\\timeform{20.0'}.  The solid black line shows $K\\propto r^{0.4}$, compared with the solid gray line indicating $K\\propto r^{1.1}$.  Recent XMM-Newton study reported entropy profiles of 31 clusters \\citep{pratt10}, and the distribution of the slope indicated a peak value of 0.98.  Therefore, the slope of 0.4 is very unusual and can be regarded as a transient value occurring in a certain phase of cluster mergers.  The slope shows a steepening and flattening across the radio relic, which itself indicates the shock heating.  A similar feature was also reported in A~3667 \\citep{akamatsu11b}.  The slope of A~3667 showed a marginal increase and a sudden drop across the radio relic region.  In A~3376, the increase of the entropy slope is more significant, which suggests a stronger shock than in A~3667\\@.  This probably corresponds to the feature that the estimated Mach number in A~3376 is factor of 2 higher that in A~3667\\@.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %Dynamical time scale %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsection{Relaxation Time Scales} \\label{sec:timescale} As shown in the previous sections, A~3376 exhibits distinct ICM properties which are significantly different from those seen in the relaxed clusters.  The density and temperature structures around the shock region, along with the irregular morphology of this cluster, suggest that A~3376 is a young system regarding the relaxation of the ICM\\@. The fact that most of the relaxed clusters now attain almost universal temperature profile indicates that the time needed for the ICM to reach thermal and dynamical equilibrium is much shorter than the cluster life.  In this section we evaluate several relevant time scales based on the present observational quantities.  We first estimate the time in which the shock front reaches the current position.  Assuming that the shock has traveled the distance between the cluster center and the relic position with a constant velocity ($v = 2000$ km s$^{-1}$), the time required to propagate this length of $\\sim 1.5$ Mpc is 0.32 Gyr.  Next, we estimate the time scale to attain thermal equilibrium after the shock heating. The shock heating should first act on ions rather than electrons, and then the thermal energy is transferred from ions to electrons through Coulomb collisions.  The equilibration time scale between ion and electron is 3 orders of magnitude longer than the electron-electron timescale.  Therefore, the cluster thermal time scale after the shock heating is limited by the electron-ion time scale.  According to \\citet{takizawa98}, the electron-ion equilibration timescale is given by \\begin{equation} t_{\\rm ie}= 2\\times 10^{8}{\\rm\\ yr} \\left(\\frac{n_{e}}{10^{-3}{\\rm\\ cm^{-3}}}\\right)^{-1} \\left(\\frac{T_{e}}{10^{8}\\rm\\ K}\\right)^{3/2} \\left(\\frac{\\ln \\Lambda}{40}\\right), \\end{equation} where $\\ln \\Lambda$ denotes the Coulomb logarithm \\citep{spitzer56}.  In the central region, this gives 0.03 Gyr, which is about 10 times shorter than the elapsed time after the shock passage. We calculated the elapsed time after the shock passage and the time for electron-ion equilibration $t_{\\rm ie}$ as a function of radius.  The resultant time scales are shown in Fig.\\ \\ref{fig:comp} (c). Here, red diamonds show $t_{\\rm ie}$, and gray dashed diamonds are the ionization time scale  $t_i$ of Fe-K$\\alpha$ assuming $n_et=3\\times10^{12}$ cm$^{-3}$~s.  Solid line shows the elapsed time after the shock passage, assuming the shock propagation from inner to outer regions.  We can see that the elapsed time is shorter than $t_{\\rm ie}$ outside of $0.6 r_{\\rm shock}$ corresponding to $17'$ (920 kpc).  This suggests that the electron temperature is likely to be lower than the ion temperature in the outer region including the radio relic \\citep{akahori08, akahori11}. Those sign already reported another merging cluster RX J1347.5-1145~\\citep{ota08}. In the region that $t_{i}$ exceed the sound crossing time $t_{sc}$ ( $t_{i} >~t_{sc}$), there will be non-equilibrium ionization state. The future high resolution X-ray spectroscopy can reveal those dynamical ionization states.  Finally, we estimate the time scale in which the sharp change of the temperature disappears.  Those high temperature region will diffuse out due to the high pressure in the post-shock region and reach equilibrium in the sound crossing time ($t_{sc}=R/v_{s}$), where $R$ is the width of the high temperature region and $v_{s}$ is the sound speed, respectively.  Using the observed temperature ($kT= 4.7$ keV which corresponds to $v_{s}=1100$~km s$^{-1}$) and the width of the high temperature region ($R=320$ kpc), we can calculate the sound crossing time.  The resultant time scale is $t_{sc}=0.28$ Gyr.  In summary of this section, the observed features of temperature change, entropy slope and the time scale consideration all point to that the outer region ($> 1$ Mpc) of A~3376 is not reaching the equilibrium. In this sense, A~3376 is still a very young system in view of the cluster evolution.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %Summary!!!! %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1112/1112.3645_arXiv.txt": {
        "abstract": "LHC-7 has narrowed down the mass range of the light Higgs boson.  This result is consistent  with the supergravity unification framework, and the current Higgs boson mass window implies a rather significant loop correction to the tree value, pointing to  a relatively heavy scalar sparticle spectrum with universal boundary conditions.  It is shown that the largest value of the Higgs boson mass is obtained on the  Hyperbolic Branch  of radiative breaking.  The  implications of  light Higgs boson in the broader mass range of 115 GeV to 131 GeV and a narrower range of 123 GeV to 127 GeV are explored  in the context of the  discovery of supersymmetry at LHC-7 and for the observation of dark  matter in direct detection experiments. ",
        "introduction": "} In models based on supersymmetry the light Higgs boson~\\cite{HiggsBoson} has a predictive mass range, and recently LHC-7 has stringently constrained the light Higgs boson to lie in the $115\\GeV$ to $131\\GeV$ range~(ATLAS) and the $115\\GeV$ to $127\\GeV$ range~(CMS) at the $95\\%$~C.L.~\\cite{dec13} with possible hints of evidence within a few  GeV of $125\\GeV$. This mass window lies in the  range predicted by supergravity unification  {(SUGRA)}~\\cite{can} (for reviews see~\\cite{Nath,IbanezRoss,BSM}). In this work we investigate supergravity model points that are consistent with the mass range given by  the new  LHC-7 data~\\cite{dec13} (for a previous work on the analysis of the Higgs boson in SUGRA and string models pointing to a heavier Higgs in the $120\\GeV$ range see~\\cite{fln-higgs}).  LHC-7 has made great strides in exploring the parameter space of supersymmetric models. Indeed,  early theoretical projections   for the  expected reach in sparticle masses and in the  $m_0-m_{1/2}$   plane for LHC-7~\\cite{arXiv:0912.4217,arXiv:1002.2430,Baer:2010tk,arXiv:1008.3423} have been met and exceeded by the $1\\fb$ and $2\\fb$  LHC-7 data~\\cite{cmsREACH,AtlasSUSY,atlas0lep,atlas165pb,atlas1fb}. { The implications of the new LHC results have been analyzed}  by a number of authors in the context of lower limits on  supersymmetric particles and in connection with dark matter~\\cite{LHC-7,Akula:2011zq,Akula:2011dd,Akula:2011ke,Akula2,Grellscheid:2011ij,Ellwanger:2011mu}. Now the most recent results from CERN~\\cite{dec13} indicate that { the two detectors,  ATLAS and CMS, have collected  as much as  $5\\fb$ of data.} One of the most interesting implications of the LHC-7 data concerns the constraints it imposes on the Higgs boson  mass.    As mentioned above we will work within the framework of a supergravity grand unification model with universal boundary conditions~\\cite{can,hlw,nac}. Here we discuss the dependence of the  light Higgs boson mass on the parameter space, i.e., on $m_0, m_{1/2}, A_0, \\tan\\beta$~\\cite{an1992}, where $m_0, m_{1/2}$ and $A_0$ are  the parameters at the GUT scale, where the GUT scale, { $M_\\mathrm{GUT} \\sim 2 \\times10^{16}\\GeV$} is defined as the scale at which the gauge couplings unify, and where $m_0$ is soft scalar mass, $m_{1/2}$,  the gaugino mass, $A_0$, the trilinear coupling and $\\tan\\beta$,  the ratio of the two Higgs VEVs in the minimal supersymmetric standard model.  An important aspect of SUGRA models is that the radiative electroweak symmetry breaking, REWSB, is  satisfied for  $A_0/m_0$ {typically in the $-5$ to $5$ range}. The renormalization group evolution then leads to a value of the trilinear coupling, $A_t$, at the electroweak scale to  also be  $\\mathcal{O}({\\rm TeV})$. The relevance of this observation is that quite generically supergravity unification { leads }  to a sizable $A_t$    which is needed to give a substantial leading order loop correction to the Higgs Boson mass  for any fixed $\\mu, \\tan \\beta$ and $m_0$, where $\\mu$ is the Higgs mixing parameter in the superpotential. Thus  a generic prediction of SUGRA  models under radiative electroweak symmetry breaking  for  a sizable $A_0/m_0$  is that there would be a substantial loop correction to the Higgs boson mass, and it is well known that the light Higgs mass at the tree level has the value   $\\mh \\leq M_Z$ and there is a significant loop correction  $\\Delta \\mh$ to lift it above $M_Z$~\\cite{1loop,hep-ph/9210242,Casas:1994us,Carena:1995wu,Espinosa:2000df,Carena:1998wq,hep-ph/0503173}.  The dominant one loop contribution arises from the top/stop sector and is given by \\beqn \\Delta \\mh^2\\simeq  \\frac{3m_t^4}{2\\pi^2 v^2} \\ln \\frac{M_{\\rm S}^2}{m_t^2} + \\frac{3 m_t^4}{2 \\pi^2 v^2}  \\left(\\frac{X_t^2}{M_{\\rm S}^2} - \\frac{X_t^4}{12 M_{\\rm S}^4}\\right)~, \\label{tloop} \\eeqn where  $v=246\\GeV$, $M_{\\rm S}$ is an average stop mass, and $X_t$ is given  by \\beqn X_t\\equiv A_t - \\mu \\cot\\beta~. \\label{xt} \\eeqn From Eq.~(\\ref{tloop}) one finds that the loop correction is maximized when \\be X_t \\sim \\sqrt 6 M_{\\rm S} ~. \\label{xt} \\ee We note that there can be  important  loop corrections also from the $b$-quark sector and a correction similar to Eq.~(\\ref{tloop}) can be written where $X_t$ is replaced by $X_b= A_b- \\mu \\tan\\beta$ along with other appropriate replacements. Thus when $\\mu \\tan\\beta$  becomes large,  the $b$-quark contribution to the loop correction, which is proportional to powers of $X_b$, becomes large and is comparable to the top contribution which implies  that a high Higgs mass can also result in stau-coannihilation  models where typically $\\mhf$ is large  and  $m_0$ is relatively small.  Further, we note that the  approximation  of Eq.~(\\ref{xt}) would not hold if the off-diagonal elements of the stop mass squared matrix are comparable to the diagonal elements  which can happen for very large $A_t$. In addition, it is well known that the two loop corrections are substantial (see e.g.~\\cite{Slavich2} for a numerical analysis). While the correction at the one loop level has the symmetry $X_t \\to - X_t$, this symmetry is lost when the two loop corrections are included and then $\\sgn\\left(A_0/m_0\\right)$ plays an important role in the corrections  to the Higgs boson mass. As seen later this observation is supported by the full numerical analysis which includes the two loop corrections. We note in passing that the theoretical predictions for the light Higgs boson mass depend sensitively on  the input parameters  which include the gauge coupling constants as well as the top mass with their experimental errors. Additionally, there are also inherent theoretical uncertainties which together with the uncertainties of the input parameters allow theoretical predictions of the light Higgs boson mass  accurate to only within an error corridor of a few GeV (see e.g.~\\cite{Slavich2}).   Since the loop corrections involve the sparticle spectrum, a large loop correction implies a relatively heavy sparticle spectrum and specifically heavy scalars. Such a possibility arises in REWSB which allows for scalars  heavier than $10\\TeV$~\\cite{Chan:1997bi}. Specifically, with scalars approaching $10\\TeV$,  the Higgs boson mass can remain heavy while the gaugino sector is free to vary. This occurs within the minimal SUGRA framework and similar situations arise in other works of radiative breaking~\\cite{Feldman:2011ud,Baer}.   Indeed, quite generally in SUGRA and string models with the MSSM field content,  the analysis of the Higgs mass with loop {corrections}  under the constraints of REWSB { gives an upper} limit on the light Higgs  boson mass of about  { $135\\GeV$ for a  wide range of input parameters}.\\footnote{ We note that  heavier Higgs boson  masses can be obtained in  a variety of different models such as hierarchical breaking models~\\cite{Wells1,adgr,Kors:2004hz} (for recent work see~\\cite{Cabrera:2011bi,arXiv:1108.6077}) or by addition of vector like multiplets~\\cite{Babu:2008ge}.} {A very} interesting aspect of the  recent LHC-7 data concerns the fact that a large portion of the Higgs boson mass window has been excluded and what remains is consistent with the  range predicted by the SUGRA models. ",
        "conclusions": "Recent data from  LHC-7 indicates  a narrow window on the light Higgs mass. This  allowed mass window is consistent with the range predicted by SUGRA models and specifically by the mSUGRA model.  Here we discussed the implications of the indicated mass range for the light Higgs mass for the sparticle mass spectrum and for dark matter. Using the allowed Higgs mass range above $115\\GeV$    the corresponding ranges for the soft masses and couplings, as well as the ratio of the vacuum expectation values of the Higgs doublets and the Higgsino mass parameter were found.  We then investigated the ranges for the sparticle masses correlated to the predicted value of the Higgs Boson mass, specifically for the chargino, the neutralino, the gluino, the stop, the stau, for the first and second generation squarks and sleptons {and for the heavier Higgs of  the minimal supersymmetric standard model, i.e., the CP odd Higgs $A^0$, the CP even Higgs $H^0$, and the charged Higgs $H^{\\pm}$.}  {Our conclusions are}   that the largest Higgs masses are realized on the Focal Surface {of the Hyperbolic Branch} of radiative electroweak symmetry breaking.   {We also point out that low values of $\\mu \\sim 150 \\GeV$ are consistent with heavy squarks and sleptons in the 10 TeV region or  larger.   We find that $\\mh~\\in (123-127)\\GeV$   does allow for light third generation stop as low as  $m_{\\tilde t_1} > 230\\GeV$, though the second generation squarks are at least $m_{\\tilde q} > 1.5\\TeV$  and second generation sleptons are at least 475 GeV.} { Thus,  the restriction of the light Higgs boson to the mass window  $\\mh~\\in (123-127)\\GeV$ provides further  constraints on the sparticle spectrum that are complimentary to the direct searches for sparticles at the LHC.  }  { Further, we find precise predictions for dark matter  if the light Higgs boson mass lies between $123\\GeV$ and $127\\GeV$.} For these light Higgs boson masses, the corresponding range of the lightest neutralino  mass  would be accessible in the next generation of direct detection dark matter experiments. The light Higgs boson in the $123\\GeV$ and $127\\GeV$ range was shown to be generic for the case of heavy scalars in minimal supergravity with {$|A_0/m_0|\\sim\\mathcal{O}(1)$.}"
    },
    "1112/1112.1689_arXiv.txt": {
        "abstract": "\\noindent A stochastic gravitational wave background (SGWB) would gravitationally lens the cosmic microwave background (CMB) photons. We correct the results provided in existing literature for modifications to the CMB polarization power spectra due to lensing by gravitational waves (GW). Weak lensing by gravitational waves (GW) distorts all the four CMB power spectra, however its effect is most striking in the mixing of power between the E-mode and B-mode of CMB polarization. This suggests the possibility of using measurements of the CMB angular power spectra to constrain the energy density ($\\Omega_{GW}$) of the SGWB. Using current data sets (QUAD, WMAP and ACT), we find that the most stringent constraints on the present $\\Omega_{GW}$ come from measurements of the angular power spectra of CMB temperature anisotropies. In near future more stringent bounds on $\\Omega_{GW}$ can be expected with improved upper limits on the B-modes of CMB polarization. Any detection of B-modes of CMB polarization above the expected signal from large scale structure(LSS) lensing could be a signal for a SGWB. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.4586_arXiv.txt": {
        "abstract": "We study shear viscosities of different species in hot and neutrino-trapped dense matter relevant to protoneutron stars. It is found that the shear viscosities of neutrons, protons and electrons in neutrino-trapped matter are of the same orders of magnitude as the corresponding shear viscosities in neutrino-free matter. Above all, the shear viscosity due to neutrinos is higher by several orders of magnitude than that of other species in neutrino-trapped matter.  Next we investigate the effect of shear viscosity in particular, neutrino shear viscosity on the thermal nucleation rate of droplets of antikaon condensed matter in protoneutron stars. The first-order phase transition from hadronic matter to antikaon condensed matter is driven by the thermal nucleation process. We compute the equation of state used for the calculation of shear viscosity and thermal nucleation time within the relativistic mean field model. Neutrino shear viscosity enhances the prefactor in the nucleation rate by several orders of magnitude compared with the $T^4$ approximation of earlier calculations. Consequently the thermal nucleation time in the $T^4$ approximation overestimates our result. Furthermore, the thermal nucleation of an antikaon droplet might be possible in neutrino-trapped matter before neutrino diffusion takes place. \\pacs{97.60.Jd, 26.60.-c,52.25.Fi,64.60.Q-} ",
        "introduction": "A first order phase transition from nuclear matter to some exotic form of matter might be possible in (proto)neutron stars. It could be either a nuclear to quark matter transition or a first order pion/kaon condensation. Consequently, it might have tremendous implications for compact stars \\cite{Gle} and supernova explosions \\cite{Sag}. Here the focus is the first order phase transition proceeding through the thermal nucleation of a new phase in particular, antikaon condensed phase in hot and neutrino-trapped matter. After the pioneering work by Kaplan and Nelson on antikaon ($K^-$ meson) condensation in dense baryonic matter formed in heavy ion collisions as well as in neutron stars \\cite{Kap}, several groups pursued the problem of antikaon condensation in (proto)neutron stars \\cite{Bro,Tho,Ell,Lee,Pra97,Gle99,Kno,Sch,Pal,Bani1,Bani2,Bani3,Bani4,Pons,Bani5}. In most cases, the phase transition was studied using either Maxwell construction or Gibbs' rules for phase equilibrium coupled with global baryon number and charge conservation \\cite{Gle92}. The first order phase transition driven by nucleation of antikaon condensed phase was considered in a few cases \\cite{Nor,Bani7}. In particular, the calculation of Ref.\\cite{Bani7} dealt with the role of shear viscosity on the the thermal nucleation of antikaon condensed phase in hot and neutrino-free compact stars \\cite{Bani7}. It is to be noted here that the first order phase transition through the thermal nucleation of quark matter droplets was also investigated in (proto)neutron stars \\cite{Nor,San,Sat,Hei,Bom,Min,Bom2} using the homogeneous nucleation theory of Langer \\cite{Nor,San,Lan}. The thermal nucleation is an efficient process than the quantum nucleation at high temperatures \\cite{Bom,Bom2}.  We adopt the homogeneous nucleation theory of Langer \\cite{Lan,Tur} for the thermal nucleation of antikaon condensed phase. Nuclear matter would be metastable near the phase transition point due to sudden change in state variables. In this case thermal and quantum fluctuations are important. Droplets of antikaon condensed matter are formed because of thermal fluctuations in the metastable nuclear matter. Droplets of the new and stable phase which are bigger than a critical radius, will survive and grow. The transportation of latent heat from the surface of the droplet into the metastable phase favours a critical size droplet to grow further. This heat transportation could be achieved through the thermal dissipation and viscous damping \\cite{Tur,Las,Raj}.  A parametrised form of the shear viscosity was used in earlier calculations of the nucleation of quark matter \\cite{Bom}. Recently, the influence of thermal conductivity and shear viscosity on the thermal nucleation time was studied in a first-order phase transition from nuclear to antikaon condensed matter in hot neutron stars \\cite{Bani7}. The shear viscosity due to neutrinos was not considered in that calculation. It would be worth studying the effect of shear viscosity on the thermal nucleation rate of droplets of antikaon condensed matter in neutrino-trapped matter relevant for protoneutron stars. Besides shear viscosities due to neutrons, protons and electrons, this involves the contribution of neutrinos to the total shear viscosity. Shear viscosities of pure neutron and neutron star matter were calculated by several groups \\cite{Flo1,Flo2,Cut,Ben,Yak,Glam}. We also performed the calculation of shear viscosity in neutron star matter using the equation of state (EoS) derived from relativistic field theoretical models \\cite{Bani6}. Transport properties of degenerate neutrinos in dense matter were estimated by Goodwin and Pethick \\cite{Good}.  We organise the paper in the following way. We describe models for EoS, shear viscosities of different species including neutrinos and the calculation of thermal nucleation rate in Sec. II. Results of this calculation are discussed in Sec. III. Sec. IV gives the summary and conclusions. ",
        "conclusions": "We have studied shear viscosities of different particle species in neutrino-trapped $\\beta$-equilibrated and charge neutral nuclear matter. We have derived equations of state of nuclear and antikaon condensed phases in the relativistic mean field model for the calculation of shear viscosity. It is noted that neutrons, protons and electrons come into thermal equilibrium in the weak interaction time scale. The shear viscosity due to neutrinos is calculated treating other particles as background and found to dominate the total shear viscosity.  Next we have investigated the first-order phase transition from neutrino-trapped nuclear matter to antikaon condensed matter through the thermal nucleation of a critical droplet of antikaon condensed matter using the same relativistic EoS as discussed above. Our emphasis in this calculation is the role of the shear viscosity due to neutrinos in the prefactor and its consequences on the thermal nucleation rate. We have observed that the thermal nucleation of a critical antikaon droplet might be possible well before the neutrino diffusion takes place. Furthermore, a comparison of our results with that of the calculation of thermal nucleation time in the $T^4$ approximation shows that the latter overestimates our results of thermal nucleation time computed with the prefactor including the neutrino shear viscosity. Though we have performed this calculation with antikaon optical potential depth of $U_{\\bar K} (n_0) = -120$ MeV, we expect qualitatively same results for other values of the antikaon optical potential depth \\cite{Bani8}."
    },
    "1112/1112.1476_arXiv.txt": {
        "abstract": "Cosmic rays fill up the entire volume of galaxies, providing an important source of heating and ionisation of the interstellar medium, and may play a significant role in the regulation of star formation and galactic evolution. Diffuse emissions from radio to high-energy \\grays{} ($> 100$ MeV) arising from various interactions between cosmic rays and the interstellar medium, interstellar radiation field, and magnetic field, are currently the best way to trace the intensities and spectra of cosmic rays in the Milky Way and other galaxies. In this contribution, I describe our recent work to model the full spectral energy distribution of galaxies like the Milky Way from radio to \\gray{} energies. The application to other galaxies, in particular the Magellanic Clouds and M31 that are detected in high-energy \\gray{s} by the Fermi-LAT, is also discussed. ",
        "introduction": "The luminosity of a star-forming galaxy like the Milky Way (MW) is dominated by the relatively narrow frequency range of the spectral energy distribution (SED) from the ultraviolet (UV) to far infrared (FIR), which is due to stellar emission and dust reprocessing in the interstellar medium (ISM). Related to the birth and death of massive stars, cosmic rays (CRs) are pervasive throughout the ISM (see, e.g.,~\\cite{Strong2007}~for a recent review). The diffuse emissions arising from various interactions between CRs and the ISM, interstellar radiation field (ISRF -- the UV--FIR component of the galactic SED), and magnetic field span radio frequencies to high-energy \\grays{} ($>$100 MeV), but at a lower level of intensity compared to the stellar and dust component. These broadband emissions are currently the best way to trace CR intensities and spectra throughout the MW and other galaxies. Gamma rays are particularly useful in this respect because this energy range gives access to the dominant hadronic component in CRs via the observation of $\\pi^0$-decay radiation produced by CR nuclei inelastically colliding with the interstellar gas. Understanding the global energy budget of processes related to the injection and propagation of CRs, and how the energy is distributed across the electromagnetic spectrum, is essential to interpret the radio/far-infrared relation (\\cite{Helou1985,Murphy2006}), galactic calorimetry (e.g.,~\\cite{Volk1989}), and predictions of extragalactic backgrounds (e.g.,~\\cite{Thompson2007,Murphy2008}), and for many other studies. ",
        "conclusions": ""
    },
    "1112/1112.1979_arXiv.txt": {
        "abstract": "The Large Area Telescope on board the \\textit{\\textit{Fermi}} satellite observed a gamma-ray flare in the Crab nebula lasting for approximately nine days in April of 2011. The source, which at optical wavelengths has a size of $\\approx$11 ly  across, doubled its gamma-ray flux within eight hours.  The peak photon flux was $(186 \\pm 6) \\times 10^{-7}$ cm$^{-2}$ s$^{-1}$ above 100 MeV, which corresponds to a 30-fold increase compared to the average value. During the flare, a new component emerged in the spectral energy distribution, which peaked at an energy of (375 $\\pm$ 26) MeV at flare maximum. The observations imply that the emission  region was likely relativistically beamed toward us and that variations in its motion are responsible for the observed spectral variability. ",
        "introduction": "The Crab nebula is the remnant of a supernova observed in 1054 AD. The explosion left behind a rotating neutron star emitting electromagnetic radiation pulsed at the rotation period, that powers a wind of relativistic particles. These particles interact with the remnant gas and magnetic field, causing the nebula to glow brightly at all wavelengths, predominantly by synchrotron radiation. The spectral energy distribution (SED) of the nebula is, accordingly, dominated by a synchrotron component extending from radio wavelengths into the gamma-ray band \\citep{hester_crab_2008}. Above 450 MeV a second component emerges, attributed to inverse-Compton scattering by the same relativistic particles \\citep{gould_1965,dejager_1992,atoyan_fluxes_1996}. The angular size of the Crab nebula is $\\approx$0.1$^\\circ$ in the optical and smaller at higher energies. This corresponds to 3.5 pc, or 11 ly, at its estimated distance of 2 kpc \\citep{trimble_distance_1973}.  Today, the pulsar and nebula (henceforth referred to together as the Crab) are considered prime examples of non-thermal sources in the Universe and serve as a laboratory for relativistic plasma physics. New puzzles for our understanding of the Crab have been posed by the detection of three bright gamma-ray flares by the AGILE satellite and the Large Area Telescope (LAT) on board the \\textit{Fermi} satellite between 2007 and 2010 \\citep{abdo_crab_2011,tavani_crab_2011}. During these flares the unpulsed component of the gamma-ray flux increased by a factor of {$\\approx$10} on time scales as short as 12 hours \\citep{balbo_crab_2011}, while the period and flux of the pulsed component remained stable.  More recently, in April of 2011 the LAT detected a fourth flare, three times brighter than any of the previous ones \\citep{buehler_crab_2011,hays_crab_2011}. The flare was swiftly confirmed by the AGILE satellite \\citep{striani_crab_2011}. Observations at lower energies, in particular by the Chandra X-ray observatory, have not yet revealed any variability correlated with the gamma-ray flares \\citep{tennant_2011}. However, analysis of these observations is ongoing and will be discussed elsewhere. Here we present the LAT gamma-ray results obtained during the flare, together with a broader analysis of the first 35 months of Crab observations by \\textit{Fermi}. ",
        "conclusions": "A year after first being reported, the gamma-ray flares from the Crab remain enigmatic. Where within the nebula does the emission come from? What produces the flux variations? How were the emitting particles accelerated? How are the flares related to the variability observed on yearly and monthly time scales? Although several ideas have been proposed, no certain answers can be given today. The observations presented here give us the most precise look into the flare phenomenon to date, during the brightest outburst detected so far. We will proceed to discuss some of the implications and challenges posed by these observations.  One striking property of the Crab nebula flares is their rapid flux variations, doubling within {t$_{d} < 8$ hours} at the rise of the 2011 April flare. Causality arguments imply that the emission region is compact, with a length {$L < c t_{d} \\approx 2.8 \\times 10^{-4}$ pc}. The emitted isotropic power at the peak of the flare of $\\approx4 \\times 10^{36}$ erg s$^{-1}$ corresponds to $\\approx$1\\% of the total spin-down power of the pulsar, the ultimate energy source of the nebula.  It is difficult to explain how this energy is focused into such a small emission volume. The focusing is generally easier to explain when the emission site is closer to the pulsar. The absence of pulsation in the flare signal implies that the emission region is at least located outside the light cylinder of the pulsar. Another possibility to explain the flare brightness is that the emission is highly anisotropic, as would be expected if the emission region moves relativistically toward us. While only mildly relativistic motion with velocities of  {$\\approx$0.5c} are observed inside the nebula \\citep{scargle_activity_1969,hester_hubble_2002,melatos_2005}, relativistic motion is expected in the pulsar wind and in the downstream medium behind the wind termination \\citep{camus_2009}. Relativistic bulk motion is particularly expected at the ``arch shock'' of the wind termination, which has been proposed as the main site of gamma-ray emission \\citep{komissarov_2011}.  The flare emission is expected to result from synchrotron radiation by relativistic electrons and positrons (henceforth referred to together as electrons) \\citep{abdo_crab_2011}. A new spectral component emerges in the SED during the flare. The hard photon spectrum ($\\gamma \\approx 1.3$) of the flaring component implies that most of the electron energy is carried by the highest energy electrons. If the electron \\emph{particle} density $n(\\epsilon)$ per energy at an energy $\\epsilon$ is characterized by a power law $n(\\epsilon) \\sim \\epsilon^{-p}$, the spectral index is $p \\le 2 \\gamma - 1 \\approx 1.6$, in a random magnetic field in which the electrons are isotropically distributed. The energy per logarithmic energy interval, which is proportional to $\\epsilon^2 \\times n(\\epsilon)$, is therefore rising with energy. We note that such a spectrum is also inferred for the radio emitting electrons in the Crab nebula, and more generally in pulsar wind nebulas \\citep{gaensel_2006,sironi_2011}. While the radio emission is produced  by a different electron population than the gamma-ray flares, efficient particle acceleration appears to be a common feature in these systems.  It is an interesting question how such a hard electron spectrum is produced.  Standard diffusive shock acceleration typically results in spectra with $p \\ge 2$ \\citep{gallant_1992,kirk_2000}. Even though it has been shown that harder spectra can be produced in certain field configurations with low-level turbulence \\citep{kirk_1989,summerlin_2011}, these conditions are not expected at termination shocks of pulsar winds. Additionally, shock acceleration appears to be inefficient at highly oblique shocks that are representative of the pulsar wind termination discontinuity \\citep{ellison_2004,summerlin_2011,sironi_2011b}. One alternative is that magnetic reconnection in the striped pulsar wind might accelerate particles \\citep{lyubarsky_2003,kirk_2004,yuan_2011,bednarek_2011}. However, simulations show that reconnection behind the pulsar termination probably does not provide the required electron energies to produce gamma-ray emission \\citep{sironi_2011}. Another interesting possibility is that acceleration is occurring directly in the electric field induced by the pulsar, as discussed by \\citet{abdo_crab_2011}.  The observation of a peak synchrotron energy of $\\approx$380 MeV is among the highest yet seen from astrophysical source today. The observation is surprising as particle acceleration in the presence of synchrotron cooling is expected to limit synchrotron emission to photon energies below $\\approx$150 MeV \\citep{guilbert_1983,de_jager_gamma-ray_1996,komissarov_2011}. Two solutions to this problem have been proposed recently in this context: \\begin{enumerate} \\item The electric field at the acceleration site is larger in magnitude than the magnetic field. This is generally an unstable state in plasma, as charges will short out the electric field; however, temporarily such a configuration is expected, e.g. in magnetic reconnection events. The gamma-ray emission might occur afterward, when the accelerated electrons enter a region of enhanced magnetic fields \\citep{uzdensky_2011,cerutti_2011}. More detailed studies are required to assess whether such a scenario is plausible in the nebula environment and can be sustained for the duration of the flares. \\item The gamma rays are emitted in a region of bulk relativistic motion, and are therefore Doppler boosted toward the observer \\citep{komissarov_2011}. For a flow moving directly toward us a Lorentz factor $\\gtrsim 2$ is sufficient to accommodate the observed peak energy. \\end{enumerate}  Variations in the Doppler boosting can naturally account for the observed flux variation \\citep{lyutikov_2011}.  The observed spectral evolution is compatible with such an interpretation: the energy flux of the emission varies approximately as a power of {$\\alpha$ = 3.42 $\\pm$ 0.86} with the cutoff energy; a correlation with {$\\alpha \\approx$ 3} is indeed expected for variations produced by changes in relativistic beaming \\citep{lind_1985}.  The flare brightness, the high frequency of the observed peak of the gamma-ray emission, and the spectral evolution during the flare all suggest the presence of relativistic beaming. We therefore conclude that, independent of the location of the emission region and the physical processes responsible for the flares, the emission region is moving relativistically toward us, and changes in its motion are likely the predominant mechanism responsible for the observed flux variations. Such a kinematic explanation does however not address the issue of how a moving source can be created dynamically and sustained radiatively in the face of strong losses. The Crab Nebula still has much more to teach us."
    },
    "1112/1112.3184_arXiv.txt": {
        "abstract": "We estimate distances to the spherical planetary nebula Abell\\,39 and the bipolar planetary nebula NGC\\,7027 by interpolating from a wide grid of photoionization models using the 3-D code, MOCASSIN. We find preliminary distances of 1.5\\,kpc and 0.9\\,kpc respectively, with uncertainties of about 30\\%. ",
        "introduction": "Accurate distances to planetary nebulae (PNe) are crucial in unraveling the connection between the physical properties of the nebulae with those of their central stars (CSs). Reliable distances facilitate the accurate estimation of  fundamental parameters such as the CS mass and luminosity, and the nebular mass and age.  We have begun a program of using photoionization modeling to refine the distances to a sample of nearby PNe (see also \\cite[Danehkar \\etal\\ 2011]{Danehkar_etal11}). In this work, we study two very different PNe as a proof of concept: Abell\\,39, a simple spherical shell with no microstructures, and NGC\\,7027, a well-known, young, luminous bipolar PN with a massive molecular envelope. Using the 3-D photoionization code (MOCASSIN; \\cite[Ercolano et al. 2003]{Ercolano_etal03}), our ultimate aim is to constrain the distance to individual PNe, utilizing the physical PN radius we have calculated and the angular size. ",
        "conclusions": ""
    },
    "1112/1112.3651_arXiv.txt": {
        "abstract": "We are building a new spectral library with the X-Shooter instrument on ESO's VLT: XSL, the X-Shooter Spectral Library.  We present our progress in building XSL, which covers the wavelength range from the near-UV to the near-IR with a resolution of $R\\sim10000$. At the time of writing we have collected spectra for nearly 240 stars.  An important feature of XSL is that we have already collected spectra of more than 100 Asymptotic Giant Branch stars in the Galaxy and the Magellanic Clouds. ",
        "introduction": "Stellar population models are powerful tools which are widely used to study galaxy evolution.  Using these models, one can determine galaxy ages, metallicities and abundances.  Spectral libraries are an integral component of stellar population models (e.g. \\cite{BC03}).  A spectral library gives the behavior of individual stellar spectra as function of effective temperature ($T _{\\mathrm{eff}}$), gravity ($\\log g$) and metallicity ([Fe/H]). A stellar population model integrates these spectra together with a set of stellar isochrones and an initial mass function to produce a model spectrum of an entire population.  In order to reproduce galaxy spectra as precisely as possible, one requires a comprehensive stellar spectral library that covers the entire desired parameter space of $T _{\\mathrm{eff}}$, $\\log g$, and [Fe/H].  Moreover, extended wavelength coverage is strongly desirable, because different stellar phases contribute their light in different bands. For instance, bright giants contribute more light in the near-infrared than the faint giant and subgiant stars, while in the optical, the situation is reversed \\cite{Frogel88}. These Asymptotic Giant Branch (AGB) stars dominate the light of intermediate-aged and old stellar populations in the near-infrared but are unimportant in the optical (e.g., \\cite{Maraston05,Conroy10}). Detecting their presence requires \\emph{broad wavelength coverage} in both the target and model spectra.  Stellar spectral libraries can be classified into empirical and theoretical libraries, depending on how the library is obtained. Both theoretical and empirical libraries have improved in recent years. The most widely used theoretical libraries in stellar population models are those of \\cite{Kurucz93}, \\cite{Munari05}, \\cite{Gustafsson08}, \\cite{Coelho05}, \\cite{Coelho07}, and \\cite{Martins05}. Theoretical libraries have the advantage of (nearly) unlimited resolution and selectable abundance patterns -- not only scaled-solar abundances but also non-solar patterns.  Unfortunately, theoretical libraries suffer from systematic uncertainties, as they rely on model atmospheres and require a reliable list of atomic and molecular line opacities \\cite{Coelho05}.  Empirical stellar libraries, on the other hand, have the advantage of being drawn from real, observed stars and therefore do not suffer this limitation; however they frequently have relatively low resolution (with a few exceptions; see below) and are unable to reproduce the indices measured in giant elliptical galaxies \\cite{Reynier89, Worthey92}, because they are based on local stars with typical Milky Way disk abundance patterns.  In this proceeding we present the X-Shooter Stellar Library, XSL.  XSL is being obtained using the new X-Shooter three-arm spectrograph on ESO's VLT \\cite{Vernet10}.  XSL has the unique advantage of simultaneously acquiring spectra covering the near-ultraviolet up to the near-infrared. Furthermore, the fact that X-Shooter is mounted on an 8.2-m telescope allows us to include faint objects in stars in the Galactic bulge and the Magellanic Clouds for the first time, along with stars in the Galactic disk and halo, at moderate spectral resolution ($R\\sim10000$).  \\subsection{Previous stellar libraries}  \\begin{table} \\caption{\\label{tablib}Some previous stellar libraries}  \\begin{center} \\lineup \\begin{tabular}{*{5}{l}} \\br Library&Resolution&Spectral range&Number &Reference\\cr &R=$\\lambda/\\Delta\\lambda$&(nm)&of stars&\\cr \\mr STELIB & 2000 & 320-930 & 249 &\\cite{LeBorgne03}\\cr ELODIE & 10000 & 390-680 & 1300 &\\cite{Prugniel01,Prugniel04,Prugniel07}\\cr INDO-US & 5000 & 346-946 & 1237&\\cite{Valdes04}\\cr MILES & 2000 & 352-750 & 985 &\\cite{miles06}\\cr IRTF-SpeX & 2000 & 800-2500& 210 &\\cite{Rayner09}\\cr NGSL & 1000 & 167-1025& 374 &\\cite{Gregg06}\\cr UVES-POP & 80000 & 307-1030& 300 &\\cite{Bagnuolo03}\\cr LW2000 & 1100 & 500-2500& 100 &\\cite{LW2000}\\cr \\br \\end{tabular} \\end{center} \\end{table}   We begin with a review of several previous empirical stellar libraries and their principal features, listed in Table~\\ref{tablib}.  In the optical, we have (among others) Lick/IDS \\cite{Worthey94}, MILES \\cite{miles06}, ELODIE \\cite{Prugniel01,Prugniel04,Prugniel07}, STELIB \\cite{LeBorgne03}, NGSL \\cite{Gregg06}, and the Pickles library \\cite{Pickles85}.  Libraries in the near-IR are a challenging task, but pioneering work has been done by Lan\\c{c}on and Wood \\cite{LW2000} (LW2000), Rayner et al.\\ \\cite{irtf09} (IRTF-SpeX), and M\\'{a}rmol-Queralt\\'{o} et al.\\ \\cite{Marmol-Queralto}. However, extended-wavelength-coverage spectral libraries at moderate resolution are still largely missing.  STELIB \\cite{LeBorgne03} has been used to construct the widely used stellar population synthesis models of Bruzual \\& Charlot 2003 (\\cite{BC03}, here after BC03).  This library has a wide wavelength range, 3200--9300 \\AA, at modest resolution ($R\\sim2000$).  The spectra of local galaxies have been reproduced by the models of BC03 through this library \\cite{Gallazzi05}. However, stars with low and high metallicites are sparse in this library, leading to problems in some regimes, like modeling old, metal-poor globular cluster spectra \\cite{Koleva08}.  The ELODIE library \\cite{Prugniel01,Prugniel04,Prugniel07} is a moderate-resolution spectral library with $R\\sim10000$. It has been applied in the P\\'{E}GASE-HR synthetic model \\cite{LeBorgne04}. The atmosphere parameter coverage of this library had been improved through its updated version, but the wavelength range is still limited in the optical region (4100--6800 \\AA).  The INDO-US library \\cite{Valdes04} includes a large number of stars (1273), covers a fair range in atmosphere parameters with moderate-resolution ($R\\sim5000$) . Atmospheric parameters of its stars are given in Wu et al. \\cite{Wuyue11}. The data have been used for automated spectral classification of high-resolution spectra over a wide wavelength range \\cite{Pickles07}.  The problem of this library is the lack of accurate spectrophotometry, making inclusion in stellar population models problematic.  The MILES library \\cite{miles06} is widely used in stellar population models. This library profits from its good physical parameter coverage, careful flux calibration, and a large number of stars (985), enabling stellar population models to predict metal-poor or metal-rich systems. The MILES library's wavelength range covers 3500--7500 \\AA\\ at modest resolution ($R\\sim2000$) \\cite{Falcon11}.  The Next General Spectral Library (NGSL\\footnote{See \\url{http://archive.stsci.edu/prepds/stisngsl/}}, \\cite{Gregg06}) is a low resolution ($R\\sim1000$) stellar library obtained with STIS on HST.  It contains $\\sim400$ stars having a wide range in metallicity and age. Since its spectra were observed in space, this library does not suffer from telluric absorption or seeing variations, and hence has excellent absolute spectrophotometry. Other features of NGSL are its broad wavelength coverage, from the space UV to 1 $\\mu$m, and high signal-to-noise.  There is also a high-resolution ($R\\sim40000$, covering 3600--11000 \\AA) extension of NGSL's southern stars taken with the UVES spectrograph of ESO's VLT (Hanuschik et al., in prep.).  The UVES Paranal Observatory Project (UVES-POP, \\cite{Bagnuolo03}) is a library of spectra of $\\sim300$ nearby bright stars taken with UVES. The spectra cover the optical region at $R\\sim80000$ and have a typical S/N ratio is 300--500 in the $V$ band. The library has been flux calibrated but has not been corrected for the severe telluric absorption in the red.  For stellar population models the sample is incomplete, since it only contains bright stars around solar metallicity.  In the near-infrared range, the IRTF Spectral Library (IRTF-SpeX, \\cite{Rayner09}) contains 200 stars observed with the cross-dispersed infrared SpeX spectrograph on IRTF at a resolving power of $R\\sim2000$. The largest spectral library of very cool supergiants and giants over the wavelength range $\\lambda=0.5$--2.5 $\\mu$m was published by Lan\\c{c}on \\& Wood (LW2000, \\cite{LW2000}), which contains $\\sim100$ stars taken at a resolution of $R=1100$, including observations over multiple phases. Unfortunately, cool giants are variable in optical and NIR, so simultaneous observations are required.  Libraries with more limited spectral coverage in the near-infrared, like the 2.3 $\\mu$m library of Marm\\'ol-Queralt\\'o et al.\\ \\cite{MQ08}, also exist, but are limited to a small number of spectral features (like the CO bandhead). ",
        "conclusions": "XSL, the X-Shooter Stellar Library, is intended to be the largest stellar library with complete wavelength coverage from 320--2480 nm and covering as large a range in stellar atmospheric parameters as possible. We will soon have a first version of XSL with spectra of $\\sim240$ stars.  The final library will contain $\\sim600$ stars at moderate resolution ($R\\sim10000$). The stellar parameter coverage will be enforced through carefully selected samples, so that not only stars with solar metallicities and abundances are included but also metal-poor or metal-rich stars from the Bulge and Magellanic Clouds. With X-Shooter's unique capability, the variable stars will be observed consistently, which will yield more reliable stellar population models.  \\subsection*"
    },
    "1112/1112.1712_arXiv.txt": {
        "abstract": "Whether analytic exact vacuum(electrovacuum) solutions of the Einstein(Einstein-Maxwell) field equations can accurately describe or not the exterior spacetime of compact stars remains still an interesting open question in Relativistic Astrophysics. As an attempt to establish their level of accuracy, the radii of the Innermost Stable Circular Orbits (ISCOs) of test particles given by analytic exterior spacetime geometries have been compared with the ones given by numerical solutions for neutron stars (NSs) obeying a realistic equation of state (EoS). It has been so shown that the six-parametric solution of Pach\\'on, Rueda, and Sanabria (2006) (hereafter PRS) is more accurate to describe the NS ISCO radii than other analytic models. We propose here an additional test of accuracy for analytic exterior geometries based on the comparison of orbital frequencies of neutral test particles. We compute the Keplerian, frame-dragging, as well as the precession and oscillation frequencies of the radial and vertical motions of neutral test particles for the Kerr and PRS geometries; then we compare them with the numerical values obtained by Morsink and Stella (1999) for realistic NSs. We identify the role of high-order multipole moments such as the mass quadrupole and current octupole in the determination of the orbital frequencies especially in the rapid rotation regime. The results of this work are relevant to cast a separatrix between black hole (BH) and NS signatures as well as probe the nuclear matter EoS and NS parameters from the Quasi-Periodic Oscillations (QPOs) observed in Low Mass X-Ray Binaries. ",
        "introduction": "One of the greatest challenges of the general theory of re\\-la\\-tivity has been the construction of solutions to the Einstein-Maxwell field equations representing the gravitational field of compact stars such as neutron stars (NSs). Stationary axially symmetric spacetimes satisfy basic properties one expects for rotating objects, namely time symmetry and reflection symmetry with respect to the rotation axis \\citep[see e.g.][]{2006CQGra..23.3251P}. The simplest stationary axially symmetric exact exterior vacuum solution describing a rotating configuration is the well-known Kerr metric \\citep{1963PhRvL..11..237K}. The Kerr metric is fully described by two free parameters: the mass $M$ and the angular momentum $J$ of the object. However, it is known from numerical models that the quadrupole moment of rotating NSs deviates considerably from the one given by the Kerr solution $Q_{\\rm Kerr}=-J^2/(M c^2)$ \\citep[see e.g.][for details]{1999ApJ...512..282L}.  In the mean time, a considerable number of analytic exterior solutions with a more complex multipolar structure than the one of the Kerr solution have been developed \\citep[see e.g.][]{1995JMP....36.3063M,2000PhRvD..62d4048M,exactsolbook}. Whether analytic exterior solutions are accurate or not to describe the gravitational field of compact stars is an interesting and very active topic of research \\citep[see e.g.][and references therein]{2002MNRAS.336..831S,2004MNRAS.350.1416B,2006PhRvD..73j4038P}.  The accuracy of analytic solutions to describe the exterior geometry of a realistic rotating compact star has been tested by comparing physical properties, e.g. the radius of the Innermost Stable Circular Orbit (ISCO) on the equatorial plane and the gravitational redshift \\citep[see][for details]{1998AstL...24..774S,2004MNRAS.350.1416B,2006PhRvD..73j4038P}. In order to do such a comparison, the free parameters (i.e. the lowest multipole moments) of the analytic exterior spacetime, are fixed to the corresponding lowest multipole moments given by numerical interior solutions of the Einstein equations, for NS realistic models \\citep[see e.g.][]{2004MNRAS.350.1416B}.  Following such a procedure, the solution of \\citet{2000PhRvD..62d4048M} has been compared by \\citet{2002MNRAS.336..831S} and by \\citet{2004MNRAS.350.1416B} with the numerical solutions for NSs calculated by \\citet{1994ApJ...424..823C} and with those derived by \\citet{2004MNRAS.350.1416B}, respectively. However, being a generalization of the solution of \\citet{1972PhRvL..29.1344T}, it cannot describe slowly rotating compact stars \\citep[see e.g.][]{2004MNRAS.350.1416B}, but the dynamics of astrophysical objects with anisotropic stresses \\citep[see][for details]{2007PhRvD..75b3008D}.  Following a similar procedure, based on tests of the ISCOs radii on the equatorial plane of the rotating neutron stars obtained by \\cite{2004MNRAS.350.1416B}, it has been shown that the six-parametric solution of \\cite{2006PhRvD..73j4038P} (hereafter PRS solution, see Sec.~\\ref{sec:2} for details) is more accurate than the model of \\cite{2000PhRvD..62d4048M}. In addition, being a generalization of the Kerr solution, this solution can be used for arbitrary rotation rates.  Besides the ISCOs radii, there are additional physical properties that can be computed with analytic and numerical models and thus useful to compare and contrast the accuracy of analytic exact models. The aim of this article is to analyze the properties of orbital frequencies of neutral test particles in the PRS and in the Kerr geometries with especial focus on the Keplerian $\\nu_{\\rm K}$, frame-dragging (Lense-Thirring) $\\nu_{\\rm LT}$, as well as the precession(oscillation) frequencies of the radial and vertical motions, $\\nu^{\\rm P}_{\\rho}$($\\nu^{\\rm OS}_{\\rho}$) and $\\nu^{\\rm P}_{z}$($\\nu^{\\rm OS}_{z}$), respectively.  The relevance of these frequencies relies on the fact that they are often invoked to explain the Quasi-Periodic Oscillations (QPOs) observed in some relativistic astrophysical systems such as Low Mass X-ray Binaries (LMXBs), binary systems harboring either a NS or a black hole (BH) accreting matter from a companion star. For instance, within the Relativistic Precession Model (RPM) introduced by \\cite{1998ApJ...492L..59S,1999ApJ...513..827M,1999ApJ...524L..63S,1999PhRvL..82...17S}, the kHz QPOs are interpreted as a direct manifestation of the modes of relativistic epicyclic motion of blobs arising at various radii $r$ in the inner parts of the accretion disk around the compact object (see Sec.~\\ref{sec:RPM}, for details).  In addition to the RPM, the Keplerian, precession and oscillation frequencies are used in other QPO theoretical models \\citep[see e.g.][for a recent comparison of the existing models]{2011ApJ...726...74L}. Due to the influence of general relativistic effects in the determination of such frequencies, an observational confirmation of any of the models might lead to an outstanding test of general relativity in the strong field regime. In this line, it is of interest to compare and contrast the orbital frequencies given by the Kerr solution and by the PRS solution (see Sec.~\\ref{sec:3}), which help to establish the differences between possible BH and NS signatures. We emphasize in this article the major role of the quadrupole moment as well as of the octupole moment of the object, whose possible measurement can be used as a tool to test the no-hair theorem of BHs \\citep[see e.g.][]{2011ApJ...726...11J} and to discriminate between the different theoretical models proposed to explain the physics at interior and exterior of the Neutron Stars. Additionally, in the case of NSs, the interpretation of QPOs as the manifestation of orbital motion frequencies might lead to crucial information of the NS parameters such as mass, angular momentum \\citep[see e.g.][]{1998ApJ...492L..59S,2010ApJ...714..748T}, and quadrupole moment \\citep[see e.g.][]{1999ApJ...513..827M}. These parameters reveal, at the same time, invaluable information about the EoS of nuclear matter.  The article is organized as follows. In Sec.~\\ref{sec:2} we recall the properties of the PRS solution. The computation of the orbital frequencies as well as the comparison of their features in the Kerr and in the PRS spacetimes, is shown in Sec.~\\ref{sec:3}. In Sec.~\\ref{sec:4} we study the accuracy of the analytic formulas of the periastron and nodal frequencies derived by \\cite{1995PhRvD..52.5707R} for stationary axially symmetric spacetimes. \\textcolor{black}{In Sections 5 and 6 we discuss the accuracy of the PRS solution in describing the frequencies of realistic NS models and its relevance in the Relativistic Precession Model, respectively.} The conclusions of this work and a discussion on possible additional effects to be accounted for in the determination of the orbital frequencies, e.g. the effect of magnetic dipole moment, are outlined in Sec.~\\ref{sec:6}.  \\section[]{The PRS analytic exact solution}\\label{sec:2} We first recall the PRS analytic model \\citep{2006PhRvD..73j4038P}, for the exterior gravitational field of a compact object\\footnote{Mathematica 8.0 scripts with the solution, some limiting cases as well as the the calculations presented in this paper are available at \\href{http://www.chem.utoronto.ca/~lpachon/scripts/nstars} {http://www.chem.utoronto.ca/$\\sim$lpachon/scripts/nstars}}. In the stationary axisymmetric case, the simplest form of the metric can be written as \\citep{1953AnP...447..309P} \\begin{equation} \\label{Papapetrou} ds^2=-f(dt-\\omega d\\phi)^2+f^{-1}\\left[e^{2\\gamma} (d\\rho^2+dz^2)+\\rho^2d\\phi^2 \\right]\\, , \\end{equation} where $f$, $\\omega$ and $\\gamma$ are functions of the quasi--cylindrical Weyl coordinates $(\\rho,z)$. Thus, the components of the metric tensor $g_{\\mu\\nu}$ are \\begin{align} g_{\\phi \\phi} &= \\frac{\\rho^2}{f(\\rho,z)} - f(\\rho,z) \\omega(\\rho,z)^2, \\\\ g_{tt} &= -f(\\rho,z), \\\\ g_{t\\phi} &= f(\\rho,z) \\omega(\\rho,z), \\\\ g_{zz} &= g_{\\rho \\rho} = \\frac{{\\rm e}^{2\\gamma(\\rho,z)}}{f(\\rho,z)} = \\frac{1}{g^{zz}} =  \\frac{1}{g^{\\rho\\rho}}. \\end{align}  Using the above line element, the Einstein-Maxwell equations can be reformulated, via Ernst's procedure in terms of two complex potentials ${\\cal E}(\\rho,z)$ and $\\Phi(\\rho,z)$ \\citep{1968PhRv..167.1175E,1968PhRv..172.1850E}. By means of Sibgatullin's integral method \\citep{1991owsg.book.....S,1993CQGra..10.1383M} this system of equations can be solved va \\begin{align} \\label{Ernst1} {\\cal E}(z,\\rho)&=\\int\\limits_{-1}^1\\frac{d\\sigma}{\\pi} \\frac{e(\\xi)\\mu(\\sigma)}{\\sqrt{1-\\sigma^2}}, \\\\ \\Phi(z,\\rho)&=\\int\\limits_{-1}^1 \\frac{d\\sigma}{\\pi} \\frac{f(\\xi)\\mu(\\sigma)}{\\sqrt{1-\\sigma^2}}, \\end{align} where $e(z):={\\cal E}(z,\\rho=0)$ and \\mbox{$f(z):=\\Phi(z,\\rho=0)$}. The unknown function $\\mu(\\sigma)$ must satisfy the singular integral equation \\begin{equation} \\dashint_{-1}^{1}\\frac{\\mu(\\sigma)[e(\\xi)+\\tilde e(\\eta)+2f(\\xi)\\tilde f(\\eta)]d\\sigma}{(\\sigma-\\tau)\\sqrt{1-\\sigma^2}}=0 \\end{equation} and the normalizing condition \\begin{equation} \\int_{-1}^1\\frac{\\mu(\\sigma)d\\sigma}{\\sqrt{1-\\sigma^2}}=\\pi, \\end{equation} where $\\xi=z+i\\rho\\sigma$, $\\eta=z+i\\rho\\tau$, $\\rho$ and $z$ being the Weyl-Papapetrou quasi--cylindrical coordinates, $\\sigma, \\tau\\in[-1,1]$, $\\tilde e(\\eta):=\\overline{e(\\bar\\eta)}$, $\\tilde f(\\eta):=\\overline{f(\\bar\\eta)}$ and the overbar stands for complex conjugation. In \\citep{2006PhRvD..73j4038P}, the Ernst potentials were chosen as \\begin{align} \\label{Potenciales eje} \\begin{split} e(z) = \\frac{z^3-z^2(m+\\textrm{i}a)-kz+\\textrm{i}s}{z^3+z^2(m-\\textrm{i}a)-kz+\\textrm{i}s}\\, , \\\\ f(z) = \\frac{q z^2 + \\textrm{i}\\mu z}{z^3+z^2(m-\\textrm{i}a)-kz+\\textrm{i}s}\\, . \\end{split} \\end{align}  We calculate the multipole moments following the procedure of \\cite{1990CQGra...7.1819H}. We denote the mass multipoles by $M_i$ while, the current (rotation) multipoles, by $S_i$. The electric multipoles are denoted by $Q_i$ and the magnetic ones by ${\\cal B}_i$. Thus, for the PRS solution we have \\begin{align} \\begin{split} \\label{multipolosP} M_0 &= m\\, ,\\quad M_2 = m k - m a^2\\, ,\\quad \\ldots \\\\ S_1 &= m a\\, ,\\quad S_3 = - m a^3 + 2 m a k  - m s\\, ,\\quad \\ldots \\\\ \\end{split} \\\\ \\begin{split} \\label{multipolosQ} Q_0 &= q\\, ,\\quad Q_2 = - a^2 q - a \\mu + k q\\, ,\\quad \\ldots \\\\ {\\cal B}_1 &= \\mu + a q\\, ,\\quad {\\cal B}_3 = - a^2 \\mu + \\mu k - a^3 q + 2 a k q - q s\\, ,\\quad \\ldots \\end{split} \\end{align} This allows us to identify $m$ as the total mass, $a$ as the total angular moment per unit mass ($a=J/m$, being $J$ the total angular moment); while $k$, $s$, $q$ and $\\mu$ are associated to the mass-quadrupole moment $M_2$, current octupole $S_3$, electric charge and magnetic dipole, respectively.  The potentials (\\ref{Potenciales eje}) can be written in an alternative way, we mean \\begin{equation} e(z)=1+\\displaystyle \\sum_{i=3}^{3} \\frac{e_{i}}{z-\\beta_{i}}\\, ,\\qquad f(z)=\\displaystyle \\sum_{i=3}^{3} \\frac{f_{i}}{z-\\beta_{i}}\\, , \\end{equation} with \\begin{align} e_{j}&= (-1)^{j}\\frac{2 m \\beta^{2}_{j}}{(\\beta_{j}-\\beta_{k})(\\beta_{j}-\\beta_{i})}\\, , \\\\ \\label{equ:f_i} f_{j}&= (-1)^{j+1}\\frac{i \\mu + d\\beta_{j} }{(\\beta_{j}-\\beta_{k})(\\beta_{j}-\\beta_{i})}\\, , \\quad i,k \\neq j\\, . \\end{align} Then, using Eqs.~(\\ref{Ernst1}) and (\\ref{Potenciales eje}), we obtain the Ernst potentials \\begin{equation}\\label{potenciales_ernst} {\\cal E}=\\frac{A + B }{A - B}\\, , \\qquad \\Phi=\\frac{C}{A - B}\\, , \\end{equation} and the metric functions in the whole spacetime \\begin{align}\\label{eq:metricfuncs} f&=\\frac{A \\bar{A}-B \\bar{B} + C \\bar{C}}{( A - B)(\\bar{A}-\\bar{B})}\\, ,\\quad e^{2\\gamma}=\\frac{A \\bar{A} -B \\bar{B} + C \\bar{C}}{\\displaystyle{K \\bar{K}\\prod_{n=1}^{6}r_n}}\\, ,\\\\ \\omega &= \\frac{{\\rm Im}[(A + B)\\bar{H}-(\\bar{A} + \\bar{B})G - C \\bar{I}]}{A \\bar{A} - B \\bar{B} + C \\bar{C}}\\, , \\end{align} where the functions $A$, $B$, $C$, $H$, $G$, $K$, and $I$ can be found in the Appendix \\ref{app:metricfuncs}.  The PRS electrovacuum exact solution belongs to the extended $N$-soliton solution of the Einstein-Maxwell equations derived by \\cite{1995PhRvD..51.4192R}, in the particular case $N=3$. In addition, the functional form of the metric functions resembles the one derived previously by \\cite{1999CQGra..16.3725B}. Besides the limiting cases discussed in \\cite{2006PhRvD..73j4038P} it is worth mentioning that, in the vacuum case $q=0$ and $\\mu=0$, for $s=0$ this solution reduces to the solution of \\cite{1995JMP....36.3063M} under the same physical conditions, namely $q=0$, $c=0$ and $b=0$ in \\cite{1995JMP....36.3063M}. ",
        "conclusions": "\\label{sec:6}  We have done an extensive comparison of the orbital motion of neutral test particles in the PRS and Kerr spacetime geometries. In particular we have emphasized on the Keplerian and frame-dragging frequencies, as well as the precession and oscillation frequencies of the radial and vertical motions.  \\textcolor{black}{We have evidentiated the differences in this respect between the Kerr and PRS solution, especially in the rapid $\\sim$kHz rotation regime. Such differences are the manifestation of the influence of the high order multipole moments such as the quadrupole and octupole.}  \\textcolor{black}{The analysis of the deviations between the Kerr and PRS features for given mass and angular momentum of a source studied in this work are useful to distinguish the signatures between BHs and NSs, which relevant to establish a separatrix for the identification of the compact objects harbored in X-Ray Binaries. In the case of BH candidates, these results might become important for testing the no-hair theorem of BHs \\citep[see e.g.][]{2011ApJ...726...11J}. Equally important, the application of the precession and oscillation frequencies to the explanation of QPOs in LMXBs possessing a NS, can unveil information on the NS parameters, leading to a possible identification of the behavior of the nuclear matter EoS at supranuclear densities. In this line, the identification of the rotation frequency of NSs in LMXBs from the pulsating X-ray flux $\\nu_{\\mathrm{burst}}$, e.g.~the case of 4U 1728--34 \\citep{1998ApJ...506L..39F}, 4U 1916--053 \\citep{2001ApJ...549L..85G} and more recently the case of IGR J17191--2821 \\citep{2010MNRAS.409.1136A}, will certainly help to constrain QPO models as well as the NS parameters.  Additional information coming from the photospheric radius expansion phenomena observed in these systems \\citep[see e.g.][for details]{2001ApJ...553L.157M} during their transient activity with Super-Eddington emission can become of paramount importance if combined with the QPO information.}  The generalization of the present work to the electrovacuum case is important to establish the influence of the magnetic dipole and quadrupole moments on the orbital motion of particles around compact objects \\citep[see e.g.][]{2010CQGra..27d5001B,2010PhRvD..82l4014S,2012CQGra..29f5012B}.  Interesting effects on the epicyclic frequencies due to the presence of the magnetic dipole have been already pointed out recently by \\cite{2010CQGra..27d5001B} and \\cite{2012CQGra..29f5012B}. These effects were predicted after neglecting the contribution of the electromagnetic field to the curvature, for $j=0$  see \\cite{2010CQGra..27d5001B} and for $j\\neq0$ \\cite{2012CQGra..29f5012B}. \\cite{2010CQGra..27d5001B} assumed the model of the star as a dipole magnetic field superposed on a Schwarzschild black hole. In the second work, they studied the case of a magnetized slowly rotating neutron stars; to build the model they  superpose an dipolar magnetic field on the Lense-Thirring geometry. The effects of the magnetic dipole on the location of the ISCO, within the PRS solution, has been investigated by \\cite{2010PhRvD..82l4014S}.  A complete analysis of the effects due to the emergence of electromagnetic structure on the orbital motion of charged particles is therefore of interest and deserve the appropriate attention. Recent observations have shown that for stars with strong magnetic fields the quadrupole and octupole magnetic terms make significant contributions to the magnetic field \\citep{2006MNRAS.370..629D}, which indicates that arbitrary higher order multipole components might be required in a realistic model. The presence of a magnetic quadrupole demands the breaking of the reflection symmetry \\citep[see] [for details]{2006CQGra..23.3251P}, by means of a slightly change to the Ernst electric potential over the symmetry axis \\begin{align} \\label{PotencialesPRSmod} \\begin{split} f(z) = \\frac{q z^2 + \\textrm{i}\\mu z + \\textrm{i} \\zeta}{z^3+z^2(m-\\textrm{i}a)-kz+\\textrm{i}s}\\, , \\end{split} \\end{align} a quadrupolar magnetic component $\\mathcal{B}_2 = \\zeta$ can be introduced to the PRS solution. Such a change generates just a redefinition of the coefficients $f_i$ in Eq.~(\\ref{equ:f_i}). In this way the PRS solution can be readily use to explore the effect of strong magnetic fields with non-dipolar structure."
    },
    "1112/1112.2711_arXiv.txt": {
        "abstract": "We present a new census of the Orion Nebula Cluster (ONC) over a large field of view ($\\gtrsim30^\\prime\\times30^\\prime$), significantly increasing the known population of stellar and substellar cluster members with precisely determined properties. We develop and exploit a technique to determine stellar effective temperatures from optical colors, nearly doubling the previously available number of objects with effective temperature determinations in this benchmark cluster. Our technique utilizes colors from deep photometry in the $I$-band and in two medium-band filters at $\\lambda\\sim753$ and $770$~nm, which accurately measure the depth of a molecular feature present in the spectra of cool stars. From these colors we can derive effective temperatures with a precision corresponding to better than one-half spectral subtype, and importantly this precision is independent of the extinction to the individual stars. Also, because this technique utilizes only photometry redward of 750~nm, the results are only mildly sensitive to optical veiling produced by accretion. Completing our census with previously available data, we place some 1750 sources in the Hertzsprung-Russel diagram and assign masses and ages down to 0.02 solar masses.  At faint luminosities, we detect a large population of background sources which is easily separated in our photometry from the bona fide cluster members. The resulting initial mass function of the cluster has good completeness well into the substellar mass range, and we find that it declines steeply with decreasing mass. This suggests a deficiency of newly formed brown dwarfs in the cluster compared to the Galactic disk population. ",
        "introduction": "\\label{section:introduction} Understanding the initial mass function (IMF) is one of the most important problems of stellar astrophysics. The IMF, together with the star formation history, dictates origin, evolution, and fate of the stellar populations, from individual clusters up to entire galaxies. The distribution of stellar masses has been studied in depth for more than half a century, starting with the pioneering work of \\citet{salpeter55}; the major open question is whether the IMF is universal or if it depends on the initial conditions of star formation \\citep[see, e.g.,][]{kroupascience,bastian2010}. Whereas it is commonly accepted that the IMF seems well reproduced by a power law for masses greater than several tenths of a solar mass, the shape and universality of the IMF in the substellar mass regime is still under investigation \\citep[e.g.,][]{Wang2011}.  Very young clusters (few Myr old) in star-forming regions provide a unique tool to investigate the IMF across the entire mass spectrum, for a number of reasons. First, they are usually young enough that neither dynamical processes nor stellar evolution have altered the mass distribution; therefore the measured distribution of stellar masses coincides with the IMF. Moreover, young low-mass stars and brown dwarfs (BDs) are in their brightest evolutionary stage when they contract towards the main sequence, therefore these objects are more easily detected and characterized when they are young.  Among the nearby Galactic star-forming regions, the Orion Nebula Cluster (ONC) is an ideal site for the study of star formation in particular for low-mass stars and BDs. This cluster counts a few thousand members, 1-3~Myr old \\citep{hillenbrand97,paperII}, spanning the entire mass spectrum ($M\\lesssim50$~M$_\\odot$). Several studies have been conducted in the past decade to measure the IMF in the ONC \\citep[e.g.,][]{hillenbrand97,hillenbrand-carpenter2000,muench2000,muench2002,lucasroche2000,lucas2005,slesnick04,paperII}; they generally find (similar to the field star population) a Salpeter-like slope above 1~M$_\\odot$, which flattens to a broad peak at 0.2-0.3~M$_\\odot$ (though the shape and position of the IMF peak is highly model dependent, \\citealt{paperII}), and the mass distribution likely decreases in the substellar mass range.  Determining masses in a young region such as the ONC is not without difficulties: besides the strong nebular emission, another major impediment is caused by differential reddening, which has strongly limited the ability to derive the stellar parameters of individual sources based on photometry alone. To overcome this shortcoming, spectroscopic surveys have been carried out \\citep{hillenbrand97,lucas2001,lucas2006,slesnick04,riddick2007,weights2009}, but they are either limited to a fraction of the members, or to the very central part of the region, the Trapezium cluster. Using near-infrared (NIR) photometry it is possible to asses stellar masses down to the planetary masses ($M<13$~M$_J$), for example by de-reddening the measured color-magnitude diagrams (CMDs) on one isochrone or even to deriving directly the IMF from the NIR luminosity functions (LFs) \\citep{hillenbrand-carpenter2000,muench2002}. This second approach, however, is not very accurate: the ONC shows a significant luminosity spread, sometimes interpreted as a real age spread \\citep{reggiani2011}, and stars of different age follow different mass-luminosity relation. Moreover the NIR excess originating in the inner circumstellar disks \\citep{meyer97} alters the observed fluxes.  Clearly, in order to improve our knowledge of the ONC IMF, a precise and systematic characterization of the stellar parameters of individual members is needed. To this purpose, the \\emph{Orion HST Treasury Program} \\citep{robberto2005treasury} has produced a high spatial resolution photometric survey of the ONC with three instruments onboard the Hubble Space Telescope, from $U-$band to the NIR, over a large field of view (FOV), with a sensitivity well into the BD mass range. To fully exploit this exceptional dataset, an accurate estimate of the effective temperature ($\\teff$) is needed to derive $A_V$ from the observed colors and therefore the stellar luminosities.  In our previous works \\citep[][hereafter Paper~I and Paper~II]{paperI,paperII} we have collected spectral types from the literature and complemented them with new optical spectroscopy. Moreover, we have presented a new observational technique, based on optical medium-band photometry at 6200\\AA---a wavelength where the spectra of cool stars show a deep, $\\teff$-dependent TiO absorption feature---to derive the spectral types of M-type sources. This allowed us to obtain the stellar parameters of $\\sim1000$ members and derive the most complete HRD of the ONC, down to the hydrogen burning limit.  In the present paper we extend our investigation to lower masses, well into the BD regime, and with higher completeness. To this purpose, we take advantage of the same observational strategy presented in Paper~I, using medium-band photometry to derive spectral types of cool sources, from which $\\teff$ and $A_V$ are derived, and place individual sources in the HRD. Here, instead of the $6200\\AA$ filter used in Paper~I, we select two bands at longer optical wavelengths, in order to increase our sensitivity for cooler sources such as BDs and further reduce the influence of extinction.  In Section \\ref{section:the_data} we present our new observations, the data reduction and calibration. In Section \\ref{section:analysis} we define two spectro-photometric indices based on our medium-band photometry. Using available stellar parameters for a small sample of ONC very-low mass stars and BDs, we define empirical transformations to convert these indices into $\\teff$ and $A_V$. We also show that this technique is not significantly affected by optical excesses from mass accretion, typical of young stars and BDs. In Section \\ref{section:HRD} we derive the new HRD for the ONC, which now includes $\\sim 1800$ sources, reaching masses as low as $0.02~M_\\odot$. We study the completeness as a function of $\\teff$ and bolometric luminosity ($L_{\\rm bol}$) accounting for photometric detection as well as differential reddening; we detect a population of candidate background contaminants, which appears well separated from the ONC members in the HRD, at lower luminosity. After excluding the contaminants, in Section \\ref{section:IMF} we derive and discuss the IMF of the ONC. ",
        "conclusions": ""
    },
    "1112/1112.0008_arXiv.txt": {
        "abstract": "We measure the real-space galaxy power spectrum on large scales at redshifts 0.5 to 1.2 using optical colour-selected samples from the CFHT Legacy Survey.  With the redshift distributions measured with a preliminary $\\sim$14000 spectroscopic redshifts from the VIMOS Public Extragalactic Redshift Survey (VIPERS), we deproject the angular distribution and directly estimate the three-dimensional power spectrum. We use a maximum likelihood estimator that is optimal for a Gaussian random field giving well-defined window functions and error estimates.  This measurement presents an initial look at the large-scale structure field probed by the VIPERS survey.  We measure the galaxy bias of the VIPERS-like sample to be $b_g=1.38 \\pm 0.05$ ($\\sigma_8=0.8$) on scales $k<0.2\\hmpc$ averaged over $0.5<{\\rm z}<1.2$.  We further investigate three photometric redshift slices, and marginalising over the bias factors while keeping other \\LCDM~ parameters fixed, we find the matter density $\\Omega_m=0.30\\pm0.06$. ",
        "introduction": "The shape of the galaxy clustering power spectrum encodes the dynamical history of the Universe under the influence of baryons, dark matter and dark energy.  On large scales the assumption of Gaussianity can be made, and the statistic summarises all of the cosmological information that is available.  Measurements of the power spectrum at $\\z\\sim0$ have led to fundamental tests of the \\LCDM~model \\citep{Tegmark04b,Efstathiou02}. The angular distribution of galaxies on the sky, although less sensitive than the full three-dimensional view, has played an important role as well.  Indeed, strong tests of the CDM model were made with the two-dimensional correlation function from the APM galaxy survey \\citep{Maddox90}.  Photometric surveys have the capability to probe significantly larger volumes at higher sampling rates than targeted spectroscopic surveys.  The advantages have become clear with the advancement of photometric redshift estimation methods.  The loss in three-dimensional precision can be compensated for with increased statistics leading to strong cosmological constraints that are comparable to the results from spectroscopic surveys.  Additionally, the projected density field is only weakly sensitive to redshift-space distortions, thus it provides a means to infer the real-space power spectrum directly.  The dependence on peculiar velocities becomes important for narrow redshift slices and can be turned into a useful measure of the growth rate \\citep{Ross11}. Measurements of the baryon acoustic feature and redshift-space distortions have now been made on photometric samples taken from the Sloan Digital Sky Survey \\citep{Padmanabhan07,Blake07,Thomas11}.   In this analysis, we present a new measurement of the real-space galaxy power spectrum using a photometric catalogue of galaxies at $0.5<\\rm{z}<1.2$ from the Canada-France-Hawaii Telescope Legacy Survey (CFHTLS) Wide survey.  The survey consists of four fields covering a total area of 133 \\sqrdeg.  The extent of the largest field, W1, is $\\sim10\\degr$ or $200~{\\rm h}^{-1}~{\\rm Mpc}$ at z=0.7, giving a maximum scale we may probe of $k_{min}\\sim0.05\\hmpc$. The dataset has been used for previous cosmological analyses, in particular for weak lensing \\citep{Fu08,Kilbinger09,Tereno09,Shan11} and galaxy correlation function measurements \\citep{Coupon11}.  A key ingredient needed to interpret the projected density field and constrain the three-dimensional power spectrum is the redshift distribution of the galaxy sample.  For this we use spectroscopy from the VIMOS Public Extragalactic Survey\\footnote{VIPERS website: \\href{http://vipers.inaf.it}{vipers.inaf.it}} (VIPERS)\\citep{VIPERS}.  VIPERS is an ongoing spectroscopic program to target $10^5$ galaxies in the redshift range 0.5-1.2 in a total area of 24\\sqrdeg~in the CFHTLS W1 and W4 fields. The accuracy of the spectroscopic measurements from VIPERS provide an unbiased estimate of the redshift distribution.  With this knowledge, we are confident that we can deproject the angular clustering signal and constrain the three dimensional power spectrum.  The primary advantage of studying the deprojected power spectrum $P_k$ is its closeness to theory.  The shape of the angular power spectrum is complicated by its dependence on survey properties, its depth and geometry.  Furthermore, in projection, scales are mixed.  Ideally we would like to separate the power on large scales in the linear regime from power on small scales that is influenced by complex astrophysical processes.   \\begin{figure} \\includegraphics[scale=1]{plots/fig1} \\caption{Galaxy count maps in the CFHTLS-Wide fields with the VIPERS-like colour selection.  We use {\\sc Healpix} cells with size 7\\arcmin.  Gaps in the survey coverage are left as blank pixels. The grid overlay has spacing of 1\\degr.\\label{fig:fields}} \\end{figure}    How to derive the three-dimensional power spectrum from the two-dimensional density field is a problem of deconvolution.  A good inversion method should be stable against noise in the data. Preliminary work done by \\citet{Baugh93,Baugh94} and \\citet{Gaztanaga98} used the Lucy deconvolution method that is known to be robust.  To further derive cosmological constraints, we must be able to estimate the covariance of the deprojection.  Methods of propagating the error from the angular correlation function to the three-dimensional power spectrum were developed by \\citet{Dodelson00} who perform the inversion with a prior on the smoothness of the power spectrum and compute a covariance matrix of the estimate.  Further work by \\citet{Eisenstein01}, \\citet{Dodelson02} and \\citet{Maller05} made use of the singular-value decomposition technique to remove modes that destabilise the inversion.  Importantly, the deprojection method should produce well-defined window functions that describe the mode mixing.  The aim is to separate the small and large scales that are mixed in projection, and the residual leakage should be understood.  A similar problem was solved with the maximum likelihood methods developed for the cosmic microwave background angular power spectrum \\citep{Tegmark97,Bond98} and then later applied to galaxy surveys \\citep{Huterer01,Tegmark02}. Applications to the deprojection of the power spectrum were presented by \\citet{Efstathiou01} and \\citet{Szalay03}.  In this work, we adopt the maximum-likelihood technique to construct an estimator for the power spectrum. The result is optimal under the assumption that the density is represented by a Gaussian random field. This is a reasonable assumption for the galaxy distribution on large scales.  Moreover, the estimator also simultaneously gives the covariance of the estimate as well as the window functions.  For small surveys, where the window functions must be handled carefully, the approach is especially useful.  We pay close attention to the window functions for the results presented here.  Maximum likelihood estimates are computationally expensive. However, because of the relatively small field sizes we consider, we can perform all computations on a consumer level four-core desktop computer.  In this article, we first introduce the CFHTLS and VIPERS datasets used.  In Sections 3 and 4 we review the angular power spectrum formalism and the maximum likelihood deprojection using a quadratic estimator.  We then apply the method to Gaussian simulations and investigate potential biases due to uncertainty in the redshift distribution and the fiducial cosmology.  Lastly, we measure the power spectrum with CFHTLS data and constrain the linear galaxy bias and matter density.  We report magnitudes using the AB magnitude convention in the CFHT $u^*g'r'i'z'$ photometric system.  We assume a flat \\LCDM~cosmology with $H_0=70.4~{\\rm km~s^{-1}~Mpc^{-1}}$, $\\Omega_m=0.272$, $\\Omega_b=0.0456$, $n_s=0.963$ and $\\sigma_8=0.8$ \\citep{Larson11}. ",
        "conclusions": "The CFHTLS-Wide fields probe a significant cosmological volume at redshifts not reached by other galaxy surveys to date.  We use the projected density field from photometric redshift samples to constrain the real-space power spectrum and derive constraints on the matter density and linear galaxy bias factors.  These results are made possible by precise knowledge of the redshift distributions provided by preliminary results from the VIPERS survey.  The primary advantage of computing the power spectrum directly from the angular distribution, instead of using conventional spherical harmonics $\\Cl$ is that we may construct window functions in Fourier space.  By optimising this, we achieve sharper constraints on the power spectrum than when we are limited to $\\ell$ bands.  This approach comes with the cost that we must adopt a fiducial power spectrum.  We showed that using the wrong fiducial power spectrum, although leading to sub-optimal weights, does not significantly bias the estimate.  This is true even on small scales, and we can effectively use this method to deconvolve small and large scales in Limber's equation. Residual systematic error on the derived power spectrum is at the 1\\% level, well below the sensitivity of the measurement.  The deprojection does strongly depend on the assumed redshift distribution of the galaxy sample as well as the cosmology used to compute the redshift-distance relation.  The cosmology dependence of the measurement makes the interpretation difficult, but to a first approximation, only the amplitude is affected; the shape of the power spectrum is recovered correctly.  Thus, a converging iterative procedure can be implemented by updating the fiducial model and repeating the analysis.  There is a degeneracy between a shift in the assumed redshift distribution and the cosmological model.  This is unavoidable when studying a field in projection.  However, the constraints on the redshift distribution can always be improved with further observation. In our analysis, from the sampling biases present in the VIPERS spectroscopy, we estimate the uncertainty in the mean redshift to be at the 1\\% level.  Thus, we do not expect a strong systematic error in the derived galaxy bias parameters.  We do note that the observed trend of low $\\chi^2$ values for the best-fit models in the higher redshift samples can arise if the covariance is overestimated.  This could be a weak hint that the true mean redshift is lower than what we assume or that a modification is needed in the fiducial cosmology.  Recently, the galaxy bias was measured from the CFHTLS-Wide fields in the context of the halo model by \\citet{Coupon11}.  Our final two photometric redshift bins, S7 and S8, correspond with samples constructed by Coupon \\etal so we are able to compare the resulting bias values.  Coupon \\etal constructed volume-limited samples using luminosity cuts resulting in a selection of brighter galaxies, thus we may expect their bias values to be larger.  The halo model constraints of \\citet{Coupon11} give for S7, $b_g=1.44\\pm0.01$, and for S8, $b_g=1.79\\pm0.03$.  These values have been scaled by 1.03 to transform from a cosmology with $\\Omega_m=0.25$ to $\\Omega_m=0.272$ which is assumed here.  Our value of $b_g$ for the S7 sample is higher, while for the S8 sample it is lower, although both are in agreement with Coupon \\etal within the 2$\\sigma$ confidence limit.  The measurements are based on different physical scales (Coupon \\etal restrict the correlation function to angular scales $<1.5\\degr$) and different model assumptions have been used.  Thus it is reasonable to consider the measurements as independent estimates.  Our results provide a preliminary look at the large-scale structure field probed by the VIPERS colour selection and demonstrate the strengths of the VIPERS sample for clustering studies at $\\z>0.5$.  We anticipate promising results with the full VIPERS spectroscopic sample."
    },
    "1112/1112.2657_arXiv.txt": {
        "abstract": "We perform a pixel-by-pixel analysis of 467 galaxies in the GOODS-VIMOS survey to study systematic effects in extracting substructure and properties of stellar populations (age, dust, metallicity and star formation history) from the pixel colors using the pixel-$z$ method. Several systematics are examined in this paper, including the effect of the input stellar population synthesis models whose SEDs are fitted to each pixel's colors, the effect of passband limitations and differences between the individual SED fits to pixels and global SED fitting to a galaxy's colors. We find that with optical-only colors ($bviz$), the systematic uncertainties due to differences among stellar population synthesis codes are well constrained. The largest impact on the stellar population age and SFR e-folding time estimates in the pixels arises from differences between the Maraston~(2005) models on one hand and the Bruzual \\& Charlot~(2003) and Charlot \\&Bruzual~(2007) models on the other, when optical-only ($bviz$) colors are used. This results in systematic differences larger than the $2\\sigma$ uncertainties in over 10 percent of all pixels in the galaxy sample. The effect of varying the number and choice of available passbands is more severe. In 26 percent of the pixels in the full sample, these limitations result in systematic biases in the age determination which are larger than the $2\\sigma$ uncertainties in the measurements. Robust results can, however, still be obtained with a minimum of 3 optical filters provided they span the 4000\\,\\AA \\ break. Near-IR data is also added to a subsample of 46 galaxies from the GOODS-NICMOS survey and systematics arising from model differences are again investigated. Differences among the models in their predicted rest-frame red/NIR colors manifest themselves as follows. For $z > 1$ galaxies the observed optical/NIR colors span the rest frame UV-optical SED, and the use of different population synthesis models does not significantly bias the estimates of the stellar population parameters compared to using optical-only colors. However, for $z < 1$, where the rest-frame NIR is still probed, there is a larger discrepancy between models when using optical-only colors compared to optical/NIR colors. This affects in particular the age determination in the pixels. With this characterization of the systematic errors, we illustrate how pixel-$z$ can be applied robustly to make detailed studies of substructure in high redshift objects such as (a) radial gradients of properties such as age, SFR and dust and (b) the distribution of these properties within subcomponents such as spiral arms and clumps. Finally, we show preliminary results from applying pixel-$z$ to galaxies in the CANDELS survey, illustrating how the new HST/WFC3 data can be exploited to probe substructure and stellar populations in $z\\sim1-3$ galaxies. ",
        "introduction": "}   \\subsection{Clumpy Disks: Resolved Stellar Populations in $z\\sim 1-3$ Objects as a Probe of High-z Disk Galaxies}  Many galaxies at $z\\sim1-3$ have morphologies that are very different from the galaxy population at $z\\sim0$ (Lotz~et~al.~2004; Lotz~et~al.~2006; Law~et~al.~2007; Scarlata~et~al.~2007; Elmegreen~et~al.~2007; Pannella~et~al.~2009). However, the morphologies of these objects only tell one part of the story. At $z>1$, rest-frame UV wavelengths, which are most impacted by dust obscuration and star formation, can be traced by high resolution optical imaging (Toft~et~al.~2007; Overzier~et~al.~2010; Cameron~et~al.~2010). The location and distribution of stellar populations within these galaxies, particularly sites of active star formation, is likely to be affected by various physical and dynamical mechanisms which are present in earlier epochs in the galaxy's evolution. The pixel-$z$ technique (Conti~et~al.~2003; Welikala~et~al.~2008; Welikala~et~al.~2009) can measure spatially resolved stellar populations for large samples of $z\\sim1-3$ galaxies using their multi-band HST images. With this approach, one can study asymmetries that exist in the star formation, age and dust distribution across these galaxies. In parallel with the detailed kinematical studies of these objects from spatially resolved spectroscopy, this will enable insights into the formation mechanisms of these galaxies.  One such mechanism that could be responsible for the formation of these clumps on kiloparsec scales is disk fragmentation in Toomre-unstable gas-rich disks (Elmegreen~et~al.~2007, Elmegreen~et~al.~2009). There is now direct kinematic evidence that high-z disks are much more turbulent than their counterparts at $z \\sim 0$ (Bournaud~et~al.~2007; Bournaud~et~al.~2008; Genzel~et~al.~2008; Genzel~et~al.~2011). For example, high-redshift disks have been found to have local intrinsic gas velocity dispersions of $20-90 \\,$km\\,s$^{-1}$ as well as high gas-to-total baryonic mass fractions (Daddi~et~al.~2010; Tacconi~et~al.~2010). Numerical simulations of gas-rich turbulent disks indicate that massive kpc-sized clumps can form in-situ through gravitational instabilities, a phenomenon known as a `clump-cluster' phase. (Noguchi~et~al.~1999; Immeli~et~al.~2004a,b; Bournaud et al.~2007; Elmegreen et al.~2008; Dekel, Sari, \\& Ceverino~2009; Agertz~et~al.~2009; Ceverino~et~al.~2010; Aumer et~al.~2010; Genzel~et~al.~2010). According to these simulations, the clumps can migrate towards the gravitational center as a result of both their mutual interactions and dynamical friction against the host disk, so that they can eventually coalesce into a young bulge on the order of a few dynamical times.  Alternatively, these high-redshift clumpy disk and irregular galaxies may form via mergers. In particular, numerical simulations of gas-rich mergers result in remnant disks that show a low velocity field asymmetry that also satisfies the criteria necessary to be classified as a high-redshift disk galaxy observationally. For example, Robertson \\& Bullock (2008) compared one of the merger remnants to the bulk properties of the well-studied $z=2.38$ galaxy BzK-15504 and showed that it has a star formation rate (SFR), gas surface density, and a circular velocity-to-velocity dispersion ratio that is in excellent agreement with BzK 15504. Such numerical simulations suggest that gas-rich mergers can play a prominent role in the formation of disk galaxies at high redshift.   Recent and upcoming larger IFU programs such as the SINS survey (F\\\"{o}rster-Schreiber et al.~2006, Genzel et al.~2006), the SAURON project (Bacon et al.~2001) and the MASSIV survey (Contini et al.~2011; Epinat et al.~2011; Queyrel et al.~2011) are expected to shed further light on these formation mechanisms by performing detailed studies of the gas kinematics of $z \\sim 2$ galaxies. Other large IFU programs include the CALIFA project (Sanchez et al.~2011), which will target approximately 600 objects with PMAS IFUs on the Calar Alto Telescope, and the SAMI program (Croom et al.~2011) which is a planned low redshift survey of serveral thousand galaxies. In addition to kinematic studies, a complementary study of the spatial distribution of the stellar populations within these galaxies can provide additional insights into whether these objects are formed through mergers or are, in fact, disks formed in-situ. F\\\"{o}rster-Schreiber et al.~(2011a, b) used a combination of near-infrared integral field spectroscopy from the SINFONI instrument on the VLT, combined with deep high resolution HST NIC2/F160W imaging of six $z \\sim 2$ star-forming galaxies to characterize the properties of kpc-scale clumps and their contribution  to the rest-frame optical emission. In our next papers, we aim to extend these previous studies and make detailed spatially-resolved measurements of the properties of large samples of disk galaxies. These studies will also make use of new fully reduced NIR data from the HST Wide Field Camera 3 (WFC3) which will become available for the GOODS-South field in early 2012.   \\subsection{Spatially Resolved Star Formation and The Role of Environment in Galaxy Evolution}   The relation between galaxy environment and galaxy properties have been extensively studied in the local Universe (Gomez et al.~2003, Blanton et al.~2005, Hogg et al.~2004). A number of studies have observed that the high-$z$ ($z\\gtrsim 1$) SFR-density relation is either reversed or weaker at $z \\sim 1$ than that seen locally (Cucciati et al.~2006; Cooper et al.~2007; Elbaz et al.~2007; Ilbert et al.~2006; Poggianti et al.~2006; Ideue et al.~2009; Salimbeni et al.~2009; Scodeggio et al.~2009; Tran et al.~2010; Gr\\\"{u}tzbauch et al.~2011). These examples are consistent with the scenario that galaxies in dense environments form stars rapidly at early times, quickly building up mass and becoming quiescent, while galaxies in less dense environments form stars at a more sedate pace but over longer timescales, a phenomenon we refer to as {\\em in-situ\\/} evolution (Wijeshinge et al.~2011). This is distinct from the \u2018\u2018infall and quench\u2019\u2019 scenario where the properties of galaxies are impacted directly by their local environment through physical mechanisms such as ram-pressure stripping and galaxy harassment. In addition, Peng at al.~(2010) makes a distinction between this type of quenching which is directly related to the environment of galaxies and {\\em mass quenching} which dominates at high masses and early cosmic times.  A few recent studies have, however,  continued to observe the same SFR-density relation upto $z\\sim1$ as seen locally (Patel et al.~2011). A few studies have investigated the relation beyond $z\\sim1$. Quadri et al.~(2011) showed that galaxies with quenched SF tend to reside in dense environments out to at least $z \\sim 1.8$ and that the the SFR-density relation holds even at fixed stellar mass. Their density estimates, however, have relatively high uncertainties as they are derived from photometric redshifts. While the SFR-density relation at high redshifts is a matter of some debate, it is clear that an {\\em in-situ\\/} evolution makes a distinct prediction compared to the \u2018\u2018infall and quench\u2019\u2019 models. The latter suggests that galaxies in dense environments should show a SFR distribution that is progressively suppressed from the outside in, as the outer regions are those which will be affected first by their rapidly changing environment. The former, on the contrary, suggests that the suppression should either happen uniformly as a galaxy ages, or that the inner regions should be suppressed first. By studying the spatial distribution of SFR in star-forming galaxies as a function of environment, we should be able to distinguish clearly between these two scenarios.  Studies by Welikala et al.~(2008; 2009)  at $z \\sim 0.1$ confirm the global SFR-density relation observed locally but also show that the SFR at $z \\sim$ 0.1 is centrally concentrated relative to the outskirts, and in more dense environments it is this central SFR that is supressed, favouring the  {\\em in-situ\\/} evolution scenario. Park et al.~(2007)  also studied the color gradients of galaxies as a function of the local galaxy density and galaxy morphology in the SDSS. They found no environmental dependence of the color gradient for early-type galaxies at a given luminosity, and only a weak dependence on environment for faint late-type galaxies which are observed to become bluer in their outskirts (relative to the galaxy center) in low density environments. It thus becomes extremely interesting to determine whether the radial variation of star formation with environment seen locally is also in place at high redshifts. The development of MOS IFU instruments and surveys with large photometric and spectroscopic samples, such as the Galaxy And Mass Assembly (GAMA) survey (Driver et al.~2011), will allow such a statistical analysis for a large sample of galaxies. There are two complementary approaches to this. While spatially resolved spectroscopy provides detailed measures of the current spatial distribution of physical properties within galaxies, pixel-$z$ adds the ability for the star formation history to be probed. While pixel-$z$ also provides the current SFR and other properties of the stellar populations in subcomponents of galaxies, these measurements will have higher uncertainties than spectroscopic measures, and cannot provide detailed kinematic information that spectroscopic observations can. However, until the advent of highly-multiplexed MOS IFUs, even IFU surveys will be limited to several thousands of targets, while pixel-$z$ can be applied to tens of thousands and up to millions of objects in order to determine the radial variation of stellar population parameters within them and the evolution of this radial variation with redshift. We therefore aim to extend studies performed by Welikala et al.~ (2008; 2009) which applied the pixel-$z$ method to $\\sim300,000$ galaxies at $z\\sim0.1$ in the SDSS and determined the radial variation of star formation for this sample, the scatter in this relation, and its dependence on galaxy environment and galaxy morphology. There are, however, inherent systematic effects that may limit the utlity of pixel-$z$ in probing spatially resolved stellar populations in galaxies. The current analysis explores these in order to quantify the impact of such systematics on the measurements made by pixel-$z$.    \\subsection{Stellar Populations Extracted from Pixel Colors: Assumptions and Limitations}  A number of studies have sought to extract stellar populations for spatially resolved colors. These compare the colors in individual pixels to a library of SEDs generated by stellar population synthesis models. The majority of these utilised small samples of galaxies at $z \\sim 0$ and $z \\sim 1$ (Abraham et al.~1999; Johnston et al.~2005; Kassin et al.~2003; Lanyon-Foster et al.~2007; Zibetti et al.~2011). Given a multi-band image of a galaxy, the pixel-$z$ method (Conti et al.~2003; Welikala et al.~2008, 2009) computes the stellar population properties such as stellar population age, SFR e-folding timescale $\\tau$, dust obscuration through $E(B-V)$ and metallicity $Z$ in each pixel in the galaxy, as well as the associated uncertainties in these parameters. It does this by taking a library of SEDs generated from stellar population synthesis (SPS) codes, redshifting each of these SEDs and convolving them with the available passbands. The simulated fluxes are then compared to the observed fluxes in each pixel and a best-fitting SED is computed for that pixel. The statistical errors in each parameter are calculated by marginalizing the likelihood function from the fit over the remaining parameters, and are conservative.  The pixel-$z$ approach makes two fundamental assumptions: \\begin{itemize} \\item{\\textbf{Independently evolving pixels: } The pixels in the galaxy image are assumed to be evolving independently of each other. This ignores any mixing of stellar populations between neighboring pixels which is expected to occur in real galaxies. Neighbouring pixels will also not truly be independent, due to a finite Point Spread Function (PSF), and this effect is also neglected.} \\item{\\textbf{Single stellar populations with exponentially declinining SFH: } The stellar population underlying each pixel is assumed to be one whose spectrum is derived from integrating over single stellar populations (SSPs) of a given age, weighted by a SFR which is characterized by an initial burst followed by an exponential decrease in SFR. The decrease is parameterized by a range of allowed values of e-folding times $\\tau$. In reality, galaxies are expected to have more complicated SFHs, including multiple bursts.} \\end{itemize}  In addition to these two fundamental assumptions, there is a limitation in the form of degeneracies among the SED models which could impact the accuracy with which pixel-$z$ determines stellar population parameters in the pixels. SPS models show that age, metallicity and dust all tend to affect spectra in similar ways (Bruzual \\& Charlot~2003). Multiple SEDs each generated for different populations (different ages, metallicites and SFHs) can, as a result, still predict very similar colors of a galaxy or its subcomponents. Some of these degeneracies arise from well known physical correlations, such as those between age and metallicity of the stellar population. Worthey~(1994) constructed detailed models of old and intermediate stellar populations and their absorption indices to show that if $\\delta_{age}$/$\\delta{Z} \\approx 3/2$ for two populations, they would appear identicial in most absorption indices. The age-metallicity degeneracy can also affect the photometric evolution of the SSPs since optical and near-infrared (NIR) colors reflect the relative contribution of hot and cool stars to the integrated light. Bruzual \\& Charlot~(2003) followed the evolution of optical and NIR colors and the stellar mass-to-light ratio $M/L$ in their model, and found that at fixed age, increasing metallicity tends to redden the colors and increase the $M/L$ ratio. This is explained by the fact that increasing the metallicity (at fixed stellar mass) causes (a) stars to evolve at lower effective temperatures and lower luminosities (Girardi et al.~2000) and (b) changes the relative number of red and blue supergiants, which strongly impacts the color evolution of the SSP. They also found that increasing metallicity at fixed age had a similar effect to increasing age at fixed metallicity. The age-metallicity degeneracy is therefore inherent in the SSP model used. There could also be other degeneracies, including those between dust obscuration and metallicity as well as between age and dust obscuration. The redshift-evolution of the obscuration-metallicity degeneracy in the pixels was characterized by Conti et al.~2003.  This work does not test the two fundamental assumptions above. As for degeneracies, exploring the different correlations between the parameters on a pixel-to-pixel basis and propagating the resulting errors are beyond the scope of this paper. Welikala et al.~2008 characterized some of these correlations using the likelihood function from the fit to the pixels in SDSS galaxies. Since the severity of the degeneracies depend on the number and type of available broadband colors, the SPS model and the chosen grid of stellar population parameters, we focus on these systematics in this paper. The validity of both the two assumptions and the impact of degeneracies will be explored in a future paper that compares the spatial distribution of stellar populations inferred from spatially resolved spectroscopy obtained from large IFU programs (Croom et al.~2011)  with those inferred from applying the pixel-$z$ technique to the same galaxies. Such a detailed comparison with IFU programs will enable a robust calibration and testing of pixel-$z$, leading eventually to a possible refinement of these assumptions. In this paper, we focus on the following systematics that could potentially bias our results:  \\begin{itemize}  \\item{\\textbf{The SPS model: } An inevitable bias is the choice of the stellar population synthesis model.  Different  stellar population synthesis models predict different colors. In particular, there are significant differences in the prescriptions for Thermally Pulsating Asymptotic Giant (TP-AGB) stars among the following models: Maraston 2005 (M05), Bruzual \\& Charlot 2003 (BC03) and Charlot \\& Bruzual 2007 (CB07). In this paper, we compare the outcomes of using these different models on a pixel-by-pixel basis, for all galaxies in the sample. Similar comparisons of population synthesis models have been performed previously on the integrated colors of galaxies but these have focused primarily on the effect on their stellar mass (Maraston et al.~2006). In this work, we focus on the effect of these differences on the age, SFR e-folding time and obscuration estimated in the pixels. } \\item{\\textbf{The available passbands: }The accuracy with which the pixel-$z$ parameters are determined in the pixels depends on the the available passbands.  This paper explores the systematic effects arising from passband limitations, using both optical and near-IR colors. } \\item{\\textbf{The parameter space searched by pixel-$z$ and the number of SEDs per model: } Pixel-$z$ searches a grid of ages, e-folding times, dust attenuation values and metallicities. In general, a large as possible range of values is chosen while physically unmeaningful values will be associated with high statistical uncertainties. Despite this, a large parameter space can result in degeneracies among the various SEDs searched by pixel-z which in turn may bias the pixel-$z$ estimates. Some priors are already built in to stop pixel-$z$ searching unphysical parameters e.g., the age of a stellar population in a pixel cannot exceed the age of the Universe at a given redshift. Here, we explore the systematic error arising from increasing the range of parameter values allowed on the pixel-$z$ grid. In particular, we focus on the impact of increasing the metallicity range. This has two predicted effects: (a) it increases the severity of the age-metallicity degeneracy and would bias the pixel-$z$ estimates in quantities such as age, $\\tau$ and dust and (b) it introduces an uncertainty associated with the models: while spectral synthesis models have been well established for solar metallicity stellar populations in optical photometry, the models are less well established for very sub-solar or super-solar populations. Increasing the metallicity range will test the impact of both these effects on the derived pixel-$z$ parameters. } \\item{\\textbf{Global SED-fitting versus pixel-$z$: } We examine any biases resulting from measuring stellar population quantities, particularly the SFR, in individual pixels compared to the same quantities derived from aperture photometry and SED-fitting. These biases arise because individual pixels can have different colors to the integrated color of the galaxy. This bias could give rise to a difference between (1) the mean properties of the galaxy obtained by integrating the fluxes in an aperture and SED-fitting to the total fluxes and (2) the mean properties derived from fitting to the individual pixels within the same aperture and summing the resulting fits.}  \\end{itemize} ",
        "conclusions": "\\label{sec:results}   \\subsection{Effect of SSP Model Differences using Optical-only Bands} \\label{subsec:model_tests}  Here, we illustrate the effect of using different stellar population synthesis models on the values of the parameters that pixel-$z$ determines using the \\textit{bviz} passbands. This is first shown qualititatively by displaying the resulting maps, and second, quantitatively, by measuring the systematic error in each parameter and following its distribution.  In the top panel of Figure~\\ref{fig:models_optical_age}, we illustrate the effect of the chosen stellar populations model on the age maps. There is a clear correlation between the pixel-$z$ maps of this galaxy and the $i$ band image. The qualitative differences from pixel-to-pixel between the various models(M05,BC03,CB07) are small but there is a difference along the right spiral arm between BC03 and CB07 on the one hand and M05 on the other. BC03 and CB07 both predict that the majority of pixels along the arm have ages $\\sim0.1$Gyr, while M05 shows a larger proportion of pixels showing younger ages ($\\sim0.01$Gyr). Otherwise, both age values and their associated uncertainties in the pixels are qualitatively similar for all 3 models. The red pixels in the inter-arm regions indicate much older populations but since these are also typically lower signal-to-noise pixels, the rms error is also significantly worse for these pixels. Pixels belonging to the sky which are artificially fit are also assigned large uncertainties.  The dust obscuration, shown in Figure~\\ref{fig:models_optical_ebv} is also correlated with the spatial distribution of the stellar age. Unlike the disagreement seen in the age values of the spiral arm between the M05 models on one hand and the BC03 and CB07 models on the other hand, the obscuration maps demonstrate a high level of agreement among the various models. In all the models, the arm regions, which were predicted to have the youngest stellar populations, also have a moderately high level of dust obscuration, with $E(B-V)\\sim 0.5$. The uncertainties in the dust obscuration in the pixels are similar among the various models as well. These uncertainties correlate strongly with morphological features in the image as in the case of the stellar population age. The knots on the spiral arms, corresponding to star-forming HII regions, are clearly displayed in the obscuration uncertainty map with the lowest errors of the entire galaxy. These statistical uncertainties are similar among all the models considered. The inner arms and the outermost edges of the galaxy, which are the regions with the lowest signal-to-noise in the entire galaxy, have, predictably, the highest rms uncertainties in their pixel-$z$ quantities.  In order to quantify these effects, we use pixels from all 467 galaxies in our sample. The systematic error due to different population synthesis models is measured according to equation 1, with the resulting distribution shown in Figure~\\ref{fig:models_test_optical}. The distributions are generally non-Gaussian, peaking around zero systematic error, and showing a relatively small fraction of outliers. The shape of the distribution is a result of the quantized nature of the problem, since the pixel-$z$ parameters lie on a discrete grid of values. The level of agreement and the fraction of outliers in this distribution for every pixel-$z$ parameter is summarized in Table 1 for each of the pairs of models being compared. The two models that have the highest level of agreement in all the parameters concerned (age, $\\tau$ and E(B-V)) are BC03 and CB07. In the $\\tau$ parameter, 94 percent of pixels within the galaxies are within $1\\sigma$ and the fraction of outliers ($>2\\sigma$) is 0.04. The high level of agreement between BC03 and CB07 is due to the fact that the continuum in the rest-frame UV is very similar between these two models. There is a larger systematic difference in the pixel-$z$ parameters between BC03 and M05, and between CB07 and M05. In particular, the $\\tau$ parameter is most sensitive to differences between these models. The fraction of pixels with a systematic difference in $\\tau$ less than $1\\sigma$, is 0.71 and 0.80 for the [BC03,M05] and [CB07,M05] model pairs respectively. The fraction of outliers in $\\tau$ ($>2\\sigma$) is 0.15 and 0.17 respectively for the same model pairs.  The impact of model differences between BC03/CB07 and M05 is much smaller for the age and obscuration. Dust obscuration is least sensitive to differences between any of the 3 models, with only $1-3$ percent of pixels having systematic differences in $E(B-V)$ which are larger than $2\\sigma$. For the inferred stellar age, the fraction of pixels having systematic differences in their stellar population age of $<1\\sigma$ is 0.86 and 0.87 for [BC03,M05] and [CB07,M05] respectively, while the fraction of outlier pixels in the age distribution for the same pairs of models is 0.11. For [BC03,CB07], 0.94 of pixels agree in their inferred stellar population age to within $1\\sigma$, while the fraction of outlier pixels for the same models is 0.04. The comparison of the stellar population synthesis models for the metallicity are not shown in Table 1 because the small number of values of metallicity being sampled means that all fits are within $2\\sigma$.  Out of the three pixel-$z$ parameters tested, it is thus in the SFR e-folding time where the largest discrepancy lies between M05 on one hand and BC03/CB07 on the other.  This discrepancy is unlikely to be due only to the differing treatment of TP-AGBs stars in these models, since if that were the case, we would expect better agreement between CB07 and M05 as contrasted against BC03. The discrepancy observed in $\\tau$ is likely related to fundamental differences between M05 on one hand and BC03/CB07 on the other when the SED fitting is performed using optical colors only. In section 3.5.2, we use some example galaxies to investigate where the discrepancy in the predicted values of $\\tau$ (between M05 and BC03/CB07) occurs spatially in the galaxies. Finally, we test the effect of redshift on these results by repeating the above analysis for subsamples of galaxies in redshift intervals of 0.5 from $z=0.5-3$ and find that our results for the optical-only passbands are not significantly impacted by the redshift of the galaxy.    \\subsection{Effect of Parameter Space and SED Degeneracies} \\label{subsec:paramspace}   In this test, we examine the effect of changing the parameter space on the pixel-$z$ estimates using the BC03 models. The number of metallicity values used in the pixel-$z$ grid of stellar population parameters is increased from 3 to 6, to include sub-solar metallicities, as detailed in Section~\\ref{sec:sspsfromcolors}. The increase in the metallicity range means we are comparing the original 2178 SEDs with an expanded set of 4356 SEDs fitted to each pixel.   The distribution of the systematic error (as a function of the statistical error) introduced by extending the allowed metallicity range into the sub-solar regime is shown in Figure~\\ref{fig:metallicity_test} for all pixels in all the galaxies in the sample. The fraction of pixels with small systematic errors ($<1\\sigma$) and the fraction of outliers in each parameter is summarized in Table 2. The largest effect of a broadening the metallicity range is, predictably, on the stellar population age but the effect is minimal: the fraction of outliers in the age systematic is only 0.06. Around 6 percent of pixels have a bias in their $\\tau$ that is larger than the $2\\sigma$ uncertainties. The obscuration is the parameter which is least affected ($< 0.01$ of pixels) by the introduction of sub-solar metallicities. This is in agreement with the findings of Conti et al.~(2003) who examined the evolution of the obscuration-metallicity degeneracy in pixels in galaxies in the HDF-North, and found that there is a strong degeneracy in galaxies at low redshifts but that the relation flattens out at $z>1$.  The conclusion of this test is that while the stellar population age and $\\tau$ are marginally affected by the increased metallicity range, as expected from the age-metallicity degeneracy, the resulting systematic differences in all the pixel-$z$ quantities are still well within the uncertainties between different stellar population synthesis models.  \\subsection{Effect of Passband Limitations (Optical-only)}  Here we test the effect of the passband limitations on the estimates of the pixel-$z$ parameters. The stellar population synthesis model is fixed to BC03 in this test and passbands are added and removed and the pixel-$z$ output is compared to that produced using the full complement of 4 optical bands. This test is initially performed for selected individual galaxies to identify the likely trends, before a full investigation on all pixels in the full galaxy sample is performed. Figure~\\ref{fig:passbands_optical_age} illustrates the test for one example galaxy, which is the same as shown in Figures~\\ref{fig:models_optical_age} and~\\ref{fig:models_optical_ebv}. It shows the impact of passband variations on the age values in the pixels and their associated uncertainties. Starting with only 2 optical bands ($bv$), the effect of different permutations of bands are tested. From the stellar population age maps, it is clear that 2 passbands are insufficient to resolve many of the features seen in the light distribution of the galaxy. For the case of 3 bands, the use of the redder $viz$ combination of filters comes closer to the result obtained with the full $bviz$ filters than does the $bvi$ set.  The $viz$ combination produces an age map that comes closer to mirroring the galaxy morphology than does the $bvi$ set, separating the arm and inter-arm regions more clearly. For this example galaxy at $z\\sim1$, most of the relevant color information for the pixel-$z$ therefore comes from the optically redder passbands.  A similar conclusion can be drawn from the map of rms values in the age parameter. With only 2 optical bands, high uncertainties throughout the disk of the galaxy reflect the inability of pixel-$z$ to find the correct SED with only 1 color, thus resulting in higher statistical errors in the ages inferred for the pixels. With 3 passbands, again it is the $viz$ combination that seems to be closest to the $bviz$ values. However, the maps of dust obscuration in Figure~\\ref{fig:passbands_optical_ebv} show that the $bvi$ combination coming closest to the result predicted from the full complement of bands. Nevertheless, the $bvi$, $viz$ and $bviz$ combinations all agree that the arms have intermediate levels of obscuration ($E(B-V)\\sim0.4$) and low associated RMS errors.  Moving to the full sample of all pixels in all galaxies, we use equation 2 to compute for each parameter and for each pixel, the systematic error (again in terms of the statisitical error in the parameter) due to the difference between a given filter permutation and the full optical filter set ($bviz$). The stellar population model is fixed to BC03 as before. Figure~\\ref{fig:passbands_optical_test} shows the distribution of this systematic error for all the pixels in the full galaxy sample. It is evident that the systematic error distribution has larger tails and a lower peak ($<1\\sigma$) than the systematic errors resulting from stellar population model differences. As for the example galaxy in Figures~\\ref{fig:passbands_optical_age}, it is clear also that 2 filters ($bv$) result in a large fraction of outliers ($>3\\sigma$). For the 3-filter permutations ($biz$, $bvi$ and $viz$), across the entire sample of galaxies, we find that there is only a small difference in the distribution of the systematic error among the filter permutations.  We show the fraction of outliers in each case in Table 3. It is clear that the SFR e-folding time ($\\tau$) and the stellar population age are the most impacted of the 3 parameters by passband limitations. In $\\tau$ and in the stellar population age, over 20 percent of the pixels become outliers ($>2\\sigma$) when one band of the full filter set $bviz$ is dropped. There is not a very significant difference between the 3 sets of filters, but omitting the $z$ band results in an outlier fraction for the $\\tau$ parameter of 0.26 compared to 0.20 for the other bands. For the stellar population age, the outlier fraction varies from $0.22-0.28$ relative to the full filter set ($bviz$). The dust obscuration is least impacted by passband changes, as long as a minimum of 3 are used. For each of the 3 filter sets, about 80 percent of pixels in the full galaxy sample have an $E(B-V)$ which agrees with the $bviz$ result to within $1\\sigma$. Finally, as in Section~\\ref{subsec:model_tests}, we test the effect of redshift on these results by repeating the above analysis for subsamples of galaxies in different redshift intervals from $z=0.5-3$ and find that our results for the optical-only passbands are not significantly impacted by the redshift of the galaxy.   \\subsection{Effect of Adding the Near-IR Data}  In this section, we explore the effect of adding NIR passbands from NIC3 to the existing optical data for a subsample of 46 galaxies. Adding the NIR data allows us to explore systematic differences between various population synthesis models more fully, since they can predict different NIR colors. In particular, it is known that BC03 and M05 differ significantly in their treatment of TP-AGB stars, consequently resulting in different predicted NIR colors. This is illustrated in Figure~\\ref{fig:spectra} which shows an example SED from all three models considered in this work, for a galaxy at $z=1.0$ containing a stellar population that is 3 Gyr old, with $\\tau=10$ Gyr, $E(B-V)=0.9$ and $Z=0.008$. There is very little difference between the shape of the continuum between the the models in the optical. However, the same galaxy in the NIR shows significant differences in the shape of the continuum among the three models. These will give rise to different predicted NIR colors in the pixels.   Figures~\\ref{fig:age_nir_example} and ~\\ref{fig:ebv_nir_example} show examples of applying pixel-$z$ to a galaxy at $z=2.5$ with $bviz$ and $J$ and $H$ band imaging from NIC3. There is a good agreement among the three models compared here in both their stellar population age maps and in the age rms map. All three models predict a very young population of stars throughout ($t<0.2$ Gyr). All models considered predict the core of the galaxy to have an older population of stars ($\\sim0.1$ Gyr) than the outskirts ($t\\sim0.01$Gyr). The one notable difference is that the M05 models predict a slightly larger proportion of very young stellar populations ($t\\sim0.001$Gyr) surrounding the nucleus. The models also agree very well on the age distribution within the companions of this galaxy (to the right). The models also agree well in the uncertainty maps in the stellar population age. Figure~\\ref{fig:ebv_nir_example} shows the obscuration maps of the same galaxy as predicted by the three models. Again, there is excellent agreement among the three models in the distribution of dust and its associated uncertainty. In general, all three models predict that the galaxy has moderate levels (E(B-V)$\\sim0.3$ magnitudes) of dust, but they also show pockets of high obscuration (E(B-V)$\\sim0.5$ magnitudes) within the main galaxy and also in its two companions. The location of these pockets are the same in three models. This example therefore implies that model differences do not introduce large systematic biases in the pixel-$z$ parameters for $z\\sim2.5$ galaxies with such irregular morphologies.  We investigate this hypothesis fully when we measure systematic differences across all pixels in the combined optical-NIR galaxy sample for all 3 models considered in this paper. For each parameter, we compare this distribution of the systematic differences with that obtained with using only the optical data. We account for the effect of redshift, particularly because for $z>1$, the 4000\\,\\AA \\ break shifts into the $J$ and $H$ passbands. We thus split the NIR sample into two: $z<1$ and $z>1$. In Figure~\\ref{fig:nir_test_zle1}, we show the distribution of the systematic difference between models for the low redshift sample. The fraction of pixels showing $<1\\sigma$ uncertainties in their pixel-$z$ parameters for each pair of models is given in Table 4 for both the optical-only colors and the optical and NIR colors. For the low redshift sample, there is a very high level of agreement between the BC03 and CB07 models in all the pixel-$z$ parameters, with over $97$ percent of pixels showing less than $1\\sigma$ deviation between the two models. With the addition of the NIR bands however, the systematic differences between BC03 and CB07 models are exposed.  This is a result of relatively small differences among the models in their predicted rest-frame UV colors but much larger differences among them in their predicted rest-frame optical (red) and rest-frame NIR colors. This affects in particular the stellar population age (with only 78 percent of pixels showing less than $1\\sigma$ difference between the two models) and $\\tau$ (88 percent of pixels showing agreement between the two models). The obscuration is less impacted by the model differences. The addition of the NIR colors also exposes differences between BC03 and M05, and this impacts the stellar population age most strongly, with the fraction of pixels showing $<1\\sigma$ difference in the age decreasing from 0.89 with optical-only colors to 0.80 with NIR colors. The differences between BC03 and M05 have less impact on $\\tau$ and the obscuration when the NIR colors are added. When comparing CB07 with M05, it is again the stellar population age in the pixels that is most affected by the addition of the NIR passbands, with only 80 percent of pixels showing an agreement in the age to within $1\\sigma$ between the two models, compared to 89 percent for the optical-only case.  For the high-redshift ($z>1$) sample, the story is somewhat different, as illustrated in Figure~\\ref{fig:nir_test_zgt1} and Table 5. There is now a much higher level of agreement between the results for the optical-only colors and the optical and NIR colors combined, for all the parameters concerned. This is due to the fact that the rest-frame UV and blue wavelength optical spectra (which are sampled by the optical and the NIR filters for $z>1$) are similar among the models. This similarity in the predicted rest-frame UV/optical (blue) colors among the models thus manifests itself as follows.  In the $z<1$ sample discussed above, the stellar population age showed a large discrepancy between the $bviz$ and $bvizJH$ samples in terms of the fraction of pixels that agree on their age for the different models being compared. For the high redshift sample, the pixels in both the $bviz$ and $bvizJH$ samples show similar levels of agreement in the age parameter between the different models compared.  For example, for the the BC03 and M05 models, the fraction of pixels with systematic differences in their age that are smaller than $1\\sigma$ is 0.95 in the optical-only case and 0.93 in the NIR case. The conclusion of these tests is that, for $z>1$ galaxies, differences among the stellar population synthesis models in their predicted NIR colors do not significantly bias the pixel-$z$ estimates of the stellar population parameters in the pixels, compared to using optical colors only.  \\subsection{Global Galaxy Properties and the Effect of Systematics }  Pixel-$z$ can measure global properties of structures composed of many pixels within the same galaxy. These will be used to determine the radial variation of stellar population parameters in galaxies as a function of the galaxy environment, and also to allow a detailed study of high redshift clumpy, disk galaxies, as outlined in Section 1.1 and 1.2. Here, we investigate the effect of systematic biases that can impact these global measurements, in particular of the SFR since this will be a direct observable in spectroscopic measurements.  \\subsubsection{Biases from using pixel-$z$ versus global SED-fitting}  In Figure~\\ref{fig:SFR_pixelz_global}, we show the level of consistency between the sum of SED fits to individual pixels in a galaxy and the global fit to the total flux in those same pixels, as described in Section 2.3.2. We see a high level of consistency between the two approaches in their estimates of the integrated SFR. We also quantify the magnitude of any deviation between the two approaches in terms of the statistical uncertainty in the SFR estimate for each galaxy from SED-fitting. We see that for the majority of galaxies in the sample, the difference between the two SFR estimates for the galaxies are within the statistical uncertainties of the SFR measurement. However, there is a small population of outliers ($>2\\sigma$) which are galaxies with low SFR ($<1 M_{\\odot}\\,yr^{-1}$) and which show a negative residual i.e., the SFR estimates from pixel-$z$, obtained by summing the results of the individual fits in each pixel, are higher than the SFR estimates from global fits to the total flux in the pixels. In addition, there is a second small population of outliers which are galaxies with high SFRs ($\\sim10 M_{\\odot}\\,yr^{-1}$). In these, the SFR estimate from the global fit is significantly higher than that estimated from the sum of the individual SED fits.   \\subsubsection{Radial gradients of stellar populations: age, dust and star formation}  In Figure~\\ref{fig:radialgradient_models}, we present the radial variation of the stellar population age, SFR e-folding time, obscuration, SFR and sSFR using the $bviz$ passbands for the same galaxy at $z\\sim1$ in Figures~\\ref{fig:models_optical_age} and~\\ref{fig:models_optical_ebv}. The radial trends in each quantity are measured according to the method described in section 2.3.3. All stellar population synthesis models agree on the qualitative trends of all the parameters with radius but there are some quantitative differences that arise, particularly between the different population synthesis models All three models agree that the stellar populations in the galaxy at all radii are young ($<0.3$Gyr) and that the stellar population age decreases rapidly from the center to the outskirts. M05 predicts a mean stellar population age that is somewhat lower in the innermost part of the galaxy compared to BC03 and CB07, and this systematic difference is of the order of the statistical uncertainty in the age measurement. BC03 and CB07 also predict a somewhat sharper drop in the stellar population age between 2 and 3 kpc, but again, the difference is well within the statistical uncertainties.  The models agree on a range of $\\tau\\sim0.5-1$ Gyr across the galaxy. Between 3 and 5 kpc, however, the M05 model prediction diverges by $1\\sigma$ from both BC03 and CB07 models. This region corresponds to the spiral arm and the clumps of star-forming HII regions within it. This implies that the different treatment of TP-AGB stars in M05 compared to the other models has its biggest effect in the spiral arms of the galaxy. The dust obscuration trends are quite similar among the models which show that the center of the galaxy is moderately obscured ($E(B-V)\\sim0.45$). The obscuration then decreases up to 2 kpc and then increases to a maximum ($E(B-V)\\sim0.6$) in the spiral arms.  The different models all find that the mean SFR in each annulus peaks in the center of the galaxy. To explain this, we note that the current SFR is determined not only by the mean age and $\\tau$ of the stellar populations, but also by the total stellar mass formed to date, which normalizes the SFR. M05 predicts a lower stellar mass throughout the galaxy than does BC03 or CB07. When the stellar mass is normalized out to give the specific SFR (sSFR), we see that the discrepancy between M05 and BC03/CB07 narrows considerably: $1\\sigma$ at most in all radial bins. Comparing the trend in sSFR with the morphology of the galaxy, we conclude that the bar structure in the center of the galaxy is actively star-forming. The models disagree, however, on how much more star-forming the bar is relative to the spiral arms at a radius of $\\sim4$kpc. BC03 and CB07 predict a slightly higher sSFR (by around $2\\times10^{-11}\\,yr^{-1}$) in the center of the galaxy compared to the spiral arms, whereas M05 predicts a somewhat lower sSFR in the center relative to the spiral arms. Nonetheless, a general picture emerges for this galaxy, which is that star formation is not simply confined to the center but takes place throughout the galactic disk. The results, albeit for a single galaxy, imply that pixel-$z$ is robust to differences between various stellar population synthesis codes when it is used to measure radial variations across galaxies.  Figure~\\ref{fig:radialgradient_passbands} shows the effect of passband limitations on measuring radial variations of stellar population parameters. The population synthesis model was fixed to BC03. Using only the $b$ and $v$ bands results in large systematic uncertainties which affects all the radial quantities but particularly the mean age and the obscuration determinations. The result of using only the $b$ and $v$ bands is that the mean age is overestimated throughout the galaxy relative to the $bviz$ filter set, by as much as $4\\sigma$. This can be explained by the fact at $z\\sim1$, the 4000\\,\\AA \\ break passes out of the bluest optical bands and, as a result, the colors predicted by the model for the pixels are redder, leading to an older stellar population for the pixels. The obscuration is underestimated throughout, and at a radius of 4 kpc the difference with respect to the $bviz$ value is as much as $4\\sigma$. The $\\tau$ and the sSFR estimates are less impacted by the loss of 2 passbands.  The 3-passband combinations $bvi$ and $viz$ both come closest to reproducing the $bviz$ radial trends in all the parameters. The systematic difference between the $bvi$ and $viz$ combinations is relatively small, but some differences do emerge nevertheless. In the center of the galaxy, the $viz$ filter combination predicts a larger (by $1\\sigma$) mean stellar population age than the bluer $bvi$ bands. Elsewhere in the galaxy, the three-filter combinations give results which are consistent (within their statistical uncertainties) with each other and also with the full $bviz$ filter set. The radial variation of the dust obscuration is consistent among the $bvi$, $viz$ and $bviz$ filter sets except in the innermost annulus where the $viz$ filters underpredict the level of obscuration relative to the other two combinations. The differences in $\\tau$ are small, although there is one important difference at $r=3$kpc where the $bvi$ set predicts a higher $\\tau$ compared to the $viz$ and $bviz$ sets. The effect of filter cominbations on the sSFR is also relatively small throughout the galaxy except in its outermost part ($r\\sim5$kpc) where the three-filter combinations overestimate the sSFR relative to the full optical set.  If we were to generalize this test, a minimum of 3 optical passbands are therefore needed for robustly computing radial trends of stellar populations in these high redshift galaxies.  \\subsubsection{Disk galaxies at $z\\gtrsim1$: spiral arms and clumps}  In Figures~\\ref{fig:spiralarms}, \\ref{fig:clumpA} and \\ref{fig:clumpB}, we illustrate the way pixel-$z$ can be used to study the stellar populations of subcomponents of galaxies such as spiral arms and clumps, and we explore how this can be biased by the choice of the SPS model. The optical filter set $bviz$ is used throughout this analysis. In the spiral arm of the disk galaxy which has been isolated, the BC03 and CB07 models predict almost identical distributions for the parameters among the pixels that make up the spiral arm. In contrast, some differences start to emerge between M05 and the other two models. The age distribution of the pixels peaks at 0.1 Gyr for the BC03 and CB07 models, and at a lower age (0.01 Gyr) for the M05 model. All three models show good agreement for the dust distribution among the pixels in the arm, which peaks at $0.5$ Gyr. There is however a systematic offset in the SFR distributions which peak at 0.003 $M_{\\odot}\\,yr^{-1}$ with M05, whereas it peaks at 0.03 $M_{\\odot}\\,yr^{-1}$ with BC03/CB07. This difference is largely due to the lower stellar mass predicted for the pixels in the arm by M05 compared to the other models. When stellar mass is normalized, we see a much closer agreement between M05 and BC03/CB07. All the models predict a sSFR distribution that is bimodal in the arm, one peaking at $10^{-9.8}yr^{-1}$ and the other at $10^{-8.3}yr^{-1}$.  We also investigate clumps in the disk and the impact of model differences on the inferred properties of clumps in the disk of the galaxy. Clump A in Figure~\\ref{fig:clumpA} is part of the outer arm of the galaxy. As in the case of the spiral arm, the distribution of the stellar population age and dust obscuration among the pixels in clump A, is quite similar among the different population synthesis models. As with the full spiral arm, a substantial difference exists in the SFR distributions. The M05 models predict the clump to have a lower SFR than BC03/CB07. The distribution peaks at $\\sim0.01 M_{\\odot}\\,yr^{-1}$ with M05 and $0.3 M_{\\odot}\\,yr^{-1}$ with BC03/CB07, due to a lower stellar mass predicted for the clump by the M05 model. It is also worth nothing that the peak of the SFR distribution within the clump occurs at a higher SFR than in the spiral arm. The sSFR distributions in the clump are far more consistent among all the models, again showing a bimodal star formation in the clump, with one population peaking at $10^{-9.8}yr^{-1}$ and the other at $10^{-8.3}yr^{-1}$ as in the case of the spiral arm. The SFR and sSFR distributions imply that the star formation process in the spiral arm might be driven by the star formation in the clump itself.  Finally, in Figure~\\ref{fig:clumpB}, we examine clump B on the left side of the galaxy. Again, model differences affect primarily the SFR, with M05 predicting the SFR distribution to peak at $\\sim0.001  M_{\\odot}\\,yr^{-1}$ while BC03/CB07 predict that it peaks around 0.03 $M_{\\odot}\\,yr^{-1}$. Irrespective of the model, however, it is evident that the mean SFR of clump B is significantly lower than clump A, indicating that clump A is the primary driver of star formation in the outskirts of the galaxy. Despite some differences between M05 and the other models, therefore, pixel-$z$ can be used to effectively and accurately probe mean properties of stellar populations within components of galaxies such as spiral arms and clumps. As discussed in section 1.1, this can be used to probe substructure in $z\\sim2$ galaxies in order to determine if these objects are largely driven by mergers or are disks formed in-situ.   \\subsection{Incorporating WFC3 to Probe Substructure: Expectations from the CANDELS Survey} \\label{subsec:candels}   Figure~\\ref{fig:CANDELS_maps} shows the results of adding the NIR WFC3 imaging in the $F125W$ and $F160W$ passbands from the CANDELS survey to the existing $bviz$ bands for an example disk galaxy at $z=0.6$, with the M05 models. The most notable difference between using the optical-only passbands and the combined optical and WFC3/NIR passbands is in the spatial distribution of the oldest ($t>4$ Gyr) stellar populations in the galaxy. In the ACS-only case, the oldest stellar populations in the galaxy are confined in highly localized regions in the galaxy, with a concentration in the center and and in pockets in the south-west and sout-east of the image.  The addition of the WFC3/NIR data, however, suggests that the oldest (and thus redder) populations are not simply confined in these pockets but are present more extensively in the disk itself. In the map of the specific star formation rate in the galaxy, both the optical-only and the combined optical and NIR cases suggest a qualitatively similar trend in the spatial distribution of the sSFR i.e., the sSFR is lowest in the center of the galaxy ($\\sim10^{-11}\\,yr^{-1}$) and increases towards the outskirts. Away from the center, however, the addition of the WFC3/NIR passbands, suggests that the disk contains a higher proportion of stellar populations with a sSFR ($\\sim10^{-10}\\,yr^{-1}$) that is lower than what is predicted by the optical-only colors. The differences observed in the maps of the stellar population age and sSFR are confirmed by the radial variation of these quantities, shown in Figure~\\ref{fig:CANDELS_radialplots}. The WFC3/NIR colors reveal an older population throughout the disk of the galaxy relative to the optical-only colors, with more than a $5\\sigma$ difference in the luminosity-weighted mean age between the $bviz$ and $bvizJH$ colors and up to a radius of 6 kpc from the galaxy center, . The radial variation of the sSFR also reflects the fact that the WFC3/NIR colors are more sensitive to the older stellar population in the disk, which were otherwise not revealed by the ACS-only passbands. The WFC3/NIR colors reveal a lower sSFR at all radii compared to the ACS-only colors, although the difference is most significant between 3 and 5 kpc from the galaxy center."
    },
    "1112/1112.0187_arXiv.txt": {
        "abstract": "Gaia is a European Space Agency (ESA) astrometry space mission, and a successor to the ESA Hipparcos mission. Gaia's main goal is to collect high-precision astrometric data (i.e. positions, parallaxes, and proper motions) for the brightest 1 billion objects in the sky. These data, complemented with multi-band, multi-epoch photometric and spectroscopic data collected from the same observing platform, will allow astronomers to reconstruct the formation history, structure, and evolution of the Galaxy.  Gaia will observe the whole sky for 5 years, providing a unique opportunity for the discovery of large numbers of transient and anomalous events, e.g. supernovae, novae and microlensing events, GRB afterglows, fallback supernovae, and other theoretical or unexpected phenomena. The Photometric Science Alerts team has been tasked with the early detection, classification and prompt release of anomalous sources in the Gaia data stream. In this paper, we discuss the challenges we face in preparing to use Gaia to search for transient phenomena at optical wavelengths. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.2182_arXiv.txt": {
        "abstract": "We present a model to self-consistently describe the joint evolution of starburst galaxies and the galactic wind resulting from this evolution. This model will eventually be used to provide a subgrid treatment of galactic outflows in cosmological simulations of galaxy formation and the evolution of the intergalactic medium (IGM). We combine the population synthesis code Starburst99 with a semi-analytical model of galactic outflows and a model for the distribution and abundances of chemical elements inside the outflows. Starting with a galaxy mass, formation redshift, and adopting a particular form for the star formation rate, we describe the evolution of the stellar populations in the galaxy, the evolution of the metallicity and chemical composition of the interstellar medium (ISM), the propagation of the galactic wind, and the metal-enrichment of the intergalactic medium. The model takes into account the full energetics of the supernovae and stellar winds and their impact on the propagation of the galactic wind, the depletion of the ISM by the galactic wind and its impact on the subsequent evolution of the galaxy, as well as the evolving distributions and abundances of metals in the galactic wind. In this paper, we study the properties of the model, by varying the mass of the galaxy, the star formation rate, and the efficiency of star formation. Our main results are the following: (1) For a given star formation efficiency $f_*$, a more extended period of active star formation tends to produce a galactic wind that reaches a larger extent. If $f_*$ is sufficiently large, the energy deposited by the stars completely expels the ISM. Eventually, the ISM is being replenished by mass loss from supernovae and stellar winds. (2) For galaxies with masses above $10^{11}\\msun$, the material ejected in the IGM always falls back onto the galaxy. Hence lower-mass galaxies are the ones responsible for enriching the IGM. (3) Stellar winds play a minor role in the dynamical evolution of the galactic wind, because their energy input is small compared to supernovae. However, they contribute significantly to the chemical composition of the galactic wind. We conclude that the history of the ISM enrichment plays a determinant role in the chemical composition and extent of the galactic wind, and therefore its ability to enrich the IGM. ",
        "introduction": "Galactic winds and outflows are the primary mechanism by which galaxies deposit energy and metal-enriched gas into the intergalactic medium (IGM).\\footnote{Some authors make a distinction between {\\it galactic winds}, which are generated over most of the lifetime of the galaxy and inject energy and metals at a steady rate, and {\\it galactic outflows}, which result from violent processes like starbursts, are short-lived, and eject material at large enough distances into the IGM to eventually reach other galaxies. In this paper, we use one or the other to designate any material that is ejected from the galaxy and deposited into the IGM.} This can greatly affect the evolution of the IGM, and the subsequent formation of other generations of galaxies. Feedback by galactic outflows can provide an explanation for the observed high mass-to-light ratio of dwarf galaxies and the abundance of dwarf galaxies in the Local Group, and can solve various problems with galaxy formation models, such as the overcooling and angular momentum problems (see \\citealt{benson10} and references therein). Galactic outflows can explain the metals observed in the IGM via the Lyman-$\\alpha$ forest (e.g. \\citealt{my87,schayeetal03,ph04,aguirreetal08,pierietal10a, pierietal10b}), the entropy content and scaling relations in X-ray clusters \\citep{kaiser91,eh91,cmt97,tn01,babuletal02,voitetal02}, and provide observational tests that can constrain theoretical models of galaxy evolution. Local examples of spectacular outflows in dwarf starburst galaxies include those of the extremely metal-poor I Zw 18 (\\citealt{pequi08,jamet10}) and NGC1569 \\citep{west09}. More massive spirals, such as NGC7213 \\citep{hameed01}, also show evidence of global outflows. For a review of the subject, see \\citet{vcbh05}.  \\subsection{Galactic Outflow Models}  Large-scale cosmological simulations have become a major tool in the study of galaxy formation and the evolution of the IGM at cosmological scales. These simulations start at high redshift with a primordial mixture of dark and baryonic matter, and a spectrum of primordial density perturbations. The algorithm simulates the evolution of the system by solving the equations of gravity, hydrodynamics, and (sometimes) radiative transfer. Adding the effect of galactic outflows in these simulations poses a major practical problem. In one hand, the computational volume must be sufficiently large to contain a ``fair'' sample of the universe, typically several tens of Megaparsecs. On the other hand, the physical processes responsible for generating the outflows take place inside galaxies, at scales of kiloparsecs or less. This represents at the very minimum 4 orders of magnitude in length and 12 orders of magnitude in mass, which is beyond the capability of current computers. Since we cannot simulate both large and small scales simultaneously, the usual solution consists of simulating the larger scales and using a {\\it subgrid physics} treatment for the smaller scale. Cosmological simulations can predict the location of the galaxies that will produce the outflow, but cannot resolve the inner structure of these galaxies with sufficient resolution to simulate the actual generation of the outflow. Instead, the algorithm will use a prescription to describe the propagation of the outflow and its effect on the surrounding material.  One possible approach consists of depositing momentum or thermal energy ``by hand'' into the system, to simulate the effect of galactic outflows on the surrounding material \\citep{std01,theunsetal02,sh03,cno05,od06, kollmeieretal06}. The algorithm determines the location of the galaxies producing the outflows and calculates the amount of momentum or thermal energy deposited into the IGM based on the galaxy properties (mass, formation redshift $\\ldots$). Then, in particle-based algorithms like smoothed particle hydrodynamics (SPH), this momentum or energy is deposited on the nearby particles, while in grid-based algorithms it is deposited on the neighboring grid points. This will result in the formation and expansion of a cavity around each galaxy, which is properly simulated by the algorithm.  A second approach consists of combining the numerical simulation with an analytical model for the outflows. Tegmark et al. (1993) have developed an analytical model to describe the propagation of galactic outflows in an expanding universe. In this model, a certain amount of energy is released into the interstellar medium (ISM) by supernovae (SNe) during an initial starburst. This energy drives the expansion of a spherical shell that propagates into the surrounding IGM, until it reaches pressure equilibrium. This model, or variations of it, has been used extensively to study the effect of galactic outflows on the IGM (\\citealt{fl01,mfr01,sb01,sfm02,sh04,lg05,pmg07}, hereafter PMG07; \\citealt{sss08,gbm09,pmp10}). In this approach, the evolution of the IGM and the propagation of the outflow are calculated separately, but not independently as they can influence one another. The presence of density inhomogeneities in the IGM can affect the propagation of the outflow, while energy and metals carried by the outflow can modify the evolution of the IGM.  There are several limitations with this second approach. In particular, it assumes that the initial starburst, which occurred during the formation of the galaxy, is the only source of energy driving the expansion of the outflow. First, the starburst lasts for a short period of time, typically 50 million years. Hence, we would expect to observe very few galaxies having an outflow. Furthermore these galaxies would be just forming and therefore would have complex and chaotic structures. Observations show instead that outflows are ubiquitous and often originate from well-relaxed galaxies (see \\citealt{vcbh05} and references therein). Second, even though the injection rate of energy is maximum during the initial starburst, the total amount of energy which is injected afterward by all generations of SNe and stellar winds could be comparable or even more important. Even if this energy is injected slowly over a large period of time, the cumulative effect could be significant. Indeed, an initial outflow caused by a starburst could be followed by a steady galactic wind that would last up to the present. The role of stellar winds has been mostly ignored in analytical models and numerical simulations of galactic outflows. Third, there is the possibility that accretion of intergalactic gas onto the galaxy might trigger a second starburst. Finally, a recent study \\citep{sbh10} suggests that there is a significant time delay between the initial starburst and the onset of the outflow, something not considered by current models.  Another important issue is the amount of metals contained in the outflow, the spatial distribution of metals in the outflow, and the relative abundances of the various elements. The metallicity of the outflow depends on the metallicity of the ISM at the time of the starburst. The metals contained in the ISM at that time can have several origins: (1) metals already present in the gas when the galaxy formed, (2) the SNe produced during the starburst, and (3) the stellar winds generated by massive stars and AGB objects. Hence, the composition of the outflows will depend on the epoch of formation of the galaxy (which determines the initial metal abundances), as well as the relative amount of metals injected into the ISM by Type~Ia SNe, core-collapse SNe (Types~Ib, Ic and II), and winds, and the timing of these various processes. As for the amount of metals injected in the IGM, models often assume that it is proportional to the mass of the galaxy, and do not provide a description of the distribution of metals in the outflow and the relative abundances of the elements (\\citealt{sb01}; PMG07).  \\subsection{Objectives}  {\\it Our goal is to develop a new galaxy evolution model to improve the treatment of galactic winds in cosmological simulations.} This model will describe not only an initial starburst and its resulting outflow, but the entire subsequent evolution of a galaxy up to the present. It will take into account the progressive injection of energy by SNe and stellar winds (which could cause a steady galactic wind that would follow the outflow and last up to the present), and the time-evolution of the metallicity and composition of the ISM that would directly affect the composition of the galactic wind. It will also provide a description of the metal content, metal distribution, and chemical composition of the galactic wind. This emphasis on the structure and composition of the galactic wind is what distinguish our model from recent semi-analytical models of galaxy formation, which tend to focus on reproducing the properties of the galaxies themselves (luminosity and mass functions, halo properties, disk sizes, $\\ldots$). For a review of the various semi-analytical models, see \\citet{baugh06}. Our approach combines a population synthesis algorithm to describe the stellar content of the galaxy, an analytical model for the expansion of the galactic wind, and a new model for the distribution of elements inside the galactic wind.  The paper is organized as follows: in \\S2, we describe the method we use to calculate the mass and chemical composition of stellar winds and SN ejecta, a key ingredient of our algorithm. In \\S3, we describe the basic equations for the evolution of the ISM. In \\S4, we describe our galaxy wind model. Results are presented in \\S5, and conclusions in \\S6. ",
        "conclusions": "We have combined a population synthesis code, interpolation tables for the mass and composition of SN ejecta, and an analytical model for galactic winds into a single algorithm that self-consistently describes the evolution of starburst galaxies. This model describes the evolution of the stellar populations in the galaxy, the evolution of the mass and chemical composition of the ISM, the propagation of the galactic wind, and the distribution and abundances of metals inside the galactic wind. In particular, the algorithm (1) provides a detailed calculation of the energy deposited into the ISM by SNe and stellar winds, which is responsible for driving the galactic wind, (2) takes into account the time-evolution of the chemical composition of the ISM, which directly affect the composition of the galactic wind, and (3) takes into account the removal of the ISM by galactic winds, which affects the metallicity of the ISM, and the metallicity of the stellar populations to follow.  \\begin{figure} \\begin{center} \\includegraphics[width=3.3in]{fig17.eps} \\caption{Metallicity profile of material ejected in the IGM at $z=0$, for a $10^9\\msun$ galaxy with an exponential SFR. The various colors represent different star formation efficiencies $f_*$. Colors and linetypes have the same meaning as in Figures~\\ref{mass2star} and~\\ref{ZISM_z}. } \\label{Z_IGM} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=3.3in]{fig18.eps} \\caption{Comoving radius of galactic wind versus redshift, for galaxies with an exponential SFR and a star formation efficiency $f_*=0.1$, for various galaxy masses $M_{\\rm gal}$ in solar masses. Solid lines: simulations with SNe only; dotted lines: simulations with SNe and stellar winds.} \\label{R_mgal} \\end{center} \\end{figure}  Our first results concern the SFR for the galaxy. For a given star formation efficiency $f_*$, a longer SFR tends to produce a galactic wind that reaches a larger extent, but this wind will be less dense. By increasing the star formation efficiency, we can produce a wind that reaches a larger extent and has a higher metallicity near its front. In some cases, the energy deposited by the stars is sufficient to completely expel the ISM. When it happens, star formation is shut down, and the galactic wind enters the post-SN phase prematurely. Hence, paradoxically, an increase in the star formation rate can sometimes result in a galactic wind that reaches a smaller extend. This happens with galaxies of masses $M_{\\rm gal}=10^8\\msun$ or less, because their shallow potential well enables the complete removal of the ISM by the galactic wind.  For galaxies with mass above $10^{11}\\msun$, the material ejected in the IGM always falls back onto the galaxy, no matter the value of $f_*$. Therefore, in the case of energy-driven galactic winds, lower-mass galaxies are more likely to be the ones responsible for enriching the IGM and potentially perturbing the formation of nearby galaxies. Below $10^{11}\\msun$, the extent of the galactic wind and its mass and metal content both increase with the mass of the galaxy at constant $f_*$. With different values of $f_*$, a less massive galaxy can sometimes produce a larger wind.  Our current model does not take into account the effect of Type Ia SNe. These are difficult to include, because of the uncertainties on the lifetime of the progenitors. The simulations presented in this paper start at redshift $z=15$, and end between redshifts $z=9$ and 6. The corresponding time periods are shorter than $1\\,{\\rm Gyr}$, which is shorter than the lifetime of several Type Ia progenitors. The energy produced by Type Ia SNe is about 20\\% of the energy produced by Type II SNe (see Fig.~10 of \\citealt{benson10}). Hence, including the Type Ia SNe would result in a slightly larger final radius for the outflow. A Type Ia SNe can produce up to 7 times more iron than a Type II SNe (see model W7 in \\citealt{nomotoetal97}), and their contribution to the iron enrichment of the ISM become important after $1\\,{\\rm Gyr}$ \\citep{wiersma10}. Hence, the abundances of iron we present in this paper are underestimated. But because of the delay, the additional iron produced would remain in the inner parts of the galactic wind.  We have assumed a minimum value of $M_i=8\\msun$ for the minimum mass of SNe progenitors. However, the correct value is actually quite uncertain. We did a few simulations with minimum masses of $6\\msun$ and $10\\msun$. Our preliminary results show differences of order 10\\% in the final radius of the outflow, and of order 20-60\\% in the total mass ejected, with the largest effect occurring when $M_i$ is reduced. We intend to study this in more detail in the future.  To conclude, properties of galactic winds depend on the host galaxy properties, such as the mass or star formation efficiency. The history of the ISM enrichment plays a determinant role in the chemical composition and extent of the galactic wind, and therefore its ability to enrich the IGM. The next step will consist of implementing this galactic outflow model into large-scale cosmological simulations of galaxy formation and the evolution of the IGM. These will be the first simulation of this kind to include a detailed treatment of the stellar winds and their impact on the chemical enrichment of the IGM"
    },
    "1112/1112.3071_arXiv.txt": {
        "abstract": "{Most galaxy evolutionary models require quasar feedback to regulate star formation in their host galaxies. In particular, at high redshift, models expect that feedback associated with quasar-driven outflows is so efficient that the gas in the host galaxy is largely swept away or heated up, hence suppressing star formation in massive galaxies. We observationally investigate this phenomenon by using VLT-SINFONI integral field spectroscopy of the luminous quasar 2QZJ002830.4-281706 at z=2.4. The spectra sample the optical emission lines redshifted into the near-IR. The [OIII]$\\lambda$5007 emission-line kinematics map reveals a massive outflow on scales of several kpc. The detection of narrow H$\\alpha$ emission reveals star formation in the quasar host galaxy, with $\\rm SFR\\sim 100~M_{\\odot}~yr^{-1}$. However, the star formation is not distributed uniformly, but is strongly suppressed in the region with the highest outflow velocity and highest velocity dispersion. This result indicates that star formation in this region is strongly quenched by the quasar outflow, which is cleaning the galaxy disk of its molecular gas. This is one of the first direct observational proofs of quasar feedback quenching the star formation at high redshift. } ",
        "introduction": "Most of the recent galaxy formation models invoke energetic outflows as a way to regulate the evolution of galaxies throughout the cosmic epochs \\citep{Silk,Bower,Springel}. In particular, quasars are expected to drive powerful outflows that eventually expel most of the gas in their host galaxies, thereby quenching both star formation and further black hole accretion \\citep[e.g. ]{Granato,DiMatteo,Menci,Bower,Hopkins,King}. According to those models, these quasar driven outflows are required to prevent massive galaxies from overgrowing, hence explaining the shortage of very massive galaxies in the local universe, and are responsible for the red color and gas poor properties of local elliptical galaxies.  Massive, large-scale outflows have been detected in the hosts of local quasars \\citep[e.g. ][]{Feruglio,Fischer,Sturm,Rupke}. However, models expect that most of the quasar feedback action occurs at high redshift, when quasars reach their peak activity (z$\\sim$2) and when star formation in the most massive galaxies is observed to decline. Evidence of outflows in luminous quasars has been detected up to very high redshift \\citep[e.g.][]{Allen,Maiolino,Alexander}. Indications that the strength of these AGN driven outflows anticorrelates with the starburst contribution to the infrared luminosity has been obtained in reddened quasars by \\cite{farrah11}. However, direct observational evidence that high-z quasar driven outflows quench star formation in their host galaxies is still missing.  Here we present VLT-SINFONI near-IR integral field spectra of the quasar 2QZJ002830.4-281706 (hereafter 2QZ0028-28). This object, at $z = 2.401$, was taken from the sample of strong [OIII]$\\lambda$5007 emitters discovered by \\cite{Shemmer}, and it is one of the most luminous quasars known. In this letter we show that the spatially resolved kinematics of the [OIII]$\\lambda$5007 line clearly reveals a prominent outflow on scales of several kpc. Even more interesting, the star formation traced by narrow H$\\alpha$ emission is suppressed in the region characterized by the strongest outflow. We suggest that this is the first observational evidence of feedback associated with a quasar-driven outflow that quenches star formation at high redshift, or one of the first.  \\begin{figure}[!] \\centering \\includegraphics[height=0.95\\linewidth,angle=90]{fig1.ps} \\vspace{0.05cm} \\caption{Upper panel: 2QZ0028-28 H band spectrum extracted from the central 0.5 arcsec, along with the various components used for the fit (see Appendix~\\ref{app1} for details). Vertical dashed lines indicate the rest frame wavelength of each line, by taking the [OIII]$\\lambda$5007 line peak for reference. Lower panel: Residuals of the fit. The green vertical lines enclose a the spectral zone affected by strong sky line residuals.} \\label{hband} \\end{figure}   \\begin{figure*}[!] \\includegraphics[height=0.4\\textwidth,angle=90]{fig2a.ps} \\hspace{1.7cm} \\includegraphics[height=0.4\\textwidth,angle=90]{fig2b.ps} \\caption{Left panel: Velocity field (first moment map) of the [OIII]$\\lambda$5007 line, showing the prominent excess of blueshifted gas with a bow-like morphology SE of the nucleus. White contours are at 330, 360, and 390 km/s. Right panel: Velocity dispersion (second moment map) of the [OIII]$\\lambda$5007 line, showing the excess of dispersion in the SE region. White contours are at 680 and 700 km/s. In both maps the black contours trace the continuum.} \\label{moments} \\end{figure*} ",
        "conclusions": "By using near-IR integral field spectroscopic observations we have revealed a powerful outflow in the host galaxy of the quasar 2QZ0028-28 at z=2.4. The outflow was revealed by the velocity field traced by the [OIII]$\\lambda$5007 line, redshifted into the H-band. We estimated that the outflow rate of ionized gas is about $\\rm 200~M_{\\odot}~yr^{-1}$, which is, however, a lower limit of the total gas outflow rate. Both the high outflow velocity ($\\rm > 1000~km/s$) and the fact that the wind is mostly traced by the [OIII] line (produced primarily in the NLR) strongly suggest that the outflow is mostly driven by the quasar. The outflow is not symmetric, the highest velocities and highest velocity dispersion are found in the region SE of the nucleus.  In the K-band, our data clearly reveal the presence of narrow H$\\alpha$ emission tracing star formation in the host galaxy, on scales of several kpc and with a rate of about $\\rm 100~M_{\\odot}~yr^{-1}$. However, star formation is not distributed uniformly in the host galaxy, but is mostly found in the regions not directly invested by the strong outflow. Instead, star formation is heavily suppressed in the SE region where the strongest outflow is detected. This observational result supports models invoking quasar feedback to quench star formation in massive galaxies at high redshift."
    },
    "1112/1112.3592_arXiv.txt": {
        "abstract": "{In this paper a new supernova catalogue containing data for 5526 extragalactic supernovae that were discovered up to 2010 December 31 is presented. It combines several catalogues that are currently available online in a consistent and traceable way. During the comparison of the catalogues inconsistent entries were identified and resolved where possible. Remaining inconsistencies are marked transparently and can be easily identified. Thus it is possible to select a high-quality sample in a most simple way. Where available, redshift-based distance estimates to the supernovae were replaced by journal-refereed distances. Examples of statistical studies that are now possible with this new catalogue are presented in this paper.} ",
        "introduction": "The observation and study of supernovae (SNe) have had significant scientific impact in the past 30 years. Most prominently, \\object{SN1987A} enabled probing neutrino properties such as mass \\citep{SNIa_neutrino_mass}, electric charge \\citep{SNIa_neutrino_charge} and magnetic moment \\citep{SNIa_neutrino_magnetic_moment}. Recently, SNe of type Ia have been used as ``standard candles'' to measure distances up to cosmological scales \\citep{SNIa_distances_calibration}. In contemporary scientific research SNe are still important. Young SNe for instance may be identified as cosmic-ray accelerators by measuring high-energetic gamma or neutrino radiation shortly after the SN explosion \\citep{pulsarmodel,andobeacom,rmwaxman}.  Zwicky started the first systematic SN search in 1932 at Caltech, but he and his collaborators detected the first SNe only after the newly built 18\\arcsec Schmidt telescope at Palomar Observatory was put into operation \\citep{zwicky_first_systematic_search}. A first overall SN list was published by \\citet{zwicky_first_catalogue} and contained 54 SNe discovered between 1885 and 1956. This Palomar Supernova Master List has been updated and occasionally republished \\citep{zwicky1965,kowal1971,sargent1974}. Two other SN lists were published contemporaneously \\citep{karpwicz1968,flin1979}.  A new and revised supernova catalogue compiled from the Palomar Supernova Master List has been published by \\citet{barbon1984}. This catalogue then became the Asiago SN catalogue \\citep[ASC;][]{barbon1989,barbon1999} and has received great recognition so far. In 1993 another catalogue, the Sternberg Astronomical Institute (SAI) SN catalogue (SSC) was first published \\citep{SSC1993,SSC2004}.  Today, running SN catalogues are easily accessible over the internet. The most important ones are the list of SNe maintained by the Central Bureau for Astronomical Telegrams (CBAT)\\footnote{http://www.cfa.harvard.edu/iau/lists/Supernovae.html}, the electronic version of the ASC\\footnote{http://web.oapd.inaf.it/supern/cat/} and the electronic version of the SSC\\footnote{http://www.sai.msu.su/sn/sncat/}. The latter two can also be accessed via VizieR\\footnote{http://vizier.u-strasbg.fr/cgi-bin/VizieR}.  The recently growing interest in SNe is demonstrated in Fig.~\\ref{fig:snperyear}. The number of detected SNe per year has grown almost exponentially after the discovery of \\object{SN1987A}, reaching a current discovery rate of a few hundred SNe per year. This gain can be mostly attributed to distant SNe, because the discovery rate of bright SNe has not significantly increased.  The need for a new supernova catalogue has arisen during a search for high-energetic neutrinos from young SNe with the AMANDA neutrino telescope \\citep{arxiv,Diplomarbeit}. In order to enhance the sensitivity of the analysis, the SNe were stacked (i.e. adding up small signals that may not be significant individually). Nearby SNe contribute most to the expected signal, but the distance derived from the redshift of the SN host galaxy gives only a very rough estimate. Hence, a new catalogue was compiled that includes journal-refereed distances of the host galaxies and therefore allows a more realistic signal estimate.  \\begin{figure}[t] \\includegraphics*[width=\\columnwidth]{SN_discoveries_per_year.pdf} \\caption{Number of detected SNe per year. The discovery year of \\object{SN1987A} is marked with a dashed line. The fraction of bright SNe, which have a magnitude at maximum $ < 15^{\\rm{m}}$, is indicated by the hatched area. Selection criteria for the data are described in Sect. \\ref{sec:studies}.} \\label{fig:snperyear} \\end{figure}  A second motivation for a new catalogue was that a comparison between the three SN catalogues (CBAT, ASC, SSC) available online revealed significant inconsistencies in the listed information. In some cases these can be attributed to simple typographical errors, while others are qualitative differences in the given information. The concept of the new catalogue is to list the undisputed information and flag differences. Therefore, the new catalogue also serves as a meta-catalogue of the current online SN catalogues. A subset of high-quality SNe with more reliable information can be selected with the meta-data. ",
        "conclusions": "In this paper a new unified catalogue of three existing supernova catalogues is presented. During the unification procedure several inconsistencies between the catalogues were identified and errors corrected. Remaining inconsistencies are transparently marked and enable the user to select high-quality subsamples. Wherever possible, redshift-independent distance estimates were added to provide a more realistic distance than the redshift only. Several examples for statistical studies that make use of the USC capabilities are shown.  The corrected and extended information contained in this unified catalogue is intended to improve the use of SN-related information, e.g. in SN-related analyses in astro-particle physics. An example is \\citet{arxiv} which, initiated this work."
    },
    "1112/1112.1845_arXiv.txt": {
        "abstract": "{As of today, over 50 planetary systems have been discovered in binary star systems, some of which have binary separations that are smaller than $20$\\,AU. In these systems the gravitational forces from the binary have a strong influence on the evolution of the protoplanetary disk and hence the planet formation process.} {We study the evolution of viscous and radiative circumstellar disks under the influence of a companion star. We focus on the eccentric $\\gamma$~Cephei and $\\alpha$~Centauri system as examples and compare disk quantities such as disk eccentricity and precession rate to previous isothermal simulations.} {We performed two-dimensional hydrodynamical simulations of the binary star systems under the assumption of coplanarity of the disk, host star and binary companion. We used the grid-based, staggered mesh code \\texttt{FARGO} with an additional energy equation to which we added radiative cooling based on opacity tables.} {The eccentric binary companion perturbs the disk around the primary star periodically. Upon passing periastron, spirals arms are induced that wind from the outer disk towards the star. In isothermal simulations this results in disk eccentricities up to $ e_\\mathrm{disk} \\approx 0.2 $, but in more realistic radiative models we obtain much smaller eccentricities of about $ e_\\mathrm{disk} \\approx 0.04 - 0.06 $ with no real precession. Models with varying viscosity and disk mass indicate that disks with less mass have lower temperatures and higher disk eccentricity.} {The fairly high disk eccentricities, as indicated in previous isothermal disk simulations, implied a more difficult planet formation in the $\\gamma$~Cephei system caused by the enhanced collision velocities of planetesimals. We have shown that under more realistic conditions with radiative cooling the disk becomes less eccentric and thus planet formation may be made easier. However, we estimate that the viscosity in the disk has to very low, with $\\alpha \\lesssim 0.001$, because otherwise the disk's lifetime will be too short to allow planet formation to occur along the core instability scenario. We estimate that the periodic heating of the disk in eccentric binaries will be observable in the mid-IR regime.} ",
        "introduction": "Presently, about 50 planets are known to reside in binary stars systems. In all systems with solar-type stars the planet orbits one of the stars while the secondary star acts as a perturber, i.e. these are in a so-called S-type configuration \\citep{1986A&A...167..379D}. Recently, planets seem to have been detected in evolved binary star systems that are in a circumbinary (P-type) configuration  \\citep[e.g.][]{2011A&A...526A..53B}. The main observational characteristic of the known planets in binary stars system have been summarized by \\citet{2004A&A...417..353E,2006ApJ...646..523R,2007A&A...462..345D}. Most planets are in binaries with very large separations with semi-major axis beyond 100\\,AU, in particular when detected by direct imaging. Observationally, there is evidence for fewer planets in binaries with a separation of less than about 100\\,AU \\citep{2007A&A...474..273E}, in accordance with the expectation that binarity constitutes a challenge to the planet formation process. Despite this, there are several systems with a quite close binary separation such as $\\gamma$~Cephei \\citep{1988ApJ...331..902C}, Gliese\\,86 \\citep{2000A&A...354...99Q,2001A&A...370L...1E}, HD\\,41004 \\citep{2004A&A...426..695Z}, and HD\\,196885 \\citep{2008A&A...479..271C}.  As known from several theoretical studies, the presence of the secondary renders the planet formation process more difficult than around single stars. Owing to the gravitational disturbance, the protoplanetary disk is hotter and dynamically more excited such that the coagulation and growth process of planetesimals as well as the gravitational instability process in the early phase of planet formation are hindered \\citep{2000ApJ...537L..65N,2004A&A...427.1097T,2006Icar..183..193T}. This is particularly true for the mentioned closer binaries with an orbital separation of about $20$\\,AU. Hence this theoretical challenge has put these tighter binaries into the focus of studies on planet formation in binary stars \\citep{2007ApJ...660..807Q,2008MNRAS.386..973P,2008A&A...486..617K} and on their dynamical stability \\citep{2004RMxAC..21..222D,2006ApJ...644..543H}. The most famous example in this category is the $\\gamma$~Cephei system, which has been known for over 20 years to contain a protoplanet \\citep{1988ApJ...331..902C,2003ApJ...599.1383H}. In recent years the orbital parameter of this system have been updated  \\citep{2007A&A...462..777N,2007ApJ...654.1095T} and the basic binary parameters are quoted in Table~\\ref{tab:standardmodel} (below) because $\\gamma$~Cephei is our standard model. At the end of the paper we will briefly discuss the $\\alpha$~Centauri system as well.  Because planet formation occurs in disks, their dynamical structure is of crucial importance in estimating the efficiency of planetary growth processes. The prime effect of the secondary star is the truncation of the disk owing to tidal torques \\citep{1994ApJ...421..651A}. The truncation radius depends on the mass ratio of the binary, its eccentricity and the viscosity in the disk. An additional effect, namely the excitation of eccentric modes in the disk, which has been noticed previously in close binary stars \\citep{1988MNRAS.232...35W,1991ApJ...381..259L,2008A&A...487..671K}, has recently drawn attention in studies of planet harboring binary stars as well \\citep{2008MNRAS.386..973P,2008A&A...486..617K}. In these studies it was shown that despite the high eccentricity of a $\\gamma$~Cephei type binary the disk became eccentric with an average eccentricity of about $e_{\\rm disk} \\approx 0.12$ and a coherent disk precession  \\citep{2008A&A...486..617K}. However, this leaves the question of numerical effects \\citep{2008MNRAS.386..973P}. Subsequently the influence of self-gravity of the disk has been analyzed by \\citet{2009A&A...508.1493M}, who concluded that the inclusion of self-gravity leads to disks with lower eccentricity on average.  Interestingly, there may be observational evidence for tidal interactions between a companion on an eccentric orbit and a circumstellar disk around the primary star \\citep{2010A&A...517A..16V}, as inferred from variable brightness on longer timescales. The possibility to detect an eccentric disk through spectral line observations has been analyzed recently by \\citet{2011A&A...528A..93R}.  In all simulations mentioned above, the disk has typically been modeled using a fixed radial temperature distribution. This simplifies the numerics such that no energy equation has to be solved but of course it lacks physical reality. In this work we will extend previous studies and consider a much improved treatment of the energy equation. We will do so within the two-dimensional, flat-disk approximation following \\citet{2003ApJ...599..548D} and \\citet{2008A&A...487L...9K}. Here, the viscous heating of the disk and the radiative losses are taken into account. We will study the structure and dynamical evolution of disks with different masses and binary parameter, and will analyze the influence of numerical aspects, in particular boundary conditions. Our first target of interest will be the well-studied system $\\gamma$~Cephei and we will present results on the $\\alpha$~Centauri system subsequently. ",
        "conclusions": "We have investigated the dynamics of a protostellar disk in binary star systems using specifically the orbital parameter of $\\gamma$~Cephei and $\\alpha$~Centauri. We assumed a coplanar system and used a two-dimensional hydrodynamical code to evolve the non-self-gravitating disk. Extending previous simulations, we included internal viscous heating given by an $\\alpha$-type viscosity prescription, and radiative cooling from the disk surface.  In a first set of simulations we investigated locally \\textit{isothermal} disks for different disk temperatures. We showed that disks in binaries of the $\\gamma$~Cephei type with a standard thickness of $H/r = 0.05$ become eccentric ($e_{\\rm disk} \\approx 0.2$) showing a coherent disk precession. This agrees well with previous simulations by \\citet{2008A&A...486..617K} and \\citet{2008MNRAS.386..973P}. Varying the temperature in the disk, we showed that the magnitude of the disk's eccentricity becomes lower when the disk thickness increases. For disks with $H/r \\gtrsim 0.065$ the mean average eccentricity has dropped below 0.08 and the disks do not show a precession anymore.  Then we studied more realistic disks with internal heating and radiative cooling, varying the disk's mass. In all cases we found relatively low eccentricities and no precession. We attribute the lack of eccentricities firstly to the increased disk height, which is, in particular for the more massive disks, higher than the standard value (see Fig.~\\ref{fig:gc-hr}). Secondly, in the full radiative models the disk's dynamical behavior is more adiabatic compared to the locally isothermal case. Then, through compressional heating ($p dV$-work), kinetic energy is transferred to internal energy, which leads to a reduced growth of disk eccentricity. We have checked that purely adiabatic models show an even lower disk eccentricity than the radiative models. Hence, the radiative case lies between the adiabatic and isothermal, as expected.  Because the disk's energy balance is determined via the viscosity, we changed the value of the parameter $\\alpha$ ranging from $0.005$ to $0.04$, all values that are consistent with the results of MHD-turbulent accretion disks. Here, we found that only the disk with the highest $\\alpha$ becomes eccentric. The reason for this rise is the larger disk radius, which leads to an enhancement of the tidal torques from the secondary. We note that the disk's outer radius in our models still lies well inside the 3:1 resonance with the binary. According to the linear instability model by \\citet{1991ApJ...381..259L}, the disk eccentricity is excited through the 3:1 resonance and hence, the disk should be sufficiently large, a condition which is fulfilled only for small mass ratios, $q=M_\\mathrm{secondary}/M_\\mathrm{primary}$. However, as shown by \\citet{2008A&A...486..617K}, disks in binary star systems with large mass ratios can turn eccentric as well, even though the disks are small, a feature confirmed in our simulations.  The inferred short lifetime of disks with standard viscosities is slightly alarming with respect to planet formation in these systems. In the core-accretion scenario planet formation proceeds along a sequence of many steps that take a few Myr. For disks to persist this long in $\\gamma$~Cephei-type binaries a very low viscosity of $\\alpha \\lesssim 10^{-4}$ seems to be required. In the gravitational instability scenario the timescale for planet formation is much shorter and hence, this scenario may be favored by our findings. Observationally, several recent studies indeed suggest that the lifetime of disks in young binary stars is significantly reduced compared to disks around single stars \\citep{2009ApJ...696L..84C,2010ApJ...709L.114D,2011arXiv1109.4141K}. Dynamically, this behavior is expected, because the perturbation of the companion star leads quite naturally to an increased mass loss of the disk, details depending on the binary separation and eccentricity.  A very critical phase in the core accretion scenario, in particular in binary stars, is the initial growth of meter to km-sized planetesimals. Here, the growth depends on the successful sticking of the two collision partners. Since the relative velocity of the bodies is increased in binary stars, planetary growth will be significantly hindered by the presence of a companion, see e.g. \\citet{2006Icar..183..193T,2011CeMDA.tmp...40T} and references therein. Here our results indicate that planetesimal growth is less negatively influenced because the disk eccentricity is reduced for more realistic radiative disks. As has been shown, eccentric disks tend to increase the mutual relative velocities of embedded objects, in particular of different sizes, because of the misaligned periastrons of the particles \\citep{2007arXiv0705.3421K,2008MNRAS.386..973P}. Hence, a radiative disk with a low viscosity could help to promote planetesimal growth. However, it remains to be seen how the inclusion of stellar irradiation (from both stars) influences the dynamics. Owing to the additional heating of the disk, we expect even more mass loss from the system and possibly higher disk eccentricities because the disks are more isothermal and will have a larger radius.  Previous studies have indicated that a in mutually inclined system planetesimal growth may be enhanced because planetesimals can be size-sorted in differently inclined planes \\citep{2009ApJ...698.2066X,2009A&A...507..505M}. However, recently \\citet{2011A&A...528A..40F} showed through full 3D hydrodynamical studies that for inclined binaries the relative velocities of different sized planetesimals increases through inclinations effects. They conclude that for inclined systems planetesimal formations can take place only for very distant binary stars with $a_{\\rm bin} \\gtrsim 60$~AU. However, the simulations considered only isothermal disks, and it remains to be seen how radiative effects influence the disk. But full 3D radiative simulations are still beyond the present computational possibilities, because thousands of orbital timescales of the disk will have to be calculated."
    },
    "1112/1112.1074_arXiv.txt": {
        "abstract": "{ Every massive globular cluster (GC) is expected to harbour a significant population of cataclysmic variables (CVs). In this review, I first explain why GC CVs matter astrophysically, how many and what types are theoretically predicted to exist and what observational tools we can use to discover, confirm and study them. I then take a look at how theoretical predictions and observed samples actually stack up to date. In the process, I also reconsider the evidence for two widely held ideas about CVs in GCs: (i) that there must be many fewer {\\em dwarf novae} than expected; (ii) that the incidence of magnetic CVs is much higher in GCs than in the Galactic field.} ",
        "introduction": "Globular clusters (GCs) are old, gravitationally bound stellar systems that typically contain $\\sim 10^6$ stars. Some of these stars are binaries, and some of these binaries are cataclysmic variables (CVs), i.e. systems in which a white dwarf (WD) accretes material from a roughly main sequence (MS) companion. These CV populations in GCs deserve special attention for at least three reasons:  {\\em 1 The Globular Cluster Perspective:} the late dynamical evolution of a GC is thought to be driven largely by its close binary (CB) population (e.g. Hut et al. 1992). However, the dominant non-interacting MS-MS binaries are difficult to detect and study in GCs. CVs can therefore be used as convenient tracers of the underlying CB populations for studies of GC dynamics and evolution.  {\\em 2 The Cataclysmic Variable Perspective:} in principle, GCs can provide us with sizeable samples of CVs at known distance and (to some extent) age. Such samples might allow critical tests of theoretical CV/binary evolution scenarios. An interesting complication is that not all CVs in GCs are likely to have formed and evolved in isolation: many are likely to have been produced, or at least affected, by dynamical interactions (see later).  {\\em 3 The Supernova Perspective:} it has been suggested that GCs might be significant Type Ia Supernova factories (Shara \\& Hurley 2002), at least in elliptical galaxies (Ivanova et al. 2006). All SN Ia progenitor populations are thought to be close relatives of CVs (e.g. double WD systems, supersoft sources, WD + red giant binaries), so CVs can again serve as a useful tracer population for these progenitors. In fact, it is even still possible that some CVs might be SN Ia progenitors themselves (Thoroughgood et al. 2001; Zorotovic, Schreiber \\& G\\\"{a}nsicke 2011). ",
        "conclusions": "I hope to have shown that GC CVs should be of great interest to anybody studying the dynamical evolution of clusters, the binary evolution of CVs or the SN Ia progenitor problem. I also hope I have made it clear that great strides have been taken in recent years in finding and understanding GC CV populations, thanks mainly to the availability of Chandra and HST. Indeed, we are finally discovering significant numbers of CVs in any given massive GC, though still not the hundreds that are theoretically predicted to lurk there. I have also discussed the theoretical and observational hints that field and GC CV populations may differ systematically, although I feel that the jury is still out regarding most of these putative differences.  The next big advances in the field are likely to come from one or both of the following directions. First, if we really wish to test for the presence of the predicted hundreds of CVs per cluster, we have to be sensitive to {\\em faint} CVs. How deep? Ideally, deep enough to discover systems like WZ Sge, at $M_{V} \\sim 13$ in quiescence. This is hard -- it requires reaching $m_v \\sim 27$ or so. But it is not impossible: Cohn et al. (2010) have essentially done this in NGC~6397 and uncovered two likely WZ-Sge-like CVs, the first such systems in any GC. An interesting short-cut might be to place limits on the number of faint CVs in GCs by considering the {\\em integrated} X-ray luminosity of individually undetected systems (e.g. Haggard, Cool \\& Davies 2009). Second, we need to obtain orbital periods for a significant sample of GC CVs, so that we can start to make meaningful comparisons to theoretically predicted period distributions."
    },
    "1112/1112.4651_arXiv.txt": {
        "abstract": "We study the impact of the late time dynamical evolution of ejecta from core-collapse supernovae on $\\nu p$-process nucleosynthesis. Our results are based on hydrodynamical simulations of neutrino wind ejecta.  Motivated by recent two-dimensional wind simulations, we vary the dynamical evolution during the $\\nu p$-process and show that final abundances strongly depend on the temperature evolution. When the expansion is very fast, there is not enough time for antineutrino absorption on protons to produce enough neutrons to overcome the $\\beta^+$-decay waiting points and no heavy elements beyond $A=64$ are produced.  The wind termination shock or reverse shock dramatically reduces the expansion speed of the ejecta.  This extends the period during which matter remains at relatively high temperatures and is exposed to high neutrino fluxes, thus allowing for further $(p,\\gamma)$ and $(n,p)$ reactions to occur and to synthesize elements beyond iron.  We find that the $\\nu p$-process starts to efficiently produce heavy elements only when the temperature drops below $\\sim 3$~GK.  At higher temperatures, due to the low alpha separation energy of $^{60}$Zn ($S_{\\alpha} = 2.7$~MeV) the reaction $^{59}$Cu$(p,\\alpha)$$^{56}$Ni is faster than the reaction $^{59}$Cu$(p,\\gamma)$$^{60}$Zn.  This results in the closed NiCu cycle that we identify and discuss here for the first time.  We also investigate the late phase of the $\\nu p$-process when the temperatures become too low to maintain proton captures.  Depending on the late neutron density, the evolution to stability is dominated by $\\beta^+$ decays or by $(n,\\gamma)$ reactions.  In the latter case, the matter flow can even reach the neutron-rich side of stability and the isotopic composition of a given element is then dominated by neutron-rich isotopes. ",
        "introduction": "\\label{sec:intro}  Neutrino-driven winds from core-collapse supernova explosions contribute to the synthesis of elements beyond iron. After the explosion, the hot proto-neutron star cools emitting neutrinos.  These neutrinos interact with the stellar matter and deposit energy in the outer layers of the proto-neutron star leading to a supersonic outflow known as neutrino-driven wind \\cite[]{duncan.shapiro.wasserman:1986}. Although neutrino-driven winds were considered the site where heavy elements are produced by the r-process \\cite[]{Woosley94}, recent simulations \\cite[]{arcones.janka.scheck:2007, Huedepohl.etal:2010, Fischer.etal:2010, Roberts.etal:2010} cannot reproduce the extreme conditions required for producing heavy r-process elements \\cite[see e.g.,][]{hoffman.woosley.qian:1997, Otsuki.Tagoshi.ea:2000, Thompson.Burrows.Meyer:2001}. The wind entropy is too low (less than $100\\,k_{\\mathrm{B}}/\\mathrm{nuc}$) and, even more significant, the ejecta is proton rich (the electron fraction $Y_e$ remains above 0.5 during seconds, see \\cite{Huedepohl.etal:2010, Fischer.etal:2010}). Even if the r-process does not take place in every neutrino-driven wind, lighter heavy elements (e.g., Sr, Y, Zr) can be synthesized in this environment as suggested by \\cite{Qian.Wasserburg:2001}. In proton-rich conditions \\cite{Froehlich06} showed that elements beyond $^{64}$Ge can be synthesized. \\cite{Wanajo.Janka.Mueller:2011} found Sr, Y, Zr in small pockets of neutron rich material ejected after the explosion of low mass progenitors. Recently, \\cite{Arcones.Montes:2011} performed a systematic nucleosynthesis study that strongly supports the production of lighter heavy elements in proton- and neutron-rich neutrino-driven winds.  In proton-rich winds, charged particle reactions (alpha and proton captures) build nuclei up to $^{56}$Ni and even up to $^{64}$Ge once the temperature drops below 3~GK. Due to their long beta-decay lifetimes and low thresholds for proton capture, the nuclei $^{56}$Ni and $^{64}$Ge act as bottlenecks that inhibit the production of heavier elements. In the $\\nu p$-process, their decay is sped up by $(n,p)$ reactions, with the neutrons produced by antineutrino absorption on the abundant free protons. This allows for the production of elements beyond iron and may explain the origin of light p-nuclei \\cite[]{Froehlich06, Pruet.Hoffman.ea:2006, Wanajo:2006, Wanajo.etal:2011}.  The synthesis of elements by the $\\nu p$-process depends thus on neutrino spectra and luminosities but also on the dynamical evolution as matter expands through the slow, early supernova ejecta.  This produces a wind termination shock or reverse shock where kinetic energy is transformed into internal energy \\cite[]{arcones.janka.scheck:2007}. The reverse shock has a big impact on the nucleosynthesis because temperature and density increase and the expansion is strongly decelerated. This hydrodynamical feature has been extensively studied for r-process nucleosynthesis \\cite[]{Qian.Woosley:1996, Sumiyoshi00, Wanajo:2007, Panov.Janka:2009, Arcones.MartinezPinedo:2011}. Recently, \\cite{Wanajo.etal:2011} have also explored the relevance of the reverse shock on the $\\nu p$-process. Motivated by their work and by new 2D hydrodynamical simulations of the neutrino-driven wind \\cite[]{Arcones.Janka:2011}, we investigate here the wind termination shock to gain further insights on the dynamical evolution relevant for the $\\nu p$-process.   Here we use a trajectory from hydrodynamical simulations (Sect.~\\ref{sec:dynamics}) combined with a complete nucleosynthesis network (Sect.~\\ref{sec:netw}).  In Sect.~\\ref{sec:results} we present our results where we analyze the impact of the wind termination temperature (Sect.~\\ref{sec:constT}), the temperature jump at the reverse shock (Sect. \\ref{sec:rs}), and the influence of late temperature evolution and the decay to stability (Sect~\\ref{sec:long}).  Our conclusions are summarized in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have studied the impact of the supernova dynamical evolution on the $\\nu p$-process and how the wind termination affects the synthesis of elements beyond iron. For a robust and strong production of nuclei with $64 \\lesssim A \\lesssim 110$ there is an optimal wind termination temperature at $\\approx$2~GK as it was found by \\cite{Wanajo.etal:2011}. Also in agreement with their work, we have shown that if the wind termination occurs at low temperatures ($T_{\\mathrm{wt}}\\approx 1$~GK), matter stays too short time in the optimal temperature range for $\\nu p$-process nucleosynthesis, $\\sim 1 -2 $~GK. Moreover, the electron antineutrino flux quickly becomes too small to produce the necessary neutrons to overcome the $\\beta^{+}$-decay waiting points. This hinders the $\\nu p$-process and consequently the efficient synthesis of heavy nuclei.  Therefore, the wind termination temperature determines the heaviest elements produced.  We have identified an end point cycle that is key at high temperatures (around 3~GK). The close NiCu cycle inhibits the production of elements heavier than 56Ni. The reactions in this closed NiCu cycle are the following: \\begin{eqnarray} && ^{56}Ni (n,p) ^{56}Co (p, \\gamma ) ^{57}Ni (n,p) ^{57}Co (p, \\gamma ) ^{58}Ni \\nonumber \\\\ & &\\longrightarrow \\, ^{58}Ni (p, \\gamma ) ^{59} Cu(p, \\alpha ) ^{56} Ni \\, . \\nonumber \\end{eqnarray} At high temperatures the reaction $^{59}$Cu$(p,\\alpha)$ $^{56}$Ni prevents the synthesis of elements beyond the iron group. When the temperature drops below $\\approx 3$~GK, $^{59}$Cu$(p,\\gamma)$ $^{60}$Zn becomes more effective than $^{59}$Cu$(p,\\alpha)$ $^{56}$Ni.  This leads to breakout from the NiCu cycle and the flow of matter can continue towards heavier nuclei.  The temperature dependence of these two reactions is critical because it sets the temperature at which the $\\nu p$-process can start to synthesize elements beyond Nickel. If this cross-over temperature were slightly higher, matter would be closer to the proto-neutron star and thus under higher neutrino flux when the path starts to move towards heavier nuclei. This will significantly increase the efficiency of the $\\nu p$-process. Therefore, our results clearly motivate further investigation of these two key reactions: $^{59}$Cu$(p,\\alpha)$ $^{56}$Ni and $^{59}$Cu$(p,\\gamma)$$^{60}$Zn. Particularly relevant are the branching ratios for the decay by alpha and gamma emission of compound states in $^{60}$Zn above the proton separation energy.   We have also explored for the first time the impact of the dynamical evolution after the wind termination on the synthesis of heavy elements by the $\\nu p$-process. Hydrodynamical simulations show that after the wind termination there is a transition from an expansion with almost constant temperature and density to a phase with almost constant velocity. The duration of the constant temperature phase, $\\Delta t$, strongly affects the final abundances of heavy nuclei. We have found that for $\\Delta t=0.0$~s the final flow of matter to stability occurs by neutron capture reactions, i.e\\ $(n,\\gamma)$ and $(n,p)$, and not by $\\beta^+$-decays.  When the phase of constant temperature is very short ($\\Delta t < 0.5$~s), the abundances of nuclei with $64<A<110$ are lower.  This leads to significantly higher neutron densities at later times.  Depending on the neutron density the matter moves to stability either by $\\beta^+$-decays ($\\Delta t = 1.0$~s) or by neutron captures.  These two different ways of reaching stability leave a distinct fingerprint in the final abundances.   We have investigated in detail the impact of dynamical evolution on the $\\nu p$-process nucleosynthesis using individual characteristic trajectories. In order to predict the complete $\\nu p$-process yields from a supernova simulation, one will need to integrate over all proton-rich ejecta.  In summary, the supernova dynamics as well as individual reactions such as the proton capture reactions on $^{59}$Cu determine how high in proton number the $\\nu p$-process can proceed. The dynamical evolution just after the wind termination can greatly affect the nucleosynthesis evolution towards stability as it determines the late neutron density.  Our results provide a link between nucleosynthesis in proton-rich winds and the dynamical evolution of the ejected matter and motivate further theoretical and experimental effort on understanding key reactions."
    },
    "1112/1112.1418_arXiv.txt": {
        "abstract": "The warm-hot intergalactic medium (WHIM) at temperatures $10^5-10^7$ K is believed to contain 30-50\\% of the baryons in the local universe. However, all current X-ray detections of the WHIM at redshifts $z>0$ are of low statistical significance ($\\lsim3\\sigma$) and/or controversial.  In this work, we aim to establish the detection limits of current X-ray observatories and explore requirements for next-generation X-ray telescopes for studying the WHIM through X-ray absorption lines. We analyze all available grating observations of Mrk~421 and obtain spectra with signal-to-noise ratio (S/N) of $\\sim90$ and 190 per 50 m\\AA\\ spectral bin from \\chandra\\ and \\xmm\\ observations, respectively. Although these spectra are two of the best ever collected with \\chandra\\ and \\xmm, we cannot confirm the two WHIM systems reported by Nicastro et al. in 2005. Our bootstrap simulations indicate that spectra with such high S/N can\\emph{not} constrain the WHIM with \\ovii\\ column densities $N_{\\rm OVII}\\approx10^{15}~{\\rm cm^{-2}}$ (corresponding to an equivalent widths of 2.5 m\\AA\\ for a Doppler velocity of $50~{\\rm km~s^{-1}}$) at $\\gsim3\\sigma$ significance level. The simulation results also suggest that it would take $>60$ Ms for \\chandra\\ and 140 Ms for \\xmm\\ to measure the $N_{\\rm OVII}$ at $\\ge4\\sigma$ from a spectrum of a background QSO with flux of $\\sim0.2$ mCrab (1 Crab = $2\\times10^{-8}~{\\rm erg~s^{-1}~cm^{-2}}$ at 0.5-2 keV). Future X-ray spectrographs need to be equipped with spectral resolution $R\\sim4000$ and effective area $A\\ge100~{\\rm cm^2}$ to accomplish the similar constraints with an exposure time of $\\sim2$ Ms and would require $\\sim11$ Ms to survey the 15 QSOs with flux $\\gsim0.2$ mCrab along which clear intergalactic \\ovi\\ absorbers have been detected. ",
        "introduction": "\\label{sec:intro}  Identifying the ``missing baryons'' is one of the major tasks of modern cosmology. Observations of the microwave background (e.g., \\citealt{kom11}) and the big bang nucleosythesis model combined with the measurement of deuterium abundance (e.g., \\citealt{bur01, ome06, pet08}) are converging on the cosmological baryon density of $\\Omega_b=0.0455\\pm0.0028$. At high redshift ($z>3$) universe, baryons exist primarily in the form of cool, photoionized intergalactic medium (IGM) that is traced by Ly$\\alpha$ absorption forest lines \\citep{rau97}. In the present-day universe, matter detected in forms of photoionized IGM, stars, galaxies, intracluster medium, etc., adds up to only $\\sim50\\%$ of the baryons \\citep{shu11}, % leaving much of the inventory yet to be found \\citep{fuk98, fuk04}. Cosmological numerical simulations for large-structure formation indicate that the missing baryons are still in the IGM, but at current epoch they have been heated by gravitational shocks and galactic feedback to temperatures of $10^5-10^7$ K when they fall into the gravitational potential wells of the dark matter cosmic web filaments (e.g., \\citealt{cen99, cen06, dav01}). The IGM at these temperatures, so-called warm-hot intergalactic medium (WHIM), mainly absorbs and emits photons in the ultraviolet (UV) and X-ray wavelength bands.  Searching for the missing baryons has been conducted extensively in the UV band. Indeed, besides the cool phase Ly$\\alpha$ absorbers (e.g., \\citealt{wey98, pen00, pen04}), the absorption lines of \\ciii, \\civ, \\nv, \\ovi, \\siiii, and \\siiv\\ have been routinely observed in spectra of background quasi-stellar objects (QSOs; e.g., \\citealt{tri00, tri08, pro04, leh06, dan05, dan08, dan06, thom08, yao11}), suggesting that a significant amount of baryons exist in the highly ionized absorbers. These high ions, \\ovi\\ in particular,  are believed to trace the WHIM at the low-temperature end ($\\sim10^5-10^6$ K). However, all these ions can also be produced through photoionization in the intergalactic environment (\\citealt{opp09}; but see also \\citealt{tep11} and \\citealt{smi12} for a contrary view) and thus some of them may not contribute to the canonical WHIM (i.e., the shock-heated IGM expected from simulations). How many of these absorbers originate in the real WHIM is still under debate \\citep{dan08, tri08}. \\neviii\\ and broad \\hi\\ Ly$\\alpha$ absorbers (BLAs) are complementary and valuable tracers of hot absorbing WHIM in UV bands. However, there are only a few detections with marginal significance of \\neviii\\ \\citep{sav05, nar09, nar11} and detections of BLAs require very high signal-to-noise-ratio (S/N) of spectra to resolve broad and weak signals from the continua \\citep{ric06, leh07, sav11, dan10, dan11}. Converting measurement of these UV absorbers to census of baryonic matter also depends on ionization fractions and metallicity of the IGM, which in many absorbers are still poorly known. Nevertheless, BLAs and \\ovi-absorbers, although partially overlapped with photoionized Ly$\\alpha$ absorbers, are estimated to contribute an additional $\\sim25\\%$ to the baryon inventory at the current epoch (e.g., \\citealt{dan08, dan11, shu11}).  X-ray observations of the intergalactic \\ovii\\ and \\oviii\\ emission/absorption features could provide essential information for establishing the existence of the WHIM and completing the baryon inventory in the local universe. At temperatures $T\\gsim10^6$ K, hydrogen and helium are nearly completely ionized \\citep{gnat07}. Without metals, the thermal gas emits/absorbs photons through bremsstrahlung, which unfortunately is optically thin even at intergalactic scales \\citep{bre07}. Because the K-shell transitions of all elements heavier than lithium and the L-shell transitions of elements heavier than iron lie in X-ray bands \\citep{pae03}, metals can greatly increase the emissivity of the gas. However, because of the density-squared dependence, the WHIM emission is expected to be weak and difficult to detect. In fact, the IGM emission has been directly observed only from the dense regions like intracluster and intragroup media (e.g., \\citealt{wang97, mul00, all02, sun09}). The WHIM emission signals were also claimed in the angular correlation of the diffuse X-ray background \\citep{sol99, urs06}, but disentangling the real WHIM signal from the unresolved point sources and local Galactic hot diffuse gas has been a major uncertainty (e.g., \\citealt{gal07}).  Unlike emission lines, absorption lines measure the column densities and thus directly sample the total mass of the intervening gas along a line of sight. Oxygen is the most abundant metal element. In contrast of \\ovi\\ as a minority ionization state, \\ovii\\ is the most abundant ion for a gas at temperatures of $T\\sim10^{5.5}-10^{6.3}$ K and \\oviii\\ takes over when $T\\gsim10^{6.3}$ K \\citep{gnat07}. Thus they have advantages over \\ovi\\ in estimating the total baryon contained in highly ionized absorbers. \\ovii\\ and \\oviii, whose production ionization potentials are 138.1 eV and 739.3 eV compared to 113.9 eV of \\ovi, are also hard to produce through photoionization in the IGM environment (e.g., \\citealt{cen06a, chen03}). Theoretical calculations indicate that the \\ovii\\ column density is about 10 times higher than \\ovi\\ in a shock-heated gas (i.e., $N_{\\rm OVII}\\gsim10N_{\\rm OVI}$) whereas $N_{\\rm OVII}\\lsim3N_{\\rm OVI}$ in a photoionized gas \\citep{fur05}. Therefore observations of the IGM \\ovii\\ and \\oviii\\ absorption lines are crucial, not only for constraining the properties of the WHIM at high temperature end but also for probing the nature of the commonly observed \\ovi-absorbers.  Unfortunately, most attempts at searching the X-ray WHIM absorption lines have been frustrated. \\ovii\\ and \\oviii\\ absorption lines at $z=0$ have been unambiguously detected in many background QSOs. However, all these X-ray absorptions are \\emph{inconsistent} with an intergalactic origin \\citep{fang06, bre07b, yao10}, but rather can be attributed to the Galactic diffuse hot gas (e.g., \\citealt{yao05, yao07, yao08, yao09, wang05}). Perhaps the most compelling intergalactic result is the detection of \\ovii\\ absorption at the redshift of the large galaxy structure Sculptor Wall ($z_{\\rm abs}=0.03$) along the QSO sight line H~2356-309 ($z_{\\rm QSO}=0.165$), albeit of low significance ($\\sim3\\sigma$; \\citealt{buo09, fang10}). However, because of the high number density of galaxies along the sight line, the absorption may sample the halo gas of one or more intervening galaxies with small impact distances ($\\lsim50$ kpc; \\citealt{wil10}), i.e., mimicking the \\ovii\\ and \\oviii\\ absorption lines at $z=0$. Therefore, it may not be representative of the typical WHIM. Furthermore, attempts at searching for the similar absorption features at the redshift of another large structure, Pisces-Cetus along the same sight line, failed \\citep{zap10}. All other claimed WHIM \\ovii\\ and \\oviii\\ absorptions have been highly debated. For instance, the detected \\oviii\\ absorption line at $z=0.0554$ in \\chandra\\ spectrum of PKS~2155-304 \\citep{fang02a, fang07} cannot be confirmed by the \\xmm\\ observations \\citep{cag04}, although at nearby redshift a small galaxy group and the corresponding UV \\ovi\\ and Ly$\\alpha$ absorption lines have been identified \\citep{shu98, shu03}. \\citet{nic05} obtained a spectrum with unprecedented S/N during the burst states of Mrk~421 and reported two WHIM detections at $z=0.011$ and $z=0.027$. Again, observations with \\xmm\\ cannot confirm the detections and reported significances have also been questioned \\citep{kaa06, ras07}. \\citet{wil10} recently found a galaxy filament at $z=0.027$, which makes the reported \\ovii\\ WHIM detections along the Mrk~421 sight line back to be a subject of discussions.  In this work, we aim to provide detection limits of current X-ray observatories, \\chandra\\ and \\xmm, in measuring the WHIM through X-ray absorption lines, and establish requirements (e.g., effective area and spectral resolution) for next-generation X-ray telescopes. We begin our investigation by scrutinizing the controversial WHIM detections along the Mrk~421 sight line. We extensively explore all the available \\chandra\\ and \\xmm\\ observations, and then run bootstrap simulations to test the reliability of any absorption feature.  The paper is organized as follows. In Section~\\ref{sec:obs} we describes the observations and our data reduction processes, and in Section~\\ref{sec:res} we report the data analysis results and compare them with previous work. In Section~\\ref{sec:sim}, we run bootstrap simulations to explore the detection limits of \\chandra\\ and \\xmm\\ and to establish requirements of next-generation X-ray telescopes. We summarize our results in Section~\\ref{sec:dis}.  Throughout the paper, we adopt the atomic data from \\citet{ver96} and solar abundances from \\citet{asp09}. We conduct our data analysis within {\\sl XSPEC} (version 12.6; \\citealt{arn96}), and report the 1$\\sigma$ confidence range or 3$\\sigma$ limit for a fitting parameter. % We take the flux normalization of $1~{\\rm Crab}=2\\times10^{-8}~{\\rm erg~ s^{-1}~cm^{-2}}$ at 0.5-2.0 keV. Thus 0.2 mCrab is $2.15\\times10^{-13}~{\\rm erg~s^{-1}~cm^{-2}}$~\\AA$^{-1}$ for a flat spectrum in wavelength space, which corresponds to $2.34\\times10^{-4}~{\\rm photons~s^{-1}~cm^{-2}}$~\\AA$^{-1}$ around the \\ovii\\ K$\\alpha$ (21.602 \\AA) at the rest-frame. We define measurement significance level (SL) of an absorption line as \\begin{equation} \\label{equ:sig} SL=\\frac{EW}{\\Delta EW}, \\end{equation} where $EW$ and $\\Delta EW$ are the equivalent width and its $1\\sigma$ uncertainty for an absorption line. ",
        "conclusions": "\\label{sec:res}  Our goal in this section is to examine the WHIM detections at redshifts $z=0.011$ and $z=0.027$. \\citet{nic05} reported the WHIM absorption from transitions of \\neix\\ K$\\alpha$, \\ovii\\ K$\\beta$, \\oviii\\ K$\\alpha$, \\ovii\\ K$\\beta$, \\nvii\\ K$\\alpha$, and \\civ\\ K$\\alpha$ (Table~2 and Figure~8 in \\citealt{nic05}), whose rest-frame wavelengths are 13.447, 18.629, 18.969, 21.602, 28.787, and 33.736 \\AA, respectively \\citep{ver96}. Therefore, we focus our attention on spectral ranges covering these transitions.  Obtaining a good continuum model is crucial to absorption-line measurement. Since all coadded \\chandra\\ spectra contain the contribution from HRC observations, to account for grating order-overlapping issues (Section~\\ref{sec:obs}), we decide to fit a broad range of the spectra from 7 to 40 \\AA, which covers the transitions to be examined. For a ease of comparison of three \\chandra\\ spectra, we jointly fit the Spectra I, II, and III. We first use a Galactic-absorption modified power-law ({\\sl wabs*powerlaw} in xspec) to fit the continuum emission, with neutral hydrogen column density $N_{\\rm H}$ linked together while the power-law index $\\Gamma$ and normalization are allowed to vary. We obtain an unacceptable fit with $\\chi^2=11646$ over 3590 degrees of freedom (DOF), and we find that some ``broad'' features cannot be accounted for by this simple model. These features are mainly due to imperfect instrumental calibration around node boundaries, CCD chip gaps plus dithering effects, the oxygen absorption edge, and imperfect cross calibration between ACIS and HRC observations \\citep{mar04, nic05}. We then use Gaussian profiles to compensate the uncounted calibration residuals. To minimize the effect of the known strong ISM absorption lines on continuum modeling, we include narrow absorption lines (with widths $\\sigma\\lsim50~{\\rm km~s^{-1}}$) of \\oi, \\oii, \\cvi, \\ovi, \\ovii, and \\neix\\ K$\\alpha$ and \\ovii\\ K$\\beta$ transitions at their rest-frame wavelengths \\citep{ver96, jue06, yao09} in our spectral fitting. The centroid energies and widths of these Gaussians are linked together while the normalizations are allowed to vary in the jointly spectral fitting. We find that, besides these narrow absorption lines produced in the ISM, we need 10 broad and one narrow Gaussian profiles to obtain an acceptable fit, and locations of these Gaussian are 10.53, 13.80, 14.60, 17.90, 18.19, 19.02, 19.10, 23.24, 23.31, 23.66, and 29.89 \\AA\\ with widths ($\\sigma$) of 0.61, 2.56, 0.04, 0.00054, 0.17, 1.36, 0.10, 1.45, 1.69, 0.41, and 0.39 \\AA. The final $\\chi^2=4448.2$ with 3512 DOF and the constrained $N_{\\rm H}=1.092\\pm0.06\\times10^{20}~{\\rm cm^{-2}}$, and $\\Gamma=2.131\\pm0.001$, $2.011\\pm0.001$, $2.067\\pm0.001$ for Spectra I, II, and III, respectively. \\emph{Because the absorption line measurement depends only on the \\emph{local} continuum, and because the centroids of these Gaussian profiles do not directly superpose on the wavelengths of the putative WHIM absorption lines, these Gaussian profiles will not have any effect on the line equivalent width measurement conducted below.} \\xmm\\ spectrum (Spectrum IV) does not have the grating-order-overlapping issue, so we use the {\\sl wabs*powerlaw} model to fit the six narrow spectral ranges as plotted in Figure~\\ref{fig:rgs}. We exclude the bad pixels in our spectral analysis, and obtain $\\chi^2/DOF=1.36, 1.49, 1.49, 1.28, 1.75, 1.56$ with DOF of 137, 75, 103, 105, 116 and 104, respectively. Figures~\\ref{fig:chan-Ni}-\\ref{fig:rgs} show these best fit continua with the ISM absorption lines removed from spectral models.  Now, let us examine the existence of the reported WHIM absorption lines. Figures~\\ref{fig:chan-Ni}-\\ref{fig:rgs} show the same portions of \\chandra\\ and \\xmm\\ spectra as shown in Figure~8 of \\citet{nic05}, in which they demonstrated their detections of the WHIM absorption lines at $z=0.011$ and $z=0.027$ from transitions of \\neix\\ K$\\alpha$, \\ovii\\ K$\\beta$, \\oviii\\ K$\\alpha$, \\ovii\\ K$\\beta$,  \\nvii\\ K$\\alpha$, and \\civ\\ K$\\alpha$. In particular, Figure~\\ref{fig:chan-Ni} reproduces their Figure~8, as they were extracted from the same \\chandra\\ observations. Several bad pixels of the RGS accidentally fall around the 21.8 \\AA\\ region and cause an instrument feature (Figure~\\ref{fig:rgs}). Although a careful calibration could yield a correct instrumental response of the RGS, in this work we do not use the \\xmm\\ spectrum to assess the WHIM detection at $z=0.011$.  Visual inspection reveals that none of the above WHIM lines reported by \\citet{nic05} is detected in the previously \\chandra\\ observations they analyzed (Figure~\\ref{fig:chan-Ni}), and in the newly available \\chandra\\ observations (Figure~\\ref{fig:chan-new}), and in the \\xmm\\ observations (Figure~\\ref{fig:rgs}).  The reported WHIM detections could be due to the improper continuum placement on the spectrum. Visual inspection also reveals that there might be an absorption feature at $\\sim21.8$ \\AA\\ in Spectrum I, which corresponds to the reported \\ovii\\ K$\\alpha$ WHIM absorption line at $z=0.011$ (panel (c) in Figure~\\ref{fig:chan-Ni}). However, such an absorption feature is not visible in Spectra II and III (panel (c) in Figures~\\ref{fig:chan-new} and \\ref{fig:chan-all}), and its significance in Spectrum I depends largely on how the local continuum is placed. Please note that our spectral profiles of the local continua were obtained by jointly fitting the Spectra I, II, and III (see above), and the best-fit model seems to account for the general spectral variation reasonably well. To quantitatively assess the existence of the \\ovii\\ K$\\alpha$ WHIM line at $z=0.011$, in the joint analysis of the Spectra I, II, and III, in addition to the 10 Gaussian profiles, we add the 11th ``broad'' Gaussian to the local continuum and another narrower Gaussian with width \\footnote{Different widths yield nearly identical results as long as the line is still unresolved by the instrument. Please see Section~\\ref{sec:exp} for justifications for the line width.} $\\sigma=50~{\\rm km~s^{-1}}$ to represent the WHIM absorption line. Again, we link the line centroids and widths together in our joint fitting but allow the normalizations to vary among the Spectra I, II, and III. The best fit yields a line centroid of 21.90 \\AA\\ and $\\sigma=0.050$ \\AA\\ for the ``broad'' component and a centroid of 21.87 \\AA\\ for the narrow component. The red curves in panel (c) of Figures~\\ref{fig:chan-Ni}-\\ref{fig:chan-all} show the modified local continua. These additional components reduce the $\\chi^2$ of 26.4 by adding nine DOF in total. However, applying the best-fit models to individual spectra yields a $\\chi^2$ change of 13.7, 4.3, and 10.4 by introducing five additional DOF to spectral fitting of Spectrum I, II, and III, respectively. All these \\chandra\\ spectra can be well described by the same continuum profile with various normalizations, which has been proved to be reasonable to other wavelength ranges (panels (a)-(b) and (d)-(f) in Figures~\\ref{fig:chan-Ni}-\\ref{fig:chan-all}). Because we find essentially no improvement in fitting the Spectrum II, adding the ``broad'' Gaussian profile to the local continuum around 21.9 \\AA\\ is statistically unnessary.  We next attempt to measure the equivalent widths (EW) of the putative WHIM absorption lines. Since we have not consistently detected any WHIM absorption of any transitions reported by \\citet{nic05}, we focus here on measuring the upper limit to the equivalent width (EW) of the \\ovii\\ K$\\alpha$ line, which is expected to be the most abundant transition in a broad temperature range of the WHIM (Section~\\ref{sec:intro}). At redshifts $z=0.011$ and $z=0.027$, the rest-frame \\ovii\\ K$\\alpha$ at 21.60 \\AA\\ is shifted to 21.84 and 22.19 \\AA, respectively, and therefore our measurements are obtained from the spectral range described by the continuum colored as blue in panel (c) of Figures~\\ref{fig:chan-Ni}-\\ref{fig:rgs}. To make a fair comparison with the results obtained by \\citet{nic05}, we also make similar measurements by using the continuum colored as red. We should emphasize that the latter measurements are only for comparison purposes, since the additional component to the local continuum that ``amplifies'' the significance of the absorption line at 21.8 \\AA\\ is statistically unnecessary. Our results are reported in Table~\\ref{tab:results}.  In summary, our analysis and results do not support the existence of the two WHIM systems at $z=0.011$ and $z=0.027$. We analyze all the available \\chandra\\ and \\xmm\\ observations of Mrk~421, and have not seen any consistent WHIM absorption of any transition at either $z=0.011$ or $z=0.027$ reported by \\citet{nic05}. We measure the EW of the \\ovii\\ K$\\alpha$ absorption line. For the system at $z=0.027$, we obtain a firm 3$\\sigma$ upper limit of $EW<1.5$ m\\AA\\ from the spectrum (Spectrum II) extracted from the newly available \\chandra\\ observations, in contrast to $EW=2.2\\pm0.8$ m\\AA\\ obtained by Nicastro et al. For the system at $z=0.011$, we confirm the existence of a small dip in the spectrum (Spectrum I) identical to that used by \\citet{nic05}, but the similar spectral feature is absent in Spectrum II. The measurement of the line heavily depends on how the local continuum is placed, and the additional spectral component that amplifies its significance is statistically unnecessary. We obtain an upper limit of $EW<1.5$ m\\AA\\ from the new spectrum, in contrast to $EW = 3.0^{+0.9}_{-0.8}$ m\\AA\\ reported by \\citet{nic05}. The findings in this work are consistent with the previous investigation made by \\citet{kaa06}, in which the authors scrutinized the same data set used by Nicastro \\etal\\ and a subset of calibration observations of \\chandra\\ and \\xmm. Therefore, we conclude that there is no WHIM line detected at either $z=0.011$ or $z=0.027$ along the Mrk~421 sight line.  There are also other important absorption lines in the spectral range of interest (Figures~\\ref{fig:chan-Ni}-\\ref{fig:rgs}). The prominent lines are \\ovii, \\neix, and \\civ\\ K$\\alpha$ and \\ovii\\ K$\\beta$ at $z\\approx0$, and the \\ovi\\ K$\\alpha$ with rest-frame wavelength 22.0 \\AA\\ at $z\\approx0$ is also visible in all four spectra albeit less significant. Again, we only measure the EWs of \\ovii\\ K$\\alpha$ and report them in Table~\\ref{tab:results}. The \\ovii\\ line at $z\\approx0$ can be well explained as absorption from the Galactic diffuse interstellar medium (ISM) \\citep{yao07}, and the \\ovi\\ line is also believed to originate from the ISM \\citep{sav05}. % % It is worth noting that the EW of \\ovii\\ K$\\alpha$ at $z\\approx0$ measured from \\xmm\\ observations is $\\gsim3\\sigma$ larger \\footnote{The differential significance is calculated as \\[|EW_1 - EW_2|/\\sqrt{(\\Delta EW_1)^2 + (\\Delta EW_2)^2}, \\] where $\\Delta EW$ is the $1\\sigma$ error of $EW$ measurement.} than that obtained from \\chandra\\ observations (Table~\\ref{tab:results}). Such a discrepancy in \\ovii\\ has also been measured by \\citet{kaa06}. The causes to this discrepancy are still under investigation, and the results will be published elsewhere. In this paper, we mainly focus on X-ray absorption measurements of the WHIM and will not discuss these non-WHIM measurements any further.   \\label{sec:dis}   1. We analyzed all available (as of 2011 April) grating data and obtained two coadded spectra with S/N$\\sim$90 per 50-m\\AA\\ resolution element from \\chandra\\ observations, and S/N$\\sim$190 from \\xmm\\ observations. Neither \\chandra\\ nor \\xmm\\ observations support the existence of the two WHIM systems previously reported by \\citet{nic05}.  2. We ran bootstrap simulations for the detecting limits of the current X-ray telescopes. We find that the reported EWs of the \\ovii\\ K$\\alpha$ at $z=0.011$ and $z=0.027$ should have been measured at $\\ge3.7\\sigma$ and $\\ge2.3\\sigma$, in contrast to the fact that we only obtained upper limits to the EWs.  3. According to the numerical simulations, the WHIM absorbers with $N_{\\rm OVII} \\gsim10^{15}~{\\rm cm^{-2}}$ could sample $>30-50\\%$ of the \\ovii-bearing baryons. To systematically survey the 15 QSO sight lines along which the IGM \\ovi\\ absorbers have been detected, future X-ray telescopes should be able to facilitate the WHIM study via the X-ray absorption line spectroscopy from background QSOs with fluxes of $\\sim0.2$ mCrab to find $\\ge10$ WHIM systems. It takes impractical long (341 Ms for \\chandra\\ and 804 Ms for \\xmm) exposure time for current X-ray telescopes to complete the survey.  4. It would require $\\sim11$ Ms on-target exposures for future X-ray spectrographs equipped with spectral resolution $R\\gsim4000$ and effective area $A\\gsim100~{\\rm cm^2}$ to finish the survey."
    },
    "1112/1112.0244_arXiv.txt": {
        "abstract": "Ultra High Energy Cosmic Rays (UHECR) at GZK cut off energy ($E\\geq 5.5 \\cdot 10^{19}$ eV) may keep sharp or diffused directionality wherever their composition is made by nucleon or light nuclei. Auger experiment UHECR (2007-2010) did show a mild clustering around Cen A. Two over three of the recent discovered AUGER multiplet (a dozen of events each) tail clustering at twenty EeV are pointing to primary sources very  near the same UHECR crowded Cen A region. These tens EeV tail is aligned with the same UHECR events. We foresaw such possibility as fragment tails of lightest UHECR nuclei. We discuss the relevance of this correlation within a model where UHECR are mostly lightest He like nuclei. UHECR fragment multiplet clustering aligned along higher Cen A events (above  $5.5 \\cdot 10^{19} $ eV  energy)  probe and reinforce our interpretation with an a priori probability to occur below $3 \\cdot 10^{-5}$. ",
        "introduction": "UHECR Astronomy and nuclear composition are more and more in severe conflict. This contradiction were already inherited in  early, apparent, discover of a Super Galactic UHECR correlation \\cite{Auger-Nov07}. In that key article the needed UHECR directionality (assumed nucleon) have been , at the same time, in disagreement with the first observed nuclei composition. Indeed UHECR particle astronomy may rise, suffering however of some directionality spread by the smearing of magnetic field bending along the UHECR flight. This bending is negligible for proton and He, but sever for iron. The bending maybe coherent (just one direction) if the magnetic field is constant or it is random , if the field directions are changing , as it happens inside the galactic plane along and across the galactic arms while pointing to Cen A.   The nucleon, the light nuclei   may keep more or less precise directionality,   a smearing astronomy, while  heaviest ones may exhibit only tiny  anisotropy if  sources are near. In addition to these signals, UHECR in flight are making  fragment secondaries nuclei as well as,  because photo-dissociation, parasite  gamma and neutrino tails somehow correlated. These UHECR compositions are leading to different secondary gamma and neutrino spectra; these different nature are making UHECR nucleon origination local and well directed (GZK cut off, tens Mpc distances) or even much local and smeared (a few Mpc) for our lightest UHECR nuclei considered in recent articles \\cite{Fargion2008}. Heavy nuclei may also have, by photon-disruption a gamma and neutrino secondary tail, probably so smeared to be mostly sink into background noise: if UHECR are only iron, as some authors still believe,  than UHECR astronomy, in particular the extragalactic one, will be  so much bent, polluted and  smeared to be  hopeless: only large scale iron anisotropy might occur by nearest galactic sources.  On the contrary extragalactic lightest nuclei UHECR astronomy may be  surrounded by a parasite (a little smeared) trace made by point source gamma, neutrinos (TeVs-PeVs) and also tens EeV energy UHECR fragments. We did suggested this possibility since few years \\cite{Fargion2011}. The lightest UHECR nuclei model \\cite{Fargion2008} is able to explain the surprising absence or paucity of events  toward Virgo (the nearest extragalactic cluster of galaxy) and the angle range $10^{o}-15^{o}$, the spread directionality (orthogonal to galactic arms) of  events around and along Cen A. A large number of authors discussed mainly the spectra of UHECR cut off as a key discriminator for UHECR modeling. However the mass composition role may confuse the real shape of any (apparent) GZK cut off. We address more on the UHECR anisotropy nature able we hope  to  correlate UHECR maps and composition, using all radio, IR, gamma MeV-GeV-TeV, maps available.   \\begin{figure}[htb] \\begin{center} \\epsfig{file=Fargion_ERICE-11_fig01.eps,scale=0.16} \\epsfig{file=Fargion_ERICE-11_fig02.eps,scale=0.15} \\caption{Left: The last 2010 UHECR event map  event map overlap with nearby infrared sky map; the clustering toward Cen A is the main UHECR signal.  Around this source the presence of a twin collinear multiplet clustering at twenty EeV shown in next figures. Note the relevance of Infrared Virgo cluster and its absence in AUGER map. Note also a tiny correlation of Vela with an unique galactic triplet event. Right: The last Planck infrared map and labels , whose spread white noise is due to galactic dust, with UHECR events. Note the partial suppression of UHECR events in those  regions where white dust is missing: this absence may hint for a galactic component of UHECR.}\\label{1r} \\end{center} \\end{figure}      \\subsection{Deuterium, proton, gamma and neutrino tails }  \\begin{figure}[!t] \\vspace{5mm} \\centering \\includegraphics[width=3.1 in]{Fargion_ERICE-11_fig03.eps} \\includegraphics[width=2.8 in]{Fargion_ERICE-11_fig04.eps} \\caption{Left: the AUGER 2010 UHECR event map and two of the three multiplet clustering toward Cen A; their sources as shown by dotted curve are  within a tiny disk area (at radius of $10^{o}$); the dotted ellipsoid area of the UHECR and multiplet clustering is also extremely correlated, aligned and small. Right: multiplet clustering toward Cen A by random bending of the galactic  spiral magnetic fields.} \\label{fig2} \\end{figure}      \\begin{figure}[!t] \\vspace{5mm} \\centering \\includegraphics[width=2.8 in]{Fargion_ERICE-11_fig05.eps} \\includegraphics[width=3.2 in]{Fargion_ERICE-11_fig06.eps} \\caption{Left: The AUGER 2010 UHECR event map on radio 408 Mhz map and two of the three multiplet clustering toward Cen A. Right:  the third Multiplet clustering points somehow toward Large and Small Magellanic clouds  as well it overlaps with the clustering along the  Magellanic stream. } \\label{fig2-3} \\end{figure}     UHECR formed (mostly) by  lightest nuclei may explain a partial clustering of events, as the one around CenA as well as a puzzling UHECR absence  around Virgo. Light nuclei are fragile and fly few Mpc before being halted by photo-disruption \\cite{Fargion2008},\\cite{Fargion2009},\\cite{Fargion2010}. Their fragments $He + \\gamma \\rightarrow D+D,D+\\gamma \\rightarrow p+n+ \\gamma  , He + \\gamma \\rightarrow He^{3}+n, He + \\gamma \\rightarrow T +p $  may nevertheless trace the same UHECR maps by a secondary  clustering at half or even fourth of the UHECR primary  energy \\cite{Fargion2011}.  Neutrinos and gamma are tracing  (both for nucleon or light nuclei) their UHECR trajectory, respectively growing at EeVs (if nucleon) or PeV-TeV (if light nuclei) energy. Gamma secondaries rays from cosmic sources are partially absorbed by microwave and infrared background making once again a very local limited UHECR-gamma astronomy. Among neutrinos $\\nu$,  muons ones $\\nu_{\\mu}$, the most penetrating and easy to detect on Earth, are unfortunately deeply polluted by a rich atmospheric  component (as smeared and as the isotropic  parent CR nucleons and nuclei ). This atmospheric  isotropy and homogeneity is probed by last TeV muon neutrino maps in a very smooth ICECUBE neutrino map. Tau neutrinos on the contrary, the last neutral lepton discovered, are almost absent in atmospheric secondaries (about five order of magnitude suppressed at TeVs). Rare $\\nu_{\\mu}\\rightarrow \\nu_{\\tau} $ neutrino oscillation at ten GeV  atmospheric windows,  may nevertheless arise; at tens TeV-PeV up to EeV  $\\nu_{\\tau} $ neutrino might be a clean signal of  UHECR-neutrino associated astronomy\\cite{FarTau}. Their tau birth in ice  may shine as a double bangs (disentangled above PeV)  \\cite{Learned}. In addition UHE tau, born tangent to the  Earth or mountain, while escaping  in air  may lead, by decay in flight, to  loud, amplified  and well detectable tau-airshower at horizons \\cite{Fargion1999},\\cite{FarTau}. Both in atmosphere fluorescence tracks or by Cherenkov blazing, or by partial skimming ground detectors. Tau astronomy versus UHECR are going to reveal most violent  sky as the most deepest probe. This tau airshowers or Earth skimming neutrino \\cite{Feng02} were considered since more than a decade and are going to be observed in AUGER or TA in a few years\\cite{FarTau},\\cite{Auger-01},\\cite{Feng02},\\cite{Auger07}.  Regarding the puzzle of UHECR let us also remind that more than a decade ago we were facing and solving a  how obsolete problem due to  AGASA and Fly Eyes events; such events were calling for sources at distances above GZK cut off. The solution was  based on the sources ejecting primary UHE neutrinos at ZeV energy scattering on relic (cosmic) ones making Z  bosons in flight and, after decay,  UHECR nucleons on Earth \\cite{Fargion1997}. ",
        "conclusions": "The history of Cosmic Rays and last UHECR discoveries (and disclaims) are exciting and surprising. The very unique correlation of UHECR with Cen A, the absence of Virgo, the hint of correlation with  Vela and a mild  connection with galactic plane or even Magellanic stream (see Fig.\\ref{1r},\\ref{fig2},\\ref{fig2-3}) might be solved by a lightest nuclei, mainly He, as a  courier, leading to a very narrow (few Mpc) sky for UHECR.  A soon answer maybe already written into predicted \\cite{Fargion2011} and now observed \\cite{Auger11} multiplet clustering (as Deuterium or proton fragments) at half UHECR edge energy aligned  around or along main UHECR group seed discussed and shown above (see Fig.\\ref{fig2}). Indeed He like UHECR  maybe bent by a characteristic as large as  $\\delta_{rm-He}  \\simeq 16^\\circ$,(while expected  proton or D fragments at half fourth these energies, along tails spread  at $\\delta_{rm-p}  \\simeq 32^\\circ$).  Future  AUGER map possibly about 2-3-4 tens EeV, UHECR fragment clustering maps along higher energy events (5-6 $10^{19}$ eV) may probe and reinforce our interpretation (already tested with present multiplet clustering at $ P (3,2) \\simeq 3 \\cdot 10^{-5}$ level). Additional clustering may occur along few galactic sources as Vela. More along galactic plane, as (inspired by Comptel gamma map correlation) with some clustering along Cassiopeia A, in future UHECR Telescope Array events  (see Fig.\\ref{fig3-4}).  The discover of expected  Neutrino astronomy at Icecube or by horizontal  tau  neutrinos  airshower  at ARGO or  Auger,TA telescopes \\cite{FarTau},\\cite{Fargion1999}\\cite{Fargion09b},\\cite{Feng02},\\cite{Auger08}, may also  shed  light on the UHECR nature, their origination  and mass composition,  opening our eyes to secret  UHECR sources."
    },
    "1112/1112.6055_arXiv.txt": {
        "abstract": "The inner regions of barred galaxies contain substructures such as off-axis shocks, nuclear rings, and nuclear spirals.  These substructure may affect star formation, and control the activity of a central black hole (BH) by determining the mass inflow rate. We investigate the formation and properties of such substructures using high-resolution, grid-based hydrodynamic simulations. The gaseous medium is assumed to be infinitesimally-thin, isothermal, and non-self-gravitating.  The stars and dark matter are represented by a static gravitational potential with four components: a stellar disk, the bulge, a central BH, and the bar. To investigate various galactic environments, we vary the gas sound speed $\\cs$ as well as the mass of the central BH $\\MBH$. Once the flow has reached a quasi-steady state, off-axis shocks tend to move closer to the bar major axis as $\\cs$ increases. Nuclear rings shrink in size with increasing $\\cs$, but are independent of $\\MBH$, suggesting that ring position is not determined by the Lindblad resonances. Rings in low-$\\cs$ models are narrow since they are occupied largely by gas on $\\xtwo$-orbits and well decoupled from nuclear spirals, while they become broad because of large thermal perturbations in high-$\\cs$ models.  Nuclear spirals persist only when either $\\cs$ is small or $\\MBH$ is large; they would otherwise be destroyed completely by the ring material on eccentric orbits.  The shape and strength of nuclear spirals depend sensitively on $\\cs$ and $\\MBH$ such that they are leading if both $\\cs$ and $\\MBH$ are small, weak trailing if $\\cs$ is small and $\\MBH$ is large, and strong trailing if both $\\cs$ and $\\MBH$ are large. While the mass inflow rate toward the nucleus is quite small in low-$\\cs$ models because of the presence of a narrow nuclear ring, it becomes larger than $0.01\\Aunit$ when $\\cs$ is large, providing a potential explanation of nuclear activity in Seyfert galaxies. ",
        "introduction": "\\label{sec:intro}  Stellar bars play an important role in the dynamical evolution of gas in galaxies.  By introducing a non-axisymmetric torque, they produce interesting morphological substructures in the gaseous medium, including a pair of dust lanes at the leading side of the bar, a nuclear ring near the center, and nuclear spirals inside the ring (e.g., \\citealt{san76,rob79,sch81,van81,ath92b,pin95,but96,mar03a,mar03b,mar06}). They also transport gas inward which can trigger starbursts in the rings (e.g., \\citealt{but86,gar91,hel94,bar95,mao01,maz08}) and if the mass inflow extends all the way to the center, they may help power active galactic nuclei (AGN) (e.g., \\citealt{shl90,reg99,kna00,lau04,van10}).  Since bar substructures represent a nonlinear response of the gas to a non-axisymmetric gravitational potential, their formation and evolution is best studied using direct numerical simulations.\\footnote{\\citet{eng00} argued that physical properties of nuclear spirals can be explained by the linear density-wave theories (see also \\citealt{mac04a}).} There have been a number of numerical studies on the gas dynamics in barred galaxies. Based on the numerical scheme employed, they can be categorized largely into two groups: (1) those using a smoothed particle hydrodynamics (SPH) technique (e.g., \\citealt{eng97,pat00,ann05,tha09}) and (2) those using a grid-based algorithm (e.g., \\citealt{ath92b,pin95,mac02,mac04b,reg03,reg04}). The numerical results from these two approaches do not always agree with each other, at least quantitatively, even if the model parameters are almost identical. For instance, \\citet{pin95} using the CMHOG code on a cylindrical grid reported that the gas near the corotation regions exhibits complex density features resulting from Rayleigh-Taylor and/or Kelvin-Helmholtz instabilities, while these structures are absent in the SPH simulations. In addition, overall shapes and structures of dust lanes and nuclear rings from CMHOG simulations are different from SPH results.  Some differences in the numerical results may be attributable to relatively large numerical diffusion of a standard SPH method and its inability to handle sharp discontinuities accurately (e.g., \\citealt{age07,pri08,rea10}).  However, after adopting and thoroughly testing the CMHOG code as part of this work, we have found it contained a serious bug in the way the gravitational forces due to the bar are added to the hydrodynamical equations.  Thus, some of the discrepancies in the flows computed by CMHOG and other codes are likely due to this bug. We discuss this bug and its affect on the results reported in \\citet{pin95} in Section 2.2.  In this paper, we revisit the gas dynamics in barred galaxies using a corrected version of the CMHOG code.  Our objectives are three-fold. First, we wish to remedy the errors in \\citet{pin95}, and to compute the formation of bar substructures with an accurate shock-capturing grid code with the correct bar potential.  Second, the morphology, shape, and strength of the bar substructures are likely to depend on the gas sound speed and the shape of the underlying gravitational potential (e.g., \\citealt{eng00}).   Thus, we report new models in which we include a central black hole (BH) that greatly affects the gravitational potential in the central regions, and we vary both the BH mass as well as the sound speed to explore the dynamics in various galactic conditions.  Third, we exploit advances in computational resources to compute models that have more than an order of magnitude higher resolution than the models in \\citet{pin95}, with a grid resolution of $0.13$ pc in the central regions.  This allows us to resolve details in the flow in the nuclear regions, in particular the formation of nuclear rings and nuclear spirals.  According to the most widely accepted theory, a nuclear ring forms near the inner Lindblad resonance (ILR) when there is only one ILR, as the gas outside (inside) ILR loses (gains) angular momentum and accumulates there, while it forms in between the inner ILR and outer ILR when there are two ILRs (e.g., \\citealt{shl90,com96,but96}). On the other hand, \\citet{reg03} argued that the ring formation is more deeply related to the existence of $\\xtwo$-orbits rather than the ILRs.  But, the arguments relying on either ILR or $\\xtwo$-orbits do not take into account the effect of thermal pressure.  Therefore, it is important to explore to what extent the concepts of ILR or $\\xtwo$-orbits are valid in describing nuclear rings, especially when the sound speed is large.  The formation, shape, and nature of nuclear spirals that may channel the gas to the galaxy centers are also not well understood. Observations using the \\emph{Hubble Space Telescope} indicate that galaxies having nuclear dust spirals are quite common (e.g., \\citealt{mar03a,mar03b}). While most of such spirals are trailing, a few galaxies including NGC 1241 and NGC 6902 reportedly possess leading nuclear spirals \\citep{dia03,gro03}. Although the linear theory suggests that leading spirals are expected when there are two ILRs (e.g., \\citealt{mac04a}), they are absent in the numerical models of \\citet{pin95} computed with the CMHOG code, while the SPH models of \\citet{ann05} with self-gravity do form leading spirals. The SPH models suffer from poor spatial resolution at the nuclear regions as most particles gather around the rings.  By running high-resolution simulations with a corrected version of CMHOG, we can clarify the issues of the nuclear spiral formation and related mass inflow rates to the galaxy center.  We in this work treat gaseous disks as being two-dimensional, isothermal, non-self-gravitating, and unmagnetized, which introduces a few caveats that need be noted from the outset. By considering an infinitesimally-thin disk, we ignore gas motions and associated dynamics along the direction perpendicular to the disk. By imposing a point symmetry relative to the galaxy center, our models do not allow for the existence of odd-$m$ modes, although this appears reasonable since $m=2$ modes dominate in the problems involving a galactic bar. In addition, we are unable to capture the potential consequences of gaseous self-gravity and magnetic stress that may not only cause fragmentation of high-density nuclear rings but also affect mass inflow rates to the galaxy center. Nevertheless, these idealized models are useful to isolate the effects of the gas sound speed and the mass of a central BH on the formation of bar substructures and mass inflows. Also, these models allow us to correct the results of previous CMHOG calculations with incorrect bar forces.   This paper is organized as follows. In Section 2, we describe the galaxy model, model parameters, and our numerical methods. In Section 3, we present the results of simulations for off-axis shocks and nuclear rings. The detailed properties of nuclear spirals are presented in Section 4. In Section 5, we study the mass inflow rates through the inner boundary obtained from our simulations. In Section 6, we conclude with a summary and discussion of our results and their astronomical implications.  \\begin{deluxetable}{lcc} \\tabletypesize{\\footnotesize} \\tablewidth{0pt} \\tablecaption{Model Parameters\\label{tbl:model}} \\tablehead{ \\colhead{$\\;\\;\\;\\;\\;$Model$\\;\\;\\;\\;\\;$} & \\colhead{$\\;\\;\\;\\;\\;$$\\cs ~ ({\\rm km\\; s^{-1}})$$\\;\\;\\;\\;\\;$} & \\colhead{$\\;\\;\\;\\MBH(0) (\\Msun)\\;\\;\\;$} } \\startdata cs05bh0 & 5  & 0  \\\\ cs05bh0t\\tablenotemark{a} & 5  & 0  \\\\ cs10bh0 & 10 & 0  \\\\ cs15bh0 & 15 & 0  \\\\ cs20bh0 & 20 & 0  \\\\ cs20bh0t\\tablenotemark{a} & 20 & 0  \\\\ \\hline cs05bh7 & 5  & 4$\\times10^7$ \\\\ cs10bh7 & 10 & 4$\\times10^7$ \\\\ cs15bh7 & 15 & 4$\\times10^7$ \\\\ cs20bh7 & 20 & 4$\\times10^7$ \\\\ cs20bh7t\\tablenotemark{a} & 20 & 4$\\times10^7$ \\\\ \\hline cs05bh8 & 5  & 4$\\times10^8$ \\\\ cs10bh8 & 10 & 4$\\times10^8$ \\\\ cs15bh8 & 15 & 4$\\times10^8$ \\\\ cs20bh8 & 20 & 4$\\times10^8$ \\tablenotetext{a}{ The BH mass is varied with time as $\\MBH(t) = \\MBH(0) + \\int_0^t \\Mdot(t^\\prime)dt^\\prime$ assuming that all the inflowing mass is added to the central BH.} \\end{deluxetable} ",
        "conclusions": "\\subsection{Summary}  We have presented detailed numerical models that explore the formation of substructures produced by the gas flow in barred galaxies. Previous models based on particle simulations (e.g., \\citealt{eng97,ann05,tha09}) did not have sufficient resolution to resolve nuclear spirals.  On the other hand, studies that used the grid-based code CMHOG (e.g., \\citealt{pin95,mac04b}) unknowingly made mistakes in the force evaluation for the bar potential, so that the results needed to be recomputed.  \\begin{figure} \\epsscale{1.2}\\plotone{fig19.eps} \\caption{ Temporal evolution of the mass inflow rates (top) and the BH mass (bottom) in Models cs05bh0t, cs20bh0t, and cs20bh7t where $\\MBH$ is allowed to vary with time. The mass inflow rates from the fixed-$\\MBH$ counterparts are compared as dotted lines.  In all models, the total increase in the BH mass over 500 Myr is not large enough to cause significant changes in $\\Mdot$. \\label{fig:mdot_MBHt}} \\end{figure}   In this paper we have corrected the errors in the original CMHOG code and run high-resolution hydrodynamical simulations. To resolve the nuclear regions, we employ a logarithmically-spaced cylindrical grid, with a zone size of $\\Delta R \\leq 6\\pc$ at $R\\leq1\\kpc$ where nuclear rings and spirals form.  We have included the potential from a central BH, and studied the flow properties as the mass of the BH $\\MBH$ and the sound speed $\\cs$ in the gas are varied. For simplicity, the effects of gaseous self-gravity and magnetic fields are not included. The main results of the  present work are summarized as follows:  1. \\emph{Off-axis Shocks} -- The imposed non-axisymmetric bar potential provides gravitational torques to the otherwise circular-rotating gas, perturbing its orbit.  The perturbed orbits crowd at the downstream sides of the bar major axis and produce overdense ridges that eventually develop into off-axis shocks. At a quasi-steady state, the off-axis shocks are overall almost parallel to $\\xone$-orbits:  they start from the bar major axis at the outer ends, are gradually displaced downstream as they move inward, and connect to the nuclear ring at the inner ends. While the positions of the off-axis shocks are almost independent of $\\MBH$, since the effect of a BH is negligible at large radii, they depend on $\\cs$ in such a way that the shocks are, on average, located closer to the bar major axis as $\\cs$ increases.  This is primarily because gas with larger $\\cs$ should be more strongly perturbed to induce shocks, which occurs deeper in the potential well and thus results in shocks on lower $\\xone$-orbits \\citep{eng97}.  The off-axis shocks are in general curved. Flow streamlines are complicated near the shocks in that they diverge before the shocks and are promptly swept inwards by inflowing gas right after the shocks.  Therefore, the usual Ranking-Hugoniot jump conditions for planar, one-dimensional shocks are not applicable to the off-axis shocks. In fact, the shock strength, as measured by the peak density $\\Sigma_{\\rm max}$ at $R\\sim(1.5-1.8)\\kpc$ from the center, is $\\Sigma_{\\rm max}/\\Sigma_0\\sim3-6$ for models with no BH and do not sensitively depend on $\\cs$. The compression factor of the off-axis shocks are insensitive to the BH mass and depends on $\\cs$ roughly as $\\alpha_{\\rm max}\\sim 7.7 (\\cs/5\\kms)^{0.92}$. The off-axis shocks have very strong velocity shear amounting to $\\sim(1-3)\\times10^3\\kms\\kpc^{-1}$. This strong shear may suppress the growth of gravitational instability in the high-density off-axis shocks when self-gravity is included.   2. \\emph{Nuclear Rings} -- When gas passes through the off-axis shocks, it loses angular momentum, flows inward, and forms a nuclear ring where the centrifugal force balances the external gravity. The ring is attached to the inner ends of the off-axis shocks and thus becomes smaller in size as $\\cs$ increases. The rings that form in our models are all located inside the (outer) ILR, but this does not imply that the ring formation is related to the ILRs. When $\\cs$ is small, the pressure perturbations on the gas orbits are so weak that the rings are quite narrow and their shape is well described by $\\xtwo$-orbits. The  mean radius is $\\Rring\\sim(0.8-0.9)\\kpc$ when $\\cs=5\\kms$ and $\\Rring\\sim(0.5-0.6)\\kpc$ when $\\cs=10\\kms$, independent of the BH mass. This suggests that the ring position is not determined by the $\\Omega-\\kappa/2$ curve, hence by the ILRs. When $\\cs \\geq 15\\kms$, on the other hand, large thermal pressure strongly affects the gas orbits in the nuclear rings. For example, some gas near the contact points between the ring and the off-axis shocks is forced out to follow relatively round orbits, while other gas is pulled in radially to make very eccentric orbits.  These diverse gas orbits near the contact points tend to spread out the ring material, making the rings much broader than in models with smaller $\\cs$.   3. \\emph{Nuclear Spirals} -- Since even the non-axisymmetric bar potential is nearly axisymmetric in the central parts, the gaseous responses inside a nuclear ring are not as dramatic as in off-axis shocks.  Nevertheless, non-axisymmetric $m=2$ perturbations are able to grow inside a ring and develop into nuclear spirals that persist for a long period of time, provided that $\\cs$ is small or $\\MBH$ is large. Although all models have weak spiral structures at early time, in models with large $\\cs$ they are soon destroyed by eccentric gas orbits as well as perturbations induced by the large pressure in the rings unless the BH mass is very large. When $\\MBH=4\\times10^8\\Msun$,  the disruptive pressure perturbations from the rings cannot penetrate the very central parts where the gas has extremely large initial angular momentum.  In this case, the nuclear spirals are protected from the surrounding and thus are long-lived.  The shape of nuclear spirals is determined by the sign of $d(\\Omega-\\kappa/2)/dR$ such that spirals that form in the regions where $d(\\Omega-\\kappa/2)/dR$ is positive (negative) are leading (trailing), confirming the theoretical expectations of \\citet{but96}. With no BH, only the model with $\\cs=5\\kms$ has persistent leading spirals in the regions where the $\\Omega-\\kappa/2$ curve is an increasing function of $R$. The leading spirals in this model are quite strong and develop into shocks with the peak density $\\pSig/\\bSig\\sim7.9$ and the compression factor $\\alpha\\sim 3.4$ at $R=0.25\\kpc$ at the end of the run. Models with $\\MBH=4\\times10^7\\Msun$ initially have hybrid features comprising of trailing spirals at $R<\\Rmin$ and leading spirals at $\\Rmin<R<\\Rmax$, where $\\Rmin$ and $\\Rmax$ refer to the radii where $\\Omega-\\kappa/2$ curve attains a local minimum and maximum, respectively.  When $\\cs\\leq10\\kms$, however, the leading parts in the hybrid spirals are destroyed by the trailing pressure waves launched by the ring, leaving only the weak trailing spirals behind. When $\\cs\\geq15\\kms$, both leading and trailing spirals are destructed completely by the pressure perturbations. In models with $\\MBH=4\\times10^8\\Msun$, $d(\\Omega-\\kappa/2)/dR<0$ in the whole nuclear regions, so that nuclear rings are always trailing. When $\\cs\\leq10\\kms$, the spirals well separated from the ring are weak and tightly wound. When $\\cs\\geq15\\kms$, on the other hand, the trailing spirals are fed by the gas in the nuclear ring and grow into shocks, the outer ends of which join the inner ends of the off-axis shocks smoothly. Although the nuclear spirals in these models have only modest density contrasts of $\\pSig/\\bSig\\sim1.7$ at $R=0.05\\kpc$, the background density resulting from the inflows of the ring material is greatly enhanced (by more than two orders of magnitude when $\\cs=20\\kms$).  4. \\emph{Mass Inflow Rates} -- Although gas experiences a significant loss in angular momentum  when it meets off-axis shocks during galaxy rotation, this does not necessarily translate into the mass inflow all the way to the galaxy center. In models with $\\cs\\leq10\\kms$, a narrow nuclear ring formed by gas with $\\xtwo$-orbits inhibits further inflows of the gas, resulting in the mass inflow rate $\\Mdot$ through the inner boundary less than $0.01\\Aunit$, regardless of the BH mass.  The mass inflow rates are greatly enhanced to $\\Mdot > 0.01\\Aunit$ in models with $\\MBH\\leq 4\\times10^7\\Msun$ and $\\cs\\geq15\\kms$ or $\\MBH=4\\times10^8\\Msun$  and $\\cs=20\\kms$ for different reasons depending on the BH mass. When $\\MBH\\leq 4\\times10^7\\Msun$, the gas orbits are affected by thermal pressure and some gas in the ring can take highly eccentric orbits, directly falling into the inner boundary.  In models with $\\MBH= 4\\times10^8\\Msun$, on the other hand, the gas orbits near the center are more or less circular, but the density in the nuclear regions is enhanced greatly because of strong nuclear spirals, increasing $\\Mdot$.  \\subsection{Discussion}  Nuclear rings play an important role in evolution of barred galaxies by providing sites of active star formation near the centers (e.g., \\citealt{but86,gar91,bar95,mao01,maz08}).  Rings certainly consist of gas that migrates from outer parts inward by losing angular momentum at off-axis shocks, but what stops further migration to form a ring remains a matter of debate. As a trapping mechanism of the ring material, \\citet{com96} proposed the non-axisymmetric bar toque that forces gas to accumulate between two ILRs or at a single ILR depending on the shape of the gravitational potential (see also \\citealt{but96}), while \\citet{reg03,reg04} favored the gas transitions from $\\xone$- to $\\xtwo$-orbits rather than orbital resonances.  However, our numerical results show that a ring is formed at early time because of the centrifugal barrier that the migrating material feels. Later on, nuclear rings slowly shrink in size as gas with lower angular momentum gas is continuously added. Although nuclear rings are located in between two ILRs in models with no BH, this is a coincidence.  Models with a central BH show that the specific ring positions are insensitive to the shape of the $\\Omega-\\kappa/2$ curves, suggesting that nuclear ring formation is not a consequence of the orbital resonances. In fact, the gas flows that produce the ring are not in force balance and have a large radial velocity, so that the concept of resonances and ILRs is not applicable to nuclear rings (e.g., \\citealt{reg03}). In addition, the notion of $\\xtwo$-orbits as the gas trapping locations is meaningful only for small $\\cs$.\\footnote{The models considered by \\citet{reg03,reg04} had the sound speed fixed to $\\cs=5\\kms$.} When the sound speed is large, thermal pressure at the contact points between the off-axis shocks and nuclear ring causes the gas orbits to deviate from $\\xtwo$-orbits considerably.  The results of our simulations suggest that not all barred-galaxies possess nuclear spirals at their centers.  Long-lasing nuclear spirals exist only when either the sound speed is small or the BH mass is large: they do not survive in models with \\emph{both} large $\\cs$ \\emph{and} small $\\MBH$. Two common views regarding the nature of nuclear spirals are low-amplitude density waves and strong gaseous shocks (see e.g., \\citealt{eng00,mac02,mac04a,mac04b,ann05,tha09}). And, our simulations indeed show that they are either tightly-wound density waves or shocks when the pith angle is relatively large ($\\ip>6^\\circ$).  One may  speculate that nuclear spirals are more likely to be shocks rather than density waves when $\\cs$ is small. In contrast to this prediction, however, nuclear spirals in models with $\\MBH=4\\times10^8\\Msun$ are density waves when $\\cs$ is small and shocks when $\\cs$ is large. This is of course because as $\\cs$ increases, waves in nuclear regions tend to be more open (with smaller $|k|$) in the beginning, and they are subsequently supplied with more gas from the rings as they grow. This is entirely consistent with the results of \\citet{ann05} who used SPH simulations to show that nuclear spirals are supported by shocks when $\\cs\\simgt15\\kms$ in models with a massive BH. We note however that weak trailing spirals seen in our Model cs10bh8 are absent in Model M2 ($\\cs=10\\kms$ and $\\MBH=4\\times10^8\\Msun$) of \\citet{ann05}, which is presumably due to an insufficient number of particles to resolve nuclear spirals in their SPH simulations.   Of 12 models with differing $\\cs$ and $\\MBH(0)$ that we have considered, only 1 model possesses leading spirals, suggesting that galaxies with leading nuclear spirals would be very uncommon in nature.  To our knowledge, only two galaxies, NGC 1241 and NGC 6902, are known to possess leading features in the nuclear regions \\citep{dia03,gro03}\\footnote{Leading arms in NGC 1241 are detected by Pa$\\alpha$ emissions tracing young stars, while those in NGC 6902 are observed in the K$^\\prime$ band tracing old populations. It is uncertain how these stellar features are related to gaseous nuclear spirals studied in this paper.}. Based on our simulations, the existence of leading spirals at centers requires two stringent conditions: (1) the gas should be dynamically cold enough to protect nuclear spirals from nuclear rings and (2) there should be a wide range of radii with $d(\\Omega-\\kappa/2)/dR<0$ in the central parts, which can be easily accomplished when there are two ILRs (or without a strong central mass concentration). The second condition is consistent with the linear theory that predicts short leading waves propagating outward from the inner ILR \\citep{mac04a}. The facts that NGC 6902 is a barred-spiral galaxy \\citep{lau04} and does not show significant $X$-ray emissions indicative of AGN activities \\citep{des09} are not inconsistent with the second requirement for the existence of leading nuclear spirals. Since NGC 1241 is a Seyfert 2 galaxy with an estimated BH mass of $\\log (\\MBH/\\Msun)=7.46$ \\citep{bia07}, however, the nuclear star-forming regions in this galaxy are unlikely to be associated with gaseous nuclear spirals studied in this work.  Finally, we discuss the mass inflow rates derived in our models in regard to powering AGNs in Seyfert galaxies.   The mass accretion rate is often measured by the Eddington ratio defined by $\\lambda \\equiv L_{\\rm bol}/L_{\\rm Edd} = 4.5\\times10^{-2} (\\epsilon/0.1) (\\Mdot/10^{-2}\\Aunit)(\\MBH/10^7\\Msun)^{-1}$, where $L_{\\rm bol}$ and $L_{\\rm Edd}$ denote the bolometric and Eddington luminosities of an AGN and $\\epsilon$ is the mass-to-energy conversion efficiency of the accreted material. Observations indicate that $\\lambda\\simlt 0.1$ for classical Seyfert 1 galaxies with broad iron emission lines (e.g., \\citealt{mey11}; see also \\citealt{pet97}), while $\\lambda\\sim 10^{-3}$ for low-luminosity Seyfert 1 AGNs (e.g., \\citealt{ho08}). In our numerical models, the mass inflow rates are larger for models with smaller $\\MBH$ and larger $\\cs$. Taking $\\epsilon\\approx0.1$ (e.g., \\citealt{yu02}), the mass inflow rates amount to $\\lambda \\simlt 10^{-4}$ for $\\MBH=4\\times10^8\\Msun$ and $\\cs\\simlt15\\kms$, $\\lambda \\sim 10^{-3}$ for $\\MBH=4\\times10^8\\Msun$ and $\\cs=20\\kms$ or for $\\MBH=4\\times10^7\\Msun$ and $\\cs=10\\kms$, and $\\lambda \\sim 0.02-0.4$ for $\\MBH=4\\times10^7\\Msun$ and $\\cs\\simgt 15\\kms$.  For classical Seyfert galaxies with strong bars, this suggests that the masses of central BHs are likely to be less than $10^8\\Msun$, which appears to be consistent with the measured values from the relation between the BH masses and the velocity dispersions of stellar bulges (e.g., \\citealt{wat08}) and reverberation mapping techniques (e.g., \\citealt{gul09,den10}). Of course, this result may depend on many factors such as the axis ratio and strength of the bar, presence of self-gravity and magnetic fields, gas cooling and heating, turbulence, etc., all of which would affect gas dynamic significantly.  Extending the present work to include these physical ingredients would be an important direction of future research."
    },
    "1112/1112.3736_arXiv.txt": {
        "abstract": "{Circumstellar disks and envelopes of low-mass young stellar objects (YSOs) contain significant amounts of ice. Such icy material will evolve to become volatile components of planetary systems, such as comets in our solar system.} {To investigate the composition and evolution of circumstellar ice around low-mass YSOs, we observed ice absorption bands in the near infrared (NIR) towards eight YSOs ranging from class 0 to class II, among which seven are associated with edge-on disks.} {We performed slit-less spectroscopic observations using the grism mode of the Infrared Camera (IRC) on board AKARI, which enables us to obtain full NIR spectra from 2.5 $\\mu$m to 5 $\\mu$m, including the CO$_2$ band and the blue wing of the H$_2$O band, which are inaccessible from the ground. We developed procedures to carefully process the spectra of targets with nebulosity. The spectra were fitted with polynomial baselines to derive the absorption spectra. The molecular absorption bands were then fitted with the laboratory database of ice absorption bands, considering the instrumental line profile and the spectral resolution of the grism dispersion element.} {Towards the class 0-I sources (L1527, IRC-L1041-2, and IRAS04302), absorption bands of H$_2$O, CO$_2$, CO, and XCN are clearly detected. Column density ratios of CO$_2$ ice and CO ice relative to H$_2$O ice are  $21-28$ \\% and $13-46$ \\%, respectively. If XCN is OCN$^-$, its column density is as high as $2-6$ \\% relative to H$_2$O ice. The HDO ice feature at 4.1 $\\mu$m is tentatively detected towards the class 0-I sources and HV Tau. Non-detections of the CH-stretching mode features around 3.5 $\\mu$m provide upper limits to the CH$_3$OH abundance of 26 \\% (L1527) and 42 \\% (IRAS04302) relative to H$_2$O. We tentatively detect OCS ice absorption towards IRC-L1041-2. Towards class 0-I sources, the detected features should mostly originate in the cold envelope, while CO gas and OCN$^-$ could originate in the region close to the protostar, where there are warm temperatures and UV radiation. We detect H$_2$O ice band towards ASR41 and 2MASSJ1628137-243139, which are edge-on class II disks. We also detect H$_2$O ice and CO$_2$ ice towards HV Tau, HK Tau, and UY Aur, and tentatively detect CO gas features towards HK Tau and UY Aur.} {} ",
        "introduction": "In molecular clouds, protostellar envelopes, and protoplanetary disks, significant amounts of oxygen and carbon are in the form of molecules in ice mantles, such as H$_2$O, CO, CO$_2$, and CH$_3$OH (e.g. Whittet \\cite{whittet93}, Murakawa et al. \\cite{murakawa00}, Gibb et al. \\cite{gibb04}). These interstellar ices are formed by the adsorption of gas-phase molecules onto grain surfaces and/or the grain-surface reactions of the adsorbed species (e.g. Aikawa et al. \\cite{aikawa05}).  Ices in circumstellar disks of low-mass YSOs are of special interest as they contribute to the raw material of planetary formation. Observations of disk ices, however, are not straightforward for the following two reasons, and thus remain limited in number. Firstly, ices exist mostly near the midplane at $r\\ge$ several AU. Observations of disk ices are basically restricted to edge-on disks, where the ice bands absorb stellar light, scattered stellar light, and/or thermal emission of warm dust in the disk inner radius (Pontoppidan et al \\cite{crbr05}). The targets are naturally heavily extincted by dust, and thus are faint. Secondly, it is unclear whether the ice absorption bands, if detected, originate in the disk or other foreground components (i.e. ambient clouds). The source CRBR 2422.8-3423 was the first edge-on object towards which detailed ice observations were performed (Thi et al. \\cite{thi02}). Pontoppidan et al. (\\cite{crbr05}) concluded, via detailed analysis and modeling of the object, that only a limited amount ($< 20$ \\%) of detected CO ice may exist in the disk, while up to 50 \\% of water and CO$_2$ may originate in the disk. They also found that the 6.85 $\\mu$m band, which is tentatively attributed to NH$_4^+$, has a prominent red wing. Since this wing is reproduced by the thermal processing in the laboratory, it indicates that along the line of sight towards CRBR 2422.8-3423 there is warm ice in the disk.  Honda et al. (\\cite{honda09}) succeeded in detecting water ice in a disk around a Herbig Ae star, which is not edge-on. They observed scattered light from a disk around HD142527 at multiple wavelength bands around $\\sim 3$ $\\mu$m. This method can be a powerful tool for investigating the spatial distribution of ices in disks. However, it would currently be difficult to detect ice features other than the 3 $\\mu$m water band via scattered light.  In the present paper, we report near infrared (NIR) ($2.5-5$ $\\mu$m) observations of edge-on disks using the grism mode of the IRC on board AKARI. We selected seven edge-on young stellar objects (YSOs); three are considered to be in a transient state of class 0-I, and four are class II (see \\S 2 for more detailed descriptions). Although the absorption by the protostellar envelope may overwhelm the absorption by the disk in class 0-I YSOs, envelope material along our line of sight (perpendicular to the rotation axis and outflow) is likely to end up in the midplane region of the forming disk, and ices with a sublimation temperature of $> 50$ K (i.e. CO$_2$ and H$_2$O) would survive the accretion shock onto the disk (Visser et al. \\cite{visser09}).  Ice observations towards low-mass YSOs have been performed using {\\it Spitzer Space Telescope (SST)} and various ground-based telescopes such as the Very Large Telescope (VLT) and Subaru. Pontoppidan et al. (\\cite{pontoppidan2003}) observed the 4.67 $\\mu$m CO ice band, which is observable from the ground, towards 39 YSOs. Using {\\it SST}, Furlan et al. (\\cite{furlan08}) and Pontoppidan et al. (\\cite{pontoppidan2008}) observed the strong CO$_2$ ice band at $15.3$ $\\mu$m towards class I sources. Observations with {\\it SST} also detected absorption bands of CH$_3$OH, HCOOH, H$_2$CO, and CH$_4$ (Boogert et al. \\cite{boogert08}, Zasowski et al. \\cite{zaowski09}, {\\rm \\\"{O}}berg et al. \\cite{oberg_ch4}). Since ground-based observations are restricted to atmospheric windows, and since {\\it SST} is restricted to the mid infrared ($5-36$ $\\mu$m), NIR observations via AKARI are important to comprehend the {\\it full} spectra of ices in circumstellar material. AKARI has a high enough sensitivity for the spectroscopic observation of low-mass YSOs; the sensitivity of the AKARI grism mode at 3 $\\mu$m is 120 $\\mu$Jy (1 $\\sigma$, 10min), which is almost comparable to the K-band sensitivity of the spectroscopic observation at 8-m ground-based telescopes (VLT and Subaru). While the L-band sensitivity on the ground is $\\sim 5$ mag or 100 times worse than at K-band, the sensitivity of the AKARI grism mode does not vary significantly at $2.5- 4$ $\\mu$m, and only gradually worsens at longer wavelengths, by up to a factor of 3-4 at 5 $\\mu$m. ",
        "conclusions": "We have observed ice absorption bands at $2.5-5$ $\\mu$m towards eight low-mass YSOs: three class 0-I protostellar cores with edge-on geometry, two edge-on class II objects, two multiple systems with edge-on class II, and one not-edge-on class II object.  Towards the class 0-I objects, L1527, IRC-L1041-2, and IRAS04302, we have detected abundant H$_2$O, CO$_2$, and CO ice in the envelope. The column density ratio of CO$_2$ to H$_2$O ice is $21-28$ \\%, which coincides with the ratio observed by {\\it SST} towards YSOs with various inclinations. The weak absorption at $\\sim 4.1$ $\\mu$m can be fitted by HDO ice; the HDO/H$_2$O ratio ranges from 2 \\% to 22 \\%. The absorption in the vicinity of the CO band (4.76 $\\mu$m) is double-peaked and fitted by combining CO ice, OCN$^-$, and CO gas. The large column density of CO ice suggests that the envelope is still very dense and cold, while OCN$^-$ and CO gas features would originate in the region close to the protostar. The column density of OCN$^-$ is as high as $2-6$ \\% relative to H$_2$O, which is much higher than previous observations. Our lines of sight (high inclinations from the rotation axis) may preferentially trace the regions with high UV irradiation, such as the surface of a forming disk and/or torus envelope.  The spectrum of IRAS04302 includes the 3.5 $\\mu$m absorption band, but the feature does not match either CH$_3$OH or CH$_4$. An OCS absorption band is tentatively detected towards IRC-L1041-2.  Towards the edge-on class II objects, ASR41 and 2MASS J1628137-243139, we have detected the H$_2$O band. The low optical depth of the water feature is due to geometrical effects (Pontoppidan et al. \\cite{crbr05}). The detected water ice mainly originates in the disk.  Both HK Tau B and HV Tau C are edge-on class II objects in multiple systems. In our spectra, which are dominated by the primaries, we have detected the absorption of H$_2$O ice and CO$_2$ ice. Ices in both the edge-on disks and foreground clouds would contribute to the absorption, although H$_2$O absorption towards HK tau could originate solely in the disk. Even if the observed features are due to the foreground clouds, the H$_2$O ice column densities in the clouds are much smaller than those observed towards HK Tau B and HV Tau C with Subaru (Terada et al. \\cite{terada07}), which confirms that the ice columns of the latter must originate in disks. The foreground H$_2$O ice and CO$_2$ ice are also detected towards UY Aur, which is not an edge-on system. We have tentatively detected CO gas towards HK Tau and UY Aur."
    },
    "1112/1112.1505_arXiv.txt": {
        "abstract": "We study the implications of asymmetric dark matter on neutron stars.  we construct a  \"mixed neutron star\" model composed of ordinary baryons and of  asymmetric dark matter baryons. We derive the general relativistic structure equations for each specie, the equation for the mass within  a given radius, and the redshift as function of radius. We present one specific numerical model as an illustrative example. In this example, the mass of the dark neutron equals half that of the ordinary neutron. The main results are: a total mass of  $3.74 M_{\\odot}$, a total mass within the neutron-sphere equaling  $1.56 M_{\\odot}$, the neutrons  mass is \\  $1.34 M_{\\odot}$, the   star radius is \\  31.9    km, the   neutron-sphere radius is \\ 11.1 km, and the redshifts from the neutron-sphere and from the star surface are \\  0.72, \\    0.25, respectively. We comment briefly on possible astrophysical implications. ",
        "introduction": "\\subsection{Background}   Cold dark matter   which is favored by most astrophysical and cosmological observations can be realized in symmetric and (or)  asymmetric scenarios. In the first class of models, dark matter is made of stable $X$  particles and an equal amount of $\\bar X$ antiparticles of mass $m_X$. In the early universe, these were   in thermal equilibrium and their residual abundance $\\Omega_X$ is fixed at the \"freeze-out\" value when the rate of the Hubble expansion overcomes that of $\\bar X-X $ annihilation rate. A prototypical example which have been most extensively studied is   SUSY(super-symmetry). Unfortunately, the searches for SUSY partners in the new large hadron collider (LHC)   have failed to detect them. More specifically, insofar as SUSY dark matter models are concerned searches for electrons, positrons or photons in clumped dark matter in and around our Galaxy, and for energetic neutrinos resulting from annihilations of Xs do not provide solid  \"indirect\" evidence for dark matter. Moreover, the ongoing direct underground searches put very strong bounds on the scattering crossections of massive Xs on nuclei.  Additional constraints are related to accretion of galactic  SUSY WIMPS onto the sun. The increased density of the captured WIMPS, accelerates the rate of particle-antiparticle annihilation. The resulting photons and electrons are   trapped in the sun but the resulting  ultra high energy  neutrinos are not.  Data from the ICE-CUBE Cerenkov radiation detector near the south pole, severely constrain  such models  \\cite{abbasi 09, abbasi 10} . In the case of asymmetric dark matter, there will be no such annihilation in the sun. Moreover, once the fraction of dark matter particles in the sun exceeds the ratio of $\\sigma_{Xn}/\\sigma_{XX}$,   scattering on the already captured X particles in the sun dominates over scattering on the baryons.  Finally,   the observed number of satellite mini-halos around  the milky way galaxy  is two orders of magnitude smaller than predicted within the symmetric dark matter framework \\cite{Klypin, Moore}.  Consequently there has been, in the very recent years, a renewed interest  in asymmetric dark matter  (ADM) which just like ordinary baryonic matter, is charge non-symmetric with say only the dark baryon (or generally only the particle)  excess remains after the annihilation of most antiparticles .While there is no single clear-cut explanation for the ordinary baryonic asymmetry, the required dark matter density is readily achieved if the mass ratio of the  X particles and baryons $m_X/m_b$ is tuned inversely with the corresponding ratio of asymmetries. An early example  of such a model has been proposed in \\cite{nuss85}.  For recent   studies of asymmetric dark matter   see   \\cite{An etal 2010,mrm11} and references therein.  \\subsection{Cosmological and astrophysical considerations}  The idea of asymmetric dark matter should reconcile with     astrophysical and cosmological data. An obvious constraint is imposed by big-bang nucleosynthesis, which implies that  the mirror neutrinos and  photons   do not contribute to the rate of the cosmic expansion at that era. Another constraint is that the mirror large scale structure-formation should precede the recombination of   ordinary matter, in order to serve as seeds gravitational potential wells  for the latter. In the model described in  \\cite{An etal 2010} the mirror neutrinos and the mirror photons are massive enough for these constraints to be satisfied. We note that even if the mirror neutrinos and photons are massless the above two constraints can be satisfied provided that  the  dark  cosmic-temperature is lower than the ordinary one.  The details of large scale structure formation will depend on the specific model for the asymmetrical dark matter. Generally, being self interacting one may expect it to form mirror structures on all astrophysical   scales.  Also, asymmetric dark matter  will  no longer provide collisionless dark halos as in the symmetric cold dark matter case, so that  the flatness of galactic rotation curves will have to be readdressed in this new context. ",
        "conclusions": "We have demonstrated    that a   mixed neutron star can, as expected, have a mass higher than ordinary neutron stars. At the same time the physical  radius, as probed by ordinary massless and massive particles, is the neutron-sphere radius which is similar   to the radius of ordinary neutron stars.  An important question, not addressed  here, is that of stability.   Since the models form a two-parameters family (the central densities) the question of stability is more complex than in the one parameter family of ordinary neutron star models. There are a number of quite interesting astrophysical implications with regard to phenomenology of compact x-ray sources. Can some of the stellar mass binary  black holes be actually mixed neutron stars? The  neutron-sphere redshift is about  50\\% higher than in the ordinary neutron star case, which may have interesting results for the temperature, radius and luminosity measured by a distant observer.  The larger neutrons binding energy would lead to a smaller value of the maximal neutrons mass, compared to an ordinary neutron star.   I wish to thank my collaborators, R. Mohapatra, S. Nussinov, D. Teplitz and V.~Teplitz as well as  W. Kluzniak and L. Zdunik for interesting discussions. Thanks are due  to   the Afeka College Research Committee  for financial support."
    },
    "1112/1112.3706.txt": {
        "abstract": "We investigate gas accretion flow onto a circumplanetary disk from a protoplanetary disk in detail by using high-resolution three-dimensional nested-grid hydrodynamic simulations, in order to provide a basis of formation processes of satellites around giant planets. % Based on detailed analyses of gas accretion flow, we find that most of gas accretion onto circumplanetary disks occurs nearly vertically toward the disk surface from high altitude, % which generates a shock surface at several scale heights of the circumplanetary disk. % The gas that has passed through the shock surface moves inward because its specific angular momentum is smaller than that of the local Keplerian rotation, while gas near the midplane in the protoplanetary disk cannot accrete to the circumplanetary disk. % Gas near the midplane within the planet's Hill sphere spirals outward and escapes from the Hill sphere through the two Lagrangian points L$_1$ and L$_2$. % % We also analyze fluxes of accreting mass and angular momentum in detail and find that the distributions of the fluxes onto the disk surface are well described by power-law functions and that a large fraction of gas accretion occurs at the outer region of the disk, i.e., at about 0.1 times the Hill radius.  The nature of power-law functions indicates that, other than the outer edge, there is no specific radius where gas accretion is concentrated. % These source functions of mass and angular momentum in the circumplanetary disk would provide us with useful constraints on the structure and evolution of the circumplanetary disk, which is important for satellite formation. ",
        "introduction": "Satellite systems around the giant planets in our solar system are commonly seen.  They are thought to have formed in circumplanetary disks, which are believed to have existed around giant planets during their gas capturing growing stage.  % MMSN model In earlier works, formation process of satellite systems have been considered based on a minimum mass subnebula (MMSN) model, in which satellites form from a disk that contains sufficient solid mass with solar composition for reproducing the current satellite systems (e.g., \\citet{Lunine82}), as an analog of the minimum mass {\\it solar} nebula model \\citep{Hayashi81}.  However, it was suggested that the MMSN model has difficulty in reproducing current satellite systems around Jupiter and Saturn \\citep{Canup02}.  One of severe problems is that the model leads to much higher temperature than that of H$_2$O ice sublimation at the current regular satellite region, which means that ice, which is the main component of the satellites, cannot be used as building material of the satellites.   % Recent models (starved disk model) In order to overcome the difficulties of the MMSN-type models which assumes a closed and static disk, alternative models have been developed. % \\citet{Canup02} proposed a model in which an accretion disk with continuous supply of gas and solid is considered as a proto-satellite disk. % This model is based on results of hydrodynamic simulations of gas capturing process of giant planets \\citep[e.g.,][]{Lubow99, DAngelo02}, which demonstrated that gas accretion from the protoplanetary disk toward the parent planets occurs through a circumplanetary disk. % In this model, the surface density and temperature of the circumplanetary disk is kept lower than assumed in the MMSN model, and thus H$_2$O ice can be used as a solid building material of satellites. % Based on comparison among time scales of different processes, they concluded that formation of the Galilean satellites can be best explained if the circumplanetary disk had one order of magnitude lower gas surface density than the MMSN disk, with slow gas accretion rate corresponding to Jupiter's growth time scale longer than a few $\\times 10^6$ yr (see also \\citet{Sasaki10}). % On the other hand, \\citet{Mosqueira03a} proposed another disk model which consists of two components, i.e., inner MMSN-type massive disk and outer low-density extended disk.  The model reproduced, for example, three inner Galilean satellites as well as only partially differentiated Callisto. % Although these models seem to reproduce the current satellite systems of the giant planets in our solar system, they needed to assume parameters for the disk structure such as surface density profile as a basis of physical processes of satellite formation.     % Hydrodynamic simulations % The structure of a circumplanetary disk is closely related to the gas accretion process of giant planets.  Gas flow around proto-giant-planets in protoplanetary disks has been studied using hydrodynamic simulations. % % 2D Earlier two-dimensional simulations showed that the gas in the Hill sphere rotates in the prograde direction \\citep{Miki82, Sekiya87, Korycansky96} and a pair of spiral shocks stands in the circumplanetary disks \\citep[e.g.,][]{Kley99, Lubow99, Tanigawa02, DAngelo02}.  However, the scale height of the proto-planetary disk for Jupiter-sized planets is comparable to the Hill radius, thus simulations with two-dimension approximation cannot capture the feature of the accretion flow. % % 3D Recent three-dimensional hydrodynamic simulations revealed that the two-dimensional picture for circumplanetary disks is not appropriate for the flow in the Hill sphere \\citep[e.g.,][]{DAngelo03,Bate03,Machida08,Paardekooper08,Ayliffe09b,Coradini10}. % One of the most important features newly found in three-dimensional calculations is that, in the region of circumplanetary disks, the gas accretes nearly vertically downward toward the midplane from high altitude.   % % Although some studies using hydrodynamic simulations were able to obtain the structures of circumplanetary disks, the direct use of results of hydrodynamic simulations for the structure, such as surface density, is problematic. % One of the reasons is that disk-like shear-dominant flow is susceptible to numerical viscosity, which is intrinsic nature of numerical hydrodynamic calculation, no matter if SPH method or mesh based method is used, thus it is difficult to obtain reliable structure of the disk directly from such simulations. % In particular, orbital radius of the rotating gas in disks basically changes only slightly with time, thus numerical error tends to accumulate easily in simulation of long-term evolution. % Also timescale required for low viscous disks such as circumplanetary disks to reach steady state is much longer than the typical dynamical time of the fluid.  While high-resolution hydrodynamic simulations is needed to resolve the structure of circumplanetary disks, their long-term evolution is difficult to follow with such time-consuming simulation. % In addition, physical (non-numerical) viscosity in the disk is not well understood, and we need to assume specific viscosity models that are hard to justify.   % Purpose of this study % As an alternative approach, in the present work, we examine gas accretion flow onto circumplanetary disks from proto-planetary disks, in order to determine gas accretion rate as a function of distance from parent planets. % Unlike the structure in a rotation disk, the accretion flow onto circumplanetary disks is not susceptible to numerical viscosity. % Also, because the circumplanetary disks are located at the downstream of super-sonic accretion flow, the accretion flow itself is hardly affected by the circumplanetary disk structure, which depends on poorly-known effective viscosity. % In \\S \\ref{sec_methods}, we describe our settings of hydrodynamic simulations.  In \\S \\ref{sec_results}, results and analyses of the simulations are shown.  We discuss implication of our analyses of the accretion flow in \\S \\ref{sec_discussion}.  We summarize our results in \\S \\ref{sec_summary}.     %\\clearpage %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec_discussion}  %======================================================================= \\subsection{Size of Circumplanetary Disks} \\label{sec_disk_size} % The size of a circumplanetary disk is closely related to the location of satellite formation.  However, it is not easy to define the outer edge of the disk, and there have been several attempts. % One natural way would be to define the disk edge based on its density distribution.  But the density is monotonically and smoothly decreasing with increasing radius toward the Hill sphere, so it is difficult to define the edge from density distribution. % % \\citet{Ayliffe09b}, who performed three-dimensional hydrodynamic simulations, suggested a criterion for the disk edge based on angular momentum distribution.  Their simulations show that specific angular momentum of the disk gas has a peak at a certain radial location, whereas specific angular momentum of a Keplerian disk increases monotonically.  Their Figure 2 suggests that the turnover point is about 1/3 of the Hill radius, and they defined the disk edge by the radial location of this peak. % \\citet{Martin11} examined periodic orbits of a particle around a planet under the influence of the gravitational force from a central star and the planet, and found that, as the size of the orbit is increased, the orbits start crossing with each other at $r \\sim 0.4r_{\\rm H}$.  They inferred that this corresponds to the location of the disk's outer edge where tidal torque of central star's gravity becomes strong, and called it tidal truncation radius ($r_{\\rm trunc}$). % The above value of $r_{\\rm trunc}$ is in agreement with the point of turnover of specific angular momentum found in their two-dimensional SPH simulation, as well as in the result of \\citet{Ayliffe09b}. % Our Figure \\ref{Fig_AM_vs_r} shows that the turnover point is $\\tilde{R} \\sim 0.3 \\tilde{r}_{\\rm H}$, which is also roughly in agreement with $r_{\\rm trunc}$.\\footnote{In this comparison, we assume that specific angular momentum in these works were not measured in inertial frame, but in the rotating frame.}   %%% We here consider an alternative criterion by examining ${\\cal V}_R$ defined by Eq.~(\\ref{vr_per_vKep}) to define the position of the outer edge. % We observed that azimuthally-averaged radial velocity in the midplane is positive (outward) in almost all the region, although the value is very small for $\\tilde{R} \\lesssim 0.2$ (Fig.~\\ref{Fig_vr-per-vKep}), and the outward velocity significantly increases at $\\tilde{R} \\gtrsim 0.2$ (see also Fig.~\\ref{Fig_streamlines_outward}). % Since gas at $\\tilde{R} \\gtrsim 0.2$ moves outward and escapes from the Hill sphere quickly, the region $\\tilde{R} \\gtrsim 0.2$ can be regarded as outside of the circumplanetary disk. % We thus define the disk edge as the radial location where ${\\cal V}_R$ starts increasing significantly and has non-negligible positive value; $\\tilde{R} \\sim 0.2$ in our case shown in Fig.~\\ref{Fig_vr-per-vKep}. % % modified 2011/11/04 Note that a two-dimensional SPH simulation shows that negative torque is exerted on the outer disk \\citep{Martin11}, which gives rise to inward flow and is contrary to our results.  This might be due to the fact that two-dimensional simulations, which create clearer spiral structure than three-dimensional ones, tend to enhance torque density. % modified 2011/12/12 In particular, the disk thickness in the outer region ($\\tilde{R} \\gtrsim 0.2$), at which outflow is observed, is very thick (thickness is comparable to radius).  Thus, it would be unlikely that our three-dimensional calculation is significantly affected by the negative torque suggested by \\citet{Martin11}. % replaced by the above at 2011/12/12 %Thus, it would be unlikely that our three-dimensional calculation is %significantly affected by the negative torque suggested by %\\citet{Martin11}.  The outer part of the circumplanetary disk is indeed %very thick (thickness is comparable to radius), thus two-dimensional %approximation is not valid. % % (connected the two paragraph at 2011/12/12) % % removed at 2011/12/12 %Similar outflow was reported by \\citet{Paardekooper08} based on %three-dimensional hydrodynamic simulations of gas flow around Earth size %planets.  They explained that accreting material from the top pushes the %envelope around the planet and the envelope material was expelled into %the equatorial plane.  Although the outflow in our simulation is much %slower than theirs, it seems to be explained by a similar mechanism, %i.e., it is caused by pushing disk material instead of an envelope. % The distribution of accreting angular momentum also seems to indicate a similar radius for the edge.  As we see in Fig.~\\ref{Fig_j_zs}, $\\bar{\\tilde{j}}_{z,\\rm s}$ and $\\tilde{j}_{\\rm Kep}$ have different dependence on $\\tilde{R}$.  If we fit each profile as a single power-law function, they cross each other at $\\tilde{R}\\sim 0.3$ in the present case. % (D).comment0705 This roughly agrees with the location of the disk edge defined above based on the significant increase of outward velocity. % The gas exterior to this radius has too large angular momentum to achieve Keplerian rotation, thus moves radially outward. % % removed at 2011/12/12 %This interpretation would be consistent with the explanation for the %outflow described in \\citet{Paardekooper08}, if we would consider %angular momentum instead of dynamical pressure.   % ell % Once angular momentum of accreting gas is obtained, one can calculate mean specific angular momentum of the accretion flow, which is sometimes used in calculating the the so-called centrifugal radius $\\tilde{r}_{\\rm c} \\equiv \\tilde{\\ell}^2/(3\\tilde{r}_{\\rm H}^3)$ to infer the disk size \\citep[e.g.,][]{Mosqueira03a,Ward10}. % Mean specific angular momentum of the accretion flow within a radius $\\tilde{R}$ is given by \\begin{equation} \\tilde{\\ell}(\\tilde{R}) = \\frac{\\dot{\\tilde{J}}_{\\rm s}(\\tilde{R})} {\\dot{\\tilde{M}}_{\\rm s}(\\tilde{R})}, \\end{equation} where \\begin{equation} \\dot{\\tilde{J}}_{\\rm s}(\\tilde{R}) = \\int_0^{\\tilde{R}} d\\tilde{R}' \\int_0^{2\\pi} \\tilde{R}'d\\phi \\tilde{f}_{\\rm s}(\\tilde{R}',\\phi) \\tilde{j}_{\\rm s}(\\tilde{R}',\\phi). \\end{equation} We can see from Figure \\ref{Fig_ell} that $\\tilde{\\ell}(\\tilde{R})$ increases nearly linearly with radius and levels around $\\tilde{R} \\sim 0.2$, where $\\tilde{\\ell} \\sim 0.2$.  This corresponds to $\\tilde{r}_{\\rm c} \\sim 0.013$, which might give us a rough estimation of the disk size. % % However, it is practically difficult to determine the upper bound of the integral range with respect to $\\tilde{R}$, which affects the value of $\\tilde{\\ell}$, because the circumplanetary disks is smoothly connected to the proto-planetary disks. % Also, unlike the case of a particle, angular momentum of a gas element is not a conserved quantity along a flow in principle, because gas is continuum medium which can transfer angular momentum through waves. This also makes it difficult to determine $\\tilde{\\ell}$ precisely, and thus $\\tilde{r}_{\\rm c}$ as well. %  %This point also suffers us when $\\bar{\\tilde{f}}_{\\rm Kep}$ and %$\\dot{\\tilde{M}}_{\\rm Kep}$ is estimated, even though, from the point of %view of total budget of angular momentum, we expect that this estimation %should not be far from real one. %% %Therefore, in order to determine the disk size or edge, we would not %have to pay attention to $\\tilde{\\ell}$ (see also \\citet{Martin11}).   %In addition, from the point of view of accreting mass or angular %momentum, typical lengths, such as centrifugal radius, should not exist %except for the outer edge of the disk because the rates of accreting %mass and angular momentum are well described by power-law functions (see %Figs.~\\ref{Fig_f_s_f_Kep} and \\ref{Fig_M_s_M_Kep}), %% %even though we do not rule out the possibility that surface density has %peaks somewhere in the circumplanetary disk because of disk internal %evolution, such as turbulent viscosity driven by MRI.    %======================================================================= \\subsection{Picture of Gas Accretion Flow onto Circumplanetary Disks} \\label{sec_picture} % Figure \\ref{Fig_CP-disk} shows a schematic picture of gas accretion flow onto a circumplanetary disk based on the results obtained in the present work.  Gas accretion occurs mostly downward from high altitude with high incident angle (see Fig.~\\ref{Fig_diskstructure}).  The accreting gas is accelerated by the planet gravity to have almost free-fall velocity of the planet (Fig.~\\ref{Fig_flux_vs_theta}b).  The value of angular momentum of the accreting gas normalized by the local Keplerian angular momentum is lower in the inner region of the disk and higher in the outer region (Fig.~\\ref{Fig_j_zs}).  The falling gas reaches the shock surface formed on the top of the circumplanetary disk.  %%% Gas near the midplane (especially where $\\tilde{R} \\lesssim 0.1$) is almost in Keplerian rotation and hydrostatic equilibrium in the $z$-direction.  In this region, radial velocity is very small and it is difficult for the gas to accrete inward through the disk midplane, as mentioned above. % On the other hand, the gas at $\\tilde{R} \\gtrsim 0.2$ in the midplane shows significant outflow, which eventually escapes from the Hill sphere (see Fig.~\\ref{Fig_streamlines_outward}). % Thus inflow from high altitude and outflow near the midplane co-exist, and they do not interfere with each other.  This means that there is a circulation across the Hill sphere and fresh (protoplanetary) gas is always supplied in the $\\tilde{R} \\gtrsim 0.2$ region. % In addition, the two Lagrangian points L$_1$ and L$_2$, which are often thought to be the most likely points of inflow, are actually the points of outflow, even in the gas accretion stage (Fig.~\\ref{Fig_flux_at_spheres}). % % added in response to the referee's comment 2). This seems consistent with \\citet{Klahr06}, who also performed three-dimensional hydrodynamic simulation with no explicit physical viscosity, but with lower resolution and with radiation.  They showed that accretion mainly occurs via the poles of the planet and no inflow along the equatorial plane, which is quite similar to ours.  But they explained that the circulation in their simulation is driven by accretion heating, which we do not consider.  Thus it is not clear if the driving mechanisms of the circulation are the same. % % added in response to the referee's comment 2). Note that \\citet{Ayliffe09b} shows that the vast majority of the mass flows into the Hill sphere near the equator, which seems to be inconsistent with our result.  There are several differences in setting between theirs and ours, so it is not easy to judge which factor causes the difference.  Disk thickness would be one of the reasons for that, even though our disk thickness is not very different from theirs. Another possibility is gap formation. % When a deep gap is formed, the contribution of the vertically accreting gas would become less significant, which reduces the difference. % However, here we think that viscosity is the main reason for the difference.  When there is explicit viscosity, the circumplanetary disk gas should transfer its angular momentum outward and most of gas would move inward accordingly, but this is not necessarily true in the inviscid limit.  Actually, inviscid simulations by \\citet{Klahr06} showed similar results with ours.  Also, inviscid limit might not be bad because circumplanetary disks are likely to be MRI inactive in most cases \\citep{Fujii11}.   %%% % new paragraph for the new Fig.17 (2011/11/04) Since dust particles tend to settle down toward the midplane, gas accretion flow from high altitude is likely dust-poor gas, which diminishes dust-to-gas ratio in the circumplanetary disk. % On the other hand, since gas in the outflow region ($\\tilde{R} \\gtrsim 0.2$) rotates significantly slower than Keplerian velocity around the planet, dust particles would migrate inward quickly.  Thus the outflow in the midplane is also likely dust-poor gas, which would be the source of solid material in the circumplanetary disk and would enhance dust-to-gas ratio in the disk. % Dust-to-gas ratio in the circumplanetary disk is one of the most important factors for satellite formation processes \\citep{CW06}, thus these two filtering effects would become important for satellite formation. % In addition, studies on size evolution of solid material in protoplanetary disks, such as \\citet{Kobayashi10}, would also be important for the filtering effects. % Further studies are needed on this issue.   %%%% %Between the shock surface and the hydrostatic region with nearly %Keplerian rotation, there is a transition layer. %% %The gas that passed through the shock surface tries to merge into the %circumplanetary disk, but its angular momentum at the shock surface is %smaller than that for Keplerian rotation at the point.  The gas %accordingly moves inward for adjustment. %% %As we see in Fig.~\\ref{Fig_j_zs}, specific angular momentum before the %shock normalized by Keplerian angular momentum decreases with decreasing %radius in $\\tilde{R} \\lesssim 0.1$, so inner region have relatively %larger inward movement in the transition region. %% %The inward stream in the transition region might play an important role %for the net mass accretion toward the planet in circumplanetary disks. %% %This is a kind of layered accretion. %% %This accretion picture should be explored more quantitatively by further %high-resolution simulations, using a code with lower artificial %viscosity such as the one with polar coordinates for numerical grids.  %%% % modified 2011.7.21 As we see in Figure \\ref{Fig_flux_vs_theta}, there is the layer just under the shock at the disk surface where gas moves inward.  This {\\it layered accretion} is explained by the fact that angular momentum of the accreting gas onto the disk surface is lower than that of the Keplerian rotation at the radius (Fig.~\\ref{Fig_j_zs}). % The inward stream in the layer might play an important role for the net mass accretion toward the planet in circumplanetary disks. % This accretion picture should be explored more quantitatively by further high-resolution simulations, using a code with lower artificial viscosity such as the one with polar coordinates for numerical grids.     %%======================================================================= %\\subsection{Difficulties in Obtaining Surface Density of Circumplanetary %  Disks} %% %% %Surface density of circumplanetary disks would be more directly needed %when one would like to consider satellite formation process, rather than %gas accretion rate onto circumplanetary disks as a function of radius. %% %Nevertheless we focused on the accretion rate because of several %reasons. % %% %First of all, effective viscosity, which is the key quantity for %determining surface density in quasi-steady state, is poorly fixed. %This is serious and essential because even the origin of viscosity is %not well understood, although turbulent viscosity driven by %magneto-rotational instability is the most likely source. %% %Even if viscosity can be fixed well, it is hard to determine surface %density by hydrodynamic simulations.  One of the reasons is that low %viscosity, which would be realized in curcim-planetary disks, is %difficult to handle precisely. %% %Also, calculation time would be too long to reach a quasi steady state %for very high-resolution hydrodynamic simulation (see the description of %surface density in \\S \\ref{sec_diskstructure}) % % %On the other hand, gas accretion profile we obtained is a more robust %quantity in the sense that this quantity basically does not chenges even %if surface density does because the accretion flow onto the surface of %circumplanetary disks is mostly supersonic, in other words, no %disturbance in the circumplanetary disks can affect the supersonic %accretion flow, even though the height of the shock surface can change %to some degree. % % %Once the gas accretion rate as a function of distance from parent %planets is obtained, one can apply it to, for example, one-dimensional %circumplanetary disk evolution models.  \\citet{Ward10} and %\\citet{Martin11} actually modeled one-dimensional circumplanetary disk %evolution by assuming distribution of the gas accretion rate, which is %the quantity we obtained in this paper. %% %Our accretion rate can basically apply those kinds of one-dimensional %models directly, even though our result only gives snapshots. % % %%% %Although specific angular momentum of the accretion flow $\\tilde{\\ell}$ %may not be important for determine the disk structure \\citep{Martin11}, %accretion rate $\\tilde{\\dot{M}}$ and evolution are definitely important. %In particular, hysteresis of the accretion rate is important for %satellite formation process via disk temperature and surface density %\\citep{Canup02}. %% %\\citet{Tanigawa07} and \\citet{Ward10} modeled the long term evolution of %the accretion rate onto giant planets by simple one-dimensional models. %However, the important period for satellite formation is the last minute %of the giant planet formation when the gas accretion is significantly %waned and about to cease.  Precise treatment of disk physical processes, %such as dissipation of protoplanetary disks by photoevapolation, viscous %evolution, and gap formation is necessary, and thus is desired to be %examined in future works.    %%======================================================================= %\\subsection{Dust to Gas Ratio of the Accreting Flow} %We have considered gas accretion flow, then we next discuss supply of %solid material into circumplanetary disks. %% %In this giant planet forming stage, large fraction of solid material in %protoplanetary disks should have been consumed to build protoplanets, %but still some dust particles or planetesimals should remain. %% %Supply of solid material to giant-planet systems is important for %satellite formation process and the fraction of heavy elements of the %giant planets in our solar system \\citep[e.g.,][]{Wuchterl00, Saumon04} %as well. % % %What we have found in this paper is that the width of gas accretion band %varies with height (see Fig.~\\ref{Fig_fate} and the corresponding text %in \\S 3.1). %% %If dust is well mixed and the ratio to gas is homogeneous in %protoplanetary disks, net dust to gas ratio of the accreting material is %the same as that of protoplanetary disks. %% %But if the ratio is not homogeneous, the net ratio differs from the %averaged one, and in reality, dust tends to settle down to the disk %midplane and thus the ratio would be enhanced around the midplane %\\citep[e.g.][]{Nakagawa86,Cuzzi93}. %% %According to Fig.~\\ref{Fig_fate}, the width is the narrowest at the %midplane and increases with height until $\\tilde{z} \\sim 0.5$, and keeps %nearly constant above the height. %% %This immediately implies that dust would be depleted in the accretion %flow if the ratio is enhanced around the midplane or the protoplanetary %disk, and that the tendency of the depletion is more noticeable for %larger size dust because sedimentation is more effective for larger %one. %% %This is the opposite sense for the larger content of heavy elements in %the current Jovian planets. % % % %Next we consider planetesimals accretion into circumplanetary disks. %% %Since motion of large planetesimals are usually decoupled from that of %gas, the motion is hardly affected by the gas and thus most of %planetesimals can pass through circumplanetary disks. %% %Even when planetesimals themselves pass through the circumplanetary %disks, dissipation due to the friction between planetesimals and %circumplanetary disk gas would give rise to ablation of the %planetesimals \\citep{Mosqueira03a}, which might be a important source %of solid in circumplanetary disks. %% %If the size of planetesimals is smaller, the planetesimals can be %directly captured by the gas drag with circumplanetary disks. %% %Even though there are some studies on the capturing process of %planetesimals during gas capturing growth phase of giant planets in the %context of accounting for composition of heavy elements in giant planets %\\citep{Zhou07, Shiraishi08}, no quantitative study have been reported %for the last stage of giant planet formation when circumplanetary disks %exist.     %======================================================================= \\subsection{Effects of Gap Formation} % One important issue to be addressed is the effect of gap, which is a lower density annulus region near the planet orbit in the protoplanetary disk. % We observed the gap as a slightly lower density band formed around $\\tilde{x}=0$ in our simulations. % However, in the last stage of giant planet formation, a giant planet would become massive enough to create a deep gap, which would be able to truncate its growth. % Since the deep gap is associated with steep density gradient at the edge, gap formation may affect accretion flow and circumplanetary disk formation. % However, at the edge of such a deep gap, the gas density changes over a radial distance comparable to the disk scale height, while the width of the accretion band onto the circumplanetary disk is much narrower than the scale height (Fig.~\\ref{Fig_fate}). % Therefore, we think that the gap formation would not affect the qualitative feature of the accretion flow. % % Another possible effect of gap formation is the disk size.  As described in \\S \\ref{sec_disk_size}, several mechanisms are proposed to explain the disk size.  If the disk size is determined by the tidal effect as suggested by \\citet{Martin11}, it should not be affected by gap formation, except when the Bondi radius is smaller than the Hill radius. % replaced by the above at 2011/12/12 %Another possible effect of gap formation is the disk size.  As described %in \\S \\ref{sec_disk_size}, several mechanisms are proposed to explain %the disk size.  If the disk size is determined by the tidal effect as %suggested by \\citet{Martin11}, it should not be affected by gap %formation. %% On the other hand, if the disk size is determined by the radially outward velocity of the gas as described in \\S \\ref{sec_disk_size}, lower density around the Hill sphere and thus stronger radial pressure gradient near the disk edge may enhance the outflow from the circumplanetary disk, and the disk size may become smaller. % In any case, effects of gap formation should be examined in future works to check the validity of the results shown in this paper.    % %If a deep gap, which should be expected to form along the orbits of %Jupiter mass planets, is developed in our simulation, resultant steeper %pressure gradient in radial direction would enhance more outward flow in %the disk.  In this sense, $\\tilde{R}_+$ should not be larger than we %obtained even when deep gap is formed.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   In order to understand the structure of circumplanetary disks, we performed high-resolution hydrodynamic simulation and analyzed gas accretion flow in detail. % We confirmed that gas accretion onto circumplanetary disks occurs in a manner that the gas is accreted from high altitude toward the disk surface downward with large incident angle, which was suggested by previous studies \\citep[e.g.,][]{DAngelo03,Bate03}, whereas \\citet{Ayliffe09b} showed that, in terms of mass flux, the downward accretion flow is not significant, which is inconsistent with our results. % replaced by the above at 2011/12/12 %We confirmed that gas accretion onto circumplanetary disks occurs in a %manner that the gas is accreted from high altitude toward the disk %surface downward with large incident angle, which was suggested by %previous studies \\citep[e.g.,][]{DAngelo03,Bate03}. %% We found that the gas that has passed through the shock surface moves inward because its specific angular momentum is smaller than that of Keplerian rotation, whereas gas accretion through the midplane does not occur. % While the net gas accretion across the Hill sphere is inward, outflow through near the two Lagrangian points is observed, although the regions around the two Lagrangian points have been, from the point of view of the potential energy, thought as a main accretion channel from protoplanetary disks. % This outflow was not observed in previous hydrodynamic simulations for Jupiter-sized planets \\citep[e.g.,][]{Bate03,Ayliffe09b}. % Outward radial velocity of the gas near the midplane significantly increases at a radial distance of about 0.2 times the Hill radius from the planet, and the gas can escape from the Hill sphere within a short period of time.   %%% We also obtained the distribution of mass and angular momentum of accreting gas onto the surface of circumplanetary disks.  We found that the accretion rates of mass and angular momentum can be well described by power-law functions. % This distribution would be useful in the study of satellite formation, for example, a radially one-dimensional viscous-evolution model for circumplanetary disks, such as \\citet{Ward10} and \\citet{Martin11}. % Recent development of viscous modeling for MRI turbulence in protoplanetary and circumplanetary disks \\citep{Okuzumi11, Fujii11} would also contribute to construct more realistic models for the circumplanetary disks. % However, in order to understand satellite formation processes, long term evolution of circumplanetary disks, which is determined by the evolution of accretion rate from protoplanetary disks to circumplanetary disks, is necessary and thus global evolution of protoplanetary disks with embedded giant planets until the complete dissipation of protoplanetary disks needs to be examined. % Together with such global models \\citep[e.g.,][]{Tanigawa07,Ward10}, we will be able to obtain long-term evolution of circumplanetary disks, which would provide better understanding of satellite formation processes.   %%% Added by Ohtsuki-san 0705 Our results demonstrate that gas accretion toward and within circumplanetary disks has a complicated vertical structure.  The width of the accretion band toward a circumplanetary disk depends on the initial height of gas elements (Fig.~\\ref{Fig_fate}).  Also, after accretion onto the disk, the gas just below the shock surface migrates inward, while the gas near the midplane moves radially outward (Fig.~\\ref{Fig_CP-disk}).  Such vertical heterogeneity of the flow may have significant influence on the dynamical evolution of solid bodies in the circumplanetary disk, which we will examine in our future work. % added in response to the referee's comment ). We also need to address the effect of viscosity and local calculation, which would affect the flow quantitatively.     %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the"
    },
    "1112/1112.0503_arXiv.txt": {
        "abstract": "Tracing the  cosmic evolution  of the Baryonic  Acoustic Oscillation (BAO) scale with galaxy two point correlation functions is currently the most promising approach to detect dark energy at early times.  A number of ongoing  and future experiments will measure  the BAO peak with  unprecedented accuracy.   We show  based  on a  set of  N-Body simulations  that  the matter  distribution  is  anisotropic out  to $\\sim150\\hMpc$,  far beyond  the  BAO scale  of $\\sim100\\hMpc$,  and discuss implications for the measurement of the BAO. To that purpose we  use  alignment correlation  functions,  i.e., cross  correlation functions  between  high  density   peaks  and  the  overall  matter distribution  measured  along  the  orientation  of  the  peaks  and perpendicular  to  it.   The  correlation  function  measured  along (perpendicular to) the orientation of high density peaks is enhanced (reduced)  by  a  factor   of  $~2$  compared  to  the  conventional correlation function and the location of the BAO peak shifts towards smaller  (larger) scales  if measured  along (perpendicular  to) the orientation of the high density peaks.  Similar effects are expected to shape observed galaxy correlation functions at BAO scales. ",
        "introduction": "\\label{sec:intro} The  Baryonic Acoustic  Oscillations (BAO) constitute a characteristic feature within the large scale structure of the Universe and can serve as  standard ruler  for  constraining the  properties  of dark  energy \\citep[e.g.,][]{Blake-Glazebrook-03, Linder-03,   Seo-Eisenstein-03, Wang-06,   McDonald-Eisenstein-07, Seo-Eisenstein-07, Seo-08, Seo-09, Kazin-Sanchez-Blanton-12}.  The BAO in the  baryon-photon fluid of  the pre-recombination era  imprint the sound horizon distance at decoupling  as a typical scale in the matter correlation   function or   power   spectrum   \\citep{Peebles-Yu-70, Sunyaev-Zeldovich-70, Eisenstein-Hu-99, Bashinsky-Bertschinger-02}. These oscillations were detected  in the cosmic  microwave background \\citep[e.g.,][]{Page-03} and  in the spatial  distribution of galaxies \\citep{Eisenstein-05, Cole-05} and have  been confirmed by a number of subsequent  studies  \\citep[e.g.,][]{Percival-07,  Cabre-Gaztanaga-09, Sanchez-09, Percival-10, Reid-10, Kazin-10}.  The next generation of large galaxy surveys, like the Panoramic Survey Telescope \\& Rapid  Response System (Pan-STARRS, \\citealt{Kaiser-02}), the   Dark  Energy  Survey   (DES,  \\citealt{DES-05}),   the  Baryonic Oscillation Spectroscopic Survey (BOSS, \\citealt{Schlegel-White-Eisenstein-09}), BigBOSS \\citep{Schlegel-11}, the   Hobby   Eberly  Telescope   Dark   Energy  Experiment   (HETDEX, \\citealt{Hill-04})    and    the    space   based    Euclid    mission \\citep{Cimatti-09},  will  cover  volumes  much  larger  than  current datasets, allowing for much more accurate determinations of the BAO.  In the two-point correlation function  the BAO are visible as a unique broad  and  quasi Gaussian  peak  \\citep{Matsubara-04}.  However,  the determination of the shape and location of the peak may be affected by sample variance \\citep{Cabre-Gaztanaga-09, Martinez-09, Kazin-10, Cabre-Gaztanaga-11}, non linear effects  and the bias of  the tracer galaxy  population \\citep{Smith-03,  Crocce-Scoccimarro-06a, Angulo-08, Crocce-Scoccimarro-08,    Seo-08}.     These    difficulties    lead \\cite{Prada-11} to suggest  to use the zero-crossing of  the two point correlation located  at $\\sim 130\\hMpc$  as standard ruler  instead of the  peak  location.   Yet,  several observations do not show  the theoretically  predicted  zero-crossing at all \\citep{  Martinez-09, Kazin-10}.  At this stage it is unclear whether this discrepancy is caused  by  systematic  effects   or  cosmic  variance  or  whether it represents  a   challenge  for  the   concordance  $\\Lambda$CDM  model \\citep{Labini-09, Sanchez-09, Kazin-10}.  One  basic assumption for  interpretation of  the BAO  measurements is that  the matter  distribution  is isotropic  at  the relevant  scales ($\\sim100\\hMpc$). In this work  we use alignment correlation functions \\citep{Paz-Stasyszyn-Padilla-08,   Faltenbacher-09}   to   show   that the amplitudes of the two-point correlation function measured along the orientations of the  high density peaks are larger  than those derived from  spherically averaged (conventional)  clustering analysis  out to scales of $\\gtrsim 150\\hMpc$ and discuss possible effects on measurements of the BAO. \\begin{figure*}[t] \\epsfig{file=fig1.eps,width=0.95\\hsize } \\caption{\\label{fig:z0} {\\it Upper  panels:} 3 dimensional two point cross  correlation functions  between FoF groups of  the indicated mass  and the  overall matter  density field  represented  by 10\\% random subset  of all simulation particles.  Solid  lines show the {\\em conventional} correlation functions.  Dotted and dashed lines show   the  {\\em  alignment}   correlation  functions   along  and perpendicular to the orientations  of the FoF groups, respectively.   {\\it   Lower   panels:}  Difference   between conventional and alignment  correlation functions above. The shaded  areas in the upper and lower panels   give  the cosmic  variance.  Poisson  errors  are negligible and are not shown here.} \\end{figure*} ",
        "conclusions": "The findings presented here are based on N-Body simulations however qualitative similar effects are expected to shape the observed galaxy correlation functions at scales up to $150\\hMpc$. Potential observational implications are:\\\\  $\\bullet$ {\\em  Large-scale anisotropy  and structure of two point correlation  functions:}   Our  results  indicate   that  matter  is distributed anisotropically out to separations far larger then the BAO scale.   At any  given  pair separation  the conventional  correlation function is the average of the alignment correlation function over the whole  range of  $\\theta$.   We find  that  the alignment  correlation function measured  along high density  peaks does not fall  below zero (for separations $\\le 150\\hMpc$)  which is counterbalanced by negative amplitudes at much  smaller separations for measurements perpendicular to  it. At the BAO scale galaxy two point correlation functions can be interpreted as the average of a more highly clustered component along the direction of high density peaks and a less clustered component perpendicular to it.\\\\ $\\bullet$  {\\em Shape and  location of the BAO peak:}  At BAO  scales the amplitudes of the correlation functions between high density peaks and the overall  matter distribution are significantly  higher if measured along the orientation of the peaks.  In this case the BAO peak is composed of a hump on top of the declining but still positive correlation function.  The  signal perpendicular to the  orientations is dominated by the BAO hump itself with negligible underlying clustering signal.  The location of the BAO peak is found at somewhat smaller (larger) scales if measured parallel  (perpendicular) to the orientation  of high density peaks. For the measurements along the density peak orientations the BAO hump is transformed into a plateau-like feature.   If the  survey  volume is  dominated  by large  filamentary structures the  shape of the conventional  correlation function should be close to  that found for the parallel signal  shown here.  This may relate to a study by \\cite{Kazin-10} who find no apparent peak in 6\\% of their mock samples.\\\\ $\\bullet$ {\\em Zero crossing and large scale power:} Several publications report the  non-detection of zero-crossing out to scales of  $250\\hMpc$ which  indicates (unexpected) large  scale power \\citep[e.g.,][]{Martinez-09,Kazin-10,Labini-09}.    We   find  similar results for the correlation  functions measured along the orientations of high  density peaks.  Furthermore, the  alignment cross correlation function  measured along  the orientations  of high  density  peaks at $z=0.6$ shows zero-crossing in contrast to the behavior at $z=0$.  The WiggleZ    redshift-space     correlation    function    at    $z=0.6$ \\citep{Blake-11}  shows  a  crossover  as  well.   Whether  these analogies are coincidental remains to be explored in future work.\\\\ $\\bullet$ {\\em Direct measurement of anisotropy:} With the  advent of enormous  cluster catalogs \\citep[e.g.,][]{Hao-10, Gilbank-11} it  should in principle be possible  to directly measure the large  scale anisotropies with alignment  correlation functions if cluster orientations  can be determined with  sufficient accuracy.  An anisotropy signal  may even be extracted simply  by using orientations of the cluster central  luminous red galaxies because the orientations of    central    galaxies   and    host    systems   are    correlated \\citep{Faltenbacher-09, Okumura-Jing-Li-09, Schneider-11}. \\\\ $\\bullet$ {\\em Improvement of BAO measurements:} If the  directional effects reported  here are observable it  would be worthwhile considering  to measure the  BAO peak perpendicular  to the orientations of  galaxy clusters (or  luminous red galaxies)  since in this direction the BAO peak is better confined."
    },
    "1112/1112.0029_arXiv.txt": {
        "abstract": "We study the evolution of the Star Formation Rate Function (SFRF) of massive ($M_\\star>10^{10} M_\\odot$) galaxies over the $0.4<z<1.8$ redshift range and its implications for our understanding of the physical processes responsible for galaxy evolution. We use multiwavelength observations included in the GOODS-MUSIC catalogue, which provides a suitable coverage of the spectral region from 0.3 to 24 $\\mu$m and either spectroscopic or photometric redshifts for each object. Individual SFRs have been obtained by combining UV and 24 $\\mu$m observations, when the latter were available. For all other sources an ``SED fitting'' SFR estimate has been considered. We then define a stellar mass limited sample, complete in the $M_\\star>10^{10} M_\\odot$ range and determine the SFRF using the $1/V_{\\rm max}$ algorithm. We thus define simulated galaxy catalogues based on the predictions of three different state-of-the-art semi-analytical models of galaxy formation and evolution, and compare them with the observed SFRF. We show that the theoretical SFRFs are well described by a double power law functional form and its redshift evolution is approximated with high accuracy by a pure evolution of the typical SFR (SFR$^\\star$). We find good agreement between model predictions and the high-SFR end of the SFRF, when the observational errors on the SFR are taken into account. However, the observational SFRF is characterised by a double peaked structure, which is absent in its theoretical counterparts. At $z>1.0$ the observed SFRF shows a relevant density evolution, which is not reproduced by SAMs, due to the well known overprediction of intermediate mass galaxies at $z\\sim2$. Semi-analytical models are thus able to reproduce the most intense SFR events observed in the GOODS-MUSIC sample and their redshift distribution. At the same time, the agreement at the low-SFR end is poor: all models overpredict the space density of SFR $\\sim 1 M_\\odot/yr$ and no model reproduces the double peaked shape of the observational SFRF. If confirmed by deeper IR observations, this discrepancy will provide a key constraint on theoretical modelling of star formation and stellar feedback. ",
        "introduction": "The evolution of the star formation rate (SFR) over the cosmic time is a fundamental constraint for every theory of galaxy formation and evolution (see e.g. \\citealt{Hopkins04, HopkinsBeacom06}). Estimating SFRs for individual galaxies is a complex task, owing to the uncertainties involved in the reconstruction of this quantity from observational data. It is widely accepted that dusty molecular clouds are the main sites for star formation: this implies that newly born stars are subject to significant dust attenuation, until they are able to escape or disrupt their parent cloud. Young OB stars emit a considerable amount of energy in the restframe ultraviolet (UV) band, which has thus been considered a key waveband for recovering the instantaneous SFR in external galaxies. However, dust absorbs UV photons, heats up and re-emits this energy as thermal radiation in the Infrared (IR) bands. For this reason, SFR estimates based on UV luminosity include correction factors to account for dust attenuation and re-emission \\citep{Calzetti94, Kennicutt98,Bell03}. Dust attenuation is particularly relevant for star forming galaxies, where the dust emission peak is the dominant component of the galactic spectral energy distribution (SED, see e.g. \\citealt{Calzetti00}), not to mention the extreme cases of sub-mm galaxies (see e.g. \\citealt{Chapman04}). The study of the cosmic IR background shows that the global energy emitted by galaxies in the IR is comparable to the direct starlight emission, detectable in the optical (see e.g. \\citealt{Lagache05}), clearly showing that a significant fraction of star formation activity is expected to be heavily extinguished and detectable only in the IR.  These uncertainties become more relevant, when SFR estimates for galaxies covering a wide redshift range are considered, since both galaxy physical properties and dust properties are expected to evolve with cosmic time \\citep{Maiolino04, Fontanot09a, Gallerani10, Fontanot11c}. Combining UV information with supplementary information from direct observations in the IR region has thus been proposed as the best tool to account for the total SFR \\citep{Kennicutt98,Bell07}. In particular, observations at $24 \\mu$m have shown to be extremely useful for estimating the global IR luminosity and the instantaneous SFR \\citep{Papovich07,Santini09}. Thanks to the advent of the {\\it Herschel} Space Observatory, we will be able to constrain total IR luminosities, by directly sampling the peak of the thermal emission. In fact, it has been recently shown by \\citet{Rodighiero10b}, that the combination of $24 \\mu$m data with {\\it Herschel} observations at longer wavelengths represents a very promising tool for a more accurate determination of SFRs of individual galaxies up to $z\\sim3$.  Nowadays, our view of the evolution of galaxy properties has substantially changed thanks to the advent of multiwavelength surveys. If the spectral sampling is fine enough, these catalogues can be used to infer the physical properties also in the absence of spectroscopic information. In recent years several groups (see e.g. \\citealt{Fontana04,Panter07}) developed a number of ``SED fitting'' algorithms: these codes compare the available photometry for individual sources with synthetic SED libraries and the best-fit model template is chosen by means of a $\\chi^2$ minimisation. The resulting estimate for redshift (the so-called photometric redshift), stellar mass ($M_\\star$), dust extinction and SFR are the most widely used results of this procedure and have been of fundamental importance for our understanding of galaxy evolution.  The common interpretation of these results is connected to the so-called ``downsizing'' scenario, in which the star formation shifts from high mass to low mass galaxies as redshift decreases (first introduced by \\citealt{Cowie96}). This picture has been recently revised by \\citet{Fontanot09b} by showing that the typical errors associated with the estimate of the physical quantities has to be taken into account, when comparing the observed evolutionary trends with the predictions of theoretical models. In particular, \\citet{Fontanot09b} concluded that the discrepancies seen between models and data for massive galaxies are not significant, if model predictions are convolved with typical observational errors, while the strongest discrepancies between model and data are seen for low-to-intermediate mass galaxies.  The importance of the study of the redshift evolution of the cosmic SFR for galaxies at a given mass range has been widely recognised by a number of authors \\citep{Noeske07, Elbaz07, Zheng07, Drory08, Dunne08, Santini09, Gilbank11}. These results and the evolution of the stellar mass density provide fundamental information about the evolution of the global process of galaxy formation. The latter approach is best complemented by the analysis of the stellar mass function (i.e. the volume density of galaxies as a function of stellar mass), which allows us to characterise galaxy evolution as a function of both redshift and stellar mass. Despite the wealth of information already available in the literature about the stellar mass function and its redshift evolution, very little is known about the corresponding Star Formation Rate Function (SFRF hereafter) and its evolution: \\citet{Bell07} studied its evolution in the $0.2<z<1.0$ redshift range in the COMBO-17 survey, \\citet{Schiminovich07} presented the distribution of SFRs in the SDSS survey and \\citet{Bothwell11} considered the distribution of SFRs in the Local Volume (defined as an 11 Mpc radius around the Milky Way) combining GALEX and {\\it Spitzer} observations. This is mainly due to the uncertainties involved in the SFR estimate, but also to the difficulties in correctly accounting for the completeness correction needed, especially at the low-SFR end.  In this paper, we focus on a new treatment of the GOODS-MUSIC data, by focusing on a mass complete sample of galaxies. In this way, we are able to fully characterise the various levels of star formation in these objects and follow their redshift evolution. Moreover, since the strongest starbursts are hosted by galaxies above the chosen mass limit, the high-SFR end of the resulting SFRF provides fundamental information about the redshift evolution of this particular galaxy class. Therefore, this paper represents a complementary study of the evolution of SFRs in galaxies of different stellar mass with respect to \\citet{Santini09}, to which we refer the reader for the analysis of the evolution of the cosmic SFR and specific SFR. A detailed comparison of the same catalogue with other estimates for the evolution of SFR as a function of redshift and stellar mass has been presented in \\citet{Fontanot09b} (see references herein). In the present paper we focus on the determination of the SFRF and compare results with the predictions of state-of-the-art semi-analytical models of galaxy formation and evolution.  This paper is organised as follows. We present the GOODS-MUSIC catalogue in sec.~\\ref{sec:dataset} and the theoretical predictions in sec.~\\ref{sec:models}; we present our results in sec.~\\ref{sec:results}; finally we discuss our conclusions in sec.~\\ref{sec:summary}. In the original GOODS-MUSIC papers a Salpeter initial mass function (IMF) is mostly used; throughout this work, we assume a Chabrier IMF and we rescale all stellar mass and SFR estimates, accordingly.  We also adopt the $\\Lambda$CDM concordance cosmological model ($H_0 = 70 \\, km/s/Mpc$, $\\Omega_m=0.3$ and $\\Omega_\\lambda=0.7$). ",
        "conclusions": "\\label{sec:summary}  In this paper we consider the GOODS-MUSIC sample to study the SFR function and its redshift evolution in the redshift range $0.4<z<1.8$, just after the peak in the cosmic SFR, a fundamental epoch to understand the physical processes responsible for the decline of star forming activity as a function of galaxy stellar mass. In order to reduce the effect of completeness corrections, we restrict our analysis to galaxies with $M_\\star > 10^{10} M_\\odot$, since at these stellar masses the GOODS-MUSIC is expected to be a complete sample. We thus compute the SFRF using SFR estimates coming from both SED fitting and/or a combination of UV light with $24 \\mu$m observations. We compare the resulting SFRF with the predictions of 3 independent semi-analytical models of galaxy formation and evolution (\\citealt{Wang08,Somerville08} and \\morgana) and we obtain the following results:  \\begin{itemize} \\item{The agreement between the observed and predicted high-SFR end (SFR$>10 M_\\odot/yr$) of the SFRF is good over the whole redshift range (once the observational errors are taken into account). This result implies that theoretical models are able to reproduce the space density evolution of the strongest starbursts. It is worth noting that the predictions of the three models are remarkably similar in this SFR range.} \\item{The SFRF of massive galaxies, as seen by the GOODS-MUSIC is characterised by a complex evolution both in the typical SFR and in the normalisation. Theoretical models are able to cope with the SFR$^\\star$ evolution, but they predict a negligible density evolution. In particular, all models overpredict the space density of SFR $ \\sim 1 M_\\odot/yr$ galaxies. We interpret this behaviour as due to the well documented excess of intermediate mass galaxies ($10^{10} M_\\odot < M_\\star < 10^{11} M_\\odot$) in SAMs \\citep{Fontanot09b}.} \\item{The observed low-SFR end of SFRF is characterised by a double peaked shape. Despite the uncertainties in the determination of such low SFRs from SED fitting methods, it is worth noting that none of the theoretical model we consider is able to reproduce this feature, and all of them predict a smooth decrease of the space density of low-SFR galaxies, with a well defined power-law slope $\\alpha$. Since this is also the SFR range where models differ most, each of them predicting a different value for $\\alpha$, we stress that future observations providing stronger constraints on the low-SFR end of the SFRF will be of fundamental importance to understand the physical mechanisms responsible for the decline of SFR since $z\\sim2$ in massive galaxies.} \\end{itemize}  The main result of this paper lies in the analysis at the high-SFR end, since this region is completely sampled by $24 \\mu$m observations, which are critical for the recovery of the SFR levels. Also the evolution of the typical SFR$^\\star$ is well sampled by means of IR data out to $z \\sim 1.0$, and allow us to put strong constraint on this relevant discrepancy seen between data and models. Unfortunately, for lower SFR levels this information is not currently available, and we have to rely entirely on the results of SED fitting. Our tests suggest that the use of SED fitting results does not introduce any obvious distortion in the SFR distribution at intermediate SFR values, but we could not completely exclude the double peaked feature in the SFRF to be due to systematics in the SED fitting algorithm. From optical photometry alone is not possible to uniquely associate SFR $ \\sim 1 M_\\odot/yr$ galaxies (the region were models and data differ most and roughly corresponding to the region between the two peaks of the SFRF) with known populations in the colour-magnitude diagram, i.e. the green valley or the blue cloud. The red $U-V$ colours, with a relatively small scatter, of galaxies belonging to the putative secondary peak provide only a qualitative support to a scenario where star formation in massive galaxies decrease at a slower pace with respect to theoretical expectations. If confirmed, this double peaked shape of the SFRF would thus represent a critical discrepancy between models and observations, since all SAMs we consider predict a smooth, featureless, single slope power-law shape for the SFRF. Modifications in the assumed star formation law, taking into account the relation between star formation and the atomic and molecular gas content of galaxies may alleviate this tension, since the properties of model galaxies with low SFR levels are very sensitive to this modeling layer \\citep[see e.g.]{Lagos11}.  On the other hand, thanks to the new facilities (e.g. {\\it Herschel}), future observations in the IR region, able to characterise the properties of the IR peak, will finally determine the SFR levels of these objects and the relevance of this feature. This is a fundamental task, since the low-SFR end of the SFRF is the region where SAM predictions differ most and therefore provides potentially strong constraints to the different approaches to star formation and stellar feedback."
    },
    "1112/1112.0359_arXiv.txt": {
        "abstract": "We present the results of the SuperLupus Survey for transiting hot Jupiter planets, which monitored a single Galactic disk field spanning 0.66 deg${^2}$ for 108 nights over three years.  Ten candidates were detected: one is a transiting planet, two remain candidates, and seven have been subsequently identified as false positives.  We construct a new image quality metric, $S_j$, based on the behaviour of 26,859 light curves, which allows us to discard poor images in an objective and quantitative manner.  Furthermore, in some cases we are able to identify statistical false positives by analysing temporal correlations between $S_j$ and transit signatures.  We use Monte Carlo simulations to measure our detection efficiency by injecting artificial transits onto real light curves and applying identical selection criteria as used in our survey. We find at 90\\% confidence level that 0.10$^{+0.27}_{-0.08}$\\% of dwarf stars host a hot Jupiter with a period of 1-10~days.  Our results are consistent with other transit surveys, but appear consistently lower than the hot Jupiter frequencies reported from radial velocity surveys, a difference we attribute, at least in part, to the difference in stellar populations probed.  In light of our determination of the frequency of hot Jupiters in Galactic field stars, previous null results for transiting planets in open cluster and globular cluster surveys no longer appear anomalously low. ",
        "introduction": "\\label{sec:intro}  The discovery of short-period, giant extrasolar planets \\citep{1995Natur.378..355M, 1996ApJ...464L.147M} provided the exciting potential for large numbers of planets to be discovered by the transit method, as these ``hot Jupiters'' have a $\\sim$10\\% geometric probability of transiting, and do so every few days.  However early predictions greatly overestimated the actual discovery rate \\citep{2003ASPC..294..361H}.  This discrepancy resulted from simplistic assumptions and a misunderstanding of the effects that systematic noise would play in lowering detection efficiency \\citep{2006MNRAS.373..231P}.  It also resulted from an over-estimation of the frequency of hot Jupiters.  More accurate, and lower, predictions were provided in \\citet{2008ApJ...686.1302B}, where it was noted that objective and quantifiable detection criteria were required for more robust inferences of planet frequencies.  By adopting such objective and quantifiable detection criteria, we propose that the frequency of hot Jupiters in the field has been overestimated by a factor of three \\citep{2011OHP...11.01008}.  The SuperLupus Survey was established to detect hot Jupiters in a field positioned just above the Galactic plane ($b=11^{\\circ}$), and also to determine the fraction of stars that host hot Jupiters in this typical Galactic field.  The survey design and data analysis was constructed so as to fulfill both of these objectives.  In Section~\\ref{observations} of this paper, we describe the SuperLupus observations and data reduction.  The photometry is described in Section~\\ref{photometry}.  The criteria for candidate selection are detailed in Section~\\ref{sec:detection}, and the 10 identified candidates are described and analyzed in Section~\\ref{sec:candidates}.  In Section~\\ref{mc_simulatinos}, we set out the details of the Monte Carlo simulations performed to calculate the detection efficiency and the effective number of stars probed for planets in the SuperLupus Survey.  In Section~\\ref{frequency}, we apply this efficiency to the actual results of the survey to determine the fraction of stars in the field that host a hot Jupiter.  Finally, in Section~\\ref{discussion} we summarize and discuss the implications of our results. ",
        "conclusions": "\\label{discussion}  The frequency of hot Jupiters in the Galaxy is an important quantity that will ultimately provide a constraint on models of planet formation and migration.  It has been suggested that the frequency of hot Jupiters in globular cluster environments is lower than that of field stars, and that this may be due to crowding or low metallicity affecting planet formation, migration, or survival \\citep{2000ApJ...545L..47G,2005ApJ...620.1043W}. Our results indicate, however, that there is little statistical disagreement between hot Jupiter frequencies in cluster and non-cluster environments, even if one of the remaining candidates turns out to be a genuine planet.  This is consistent with the work presented in \\citep{2011...arXiv:1009.3013v1} which concludes that there is no evidence to support that open clusters have a lower frequency of hot Jupiters.  Initial estimates for planet yields from transit surveys turned out to be far in excess of the actual discovery rate \\citep{2003ASPC..294..361H}.  One of the many factors that led to this over-estimation was the adoption of the hot Jupiter frequency derived from early radial velocity surveys, which as we have shown is higher than is found from transit surveys. The {\\it Kepler\\/} result of 0.37$\\%$ \\citep{2011arXiv1103.2541H} obviously provides a robust statistic for hot Jupiter frequencies due to the high detection efficiency of that survey.  However the result should be adopted cautiously when calculating expected yields from typical transit surveys, as it is based on a sample of ``solar type'' stars drawn from Kepler target stars, rather than the ensemble field stars monitored by most blind transit surveys."
    },
    "1112/1112.4670_arXiv.txt": {
        "abstract": "We constrain properties of cluster haloes by performing likelihood analysis using lensing shear and flexion data. We test our analysis using two mock cluster haloes: an isothermal ellipsoid (SIE) model and a more realistic elliptical Navarro-Frenk-White (eNFW) model. For both haloes, we find that flexion is more sensitive to the halo ellipticity than shear.  The introduction of flexion information significantly improves the constraints on halo ellipticity, orientation and mass. We also point out that there is a degeneracy between the mass and the ellipticity of SIE models in the lensing signal. ",
        "introduction": "The properties of galaxy and cluster haloes are of great interest in cosmology, and can be powerful tests of the cosmological paradigm and the nature of dark matter. Two important parameters that describe a dark matter halo are its mass and shape, which are related to many physical processes, such as the growth and merging history \\citep{1993MNRAS.264..201K,2005Natur.435..629S}.  Models of dark matter haloes beyond the spherical approximation are favored by many numerical simulations \\citep{2002ApJ...574..538J, 2004IAUS..220..421S,2004ApJ...611L..73K, 2006MNRAS.367.1781A} and observations \\citep{2000A&A...364..377R, 2004ApJ...601..599L, 2005ApJ...625..108D, 2006ApJ...645..170S, 2010ApJ...718..762W}. Furthermore, numerical simulations with different assumptions predict different properties of dark matter haloes \\citep[e.g.][]{2002sgdh.conf..109B, 2005ApJ...627..647B,2007MNRAS.380...93W}.  Current models based on N-body simulations, semi-analytic models or hydrodynamic simulations can predict several halo properties, but several ingredients of these models remain uncertain. A precise understanding of halo properties such as mass and ellipticity is important to confirm and improve the existing models of galaxy formation and probe the physical nature of dark matter.  Gravitational lensing is a powerful tool to study mass distributions, independent of the nature or dynamical state of the matter \\citep[see][for reviews]{1999ARA&A..37..127M,2006glsw.book..269S, 2008PhR...462...67M}. Galaxy-Galaxy Lensing (GGL) is concerned with the mass associated with galaxies and dark matter haloes in which galaxies reside \\citep{1984ApJ...281L..59T, 1996ApJ...466..623B, 1998ApJ...503..531H, 2004ApJ...606...67H, 2004AJ....127.2544S, 2006MNRAS.368..715M}. The distortion caused by a single galaxy cannot be detected, but the statistics of many foreground-background pairs yield a detectable signal for a population of galaxies. \\citet{1996ApJ...466..623B} discovered a significant GGL shear signal. \\citet{1997ApJ...474...25S} developed a maximum likelihood analysis that can constrain the halo properties of the lens galaxy populations through GGL, allowing to estimate the mean velocity dispersion and the characteristic scale for a non-singular isothermal sphere halo model.  Flexion as the gradient of the projected mass density, is sensitive to the small-scale variations of mass distributions \\citep{2002ApJ...564...65G,2005ApJ...619..741G,2006MNRAS.365..414B}. Different techniques have been developed to measure flexion \\citep[see][for examples]{2006ApJ...645...17I, 2007ApJ...660..995O, 2008A&A...485..363S,2011MNRAS.416.1616F}. Recently, \\citet{2011MNRAS.412.2665V} applied the shapelets technique on the COSMOS survey, and \\citet{2011ApJ...736...43C} introduced a new method (so-called analytic image model) to analyse lensing flexion images. It has been noted that flexion can contribute to cosmology in several aspects, such as exploring the mass distribution of dark matter haloes of galaxies and clusters, especially substructures \\citep{2009MNRAS.395.1438L, 2010MNRAS.409..389B, 2010arXiv1008.3088E}. \\citet{2011MNRAS.412.1023H} also propose to reduce the distance measurement errors of standard candles using lensing shear and flexion maps.  In \\citet{2009MNRAS.400.1132H, 2011A&A...528A..52E, 2011MNRAS.417.2197E}, the ellipticity of a galaxy halo has been studied with flexion. It was found that the constraints from flexion are tighter than those from shear. In \\citet{2005ApJ...619..741G}, Galaxy-Galaxy lensing Flexion (GGF) has been studied using the distribution function of orientations between the line connecting foreground and background pairs and the flexion of the background galaxies. A GGF signal has been detected by \\citet{2007ApJ...666...51L} using images taken by HST ACS in the cluster Abell 1689. Moreover, combining shear and flexion information provides tighter constraints on halo properties by studying mass distribution on different scales \\citep{2010arXiv1008.3088E,2010MNRAS.404..858S,2011MNRAS.412.1023H}.  In this paper, we combine shear and flexion data to constrain the properties of dark matter haloes. The tangential shear is mainly sensitive to the mass of the halo, whereas flexion is sensitive to the halo ellipticity.  Moreover the usable number density of flexion data is relatively low, since the flexion signal drops faster than the shear signal with the angular distance to the centre of the halo. It is thus not sufficient to constrain the halo properties with shear or flexion alone. Therefore, we propose to take advantage of both shear and flexion in our analysis. In particular, we use the angular positions, tangential shear and flexion of galaxies in our likelihood functions. A singular isothermal ellipsoid model and an elliptical NFW model for a galaxy or cluster halo with 3 parameters (mass, ellipticity and orientation) are adopted in this paper. In Sec. 2, we recall the basic lensing equations. Our likelihood function is introduced in Sec. 3. We perform numerical tests of our method and results in Sec. 4 and present our conclusions in Sec. 5.  Throughout this paper, we adopt a $\\Lambda$CDM model with $\\Omega_{\\Lambda}=0.75$, $\\Omega_{\\rm m}=0.25$, and a Hubble constant $H_0 = 73$ km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "conclusions": "In this paper, we study the potential of weak lensing flexion in the study of galaxy cluster haloes. We use mock data including shear, $\\cal F$ flexion, $\\cal G$ flexion and redshift information.  We find that the inclusion of flexion significantly improves the estimate of foreground haloes parameters, although the details are model-dependent. In particular, in the case of a SIE halo, the presence of a mass-ellipticity anti-correlation implies that analyses where the halo is incorrectly assumed to be spherical will overestimate the halo mass.  On the other hand, we did not find significant correlation between the halo mass and ellipticity in the eNFW model.  The noise in the mock data is determined by the dispersion of the intrinsic shear and flexion distributions, and by the density of background galaxies.  After applying stringent cuts in the data, we are left with a galaxy density of roughly 10 arcmin$^{-2}$.  Our approach assumes that the flexion estimators are linear in the flexion observables and the point spread function can be removed without producing a bias. In reality, the noise of flexion estimators can be complicated and non-Gaussian. An accurate study of flexion noise is important to evaluate the estimated error.  The analysis considers a single cluster halo. This approach is not possible if the number density of background galaxies is low. In that case, stacking of several halo fields can be used to increase the number of background images and thus the signal-to-noise. That approach requires the alignment of the major axis of various foreground galaxies and selecting haloes with similar properties, for example similar shapes of their central galaxies. Such stacking analysis can constrain the halo shapes more tightly, and as a function of other halo properties, e.g. mass.  Our results emphasize that a combined weak lensing analysis will be a useful technique for precise measurements of the properties of galaxy or cluster haloes from future weak lensing surveys, such as EUCLID."
    },
    "1112/1112.3116_arXiv.txt": {
        "abstract": "Recent advances in N-body simulations of dark matter haloes have shown that three-parameter models, in particular the Einasto profile characterized by $d\\ln\\rho(r)/d\\ln r\\propto r^\\alpha$ with a shape parameter $\\alpha \\lesssim 0.3$, are able to produce better fits to the 3D spatial density profiles than two-parameter models like the Navarro, Frenk and White (NFW), and Moore et al. profiles.  In this paper, we present for the first time an analytically motivated form for the 2D surface mass density of the Einasto family of dark matter haloes, in terms of the 3D spatial density parameters for a wide range of the shape parameter $0.1$ $\\leq$ $\\alpha$ $\\leq$ $1$. Our model describes a projected (2D) Einasto profile remarkably well between $0$ and $(3-5)$ $r_{200}$, with errors less than $0.3$ per cent for $\\alpha$ $\\leq 0.3$ and less than $2$ per cent for $\\alpha$ as large as 1. This model (in 2D) can thus be used to fit strong and weak lensing observations of galaxies and clusters whose total spatial(3D) density distributions are believed to be Einasto-like. Further, given the dependence of our model on the 3D parameters, one can reliably estimate structural parameters of the spatial (3D) density from 2D observations.  We also consider a Sersic-like parametrization for the above family of projected Einasto profiles and observe that fits with a Sersic profile are sensitive to whether one fits the projected density in linear scale or logarithmic scale and yield widely varying results. Structural parameters of Einasto-like systems, inferred from fits with a Sersic profile, should be used with caution. ",
        "introduction": "Gravitational lensing signatures are a response to the projected surface mass column density of matter $\\Sigma(\\vec R)$ along the line of sight in galaxies and clusters. Upon a suitable deprojection and circular averaging, an estimate of $\\Sigma (R)$ can, in principle, be used to trace the spherically averaged 3D density profile $\\rho(r)$.  In the past few years, N-body simulations have shown [\\cite{P03}, \\cite{Nav04} (Nav04), \\cite{M06} (M06), \\cite{S09}(S09)] that three-parameter models, especially the Einasto \\citep{Einasto} profile and the \\cite{PS97} de-projected Sersic profiles (PS97), are able to produce better fits to the 3D density profiles of galaxy and cluster-sized dark matter haloes than two-parameter models (\\cite{NFW} $\\&$ (1996) (NFW), \\cite{Moore99}). While the PS97 profile has a well known 2D sky projected form - the Sersic \\citep{Sersic} profile, there has been no such analytical counterpart for the Einasto profile.  In this paper, we present a very good approximation for the 2D projection of the Einasto family of 3D profiles. Thus, if the 3D total mass density is believed to be Einasto-like, our model can be used to parametrically describe the projected 2D surface mass densities of galaxies and clusters in the weak and strong lensing regimes. However, note that even upon radial averaging to smooth out the substructure, not all haloes subscribe to an Einasto profile. For the rest of this paper, we will limit the discussion to the 2D projection of the special case where the 3D profile is Einasto-like.  In 3D, the functional form of the Einasto \\citep{Einasto} profile is given by: \\begin{align}\\label{einasto} \\ln[\\frac{\\rho(r)}{\\rho_s}]= -b[(\\frac{r}{r_s})^{\\frac{1}{n}} -1] \\end{align} where, $\\rho(r)$ is the 3D (spatial) density at $r$, $n$ (or $\\alpha=\\frac{1}{n}$) is the shape parameter, $b$ is a function of $n$, $\\rho_s$ the spatial density at a scale radius $r_s$ and $\\rho(0) = \\rho_s e^b$.  In 2D, the Sersic \\citep{Sersic} profile, which has been used to describe the projected surface brightness profiles of galaxies, is similar in form to the Einasto profile in 3D \\eqref{einasto} and is given by: \\begin{align}\\label{sersic} \\ln [\\frac{\\Sigma_S(R)}{\\Sigma_{R_E}}] = -q[(\\frac{R}{R_E})^{\\frac{1}{m}} -1] \\end{align} such that $\\Sigma_S(0)$$=$$\\Sigma_{R_E}~e^q$ where, $R$ is the projected distance in the plane of the sky, $\\Sigma_{R_E}$ is the line of sight projected surface brightness at a projected scale radius $R_E$, which can be defined to be the half-light radius of a Sersic profile under the condition: $q=2m-0.3333+0.009876/m$ (PS97) with $m$ (or $\\lambda$$=$$\\frac{1}{m}$) characterizing the shape of the Sersic profile.  The parameter $\\alpha$$=$$\\frac{1}{n}$ in \\eqref{einasto} defines the shape of the Einasto profile. In the first ever fits to N-body haloes with a Einasto profile, Nav04 found an average value of $0.172$$\\pm$$0.032$ for a wide range of halo masses from dwarfs to clusters. For galaxy-sized haloes \\cite{Pr06} found $0.133$$\\leq$$\\alpha$$\\leq$$0.167$. \\cite{HW08} observed an evolution of $\\alpha$ with mass and redshift ($z$) in the Millennium Simulation ($MS$) of \\cite{Spring05} and found $\\alpha$$\\sim$$0.17$ for galaxy and $\\sim$ $0.23$ for cluster-sized haloes. A similar trend is supported by \\cite{G08} where $\\alpha$$\\sim$$0.3$ for the most massive clusters in $MS$. Hence, although in this paper we discuss our approximation to a projected Einasto profile for $0.1$$\\leq$$\\alpha$$\\leq$$1$, particular attention is drawn to the domain $\\alpha$$\\lesssim$$0.3$; where as we shall show in \\S 3, the errors due to our approximation out to $30$$r_{-2}$ ($\\sim$$(3-5)$$r_{200}$) are $<$ $0.3\\%$.  The parameters $r_s$, $b$ and $n$ in the Einasto profile are not independent. It can be seen, that in terms of a dimensionless length $X=\\frac{r}{r_s}$, the logarithmic slope of the density profile is given by: \\begin{align} \\beta = \\frac{d \\ln (\\rho/\\rho_s)}{d \\ln X}= -\\frac{b}{n} X^{1/n} \\end{align} Nav04 chooses to define $r_s$ such that  $b/n=2$ (the isothermal value of $\\beta=-2$) at $r=r_s$ and hence label $r_s$ as $r_{-2}$ and $\\rho_s$ as $\\rho_{-2}$. Another approach is to use the convention of M06, requiring $r_s$ to include half the total mass. They quote a numerical estimate of $b$$=$$3n-0.333+0.0079/n$. In this paper we will follow the Nav04 parametrization.  Since the Sersic profile describes the 2D surface brightness profiles of galaxies reasonably well, a natural question is: does the 2D surface mass density of Einasto-like 3D dark matter haloes also follow a Sersic-like description? For the Nav04 simulations, \\cite{M05} (M05) have shown that a Sersic function does produce fits with acceptable errors in the range of $r_{conv}$ to $r_{200}$, where $r_{conv}$ is the minimum radius of convergence and $r_{200}$ is the virial radius in the N-body simulations of P03 and Nav04.  In this paper, we focus on the analytical description of the projected Einasto profile. In order to see how well a projected Einasto profile is described by a Sersic profile, we numerically project \\eqref{einasto} and find that a Sersic profile produces acceptable fits only in a limited range of the projected radius $R$. We shall show in \\S 3 that this range also depends on whether one fits the numerically projected Einasto profile in linear scale (large errors with increasing R) or log scale (large errors with decreasing R) of density. It is thus evident that the Sersic profile does not give an adequate description of a projected Einasto profile for all $R$.  The choice of Sersic fit (log or linear) may, for example, have possibly strong implications in the strong and weak lensing regimes respectively, yielding incorrect results. It is with this perspective, and the observation that the projection of an Einasto profile (i.e. integral of a Sersic-like function) is not a Sersic-like function but rather a Gamma-like function, we present a derivation of the surface mass density (\\S 2.1) for the Einasto profile.  Throughout this paper we will describe 3D spatial parameters as $r$, $r_s$ (or $r_{-2}$), $\\rho_s$ (or $\\rho_{-2}$), $b$ and $n$ (or $\\alpha = 1/n$))and 2D projected parameters as $R$, $R_E$, $q$ and $m$ (or $\\lambda = 1/m$). We will also refer to the Sersic profile as $\\Sigma_S$, our approximation to a projected 2D Einasto profile as $\\Sigma_E$ and a numerically projected Einasto profile as $\\Sigma_N$. Further, for this paper, $r_{conv}$ and $r_{200}$ have no physical meaning as such. We note from the Nav04 simulations, that on an average $r_{conv} \\sim 5$ per cent of $r_{-2}$, and $r_{200} \\sim 8-10$ per cent of $r_{-2}$ for galaxies and $\\sim 4$ per cent of $r_{-2}$ for clusters. In this paper we will, for the sake of discussion, refer to $r_{conv} \\sim$  $0.05$ $r_{-2}$ and $r_{200}$ as $\\sim$ $6$ $r_{-2}$.  In \\S 2 we derive a semi-analytical form \\eqref{sigmae} for $\\Sigma_E$, which for the Nav04 parametrization of $b=2n$, simplifies to \\eqref{sigmae2}. In \\S 3 we discuss details of estimating the fit parameters in our expression for $\\Sigma_E$, followed by a discussion of errors in our approximation. We then present a comparison between $\\Sigma_N$ and the best-fitting Sersic profile $\\Sigma_S$. In \\S 3.4, we note the remarkable accuracy with which \\eqref{sigmae2}, along with \\eqref{zeta2param} and \\eqref{muparam}, can be used to extract the 3D parameters of Einasto-like systems from its 2D form. ",
        "conclusions": "Non-parameteric estimates of density profiles in N-body simulations (Nav04,M06) favour Einasto-like profiles, since they provide better fits than the two-parameter NFW and Moore profiles. \\cite{M06} have also shown that a de-projected Sersic profile fits the 3D halo mass distribution almost as well as the Einasto profile, and a Sersic profile provides good fits to non-parametric estimates of surface mass densities (M05) of the Nav04 N-body haloes.  We have observed that fits with a Sersic function ($\\Sigma_S$) to a numerically projected Einasto profile ($\\Sigma_N$) are sensitive to whether one fits using linear density $\\Sigma_S$ (errors increasing for large R) or log density $ln(\\Sigma_S)$ (errors increasing for small R) yielding widely varying results. Consequently, the Sersic profile does not give an adequate description of the projected Einasto profile.  Sersic profile fits to the surface mass density of N-body haloes (M05), whose 3D spatial densities are well fit by Einasto profiles with $0.12$$\\leq$$\\alpha$$\\leq$$0.22$, have been obtained from the limited radial range of $r_{conv}$ to $r_{200}$. For the haloes in M05 and Nav04, this range is generally less than two decades in radius. Hence, if the 3D distribution is indeed Einasto-like, interpreting structural properties from fits with a Sersic profile, especially $m$ in \\eqref{sersic} as the shape parameter, $\\Sigma_S(0)$ as the central density and $R_E$ as the half-mass radius, can be misleading.  In this paper, we have provided an analytical approximation \\eqref{sigmae2} to the projected surface mass density of Einasto-like 3D density distributions. The fit errors are well contained to $<$$2\\%$ for the projected radial range $0$$\\leq$$R$$\\leq (10-30)$ $r_{-2}$ equivalent to $(3-5)$$r_{200}$ and shape parameter, $1$$\\leq$$n$$\\leq$$10$, or $0.1$$\\leq$$\\alpha$$\\leq$$1$. This model can therefore be used both as a fitting function for 2D observations and also to extract the 3D parameters of Einasto-like profiles. Since $\\Sigma_E$ fits a projected Einasto profile in a wide radial range, it can be used for fitting strong and weak lensing observations in systems whose total 3D density distribution is believed to be Einasto-like. One can also numerically integrate \\eqref{sigmae2} to get reliable estimates of the mass enclosed.  Finally, we note that the form similarity of \\eqref{einasto} and \\eqref{sersic}, i.e. fitting functions that describe the 3D mass density of dark matter haloes and the 2D light distributions of galaxies, respectively, could be largely coincidental and should be used with caution when drawing conclusions about the similarity of dynamical evolution that lead to the formation of the stellar components of ellipticals and dark matter haloes."
    },
    "1112/1112.1922_arXiv.txt": {
        "abstract": "*{}  \\abstract{Long tails and streams of stars are the most  noticeable upshots of galaxy collisions. Their origin as gravitational, tidal, disturbances has however been recognized only less than fifty years ago and more than ten years after their first observations.  This Review describes how the idea of galactic tides emerged, in particular thanks to the advances in  numerical simulations, from the first ones that included tens of particles to the most sophisticated  ones with tens of millions of them and state-of-the-art hydrodynamical prescriptions.  Theoretical aspects pertaining to the formation of tidal tails are then presented. The third part of the review turns to observations and underlines the need for collecting deep multi-wavelength data to tackle the variety of physical processes exhibited by collisional debris. Tidal tails are not just stellar structures, but turn out to contain all the components usually found in galactic disks, in particular atomic / molecular gas and dust. They host star-forming complexes and are able to form star-clusters or even second-generation dwarf galaxies. The final part of the review discusses what tidal tails can tell us (or not) about the structure and content of present-day galaxies, including their dark components,  and explains how tidal tails may be used to probe the past evolution of galaxies and their mass assembly history. On-going deep wide-field surveys disclose many new low-surface brightness structures in the nearby Universe, offering great opportunities for attempting galactic archeology with tidal tails.} ",
        "introduction": " ",
        "conclusions": "It is an undeniable fact that tidal forces and the formation of tidal tails are overall a second order  process in galaxy evolution. The fraction of stars expelled in the intergalactic medium is low, at most a few percent in major mergers. The fraction of gas is more important, but the bulk of the gaseous reservoir is funneled into the central regions. Collisional debris may host star-forming regions, but their contribution to the total star formation rate is minimum. Clearly, most of the activity occurs in the more central and nuclear regions where starbursts and/or AGN fueling is triggered. However, one of the aims of the present Review is to emphasize the idea that tidal debris can provide insightful information about the properties of galaxies, the same way as garbage in trash cans tells us much about the way of life of their owners.  The presence of tidal features is an unambiguous proof that a major/minor merger occurred in the recent past, and that at least one of the colliding galaxies had a stellar and/or gaseous disk. The converse  is not true though,  as not all collisions produce prominent tidal features. Determining when the merger took place is less strait-forward. However, numerical simulations done in  cosmological context will soon be able to constrain the survival time of collisional debris and thus give predictions on their age. Comparisons between observations and simulations should then allow us to reconstruct the mass assembly of galaxies.  Current generation of wide-field-of-view cameras and the on-going extremely deep surveys of the nearby Universe detect numerous new tidal features of very low surface brightness, offering interesting prospects to galactic archeology. At high redshift, the census of tidal perturbations is much more complex, not only because of dimming and band shifting issues, but also because distant galaxies are much more gas--rich and therefore are intrinsically irregular. This makes the separation between secular and external effects rather ambiguous.  Multi-wavelength surveys  have revealed the presence in collisional debris  of all the constituents of regular galaxies  though with different proportions: young and old stars, atomic gas, molecular gas, even possibly dark gas, heavy elements and dust. Star formation seems to proceed there in a similar way as in isolated spiral disks, despite the very different environment at large scale.  Tidal tails may in principle even be used to probe some fundamental aspects of physics, including, of course, the properties of tidal forces but also the laws of gravitation, as shown by recent experiments with modified gravity. The fact that tidal forces can be compressive and for instance  contribute to the stability of star clusters whereas they are usually associated with destruction processes has only recently been understood. The shape of the tidal tensor explains why massive tidal dwarf galaxies may only be formed within an extended dark matter halo. A theoretical study on the nature and the role of tidal forces in galaxies remains largely to be done and might provide further surprises."
    },
    "1112/1112.4500_arXiv.txt": {
        "abstract": "The electromagnetic modes possibly unstable in  strongly magnetized plasmas are identified. The regime where  this instability might stand out compared to the incoherent electron-cyclotron radiation is explored. These modes are relevant to the inertial confinement fusion and the gamma ray burst. ",
        "introduction": "The magnetic field is ubiquitous in plasmas. The plasma interaction with the magnetic field often determines its dynamical behavior~\\citep{Stone, Biskamp, Fishman, Nakar, Innes}. In particular, the magnetic field could convert the electron kinetic energy into the collective photons, as in the magnetron or the gyrotron~\\citep{gyrotron3, tgyro} and the free electron laser (FEL)~\\citep{Fisch, sonlandau, sonbackward}. In the FEL, a relativistic electron beam gets periodically  accelerated amplifying  coherent electromagnetic (E\\&M) waves. The resulting lasers have various applications. One natural question would be whether the electron thermal gyro-motion in the magnetic field, not the specifically designed population-inverted plasma state, can excite an analogous process.  In this letter, we examine the instability of collective E\\&M waves arising from the electron thermal gyro-motion. The analysis of the electron motion in the presence of a strong magnetic field leads to a theoretical framework similar to that of the Landau damping~\\citep{sonlandau, songamma}, which reveals that there could exist numerous unstable E\\&M modes in  relativistic plasmas. Our theory is similar to the well-known electron maser theory~\\citep{chu, Treumann}, but our theory deals with the long time limit while the previous researches mainly focus on the short time limit, which will be discussed in detail at the end of Sec.~(4). We identify the regime where the coherent radiation from this instability might be  more intense than the incoherent cyclotron radiation. Various implications of our study on the astrophysical and laboratory plasmas are discussed, including the short gamma ray burst~\\citep{Nakar}, the non-inductive current drive~\\citep{Fisch, sonprl} and the soft x-ray generation in the inertial confinement fusion~\\citep{Tabak}.  This paper is organized as follows. In Sec.~(2), the theory of the instability is developed based on the Landau damping theory for the non-relativistic plasmas. In Sec.~(3), the fully relativistic theory is developed. This section is the major result of this paper. In Sec.~(4), we identify the instability inherent in the relativistic plasmas.  In Sec.~(5), we discuss the implications of our results. ",
        "conclusions": "To summarize, the coherent E\\&M  instabilities of strongly magnetized relativistic electrons are analyzed in the context of the Landau damping theory. Our main results are Eqs.~(\\ref{eq:landau2}) and (\\ref{eq:landau3}). Our theory based on the long-time scale is relevant with the hard x-ray and the gamma ray generation in the strongly magnetized  dense plasmas. In an astrophysical plasma, a strong and large spatial scale magnetic field in the range of $10^{10} - 10^{15}$ gauss is often encountered. Its radiation may have to be re-examined in terms of the possible unstable modes identified in our analysis. More specifically, the theory would be relevant with the generation of the soft and the hard x-ray, and the gamma ray~\\citep{gamma}. The complications when applying the above theory to astrophysical plasmas are the relativistic effect and the electron quantum diffraction~\\citep{sonpla, sonprl, sonbackward, sonlandau}. One more relevant phenomenon is the non-inductive current drive~\\citep{Fisch, sonprl}. During the E\\&M wave amplification, the electrons giving the energy to the E\\&M wave lose their momentum at an increased rate, as the collision frequency increases compared to the case of no interaction with the E\\&M wave. This leads to the well-known non-inductive current drive~\\citep{Fisch, sonprl}.  The authors are thankful to Dr.~I.~Dodin and Prof.~N.~J.~Fisch for many useful discussions and advice."
    },
    "1112/1112.2921_arXiv.txt": {
        "abstract": "We measured the magnetic field of the bright, evolved hot subdwarf HD\\,76431 by means of high-resolution spectropolarimetry. In contrast to previous measurements we were not able to detect a significant magnetic field. We discuss the possibility that this field may be variable. Our search for a possible companion star to HD\\,76431 led to inconclusive results. ",
        "introduction": "HD\\,76431 (PG\\,0853+019, $m_{\\rm B}=8.9\\,{\\rm mag}$) is a bright B-type star located at high Galactic latitude. According to its atmospheric parameters \\citep[$T_{\\rm eff}=31\\,000\\,{\\rm K}$, $\\log{g}=4.51; $][]{ramspeck01} it is located below the main sequence in the $T_{\\rm eff}$-$\\log{g}$-diagram (see Fig.~\\ref{fig:tefflogg}). \\citet{ramspeck01} measured a subsolar He abundance ($\\log{y}=-1.51$) and concluded that HD\\,76431 is most likely a post-EHB object.  \\begin{figure} \\begin{center} \\includegraphics[width=10cm]{tefflogg_hd76431.eps} \\label{fig:tefflogg} \\caption{$T_{\\rm eff}-\\log{g}$ diagram. The grey circles mark sdBs from the SPY project \\citep{lisker05}. EHB evolutionary tracks are taken from \\citet{dorman93}. The location of the main sequence is indicated by the long-dashed horizontal lines. The atmospheric parameters of HD\\,76431 are taken from Ramspeck et al. (2001).} \\end{center} \\end{figure}  The role of magnetic fields in stellar evolution remains unclear. One way to address this issue is the analysis of the magnetic properties of evolved stars. While the strong MG-fields of magnetic white dwarfs are easily detectable due to the prominent Zeeman splitting of their spectral lines, weaker fields can only be detected with spectropolarimetry. Such fields of the order of a few kG have been detected in white dwarfs and central stars of planetary nebulae.  A first attempt to determine the magnetic fields of bright hot subdwarf stars using this method was made by \\citet{elkin98}, who also targeted HD\\,76431, but did not detect a magnetic field ($B_{\\rm e}=-50\\pm130\\,{\\rm G}$). \\citet{otoole05} on the other hand reported the detection of a magnetic field in HD\\,76431 with a mean longitudinal field strength of $-1096\\pm91\\,{\\rm G}$ using the medium resolution instrument FORS1 mounted at the ESO-VLT. Here we present the results of a spectropolarimetric monitoring campaign of this star using high-resolution data. ",
        "conclusions": "In contrast to the detection reported by \\citet{otoole05} no significant magnetic field could be found in HD\\,76431. This result is consistent with the non-detection by \\citep{elkin98}. Given that the measurement  obtained by \\citet{otoole05} is correct, the magnetic field of this star may be variable on a timescale of years. The reason for this variation may be slow rotation of the star. The sharp metal lines indicate a projected rotational velocity $<10\\,{\\rm km\\,s^{-1}}$.  However, Petit et al. (these proceedings) also report the non-detection of a magnetic field in HD\\,76431 based on high-resolution spectropolarimetry. Furthermore, the authors reanalysed the FORS1-data on which the results of \\citet{otoole05} are based and detected no significant magnetic field either. They conclude that the detection of \\citet{otoole05} was spurious and caused by instrumental effects. The question, whether the subdwarf has a companion, remains unclear."
    },
    "1112/1112.0115_arXiv.txt": {
        "abstract": "{ Physical and phenomenological properties (radius, luminosity, shape of the light curve,  etc.) of Cepheids strongly depend on the pulsation period, with the exception of the pulsation amplitude. A possible factor causing a wide range of pulsation amplitudes might be the different atmospheric metallicities of individual Cepheids. } { We studied the influence exerted by the atmospheric iron content, [Fe/H], on the pulsational amplitude of Galactic Cepheids.} { We searched for correlations between the [Fe/H] value and both the observed amplitudes and amplitude related parameters. } { The amplitude of the Cepheid pulsation slightly decreases with increasing iron abundance. This effect is more pronounced for the radial velocity variations and for the shorter pulsation periods. The wavelength dependence of photometric amplitudes is also found to be sensitive to the metallicity. Some of these effects are not consequences of differential line blanketing. Based on the calibrations of the metallicity sensitivity relationships, we derived photometric iron abundance for 21 Galactic Cepheids. The dichotomic behaviour dividing Galactic Cepheids that pulsate in the fundamental mode into short- and long-period groups at the period of 10\\fd47 can be noticed in some diagrams that show metallicity-related dependences. } { We confirm that variety in atmospheric metallicity in Cepheids contributes to the finite range of pulsation amplitudes at a given period. Effects of metallicity on the amplitudes revealed from observational data and the occurrence of the dichotomy also derived from phenomenological data have to be confirmed by appropriate theoretical models of stellar structure and pulsation.} {} ",
        "introduction": "\\label{sect_1}  Classical Cepheids are radially pulsating supergiant stars. This type of pulsation, maintained by the $\\kappa$-mechanism in the outer layers of the star, is stable in a narrow range of the effective temperature, in the nearly vertical {\\em instability strip} in the Hertzsprung-Russell diagram. The fact that this pulsation is a free radial oscillation of the star results in a relationship between the period, $P$, and luminosity, $L$, of Cepheids.  Owing to the existence of the period-luminosity ($P$-$L$) relationship, Cepheids are primary distance indicators in astronomy. The precise calibration of this relationship has been in the forefront of Cepheid-related studies for decades. A recently emerged aspect of these studies is the role of metallicity in shaping the $P$-$L$ relationship. In spite of thorough theoretical and observational investigations, the results on the metallicity dependence are controversial (see summaries by \\cite{Retal08}; \\cite{M09}; \\cite{FM10}).  Along with other Cepheid-related relationships, the $P$-$L$ relationship is roughly linear if plotted against $\\log P$. The only exception is the period-amplitude ($P$-$A$) relationship. The wide range of possible photometric and radial velocity amplitudes at a given pulsation period may be caused by various effects (pulsation mode and energy, companion star(s), metallicity, etc.).  Motivated by the fact that the dependence of the pulsation amplitude on $\\log P$ is neither linear nor unique, we studied the $P$-$A$ relationship of Galactic Cepheids in Paper~I (\\cite{KSz09}). We revised the $P$-$A$ graphs for the $U$, $B$, $V$, $R_{\\rm C}$, and $I_{\\rm C}$ photometric bands. In addition, a $P$-$A$ graph was constructed using the observed peak-to-peak amplitudes of the pulsational radial velocity variations updating the much earlier diagrams (including that by \\cite{CS84} to which we inadvertently missed to refer in Paper~I). One purpose of constructing new $P$-$A$ diagrams has been to study the possible effect of metallicity on the pulsational amplitude, which may contribute to the observed finite range of actual amplitudes at a given pulsation period.   \\begin{table*}[t] \\caption{Average [Fe/H] values for various groups of Cepheids} \\begin{tabular}{lcccccccccccccc} \\hline \\noalign{\\smallskip} Sample && \\multicolumn{3}{c}{Solitary Cepheids} && &\\multicolumn{3}{c}{Binary Cepheids}& &&\\multicolumn{3}{c}{Total sample}\\\\ \\noalign{\\smallskip} \\cline{3-5} \\cline{8-10} \\cline{13-15} && [Fe/H] & $\\sigma$ & $N$ &&& [Fe/H] & $\\sigma$ & $N$ &&&[Fe/H] & $\\sigma$ & $N$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Fundamental mode (F) Cepheids, $\\log P > 1.02$ && 0.167 & 0.164 & 49 &&& 0.184 & 0.154 & 34 &&& 0.174 & 0.160 & 83\\\\ Fundamental mode (F) Cepheids, $\\log P < 1.02$ && 0.091 & 0.158 & 114 &&& 0.103 & 0.100 & 78 &&& 0.096 & 0.137 & 192\\\\ First overtone (1OT) Cepheids  && 0.108 & 0.142 & 28 &&& 0.068 & 0.118 & 24 &&& 0.089 & 0.133 & 52\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\label{tab_aveFe} \\end{table*}  An influence of the chemical composition on the amplitude of pulsation is expected based on different patterns of the $P$-$A$ plots for Cepheids in our Galaxy and the Magellanic Clouds. van Genderen (\\cite{vG78}) was the first who attributed the differences in the largest pulsation amplitudes of Cepheids to the different average metallicities of the respective host galaxies. If atmospheric metallicity has an influence on largest possible amplitude at a given pulsation period, then it should exert influence on the actual amplitude of individual Cepheids.  Here we study the effect of the atmospheric iron content, [Fe/H], on peak-to-peak amplitudes and amplitude-related parameters introduced in Paper~I for classical Cepheids of our Galaxy. Input values have been taken from Paper~I (where the sources of the amplitude and spectroscopic [Fe/H] data as well as the method of their homogenization is also described); in addition, [Fe/H] values published recently by Luck \\& Lambert (\\cite{LL11}) and Luck et~al. (\\cite{Letal11}) have also been taken into account (after applying the same correction as for the [Fe/H] data tabulated in Paper~I). The present sample consists of 329 Galactic Cepheids with known spectroscopic [Fe/H] value.  This paper is organized as follows. Section~\\ref{sect_per} deals with the period dependence of the iron content of Galactic Cepheids. Correlations between the [Fe/H] value and various amplitudes and amplitude-related parameters are presented in Sect.~\\ref{sect_amp}. Consequences of our results are discussed in Sect.~\\ref{sect_disc}, while the conclusions are drawn in Sect.~\\ref{sect_concl}. ",
        "conclusions": "\\label{sect_concl}  We confirmed that the observed variety in the atmospheric iron content of Cepheids contributes to the wide range of the actual pulsation amplitudes: more metal-deficient Cepheids tend to pulsate with larger amplitudes. We were able to reveal the metallicity dependence of several parameters ($q$, $k$, $m$, $\\Delta k$, and $\\Delta m$) of Galactic Cepheids. In physical terms, these relationships mean that the amplitude of the radial velocity variations, the ratio of photometric to radial velocity amplitudes, and the wavelength dependence of the photometric amplitudes depend on the iron abundance of the pulsating stellar atmosphere. In some cases, the metallicity sensitivity of the amplitude parameters is not the consequence of the differential line blanketing with the result that stars of higher metallicity are fainter towards shorter wavelengths.  These correlations, although not strong, can still be used for deriving the [Fe/H] value of individual Cepheids from photometric observational data. Improvement of these correlations is mainly hindered by the current uncertainties in the spectroscopic abundance determination: individual [Fe/H] values have an error of $0.05 - 0.10$. From multicolour photometric data, we derived a photometric [Fe/H] value for 21 Galactic long-period Cepheids, one of which (GY~Sge) is lacking prior spectroscopic abundance determination.  Availability of direct spectroscopic abundance values and multicolour light curves of Cepheids in the Magellanic Clouds would extend the newly found correlations to more negative [Fe/H] values because the interval of atmospheric iron content of the Galactic Cepheids (typically $-0.25 < {\\rm [Fe/H]} < +0.25$) is not particularly wide. Quite a few Cepheids more deficient in iron are known in the outer spiral arm of the Milky Way system but these Cepheids are so distant that their available spectroscopic [Fe/H] values are too uncertain for calibration purposes. Magellanic Cepheids are much more metal-deficient: the mean [Fe/H] is $-0.34 \\pm 0.03$ for the LMC and $-0.64 \\pm 0.04$ for the SMC (Keller \\& Wood \\cite{KW06}). The available observational data on the abundance of individual Magellanic Cepheids are scanty: photometric metallicity values determined by Harris (\\cite{H81}) using the Washington system resulted in metallicities of limited accuracy, and spectroscopic metallicities for a sample consisting of 12-12 Cepheids in each cloud (\\cite{Retal05}).  The finite interval of actual [Fe/H] values of Cepheids situated in different parts of the same galaxy is a drawback when attempting to decrease the scatter of the $P$-$L$ relationship for Cepheids in individual galaxies and in settling the question of the metallicity sensitivity of the $P$-$L$ relationship. The wide variety of iron abundance values of Cepheids in the same host galaxy is a reason to be cautious when assigning an `average' metallicity to all Cepheids in a given galaxy. As a rule, the determination of the heavy element abundance of external galaxies (except the nearest ones) is based on the chemical properties of their bright HII regions (\\cite{Zetal94}). Metallicity of these regions, however, corresponds to youngest population of stars involving long-period Cepheids. Short-period Cepheids are somewhat older (and evolve slower) and may have formed from a less processed interstellar matter.  Both the effect of metallicity (of non-blanketing origin) on the amplitudes revealed from observational data and the occurrence of the dichotomy at the limit of 10\\fd47 (including its relation to the $2P_{\\rm 2OT} = P_{\\rm F}$ resonance phenomenon) have to be investigated through appropriate theoretical models of stellar pulsation."
    },
    "1112/1112.5730_arXiv.txt": {
        "abstract": "The differential intensities of Cosmic Rays at Earth were calculated using a 2D stochastic Montecarlo diffusion code and compared with observation data. We evaluated the effect of stretched and compressed heliospheres on the Cosmic Ray intensities at the Earth. This was studied introducing a dependence of the diffusion parameter on the heliospherical size. Then, we found that the optimum value of the heliospherical radius better accounting for experimental data. We also found that the obtained values depends on solar activity. Our results are compatible with Voyager observations and with models of heliospherical size modulation. ",
        "introduction": "HelMod Code\\cite{HelMod2011} solves the bi-dimensional Parker's particle transport equation\\cite{parker1965}\\!. A Monte Carlo technique is applied on a set of Stochastic Differential Equations (SDEs) fully equivalent to the Parker's equation\\cite{AstraArticle2011}\\!. The model takes into account particle drift effects and latitudinal dependence of the solar wind speed and of the Interplanetary Magnetic Field (IMF). It is described in details in Ref.~\\refcite{HelMod2011}. In the model, the IMF from Parker\\cite{parker1958} is modified introducing a small latitudinal components as described in Ref.~\\refcite{langner2004}. For periods of low solar activity, we take a solar wind speed gradually increasing from the Earth position up to a maximum value near the heliospherical poles ($\\simeq 760$\\,km/s)~\\cite{McComas2000}\\!. For periods approaching the solar maximum we assume a solar wind speed independent on the latitude.  The symmetric part of the diffusion tensor, in a reference frame with one axis aligned with the Parker's magnetic-field, is purely diagonal containing transverse ($K_{\\perp \\theta}$ and $K_{\\perp r}$) and parallel ($K_{||}$) components~\\cite{potgieter2000}\\!. The diffusion coefficients are given by~\\cite{potgieter1994}  \\begin{eqnarray}\\label{duffusion_comp} \\begin{array}{l} K_{||} = \\beta \\, K_0(t) \\,K_{P}(P)\\,\\left[\\frac{B_{\\oplus}}{3B} \\right] \\ , \\\\ K_{\\perp r} = \\rho_k \\ K_{||} \\ , \\\\ K_{\\perp \\theta} = \\iota(\\theta) \\ \\rho_k \\ K_{||} \\ , \\end{array} \\end{eqnarray}  \\noindent where $\\beta = v/c$, $v$ the particle velocity and $c$ the speed of light; the diffusion parameter $K_0$ accounts for the dependence on the solar activity; $ B_{\\oplus}$ is the measured value of IMF at the Earth position - typically $ \\approx 5$\\,nT, but changing with time - obtained from Ref.~\\refcite{SW_web}; $B$ is the magnitude of the large scale IMF as a function of heliocentric coordinates; finally, the term $K_{P}$ takes into account the dependence on the rigidity $P$ of the GCR particle usually expressed in GV.~In the present model $K_P \\approx P$ (e.g., see Ref.\\cite{perko1987}). Furthermore \\ $\\rho_k = 0.05$ \\ and, as described in Ref.~\\refcite{HelMod2011},  \\begin{eqnarray}\\label{enhaced_K} \\iota(\\theta)  = \\left\\{ \\begin{array}{rll} 10 \\ ,  & \\textrm{ in the polar regions},\\\\ 1 \\ , & \\textrm{ in the equatorial region}. \\end{array}\\right. \\end{eqnarray}  After the transformations from 3D field-aligned into 2D heliospherical coordinates~\\cite{burg2008}\\!, the symmetric components of the diffusion tensor contains both diagonal ($K_{r r}$ and $K_{\\theta \\theta}$) and off-diagonal terms ($K_{r \\theta}$ and $K_{\\theta r}$), resulting by a proper combination of $K_{\\perp \\theta}$, $K_{\\perp r}$ and $K_{||}$~\\cite{HelMod2011}\\!. ",
        "conclusions": "We presented the HelMod 2D Monte Carlo code for the study of Cosmic Rays propagation in the inner heliosphere. Both heliospherical shape and size are supposed to be relevant for the modulation process. We introduced a dependence of the diffusion parameter on the heliospherical size, which accounts for the variation with time and solar activity. We compare modulated spectra with experimental data covering the solar cycle 23. Then we found, for our 2D model, the best value of the heliospherical radius, which changes with time. Most of the solar modulation occurs in the inner heliosphere and differences in the heliospherical radius are effective only at energy below a few hundred MeV. Our results are not in contradiction with Voyager observations and models of TS distance as a function of solar activity. We found that LIS form Ref.~\\refcite{Burger2000} fits better observation data at low energy."
    },
    "1112/1112.0609_arXiv.txt": {
        "abstract": "A quantitative study of the observable radio signatures of the sausage, kink, and torsional MHD oscillation  modes  in flaring coronal loops is performed. Considering first non-zero order effect of these various MHD oscillation modes on the radio source parameters such as magnetic field, line of sight, plasma density and temperature, electron distribution function, and the source dimensions, we compute time dependent radio emission (spectra and light curves). The radio light curves (of both flux density and degree of polarization) at all considered radio frequencies are than quantified in both time domain (via computation of the full modulation amplitude as a function of frequency) and in Fourier domain (oscillation spectra, phases, and partial modulation amplitude) to form the signatures specific to a particular oscillation mode and/or   source parameter regime. We found that the parameter regime and the involved MHD mode can indeed be distinguished using the quantitative measures derived in the modeling. We apply the developed approach to analyze radio burst recorded by Owens Valley Solar Array and report possible detection of the sausage mode oscillation in one (partly occulted) flare and kink or torsional oscillations in another flare. ",
        "introduction": "Oscillations and quasi-periodic pulsations (QPPs) in radio emission from solar flares have been observed for half a century \\citep{Young61}. Pulsation phenomena can be formally categorized by wavelength of the emission and by frequency/period of the QPPs. For example, at meter  and decimeter wavelengths, rapid QPPs occur with periods of tens to hundreds of milliseconds. Typically, these QPPs have rapidly varying amplitudes, variable periods, and narrow bandwidths $\\delta f/f\\lesssim 1$, \\citep[e.g.][]{Young61, Zlobec_etal_1987, Stepanov_Yurovsky_1990, Kurths_etal_1991, Yurovsky_1991, Fleishman94, Aschwanden_etal_1995, Fl_etal_2002, PRL, Magdalenic_etal_2002, Benz_etal_2005, Meszarosova_etal_2005, Chen_Yan_2007, Benz_etal_2011, Chen_etal_2011}. Longer-period QPPs, with periods $\\tau\\gtrsim10$~s  are observed in microwaves; such QPPs are often accompanied by QPPs in hard X-rays \\citep[HXR; e.g.,][] {Parks_Winckler_1969, Janssens_etal_1973, Zaitsev_Stepanov_1982b, Kane_etal_1983, Nakajima_etal_1983, Asai_etal_2001, Grechnev_etal_2003, Nakariakov_etal_2003, Stepanov_etal_2004, Melnikov_etal_2005, Fl_etal_2008}.  Numerous models have been proposed to account for the various types of oscillations, QPPs, and related phenomena observed at radio wavelengths \\citep[e.g.,][and references therein]{Aschwanden87, Nakariakov_etal_2003, Fl_etal_2008}. Generally speaking, for radio emissions produced by coherent radiation processes (such as fast millisecond pulsations), QPPs are believed to result from nonlinear, self-organizing wave-wave or wave-particle interactions \\citep{Zaitsev_Stepanov_1975b, Meerson_etal_1978, Bardakov_Stepanov_1979, Aschwanden88,Fleishman94,Korsakov98}. In contrast, for radio emission produced by incoherent gyrosynchrotron (GS) radiation from energetic electrons, QPPs are believed to result from modulation of the source parameters (such as the energetic electron distribution, the magnetic field strength, the line of sight, etc.) via magnetohydrodynamic (MHD) oscillations (kink, sausage, or torsional modes) \\citep{Aschwanden04} and/or modulation of electron distribution function via time-dependent acceleration and injection. Most of the reported so far QPPs of GS microwave emission \\citep[e.g.,][]{Kane_etal_1983, Nakajima_etal_1983, Fl_etal_2008} are consistent with modulation of fast electron acceleration/injection, while there is no compelling case in which modulation of GS emission by an MHD oscillation mode would be firmly confirmed. As has been shown by \\citet{Fl_etal_2008}, a well spectrally resolved radio spectrum is crucial to distinguish between MHD oscillations and modulation of fast electron injection/acceleration. Another component needed to firmly detect a particular MHD oscillation mode from the observed radio data is detailed quantitative modeling of observable measures, such as oscillation phase or modulation power, of the GS flaring emission modulated by one or another MHD mode, which is yet unavailable.   This paper attempts to partly remedy the situation by analyzing  dependence of GS radiation on modulation of the source parameters by various MHD modes, namely, large-scale kink, sausage, and torsional oscillation modes, while we do not consider a special case of the longitudinal slow mode (essentially, the sound wave) where only the plasma density oscillates since this mode can only affect low-frequency radiation from a dense plasma, which has already been studied by \\citet{Nakariakov_Melnikov_2006}. For each mode, time-domain and Fourier analyses are performed on the radio emission and polarization to characterize the properties of the pulsations, and the results are evaluated to identity a signature by which each oscillation mode can be distinguished in observational data. In contrast to approximate sausage-mode modulation modeling performed by  \\citet{Fl_etal_2008} based on relatively crude Dulk-Marsh approximation, in this study the microwave (GS plus free-free) emission is modeled using precise fast GS codes recently developed by \\cite{Fl_Kuzn_2010}.   The results of the modeling are then used to search for the corresponding QPP signatures in the total power microwave database available from the Owens Valley Solar Array (OVSA). Specifically, we looked through well-calibrated data recorded during 2001-2002 when microwave emission from 412 solar flares was recorded \\citep{Nita_etal_2004}. We found the the vast majority of the radio bursts does not display unambiguous QPPs in the dynamic spectra; only about 10\\% of all events show time variations at the light curves that can be classified as QPPs. From them we selected a few events in which the presence of QPPs is apparent from their dynamic spectra; for those selected events we computed all possible quantitative measures to be compared with the corresponding measures derived from the modeling in the attempt to detect an indication of one or another MHD loop oscillation mode. In most of the case we were unable to firmly identify the genuine cause of the QPPs; however, in two best cases, presented in the paper, the observed QPPs indeed could be produced by an MHD mode: the sausage mode in an (occulted) flare of 28 August 2002 21:41 UT and kink or torsional mode in a flare of 20 July 2002 21:27 UT. % We note that some MHD oscillation modes may or may not produce  corresponding oscillations in HXR and/or soft X-rays (SXR), which we also tested for the mentioned events using HXR data available from the Reuven Ramaty High Energy Solar Spectroscopic Imager (RHESSI) and SXR data from the Geostationary Operational Environmental Satellites (GOES). % ",
        "conclusions": "We have described a detailed modeling of the MHD loop oscillation effect on the GS microwave emission from solar flares. Unlike most of other studies concentrating on observed QPP periods and their possible correspondence to one or another MHD mode, we address the question how the oscillation phase and corresponding modulation amplitude behave as a function of the radio frequency, given that these QPP measures were shown to be highly efficient to identify the true cause of QPPs \\citep{Fl_etal_2008}. We found that the QPP relative phase behavior is distinctly different for the cases of relatively strong magnetic field when the phase repeatedly changes at each of the local peaks at low gyroharmonics  and of  either dense or tenuous flaring loops  with a lower field, displaying no or $\\pi$ phase shift around the (single) spectral peak, respectively. Furthermore, the sausage mode oscillations are distinguishable from the kink and torsional mode oscillations using the frequency behavior of the modulation amplitude. The kink and torsional modes are distinguishable from each other in the case of tenuous source, when oscillations of the source visible area dominate the kink mode effect on the radio QPPs in the optically thick region. However, the kink and torsional mode oscillations in a dense loop are indistinguishable from each other as far as linear order effect (which is identical for these modes) dominates the radio emission response on these MHD oscillations.  We have used the modeling results to analyze the OVSA microwave burst database and search for the expected MHD oscillation signatures in some bursts displaying QPPs. Although we did not find many promising events-candidates, we did identified two reasonably good examples in which detection of one or another MHD oscillation mode is likely. Specifically, in one of them the radio QPPs might be caused by the sausage mode, while in the other one by either kink or torsional mode. Comparison of the radio QPPs with SXR and HXR QPPs favors the MHD loop oscillations for these two events compared with the fast electron injection/acceleration oscillations detected in another event using a similar analysis \\citep{Fl_etal_2008}. However, this conclusion must be taken with some caution (given incompleteness and insufficient accuracy of the data) and must be further evaluated and tested on other events with radio QPPs; we plan to perform a more comprehensive statistical study based on full available OVSA database elsewhere including analysis of the polarization data when available, because as evident from the above modeling, such a study would benefit greatly from  accurate polarization data.  We anticipate that similar data analysis will be even more important for soon-to-be-available imaging spectroscopy data after on-going OVSA expansion (EOVSA) has been completed. In particular, all EOVSA antennas will be equipped with dual-polarization feeds for simultaneous right-hand and left-hand circular polarization measurements. Furthermore, the expansion will greatly improve the image quality and  the frequency resolution, which are all needed to perform a more meaningful analysis and yield more reliable  conclusions on the cause of radio QPPs in each particular case."
    },
    "1112/1112.0323_arXiv.txt": {
        "abstract": "I summarise recent results on multi-wavelength properties of distant lensed galaxies, with a particular focus on {\\it Herschel}. Submm surveys have already resulted in a breakthrough discovery of an extremely efficient selection technique for strong gravitational lenses. Benefitting from the gravitational magnification boost, blind mm-wave redshifts have been demonstrated on IRAM, SMA and GBT, and follow-up emission line detections have been made of water, [O{\\sc iii}], [C{\\sc ii}] and other species, revealing the PDR/XDR/CRDR conditions. I also discuss HST imaging of submm lenses, lensed galaxy reconstruction, the prospects for ALMA and e-Merlin and the effects of differential magnification. Many emission line diagnostics are relatively unaffected by differential magnification, but SED-based estimates of bolometric fractions in lensed infrared galaxies are so unreliable as to be useless, unless a lens mass model is available to correct for differential amplification. ",
        "introduction": "Gravitational lens magnification is a fast route to faint object and faint population studies. This technique has been exploited to reveal or constrain the populations responsible for integrated far-IR and submm extragalactic backgrounds (e.g. \\cite[Smail et al. 1997]{Smail97}, \\cite[Knudsen et al. 2008]{Knudsen08}).  Lensing is also one of the few available means of inferring the distribution of dark matter in galaxies (e.g. \\cite[Gavazzi et al. 2007]{Gavazzi07}). Rare lens configurations even have the promise of yielding $\\Omega_{\\rm M}$ and $w$ to $10\\%$; this requires $50$ double Einstein rings, requiring in turn a parent sample of thousands of lenses (\\cite[Gavazzi et al. 2008]{Gavazzi08}). Indeed many current and future applications of lensing are limited by sample size (\\cite[Treu 2010]{Treu10}). Furthermore, most lens surveys have selected strongly for foreground lenses at relatively low redshifts, e.g. $z<0.2$, restricting studies of the evolving dark matter distribution in galaxies. These are strong drivers for both larger and higher redshift samples of lenses.  The past year has seen an explosion of interest in submm-selected strong gravitational lenses, mainly driven by {\\it Herschel} open and guaranteed time surveys and their follow-ups. New techniques pioneered with {\\it Herschel} are capable of detecting lenses in much higher numbers and to much higher redshifts than previous optical/near-IR surveys such as SLACS or SLQS (\\cite[Bolton et al. 2006]{Bolton06}, \\cite[Oguri et al. 2006]{Oguri06}). This paper will briefly review some recent key infrared lensing results with a particular focus on the multiwavelength spectral energy distributions and on new data from {\\it Herschel}. ",
        "conclusions": "Galaxy cluster lenses are a rapid route to probing the populations dominating the cosmic infrared backgrounds, both through direct detections and stacking analyses. The bulk of the observed or predicted cosmic infrared backgrounds at $15\\,\\mu$m, $100\\,\\mu$m, $160\\,\\mu$m and $850\\,\\mu$m have now been resolved into individual sources through ultra-deep imaging in the fields of foreground gravitationally-lensing galaxy clusters. Meanwhile, there is a clear need for larger numbers of strong galaxy-galaxy lens systems, particularly at higher lens redshifts. Following the spectacular confirmation of the $\\sim100\\%$ gravitational lens selection efficiency from submm-selected lenses, there is now a tremendous scope for such lens discovery in the submm with {\\it Herschel}. A further breakthrough has been ``blind'' submm/mm-wave redshift determination, i.e. without recourse to the torturous multi-stage multi-wavelength cross-identification and long-term 8m/10m-class optical/near-IR spectroscopy. Emission line diagnostics already yielding constraints on the physical conditions in these lensed galaxies. However, differential magnification effects cannot be neglected. In particular, the observed CO SLED in any lensed system can easily be un-representative of the SLED in the underlying source. Furthermore, the uncorrected bolometric fractions of strongly-lensed galaixes (magnifications $\\mu>10$) inferred from broad-band SEDs are so unreliable as to be useless.  SS thanks friends and colleagues the H-ATLAS (\\cite{Eales10}) and HerMES consortia, whose results are quoted in this paper, STFC (grant ST/G002533/1) for financial support, and the conference organisers for the kind invitation."
    },
    "1112/1112.0053_arXiv.txt": {
        "abstract": "Be stars possess gaseous circumstellar decretion disks, which are well described using standard $\\alpha$-disk theory. The Be star 28\\,CMa recently underwent a long outburst followed by a long period of quiescence, during which the disk dissipated. Here we present the first time-dependent models of the dissipation of a viscous decretion disk. By modeling the rate of decline of the $V$-band excess, we determine that the viscosity parameter $\\alpha=1.0\\pm0.2$, corresponding to a mass injection rate $\\dot{M}=(3.5\\pm 1.3) \\times 10^{-8}\\ M_\\sun\\,\\mathrm{yr}^{-1}$. Such a large value of $\\alpha$ suggests that the origin of the turbulent viscosity is an instability in the disk whose growth is limited by shock dissipation.  The mass injection rate is more than an order of magnitude larger than the wind mass loss rate inferred from UV observations, implying that the mass injection mechanism most likely is not the stellar wind, but some other mechanism. ",
        "introduction": "\\label{intro}  It is now generally accepted that Classical Be stars possess gaseous circumstellar disks that are responsible for their excess emission \\citep[for reviews, see][]{por03,baa01}.  Furthermore spectroastrometry and optical interferometic measurements indicate that this material is in Keplerian rotation \\citep{oud11,kra11}.  Consequently, as detailed in the review of \\citet{car11}, the steady-state viscous decretion disk model \\citep{lee91} is capable of explaining many of the (time-averaged) observational properties of Be star disks.  Most likely these disks are fed by mass-loss from their rapidly rotating central stars.  Both observational \\citep{riv98,nei02} and theoretical \\citep{and86,cra09} studies suggest that non-radial pulsations may play an important role.  However, to create a viscous decretion disk, this material must be placed into orbit.  Since the stars most likely are rotating at less than their critical speed \\citep[e.g.,][]{cra05}, the mass loss mechanism must add some angular momentum before injecting the material into to the disk.  From a theoretical point of view, this requirement is problematic, so, as yet, no entirely satisfactory mechanism has been proposed \\citep[see][for a detailed assessment]{owo06}.  Regardless of how material is injected into the disk, if this material is in orbit, turbulent viscosity will redistribute it on a viscous diffusion timescale, producing an outflowing decretion disk, as long as the disk is fed continually.  However, Be stars exhibit a wide range of variability, ranging from emission line profile variations \\citep[e.g, $V/R$ variability, indicative of density waves within the disk; ][]{ste09,car09}, to photometric and polarimetric variability, indicating that frequently the disks are not fed at a constant rate.  In some cases, the disk can disappear for decades before being rebuilt \\citep[see review by][]{und82,wis10}.  Oftentimes, this rebuilding phase occurs in a series of outbursts, during which the disk apparently is fed for a short period of time followed by a period of quiescence, during which the nascent disk can partially reaccrete onto the star.  Similar outbursts can occur while the disk is disappearing.  Once such example is the Be star 28\\,CMa.  Over the past 40 years, 28\\,CMa has exhibited quasi-regular outbursts, every 8 years or so, when the star brightens by about half a magnitude in the $V$-band.  This usually occurs via multiple individual injections (e.g., the rapid rises seen in Fig. 1), during which the star brightens at a rate of $0.02\\pm0.01\\,\\mathrm{mag\\,d}^{-1}$, due to free-bound and free-free radiation from the disk.  The outbursts last for about two years \\citep[see][for a detailed discussion of the long-term photometric  variations of 28\\,CMa]{ste03}.  A particularly well-monitored outburst occurred from 2001 to 2004, and this outburst was followed by a long quiescent phase from 2004 to 2009.  Furthermore it appears that during this quiescent phase, very little additional material was injected into the disk (only a few minor events).  Thus nature has provided us the perfect experiment to study how Be star disks dissipate.  If we accept as a premise that the viscous decretion disk model is correct (and that no other mechanisms participate in the disk mass budget), then the timescale of the post-outburst photometric decline is controlled solely by the disk viscosity.  In this paper, we model the $V$-band light curve with the goals of: 1) directly measuring the disk viscosity parameter, $\\alpha$, and 2) further testing the viscous decretion disk model by comparing its time-dependent predictions of the disk dissipation against the detailed shape of the photometric light curve.  \\begin{figure*} \\centerline{\\includegraphics[width=0.85\\linewidth]{f1.eps}} \\caption{$V$-band Light Curve.  Visual observations of 28\\,CMa (grey triangles) are shown in comparison to our model fits for different values of $\\alpha$, as indicated.  Phase~I (MJD = 52900--53070) is the initial decline, which was used to determine the value of $\\alpha$. Phase~II (MJD = 53070--54670) is the slow disk-draining phase. The inset shows the reduced chi-squared of the Phase~I fit for different values of $\\alpha$. The horizontal dotted line indicates the 90\\% confidence level. } \\label{fig:LightCurve} \\end{figure*} ",
        "conclusions": "Light curve fitting to the dissipation of a disk can be a powerful tool for measuring $\\alpha$. Further it provides additional quantitative tests of the viscous decretion disk model.  Here we have shown that the viscous decretion disk model can reproduce the detailed time-dependence of the dissipation of a Be star disk.  Using the $V$-band excess to measure how the disk density decreases with time, we have been able for the first time to directly measure the value of the disk viscosity parameter.  We find $\\alpha = 1.0\\pm0.2$, which provides an important clue about the origin of the turbulent viscosity, suggesting that it likely is produced by an instability in the disk whose growth is limited by shock dissipation.  Finally, knowing the value of $\\alpha$, we % calculate that the mass decretion rate of 28\\,CMa was ${\\dot M} = (3.5\\pm1.3) \\times 10^{-8}\\ M_\\sun\\ \\mathrm{yr}^{-1}$.  Such a large required mass injection rate places strong constraints on any proposed stellar mass loss mechanism."
    },
    "1112/1112.0862_arXiv.txt": {
        "abstract": "{} {\\zp, one of the closest and brightest massive stars, was the first early-type object observed by the current generation of X-ray observatories. These data provided some surprising results, confirming partly the theoretical predictions while simultaneously unveiling some problematic mismatches with expectations. In this series of papers, we perform a thorough study of \\zp\\ in X-rays, using a decade of \\xmm\\ observations.} {\\zp\\ was observed 18 times by \\xmm, totaling 1Ms in exposure. This provides the highest-quality high-resolution X-ray spectrum of a massive star to date, as well as a perfect dataset for studying X-ray variability in an ``archetype'' object. } {This first paper reports on the data reduction of this unique dataset and provides a few preliminary results. On the one hand, the analysis of EPIC low-resolution spectra shows the star to have a remarkably stable X-ray emission from one observation to the next. On the other hand, the fitting by a wind model of individual line profiles recorded by RGS confirms the wavelength dependence of the line morphology.} {} ",
        "introduction": "With its very early spectral type (O4Infp, \\citealt{wal72}) and a distance of only 335\\,pc \\citep{van07,mai08}, the star Naos, better known as \\zp\\ (or HD\\,66811), is one of the closest and brightest massive stars. It is therefore one of the most studied objects amongst the O-star population. However, despite the intense work, many open questions remain on its nature.  Indeed, \\zp\\ displays several intriguing properties. First, its visible spectrum shows clear signs of helium overabundance and chemical enrichment by CNO-processed material (e.g. \\citealt{pau01}) as well as fast rotation (more than 200\\kms\\ for $v \\sin(i)$, \\citealt{pen96,how97}). Second, it is a known runaway (e.g. from Hipparcos data, \\citealt{mof98}). These properties have led to speculations on its evolutionary status. On the one hand, the chemical enrichment and fast rotation could result from mass and angular momentum exchange through Roche lobe overflow in a binary. \\zp\\ could therefore have been the secondary component of such a system, the supernova explosion of its companion having ejected it from its birth place a few millions years ago \\citep{van96}. On the other hand, \\zp\\ {displays a similar Hipparcos parallax} as stars of the Vela R2 association\\footnote{ This conclusion was based on the original release of the Hipparcos catalog, hence the use of the `old' distance of 430\\,pc in the Schaerer et al. paper, but the parallax similarity remains when using the new reduction of Van Leeuwen (and thus the `new' distance of 335\\,pc).} \\citep{sch97}, and dynamical interactions within this association could have led to the ejection of the (single) O-star \\citep{van96}. In this scenario, the chemical enrichment of \\zp\\ would be explained by the intense rotational mixing occurring in the fast-rotating main-sequence progenitor \\citep{mey00}. In addition, \\zp\\ displays double-peaked emission lines, suggested to arise in a rotating wind \\citep{con74,pet96}, and a compression of the wind in the equatorial plane was detected by \\citet{har96}.  Due to its brightness, \\zp\\ was one of the first massive stars observed with high-resolution in X-rays \\citep{kah01,cas01}. At first, its X-ray lines appeared to match expectations as they did show the broad, blueward-skewed profiles expected for the wind embedded shock model \\citep{owo01}. However, the devil was in the details. When quantitatively fitting the line profiles, \\citet{kra03} found a much lower wind attenuation than expected on the basis of the mass-loss rate determined from optical and UV observations (see also \\citealt{osk06}). They also found that the typical optical depths $\\tau_*$, used in the wind-shock models, seemed independent of wavelength, which can only be explained by invoking porosity \\citep{fel03, osk06}. To improve the fitting of the X-ray line profiles, \\citet{leu07} included the effect of resonance scattering: better fits were indeed obtained, without the need of a large reduction in the mass-loss rate, but they also showed that some unexplained discrepancies remain. Re-analyzing the {\\it Chandra} data of \\zp, \\citet{coh10} argue in favor of a reduced mass-loss rate, without the need of any porosity as their new derivation of the optical depths implies an increase with wavelength, as expected from the bound-free absorption opacity of the (cool) wind. Except for \\citet{leu07}, all above studies relied on a single 68\\,ks {\\it Chandra} observation or a 57\\,ks \\xmm\\ exposure taken in 2000.  Both facilities have their advantages: while \\xmm\\ globally has a higher sensitivity, {\\it Chandra} has a lower background, and a higher spectral resolution and sensivity at short wavelengths for its grating spectra.  Today, however, much more data are available (see below).  Considering the uncertainties in the line profile results and the lack of new variability studies, we decided to re-investigate \\zp\\ using the best dataset available at the present time: 18 \\xmm\\ exposures, corresponding to an exposure of 1Ms totaling $>$700\\,ks of useful time (i.e. an improvement by an order of magnitude compared to most previous studies). This dataset thus provides the most detailed X-ray view of an O star to date. The results that we obtained will be presented in a series of papers. The first one will present the data, their reduction, and a few first results; the second one will focus on the X-ray variations of \\zp, using EPIC and RGS data; the last one will present a global analysis of the high-resolution X-ray spectrum, using the merged high-resolution data.  This first paper is organized as follows. The dataset and its reduction are presented in Sect. 2, the spectral fits are presented in Sect. 3, the individual line profile fitting in Sect. 4, and the results are summarized and discussed in Sect. 5. ",
        "conclusions": "In the past decade, about 1Ms of data were obtained by \\xmm\\ on \\zp. Of these, about 30\\% is strongly affected by flares, reducing the useful exposures to 579ks for EPIC-MOS, 477ks for EPIC-pn, and 751ks for RGS. A variety of modes was used, the most reliable being the SW+Thick Filter mode; the use of the medium filter yielded piled-up data, while the data taken in LW+Thick Filter mode do not appear significantly different from those obtained with SW+Thick Filter mode. Attention was paid to this problem, notably by choosing similar extraction regions for all datasets and using only single events.  Broad-band EPIC data taken with the Thick filter were analyzed using absorbed optically-thin thermal emissions. Four temperatures are needed to reproduce in a reasonable way the data. Note that the fits are not formally acceptable, but that (1) instruments do not always agree with one another and (2) the reduced noise amplifies the limitations due to the imperfect atomic parameters and standard line profiles. Nevertheless, the EPIC spectra appear remarkably stable over the decade of observations, with only 3\\% dispersions around the average fluxes. A detailed variability study, based on lightcurves, will be presented in Paper II.  The combined high-resolution RGS spectrum confirms that the X-ray line profiles vary with wavelength. Fitting individual line profiles using a wind model yields similar onset radius for the X-ray emission, but wind continuum opacities depending on wavelength. This is simply due to the fact that the cool absorbing clumps in the wind are not fully optically thick at all wavelengths, though further modelling is needed in order to adequately reproduce the opacity variations."
    },
    "1112/1112.0786_arXiv.txt": {
        "abstract": "Astrophysical sources of TeV gamma rays are usually established by Cherenkov telescope observations. These counting type instruments have a field of view of few degrees in diameter and record large numbers of particle air showers via their Cherenkov radiation in the atmosphere. The showers are either induced by gamma rays or diffuse cosmic ray background. The commonly used test statistic to evaluate a possible gamma-ray excess is Li and Ma (1983), Eq.\\ 17, which can be applied to independent on- and off-source observations, or scenarios that can be approximated as such. This formula however is unsuitable if the data are taken in so-called \"wobble\" mode (pointing to several offset positions around the source), if at the same time the acceptance shape is irregular or even depends on \\bref{operating parameters such as the pointing direction or telescope multiplicity}. To provide a robust test statistic in such cases, this paper explores a possible generalization of the likelihood ratio concept on which the formula of Li and Ma is based. In doing so, \\bref{the multi-pointing nature of the data and the typically known instrument point spread function are fully exploited} to derive a new, semi-numerical test statistic. Due to its flexibility and robustness against systematic uncertainties, it is not only useful for detection purposes, but also for skymapping and source shape fitting. \\bref{Simplified Monte Carlo simulations are presented} to verify the results, and \\bref{several applications and further generalizations of the concept are discussed}. ",
        "introduction": "\\label{sec:intro}  The field of very-high-energy (VHE, $>100\\gev$) gamma-ray astronomy is currently in its third generation of instruments, with an advanced future project, CTA \\citep{cta}, already being in its preparatory phase. The currently operated systems H.E.S.S. \\citep{hess}, MAGIC \\citep{magicstereoperformance} and VERITAS \\citep{veritas} have established more than 100 VHE sources in the sky\\footnote{http://tevcat.uchicago.edu/}. The acceptence of these telescope systems covers few degrees in diameter, and declines smoothly with increasing distance from the pointing position. Establishing a gamma-ray signal from an astrophysical source requires to significantly \\bref{prove} a gamma-ray excess over background events that typically remain dominant even after selection cuts. This background is mostly composed of diffuse hadrons, part of which appears almost identical \\bref{to} electromagnetic showers and has to be considered to be irreducible \\citep{irreducible}. Besides these \"gamma-like hadrons\", the irreducible background also contains smaller fractions of diffuse electrons and gamma rays. \\bref{This irreducible background, and the statistical and systematic uncertainties that come with it, are one of the main limiting factors of TeV astronomy; usually, an observational campaign for a given source either reveals one source or none, and only in few cases, or if the effort of a large scan is undertaken, additional unexpected sources are detected. Therefore, a statistical source detection technique that is both sensitive to weak sources and stable against systematic effects is of crucial importance to the field.}  The standard test statistic to evaluate an excess of gamma rays from a given sky direction is Li and Ma \\cite{lima}, Eq.\\ 17, which hereafter will be refered to as $S\\tin{LM}$. It was established among several alternatives they evaluated in their paper, based on the fact that this likelihood ratio test statistic was the only one that yielded a satisfyingly Gaussian null-hypothesis distribution. The formula they presented was designed to compare the event numbers of an on- and off-source observation ($N\\tin{on}$, $N\\tin{off}$), and allows for a scaling factor $\\alpha$ between the effective observation times ($t\\tin{on} = \\alpha\\, t\\tin{off}$) to account for unequal exposures.  In modern observation practice, most Cherenkov telescope data are \\textit{not} taken in On/Off-mode (see Fig.~\\ref{fig:on_off_wobble}, left), because this strategy implies a lot of observation time dedicated to empty sky regions. Also, it is prone to systematic differences between the on- and off-data caused by instabilities in electronic or atmospheric \\bref{operating} conditions, especially if the off-data \\bref{could not be scheduled contemporaneously enough with the on-source observation.} Therefore, usually the \"wobble\" technique \\citep{wobble} is applied, in which the data  are taken at two or more observation positions offset in different directions from the main target coordinates (see Fig.~\\ref{fig:on_off_wobble}, right). In this way, each wobble set provides both on- and off-data at the same time. In its original idea, the exposure shape is considered to have some circular symmetry, and for each wobble set the off-data can be taken from the same observation, at sky positions of similar distance to the telescope pointing position (\"reflected regions\" \\citep{reflectedregion}). In that case, the off-exposure ratio $\\alpha$ is the same for all wobble data sets, and $S\\tin{LM}$ can be applied after summing up the on- and off-events of all wobble sets. This constant $\\alpha$ can also be achieved approximately if the off-data are taken from \\textit{hadron-like} background events (\"template background\", \\citep{templatebackground}), or a ring area around the source position (\"ring background\" \\citep{hessskymapping}), both of which require \\bref{also symmetry assumptions or} an efficiency correction through Monte Carlo simulations of the isotropic background.  \\begin{figure} \\centering \\includegraphics[width=13cm]{on_off_wobble.eps} \\caption{ Observation schemes with asymmetric acceptance shapes (green areas). The sky position to be evaluated for a signal is marked with a yellow star, the corresponding off-regions with grey stars. Left: Original On/Off scheme. The off-data provides $N\\tin{off}$, the on-data $N\\tin{on}$, and it needs one exposure ratio $\\alpha=t\\tin{on}/t\\tin{off}$ between those. Right: Wobble scheme with four observation positions. The off data of each wobble set can be taken from all other wobble sets, resulting in four $\\alpha$ parameters and their corresponding $N\\tin{on}$, $N\\tin{off}$. } \\label{fig:on_off_wobble} \\end{figure}  In the general case, though, the acceptance symmetry might not hold and a Monte Carlo  correction may involve too high uncertainties. This occurs for instance in very \\bref{low-energy observations}, where camera-based acceptance inhomogeneities (dead pixels, trigger fluctuations) are both difficult to model in simulations and furthermore lead to \\bref{features in the acceptance shape that can depend on the Alt/Az pointing direction or other operation parameters of the system}. This is particulary troublesome in two-telescope systems like MAGIC \\citep{magicstereoperformance}, where already the basic geometrical overlap of the two fields of view implies an elongated, Alt/Az-dependent exposure \\bref{shape}. Under these conditions, the \"reflected regions\" approach does not hold, and the off-data for each wobble set have to be taken from the other wobble sets (see Fig.~\\ref{fig:on_off_wobble}, right). This approach can provide \\bref{a sensitive measurement}, because with several wobble sets, the off regions are numerous and well-populated, but it results in a different $\\alpha$ for each wobble set, which is not supported by $S\\tin{LM}$. On top of that, if this procedure has to be done \\bref{separately for different types of data (be it for instance different Azimuth angles or telescope multiplicities)}, it leads to many more $\\alpha$-parameters, and in general, some off-events may happen to be oversampled if they lie in more than one off-region, which is also not considered in $S\\tin{LM}$. As a consequence, while the background density can still be modeled under certain assumptions and some numerical effort \\citep{icrcadvanced}, the test statistic $S\\tin{LM}$, if still applied in some way, becomes very approximative. This may be dealt with in practice by making additional high requirements to the signal-to-background ratio of a detection \\citep{magicstereoperformance}, which however \\bref{limits the effective sensitivity of the instrument}.  Therefore, unlike previous efforts \\citep{hessskymapping,icrcadvanced}, this work will not pursue to extract the variables needed for $S\\tin{LM}$ through the complex task of explicit background modeling. Instead, \\bref{the likelihood ratio concept behind that test statistic will be generalized, and new formulae will be derived} that can directly be applied to multi-wobble Cherenkov telescope data. To do so, \\breftwo{no} Monte Carlo simulations or exposure symmetry assumptions will be required. It will only be assumed that the telescope acceptance shape is the same for different wobble data sets if they are taken \\bref{under similar operating conditions}\\footnote{Note that if this basic assumption does not hold, any other symmetry assumption also breaks.}.  Besides that, \\bref{this work will also address the disadvantage of $S\\tin{LM}$ that it} depends on the size of the signal region in which $N\\tin{on}$ is calculated. This area is usually defined through an integration radius (\"$\\theta^2$-cut\", with $\\theta$ being the angular distance between reconstructed gamma direction and source position). Its optimization either depends on the background density and an assumption of the source strength, or may involve several trials. \\bref{In this paper,} these assumptions \\bref{are reduced} by accomodating the knowledge of the point spread function (PSF) in the formulae, which makes them independent of the source strength or background density.  Likelihood ratios are frequently used to convert a complex likelihood maximization problem to a test statistic that follows a $\\chi^2$-distribution. This possibility was first proven in \\cite{llhratio}, and was suggested for astronomical purposes in \\citep{llhratioastronomy}. The technique is now widely used in counting type experiments, mainly in X-ray and gamma-ray astronomy, and can be applied both for detection and optimization purposes \\citep{llhrvschi2}. \\bref{It should be pointed out} that criticism and potentially more accurate or more general alternatives to the likelihood ratio concept exist \\citep{bayesianlima,limacritics,fisherexactlima}, but are not subject of this work.  The structure of the paper is to first define the coordinate systems and naming conventions needed for the calculations. Then, a likelihood function is set up and maximized to gain all relevant parameters \\breftwo{(Sec.~\\ref{sec:llh})}. Based on that, the test statistic is formulated in Sec.~\\ref{sec:ts} and its intrinsic inclusion of $S\\tin{LM}$ is demonstrated. In Sec.~\\ref{sec:generalizations}, some further possible generalizations and applications of the formulae are discussed and a recipe for \"Likelihood Ratio Skymapping\" is suggested. Finally, some example toy simulations are shown in Sec.~\\ref{sec:mc} to verify the method. ",
        "conclusions": "This paper derives a new generalized test statistic (\\eref{eq:significance}) that can be used for significance calculation in VHE astronomy. The advantages over the existing test statistics are that it flexibly takes into account any number of data subsets from different wobble coordinates and \\bref{operating conditions of the system}, even if the acceptance shape is very irregular \\bref{and different between these operating conditions}. Also, it \\bref{takes} advantage of a known gamma-ray PSF while not requiring the optimization of an integration radius (\"$\\theta^2$-cut\").  The test statistic can be applied to any position in the field of view, so it is very suitable for skymapping purposes. The advantages of this approach is that the test statistic only makes minimal assumptions on the acceptance field of view and does not require any exposure symmetry or Monte Carlo simulations. It is hence unaffected by many systematic uncertainties. \\bref{A} \"Likelihood Ratio Skymapping\" procedure \\bref{is suggested} in Sec.~\\ref{sec:procedure}.  The log-likelihood function (\\eref{eq:lllhfphi}) can also be applied to fit the shape and position parameters of the source. \\bref{The} formulae are furthermore extendable to accomodate established sources in the field of view, a non-homogeneous PSF shape or gamma-to-hadron acceptance ratio, an unbinned analysis approach or the \"orbit\" observation mode. \\bref{If the background event density is sufficiently high, the relative} \\breftwo{excess} parameter $\\phi$ is well-suited to calculate a gamma-ray flux map of the field of view.  \\bref{In several simulated scenarios it is verified} that the test statistic can not only handle the difficult situations it is designed for, but \\bref{also seems to be systematically higher in the signal case, and therefore more sensitive,} than the commonly used test statistic after Li and Ma \\cite{lima}, Eq.\\ 17."
    },
    "1112/1112.2260_arXiv.txt": {
        "abstract": "In order to understand environmental effects on star formation in high-redshift galaxies, we investigate the physical relationships between the star formation activity, stellar mass, and environment for $z\\simeq1.2$ galaxies in the 2 deg$^2$ COSMOS field$^*$.  We estimate star formation using the \\OII$\\lambda 3727$ emission line and environment from the local galaxy density.  Our analysis shows that for massive galaxies ($M_* \\gtrsim 10^{10} M_{\\sun}$), the fraction of \\OII\\ emitters in high-density environments ($\\Sigma_{\\rm 10th}\\gtrsim 3.9\\ {\\rm Mpc^{-2}}$) is $1.7\\pm 0.4$ times higher than in low-density environments ($\\Sigma_{\\rm 10th} \\lesssim 1.5\\ {\\rm Mpc^{-2}}$), while the \\OII\\ emitter fraction does not depend on environment for low-mass $M_* \\lesssim 10^{10} M_{\\sun}$ galaxies.  In order to understand what drives these trends, we investigate the role of companion galaxies in our sample.  We find that the fraction of \\OII\\ emitters in galaxies with companions is $2.4\\pm0.5$ times as high as that in galaxies without companions at $M_*\\gtrsim 10^{10} M_{\\sun}$.  In addition, massive galaxies are more likely to have companions in high-density environments.  However, although the {\\it number} of star forming galaxies increases for massive galaxies with close companions and in dense environments, the {\\it average} star formation rate of star forming galaxies at a given mass is independent of environment and the presence/absence of a close companion.  These results suggest that interactions and/or mergers in high-density environment could induce star formation in massive galaxies at $z\\sim1.2$, increasing the fraction of star-forming galaxies with $M_* \\gtrsim 10^{10} M_{\\sun}$. ",
        "introduction": "A key question in understanding the formation and evolution of galaxies is to find what physical parameters are most sensitive to the star formation process in galaxies; e.g., how the star formation activity depends on the environment, and how the relation between star formation and environment changes during the course of galaxy evolution over 10 Gyrs.  Observational properties of galaxies in the local universe have been extensively studied in the past.  Typically, actively star-forming galaxies in the local universe have lower masses than passive galaxies and the most massive galaxies tend to be inactive for star formation (so-called mass-downsizing; Cowie et al. 1996).  In addition, the star formation activity strongly depends on environment: star formation rate (SFR) decreases with increasing galaxy density (e.g., Lewis et al. 2002; Gomez et al.2003; Kauffmann et al. 2004; Mahajan et al. 2010) and the fraction of star-forming galaxies also decreases with increasing galaxy density (e.g., Carter et al. 2001; Balogh et al. 2004; Mateus \\& Sodr\\'e 2004).  The fraction of early-type (passive) galaxies is higher in higher-density regions while the fraction of late-type (star-forming) ones is lower in such environment, called the ``morphology-density relation'' and consistent with the above results (e.g., Dressler 1980; Dressler et al. 1997; Goto et al. 2003; Capak et al. 2007).  These findings indicate that the star formation activity in galaxies is strongly related to both the stellar mass and their environment.  On the other hand, it is well known that the star formation rate density (SFRD) steeply increases in the first $\\sim 2$ Gyrs, peaks at $z\\sim 1-3$, and then decreases by an order of magnitude toward the present day (e.g., Madau et al.1996; Lilly et al. 1996; Shioya et al. 2008; Bouwens et al. 2009).  These results suggest that the redshift range of $z \\sim 1-3$ is the key epoch for the most active cosmic star formation in the history of universe.  Therefore, studies of galaxies at $1 < z < 3$ are important for understanding galaxy evolution.  In recent years, wide and deep surveys have allowed us to study star formation activity in $z \\gtrsim 1$ galaxies.  For example, Elbaz et al. (2007) and Cooper et al. (2008) studied the relation between the star formation rate (SFR) in galaxies and the environment at $z\\sim 1$, in the Great Observatories Origins Deep Survey (GOODS; Giavalisco et al. 2004) and the DEEP2 Galaxy Redshift Survey (DEEP2; Davis et al. 2003).  These studies showed that the average star formation rate increases with increasing galaxy density (Elbaz et al. 2007; Cooper et al. 2008).  We also investigated the relation between the fraction of star-forming galaxies with \\OII\\ $\\lambda 3727$ emission (hereafter ``\\OII\\ emitters'') and the local galaxy density at $z\\simeq 1.2$ in the Cosmic Evolution Survey (COSMOS; Scoville et al. 2007a; Koekemoer et al. 2007) and found that the fraction of star-forming galaxies in high-density regions is as high as that in low-density regions (Ideue et al. 2009; hereafter ``Paper {\\sc i}'').  Similar trends were also reported by other observational studies in groups of galaxies (Iovino et al. 2010; Li et al. 2011; Sobral et al. 2011) and cluster environments (Hayashi et al. 2010; Hilton et al. 2010) at $z \\gtrsim 1$.  Since the fraction of star-forming galaxies decreases with increasing galaxy density in the local universe as mentioned above, these results suggest that the relationship between the star formation activity in galaxies and their environment dramatically changed from $z\\sim1$ to the present day.  On the other hand, many studies suggest that the star formation activity or the fraction of star-forming galaxies also strongly depends on stellar mass even at $z\\gtrsim 1$ (e.g., Noeske et al. 2007; Elbaz et al. 2007; Ilbert et al. 2010).  Therefore, it is important to investigate the relationships between the star formation activity, the stellar mass of galaxies, and the environment at $z\\sim 1$, in order to reveal the origin of the observed difference in the environmental dependence of the star formation activity between $z\\sim0$ and $z\\sim1$.  In this paper, we focus on the relationships between the star formation, stellar mass, and environment of galaxies at $z \\simeq 1.2$ in the COSMOS field. We have already carried out Subaru imaging observations of the COSMOS field (Taniguchi et al. 2007). We can obtain the sample of \\OII\\ emitters in the COSMOS field by using a narrowband filter NB816 ($\\lambda_c = 8150 $ \\AA\\ and $\\Delta \\lambda {\\rm (FWHM)} = 120 $ \\AA, see Takahashi et al. 2007) and estimate their SFR by using the \\OII\\ luminosity.  The \\OII\\ emission line provides a good estimator of the star formation in galaxies at intermediate redshift (e.g., Kennicutt 1998; Jansen et al. 2001).  The stellar mass in galaxies in the COSMOS field was obtained from the fitting of their spectral energy distributions (SEDs; Ilbert et al. 2010).  Thanks to the very large survey area ($\\sim$ 2 deg$^2$) of the COSMOS field, we are able to obtain an unbiased picture of the star formation activity in galaxies, avoiding the cosmic variance effect.  Indeed, the COSMOS field includes various regions with a wide range of galaxy density (e.g., Scoville et al. 2007a, 2007b; Takahashi et al. 2007), which enables us to investigate systematically the environmental effects on galaxy properties.  This paper is organized as follows. In section 2, we describe the sample selection and possible contamination of the \\OII\\ emitter sample.  Descriptions of the measurements of local galaxy density, stellar mass, and SFR are presented in section 3.  We investigate the dependence of star formation on environment and stellar mass, and the effect of close companions on the star formation activity, for our sample at $z\\simeq 1.2$ in section 4.  In section 5, we summarize our findings and discuss implications of our results on the evolution of galaxies from $z\\sim 1.2$ to $z\\sim 0$.  Throughout this paper, magnitudes are given in the AB system. We adopt a flat universe with the following cosmological parameters; $\\Omega_{\\rm matter} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_0 = 70\\; {\\rm km\\;s^{-1}\\;Mpc^{-1}}$. ",
        "conclusions": "In this paper, we have investigated the relations between the star formation, stellar mass, and environment at $z\\simeq 1.2$ in the COSMOS field.  As shown in section 4, our results are different from the observational properties of galaxies in the local universe.  Based on our observational results, we discuss possible origins of the difference between $z\\sim 0$ and $z\\sim 1$ from a viewpoint of the evolution of galaxies.  \\subsection{The Observational Properties of Galaxies at $z \\simeq 1.2$ in the COSMOS Field}  The findings of this study are summarized by the following:  \\begin{enumerate}  \\item At $M_* \\gtrsim 10^{10} M_{\\sun}$, the fraction of \\OII\\ emitters in high-density environments is $1.7\\pm0.4$ times as high as in low- and intermediate-density environments.  The fraction of \\OII\\ emitters does not depend on environment at $M_* \\lesssim 10^{10} M_{\\sun}$ (Figure \\ref{fracoii-mass-env}).  \\item The fraction of \\OII\\ emitters in galaxies with a companion (likely interacting galaxies) is $2.4\\pm0.5$ times higher than that in those without a companion (likely isolated galaxies) over $M_* \\sim 10^{10}-10^{11.5} M_{\\sun}$ (Figure \\ref{fracoii-mass-compnoncomp}).   \\item The fraction of galaxies with a companion for the \\OII\\ emitters at high mass ($M_{*}=10^{10}-10^{11.5}M_{\\sun}$) is $1.5\\pm0.3$ times higher than that at low mass ($M_{*}<10^{10}M_{\\sun}$), while that for the non-\\OII\\ emitters at high mass is $0.7\\pm0.2$ times that at low mass (Figure \\ref{fraccomp-mass-oiinonoii}).   \\item The fraction of \\OII\\ emitters with a companion is higher in the higher-density environments (Figure \\ref{fraccomp-mass-env}). The fractions in the low-, intermediate-, and high-density environment at high mass ($M_{*}=10^{10}-10^{11.5}M_{\\sun}$) are $0.2\\pm0.1$, $0.3\\pm0.1$, and $0.4\\pm0.1$, respectively.   \\item The average SFR of \\OII\\ emitters strongly correlates with stellar mass in all the environments. For example, the linear fit gives the relation in the high-density environment, $\\log {\\rm <SFR>}=(0.59\\pm0.08)\\times\\log M_{*}-(4.56\\pm0.78)$. The SFR at a given mass is independent of the environment and the presence of a close companion (Figures \\ref{sfr-mass-env} and \\ref{sfr-mass-compnoncomp}).  \\end{enumerate}  From the first item described above, it is expected that high-mass \\OII\\ emitters contribute significantly to the active star formation in high-density regions at $z\\simeq 1.2$.  Moreover, items 2, 3, and 4 suggest that interactions and/or mergers could trigger star formation in massive galaxies preferentially in high-density environments, while the star formation in low-mass galaxies appears to be independent of environment and does not seem to be affected by the interactions and/or mergers.  In other words, $EW$(\\OII) of some inactive high-mass galaxies may become higher due to the active star formation induced by galaxy interactions, and thus we observe such high-mass galaxies as \\OII\\ emitters.  On the other hand, considering item 5, it appears that the strength of star formation activity in the star-forming galaxies is not influenced by environmental effects, although the environment and interactions affect the fraction of star-forming galaxies.  The environmental dependence of the \\OII\\ fraction is seen only at high mass.  It is interesting to note that the influence of interactions and/or mergers on star formation depends both on environment and stellar mass.  This can be interpreted such that massive galaxies need external triggers (i.e., interactions and/or mergers) for active star formation, while the less massive galaxies are form stars regardless of the presence of a close companion.  One possible triggering mechanism for less massive galaxies is minor merger events (e.g., Taniguchi \\& Wada 1996).  However, such satellite galaxies would be too faint to be detected in our imaging survey.  As an another possible interpretation, it is expected that less massive galaxies contain much gas and can naturally form stars since evolution of those galaxies is slow compared to massive galaxies, which is expected from mass-downsizing (e.g., Cowie et al. 1996).  Therefore, we conclude that the environmental dependence of the \\OII\\ fraction (i.e., the fraction of star-forming galaxies) is seen only at high mass, because 1) interactions and/or mergers induce the star formation only in massive galaxies and 2) the probability of interactions and/or mergers depends on local galaxy density.  Peng et al. (2010, 2011) recently found that the environment quenching (or the satellite quenching) depends on the local galaxy density but not on the stellar mass of galaxies.  We observe a different environmental effect from that found by Peng et al. (2010, 2011), since our results suggest that the influence of interaction and/or mergers depend on both environment and stellar mass.   \\subsection{Implications to the evolution of galaxies from $z\\sim 1$ to $z\\sim 0$}  In the local universe, it is observed that the fraction of red galaxies in high-density environments is higher than that in low-density environments, and this difference becomes larger at lower mass (e.g., Baldry et al. 2006).  Namely, the fraction of blue galaxies in high-density environments is lower than that in low-density environments, especially at low mass.  Iovino et al (2010) also found a similar trend at $z<0.5$.  Furthermore, they found that the fraction of blue galaxies at $z=0.6-0.8$ does not depend on the environment in $M_* > 10^{10}$.  These are in contrast to the environmental dependence of the \\OII\\ fraction in massive galaxies at $z\\simeq 1.2$ mentioned above.  Now a question arises: what is the origin of the difference between $z\\sim 1$ and $z\\sim 0$?  Our results suggest that interactions and mergers are likely to induce star formation in massive galaxies in high-density environments at $z\\simeq 1.2$, resulting in the observed high \\OII\\ fraction of massive galaxies in high-density environments. These inactive massive galaxies where star formation would be induced by interactions are likely to have some cold gas for star formation. We have also found that the average SFR at a given mass is similar among the different environments and is independent of the presence of a close companion (Figures \\ref{sfr-mass-env} and \\ref{sfr-mass-compnoncomp}), the strength of the star formation induced by interactions in massive galaxies does not seem to be significantly different from other massive star-forming galaxies on average.  Therefore we expect that there are already some inactive high-mass galaxies with sufficient cold gas for SFR at a given mass at $z\\simeq 1.2$.  On the other hand, the fraction of gas-poor galaxies is higher in higher density regions at low-$z$ (e.g., Giovanelli \\& Haynes 1985; Bertram et al. 2006).  Therefore it is unlikely that a significant starburst occurs in interactions or mergers between gas-poor galaxies (i.e., dry mergers; e.g., Tran et al. 2005; Cattaneo et al. 2008), even though interactions and/or mergers still occur in high-density regions at low-$z$ (e.g., Patton et al. 2011).  In this context, one important difference in galaxies between $z\\sim 1.2$ and low $z$ is the cold gaseous content.  Since galaxies at $z\\sim 1.2$ are in early stages of their evolution, it is likely that their cold gaseous content is, on average, larger than that in galaxies at low $z$ even if such galaxies are located in relatively high-density regions.  It is thus expected that interactions and mergers at $z\\simeq 1.2$ could cause intense star formation more frequently since it is known that interactions and mergers between two gas-rich galaxies (wet mergers) can induce active star formation and rapidly consume the cold gas (e.g., Barton et al. 2000; Woods et al. 2006).  In fact, Lin et al. (2008, 2010) suggest that the fraction of dry mergers gradually increases from $z\\sim 1$ to $z\\sim 0$, while the fraction of wet mergers decreases at the redshift range.  If there are many inactive high-mass galaxies with sufficient cold gas for star formation at $z\\simeq 1.2$ as mentioned above, there may be some mechanisms responsible for quenching star formation in galaxies while retaining their gas.  For example, it might be expected that high-mass galaxies have experienced starbursts in the past, and then their gas becomes hot and their star formation stops in spite of the existence of the (hot) gas.  Later the hot gas cools down into cold gas, and then high-mass galaxies actively form stars when interactions and/or mergers occur.  As another scenario, high-mass galaxies might have consumed most of gas in the past, which is expected from the mass-downsizing of galaxy formation (e.g., Cowie et al. 1996), and the gas density was already too low to cause intense star formation (Kennicutt 1989).  However, if interactions can sufficiently perturb the remaining gas, the star formation can be triggered and they can be observed as high-mass star-forming galaxies in high-density environment.  In this case, such galaxies immediately consume their gas in the star formation and these massive galaxies in high-density environment are observed to be passive at $z \\lesssim 1$.  From these considerations, we suggest the following scenario; The high-mass star-forming galaxies for which star formation are induced by interactions and/or mergers contribute significantly to the active star formation in high-density regions at $z\\simeq 1.2$.  These massive star-forming galaxies in high-density regions at $z\\simeq 1.2$ could quickly consume most of the accreted cold gas.  If this is the case, the star formation activity may not be enhanced when interactions and/or mergers occurred in high-density environments at lower redshift.  In this context, the quenching of star formation in massive galaxies in high density environments is expected to lead to the shift of major star formation in the universe from high-density regions to low-density ones at $z\\lesssim 1$.    We would like to thank the anonymous referee for her/his very useful comments and suggestions.  We also thank all members of the COSMOS team. This work was financially supported in part by the Japan Society for the Promotion of Science (Nos. 17253001, 19340046, 23244031, and 23654068). Y. I. is financially supported by the Japan Society for the Promotion of science (JSPS) through JSPS Research Fellowship for Young Scientists."
    },
    "1112/1112.0337_arXiv.txt": {
        "abstract": "We investigate whether any multi-planet systems among {\\em Kepler} candidates (2011 February release) can harbor additional terrestrial-mass planets or smaller bodies.  We apply the ``packed planetary systems'' hypothesis that suggests all planetary systems are filled to capacity, and use a Hill stability criterion to identify eight 2-planet systems with significant gaps between the innermost and outermost planets. For each of these systems, we perform long-term numerical integrations of 10$^7$ years to investigate the stability of 4000$-$8000 test particles injected into the gaps. We map out stability regions in orbital parameter space, and therefore quantify the ranges of semi-major axes and eccentricities of stable particles. Strong mean-motion resonances can add additional regions of stability in otherwise unstable parameter space.  We derive simple expressions for the extent of the stability regions, which is related to quantities such as the dynamical spacing $\\Delta$, the separation between two planets in units of their mutual Hill radii.  Our results suggest that planets with separation $\\Delta < 10$ are unlikely to host extensive stability regions, and that about 95 out of a total of 115 two-planet systems in the {\\em Kepler} sample may have sizeable stability regions.  We predict that {\\em Kepler} candidate systems including KOI 433, KOI 72/Kepler-10, KOI 555, KOI 1596, KOI 904, KOI 223, KOI 1590, and KOI 139 can harbor additional planets or low-mass bodies between the inner and outer detected planets. These predicted planets may be detected by future observations. ",
        "introduction": "Early studies of extrasolar planetary systems showed residual velocity trends in Keplerian orbit fits to radial velocity data \\citep[e.g.,][]{marc98,butl98,marc99,vogt00,fisc01}, suggesting that these systems may host additional, undetected planets. \\citet{fisc01} noted that about half of the stars in their sample of 12 systems showed residual trends greater than the expected scatter due to measurement uncertainties and stellar noise. Most of these systems were later confirmed to harbor additional planet(s).  In more recent years, the study and prediction of undiscovered planets have been aided by long-term N-body simulations. These numerical investigations searched for stability zones in multi-planet systems by integrating hundreds to thousands of test bodies, which were injected into empty regions between known planets \\citep[e.g.,][]{meno03,barn04,raym05,ji05,rive07,raym08}. For example, a putative Saturn-mass planet in HD~74156 was first predicted by \\citet{raym05} through numerical simulations that showed a stable region between planets b and c.  The planet was later discovered by \\citet{bean08}, although there have been questions about the validity of the detection~\\citep{witt09,mesc11}. This prediction was motivated by the ``packed planetary systems'' (PPS) hypothesis.  The PPS hypothesis is the idea that planetary systems are formed ``dynamically full'' and filled to capacity, and any additional planets will cause the systems to be unstable \\citep[e.g.,][]{barn04, raym05, raym06, barn07}. Consequently, planetary systems with stable stability zones between the innermost and outermost planets are suggestive of additional, yet-undetected planets. Reasons for the non-detections of hypothetical planets include lack of sufficient data, such as non-transiting planets that require more data to detect them via transit timing variations, or planetary masses/radii that are below detection limits. The orbital properties of predicted planets can be identified through long-term numerical simulations. Support for the PPS hypothesis comes from early observations of packed multi-planet systems that led to this hypothesis \\citep[e.g.,][]{butl99,marc01b,marc01a,fisc02,mayo04}, apparent consistency between the planet-planet scattering model and packed systems \\citep{raym09}, the remarkably dense and packed Kepler-11 system \\citep{liss11a}, theoretical work \\citep[e.g.,][]{cham96,smit09}, and other investigations \\citep[e.g.,][]{rive00,gozd06}.  In the present study, we apply the PPS hypothesis to multi-planet candidate systems discovered by the {\\em Kepler} team during the mission's first four and a half months of data \\citep{boru11}. The {\\em Kepler} mission is a transit survey designed to search for Earth-sized planets \\citep{boru10,koch10,jenk10,cald10}, and is sensitive to terrestrial-class and larger planets located at a large range of separations from their host star. {\\em Kepler} can detect multiple transiting systems for densely-packed planets with nearly coplanar configurations or with serendipitous geometric alignment, and the dynamics and statistics of {\\em Kepler} multi-planet systems are providing a wealth of information about planetary systems \\citep[e.g.,][]{stef10,lath11,liss11a,liss11b}.  Given that planetary systems have been discovered with densely packed planets, we seek to test the PPS hypothesis and predict additional planets in {\\em Kepler} candidate multi-planet systems. In Section \\ref{datamethods}, we discuss {\\em Kepler}'s sample of multi-planet systems as well as our methodology for evaluating each planetary system's level of dynamical packing. We also explain our methods for running numerical simulations and our choice of initial conditions. In Section \\ref{results}, we present the results from long-term numerical integrations and illustrate them using stability maps. Section \\ref{mountain} discusses the dynamical interpretation of our work, in particular the relationship between planetary spacing and the extent of an inter-planet stability region.  We then summarize the restrictions and scope of our study (Section \\ref{scope}) and state our conclusions (Section \\ref{conclusion}). ",
        "conclusions": "\\label{conclusion}  The ``packed planetary systems'' model advocates the idea that all planetary systems are formed to capacity. To test this hypothesis, we investigated the packing density of {\\em Kepler} candidate two-planet systems from the first four and a half months of the mission. Through numerical calculations, we determined whether regions of stability exist between known planets with wide separations, i.e.,~in systems that seemed the most unpacked based on how well they satisfy Hill stability. Discovery of a stable region suggests that a low-mass body could be present in the gap, which would then bring the system to a more ``packed'' state. With time, such predictions will be shown to be correct or incorrect, allowing us to gauge the success of this model.  We performed detailed numerical simulations of eight, two-planet {\\em Kepler} systems, selected using an analytical $\\beta/\\beta_{\\rm crit}$ stability criterion. In addition to the known planets, we included 4000$-$8000 test particles per planetary system, allowing both circular and non-circular, and coplanar and non-coplanar orbits. These test particles are good proxies for low-mass bodies such as terrestrial planets as well as small bodies such as asteroids or dwarf planets. We integrated all bodies for 10$^7$ years; we defined stable particles as those that remained stable during the length of the integration and unstable particles as particles that experienced a collision or ejection.  Our results (Figures \\ref{maps1} to \\ref{maps3}) indicated that all of the planetary systems investigated here (KOIs 433, 72, 555, 1596, 904, 223, 1590, and 139) can pack additional, yet-undetected bodies in the identified stable locations. We also discussed relationships relating dynamical spacing between known planets and the extent of the inter-planet stability region. We derived an analytical relationship relating the largest possible eccentricity of a stable test particle to the semi-major axes and Hill radii of the two planets surrounding the particle. We also demonstrated that $\\Delta$, the separation between two planets in units of their mutual Hill radii, can be a reasonable predictor of whether or not stability regions can exist between planets.  The cut-off occurs at a critical $\\Delta$ between 10 and 15.  We suggest that planets with separation $\\Delta < 10$ are unlikely to host extensive stability regions.  Based on this $\\Delta=10-15$ cut-off, we suggest that about 95 out of a total of 115 two-planet systems in the February 2011 {\\em Kepler} sample may have sizeable stability regions."
    },
    "1112/1112.1107_arXiv.txt": {
        "abstract": "We build a hydrodynamical model for computing and understanding the Sun's large-scale high latitude flows, including Coriolis forces, turbulent diffusion of momentum and gyroscopic pumping. Side boundaries of the spherical 'polar cap', our computational domain, are located at latitudes $\\geq 60^{\\circ}$. Implementing observed low latitude flows as side boundary conditions, we solve the flow equations for a cartesian analog of the polar cap. The key parameter that determines whether there are nodes in the high latitude meridional flow is $\\epsilon=2 \\Omega n \\pi H^2/\\nu$, in which $\\Omega$ is the interior rotation rate, n the radial wavenumber of the meridional flow, $H$ the depth of the convection zone and $\\nu$ the turbulent viscosity. The smaller the $\\epsilon$ (larger turbulent viscosity), the fewer the number of nodes in high latitudes. For all latitudes within the polar cap, we find three nodes for $\\nu=10^{12}{\\rm cm}^2{\\rm s}^{-1}$, two for $10^{13}$, and one or none for $10^{15}$ or higher. For $\\nu$ near $10^{14}$ our model exhibits 'node merging': as the meridional flow speed is increased, two nodes cancel each other, leaving no nodes. On the other hand, for fixed flow speed at the boundary, as $\\nu$ is increased the poleward most node migrates to the pole and disappears, ultimately for high enough $\\nu$ leaving no nodes. These results suggest that primary poleward surface meridional flow can extend from $60^{\\circ}$ to the pole either by node-merging or by node migration and disappearance. ",
        "introduction": "It has been demonstrated both observationally and theoretically that meridional circulation plays a crucial role in the workings of the solar cycle. In flux-transport dynamos, the meridional circulation is responsible for determining the cycle period \\citep{ws91,dc99,krs01} and also plays an important role in setting the cycle 'shape', that is, its rise and fall patterns. It follows that it is very important for understanding and predicting solar cycles that we have the best possible information about meridional circulation on the Sun. The accurate measurements of surface flow-patterns as a function of latitude as well as time would be very useful for simulating the evolutionary pattern of the Sun's large-scale fields, particularly the polar fields \\citep{bsss04,wls05,ddg08}.  \\begin{figure}[hbt] \\epsscale{1.0} \\plotone{f1lores.eps} \\caption{The Sun's characteristic conveyor-belts during cycles 22 (left frame) and 23 (middle frame), derived from surface observations and mass-conservation, are shown. The speed associated with the colors is shown in the colorbar (right frame). The maximum surface flow-speeds at $25^{\\circ}$ latitude were almost the same in both the conveyor-belts, but the flow turning down towards the equator around $60^{\\circ}$ latitude during cycle 22 made the length of the primary conveyor-belt shorter than that during cycle 23, in which flow went all the way to the pole before turning equatorward.} \\label{mc2223} \\end{figure}  While meridional flow at all latitudes is important for the solar dynamo, its pattern near the poles of the Sun may be particularly significant. Recently \\citet{dgdu10} showed the variation in solar cycle duration could be explained by variations in the latitudinal extent of the Sun's primary 'conveyor belt' or meridional circulation. Observations of the Sun's surface Doppler plasma flow from Mount Wilson Observatory data indicate that in both cycles 22 and 23, the maximum poleward surface flow-speed in the primary belt was the same (\\citet{ub05}; also see the figure 1 of \\citet{dgdu10}). But in cycle 22, the primary circulation cell flowed poleward only to about $60^{\\circ}$ latitude, thus making a shorter path for the magnetic flux transport via the conveyor belt and resulting in a cycle duration of $\\sim 10.5$ years (see the left frame of Figure 1). On the other hand, in cycle 23, the poleward surface flow went all the way to the poles (see the middle frame in Figure 1), leading to a longer path via the conveyor-belt cycle of $\\sim 12.5$ years.  In the past, some flux-transport dynamo calculations \\citep{bebr05} and surface-transport calculations \\citep{jcss08} have dealt with possible multi-cell meridional flow scenarios, in the context of understanding the role of these flows in solar cycle features. We now see that these studies are not just the merely playing with models; such scenarios could happen in reality. The change in the surface poleward flow-pattern -- its reversing around $60^{\\circ}$ as it did in cycle 22 and maintaining poleward flow all the way to the pole as in cycle 23 -- have significantly impacted the duration of cycle 23 and the length of the minimum that followed it, compared to that in cycle 22.  The plasma velocity can also be determined by helioseismic analysis, either from ring diagrams or from time-distance diagrams \\citep{gdsb97, hetal02,zk04,gkhhk08,gbs10}. Alternatively, features seen on the images such as magnetic structures and supergranule cells, can be tracked with cross-correlation analysis to yield a drift velocity for that feature \\citep{khh93,sd96,sks06,szk07,sksb08}. \\citet{u10} reanalyzed the Mount Wilson surface Doppler data for cycles 22 and 23, from 1986 through 2009. \\citet{u10} computed meridional flow profiles up to at least $80^{\\circ}$ latitude, found smooth evolution of the signal from one year to the next, and confirmed the difference in high latitude flow patterns between cycles 22 and 23, in particular, the existence of a second reversed cell poleward of about $60^{\\circ}$ during most of cycle 22, and a single cell with poleward surface flow all the way to the poles during the major part of cycle 23.  Differential rotation throughout the solar convection zone is fairly well known from helioseismic measurements \\citep{t04}. It is difficult to measure at high latitudes because of foreshortening and other effects. Most methods measure the linear rotational velocity rather than the angular rotation rate, so it is particularly difficult to calculate the angular measure with the short moment-arm near the poles. It appears that on average the angular measure of differential rotation at the surface declines monotonically to the poles \\citep{b00}, though the existence of a 'polar vortex' has been suggested by theory \\citep{g79}. Although our focus on this paper is more on meridional flow at high latitudes than on differential rotation there, our model will calculate both quantities, so we will need to compare results from the model with both flows.  Given the important consequences of a second, reversed meridional cell in high latitudes, it is important to develop hydrodynamical theories that could indicate what to expect to occur in the Sun. What physics determines the presence or absense of a second cell? Can we develop a simple physical argument that we should expect to see a single primary cell, or two or more cells? This is the question we attempt to answer in this paper. The theory we develop filters out all convective instability and concentrates on forcing high latitude meridional circulation mechanically by the primary meridional flow, and possibly by differential rotation, from lower latitudes. We recognize that in the Sun there could be both this mechanical forcing and axisymmetric convection in high latitudes.  Meridional circulation is produced in virtually all fluid dynamical models used to simulate and understand the origins of the differential rotation of the Sun. These fluid dynamical models fall generally into two classes: mean field models that are axisymmetric and include parameterizations of turbulent transport of momentum \\citep{drsc89,r05}, and global 3D numerical models that simulate global convection in a deep rotating spherical shell \\citep{mbdt08}. For both types of models there are numerous results that show a wide variety of meridional flow patterns, but recently both approaches have been converging toward a common result of a dominant meridional flow cell that has poleward flow near the outer boundary, and equatorward return flow near the bottom \\citep{r05,mbdt08}.  For some parameter choices the mean field models give a second, reversed cell at high latitudes \\citep{r05} and the 3D global convection models can give multiple high latitude cells \\citep{mbdt08}. Both classes of models also show that the percentage fluctuations in the amplitudes of meridional flow with time are much larger than for the differential rotation. This appears to be due to time-fluctuations in the Reynolds stresses and the fact that the meridional flow is a result of a slight imbalance between large forces, while the forces responsible for differential rotation are small enough to make the differential rotation change only very slowly.  These models are all global, and focused primarily on the problem of understanding the details of differential rotation with depth and latitude in the Sun, as well as such features as torsional oscillations. To focus primarily on the fluid dynamics of high latitudes, it is possible to build simpler models, particularly at first, and then build up to more realistic models more comparable to the global models just described. This is the approach we take below. ",
        "conclusions": "We have developed a relatively simple hydrodynamical model of the circulation at high latitudes in the solar convection zone and photosphere that contains only three forces: pressure gradients, viscous and Coriolis forces. The model equations are solved in a cartesian 'analog' of a spherical polar cap that leaves out curvature effects as well as the large density increase with depth. This system is assumed to be stress-free at the top and bottom, corresponding to the top and bottom of the solar convection zone. It is forced with meridional flow and differential rotation, guided by observations, imposed at the low latitude boundary of the cap, placed at latitudes between $60^{\\circ}$ and $70^{\\circ}$ latitude. While the inclusion of such a boundary is artificial, it is intended to separate the physics of low and mid-latitudes, responsible for the primary meridional circulation cell that has poleward flow at the top, from the physics active at high latitudes, which should be the primary determinant of the circulation found there.  Our general results are that, as the turbulent viscosity is increased, the number of nodes decreases. As the meridional flow is increased for a given turbulent viscosity, the latitude of a node increases; or, in other words, the node migrates poleward. The first of these results is explained by the viscous forces increasingly overpowering the deflecting effect of the coriolis force to allow the poleward flow of the primary cell to reach a higher latitude. The second is due to the increased poleward momentum of the poleward moving particles that allows them to reach a higher latitude before sinking down to feed the return flow. Overall, we find that our general results are not particularly sensitive to the differential rotation imposed at the boundary, provided it is plausible for the Sun. Unlike for the meridional circulation, changes in the differential rotation of the Sun at all latitudes are very small percentages of the mean differential rotation.  Most interestingly, we find that the decrease in the number of nodes as the turbulent viscosity is increased is not monotonic. In particular, in the neighborhood of $\\nu=10^{14} {\\rm cm}^2\\,{\\rm s}^{-1}$ in this model, we find that two nodes merge as the meridional flow is increased, leaving no nodes for higher meridional flow at the boundary. This is true even though, for still higher viscosity, there remains one node. Two nodes implies the presence of both a reversed (clockwise) cell and, on its poleward side, a second, weak, counterclockwise cell that has a poleward flow near the top, just as in the primary cell. With the merger of the two nodes, the reverse cell is squeezed out, merging the primary cell with its polar counterpart.  We speculate that it is this phenomenon that could be responsible for the primary cell reaching all the way to the poles during much of cycle 23, whereas there was a node near $65^{\\circ}$ in both cycles 22 and the early stages of cycle 24. At present, observations can not tell us whether there is a second node near $80^{\\circ}$ as our model predicts. Better observations in the future of velocities at the highest latitudes would obviously be valuable for testing this theory and that of the spherical polar cap that will be developed later. Better information about variations of meridional flow with depth at high latitudes would also be very useful.  Our solutions are for steady flow, but we know that the meridional flow at high latitudes on the Sun changes with time. Given our results, we can expect that a time dependent theory could determine whether changes in meridional flow at the boundary would lead to changes in meridional flow at higher latitudes that agree with observed changes in flow at the highest latitudes. If the Sun merges two nodes with a meridional flow increase, does the model do the same? Does the model predict a merger as a result of an increase in meridional flow that is not observed on the Sun? Questions such as these can only be answered with a time dependent model.  Our results suggest that it may not be necessary to invoke any additional physics to explain changes in the meridional flow cells with time in high latitudes. That does not mean, however, that the presence or importance of such physics can be ruled out at this time.  There are several additional effects that should be included in the high latitude meridional circulation model to make it more realistic. Even keeping the same physics, results from this model using spherical geometry could change significantly. With spherical geometry the meridians converge to the pole, making it harder for as much mass flux to reach the pole as does in the straight channel. In addition, the spherical problem does not separate easily in radius and colatitude due to the Coriolis forces. This has the effect of linking different latitudes and different depths of the flow in ways not present in the cartesian analog. All of these spherical effects should influence the structure of the flow, including the location of nodes in the streamfunction, and how many nodes there are for a given turbulent viscosity.  Within spherical geometry, the flow patterns will change substantially when the density increase with depth through the convection zone is included; the patterns of mass flux, or $\\rho v$, could look somewhat like those without the density variation, but the velocity itself should decline substantially with depth. This effect could also change the location and number of nodes present. In this version of the model, effects of the variation of the turbulent viscosity with depth should also be studie; we have already seen that the number and latitude of nodes in the meridional flow is sensitive to the turbulent viscosity used. Allowing this quantity to vary with radius in the spherical case could change the results substantially, as could allowing for the density increase with depth.  In its present form, the model is for axisymmetric motions, and, beyond the turbulent diffusion, our model contains no effect of organized global scale Reynolds stresses. Yet these stresses are surely important in driving the global differential rotation as well as playing a role in the maintenance of the primary meridional cell. We do not know to what latitude these Reynolds stresses reach, but they could be active within the polar cap. Their possible effects on the polar meridional flows should also be considered.  Our model currently also does not include explicitly any thermodynamics. Allowing for departures of the temperature from the adiabatic gradient could be important, particularly at the top and the bottom of the convection zone. Finally, our model is hydrodynamic, so no effects of magnetic fields are included. But the Sun is a dynamo which generates fields, some quite strong, throughout the convection zone, so the effects of these should also be taken into account. One possible effect is that different amplitudes of polar fields might influence the amplitude of meridional flow in high latitudes, as well as the locations of its nodes."
    },
    "1112/1112.3102_arXiv.txt": {
        "abstract": "Disk accretion onto weakly magnetized astrophysical objects often proceeds via a boundary layer that forms near the object's surface, in which the rotation speed of the accreted gas changes rapidly. Here we study the initial stages of formation for such a boundary layer around a white dwarf or a young star by examining the hydrodynamical shear instabilities that may initiate mixing and momentum transport between the two fluids of different densities moving supersonically with respect to each other. We find that an initially laminar boundary layer is unstable to two different kinds of instabilities. One is an instability of a supersonic vortex sheet (implying a discontinuous initial profile of the angular speed of the gas) in the presence of gravity, which we find to have a growth rate of order (but less than) the orbital frequency. The other is a sonic instability of a finite width, supersonic shear layer, which is similar to the Papaloizou-Pringle instability. It has a growth rate proportional to the shear inside the transition layer, which is of order the orbital frequency times the ratio of stellar radius to the boundary layer thickness. For a boundary layer that is thin compared to the radius of the star, the shear rate is much larger than the orbital frequency. Thus, we conclude that sonic instabilities play a dominant role in the initial stages of nonmagnetic boundary layer formation and give rise to very fast mixing between disk gas and stellar fluid in the supersonic regime. ",
        "introduction": "Accretion onto astrophysical objects possessing a material surface (as opposed to accretion onto black holes) always involves a non-trivial interaction of the incoming gas with the object's outer layers. Examples of such objects include white dwarfs in cataclysmic variables (CVs), young stars gaining material from a protoplanetary disk, and accreting neutron stars. In all these systems, one must understand how the incoming gas shares its angular momentum with the accreting object and how it mixes with the previously accreted material.  If the magnetic field of the central object is strong enough, it can disrupt the disk at some distance above the surface. Inside this region, accretion proceeds along field lines \\citep{GhoshLamb,Koldobaetal}. The critical value of the magnetic field at the surface of the central object for magnetic disruption is given by \\ba B_{*,\\text{crit}} = 5 \\times 10^4 \\left(\\frac{\\beta}{.5} \\right)^{-7/4} \\left(\\frac{M_*}{.6 M_\\odot}\\right)^{1/4} \\left(\\frac{\\dot{M}}{10^{-8} M_\\odot \\ \\text{yr}^{-1}}\\right)^{1/2}\\left(\\frac{R_*}{9 \\times 10^3 \\ \\text{km}}\\right)^{-5/4} \\text{G}. \\label{Bdisrupt} \\ea Here, we have used parameters for the mass and radius that are typical for CVs in outburst \\citep{Bergeronetal,Hachisu,Patterson}. The parameter $\\beta$ is a dimensionless factor of order unity that depends on the model adopted for the disruption of the disk by the stellar magnetic field \\citep{GhoshLamb}.  For $B_* < B_{*,\\text{crit}}$, the accretion disk will extend all the way to the surface of the star, for which there is good observational evidence in neutron star low mass X-ray binaries \\citep{GilfanovRevnivtsevMolkov} and in dwarf nova systems in outburst \\citep{WheatleyMaucheMattei}. Since the rotation rate of the star, $\\Omega_*$, is slower than the Keplerian rotation rate at the stellar surface, $\\Omega_K(R_*)$, a boundary layer (BL) will exist inside of which $d \\Omega / d R > 0$ and the rotation profile of the star attaches smoothly to that of the disk. If the accretion rate is high enough, it is also possible for material to spread meridionally to high latitudes forming a spreading layer (SL) \\citep{InogamovSunyaev, PiroBildsten}. For both a BL and a SL, the intense energy release in a localized region near the stellar surface leads to easily observable signatures such as hard spectral components, variability of emission, and so on.  A considerable amount of effort has been previously devoted to understanding the structure of well-developed, steady-state BLs, which have been evolving for long enough to establish a smooth rotation profile in their interior. An outstanding question in the study of steady-state BLs is identifying the mechanism of angular momentum transport in the layer. Several potential mechanisms have been explored over the years, among them shear instabilities \\citep{KippenhahnThomas}, baroclinic instabilities \\citep{Fujimoto, Hanawa}, and the Tayler-Spruit dynamo \\citep{Spruit, PiroBildsten1}. Despite this effort, no clear answer exists at present regarding the nature of the angular momentum transport and mixing in well-developed BLs.   \\subsection{Initiation of Boundary Layers}  An interesting aspect of the BL problem that has received less attention in the past is the issue of BL {\\it initiation}, i.e. the initial stage of BL formation which must occur when the accreted material first touches the surface of the star. In this case the physical setup is going to be quite different from the steady-state BL primarily because of the much larger velocity shear, justifying the study of BL initiation in the limit of an essentially discontinuous rotation profile at the interface between the star and accreting material.  Even though the initiation stage is just a transient phase in the BL evolution, one still needs to understand it to get a full picture of the BL phenomenon. Also, one need not think that the BL initiation is a unique event for every accreting object --- it may, in fact, be repetitive. The most common situation in which this BL initiation is recurrent involves the dramatic increase of the mass accretion rate (due to some sort of outburst triggered by an instability) through the disk which is magnetically disrupted outside the star under normal circumstances. According to Equation (\\ref{Bdisrupt}), a sudden increase of $\\dot M$ by several orders of magnitude as typical for some accreting systems (e.g. dwarf novae, FU Ori outbursts, etc.) can rapidly compress the stellar magnetosphere to the point at which disk material starts touching the stellar surface and a BL starts to form.  There is good observational evidence for this kind of recurrent behavior. For instance, \\citet{LivioPringle} have argued that the observed lag in the rise time of the UV emission relative to the optical in a CV system transitioning to outburst can be explained if the disk is magnetically disrupted in quiescence but not in outburst. Thus, the hottest, innermost part of the disk is evacuated in quiescence, and UV radiation is not emitted immediately in the transition to outburst, since the empty region requires time to become filled.  FU Ori stars are another type of system in which the disk can extend all the way to the stellar surface, due to the high accretion rate - $\\dot{M} \\sim 10^{-4} M_\\sun \\text{yr}^{-1}$ \\citep{KHH}. In these systems, the BL can puff up to become of order the stellar radius and the distinction between the boundary layer, the disk, and the star becomes blurred \\citep{PNFuOri}.  The goal of our present work is to make a first step towards understanding the {\\it formation} of the BL. As in the case of a well-developed BL, the major issue for this initial stage lies in details of the angular momentum transport and mixing. However, because of the enormous shear present at the star-disk interface at this stage, it is highly likely that purely hydrodynamical shear instabilities would dominate the transport of mass and momentum rather than anything else.  For that reason, in this work we primarily focus on exploring the properties of various shear-driven instabilities under the conditions typical at the BL formation stage. {First, we seek to identify a particular variety of the shear instability that most efficiently initiates mixing between the two fluids, i.e. has the fastest growth rate. Second, we examine the conditions needed for its operation, such as the density contrast between fluids, initial velocity profile, etc.}  There are several physical ingredients which can potentially affect operation of shear instabilities: stratification, rotation, magnetic fields, turbulence, radiative transfer, and the supersonic nature of the flow. The latter aspect of the problem is very important and is inevitable during the BL initiation phase, when disk material rotating at the Keplerian speed comes into contact with the more slowly spinning stellar surface. The differential azimuthal velocity between the two interacting flows is then bound to be a significant fraction of the Keplerian velocity at $R_*$ and should highly exceed the sound speed both in the disk and in the outer layers of the star.   \\subsection{Shear Instabilities in Compressible Fluids}  The study of shear instabilities in compressible fluids is of fundamental physical significance and has received much attention in the past. \\citet{Landau}, \\citet{Hatanaka}, \\citet{Pai}, \\citet{MilesKH}, and \\citet{Gerwin} have all studied the problem in the vortex sheet approximation, where one half plane of compressible fluid moves at constant velocity over another. When the two fluids move at a low Mach number relative to each other, one finds that infinitesimal perturbations are governed by the classical Kelvin-Helmholtz (KH) dispersion relation for two incompressible fluids. However, \\citet{MilesKH} showed that above a critical Mach number, the vortex sheet becomes marginally stable to two-dimensional disturbances along the direction of the flow. This is a surprising result considering that one might have expected increasing the shear to lead to an increased growth rate for the instability rather than stabilization. Understanding the implications of this result for the momentum and mass transport in the astrophysical BLs is one of the goals of our study.  Later, \\citet{Blumenetal}, \\citet{Ray}, \\citet{ChoudhuryLovelace} and \\citet{Glatzel} studied the stability of a fluid with a continuous and monotonically varying velocity profile. They found that unlike the vortex sheet, continuous velocity profiles with a supersonic velocity difference across them were unstable even at high Mach number. \\citet{Glatzel} showed that the instability was similar to the Papaloizou-Pringle instability which operates in hydrodynamical disks with radial boundaries \\citep{PP}. Thus, for high Mach number flow in a compressible fluid, a finite thickness velocity profile exhibits fundamentally different behavior from the vortex sheet, since it can support unstable modes. Moreover, the growth rate of the unstable modes scales as $\\propto 1/L$, where $L$ is the thickness of the shear layer. Thus, the thinner the shear layer, the faster the instability operates! This is exactly the opposite of what one might have expected from the vortex sheet stability criterion and we will provide an explanation for this in \\S \\ref{finite_width_SL}.  Previous studies of possible instabilities inside astrophysical BLs have been primarily concerned with the sub-sonic regime when the flow can be considered as almost incompressible \\citep{KippenhahnThomas, Fujimoto1}. This regime may indeed apply in well-developed BLs with smooth shear profiles, even though at the present level of knowledge one can hardly exclude the possibility of the existence of localized regions in steady-state BLs where compressibility is still important. Our primary goal here is to extend these studies into the regime of highly compressible, supersonic flows and to explore the implications for astrophysical objects. In the course of studying supersonic shear instabilities, we will sometimes also incorporate stratification in our calculations and explore the role of rotation.  The paper is organized as follows. In \\S \\ref{goveqn} we present our governing set of equations and describe the formalism we use to study the stability problem. Then in \\S \\ref{compress_KH} we study the case of the compressible vortex sheet with no gravity and show that our formalism reproduces the dispersion relation obtained by previous authors \\citep{Landau, Hatanaka, Pai, MilesKH, Gerwin}. In \\S \\ref{compress_grav}, we introduce gravity as a small parameter and perform a perturbation expansion to study what influence this has on the stability of the vortex sheet. Finally, in \\S \\ref{finite_width_SL} we study the stability of a finite width shear layer with a linear velocity profile. We derive new dispersion relations for this case and test the growth rate of the fastest growing mode using the Godunov code Athena \\citep{Stoneetal}. ",
        "conclusions": "\\label{discon}   We have studied supersonic shear instabilities that could drive the turbulence in the BLs of stars for which the disk is undisrupted by a magnetic field. Our study is aimed mainly at identifying the instabilities that lead to the formation of the BL when the disk just touches the surface of the star. The main result of our work is the identification of two types of instabilities that could operate in the BLs of such systems and had not been previously discussed in this context.  The first is an instability of a vortex sheet at high Mach number caused by gravity. Although the vortex sheet is stable above a critical Mach number, the addition of a small amount of gravity destabilizes it. We have found that the eigenfrequencies of the dispersion relation in the limit $G \\rightarrow 0$ acquire a purely imaginary term, which is proportional to the small parameter $G$. This has the effect of making the upper branch unstable. We now consider whether the instability of the upper branch in the limit $G \\rightarrow 0$ is likely to be relevant during the initiation of the BL, when one may expect the vortex sheet approximation to be valid. Redimensionalizing Equation (\\ref{uvp1}), we obtain \\ba \\label{gfreq} \\omega_1 \\sim \\Omega_K \\eps^{1/2} i, \\ea where $\\Omega_K$ is the Keplerian velocity at the surface of the star. During the initiation of the BL, we expect disk material to be less dense than stellar material, which means that $\\eps \\lesssim 1$. It then follows from Equation (\\ref{gfreq}) that the characteristic growth rate of the instability is $\\lesssim \\Omega_K$ and is independent of the wavenumber.   If the BL is radially thin, as would be expected during the initiation phase, then the growth rate given by Equation (\\ref{gfreq}) is small compared to the shear rate $S \\sim \\Omega_K R_*/\\delta_{BL} \\gg \\Omega_K$. Considering that $S$ is the characteristic growth rate for shear instabilities, if there are other mechanisms for instability with growth rates proportional to $S$, they will quickly become dominant.  In particular, we have demonstrated that the growth rate of the sonic instability of a finite width shear layer at high Mach number is proportional to $S$. The sonic instability is similar in nature to the Papaloizou-Pringle instability in that both are global instabilities and cannot be derived from a local analysis. There are two destabilizing mechanisms for the sonic instability. The first corresponds to emission of radiation and gives instability over a broad range of wavenumbers. The second corresponds to over-reflection of trapped modes and results in sharply peaked, disconnected regions of instability in $k$-space. Because sonic instabilities operate on a much faster timescale than the gravity mechanism for a vortex sheet, this makes them an appealing candidate for the initial stages of boundary layer formation, when one might expect large shears to be present.  We mention that \\citet{AlexakisMiles} have considered the Miles instability \\citep{Mileswave} in the context of boundary layer formation. Specifically, they invoked the Miles instability to generate mixing of WD and stellar material and explain the enrichment in heavy elements observed in nova explosions. The Miles instability was proposed as a mechanism for generating waves over water at low wind speeds and operates through a resonant interaction between the wind and the water wave. Due to this interaction, a component of the pressure perturbation is created that is in phase with the velocity of the air-water interface, and much like pumping a swing, swings up the interface to large amplitudes. However, the Miles instability was initially introduced as a way of explaining the formation of waves on water at weak wind speeds, for which the air-water interface is stable to the KH instability. During boundary layer formation, however, large shears are generated, so we are not in the weak wind regime, and the Miles instability is likely to be swamped by sonic instabilities.  There are two main astrophysical implications of our findings. First, we demonstrate that the initiation of the BL is likely to take only very short amount time after the first material from the disk arrives to the stellar surface. This is because the sonic instabilities that we explored in this paper have an extremely short growth rate, which is very weakly dependent on the density contrast between the disk and stellar material.\\footnote{This statement is not true for the supersonic KH instability with gravity investigated in \\S \\ref{compress_grav}, see Equation (\\ref{uvp1}).} Thus, mixing of the two fluids starts almost immediately after they come into contact.  Second, given the efficiency with which the sonic instability operates, it is likely that it may play important role also in more developed phases of the BL evolution. As long as the BL possesses some effective \"boundaries\" (e.g. sharp changes in the velocity of density behavior) and the gas flow within it is supersonic the purely hydrodynamic sonic instabilities are going to operate in it potentially providing means for continuing mixing and angular momentum transport inside the layer. Future numerical calculations capable of following the nonlinear development of sonic instabilities should be able to address this issue."
    },
    "1112/1112.4514_arXiv.txt": {
        "abstract": "We present the discovery of the \\starname\\ planetary system, which we initially identified through the detection of five distinct periodic transit signals in the \\kepler\\ light curve of the host star \\tmid. From high-resolution spectroscopy of the star, we find a stellar effective temperature $\\teff=5455\\pm100$~K, a metallicity of \\feh$=0.01\\pm0.04$, and a surface gravity of $\\logg=4.4\\pm0.1$. We combine these estimates with an estimate of the stellar density derived from the transit light curves to deduce a stellar mass of $\\mstar=0.912\\pm0.034~\\msun$ and a stellar radius of $\\rstar=0.944^{+0.060}_{-0.095}~\\rsun$.  For three of the transit signals, we demonstrate that our results strongly disfavor the possibility that these result from astrophysical false positives.  We accomplish this by first identifying the subset of stellar blends that reproduce the precise shape of the light curve and then using the constraints on the presence of additional stars from high-angular resolution imaging, photometric colors, and the absence of a secondary component in our spectroscopic observations. We conclude that the planetary scenario is more likely than that of an astrophysical false positive by a factor of $2\\times10^5$ (\\planetb), $1\\times10^5$ (\\planetc), and $1.1\\times10^3$ (\\planetd), sufficient to validate these objects as planetary companions.  For \\planetc\\ and \\planetd, the blend scenario is independently disfavored by the achromaticity of the transit: From \\spitzer\\ data gathered at 4.5~\\micron, we infer a ratio of the planetary to stellar radii of $0.075\\pm0.015$ (\\planetc) and $0.065\\pm0.011$ (\\planetd), consistent with each of the depths measured in the \\kepler\\ optical bandpass. We determine the orbital periods and physical radii of the three confirmed planets to be $3.70$~d and $1.91^{+0.12}_{-0.21}~R_{\\Earth}$ for \\planetb, $10.85$~d and $3.07^{+0.20}_{-0.31}~R_{\\Earth}$ for \\planetc, and $77.61$~d and $2.75^{+0.17}_{-0.30}~R_{\\Earth}$ for \\planetd.  From multi-epoch radial velocities, we determine the masses of \\planetb\\ and \\planetc\\ to be $8.7\\pm2.2~M_{\\Earth}$ and $16.1\\pm3.5~M_{\\Earth}$, respectively, and we place an upper limit on the mass of \\planetd\\ of $20.1~M_{\\Earth}~(2~\\sigma)$. ",
        "introduction": "Systems with multiple exoplanets, and transiting exoplanets, each bolster confidence in the reality of the planetary interpretation of the signals and offer distinct constraints on models of planet formation.  The first extrasolar planets were found around a pulsar \\citep{wols92}, and it was the multi-planetary nature --- in particular the gravitational perturbations between the planets \\citep{rasio92, wols94} --- which solidified this outlandish claim.  Around Sun-like stars as well, the origin of radial velocity signals continued to be questioned by some, at the time multiple planets were found around ups Andromeda \\citep{butler99}. The orbital configuration of planets relative to each other has shed light on a host of physical processes, from smooth radial migration into resonant orbits \\citep{lee02} to chaotic scattering into secular eccentricity cycles \\citep{malh02,ford05}.  Now with ever-growing statistics of ever-smaller Doppler-detected planets in multiple systems \\citep{mayor11}, the formation and early history of planetary systems continues to come into sharper focus.  Concurrently, transiting exoplanets have paid burgeoning dividends, starting with the definitive proof that Doppler signals were truly due to gas-giant planets orbiting in close-in orbits \\citep{charbonneau00, henry00}. Transit lightcurves offer precise geometrical constraints on the orbit of the planet \\citep{winn10}, such that radial velocity and photometric measurements yield the density of the planet and hence point to its composition \\citep{adams08, miller11}.   Transiting configurations also enable follow-up measurements \\citep{charbonneau02, knutson07, triaud10} which inform on the mechanisms of planetary formation, evolution, and even weather.  These two research streams, multiplanets and transiting planets, came together for the first time with the discovery of Kepler-9 \\citep{holman10, torres11}. This discovery was enabled by data from the \\kepler\\ Mission \\citep{borucki10, koch10a}, which is uniquely suited for such detections as it offers near-continuous high-precision photometric monitoring of target stars.  Based on the first 4 months of \\kepler\\ data, \\citet{borucki11} announced the detection of 170 stars each with 2 or more candidate transiting planets; \\citet{steffen10} discussed in detail 5 systems each possessing multiple candidate transiting planets. A comparative analysis of the population of candidates with multiple planets and single planets was published by \\citet{latham11}, and \\citet{lissauer11a} discussed the architecture and dynamics of the ensemble of candidate multi-planet systems.  The path to confirming the planetary nature of such \\kepler\\ candidates is arduous.  At present, three stars (in addition to Kepler-9) hosting multiple transiting candidates have been presented in detail and the planetary nature of each of the candidates has been established: These systems are Kepler-10 \\citep{batalha11, fressin11a}, Kepler-11 \\citep{lissauer11b}, and Kepler-18 \\citep{cochran11}.  Transiting planets are most profitable when their masses can be determined directly from observation, either through radial velocity (RV) monitoring of the host star or by transit timing variations (TTVs), as was done for Kepler-9bc, Kepler-10b, Kepler-11bcdef, and Kepler-18bcd. When neither the RV or TTV signals is detected, statistical arguments can be employed to show that the planetary hypothesis is far more likely than alternate scenarios (namely blends of several stars containing an eclipsing component), and this was the means by which Kepler-9d, Kepler-10c, and Kepler-11g were all validated.  While such work proves the existence of a planet and determines its radius, the mass and hence composition remain unknown save for speculation from theoretical considerations.  This paper presents the discovery of a new system, \\starname, with five candidate transiting planets. We validate three of these by statistical argument; we then proceed to use RV measurements to determine the masses of two of these, and we place an upper limit on the mass of the third. We do not validate in this paper the remaining two signals (and hence remain only candidates, albeit very interesting ones, owing to their diminutive sizes), rather the validation of these two remaining signals is addressed in a separate effort \\citep{fressin12}.  The paper is structured as follows: In \\S2, we present our extraction of the \\kepler\\ light curve (\\S2.1), our modeling of these data (and RVs) to estimate the orbital and physical parameters of the planets and star (\\S2.2), as well as limits on the motion of the photocentroid during transit (\\S2.3) and a study of the long-term astrophysical variability of the star from the \\kepler\\ light curve (\\S2.4).  In \\S3 we present follow-up observations that we use to argue for the planetary interpretation, including high-resolution imaging (\\S3.1) and \\spitzer\\ photometry (\\S3.2), and the spectroscopy we use to characterize the star and determine the radial velocity signal (\\S3.3). In \\S4, we present our statistical analysis that validates the planetary nature of the three largest candidate planets in the system. In \\S5 we consider the dynamics of the system, and in \\S6 we discuss the constraints on the composition and formation history of the three planets.  \\subsection{Nomenclature} Throughout the course of the \\kepler\\ Mission, a given star is known by many different names \\citep[see][for an explanation of \\kepler\\ naming conventions]{borucki11}, and we pause here to explain the relationship of these names in the current context.  The star that is the subject of this paper is located at $\\alpha=$\\kicra, $\\delta=$\\kicdec\\ (J2000).  It was already known as \\tmid, and in the \\kepler\\ input catalog it was designated \\kicid.  After the identification of candidate transiting planets it became a \\kepler\\ Object of Interest (KOI) and was further dubbed \\koi, and it appeared as such in the list of candidates published by \\citet{borucki11}. Some authors have elected to denote KOIs using a different nomenclature, in which case \\koi\\ would be identified as KOI-70. After the confirmation of the planetary nature of three of these candidates it was given its final moniker \\starname. This paper describes that process of confirmation, but for simplicity we refer to the star as \\starname\\ throughout.  The three confirmed exoplanets were initially assigned KOI designations representing the chronological order in which the transiting signals were identified, but to avoid confusion we will refer to them henceforth by their \\starname\\ designations in which they are ordered by increasing orbital period $P$; \\planetb\\ (\\koib, $P=3.70$~d), \\planetc\\ (\\koic, $P=10.85$~d), and \\planetd\\ (\\koid, $P=77.61$~d).  We will refer to the two remaining candidates as \\koie\\ and \\koif, but note (as described below) that the period of \\koie\\ ($P=6.10$~d) is intermediate between those of \\planetb\\ and \\planetc, and the period of \\koif\\ ($P=19.58$~d) is intermediate between those of \\planetc\\ and \\planetd. ",
        "conclusions": ""
    },
    "1112/1112.3875_arXiv.txt": {
        "abstract": "Thanks to {\\tt SAURON\\/} integral-field observations we uncovered the Planetary Nebulae (PNe) populations inhabiting the central and nuclear regions of our galactic neighbours M32 and M31, respectively, and discuss the significant differences between their corresponding PNe luminosity functions in light of the properties of their parent stellar populations. In particular, we conclude that the lack of bright PNe in the nuclear regions of M31 is likely linked to the nearly Solar value for the stellar metallicity, consistent with previous suggestions that a larger metallicity would bias the Horizontal-Branch (HB) populations toward bluer colors, with fewer red HB stars capable of producing PNe and more blue HB stars that instead could contribute to the far-UV flux that is observed in metal-rich early-type galaxies and, incidentally, also in the nucleus of M31. ",
        "introduction": "Planetary Nebulae (PNe) in external galaxies are mostly regarded either as tracers of the gravitational potential (e.g., Romanowsky et al. 2003) or as indicators for the distance of their galactic hosts (e.g., Ciardullo et al. 1989), with the latter advantage owing to the nearly universal -- though not fully understood -- shape of the PNe luminosity function (PNLF, generally in the \\Oiii$\\lambda5007$ emission). Yet extra-galactic PNe can also be used as probes of their parent stellar population (see, e.g., Ciardullo 2006) and understanding in particular the origin of the PNLF is a puzzle that, once solved, promises to reveal new clues on the late stages of stellar evolution and on the formation of PNe themselves (e.g., Ciardullo et al. 2005; Buzzoni, Arnaboldi \\& Corradi 2006).  PNe originates from horizontal-branch (HB) stars that climb back the asymptotic giant branch (AGB) at the end of their helium-burning phase, when these stars leave the AGB and quickly cross the Hertzprung-Russell diagram on their way towards the cooling track of white dwarves (WD). For a population with a given age and metallicity, HB stars have nearly the same helium core mass ($\\sim 0.5\\,M_{\\odot}$) but a range of hydrogen shell mass ($\\sim 0.001 - 0.3 \\,M_{\\odot}$), with the reddest stars having also the largest H-shells and originating from the most massive main-sequence progenitors. Only HB stars with a considerable H-shell ascend toward the AGB and eventually lead to the formation of a PN, whereas the bluest HB stars with little envelope mass head straight toward the WD cooling curve by evolving first to higher luminosities and effective temperatures (the so-called AGB-manqu\\'e phase).  According to this simple picture, galaxies with on-going star formation should show brighter PNe than quiescent systems where massive stars have long disappeared (e.g., Marigo et al. 2004), but in fact the PNLF of young and old galaxies are relatively similar. In particular, all extra-galactic PNe surveys indicate a common and bright cut-off for the PNLF, which led Ciardullo et al. (2005) to suggest a binary evolution for the progenitors of the brightest PNe that would be common to different kind of galaxies. If galaxies seem to invariably host very bright PNe, their specific content of PNe - that is the number of PNe normalised by a galaxy bolometric luminosity - appears to vary with the metallicity of the stellar population. More specifically, Buzzoni, Arnaboldi \\& Corradi (2006) found that more metal rich galaxies show comparably fewer PNe, which also corresponds to larger far-UV fluxes. Interestingly, this may indicate that at a given mean stellar age, a larger metallicity would bias the HB population towards fewer stars with massive H-shell capable to lead to the formation of PNe, with a larger fraction of blue HB stars that contribute instead to the overall far-UV flux of their host galaxy (i.e. the so-called UV-upturn, Burnstein et al. 1988) as they follow their AGB-manqu\\'e evolution.  Within this context, we note that whereas our knowledge of both the shape and normalisation of the PNLF comes chiefly from the peripheral PNe populations of galaxies, both measurements for the stellar metallity and the UV spectral shape of galaxies pertain to their optical regions. This is because narrow-band imaging or slitless spectroscopy - the most common techniques employed to find extragalactic PNe - find it hard to detect PNe against a strong stellar background, whereas measuring the strength of stellar absorption lines or imaging the far-UV flux of galactic halos is prohibitively expensive in terms of telescope time. Such a dramatic spatial inconsistency needs to be resolved if we ought to understand the link between PNe and the properties of their parent stellar populations, in particular if we consider that such a connection may already not be entirely within our grasp, as Hubble Space Telescope (HST) observations for the UV color-magnitude diagram of M32 suggests (Brown et al. 2008). ",
        "conclusions": ""
    },
    "1112/1112.4387_arXiv.txt": {
        "abstract": "We explore a scenario in the Next-to-Minimal-Supersymmetric-Standard-Model (NMSSM) with both a light ${\\cal O}(10)$ GeV neutralino and a CP-odd Higgs boson with significant coupling to down-type fermions, evading all current B physics, LEP and WMAP bounds. Motivated by a possible slight lepton universality breaking hinted in $\\Upsilon$ decays, we consider the effect of the mixing of $\\eta_b$ resonances with the pseudoscalar Higgs on the spin-dependent scattering neutralino cross section off nucleons. We conclude that this mechanism could be relevant provided that non-perturbative effects enhance the effective $\\eta_b$-nucleon coupling, taking over velocity/$q^2$ suppression factors, perhaps giving a new insight into the current controversial situation concerning direct search experiments of dark matter. ",
        "introduction": "Evidence has been accumulated both from astrophysics and cosmology that about 1/4 of the energy budget of the present universe consists of the so-called (cold) dark matter (DM), namely, a component which is non-relativistic and neither feels the electromagnetic nor the the strong interaction. It is fair to say that the most popular DM candidate for a WIMP (Weakly Interacting Massive Particle) is the lightest supersymmetric particle (LSP) in supersymmetric models with $R$-parity conservation. Leaving aside the axion and the axino, the superpartners with the right properties for playing the role of a WIMP in the universe are the gravitino and the lightest neutralino ($\\chi$) - by far, the most discussed case in the literature.  Although the LHC is running smoothly and collecting large amounts of data useful to look for physics beyond the Standard Model (SM), other complementary facilities are certainly needed, especially concerning DM detection. In fact, DAMA/LIBRA, CoGeNT, and more recently CRESST experiments have reported the observation of events in excess of the expected background, hinting at the existence of a light WIMP~\\cite{Bernabei:2010mq,Aalseth:2011wp,Angloher:2011uu}. However, exclusion limits set by other direct searches, such as Xenon10~\\cite{arXiv:1104.3088} and Xenon100~\\cite{arXiv:1104.2549}, are in tension with the above claims.  In the NMSSM, a light neutralino (as a DM candidate) can efficiently annihilate through the resonant s-channel via a light pseudoscalar Higgs mediator satisfying the requirements from the relic density~\\cite{Gunion:2005rw}. However, following a scan of the NMSSM parameter space, the authors of \\cite{Das:2010ww} obtained upper limits on the spin-independent (SI) $\\chi$-nucleon cross section which are substantially below the requirements of DAMA and CoGeNT. The spin-dependent (SD) cross section (via $Z$-exchange) was also found several orders of magnitude below current experimental bounds. On the other hand, the authors of \\cite{Gunion:2010dy} were able to achieve a somewhat larger SI cross section in a similar scenario. Admittedly, such cross sections can be further enhanced by increasing the $s$-quark content of the nucleon, but the agreement with the low range of DAMA results turns out to be only marginally acceptable.  In this work we revisit SD $\\chi$-nucleon scattering via pseudocalar-exchange in the NMSSM, usually neglected in most analyses \\cite{Das:2010ww}, which however might be enhanced due to a non-perturbative mechanism as later argued.     \\begin{table*}[hbt] \\caption{Phase-space corrected leptonic branching fractions, $\\hat{\\cal B}\\left(\\Upsilon(nS) \\to \\ell \\ell\\right)$ (in \\%), and error bars (summed in quadrature) of $\\Upsilon(1S)$, $\\Upsilon(2S)$, and $\\Upsilon(3S)$ resonances \\cite{pdg}. Error bars of the ratios $\\hat{R}_{\\tau/\\ell}(nS)$ are likely overestimated because of expected correlations between the numerator and denominator experimental uncertainties.}   \\label{FACTORES}   \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|} \\hline & $\\hat{\\cal B}\\left(e^+e^-\\right)$ & $\\hat{\\cal B}\\left(\\mu^+\\mu^-\\right)$ & $\\hat{\\cal B}\\left(\\tau^+\\tau^-\\right)$ & $\\hat{R}_{\\tau/e}(nS)$ & $\\hat{R}_{\\tau/\\mu}(nS)$\\\\ \\hline $\\Upsilon(1S)$ & $2.48 \\pm 0.07$ & $2.48 \\pm 0.05$ & $2.62 \\pm 0.10$ & ${\\bf 0.057 \\pm 0.050}$ & ${\\bf 0.057 \\pm 0.046}$ \\\\ \\hline $\\Upsilon(2S)$ & $1.91 \\pm 0.16$ &  $1.93 \\pm 0.17$ & $2.01 \\pm 0.21$ & ${\\bf 0.052 \\pm 0.141}$ & ${\\bf 0.041 \\pm 0.141}$ \\\\ \\hline $\\Upsilon(3S)$ & $2.18 \\pm 0.21$ & $2.18 \\pm 0.21$ & $2.30 \\pm 0.30$ & ${\\bf 0.056 \\pm 0.171}$ & ${\\bf 0.056 \\pm 0.171}$ \\\\  \\hline \\hline $\\psi(2S)$ & $0.773 \\pm 0.017$ & $0.77 \\pm 0.08$ & $0.772 \\pm 0.100$ & ${\\bf -0.001 \\pm 0.100}$ & ${\\bf 0.002 \\pm 0.100}$ \\\\  \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "In this paper, we have considered a particular scenario within the NMSSM with both a light neutralino and a light CP-odd Higgs boson, the latter sizably mixing with pseudoscalar $\\eta_b$ resonances. Implicit in our work is the idea that non-perturbative effects (e.g. instanton-induced interaction) may lead to a non-negligible pseudoscalar contribution to the $\\chi$-nucleus scattering, thereby introducing a momentum-dependent form factor in the cross section which might be helpful (see e.g. Ref.~\\cite{Feldstein:2009tr}) to interpret the results of direct DM search experiments, with variable sensitivity along the nuclear recoil energy range.  To conclude we stress that an accurate experimental test of lepton universality in $\\Upsilon$ decays, the discovery of the $\\eta_b(2S)$ resonance together with the measurements of ${\\cal B}[\\eta_b(nS) \\to p\\bar{p}]$ at a (Super) B factory \\cite{Bona:2007qt} could be relevant for a better understanding of DM searches and related astrophysical questions.  \\subsection*"
    },
    "1112/1112.0498_arXiv.txt": {
        "abstract": "{We describe the growth of gas giant planets in the core accretion scenario.} { The core growth is not modeled as a gradual accretion of planetesimals but as episodic impacts of large mass ratios, i.e.\\ we study impacts of 0.02 - 1 \\mearth onto cores of 1-15 \\mearth. Such impacts could deliver the majority of solid matter in the giant impact regime. We focus on the thermal response of the envelope to the energy delivery. Previous studies have shown that sudden shut off of core accretion can dramatically speed up gas accretion. We therefore expect that giant impacts followed by periods of very low core accretion will result in a net increase in gas accretion rate. This study aims at modelling such a sequence of events and to understand the reaction of the envelope to giant impacts in more detail. } {To model this scenario, we spread the impact energy deposition over a time that is long compared to the sound crossing time, but very short compared to the Kelvin-Helmholtz time. The simulations are done in spherical symmetry and assume quasi-hydrostatic equilibrium.} {Results confirm what could be inferred from previous studies: gas can be accreted faster onto the core for the same net core growth speed while at the same time rapid gas accretion can occur for smaller cores -- significantly smaller than the usual critical core mass. Furthermore our simulations show, that significant mass fractions of the envelope can be ejected by such an impact. } { Large impacts are an efficient process to remove the accretion energy by envelope ejection. In the time between impacts, very fast gas accretion can take place. This process could significantly shorten the formation time of gas giant planets. As an important side-effect, the episodic ejection of the envelope will reset the envelope composition to nebula conditions. } ",
        "introduction": "\\label{sec:introduction} We study the formation of gas giant planets in the core accretion scenario \\citep{mizuno1980,1996Icar..124...62P,1986Icar...67..391B}. In this scenario, a planetary embryo grows by accreting from a swarm of planetesimals. At some point the embryo becomes massive enough to gravitationally attract a gaseous envelope. The growth process, both in terms of solid and gas accretion, is controlled by the planetesimal accretion, which is typically modeled as a gradual accretion of small planetesimals. However, in the giant collision phase under certain conditions, the accretion process could be dominated by relatively large impacts  \\citep{1969Icar...10..109S}. This is confirmed by Monte-Carlo planet formation models (T. Schr\\\"oter et al., in preparation) and recent results from N-body simulations \\citep{Raymond:2005p21625,Nimmo:2006p21624}. In such cases, while the collisions are less frequent, each one increases the mass of the protoplanet by a significant amount (typically of order 10\\%).  This possibility led us to investigate the importance of the nature of the solid accretion process in the overall growth of giant planets. In particular, we want to investigate if episodic but large impacts result in changes  in the mass and structure of the envelope when compared to gradual core growth?  A body of work investigates the importance of impacts and/or core luminosity on the evolution of the envelope of giant planets. One such study concerned itself with the possibility of stripping the envelope of Uranus by an impact induced shockwave \\citep{1990Icar...84..528K}. However, the authors did not follow the long-term evolution of the post impact planet and did not consider the possibility of subsequent re-accretion of gas. \\citet{2006ApJ...650.1150I} study the collision of two giant planets to explain the low envelope mass of HD~149026b.  Another study \\citep{2007A&A...466..717A} tried to assess  the effect of a large impact on the long-term luminosity evolution of a giant planet with an eye on its potential detection. Other studies investigated the influence of the thermal energy content of the solid core as well as the energy provided by its contraction on the overall evolution of the luminosity \\citep{2008A&A...482..315B}. Further studies  \\citep{2007MNRAS.381..640P} concern the dynamic response of proto-planetary envelopes to a perturbation within an ideal gas approach. Recently, \\citet{2010ApJ...720.1161L} have studied the merger of planetary embryos focussing on the re-distribution of heavy elements following the merger.  Most closely related to the problem at hand is the sudden core luminosity shut-off scenario: the evolution of the planet when the core luminosity is suddenly shut off. This has been studied in detail by \\citet{2000ApJ...537.1013I} and \\citet{Hubickyj2005415}. It is thus expected, that the periods of low core accretion in-between impacts will lead to massive gas accretion.  The effect of sporadic, relatively massive impacts during the growth phase of the core on the gas accretion has, however, not been studied in detail. Here we attempt to determine the thermal response of a gaseous envelope upon a sudden energy input delivered by a large impact to the core, and how such episodic events could modify the build-up of the envelope when compared to the nominal case of gradual accretion. ",
        "conclusions": "% \\label{sec:conclusion}  We have analyzed the effect of relatively massive impacts onto the cores of giant planets in the growth phase. Previous studies \\citep{2000ApJ...537.1013I,Hubickyj2005415} have shown how the envelope reacts to a shut-off of core luminosity with rapid envelope accretion. Therefore, intermittent giant impacts with long periods of very low core accretion in-between are expected speed up gas accretion. The aim of this study was to study the reaction of the envelope to the giant impact and the subsequent period of low core luminosity in detail. Wether or not a net acceleration of gas accretion takes place and how strong this effect is, depends on the timescale of the reaction to the impact in relation to the time in-between impacts.    There are two major effects: 1) Due to a very large core luminosity after the impact, a large fraction of the envelope is ejected. The remaining tiny envelope allows a very fast energy transport and in consequence the huge impact energy is used up or radiated away very fast. The huge luminosity in combination with a small envelope leads to a small Kelvin-Helmholtz time. Therefore the envelope 'forgets' that the impact has taken place in very short time. Afterwards, very little solid accretion is necessary for a given net core growth speed and the subsequent evolution up to the next impact can be understood by the well-studied sudden shut-off of core luminosity. However, due to the episodic impacts, further core growth is still possible. 2) This effect is enhanced by a second: The very large luminosity during the impact reconfigures the envelope structure so that it is in hydrostatic equilibrium for very high energy fluxes. Shutting off core accretion after this reconfiguration triggers higher gas accretion rates than the shut-off without prior impact. This can also be understood in terms of the negative gravothermal specific heat of self-gravitating non-degenerate gases. In this state the impacts actually lower the central temperature which explains why subsequent gas accretion can be faster after the impact. This second effect helps reduce the time it takes to accrete once more the envelope gas that has been ejected by the impact. Once the planet has again the same envelope mass as before the impact but a much lower core luminosity, the evolution follows the shutdown scenario.   Together, the alternation between very high and low energy input allows more gas to be accreted even though the impact initially ejects some (or all) of the envelope gas. The subsequent high gas accretion rate quickly reforms the original envelope and continues to accrete faster than the gradually growing case. In this way, even a rapid succession of impacts can lead to a faster envelope growth.  A further interesting consequence of the envelope loss caused by the impact relates to the dust opacity: every time the envelope is ejected, new fresh gas is accreted from the nebula, therefore resetting any former modifications to the opacity caused by e.g. dust growth and settling.  In summary, we find that episodic large impacts significantly speed up gas accretion as was expected from shut-off calculations. The new result of this study is the fact that 1) almost the entire envelope of the planet is ejected as a consequence of the impact energy and 2) it takes only a very short time to accrete the lost gas after the ejection. In fact, this is so fast that the ejection has practically no effect on the long-term evolution - envelope masses are equivalent with the shutdown-case using the increased post-impact core mass.    We can therefore conclude the following for the formation of giant planets: If planetesimals are accreted in a regime where mass ratios are large, i.e. most mass is delivered by massive impacts, this will accelerate envelope build-up. Furthermore, the ratio of envelope to core mass will be significantly enhanced and smaller cores can begin rapid gas accretion while the core is growing by large impacts. The speedup with consecutive impacts, can be understood in principle from these calculations. We will study a full evolution calculation based on episodic core growth in a future publication."
    },
    "1112/1112.1065.txt": {
        "abstract": "{Interacting galaxies are well-known for their high star formation rates and rich star cluster populations, but it is also recognized that the rapidly changing tidal field can efficiently destroy clusters. We use numerical simulations of merging disc galaxies to investigate which mechanism dominates. The simulations include a model for the formation and evolution of the entire star cluster population, accounting for the evaporation of clusters due to two-body relaxation and tidal shocks. We find that the dynamical heating of stellar clusters by tidal shocks is about an order of magnitude higher in interacting galaxies than in isolated galaxies. This is driven by the increased gas density, and is sufficient to destroy star clusters at a higher rate than new clusters are formed: the total number of stellar clusters in the merger remnant is 2--50\\% of the amount in the progenitor discs, with low-mass clusters being disrupted preferentially. By adopting observationally motivated selection criteria, we find that the observed surplus of star clusters in nearby merging galaxies with respect to isolated systems is caused by the observational bias to detect young, massive clusters, and marks a transient phase in galaxy evolution. We provide a general expression for the survival fraction of clusters, which increases with the gas depletion time-scale, reflecting that both the formation and the destruction of clusters are driven by the growth of the gas density. Due to the preferential disruption of low-mass clusters, the mass distribution of the surviving star clusters in a merger remnant develops a peak at a mass of about $10^{3}~\\msun$, which evolves to higher masses at a rate of 0.3--0.4~dex per Gyr. Briefly after a merger, the peak mass depends weakly on the galactocentric radius, but this correlation disappears as the system ages due to the destruction of clusters on eccentric orbits. We discuss the similarities between the cluster populations of the simulated merger remnants and (young) globular cluster systems. Our results suggest that the combination of cluster formation and destruction should be widespread in the dense star-forming environments at high redshifts, which could provide a natural origin to present-day globular cluster systems.} ",
        "introduction": "\\label{sec:intro} Merging and interacting galaxies host huge starbursts and large populations of young massive stellar clusters \\citep[e.g.][]{holtzman92,schweizer96,whitmore99}. A galaxy interaction triggers inflows of interstellar gas towards the galaxy centres, where it fuels a burst of star formation \\citep{hernquist89,mihos96,barnes96}. Merger-induced starbursts play a central role in the history of the universe, as galaxies are thought to have formed through hierarchical merging \\citep[e.g.][]{white78,white91,cole00}. Some fraction of this star formation takes place in compact stellar clusters \\citep{elmegreen83,whitmore99,bressert10} with masses in the range $10^2$--$10^8~\\msun$ \\citep{portegieszwart10}. The clusters that remain after a merger are often used as fossils to trace the formation history of the galaxy \\citep{larsen01}.  During the past two decades, observations with the Hubble Space Telescope have revealed that many nearby ongoing galaxy mergers host exceptionally rich star cluster populations with cluster masses exceeding $10^7~\\msun$ \\citep{schweizer82,holtzman92,miller97,schweizer98,bastian06}, which are formed due to the perturbation of the interstellar medium (ISM) \\citep{schweizer87,ashman92}. The multitude of star clusters suggests that they are useful tracers of past galaxy mergers, especially because they are easily observed up to distances of several tens of megaparsecs. The observed clusters ($>10^4~\\msun$) are distributed according to a power law with index $-2$ down to the detection limit \\citep{zhang99}. These clusters are thought to be just the `tip of the iceberg', since the initial cluster mass function (ICMF) appears to continue beyond the detection limit and down to {a certain} physical lower mass limit \\citep[see e.g.][]{portegieszwart10}.  However, high gas densities and tidal shocks, both of which are prevalent in coalescing galaxies, are known to have a disruptive effect on star clusters \\citep{spitzer58,weinberg94b,gieles06}. The destruction rate of star clusters decreases with increasing cluster mass and density \\citep{spitzer87,lamers05}.\\footnote{Unless the environment in which they reside is so disruptive that it can efficiently destroy a cluster regardless of its mass. In that case, the recently argued scenario in which cluster disruption is mass-independent \\citep{whitmore07} can arise \\citep{elmegreen10b,kruijssen11}. This would then not be universal, but depends on the environmental conditions.} This indicates that the effects of star cluster disruption could be masked by observational selection effects and go unnoticed in observations, i.e. the brightest and therefore most massive clusters are easiest to detect but also least affected by disruption.  If the ICMF is universal, i.e. all stellar clusters are formed according to a power law with index $-2$ throughout space and time \\citep[e.g.][]{kruijssen11d}, then the important role of cluster disruption is supported by the old (`globular') star cluster systems that are observed in nearby spiral and giant elliptical galaxies, which are strongly lacking low-mass clusters with respect to the young populations in presently merging galaxies \\citep{vesperini01,fall01,elmegreen10}. For a power law ICMF, the size-of-sample effect would also require that the most massive clusters are formed in the largest bursts of star formation, implying that globular clusters originate from starburst environments. The question thus arises whether or not the disruption of star clusters dominates over their formation in starburst galaxies. This is not easily determined on analytical grounds.  Globular cluster systems are present over most of the galaxy mass range \\citep[e.g.][]{peng08}, and as such it is evident that they were not only formed in interactions between massive spiral galaxies; the presence of globular clusters in dwarf galaxies suggests that these also endured starbursts during their early evolution. While the globular clusters of dwarf galaxies are generally metal-poor, the colour distribution of globular clusters is often bimodal in massive spiral galaxies and giant ellipticals \\citep{searle78,forbes97,kundu01,peng06}. This colour bimodality may translate into a metallicity bimodality, although it has recently been suggested that it is a relic of a non-linear relation between colour and metallicity \\citep{yoon06,chies11,yoon11}. {Regardless of whether the metallicity distribution is bimodal}, a popular explanation for the broad range in metallicities \\citep{muratov10} is that the metal-poor clusters preferentially originate from accreted dwarf galaxies \\citep{prieto08}, while the metal-rich population was mainly formed in-situ, either by disc instabilities \\citep{shapiro10} or in galaxy mergers \\citep{ashman92}.  The formation of globular clusters has been investigated in several theoretical and numerical studies \\citep{harris94,elmegreen97,bekki02,li04,bournaud08}. As expected from {power law statistics}, these studies all point to dense, gas-rich environments, which are typically correlated with high star formation rate densities. However, the present-day population of globular clusters is not recovered in these studies, because they only concern cluster formation and contain either no description for the further evolution of clusters or a very simplified one. {Separate studies, both analytical and numerical, have shown that} the evolution of (globular) clusters is related to the galactic environment \\citep{spitzer87,baumgardt03,lamers05a,gieles06,elmegreen10b,kruijssen11}. To obtain a more complete, quantitative understanding of the origin of present-day globular clusters, it is necessary to consider their formation and further evolution simultaneously.  At present, the most commonly used method to model the evolution of star clusters is through $N$-body simulations \\citep{vesperini97b,portegieszwart98,baumgardt03,gieles08,praagman10,renaud11}. However, this method is computationally too expensive to follow the formation and evolution of the entire cluster population. \\citet{kruijssen11} therefore introduced a method in which numerical simulations of galaxies are supplemented with a semi-analytic model for the formation and evolution of star clusters, of which the results are consistent with (observed and simulated) formation and destruction rates from the literature. {This model enables us to track the formation and evolution of the entire star cluster population throughout the assembly histories of galaxies.}  As a first effort to understand the (im)balance between the formation and destruction of star clusters in starburst environments, we use the method from \\citet{kruijssen11} to model the star cluster populations of galaxy mergers. This allows us to quantify the net effect of a galaxy merger on its cluster population. With this setup, we aim to investigate: \\begin{itemize} \\item[(1)] {the relative importance of} cluster formation and destruction in interacting galaxies; \\item[(2)] whether galaxy mergers can produce the progenitors of present-day metal-rich globular clusters. \\end{itemize} In Sects.~\\ref{sec:clform} and~\\ref{sec:clevo} we summarise our model, while the initial conditions of the simulations are presented in Sect.~\\ref{sec:init}. The evolution of the star cluster population in galaxy mergers is assessed in Sect.~\\ref{sec:evo}, where we also address their sensitivity to model parameters. We end this paper with a summary of our conclusions. ",
        "conclusions": "\\label{sec:concl} We have performed a numerical study of major mergers of comparable-mass disc galaxies, complemented with a sub-grid model for the ongoing formation and evolution of their star cluster populations. The simulations have been used to address the relative contributions of cluster formation and disruption over the course of a galaxy merger, and to investigate the potential formation of the metal-rich part of a globular cluster system. The main results from our model are as follows.  \\begin{itemize} \\item[(i)] During a galaxy merger, the total number of star clusters decreases. The increase of the star formation rate during merger-induced starbursts is compensated by a stronger increase of the cluster disruption due to tidal shock heating by dense gas. \\item[(ii)] Although during certain episodes the destruction rate is high enough to disrupt clusters independently of their mass, over the entire course of a merger low-mass clusters are most strongly affected by the destruction. When considering increasingly massive clusters, their number decreases by a smaller amount during a merger. If the cluster sample is limited to massive and young clusters to mimic observational selection effects, the net destruction cannot be detected and changes to a transient increase of the number of clusters during the starbursts, in agreement with observational results. \\item[(iii)] The relative decrease of the number of clusters is stronger for higher peak star formation rates, because the enhanced formation and destruction of clusters are both caused by the high gas density. This trend is weaker for higher masses and may be reversed above $M\\sim 1$--$3\\times10^5~\\msun$, where a stronger starburst may produce {\\it more} clusters than a weak starburst. In Eq.~\\ref{eq:tdepl}, we provide a generalised expression for the survival fraction of clusters as a function of the gas depletion time-scale, which reflects the intensity of the starburst. \\item[(iv)] {The peaks in the cluster age distribution and star formation history can be offset with respect to each other due to the elevated cluster disruption rate at the height of a starburst. This offset can be as large as 200~Myr \\citep{kruijssen11}, which implies that while the cluster age distribution can be used to reveal the occurrence of a starburst, it cannot necessarily be used to determine its time or duration.} \\item[(v)] The orbital kinematics of the star clusters in a merger remnant are isotropic within galactocentric radii of $\\sim 50$--60~kpc due to the destruction of clusters on highly eccentric orbits. This value is similar to the result for the accretion of globular clusters from satellite dwarf galaxies \\citep{prieto08}, which shows that it may not be possible to distinguish between in-situ and ex-situ cluster formation based on solely the orbital (an)isotropy of the cluster population. \\item[(vi)] The preferential destruction of low-mass clusters causes the power law initial cluster mass function to develop a peak at a mass of about $10^{2.5}~\\msun$ during the final coalescence of the galaxies. This is a lower limit, as the precise value depends on the relation between cluster mass and radius, with the post-merger peak mass potentially reaching up to $10^4~\\msun$ if the cluster radii are completely unrelated to their masses. The peak mass only weakly correlates with galactocentric radius due to the destruction of clusters on radially anisotropic orbits, and (for the adopted mass-radius relation) increases by about 0.3--0.4~dex per Gyr after the completion of a merger. Young to intermediate-age ($\\sim 2$~Gyr old) merger remnants should display a peak in the star cluster mass distribution at about $10^3~\\msun$ due to the destruction of low-mass clusters (see Figs.~\\ref{fig:histgc} and~\\ref{fig:gcmfsens}). \\item[(vii)] After a merger is completed, the star cluster population is similar to what a young globular cluster system would look like. Firstly, the ejection of clusters from star-forming regions into the stellar halo produces a spatial distribution that is comparable to that of globular clusters. Secondly, the peaked cluster mass distribution is intermediate to that of young massive clusters and old globular clusters. Thirdly, the high star formation rate during a merger is capable of producing clusters that are massive enough to survive for a Hubble time. \\end{itemize}  Interestingly, the high disruption rate after (globular) cluster formation could lead to a mass distribution with a peak mass of $10^3$--$10^4~\\msun$ on such a short time-scale that only little further disruption is required to obtain the current peak mass of the globular cluster mass distribution. This would imply that even the subset of clusters on the widest orbits around their host galaxies would be able to reach it before the present day. If the ICMF of globular clusters had a Schechter-type truncation at the high-mass end \\citep{kruijssen11d}, any further disruption would not yield an additional increase of the peak mass because it then saturates at about 10\\% of the truncation mass\\footnote{This percentage applies if the ICMF of globular clusters followed a power law with index $-2$ below the truncation, and the mass dependence of the disruption time-scale is $\\gamma\\sim0.7$ \\citep{gieles09,kruijssen09b}.} \\citep{gieles09}. This would thus lead to a `universal' globular cluster mass function, independent of galactocentric radius and current galactic environment. The current peak mass of globular cluster systems throughout the universe indeed happens to be $\\sim 2\\times10^5~\\msun$ \\citep[e.g.][]{jordan07}, around 10\\% lower than the estimated truncation mass of their ICMF \\citep{kruijssen09b}.  The increased cluster disruption rate in galaxy mergers is driven by the high gas densities that also cause the burst of star formation. This indicates that the mechanism of enhanced disruption is not necessarily constrained to major mergers, and can be generalised to any environment with a high gas density and a correspondingly high SFR. While major mergers may provide an efficient formation channel for globular cluster populations, they are not a prerequisite. Any extremely high-density, dynamically active, star-forming environment -- be it in a starburst dwarf galaxy, during bulge assembly, in an unstable high-redshift disc or in a major merger -- would cause the enhanced disruption of clusters at young ages. The clusters that eventually survive are characterised by a more quiescent evolution due to cluster migration and natural selection \\citep{kruijssen09b,elmegreen10,elmegreen10b,kruijssen11}. Indeed, the wide variety of galaxy types with remarkably similar globular cluster mass distributions is hard to explain if cluster disruption is governed by the present-day environment, and suggests that the bulk of the disruption occurred at the epoch of globular cluster formation, when the host galaxies were likely more similar. A generalisation to all dense environments is supported by dwarf galaxies like Fornax, which has not undergone a major merger and yet harbours a handful of globular clusters \\citep{shapley39,hodge61} {that presumably formed in a starburst during the early formation of the galaxy}. If such a generalisation to all dense environments indeed holds, it would suggest that globular cluster populations may be the inevitable outcomes of the large starbursts occurring in the early universe."
    },
    "1112/1112.4933_arXiv.txt": {
        "abstract": "We report on our second-year campaign of X-ray follow-up observations of unidentified \\F $\\gamma$-ray sources at high Galactic latitudes ($|b|>10^{\\circ}$) using the X-ray Imaging Spectrometer onboard the \\SU X-ray Observatory. In this second year of the project, seven new targets were selected from the First \\F Catalog, and studied with $20-40$\\,ks effective \\SU exposures. We detected an X-ray point source coincident with the position of the recently discovered millisecond pulsar PSR~J2302+4442 within the $95\\%$ confidence error circle of 1FGL~J2302.8+4443. The X-ray spectrum of the detected counterpart was well fit by a blackbody model with temperature of $kT \\simeq 0.3$\\,keV, consistent with an origin of the observed X-ray photons from the surface of a rotating magnetized neutron star. For four other targets which were also recently identified with a normal pulsar (1FGL~J0106.7+4853) and millisecond pulsars (1FGL~J1312.6+0048, J1902.0$-$5110, and J2043.2+1709), only upper limits in the $0.5-10$\\,keV band were obtained at the flux levels of $\\simeq 10^{-14}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$. A weak X-ray source was found in the field of 1FGL~J1739.4+8717, but its association with the variable $\\gamma$-ray emitter could not be confirmed with the available \\SU data alone. For the remaining \\F object 1FGL~J1743.8$-$7620 no X-ray source was detected within the LAT $95 \\%$ error ellipse. We briefly discuss the general properties of the observed high Galactic-latitude \\F objects by comparing their multiwavelength properties with those of known blazars and millisecond pulsars. ",
        "introduction": "Since its successful launch in 2008 June, the Large Area Telescope (LAT) onboard the {\\it Fermi} Gamma-ray Space Telescope \\citep{atw09} has enabled many important breakthroughs in the understanding of the origin of high energy $\\gamma$-ray emissions of various classes of astrophysical objects. The number of detected $\\gamma$-ray sources increased dramatically, from 271 objects listed in the 3rd EGRET Catalog \\citep[3EG;][]{har99}\\footnote{See also \\citet{cas08} for the revised catalog of EGRET $\\gamma$-ray sources.} to 1873 in the Second \\F Catalog \\citep[2FGL;][]{2FGL}. About 800 $\\gamma$-ray sources included in 2FGL have been identified as blazars \\citep{2LAC}, i.e., jetted active galactic nuclei (AGN) characterized by strong relativistic beaming. Other associations included pulsars \\citep[e.g.,][]{LATPSR}, high-mass X-ray binaries \\citep[e.g.,][]{HMXB}, radio galaxies \\citep[e.g.,][]{MAGN}, pulsar wind nebulae \\citep[e.g.,][]{pwn}, supernova remnants \\citep[e.g.,][]{snr}, globular clusters \\citep[e.g.,][]{glbcls}, starburst galaxies \\citep[e.g.,][]{sburst}, and distinct objects like the Large Magellanic Cloud \\citep{lmc}. However, no obvious counterparts at longer wavelengths have been found for as much as $38\\%$ of \\F objects so that several hundreds of GeV sources currently remain \\emph{unassociated} with any known astrophysical systems. Fortunately, an improved localization error for the \\F \\citep[typical $95\\%$ confidence radii $r_{\\rm 95} \\sim 0^{\\circ}.1-0^{\\circ}.2$, and even $0^{\\circ}.005-0^{\\circ}.01$ for the brightest sources;][]{2FGL}, when compared to that of EGRET (typical $r_{\\rm 95} \\simeq 0^{\\circ}.4-0^{\\circ}.7$), allows for much more effective follow-up studies at radio, optical, and X-ray frequencies, which can help to unravel the nature of the unidentified $\\gamma$-ray emitters.  In this context, X-ray follow-up observations of unidentified \\F objects are of particular importance, since some classes of astrophysical sources of $\\gamma$-rays such as AGN are strong X-ray emitters as well, while the others like most of $\\gamma$-ray emitting pulsars are faint X-ray sources. Note that assuming the keV--to--GeV emission continuum in the form of a broad-band power-law ($F_{\\nu} \\propto \\nu^{-\\alpha_{{\\rm x} \\gamma}}$), which could be a relatively good zero-order approximation in the case of blazar sources but not necessarily in the case of other classes of $\\gamma$-ray emitters,  the monochromatic X-ray flux energy density scales as $[\\nu F_{\\nu}]_{\\rm 1\\,keV} = \\left({\\rm 1\\,keV}/{\\rm 0.1\\,GeV}\\right)^{1-\\alpha_{{\\rm x} \\gamma}} \\times [\\nu F_{\\nu}]_{\\rm 0.1\\,GeV} \\simeq 3 \\times 10^{-3} \\times [\\nu F_{\\nu}]_{\\rm 0.1\\,GeV}$ for a relatively flat spectral index of $\\alpha_{{\\rm x} \\gamma} \\simeq 0.5$. Hence, if an X-ray counterpart of a bright \\F source is characterized by, e.g., $[\\nu F_{\\nu}]_{\\rm 0.1\\,GeV} \\simeq 10^{-11}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ and the X-ray--to--$\\gamma$-ray power-law emission continuum with the slope $\\alpha_{{\\rm x} \\gamma} \\geq 0.5$, such source can be expected to be detectable with modern X-ray instruments such as {\\it Chandra}, XMM-{\\it Newton}, {\\it Swift}, and \\SU within reasonable exposure times. In particular, a point source search to the level of $[\\nu F_{\\nu}]_{\\rm 1\\,keV} \\sim (10^{-14} - 10^{-13})$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ is easily attainable with the X-ray Imaging Spectrometer \\citep[XIS;][]{koy07} onboard \\SU \\citep{mit07} with relatively short exposures of few tens of ksec \\citep[e.g.,][]{aka11}. In the case of a positive detection, correlated flux changes at X-ray and $\\gamma$-ray frequencies provide an identification. The lack of correlated variability, or non-detection of an X-ray counterpart, provide on the other hand only circumstantial evidence regarding the nature of a studied target. Yet in many cases such evidence may be crucial, since the non-detection of an X-ray counterpart despite a long, dedicated observation of a bright \\F object may disprove a potential association with given classes of astrophysical sources. That is because, as mentioned above, only a few established high-energy emitters are that bright in $\\gamma$-rays but very faint in X-rays \\citep[e.g., Geminga pulsar; see the discussion in][]{tho04,mat07}.  Thus motivated we started a project to investigate the nature of unidentified high Galactic-latitude \\F objects through deep X-ray follow-up observations with \\SU XIS. The results of the first-year campaign conducted over the span of \\SU AO4 were presented in \\citet{mae11}. The AO4 program included four steady/weakly variable \\F sources from the initial \\F Bright Source List \\citep[0FGL;][]{0FGL} and can be summarized as follows. The X-ray counterpart for one of the brightest unassociated \\F objects, 1FGL~J1231.1$-$1410 (also detected by EGRET as 3EG~J1234$-$1318 and EGR~J1231$-$1412), was found. The X-ray spectrum of the counterpart was well fit by a blackbody model with a temperature of $kT \\simeq 0.16$\\,keV plus an additional power-law component dominating above 2\\,keV photon energies. This power-law component was confirmed in subsequent {\\it Swift} and XMM-{\\it Newton} exposures. Considering a recent identification of 1FGL~J1231.1$-$1410 with the millisecond pulsar (MSP) PSR~J1231$-$1411 \\citep{ran11}, in \\citeauthor{mae11} we concluded that the detected thermal X-ray photons originate from the surface of a rotating magnetized neutron star, while the non-thermal X-ray component is most likely produced within the pulsar magnetosphere. In the case of 1FGL~J1311.7$-$3429, two possibly associated X-ray point sources were discovered, one of which is now excluded from the smaller error ellipse of the GeV emitter as catalogued in the 2FGL \\citep{2FGL}. The identification of the remaining X-ray counterparts with the respective $\\gamma$-ray objects remain uncertain despite a robust determination of the the spectral and variability properties of the X-ray sources. In the case of 1FGL~J1333.2+5056, we found several weak X-ray sources within the \\F error circle, and speculated on the AGN nature of the target. Finally, one X-ray point source was detected at the edge of the error ellipse of 1FGL~J2017.3+0603. The physical connection was however viewed as unlikely, since the X-ray source did not coincide with the location of the MSP PSR~J2017+0603 discovered by the Nan\\c{c}ay radio telescope which constituted a more highly probable association with the \\F object \\citep{cog11}. The MSP identification was later indeed confirmed by the detection of the pulsed emission in the \\F data, with the same period as the radio pulsations.  In this paper, we report the results of our second-year campaign, conducted over the span of \\SU AO5 (2010 April to 2011 March) which included observations of seven \\F sources located at high Galactic latitudes ($|b|>10^{\\circ}$). The targets were selected from the First \\F Catalog of point sources \\citep[1FGL;][]{1FGL} as objects \\emph{unidentified} at the time of writing of the \\SU AO5 proposal. Since then however, four of the selected targets have been associated with MSPs: 1FGL~J1902.0$-$5110 with PSR~J1902$-$5105 \\citep{cam11}, 1FGL~J2043.2+1709 with PSR~J2043+1711 \\citep{gui11}, 1FGL~J2302.8+4443 with PSR~J2302+4442 \\citep{cog11}, and 1FGL~J1312.6+0048 with PSR~J1312+00 \\citep{2FGL}. And moreover, 1FGL~J0106.7+4853 have very recently been associated with a normal pulsar PSR~J0106+4855 \\citep{ple11}. The XMM-{\\it Newton} and {\\it Swift} satellites detected the weak X-ray counterpart of the MSP in 1FGL~J2302.8+4443 \\citep{cog11}. In the following, our new \\SU observations and data reduction procedure are described in section~\\ref{sec:obs}. The analysis results are given in section~\\ref{sec:results}, and discussed further in section~\\ref{sec:disc_and_conc} in the context of multiwavelength studies of unidentified \\F objects. ",
        "conclusions": "\\label{sec:disc_and_conc}  In this paper, we report on the results of X-ray follow-up observations of seven bright \\F sources at high Galactic latitudes ($|b|>10^{\\circ}$) using \\SU XIS. We discovered the X-ray counterpart of 1FGL~J2302.8+4443 coinciding with the position of the MSP PSR~J2302+4442 recently claimed to be associated with the $\\gamma$-ray emitter. We did not however, detect X-ray counterparts for the other four \\F objects similarly identified with a normal pulsar (1FGL~J0106.7+4853) and MSPs, namely for 1FGL~J1312.6+0048, 1FGL~J1902.0$-$5110 and 1FGL~J2043.2+1709. (In a few cases the X-rays sources have been detected within the 2FGL error ellipses, but none at the positions of the pulsars.) A relatively weak X-ray source was found inside the 2FGL error region of 1FGL~J1739.4+8717. Finally, no candidate for the X-ray counterpart was detected for the remaining object 1FGL~J1743.8$-$7620.  Including our previous observations of 1FGL~J1231.1$-$1410, 1FGL~J1311.7$-$3429, 1FGL~J1333.2+5056, and 1FGL~J2017.3+0603 reported in \\citet{mae11}, our sample of high Galactic-latitude \\F objects initially selected as unidentified and studied with \\SU consists now of eleven targets. For eight of these, we have detected single or multiple X-ray sources within the LAT error ellipses. Over the time period when the \\SU observations were being obtained, six targets from the $\\gamma$-ray sample were found to be associated with MSPs, one target (1FGL~J0106.7+4853) has been associated with a normal pulsar, and one source (1FGL~J1333.2+5056) has been classified as an AGN candidate, all in agreement with the gathered X-ray data. Still, four objects from the list remain unidentified. As argued below, one of these four, 1FGL~J1739.4+8717, is quite likely a high-redshift blazar.  The source 1FGL~J1739.4+8717 was characterized by an enhanced flux level within the LAT photon energy range during the the first seven months of \\F operation\\footnote{\\texttt{http://heasarc.gsfc.nasa.gov/FTP/fermi/data/lat/catalogs/source/lightcurves/2FGLJ1738.9+8716.png}}. After that time, the activity of the $\\gamma$-ray emitter decreased. The photon index in the LAT energy band is $\\Gamma_{\\gamma}$ = 2.1 $\\pm$ 0.1 \\citep[where dN/dE $\\propto$ E$^{-\\Gamma_{\\gamma}}$ is the differential photon flux; ][]{2FGL}, which is a typical value for the $\\gamma$-ray spectra of BL Lac type blazars \\citep[see][]{2LAC}. Importantly, as written in section~\\ref{sec:J1739} one relatively bright radio source, NVSS\\,J173722+871744, is located inside the 2FGL error region of 1FGL~J1739.4+8717 (see Figure\\,\\ref{fig:J1739_Ximage}). With a typical radio spectral index for blazar sources and radio-to-optical and optical-to-X-ray spectral indices that are consistent with blazar broadband spectrum, it is quite likely that 1FGL~J1739.4+8717 is indeed associated with a distant blazar currently characterized by an activity level low enough so that its X-ray emission was below the detection limit of the XIS instrument ($\\sim 10^{-15}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$) at the time of the performed \\SU observations.  In general, unidentified sources constituted a large fraction of the population of $\\gamma$-ray emitters detected by EGRET  ($\\sim 60\\%$ in 3EG), and at present about 31 $\\%$ of \\F sources in 2FGL catalog remain unassociated (specifically 273 sources at high Galactic latitudes ($| b | > 10^{\\circ}$) and 303 sources at low Galactic latitudes $| b | < 10^{\\circ}$). Those located at the lowest Galactic latitudes ($|b| < 5^{\\circ}$) are most widely expected to be associated with local systems such as molecular clouds, supernova remnants, massive stars, high-mass X-ray binaries, radio quiet pulsars, and pulsar wind nebulae \\citep[e.g.,][]{kaa96,yad97,rom99}. In particular, half a dozen of the brightest 3EG sources in the Galactic plane were identified as young pulsars \\citep{tho99}, despite the relatively poor localization of the EGRET sources and the source confusion complicated substantially the identification procedure. On the other hand, most of the unassociated 3EG sources at high Galactic latitudes ($|b| > 10^{\\circ}$) were later identified as blazars \\citep{sow03,sow04}. Pulsars were therefore expected to be found mainly among GeV emitters at low Galactic latitudes, while blazars were supposed to constitute the main population of GeV emitters at high Galactic latitudes. But the $\\gamma$-ray--bright pulsars were also expected to be found at intermediate Galactic latitudes \\citep[$5^{\\circ} < |b| < 73^{\\circ}$;][]{cra06}. The identification of a number of \\F objects located above the Galactic plane ($| b | > 10^{\\circ}$) with such systems (predominantly with MSPs) confirmed these expectations \\citep[see, e.g.,][and the discussion in \\citealt{mae11}]{ran10}.  In Figure\\,\\ref{fig:ratio} we plot the X-ray--to--$\\gamma$-ray energy flux density ratios ($F_{2-10\\,{\\rm keV}}/F_{0.1-100\\,{\\rm GeV}}$) versus radio--to--$\\gamma$-ray energy flux density ratios ($F_{1.4\\,{\\rm GHz}}/F_{0.1-100\\,{\\rm GeV}}$) for the \\F objects from our \\SU sample discussed here (blue circles) with the radio data available in the literature \\citep[see Table\\,\\ref{tab:fits}]{1FGL,con98}. These can be compared with the analogous ratios evaluated for bright \\F objects identified with blazars and MSPs (denoted in the figure by red crosses and squares, respectively). We remind the reader that the blazar class includes flat spectrum radio quasars (FSRQs) and BL Lacertae objects (BL Lacs). In addition, in the figure we plot the two targets discussed in \\citet{mae11}, namely 1FGL~J1333.2+5056 most likely associated with an AGN, and a peculiar object 1FGL~J1311.7$-$3429 (pink stars). As shown, the blazar and MSP populations are clearly separated in the constructed flux ratio plane. Also, four objects from our sample which have recently been associated with a normal pulsar (1FGL~J0106.7+4853) and MSPs (1FGL~J1312.6+0048, 1FGL~J2043.2+1709, and 1FGL~J2302.8+4443) occupy the same region in the analyzed parameter space as the previously known MSPs detected in the GeV range\\footnote{In the case of 1FGL~J1312.6+0048, the X-ray--to--$\\gamma$-ray energy flux density ratio shown in Figure\\,\\ref{fig:ratio} is evaluated assuming the association of the $\\gamma$-ray sources with the MSP PSR~J1312+00. That is, the \\SU XIS upper limit derived at the position of the pulsar is considered, and not the X-ray flux of the \\SU source detected within 2FGL error region.}. On the other hand, in the case of 1FGL~J1739.4+8717 the evaluated energy flux density ratios --- which are very similar to those characterizing 1FGL~J1333.2+5056 --- are consistent with the blazar identification proposed above, if only NVSS\\,J173722+871744 is considered as the true counterpart of the $\\gamma$-ray emitter.  In all, we conclude that the gathered \\SU XIS data together with the broad-band properties of the analyzed \\F objects are in agreement with the identification of most of them as MSPs. Yet a few cases in the analyzed sample (1FGL~J1739.4+8717, 1FGL~J1333.2+5056) constitute quite probable associations with AGN (high-redshift blazars). Finally, the nature of the remaining few targets (like 1FGL~J1311.7$-$3429) is still an open question, although, as inferred from Figure \\ref{fig:ratio}, the MSP identification seems more viable than the blazar one. In the near future, we are further continuing our X-ray studies during the \\SU AO6 cycle, focusing on both new targets (1FGL~J0103.1+4840, 1FGL~J1946.7$-$5404, and 1FGL~J2339.7$-$0531), but also performing ultra-deep exposures on particularly intriguing sources like 1FGL~J1311.7-3429.  Finally, let us comment in more detail on the case of 1FGL~J2302.8+4443. This object, as already mentioned above, has recently been associated with the millisecond pulsar PSR~J2302+4442 discovered by the Nan\\c{c}ay radio telescope \\citep{cog11}. The rotation period of the pulsar is $P \\simeq 5.19$\\,ms, the spin-down luminosity is $\\dot{E} \\simeq 3.74 \\times 10^{33}$\\,erg\\,s$^{-1}$, and the characteristic age can be evaluated as $\\tau \\simeq 6.2$\\,Gyr. \\citet{cog11} reported also on the detection of the X-ray counterpart of the pulsar with XMM-{\\it Newton}, with the unabsorbed $0.5-3$\\,keV flux of $3.1^{+0.4}_{-0.4} \\times 10^{-14}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$. This is consistent with our \\SU detection (the re-calculated flux in the same photon energy range $2.9^{+1.1}_{-1.2} \\times 10^{-14}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$). Anticipating the distance of the pulsar $d \\simeq 1.18$\\,kpc which was inferred from the NE2001 model of Galactic free electron density \\citep{cor01}, the X-ray luminosity of PSR~J2302+4442 can therefore be evaluated roughly as $L_{\\rm x} \\sim 3 \\times 10^{30}$\\,erg\\,s$^{-1}$. This, together with $\\dot{E}$ provided above, is then in good agreement with the scaling relation between the X-ray and spin-down luminosities $L_{\\rm x} \\sim 10^{-3} \\times \\dot{E}$ established for MSPs \\citep{bec97,gae06,zha07}, although the inferred distance of PSR~J2302+4442 is indicated to be smaller by a factor of four considering an unphysically high $\\gamma$-ray efficiency and, instead, assuming average efficiency of $\\gamma$-ray MSPs $\\sim$ 10 \\% \\citep{cog11}.  The spectrum of the X-ray counterpart of PSR~J2302+4442/1FGL~J2302.8+4443 was well fit by a blackbody model, and this is again in agreement with the idea that the observed X-ray photons originate from thermal emission from the surface of a rotating magnetized neutron star. Because of the limited photon statistics in the higher energy range of XIS, we could not however confirm the presence of a non-thermal component above 2\\,keV \\citep[clearly detected in the case of 1FGL~J1231.1$-$1410 by][]{mae11}. Interestingly, for the isolated pulsars as old as PSR~J2302+4442 (characteristic age of about 6\\,Gyr), the surface temperature of the neutron star is expected to be $T \\lesssim 10^{5}$\\,K \\citep{nam87,pag92}. Our detection of the X-ray counterpart indicated $kT \\simeq 0.31 \\pm 0.03$\\,keV, i.e. $T \\simeq 3.6 \\times 10^{6}$\\,K instead. Some reheating process is therefore required, possibly related to the impact of relativistic particles on polar caps \\citep[and references therein]{bec97}.   C.C. Cheung's work at NRL is sponsored by NASA DPR S-15633-Y. \\L .S. is grateful for the support from Polish MNiSW through the grant N-N203-380336. We thank the anonymous referee for a careful reading of the manuscript and useful suggestion which helped to improve the paper."
    },
    "1112/1112.4334_arXiv.txt": {
        "abstract": "We study the effect of the non-Gaussianity induced by gravitational evolution upon the statistical properties of absorption in quasar (QSO) spectra. Using the generic hierarchical ansatz and the lognormal approximation we derive the analytical expressions for the one-point PDF as well as for the joint two-point probability distribution (2PDF) of transmitted fluxes in two neighbouring QSOs. These flux PDFs are constructed in 3D as well as in projection (i.e. in 2D).  The PDFs are constructed by relating the lower-order moments, i.e. cumulants and cumulant correlators, of the fluxes to the 3D neutral hydrogen distribution which is, in turn, expressed as a function of the underlying dark matter distribution. The lower-order moments are next modelled using a generating function formalism in the context of a {\\em minimal tree-model} for the higher-order correlation hierarchy. These different approximations give nearly identical results for the range of redshifts probed, and we also find a very good agreement between our predictions and outputs of hydrodynamical simulations. The formalism developed here for the joint statistics of flux-decrements concerning two lines of sight can be extended to multiple lines of sight, which could be particularly important for the 3D reconstruction of the cosmic web from QSO spectra (e.g. in the BOSS survey). These statistics probe the underlying projected neutral hydrogen field and are thus linked to ``hot-spots'' of absorption.  The results for the PDF and the bias presented here use the same functional forms of scaling functions that have previously been employed for the modelling of other cosmological observation such as the Sunyaev-Zel'dovich effect. ",
        "introduction": "Ongoing Cosmic Microwave Background (CMB) experiments such as Planck\\footnote{http://www.rssd.esa.int/Planck}, the Atcama-Cosmology Telescope\\footnote{http://www.physics.princeton.edu/act/} (ACT) and the South Pole Telescope\\footnote{http://pole.uchicago.edu/} (SPT) will pinpoint the cosmological parameters that describe the background geometry and dynamics of the Universe in an unprecedented detail. Along with very precise constraints on the structure of the Universe on the largest scales, smaller-scales observables will be crucial in order to further constrain the cosmological concordance model or find possible deviations from it.  In particular, galaxy clustering (BOSS, 6dF, etc.) surveys and future weak lensing and clustering observations (e.g. EUCLID) could probe smaller scales and new redshif regimes.  Spectroscopic surveys such as BOSS\\footnote{http://cosmology.lbl.gov/BOSS/} (and BIG BOSS) will also trace the large scale distribution of the baryonic matter in the Universe through the study of the flux distribution of the Lyman-${\\alpha}$ absorption systems in a very large number of quasar (QSO) spectra.  The Lyman-$\\alpha$ ``forest'', the many absorption features in QSO spectra, produced by intervening neutral hydrogen in the intergalactic medium (IGM) along the line-of-sight, is well known to be an important cosmological probe (for a recent review see Meiksin 2009). In the standard cosmological paradigm, the IGM consists of mildly non-linear gas, making up the cosmic web, that traces the dark matter and is photo-heated by a Ultra Violet (UV)-background.  The Lyman-$\\alpha$ forest is thus the main probe of the IGM and it has been shown to arise naturally in hierarchical structure formation scenarios. Astrophysical effects produced by feedback from galaxies and/or AGNs do not seem to strongly affect the vast majority of the baryons in the cosmic web \\citep{Mc05,T02}, thereby this can be used as a dark matter tracer.  The relation between the Lyman-${\\alpha}$ forest flux and the underlying matter field is a nonlinear one and it is generally expected that statistics of Lyman-$\\alpha$ are biased relative to the underlying dark matter distribution. The Lyman-${\\alpha}$ forest has been studied using variety of analytical techniques such as the Zel'dovich approximation, which is valid in the quasilinear regime and often used in modelling of the nonlinear gravitational clustering \\citep{DS97, HGZ97,McGill90,MM02}.  In addition, the lognormal approximation \\citep{CJ91} is frequently used to model the statistics of Lyman-$\\alpha$ forest \\citep{Bi93,GH96,BD97,RPT01,Viel02}.  Models based on the hierarchical or scaling {\\em ansatz} \\citep{BS89,BeS92,BS99} for higher-order correlation functions have also been investigated in order to model the statistics of Lyman-$\\alpha$ forest \\citep{VSS99}.  In addition to analytical modelling, hydrodynamical simulations have also played a very important role in this field (e.g. \\cite{Cen94,GH98,Cr98,Cr99,MW01}) and support the simple analytical picture. Thus, analytical schemes, including the ones that we develop here, can be calibrated using numerical simulations and this in turn allows to explore a large parameter space efficiently. However, numerical simulations are required to resolve the Jeans scale of the photo-ionized warm IGM, and this requirement typically means small box sizes that sample larger scales modes rather poorly.  Indeed several semi-numerical prescriptions, which are not entirely based on hydrodynamical simulations, have also been developed to model Lyman-$\\alpha$ flux and recover the correlation function from observed data sets \\citep{Sls11}. \\begin{figure} \\begin{center} {\\epsfxsize=10 cm \\epsfysize=5.1 cm {\\epsfbox[29 436 581 714]{flux.eps }}} \\end{center} \\caption{The PDF $p(F_{\\alpha})$ and bias $b(F_{\\alpha})$ of the flux $F_{\\alpha}$ is plotted as a function of the flux $F_{\\alpha}$. The PDF (left-panel) and the bias (right-panel) both are constructed using a lognormal model for the underlying mass distribution. The resulting PDFs for $\\tilde\\delta_{}$ are next transformed into the flux PDFs using the fluctuating Gunn-Peterson approximation Eq.(\\ref{eq:gp}). We compare the results from the lognormal approximation (solid lines) against the one based on the Gaussian approximation (dashed lines). The values of $A(z)$ and $\\beta$ are constructed using the functional fit in Eq.(\\ref{eq:fit}) given by \\citep{Kim07} see text for more details. The bias changes signs at an intermediate flux value which depends on cosmology and hydrodynamical parameters that define the equation of state of the photo-ionized medium $A(z)$ and $\\beta$. A fiducial value for the variance $\\sigma=3$ was assumed for this plot. } \\label{fig:demo} \\end{figure}  The two most commonly used approaches in Lyman-${\\alpha}$ studies are based either on decomposing the information encoded in the transmitted flux via Voigt profile fitting or treating the flux as a continuous field.  In the first approach, the shapes and clustering properties of absorption lines fitted by Voigt profiles have been investigated in variety of studies involving the temperature of the IGM \\citep{STLE99,MMR01}, in order to constrain the reionization history \\citep{TSS02,HH03}, to measure the matter power spectrum and cosmological parameters \\citep{Croft99,VHS04,MSB06,MM99,Rol03,Copp06,Gui07,VH06}.  By using the second set of methods, statistical properties of the flux such as the mean flux level, flux PDF, flux power spectrum \\citep{VHS04,Viel08,SSM06} and flux bi-spectrum are typically employed to explore flux statistics. For example, it has been shown that the mean-flux level can be used to constrain the amplitude of intergalactic UV background \\citep{T04,Bolt05} while the flux PDF \\citep{Mc09, Bolt09} is sensitive to the thermal evolution of the IGM (see also \\cite{L06}). The flux power spectrum on the other hand can be used to constrain the cosmological parameters and nature of dark matter \\citep{Croft02}.  The flux-bispectrum \\citep{Viel09} contains useful information about the primordial, as well as gravity-induced (i.e. secondary) non-Gaussianity. The data typically used in these investigations consists mainly of two different sets of QSO spectra: the SDSS low resolution low signal-to-noise spectra and UVES/VLT or HIRES/KECK high resolution spectra. The number of SDSS spectra is about a factor $\\sim 200$ larger than that of high resolution samples though the later probes the smaller scales with greater accuracy.  More recently, it has been argued that BOSS-like QSO spectroscopic surveys could detect Baryon Acoustic Oscillations (BAO) signatures at high redshift \\citep{ME07} and a sample of QSO pairs can constrain the geometry of the high redshift universe \\citep{Mc03}. Furthermore, analysis of coincident absorption lines in QSO pairs can also allow departures from the Hubble flow and non-gravitational effects to be measured \\citep{Rauch05}.  The SDSS-III/BOSS survey aims at identifying and observing more than 160,000 QSO over $\\sim 10,000$ square degrees within a redshift range $z=(2.15-3.5)$. This survey is primarily design to studying baryonic acoustic oscillations by performing a full 3D sampling of the matter density.  Such studies will also provide an unprecedented opportunity to study the clustering statistics using projected Lyman-$\\alpha$ flux decrements of QSOs.  In this paper, we will consider the Lyman-$\\alpha$ flux decrement in two dimensions (2D) which is related the projected density of neutral hydrogen. The statistical study of projected Lyman-$\\alpha$ flux decrement can be performed using one- or two-point PDF or their lower order moments. The PDFs contain information of cumulants or their correlators to an arbitrary order and can be constructed using well-established machinery of hierarchical ansatz \\citep{MCM1,MCM2,MCM3}. Several authors have recently studied the lower-order cumulants of Lyman-$\\alpha$ forests and cross-correlated them against weak lensing convergence as well as to the CMB sky \\citep{VDSS09,VVDS11}.  The results presented here are complementary to such studies as we take into account the lower order moments to an arbitrary order not just for the one-point cumulants but also for their two-point counterparts or {\\em cumulant correlators}. The cumulant correlators are the two-point analogues of one-point cumulants and are already in use in different areas of cosmology, e.g.  in analysis of galaxy surveys \\citep{SS97}. The lowest in the two-point hierarchy is the two-point correlation function. In the context of Lyman-$\\alpha$ studies, the two-point correlation function has already been introduced in studies involving two neighbouring line of sights \\citep{Viel02,Dod02}.  Our study generalizes these results to probe non-Gaussian correlation functions involving multiple line of sight. We model the statistics of underlying neutral hydrogen distribution using lognormal distribution as well as an extension of perturbation theory approach \\citep{VaMu04} in 3D.  Predictions from these models are then tested using hydrodynamical simulations at three different redshifts ($z=2,3,4$). Next we use these results to build and test statistical description of the projected flux distribution. The results presented here can be generalized to cross-correlation studies involving external data sets and Lyman-$\\alpha$ flux distribution. The correlation functions (CCs) are equivalent to their Fourier (harmonic) space counterpart, the multispectra, recently introduced by \\citep{MuHe09}.  The plan of this paper is as follows. In \\textsection\\ref{sec:form} we introduce the notations and define the relevant quantities such as the transmitted flux and its relation to the underlying density contrast. In \\textsection\\ref{sec:sims} we give details of the simulations that were used in our study.  In \\textsection\\ref{sec:3D} details of modelling of the 3D flux are presented.  In \\textsection\\ref{sec:lower} we derive the lower order statistics for the flux in terms of that of the underlying density contrast. In \\textsection\\ref{sec:hier} we provide a very brief introduction to the hierarchical ansatz which we use to model the statistics of underlying density contrast.  In \\textsection\\ref{sec:pdf} we provide derivation of the PDF and the bias associated with the flux distribution and finally \\textsection\\ref{sec:conclu} is left for discussion of results.  We also provide a brief appendix introducing the lognormal approximation as well as the hierarchical ansatz. ",
        "conclusions": "\\label{sec:conclu} The diffuse intergalactic medium acts as a significant reservoir of baryons at low to intermediate redshift ($z<5$), which can be probed via the absorption lines in the spectrum of distant QSOs. The Lyman-$\\alpha$ absorption lines along the line of sights are due to a vast range of completely different classes of object. These objects include under-dense neutral hydrogen clouds, halos of large and over-dense systems which may have already reached virial equilibrium, as well as UV heated systems that are strongly coupled to their environments. Previous studies have mainly dealt with the problem of detailed modelling of number of these objects as a function of their clustering and internal properties. Using self-consistent scaling models, which have a long history and were initially employed in galaxy clustering statistics, several authors have computed the column density distribution of Lyman-$\\alpha$ absorption systems not only for the low column density Lyman-$\\alpha$ forest systems but also for Lyman-limit systems and the damped absorption systems.  In addition to the hierarchical modelling, lognormal approximation is also quite successful in reproducing the clustering statistics of Lyman-{$\\alpha$} absorption lines. In this study we showed the approach taken by the lognormal approximation and the hierarchical ansatz generates near identical results. While previous studies focussed only on one-point statistics we extend these results to the two-point distributions and their lower order moments.  In addition to the statistics of absorption lines, the statistics of the transmitted flux play an important complementary role in Lyman-$\\alpha$ studies. The flux can be treated as a continuous field. Various statistics which are often employed in analysing the flux include the LOS power spectrum measurements, estimation of bispectrum and more recently the entire PDF. The flux PDF contains the information regarding cumulants to an arbitrary order. Clearly the flux statistics and the column density distributions are related. One major goal in this study was to unify these two pictures in the context of the hierarchical ansatz or scaling models as well as the lognormal approximation.  The hierarchical model is primarily valid at smaller scales where the correlation functions assumes a hierarchical form. Gravitational clustering is known to develop such a form of hierarchy both in the perturbative (quasi-linear) as well as in the highly non-linear regime. The hierarchical ansatz can be used to describe the mass functions and the bias associated with collapsed objects using the scaling function $h(x)$ and $b(x)$; the variable $x$ is a scaling variable. Previous studies that have employed the hierarchical ansatz have shown that the column density distribution of absorption systems can be described using the same functional form for $h(x)$ and $b(x)$ that are often used for wide ranging studies from the galaxy clustering to thermal Sunyaev-Zel'dovich effects or X luminosity of clusters of galaxies. We showed that the same analytical framework can also be used to understand the statistics of QSO flux measurements, thereby providing a unified statistical approach to what at first sight appear to be very different observables.  We have approached the modelling of the flux PDF in two different ways. For the 3D analysis, the PDF of the neutral hydrogen density contrast $\\hat\\delta$ is modelled according to the hierarchical ansatz as well as lognormal distribution. We find very good agreement with the numerical simulations for the entire range of redshifts that we have considered. The predictions from both these models are almost identical and differ only marginally in the less interesting under-dense regions. We also employ a modified version of scale invariant hierarchical approximation which was developed recently \\citep{VaMu04}. This particular approximation provides a very accurate model for the clustering of $\\hat\\delta$. Using the fluctuating Gunn-Peterson approximation we map the $\\hat\\delta$ PDF to that of the transmitted flux $F_{\\alpha}$. We find reasonable approximation for the allowed range of parameters $A(z)$ and $\\beta$ that define the Gunn-Peterson approximation. Next we consider the projected or the 2D distribution of flux. For projected statistics, we start with linking the cumulants and cumulant correlators for the flux and the underlying neutral hydrogen density contrast $\\hat \\delta$. Then $\\hat \\delta$ is taken to be a tracer of the underlying density contrast $\\delta$, modelled statistically using the hierarchical ansatz. Next it is shown that, under certain simplifying assumptions, the PDF of $\\delta$ and PDF of a suitably defined reduced flux decrement are linked through a very simple relation. Tests against numerical simulations shows good agreement. Results were obtained for the entire bias functions that act as a generating function for the cumulant correlators. The scaling function $h(x)$ and the bias function $b(x)$ are also known to describe the number density and bias of over-dense objects respectively.  The formalism developed here can also be used to probe the higher order cross-correlation statistics involving Lyman-$\\alpha$ flux decrement and weak lensing convergence of CMB maps or those extracted from convergence maps constructed using the weak gravitational lensing of optical galaxies."
    },
    "1112/1112.1382_arXiv.txt": {
        "abstract": "{Radial velocity (RV) measurements from near-infrared spectra have become a potentially powerful tool to search for planets around cool stars and sub-stellar objects. As part of a large survey to characterize M-dwarfs using NIRSPEC at Keck~II, we obtained spectra of eight late M-dwarfs (spectral types M5.0-M8.0) during two or more observing epochs per target. These spectra were taken with intermediate spectral resolving powers ($R\\sim20,000$)  in the $J$-band.  } {We search for relative radial velocity variability in these late M-dwarfs and test the NIRSPEC capability of detecting short period brown dwarf and massive planetary companions around low-mass stars in the $J$-band ($\\approx 1.25~\\mu$m).  Additionally, we reanalyzed the data of the M8-type star vB10 (one of our targets) presented in Zapatero Osorio et al. (2009), which were obtained with the same instrumentation as our data.} {To achieve a precise RV measurement stability, the NIRSPEC spectra are self-calibrated by making use of the telluric absorption lines, which are present in the observed spectra and used as a long-term stable reference. In the modeling process a multi-parameter $\\chi^2$-optimization is employed to generate an accurate description of the observation. The telluric lines allow us to model the instrumental profile of the spectrograph and the determination of the Doppler shift of the stellar absorption lines. } {For the entire M-dwarf sample, we do not find any evidence of relative RV variations induced by a short period brown dwarf or massive planetary companion. The typical RV precision of the measurements is between 180 and 300~m~s$^{-1}$, which is sufficient to detect hot Neptunes around M-dwarfs. Also, we find that the spurious RV shift in Zapatero et al.~(2009) of the star VB10 was caused by asymmetries in the instrumental profile between different observing epochs, which were not taken into account in their analysis.} {} ",
        "introduction": "The search for extrasolar planets  has led to more than 700 confirmed discoveries\\footnote{The Extrasolar Planets Encyclopedia; http://www.exoplanet.eu; 2011-Nov-15} by using all detection techniques. Up to now, most of them have been detected by means of the radial velocity (RV) technique using high-resolution spectrographs ($R=\\lambda / \\Delta \\lambda \\ge 40,000$) at optical wavelengths. Most discoveries are giant gaseous planets (typically hot Neptunes and Jupiters) of short periods (of a few days) around stars of spectral types F, G and K.  As potential hosts to rocky planetary companions, M-dwarfs have become increasingly popular as targets for RV searches (e.g. Endl et al. 2006, Charbonneau et al. 2009, Mayor et al. 2009, Zechmeister et al.~2009). Very cool stars such as M-dwarfs are the most abundant type ($\\sim70 \\%$) of stars in the solar neighborhood and the Milky Way in general (Henry at al.~1997). The effective temperatures and masses of M-dwarfs, respectively, are in the range 3700 to 2200~K and 0.5 to 0.07 solar masses for the M0 to M9.5 spectral types. They exhibit prominent absorption features corresponding to strong neutral atoms, H$_2$0, FeH, VO, CO, and TiO. Owing to the low masses of these objects, the reflex motion of the host star due to the gravitational pull of the extrasolar planet is higher and more easily detectable than for more massive host stars. Since M-dwarfs are very cool stars in comparison with solar-type stars, short period planets would more likely be situated in the habitable zone.      M-dwarfs emit most of their energy around $1.1-1.3~\\mu {\\rm m}$, in the near-infrared (NIR), while they appear very faint at optical wavelengths. First attempts to measure RV variations among very cool M-dwarfs at NIR wavelengths were done by Mart\\'in et al.~(2006). They achieved a RV precision of around 300 m~s$^{-1}$ for the M9.5-dwarf LP944-20 by using the spectrograph NIRSPEC, which is mounted on the Keck II telescope in Hawaii (McLean et al.~1998).  Recently, several research groups have reported high-precision RV measurements taken in the NIR with CRIRES (K\\\"aufl et al.~2004), mounted at the UT1/VLT in the Paranal Observatory of ESO in Chile. Bean et al.~(2010a) conducted high-resolution spectroscopic data of over 60 M-dwarfs (spectral types M4-M9) and used a NH$_3$ gas cell spectrum as a stable reference, and report an RV precision of better than $5~{\\rm m~s^{-1}}$. Figueira et al.~(2010) took observations of the planetary candidate TW~Hya and achieved a RV precision better than $10~{\\rm m~s^{-1}}$ by adopting telluric lines as a stable reference. Blake et al.~(2010) report RV measurements of 59 M- and L-dwarfs using the Keck/NIRSPEC spectrograph, with the aim to detect low-mass companions. They made use of strong CO absorption features around 2.3~$\\mu$m in M- and L-dwarfs and achieved  RV precisions between 50 and 200~m~s$^{-1}$. Tanner et al.~(2010) report preliminary results of a late M-dwarf survey by using Keck/NIRSPEC with RV precisions between 150-300~m~s$^{-1}$.  In 2009, Pravdo \\& Shaklan~(2009) announced a massive planet around the M-dwarf vB10 discovered by means of astrometrical data. Zapatero Osorio et al.~(2009; hereafter ZO09) made use of our NIRSPEC data set and found evidence for RV variations, which supported the planet hypothesis. They achieved a RV precision of about 300~m~s$^{-1}$. However, this planet was later refuted by different groups: Bean et al. (2010b), who took high-resolution spectra ($R=\\lambda/\\Delta\\lambda\\sim100,000$) with CRIRES and who achieved a RV precision of $\\sim 10$~m~s$^{-1}$, and by Anglada-Escud\\'e et al. (2010). Additionally, Lazorenko et al.~(2011) carried out an astrometric survey using the FORS2 camera of the ESO/VLT on Cerro Paranal, Chile, but found no evidence for the existence of a massive planet orbiting vB10. As part of this work, we aimed at finding out what had caused the spurious RV variations in the data analysis of ZO09.  Here, we report relative RV measurements of 8 late M-dwarfs with NIRSPEC, and we support the capability of this instrument to detect giant planetary companions with short orbital periods. In Section 2 we describe our M-dwarf sample, our observations and data reduction. In Section 3 we outline the details of the data analysis, followed by the results and discussion (Section~4). ",
        "conclusions": "We analyzed the data sets with our relative radial velocity measurement approach and determined the relative RV measurements with respect to the selected reference epoch. For any of the eight M-dwarfs in our sample, we have not found significant evidence of relative RV variations at the level of 3$\\sigma$ (Table~\\ref{xoxo:T3}), where $\\sigma$ stands for the measurements uncertainty. The RV precisions are in the order of 180-300~m~s$^{-1}$, except for the observations in July 2008, which were taken at low SNR.  We investigated the period and mass range of companions which could be detected with such RV precisions. We determined the minimum mass of the planet by employing a Monte-Carlo analysis, thereby probing planetary orbits with different parameters and investigating how many of these orbits could be recovered for the five measurements of vB10. We considered only the case of a circular orbit and the mass of vB10, which is $m_\\star=0.078~{\\rm M}_\\odot$. Fig.~\\ref{xoxo:F1} shows the 3$\\sigma$ detection limit.  We find that for companions with only a few days period, even planets with minimum masses of $m_{\\rm{p}}\\sin i \\ge 0.3~M_{\\rm Jup}$ can be detected with a RV precision of $\\sim220$~m~s$^{-1}$.  \\begin{figure} \\includegraphics[angle=270,scale=.4]{out.dat.ps} \\caption{Monte-Carlo analysis for the five vB10 measurements. For the mass of vB10, we adopted 0.078~M$_\\odot$ from Pravdo \\& Shaklan (2009). It is shown that with a RV precision of $\\sim220$~m~s$^{-1}$ even hot Jupiters with minimum masses $m_{\\rm p} \\sin i> 0.3~$M$_{\\rm Jup}$ could be detected around late M-dwarfs with 3$\\sigma$ confidence. We note that for a larger number of measurements the number of aliasing peaks can be significantly decreased.\\label{xoxo:F1}} \\end{figure}   In Fig.~4 we show our relative RVs of vB10 and the measurements by ZO09. We note that for a proper comparison, we adopted the same reference epoch as in ZO09. The agreement between ZO09 and our measurements is within 1$\\sigma$ of the quoted uncertainties for all epochs except for the 2001 epoch (BJD = 2452076). We provide next an explanation for the discrepancy of this one measurement. Similar to our data analysis, ZO09 used the telluric lines present in the target spectra as a stable reference, but contrary to our analysis they did not account for any IP variations in their analysis, but calculated the RVs by cross correlation. In our analysis, we do not see any RV shift exceeding the RV-precision for any measurement. We get evidence that the different instrumental setting used on 2001-Jun-15 (0.576\" slit instead of the standard setting of 0.432\")  produced an asymmetric instrumental profile (Fig.~\\ref{xoxo:F2}), which led to a significant RV shift when a simple cross correlation is adopted for the RV determination. Our results clearly demonstrate the importance of modeling the IP especially when observations are carried out with different instrumental settings.  \\begin{figure} \\includegraphics[angle=270,scale=.355]{xoxo-fg2.ps} \\caption{Our RV measurements of vB10 (crosses) vs. the measurements of ZO09 (open circles). The RV precision given by our analysis is about $~220$~m~s$^{-1}$, except for the last epoch in 2008, which was hampered by bad weather. We prove that the RV shift in the work of ZO09 in the first 2001 epoch (BJD$=245~2076$) originates from unaccounted asymmetries in the IP rather than from a planetary companion. \\label{xoxo:F2A}} \\end{figure}  \\begin{figure} \\includegraphics[angle=270,scale=.36]{new-ip-nogrid6.ps} \\caption{Instrumental profiles (IPs) of NIRSPEC for two different observing epochs of vB10. On 2001-06-15, a broader slit was used ($0.576\"$; solid line) than for the reference epoch ($0.432\"$; 3007-06-25; dashed line). To visualize the asymmetries between both IPs, we calculated the ratio between both IPs and scaled the resulting function for better visibility (dotted line). ZO09 did not account for these asymmetries between both IPs, which led to a spurious RV shift of about 1~km~s$^{-1}$ for the 2001-06-15 measurement in their data analysis.\\label{xoxo:F2}} \\end{figure}   \\begin{table} \\begin{center} \\caption{Relative radial velocity measurements.} \\label{xoxo:T3} \\begin{tabular}{lcccc} \\hline\\hline Obs. Date & BJD & rel. RV \\\\ & 245~0000+ & (m~s$^{-1}$) \\\\ \\hline \\multicolumn{2}{l}{{\\bf 2MJ2331-2749}:} \\\\ 2007-Jun-24 & 4276.10874 & $194\\pm224$\\\\ 2007-Jun-25 & 4277.09424 & {\\it ref. epoch}\\\\ \\hline \\multicolumn{2}{l}{{\\bf GJ 406}:} \\\\ 2007-Apr-30 & 4220.79751 & {\\it reference epoch} \\\\ 2007-Dec-23 & 4458.14962 & $-77\\pm 238$ \\\\ \\hline \\multicolumn{2}{l}{{\\bf GJ 905}:} \\\\ 2007-Jun-25 & 4277.11249 & $-233\\pm 201$\\\\ 2007-Oct-27 & 4400.95657 & {\\it reference epoch} \\\\ \\hline \\multicolumn{2}{l}{{\\bf GJ 1156}:} \\\\ 2007-Apr-30 & 4220.81513 & $141\\pm185$\\\\ 2007-Dec-22 & 4457.16844 & {\\it reference epoch} \\\\ \\hline \\multicolumn{2}{l}{{\\bf LHS 1363}:} \\\\ 2007-Oct-26 & 4400.01234 & $-27\\pm196$\\\\ 2007-Oct-27 & 4401.01511 & {\\it reference epoch}\\\\ \\hline \\multicolumn{2}{l}{{\\bf LP 412-31}:} \\\\ 2007-Oct-26 & 4400.03654 & {\\it reference epoch} \\\\ 2007-Oct-27 & 4401.04491 & $298\\pm260$\\\\ \\hline \\multicolumn{2}{l}{{\\bf RXJ2208.2}:} \\\\ 2007-Jun-24 & 4276.07879 & $-189\\pm 307$\\\\ 2007-Jun-25 & 4277.07266 & {\\it reference epoch}\\\\ \\hline \\multicolumn{2}{l}{{\\bf vB10}:} \\\\ 2001-Jun-15 & 2076.08951 & $19\\pm230$ \\\\ 2001-Nov-02 & 2215.70560 & $69\\pm223$ \\\\ 2001-Nov-02 & 2215.74225 & $-3\\pm217$ \\\\ 2007-Jun-25 & 4277.05865 & {\\it reference epoch} \\\\ 2008-Jul-28 & 4675.75669 & $131\\pm497$ \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  We compare our results to the work of Blake et al.~(2010), who searched for companions to M- and L-dwarf wby using NIRSPEC at a spectral resolving power of $\\sim25,000$ in the $K$-band. They adopted one spectral order covering the wavelength range from 2.285 to 2.318~$\\mu$m to measure the dense and strong CO-absorption line pattern present in those dwarfs. As a stable wavelength reference, they made use of the CH$_4$ telluric absorption lines present in the observations. Similarly to us, they employed a self calibrating approach, with the difference that they adopted theoretical models for M- and L-dwarfs, which well-described the observations.  Blake et al. obtained measurements with SNRs in the range of 50 to 100 in the pseudo stellar continua, and they report RV precisions of 100-300 m/s for slowly rotating late-M and L dwarfs. The uncertainty of Blake et al. in the K-band is in agreement with our derivation of 180-300 m/s in view of our SNRs. However, our wavelength coverage is about twice that in Blake et al. According to the relative RV precision formulae, we should have obtained better velocity precision in terms of wavelength coverage, which is not the case. We conclude that both the larger number of deep lines (more than 30 lines with a line depth larger than 50\\%) in the CO-band region as compared to the J-band (only a few lines with a depth larger than 25\\%) as well as the use of theoretical template spectra instead of deconvolved stellar spectra  appear to compensate for the shorter wavelength coverage in a similar factor (c.f. Equation~6 in Butler et al. 1996).  We note that Reiners et al.~(2010) and Rodler et al.~(2011) carried out theoretical RV precision studies of M- and L-dwarfs, by adopting theoretical models of M-dwarfs (e.g. del Burgo et al.~2009). As result, they find that the highest RV precision for M-dwarfs is attained in the $Y$ band around $1~\\mu$m, rather than in the $J-$ , $H-$ or $K$-band. For L-dwarfs, however, Rodler et al.~(2011) reported that the highest RV precision is attained in the $J$-band.   We conclude that for an accurate relative RV determination with NIRSPEC, a self-calibrating approach, which accounts for changes in the instrumental setting, produces the best measurements in terms of RV precision. Although with our RV precision we would be able to detect massive hot Neptunes around late M-dwarfs, we have not found any brown dwarf or massive planetary companion in our survey. Additionally, the re-analysis of the data of the M8-dwarf vB10 presented in ZO09 now clearly confirms the non-existence of a massive planet orbiting that dwarf and agrees with the results by other research groups (e.g. Anglada-Escud\\'e et al.~2010; Bean et al.~2010b; Lazorenko et al.~2011)."
    },
    "1112/1112.3542_arXiv.txt": {
        "abstract": "{Convergent point (CP) search methods are important tools for studying the kinematic properties of open clusters and young associations whose members share the same spatial motion.} {We present a new CP search strategy based on proper motion data. We test the new algorithm on synthetic data and compare it with previous versions of the CP search method. As an illustration and validation of the new method we also present an application to the Hyades open cluster and a comparison with independent results.} {The new algorithm rests on the idea of representing the stellar proper motions by great circles over the celestial sphere and visualizing their intersections as the CP of the moving group. The new strategy combines a maximum-likelihood analysis for simultaneously determining the CP and selecting the most likely group members and a minimization procedure that returns a refined CP position and its uncertainties. The method allows one to correct for internal motions within the group and takes into account that the stars in the group lie at different distances.  } {Based on Monte Carlo simulations, we find that the new CP search method in many cases returns a more precise solution than its previous versions. The new method is able to find and eliminate more field stars in the sample and is not biased towards distant stars.  The CP solution for the Hyades open cluster is in excellent agreement with previous determinations. } {} ",
        "introduction": "Ever since their discovery, the existence of stellar groups with common space motion in the solar neighborhood has been an intriguing issue whose understanding is still far from complete. The origin and evolution of these comoving groups of stars, usually referred to simply as moving groups, is explained by different scenarios including cluster disruption, dynamical effects, and accretion events \\citep{Eggen(1996),Dehnen(2000),Fux(2001),Navarro(2004)}.  Moving groups, which are observed as overdensities in the velocity space and exhibit a low internal velocity dispersion, typically a few km/s or less \\citep{Mathieu1986}, allow study of the large-scale structure and dynamics of the Milky Way \\citep{Antoja2008}. Because of perspective effects, the proper motions of comoving stars\\footnote{In fact, the proper motions of a group of stars converge to a vertex either if their space motions are parallel or if they are expanding uniformly from a moving point. These two dynamical states are strictly equivalent as far as proper motions are concerned and radial velocity information is needed to distinguish between them \\citep{Blaauw1964}.} appear to converge to a single point in the celestial sphere referred to as the \\textit{convergent point} (hereafter CP) of the moving cluster. The CP is important not only for determining which stars are actual members of a moving cluster, but also for deriving individual kinematic distances of moving group members, provided that their radial velocities are known. This is very valuable when the trigonometric parallax from the ground is not accessible and \\rm{\\sc{Hipparcos}} parallaxes are not available \\citep[for recent applications of this strategy see][]{Mamajek2005, Bertout2006}.   The first algorithms implementing a method for calculating the CP coordinates come from Charlier and Bohlin in 1916 \\citep[see][]{Smart1968}. Each of them derived an equation independently using the position angle of stars to determine the CP position on the celestial sphere. However, the constant least-square coefficients involved in solving these equations were subject to measurement errors, thus leading to systematic errors in the derived CP coordinates. Later, \\citet{Seares1945}, \\citet{Petrie1949} and \\citet{Roman1949} proposed different ways of correcting the Charlier and Bohlin equations. Their strategies represent, in a first approximation, different ways of using the position angle of stars to determine the CP solution.  Another approach was proposed by \\citet{Brown1950}. He introduced a reference frame with the origin approximately at the center of the cluster and a fundamental plane defined by drawing a great circle through the origin in the direction of the average proper motion of the cluster. The proper motion of each star is then resolved into two components, one parallel to the reference plane and the other one perpendicular to the same plane, with the former expected to be much larger than the latter. The coordinates of the CP were then derived by applying the method of maximum likelihood.  Based on this method, \\citet[][hereafter J71]{Jones1971} presented a  twofold algorithm dedicated to simultaneously  selecting the moving group members and calculating the CP position. This method was later improved and reformulated by \\citet[][hereafter B99]{deBruijne1999} to make full use of the \\textit{Hipparcos} data and allow for internal motions within the moving group.  There are many methods of finding the CP coordinates of moving groups. Several of them use, as J71 and B99 do, the observed proper motions of stars and differ in the details of the search strategy \\citep[see for example][]{Makarov2001, Makarov2007a, Makarov2007b}. Other methods use different observed quantities, such as parallaxes and radial velocities \\citep[]{Chen1997, Chereul1999, Hoogerwerf1999,Asiain1999}. The choice of which method to use depends essentially on the observational information that is available. The new CP search method that we propose in this article uses the observed proper motions to select the stars that belong to the moving group and determine the CP position. It builds on the works of J71 and B99, so we give in Sect.~2 some details on these methods that will be useful in presenting our own work.  In the rest of this paper we present and test a new version of the CP search method. Section~3 describes our new algorithm while Sect.~4 deals with the construction of synthetic data from moving group simulations dedicated to testing and investigating the performance of our algorithm in comparison with previous ones. Section~5 presents an application of our method to the Hyades open cluster using \\rm{\\sc{Hipparcos}} data, and finally Sect.~6 summarizes the results of this work. ",
        "conclusions": "We presented a new CPSM based on the idea of representing the stellar proper motions by great circles on the celestial sphere. The new CP method combines  (i) maximum-likelihood analysis to simultaneously determine the CP and select moving group members with (ii) a direct minimization routine used to return a more refined CP position and its uncertainties. Our new method allows us to correct for internal motions within the group by applying an individual correction for each star that depends on its proper motion. This procedure takes into account that the stars in the group are not located at the same distance.  We performed extensive Monte Carlo simulations to test and compare our new CPSM with the classic CPSM regarding the selection of moving members and the accuracy of the CP solution. We investigated the effects of (a) the velocity dispersion of the cluster, (b) observational errors on proper motion, (c) cluster distance, (d) number of group members, and (e) angular distance to the CP on the CP solution. Our new CPSM returned a more precise CP solution than the classic CPSM for 95\\% of the simulations. We verified that in the absence of velocity dispersion (ideal moving group) both methods exhibit the same performance. We also found that the new CPSM finds and eliminates more field stars than the classic CPSM. This situation is even more evident in the presence of background stars that generally have smaller proper motions. The new CPSM is able to retrieve more than 80\\% of all cluster members with a contamination around 20\\% of the total number of field stars at distances of 100~pc and 200~pc. At larger distances the efficiency of rejecting field stars decreases, but it is still higher than the efficiency of the classic CPSM. The definition of the $X^{2}$ function in the new CP method considers both the amplitude and direction of the stellar proper motion vector instead of only one directional component that removes the bias towards distant stars that is observed in the classic CPSM. Additional information (e.g. parallaxes and radial velocities) is required to eliminate the remaining field stars in the sample that were not rejected by the CPSM.  The new CPSM is shown to work well when applied to the Hyades open cluster, and its results are agree well with previous determinations of the CP position. The individual parallaxes derived using our CP solution are fully consistent with the trigonometric parallaxes given in the \\rm{\\sc{Hipparcos}} catalog. Although our results agree well with data from both reductions, we find that our external precision is slightly better for HIP07. Our results are also in good agreement with the secular parallaxes found in the literature, which are more precise than the trigonometric parallaxes given in the \\rm{\\sc{Hipparcos}} catalog.  The original implementation of the new CPSM was adapted to take the full covariance matrix of the \\rm{\\sc{Hipparcos}} catalog into account. The same procedure can also be applied to the future data of the \\textit{Gaia} mission, which will also be published with the full covariance matrix. The new CPSM will be used in forthcoming papers to investigate the kinematic properties of several young associations."
    },
    "1112/1112.4154_arXiv.txt": {
        "abstract": "{We study the structure of neutron stars in $R+\\beta R^{\\mu \\nu} R_{\\mu \\nu}$ gravity model with perturbative method.  We obtain mass--radius relations for six representative equations of state (EoSs). We find that, for $|\\beta| \\sim 10^{11}$ cm$^{2}$, the results differ substantially from the results of general relativity. Some of the soft EoSs that are excluded within the framework of general relativity can be reconciled for certain values of $\\beta$ of this order with the 2 solar mass neutron star recently observed. For values of $\\beta$ greater than a few $10^{11}$ cm$^2$ we find a new solution branch allowing highly massive neutron stars. By referring some recent observational constraints on the mass--radius relation we try to constrain the value of $\\beta$ for each EoS. The associated length scale $\\sqrt{\\beta}\\sim 10^6$ cm is of the order of the the typical radius of neutron stars implying that this is the smallest value we could find by using neutron stars as a probe. We thus conclude that the true value of $\\beta$ is most likely much smaller than $10^{11}$ cm$^{2}$.} ",
        "introduction": "Einstein's theory of gravity, general relativity, enjoyed impressive observational and experimental support in the past century. Main successes of this theory were explanations it brought to phenomena such as precession of the perihelion of Mercury, bending of light and the gravitational redshift of light near massive bodies. Although all of these impressive tests were done solely in the solar system, the standard model of cosmology assumes the validity of general relativity in all scales up to the very large scale structure of the universe. However, as the 20th century was ending, data from distant supernovae Ia \\cite{per99,rie98,rie04} were interpreted as evidence of late time acceleration in the expansion rate of the universe. If we continue to assume the validity of general relativity in all scales, then the best fit to observational data requires us to introduce a non--vanishing positive cosmological constant into the theory. This is the simplest way to follow without changing the basic paradigm. Yet there are several theoretical problems related with the existence of cosmological constant (see for example \\cite{Weinberg,Peebles,Nobbenhuis,Bousso}). One of the most important among these problems is the lack of a quantum theoretical method to calculate its inferred value from cosmological data \\cite{car01}. There are several proposals in order to avoid the problems of cosmological constant with alternative routes of explanations \\cite{uzan,tsu10}. These proposals can be collected into a few groups: to explain late time accelerated expansion one can either modify general relativity by modifying the Einstein--Hilbert action, or add new gravitational degrees of freedom other than the metric to the theory of gravity, or change how the matter fields and perhaps the cosmological constant gravitates \\cite{man05,clif11}.  A modified gravity theory, which is proposed to solve late time cosmic acceleration problem, should also be able to pass several tests before it can be considered a viable theory of gravity. First of all, in the weak gravity regime, such a theory should  be compatible with the solar system tests and table--top experiments. In cosmological scales, other than producing the late time accelerated expansion, it should be free of gravitational instabilities, and obey constraints of the standard model of cosmology. Such a theory is also expected to do well in strong gravity regime: for example it should have solutions for neutron stars with mass--radius relation inside the current observational bounds. There are alternatives to and generalizations of general relativity theory which have the same predictions in the weak-field regime as the general relativity, and also provide an explanation of the evolution of the universe in the large scale. Thus, the difference between general relativity and alternatives might become prominent in the strong-gravity regime \\cite{psa08}.  One important family of modifications of Einstein--Hilbert action is the $f(R)$ theories of gravity (see reviews \\cite{Odintsov-rev,Capozziello-rev,Sotiriou-rev,deFelice-rev} and references therein). In such theories one uses a function of curvature scalar as the Lagrangian density. The $f(R)$ term must have a lower order expansion in Ricci scalar in order to include general relativity as, perhaps, a weak-field limit of it. Such models of gravity can be made to pass Solar System tests, explain the late-time accelerated expansion of the universe and also work well in the strong-field regime. In a previous work \\cite{ara11} a simple version of $f(R)$ gravity theory, in which $f(R) = R +\\alpha R^2$, is studied by two of the present authors. It is curious that predictions of such theories can be shown to be equivalent to scalar-tensor theories of gravity \\cite{mag94}, which have been analyzed throughly since the seminal paper of Brans and Dicke \\cite{bd}. However, if we modify the Einstein--Hilbert term by contractions of Ricci and Riemann tensors, $R_{\\mu \\nu }R^{\\mu \\nu }$ and $R_{\\mu \\nu \\rho \\sigma}R^{\\mu \\nu \\rho \\sigma}$, then we have an alternative gravity theory with independent predictions.  Inspiration for such an alternative gravity theory may come from string theory. Absence of ghosts in low energy string theory in flat backgrounds requires the quadratic corrections to Einstein's gravity to be of the Gauss-Bonnet (GB) form \\cite{zwi85}: \\begin{equation} S=\\int d^{4}x\\sqrt{-g}\\left[ R+\\gamma ( R^2 -4 R_{\\mu \\nu }R^{\\mu \\nu } +R_{\\mu \\nu \\rho \\sigma}R^{\\mu \\nu \\rho \\sigma} )\\right]\\ . \\label{zwei} \\end{equation} However, the Gauss--Bonnet term will not contribute to the equations of motion, because it is equivalent to a total derivative in four dimensions. This means that variations of $R_{\\mu \\nu \\rho \\sigma}R^{\\mu \\nu \\rho \\sigma}$ can be expressed in terms of variations of $R^2$ and $R_{\\mu \\nu }R^{\\mu \\nu }$. Contraction of Riemann tensors, could have relevance only in the cases related to quantum gravity \\cite{psa09,par99}. Therefore for the classical physics applications one can take the action of our string-inspired gravity as \\begin{equation} S=\\int d^{4}x\\sqrt{-g}\\left[ R+\\alpha R^2 +\\beta R_{\\mu \\nu }R^{\\mu \\nu }\\right]\\ , \\label{rqugr} \\end{equation} where $\\alpha $ and $\\beta $ are free-parameters, taken independent of each other. However, special relations between $\\alpha $ and $\\beta $ exists and such cases correspond to some special theories: for example, $\\beta=-3\\alpha$ correspond to the Weyl tensor squared modification of general relativity \\cite{s08,fs09}, and $\\beta=-4\\alpha$ has unique energy properties \\cite{bay02}.  As mentioned above, the absence of ghosts in the low energy string theory in flat backgrounds requires the quadratic corrections to Einstein's gravity to be of the form given in equation (\\ref{zwei}). This means that the alternative gravity theory defined by (\\ref{rqugr}) cannot be free of ghosts and this brings up the question of stability of the neutron star solutions discussed in this paper. However, all such modified gravity theories, bar the theory defined with (\\ref{zwei}), suffer from the same problem, the solution of which is beyond the scope of this paper. Our aim in this paper is to see the possibility of neutron star solutions in a specific modified gravity theory defined by (\\ref{rqugr}) with $\\alpha =0$, to analyze solutions with mass-radius relations obeying the observational constraints, and to see if there are problems in this gravity model other than the well known stability issue. When combined with the results of \\cite{ara11} this analysis will allow for a comparison of the effects of $R^2$ and $R_{\\mu \\nu }R^{\\mu \\nu }$ terms in (\\ref{rqugr}) on the structure of neutron stars.  We further note that the mass dimensions of quadratic corrections to Einstein--Hilbert action in (\\ref{rqugr}) is order $[L]^{-4}$. It is possible to add terms with the same mass dimensions to the above action. In fact referring to the quantum gravity arguments in \\cite{par99}, we could also add the following dimension $[L]^{-4}$ terms to the action: \\begin{equation} \\square R\\quad \\mathrm{and}\\quad \\nabla ^{\\mu }\\nabla ^{\\nu }R_{\\mu \\nu }\\ . \\end{equation} However, as it is pointed out in \\cite{emi10} these terms do not contribute to the field equations and therefore in the context of the present paper they are also redundant.  The effect of the value of $\\alpha $, while $\\beta =0$, on the mass-radius (M-R) relation of neutron stars has been studied, for a polytropic EoS in \\cite{coo10}, and for realistic EoS in \\cite{ara11}. The latter work constrained the value of $\\alpha $ to be $|\\alpha |\\lesssim 10^{10}$ cm$^{2}$. In an other study, Santos analyzes neutron stars in this model of gravity for a single EoS of ideal neutron gas \\cite{emi11} and for a specific choice of $\\beta=-2\\alpha$, $\\sqrt{\\alpha}=0.96$ km. In that work, the author shows, for this restricted choice of EoS and parameters, that stable configurations of neutron stars are possible even for arbitrarily large baryon numbers of the neutron star.  Our approach in this work is different than \\cite{emi11} in two ways: (i) Rather than choosing a fixed specific value for $\\beta$, while $\\alpha =0$, we study its effect, as a free parameter, on the M-R relation and constrain its value by referring to observations, (ii) As the interaction between nucleons can not be neglected, the EoS of ideal neutron gas can not provide realistic M-R relations. We thus employ six different realistic EoSs corresponding to a variety of assumptions for the interaction between nucleons rather than the very restrictive EoS of ideal neutron gas.  In the present work, we adopt an alternative theory of gravity, in which the Einstein-Hilbert action is modified with the term $R_{\\mu\\nu}R^{\\mu\\nu}$ only. In \\S 2, we assume a perturbative form of our alternative gravity model and obtain the field equations. The reason of perturbative approach is that the equations of motion derived from this alternative gravity theory are fourth order differential equations, and their treatment in four dimensions is problematic. To obtain the modified Tolman--Oppenheimer--Volkoff (TOV) equations from the field equations we also assumed perturbative forms of metric and hydrodynamical functions. In \\S 3, we solve the structure of neutron stars in this gravity model for six representative equations of state describing the dense matter of neutron stars. We plot the mass-radius relations of neutron stars for $\\beta$ changing in the range $\\sim \\pm 10^{11}$ cm$^2$. This way the value of the perturbative parameter $\\beta$ is constrained by the recent measurements of the mass-radius relation \\cite{oze10} and the observed 2 solar mass neutron star \\cite{dem10}. We identify that $\\beta \\sim \\pm 10^{11}$ cm$^2$ produces results that can have observational consequences. Lastly, the results of the numerical study and the significance of the scale of the perturbation parameter is discussed in the conclusions. ",
        "conclusions": "In this paper we studied the structure of neutron stars (NSs) in a generalized theory of gravity motivated by string theory. In the first section we discussed the motivations for modifying gravity. In the second section we derived the field equations of this gravity model from its action. In the third section we obtained the hydrostatic equilibrium equations in spherical symmetry from the field equations by using a perturbative approach in which general relativity (GR) stands for the zeroth order gravity model. In the fourth section we solved the hydrostatic equilibrium equations for NSs by using numerical methods. In order to solve the equations we have used realistic equations of state (EoSs) that describe the dense matter inside NSs and obtained the mass-radius (M-R) relations depending on $\\beta$, the free parameter of the generalized gravity model considered. These M-R relations are then compared with the recent observational measurements of mass and radius of NSs to constrain the value of $\\beta$.  We have shown that observationally significant changes on the M-R relation are induced for $\\beta \\sim 10^{11}$ cm$^{2}$. An order of magnitude smaller values for $\\beta $ gives results that are not significantly different from what GR predicts. An order of magnitude greater values, on the other hand, leads to results that can not be associated with known properties of NSs. For such values, the perturbative approach breaks down too. For this selection of EoSs we see that none of them are consistent with the observations for $\\beta<-5\\times 10^{11}$ cm$^2$ and $\\beta>4\\times 10^{11}$ cm$^2$. The only exception is MPA1 which accepts large positive values like  $\\beta_{11}\\sim 10$. We thus conclude from this analysis that $|\\beta |\\lesssim 5\\times 10^{11}$ cm$^{2}$ is an upper limit brought by observations of NSs \\cite{oze10,dem10}.  We note that, for the the gravity model $R+\\alpha R^2$, the constraint on $\\alpha$ obtained by \\cite{ara11} by using NSs confronting the M-R relation with the same observations, is of the order $10^{10}$ cm$^2$, two orders of magnitude smaller than the constraint obtained on $\\beta$ in the present work, although both $\\alpha$ and $\\beta$ have the same dimensions ($[L]^2$). This indicates that the $R^2$ term in eq. (\\ref{rqugr}) is more effective on the structure of neutron stars than the $R_{\\mu \\nu }R^{\\mu \\nu }$ term.  The last observation is very interesting from a theoretical point of view. As mentioned in the introduction, theory described by (\\ref{rqugr}) is equivalent to Weyl tensor squared modification of general relativity for $\\beta = -3\\alpha$ \\cite{s08,fs09}. That is, from the point of view of Weyl tensor squared modification, $\\alpha$ and $\\beta$ are in the same order. Our perturbative analysis in this paper also in some way forces us to make a similar statement: if $\\beta$ is taken in the same order as of $\\alpha$, then there would not be the problem of the break-down of the perturbative approach. What extra we learn from this analysis is that even though the $R^2$ and $R_{\\mu \\nu }R^{\\mu \\nu }$ terms seem to be on equal footing in the expansion of a Weyl tensor square term, \\begin{equation} \\label{Weyl} \\frac12 C_{\\mu \\nu \\rho \\sigma}C^{\\mu \\nu \\rho \\sigma} = -\\frac13 R^2 + R_{\\mu \\nu }R^{\\mu \\nu } + \\mathrm{topological\\, term}, \\end{equation} the contributions of them to neutron star physics are not equal. This is pointed out to us by the analysis in this paper and it is an important statement just by itself. It would be significant if this observation is confirmed and understood analytically. We plan to do this in a future publication.  We find that, some of the EoSs, which do not give M-R relations consistent with the observations within the framework of GR, can be reconciled with these observations via the free parameter $\\beta $ in the generalized gravity model considered in this work. This then brings up the question of degeneracy between the EoSs and the free parameter $\\beta $. This degeneracy does not effect the constraint $|\\beta|\\lesssim 5\\times 10^{11}$ cm$^{2}$ which is a bound for all EoSs we considered except for MPA1.  We finally comment that the constraint we obtained is actually the strongest constraint we could obtain by using NSs as the experimental apparatus. Another estimate of the nominal value $\\beta _{0}$ of the previous section is as follows: Typical radius of a NS is $R_{\\ast}\\sim 10$ km, the only length scale in the system. This corresponds to an estimate of curvature $R^{\\mu \\nu }R_{\\mu \\nu }\\sim R_{\\ast}^{-2}\\sim 10^{-12}$ cm$^{-2}$ and so the new perturbative term $\\beta R^{\\mu \\nu }R_{\\mu \\nu }$ will become of order $R_{\\ast}$ and lead to variations on the structure of neutron stars for $\\beta \\sim R_{\\ast}^{2}\\sim 10^{12}$ cm$^{2}$. As we obtain such variations in this limit, the value that we should obtain by using NSs, we infer that the actual value of $\\beta $ is likely much smaller than this. As we mentioned before, deviations from GR are not significant for NSs for values much less than this value."
    },
    "1112/1112.4648_arXiv.txt": {
        "abstract": "{} {We intend to compile a new galaxy group and cluster sample of the latest available SDSS data, adding several parameter for the purpose of studying the supercluster network, galaxy and group evolution, and their connection to the surrounding environment.} {We used a modified friends-of-friends (FoF) method with a variable linking length in the transverse and radial directions to eliminate selection effects and to find reliably as many groups as possible. Using the galaxies as a basis, we calculated the luminosity density field.} {We create a new catalogue of groups and clusters for the SDSS data release~8 sample. We find and add environmental parameters to our catalogue, together with other galaxy parameters (e.g., morphology), missing from our previous catalogues. We take into account various selection effects caused by a magnitude limited galaxy sample. Our final sample contains 576493 galaxies and 77858 groups. The group catalogue is available at \\texttt{http://www.aai.ee/$\\sim$elmo/dr8groups/} and from the Strasbourg Astronomical Data Center (CDS).} {} ",
        "introduction": "Observations of the local Universe have shown that basically all galaxies are located in groups -- it is their natural environment. Groups and clusters of galaxies form the basic building blocks of the Universe. Therefore, it is essential to extract groups of galaxies from galaxy surveys, and their study can provide new understanding of the evolution of galaxies, of the large-scale structure, and of the underlying cosmological model.  In our previous papers \\citep{Tago:08,Tago:10} we have extracted groups from the SDSS DR5 and DR7 samples, respectively. In these papers we have given an extensive review of papers dedicated to group search methods and of the published group catalogues. In this introduction we present only a short review of the studies of galaxy groups.  During the last decade, several group catalogues that use spectroscopic redshifts have been published, either based on the 2dFGRS \\citep{Eke:04, Yang:05, Tago:06}, or on earlier releases of the SDSS \\citep{Einasto:03, Merchan:05, Zandivarez:06, Berlind:06, Berlind:09, Yang:07,Koester:07}. However, similar algorithms used to compile these catalogues have yielded groups of galaxies with rather different statistical properties.  Several authors have recently compiled group catalogues up to the redshift 0.6: GAMA (Galaxy And Mass Assembly) by \\citet{Robotham:11} is a galaxy group catalogue based on the SDSS target catalogue; \\citet{Farrens:11} derived a catalogue, based on the LRGs and QSOs in the 2dF-SDSS surveys. Using photometric redshifts several group/cluster catalogues have been compiled \\citep[e.g.][]{Gal:09,Szabo:11}. \\citet{Hao:10} applied a Gaussian mixture BCG algorithm to the SDSS DR7 data and assembled a photometric group catalogue up to the redshift 0.55. Using spectroscopic redshifts, \\citet{Knobel:09} compiled the deepest group catalogue so far, reaching up to the redshift 1 in the zCOSMOS field. In addition, catalogues of rich galaxy clusters have been created by \\citet{Miller:05}, \\citet{Aguerri:07}, and \\citet{Popesso:07}. We discuss and compare some of these catalogues in a separate paper.   The papers dedicated to group and cluster search use a wide range of both sample selection methods as well as cluster search methods and parameters. The choice of the methods and parameters depends on the goal of the catalogue. For example, while \\citet{Weinmann:06} searched for compact groups in the SDSS DR2 sample, applying strict criteria in the friend-of-friend (FoF) method, then \\citet{Berlind:06} applied the FoF method to the volume-limited samples of the SDSS with the goal to measure the group multiplicity function and to constrain dark matter haloes. Hence, the parameters of the algorithm depend on the goal of the study.  Our goal is to generate an up-to-date catalogue of groups and clusters for large-scale structure studies. The catalogue is based on the SDSS data release~8 (DR8). Since the SDSS spectroscopic main sample basically did not change from DR7 to DR8, we use exactly the same group finding algorithm and parameters as described in \\citet{Tago:10}, yielding a group sample with similar properties. The photometry of the galaxies in DR8 has been reprocessed, yielding more accurate luminosities. This is important for detailed photometric galaxy modelling (Tempel et al. in prep).  Compared to our previous catalogues, we added several additional descriptors, including environmental parameters (both local and global) and morphology. For the DR7, such data have been already used in several papers: the environment and morphology have been used to study the environmental effects on galaxy evolution \\citep{Tempel:11a}; global environments have been used to extract superclusters from the cosmic network \\citep{Liivamagi:10}. The present catalogue, based on the DR8, has already been used to compare the local and global environments of galaxies (Lietzen et al. in prep), to study the structure of rich groups \\citep{Einasto:12}, and to study the photometric structure of galaxies (Tempel et al. in prep).  The paper is organised as follows. The data used are described in Sect.~\\ref{sect:data}. Section~\\ref{sect:cat} gives a brief overview of the group finding algorithm used, together with a short comparison with our DR7 catalogue. In Sect.~\\ref{sect:dens} we describe our method to calculate the luminosity density field. In Sect.~\\ref{sect:param} we describe the additional galaxy and group parameters. All the parameters in the resulting catalogue are described in Appendix~\\ref{app:cat}. The catalogue can be downloaded from \\texttt{http://www.aai.ee/$\\sim$elmo/dr8groups/} or from the Strasbourg Astronomical Data Center (CDS)\\footnote{Galaxy and group/cluster tables will be available at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/}.  Throughout this paper we assume the following cosmology: the Hubble constant $H_0 = 100\\,h\\,\\mathrm{km\\,s^{-1}Mpc^{-1}}$, the matter density $\\Omega_\\mathrm{m}=0.27$ and the dark energy density $\\Omega_\\Lambda=0.73$. ",
        "conclusions": "We have constructed a group catalogue for the SDSS DR8 sample, following the same procedure that we used for the DR7 sample \\citep{Tago:10}. Since the spectroscopic data is the same in the DR7 and DR8, the new catalogue is similar to the previous DR7 group catalogue. The improvements are in the area of initial galaxy selections and from the SDSS side, the photometry of galaxies.  In addition to the properties, presented in our previous catalogue, we added some new qualitative information. Most importantly, we tracked the missing galaxies in the SDSS (due to fibre collisions) and calculated the environmental density parameters for each galaxy and group. We also added our galaxy morphology as derived in \\citet{Tempel:11a}.  In the study of galaxy groups, the most important problem at present is the dynamical status of groups of galaxies. Recently \\citet{Plionis:06} and \\citet{Tovmassian:09} studied shapes and virial properties of groups and found a strong dependence on richness and concluded that groups are not in dynamical equilibrium but rather are at various stages of their virialisation process.  The dynamical status of groups is also characterised by the existence of subgroups  and substructures seen in many studies \\citep{Burgett:04, Coziol:09, Einasto:10}. These studies also support the non-virialised nature of groups of galaxies. With both observational and simulated data \\citet{Niemi:07} showed that about 20\\% of nearby groups are not bound and are groups merely in a visual sense.  The controversial results  obtained by various authors may be an indication that our present knowledge of groups of galaxies is poor. An optimistic viewpoint is that the study of a broad and inhomogeneous class of galaxy systems -- groups of galaxies -- will help step by step to solve important problems of galaxy formation and evolution, and of the evolution of the large scale structure."
    },
    "1112/1112.4012_arXiv.txt": {
        "abstract": "We have studied the Residual Blaze Functions (RBF) resulting from division of individual echelle orders by extracted flat-field in spectra obtained by slit-fed OES spectrograph of 2m telescope of Ond\\v{r}ejov observatory, Czech Republic. We have eliminated the dependence on target and observation conditions by semiautomatic fitting of global response function, thus getting the instrument-only dependent part, which may be easily incorporated into data reduction pipeline. The improvement of reliability of estimation of continuum on spectra of targets with wide and shallow lines is noticeable and the merging of all orders into the one long spectrum gives much more reliable results. ",
        "introduction": "The requirement of a modern astrophysical analysis, to study the wide range of the object spectrum in the greatest details, or to observe its variability in many spectral lines at the same time, lead naturally to the usage of cross-dispersed echelle spectrographs.  Although the reduction of the echelle raw frames is quite straightforward and a number of automatic reduction pipelines has been in use for many instruments, there remains still unsolved the problem of the reliable normalization of individual extracted spectral orders before merging into one long continuous spectrum.  Due to its inherent nature the echelle blaze function has to be removed from all orders with the immense precision, otherwise the strong ripple patterns appear on the merged spectrum. They   make difficult the fitting of the continuum an hence influence the reliability of the measurement of equivalent widths of wide shallow lines, the placement of wings of hydrogen lines of hot rapidly rotating stars, or the studies of diffuse interstellar bands. ",
        "conclusions": "We have shown a one possible way how to tackle the problem of precise unblazing of echelle spectra using the separation of correction function to instrument-only dependent part (IRBF) changing on the scale of one echelle order and observation-dependent part (GSF) smoothly changing over the whole observed spectrum.  Even if this is not the final solution, and still some quickly changing variations remain, the resulting corrected echelle orders may be merged in the one long spectrum much more easily avoiding strong ripples. If the estimate of GSF is known for given observation, the spectra may be  normalized to continuum almost automatically.  Our method is just a first simple approximation of the general correction procedure, it may be instrument dependent (better suited to slit-fed spectrographs) and so the future investigation of the problem is still highly desirable."
    },
    "1112/1112.5201.txt": {
        "abstract": "%\\begin{abstract} In this work we present a model of dark matter based on scalar-tensor theory of gravity. With this scalar field dark matter model we study the non-linear evolution of the large scale structures in the universe. %using a dark matter %model steaming from a scalar-tensor theory of gravitation. The equations that govern the evolution of the scale factor of the universe are derived together with %and also the appropriate Newtonian equations to follow the non-linear evolution of the structures. Results are given in terms of the power spectrum that gives quantitative information on the large-scale structure formation. The initial conditions we have used are consistent with the so called concordance $\\Lambda$CDM model. %We also show how to estimate %the scalar field dark matter parameters using a sample of observed rotation curves. % ",
        "introduction": "\\label{sec:Intro}  %The Big Bang theory The standard model of cosmology is supported by three main astronomical observations: the surveys of supernovae Ia, the cosmic microwave background radiation (CMB), and the primordial nucleosynthesis. These observations together with other % %M modern cosmological observations, % like galaxies surveys (SDSS, 2dF), galaxy rotation curves, the Bullet Cluster observation, studies of clusters of galaxies, % establish that the universe behaves as dominated by dark matter (DM) and dark energy. However, the direct evidence for the existence of  these invisible components remains lacking. Several theories that would modify our understanding of gravity have been proposed in order to explain the large scale structure formation in the universe and the galactic dynamics. The best model we have to explain the observations is the $\\Lambda$CDM model, i.e., the model of cold dark matter (CDM) --non-relativistic particles %at the epoch of last scattering of unknown origin-- with cosmological constant ($\\Lambda $), in particular, this model explains very well the universe on scales of galaxy clusters and up \\citep{Breton2004}.  The $\\Lambda$CDM model has become the theoretical paradigm leading the models of the universe to explain the large scale structure (LSS) formation  and several other observations. Where ``large'' means scales larger than $1$ Mpc --about the size of the group of galaxies that our Milky Way belongs. Together with the cosmic inflation theory, this model makes a clear prediction about the necessary initial  conditions that the universe has to have in order to have the structures we observe and that those structures build hierarchically due to a gravitational instability. % One of its main predictions is that the density profile of galaxies, clusters of galaxies, and so on, is of the form \\citep{Navarro96,Navarro97}, \\begin{displaymath} \\rho_{NFW}(r) = \\frac{\\rho_0}{(r/r_0)(1+r/r_0)^2}, \\end{displaymath} a density profile known as Navarro-Frenk-White profile (NFW). Parameters $\\rho_0$ and $r_0$ must be fitted, for example, using rotation curves of galaxies.  The $\\Lambda$CDM model and its success in explaining several  observations --this is why this model is also known as the concordance model--",
        "conclusions": "\\label{sec:Conclusions}  We have used a general, static STT, that is compatible with local observations by the appropriate definition of the background field constant, i. e., $\\langle \\phi\\rangle = (1+\\alpha)/G_N$, to study the LSS formation process. The initial condition for the several cases (different values of parameter $\\alpha$) was built in such a way that the geometry of the model universe were flat. Quantitatively, this demands that our models have $\\Omega/(1+\\alpha)=1$ and this changes the amount of dark matter and energy of the models in order to have a flat cosmology.  Using the resulting modified dynamical equations, we have studied the LSS formation process of a $\\Lambda$CDM universe. We varied the amplitude and sign of the strength of the SF (parameter $\\alpha$) in the interval $(-1,1)$ and performed several 3D-simulations with the same initial conditions. From our simulations we have found that the inclusion of the SF changes the local dynamical properties of the most massive groups, however, the overall structure is very similar, as can be seen in Fig. \\ref{fig:LSSF}.  The general gravitational effect is that the interaction between dark matter particles given by the potential $\\Phi_N$ (see equation (\\ref{eq:pointMassPotential})) changes by a factor $F_{SF}$, equation (\\ref{eq:FSFFactor}), in comparison with the purely Newtonian case. Thus, for $\\alpha > 0$ the growth of structures speeds up in comparison with the Newtonian case (without SF). For the $\\alpha <0$ case the effect is to diminish the formation of structures.  It is important to note that particles in our model are gravitating particles and that the SF acts as a mechanism that modifies gravity. The effective mass of the SF ($m_{SF}=1/\\lambda$) only sets an interaction length scale for the Yukawa term.  In this work we only varied the amplitude of the SF --parameter $\\alpha$-- leaving the scale length, $\\lambda$, of the SF unchanged. However, in other studies we have done \\citep{mar2009b} we have found that increasing $\\lambda$ enhances the structure formation process for $\\alpha$ positive, and decreasing $\\lambda$ makes the structures grow at a slower rate.  We have computed the mass power spectrum in order to study the LSS formation process. The theoretical scheme we have used is compatible with local observations because we have defined the background field constant $<\\phi>  =  G_{N}^{-1} (1+\\alpha)$ or equivalently that the local gravitational constant is given by $ (1 + \\alpha) \\langle \\phi \\rangle^{-1} $, instead of being given by $\\langle \\phi \\rangle^{-1}$. % A direct consequence of the approach is that  the amount of matter (energy) has to be increased for positive values of $\\alpha$ and diminished  for negative values of $\\alpha$ with respect to the standard $\\Lambda$CDM model in order to have a flat cosmological model. % Quantitatively, our model demands to have $\\Omega/ (1+\\alpha) =1$ and this changes the amount of dark matter and energy of the model for a flat cosmological model, as assumed. The general gravitational effect is that  the interaction including  the SF changes by a factor $F_{SF}(r,\\alpha,\\lambda) \\approx 1+\\alpha \\, \\left( 1+\\frac{r}{\\lambda} \\right)$ for $r<\\lambda$ in comparison with the Newtonian case. Thus, for $\\alpha >0$ the growth of structures speeds up in comparison with the Newtonian case.  For the   $\\alpha <0$ case the effect is to diminish the formation of structures.  For $r> \\lambda$ the dynamics is essentially Newtonian.  Comparison with the power spectrums from galaxies in the SDSS catalog and that inferred from Lyman-$\\alpha$ forest observations tell us that $\\Lambda$CDM predicts more structure formation in the regime of smaller scales ($k>0.4$ $h/$Mpc). Whereas,  the model with SF with $\\alpha=-1/2$ and $\\lambda=5$ Mpc$/h$, follows the general trend of the observed power spectrum. In this way we are able to construct a model that predicts the observed structure formation in the regime of small scales, with lesser number of halo satellites than the $\\Lambda$CDM model.   % %\\begin{acknowledgement} %This work received financial support by %CONACYT grant number \\ CB-2007-01-84133. %\\end{acknowledgement}   %%%%%%%%%%%%%%%% REFERENCIAS %%%%%%%%%%%%%%%%  %"
    },
    "1112/1112.1073.txt": {
        "abstract": "The CO-\\htwo \\ conversion factor in galaxies is typically described as bimodal: one value for discs and quiescent regions, and another (lower) value for mergers and starbursts.  In this proceeding, I will describe both empirical observational evidence that the conversion factor varies with physical environment, as well as a theoretical model which aims to understand the physical processes which drive these variations.  I present a functional form for \\xco \\ which can be applied to observations ranging in scale from $\\sim 70$ pc to galaxy-wide scales, and show the consequences of the application of this model to the Kennicutt-Schmidt star formation law. ",
        "introduction": "Deriving an \\htwo \\ gas mass from a giant molecular cloud (GMC) or galaxy full of clouds involves the usage of tracer molecules such as $^{12}$CO (hereafter, CO).  This is because \\htwo \\ has no permanent dipole moment, and its first quadrupole moment lies $\\sim 500$ K above ground, significantly warmer than the typical ISM temperature in a GMC.  Converting from a CO emission line strength to an underlying \\htwo \\ gas mass  involves the usage of a CO-\\htwo \\ conversion factor.  The CO-\\htwo \\ conversion factor goes by two names in the literature: \\xco\\footnote{\\xco \\ is also referred to as the $X$-factor.  I will use \\xco \\ or the $X$-factor arbitrarily in this proceeding.} \\ and \\alphaco. The former has units of \\xcounits, while the latter \\msun ${\\rm pc}^{-2}$(K-\\kmsend)$^{-1}$.  The two are equivalent, and simply related via \\xco = $6.3 \\times 10^{19} \\alpha_{\\rm CO}$ in the aforementioned units.  Since both \\xco \\ and \\alphaco \\ are used in the literature, I will present the results of this proceeding in terms of both units.  Observationally determining an $X$-factor requires deriving an independent measurement of \\htwo \\ gas mass, and comparing this to an observed CO line flux.  Via a variety of methods \\citep[some of these are summarised in the seminal conference proceeding by][]{sol91}, it has been found that the conversion factor is roughly constant in the Galaxy and nearby galaxies (when the metallicity is of order solar) with value \\xco= $2-4 \\times 10^{20} \\xcounits$, or \\alphaco $\\approx 3-6$ \\alphacounits \\ \\citep[for a reference list of these observations, please see the introduction of][]{nar11c}.  However, despite the seeming constancy of the $X$-factor in relatively ``normal'' GMCs, observations of nearby ultraluminous infrared galaxies (ULIRGs) show that the usage of a Milky Way \\xco \\ in these environments would cause the inferred \\htwo \\ gas mass to exceed the measured dynamical mass \\citep{dow98}.  Hence, in ULIRGs, the $X$-factor must be lower than within the Galaxy.  In Figure~\\ref{figure:observational_xco}, we show the distribution of \\xco \\ and \\alphaco \\ values from the \\citet{dow98} survey, as well as a shaded region for the typical range of Galactic values.  There is a dispersion amongst the ULIRG $X$-factors, though on average they are lower than the Galactic mean.  Despite the observed dispersion in ULIRG conversion factors, the fact that they are lower on average than the Milky Way value has led the community to largely adopt a bimodal picture of \\xco: a value of $\\alphaco \\approx 4$ for disc galaxies, and $\\alphaco \\approx 0.8$ for starbursts and mergers.  The ramifications of a bimodal form of the CO-\\htwo \\ conversion factor are significant.  As an example, \\citet{dad10b} and \\citet{gen10} demonstrated the effects of assuming a bimodal \\xco \\ on the Kennicutt-Schmidt (KS) star formation rate-gas density relation.  While in $\\Sigma_{\\rm SFR}-\\Sigma_{\\rm CO}$ space, galaxies at low and high-\\z \\ lie on an arguably unimodal relation, the introduction of a bimodal \\xco \\ to convert the CO luminosity to an \\htwo \\ gas mass causes the KS relation to become bimodal.  That is, mergers and discs lie on different track on the KS relation when utilising a different \\xco \\ for each.  This is shown explicitly in Figure~\\ref{figure:ks}, where I have compiled the galaxies presented in \\citet{dad10b} and \\citet{gen10}, and plotted (in the left panel) the $\\Sigma_{\\rm SFR}-\\Sigma_{\\rm CO}$ relation, and (in the middle panel), the $\\Sigma_{\\rm SFR}-\\Sigma_{\\rm H2}$ relation which utilises a bimodal $X$-factor.  The mergers (squares) and discs (circles) occupy different tracks.  This has given rise to a terminology in the literature, as well as at conferences (including this one!), that mergers and discs have different ``modes'' of star formation.  Mergers, according to this relation, form stars more efficiently (i.e. needing less gas to sustain a given SFR) than discs.  We will neglect the right panel of Figure~\\ref{figure:ks} for now, and will return to it shortly. ",
        "conclusions": "The principal result from this work is that the CO-\\htwo \\ conversion factor in galaxies varies with the physical conditions in the galaxy. It is not  bimodal with the global morphology of the galaxy.  Rather, the physical environment sets the CO-\\htwo \\ conversion factor.  The dominant drivers in the conversion factor are the gas temperature, velocity dispersion and metallicity.  While the former two are difficult to observe, they can be parameterised by the observable CO intensity.  With this, the CO-\\htwo \\ conversion factor can be described as function of CO intensity and gas phase metallicity.  We show that this function can be derived both from empirical observational results, as well as numerical modeling.     %\\begin{figure} %    \\centering %    \\includegraphics[angle=90,scale=0.4]{observational_xco.ps} %\\caption{\\label{figure:observational_xco}} %\\end{figure} %%%%%%%"
    },
    "1112/1112.1637_arXiv.txt": {
        "abstract": "We outline the scientific motivation behind a search for gravitational waves associated with short gamma ray bursts detected by the InterPlanetary Network (IPN) during LIGO's fifth science run and Virgo's first science run. The IPN localisation of short gamma ray bursts is limited to extended error boxes of different shapes and sizes and a search on these error boxes poses a series of challenges for data analysis. We will discuss these challenges and outline the methods to optimise the search over these error boxes. ",
        "introduction": "Short hard gamma ray bursts (short GRBs) are believed to be produced by mergers of either two neutron stars or a neutron star and a stellar mass black hole \\cite{Nakar:2007, ShibTan06, Berger:2010qx}. These events are ideal sources for strong gravitational wave emission \\cite{ACST94, Kiuchi:2010ze}. If an observation of both gamma rays and gravitational waves (GW) originating from the same event could be achieved, it will increase confidence and allow for better science output. It is thus of great importance to constantly monitor and record GRBs to allow a GW search to be performed around the burst times. Systematic analyses of GW data around short GRB times have been done in the past and the most recent publications from the LIGO-Virgo group contain results from the Swift-observed GRBs during LIGO's fifth science run (S5) and Virgo's first science run (VSR1) \\cite{Abadie:2010uf, Collaboration:2009kk}. This paper reports the methodology and motivation for a proposed GW search around the times of 20 additional GRB during S5/VSR1. These bursts were observed by the InterPlanetary Network (IPN) \\cite{Hurley:2002wv, HurleyHTML}, a group of satellites orbiting the Earth and Mars and operating, among other equipment on board, gamma ray detectors. These bursts were detected between 2006 and 2007 and have been localized, in both time and sky location, such that a targeted GW search is possible.  The InterPlanetary Network \\cite{Hurley:2002wv, Hurley:1999ym} employs several space missions and synthesizes data obtained from the detection of the same burst by different spacecraft equipped with gamma ray detectors. The IPN has been operating for three successive generations; presently the third IPN (IPN3) began its operation in November 1990. Currently the spacecraft gathering data are Konus-WIND, Suzaku, INTEGRAL,  RHESSI, Swift, Fermi/GBM (in Earth orbit), MESSENGER (in Mercury orbit) and Mars Odyssey (in Mars orbit) \\cite{HurleyHTML}. When the duty cycles and effective fields of view of all the missions in the network are considered, the IPN is an all-time, isotropic GRB monitor.  In this paper, we first review the motivation for performing a joint GW-GRB search as an overview and then provide details of the IPN short GRBs that were identified during the LIGO-Virgo S5-VSR1 science runs.  We discuss the necessary changes to the current ongoing analysis (for the S6/VSR2-3 science runs Swift and Fermi-observed GRBs, \\cite{lvc:s6grb, Harry:2010fr}) that are needed to implement a search on GRBs identified by the IPN network which may be less well localized in both sky position and time than the corresponding bursts identified by the Swift satellite and used in previous analyses \\cite{Abadie:2010uf, Collaboration:2009kk}. ",
        "conclusions": "We have presented the methodology for a search for gravitational waves around the times of short GRBs, detected by IPN during S5 and VSR1. This search has all the needed tools and will commence in a short time. The work in this paper was presented in the form of a poster at the 10th Amaldi Conference for Gravitational Waves organised in Cardiff, UK in July 2011. These types of searches are very promising for the future detection of gravitational waves and combined with the prospect of detecting other electromagnetic counterparts from GRBs, e.g. radio or optical pre- or afterglows (summarised in \\cite{Predoi:2009af, Coward:2011yr}), may open the doors for true multi-messenger astronomy with gravitational waves."
    },
    "1112/1112.2772_arXiv.txt": {
        "abstract": "{ The stellar astronomy has always been considered  the fundamental source of knowledge about the basic building blocks of the universe --- the stars. It has proved correctness of many physical theories --- like e.g.  the idea of nuclear fusion in stellar cores, the exchange of mass in interacting binaries or models of stellar evolution towards white dwarfs or neutron stars.  Despite its well acknowledged importance  it seems to be loosing its interestingness for students, for telescope allocation committees at large observatories, as well as for granting agencies.  In the domain of big telescopes it has been gradually overtaken by the extra-galactic research and cosmology, surviving however at smaller observatories and among most advanced amateur astronomers.  We try to analyse the main obstacles lowering the efficiency of research in contemporary stellar astronomy. We will shortly tackle several  problems induced by paradigmatic changes in handling the extraordinary amount of data provided by current instruments as well as by introduction of economical criteria  and  factory-like management into the modern astronomy.  Finally we speculate  the reasons of a marginal role of Virtual observatory  in contemporary stellar research and give some ideas of possible improvements. }  \\FullConference{Accelerating the Rate of Astronomical Discovery, sps5\\\\ August 11-14, 2009\\\\ Rio de Janeiro, Brazil}  \\begin{document} ",
        "introduction": "The technological potential of the current astronomical research is enormous. The opening of almost all windows of electromagnetic spectrum facilitated by advances in detector technology together with specialized astronomical satellites as well as 8 to 10m class telescopes harnessed by extremely sensitive instrumentation provides the astronomical community with overwhelmingly massive amount of data about the very deep structure of our universe as well as detailed multi-wavelength information about the billions of even very distant objects. The most powerful GRIDs of computers as well as clever infrastructure of Virtual Observatory are supposed to easily handle the data avalanche and so amount of principal astronomical discoveries should follow the technological Moore's law. But is it really so ?   Unfortunately, there is a danger of loosing the important dimension --- the time --- from observations of the most powerful telescopes and their archives. The astronomy of past was  based on long-term monitoring of selected objects and most physical analysis was concerning their  time-dependent nature (e.g. the search of periods, phase dependent changes of line profiles or changes of intensity or spectra during the outbursts of novae or cataclysmic variables). The largest amount of data coming from large facilities is currently only some kind of snapshots of the Universe, being forced by the way the telescope time is allocated  and by pressure on rapid publication of results.  All of this may lead  to the paradoxical situation when the next generation of astronomers will have from the then available data the feeling of the universe like something static and unchangeable as it was in the ancient times before the Galileo's invention of telescope. ",
        "conclusions": "Stellar astronomy seems to be  now in stagnation period. The interest of professional astronomers is turned towards very distant extragalactic objects. The still existing professional stellar astronomy has been driven towards producing only  snapshots  than  systematic monitoring.  The stars are interesting for many astronomers only thanks to their planets, whereas most of the physical problems of the stellar structure and stellar evolution as well as nature of some phenomena remain still unsolved.  Unfortunately, the stellar research seems  not to be very appealing for students either, despite the great importance of stellar astronomy for understanding of galactic structure and evolution.  On the other hand there is a lot of enthusiastic people among amateur astronomers, who continue to monitor the stellar variability on the long time scales and report outburst of cataclysmic variables, but they can hardly take spectra of most interesting stars, mainly due to their limited capabilities.  The solution of the problems described might be the Virtual observatory, allowing the integration of dispersed observations from different (even small) teams.  The VO principles might attract the number of young people fascinated by new software technologies.  Unfortunately, there is a lack of stellar spectra in VO archives as well as little support of specific techniques important in stellar astronomy in VO tools with respect to existing legacy applications.  There is also very little VO awareness among stellar astronomers and their students.  Thus the possible solution of the current stellar crisis could be expressed by several goals: \\begin{itemize} \\item Convince people to put more  stellar spectra into VO archives \\item Circumvent the data jealousy by convincing people about advantage of openness \\item Get resources needed for development of new  more versatile analysis tools with VO interface (or for adding  the VO interface to legacy applications) \\item Rise the level of VO awareness by systematic VO education or rather \"VO evangelization\" \\end{itemize} We can hopefully still preserve the image of the dynamic eternally changing universe  for  our grandsons instead of hinting them the idea of the Immutable Heavens again."
    },
    "1112/1112.0777_arXiv.txt": {
        "abstract": "We report the discovery of two consecutive supernovae (SNe), \\object{2010cu} and \\object{2011hi}, located at 0.37\\arcsec\\ (180 pc) and 0.79\\arcsec\\ (380 pc) projected distance respectively from the centre of the \\textit{K}-band nucleus of the luminous infrared galaxy \\object{IC 883}. The SNe were discovered in an ongoing near-infrared \\textit{K}-band search for core-collapse SNe in such galaxies using the ALTAIR/NIRI adaptive optics system with laser guide star at the Gemini-North Telescope. These are thus the closest SNe yet discovered to a LIRG nucleus in optical or near-infrared wavelengths. The near-infrared light curves and colours of both SNe are consistent with core-collapse events. Both SNe seem to suffer from relatively low host galaxy extinction suggesting that regardless of their low projected galactocentric distances, they are not deeply buried in the nuclear regions of the host galaxy. ",
        "introduction": "Luminous ($L_{{\\rm IR}} > 10^{11} L_{\\odot}$) and ultraluminous ($L_{{\\rm IR}} > 10^{12} L_{\\odot}$) infrared (IR) galaxies (LIRGs and ULIRGs, respectively), have high star formation (SF) rates. The fraction of SF in U/LIRGs in the local universe is small compared to that in normal spiral galaxies; however, at high redshift they become the dominant sources of SF \\citep[e.g.,][]{magnelli11}. Stars more massive than $\\sim8M_{\\sun}$ explode as core-collapse supernovae (CCSNe), and due to their short life cycles they are a very useful tool for tracing ongoing SF rates, independent of the conventional way of using the galaxy IR luminosity. Recently, \\citet{anderson11} have presented evidence that the SN population in the interacting LIRG system \\object{Arp 299} differs from those observed in normal spiral galaxies with the Ib and IIb SNe being more numerous and centrally concentrated than more common Type II SNe. Until now, this kind of study has been impossible due to insufficient numbers of SNe discovered in U/LIRGs, providing a strong motivation for SN searches in galaxies with high SF rates. The extinction-free searches at radio wavelengths using very long baseline interferometry (VLBI) have been successful in discovering SNe and SN remnants in U/LIRGs, e.g., in \\object{Arp 220} \\citep[][]{parra07, batejat11} and \\object{Arp 299} \\citep{perez-torres09, ulvestad09}. The reason for the low optical discovery rate of CCSNe in U/LIRGs is the high dust extinction combined with spatial concentration of SF in the crowded nuclear regions within the central few hundreds of parsecs \\citep[see e.g.,][]{soifer01}. Therefore imaging in the near-IR where the extinction is strongly reduced, and high spatial resolution is achievable with ground-based adaptive optics (AO) or space-based imaging is required. Recently such methods have been shown to be very efficient with the discovery of several SNe in the LIRG nuclear/circumnuclear regions. \\object{SN 2004ip} was discovered with the Very Large Telescope using the natural guide star NACO AO system \\citep{mattila07}. \\object{SNe 2008cs} and \\object{2004iq} reported in \\citet{kankare08} were discovered using the Gemini-North Telescope with the laser guide star (LGS) assisted ALTAIR AO system and the \\textit{Hubble Space Telescope} (\\textit{HST}) NICMOS archive data, respectively.  In this Letter, we report the consecutive discovery of two SNe in the same host galaxy with the Gemini-North Telescope\\footnote{Based on observations obtained at the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the Science and Technology Facilities Council (United Kingdom), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), Minist\\'{e}rio da Ci\\^{e}ncia e Tecnologia (Brazil) and Ministerio de Ciencia, Tecnolog\\'{i}a e Innovaci\\'{o}n Productiva (Argentina).}, \\object{SN 2010cu} and \\object{SN 2011hi}. \\citet{sanders03} report the host LIRG \\object{IC 883} (also known as e.g., \\object{UGC 8387}, \\object{Arp 193}, and \\object{IRAS 13183+3423}) to have an IR luminosity of $L_{IR} = L[8-1000 \\mu m] = 4.7 \\times 10^{11} L_{\\odot}$ and a distance of 100 Mpc (H$_{0}$ = 75 km~s$^{-1}$Mpc$^{-1}$) giving a projected linear distance scale $1\\arcsec = 485$~pc. The IR luminosity of \\object{IC 883} indicates a CCSN rate of $\\sim1.3$ yr$^{-1}$ derived with the empirical relation of \\citet{mattila01}, consistent with the discovery of two SNe within two years. In \\citet{romero-canizales11b} we report the details of our radio follow-up of \\object{SNe 2010cu} and \\object{2011hi} and investigate the innermost nuclear regions of \\object{IC 883} through high-angular resolution, high-sensitivity radio observations. ",
        "conclusions": "The adaptive optics assisted discovery of two consecutive SNe located at $<400$~pc from the LIRG nucleus demonstrates the importance of high spatial resolution in searches for CCSNe in highly crowded starburst regions in LIRGs. We note that of the five SNe discovered using high resolution near-IR observations in our sample of LIRGs (see Sect.~1), four are within 1~kpc from the host nucleus.  Based on the near-IR light curve fitting to templates we find \\object{SN 2010cu} to be consistent with either a Type IIP or a near-IR bright Type IIn/L SN with a host galaxy extinction of $A_{V}\\lesssim1$~mag. Taking into account also the extinction map \\object{SN 2011hi} is found to be most consistent with a Type IIn/L SN during the near-IR excess phase with host galaxy extinction of $A_{V}\\sim0$~mag. A Type Ib/c classification was found unlikely for both the SNe. Future near-IR and radio follow-up observations of \\object{SN 2011hi} will be able to better constrain the nature and line-of-sight extinction of this potentially very interesting event  The surprisingly low derived line-of-sight extinctions appear to be in conflict with the expectation of the occurrence of highly obscured CCSNe in the nuclear region of LIRGs. For example, in \\citet{kankare08} $A_{V}=16$~mag of host galaxy extinction was derived with near-IR template light curve fitting for \\object{SN 2008cs}, located at a projected distance of 1.5~kpc from the host LIRG nucleus. Based on this it seems more likely that \\object{SNe 2010cu} and \\object{2011hi} are located in the outer parts of the LIRG system which we see projected close to the LIRG nucleus. Additionally the majority of LIRGs at the high IR luminosity end are complex interacting systems and their dust extinction distributions can be very complex compared to normal spiral galaxies with symmetric exponential disks, see e.g., \\citet{alonso-herrero06} and \\citet{vaisanen08}. Despite the low extinctions such SNe occurring close to the LIRG nuclei remain undiscovered by the current SN searches and will therefore contribute to the 'SN rate problem' recently suggested by \\citet{horiuchi11}."
    },
    "1112/1112.4859_arXiv.txt": {
        "abstract": "{} {We develop a statistical analytical model that predicts the occurrence frequency distributions and parameter correlations of avalanches in nonlinear dissipative systems in the state of a slowly-driven self-organized criticality (SOC) system.} {This model, called the fractal-diffusive SOC model, is based on the following four assumptions: (i) The avalanche size $L$ grows as a diffusive random walk with time $T$, following $L \\propto T^{1/2}$; (ii) The energy dissipation rate $f(t)$ occupies a fractal volume with dimension $D_S$, (iii) The mean fractal dimension of avalanches in Euclidean space $S=1,2,3$ is $D_S \\approx (1+S)/2$; and (iv) The occurrence frequency distributions $N(x) \\propto x^{-\\alpha_x}$ based on spatially uniform probabilities in a SOC system are given by $N(L) \\propto L^{-S}$, with $S$ being the Eudlidean dimension. We perform cellular automaton simulations in three dimensions ($S=1,2,3$) to test the theoretical model.} {The analytical model predicts the following statistical correlations: $F \\propto L^{D_S} \\propto T^{D_S/2}$ for the flux, $P \\propto L^{S} \\propto T^{S/2}$ for the peak energy dissipation rate, and $E \\propto F T \\propto T^{1+D_S/2}$ for the total dissipated energy; The model predicts powerlaw distributions for all parameters, with the slopes $\\alpha_T=(1+S)/2$,  $\\alpha_F=1+(S-1)/D_S$, $\\alpha_P=2-1/S$, and $\\alpha_E=1+(S-1)/(D_S+2)$.\tThe cellular automaton simulations reproduce the predicted fractal dimensions, occurrence frequency distributions, and correlations within a satisfactory agreement within $\\approx 10\\%$ in all three dimensions.} {One profound prediction of this universal SOC model is that the energy distribution has a powerlaw slope in the range of $\\alpha_E=1.40-1.67$, and the peak energy distribution has a slope of $\\alpha_P=1.67$ (for any fractal dimension $D_S=1,...,3$ in Euclidean space $S=3$), and thus predicts that the bulk energy is always contained in the largest events, which rules out significant nanoflare heating in the case of solar flares.} ",
        "introduction": "The statistics of nonlinear processes in the universe often shows powerlaw-like distributions, most conspicously in energetic dynamic phenomena in astrophysics (e.g., solar and stellar flares, pulsar glitches, auroral substorms) and in catastrophic events in geophysics (e.g., earthquakes, landslides, or forest fires). The most widely known example is the distribution of earthquake magnitudes, which has a powerlaw slope of $\\alpha \\approx 2.0$ for the differential frequency distribution (Turcotte 1999), the so-called Gutenberg-Richter (1954) law.  Bak, Tang, and Wiesenfeld (1987, 1988) introduced the theoretical concept of self-organized criticality (SOC), which has been initially applied to sandpile avalanches at a critical angle of repose, and has been generalized to nonlinear dissipative systems that are driven in a critical state. Comprehensive reviews on this subject can be found for applications in geophysics (Turcotte 1999), solar physics (Charbonneau et al.~2001), and astrophysics (Aschwanden 2011).  Hallmarks of SOC systems are the scale-free powerlaw distributions of various event parameters, such as the peak energy dissipation rate $P$, the total energy $E$, or the time duration $T$ of events. While the powerlaw shape of the distribution function can be explained by the statistics of nonlinear processes that have an exponential growth phase and saturate after a random time interval (e.g., Willis and Yule 1922; Fermi 1949; Rosner and Vaiana 1978; Aschwanden et al.~1998; Aschwanden 2004, 2011), no general theoretical model has been developed that predicts the numerical value of the powerlaw slope of SOC parameter distributions. Simple analytical models that characterize the nonlinear growth phase with an exponential growth time $\\tau_G$ and the random distribution of risetimes with an average value of $t_S$, predict a powerlaw slope of $\\alpha_P = 1 + t_S/\\tau_G$ for the energy dissipation rate (e.g., Rosner and Vaiana 1978; Aschwanden et al.~1998), but cellular automaton simulations suggest a much more intermittent energy release process than the idealized case of an avalanche with a single growth and decay phase. An alternative theoretical explanation for a slope $\\alpha_E=3/2$ was put forward by a dimensional argument (Litvinenko 1998), which can be derived from the definition of the kinetic energy of convective flows, but this model entails a specific physical mechanism that has not universal validity for SOC systems.  In this Paper we propose a more general concept where the powerlaw slope of the occurrence frequency distribution of SOC parameters depends on the fractal geometry of the energy dissipation domain. We aim for a universal statistical model of nonlinear energy dissipation processes that is independent of any particular physical mechanism. The fractal structure of self-organized processes has been stressed prominently from beginning (Bak, Tang, \\& Wiesenfeld 1987, 1988; Bak and Chen 1989), but no quantitative theory has been put forward that links the fractal geometry to the size distribution of SOC events. Fractals have been studied independently (e.g., Mandelbrot 1977, 1983, 1985), while the fractal geometry of SOC avalanches was postulated (e.g., see textbooks of Bak 1996; Sornette 2004; Aschwanden 2011), but no general self-consistent model has been attempted.  This paper presents an analytical theory that derives a theoretical framework to quantitatively link the concept of fractal dimensions to the occurrence frequency distributions of SOC avalanche events (Section 2), tests of the analytical theory with numerical simulations of cellular automaton models in three Euclidean dimensions (Section 3), a comparison and application to solar flares (Section 4), and a summary of the model assumptions and conclusions (Section 5). ",
        "conclusions": "We developed an analytical theory for the statistical distributions and correlations of observable parameters of SOC events, which includes the avalanche length scale $L$, the time duration $T$, the peak $P$ and energy dissipation rate $F$. The basic assumptions of our analytical model, which we call the {\\sl fractal-diffusive SOC model}, are the following:  \\begin{enumerate} \\item{{\\bf Diffusive expansion of SOC avalanches:} The radius $r(t)$ or spatial length scale $L$ of an avalanche grows with time like the average of a diffusive random walk, which predicts a statistical correlation $L \\propto T^{1/2}$ between the length scale $L$ and time duration $T$ of the avalanche.}  \\item{{\\bf Fractal Energy Dissipation Rate:} The complexity of random next-neighbor interactions in a critical SOC state can be characterized approximately with a fractal geometry. The volume (or area) of the instantaneous energy dissipation rate is assumed to have a fractal dimension $D_S$. The predicted statistical correlations are: $F \\propto T^{D_S/2}$, $P \\propto T^{S/2}$, and $E \\propto T^{1+D_S/2}$.}  \\item{{\\bf Mean Fractal Dimension:} The mean fractal dimension $D_S$ for different Euclidean space dimensions $S=1,2,3$ can be estimated from the arithmetic mean of the minimum dimension for a propagating avalanche, $D_{S,min} \\approx 1$, and the maximum (Euclidean) dimension $D_{S,max} = S$, which yields $D_S \\approx (1+S)/2$.}  \\item{{\\bf Occurrence Frequency Distributions:} Equal probability of avalanches with size $L$ at various spatial locations in a uniform, slowly-driven SOC system predicts a probability distribution of $N(L) \\propto L^{-S}$. A direct consequence of this assumption, together with the other assumptions made above, yields a powerlaw function $N(x) \\propto x^{-\\alpha_x}$ for the occurrence frequency distributions of all parameters, which is the hallmark of a SOC system. The predicted powerlaw indices are: $\\alpha_T=(1+S)/2$, $\\alpha_F=1+(S-1)/D_S$, $\\alpha_P=2-1/S$, and $\\alpha_E=1+(S-1)/(D_S+2)$. Specifically, for applications to 3-D phenomena, absolute values are predicted for the powerlaw slopes $\\alpha_L=3$ and $\\alpha_T=2$, and $\\alpha_P=1.67$, while a range of $\\alpha_E=1.4,...,1.67$ is expected for any fractal dimension in the range of $1 \\le D_3 \\le 3$.} \\end{enumerate}  We have validated our theory by a detailed comparison with a set of SOC simulations using a specific form of a cellular automaton avalanche model (connectivity, stability threshold, redistribution rule, etc), and found a good agreement between theory and numerical simulations, in the order of $\\approx 10\\%$ for the powerlaw slopes ($\\alpha_L, \\alpha_T, \\alpha_E, \\alpha_P$), the power indices of correlated parameters ($\\beta_{TL}, \\beta_{TP}, \\beta_{TE}$), and the fractal dimensions $D_S$, for all three Euclidean space dimensions $S=1,2,3$. Yet, at its most general level our theory is saying that the self-similarity of energy release statistics in such models is a direct reflection of the fractal nature of avalanches. The SOC flare model recently proposed by Morales \\& Charbonneau (2008, 2009) offers an interesting test of this conclusion. Their model, defined on a set of initially parallel magnetic flux strands, contained in a plane and subjected to random sideways deformation, with instability and readjustment occurring when the crossing angle of two flux strands exceeds some threshold angle. This model is thus strongly anisotropic, with pseudo-local stability and redistribution, in the sense that these operators now act on nearest-neighbors nodes located along each flux strand, rather than in the immediate spatial vicinity of the unstable sites. This is very different from the isotropic Lu et al.~(1993)-type SOC model used here for validation. Yet, the results compiled in Table 2 of Morales \\& Charbonneau (2008) for their highest resolution simulations reveal that the theoretical occurrence frequency distributions obtained herein, do hold within the stated uncertainties on the power-law fits. Likewise, the power indices of the correlated parameters also agree with theory within the inferred uncertainties. This provides additional empirical support to our conjecture that the fractal-diffusive SOC model does represent a robust characterization of avalanche energy release in SOC systems in general. However, it should be remembered that the inferred scaling laws are only valid for a slowly-driven SOC system, while alternative SOC systems with time-variable drivers or non-stationary input rates exhibit modified occurrence frequency and waiting time distributions (Charbonneau et al.~2001; Norman et al.~2001), which was also found in solar observations extending over multiple solar cycles (Crosby et al.~1993; Biesecker 1994; Bai 1993; Aschwanden 2010a).  What other predictions can be made from our analytical model? For SOC processes in 3-D space, which is probably the most common application in the real world, the mean fractal dimension is predicted to be $D_3\\approx 2.0$, which can be tested by measurements of fractal dimensions in observations. The 20 largest solar flares observed with TRACE have been analyzed in this respect and an area fractal dimension of $D_2=1.89\\pm0.05$ was found at the flare peaks, which translates into a value of $D_3=2.10\\pm0.14$ if we use an anisotropic flare arcade model (Aschwanden and Aschwanden 2008a). The distribution of flare energies is predicted to have a powerlaw slope of $\\alpha_E=1.50$, which closely matches the observed statistics of solar flare hard X-ray emission ($\\alpha_E \\approx 1.49-1.56$). Since this value is undisputably below the critical limit $\\alpha=2$ of the energy integral, the total released energy is contained in the largest flares and thus rules out any significant nanoflare heating of the solar corona. Another prediction, that we did not test here with solar flare data, is the diffusive flare size scaling. Straightforward tests could be carried out by gathering statistics of the flare size evolution during individual flares (which are predicted to scale as $x(t) \\propto t^{1/2}$), as well as from the statistics of a large sample of flares, which is predicted to show a correlation $L \\propto T^{1/2}$. The application of our fractal-diffusive SOC model to solar flares implies that the subsequent triggering of local magnetic reconnection events during a flare occurs as a diffusive random walk. A similar finding of diffusive random walk was also found in the turbulent flows of magnetic bright points in the lanes between photospheric granular convection cells (Lawrence et al.~2001). The spatio-temporal scaling of the diffusive random walk predicts also the size, duration, and energy of the largest flare, which is likely to be constrained by the size $L_{AR} \\propto T_{max}^{1/2}$ of the largest active region."
    },
    "1112/1112.3965_arXiv.txt": {
        "abstract": "{ We present \\swift\\ follow-up observations of a sample of 35 unclassified faint X-ray sources drawn from the \\asca\\ Galactic centre and Galactic plane surveys. Our short, pointed XRT observations allow detections down to a limiting 0.3--10 keV flux of $F_X \\sim10^{-13}~\\flux$, which translates into a luminosity of $L_X\\sim10^{33}~\\lum$ for an assumed distance of $D=8$~kpc. The brightest source in our sample reaches a maximum 0.3--10 keV luminosity of $L_X\\sim2\\times10^{36}~(D/8~\\mathrm{kpc})^2~\\lum$ during our \\swift\\ observations. We detect 16 ($46\\%$) of the \\asca\\ sources with the XRT, while 19 were not detected during our program. Since we are probing the faint end of the \\asca\\ source populations, we expect a large fraction of the non-detections to be due to the Eddington bias. This is strengthened by the fact that we find the observed XRT count rates to be predominantly lower than expected based on the reported \\asca\\ intensities. Nevertheless, investigation of the \\asca\\ properties and any possible long-term X-ray variability leads us to conclude that the non-detections likely include two spurious \\asca\\ detections and three objects that are highly variable or transient X-ray sources. For the 16 XRT-detected sources we obtain positional accuracies of $\\sim2-4''$, which significantly improves upon their \\asca\\ uncertainties of $\\sim1'$. We use the X-ray spectra and variability to characterise these objects. Most appear to be faint, persistent X-ray emitters that have highly absorbed spectra. Based on their X-ray properties we identify three accreting compact objects: one confirmed (\\ucxb) and one candidate (\\porb) X-ray binary, and one possible magnetically accreting white dwarf (\\polar). Furthermore, we use the improved positions of XRT-detected sources to search for counterparts in simultaneously obtained \\swift/UVOT images and possible associations with catalogued sources at various wavelengths. This reveals three possible main sequence stars amongst our sample. The other sources remain unclassified, but our improved XRT positions provide good prospects for dedicated follow-up observations that have the potential to elucidate the nature of these faint, unclassified \\asca\\ sources. } ",
        "introduction": "The Advanced Satellite for Cosmology and Astrophysics (\\asca) was launched in 1993 and operated until 2000. It was the first mission that combined imaging capabilities with a relatively broad energy passband (0.7--10 keV). During its lifetime, the \\asca\\ satellite pointed towards the Galactic plane and the inner bulge of the Galaxy many times. In particular, two dedicated surveys were performed that detected about 200 X-ray sources with 0.7--10 keV fluxes of $F_X\\gtrsim 3\\times 10^{-13}~\\flux$ \\citep[][]{sugizaki01,sakano02}. Both surveys used the GIS (gas imaging spectrometer) detector, which had a spatial resolution of $\\sim30\\arcsec$ and a field of view (FOV) of $\\sim50'$.  The \\asca\\ Galactic plane survey entailed a systematic study of X-ray sources located in the inner Galactic disk \\citep{sugizaki01}. The region of $| l | \\lesssim45^{\\circ}$, $| b |\\lesssim 0.4^{\\circ}$ was covered between 1996 and 1999 using successive 10-ks observations. Investigation of three different energy bands (0.7--2, 2--10 and 0.7--7 keV) resulted in the identification of 163 individual X-ray sources detected with a significance of $\\gtrsim 4 \\sigma$. \\asca\\ targeted a region of $5^{\\circ} \\times 5^{\\circ}$ around the centre of our Galaxy between 1993 and 1999 \\citep{sakano02}. Searching the 0.7--3 and 3--10 keV images yielded a total of 52 point sources that were detected with a significance of $\\gtrsim5\\sigma$. In both \\asca\\ surveys, spectral analysis was performed for sources that were detected with a significance of $\\gtrsim5\\sigma$. The GIS data were fitted to a simple absorbed powerlaw model to characterise the X-ray spectra \\citep[][]{sugizaki01,sakano02}. Source positions could be determined with a $90\\%$ confidence positional accuracy of $\\sim1'$ \\citep{ueda1999,sugizaki01,sakano02}.  After cross-correlation with the \\einstein\\ and \\rosat\\ bright source catalogues, about 2/3 of the \\asca-detected X-ray sources remained unclassified \\citep{sugizaki01,sakano02}. These objects have 0.7--10 keV X-ray fluxes in the range of $F_X \\sim 10^{-13} - 10^{-11}~\\flux$, which translates into a luminosity of $L_X \\sim 10^{33}-10^{35}~\\lum$ for a distance of $D=8$~kpc. This is right in between the well-known bright X-ray point sources ($L_X>10^{36}~\\lum$) that are mostly active X-ray binaries \\citep[e.g.,][]{chen97}, and dim X-ray sources ($L_X<10^{33}~\\lum$) that consist of a collection of source types such as accreting white dwarfs, dormant X-ray binaries, active stars or young pulsars \\citep[e.g.,][]{muno2009}.  The luminosity range traced by the unclassified \\asca\\ sources suggests that a variety of source types might be amongst them. For example, it covers the intensities typically seen for strongly magnetised neutron stars (magnetars), bright magnetically accreting white dwarfs (polars and intermediate polars), sub-luminous X-ray binaries and X-ray emitting massive stars \\citep[e.g.,][]{maeda1996,ezuka1999,wijnands06,chelovekov07,gelfland2007,heras08,kaur2010,degenaar09_gc,anderson2011}. In addition, the sample may contain foreground stars or background active galactic nuclei (AGN). The main difficulty in classifying the \\asca\\ sources is their relatively large positional uncertainties of $\\sim$1$'$. This inhibits follow-up observations aiming to search for counterparts at other wavelengths.  We launched a program to observe unclassified \\asca\\ sources with the \\swift\\ satellite \\citep[][]{gehrels2004}. The primary instrument for our observations is the X-ray telescope \\citep[XRT;][]{burrows05}. The XRT is sensitive in the 0.3--10 keV energy range and has a FOV of $\\sim23' \\times 23'$. Furthermore, we used data obtained with the Ultra-Violet/Optical Telescope \\citep[UVOT;][]{roming05} that has a FOV of $\\sim17' \\times 17'$ and can be operated using a variety of filters in a wavelength range of $\\sim1500-8000$ \\AA{} \\citep[][]{poole2008}. In addition to these narrow-field instruments, \\swift\\ is equipped with the Burst Alert Telescope \\citep[BAT;][]{barthelmy05}: a hard X-ray (15--150 keV) monitoring instrument that has a wide FOV ($2$~steradian). The combination of its pointing flexibility, good sensitivity in the soft X-ray band (below 10 keV) and multi-wavelength coverage makes \\swift\\ an ideal tool to perform follow-up observations of X-ray sources discovered by other missions \\citep[e.g.,][]{rodriguez09,starling2011,kennea2011}.  By observing unclassified \\asca\\ sources with \\swift, our primary aim is to improve their positional uncertainties to a few arcseconds, so that dedicated searches for counterparts at other wavelengths become feasible. In addition, we strive to obtain X-ray spectral information and to study possible variability in their X-ray emission. This approach provides a basis for further investigation of the nature of individual sources, as well as the population of faint \\asca\\ sources as a whole.\\footnote{See also the {\\it ChIcAGO} project, which targets a (different) sample of unclassified \\asca\\ sources using \\chan\\ and multi-wavelength follow-up observations \\citep[][]{anderson2011}.} ",
        "conclusions": "\\label{sec:discuss} We carried out \\swift\\ follow-up observations of a sample of 35 unclassified faint X-ray sources detected during \\asca\\ surveys of the Galactic plane and centre between 1993 and 1999 \\citep[][]{sugizaki01,sakano02}. A total of 16 sources from our sample were detected with the XRT, which allows us to improve their positional accuracy from $\\sim$1$'$ to $\\sim$2--4$''$. The mean offset of the improved XRT position from the \\asca\\ coordinates is $47.9''$ ($0.8'$), which is relatively close to the reported 90\\% confidence \\asca\\ positional uncertainties.  We investigated the X-ray spectra of the detected sources and studied any possible long-term X-ray variability of all our 35 targets. Our program was conducted between 2006 and 2010, hence $\\sim7-17$ years since the \\asca\\ detections. In general, it is not possible to classify the sources based on their X-ray spectral properties, particularly given the fact that the total number of counts detected during our observations was often low (see Table~\\ref{tab:spec}). Using the improved \\swift/XRT position information, we therefore searched optical and infrared source catalogues to assess any possible associations. Below we discuss the possible nature of the sample of unclassified \\asca\\ sources covered by our program. Proposed classifications and possible associations are summarised in Table~\\ref{tab:class}.   \\subsection{The nature of the \\asca\\ sources detected with \\swift}\\label{subsec:nat_detect}  \\subsubsection{Accreting compact objects}\\label{subsubsec:compactobjects}  {\\bf \\ucxb} was observed with \\swift/XRT on two different epochs in 2007 July and 2008 August--July, during which it is detected at an average 2--10 keV luminosity of $L_X \\sim 6\\times10^{34}~(D/9.2~\\mathrm{kpc})^2~\\lum$. This is considerably lower than typically observed for neutron star LMXBs and classifies the source as a sub-luminous X-ray binary \\cite[][]{muno05_apj622,wijnands06,degenaar09_gc,campana09}. The \\swift/XRT data show that the quiescent phase signalled by \\chan\\ in 2008 May must have been short \\citep[$\\lesssim11$~months, see][]{bassa08,jonker08}. Its ability to spend long times accreting at a low X-ray luminosity, combined with its thermonuclear X-ray burst properties and low absolute optical magnitude ($M_i \\gtrsim 2$), led to the suggestion that \\ucxb\\ could be an ultra-compact X-ray binary \\citep[][]{bassa08}. Such systems harbour hydrogen-poor donor stars in small orbits of $\\lesssim 90$~min \\citep[][]{nelson86}.  The XRT spectrum of \\ucxb\\ is very soft with a photon index of $\\Gamma=2.7$. Such a high spectral index is quite unusual, although other LMXBs accreting at similarly low luminosities have also been found to display very soft X-ray spectra \\citep[e.g.,][]{zand05,delsanto07,degenaar09_gc,degenaar2010_burst,armas2011}. An in-depth investigation of high-quality \\xmm\\ data of the sub-luminous (candidate) LMXB \\xte\\ revealed that the X-ray spectrum contained a soft (thermal) component in addition to the hard powerlaw emission \\citep[][]{armas2011}. In contrast, the XRT spectrum of \\ucxb\\ can be adequately fit by an absorbed powerlaw alone and does not require the addition of a soft spectral component. \\\\  \\noindent {\\bf \\porb} is the brightest source in our sample and was detected with the XRT at an average 2--10 keV luminosity of $L_X \\sim 2\\times10^{35}~(D/8~\\mathrm{kpc})^2~\\lum$. It has a highly absorbed ($N_H \\sim 8 \\times 10^{22}~\\nh$) X-ray spectrum, that can be described by a simple powerlaw model with an index of $\\Gamma=2.3$. The XRT light curve shows that the source intensity varies considerably by a factor of $\\sim10$ on a timescale of days. Given its X-ray spectral characteristics, its average intensity and the timescale of the observed variability we tentatively identify \\porb\\ as an X-ray binary.  During our program \\porb\\ spanned a luminosity range of $L_X\\sim1\\times10^{34}-5\\times10^{35}~(D/8~\\mathrm{kpc})^2~\\lum$. Such large variations in X-ray intensity (a factor $\\sim50$) are not uncommon for X-ray binaries and are usually ascribed to considerable variations in the mass-accretion rate onto the compact primary. For example, strong X-ray variability is a characteristic feature of Supergiant Fast X-ray Transients (SFXTs), in which the compact primary accretes matter from the varying wind of a massive ($\\gtrsim10~\\Msun$) supergiant companion star \\citep[e.g.,][]{sidoli08,romano2011}. Similar to what we observe for \\porb, these kind of systems are typically highly absorbed. The comparison breaks down, however, when considering the X-ray spectral shape: SFXTs typically have much harder X-ray spectra ($\\Gamma \\sim 1$) in the luminosity range that we observe for \\porb\\ \\citep[][]{romano2011}.  Variations of similar magnitude and on a similar timescale as seen for \\porb\\ have also been observed from LMXBs in which the companion star overflows its Roche lobe and accretion is governed via an accretion disk \\citep[e.g.,][]{wijnands2001_1808,linares2008,degenaar2011_1701,fridriksson2011}. Although the average 2--10 keV X-ray luminosity of this source is lower than is typically seen for active LMXBs \\citep[e.g.,][]{chen97}, there is a growing group of sub-luminous sources that accrete at similar X-ray intensities to \\porb\\ \\citep[e.g., \\ucxb\\ discussed above, but see also][]{cornelisse02,muno05_apj622,wijnands06,zand05,zand07,delsanto07,degenaar09_gc,campana09,degenaar2010_burst}.  Based on its X-ray spectral characteristics we consider it more likely that \\porb\\ belongs to the class of LMXBs and thus harbors a low-mass ($\\lesssim1~\\Msun$) companion star as opposed to a more massive donor. Dedicated follow-up observations at longer wavelengths have the potential to rule out a high-mass X-ray binary nature and thus provide a test for our proposed classification.\\\\   \\noindent {\\bf \\polar} stands out by displaying a much less absorbed ($N_H\\lesssim3\\times10^{21}~\\nh$) and much flatter ($\\Gamma \\sim 0$) X-ray spectrum than observed for the majority of X-ray sources in our sample (see Table~\\ref{tab:spec}). The hard X-ray spectrum of \\polar\\ (see Fig.~\\ref{fig:wd}) resembles that of magnetically accreting white dwarfs \\citep[polars and intermediate polars;][]{ezuka1999,muno2003,hong2009}. Although wind-accreting neutron stars in high-mass X-ray binaries (HMXB) may show similarly hard X-ray spectra, the low hydrogen absorption column density inferred from fitting the X-ray spectrum and the lack of any UVOT or catalogued optical/infrared counterpart within the $2.8''$ XRT positional uncertainty, renders a HMXB nature less likely.  Based on its X-ray spectral properties and lack of a counterpart at longer wavelengths we tentatively classify \\polar\\ as a magnetically accreting white dwarf. The limited statistics of the XRT data does not allow us to investigate the timing properties of the source. Deeper X-ray observations may test our proposed source classification by searching for common features of magnetically accreting white dwarfs, such as prominent iron emission lines in the X-ray spectrum or coherent X-ray pulsations with a periodicity on the order of minutes to hours \\citep[e.g,][]{ritter2003,muno2003,hong2009,kaur2010}. Follow-up observations at other wavelengths (optical/infrared) can further aid in the classification of this source.  \\subsubsection{Main sequence stars}\\label{subsubsec:ms_stars}  AX J1738.4--2902 is positionally coincident with an optical/infrared catalogued source. We also detect a likely counterpart at this position in our UVOT images. Based on the fact that we a find very soft X-ray spectrum with little absorption, we tentatively associate AX J1738.4--2902 with the catalogued optical/infrared object. As such, this \\asca\\ source is identified as a K0V pre-main sequence star (see Sect.~\\ref{subsec:oir}).  For another 7 of our XRT detections we find possible associations with X-ray/UV/optical/infrared/radio sources (see Tables~\\ref{tab:oir} and~\\ref{tab:class}). In Fig.~\\ref{fig:ir} we plot the approximate {\\it 2MASS} colours for the 6 possible infrared counterparts to detected \\asca\\ sources in our sample (magnitudes are not corrected for extinction).\\footnote{Taking into account extinction would push the data points down and towards the left in the colour-colour diagram (Fig.~\\ref{fig:ir}).} Circles represent sources which also have optical information, whereas squares denote those which are only detected in the infrared bands (see Table~\\ref{tab:oir}). The dashed line in Fig.~\\ref{fig:ir} indicates the expectation for A- through M-type main-sequence stars \\citep[][]{bessell1988,tokunaga2000}. The labels to the data points correspond to the source numbers used in Table~\\ref{tab:oir}.  We can see in Fig.~\\ref{fig:ir} that three sources lie along the main-sequence line: AX J1504.6--5824, AX J1720.8--3710 and AX J1738.4--2902. Our proposed counterpart for the latter is indeed classified as a K0V star (see above). If AX J1504.6--5824 and AX J1720.8--3710 are related to the infrared objects located within their XRT error circles, both are likely main sequence stars as well. The low absorption column density ($N_H\\lesssim3\\times10^{21}~\\nh$) inferred from fitting the X-ray spectrum of AX J1720.8--3710 makes this association likely. The statistics of the spectral data are limited, however, and the offset from the XRT position to the {\\it 2MASS} source is relatively large (see Tables~\\ref{tab:spec} and~\\ref{tab:oir}). In the case of AX J1504.6--5824, there is also a considerable offset between the XRT position and the UV/optical/infrared object, whereas the X-ray spectrum appears relatively hard and highly absorbed (Tables~\\ref{tab:spec} and~\\ref{tab:oir}). This casts doubt on the possible classification as a foreground star.  We emphasise that the probability of a chance detection of an optical/infrared source within the XRT error circles is high (typically $\\sim50\\%$) as the source fields are crowded (see Sect.~\\ref{subsec:oir}). It may not be a coincidence that the sources for which we find possible associations are often the ones that were detected at low significance in the XRT data and consequently have relatively large positional errors compared to the brighter XRT detections (see Table~\\ref{tab:pos}). Our proposed possible associations may be confirmed or rejected by obtaining even more accurate X-ray positions (which is currently only possible with \\chan), or by conducting dedicated follow-up observations at optical/infrared wavelengths.   \\begin{figure} \\begin{center} \\includegraphics[width=8cm]{2MASS_ASCA_custom.eps} \\end{center} \\caption[]{Approximate {\\it 2MASS} colours for possible near-infrared counterparts of detected sources in our sample (magnitudes are not corrected for extinction). Circles represent sources that also have optical information, whereas squares indicate objects that are only detected in the infrared bands. The data points are labelled to indicate the corresponding source numbers (see Table~\\ref{tab:oir}). The dashed line indicates the expectation for A- through M-type main-sequence stars. Correcting for extinction would move the data points down and towards the left in this diagram. } \\label{fig:ir} \\end{figure}   \\begin{figure} \\begin{center} \\includegraphics[width=8cm]{rates_sigma2.eps} \\end{center} \\caption[]{Observed versus expected XRT count rates for both detected and undetected sources in our sample (represented by squares and upper limit symbols, respectively). The solid line corresponds to a ratio of observed over expected count rates that equals unity, whereas the dashed lines represents $1\\sigma$ deviations from this relation. } \\label{fig:rates} \\end{figure}    \\begin{figure} \\begin{center} \\includegraphics[width=8cm]{asca_par_histo_custom.eps} \\end{center} \\caption[]{The source detection significances of the \\asca\\ Galactic plane survey are shown for the 3 energy bands investigated in that work \\citep[][]{sugizaki01}. \\swift/XRT non-detected sources are represented by the solid black line, the firmly detected sources are coloured red and tentative detections are indicated by hatched lines. These three different groups are overlaid in the graph (i.e., it is not cumulative). } \\label{fig:ascapar} \\end{figure}      \\subsection{The nature of non-detected sources: a comparison with the \\asca\\ results}\\label{subsec:nat_nondetect}  In Fig.~\\ref{fig:rates}, we compare the observed and expected XRT count rates for our sample sources. This figure shows that to within 1$\\sigma$ (dashed lines) only one source was brighter than predicted, while for 46\\% of our sample the count rates lay below expectations. Among these fainter than expected sources are many of the XRT non-detections, implying either transient or highly variable sources, significant Eddington bias effects, extended rather than point-like sources or spurious detections in the \\asca\\ surveys. The Eddington bias, which boosts the true count rate of sources at or near the detection limit \\citep[][]{eddington1913}, is likely to be at work in this sample as we are probing the faint end of the \\asca\\ catalogues. This idea is strengthened by the fact that we find the observed XRT count rates to be predominantly lower than expected.  Without knowing the intrinsic source population we cannot estimate the magnitude of this effect, but we note that work on the \\xmm\\ hard band survey (2--10 keV) has shown that for the \\xmm\\ slew survey, the true count rate distribution peaks a factor of 2 below the observed count rate for detections with low numbers of counts \\citep[][Warwick, R. et al., submitted to A\\&A]{starling2011}. Three of our XRT non-detected sources have count rate upper limits that are a factor 5--10 lower than predicted from the \\asca\\ observations (see Table~\\ref{tab:nondet}), which strongly suggests these are not faint, steady sources with count rates boosted by the Eddington bias, but are either variable (possibly transient) sources or spurious detections. Of these three, AX J1833.9--0822 was a firm 6$\\sigma$ detection above 2 keV with \\asca\\ and AX J1836.3--0647 was a 5$\\sigma$ detection in the 0.7--7 keV band dominated by its soft flux, while no similar information is available for AX J1751.1--2748. These detection significances disfavour a non-astrophysical, spurious nature. Various classes of X-ray sources (e.g., X-ray binaries, accreting white dwarfs, magnetars, stars, AGN) vary in intensity by a factor of a few. X-ray sources are generally denoted as transient when the variations are as large as $>2$ orders of magnitude. We thus tentatively identify AX J1751.1--2748, AX J1833.9--0822 and AX J1836.3--0647 as strongly variable or transient X-ray sources, although we cannot classify them further.  While it is not possible to classify the sources \\swift\\ has not detected, we may learn something of their nature from further examining the details of their \\asca\\ measurements. A total of 26 objects in our sample were drawn from the \\asca\\ Galactic plane survey, which were evaluated in 3 energy bands (0.7--2 keV, 2--10 keV and 0.7--7 keV) and assigned a detection significance \\citep[][]{sugizaki01}. We plot these values in Fig.~\\ref{fig:ascapar}, distinguishing between the XRT non-detections (solid black lines), detections (coloured red) and tentative detections (hatched lines). It can be seen that just 3 sources were significantly detected in the soft energy band (0.7--2 keV), and these were also detected both in other \\asca\\ bands and later with the XRT. This is unsurprising given the high absorbing columns in these directions (see Tables~\\ref{tab:spec} and \\ref{tab:nondet}). This causes the broad band (0.7--7 keV) to largely follow the results of the hard band (2--10 keV), where XRT-detected sources generally represent the higher-significance \\asca\\ detections while XRT non-detections are found below 5--6$\\sigma$. Our tentative XRT detections span a wider range of \\asca\\ detection significances, and are concentrated in the 0.7--7 keV band over the 4--5$\\sigma$ range.  It is of note that Galactic plane survey sources with $>$4$\\sigma$ detections in all 3 \\asca\\ energy bands number just two, AX J1538.3--5541 and AX J1651.0--4403, both of which are well detected with \\swift/XRT (see Sect.~\\ref{subsubsec:porb} for further details of AX J1538.3--5541). There are 3 hard-band only sources: 1 was tentatively detected with the XRT and the other 2 were not. Broad-band only detections generally have the fewest total counts and are therefore most likely to be spurious. We have two, AX J1835.1--0806 and AX J1857.3+0247. Both went undetected with the XRT adding further uncertainty to assignment of an astrophysical nature to these sources. We thus tentatively mark the two as spurious \\asca\\ detections. However, we note that their XRT count rate limits are not more than a factor of 2 below the expected rates. This is certainly possible within the Eddington bias effect and a large number of sources are variable at this level. Furthermore, our expected XRT count rates are overestimated if the sources are extended rather than point-like. As mentioned in Sect.~\\ref{subsec:hard}, AX J1857.3+0247 may possibly be related to an extended X-ray source.  Spectral analysis was attempted for those sources with \\asca\\ detections above 5$\\sigma$ in both \\citet{sugizaki01} and \\citet{sakano02}, over the full energy range of 0.7--10 keV and adopting an absorbed power law model, although only \\citet{sugizaki01} provide error estimates on the parameters. For the five firm XRT detections with well-measured \\asca\\ spectral parameters, $N_H$ and $\\Gamma$ are consistent within the 90\\% uncertainties with one exception: AX J1538.3--5541. For this source $N_H$ is consistent between the \\asca\\ and \\swift/XRT averaged spectra, while the photon index $\\Gamma$ is just consistent at the 3$\\sigma$ level possibly having been harder during the \\asca\\ observation. We found this source to be highly time variable from multiple XRT observations, and possible spectral evolution is investigated in Sect.~\\ref{subsubsec:porb}. The overall consistency between spectral measurements validates our use of the \\asca\\ spectra to make count rate predictions for XRT as used in Tables~\\ref{tab:spec} and~\\ref{tab:nondet}.   In conclusion, the 31\\% of sample sources firmly detected with {\\it Swift}/XRT were generally found with greater confidence in the \\asca\\ Galactic plane and Galactic centre surveys than the XRT non-detected sources. Very few sources were detected below 2 keV with \\asca, which is likely due to high foreground absorbing columns (see Tables~\\ref{tab:spec} and~\\ref{tab:nondet}). However, the greatest deviation from the expected count rate among the XRT non-detections corresponds to three sources that were largely well-detected with \\asca, strengthening the probability that they are astrophysical sources of variable or transient nature. We identify two sources that most likely have been spurious \\asca-detections given that these were weakly detected only in the full \\asca\\ energy band and not detected with XRT. For sources well-detected with both \\asca\\ and XRT we find their spectra to be unchanged within the given errors between the two epochs.   \\subsection{Summary}\\label{subsec:summary} Our sample of 35 unclassified \\asca\\ sources includes one confirmed and one candidate X-ray binary, and likely one magnetically accreting white dwarf. We assign these classifications based on their X-ray spectral and temporal properties. Furthermore, we used the improved XRT positions of the 16 sources detected in our sample to search for possible associations with catalogued sources at X-ray, UV, optical, infrared and radio wavelengths. This results in the identifications of possible counterparts for 8 of our targets, amongst which are 3 likely main sequence stars (see Table~\\ref{tab:class}). The probability of an optical/infrared chance detection is, however, considerable in these crowded fields ($\\sim50\\%$).  By assessing any possible long-term variability and the reported \\asca\\ properties of the 19 sources that were not detected by XRT, we identify two potentially spurious \\asca-detections amongst them and three X-ray sources that appear to be variable or transient (see Table~\\ref{tab:class}). A substantial fraction of the XRT non-detections are expected to be due to the Eddington bias, and might thus involve weak, persistent X-ray sources with intensities that lie near the detection limit of our \\swift/XRT observations.  With our improved $\\sim2-4''$ X-ray positions obtained for the 16 XRT-detected sources dedicated follow-up observations at different wavelengths become feasible. These have the potential to further classify these faint X-ray sources. As can be seen in Table~\\ref{tab:class}, our study shows that the unclassified \\asca\\ sources harbour a variety of astrophysical objects.    \\begin{table*} \\begin{threeparttable}[t] \\begin{center} \\caption[]{{Proposed source associations/classifications.}} \\begin{tabular}{l l l} \\toprule $\\#$ & Name & Comments   \\\\ \\midrule \\multicolumn{3}{l}{{\\bf Firm \\swift/XRT detections}} \\\\ 2 & AX J1504.6--5824 & Possible association with an optical/infrared object: main sequence star \\\\ 5 & AX J1538.3--5541 & Candidate X-ray binary, most likely an LMXB \\\\ 7 & \\polar & Candidate accreting magnetised white dwarf\\\\ 16 & AX J1738.4--2902 & Likely associated with 1RXS J173826.7--290140/2MASS J17382620--2901494: pre-main sequence star (K0V) \\\\ 19 & AX J1742.6--2901 & Possible association with 2RXP J174241.8--290215, optical/infrared/radio object \\\\ 22 & AX J1754.2--2754 & Confirmed neutron star LMXB, proposed ultra-compact X-ray binary \\\\ 29 & AX J1846.0--0231 & Possible association with an infrared object\\\\ \\midrule \\multicolumn{3}{l}{{\\bf Tentative \\swift/XRT detections}} \\\\ 11 & AX J1717.2--3718 & Possible association with an infrared object \\\\ 13 & AX J1720.8--3710 & Possibly related to 1RXS 172051.9--371033, UV/optical/infrared object: main sequence star \\\\ 17 & AX J1739.5--2730 & Possible association with an optical object \\\\ 30 & AX J1846.1--0239 & Possible association with an optical/infrared object \\\\ \\midrule \\multicolumn{3}{l}{{\\bf \\swift/XRT non-detections}} \\\\ 20 & AX J1751.1--2748 & Strongly variable/transient X-ray source \\\\ 25 & AX J1833.9--0822 & Strongly variable/transient X-ray source \\\\ 27 & AX J1835.1--0806 & Potential spurious \\asca\\ detection \\\\ 28 & AX J1836.3--0647 & Strongly variable/transient X-ray source \\\\ 34 & AX J1857.3+0247 & Potential spurious \\asca\\ detection \\\\ \\bottomrule \\end{tabular} \\label{tab:class} \\begin{tablenotes} \\item[]% \\end{tablenotes} \\end{center} \\end{threeparttable} \\end{table*}"
    },
    "1112/1112.3648_arXiv.txt": {
        "abstract": "{\\it Chandra} spectroscopy of transient stellar-mass black holes in outburst has clearly revealed accretion disk winds in soft, disk--dominated states, in apparent anti-correlation with relativistic jets in low/hard states.  These disk winds are observed to be highly ionized, dense, and to have typical velocities of $\\sim$1000~km/s or less projected along our line of sight.  Here, we present an analysis of two {\\it Chandra} High Energy Transmission Grating spectra of the Galactic black hole candidate IGR J17091$-$3624 and contemporaneous EVLA radio observations, obtained in 2011.  The second {\\it Chandra} observation reveals an absorption line at 6.91$\\pm$0.01~keV; associating this line with He-like Fe XXV requires a blue-shift of $9300^{+500}_{-400}$~km/s (0.03$c$, or the escape velocity at 1000~R$_{Schw}$).  This projected outflow velocity is an order of magnitude higher than has previously been observed in stellar-mass black holes, and is broadly consistent with some of the fastest winds detected in active galactic nuclei.  A potential feature at 7.32 keV, if due to Fe XXVI, would imply a velocity of $\\sim 14600$~km/s (0.05$c$), but this putative feature is marginal.  Photoionization modeling suggests that the accretion disk wind in IGR J17091$-$3624 may originate within 43,300 Schwarzschild radii of the black hole, and may be expelling more gas than accretes. The contemporaneous EVLA observations strongly indicate that jet activity was indeed quenched at the time of our {\\it Chandra} observations.  We discuss the results in the context of disk winds, jets, and basic accretion disk physics in accreting black hole systems. ",
        "introduction": "A detailed observational account of how black hole accretion disks drive winds and jets remains elusive, but the combination of high resolution X-ray spectroscopy, improved radio sensitivity, and comparisons across the black hole mass scale hold great potential. The broad range in X-ray luminosity covered by transient stellar-mass black holes makes it possible to trace major changes in the accretion flow as a function of the inferred mass accretion rate; this is largely impossible in supermassive black holes.  Disk winds and jets, for instance, appear to be state-dependent and mutually exclusive in sources such as H 1743$-$322 \\citep{Miller06, Blum10}, GRO J1655$-$40 \\citep{Miller08,Luketic10,Kallman09}, and GRS 1915$+$105 \\citep{Miller08, Neilsen09}.  This may offer insights into why many Seyfert AGN, which are well known for their disk winds, are typically radio--quiet \\citep[though not necessarily devoid of jets; see][]{King11, Jones11,Giroletti09}.  The proximity of Galactic black hole binaries (BHB) ensures a high flux level and spectra with excellent sensitivity in the Fe K band.  This is of prime importance because He-like Fe XXV and H-like Fe XXVI lines can endure in extremely hot, ionized gas \\citep[see, e.g.][]{Bautista01}, and therefore trace the wind region closest to where it is launched near the black hole. Studies of some stellar-mass black hole disk winds find that the gas is too ionized, too dense, and originates too close to the black hole to be expelled by radiative pressure or by thermal pressure from Compton heating of the disk, requiring magnetic pressure \\citep{Miller06, Miller06b, Kubota07}.  Winds that may originate close to the black hole and carry high mass fluxes are also observed in AGN \\citep[e.g.,][]{Kaspi02, Chartas02,King11b,Tombesi10}.  In this Letter, we present evidence of a particularly fast disk wind in the black hole candidate IGR J17091$-$3624. The current outburst of IGR J17091$-$3624 was first reported on 2011 January 28 \\citep{Krimm11}. Our observations caught IGR J17091$-$3624 in the high/soft state, but it is important to note that the source has also showed low/hard state episodes with flaring and apparent jet activity in radio bands  \\citep{Rodriguez11}. X-ray flux variations in IGR J17091$-$3264 bear similarities to the microquasar GRS 1915$+$105 \\citep[e.g.,][]{Altamirano11}. ",
        "conclusions": "At ionizations above $10^{3}$, radiation pressure is inefficient, and it is not able to drive these winds \\citep[e.g.,][]{Proga00b}.  Thus, although the UV components of disk winds in AGN are driven at least partially by radiation pressure, the wind in IGR J17091$-$3624 likely cannot be driven in this way.  A thermal wind can arise at radii greater than $0.2~R_{C}$ \\citep{Woods96}, where $R_{C} = (1.0\\times 10^{10})\\times (M_{BH}/M_{\\odot})/T_{C8}$, where $T_{C8}$ is the Compton temperature of the gas in units of $10^{8}$~K.  The spectrum observed in the second observation gives $R_{C} \\simeq 5\\times 10^{12}$~cm. Therefore, if we assume our conservative estimate of the launching radius, it is possible for IGRJ17091$-$3624 to have a thermally driven wind. However, if the wind originates closer to the black hole, then it is likely that magnetic processes -- either pressure from magnetic viscosity within the disk \\citep[e.g.,][]{Proga03} or magneto-centrifugal acceleration \\citep[e.g.,][]{Blandford82} -- must play a role in launching the wind observed in IGR J17091$-$3624.  Fast X-ray disk winds are not only seen in BHB like IGR J17091$-$3624, but also in AGN and quasars \\citep[e.g.,][]{King11b, Chartas02}. The fastest UV winds observed in AGN are pushed to high velocities by radiation pressure.  It remains to be seen whether a common driving mechanism works across the black hole mass scale to drive fast, highly ionized X-ray disk winds.  \\cite{Chartas02} show that in the quasar APM 08279+5255 there are broad absorption features, which are likely highly relativistic Fe XXV and/or Fe XXVI lines. In these regards, it bears some similarities to the most extreme winds in BHB's.  Observations of BHB point to an anti-correlation of wind and jet outflows from accretion disks \\citep{Miller06,Miller08, Blum10, Neilsen10}.  Winds appear to only be detected, or at least are considerably stronger, in soft, disk--dominated states, and absent in hard states where compact, steady jets are ubiquitous \\citep{Fender06}.  In H 1743$-$322, in particular, there is evidence that the absence of winds in hard states is {\\it not} an artifact of high ionization hindering the detection of absorption lines, but instead represents a real change in the column density (and thus the mass outflow rate) in any wind \\citep{Blum10}.  It appears that our coordinated {\\it Chandra} and EVLA observations of IGR J17091$-$3624 support this anti-correlation.  The EVLA observations place very tight limits on the radio flux when the disk wind is detected, orders of magnitude below the level at which IGR J17091$-$3624 was detected in radio during its low/hard state only a few months prior \\citep{Rodriguez11}.  \\cite{Neilsen09} suggested that the production of winds may be responsible for quenching jets in GRS 1915$+$105.  It might then be the case that jets should be observed whenever winds are absent.  In our first observation of IGR J17091$-$3624, however, neither a wind nor a jet is detected, with tight limits.  Instead, the apparent dichotomy between winds and jets may signal the magnetic field topology in and above the disk is state-dependent.  This is broadly consistent with multi-wavelength studies suggesting synchrotron flares above the disk, but only in the hard state \\citep[e.g. GX 339$-$4, XTE J1118$+$480,][]{DiMatteo99, Gandhi10}.  It is interesting to speculate that the magnetic field might be primarily toroidal in the soft state, where a Shakura-Sunyaev disk is dominant, but primarily poloidal in the hard state, when the mass accretion rate is lower \\citep[e.g.,][]{Beckwith08}. The type of outflow that is observed may also depend greatly on how much mass is loaded onto magnetic field lines; that could depend on variables including the mass accretion rate through the disk.  \\vspace{0.2in} We would like to thank the anonymous referee. We thank Michael Nowak for his instrumental help as well. ALK gratefully acknowledges support through the NASA Earth and Space Sciences Fellowship.  JMM gratefully acknowledges support through the {\\it Chandra} Guest Observer program. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.  \\begin{deluxetable}{l l l l l} \\tablecolumns{5} \\tablewidth{0pc} \\tabletypesize{\\scriptsize} \\tablecaption{Spectral Modeling Parameters of the 2$^{nd}$ HEG observation} \\tablehead{Parameter  & Model 1 & Model 2 & Model 3 & Model 4 \\\\ \\hline & diskbb + po & (diskbb + po) & comptt & (comptt) \\\\ & + Gauss + Gauss & $\\times$ Xstar $\\times$ Xstar  & + Gauss + Gauss  &  $\\times$ Xstar $\\times$ Xstar } \\startdata N$_H$ (10$^{22}$ cm$^{-2}$) & 1.14$\\pm0.06$ & 1.13$\\pm0.06$ & 0.475 $^{+0.017}_{-0.018}$ & 0.558$^{+0.025}_{-0.028}$  \\\\ -\\\\ T$_{in}$ (keV) & 1.53 $\\pm0.09$ & 1.51 $^{+0.11}_{-0.09}$& - & -\\\\ Norm & 13.1$^{+3.6}_{-2.5}$ & 13.8$^{+4.0}_{-1.7}$ & - &- \\\\ $\\Gamma$ & 1.93 $^{+0.15}_{-0.16}$ & 1.91 $\\pm0.17$& - & -\\\\ Norm &0.35 $^{+0.07}_{-0.08}$ & 0.34 $\\pm0.08$ & - & -\\\\ -\\\\ T$_0$ (keV) & - &- & 0.58$\\pm0.01$ & 0.59 $\\pm0.01$  \\\\ kT (keV) & -& -& 10.5$^{+30}_{-1.7}$ & 9.8 $\\pm0.02$ \\\\ $\\tau_{plasma}$ & - & - & 2.24$\\pm0.01$ &2.28 $\\pm0.01$ \\\\ Norm & -&-& 0.0584$\\pm0.0001$ & 0.062$^{+0.002}_{-0.05}$ \\\\ -\\\\ E$_{Fe XXV}$ (keV) & 6.91$\\pm0.01$  & -& 6.91$^{+0.02}_{-0.01}$&-\\\\ FWHM (keV) & 0.091$^{+0.022}_{-0.049}$ & - & 0.13$^{+0.19}_{-0.04}$ &-\\\\ EW (keV)& 0.021$^{+0.005}_{-0.002}$ &  -&$0.040^{+0.007}_{-0.009}$&-\\\\ Norm ($10^{-4}$) & 3.5 $^{+0.8}_{-0.6} $& -&6.0$^{+1.1}_{-1.3}$ &-\\\\ v (km/s) & 9300$^{+500}_{-400}$ & -& 9300$^{+400}_{-800}$&-\\\\ -\\\\ E$_{Fe XXVI}$ (keV) &7.32$^{+0.02}_{-0.06}$& -& 7.30$\\pm$0.02 &-\\\\ FWHM (keV) & 0.081$^{+0.079}_{-0.027}$ & - & 0.25 $^{+0.13}_{-0.01}$&-\\\\ EW (keV) & 0.032 $^{+0.018}_{-0.004}$&  -&$0.089^{+0.013}_{-0.014}$&-\\\\ Norm ($10^{-4}$) & 3.4$^{+1.9}_{-0.4}$ & -&11.8$^{+1.7}_{-1.6}$ &-\\\\ v (km/s) &14600$^{+2500}_{-800}$& -& 13800$\\pm800$&-\\\\ -\\\\ N (10$^{22}$ cm$^{-2}$) & - & 0.47$^{+0.17}_{-0.19}$ & - & 0.45$^{+0.33}_{-0.17}$ \\\\ $\\log \\xi$(ergs cm s$^{-1}$) & - &  3.3 $^{+0.2}_{-0.1}$ &-&3.4 $^{+0.2}_{-0.1}$ \\\\ v (km/s) & - & 9600$^{+400}_{-500}$ & - & 9600$\\pm300$ \\\\ -\\\\ N (10$^{22}$ cm$^{-2}$) & - & 1.66 $^{+1.18}_{-0.83}$ & - & 1.97$^{+1.26}_{-0.51}$ \\\\ $\\log \\xi$(ergs cm s$^{-1}$) & - &  3.9$^{+0.5}_{-0.3}$ &-&3.7 $^{+0.3}_{-0.1}$ \\\\ v (km/s) & - & 15400$\\pm400$& - & 15400 $^{+400}_{-300}$\\\\ -\\\\ $\\chi^2/\\nu$ & 2725/3408 = 0.80 & 2731/3408 = 0.80 & 2793/3408 = 0.82 & 2761/3408 = 0.81 \\\\ \\enddata \\label{lines} \\tablecomments{This Table lists the line detections using Gaussian functions as well as more self-consistent, photoionization components created with XSTAR, assuming two different continuum models. {\\it TBabs} is applied to all the models and the errors are 1$\\sigma$ confidence level.} \\end{deluxetable}     \\clearpage"
    },
    "1112/1112.3681_arXiv.txt": {
        "abstract": "{Dark matter particles may be captured by a star and then thermalized in the star's core. At the end of its life a massive star collapses suddenly and a compact object is formed. The dark matter particles redistribute accordingly. In the inelastic dark matter model, an extended dense dark matter mini-halo surrounding the neutron star may be formed. Such mini-halos may be common in the Galaxy. The electron/positron flux resulting in the annihilation of dark matter particles, however, is unable to give rise to observable signal unless a nascent mini-halo  is within a distance $\\sim ~{\\rm a~few}~0.1$ pc from the Earth.} ",
        "introduction": "Dark matter (DM) is a form of matter necessary to account for gravitational effects observed in very large scale structures such as anomalies in the rotation of galaxies and the gravitational lensing of light by galaxy clusters that cannot be accounted for by the quantity of observed matter \\cite[e.g.,][]{lsp,lkp}. The existence of dark matter is a piece of solid evidence for new physics that is beyond the standard model and the relevant researches are the focus of modern physics and astrophysics. Among the hypothetical particles proposed so far, the weakly interacting massive particles (WIMPs) are the leading candidate \\cite{lsp,lkp}. The direct-detection searches look for the signal of WIMP-nuclei scattering in underground detectors. The indirect-detection experiments, instead, aim to catch the annihilation/decay products of DM particles. Currently there are some anomalous signals in favor of WIMPs but the evidence is not conclusive \\cite{atic,pamela,fermi,fan}.  The DM particles may be captured by the stars. The capture of WIMPs by the sun, earth and other astrophysical objects has been extensively investigated \\cite{gould1,gould2,press85}. The captured particles will eventually be thermalized and centered in the core of the star, namely, a DM core forms in the center of a star \\cite{peter47}. After the death of a (massive) star, a compact object is formed. The collapse of the star lasts for a very short time and the density distribution of the DM core will be modified. If the compact object is small enough, a significant part of the DM core may be left outside the compact object, i.e., a DM mini-halo is formed. As the outer shell of the stellar remnant eventually disappears, the annihilation products of the DM matter, such as electrons and photons, are visible for the observers. In this work we investigate the formation and evolution of such a mini-halo and calculate the flux of the annihilated electrons/positrons. This work is structured as follows. In section 2 we describe the distribution of WIMPs inside the star. In section 3 we investigate the redistribution of WIMPs in the collapse process. In section 4 we calculate the evolution of the DM mini-halo and in section 5 we estimate the electron/positron flux detectable on the earth. We summarize our results in section 6. ",
        "conclusions": "In this work we have investigated the formation, evolution and the annihilation of the particles in dark matter mini-halo surrounding the compact object. In our model, the dark matter particles are captured by the stellar matter and then distributed in the core of the star. At the end of its life the star uses up the fuel and a compact object, either a white dwarf, a neutron star or a black hole,  is formed. The dark matter particles will redistribute around the newly formed compact object. In the elastic dark matter model and if the elastic cross section is large enough in the inelastic dark matter model, the dark matter particles will be thermalized and no prominent dark matter mini-halo can be formed. In the inelastic dark matter model, however, if the elastic cross section is very small (for example $\\lesssim 10^{-47}~{\\rm cm^{2}}$) an extended dense dark matter mini-halo surrounding neutron star may be formed. Though such mini-halos may be very common in the Galaxy, the electron/positron flux resulting in the annihilation of these dark matter particles is negligible compared with the observational data unless a nascent mini-halo is within $\\sim$ a few $0.1~{\\rm pc}$ to the Earth."
    },
    "1112/1112.4774_arXiv.txt": {
        "abstract": "The Hubble radius is a particular manifestation of the Universe's gravitational horizon, $R_{\\rm h}(t_0) \\equiv c/H_0$, the distance beyond which physical processes remain unobservable to us at the present epoch. Based on recent observations of the cosmic microwave background (CMB) with WMAP, and ground-based and HST searches for Type Ia supernovae, we now know that $R_{\\rm h}(t_0)\\sim$$13.5$ Glyr. This coincides with the maximum distance ($ct_0\\approx 13.7$ Glyr) light {\\it could} have traveled since the big bang. However, the physical meaning of $R_{\\rm h}$ is still not universally understood or accepted, though the minimalist view holds that it is merely the proper distance at which the rate of cosmic recession reaches the speed of light $c$. Even so, it is sometimes argued that we can see light from sources beyond $R_{\\rm h}$, the claim being that $R_{\\rm h}$ lies at a redshift of only $\\sim$$2$, whereas the CMB was produced at a much greater redshift ($\\sim$$1100$). In this paper, we build on recent developments with the gravitational radius by actually calculating null geodesics for a broad range of FRW cosmologies, to show---at least in the specific cases we consider here, including $\\Lambda$CDM---that no photon trajectories reaching us today could have ever crossed $R_{\\rm h}(t_0)$. We therefore confirm that the current Hubble radius, contrary to a commonly held misconception, is indeed the limit to our observability. We find that the size of the visible universe in $\\Lambda$CDM, measured as a proper distance, is approximately $0.45ct_0$. ",
        "introduction": "The standard model of cosmology is confronted with several unpalatable coincidences, suggesting that we do not yet have a fully consistent picture of the Universe's dynamical expansion (see, e.g., Melia \\& Shevchuk 2011). Part of the problem is that cosmological observations can only be interpreted from within the context of a pre-assumed model, and the data can be quite compliant.  The $\\Lambda$CDM (Cold Dark Matter with a cosmological constant $\\Lambda$) model has been without peer in cosmology (see, e.g., Spergel et al. 2003, and Tegmark et al. 2004). For example, this model has been used with complementary measurements of the cosmic microwave background (CMB) radiation to determine that the Universe is flat, so its energy density $\\rho$ is at (or very near) the ``critical\" density $\\rho_{\\rm c}\\equiv 3c^2H^2/ 8\\pi G$. But among the many peculiarities of this description of the universe is the inference, based on current observations, that the density $\\rho_{\\rm d}$ of dark energy must itself be of order $\\rho_{\\rm c}$. (In the context of $\\Lambda$CDM, the best fit to the WMAP data indicates that dark energy represents approximately $73\\%$ of the total $\\rho\\approx \\rho_{\\rm c}$; see Spergel et al. 2003.) Dark energy is often thought to be the manifestation of the aforementioned cosmological constant, $\\Lambda$, though no reasonable explanation has yet been offered as to why such a fixed, universal density ought to exist at this scale. It is well known that if $\\Lambda$ is associated with the energy of the vacuum in quantum theory, it should have a scale representative of phase transitions in the early Universe---120 orders of magnitude greater than $\\rho_{\\rm c}$.  Many workers have attempted to circumvent these difficulties by proposing alternative forms of dark energy, including Quintessence (Ratra \\& Peebles 1988; Wetterich 1988), which represents an evolving canonical scalar field with an inflation-inducing potential, a Chameleon field (see, e.g., Mota \\& Barrow 2004; Khoury \\& Weltman 2004; Brax et al. 2004) in which the scalar field couples to the baryon energy density and varies from solar system to cosmological scales, and modified gravity, arising out of both string motivated, or General Relativity modified actions (Capozziello et al. 2003; Nojiri \\& Odintsov 2003; Carroll et al. 2004), which introduce large length scale corrections modifying the late time evolution of the Universe. The actual number of suggested remedies is far greater than this small, illustrative sample.  An equally perplexing puzzle with $\\Lambda$CDM has been dubbed the ``coincidence problem,\" arising from the peculiar near-simultaneous convergence of the matter energy density $\\rho_{\\rm m}$ and $\\rho_{\\rm d}$ towards $\\rho_{\\rm c}$ in the present epoch. Though $\\rho_{\\rm m}$ and $\\rho_{\\rm d}$ are expected to change at different rates as the Universe expands (particularly if dark energy is a cosmological constant) they are nearly equal in the present epoch, implying that we live at a special time in cosmic history.  In a recent paper (Melia \\& Shevchuk 2011), we proposed an explanation for yet another disturbing coincidence, having to do with the apparent equality of our gravitational (or Hubble) radius $R_{\\rm h}$ with the distance $ct_0$ light could have traveled since the big bang (in terms of the presumed current age $t_0$ of the Universe). This equality has received some scrutiny in recent years (Melia 2003, 2007, 2009, Melia \\& Abdelqader 2009, van Oirschot et al. 2010; see also Lima 2007 for a related, though unpublished, work).  Unfortunately, there is still some confusion regarding the properties of $R_{\\rm h}$ due to a misunderstanding of the role it plays in our observations. For example, it is sometimes suggested  (see, e.g., Davis \\& Lineweaver 2004; van Oirschot et al. 2010) that sources beyond $R_{\\rm h}(t_0)$ are observable today, which is certainly not the case. We will therefore begin by elaborating upon what the gravitational radius $R_{\\rm h}$ is---and what it is not. Though first defined in Melia (2007), an unrecognized form of $R_{\\rm h}$ actually appeared in de Sitter's (1917) own account of his spacetime metric. And we will advance the discussion further by actually calculating photon trajectories for various well-studied Friedmann-Robertson-Walker (FRW) Cosmologies, demonstrating that the null geodesics reaching us at $t_0$ have never crossed $R_{\\rm h}(t_0)$. Some come close, and in one case---the de Sitter model---they approach $R_{\\rm h}$ asymptotically as $t$ recedes to our infinite past. Our conclusion in this paper will be that $R_{\\rm h}(t_0)$ is a real limit to our observability at the present time $t_0$. ",
        "conclusions": "Throughout this paper, we have made a conscious effort to discuss the properties of null geodesics in FRW cosmologies without resorting to conformal diagrams. This approach, also used to great effect by Ellis and Rothman (1993), makes it easier to think in terms of familiar quantities (proper distances and proper time) that are not always straightforward to interpret otherwise. Misconceptions often arise from the misinterpretation of coordinate-dependent effects. In their paper, Ellis and Rothman clearly delineated true horizons from apparent horizons, and extended the definitions, first introduced by Rindler (1956), in a clear and pedagogical manner. In this paper, we have paid particular attention to the gravitational horizon, also manifested as the Hubble radius, which is time-dependent and may or may not turn into an event horizon in the asymptotic future, depending on the equation of state.  It is useful at this point to be absolutely clear about how far light has traveled in reaching us. It is quite evident from our results that sources whose light we see today were at a proper distance $R(t_e)\\rightarrow 0$ when $t_e\\rightarrow 0$. For the FRW cosmologies we have considered here (which include $\\Lambda$CDM) the light reaching us today---including that from the recombination region associated with the CMB---has traveled a {\\it net} proper distance of at most $\\sim$0.3--$0.4ct_0$. It is therefore not correct to claim that the size of the visible universe in these cosmologies is $ct_0$ (or even greater in some interpretations). Because all causally connected sources in an expanding universe began in a vanishingly small volume as $t\\rightarrow 0$, the maximum proper distance from which we receive light today must necessarily be less than $ct_0$, since presumably there were no pre-existing sources at a non-zero proper distance prior to $t=0$ with which we were in causal contact.  It is true, however, that more and more sources become visible to us as time advances, since for $t>t_0$, the geodesic curves terminating in our future all rise above their current counterparts shown in figures~1,2,3,4, and 5d. In our future, we will see light from sources that radiated at proper distances greater than those shown here. Of course, $R_{\\rm h}(t)$ will also continue to increase, and it is not difficult to convince oneself from Equations~(5) and (8) that the limits of observability will always be $R_{\\rm h}(t)$, since $R_\\gamma(t_e)/R_{\\rm h}(t)$ is smaller than $1$ for all $t_e<t$. To see this, let us first find the emission time $t_{e,\\; {\\rm max}}$ at which $R_\\gamma(t_e)$ attains its maximum value.  Clearly, this happens when $\\dot{R_\\gamma}=0$, which means that $R_\\gamma(t_{e,\\;{\\rm max}}) =R_{\\rm h}(t_{e,\\;{\\rm max}})$. But for all $w>-1$, Equation~(5) shows that $\\dot{R_{\\rm h}}>0$, and therefore $R_{\\rm h}(t)>R_{\\rm h}(t_{e,\\;{\\rm max}})$ for $t>t_{e,\\;{\\rm max}}$, which also means that $R_{\\rm h}(t)> R_\\gamma(t_{e,\\;{\\rm max}})$ for all $t>t_{e,\\;{\\rm max}}$."
    },
    "1112/1112.2708_arXiv.txt": {
        "abstract": "%  Kuiper Belt object 2007 TY430 is the first wide, equal-sized, binary known in the 3:2 mean motion resonance with Neptune. The two components have a maximum separation of about 1 arcsecond and are on average less than 0.1 magnitudes different in apparent magnitude with identical ultra-red colors ($g-i=1.49\\pm0.01$ mags).  Using nearly monthly observations of 2007 TY430 from 2007-2011, the orbit of the mutual components was found to have a period of $961.2\\pm4.6$ days with a semi-major axis of $21000\\pm160$ km and eccentricity of $0.1529\\pm0.0028$.  The inclination with respect to the ecliptic is $15.68\\pm0.22$ degrees and extensive observations have allowed the mirror orbit to be eliminated as a possibility.  The total mass for the binary system was found to be $7.90\\pm0.21 \\times10^{17}$ kg. Equal-sized, wide binaries and ultra-red colors are common in the low inclination ``cold'' classical part of the Kuiper Belt and likely formed through some sort of three body interactions within a much denser Kuiper Belt.  To date 2007 TY430 is the only ultra-red, equal-sized binary known outside of the classical Kuiper belt population.  Numerical simulations suggest 2007 TY430 is moderately unstable in the outer part of the 3:2 resonance and thus 2007 TY430 is likely an escaped ``cold'' classical object that later got trapped in the 3:2 resonance.  Similar to the known equal-sized, wide binaries in the cold classical population, the binary 2007 TY430 requires a high albedo and very low density structure to obtain the total mass found for the pair.  For a realistic minimum density of 0.5 g/cm$^{3}$ the albedo of 2007 TY430 would be greater than 0.17.  For reasonable densities, the radii of either component should be less than 60 km, and thus the relatively low eccentricity of the binary is interesting since no tides should be operating on the bodies at their large distances from each other.  The low prograde inclination of the binary also makes it unlikely the Kozai mechanism could have altered the orbit, making the 2007 TY430 binary orbit likely one of the few relatively unaltered primordial binary orbits known.  Under some binary formation models, the low inclination prograde orbit of the 2007 TY430 binary indicates formation within a relatively high velocity regime in the Kuiper Belt. ",
        "introduction": "Small solar system bodies, such as asteroids and Kuiper Belt Objects (KBOs), may be leftover remnants of the planetesimals that went into the formation of the terrestrial and giant planets.  These small bodies' orbital and physical properties give insight into the Solar System's origins and evolution.  KBO binaries are particularly informative about the conditions at the time when they formed, which must have been at an earlier epoch of Solar System history.  The Kuiper Belt is dynamically structured with three main dynamical classes (Figures~\\ref{fig:kboeabinary} to~\\ref{fig:kboiabinaryblowup}).  1) Classical KBOs have semi-major axes between about 40 and 50 AU with moderate eccentricities and inclinations.  These objects may be regarded as the population originally predicted for the Kuiper Belt (Fernandez and Ip 1981), but they have higher eccentricities and inclinations than expected (Jewitt et al. 1998).  The dynamics of the classical KBOs have shown that the outer Solar System has been highly modified through the evolution of the planets (Hahn and Malhotra 2005).  There appears to be two subsets of classical KBOs.  The low inclination ``cold'' classical population has generally smaller and redder objects compared to the high inclination ``hot'' classical population (Tegler and Romanishin 2000; Brown 2001; Levison and Stern 2001; Trujillo and Brown 2002; Peixinho et al. 2008).  These two classical populations likely represent different formation regions from the outer solar system.  Through migration and scattering of the planets very early in the Solar System the two classical populations came to reside where we see them now (Gomes 2003; Levison and Morbidelli 2003; Batygin et al. 2011). 2) Scattered disk objects have large eccentricities with perihelia near the orbit of Neptune ($q \\sim 30-45$ AU).  The scattered disk objects are likely to have been moved to their current orbits through interactions with Neptune (Gomes et al. 2008).   3) Resonant KBOs are in mean motion resonances with Neptune and generally have higher eccentricities and inclinations than classical KBOs.  Most of the resonance objects were likely captured into their resonances from the outward migration and circularization of Neptune's orbit (Malhotra 1995; Levison et al. 2008).  Where the resonance objects came from and how and when this capture occurred is still highly debated.  The 3:2 resonance, called Plutinos since Pluto is in this resonance, appears to be the most populated resonance (Jewitt et al. 1998; Chiang and Jordan 2002; Sheppard et al. 2011).  The other resonances with a significant number of known objects are the outer mean motion resonances 5:3, 7:4, 2:1 and 5:2 (Elliot et al. 2005; Hahn and Malhotra 2005; Gladman et al. 2008).  Binary small bodies are a particularly good way to obtain clues about the Solar System's past.  Remarkably, the small bodies in the various Solar System reservoirs appear to have significantly different binary characteristics (Richardson and Walsh 2006; Noll et al. 2008a).  In the main asteroid belt, most secondaries are close to and significantly smaller than the primaries (Merline et al. 2002; Margot 2002; Marchis et al. 2006b).  This is indicative of formation through direct collisions.  The few binaries known in the Jupiter Trojan population hint at these objects being of very low density, implying they may be similar to comets and Trans-Neptunian Objects (TNOs) (Marchis et al. 2006a).  Starting with the discovery of the first KBO binary other than Pluto, it has become apparent that most of the low inclination classical Kuiper Belt binaries have large separations and similar sized components (Veillet et al. 2002; Noll et al. 2008a). Because of the large angular momentum in these equal-sized, wide binary systems, they cannot be formed by direct collisions.  Direct formation by gravitational collapse, the dynamical friction of a sea of KBOs, three body interactions, some other collisionless interactions, or a combination of these are the best viable scenarios for binary formation in the classical Kuiper Belt (Nesvorn{\\'y} et al. 2010; Goldreich et al. 2002; Weidenschilling 2002; Funato et al. 2004; Astakhov et al. 2005; Lee 2007; Schlichting and Sari 2008a; Gamboa Suarez et al. 2010).  The cold classical disk also has a high fraction of binaries, irrespective of separation and difference in magnitude of the individual components, compared to the hot classical, resonant and scattered populations, about 30 percent versus 5 percent binary fraction, respectively (Stephens and Noll 2006; Noll et al. 2008b).  Unlike the low inclination classical KBOs, many of the other known KBO binaries do not have such equal-sized and wide components.  The different types of binary formation mechanisms proposed will each tend to create a particular subset of satellite semi-major axes, eccentricities, inclinations and secondary diameters (Kern and Elliot 2006a; Noll et al. 2008a).  Thus, understanding the orbital elements and physical characteristics of the components of a binary is essential to constraining how the binary may have formed. Understanding the formation mechanism of binaries gives insight into the original solar nebula and the collisional environment in the distant past.  The two distinct binary populations seen in the cold and hot dynamical classes are strong evidence of a different history. For example, Parker and Kavelaars (2010) use the fragile nature of the wide binaries in the cold classical region as evidence to show that cold classical KBOs were not likely to have been implanted dynamically, as suggested by Levison et al. (2008). Binaries also likely played an important role in planet formation by acting as a heat source in the planetesimal disk, tranforming gravitational potential energy into kinetic energy of planetesimals in the disk (Perets 2011).  2007 TY430 is a unique binary KBO because it is the first known equal-sized, wide binary in the 3:2 mean motion resonance with Neptune, which is the innermost well populated Neptune resonance.  The other known 3:2 resonance binary objects (including Pluto, Orcus, (47171) 1999 TC36 and (208996) 2003 AZ84) have significantly smaller and relatively closer secondaries (magnitude differences greater than 2 mags) that are more indicative of formation through direct collisions (Stern 2002; Brown et al. 2006; Canup 2011).  Because 2007 TY430 indicates a different binary formation mechanism operating within the 3:2 resonance population, 2007 TY430 was extensively observed over the past several years in order to determine the orbital and physical properties of its two components. ",
        "conclusions": "The ultra-red color and wide equal-sized binary nature of 2007 TY430 makes it very similar to the unique characteristics of the low inclination classical Kuiper Belt.  This strongly favors the argument that 2007 TY430 originated as a typical low inclination ``cold'' classical Kuiper Belt object (assuming that the cold classical population forms in situ).  Weak chaos over long timescales (or stronger chaos on faster timescales) slowly removed 2007 TY430 from the cold classical region into a dynamically excited orbit in the scattered disk (see Volk and Malhotra 2011).  Eventually 2007 TY430 became trapped in the unstable outskirts of the 3:2 resonance where it is seen today. During this process, the binary remained stable.  This would suggest that 2007 TY430's formation environment is likely the cause of its ultra-red color and wide binary nature and not its current orbital characteristics. This bolsters the argument of Benecchi et al. (2009) who suggest that similar colors for binary components indicate common formation conditions, regardless of present location.  \\subsection{Stability of 2007 TY430}  The Hill sphere of an object orbiting the Sun is the region where satellites are generally stable around the object, given by:  \\begin{equation} r_{H} = a_{p} \\left[\\frac{(m_{p1}+m_{p2})}{3M_{\\odot}}\\right]^{1/3} \\label{eq:hill} \\end{equation}  where $a_{p}$ is heliocentric distance of the object in AU, $M_{\\odot}$ is the mass of the Sun and $m_{p1}$ and $m_{p2}$ are the mass of the first and second components of the binary, respectively. The Hill radius for 2007 TY430 is about $3.015\\times 10^{5}$ km.  Using the heliocentric semi-major axis of 2007 TY430 from Table 1 and a total mass of $7.90\\pm0.21 \\times10^{17}$ kg for the binary, gives $a_{bin}/r_{H}$ of about 0.07 for 2007 TY430.  While this is still well within the stable region of orbital phase space, it is among the widest known binaries (Nesvorn{\\'y} et al. 2003, Nicholson et al. 2008, Noll et al 2008a).  Widely separated binaries were probably much more common in the past as they are unstable over the age of the solar system from dynamical perturbations (Petit and Mousis 2004; Noll et al. 2006, 2008a; Petit et al. 2008; Parker and Kavelaars 2010).  Its possible the cold classical KBOs have had lower numbers of collisions or fewer planet interactions allowing the largest population of surviving wide, equal-sized binaries in the Kuiper belt.  2007 TY430 has similar separation and component sizes as the cold classical KBOs 1998 WW31 and 2000 CF105 (Table 5).  Petit and Mousis (2004) determined these two binary KBOs were only stable for about 1-2 Gyrs.  The Petit and Mousis (2004) result is likely a lower limit on the binary lifetimes since they used a steeper size frequency distribution for the smaller Kuiper belt objects than currently believed to exist, allowing for more small impactors to disrupt the binaries than is currently likely (Fraser et al. 2008; Fuentes and Holman 2008; Fuentes et al. 2009; Fraser 2009).  The 2007 TY430 binary probably has a similar timescale of stability, though 2007 TY430 is in the Plutino population and not the low inclination classical belt.  The timescales of destruction for wide binary TNOs by Parker and Kavelaars (2010) also suggest 2007 TY430 would not have a high probability for surviving for the age of the Solar System, and thus 2007 TY430 could just be one of a few binaries left of a once much larger population of binaries.  \\subsubsection{2007 TY430: The Case for a Primordial Binary Orbit}  The inclinations of many TNO binary components are generally high (Naoz et al. 2010).  The Kuiper Belt binary 2001 QW322 was found to have a fairly low eccentricity ($e<0.4$) yet a very large inclination ($i\\sim125$ degrees) and semi-major axis ($\\sim 120000$ km) (Petit et al. 2008).  Because of 2001 QW322's large inclination, its relatively low eccentricity and large semi-major axis can be explained by the Kozai mechanism (Perets and Naoz 2009).  Most binary minor planets do not appear to have primordial orbits as they could have been affected by either tidal forces or have high eccentricities or high inclinations at which the Kozai mechanism may operate (Ragozzine 2009, Naoz et al. 2010).  The binary orbit of 2007 TY430 is compared to other known equal-sized binary orbits in Figures ~\\ref{fig:binarymutualrhi} and ~\\ref{fig:binarymutualrhe}.  2007 TY430 does not have orbital elements similar to other TNO binaries shown in a similar Figure 3 of Naoz et al. (2010), which compares the normalized separations of the binary components to their eccentricities.  The $log(a_{bin}/r)$ for 2007 TY430 is slightly greater than 2.5, but still in a stable region of the Naoz et al. Figure 3.  Interestingly, when including recent results from Parker et al. (2011), it appears most of the equal-sized binaries with very large semi-major axes appear to have moderate to low eccentricities and low inclinations, making them unsusceptible to the Kozai resonance and tides.  This is in contrast to many of the lower semi-major axis, equal-sized binaries that appear susceptible to the Kozai mechanism because of their large mutual inclinations and higher eccentricities.  This result is likely because the large semi-major axis objects would become unstable to perturbations if their eccentricities or inclinations were too large.  2007 TY430 is the lowest eccentricity, wide component, equal-sized binary known.  Because the mutual orbital elements of the 2007 TY430 binary has low inclination, low eccentricity and large semi-major axis (Figures~\\ref{fig:binarymutualrhi} and ~\\ref{fig:binarymutualrhe}), the Kozai mechanism and tidal interactions have not likely modified the primordial orbit of 2007 TY430 (Chyba et al. 1989; Murray and Dermott 1999; Perets and Naoz 2009).  Modification of the 2007 TY430 binary could have occurred over the age of the solar system from direct collisions on either of the two components, from relatively massive bodies passing within the Hill sphere of the two components or interactions with the giant planets (Petit and Mousis 2004; Nesvorn{\\'y} et al. 2011).  The high eccentricities and inclinations of other TNO binary components suggest formation in a dense collisional environment.  In this environment, gravitational encounters can create these high eccentricities and inclinations along with secular Kozai effects and tidal evolution (Naoz et al. 2010).  Other binaries observed to date could have formed in a similar manner as 2007 TY430, but later three body encounters or direct collisions after binary formation changed the binary components orbits to be highly inclined and eccentric.  It is thus possible that 2007 TY430 simply has escaped any significant collisions or gravitational encounters while other known binaries have not or formed in a different collisional environment.  This could be because 2007 TY430 diffused out of the low inclination classical belt early on and the Plutino orbit made it less susceptible to collisions. Thus 2007 TY430's binary orbit could be more primordial than many of the other known equal-sized wide binary objects observed to date.  \\subsection{Formation of the 2007 TY430 binary system}  The proposed binary formation by Funato et al. (2004) from exchanged binaries is unlikely since TNOs appear to lack the many high-eccentricity binaries suggested by this mechanism (Naoz et al. 2010), of which 2007 TY430 is another example of a low eccentricity binary.  Weidenschilling (2002) proposed a hybrid mechanism of a direct collision between two bodies while in the Hill Sphere of a third body, which can directly form large separation binaries.  This mechanism would only be likely to occur during the formation of the Kuiper Belt, as many more large bodies than currently observed are required as well as low velocities are needed.  A similar mechanism, the L$^{3}$ mechanism proposed by Goldreich et al. (2002), has a third body strongly interacting with two other bodies while they each are within the other body's Hill sphere.  The L$^{3}$ mechanism should dominate over the Weidenschilling (2002) mechanism because strong gravitational interactions should occur much more frequently than the actual collisions of objects.  Chaos assisted capture (Astakhov et al. 2005, Lee et al. 2007), is effectively the same as the L$^{3}$ mechanism; in chaos assisted capture, two bodies are temporarily trapped in their mutual Hill Spheres and can become permanently trapped when a third 'intruder' body is scattered by the pair.  This mechanism would create wide separation binaries with equal-size and moderate eccentricity.  A second binary formation mechanism proposed by Goldreich et al. (2002), the L$^{2}$s mechanism, involves two objects forming an unbound transient binary that becomes bound with the aid of the dynamical friction from a sea of small objects.  Any binary formation that involves dissipation of energy in a smooth and gradual manner, like the L$^{2}$s mechanism, will likely form retrograde binaries (Schlichting and Sari 2008b).  Because of 2007 TY430's prograde orbit, it is unlikely to have formed by the L$^{2}$s mechanism.  Unlike the L$^{2}$s mechanism, the L$^{3}$ mechanism of Goldreich et al. (2002) should form equal populations of prograde and retrograde binaries (Schlichting and Sari 2008).  So 2007 TY430 could have formed through the L$^{3}$ mechanism which would only be likely if the Kuiper Belt objects had relatively large velocities, called super-Hill velocities (Schlichting and Sari 2008a, 2008b).  The Hill velocity (Rafikov 2003; Goldreich et al 2004; Murray-Clay and Chiang 2006; Lee et al. 2007) is defined as:  \\begin{equation} v_H = \\left[\\frac{G(m_{p1}+m_{p2})}{r_{H}}\\right]^{1/2} \\label{eq:hill} \\end{equation}  where $G$ is the gravitational constant.  For 2007 TY430, the Hill velocity is about 0.4 m s$^{-1}$ which is much less than the Keplerian velocity of its Plutino type orbit that is about 4 km s$^{-1}$.  Formation through the L$^{3}$ mechanism would require large velocities between KBOs that several authors do not think were prevalent in the early Kuiper Belt (Goldreich et al. (2002), Schlichting and Sari (2008b), Murray-Clay and Schlichting (2011)).  Goldreich et al. (2002) suggest the velocities of KBOs were about a third of the Hill velocity, which would make retrograde orbits through the L$^{2}$s mechanism the dominant binary formation mechanism.  This is because binary formation efficiency from gravitational encounters decreases significantly with higher relative velocities since the sphere of influence or time an object affects another is decreased (Noll et al. 2008a; Schlichting and Sari 2008b).  Murray-Clay and Schlichting (2011) argue that sub-Hill velocities and the binary formation from dynamical friction (L$^{2}$s mechanism) is the best mechanism for the equal-sized binary formations.  2007 TY430 does not fall in along these formation lines because it is not retrograde, meaning large velocities during formation were likely and thus dynamical friction unlikely to be the capture mechanism.  A weakness to the high velocity scenario is that Schlicting and Sari (2008b) suggest that the velocity would need to be finely tuned and stay in such a state for a considerable amount of time in order for a high velocity regime to form these prograde, equal-sized binaries.  To date, not enough equal-sized binary orbital inclinations are known to draw a strong conclusion.  Only 17 equal-sized KBO binaries have known inclinations (Figure~\\ref{fig:binarymutualrhi}).  Of these five are retrograde and twelve are prograde, but three of the prograde objects have $\\Delta Mag>1$ mag and thus are not quite equal-sized binaries.  Several authors have suggested that the observed inclinations of TNO binaries are consistent with them being randomly distributed (Noll 2003, Chiang et al. 2006, Grundy et al. 2011).  This suggests the Kuiper Belt disk was in the dispersion-dominated (dynamically hot) regime during binary formation (see Stewart and Ida 2000; Collins and Sari 2006; Schlichting and Sari 2008b).  Lee et al. (2007) point out there are more prograde than retrograde binaries and determine this as a sign of how the binaries formed.  Schlichting and Sari (2008b) suggest this sign is showing the velocity regime in which the binaries formed.  Schlichting and Sari (2008b) predicted that over 97 percent of binaries with comparable masses will have retrograde orbits if the relative velocities of the Kuiper Belt objects were low.  Since it appears to be well below 50 percent, the relative velocities must have been significantly higher, as discussed above for the formation of 2007 TY430.  Binary formation from direct gravitational collapse of an over density of concentrated cm to meter sized solids was suggested by Nesvorn{\\'y} et al. (2010).  This mechanism appears to be able to produce a wide range of distant equal-sized binaries with a large range of eccentricities with most having prograde inclinations.  This mechanism could explain 2007 TY430, but may have trouble producing the retrograde binaries found to date.  \\subsection{Plutino Resonance and the Cold Classical Belt}  The Neptune mean motion resonance populations could have been emplaced by one of two favored mechanisms.  1) The slow migration outwards of Neptune by over 10 AU from its formation location from planetesimal scattering would have allowed Neptune to sweep many objects into the resonances we see today (Malholtra 1995; Hahn and Malhotra 2005). This scenario would likely mean that the inner 3:2 resonance would have a significantly different population of objects than the more outer resonances that would have swept through the classical Kuiper Belt population.  This scenario would also suggest the resonance populations should have a cold component representative of the cold classicals along with a hot component representing the objects from closer in.  2) The chaotic population of the resonances could have occurred if Neptune was scattered with a relatively large eccentricity out from near the current Saturn region to near its present semi-major axis (Levison and Morbidelli 2003; Tsiganis et al. 2005; Levison et al. 2008).  As Neptune circularized its orbit through interactions with planetesimals, objects in the resonance areas would become trapped.  In this scenario, the resonance populations would be more uniform and not have a cold component, unlike the slow migration scenario.  The resonance populations should have a cold component if Neptune experienced extensive slow, smooth migration.  This cold component should be similar in characteristics to the low inclination classical Kuiper Belt, that is, equal-sized binaries should be common as well as ultra-red material (Murray-Clay and Schlichting 2011).  This cold component should also have lower eccentricities compared to other objects, meaning there could be a correlation between binaries, color and eccentricity as well as with inclination in the resonance populations.  In the Levison et al. (2008) scenario, Chaotic scattering of Neptune and then circularization, the resonance populations should match the scattered disk in characteristics.  No obvious cold component is observed in the 3:2 resonance population and the statistics are too low for the other resonance populations to determine if there is a cold component (Murray-Clay and Schlichting 2011).  The 3:2 having no cold component is still consistent with slow, smooth migration since the 3:2 does not sweep through the cold classicals (Murray-Clay and Schlichting 2011).  Thus the 3:2 resonance is not as strong as a marker as the other resonances, which should have cold components as they swept through the classical region of the Kuiper Belt.  It is possible that sweeping of the $\\nu_{8}$ resonance through the 3:2 resonance and inner classical Kuiper belt cleared out much of these regions (Petit et al. 2011).  If this is true, the $\\nu_{8}$ resonance likely did not sweep through until after the formation of any equal-sized binaries since the binaries are unlikely to have formed as the density of such objects would have been too low.  2007 TY430 has a moderately inclined and eccentric heliocentric orbit, and thus would be considered part of the hot component of the 3:2 resonance on dynamical grounds. Although 2007 TY430 could easily have been emplaced recently by dynamical scattering and resonance sticking, it is interesting to speculate on the meaning of this cold-classical like binary if it was placed in the resonance primordially. As discussed above, if 2007 TY430 was captured from an originally low eccentricity, low inclination orbit by slow, smooth migration, it would have originated around 31 AU by conservation of Brouwer's integral. This would suggest that the primoridal cold classical region extended much further inward than is seen today, but the observational evidence for a cold region of the 3:2 resonance is insufficient (Murray-Clay and Schlichting 2011).  If its dynamical history was well-known, the relatively high eccentricity and inclination of 2007 TY430 could have been seen as an indication that the slow, smooth migration model is insufficient.  More binary objects need to be discovered in the 3:2, but the existence of 2007 TY430 in the hot region of the 3:2 and no other known equal-sized 3:2 resonance objects, would favor chaotic scattering by Neptune over slow smooth migration.  \\subsection{Other Known Equal-Sized Wide Binaries}  Table 5 shows the other known equal-sized TNO binaries.  Here, equal-sized is defined as the components having less than 1 magnitude difference between components.  This magnitude difference corresponds to mass ratios less than 4, assuming similar albedos and densities for the two components and the lack of significant rotational brightness variations.  As shown in Table 5 and Figures~\\ref{fig:kboeabinary} to~\\ref{fig:kboiabinaryblowup}, the equal-sized binary population is mostly in the low inclination classical belt. Interestingly, most of the equal-sized binary TNOs not in the low inclination classical belt appear to be in mean motion resonances.  We see no obvious increase in the binary population between the semi-major axes of 43.5 and 44.5 AU, which has been called the ``Kernel'' area by Petit et al. (2011) as this region appears to be somewhat more dense than the rest of the classical Kuiper belt (Figures~\\ref{fig:kboeabinaryblowup} and~\\ref{fig:kboiabinaryblowup}).  It does appear that there is an absence of equal-sized binaries around 43.5 AU, which is near the 7:4 resonance, as well as very few with eccentricities greater than 0.10.  The colors of the equal-sized binaries are not all ultra-red, as many are near neutral in color (Tegler and Romanishin 2000; Trujillo and Brown 2002; Gulbis et al. 2006; Peixinho et al. 2008; Petit et al. 2008; Benecchi et al. 2009; Sheppard 2010).  As Figure~\\ref{fig:binarycolors} shows, the only equal-sized binary that is ultra-red in color ($S>25$ as defined in Sheppard (2010)) and not in the classical population is 2007 TY430 (colors are from the database described in Hainaut and Delsanti 2002).  In all, only 1 of 6 equal-sized binaries with known colors outside of the classical population has an ultra-red color, which is 2007 TY430 (Table 5).  In contrast 13 of 17 equal-sized binaries in the classical population have spectral gradients near or within the ultra-red color region ($S>20$).  Numerically integrating the non classical equal-sized binaries from Table 5 (2000 QL251, 2000 FE8, 2006 SF369, 1998 WV24, 2000 CM114, 2001 QC298) with ten clones each for 500 MYr found all but 2000 CM114 were fairly stable. Some of the clones of 2000 QL151 ($6/10$, 1 a little chaotic) and 2006 SF369 ($10/10$, 3 a little chaotic) were experiencing Kozai oscillations (not unusual for resonant objects). Thus only 2000 CM114 is likely to have come from the classical region, but its inclination is currently very large and it is not ultra-red in color. Since all of the equal-sized binaries outside the classical belt, except for 2007 TY430, do not have ultra-red colors, it suggests that equal-sized binaries formed in multiple locations and not just in the cold classical belt.  As shown in Figure~\\ref{fig:binarycolors}, the inner classical Kuiper belt is likely an extension of the main classical Kuiper belt as both have low inclination, ultra-red, equal-sized binaries.  If the ultra-red, equal-sized binary of 1999 OJ4 formed in the inner classical belt, where we see it today, it must have formed before any clearing through the $\\nu_{8}$ resonance sweeping mechanism since the desnity of KBOs would have been to low for likely equal-sized binary formation after any significant clearing of the region."
    },
    "1112/1112.0011_arXiv.txt": {
        "abstract": "In this proceeding, we show that when we combined WMAP and the most recent results of XENON100, the invisible width of the Higgs to scalar dark matter is negligible($\\lesssim 10 \\%$),  except in a small region with very light dark matter ($\\lesssim 10$ GeV) not yet excluded by XENON100 or around 60 GeV where the ratio can reach 50\\% to 60\\%. The new results released by the Higgs searches of ATLAS and CMS set very strong limits on the elastic scattering cross section. ",
        "introduction": "Two of the most important issues in particle physics phenomenology are the nature of the dark matter and the mechanism to realize spontaneously the electroweak symmetry breaking of the Standard Model (SM). The observations made by the WMAP collaboration show that the matter content of the universe is dark, making up about 85 \\% of the total amount of matter whereas the XENON collaboration recently released its constraints on direct detection of Dark Matter . These constraints are the most stringent in the field nowadays, and begin to exclude a significant part of the parameter space of the Weakly Interacting Massive Particle (WIMP) paradigm. On the other front, the accelerator collaborations ATLAS, CMS and D0/CDF \\cite{ATLAS} have obtained important results concerning the Higgs searches. . It is obvious that the Higgs hunting at LHC is intimately linked with measurement of elastic scattering on nucleon, especially in Higgs-portal like models where the Higgs boson is the key particle exchanged through annihilation/scattering processes. It has already been showed recently that a combined LEP/TEVATRON/XENON/WMAP analysis can restrict severely the parameter space allowed in generic constructions \\cite{Mambrini:2011pw}. In this work, we apply such analysis in the specific context of a scalar singlet dark matter extension of the Standard Model and show that most of the region allowed by WMAP will be excluded/probed by LHC and XENON100 by the end of next year. ",
        "conclusions": ""
    },
    "1112/1112.2478_arXiv.txt": {
        "abstract": "% We consider the evolution of an initially FLRW universe after the formation of a long, straight, cosmic string with arbitrary tension and mass per unit length. The birth of the string sources scalar and tensor-type perturbations in the background metric and both density and velocity perturbations in the background fluid, which compensate for the string mass and maintain energy conservation. The former generate the deficit angle within the light cone of the string and a gravitational shock front at the cosmological horizon, whereas the latter are confined within the sound cone. We study the properties of the metric within each region of the resulting spacetime and give the explicit coordinate transformations which demonstrate non-violation of causality. This paper generalizes the work of previous studies for the Nambu-Goto string. ",
        "introduction": "Cosmic strings are linear concentrations of energy which may have formed during symmetry-breaking phase transitions in the early Universe \\cite{Nielsen_Olesen, hep-ph/9411342, VS1994,hep-th/0508135v2, Preskill}. Though the string width is determined by the inverse of the symmetry-breaking energy scale, this is small compared to cosmological distances, and they may be approximated as one-dimensional objects for many purposes \\cite{Goto1,Anderson2003}. Specifically, strings may have been produced at the epoch of electro-weak symmetry-breaking \\cite{Nambu1}, or the GUT scale, but their formation is also a generic feature of the phase transitions inherent in many extensions of the standard model,  (c.f. \\cite{hep-ph/0308134} and references therein). In field theory they are a type of topological defect, analogous to the magnetic flux tubes and other vortex-type defects created in certain condensed matter systems \\cite{Abrikosov1, Zurek1, Zurek2, Zurek3,  Bowick1, Williams1, Hendry1, Chuang1, Annett1}, and are produced via the Kibble mechanism \\cite{ICTP/75/5}, if the vacuum manifold ($\\mathcal{M}$) possesses a nontrivial first homotopy group (e.g. $\\pi_1(\\mathcal{M})=\\mathbb{Z}$, in which each element, a nonzero integer, corresponds to an allowed winding number for the string vortex cross-section). In recent years, the formation of ``cosmic\", i.e. horizon-sized, fundamental strings ($F$-strings) and one-dimensional $D$-branes ($D$-strings) has also been extensively studied in string theory \\cite{Witten:1985, Polchinski_Intro, CMP_FD1,CMP_FD2,0911.1345v3, 0811.1277v1, hep-th/0505050v1, astro-ph/0410073v2, Sakellariadou:2009, Rajantie:2007, Copeland:2005}, particularly in the context of brane inflation (see \\cite{Cline1, Gasperini1, Tye1, Carroll:TASI, Quevedo:Lectures, Danielsson1} for reviews), in which such defects can be copiously produced \\cite{Sarangi1, Jones_etal1, Pogosian_Obs1}. However, regardless of the precise details of individual models, the main phenomenological, and observationally relevant, parameters that characterize the string are the energy of the per unit length, which we denote $U$, and the tension, $\\mu = qU$. \\footnote{This is not strictly true in the case of models containing compact extra-dimensions in which the intercommuting probability for strings which ``cross\" in the noncompact space can be much less than unity ($P<<1$). In this case the dynamics of string networks and their implications for cosmological observations can be substantially altered \\cite{Sarangi1,Jones_etal1,Pogosian_Obs1,Avgoustidis:2004,Copeland:2005}. Extra-dimensional scenarios, especially those motivated by, or directly embedded into, string theory also allow for the formation of cosmic ``necklaces\" of various kinds, which introduce significant alterations to the network dynamics along with additional phenomenological parameters, such as the bead mass \\cite{Leblond:2007, Dasgupta1, Martins1, Matsuda2, Lake1, Lake3}. The analogues of a superstring necklaces in field theory are series of monopoles connected by string-segments, which may also be formed in some symmetry-breaking models \\cite{Siemens1, Blanco-Pillado1}. The authors hope to address the important question of the inborn metric for a cosmic necklace in a future article.} The Nambu-Goto (NG) action represents the simplest model and enforcing Lorentz symmetry along the string corresponds to setting $q=1$. If the NG string is wiggly at microscopic scales, it may be approximated macroscopically by a simpler embedding with an effective mass per unit length which differs from the intrinsic tension. Thus in the coarse-grained limit, small-scale structure, even on pure NG strings, can lead to a value of $q$ different from unity \\cite{vilenkin,carter}. For current carrying strings, the mass-density and tension may differ, as originally shown by Witten \\cite{witten,FPRINT-92-39}, so that $q \\neq 1$ in general. \\\\ \\indent The gravitational field around a static, straight string exhibits an interesting feature not seen in fields surrounding spherically-symmetric distributions of matter. In a perturbative analysis of a string in the wire approximation (to first order in $G\\mu$), Vilenkin \\cite{167218} showed that test particles within the string light-cone experience a gravitational acceleration along the radial direction equal to $2GU(q-1)/r$, where $r$ is a distance from the string, and that the surrounding space has a conical structure with deficit angle $4\\pi (1+q)GU$. For the NG string therefore, the gravitational force vanishes and the spacetime becomes locally Minkowski. These results were found to hold in a nonperturbative analysis (i.e. to all orders in $G\\mu$), independently, by Gott \\cite{Gott1} and Hiscock \\cite{Hiscock1}, even for a string of finite width. One observational consequence is the formation of double images from bright sources lying behind the string in relation to an observer. The images have exactly the same shape and brightness and for NG strings their angular separation on the sky is of the order of the deficit angle, though for $q>>1$ it may be substantially smaller \\cite{Pe94,Uz01}. This feature has been used to search for long strings \\cite{Gott1, Vilenkin_lensing1,Hogan_lensing1,Shlaer_lensing1,Huterer_lensing1,Hindmarsh_lensing1, Dyda_lensing1,Gasperini_lensing1, Morganson_lensing1, deLaix_lensing1} and loops \\cite{Cowie_lensing1, deLaix_lensing2, Mack_lensing1} via astronomical observations. Bounds on phenomenological parameters such as the string tension and number density have also been obtained using the predicted effects of strings on CMB anisotropy \\cite{Zeldovich_CMB1, Kaiser_CMB1,Stebbins_CMB1,Gangui_CMB1,Allen_CMB1,Perivolaropoulos_CMB1,Benabed_CMB1,Landriau_CMB1, Landriau_CMB2, Pogosian_CMB1, Wyman_CMB1, Battye_CMB1, Battye_CMB2, Contaldi_CMB1, Bevis_CMB1, Bevis_CMB2, Bevis_CMB3, Bevis_CMB4,Suyama_CMB1,Suyama_CMB2}, the expected emission of high-energy rays from cusps and kinks \\cite{Bhattacharjee_CosRay1, Sigl_CosRay1, Berezinsky_CosRay1, Wichoski_CosRay1} and possible signatures in the 21cm line \\cite{Khatri_21cm}.  Interestingly, it has recently been proposed that cusp-emission from superconducting strings may be responsible for anomalous gamma-ray burst observations \\cite{Gr09,Ta09,Sa09} and there remains ongoing controversy on that point (c.f. \\cite{Ba87,Pa88,Br93,Pl94,Be01,Be04,Ch10,Wa11,Ch11} and references therein). With the advent of the Planck experiment \\cite{Planck}, it is also hoped that the polarization of CMB $B$-modes caused by string networks in the early universe could confirm, or rule out, the existence of GUT-scale strings in the near future \\cite{Seljak_Pol1,Pogosian_Pol1,Pogosian_Pol2,Bevis_Pol1,Garcia_Pol1}. Finally, gravitational radiation from strings has been extensively studied \\cite{Vachaspati_GravRad1,Hindmarsh_GravRad1,Darmour_GravRad1,Darmour_GravRad2,Darmour_GravRad3} and future gravitational wave detections from experiments such as LISA/NGO \\cite{LISA/NGO} or LIGO \\cite{LIGO} may place even stronger bounds on string parameters. \\\\ \\indent It is interesting to consider how the conical structure of the spacetime surrounding the string evolves from the initial Friedmann-Lema\\^\\i tre-Robertson-Walker (FLRW) background, which exists before the symmetry-breaking epoch, and a detailed analysis of this problem was originally given by Magueijo \\cite{Magueijo:1992tt,mag-unpublish} for the case of an NG string. He introduced a phenomenological string energy-momentum tensor which is localized at the origin and appears only after the symmetry-breaking time. The generation of string mass-energy is compensated by a decrease in the mass-energy of the perfect fluid that dominates the Universe, so that conservation of energy and momentum are not violated. \\footnote{In previous analyses this was not always the case, though particle production by cosmic strings was studied in detail in the following sources, \\cite{Parker:1987qx,Sa88,Da88}.} Inclusion of the background fluid is therefore compulsory and the problem of string generation must be considered on an (initially) FLRW background. In his analysis, Magueijo obtained solutions of the linearized Einstein equations on the FLRW background metric by assuming the dimensionless parameter $GU$ to be small and constant. This is consistent with known observational bounds on the string tension in various cosmological models, which give an upper limit of $GU \\leq \\mathcal{O}(10^{-6})$ for field-theoretic strings \\cite{VS1994, Bevis_CMB1,Bevis_CMB2,Bevis_CMB3,Bevis_CMB4} and $10^{-11} \\leq GU \\leq 10^{-6}$ for the warped tension of cosmic $F$/$D$-strings \\cite{Jones_etal1,Pogosian_Obs1} (see also \\cite{Hindmarsh_rev} for a review of observational constraints up to 1990). He showed that, at late times and in the region close to the string, the spacetime takes the form of that surrounding an eternal string having deficit angle $8\\pi GU$ \\cite{167218,Gott1,Hiscock1} and that a gravitational shock propagates along the lightcone, with the spacetime outside the horizon remaining unperturbed so that the FLRW metric remains valid, though detailed calculations confirming this picture remain unpublished \\cite{mag-unpublish}. \\\\ \\indent In this paper, we reconsider this issue and also extend the original analysis by considering the more general case of $q \\ne 1$. This allows us to include the physically interesting class of ``conical\" spacetimes which are not locally flat and to consider the gravitational acceleration of test particles in the vicinity of a noneternal string. We follow the same approach taken in \\cite{Magueijo:1992tt} to evaluate the metric perturbations induced by string formation and find that the general picture remains qualitatively similar. However, as shown in detail in Sec.~IV, our solutions do not exactly match those obtained previously, even for $q=1$. We believe that this is because of the choice of boundary condition for the Poisson equation governing the gravitational potential given in \\cite{Magueijo:1992tt} is not consistent with trace part of the Einstein equations. We also provide concrete coordinate transformations both for inside and outside the sound horizon of the perfect fluid, as well as in the region between the sound horizon and the edge of the light cone, which result in an FLRW metric beyond the causal horizon and the metric for an eternal string deep inside the sound horizon, respectively, in accordance with previous results \\cite{167218,Gott1,Hiscock1,Magueijo:1992tt,mag-unpublish}. \\\\ \\indent The structure of the paper is as follows: In Sec.~II we define the perturbation variables and decompose them into scalar and tensor components. In Sec.~III we derive the master equations for each type, while analytic solutions are given in Sec.~IV together with a comparison of the corresponding solutions obtained in \\cite{Magueijo:1992tt}. Section V is devoted to studying the properties of the resulting class of spacetimes and a brief summary of our results and discussion of the prospects for future work is given in the conclusion, Sec.~VI. ",
        "conclusions": "We derived the analytic form of the metric perturbations on an FLRW background, excited by the formation of a straight cosmic string under the wire approximation. Our results are valid for a general string with differing tension and energy per unit length. For the Nambu-Goto string, our metric does not coincide with that given in previous work due to a differing choice of boundary conditions for the Poisson equation. We found that, as a result of this choice, the solutions given in \\cite{Magueijo:1992tt}  for the gravitational potential and curvature perturbation variables do not satisfy the trace part of the Einstein equations. \\\\ \\indent By performing the appropriate coordinate transformation, we explicitly verified that the spacetime outside the light cone remains unperturbed and governed by the FLRW metric. Whereas it is nontrivial to demonstrate nonviolation of causality in the original metric, it is manifestly respected in the new coordinate system. At the cosmological horizon, we found that there exists a gravitational shock giving rise to a divergent tidal force along both the angular direction and the direction of the string axis, $z$. The shock distorts a small ring of test particles placed in the $\\theta-z$ plane by stretching it along $\\theta$-direction and shrinking it along the $z$-direction. The degree of divergence is weak, so that the resultant distortion remains finite. Inside the cosmological horizon, the metric perturbations exhibit self-similarity at late time, i.e. they depend only on a ratio of the radial coordinate to the conformal time. This is reasonable since the Hubble distance is the only characteristic length scale in the system. Deep inside the horizon, we found that the metric tends to that for an eternal string at late times, which has been used to search for the signatures of cosmic strings in astronomical observations."
    },
    "1112/1112.5599_arXiv.txt": {
        "abstract": "{We constrain the energy at which the transition from Galactic to extragalactic cosmic rays occurs by computing the anisotropy at Earth of cosmic rays emitted by Galactic sources. Since the diffusion approximation starts to loose its validity for $E/Z \\gtrsim 10^{16-17}$\\,eV, we propagate individual cosmic rays using Galactic magnetic field models and taking into account both their regular and turbulent components. The turbulent field is generated on a nested grid which allows spatial resolution down to fractions of a parsec. Assuming sufficiently frequent Galactic CR sources, the dipole amplitude computed for a mostly light or intermediate primary composition exceeds the dipole bounds measured by the Auger collaboration around $E \\approx 10^{18}$\\,eV. Therefore, a transition at the ankle or above would require a heavy composition or a rather extreme Galactic magnetic field with strength $\\gtrsim 10\\,\\mu$G. Moreover, the fast rising proton contribution suggested by KASCADE-Grande data between $10^{17}$\\,eV and $10^{18}$\\,eV should be of extragalactic origin. In case heavy nuclei dominate the flux at $E \\gtrsim 10^{18}$\\,eV, the transition energy can be close to the ankle, if Galactic CRs are produced by sufficiently frequent transients as e.g.\\ magnetars. } ",
        "introduction": "\\label{Introduction}  The question at which energy the transition from Galactic to extragalactic cosmic rays (CRs) takes place is one of the major unresolved issues of cosmic ray physics. Two promising possibilities are to associate the transition with one of the two evident features of the cosmic ray spectrum: The second knee around $E\\simeq 5\\times10^{17}$\\,eV or the ankle at $E\\simeq3\\times 10^{18}$\\,eV. Since the chemical composition of galactic and extragalactic CRs should differ in general, both because of propagation effects and of the different nature of their sources, the transition may be detected experimentally studying the chemical composition of CRs as function of energy.  In the case of a transition around the second knee, Galactic CR sources such as e.g.\\ supernova remnants would accelerate CRs up to the rigidity-dependent knee, which is close to $10^{17}$\\,eV for iron. If the extragalactic CR flux dominating at higher energies would consist mainly of protons, the ankle could be explained as a dip in the extragalactic CR spectrum due to the pair-production losses of protons on cosmic microwave background (CMB) photons $p + \\gamma_{\\rm CMB} \\rightarrow p + e^+ + e^-$~\\cite{dip}. Below $\\sim10^{17-18}$\\,eV, the extragalactic CR flux may be suppressed because of CR propagation in extragalactic magnetic fields~\\cite{Lemoine:2004uw,Kotera:2007ca}. On the other hand, the scenario of Ref.~\\cite{Allard:2005cx} would favour a transition at the ankle. The composition of the CR flux at high energies is the subject of current debate due to the facts that hadronic physics must be extrapolated from lower energies and that the complex experimental analyses for different experiments are not yet completely reconciled. The scenario of Ref.~\\cite{dip} is supported by the composition measurements of HiRes~\\cite{HiRes} and the first results of the Telescope Array~\\cite{TA}, which are consistent with a light composition around the ankle and above. On the other hand, recent results from the Pierre Auger Observatory~\\cite{Abraham:2010yv,auger:2011pe} indicate a composition becoming heavier with increasing energy above the ankle, and the Yakutsk EAS array muon data suggests a non negligible fraction of heavy nuclei above $\\simeq10^{19}$\\,eV~\\cite{Yakutsk}. Moreover, the measurements of the KASCADE-Grande~\\cite{KASCADE-Grande} collaboration are consistent with a dominantly heavy composition up to $10^{18}$\\,eV. However, the KASCADE-Grande data indicate a fast rising proton contribution above $10^{17}$\\,eV.  Thus at present the experimental data on the CR composition do not allow us yet to determine the transition energy between Galactic and extragalactic CRs. In this paper we suggest to use instead experimental limits on the anisotropy of the arrival directions of UHECRs to constrain the maximal contribution of Galactic CRs at $E\\gtrsim 10^{18}$\\,eV. At energies below $10^{17}$\\,eV, the diffusive propagation of Galactic cosmic rays and their resulting anisotropy at Earth was studied in details in Refs.~\\cite{Blasi:2011fi,Blasi:2011fm}.  Since the propagation of CRs in the Galactic magnetic field (GMF) is not longer diffusive at $E\\gtrsim 10^{17}$\\,eV, we directly propagate UHECRs in the GMF using the numerical code developed in Refs.~\\cite{Giacinti:2010dk,Giacinti:2011uj}. We present also a way to generate the turbulent field on a nested grid without limitation on its spatial resolution. This method allows us to include magnetic field fluctuations spanning the required large dynamical range of scales, from negligible compared to the CR Larmor radii up to 300\\,pc. As main result of this work we show that the existing limits on CR anisotropies strongly restrict the contribution of the CNO element group to the Galactic CR component above $E\\gtrsim1$\\,EeV, while the contribution of iron is restricted above $E\\gtrsim3$\\,EeV.  Details of the method to generate turbulent magnetic fields are discussed in the Section~\\ref{Method_TF}. In Section~\\ref{Method_Anisotropy}, we review the GMF models used and discuss how the CR anisotropy is calculated. Results of numerical simulations are presented in the Sections~\\ref{Anisotropy} and \\ref{Spectrum} for anisotropies and the spectrum of UHECR. ",
        "conclusions": "\\label{Conclusions}  In this work we studied the consistency of a transition from Galactic to extragalactic CRs with existing anisotropy limits as a function of energy above $E=10^{18}$\\,eV. The diffusion approximation predicts a dipole anisotropy $\\delta=-3D_{ij}\\nabla_j\\ln(n)$ increasing with energy, since both the diffusion tensor $D_{ij}$ and the relative CR gradient $\\nabla_j\\ln(n)$ increase with energy. However, this approximation becomes unreliable at ${\\cal O}(E/Z)\\sim 10^{16}$\\,eV, and therefore we studied the propagation of  CRs in the Galactic magnetic field directly by backtracking trajectories. We simulated the turbulent magnetic field on nested grids which allows one to include turbulent field modes $\\Bk$ with arbitrary small wave-lengths. For the regular Galactic magnetic field we used up-to-date models from Ref.~\\cite{Pshirkov:2011um}. Because the global structure of the GMF is still rather uncertain, we studied the dependence of the resulting anisotropy on the magnetic field parameters such as its strength $B_0$, scale height $z_0$, correlation length $L_{\\rm c}$ and exponent $\\alpha$ of its power-spectrum. We also examined the dependence of our results on the width and height of the disk in which sources are located.  The main results of this study are presented in the Figs. 3--5. They show that the anisotropy mostly depends on the amplitude $B_0$ of the magnetic field in the disk. As our main conclusion from this study, we found that existing anisotropy limits are not compatible with light (proton) and intermediate (CNO) nuclei of Galactic origin as dominant contribution to the CR flux above 1\\,EeV. By contrast, Galactic iron nuclei as CR primaries are consistent with the existing limits even up to 10--20\\,EeV, if the strength of the turbulent field is as large as $B_{\\rm rms}\\sim 8\\,\\mu$G. This finding implies that determining the chemical composition of the CR flux around $10^{18}$\\,eV settles also the question of the transition energy between Galactic and extragalactic component: As light nuclei at this energy are not sufficiently isotropized, they have to be extragalactic. Therefore the fast increasing proton contribution indicated by the KASCADE-Grande collaboration between $10^{17}$\\,eV and $10^{18}$\\,eV suggests the beginning of an extragalactic component.  We also studied qualitatively the dependence of the anisotropy on the effective density of sources, see Figs.~6--8. The average escape time of iron nuclei with 10\\,EeV energy from the Galaxy is $\\sim 10^{5}$\\,yr. Assuming for magnetars a rate of $10^{-3}/$yr, the effective density of magnetars as sources of CR at 10\\,EeV is $\\sim 100$/Galaxy. Thus magnetars satisfy the anisotropy constraint and can be natural candidates for the sources of the high-energy end of the Galactic CR flux in the scenario where the transition from Galactic to extragalactic cosmic rays occurs at the ankle, provided they are able to accelerate iron up to few$\\;\\times 10^{18}$\\,eV.  In summary, we conclude that models with a transition from Galactic to extragalactic cosmic rays around the ankle are consistent with the existing anisotropy limits if the composition of Galactic cosmic rays at $E \\gtrsim 10^{18}$\\,eV is dominated by heavy nuclei. In contrast, if the chemical composition at these energies turns out to be light or intermediate, a transition at the ankle would be very strongly disfavoured."
    },
    "1112/1112.4255.txt": {
        "abstract": "In this paper we revisit the dynamical dark energy model building based on single scalar field involving higher derivative terms. By imposing a degenerate condition on the higher derivatives in curved spacetime, one can select the models which are free from the ghost mode and the equation of state is able to cross the cosmological constant boundary smoothly, dynamically violate the null energy condition. Generally the Lagrangian of this type of dark energy models depends on the second derivatives linearly. It behaves like an imperfect fluid, thus its cosmological perturbation theory needs to be generalized. We also study such a model with explicit form of degenerate Lagrangian and show that its equation of state may cross $-1$ without any instability. ",
        "introduction": "The recent data from type Ia Supernovae (SNIa) and cosmic microwave background (CMB) radiation and so on have provided strong evidences for a spatially flat and accelerated expanding universe. In the context of Friedmann-Robertson-Walker (FRW) cosmology with Einstein gravity, this acceleration is attributed to the domination of a component with negative pressure, called dark energy. So far, the nature of dark energy remains a mystery. Theoretically, the simplest candidate for such a component is a small positive cosmological constant, but it suffers the difficulties associated with the fine tuning and the coincidence problems. Therefore, many physicists are attracted by the idea of dynamical dark energy models, such as quintessence \\cite{Ratra:1987rm,Wetterich:1987fm,Caldwell:1997ii}, phantom \\cite{Caldwell:1999ew}, k-essence \\cite{ArmendarizPicon:2000dh}, quintom \\cite{Feng:2004ad} and so on (see Refs. \\cite{Copeland:2006wr, Frieman:2008sn, Caldwell:2009ix, Cai:2009zp, Li:2011sd} for recent reviews).  Although the recent fits to the data \\cite{Li:2011dr}, in combination of the 7-year WMAP \\cite{Komatsu:2010fb}, the Sloan Digital Sky Survey \\cite{Reid:2009xm}, and the recently released \u00a1\u00b0Union2\u00a1\u00b1 SNIa data \\cite{Riess:2011yx}, show remarkably the consistence of the cosmological constant, it is worth noting that a class of dynamical models with the equation-of-state (EoS) across $-1$, dubbed as {\\it quintom} which dynamically violates the null energy condition (NEC), is mildly favored \\cite{Feng:2004ad}. However, it was noticed that consistent single field models realizing the quintom scenario are difficult to be constructed. For example, the EoS of quintessence is limited to be in the region $-1\\leq w \\leq 1$, while for phantom $w$ is always smaller than $-1$. It was proved that in models described by a single perfect fluid or a single scalar field with a Lagrangian of k-essence form \\cite{ArmendarizPicon:1999rj}, the cosmological perturbations encounter a divergence when the background EoS crosses $-1$ \\cite{Feng:2004ad, Vikman:2004dc, Hu:2004kh, Caldwell:2005ai, Zhao:2005vj}. This statement was explicitly proven in Ref. \\cite{Xia:2007km} as a ``No-Go\" theorem for dynamical dark energy models.  To realize a viable quintom scenario, one usually needs to add more degrees of freedom into the dark energy budget. The simplest and also the first quintom model was constructed by a combination of a canonical scalar and a phantom scalar field \\cite{Feng:2004ad}. However theorists are still interested in pursuing single field quintom models. The first single scalar field quintom model was realized by introducing higher order derivative terms \\cite{Li:2005fm}, also see \\cite{Zhang:2006ck} for generalization. It was also obtained in the frame of nonlocal string theory \\cite{Aref'eva:2005fu} and decaying tachyonic branes \\cite{Cai:2007gs}. However, there is a quantum instability due to an unbounded vacuum state in such type of models \\cite{Carroll:2003st, Cline:2003gs}. A possible approach to stable violations of NEC is the ghost condensation of Ref. \\cite{ArkaniHamed:2003uy}, in which the negative kinetic modes are bounded via a spontaneous Lorentz symmetry breaking, although it might allow for superluminal propagation of information in some cases \\cite{Dubovsky:2005xd}. Various theoretical realizations of quintom scenarios and their implications for early universe physics were reviewed in Ref. \\cite{Cai:2009zp} (see also \\cite{Qiu:2010ux}).  Recently, a scalar field model which stably violates the NEC has been studied extensively, which is the so-called Galileon \\cite{Nicolis:2008in}. It was viewed as a local infrared modification of General Relativity, generalizing an effective field description of the DGP model \\cite{Dvali:2000hr}. The key feature of these models is that they contain higher order derivative terms in the action while the equation of motion remains second-order in order to avoid the appearance of ghost modes, realizing the idea pioneered by Horndeski thirty years ago \\cite{Horndeski:1974}. Later on, various phenomenological studies of this type of models were performed, namely, see Refs. \\cite{deRham:2010eu, Deffayet:2010qz, Pujolas:2011he, Deffayet:2009wt, Deffayet:2011gz, Gao:2011qe,Chow:2009fm, Silva:2009km, DeFelice:2010pv, Kobayashi:2010cm, Hinterbichler:2010xn, Creminelli:2010ba}. Motivated by the feature of the Galileon model, in this paper we revisit quintom dark energy models containing higher derivative terms. We start from a general covariant Lagrangian of single scalar field involving higher derivative terms and pursue how to keep the model free from extra degree of freedom. We show that it is able to eliminate the ghost mode by imposing a degenerate condition that the Lagrangian only depends on the second derivative terms linearly. Based on the degenerate Lagrangian, we build an explicit dark energy model and study its dynamics of its homogeneous background and its perturbations. Our numerical calculations show that the EoS is able to cross the cosmological constant boundary smoothly and the perturbation modes are well controlled when the crossing takes place. We understand the reason of realizing the single field quintom without ghosts is that, such a single scalar field model is no longer be able to correspond to a perfect fluid. Thus, this model does not conflict with the ``No-Go\" theorem for quintom dark energy model building as proposed in \\cite{Xia:2007km}.  This paper is organized as follows. In Section II, we simply review the difficulty of constructing a single field dark energy model which gives rise to quintom scenario. In Section III, we start with a single field action involving higher derivative terms, and present the general analysis. We also discuss under what condition the higher derivative terms do not bring a ghost mode to the effective field description in Section IV. Section V is devoted to the study of the background and perturbation dynamics of single field dark energy model with degenerate higher derivatives in the frame of flat FRW universe. Specifically, we present an explicit model of quintom dark energy with degenerate higher derivatives. We perform numerical computation to illustrate such a model can realize the EoS across $-1$ smoothly. Section VI is the summary. ",
        "conclusions": "The model building of dark energy which is possible to cross the cosmological constant boundary, i.e., the quintom dark energy has attracted a lot of attention in the literature (for example see Refs. \\cite{Quintom_sum}). A very interesting question is how to build a single scalar model which gives rise to the crossing without any instabilities. The past studies have shown that a single scalar field satisfying a generic k-essence Lagrangian cannot give rise to the stable violation of NEC. However, it becomes possible when higher derivative operators are introduced. The non-degenerate higher derivative model can eliminate the classical instabilities but is still plagued by the ghost mode. In this paper we considered the degenerate higher derivative model inspired by the Galileon theory and its generalizations. We started with a single scalar of which the Lagrangian contains generic higher derivative operator, and then suggested a degenerate condition to eliminate the extra degrees of freedom brought by these higher derivative operators. Using the degenerate condition in the curved spacetime, the Lagrangian can be reduced to the form (\\ref{kg}). We studied the cosmological application of this model, including the background and perturbation analysis. Particularly, we chose an explicit example and performed a detailed numerical computation. Our numerical results verified that this model is able to realize the EoS to cross the cosmological constant boundary and the perturbation evolves smoothly without any pathologies.  A quintom model which successfully violates the NEC without leading to quantum instability also has important implications to the early universe physics. In Refs. \\cite{Cai:2007}, the authors found that a quintom model can avoid the big bang singularity widely existing in standard and inflationary cosmologies. In this picture, the moment of initial singularity can be replaced by a big bounce. Based on this scenario, the corresponding perturbation theory has been extensively developed, namely, on adiabatic perturbations \\cite{Cai:2008qb, Cai:2008qw, Cai:2009hc}, non-Gaussianities\\cite{Cai:2009fn}, entropy fluctuations\\cite{Cai:2008qw, Cai:2011zx}, and the related preheating phase\\cite{Cai:2011ci}. Recently, it was observed that the Galileon model with the EoS across $-1$ exactly leads to a bouncing solution in the frame of the flat FRW universe\\cite{Qiu:2011cy}. Consequently, we expect that bouncing cosmologies can be realized in a generic quintom model with degenerate higher derivative operators."
    },
    "1112/1112.5882_arXiv.txt": {
        "abstract": "{DBV stars  are pulsating white dwarfs  with atmospheres rich in  He.  Asteroseismology of  DBV stars  can provide  valuable clues about  the  origin, structure  and  evolution of  hydrogen-deficient white dwarfs,  and may  allow to study  neutrino and  axion physics. Recently, a  new DBV star, KIC  8626021, has been  discovered in the field of  the \\emph{Kepler} spacecraft. It is  expected that further monitoring of this star in the next years will enable astronomers to determine  its   detailed  asteroseismic  profile.}{We   perform  an asteroseismological analysis  of KIC 8626021  on the basis  of fully evolutionary DB  white-dwarf models.}  {We employ a  complete set of evolutionary  DB white-dwarf  structures  covering a  wide range  of effective temperatures and stellar  masses.  They have been obtained on the basis of a  complete treatment of the evolutionary history of progenitors stars.  We  compute $g$-mode adiabatic pulsation periods for  this  set  of  models  and  compare  them  with  the  pulsation properties  exhibited by  KIC  8626021.}{On the  basis  of the  mean period spacing of the star, we found that the stellar mass should be substantially    larger   than    spectroscopy    indicates.    From period-to-period   fits  we   found  an   asteroseismological  model characterized  by  an effective  temperature  much  higher than  the spectroscopic     estimate.}{In    agreement    with     a    recent asteroseismological  analysis  of this  star  by  other authors,  we conclude that KIC  8626021 is located near the blue  edge of the DBV instability strip, contrarily  to spectroscopic predictions. We also conclude that the mass of KIC 8626021 should be substantially larger than thought.} ",
        "introduction": "\\label{intro}  V777  Her (or  DBV)  stars  are $g$-mode  variable  white dwarfs  with He-rich  atmospheres  (DB) and  effective  temperatures  in the  range $21\\,500 \\lesssim  T_{\\rm eff} \\lesssim  29\\,000$ K that  pulsate with periods between 100  and $1100$ s (Winget \\&  Kepler 2008; Althaus  et al. 2010).  They are  the hotter cousins of the  ZZ Ceti  (or DAV)  stars, which  are pulsating  H-rich atmosphere white dwarfs that define an instability strip centered at $T_{\\rm eff} \\approx 12\\,000$ K.  As an  excellent demonstration of the validity of the  stellar pulsation  theory,  the  existence of  the  DBV class  of compact pulsators was predicted  on theoretical grounds (Winget et al. 1982a)  before  it  were  confirmed  observationally  (Winget  et  al. 1982b).   Pulsations  in  V777 Her  are  thought  to  be driven  by  a combination of the $\\kappa-\\gamma$  mechanism acting in the He partial ionization  zone  ---and  thus  setting  the  blue  edge  of  the  DBV instability  strip  (Winget et  al.   1983;  Bradley  \\& Winget  1994; C\\'orsico  et  al. 2009),  and  the  ``convective driving''  mechanism (Brickhill 1991; Goldreich \\& Wu 1999) which is thought to be dominant once the outer convection zone has deepened enough.  White-dwarf asteroseismology  ---the comparison between  the pulsation periods  of white  dwarfs  and the  periods  computed for  appropriate theoretical  models---  allows us  to  infer  details  of the  origin, internal  structure and  evolution  of white  dwarfs.  In  particular, estimates of the stellar mass,  He and H layer mass, core composition, magnetic field, rotation  rate, seismological distance, and properties of  the  outer convection  zone  can  be  inferred from  the  observed pulsations of  DAV and DBV stars  (Winget \\& Kepler  2008; Althaus et al. 2010).  Finally, an eventual measurement of the temporal  changes in the observed stable  periods of DBVs could allow  to  study neutrino  emission  (Winget  et  al. 2004)  and  also \\emph{axion} emission (see Isern et al. 2010 and references therein).  Unlike the ZZ Ceti variables,  that constitute the most numerous class of pulsating white dwarfs (about 150 objects are known today; Castanheira et  al. 2010),  V777 Her are  quite rare and  difficult to find,  and until  recently, only  20 stars  of this  class  were known (Beauchamp et al. 1999; Nitta et al. 2009; Kilkenny et al. 2009).  The list   is  now   a  bit   enlarged   with  KIC   8626021  ($=$   GALEX J192904.6$+$444708,  $T_{\\rm eff}=  24\\,950 \\pm  750$ K  and  $\\log g= 7.91\\pm 0.07$ dex), a very recently discovered DBV star located in the field  of view  of the  \\emph{Kepler  Mission} by  {\\O}stensen et  al. (2011) (hereinafter  {\\O}EA11)\\footnote{Even more recently,  Hermes et al.  (2011) have reported the discovery of the first DAV star in the \\emph{Kepler  Mission} field.}.   This star  exhibits at  least five periodicities in  the range $197-376$  s and amplitudes between  1 and 5.2 mma,  with the three  strongest modes showing a  triplet structure due possibly to  rotation.  It is expected that  additional modes with higher radial orders (longer periods)  will be detected in future runs of \\emph{Kepler}, thus increasing the potential of asteroseismology to infer its interior  structure.  Given the values of  $T_{\\rm eff}$ and $\\log g$ quoted above and the DB models of Althaus et al. (2009a), KIC 8626021 should have a mass of $M_*= 0.56 \\pm 0.03 M_{\\odot}$.  Motivated by  the exciting discovery of the  first pulsating white dwarf   in  the   \\emph{Kepler}   field  of   view,   we  present   an asteroseismological analysis of KIC 8626021 on the basis of the DB white-dwarf models presented in Althaus et al.  (2009a).  This grid of evolutionary models was computed for a wide range of stellar masses on the  basis of  a complete  treatment  of the  evolutionary history  of progenitors stars,  including those  stages relevant for  the chemical profiles of the white dwarf, such  as the core H and He burning phases, the  thermally pulsing asymptotic giant branch  phase, and the born-again episode that is responsible for the H deficiency.  By considering the mean period spacing exhibited by the star, we found that   KIC  8626021   should  be   substantially  more   massive  than spectroscopy  indicates.  We  also found  that period  to  period fits favour an  asteroseismological solution characterized  by an effective temperature much higher than the spectroscopic estimate.  While   writing  this  paper,   Bischoff-Kim  \\&   {\\O}stensen  (2011) (hereinafter BK{\\O}11) announced the  results of their own independent asteroseismological analysis on KIC 8626021. They found that this star is actually a hot DBV.  Using entirely different models and methods, we arrive at the same  conclusion. The fundamental difference between our approach  and that  of BK{\\O}11  is  that we  run evolutionary  models starting  on  the  Zero  Age  Main Sequence,  accounting  for  nuclear burning, time  dependent diffusion of the elements  and other physical processes  that  take  place in  the  course  of  the evolution  of  a star.  BK{\\O}11 perform ``fast  white dwarf  asteroseismology'', where they guess and parameterize the internal chemical composition profiles and build static models. That  method allows them a fuller exploration of parameter space. For KIC  8626021, they find very good fits (hotter than  what the  preliminary spectroscopy  suggests) but  also conclude that their internal Oxygen  composition profiles are in poor agreement with stellar  evolution calculations. They suggest  further studies to see what  the physical  parameters of models  that agree  with stellar evolution calculations would be. In  essence, this is the kind of study we present here. While our best fit models are evidently different, we concur  with the  conclusion  that  KIC 8626021  is  hotter (and  more massive)  than  suggested  by  the initial  spectroscopic  study.  The conclusion that  this star  is residing  at the blue  edge of  the DBV instability strip appears to be robust, and calls for the necessity of a  new  and  improved  spectroscopic determination  of  its  effective temperature.  In Section \\ref{modeling} we present some details about our models and methods.  In  Section  \\ref{avgspacing},  we use  the  average  period spacing  of KIC  8626021's period  spectrum to  draw  some conclusions about its  mass and effective  temperature. We contrast our  models to the grid of models computed by BK{\\O}11. In Section \\ref{best-fit}, we present our best fit models. Again, we contrast our results with those of BK{\\O}11. In Section \\ref{discussion} we discuss the discrepancies between the results coming from spectroscopy and from asteroseismology, and the differences between our asteroseismological results and those of BK{\\O}11. We conclude in Section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  In  this paper  we have  presented  a detailed asteroseismic analysis of  KIC 8626021, the  first pulsating  DB white dwarf  star discovered  by the \\emph{Kepler  Mission},  on the  basis  of  the  full evolutionary  DB white-dwarf  models presented in  Althaus et  al.  (2009a)  which were computed for a  wide range of stellar masses  and He envelopes.  These DB  white  dwarf  models  are  characterized  by  consistent  chemical profiles for both the core  and the envelope.  These chemical profiles are the result  of the computation of the  full and complete evolution of the progenitor stars from the zero age main sequence, including the core H and He  burning phases, the thermally pulsing asymptotic giant branch phase, the born-again episode that is responsible for the H  deficiency, and from  time-dependent element  diffusion predictions during the white-dwarf stage.  By considering the mean period spacing exhibited by the star, we found that  KIC  8626021 should  have  a stellar  mass  in  the range  $0.60 \\lesssim M_*/M_{\\odot} \\lesssim 0.87$, substantially larger than those derived  by   previous  spectroscopic  ($M_*   \\sim  0.56  M_{\\odot}$; {\\O}EA11)  and  asteroseismic ($M_*  \\sim  0.57 M_{\\odot}$;  BK{\\O}11) studies.  We  also  found  that  period-to-period  fits  point  to  an asteroseismological  model  with  an  effective temperature  of  $\\sim 27\\,300$  K,  in  strong  conflict  with  the  spectroscopic  estimate ($T_{\\rm eff}  \\sim 24\\, 900$ K).   Our results are  in agreement with the  recent asteroseismic  analysis  of BK{\\O}11  on  KIC 8626021,  in particular  regarding  its  effective  temperature.   In  fact,  these authors  conclude that  KIC  8626021  \\emph{must be}  a  hot DBV  with $T_{\\rm eff} \\sim 29\\,200$ K. It would be interesting to see what a spectroscopic analysis based on higher signal-to-noise spectra will tell about the surface gravity and effective temperature of the star.  If KIC 8626021 is a hot  DBV, as first found by BK{\\O}11 and confirmed now  by our  results,  then  further monitoring  of  KIC 8626021  with \\emph{Kepler} in the  next years probably will allow  a measurement of $\\dot{\\Pi}$, which in turn  could open the possibility to constrain the plasmon neutrino emission rate (Winget et al. 2004; BK{\\O}11). This endeavour is evidently dependent on the models, which can always be improved. Uncertainties in the models can also be assessed (Bischoff-Kim et. al. 2008)."
    },
    "1112/1112.0005_arXiv.txt": {
        "abstract": "{ Gravitationally lensed quasars can be used as powerful cosmological and astrophysical probes. We can (i) infer the Hubble constant $H_{0}$ based on the so-called time-delay technique,  (ii) unveil substructures along the line-of-sight toward distant galaxies, and (iii) compare the shape and the slope of baryons and dark matter distributions in the inner regions of galaxies. To reach these goals, we need high-accuracy astrometry of the quasar images relative to the lensing galaxy and morphology measurements of the lens. In this work, we first present new astrometry for 11 lenses with measured time delays, namely, JVAS~B0218+357, SBS~0909+532, RX~J0911.4+0551, FBQS~J0951+2635, HE~1104-1805, PG~1115+080, JVAS~B1422+231, SBS~1520+530, CLASS~B1600+434, CLASS~B1608+656, and HE~2149-2745. These measurements proceed from the use of the Magain-Courbin-Sohy (MCS) deconvolution algorithm applied in an iterative way (ISMCS) to near-IR HST images.  We obtain a typical astrometric accuracy of about 1-2.5 mas and an accurate shape measurement of the lens galaxy. Second, we combined these measurements with those of 14 other lensing systems, mostly from the COSMOGRAIL set of targets, to present new mass models of these lenses. The modeling of these 25 gravitational lenses led to the following results: 1) In four double-image quasars (HE0047-1746, J1226-006, SBS~1520+530, and HE~2149-2745), we show that the influence of the lens environment on the time delay can easily be quantified and modeled, hence putting these lenses with high priority for time-delay determination. 2) For quadruple-image quasars, the difficulty often encountered in reproducing the image positions to milli-arcsec accuracy (astrometric anomaly problem) is overcome by explicitly including the nearest visible galaxy/satellite in the lens model. However, one anomalous system  (RXS~J1131-1231) does not show any luminous perturber in its vicinity, and three others (WFI~2026-4536, WFI~2033-4723, and B2045+265) have problematic modeling. These four systems are the best candidates for a pertubation by a dark matter substructure along the line-of-sight. 3) We revisit the correlation between the position angle (PA) and ellipticity of the light and of the mass distribution in lensing galaxies. As in previous studies, we find a significant correlation between the PA of the light and of the mass distributions. However, in contrast with these same studies, we find that the ellipticity of the light and of the mass also correlate well, suggesting that the overall spatial distribution of matter is not very different from the baryon distribution in the inner $\\sim$5\\,kpc of lensing galaxies. This offers a new test for high-resolution hydrodynamical simulations. } ",
        "introduction": "The measurement of the time delay between the lensed images of a gravitationally lensed quasar offers one of the most elegant ways to measure the \\textit{Hubble constant}, $H_{0}$ \\citep{Refsdal1964}. This cosmological application of strongly lensed quasars motivated many of the early gravitational lensing studies and time-delay measurement campaigns. Unfortunately,  the precision on $H_{0}$ obtained with the time delay method has been, until now, challenged by systematic errors. Recently, \\citet{Suyu2009, Suyu2010} showed that a detailed study of the lensed quasar CLASS~B1608+656, combining different high-resolution data and advanced lens modeling techniques, could lead to an estimate of $H_{0}$  with very small random and systematic uncertainties: $H_{0}= 70.6 \\pm 3.1 \\ \\rm km \\ \\rm s^{-1} \\rm Mpc^{-1}$. The determination of $H_{0}$ based on a large sample of lensed quasars can also reach a precision competitive with the one of standard methods \\citep{Oguri2007,Coles2008}.  In order to use the time-delay method as a precision-cosmology tool, it is necessary to derive accurate time-delay measurements and to properly understand the uncertainties associated to the lens environment and to the lens mass distribution. The COSMOGRAIL project, i.e.  the COSmological MOnitoring of GRAvItational Lenses, aims to provide the community with exquisite lightcurves and accurate time-delay measurements for a sample of more than 20 lensed quasars \\citep{Cosmograil1, Cosmograil5, Cosmograil7, Courbin2011}. To understand degeneracies affecting lens models, we started to model all systems monitored by COSMOGRAIL as well as those monitored by earlier campaigns in a uniform way. Because the relative astrometry of the lensed images is the main constraint on the lens models, we first worked on deriving an updated astrometry of these systems by applying the \\textit{Iterative Strategy coupled with the MCS{\\footnote{MCS is an abbreviation build based on the name of the three designers of the method, Magain, Courbin, and Sohy.}} deconvolution algorithm} \\citep{MCS98}, i.e. ISMCS \\citep{Chantry2007} to HST images. This allows us to obtain accurate relative astrometry and morphological information of the lensing galaxy (i.e. ellipticity and position angle PA of its major axis). We presented the first part of this work in \\citet[][ hereafter Paper I]{Chantry2010}, where seven systems currently monitored by COSMOGRAIL were analyzed. We present here the deconvolution of a sample of an additional eleven lensed quasars with published time delays: JVAS~B0218+357, SBS~0909+532, RX~J0911.4+0551, FBQS~J0951+2635, HE~1104-1805, PG~1115+080, JVAS~B1422+231, SBS~1520+530, CLASS~B1600+434, CLASS~B1608+656, and HE~2149-2745.  High-accuracy relative astrometry of lensed quasars allows one to test the ability of a smooth mass model of the galaxy to reproduce the observed image configuration. The failure of a smooth lens model to do this, which means that we identify an astrometric anomaly, may be caused by a perturbation of the potential by dark matter clump(s) along the line-of-sight, and likely located in the halo of the main lens. During the past five years, the occurrence and amplitude of astrometric anomalies caused by dark matter clumps has been questioned and their use to study the amount of substructures in galaxies has been investigated \\citep[][ Paper I and references therein]{Zackrisson2009}. We use the seven quads studied here and in Paper I plus seven other quads with astrometry derived with ISMCS or VLA+HST imaging, to revisit the evidence for astrometric anomalies. We also  quantitatively address the role of the nearby lens environment in producing these anomalies.  The deconvolution method we used also provides accurate morphological information on the lensing galaxies. With these data, we compared the mass and the light distribution in these systems. This problem has first been approached by \\citet{Keeton1998} who compared the mass and the light distribution for 14 multiply-imaged quasars (seven doubles and seven quads). More recently, \\citet{Ferreras2008} and \\citet{Treu2009} studied a sample of nine (resp. 25) lensing galaxies from the SLACS sample \\citep{Bolton2006}. Although these works probe different populations of lenses, the ``SLACS lenses'' lying in typical elliptical galaxy environments while lensed quasars often lie in richer environments \\citep{Huterer2005, Oguri2005b, Suyu2009}, they find that the PA of the light and of the mass do agree within 10 degrees. Conversely, no correlation has been found between the ellipticity of the light and of the mass distributions in \\citet{Keeton1998} and in \\citet{Ferreras2008}.  The outline of the paper is the following. Section \\ref{material11} presents the eleven lensed systems deconvolved with the ISMCS technique. The deconvolution method and the derived astrometry are described in Sect. \\ref{dec11}. Section~\\ref{sec:model} explains the mass modeling strategy. We discuss the lens models for doubly imaged quasars in Sect.~\\ref{sec:doubles}. In Sect.~\\ref{sec:quads} we present an enlarged sample of quadruply imaged systems for which we provide additional mass models and a comparison of the shape of the mass and of the light distribution. Finally, our conclusions are presented in in Sect. \\ref{Conc11}. ",
        "conclusions": "\\label{Conc11}  We applied the iterative strategy coupled with the MCS deconvolution algorithm, i.e. ISMCS, to high-resolution near-IR images of a sample of eleven lensed quasars. The ISMCS allowed us to obtain accurate astrometry, in most cases with 1 to 2.5 mas error bars, and shape parameters for the light distribution of the lensing galaxy. In three cases, the deconvolution process was not entirely satisfactory, i.e. HE~1104-1805, PG~1115+080 and CLASS~B1608+656, but still provided better results than simple PSF subtraction. For one special case, i.e. JVAS~B0218+357, the small angular separation and the poor resolution of NICMOS prevents us from reaching the desired accuracy. For one system, SBS~0909+532, we were able to detect the lensing galaxy for the first time, which is angularly small and relatively bright. Two objects reveal partial Einstein rings (PG~1115+080, CLASS~B1608+656).  We provided, for these systems, simple mass models (isothermal and de Vaucouleurs mass distributions) that give an overview of the ability of smooth mass models to reproduce the known lensed quasars. These models give some hints about the effect of the radial profile on the flux ratios and on the time delays. For the sample of ten doubly imaged quasars presented here and in Paper I  \\citep[i.e.][]{Chantry2010}, we have also studied the role of intrinsic ellipticity and external shear in the models. For half of these systems{\\footnote{HS~2209+1914 is more uncertain because of the lack of detailed environmental studies}} (HE0047-1746, J1226-006, SBS~1520+530, HE~2149-2745, and HS~2209+1914) one should be able to unfold intrinsic and extrinsic shear and therefore minimize systematic biases on $H_0$ based on the time-delay method. This identification of the most promising systems is partly subjective because detailed environmental studies of the whole sample are lacking. Exhaustive studies of lens environment are fortunately ongoing \\citep{Williams2006, Momcheva2006, Wong2011} and should allow one to address this question more quantitatively.  In addition to the ten doubly imaged quasars, we have performed modeling with singular isothermal ellipsoid (SIE) of a sample of 14 quadruply imaged systems for which we have relative astrometry with accuracy typically better than 2 mas. This sample allowed us to investigate two questions:  \\begin{enumerate} \\item {\\it {How efficient are SIE models in reproducing the image and galaxy position of the observed quads and how does the local environment of the lens influence this result ? }} We found that the environment of most of the quads is rich and that at least the nearest galaxy or a nearby group has to be included explicitly in the model to recover the image and galaxy positions. Following this procedure leads to a formally perfect fit with an SIE model of the main lens for 11 of the 14 systems. When this perturber is not included, models generally tend to predict the position of the main lensing galaxy significantly offset with respect to the observed position, the lensed image positions being less affected. For the two systems RXS~J1131-1231 and WFI~2033-4723, we were not able to obtain an acceptable fit. The problems encountered with WFI~2033-4723 might be caused by other nearby galaxies not included explicitly in our model.  RXS~J1131-1231 remains the best candidate of our sample for which astrometric anomalies may be caused by a dark matter substructure. Two other systems (WFI~2026-4536 and B2045+265) show a significant mismatch between the $\\theta_e$ of the light and of the mass. The origin of this discrepancy is unclear and might also be the signature of dark matter substructures. We mention that astrometric anomalies may become proeminent when higher accuracy images are available. This is the case for B0128+437, which is perfectly reproduced by our model but shows astrometric anomalies once sub-mas astrometry of the lensed images is considered \\citep{Biggs2004}. \\\\ \\item {\\it {How do the morphology of the mass and of the light profile compare ?}} Our analysis of 12 out of 14 quads with measured structural parameters allowed us to find that the PA of the mass and of the  light distribution are aligned within typically 10 degrees, in agreement with previous studies \\citep{Keeton1998, Ferreras2008, Treu2009}. We also discovered that the ellipticity of the light generally agrees well with that of the total mass distribution (8 out of 12 lenses show a good correlation). Overall, this indicates that the light provides a good proxy to the total mass distribution in these lensing galaxies{\\footnote{The radial distribution of the total mass disagrees however with that of the light, as illustrated by the inability of deVaucouleurs mass profiles to model quads.}}, a result particularly interesting in the light of the rich environment of the lenses. Understanding why these correlations take place is important for both our comprehension of dark matter and galaxy formation and for time-delay studies. Therefore, an increase in the sample of lenses with shape measurements and a comparison with numerical simulations, including problems caused by selection biases \\citep{Vandeven2009, Mandelbaum2009}, are necessary. This will provide important clues on the role of dark matter in shaping galaxies. \\\\  \\end{enumerate}"
    },
    "1112/1112.3006_arXiv.txt": {
        "abstract": "Cosmological data have provided new constraints on the number of neutrino species and the neutrino mass. However these constraints depend on assumptions related to the underlying cosmology. Since a correlation is expected between the number of effective neutrinos $N_{eff}$, the neutrino mass $\\sum m_\\nu$, and the curvature of the universe $\\Omega_k$, it is useful to investigate the current constraints in the framework of a non-flat universe. In this paper we update the constraints on neutrino parameters by making use of the latest cosmic microwave background (CMB) data from the ACT and SPT experiments and consider the possibility of a universe with non-zero curvature. We first place new constraints on $N_{eff}$ and $\\Omega_k$, with $N_{eff} = 4.03 \\pm 0.45$ and $10^3 \\, \\Omega_k = -4.46 \\pm 5.24$. Thus, even when $\\Omega_k$ is allowed to vary, $N_{eff} = 3$ is still disfavored with 95\\% confidence. We then investigate the correlation between neutrino mass and curvature that shifts the $95 \\%$ upper limit of  $\\sum m_\\nu < 0.45$ eV to $\\sum m_\\nu < 0.95$ eV. Thus, the impact of assuming flatness in neutrino cosmology is significant and an essential consideration with future experiments. ",
        "introduction": "Throughout the previous decades experimental cosmology has benefited from accurate measurements of the cosmic microwave background (CMB). The data have determined constraints on several cosmological parameters to remarkable accuracy and the ability to constrain new physics with the CMB continues to improve. Future CMB experiments might even be able to measure B-mode polarization and distinguish between neutrino hierarchy models. However, when constraining new parameters one must be careful when constraints depend on assumptions about the underlying cosmology. For example, a correlation between the neutrino properties and the curvature of the universe is clearly expected since, a higher number of neutrino species  or large mass would introduce pre-recombination effects, shifting the positions of the peaks in the angular CMB spectrum (cf.~\\cite{cormelk,Mantz:2009rj}).  Here we present an update on the constraints of the number of neutrino species $N_{eff}$ and the sum of neutrino masses $\\Sigma m_{\\nu}$ in the framework of non-flat universes with  $\\Omega_k \\neq 0$  combining the Wilkinson Microwave Anisotropy Probe (WMAP) 7-year~\\cite{Komatsu:2010fb}, South Pole Telescope (SPT)~\\cite{Keisler:2011aw} and Atacama Cosmology Telescope (ACT)~\\cite{Dunkley:2010ge} datasets. %    The paper is organized as follows. In Section~\\ref{sec:theory} we % give theoretical arguments for why the $N_{eff}$ and $\\Omega_k$ % parameters should be correlated. In Section~\\ref{sec:method} we discuss our method of constraining the parameters $N_{eff}$, $\\Sigma m_{\\nu}$, and $\\Omega_k$. We present the results of the analysis in Section~\\ref{sec:imp}. Finally, in Section~\\ref{sec:conclusion} we conclude and discuss the implications of assuming flatness in neutrino cosmology. ",
        "conclusions": "\\label{sec:conclusion}  The resolution of the high effective neutrino number in cosmology remains an open question. However, additional neutrinos % may be due to parameter degeneracy or other issues in statistical analysis rather than new physics. % The focus of this paper has been an argument for correlation between the number of effective neutrinos $N_{eff}$ and the curvature of the Universe $\\Omega_k$, which arises from the effect of these parameters on distance measurements. The qualitative argument is confirmed by a statistical analysis of CMB anisotropy measurements using CosmoMC.  In this paper we have shown that there is a correlation between $N_{eff}$ and $\\Omega_k$ that gets stronger when SPT and ACT datasets are added to WMAP alone.  However, even when $\\Omega_k$ is % allowed to vary, $N_{eff} = 3$ is still disfavored by the data with 95\\% confidence. Although the correlation favors a closed universe with $\\Omega_k < 0$, % if CMB data were to favor open models then the neutrino number would decrease as predicted. Perhaps % the same element of the data that favors a closed universe may also be responsible for the trend toward a higher $N_{eff}$. More importantly, % we find a strong correlation between curvature and % the sum of the neutrino masses. %  Future experiments will provide further insight into both $N_{eff}$ and $\\sum m_\\nu$~\\cite{Carbone:2010ik}. Our results are consistent with the current understanding of the data available. The strongest constraints on these parameters from the statistical analysis assuming a flat universe are $N_{eff} = 3.89 \\pm 0.41$ and $\\sum m_\\nu < 0.45$ eV with 95\\% confidence level using WMAP7+ACT+SPT+BAO+$H_0$. The constraints are weakened by % degeneracy with the curvature parameter $\\Omega_k$. However, this still represents the continued effort toward significant improvements on parameter constraints in cosmology. Although the sum of the neutrino masses is significantly improved from the WMAP 7-year result of $\\sum m_\\nu < 0.57$ eV, the constraint is far from being sensitive enough to rule out one of the mass hierarchies. Furthermore, we have shown that the mass uncertainty more than doubles when $\\Omega_k \\neq 0$. Based on our results and the estimated quality of data for Planck and other experiments, it should be possible to determine the existence or nonexistence of sterile radiation to much greater confidence in the near future."
    },
    "1112/1112.2882_arXiv.txt": {
        "abstract": "{ \\object{Comet P/2010~A2 LINEAR} is an object on an asteroidal orbit within the inner Main Belt, therefore a good candidate for membership with the Main Belt Comet family. It was observed with several telescopes (ESO New Technology Telescope, La Silla, Chile; Gemini North, Mauna Kea, Hawai`i; University of Hawai`i 2.2~m, Mauna Kea, Hawai`i) from 14~Jan. until 19~Feb. 2010 in order to characterize and monitor it and its very unusual dust tail, which appears almost fully detached from the nucleus; the head of the tail includes two narrow arcs forming a cross. No evolution was observed during the span of the observations. Observations obtained during the Earth orbital plane crossing allowed an examination of the out-of-plane 3D structure of the tail. The immediate surroundings of the nucleus were found dust-free, which allowed an estimate of the nucleus radius of 80--90~m, assuming an albedo $p=0.11$ and a phase correction with $G=0.15$ (values typical for S-type asteroids). A model of the thermal evolution indicates that such a small nucleus could not maintain any ice content for more than a few million years on its current orbit, ruling out ice sublimation dust ejection mechanism. Rotational spin-up and electrostatic dust levitations were also rejected, leaving an impact with a smaller body as the favoured hypothesis. This is further supported by the analysis of the tail structure. Finston-Probstein dynamical dust modelling indicates the tail was produced by a single burst of dust emission. More advanced models (described in detail in a companion paper), independently indicate that this burst populated a hollow cone with a half-opening angle $\\alpha\\sim 40\\degr$ and with an ejection velocity $v_{\\rm max}\\sim 0.2$~m~s$^{-1}$, where the small dust grains fill the observed tail, while the arcs are foreshortened sections of the burst cone. The dust grains in the tail are measured to have radii between $a=1$--20~mm, with a differential size distribution proportional to $a^{-3.44\\pm0.08}$. The dust contained in the tail is estimated to at least $8\\times 10^8$~kg, which would form a sphere of 40~m radius (with a density $\\rho=3\\,000$~kg~m$^{-3}$ and an albedo $p=0.11$ typical of S-type asteroids). Analysing these results in the framework of crater physics, we conclude that a gravity-controlled crater would have grown up to $\\sim 100$~m radius, i.e. comparable to the size of the body. The non-disruption of the body suggest this was an oblique impact. ",
        "introduction": "Habitability within our solar system is determined by the distribution of water and volatiles, yet the origin of this distribution is currently a fundamental unresolved planetary science issue \\citep{KC03_habitable}. There are three leading scenarios for the origin of terrestrial planetary water, including: ($i$) nebular gas adsorption on micron-sized dust grains inside the snow line \\citep[][the distance from the Sun outside of which water the temperature is low enough for water ice to condense]{Mur+08_water}, ($ii$) chemical reactions between an early hydrogen envelope and oxides in a magma ocean \\cite{GI07_DHratio, GI08_ocean}, or ($iii$) delivery by volatile-rich planetesimals (asteroids or comets) formed beyond the snow line \\citep{Mor+00_water}. The first two processes probably contributed to, but may be unable to account for all of Earth's water \\citep{Mot+07_water}. Within the broader context of icy bodies, main belt comets (MBCs), a newly discovered class of ``comets having stable orbits completely confined to the main asteroid belt'' \\citep{HsiehJ06_MBC}, present a subclass of particular significance to the history of terrestrial water and other important volatiles.  In addition to possibly P/2010~A2 (LINEAR), there are currently five known MBCs: 133P/Elst-Pizarro, 176P/LINEAR, 238P/Read, C/2008 R1 (Garradd) and P/2010 R2 (La Sagra).  The main belt comets are defined \\citep{HsiehJ06_MBC} by ($i$) a semi-major axis that is less than Jupiter's, ($ii$) Tisserand parameters significantly greater than 3 ---meaning they are dynamically decoupled from Jupiter, like ordinary asteroids \\citep{Vaghi73, Kresak80}--- and ($iii$) mass loss with a cometary appearance. As comets in near-circular orbits within the asteroid belt, these objects likely still harbor nebular water frozen out from beyond the primordial ``snow line'' \\citep{Encrenaz08, GaraudL07_snow} of the young solar system.  Dynamical simulations suggest they probably formed in-situ \\citep{Haghighipour09_MBC} at a different temperature from either comets or the asteroidal reservoir that has been sampled through our meteorite collections.  Thus, understanding their chemistry can provide unique insight into the distribution of volatiles in the early stages of planet formation.  P/2010 A2 (LINEAR) was discovered by the LINEAR project on 7~Jan., 2010 \\citep{CBET2114} and was described as ``a headless comet with a straight tail, and no central condensation'' \\citep{IAUC9105}. On 11--12 Jan., observers at the WYNN 3.5-m and at the 2.5-m Nordic Optical Telescope reported that the object had a asteroidal-like body $\\sim$ 150--200m in diameter which was connected to the tail by an unresolved light bridge \\citep{IAUC9109}. Based on the orbital semi-major axis and on the Tisserand parameter, $T_J=3.6$, Jewitt \\citep{IAUC9109} concluded that P/2010~A2 was a new MBC. With the smallest perihelion distance of any of the MBCs ($q=2.29$~AU), these elements suggested a membership in the Flora family. Flora is a large asteroid family of $\\sim$500 members which can be broken up into many sub-families or clans and which are broadly compositionally consistent with space-weathered S-type asteroidal spectra \\citep{Flo+98}.  Jewitt \\citep{IAUC9109} and Licandro \\citep{CBET2134} suggested that the location of the nucleus outside the coma might be the consequence of an impact.  Is P/2010~A2 a genuine, sublimating comet (whose activity was possibly triggered by an impact), or is the tail the signature of an impact (or another alternative process), with a dispersion of the dust ejecta but no sublimation? If it can be demonstrated that the object is a comet, it would indicate that at least one object in the inner asteroid belt still contains volatile material. The water snow line in the proto-solar nebula is estimated to have been around 2.5~AU from the Sun \\citep{Jon+90} or possibly even as close as 1.5~AU \\citep{Lec+06,Mac+10,Min+11}. Initially parent bodies beyond the snow line are believed to be mixtures of water ice and silicates, which are then heated due to the radioactive decay of $^{26}$Al within 1--2 million years after nebular collapse \\citep{Krot+06}. \\citet{Grimm+89} find that once the water is liquid, it is consumed by hydration reactions, preferentially in the interior \\citep{Cohen+00,Wilson+99}, possibly leaving ice in the outer layers. In the inner asteroid belt, it is likely that most of that water has been converted in hydrated material observed on S-type asteroids \\citep{Riv+02}. MBCs may be the frozen components of the outer edges of asteroid parent bodies that have survived the age of the solar system.  Four of the MBCs are located in the outer main belt, where asteroids with hydrated material are less common, suggesting water ice could still exist.  The orbital elements of the MBCs are listed in Table~\\ref{tab:ele} and displayed in Fig.~\\ref{fig:ele}. 133P/Elst-Pizarro, the first MBC discovered, is a member of the Themis family; 176P/LINEAR was found in a survey targeting objects with orbits similar to that of 133P/Elst-Pizarro \\citep{HHH09}.  238P/Read, however, was discovered serendipitously (i.e. not within a MBC-dedicated program) on a similar orbit. A fourth MBC, P/2008~R1 Garradd, was discovered serendipitously in the central region of the main belt. The fifth MBC, P/2010~R2 La Sagra, was also discovered serendipitously with a semi-major axis similar to that of 133P/Elst-Pizarro \\citep{Mar+10}. More recently, asteroid (596)~Scheila presented a comet-like appearance \\citep{Lar10}. \\citet{Bod+11} and \\citet{Jew+11} conclude however that this dust cloud was likely caused by the impact of a small asteroid on Scheila. The survival of an icy body in the inner asteroid belt would therefore give exciting constraints on the evolution of water in the belt. It would also raise fundamental questions on how to shield water ice in that area of the solar system in a rather small body. \\begin{figure} \\caption{(not available on astroPh) The orbital elements of P/2010~A2 (green square to the left), of the 5 known MBCs (circles, 133P and 238P are indistinguishable) and of the numbered main belt asteroids (small dots). When available, the proper elements were used (source: the Asteroid Dynamics Site, http://hamilton.dm.unipi.it/astdys/). The inclination of P/2010~R2 La Sagra is off scale; its $a$ is marked by a triangle on the top of the plot; $i=21.39$ . } \\label{fig:ele} \\end{figure}   \\begin{figure}[h] {\\bf a.} \\includegraphics[width=8.8cm]{fig_2010-01-14_NTT_OH.eps}\\\\ {\\bf b.} \\includegraphics[width=8.8cm]{fig_2010-01-15_NTT_HH.eps}\\\\ {\\bf c.} \\includegraphics[width=8.8cm]{fig_2010-01-16_NTT_HH.eps}\\\\ {\\bf d.} \\includegraphics[width=8.8cm]{fig_2010-01-17_NTT_NN.eps}\\\\ {\\bf e.} \\includegraphics[width=8.8cm]{fig_2010-01-19_GN_OH.eps} \\caption{P/2010~A2, images from UT~14.2 (a), 16.2 (b), 17.2 (c), 18.2 (d) Jan. 2010 using the NTT, UT~19.5 Jan. 2010 using GN (e). The linear gray scale covers the range of (0--3) $\\times 10^{-8} Af$, a proxy for the amount of dust present (see Section~\\ref{sec:dust}). Each panel includes a 40$\\times 30''$ zoomed portion to show details of the inner structure. The positions of N and E are indicated, as are the anti-solar direction and the heliocentric velocity vector.  } \\label{fig:image}% \\end{figure} \\addtocounter{figure}{-1} \\begin{figure}[h] {\\bf f.}  \\includegraphics[width=8.8cm]{fig_2010-01-22_UH_BY.eps}\\\\ {\\bf g.}  \\includegraphics[width=8.8cm]{fig_2010-01-23_UH_BY.eps}\\\\ {\\bf h.}  \\includegraphics[width=8.8cm]{fig_2010-01-25_UH_BY.eps}\\\\ {\\bf i.}  \\includegraphics[width=8.8cm]{fig_2010-02-02_GN_OH.eps}\\\\ {\\bf j.}  \\includegraphics[width=8.8cm]{fig_2010-02-19_UH_JP.eps} \\caption{continued. P/2010~A2, images from UT~22.3 (f), 23.4 (g), 25.4 (h) Jan. 2010 using UH~2.2-m, UT~2.3 Feb. 2010 using GN (i), UT~19.5 Feb. 2010 using UH~2.2-m (j).} \\end{figure}  \\citet{Mor+10} presented observations obtained with the GTC, WHT and NOT on La Palma, which they modeled with an extended period of water-driven cometary activity. \\citet{Jew+10} acquired a series of HST images; from the orientation and geometry of the tail, they favor the disruption of an asteroid (either by collision or spin up) in Feb.--Mar. 2009. Finally, \\citet{Sno+10} secured observation from the Rosetta spacecraft. Thanks to the position of the space probe, they could observe the object with a very different geometry (but with a much more modest resolution), from which they concluded that the observed dust tail was caused by an impact.     In order to investigate the process that generated the observed dust tail, we acquired deep images of P/2010~A2 over various epochs. The observations are presented in Section~\\ref{sec:obs}. The analysis and modelling of the nucleus and dust are described in Section~\\ref{sec:ana}.  In Section~\\ref{sec:con}, we discuss the conclusions and summarize the results.  A companion paper \\citep{PAPER2} is devoted to the details of the dust models developed for this study, and a second follow-up article \\citep{PAPER3} focuses on the characteristics of the impact process and of the nucleus based on the interpretation of the dust as an impact plume. ",
        "conclusions": "\\label{sec:con}   The dust tail of P/2010~A2 was observed and analyzed so to cast some light on the possible cometary nature of that object, as its appearance and orbit  could make it a member of the small family of Main Belt Comets. The main results and conclusions of this study are listed hereafter.   \\begin{itemize} \\item Observations were acquired with GN, NTT, UH~2.2-m from 14 Jan. until 19 Feb. 2010. % The tail did not show any measurable evolution over that time span, which included the crossing of the orbital plane, which indicates the tail is an out-of-plane 3D structure, suggesting the dust was released with a significant orthogonal velocity. \\item The nucleus, detached from the tail, appears not to be surrounded by dust (at most 3\\% of the light corresponds to near-nucleus dust); dynamical dust models also indicate that no recent dust is present, suggesting no cometary activity at the time of the observations. The absolute magnitude of the comet was estimated to SDSS $r'(1,1,0)=21.74\\pm0.04$ on 19~Jan. 2010 and $21.55\\pm0.05$ on 2~Feb. Assuming an albedo $p=0.11$, a value typical for S-type asteroids, this converts into a nucleus of radius $r=80$--90~m. Using a density of $\\rho=3$\\,000~kg~m$^{-3}$ (also typical of a S-type asteroid), the escape velocity of the body is $v_e=0.10$--0.12~m~s$^{-1}$. \\item A thermal model of the nucleus indicates that water ice (as well as any more volatile ice) would not survive more than a few million years in the object on its current orbit. As we have no reason to suspect a recent insertion on the current orbit, this rules out ice sublimation as the source of the observed dust. Rotational spin up and electrostatic lifting were also rejected as possible source for the dust tail. \\item A Finston-Probstein-type dynamical dust modelling of the morphology, light distribution and detachment of the tail indicates that a short burst of dust emission (a day long or less) represents best the observed tail, with an ejection velocity in the range $v_e=0.20$--0.30~m~s$^{-1}$. The ejection date is in agreement with other studies of this object \\citep{Jew+10, Sno+10}; we continue this study using Snodgrass' date, 10 Feb. 2009. The model, together with direct measurements of the tail within the camera's field of view, indicate it is constituted of dust grains with radii $a$ ranging between a few tenths of millimetres to 20~mm, with a number of particles proportional to $a^{-3.44\\pm0.08}$ {(differential size distribution index).} The total dust content of the tail is estimated to at least $8\\times 10^8$~kg, which could be packed in a sphere with a radius of 40~m (assuming the same density $\\rho=3$\\,000~kg~m$^{-3}$ as above). \\item The position and shape of the tail's envelope was modeled as originating from dust emitted on a cone with a half-opening angle $\\alpha \\sim 40\\degr$ pointing below the orbital plane, toward the forward direction, with a ejection velocity $v_e \\le 0.20$~m~s$^{-1}$. The X-shaped arcs were modeled as a collection of large particles, with $v_e \\ge 0.20$~m~s$^{-1}$; they were independently found to be emitted on a section of the same hollow cone. The are seen as an arc as a consequence of the geometric foreshortening of the cone section (after one year of evolution); other sections of the cone, which are not foreshortened in the same way, contribute to the overall tail. The Rosetta observations \\citep{Sno+10} did not see other sections of the cone as arcs because of the coarser resolution of its camera. \\item Because ice sublimation, electrostatic lifting and rotational spin up were rejected as possible causes for the dust emission, and because of the very short release time and the geometry of the dust emission on a hollow cone with an half-opening angle $\\alpha = 40\\degr$, we conclude that the dust release was caused by an impact by a small object. \\item Considering the volume of ejected dust in the framework of a gravity-dominated crater formation process, a crater of $\\sim 100$~m radius was produced in a total time of the order of seconds. While that crater is large compared to the body, other bodies show craters as large as themselves. An oblique impact ---which is statistically the most likely--- would explain how P/2010~A2 escaped complete, catastrophic disruption.   \\end{itemize}  In summary, the tail of P/2010~A2 can be explained by an impact on a $\\sim 80$--90~m asteroid. The impact released dust in a hollow cone, which evolved for about one year to form the observed dust distribution.  The fine-grained dust formed a long tail that was extended by solar radiation pressure, and the large grains (un-affected by the solar radiation pressure) retained the conical shape. The viewing geometry revealed sections of that cone as the arcs in the head of the dust cloud. Therefore, what we witnessed was likely an event of collisional grinding within the Flora collisional family, and not a genuine, water-ice driven Main Belt Comet. With increasingly sensitive sky surveys covering larger fractions of the sky during each dark period, we can expect to detect more and more of these collisions, as illustrated by the recent discovery of the dust cloud around (596)~Scheila. While the original goal of this project was to probe the water ice content in the asteroid main belt, its end result was the observation of a natural impact on a small asteroid."
    },
    "1112/1112.0555_arXiv.txt": {
        "abstract": "The ultrahigh energy neutrino cross section is a crucial ingredient in the calculation of the event rate in high energy neutrino telescopes. Currently there are several approaches which predict different behaviors for its magnitude for ultrahigh energies. In this contribution is presented a summary of current predictions based on the non-linear QCD evolution equations, the so-called perturbative saturation physics. In particular, predictions are shown based on the parton saturation approaches and the consequences of  geometric scaling property at high energies are discussed. The scaling property allows an analytical computation of the neutrino scattering on nucleon/nucleus at high energies, providing a theoretical parameterization. ",
        "introduction": "The investigation of ultrahigh energy (UHE) cosmic neutrinos provides an opportunity for study particle physics beyond the reach of the LHC \\cite{neu_review}. As an example, nowadays the Pierre Auger Observatory  is sensitive to neutrinos of energy $\\ge 10^8$ GeV \\cite{pao}. A crucial ingredient in the calculation of attenuation of neutrinos traversing the Earth and the event rate in high energy neutrino telescopes is the  high energy neutrino-nucleon cross section, which  provides a probe of Quantum Chromodynamics (QCD) in the kinematic region of very small values of Bjorken-$x$. The typical $x$ value probed is $x \\approx m_W^2/2m_NE_{\\nu}$, which implies that for $E_{\\nu} \\approx 10^8 - 10^{10}$ GeV one have $x \\approx 10^{-4} - 10^{-6}$ at $Q^2 \\approx 10^4$ GeV$^2$. This kinematical range was not explored by the HERA measurements of the structure functions \\cite{hera}.  The description of QCD dynamics in such very high energy limit is a subject of intense debate \\cite{hdqcd}. Theoretically, at high energies (small Bjorken-$x$)  one expects the transition of the regime described by the linear dynamics, where only the parton emissions are considered, to a new regime where the physical process of recombination of partons becomes important in the parton cascade and the evolution is given by a non-linear evolution equation.  This regime is characterized by the limitation on the maximum phase-space parton density that can be reached in the hadron wavefunction (parton saturation), with the transition being specified  by a typical scale, which is energy dependent and is called saturation scale $Q_{\\mathrm{s}}$. Moreover,  the growth of the parton distribution is expected to saturate, forming a  Color Glass Condensate (CGC), whose evolution with energy is described by an infinite hierarchy of coupled equations for the correlators of  Wilson lines \\cite{hdqcd}. In the mean field approximation, the first equation of this  hierarchy decouples and boils down to a single non-linear integro-differential  equation: the Balitsky-Kovchegov (BK) equation \\cite{BAL,KOVCHEGOV}. Experimentally, possible signals of parton saturation have already been observed both in  $ep$ deep inelastic scattering at HERA and in deuteron-gold collisions at RHIC \\cite{hdqcd}.   Currently, there are predictions of the neutrino nucleon cross sections with structure functions constrained by HERA data are based on linear  dynamics \\cite{ccs}, using DGLAP or an unified DGLAP/BFKL evolution, or phenomenological models that resembles the expected behavior predicted by  the non-linear QCD dynamics \\cite{bhm} (i.e., the proton structure function saturating the Froissart bound at asymptotic energies, $F_2^p \\propto \\ln^2 (1/x)$). As a general feature, the nonlinear QCD dynamics predicts sizable suppression of UHE neutrino cross section in comparison with standard approaches. Here, we summarize the main results of works presented in Refs. \\cite{magno1,magno2}. In Ref. \\cite{magno1} the contribution of non-linear effects  was estimated considering distinct phenomenological models based on saturation physics. An update on those calculations have been done recently \\cite{victor}. In Ref. \\cite{magno2}, the geometric scaling property (which is a natural consequence of the asymptotic solutions of the nonlinear QCD evolution equations) is considered to obtain an analytical parameterization for the UHE neutrino cross sections. In what follows we introduce the theoretical and phenomenological tools and present the main results and predictions ",
        "conclusions": ""
    },
    "1112/1112.5236_arXiv.txt": {
        "abstract": "We analyze \\textit{in situ} measurements of solar wind velocity obtained by the \\textit{Advanced Composition Explorer} (ACE) spacecraft during the solar activity cycle $23$. We calculated a robust complexity measure, the permutation entropy ($S$) of solar wind time series at different phases of a solar activity cycle. The permutation entropy measure is first tested on the known dynamical data before its application to solar wind time series. It is observed that complexity of solar wind velocity fluctuations at $1$ AU shows hysteresis phenomenon while following the ascending and descending phases of the activity cycle. This indicates the presence of multistability in the dynamics governing the solar wind velocity over a solar activity cycle. ",
        "introduction": "\\label{S-intro} The corona changes its shape enormously during a solar activity cycle resulting in temporal and structural properties of the solar wind velocity variation in a solar cycle \\cite{schw07}. \\textit{In situ} solar wind plasma observations show that its local properties at $1$~AU are modulated by the solar activity cycle \\cite{hapg91,rich08}. Around solar activity minimum, the structures of the corona and the solar wind are rather simple and remain so for several months \\cite{schw07}. At solar activity maximum, slow solar wind dominates at all helio latitudes \\cite{mcco00}. Recently it was reported that slow solar wind velocity just before the maximum of the solar activity cycle is least correlated to data obtained from the rest of the solar activity cycle \\cite{suya11}.   Observations of solar wind velocity made by different spacecraft have been analyzed in considerable detail and reported by several authors. Values of complexity measures such as entropy (\\opencite{mace97}, \\citeyear{maco98}; \\opencite{mace00}; \\opencite{reda01}), correlation dimension \\cite{mace97,gupt08} and Lyapunov exponents \\cite{mace98,reda01,gupt08} show that solar wind velocity fluctuations are a consequence of complex nonlinear dynamical processes. Inherent changes in the dynamics governing the solar wind velocity at $0.3$ AU have been observed \\cite{gupt08}. \\inlinecite{mila04} analyzed magnetic and bulk velocities measured by ACE and found an anisotropy in the velocity, magnetic, and cross helicity correlation functions and power spectra. \\inlinecite{mcco00} analyzed Ulysses observations to demonstrate that the mid-latitude solar wind structure becomes increasingly complex as solar activity increases. \\inlinecite{cons09} investigated the emergence of spatio-temporal complexity in the $11$-year solar cycle monitored by sunspot activity. They showed that spatio-temporal or dynamical complexity is an intrinsic property of the solar cycle. Using information entropy approach to the sunspot number time series, they showed how the dynamical complexity increases during the maximum phase of the solar cycle.   Hysteresis occurs in several phenomena in physics, chemistry, biology, and engineering. It is a nonlinear phenomenon observed in systems from diverse areas of science, \\textit{e.g}., electromagnetism, electro-plasticity, superconductivity, and granular motion \\cite{bert98,guye99,katz02,zhar02}. This phenomenon occurs when a nonlinear system has at least two coexisting stable states in the hysteresis region where the system is found to depend on the history of the dynamics. Here, in dynamical systems, history corresponds to the  system's initial conditions. For example, hysteresis is observed in van der Pol system, Duffing system \\cite{thom86}, and Lorenz system \\cite{alfs85}. The hysteresis has been observed in coupled nonlinear systems as well  \\cite{pras05}.   Hysteresis phenomenon has been noticed in various solar indices. \\inlinecite{bach94} observed the presence of hysteresis patterns among many pairs of activity indices during solar cycle $21$ and $22$. They found that this hysteresis can be expressed approximately as a hierarchy of delay times behind the leading index, the sunspot number. \\inlinecite{reye98} analyzed the low-degree $p$-mode frequency shifts and solar activity indices (radio flux at $10.7$ cm and magnetic index) over solar cycle $22$ and observed a hysteresis phenomenon. \\inlinecite{more00} suggested that high latitude fields are necessary to produce a significant difference in hysteresis between odd and even-degree $p$-modes frequencies. \\inlinecite{trip00} reported that the intermediate degree $p$-mode frequencies of solar cycle $22$ show a hysteresis phenomenon with the magnetic indices whereas no such effect exists for the radiative indices. \\inlinecite{ozgu01} showed the presence of hysteresis between the solar flare index and some solar activity indicators such as total sunspot area, mean magnetic field, and coronal index during solar cycles $21$ and $22$. They found that these indices follow different paths for ascending and descending phases of the solar cycles while saturation effect exists at the extreme phases.   In the present work, we attempt to understand the dynamics of ascending and descending phases of a solar cycle. We use permutation entropy $(S)$ of hourly averaged solar wind velocity time series to capture the complexity trend over a cycle. We use the data obtained from ACE during $1998-2010$. This period belongs to solar activity cycle $23$. We use permutation entropy $(S)$ to detect the hysteresis in a dynamical system and calculate it for time series obtained from simulated as well as solar wind velocity data.   In the next section, we review the algorithm to calculate the permutation entropy $(S)$ of a time series. In Section~\\ref{S-dynamical_hysteresis} we describe how $S$ detects the multistability present in the modeled dynamical system. In Section~\\ref{S-solar_wind} we analyze solar wind data using permutation entropy. This is followed by  conclusions in Section~\\ref{S-conc}. ",
        "conclusions": "The long term sunspot time series shows an average cycle length of $11$ years. Although nearly periodic, the period as well as amplitude of the cycle varies irregularly. Apart from the sunspot data, the intrinsic irregularity of the solar cycle is seen in other observable variables like surface flows, solar irradiance (solar constant), the solar wind, and so on. The Maunder Minimum and several ancient periods from solar proxy-data suggest that the Sun exhibits quasi-periodic or intermittent behavior \\cite{femi97}. Owing to changes in magnetic activity, many aspects of the solar wind change over a solar cycle, including the speed, the density, the dynamic pressure, the composition, and the temperature \\cite{rich08}.   In this paper, we present the analysis of solar wind velocity data during the solar activity cycle $23$. We obtained $18$ different time series of hourly averaged solar wind velocity, measured by ACE spacecraft at $1$ A.U. To quantify the randomness of these time series, we use a robust, conceptually simple and computationally efficient measure called permutation entropy. A smaller value of permutation entropy indicates a more regular time series. We observe that as the solar cycle $23$ progresses towards maximum, the permutation entropy increases, saturates around the peak of activity and then decreases as the activity of cycle $23$ subsides. We also note (Figure. \\ref{fig5}) that the value of permutation entropy follows different paths in the ascending and descending phases of the solar activity cycle. In addition, while the ascent is fluctuating, the descent is smooth. This hysteresis phenomenon shows the multistability in the dynamics of solar wind, over the solar activity cycle. The behavior is similar to the one observed for hysteresis phenomenon of other solar indices \\cite{bach94,reye98,ozgu01} and confirms (\\textit{cf}. \\opencite{cons09}) that spatio-temporal or dynamical complexity is an intrinsic property of the solar cycle.  \\label{S-conc} \\begin{acks} The authors thank the ACE Science Center and instrument teams for making available the ACE data used here. VS and AP thank CSIR for SRF and DST Govt. of India for financial supports respectively. \\end{acks}"
    },
    "1112/1112.4837_arXiv.txt": {
        "abstract": "{ We propose a fully nonlinear framework to construct consistency relations for testing generic cosmological scenarios using the evolution of large scale structure. It is based on the covariant approach in combination with a frame that is purely given by the metric, the normal frame. As an example, we apply this framework to the $\\Lambda$CDM model, by extending the usual first order conditions on the metric potentials to second order, where the two potentials start to differ from each other. We argue that working in the normal frame is not only a practical choice but also helps with the physical interpretation of nonlinear dynamics. In this frame, effective pressures and anisotropic stresses appear at second order in perturbation theory, even for ``pressureless'' dust. We quantify their effect and compare them, for illustration, to the pressure of a generic clustering dark energy fluid and the anisotropic stress in the DGP model. Besides, we also discuss the effect of a mismatch of the potentials on the determination of galaxy bias. } ",
        "introduction": "\\label{sec:intro}  Current cosmological observations provide evidence for a recent onset of a phase of accelerated expansion of the Universe \\cite{Riess:1998cb, Perlmutter:1998np, Komatsu:2010fb,Suzuki:2011hu,Sherwin:2011gv}. Under the assumption of large-scale homogeneity and isotropy, the observations can only be described in the framework of General Relativity (GR) by invoking some form of Dark Energy (DE) with negative pressure. Apart from the Standard Model (SM) of particle physics, the preferred model of the content of the Universe also assumes the presence of Cold Dark Matter (CDM) and DE in the form of a cosmological constant ($\\Lambda$) and is therefore named $\\Lambda$CDM. The dynamics of the Universe is based on the Einstein equations and on the assumption that the average geometry is described by a Friedmann-Lema\\^itre-Robertson-Walker (FLRW) metric. The formation of structures is then modelled using cosmological perturbation theory \\cite{Bardeen:1980kt, Kodama:1985bj, Mukhanov:1990me, Durrer:1993db, Ma:1995ey, Bernardeau:2001qr} which describes the evolution of small fluctuations in the energy-momentum content and the metric.  The $\\Lambda$CDM model comes with a minimal set of parameters that are increasingly well constrained by the data and, to date, no significant departures from its predictions on cosmological scales have been found \\cite{Komatsu:2010fb, Bean:2010zq, Zhao:2010dz}. However, the difficulty in explaining the required tiny value of $\\Lambda$ has motivated many alternative scenarios with dynamical DE fields or modifications to GR aka Modified Gravity (MG)~\\cite{Wetterich:1987fm, Ratra:1987rm, ArmendarizPicon:1999rj, Dvali:2000hr, Deffayet:2001pu, Copeland:2006wr, Durrer:2008in, Tsujikawa:2010zza, Clifton:2011jh, Capozziello:2011et,Kunz:2012aw}. Given the existing tight constraints, it is essential  to understand very accurately what the theoretical predictions are and how to interpret the observables, in order to maximise the benefits of the future data. A possible falsification of $\\Lambda$CDM would then be the first step towards a complete understanding of the dark energy. In practice, we can choose between comparing $\\Lambda$CDM to a large representative number of alternative scenarios, parameterise departures from it in a sufficiently flexible (but still economical) way \\cite{Caldwell:2007cw, Amendola:2007rr, Hu:2007pj, Song:2008vm, Bertschinger:2008zb, Song:2010rm, Pogosian:2010tj, Daniel:2010yt, Song:2010fg, Hojjati:2011ix, Zhao:2011te, Dossett:2011zp, Dossett:2011tn, Baker:2011jy, Zuntz:2011aq}, or identify key relations whose breakdown would hint towards a falsification of the standard paradigm \\cite{Huterer:2006mva, Song:2008vm, Song:2008xd}. All three strategies are useful in their own right: the third option is initially most powerful for falsifying or confirming $\\Lambda$CDM, while the second can be helpful as a guidance for the first, more traditional option. In any case, exquisite precision in the model predictions and understanding of the observables are needed to make decisive statements.  In this work we propose and start to explore a program for constructing fully nonlinear consistency relations between geometry and matter content, not assuming GR but allowing for generic modifications of the field equations. This can be seen as a (nonlinear) step beyond the popular modified growth parameterisations \\cite{Caldwell:2007cw, Amendola:2007rr, Hu:2007pj, Song:2008vm, Bertschinger:2008zb, Song:2010rm, Pogosian:2010tj, Daniel:2010yt, Song:2010fg, Hojjati:2011ix, Zhao:2011te, Dossett:2011zp, Dossett:2011tn, Baker:2011jy, Zuntz:2011aq}. Such relations can either be implemented as consistency checks to falsify $\\Lambda$CDM (or some other paradigm), or can be parameterised to explore a possible departure from $\\Lambda$CDM. To construct the consistency relations we make use of the fact that at late times matter (baryonic and dark) is well described by a pressureless perfect fluid (dust) on cosmological scales. The absence of isotropic and anisotropic pressures in the matter rest frame directly implies two consistency relations between the geometry and the physics that gives rise to the late time accelerated expansion. However, we point out some important subtleties that arise due to the choice of frame and that can become relevant on nonlinear scales.  We first formulate the consistency relations in the exact nonlinear theory employing the covariant approach \\cite{Ellis:1971pg, Ellis:1989jt, Bruni:1992dg, Ehlers:1993gf, Maartens:1998xg,Clarkson:2010uz}. Then, we discuss how the conditions come about after matter-radiation equality in first and second order cosmological perturbation theory, using the longitudinal (or conformal Newtonian) gauge. At first order, the resulting conditions are well known. They simply tell us that the two metric potentials are equal, $\\phi_1=\\psi_1$, and their evolution is governed by a simple second order differential equation that can be solved exactly: the Bardeen equation \\cite{Bardeen:1980kt}. At second order, the total measured pressure and anisotropic stress depend on the observer frame that is used to define these quantities, and this is where the subtleties come into play \\cite{Hwang:2005hd}. The pressure measured in the frame comoving with the matter is zero by definition (in analogy to the Lagrangian picture in the Newtonian approximation, see for instance \\cite{Hwang:2006iw,Villa:2011vt}). However, in a non-comoving frame, effective pressures and anisotropic stresses are induced by the relative velocity field, even if the source is a pressureless perfect fluid \\cite{Hwang:2005hd}. As we shall discuss in detail, there are good reasons to work in a specific non-comoving frame, called the \\emph{normal frame}. This frame is orthogonal to the surfaces of constant time and is given purely by the geometry, i.e.\\ the (perturbed) metric. An added advantage of this framework is that the density and pressure of matter (and of any other component) in the normal frame are then immediately related to the \\emph{geometrical pressures} that can be read off from the Einstein tensor. This means that projecting on the normal frame lets us cleanly separate the geometrical fluctuations from those of the stress-energy content. Consequently, the consistency relations in the normal frame take a particularly simple form. The change of frame from comoving to non-comoving induces effective matter pressures that we quantify. We find them to be small in comparison to typical non-standard sources of pressure and anisotropic stress in models beyond $\\Lambda$CDM.  The outline of the paper is as follows. First, in section \\ref{sec:covariant}, we review the effects of changing the observer frame from comoving to non-comoving, and discuss why the normal frame is particularly useful. In section \\ref{sec:metric_linear} we state the consistency relations in the exact nonlinear case, discuss the connection between the covariant and the metric perturbations approach and review the first order consistency relations. In section \\ref{sec:nonlinear} we derive the second order consistency relations and compute the effective matter pressure and anisotropic stress. Section \\ref{sec:discussion} is devoted to the discussion of a number of typical cases where the effective matter stresses could potentially lead to wrong physical interpretations  of the data. We consider the cases of a general clustering DE fluid and the DGP (Dvali-Gabadadze-Porrati \\cite{Dvali:2000hr}) model as a prototypical example of MG, because for both their respective intrinsic pressure and/or anisotropic stress perturbations could, in principle, be confused with the contributions from matter. Further, we discuss the impact on the determination of the galaxy bias from weak lensing and galaxy clustering. We summarise our work and conclude in section \\ref{sec:conclusions}. We include three appendices. In appendix \\ref{app:conventions} we state our conventions on notation and units. Appendix \\ref{app:covariant} briefly reviews the covariant approach and appendix \\ref{app:metric} gives more details on the metric perturbation calculations. ",
        "conclusions": "\\label{sec:conclusions}  With current cosmological observations tightening constraints on the $\\Lambda$CDM scenario and not (yet) hinting at strong deviations from it, or from GR, it becomes increasingly important to devise ways of falsifying or confirming the standard paradigm. To be able to do so, a very precise understanding of the observations and the corresponding predictions is needed not only on linear but also on nonlinear scales. In recent years, much progress has been made in nonlinear Newtonian perturbation theory to understand and complement numerical analyses. Simulation efforts in non-standard scenarios are also underway. On the other hand, relativistic perturbation theory predictions and consistency relations are mostly restricted to linear theory. But to be able to confidently falsify or confirm the standard paradigm it is imperative to understand the implications of dynamics in curved space on weakly nonlinear scales where Newtonian theory may not be sufficient, even more so if modifications to GR are allowed.  A framework based on casting beyond-$\\Lambda$CDM structure growth in terms of two functions, that can either be parameterised in generic ways or predicted from a given theory, has recently received considerable attention. This tool is often referred to as Parameterised-Post-Friedmann approach or modified growth parameterisations. In this paper, we propose an alternative framework deriving fully nonlinear consistency relations between the geometry, the matter content and any possible dark component. Since our framework requires neither the assumption of GR nor minimal interactions, but only relies on the symmetries of the Einstein tensor, it can be applied to different scenarios and models. The framework derives from the covariant approach, employing a very practical choice of 4-vector field (the normal frame) to decompose the Einstein tensor, the energy-momentum tensor and their difference, $X^\\mn\\equiv M_P^2 G^\\mn - T^\\mn$, that describes the physics beyond GR and/or beyond the standard constituents of the cosmological model. The normal frame, $n^\\mu$, is defined to be orthogonal to the surfaces of constant time. Because its definition is purely geometric it is only given by the metric itself and so are the irreducible components of $M_P^2G^\\mn$: the \\emph{geometrical} dynamic quantities. This is the reason why the normal frame is such a convenient choice, as it naturally allows us to write consistency relations that separate into a geometrical part, a matter part and a dark energy part: \\be \\nn p_X = p_G - \\frac{1}{3}v_m^2\\rho_m \\,, \\qquad \\pi_X^\\mn = \\pi_G^\\mn - \\rho_m v_m^{\\la\\mu}v_m^{\\nu\\ra} \\ee where the $G$-variables are only given by the metric, see section \\ref{sec:fullconsistency}. These fully nonlinear consistency relations can either be taken as conditions for $\\Lambda$CDM+GR to hold by requiring $p_X=-M_P^2 \\Lambda$, $\\pi_X^\\mn=0$, or the $X$-quantities can be parameterised and reconstructed from data. The advantage of this very general framework is that it is a top-down approach that a priori takes all non-linearities and all scalar, vector and tensor degrees of freedom into account, and can be simplified in perturbation theory at will.  As a consequence of selecting the normal frame, the dynamical quantities on the matter side (the energy density, the pressure, the energy flux and the anisotropic stress) do not take the values that one would expect, or find when measured in the matter rest frame. The transformations from the rest frame to the normal frame can be written fully nonlinearly. These show that a non-vanishing effective pressure and anisotropic stress are measured in the normal frame, even if the matter is assumed to be a perfect pressureless fluid in its rest frame, see section \\ref{sec:framechange}. This is the origin of the terms $\\sim v_m^2$ in the consistency relations. However, these effects are only relevant at second order in perturbation theory and they can be accounted for, knowing that they are present. Furthermore, we point out that cosmological observations do not actually measure the density field in its comoving frame, as we do observe the effect of the velocity field of the sources through the projection onto the photon wave vector. This means that  there are non-comoving effects that need to be taken into account. More importantly, apart from gravitational nonlinearities, we identify the pressure and anisotropic stress measured in the normal frame as the physical quantities that enter as sources in the Bardeen equation for the second order Weyl potential and the constraint for the second order gravitational slip. This is a consequence of the fact that the normal frame is purely defined in terms of the metric. Finally, let us recall that the normal frame helps to recover the Eulerian picture of Newtonian perturbation theory from GR. Although GR is a theory of curved spacetime, the normal frame is related to the Eulerian picture in Newtonian nonlinear dynamics.  We compute the consistency relations for $\\Lambda$CDM in first and second order scalar perturbation theory in the generalised longitudinal gauge (that is the gauge where $n^\\mu$ has zero spatial components and does not see shear or vorticity). At first order we recover the usual conditions: the gravitational slip vanishes, $\\Pi_1\\equiv\\phi_1-\\psi_1=0$, and the Weyl potential, $\\vf_1\\equiv\\frac{1}{2}(\\phi_1+\\psi_1)$, follows the standard Bardeen equation, see section \\ref{sec:linear}. We solve the Bardeen equation analytically to provide the basis for a careful study of the matter pressure and anisotropic stress. We compute $\\frac{1}{3}v_m^2\\rho_m$ and $\\rho_m v_m^{\\la\\mu}v_m^{\\nu\\ra}$ at second order and find these terms to be small on large scales. They become relevant only on scales $k\\gtrsim 0.1\\, h/\\Mpc$. The second order anisotropic stress consistency relation reveals a non-vanishing gravitational slip due to the effective matter anisotropic stress and non trivial combinations of products of spatial derivatives of $\\vf_1$, see equation \\eqref{eq:k4pi} and also refs. \\cite{Acquaviva:2002ud,Vernizzi:2004nc}. Here we clarify the origin of the different contributions to $\\Pi_2$. The pressure consistency, on the other hand, yields the second order analogue of the Bardeen equation for $\\vf_2$, see equation \\eqref{eq:bardeen2}. It acquires different source terms which we analyse, compute and compare and we find subtle cancellations between the contributions from the effective matter stresses and $\\partial^2\\vf_1^2$--type terms. The fact that the source term of the second order Bardeen equation depends non trivially and sensitively on the pressure and anisotropic stress perturbations of a possible dark energy makes the weakly non-linear evolution a promising tool to test the $\\Lambda$CDM paradigm. Combining different cosmological probes enables us to reconstruct the consistency relations, as all observations probe the metric potentials, either directly or indirectly via the Poisson equation, see also e.g.\\ \\cite{Song:2008xd}. Let us emphasise that the effective matter pressure and anisotropic stress are seen by the second order metric perturbations in any case, irrespectively of the yet to be solved issue of which precise frame is the one that corresponds to us as observers.  Finally, we discuss a set of applications: we compare the effective matter pressure with the pressure perturbations in clustering dark energy, concluding that the matter pressure is negligible on large scales and becomes relevant on weakly nonlinear scales. Thus, the matter pressure only needs to be taken into account when deriving constraints on the dark energy sound speed from small scale observations. Furthermore, we compare the matter anisotropic stress to typical anisotropic stresses arising in MG scenarios, specifically the DGP model. Anisotropic stress is normally considered a smoking gun for the detection of MG. We find that, typically, the matter anisotropic stress is by far subdominant and not an issue. Finally, we assess the impact of the anisotropic stress on the determination of galaxy bias from combining WL and galaxy clustering data. We find that the gravitational slip affects the Lagrangian galaxy bias and the velocity bias in opposite ways, but only marginally. Using multiple tracers at several redshifts and different combinations of observables, we expect that one could be able to filter out or strongly constrain the gravitational slip."
    },
    "1112/1112.1415_arXiv.txt": {
        "abstract": "We examine to what extent the inferred surface temperature of magnetars in quiescence can constrain the presence of a superfluid in the neutron star core and the role of magnetic field decay in the core. By performing detailed simulations of neutron star cooling, we show that extremely strong heating from field decay in the {\\it core} cannot produce the high observed surface temperatures nor delay the onset of neutron superfluidity in the core. We verify the results of Kaminker et al., namely that the high magnetar surface temperatures require heating in the neutron star {\\it crust}, and crust heating is decoupled from cooling/heating in the core. Therefore, because crust heating masks core heating, it is not possible to conclude that magnetar cores are in a non-superfluid state purely from high surface temperatures. From our interior temperature evolutions and after accounting for proton superconductivity in the core, we find that neutron superfluidity in the core occurs less than a few hundred years after neutron star formation. This onset time is unaffected by heating due to core field decay at fields $\\lesssim 10^{16}\\mbox{ G}$. Thus all known neutron stars, including magnetars, without a core containing exotic particles, should have a core of superfluid neutrons and superconducting protons. ",
        "introduction": "\\label{sec:intro}  Neutron stars (NSs) begin their lives very hot (with temperatures $T>10^{11}\\mbox{ K}$) but cool rapidly through the emission of neutrinos. Neutrino emission processes depend on uncertain physics at the supra-nuclear densities ($\\rho>\\densnuc\\approx 2.8\\times 10^{14}\\mbox{ g cm$^{-3}$}$) of the NS core (see \\citealt{tsuruta98,yakovlevpethick04,pageetal06}, for review). Current theories indicate that the core may contain a neutron superfluid and proton superconductor or even exotic particles, such as hyperons and deconfined quarks (see, e.g., \\citealt{lattimerprakash04,haenseletal07}, for review). The recent observation of rapid cooling \\citep{heinkeho10,shterninetal11} of the NS in the Cassiopeia~A supernova remnant provides the first constraints on the critical temperatures for the onset of superfluidity of core neutrons $\\Tcnt$ (in the triplet state) and protons $\\Tcp$ (in the singlet state), i.e., $\\Tcnt\\approx(5-9)\\times 10^8$~K and $\\Tcp\\sim (2-3)\\times 10^9$~K \\citep{pageetal11,shterninetal11}.  Anomalous X-ray pulsars and soft gamma-ray repeaters form the magnetar class of NSs, i.e., systems which possess superstrong magnetic fields ($B\\gtrsim 10^{14}\\mbox{ G}$)\\footnote{However, there exists an apparently low surface magnetic field magnetar \\citep{reaetal10}.}, and their strong fields likely power the activity seen in these objects (see \\citealt{woodsthompson06,mereghetti08}, for review). One notable property of magnetars is that their observed surface temperatures $\\Ts$ in quiescence are significantly higher than those of other NSs of a similar age.  In fact, they are too high for NSs that cool passively, i.e., without an additional source of internal heat (accretion heating can be excluded by, e.g., non-detections of binary companion or disk emission). An interesting problem concerns the heat generated from magnetic field decay, which has been proposed to be the source for the high temperatures of magnetars \\citep{thompsonduncan96,heylkulkarni98,colpietal00,aguileraetal08b}. This heat can strongly influence the time/age at which the core becomes superfluid if heating/field decay occurs in the core \\citep{thompsonduncan96,arrasetal04,dallossoetal09}. The problem is important since the presence of superfluid components has a strong impact on magnetar interior dynamics, such as the mechanism for producing glitches (see, e.g., \\citealt{sauls89}), fluid oscillations (see, e.g., \\citealt{passamontiandersson11}), magnetic field and rotational evolution (see, e.g., \\citealt{glampedakisandersson11}), and magnetohydrodynamical equilibrium \\citep{glampedakisetal12,landeretal12}.  We address this problem by conducting detailed calculations of the thermal evolution of a NS with various prescriptions for an internal heat source that can be associated with magnetic field decay. We show that, regardless of the magnetic field strength and detailed mechanism for field decay in the core, the heat generated by the decay is insufficient to power the surface emission of magnetars. Furthermore, by accounting for the effects of proton superconductivity on field decay, we find that this core heating is not strong enough to balance neutrino cooling and cannot delay the onset of core neutron superfluidity for fields $\\lesssim 10^{16}\\mbox{ G}$ (onset at age $\\lesssim\\mbox{a few}\\times 100\\mbox{ yr}$). Thus the cores of all currently known magnetars should be in a superfluid state.  We briefly describe past works and note key findings and assumptions made in these works that we improve upon here. \\citet{arrasetal04} consider coupled magnetic field decay and thermal evolution of magnetars. The internal field and superfluid temperatures are assumed to be $>10^{15}\\mbox{ G}$ and $\\Tcp=5\\times 10^9\\mbox{ K}$ and $\\Tcnt=(5-9)\\times 10^8\\mbox{ K}$, respectively. \\citet{arrasetal04} find that magnetic field decay can delay the transition to core neutron superfluidity and maintain a relatively high surface temperature to ages $\\approx 10^3-10^5\\mbox{ yr}$, depending on $\\Tcnt$. However, their calculation only considers volume-averaged quantities and thus assumes that the NS interior is isothermal (see Section~\\ref{sec:nscooleq}), which cannot be the case if there is a localized heating source such as field decay in the crust or core due to the large but finite thermal conductivity of NS matter.  \\citet{kaminkeretal06} calculate the evolution of the temperature profile [i.e., $T(\\rho)$] by solving the energy balance and flux equations [see eqs.~(\\ref{eq:energybalancegr}) and (\\ref{eq:heatfluxgr})]. They demonstrate that, in order to explain the observed high surface temperatures, magnetars require a heat source and, most importantly, this heat source (e.g., from field decay; \\citealt{ponsetal09}) must be located in the outer crust; if the heat source is located too deep in the NS interior (e.g., in the core), then neutrino emission efficiently removes the heat locally, and the surface temperature cannot be increased sufficiently to match the observed values. Furthermore, the outer crust is thermally decoupled from the core so that heating of the crust does not affect (neutrino) cooling of the core. Because of this thermal decoupling, the results and conclusions of \\citet{kaminkeretal06} are not particularly sensitive to the state of matter in the core, e.g., superfluidity of the core nucleons. Though no quantitative results are shown, \\citet{kaminkeretal06} state that cooling calculations with the effects of inner crust superfluidity and neutrino emission by Cooper pair formation produce different temperature profiles, but these effects do not change the surface temperature.  In follow-up work, \\citet{kaminkeretal09} improve their calculations and examine the effects of light-element accreted envelopes and anisotropic heat conduction due to the magnetic field in the envelope and outer crust. While accreted envelopes can give higher surface temperatures for the same core temperatures, core heating is still unable to produce surface temperatures that are high enough to explain magnetar observations.  \\citet{dallossoetal09} criticize the work of \\citet{kaminkeretal06}, arguing that the phenomenological heating model considered is independent of magnetic field and decays on the wrong timescale [see eq.~(\\ref{eq:crustheat})] and that the correct heating should depend inversely on temperature. This last point implies that more heat is generated at lower temperatures and can lead to an equilibrium between neutrino cooling and heating from field decay (see also \\citealt{thompsonduncan96}). The equilibrium temperature (above $10^9\\mbox{ K}$) can be maintained for $\\approx 10^4\\mbox{yr}$.  The results of \\citet{dallossoetal09} suggest that the core of magnetars are not superfluid until after this time, since the critical temperature for the onset of neutron superfluidity is $(5-9)\\times 10^8\\mbox{ K}$ \\citep{pageetal11,shterninetal11}. However no cooling calculation is performed by \\citet{dallossoetal09}.  Here we perform NS cooling simulations to determine the role of core heating by magnetic field decay. We do not examine in detail field evolution and heating in the crust (see, e.g., \\citealt{ponsetal09}). Rather we use the phenomenological model of \\citet{kaminkeretal06,kaminkeretal09} to demonstrate the effect of {\\it crust heating} on the NS cooling behavior. This will be adequate for our purposes since the goal here is to assess the importance of {\\it core heating} in magnetars. As we will show, as long as the magnetic field decay time is longer than the cooling time, the core heating we use is the maximum one can use. While this maximal heating can delay the onset of neutron superfluidity [on the timescale estimated by \\citet{thompsonduncan96,dallossoetal09}], it cannot reproduce the high surface temperatures seen for magnetars. We arrive at this same conclusion using the prescription for heating and field decay given in \\citet{dallossoetal09}. Furthermore, the above works neglect the effects of proton superconductivity; when this is taken into account, there is no delay in superfluidity onset for any reasonable core magnetic field strength. In Section~\\ref{sec:model}, we describe the thermal evolution equations and input physics, including superfluid properties and prescriptions for internal heat sources. Section~\\ref{sec:results} presents the results of our calculations. We summarize our results and discuss their implications in Section~\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  We performed detailed calculations of neutron star cooling, including the effects of superfluidity and additional heating (due to magnetic field decay) in the crust and core. We find that magnetic field decay in the neutron star core cannot be the sole source powering the high observed surface temperature of magnetars in quiescence; the high temperatures require outer crust heating \\citep{kaminkeretal06,kaminkeretal09}. Because of crust heating and effective thermal decoupling between the outer crust and core, the state of matter in the core cannot be deduced from these surface temperature measurements. By computing the evolution of the temperature profile $T(\\rho)$, we determine the time when core neutrons first become superfluid, i.e., when $T<\\Tcnt(\\rho)$. We find that heating by field decay in the core (with fields $\\lesssim 10^{16}\\mbox{ G}$) cannot balance neutrino cooling and thus cannot maintain relatively high core temperatures (c.f., \\citealt{thompsonduncan96,dallossoetal09}). As a result, onset of superfluidity for neutrons in the core cannot be delayed, and neutron stars possess superfluid and superconducting cores after a few hundred years; this does not strongly depend on the nuclear EOS.  Since core heating is not significant, the temperature profiles $T(\\rho)$ and surface temperature evolution $\\Ts^\\infty(t)$ for magnetars are just those with crust heating (see Figs.~\\ref{fig:tempprofcrust} and \\ref{fig:timets}, respectively), and crust heating does not affect the onset of core neutron superfluidity because of thermal decoupling between the outer crust and core.  Magnetar activity may be driven by field decay in the core \\citep{thompsonduncan95,thompsonduncan96} or by processes in the crust \\citep{thompsonduncan93,thompsonduncan95,glampedakisetal11,priceetal12}. Although the former may still be true, it must be accompanied by heating in the (outer) crust. On the other hand, field evolution in the crust easily couples to surface emission. Thus in order to understand the high surface temperature of magnetars, detailed studies should focus on magnetic field evolution and heating in the crust (see, e.g., \\citealt{ponsetal09,cooperkaplan10,priceetal12}). We note that several magnetars may have very similar X-ray luminosities \\citep{durantvankerkwijk06}; this could suggest that their crustal field strengths are similar and the field decay timescale is longer than the age of the oldest of these sources.  Finally, the (normal versus superfluid) state of the core has important consequences for magnetic and rotational evolution of magnetars \\citep{glampedakisandersson11}, as well as their glitching behavior and possible stellar oscillations; studies of these effects may be effective probes of the neutron star core. We showed that the core can be treated as being in a superfluid and superconducting state after the neutron star is a few hundred years old."
    },
    "1112/1112.1765_arXiv.txt": {
        "abstract": "The current direct observations of brown dwarfs and exoplanets have been obtained using instruments not specifically designed for overcoming the large contrast ratio between the host star and any wide-separation faint companions.  However, we are about to witness the birth of several new dedicated observing platforms specifically geared towards high contrast imaging of these objects.  The Gemini Planet Imager, VLT-SPHERE, Subaru HiCIAO, and Project 1640 at the Palomar 5m telescope will return images of numerous exoplanets and brown dwarfs over hundreds of observing nights in the next five years.  Along with diffraction-limited coronagraphs and high-order adaptive optics, these instruments also will return spectral and polarimetric information on any discovered targets, giving clues to their atmospheric compositions and characteristics. Such spectral characterization will be key to forming a detailed theory of comparative exoplanetary science which will be widely applicable to both exoplanets and brown dwarfs. Further, the prevalence of aperture masking interferometry in the field of high contrast imaging is also allowing observers to sense massive, young planets at solar system scales ($\\sim$3-30 AU)--- separations out of reach to conventional direct imaging techniques. Such observations can provide snapshots at the earliest phases of planet formation---information essential for constraining formation mechanisms as well as evolutionary models of planetary mass companions.  As a demonstration of the power of this technique, I briefly review recent aperture masking observations of the HR 8799 system.  Moreover, all of the aforementioned techniques are already extremely adept at detecting low-mass stellar companions to their target stars, and I present some recent highlights. ",
        "introduction": "In recent years, astronomers have identified more than 400 planets outside our solar system, launching the new and thriving field of exoplanetary science (Marcy \\etal~2005).  The vast majority of these objects have been discovered indirectly by observing the variations induced in their host star's light.  The radial velocity surveys can  provide orbital eccentricity, semi-major axes, and lower limits on the masses of companion planets while observations of transiting planets can provide fundamental data on planet radii and limited spectroscopy of the planets themselves.  However, studying those objects out of reach to the radial velocity and doppler methods will reveal completely new aspects of exoplanetary science in great detail.    {\\it Direct imaging} of exoplanets provides a complementary set of parameters such as photometry (and hence luminosity), as well as detailed spectroscopic information.   \\begin{figure} \\centering \\resizebox{13.5cm}{!}{\\includegraphics[angle=-90]{f1.ps}} \\caption[]{{\\it Left:} A schematic demonstration taken from \\cite{oh09} demonstrating the challenges associated with high contrast imaging. Hypothetical orbits of a Jupiter and an Earth have been overlaid on the Point Spread Function of a nearby star. Planetary mass companions typically have angular separations of a fraction of an arcsecond, and contrast ratios spanning several orders of magnitude ($\\sim$10$^4$-10$^7$).  {\\it Right:} The Mass vs. Orbital separation for planetary mass companions detected through transits (filled circles), radial velocity measurements (squares), solar system planets (triangles), as well as those objects detected through direct imaging (open circles). } \\label{highcontrast} \\end{figure}  Moreover, the direct imaging of planetary mass companions will allow researchers to more fully probe the mass-separation parameter space (Figure~\\ref{highcontrast}, right panel) occupied by these objects (Oppenheimer \\& Hinkley 2009). At young ages (\\S 3), the orbital placement of planetary mass companions serves as a birth snapshot, lending support to models (Pollack \\etal~1996) that may be more efficient at building a massive core at only a few AU, or models that allow for the formation of massive objects at tens of AU though fragmentation of the gaseous disk (Boss 1997).  Finally, imaging of multi-planet systems at any age will serve as dynamical laboratories for studying the planetary architectures.  \\begin{figure} \\centering \\resizebox{14.5cm}{!}{\\includegraphics[angle=90]{f2.ps}} \\caption[]{{\\it Top Panel:} Example Spectra for M, L, and T-dwarfs, as well as Jupiter with prominent absorption bands marked, taken from Marley \\etal~(2009).  {\\it Lower Panels:}  The points with error bars are simulated measurements gathered with a low resolution ($\\lambda/\\Delta\\lambda$$\\sim$45) spectrograph, shown along with higher resolution spectra from NASA IRTF (black curves) for the labelled spectral types. Such low spectral resolution measurements will be obtained by GPI, SPHERE and Project 1640, and can easily sample broad absorption features present in late type stars and exoplanets.  Lower panels are taken from Rice \\etal~(2011) in prep. Images used courtesy of Michael Cushing, Mark Marley, and Emily Rice. } \\label{spectra} \\end{figure}   Perhaps just as essential, spectroscopy provides clues to the atmospheric chemistry, internal physics, and perhaps may even shed light on non-equilibrium chemistry associated with these objects.  More robust classification schemes for planets in general will arise from observing as many planets as possible at different ages, in different environments, and with a broad range of parent stars. Figure~\\ref{spectra} provides an illustration of the diversity of the spectra of late-M, L and T dwarfs as well as a spectrum of Jupiter showing broad absorption bands due to e.g. H$_2$O, CH$_4$, and NH$_3$.  Spectral characterization of such objects can be accomplished even with the somewhat low spectral resolution ($\\lambda/\\Delta\\lambda\\sim$30-50) of the instruments described in this review. These pieces of information not only reveal the detailed properties of the objects themselves, they serve as key benchmarks for competing evolutionary models describing the thermal and atmospheric properties of these objects.     It is often overlooked, but should be emphasized, that the recent spectacular images (Marois \\etal~2008, Kalas \\etal~2008, and Lagrange \\etal~2010)  of Jovian planets were obtained with Adaptive Optics, hereafter ``AO'', systems and infrared cameras not specifically designed for the task of overcoming the high contrast ratio between the planets and their host stars. These were obtained using existing instrumentation, but with observing and data reduction strategies customized for high contrast imaging. Further, the recent L-band images of HR 8799 and $\\beta$ Pic b (Marois \\etal~2010, Lagrange \\etal~2010) were obtained without the use of a  coronagraph! These studies have demonstrated that direct imaging of planetary mass companions as well as disks (e.g. Oppenheimer \\etal~2008, Hinkley \\etal~2009) is now a mature technique and may become routine using ground-based observatories. More so, the handful of coming instruments dedicated to detailed spectroscopic characterization of planetary mass companions may make these kinds of discoveries routine, initiating an era of comparative exoplanetary science.  \\subsection{The Challenge of High Contrast Imaging}  The major obstacle to the direct detection of planetary companions to nearby stars is the overwhelming brightness of the host star.  If our solar system were viewed from 20 pc, Jupiter would appear $10^8-10^{10}$  times fainter than our Sun in the near-IR (Barraffe \\etal~2003) at a separation of 0.25$^{\\prime\\prime}$, completely lost in its glare (See Figure~\\ref{highcontrast}, left panel). The key requirement is the suppression of the star's overwhelming brightness through precise starlight control (Oppenheimer \\& Hinkley 2009).  A promising method for direct imaging of stellar companions involves two techniques working in conjunction.  The first, high-order AO, provides control and manipulation of the image by correcting the aberrations in the incoming stellar wave front caused by the Earth's atmosphere.  AO has the effect of creating a nearly diffraction-limited point spread function, with the majority of the stellar flux concentrated in this core.  Second, a Lyot coronagraph (Sivaramakrishnan \\etal~2001) suppresses this corrected light.  Together, these two techniques can obtain contrast levels of $10^4$-$10^5$ at 1$^{\\prime\\prime}$ (Leconte \\etal~2010).  Improvements in coronagraphy, specifically the apodization of the telescope pupil (Soummer 2005), as well as post-processing to suppress speckle noise (Hinkley \\etal~2007, Crepp et al 2011), can significantly improve the achieved contrast, especially at high Strehl ratios.  Below we briefly describe some of the instrumentation being built at the time of this writing, or currently in place on large telescopes. ",
        "conclusions": ""
    },
    "1112/1112.6373_arXiv.txt": {
        "abstract": "We have performed extensive two-dimensional magnetohydrodynamic simulations to study the amplification of magnetic fields when a supernova blast wave propagates into a turbulent interstellar plasma. The blast wave is driven by injecting high pressure in the simulation domain. The interstellar magnetic field can be amplified by two different processes, occurring in different regions. One is facilitated by the fluid vorticity generated by the ``rippled\" shock front interacting with the background turbulence. The resulting turbulent flow keeps amplifying the magnetic field, consistent with earlier work \\citep{Giacalone2007}. The other process is facilitated by the growth of the Rayleigh-Taylor instability at the contact discontinuity between the ejecta and the shocked medium. This can efficiently amplify the magnetic field and tends to produce the highest magnetic field. We investigate the dependence of the amplification on numerical parameters such as grid-cell size and on various physical parameters. We show the magnetic field has a characteristic radial profile that the downstream magnetic field gets progressively stronger away from the shock. This is because the downstream magnetic field needs a finite time to reach the efficient amplification, and will get further amplified in the Rayleigh-Taylor region. In our simulation we do not observe a systematic strong magnetic field within a small distance to the shock. This indicates that if the magnetic-field amplification in supernova remnants indeed occurs near the shock front, other processes such as three-dimensional instabilities, plasma kinetics and/or cosmic ray effect may need to be considered to explain the strong magnetic field in supernova remnants. ",
        "introduction": "Powerful shocks associated with supernova remnants (hereafter SNRs) sweeping through the interstellar medium (ISM) are remarkable high-energy phenomena in astrophysics. It is widely believed that these high-Mach number shocks are the sources of galactic cosmic rays with energies up to at least $10^{15}$ eV. SNRs are also sources of strong radio and/or X-ray emissions. In these high-energy processes, the magnetic field is of great importance. Moreover, it provides information on energetic charged particles, which are presumably accelerated by the supernova shocks.  The ISM is known to be turbulent. Measurements of the ISM radio-wave scintillation have established the existence of large-scale density turbulence which has a Kolmogorov-like power spectrum spanning more than ten decades of spatial scale with an outer scale of several parsecs \\citep[e.g.,][]{Lee1976ApJ,Armstrong1981,Rickett1990,Armstrong1995,Minter1996}. This has been called ``the big power law in the sky.\" \\citep[see,][]{Spangler2007} The galactic magnetic field is observed to be a few micro-Gauss and has uniform and fluctuating components that are roughly in equipartition \\citep[e.g.,][]{Beck1996,Minter1996,Han2004}. The turbulent magnetic field can interact with the shock waves, distorting their surfaces leading to shock ripples \\citep{Neugebauer2005} and the enhanced downstream magnetic fluctuations \\citep{Lu2009}. It is also important for efficient particle acceleration \\citep{Giacalone2005,Jokipii2007,Guo2010a}. Turbulence in the upstream medium has also been considered \\citep[][]{Balsara2001} to explain the irregular and patchy emission morphology observed in SNRs \\citep[e.g.,][]{Anderson1996}.  Recently, it has been inferred from observations that the magnetic field in young SNRs is strongly enhanced to a magnitude much greater than the compression given by the shock jump condition. For example, by assuming that the so-called X-ray ``thin rims\" seen in several young SNRs \\citep[e.g.,][]{Bamba2005} are caused by shock accelerated electrons rapidly losing energy in strong magnetic field through synchrotron radiation, the associated magnetic field may be more than $50-100\\mu G$ in order to explain the thickness of the ``thin rims\" \\citep{Berezhko2003,Volk2005,Ballet2006,Parizot2006}. For SNR shocks with higher shock speeds propagating in more inhomogenous media such as Cas A and Tycho, the downstream magnetic fields  are inferred to be more enhanced. ``Thin rims\" are also seen in radio emissions \\citep{Reynoso1997}, which cannot be explained by the electrons losing energy in strong magnetic field. This indicates some other mechanism, e.g., decay of magnetic fluctuations may need to be considered \\citep{Pohl2005}. Further downstream of the shock, the magnetic field may possibly be even higher than the region right behind the shock \\citep{Vink2003}. The morphology of X-ray emission in SNRs shows filamentary structure and rapid time variation, which indicates that the magnetic field could be as high as $1 mG$ \\citep{Uchiyama2007}, over two orders of magnitude higher than the background fluid. It should be pointed out this rapid variation of synchrotron emission could be well reproduced in the case of strong magnetic fluctuations \\citep{Bykov2008}. Since the effect of turbulence is not fully understood at this point, the actually magnetic-field amplification factors in young SNRs remain uncertain.  \\citet{Bell2001} and \\citet{Bell2004} proposed that cosmic rays accelerated by the SNR forward shock waves provide a current that leads to an instability that can amplify the magnetic field close to the shock front. Numerical simulations show the evidence of this instability, however, it is found to easily saturate and the amplification factor may be limited \\citep[e.g.,][]{Riquelme2009}. Recently, \\citet{Giacalone2007} proposed an alternative mechanism, in which the interaction between the warped shock front and large-scale density fluctuations produces fluid vorticity downstream of strong shocks. That fluid vorticity can stretch, distort and amplify the magnetic field. The magnetic-field amplification in this mechanism relies on the dynamics of a magnetized fluid rather than the cosmic-ray kinetic physics. It is interesting to note, however, \\citet{Balsara2001} performed three-dimensional MHD blast wave simulation with moderate resolution which does not show magnetic-field enhancement larger than $50 \\mu G$, whereas two-dimensional Cartesian geometry simulations with high numerical resolution give strong amplification \\citep{Giacalone2007,Inoue2009} with maximum values larger than $100\\mu G$. This discrepancy warrants further investigation.   In addition, a critical constraint from several observed ``thin rims\" is that the required magnetic-field amplification should occur within a narrow distance $\\lesssim 0.01$ pc of the supernova shocks \\citep[for a review, see][]{Reynolds2011}. This is supported by some coincidences of X-ray ``thin rims\" and shock locations inferred from H$\\alpha$ observations \\citep{Winkler2003} and radio polarimetry \\citep{Gotthelf2001}. This places an important constraint on various field amplification models, although it should be noted that there is no one-to-one correspondence between X-ray ``thin rims\" and inferred shock locations. To our knowledge, this constraint has not been taken into account previously in comparing models and simulations with observations.  In this study, we perform a series of two-dimensional ideal MHD simulations with high spatial resolution to study strong supernova blast shock waves propagating into the ISM containing pre-specified large-scale density and magnetic fluctuations. We investigate how the amplification depends on a variety of parameters, including the explosion energy, the level of background turbulence, and the numerical resolution, etc. The paper is organized as follows. In Section 2, we describe our numerical model and simulation setup. We present the simulation results in Section 3, and in Section 4 we summarize and discuss our results. ",
        "conclusions": "The inferred strong magnetic field in young SNRs is a significant result and can be important in the high-energy processes including particle acceleration and thermal/nonthermal emissions. The origin of this process, however, is still under debate. In this work we study the interaction between a supernova blast wave with a turbulent upstream medium which contains density and magnetic-field fluctuations. The vorticity produced at the rippled shock front can stretch and distort the magnetic field lines, and this leads to a strong magnetic-field amplification downstream \\citep{Giacalone2007,Inoue2009}. Using two-dimensional MHD simulations of a blast wave, we confirm the key features of this process. Based on our simulations, we conclude that the increase of magnetic field is dependent on shock speed and background density turbulence amplitudes. Furthermore, the numerical resolution used in the simulations can play an important role as well. Previous work \\citep{Balsara2001} using three-dimensional MHD simulation with moderate resolution shows no magnetic-field amplification beyond $50\\mu G$. Here we show the magnetic evolution downstream is sensitive to the resolutions used in the simulation. For high resolutions, the simulations allow rapid growth at small scales, this leads to efficient field amplification.  Furthermore, we find that there are two different processes and spatial regions where magnetic field is amplified. One is associated with the shock amplification immediately downstream and the other is associated with the RTI at the interface between the ejecta and the shocked medium.  However, in our simulations, we did not observe a systematic strong magnetic field within a thin region immediate downstream of the supernova shock. For example, using the results of the highest resolution case, within $0.15$ pc downstream of supernova shock, we observe only about $0.8\\%$ region which has magnetic field larger than $30 \\mu G$. This lack of strong magnetic field can be understood as the downstream dynamo process requires an efficient stretching to produce strong magnetic field. The time scale for the growth of magnetic field depends on the eddy turnover time. Only after a certain time can the field get sufficient amplification. If the thin rims ($0.01 - 0.1$ pc) observed in young SNRs are indeed caused by the electrons losing energy in strong fields ($\\sim$ several hundred $\\mu G$), some other processes such as three-dimensional instabilities, plasma kinetics or the effect of cosmic rays might be needed to explain the magnetic-field amplification in young SNRs.  We note that in observation there is no one-to-one correspondence between X-ray ``thin rims\" and inferred shock locations. In fact many ``thin rims\" and filamentary structures seen in X-ray observation can hardly be related to shock front due to observation limitation. Further understanding about the relationship between the small scale X-ray structure and shock locations is needed to further constraint and distinguish the different mechanisms for amplification of magnetic field.  We also note that the two-dimensional simulation in this study could be significantly different from three-dimensional simulation. It is known that in two-dimensional simulation the dynamics of MHD flow is very different from that for three-dimensional simulation. For example the inverse cascade of enstrophy can causes strong intermittency due to two-dimensional effect \\citep{Biskamp2003}. There are several recent studies show that magnetic field amplification behind high-mach number shock \\citep{Inoue2011a,Inoue2011b} and in Rayleigh-Taylor region \\citep{Stone2007} can still operate in three-dimensions, which confirm the results found in two-dimensional simulations. Also, in three-dimensional simulation the MHD flow could develop other types of instabilities with larger growth rates. Further three-dimensional MHD simulation with high resolution will be useful in confirming the conclusions of this paper."
    },
    "1112/1112.1881_arXiv.txt": {
        "abstract": "{In most of particle dark matter (DM) models, the DM candidate injects sizable fluxes of high-energy electrons and positrons through its annihilations or decays. Emitted in regions with magnetic field, they in turn give raise to a synchrotron radiation, which typically covers radio and infrared bands. We discuss the possibility of detecting signatures of Galactic and extra-galactic DM in the total intensity and small-scale anisotropies of the radio background.}  \\FullConference{XXIst International Europhysics Conference on High Energy Physics\\\\ 21-27 July 2011\\\\ Grenoble, Rhones Alpes France}  \\begin{document} ",
        "introduction": "From the title of my talk, one may wonder how it can be possible to infer something completely unknown, as non-gravitational signals of dark matter (DM), by means of something quite well-known, as radio-astronomy. Discovery of new physics normally goes through designing new experiments. Here the idea is very simple, namely, that the technique of radio observations is potentially very promising in the quest for particle DM, but requires improved flux and angular sensitivities with respect to current capabilities. The development of ASKAP, EVLA, MeerKAT, and, in particular, SKA makes the near future promisingly bright in this respect. Those new radio telescopes will certainly discover signals of previously unknown physical mechanisms, and we discuss the possibility that an emission induced by particle DM will be among them.  Weakly interacting massive particles (WIMPs) are the most investigated class of DM candidates in the literature (for a review, see, e.g.,~\\cite{Bertone:2010zza}). One of the routes to test the hypothesis of WIMP DM stems from the the bases of the framework themselves. Indeed, given the weak interaction, there is a (\"weak\" but finite) probability that WIMPs in DM halos annihilate in pairs or decay into detectable species. In particular, and with the exception of WIMP models annihilating/decaying into neutrinos only, a sizable branching ratio of annihilation/decay into electrons and positrons is a general feature of WIMP models (see, e.g., Fig.~4 in \\cite{Regis:GC}). Interactions of high-energy $e^+/e^-$ with the interstellar magnetic field in astrophysical structures give rise to magnetic bremsstrahlung called synchrotron radiation. In the monochromatic approximation of synchrotron power, the energy of an electron emitting at frequency $\\nu$ is given by $E\\simeq 15\\sqrt{\\nu_{GHz}/B_{\\mu G}}$ GeV (where $\\nu_{GHz}$ is the frequency in GHz and $B_{\\mu G}$ is the magnetic field in $\\mu G$). It follows that, assuming a magnetic field of few $\\mu G$ (as typical for galaxies), emissions at radio frequencies are mostly generated by electrons with energy around 1-10 GeV. Therefore, assuming a DM mass $\\gtrsim 10$ GeV, a sizable DM-induced radio synchrotron emission is a general prediction of WIMP models (with, of course, spectrum and absolute flux depending on the specific model). ",
        "conclusions": ""
    },
    "1112/1112.6328_arXiv.txt": {
        "abstract": "We discuss the possible impact of strange quark matter on the evolution of core-collapse supernovae with emphasis on low critical densities for the quark-hadron phase transition. For such cases the hot proto-neutron star can collapse to a more compact hybrid star configuration hundreds of milliseconds after core-bounce. The collapse triggers the formation of a second shock wave. The latter leads to a successful supernova explosion and leaves an imprint on the neutrino signal. These dynamical features are discussed with respect to their compatibility with recent neutron star mass measurements which indicate a stiff high density nuclear matter equation of state. ",
        "introduction": "A core-collapse supernova (SN) explosion marks the disruption of a massive star by an energetic shock wave followed by the formation of a neutron star or a black hole. In the first hundreds of milliseconds (ms) of a supernova, temperatures in the range of tens of MeV, densities beyond nuclear matter saturation density $n_0 \\sim 0.145\\:$fm$^{-3}$ and proton fractions $Y_p \\leq 0.3$ are reached. With such properties, supernovae (SNe) are astrophysical laboratories for dense nuclear matter which, in the phase diagram of strongly interacting matter, possess overlap regions with heavy-ion experiments, such as in the future FAIR facility at GSI, Darmstadt (Germany) and the NICA facility at the JINR in Dubna (Russia). The modeling of core-collapse SNe represents a great computational challenge. It requires as input amongst others a nuclear matter equation of state (EoS), which  provides information about the thermodynamic properties and compositions of matter for a large range of baryon number densities $n_b$, temperatures $T$, and isospin states, characterized by $Y_p$. Being discussed to populate neutron star interiors, hyperons and quark matter can also be included in SN equations of state (EoSs). At present, both components are tested on their impacts in SN simulations. Hereby, the applied quark and hyperon EoSs should be compatible with observed neutron star properties, e.g. pulsar masses. As we will argue in the next section, the latter indicate a stiff nuclear matter EoS at high density. \\newline In the following we will give a short summary on recent pulsar mass measurements and their implications for quark matter in neutron star interiors. We will proceed with an overview of quark matter studies in core-collapse SNe. In the last section, we will focus on the impact of low density quark-hadron phase transitions on the gravitational collapse of light and intermediate mass progenitor stars. ",
        "conclusions": "The onset of strangeness in core-collapse supernovae (SNe) can be studied by implementing hyperons and strange quark matte in SN equations of state (EoSs). Hereby, the chosen parameter sets must fulfill restrictions from the recent finding of a two solar mass pulsar PSR J1614-2230, which is only compatible with a stiff high density nuclear matter EoS. Within currently applied quark EoS models, high critical densities were shown to shorten the time to black hole formation in the gravitational collapse of massive progenitor stars. For the onset of quark matter with a soft EoS at critical densities $n_{crit}$ around nuclear matter saturation density $n_0\\sim 0.145$fm$^{-3}$, the proto neutron star collapses to a more compact hybrid star configuration within hundreds of ms after core-bounce. This launches a shock wave which leads to the SN explosion and releases a neutrino burst, dominated by $\\bar{\\nu}_e$. The properties of the neutrino burst are dependent on $n_{crit}$ as well as the model of the progenitor star. A first study with a stiff quark EoS shows no impacts on the SN dynamics as the critical densities are too high to be reached in the early post-bounce phase. However, studies are on the way, in which we apply a stiff quark EoS parameter set where SN matter enters the quark-hadron mixed phase at $n_{crit} \\sim n_0$.\\\\ \\newline \\textit{Acknowledgements}\\\\ The project was funded by the Swiss National Science Foundation (SNF) under project numbers PP00P2 - 124879/1, 200020 - 122287. T.F is support by HIC for FAIR and by the SNF under project~no.~PBBSP2-133378. The work of G.P. is supported by the DFG under Grant No. PA 1780/2-1 and J.S.-B. is supported by the DFG through the Heidelberg Graduate School of Fundamental Physics. M. H. is supported by the SNF under project number no. 200020-132816/1. M. H. is also grateful for participating in the EuroGENESIS collaborative research program of the European Science Foundation (ESF) and the ENSAR/THEXO project. I.S. is supported by the AvH foundation via a Feodor Lynen fellowship and wishes to acknowledge the HPCC of MSU and the iCER. The authors are additionally supported by CompStar, a research networking program of the ESF, and the scopes project funded by the SNF grant. no. IB7320-110996/1."
    },
    "1112/1112.0574_arXiv.txt": {
        "abstract": "We report \\emph{Warm Spitzer} full-orbit phase observations of WASP-12b at 3.6 and 4.5 $\\mu$m. This extremely inflated hot Jupiter is thought to be overflowing its Roche lobe, undergoing mass loss, accretion onto its host star, and has been claimed to have a C/O ratio in excess of unity.   We are able to measure the transit depths, eclipse depths, thermal and ellipsoidal phase variations at both wavelengths. The large amplitude phase variations, combined with the planet's previously-measured day-side spectral energy distribution, is indicative of non-zero Bond albedo and very poor day--night heat redistribution. The transit depths in the mid-infrared ---$(R_{p}/R_{*})^{2} = 0.0123(3)$ and $0.0111(3)$ at 3.6 and 4.5 $\\mu$m, respectively--- indicate that the atmospheric opacity is greater at 3.6 than at 4.5~$\\mu$m, in disagreement with model predictions, irrespective of C/O ratio. The secondary eclipse depths are consistent with previous studies: $F_{\\rm day}/F_{*} = 0.0038(4)$ and 0.0039(3) at 3.6 and 4.5 $\\mu$m, respectively. We do not detect ellipsoidal variations at 3.6 $\\mu$m, but our parameter uncertainties ---estimated via prayer-bead Monte Carlo--- keep this non-detection consistent with model predictions. At 4.5 $\\mu$m, on the other hand, we detect ellipsoidal variations that are much stronger than predicted. If interpreted as a geometric effect due to the planet's elongated shape, these variations imply a 3:2 ratio for the planet's longest:shortest axes and a relatively bright day--night terminator. If we instead presume that the 4.5 $\\mu$m ellipsoidal variations are due to uncorrected systematic noise and we fix the amplitude of the variations to zero, the best fit 4.5 $\\mu$m transit depth becomes commensurate with the 3.6 $\\mu$m depth, within the uncertainties. The relative transit depths are then consistent with a Solar composition and short scale height at the terminator.   Assuming zero ellipsoidal variations also yields a much deeper 4.5 $\\mu$m eclipse depth, consistent with a Solar composition and modest temperature inversion. We suggest future observations that could distinguish between these two scenarios. ",
        "introduction": "Thermal phase variations are a powerful way to constrain the climate on exoplanets. Such observations have been made for non-transiting short-period planets \\citep{Cowan_2007, Crossfield_2010}, but are most potent when combined with transit and eclipse observations for edge-on systems, because of the additional knowledge of the planet's inclination, mass and radius \\citep{Knutson_2007a, Knutson_2009a, Knutson_2009b}. Secondary eclipse depths provide a constraint on the planet's day-side temperature. Thermal phase variations probe the day--night temperature contrast and hence the planet's heat redistribution efficiency. If the observational cadence and signal-to-noise ratio are sufficiently high, phase variations are also sensitive to the offset between the noon meridian and the planet's hottest local stellar time, hence constraining wind speed and direction.  By considering eclipse depths at a variety of wavelengths for a sample of 24 transiting planets, \\cite{Cowan_2011b} estimated their day-side effective temperatures, hence placing a joint constraint on the Bond albedo and heat recirculation efficiency of these planets. That study found that typical hot Jupiters exhibit a variety of albedo/recirculation efficiencies, but planets with substellar equilibrium temperatures greater than $T_{0}\\approx2700$~K all seem to have lower albedo and/or recirculation efficiency. In other words, the hottest transiting giant planets have a qualitatively different climate than the merely hot Jupiters, but it is not known whether this is due to a difference in albedo, circulation, or both. Direct measurements of hot Jupiter geometric albedos from optical secondary eclipse observations span more than an order of magnitude and do not resolve this degeneracy.  In this paper, we break the albedo-recirculation degeneracy for WASP-12b \\citep{Hebb_2009}, one of the very hottest known exoplanets, with a day-side temperature of $\\sim3000$~K: the amplitude of thermal phase variations is a direct measure of the planet's day-night temperature contrast and hence heat transport efficiency.  If the night-side temperature is high, then the planet's albedo must be exceedingly low to be consistent with its high day-side temperature.  If, on the other hand, the night-side temperature is low, then the planet has an albedo in the tens of percent.  WASP-12b has been a fascinating planet since its discovery. The discrepant timing of its secondary eclipse indicated that the planet had a slight eccentricity \\citep{Lopez-Morales_2010}, but subsequent eclipse \\citep{Campo_2011, Croll_2011} and radial velocity \\citep{Husnoo_2011} observations have all but ruled this out. Nevertheless, the planet's short-period orbit (1.1 days; just outside its star's Roche limit) and inflated radius (1.8 $R_{J}$) led to the prediction that it is tidally distorted \\citep{Ragozzine_2009, Leconte_2011, Budaj_2011}, and undergoing Roche-lobe overflow followed by accretion onto its host star \\citep{Li_2010, Lai_2010}. The putative early ingress of an ultraviolet transit observed by \\cite{Fossati_2010} seems to support this prediction, but may also be explained in terms of a leading bow shock from material streaming off the planet \\citep{Vidotto_2010, Llama_2011}.  More recently, \\cite{Madhusudhan_2011} used the wavelength-dependance of mid-infrared eclipse depths of \\cite{Croll_2011} and \\cite{Campo_2011} to constrain the atmospheric composition of WASP-12b, and found it has a carbon to oxygen ratio (C/O) greater than unity, unlike Solar System planets, or the assumed composition of extrasolar planets. Those findings rested heavily on the relative eclipse depths at 3.6 and 4.5 $\\mu$m. Our observations of eclipses and transits at these wavebands should be able to reinforce or rule out the high C/O scenario. ",
        "conclusions": "We obtained \\emph{Warm Spitzer} full-orbit phase observations of WASP-12b at 3.6 and 4.5 $\\mu$m, allowing us to measure the transit depths, eclipse depths, thermal and ellipsoidal phase variations at both wavelengths. We are able to push \\emph{Warm Spitzer} photometry to within 10--20\\% of the Poisson limit, but there are two important caveats:  A) Removing intra-pixel sensitivity variations (IPSVs) from the data is inherently a model-dependent endeavor. This means that we must specify not only an IPSV model, but also an astrophysical model before getting close to the quoted precision. The simultaneous fit to astrophysical and systematic effects makes it difficult to produce a ``clean'' lightcurve independent of astrophysical assumptions.  For example, we obtain very different thermal phase variation parameters depending on how we correct for systematics, and it is difficult to distinguish between these scenarios based solely on goodness-of-fit. Instead, we must resort to a number indirect clues as to which IPSV-removal scheme is more trustworthy.  B) There is still red noise in our residuals, no matter how we remove IPSVs.  This remaining red noise is the dominant source of uncertainty for all of our astrophysical parameters.  We find that WASP-12b exhibits large-amplitude thermal phases ---indicative of poor day--night heat transport and a moderate Bond albedo--- but also an unexpectedly large phase offset at 3.6 $\\mu$m.  We do not detect ellipsoidal variations at 3.6 $\\mu$m, while we detect an unexpectedly strong signal at 4.5 $\\mu$m.  This leads us to two possible hypotheses:  1) If we take the 4.5~$\\mu$m ellipsoidal variations at face value, we find: deeper transits at 3.6 $\\mu$m as compared to 4.5 $\\mu$m, inconsistent with either Solar or enhanced CO models; eclipse depths consistent with previous studies. If the 4.5 $\\mu$m ellipsoidal variations are astrophysical in nature, it indicates that the planet is far more distorted than predicted, and exhibits a bright terminator. In this scenario, the 3.6 $\\mu$m ellipsoidal variations are attenuated due to detector systematics, possibly throwing off the 3.6 $\\mu$m transit depth as well.  2) If instead we presume that the 4.5 $\\mu$m ellipsoidal variations are caused by detector systematics and set them to zero ---the null hypothesis--- we find: transit depths consistent with a Solar composition and short atmospheric scale height at the planet's terminator; eclipse depths consistent with a Solar composition and a modest temperature inversion; ellipsoidal variations in line with predictions.  The null hypothesis is attractive in its simplicity, but requires that we were very unlucky; follow-up \\emph{Warm Spitzer} observations would have different systematics (the PSF would fall on different regions of the pixels) and could settle the question of ellipsoidal variations.  It is likely that near infrared transit spectroscopy could break the composition degeneracy, or at least determine the atmospheric structure of WASP-12b; if the planet has a short scale-height at the terminator it will lend credence to the null hypothesis.  Further optical transit photometry will be useful in pinning down the transmission spectrum and refining geometrical parameters; if $a/R_{*}<3$, then the planet could very well be more distorted than predicted, making the large ellipsoidal variations more plausible.  Optical eclipse measurements from the ground or from space might confirm the moderate albedo of the planet.  The planet is hypothesized to be losing mass to its host star.  If this is indeed the case, the presence of an accretion disk, accretion stream and impact hot spot may necessitate a more holistic model to properly interpret observations."
    },
    "1112/1112.6214_arXiv.txt": {
        "abstract": "The number of methods used to study the properties of galaxies is increased, and testing these methods is very important. Galactic globular clusters (GCs) provide an excellent medium for such test, because they can be considered as simple stellar populations. We present ages and metallicities for 40 Galactic GCs as determined from three publicly available techniques, including colour, Lick-index and spectrum-fitting methods, based on Bruzual \\& Charlot evolutionary population synthesis (EPS) models. By comparing with the ages obtained from colour-magnitude diagrams (CMDs) and metallicities obtained from spectra of stars, we are able to estimate the ability of these methods on determination of GCs$^{'}$ parameters, which is absolutely necessary. As a result, we find that: (i) for the metallicity, our derived metallicities agree with those derived from the spectra of stars, Lick-index method is suitable to study metallicity for the stellar population systems in the range of $-1.5\\la$[Fe/H]$\\la-0.7$ and spectrum-fitting method is suitable to study metallicity for the stellar population systems in the range of $-2.3\\la$[Fe/H]$\\la-1.5$; (ii) for the age, these three methods have difficulties in age determination, our derived ages are smaller (about 2.0\\,Gyr, on average) than the results of CMDs for all these three methods. We use Vazdekis and Maraston models to analyze whether our results are dependent on EPS models, and find that the tendency of these two models is the same as that of Bruzual \\& Charlot models. Our results are independent of the EPS models. In addition, our test is based on the old GCs and our conclusions may hold for old stellar population systems.   Besides the age-metallicity degeneracy, we examine the possible effects of other factors (horizontal branch morphology, blue straggler stars, binary interactions and $\\alpha-$enhancement) and give quantitative analysis of the influences of these factors on age determinations (except for $\\alpha-$enhancement). For colour and spectrum-fitting methods, the age can be underestimated about $0.0-3.0$\\,Gyr, $0.0-2.0$\\,Gyr, and $0.0-3.0$\\,Gyr due to influences of horizontal branch, blue straggler and binary stars, respectively. And for Lick-index method, the lower limit of maximal change of age is 6.0\\,Gyr, 5.0\\,Gyr and 3.0\\,Gyr due to influences of horizontal branch, blue straggler and binary stars, respectively. ",
        "introduction": "\\end{flushleft} \\end{figure} \\begin{figure} \\begin{flushleft} \\includegraphics[bb=80 35 470 785,height=8.5cm,width=4cm,clip,angle=270]{zhy2.ps} \\caption{H$\\beta$ versus three indices (Mg\\,$b$, $<$Fe$>$ and [MgFe]$^{'}$) for the models of Thomas, Maraston \\& Johansson (2011a), the ages are from $1$ to $15$ Gyr and [Fe/H] are from $-2.3$ to $0.67$ that are labeled in the figure. The gray solid and red dashed lines stand for [$\\alpha$/Fe]$=0.0$ and 0.5, respectively.}\\label{entropy-rlof} \\end{flushleft} \\end{figure}  Fig.\\,1 shows an example of the method of using contour map of $\\chi^2$ to obtain age and metallicity for NGC\\,6235. The star stands for the $\\chi^2_{min}$ and corresponds to the obtaining $\\tau$ and [Fe/H]. The shaded area maps 1$\\sigma$ confidence region, and the $\\Delta$$\\tau$ and $\\Delta$[Fe/H] are obtained based on this shaded area.  \\subsection{Lick-index method}  Lick/IDS indices have been used to break the age-metallicity degeneracy (Trager et al. 2000; Terlevich \\& Forbes 2002; Gallazzi et al. 2005) of SP systems. In this paper, we select the Balmer (H$\\beta$, H$\\delta_{A}$, H$\\delta_{F}$, H$\\gamma_{A}$ and H$\\gamma_{F}$), Mg\\,\\emph{b}, Fe5270 and Fe5335 indices to study the parameters of GCs (Balmer indices are age-sensitive indices, Mg\\,\\emph{b}, Fe5270 and Fe5335 are metal-sensitive indices). The H$\\delta$ and H$\\gamma$ are high order Balmer indices which measured with a narrower ($\\sim$\\,20\\,${\\rm\\AA}$; marked by subscript F) and a wider ($\\sim$\\,40\\,${\\rm\\AA}$; marked by subscript A) central bandpass \\citep{wort97}. Among these five Balmer indices, H$\\beta$ and H$\\gamma_{F}$ are little sensitive to $\\alpha/$Fe and the other three indices are sensitive to $\\alpha/$Fe. And in Section 5.4, we will discuss the influence of $\\alpha$-enhancement on age determination for Lick-index method. For the three metal-sensitive indices, we adopt the following definition: [MgFe]$^{'}$\\,=\\,[Mg\\,\\emph{b}(0.72\\,$\\times\\,$Fe5270\\,+\\,0.28\\,$\\times$\\,Fe5335)]$^{1/2}$, which defined by Thomas, Maraston \\& Bender (2003). The [MgFe]$^{'}$ has been found to be good tracer of metallicity. Adopting the models of Thomas, Maraston \\& Johansson (2011a) we present this character in Fig.\\,2, the grids of H$\\beta$ versus Mg\\,$b$, $<$Fe$>$ ($<$Fe$>=$1/2(Fe5237$+$Fe5335)) and [MgFe]$^{'}$ with two types of $\\alpha$/Fe ([$\\alpha$/Fe]$=0.0$ and $0.5$), the gray solid and red dashed lines stand for [$\\alpha$/Fe]$=0.0$ and $0.5$, respectively. From them we can see that the influence of $\\alpha$/Fe ratio is obvious for H$\\beta$ versus Mg\\,$b$ and $<$Fe$>$ grids, and is small for H$\\beta$ versus [MgFe]$^{'}$ grid. The Mg\\,$b$ increases and $<$Fe$>$ decreases with increasing $\\alpha$/Fe, and the combination index [MgFe]$'$ counteract the effect of $\\alpha$/Fe ratio on Mg$b$, Fe5237 and Fe5335. This certify that the [MgFe]$'$ is independent of $\\alpha$/Fe and can be considered as good tracer of metallicity \\citep{thom03}.   \\begin{figure*} \\includegraphics[bb=20 20 600 780,height=16cm,width=9.5cm,clip,angle=270]{zhy3.ps} \\caption{The distribution of GCs on the Balmer indices and [MgFe]$^{'}$ planes. Model grides from BC03 are shown: age is constant for solid line (top to bottom, 1, 2, 4, 5, 8, 10, 12 and 15 Gyr), and metallicity is constant for dashed line (left to right: [Fe/H]$=-2.3, -1.7, -0.7, -0.4$, $0.0$ and $+0.4$). Line strengths of Galactic GCs are plotted on SSP grides. The red open squares, blue open circles and purple open pentacles stand for GCs with red, blue and unknown HBR, respectively. In the top left plane (H$\\beta$ versus [MgFe]$^{'}$ plane), we give these values of age and metallicity.}\\label{entropy-rlof}  \\end{figure*} In Fig.\\,3, we show the distribution of GCs on five Balmer (H$\\beta$, H$\\delta_{A}$, H$\\delta_{F}$, H$\\gamma_{A}$ and H$\\gamma_{F}$) indices versus [MgFe]$^{'}$ planes. We can see that the majority of GCs lie beyond the grid covered by H$\\beta$ (this problem does not involve all Balmer indices, Thomas et al. 2011a), which would be inconsistent with the typical universe age 13.6\\,Gyr. This is a general problem (see Fig.\\,2 of Mendel et al. 2007) for all studies of using H$\\beta$ index to study the GCs based on existing models. This problem is also clearly demonstrated by the $'$zero-point$'$ offset (some GCs lie beyond the grid of H$\\beta$ and most GC ages are greater than 14\\,Gyr). Although these problems have been pointed by some groups (Vazdekis et al. 2001; Schiavon et al. 2002; Cenarro et al. 2008), but few works have been done on studying the specific sources of this problem or quantified their effects on the derived ages and metallicities. Similar to the technique described in Puzia et al. (2005), the ages and metallicities for individual GCs are computed as the weighted mean of the parameters derived from five Balmer indices versus [MgFe]$^{'}$ grids.  Due to the influence of hot blue HB stars, the isochrones for old ($\\tau$\\,$>$\\,8.0\\,Gyr) stellar populations with metallicity below [Fe/H]$\\,=-0.7$ tend to overlap (Maraston \\& Thomas 2000; Mendel et al. 2007), this nature is reflected by H$\\beta$ in panel (a) of Fig.\\,3. This brings an ambiguity in determining ages and metallicities of GCs and unnaturally extends the age distributions for old ages for Lick-index method. Note that, for the GCs lying beyond the Lick indices grids, we assume the values of GCs are 15.0\\,Gyr.  \\subsection{spectrum-fitting method} With the increase of spectroscopic data, people use spectra to study the properties of SP systems. So far there are many kinds of full spectrum-fitting techniques used in such study (e.g. Cid Fernandes et al. 2005; Mathis, Charlot \\& Brinchmann 2006; Tojeiro et al. 2007; K08).  We employ a modified version of the \\tiny{STARLIGHT} \\normalsize{code} by Cid Fernandes et al. (2005) to analyze the observed spectra of Galactic GCs. The \\tiny{STARLIGHT} \\normalsize{}code is originally used to study the properties of galaxy, and is achieved by fitting the observed spectrum \\emph{F$\\rm_{O}$} with a model spectrum \\emph{F$\\rm_{M}$} that mixed by \\emph{N$_{\\star}$} SSPs with different ages and metallicities from the BC03 models. The code is carried out with a simulated annealing plus Metropolis scheme (Cid Fernandes et al. 2001), which searches for the minimum  \\begin{equation} \\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\chi^2= \\sum\\limits_{\\lambda}\\, [(F\\rm_{O}-\\emph{F}\\rm_{M})\\omega_{\\lambda}]^2, \\end{equation} where $\\omega$$^{-1}_{\\lambda}$ is the error in \\emph{F$\\rm_{O}$}. The line-of-sight stellar motions are modeled by a Gaussian distribution centered at velocity \\emph{v}$_{\\star}$ and with dispersion \\emph{$\\sigma$}$_{\\star}$. The output includes the distributions of stellar age, metallicity, extinction, velocity dispersion and stellar mass. ",
        "conclusions": "We use colour, Lick-index and spectrum-fitting methods based on BC03 model to obtain the ages and metallicities for Galactic GCs, and compare them with those determined from spatially resolved observations (such as CMDs, spectra of stars). We also compare the ages and metallicities derived from other two models (Vazdekis and Maraston models) to test whether our results are dependent on EPS models. Before displaying the results of this work, we should introduce some pretreatments of model, sample data and methods. The details are given as follows.  \\subsection{Pretreatments}   \\begin{itemize} \\item For the EPS model, as said above, the BC03 model provides six metallicities, we linearly interpolate colours, Lick indices and spectra of SSPs that span [Fe/H] from $-$2.3 to $+$0.4 in increments of 0.1 dex.   \\item In order to discuss the influence of HB morphology on age determinations in Section 5.1, we divide the GCs into three groups according to the value of HB ratio (HBR \\footnote{ HBR\\,$=(B-R)/(B+R+V)$, B is the number of blue HB stars, R is the number of red HB stars and V is the number of RR Lyrae variable.}) from the Harris$'$ catalogue. One is the blue HBR type (HBR\\,$>$0), the second is red HBR type ( HBR\\,$<$0) and the third is unknown HBR (HBR is not given).  \\item For the colour method, we do extinction corrections for observed magnitudes of GCs by adopting the extinction curves of \\citet{schl98} and the $E(B-V)$ of Harris$'$ catalogue (the 8th column of Table A1).  \\item For the Lick-index method, in this work we calculate the Lick indices from the spectra of Galactic GCs by degrading to wave-dependent resolution of the Lick system (see Section 4.4 of BC03), without adopting the results by using fitting functions.  \\begin{figure*} \\includegraphics[bb=20 20 580 730,height=16.0cm,width=9.5cm,clip,angle=270]{zhy4.ps} \\caption{Top panels: the comparisons between the metallicities obtained by three methods and the spectra of stars for Galactic GCs. Symbols have the same meanings in Fig.3. Bottom panels are the comparisons between the ages obtained by three methods and the CMDs. Left, middle and right panels are comparison results for methods of colour, Lick-index and spectrum-fitting, respectively. }\\label{entropy-rlof}  \\end{figure*}  \\item For the spectrum-fitting method, we adopt a modified version of the \\tiny{STARLIGHT}\\normalsize{} code. Because it is generally accepted that the GC can be represented by a SSP, only one component in the base. We fit the observed GC spectrum \\emph{F$\\rm_{O}$} with each of the SSP spectrum \\emph{F$_{i}$}, do not fit the observed spectrum with a linear combination of \\emph{N$_{\\star}$} SSPs. In this study, following \\citet{cid05}, we find the best fitting SSP, which matches a given observed spectrum of GC, through a standard $\\chi^2$ minimization procedure:  \\begin{equation} \\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\chi^2_{i}= \\sum\\limits_{\\lambda}\\, [(F\\rm_{O}-\\emph{F$_{i}$})\\omega_{\\lambda}]^2, \\end{equation} instead of equation (2), where \\emph{F$_{i}$} corresponds to the spectrum of i-th SSP model. And in this work we construct a base of $18$\\,ages\\,$\\times\\,24$\\,metallicities\\,$=432$\\,SSP model spectra. The use of \\tiny{STARLIGHT}\\normalsize{} to study the integrated spectra of clusters has also been extensively discussed by Cid Fernandes \\& Gonz$\\acute{a}$lez Delgado (2010).  In the spectrum-fitting process, we use the spectra in the 3700$-$5700\\,${\\rm\\AA}$ range because the red and blue parts of the observations have poor quality. Errors on the age and metallicity estimation are also determined by the $\\chi^2$ contours.  \\end{itemize}   \\subsection{Comparison of parameters with other studies }  Using the methods and pretreatments outlined in Sections\\,3 and 4.1, we obtain ages, metallicities and corresponding errors for GCs. In order to distinguish the ages and metallicities derived from different methods, we define some parameters in this paper. $\\tau\\rm_{C}$ and [Fe/H]$\\rm_{C}$ stand for the age and metallicity determined from colour method, $\\tau\\rm_{I}$ and [Fe/H]$\\rm_{I}$ represent those obtained from Lick-index method, and $\\tau\\rm_{S}$ and [Fe/H]$\\rm_{S}$ stand for those determined from spectrum-fitting method. The parameters of literatures are represented by $\\tau\\rm_{L}$ and [Fe/H]$\\rm_{L}$. In Table A2 we list the GC IDs in the first column. In the 2nd to 7th columns, $\\tau\\rm_{C}$, [Fe/H]$\\rm_{C}$, $\\tau\\rm_{I}$, [Fe/H]$\\rm_{I}$, $\\tau\\rm_{S}$ and [Fe/H]$\\rm_{S}$ are given with errors, respectively. And in 8th and 9th columns we list the $\\tau\\rm_{L}$ and [Fe/H]$\\rm_{L}$.    In the top panels of Fig.\\,4, we give the comparison between our GC metallicities and those given in the literatures that obtain metallicities from the spectra of stars. And (a), (b) and (c) panels are the comparison results for colour, Lick-index and spectrum-fitting methods, respectively. The vertical bars are the errors in metallicities derived from three methods. From them we see that the metallicities obtained by three methods are in agreement with literature value for the whole sample. Meanwhile we also find that all GCs with blue HBR have lower metallicity than those with red HBR for these three methods, which agrees with the study of Lee, Yoon \\& Lee (2000). Our specific analysis are given as follows.\\\\ (a) For colour method, [Fe/H]$\\rm_{C}$ agrees with [Fe/H]$\\rm_{L}$ for the whole sample, but there exists discreteness. We can see that [Fe/H]$\\rm_{C}$ is a bit smaller than [Fe/H]$\\rm_{L}$ in the range of [Fe/H]$\\rm_{L}$$\\la-$1.0 and [Fe/H]$\\rm_{C}$ agrees with [Fe/H]$\\rm_{L}$ for the range of [Fe/H]$\\rm_{L}$$>-$1.0. But for the value of [Fe/H]$\\rm_{L}$ is about $-$1.1, there exist three GCs (NGC 6171, 6342 and 6652) having lager [Fe/H]$\\rm_{C}$ than [Fe/H]$\\rm_{L}$. From the [Fe/H]$\\rm_{C}$ and $\\tau\\rm_{C}$ of these GCs listed in Table A2, we see that $\\tau\\rm_{C}$ is smaller than $\\tau\\rm_{L}$, so these GCs can be affected by age-metallicity degeneracy.\\\\ (b) For the Lick-index method, the [Fe/H]$\\rm_{I}$ is perfectly in agreement with [Fe/H]$\\rm_{L}$ in the range of $-1.5\\la$\\,[Fe/H]\\,$\\la-0.7$. But there have difficulties in metallicity determination for the range of [Fe/H]\\,$\\la-1.5$, because it tends to overlap for Balmer and [MgFe]$^{'}$ indices at low-metallicity range for BC03 model. And we can see this phenomenon clearly in the panel (a) of Fig.\\,3. This indicates that Lick-index method is suitable to study metallicity for SP systems in the range of $-1.5\\la$\\,[Fe/H]\\,$\\la-0.7$ and has some difficulties in studying SP systems for the range of [Fe/H]\\,$\\la-1.5$.\\\\ (c) For the spectrum-fitting method, we find that [Fe/H]$\\rm_{S}$ matches [Fe/H]$\\rm_{L}$ in the range of $-2.3\\la$\\,[Fe/H]\\,$\\la-1.5$, and [Fe/H]$\\rm_{S}$ is smaller than [Fe/H]$\\rm_{L}$ for the range of [Fe/H]\\,$\\ga-1.5$. All these results imply that the spectrum-fitting method may be suitable to study metallicity for SP systems in the range of $-2.3\\la$\\,[Fe/H]\\,$\\la-1.5$ and [Fe/H]$\\rm_{S}$ may be smaller than [Fe/H]$\\rm_{L}$ for the range of [Fe/H]\\,$\\ga-1.5$.   On the whole, our metallicities obtained from three methods match those determined from spectra of stars in the entire metallicity range spanned by the GCs. The Lick-index method is suitable to study the metallicity in the range of $-1.5\\la$\\,[Fe/H]\\,$\\la-0.7$ and spectrum-fitting method is suitable to study the metallicity in the range of $-2.3\\la$\\,[Fe/H]\\,$\\la-1.5$.    In the bottom panels of Fig.\\,4, we compare our derived GC ages of three methods with the values of literature that determine ages from CMDs. And (d), (e) and (f) panels are corresponding to colour, Lick-index and spectrum-fitting methods, respectively. From them we can see there exists large discrepancy in age determinations, and the errors are also large for these three methods. Our specific analysis are given as follows.\\\\ (d) For the colour method, most of the GCs with blue and unknown HBR have relative lower $\\tau\\rm_{C}$ than $\\tau\\rm_{L}$, but there has no obvious tendency for those with red HBR. The age of GCs with blue HBR can be influenced by HB stars for the existence of HB stars making GCs look younger. Five GCs are found to possess extreme low $\\tau\\rm_{C}$, including two with unknown HBR (NGC 6388 and 6441) and three with red HBR (NGC 6171, 6342, and 6652). Among these GCs, three have relative lager [Fe/H]$\\rm_{C}$ than [Fe/H]$\\rm_{L}$, which can be affected by the age-metallicity degeneracy (see Fig.\\,4a). Meanwhile from the 4th and 5th columns of Table A1, we know that the other two GCs are close to the Galactic bulge, which are badly contaminated by the field stars or affected by differential reddening, so we can not get reliable parameters for them.\\\\ (e) For the Lick-index method, we can see there exist some GCs with extreme large $\\tau\\rm_{I}$ (about 15.0\\,Gyr), and except these GCs most of GCs have smaller $\\tau\\rm_{I}$ which are influenced by HB stars, especially for those with blue HBR. About eight GCs have large $\\tau\\rm_{I}$ of 15.0\\,Gyr which lie beyond the grids, including two with red HBR (NGC 1851 and 6637), one with unknown HBR (NGC 5946) and five with blue HBR (NGC 1904, 2298, 3201, 6254 and 7078). Similar to above, Lick-index method has some difficulties in parameter determinations for the range of [Fe/H]\\,$\\la-1.5$, and most of these eight GCs are metal-poor ([Fe/H]$\\rm_{I}$$\\la-$1.5), so we obtain extreme large $\\tau\\rm_{I}$ for them.\\\\ (f) For the spectrum-fitting method, most GCs with blue HBR have relative smaller $\\tau\\rm_{S}$ than $\\tau\\rm_{L}$, but there has no obvious tendency for those with red and unknown HBR. From panel (c), we know that [Fe/H]$\\rm_{S}$ is consistent with [Fe/H]$\\rm_{L}$ in the range of [Fe/H]\\,$\\la-1.5$ and most of GCs with blue HBR are metal-poor, so the $\\tau\\rm_{S}$ of GCs with blue HBR are affected by HB stars rather than age-metallicity degeneracy. For the GCs with red and unknown HBR, one part has large $\\tau\\rm_{S}$ which can be influenced by age-metallicity degeneracy, and the other part has small $\\tau\\rm_{S}$ which may be influenced by HB stars and BSs. \\begin{table*} \\begin{center} \\caption[]{Comparison of main model ingredients in each EPS model. The first, second and third rows list the stellar evolution track, stellar spectral library and IMF of the EPS models. The 4th and 5th rows show the metallicity ($Z$) and age range that these models cover, and the number of metallicities is also given in the parenthesis of the 4th row.  } \\tabcolsep=0.15in  \\label{Tab:frac} \\renewcommand{\\arraystretch}{1.5} \\begin{tabular}{|l|l|l|l|r|r|r|r|r|r|c|c|c|c|c|c|c|c|c|c|} \\hline\\hline\\noalign{\\smallskip} Models &BC03   &Vazdekis &Maraston \\\\ \\hline\\noalign{\\smallskip} Stellar evolution track   &Padova 1994  &Padova 2000  &Cassisi+Geneva \\\\ Stellar spectral library   &STELIB &MILES  &MILES\\\\ IMF                       &Chabrier &Salpeter  &Chabrier\\\\ $Z$(number)               &$0.0001-0.0500(6)$&$0.0001-0.0300(7)$&$0.0001-0.0400(5)$\\\\ Age(Gyr)                  &$0.0001-20.0000$  &$0.0630-17.1800$  &$0.0060-15.0000$\\\\    \\hline \\end{tabular} \\end{center} \\end{table*} Note that, relevant to this test is that CMD-derived ages also can carry their own problems for this test, because the CMD-derived ages depend on the adopted tracks and on whether element ratios be taken into account, which has also been discussed in \\citet{mara11}.   On the whole, all of these three methods have difficulties in age determination and we can not directly say which method is better on age determination. Except the influences of CMD-derived ages and age-metallicity degeneracy in the entire age range, there exist some uncertainties in age determinations, some factors (e.g. HB stars, BSs, binary stars and $\\alpha$-enhancement) can affect the age determinations and we will discuss them in the next Section.  \\begin{figure*} \\includegraphics[bb=20 15 580 770,height=16.0cm,width=9.5cm,clip,angle=270]{zhy5.ps}  \\caption{Comparisons between the results from three EPS models and the literatures. Top and bottom panels are the comparisons for metallicity and age, respectively. The open squares, pentacles and circles stand for the results of BC03, Vazdekis and Maraston EPS models, respectively. Left, middle and right panels are comparison results for methods of colour, Lick-index and spectrum-fitting, respectively. The dispersion (rms) of each model is also listed in the lower right corner of each panel. }\\label{entropy-rlof}  \\end{figure*}  \\subsection{Consistency checks} We use Vazdekis\\footnote{http://miles.iac.es/} \\citep[the updated v9.1 version,][]{vazd10,falc11} and Maraston \\citep{mara11} models to investigate whether our results are dependent on EPS models. The Vazdekis models use the Padova 2000 theoretical isochrones \\citep{gira00}, the MILES library \\citep{sanc06} and four types of IMF with stellar mass limits (Vazdekis et al. 2003) of 0.1 and 100 M$_{\\odot}$. This version covers the wavelength range of 3540$-$7410\\,${\\rm\\AA}$ with a spectral resolution (FWHM) of $\\sim$2.5\\,${\\rm\\AA}$ \\citep{falc11}, seven metallicities $0.0001\\leqslant Z \\leqslant0.03$ and 50 ages across the range of 0.0630$-$17.1800\\,Gyr. In this work, we select the unimodal Salpeter IMF \\citep{salp55} to study the parameters of GCs for spectrum-fitting and Lick-index methods. For the colour method, because the MILES library has relatively limited wavelength coverage and does not extend into I-band, we select the colours from previous models of Vazdekis \\citep{vazd99} to study the parameters of GCs. This does not influence our results (private discussion).  The Maraston models assume a stellar evolution prescription, which consists of the isochrones and stellar tracks by Cassisi, Castellani \\& Castellani (1997) for ages larger than $\\sim$30.0Myr and by Ganeva \\citep{scha92} for younger populations. Sets of models have been also computed with Padova stellar evolutionary models \\citep{gira00}. Four different libraries of flux-calibrated empirical stellar spectra and three types of IMF have been considered. We select the MILES library and the \\citet{chab03} IMF to study the parameters of GCs. These updated models cover the wavelength range of 3500$-$7430\\,${\\rm\\AA}$ with a spectral resolution (FWHM) of $\\sim$2.54\\,${\\rm\\AA}$ \\citep{beif11} and five metallicities $0.0001\\leqslant Z \\leqslant0.04$ with different ages across the range of 0.0060$-$15.0000\\,Gyr (Maraston \\& Str\\\"{o}mb\\\"{a}ck 2011, and references therein). Just as said above, we use the updated models of \\citet{mara11} based on MILES library for spectrum-fitting and Lick-index methods. And for the colour method we select the models of \\citet{mara05}, which based on the $BaSeL$ library, to study the parameters of GCs. In Table 1 we list the comparison of main model ingredients in each EPS model. The first, second and third rows list the stellar evolution track, stellar spectral library and IMF of the EPS models that we used in this study. The 4th and 5th rows show the metallicity ($Z$) and age range that these models cover, and the number of metallicities is also given in the parenthesis of the 4th row.   We use these two EPS models to obtain the parameters of GCs based on three methods described above and compare the results of these two models with that of the literatures. In Fig.\\,5, we show the comparisons between the parameters of GCs obtained from tree methods based on three EPS models (BC03, Vazdekis and Maraston) and the literatures, in which the open squares, pentacles and circles stand for the results of BC03, Vazdekis and Maraston EPS models, respectively. The dispersion (rms) is also listed in the lower right corner of each panel, rms$\\rm_{bc03}$, rms$\\rm_{vazd}$ and rms$\\rm_{mara}$ stand for the rms of BC03, Vazdekis and Maraston EPS models, respectively. We find that the metallicities obtained from three methods of these two EPS models (Vazdekis and Maraston) have an agreement with those of the literatures in the entire metallicity range. From the top panels (a, b and c panels), we find that the comparisons between metallicities obtained by three methods and the spectra of stars for these two models are similar to that of BC03 models. For Lick-index method in panel (b), we see that the rms$\\rm_{vazd}$ and rms$\\rm_{mara}$ are smaller than rms$\\rm_{bc03}$, this may be due to the adoption of MILES for these two models and STELIB for BC03 models \\citep{mara11}. For the spectrum-fitting method in panel (c), the [Fe/H]$\\rm_{S}$ of Vazdekis model is larger than that of the other two models, so the rms$\\rm_{vazd}$ is smaller than rms$\\rm_{bc03}$ and rms$\\rm_{mara}$. From these rms, we can know that the dispersion of colour method is larger than that of other two methods. All these can not influence our conclusions, Lick-index method is suitable to study metallicity in the range of $-1.5\\la$\\,[Fe/H]\\,$\\la-0.7$ and spectrum-fitting method is suitable to study metallicity in the range of $-2.3\\la$\\,[Fe/H]\\,$\\la-1.5$. From bottom panels, we can see that all these three methods have difficulties in age determinations for these two EPS models and the rms is larger than that of metallicity. There also exist some GCs with extreme large $\\tau\\rm_{I}$ (about 15.0\\,Gyr) for these two models, and most of GCs have relative smaller $\\tau\\rm_{C}$, $\\tau\\rm_{I}$ and $\\tau\\rm_{S}$ than $\\tau\\rm_{L}$, which is similar to that of BC03 models. The comparisons of derived metallicities and ages of individual GCs show that these three models make different parameter  predictions. However, the whole tendency for these three methods of these three EPS models is nearly the same. This indicates that our conclusions are independent of the EPS models.  \\begin{figure} \\includegraphics[bb=95 45 590 690,height=8.0cm,width=6.0cm,clip,angle=270]{zhy6.ps} \\caption{Taking NGC 7078 as an example to illustrate the method of selecting HB stars and BSs from the CMD. The HB stars are marked with red asterisks and BSs are marked with blue circles.}\\label{entropy-rlof} \\end{figure}   \\begin{figure*} \\includegraphics[bb=35 15 555 780,height=16.cm,width=9.5cm,clip,angle=270]{zhy7.ps} \\caption{Theoretical Balmer index$-$\\emph{T}$\\rm_{eff}$ relations obtained from the B\\tiny{LUERED} \\normalsize{} library for different log \\emph{g} (solid, dashed, dotted, and dash-dot-dot-dot lines stand for log \\emph{g}$=3.0, 4.0, 4.5$ and $5.0$, respectively) at [Fe/H]$=-1.0$.}\\label{entropy-rlof} \\end{figure*}  We have investigated the utility of three methods (colour, Lick-index and full spectrum-fitting) with the same EPS models (BC03 models) for estimating ages and metallicities of Galactic GCs. We also compared our results with those estimated in the literatures from other methods, and the main results of this study are as follows.  (1) Our results show that Galactic GCs are almost old metal-poor SP systems.  (2) Metallicities determined from these three methods are in agreement with those of literatures in the entire range spanned by GCs. The Lick-index method is suitable to obtain metallicity for SP systems in the range of $-1.5\\la$[Fe/H]$\\la-0.7$ and spectrum-fitting method is suitable to study metallicity for SP systems in the range of $-2.3\\la$[Fe/H]$\\la-1.5$.  (3) There exist some discrepancies between our ages and the literatures, we can not directly say which method is more suitable for age determination. Through comparing with the literatures we find our results are younger than the literatures on average, especially for GCs with blue HBR. Many factors can affect the age determinations and make GCs become younger (such as the HB morphology, BSs, binary stars, $\\alpha-$enhancement).  (4) We use Vazdekis and Maraston models to investigate whether our results are dependent on EPS models. The comparisons of derived metallicities and ages for individual GCs show that three models make different parameter predictions. However, the whole tendency for these three methods of these three EPS models is nearly the same. All these indicate that our above conclusions are independent of the EPS models.  (5) In this work we use the old Galactic GCs to test these three methods, and our results may hold for old SP system.  (6) We also study the quantitative influences of HB, BSs and binary stars on age determinations for three methods. The existence of all these stars can make the GCs look younger. For the colour and spectrum-fitting methods, the age can be underestimated about 0.0$-$3.0\\,Gyr, 0.0$-$2.0\\,Gyr and 0.0$-$3.0\\,Gyr due to influences of HB, BSs and binary stars, respectively. And for Lick-index method, the lower limit of maximal change of age is 6.0\\,Gyr, 5.0\\,Gyr and 3.0\\,Gyr due to influences of horizontal branch, blue straggler and binary stars, respectively.  (7) For the Lick-index method, we also investigate the influence of $\\alpha-$enhancement on age determination, and find that the $\\alpha-$enhancement has little effect on age determination for Galactic GCs."
    },
    "1112/1112.4211_arXiv.txt": {
        "abstract": "We are extending our search for faint PNe in the LMC to include the outer 56 deg$^{2}$ area not covered in the original UKST survey of the central 25deg$^{2}$ region. Candidate PNe have been selected using the Magellanic Cloud Emission Line Survey (MCELS) and the first round of observations has yielded 93 new LMC PNe while confirming the 102 previously known PNe in the outer LMC. We plan to continue our spectroscopic object identification program until we cover all our remaining candidates in the survey area. These observations, providing medium and high resolution spectra from $\\lambda$3650\\AA~to $\\lambda$6900\\AA~will additionally be used to measure fluxes for a series of research projects including luminosity functions, abundances and LMC kinematics. ",
        "introduction": "In 2006 Reid and Parker added 460 planetary nebulae (PNe) to the 169 previously known across the central 25deg$^{2}$ of the LMC. These candidates were assigned a probability rating which relied on spectroscopic and optical confirmation. With the assistance of increased high resolution spectroscopy and NIR imaging from Spitzer Space Telescope SAGE data (see \\cite[Meixner et al. 2006]{Meixner2006}) we have been able to either confirm or re-classifying most of the lower classed ``possible PNe''. The resulting large number of PNe now known in the central LMC (\\cite[Reid \\& Parker, 2006a,b]{Reid2006b}) have yielded significant advances in our knowledge of the central LMC's kinematical sub-structure (rotation, inclinations, transverse velocity) and raised interesting questions regarding the kinematical structure of the outer regions. Our access to the MCELS survey (\\cite[Smith et al. 1998]{Smith1998}) provides the opportunity to achieve equivalent results across a much larger 84deg$^{2}$ of the LMC, critical for a complete determination of kinematics, abundance gradients and central star properties (eg. \\cite[Reid \\& Parker, 2010b]{Reid2010b}) . We have already used AAOmega to spectroscopically verify 93 new PNe to add to the 102 previously known in the outer LMC (eg. Fig.\\,\\ref{fig1}). New PN radial velocities are being compared to other tracers and the HI gas disk. These are being added to existing kinematic data to create gradients and verify models. With a near complete LMC PNe census to V=22, an unbiased LMC PNe luminosity function (PNLF) is being built in order to identify any population sub-trends or `dips' while providing an accurate bright-end cut off which is used as a standard candle. ",
        "conclusions": ""
    },
    "1112/1112.4482_arXiv.txt": {
        "abstract": "The usage of the high-level scripting language Python has enabled new mechanisms for data interrogation, discovery and visualization of scientific data.  We present \\yt{}\\footnote{Available at \\texttt{http://yt-project.org/}}, an open source, community-developed astrophysical analysis and visualization toolkit for data generated by high-performance computing (HPC) simulations of astrophysical phenomena. Through a separation of responsibilities in the underlying Python code, \\yt{} allows data generated by incompatible, and sometimes even directly competing, astrophysical simulation platforms to be analyzed in a consistent manner, focusing on physically relevant quantities rather than quantities native to astrophysical simulation codes.  We present on its mechanisms for data access, capabilities for MPI-parallel analysis, and its implementation as an \\textit{in situ} analysis and visualization tool. ",
        "introduction": "\\label{sec:introduction} In the last decade, multiphysics astrophysical simulations have increased exponentially in both sophistication and size \\cite{ 2005Natur.435..629S, % 2008JPhCS.125a2008K, % 2009JCoPh.228.6833R, % 2008MNRAS.390.1326O, % 2007arXiv0705.1556N, % 2010arXiv1008.4368K, % 2010arXiv1008.2801A, % 2010arXiv1002.3660K, % 2000ApJS..131..273F, % 2007ApJ...659L..87A}; % however, the software tools to mine those simulations have not kept pace. Typically, methods for examining data suffer from a lack of agility, discouraging exploratory investigation.  To accommodate this, massively parallel visualization tools such as VisIT and ParaView \\cite{visit_paper, paraview_paper} have been repurposed as domain-specific astrophysical tools. This repurposing, while effective, does not benefit from domain-specific analysis or data structures.  The lack of domain-specific quantitative analysis tools designed for astrophysical data leads to the development of specialized tools by individual researchers or research groups, most of which are never shared outside the research group.  This can substantially inhibit collaboration between different groups--even those using the same simulation code.  This fractionation of the astrophysical community demonstrates a clear need for a flexible and cross-code software package for quantitative data analysis and visualization.  In this paper we present \\yt{} \\cite{yt_full_paper}, a data analysis and visualization package that works with several astrophysical simulation codes.  \\yt{} is developed openly and is freely available at \\texttt{http://yt-project.org/}.  It has been designed to be a common platform for simulation analysis, so that scripts can be shared across groups and analysis can be repeated by independent scientists\\footnote{A platform for sharing scripts is provided with yt, with command-line helpers, at \\texttt{http://hub.yt-project.org}.}.  By making this tool available, we hope not only to encourage cross-group collaboration and validation of results, but to remove or at least greatly lower the barrier to entry for exploratory simulation analysis.  \\yt{} provides mechanisms for conducting complete analysis pipelines resulting in publication quality figures and data tables, as well as the necessary components for constructing new methods for examining data.  The concepts for data handling and representation in \\yt{} are certainly not new, but their application to astrophysical data enables complex, detailed analysis pipelines to be shared between individuals studying disparate phenomena using disparate methods.  This enables and even encourages reproducibility and independent verification of results.  We have built this analysis and visualization code in Python, using NumPy \\cite{numpy_paper} for fast mathematical operations, \\texttt{mpi4py} for MPI-parallelism \\cite{MPI4PY:PAPER1, MPI4PY:PAPER2}, and optionally Matplotlib for 2D visualization \\cite{matplotlib_paper}.  Additionally, several core library routines in \\yt{} such as the AMR volume rendering, multi-dimensional binning, and file access routines, are written in Cython.  In addition to utilizing community-developed Python modules, \\yt{} is itself a Python module suitable for direct scripting or access as a library.  A community of users and developers has grown around the project, with over 20 committers in the history of the project, and it has been used in numerous published papers and posters. (See, for example, \\cite{ 2010ApJ...715.1575S, 2010ApJ...721.1105B, 2009Sci...325..601T, 2011ApJ...738...54K, 2011arXiv1108.4427Z, 2011ApJ...737...63A, 2011MNRAS.414.2297I, 2011ApJ...735...49M, 2011ApJ...731...59C}.)  In order to accomodate the diverse computing environments on which astrophysical simulations are run, \\yt{} was designed to use primarily off-screen rendering and scripting interfaces, although several smaller tools are provided for specific, interactive visualization tasks.  The former method is well-suited to remote visualization and can be run via a job execution queue on a batch-compute cluster, such as those on which the underlying simulation are run.  \\yt{} is subdivided into several sub-packages for data handling, data analysis, and plotting.  This modularity encourages the creation of reusable components for multi-step analysis operations.  While work continues on the exploratory, post-processing methods of data analysis and visualization that \\yt{} was originally designed for, current development has focused on analysis during the course of a simulation, or so-called \\textit{in situ} analysis.  This allows for high-cadence analysis to be conducted without writing data to disk.  Future simulations, such as those to be conducted on petascale machines, will require analysis and visualization during the simulation rather than exclusively as a post-processing technique.  In this paper, we will describe the mechanisms that \\yt{} provides for accessing data (\\S \\ref{sec:mechanisms}), methods of interacting with \\yt{} (\\S \\ref{sec:interacting}), the visualization techniques offered by \\yt{} (\\S \\ref{sec:visualization}),  the parallelism strategy for data analysis and generation of visualizations (\\S \\ref{sec:parallelism}), and end with a discussion of the process of embedding \\yt{} in running simulation codes and how this will be inverted in the future (\\S \\ref{sec:embedding}). ",
        "conclusions": "\\label{sec:conclusions}  The \\yt{} project is fully free and open source software, released under the GNU General Public License, with no dependencies on external code that is not also free and open source software.  The development process occurs completely in the open at \\texttt{http://yt-project.org/}, with publicly-accessible source control systems, bug tracking, mailing lists, and regression tests.  Building a community of users has been a priority of the \\yt{} development team, both to encourage collaboration and to solicit contributions from new developers; both the user and developer communities are highly distributed around the world. \\yt{} is developed using Mercurial\\footnote{\\texttt{http://mercurial.selenic.com/}}, a distributed version control system that enables local versioned development and encourages users to make and contribute changes upstream.  Many of the operations conducted in \\yt{}: fluid analysis, phase diagrams, volume rendering, parallelism, and in situ analysis could feasibly be applied to domains other than astrophysics.  We intend to generalize the underlying code base such that it can be applied to many other data formats in astrophysics, and ultimately we hope to provide these tools and techniques to domains other than astrophysics.  Our first steps toward this, providing a generic and arbitrary data loader, have shown that it is feasible.  Future versions of \\yt{} will generalize fields and particle handling, and should make this process much easier.  The creation of a freely available, publicly inspectable platform for simulation analysis allows the community to disentangle the coding process from the scientific process.  Simultaneously, by making this platform public, inspectable and freely available, it can be openly improved and verified.  The availability and relatively approachable nature of \\yt{}, in addition to the inclusion of many simple analysis tasks, reduces the barrier to entry for young scientists.  Rather than worrying about the differences between Enzo and FLASH hierarchy formats, or row versus column ordering, or HDF4 versus HDF5 versus unformatted fortran data formats, researchers can focus on understanding and exploring their data.  More generally, however, by orienting the development of an analysis framework as a community project, the fragmentation of methods and mechanisms for astrophysical data analysis is greatly inhibited.  Future generations of simulations and simulation codes will not only benefit from this collaboration, but they will require it."
    },
    "1112/1112.3094_arXiv.txt": {
        "abstract": "There is great interest in numerical relativity simulations involving matter due to the likelihood that binary compact objects involving neutron stars will be detected by gravitational wave observatories in the coming years, as well as to the possibility that binary compact object mergers could explain short-duration gamma-ray bursts.  We present a code designed for simulations of hydrodynamics coupled to the Einstein field equations targeted toward such applications.  This code has recently been used to study eccentric mergers of black hole-neutron star binaries.  We evolve the fluid conservatively using high-resolution shock-capturing methods, while the field equations are solved in the generalized-harmonic formulation with finite differences.  In order to resolve the various scales that may arise, we use adaptive mesh refinement (AMR) with grid hierarchies based on truncation error estimates.  A noteworthy feature of this code is the implementation of the flux correction algorithm of Berger and Colella to ensure that the conservative nature of fluid advection is respected across AMR boundaries. We present various tests to compare the performance of different limiters and flux calculation methods, as well as to demonstrate the utility of AMR flux corrections. ",
        "introduction": "\\label{intro}  The interface between strong field gravity and matter dynamics promises to be one of the important frontiers in the coming years. A new generation of gravitational wave detectors (LIGO~\\cite{LIGO}, GEO~\\cite{GEO},TAMA~\\cite{TAMA}, and VIRGO~\\cite{VIRGO}) are now operational, and within the next few years are expected to reach sensitivities that will allow observations of the Universe in gravitational radiation for the first time.  The prime targets of these observations are compact object (CO) binaries composed of combinations of black holes (BHs) and neutron stars (NSs). Modeling of such sources is a crucial ingredient to realize the promise of gravitational wave astronomy.  Even if an event is detected with a high signal-to-noise ratio, reconstructing the dynamics of the system that produced the signal cannot be done directly but instead will require template banks of theoretical waveforms informed by numerical simulations.  Compact object mergers involving NSs are expected to be significant sources of not only gravitational radiation, but also possible progenitors for short-gamma-ray bursts (SGRBs)~\\cite{npp92,2005Natur.437..851G,2005ApJ...630L.165L} and other electromagnetic and neutrino counterparts~\\cite{Metzger:2011bv}. Efforts are already underway to use potential gravitational wave sources as triggers for searches for electromagnetic transients~\\cite{LIGO_EM,LIGO_EM2}. Observations would help constrain evolutionary models for the progenitor stars and their environments. Perhaps most intriguingly, the observations would give clues to the equation of state (EOS) of matter at nuclear densities (as in NS interiors), which cannot be probed in laboratories on Earth and is not fully understood at the theoretical level (for a broad discussion see for example~\\cite{glendenning}). The reason that the gravitational wave signature could contain information about the matter EOS (and other details about the internal structure of neutron stars) is that the EOS in general has a significant effect on the bulk motion of matter, and it is this bulk motion that is the mechanism by which gravitational waves are produced. Several studies to date have looked into this issue, suggesting the imprint of the EOS on the gravitational waves will be strong enough to detect~\\cite{Kiuchi:2011re,Pannarale:2011pk,jocelyn,Oechslin:2007gn,Tsui:2006tr,Kokkotas:2005vr,Shibata:2005xz,Shibata:2005ss,Faber:2004fs,Bejger:2004zx,Lai:1996sv,Xing:1994ak} (though, in some cases, the expected frequencies are higher than the range to which the current generation of ground-based detectors are most sensitive, thus limiting the information which can be extracted). While CO binaries containing NSs are a particularly interesting class of sources involving general relativistic (GR) hydrodynamics, they are by no means the only such systems. Examples of additional systems that have already been considered include BH accretion tori~\\cite{montero,korobkin,2011PhRvD..84b4024F,2012ApJ...744...45B} and NS-white dwarf mergers~\\cite{vasileos}.  Thoroughly modeling systems like those described above would require evolution of the spacetime, the photon and neutrino radiation fields, and the magnetized, relativistic fluid. Even a minimalistic treatment, with the Einstein equations coupled to the equations of relativistic hydrodynamics, represents a complex, nonlinear system of partial differential equations.   Numerical simulations are thus essential for exploring such strong field, dynamical systems. There is a long history of adapting successful techniques for simulating Newtonian hydrodynamics to relativistic and general relativistic fluids which we will not attempt to summarize (see~\\cite{grhydro_lr} for a review of general relativistic hydrodynamics). Instead, we will briefly attempt to place the code described in the present paper in the context of other recent codes developed for fluids on evolving spacetimes. \\footnote{Note that our focus is restricted to codes which handle dynamically evolving gravitational fields.  Such codes, however, frequently owe much to earlier, fixed-background evolution codes (see~\\cite{grhydro_lr}).  In addition, advancements such as GR-hydro with multipatch grids~\\cite{thor} and with GPUs~\\cite{horizon} have recently been made with fixed-background codes.}  Several of these codes~\\cite{mhdcode1,SACRA,Giacomazzo:2007ti,Jena_nsns,Bode:2009mt} solve the field equations in the BSSN formulation~\\cite{sn,bs}.  The remainder~\\cite{matt,Anderson:2006ay} use the generalized-harmonic formulation~\\cite{garfinkle,gh3d} which we also employ; unlike our code, however, these groups convert to a fully first-order formulation~\\cite{new_lindblom_et_al}.  Most groups use finite-difference methods for the metric evolution and a conservative, high-resolution shock-capturing (HRSC) scheme for the hydro evolution; these unigrid algorithms are then interfaced with some sort of adaptive mesh refinement (AMR). A notable exception for the metric evolution is~\\cite{matt}, which employs pseudospectral methods for the metric and then interpolates to a finite-volume grid for the fluid.  Some groups have implemented the MHD equations in full GR; since these codes all make use of conservative HRSC methods, they may be principally differentiated by how they meet the challenge of preserving the $\\nabla \\cdot {\\mathbf B} = 0$ constraint. (A straightforward finite-difference evolution of the magnetic field would generically lead to magnetic monopoles and, hence, unphysical behavior.) {\\tt WhiskyMHD} employs constrained transport~\\cite{Giacomazzo:2007ti} for this purpose, which preserves the constraint to machine accuracy, whereas the code of~\\cite{chawla} uses hyperbolic divergence cleaning.  Constrained transport, however, requires special interpolation at refinement level boundaries in order to preserve the constraint. The Illinois group found that a vector-potential formulation of the MHD equations works well when coupled to AMR~\\cite{illinoisNewMHD}. This is because the constraint is preserved by construction with the vector-potential, even with the restriction and prolongation operations of AMR (see also~\\cite{illinoisMHD2} for a thorough examination of the electromagnetic gauge condition).  Studies indicate that magnetic fields do not significantly affect the gravitational dynamics of CO mergers (see {\\em e.g.} \\cite{chawla}), but they could be critical for understanding EM counterparts including the possible formation of a SGRB engine. A new method to treat the MHD equations was recently presented in~\\cite{lehner_new}, where ideal MHD is used in high matter density regions ({\\em e.g.} inside a NS), while the force-free approximation is used elsewhere ({\\em e.g.} the magnetosphere of a NS). The authors applied the method to study the collapse of magnetized hypermassive NSs (which could be formed via binary NS mergers) and suggested that intense EM outbursts could accompany such events.  Besides MHD, the other major advances in the physical model for numerical relativity codes have been in the arena of microphysics. While the $\\Gamma=2$ EOS was the community standard for quite some time, most codes now allow for a nuclear theory-based EOS~\\cite{matt2,shibataBNS} and/or use various parametrized, piecewise polytropic EOSs inspired by the range of plausible nuclear EOSs~\\cite{shibataBHNS4,illinoisWDNS}.  These advances in EOS description primarily affect the cold NS structure, but the group developing the {\\tt SACRA} code has also begun to account for neutrino transport via a simplified leakage scheme~\\cite{Kiuchi:2011re,2011PhRvL.107e1102S}.  The same group has also made available a formulation for a more accurate truncated moment scheme with a variable Eddington factor closure~\\cite{truncated_moment}, which shows much promise for numerical relativity simulations with neutrino physics beyond the leakage approximation.  Another category of GR hydrodynamics codes employs the conformal-flatness approximation, which is particularly useful when supernova simulations are the target application.  An example is CoCoNuT/VERTEX, which incorporates relativistic hydrodynamics, conformally flat gravity, and ray-by-ray neutrino transport~\\cite{coconut}. The code of~\\cite{cerda_duran} employs a similar scheme for hydrodynamics and gravity but adds a test magnetic field; this code has been used to study the magnetorotational instability in supernovae.  Newtonian (and semi-Newtonian)~\\cite{lee_kluzniak,magma}, conformally flat~\\cite{faber,oechslin}, and fixed-background~\\cite{laguna} SPH codes represent an important, orthogonal approach to studying CO interactions. SPH has an advantage over Eulerian schemes when a large range of spatial scales is involved.  Such a situation may arise in CO mergers when material is stripped from a star in a tidal interaction and forms an extended tail.  On the other hand, Eulerian codes are the standard approach when strong shocks are present, as would arise in binary NS mergers or disk circularization.  (Recent progress has been made, however, in applying SPH to situations with relativistic shocks~\\cite{Rosswog_SRHD}.)  In addition, SPH has not (to our knowledge) yet been coupled to a code which evolves the full Einstein equations.  Nonetheless, comparisons between Eulerian and SPH results could prove very useful on a problem-by-problem basis to characterize the errors in both methods.  Though current efforts in GR simulations involving matter tend to focus on increasingly complex physical models, there remain many unanswered questions in the astrophysics of compact objects that can be addressed with a code which solves the Einstein equations coupled to perfect fluid hydrodynamics.  We have thus focused our code development on hydrodynamics in full GR, while maintaining a flexible infrastructure to accommodate additional physics modules in the future.  We evolve the field equations in the generalized-harmonic formulation using finite differences.  The fluid is evolved conservatively using one of several different shock-capturing techniques we test here. We have also implemented the hydrodynamical equations in a manner that is independent of EOS.  We make use of AMR by dynamically adapting the mesh refinement hierarchy based on truncation error estimates of a select number of the evolved variables.  We also utilize Berger and Colella~\\cite{bc89} style flux corrections (also known as ``refluxing'') in order to make the use of AMR compatible with the conservative nature of the hydrodynamic equations. Though AMR flux corrections have been implemented in other astrophysical hydrodynamics codes (such as Athena~\\cite{ATHENA}, CASTRO~\\cite{CASTRO}, {\\tt Enzo}~\\cite{EnzoMHD}, and FLASH~\\cite{FLASH}), to our knowledge this algorithm has not been used previously for hydrodynamics simulations in full general relativity.\\footnote{Note that ``flux correction'' here refers to the enforcement of conservation at AMR boundaries, not the recalculation of fluxes with a more dissipative scheme to preserve stability as in Athena~\\cite{athena_rmhd}.} A further noteworthy feature of our implementation is that we store corrections to the corresponding fluid quantity integrated in the volume of a given cell instead of the flux, allowing for easy implementation within a computational infrastructure that supports cell-centered but not face-centered distributed data structures. The code described here has recently been applied to studying BH-NS mergers with eccentricity as may arise in dense stellar systems such as galactic nuclear clusters and globular clusters~\\cite{ebhns_letter,ebhns_paper}.  In the remainder of this paper we outline our computational methodology for simulating hydrodynamics coupled to the Einstein field equations and describe tests of this methodology. In Sec.~\\ref{methods} we review the generalized-harmonic approach to solving the field equations and present our methods for conservatively evolving a perfect fluid coupled to gravity, including our method for inverting the conserved quantities to obtain the primitive fluid variables and the implementation of flux corrections to enforce the conservation of fluid quantities across AMR boundaries.  In Sec.~\\ref{tests} we present simulation results which test these methods, highlight the strengths and weaknesses of various shock -capturing techniques, and demonstrate the utility of the flux correction algorithm. ",
        "conclusions": "Numerous scenarios that fall within the purview of general relativistic hydrodynamics are still mostly unexplored---especially CO mergers involving neutron stars. There is a rich parameter space, of which large areas remain uncharted due to uncertainty or potential variability in BH and NS masses, BH spin and alignment, the NS EOS, and other aspects. Beyond the pure hydrodynamics problem, the roles of magnetic fields and neutrino physics are just beginning to be explored by various groups, and we expect to add support for such physics to our code in the future. The potential applications of robust and flexible numerical algorithms for evolving hydrodynamics together with the Einstein field equations are manifold.  With this in mind, we have implemented methods for conservatively evolving arbitrary EOSs, in particular for converting from conserved to primitive variables without knowledge of derivatives; and we have implemented numerous reconstruction and flux calculation methods that can be used interchangeably to meet problem specific requirements.  Though accurate treatment of shocks may not be crucial for BH-NS mergers (where shocks are not expected to be dynamically important), the same is not true of NS-NS binaries, especially eccentric ones where the stars may come into contact during nonmerger close encounters~\\cite{Roman}. We have also taken care to implement a flux correction algorithm that preserves the conservative nature of hydrodynamical advection across AMR boundaries.  Though strict conservation is not, strictly speaking, essential (since any nonconservation would be at the level of truncation error), it is an especially appealing property when studying, for example, CO mergers as potential SGRB progenitors. After merger, material that did not fall into the black hole --- typically on the order of a few percent of the original NS mass --- will fill a large volume making up an accretion disk and potentially unbound material.  Though accurately tracking this material is not important for the gravitational dynamics, it is critical for characterizing potential EM counterparts to the merger."
    },
    "1112/1112.1058_arXiv.txt": {
        "abstract": "{} {We present the analysis of a large sample of gamma-ray burst (GRB) X-ray light curves in the rest frame to characterise their intrinsic properties in the context of different theoretical scenarios.} {We determine the morphology, time scales, and energetics of $64$ long GRBs observed by \\emph{Swift}/XRT \\emph{without} flaring activity. We furthermore provide a one-to-one comparison to the properties of GRBs \\emph{with} X-ray flares.} {We find that the steep decay morphology and its connection with X-ray flares favour a scenario in which a central engine origin. We show that this scenario can also account for the shallow decay phase, provided that the GRB progenitor star has a self-similar structure with a constant envelope-to-core mass ratio $\\sim 0.02-0.03$. However, difficulties arise for very long duration ($t_p\\gtrsim10^4$ s) shallow phases. Alternatively, a spinning-down magnetar whose emitted power refreshes the forward shock can quantitatively account for the shallow decay properties. In particular we demonstrate that this model can account for the plateau luminosity vs. end time anticorrelation.} {} ",
        "introduction": "Since the launch of the \\emph{Swift} satellite in 2004 \\citep{2004ApJ...611.1005G}, the evolution of the gamma-ray burst (GRB) X-ray light curves has appeared to be more complex than previously thought. The different shapes of the light curves \\citep{2006ApJ...642..389N} and the existence of X-ray flares superimposed on the smooth continuum \\citep{2007ApJ...671.1903C,chinca10} up to very late times \\citep{2008A&A...487..533C,2011A&A...526A..27B} have prompted a large number of studies conducted on different samples of X-ray light curves in order to single out their main properties \\citep{2006ApJ...642..354Z,2006ApJ...647.1213O,2007ApJ...662.1093W,2007ApJ...666.1002Z,2007ApJ...670..565L,2008ApJ...675..528L,2009ApJ...707..328L,2008ApJ...689.1161G,2009ApJ...698...43R,2009MNRAS.397.1177E}.  Analysis of the observational properties of X-ray flares \\citep{2007ApJ...671.1903C,chinca10,giantflares10,2011A&A...526A..27B,2011MNRAS.410.1064M} allowed us to establish a connection with the prompt emission. In particular, the extension of the prompt lag-luminosity relation to X-ray flares strongly suggests that X-ray flares and prompt emission pulses are produced by the same mechanism \\citep{giantflares10}. The nature of the underlying X-ray \\emph{continuum} is still debated. It was immediately clear soon after the first set of observations by \\emph{Swift}/XRT \\citep{2005SSRv..120..165B} and redshift measurements that the light curves differed from a simple power-law behaviour and that some of the characteristics repeat in a systematic way  \\citep[see][]{2006ApJ...642..389N,2006ApJ...642..354Z,2006ApJ...647.1213O}. After six years of data collection, we now have a much larger sample for investigating the morphology of the X-ray light curves, and a rather large subsample of redshift measurements for this analysis in the rest frame. The rest frame allows us to understand the intrinsic properties, the timescales involved, and the energetics. In fact, the parameters derived are directly related to the physics of the phenomenon, and the distributions are not affected by the distribution of redshift. This is the approach we adopt in this work.  We select light curves observed by XRT without flares, in order to avoid any kind of contamination from other emission components. When necessary we compare our results with the sample analysed in \\citet{2011MNRAS.410.1064M} of GRBs with flares, providing for the first time a one-to-one comparison of light curves \\emph{with} and \\emph{without} flares.  The initial steep decay of the afterglow is commonly due to the tail of the prompt emission \\citep{2000ApJ...541L..51K}, and the plateau phase is often connected with external shock emission with an injection of energy. However, the standard model for GRBs suffers from several problems \\citep[for a review see][]{2009arXiv0911.0349L}. Here we aim at exploring alternative mechanisms to shape the light curve, such as the accretion of the stellar material left behind by the collapse of the massive star \\citep{2008MNRAS.388.1729K} or the rotational energy of a newly-born magnetar \\citep[see e.g.][and references therein]{2010arXiv1012.0001M}. We show indeed that an accretion model would explain the X-ray light curves well enough, as well as the properties of flares as discussed by \\citet{2011MNRAS.410.1064M}. The plateau phase agrees with the assumption that new energy is injected from a spinning down neutron star into the forward shock. In particular within this framework we demonstrate that this model can account for the anticorrelation between the plateau luminosity and its end time found by \\citet{2008MNRAS.391L..79D,2010ApJ...722L.215D}.  In Sect.~\\ref{data_analysis} we describe the light curve extraction and the luminosity calibration procedures. In Sect.~\\ref{sample} we describe the sample selection criteria, the fitting procedure and the main results about the morphology and spectral properties of the sample. In Sect.~\\ref{energy} we present the analysis of the energy of the sample compared with the prompt emission energy. In Sect.~\\ref{flares} we compare the present sample with the light curves with flares analysed in \\citet{2011MNRAS.410.1064M}. In Sect.~\\ref{discussion} we discuss our findings and in Sect.~\\ref{conclusions} we draw our main conclusions. We adopt standard values of the cosmological parameters: $H_\\circ=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.27$, and $\\Omega_{\\Lambda}=0.73$. Errors are given at $1\\, \\sigma$ confidence level unless otherwise stated. Times are in rest frame unless otherwise stated. ",
        "conclusions": "We analysed $64$ long GRB X-ray light curves observed by XRT with redshift measurement that do not exhibit flaring activity. This allowed us to characterise the morphology and energetics of the sample in order to constrain the mechanism that produces the X-ray continuum. The light curves of the sample can be divided into four morphological Types, with similarities between different parts. When necessary, we compared our sample with the one considered in \\citet{2011MNRAS.410.1064M} of $44$ GRBs that exhibit flaring activity. We found that  \\begin{itemize} \\item a large fraction of the light curves are on average consistent with the high latitude emission from a structured jet, but all of them show a strong or moderate spectral evolution. The accretion model accounts for the observed range of temporal decays and the spectral evolution of the steep decay phases of Types Ib and II; \\item the injection phase observed in Type Ia and II can be interpreted as power emitted from spinning-down ultramagnetised neutron star that refreshes the forward shock. We applied such a model to our sample and found that it fits the data well enough; \\item this scenario accounts for the observed $L_p-t_p$ anticorrelation \\citet{2008MNRAS.391L..79D,2010ApJ...722L.215D}: we reproduced the normalisation, slope, and scatter of the anticorrelation starting from plausible values of the parameters. The consistency of Type Ia shallow decay with the anticorrelation within $99\\%$ confidence level suggests that they can be reproduced by this model with the same distribution of $B$ and $P$ as for Type II; \\item a second phase of accretion of a progenitor with a structure core envelope can originate the shallow phase in Types Ia and II. The correlation between the prompt emission total isotropic energy with the shallow decay X-ray energy implies a self-similar structure for the progenitor star, with $\\sim 0.02-0.03$. However in this scenario, difficulties arise for the very long duration ($t\\gtrsim 10^4$ s) of Type Ib shallow decay; \\item $70\\%$ of the light curves with flares are Type II, while only $23\\%$ of the non-flare sample are, confirming our previous finding \\citep{giantflares10}. This suggests that Type II light curves are related to the suitable conditions for the occurrence of flares. The accretion model has the merit of proposing a unique mechanism to explain the steep decay and the flaring activity; \\end{itemize}  The magnetar model is able to explain the different morphologies of the X-ray continuum and the properties of the shallow decay. The main limit in the accretion model is the difficulty of achieving very long timescales as observed in Type Ib shallow decay. The behaviour of the steep decay and the suggested connection with the flaring activity \\citep[see also][]{2009MNRAS.395..955B,giantflares10,2011MNRAS.410.1064M} indicates that the steep decay may not be viewed as simply the tail of the prompt emission."
    },
    "1112/1112.6164_arXiv.txt": {
        "abstract": "We study the role of viscosity and the effects of a magnetic field on a rotating, self-gravitating fluid, using Newtonian theory and adopting the ideal magnetohydrodynamic approximation. Our results confirm that viscosity can generate vorticity in inhomogeneous environments, while the magnetic tension can produce vorticity even in the absence of fluid pressure and density gradients. Linearizing our equations around an Einstein-de Sitter cosmology, we find that viscosity adds to the diluting effect of the universal expansion. Typically, however, the dissipative viscous effects are confined to relatively small scales. We also identify the characteristic length bellow which the viscous dissipation is strong and beyond which viscosity is essentially negligible. In contrast, magnetism seems to favor cosmic rotation. The magnetic presence is found to slow down the standard decay-rate of linear vortices, thus leading to universes with more residual rotation than generally anticipated. ",
        "introduction": "Most, if not all, astrophysical systems rotate and it is conceivable that vorticity can play an important role in a number of situations. So far, however, the majority of the related studies has been confined to relatively small-scale systems. Although the question of universal rotation and how large this may be has been posed more than once (see, e.g.,~\\cite{BJS,KO,KHB,GSFB,JBEGH,SC} and references therein), in cosmology the role of vorticity is usually bypassed. One of the reasons for doing so is that typical inflationary scenarios are irrotational, since the vacuum fluctuations do not generate vorticity~\\cite{LL}. Another, is that linear rotational perturbations are known to decay as the inverse square of the cosmological scale factor, which implies strong dilution for any vortex-like distortions that could happen to exist. Dissipative effects during the subsequent radiation era, like Silk damping and neutrino free-streaming, may also have a negative impact on large-scale rotation, though (to the best of our knowledge) studies of this nature are not yet available in the literature.  In cosmology, where the terms vorticity and rotation are used interchangeably, vorticity is usually generated from nonlinear effects. This is one additional reason for not taking vorticity into account in many cosmological perturbation studies. Inhomogeneous non-barotropic (e.g., non-adiabatic) media, for example, can generate rotation at the second perturbative level~\\cite{CMM,CM}, and so can fluid viscosity~\\cite{MB,D-SB,Fetal}. Nevertheless, resorting to nonlinear effects is not always necessary, since there are agents that can induce rotation still at the linear level. For instance, a coherent magnetic field (`$B$-field') can act as a source of cosmic vorticity~\\cite{W,TM}. The opposite is also theoretically possible. Thus, following~\\cite{H}, a number of authors have explored the likelihood that rotational distortions can lead to cosmic magnetogenesis (see, e.g.,~\\cite{Metal,MMNR,KMST,GR} and references therein). In addition, both the fluid viscosity and the $B$-field can affect the evolution of rotational perturbations and thus change the standard picture, which usually assumes non-magnetized, perfect media. Therefore, in principle at least, the effects of the aforementioned sources could lead to more residual vorticity than previously anticipated.  The present paper uses Newtonian gravity to study the effects of viscosity and those of a magnetic field on rotation. We adopt the covariant, Lagrangian approach to fluid dynamics, which has been proved helpful in several (mainly relativistic) cosmological studies (see~\\cite{E1} for an introduction to the Newtonian version of the formalism). First, we look into the effects of fluid viscosity and then consider those of a coherent magnetic field. Our results, confirm those of earlier studies claiming that both of the aforementioned agents can act as sources of cosmic rotation. We show, in particular, that viscosity can only generate vorticity in media with an inhomogeneous density distribution. This is not the case for the $B$-field, which does not (in principle at least) require the presence of density perturbations. Assuming the Newtonian analogue of the Einstein-de Sitter universe as our background model, we find that viscosity can only generate vorticity at second order. Magnetism, on the other hand, can do the same at the linear perturbative level. Moreover, the irreducible decomposition of variables and equations that the covariant formalism introduces, allows us to identify the agents leading to vorticity generation, as well as the way they do so. Viscosity, for example, sources rotational distortions via the coupling of the shear with certain types of density perturbations. Alternatively, the magnetic field acts primarily through the tension component of the Lorentz force.  Our next step is to investigate how viscous and/or magnetized media can affect the linear evolution of rotational perturbations. In the first case, we find that viscosity always leads to vorticity dissipation, though the effects are generally scale dependent. We identify, in particular, a characteristic length beyond which viscosity is essentially negligible and below which its role becomes progressively more important. In particular, well inside this ``viscous scale'' vorticity dissipates very quickly, while far beyond that scale the fluid kinematics are dominated by the background Hubble expansion. There, one recovers the standard inverse-square decay law for rotational perturbations. The size of the aforementioned viscous scale depends on the specifics of the cosmic fluid, such as the average velocity and the mean free-path of the species involved. When dealing with non-relativistic matter, however, the typical domain affected by viscosity is very small (in cosmological terms).  In contrast to viscosity, the presence of magnetic fields seems to have a favorable effect on linear rotational perturbations. In addition to generating vorticity at the linear order, the $B$-field slows down the standard decay rate of rotation in perturbed Friedmann-Robertson-Walker (FRW) universes. As mentioned earlier, vorticity typically depletes as $a^{-2}$, with $a$ representing the cosmological scale factor. When the linear magnetic effects are accounted for, however, the vorticity decay-rate changes to $a^{-3/2}$. The implication of this effect, which comes entirely from the magnetic tension, is that magnetized universes, can have more residual vorticity than their non-magnetic counterparts. An analogous, conclusion was also reached indirectly, through the relativistic analysis of linear density inhomogeneities in magnetized cosmologies~\\cite{TM}. To the best of our knowledge, this is the first direct indication that magnetism ``favors'' both the generation and the survival of cosmic rotation. ",
        "conclusions": "Realistic media do not behave like perfect fluids and magnetic fields, of various scale and strength, appear to be everywhere in the universe. At the same time,  cosmic rotation still poses a number of fundamental questions that remain as yet unanswered. In this respect, it would be interesting to see whether viscosity and magnetism have something new to say about vorticity in general. So far, both agents have been known to act as sources of rotation, though their effects on pre-existing vorticity are not clear yet.  The present article adopts the Newtonian approach to investigate the rotational behavior of a viscous fluid and subsequently that of a magnetized medium. At first, our analysis is fully nonlinear and confirms that both viscosity and magnetism can generate vorticity. We also show that the former of these two agents does so through the coupling of the shear with certain types of density inhomogeneities. Magnetism, on the other hand, does not necessarily need an inhomogeneous matter distribution to source rotational distortions. The $B$-field generates vorticity by means of its Lorentz force, especially via the tension component of the latter.  Assuming the Newtonian analogue of the Einstein-de Sitter universe and then linearizing our equations around this background, we looked into the effects of viscosity and magnetism on rotational cosmological perturbations. We found that the former effects are negative. Viscosity generally leads to vorticity dissipation, though (for non-relativistic species) its role is typically confined to relatively small lengths. Nevertheless, further study is required to establish the actual range of the dissipative effects. Here, we have taken the first step in this direction by identifying the characteristic ``viscous scale'', below which viscosity dominates and beyond which the universal expansion takes over.  In contrast, the magnetic presence seems to favor rotation. More specifically, after linearizing our equations around the aforementioned Einstein-de Sitter background, we found that the $B$-field slows down the standard decay rate of rotational cosmological perturbations. In particular, vorticity no longer obeys the standard $a^{-2}$ decay-law, but drops as $a^{-3/2}$ ($a$ is the cosmological scale factor). Also, the magnetic effects are scale-independent, which means that their domain of influence is decided by the coherence length of the $B$-field. Analogous results were also obtained when studying vortex-like perturbations in the density distribution of a magnetized medium. Overall, it appears as though magnetism favors both the emergence and the survival of cosmic rotation. Put another way, magnetized universes should contain more residual vorticity than their magnetic-free counterparts.  In closing, we should also note that it would be worth-studying the above described effects using a relativistic approach. This will enable one to involve highly energetic particles, look deep into the pre-decoupling universe and also allow for spacetimes more complex than the Einstein-de Sitter universe."
    },
    "1112/1112.4343_arXiv.txt": {
        "abstract": "{The composition of Galactic Cosmic Rays (GCR) presents strong similarities to the standard (cosmic) composition, but also noticeable differences, the most important being the high isotopic ratio of \\neo, which is $\\sim$5 times higher in GCR than in the Sun. This ratio provides key information on the GCR origin.} {We investigate the idea that GCR are accelerated by the forward shocks of supernova explosions, as they run through the presupernova winds of the massive stars and through the interstellar medium.} { We use detailed wind and core yields of rotating and non-rotating models of massive stars with mass loss, as well as simple models for the properties of the forward shock and of the circumstellar medium.} { We find that the observed GCR \\neo \\ ratio can be explained if GCR are accelerated  only during the early Sedov phase, for shock velocities $>$1500-1900 km/s. The acceleration efficiency is found to be of the order of 10$^{-6}$-10$^{-5}$, i.e. a few particles out of a million encountered by the shock escape the SN at GCR energies. We also show quantitatively that the widely publicized idea that GCR are accelerated in superbubbles fails to account for the high \\neo \\ ratio in GCR.} {} ",
        "introduction": "Supernova (SN) shocks are generally thought to be the main accelerator of the bulk of Galactic Cosmic Rays (GCR). Indeed, the power of GCR in the Milky Way is estimated to 10$^{41}$ ergs/s, corresponding to 10-20 \\% of the kinetic power of Galactic supernovae (assuming canonical values of 2 SN per century, each one releasing 10$^{51}$ ergs of kinetic energy).  The site of the acceleration of GCR remains debatable today, despite more than five decades of theoretical and observational studies (e.g. Strong et al. 2007 and references therein). Over the years, it has been suggested that GCR are accelerated in 1) SN remnants (either by the forward or the reverse shock or both), 2) the interstellar medium (ISM), 3) the winds of massive stars, 4) the interiors of superbubbles, excavated by the massive star winds and the subsequent SN explosions of an OB association.  Each one of the proposed sites has its own advantages and shortcomings, regarding the energetics and/or the composition of accelerated matter. For instance, it has been argued that the hot, low density environment of a superbubble minmizes radiative losses of SN shocks and energy losses of accelerated particles, thus allowing the latter to reach substantial energies, up to the \"knee\"of the GCR spectrum (e.g. Parizot et al. 2004). On the other hand, reverse SN shocks running into the SN interior carry insufficient energy to explain the bulk GCR energetics (Ramaty et al. 1997). Moreover, they should accelerate $^{59}$Ni, a product of explosive nucleosynthesis which is unstable to e$^-$-capture (with a lifetime of 10$^4$ yr) and which has not been detected in GCR (Wiedenbeck et al. 1999), while a rapid acceleration would render it practically stable and thus detectable.  It was realized early on that the {\\it elemental} composition of GCR differs significantly from the one of the ISM. Those differences may provide valuable information on the origin of GCR particles and, perhaps, on the site - and even the mechanism - of acceleration. Volatiles behave  differently from refractories: the former display a mass-dependent enrichment with respect to H, which reaches a factor of 10 for the heaviest of them; the latter are all overabundant (w.r.t. H) by a factor of 20, while C and O display intermediate overabundances, by factors of 9 and 5, respectively (e.g. Wiedenbeck 2007 and references therein).  This complex pattern is now thought to result not from ionization effects (as suggested in Cass\\'e and Goret 1978, and further developped by Meyer 1985) but rather from effects related to elemental condensation temperature (Meyer et al. 1997): refractories are locked in dust grains, which are sputtered by repeated SN shocks and the released ions are easily picked-up and accelerated (Ellison et al. 1997). This, quite elaborate, scheme, which builds on earlier ideas by e.g. Cesarsky and Bibring (1981), accounts quantitatively for most of the observed features of GCR source composition; still, it leaves unanswered the key issue about the acceleration site of GCR (and how it affects the composition of accelerated matter).  The most conspicuous feature of GCR source composition   is undoubtely the high isotopic \\neo \\ ratio. Its  value was measured since the late 1970ies (Garcia-Munoz, Simpson and Wefel 1979, Wiedenbeck and Greiner 1981). The most accurate  measurement today, obtained from analysis of the CRIS instrument,  leads  to a best estimate  (Binns et al. 2008) of 0.387 $\\pm$ 0.007 (statistical) $\\pm$ 0.022 (systematic). This is 5.3$\\pm$0.3 times the value of the (\\neo)$_{\\odot}$ \\ ratio in the solar wind. Contrary to the case of the elemental source GCR abundances, which may be affected by various physico-chemical factors (first ionization potential, condensation temperature, etc.) isotopic ratios can only be affected by nucleosynthetic processes and thus provide crucial information on the origin of cosmic ray particles. It should be noticed that up to now there is now clear evidence for any other GCR isotopic ratio to differ from solar, with the potential exception of $^{58}$Fe/$^{56}$Fe, which is estimated to 1.5$\\pm$0.3 times solar (Binns et al. 2008).  Soon after the discovery of the anomalous GCR \\neo \\ ratio, Cass\\'e and Paul (1982) suggested that it could be explained by a mixture of $\\sim$2\\% of material from the wind of a WC star to 98\\% of material with standard composition. In early He-burning, $^{14}$N (produced through the CNO cycle in the previous H-burning phase) is transformed almost totally in \\neb \\ through $^{14}$N($\\alpha,\\gamma$)$^{18}$F($\\beta^+$)$^{18}$O($\\alpha,\\gamma$)$^{22}$Ne. He-burning products (like $^{12}$C  and \\neb) are expelled by the stellar winds of  massive stars during their WC phase. The observed GCR \\neo \\ ratio is obtained by assuming dilution of WC material with matter of standard composition. Subsequent studies put the aforementionned idea on a quantitative basis, with the use of detailed models of the evolution and nucleosynthesis of massive, mass losing stars (Maeder 1983, Prantzos 1984, Meyer 1985, Prantzos et al. 1987). In those studies, the acceleration site of GCR was considered as  decoupled from the nucleosynthesis site, and unrelated to the fraction of admixtured WC material.  Higdon and Lingenfelter (2003) evaluated quantitatively the \\neo \\ ratio within a superbubble, created by the collective action of stellar winds and SN shockwaves. They adopted  stellar wind yields for \\nea \\ and \\neb \\ from the models of Schaller et al. (1992) and SN yields from the models of Woosley and Weaver  (1995). They found that the \\neo \\ ratio in the superbubble decreases with time (since \\neb \\ from the winds dominates the evolution of \\neo \\ at early times) and that its time average value is compatible with the GCR source \\neo \\ inferred from observations. In a subsequent paper, Lingenfelter and Higdon (2007) recognised that the Schaller et al. (1992) yields of \\neb \\ were highly overestimated\\footnote{The reason was the excessively high mass loss rates adopted in that work.} and, consequently, \"... new detailed calculations of the expected GCR isotopic ratio are called for...\", but they did not attempt such a re-evaluation. In the meantime, Binns et al. (2005), using updated wind yields of massive stars with rotation (from the Geneva group, see Sec. 2.2), found good agreement between the observed \\neo \\ ratio and an admixture of $\\sim$20\\% material from WR stars  with 80\\% material of standard composition. According to  Binns et al. (2008), since WR stars are evolutionary products of OB stars, such an agreement  \"...suggests that OB associations within superbubbles are the likely source of at least a substantial fraction of GCR\".  Howewer,  theoretical studies in the past 10 years  are based mostly on the  paradigm of GCR being accelerated in SN remnants, not in superbubbles, e.g. \\cite{Ptuskin05, Berezhko06,Berezhko09,Ptuskin10,Caprioli10,Schure10,Ellison11} \\ and references therein. The kinetic energy of the bulk motion of the forward shock of the SN explosion is converted to GCR  energy through diffusive shock acceleration. The process is highly non-linear and involves the dynamical reaction of both the accelerated particles and of the magnetic field on the system. Those studies usually take into account the fact that the SN explosion often occurs within the cavity excavated in the interstellar medium (ISM) by the wind of the massive star prior to the explosion (\\cite{Biermann01}); however, the structure of the circumstellar environment in that case is quite complex and simplified models are used for its description. Although \\cite{Caprioli11b} considered the composition of GCR (H,He, CNO, MgSiAl, Fe)  resulting from such an acceleration site, none of those studies considered the \\neo \\ ratio.   In this work we study the \\neo \\ ratio of GCR accelerated by the forward shocks of   SN  explosions, as they run through the presupernova winds of  massive stars and through the interstellar medium. We consider the whole mass spectrum of massive stars (from $\\sim$10 to 120  \\ms), including  stars with either small or large mass losses prior to their explosions. We consider stellar properties (masses of winds, ejecta, yields etc.) from recent models with mass loss and or without rotation (from Hirschi et al. 2005 and Limongi and Chieffi 2006, respectively), the former having larger \\neb \\ enhancements in their winds. We adopt a simplified prescription (suggested in Ptuskin and  Zirakashvili 2005 and reformulated in Caprioli 2011) to describe the structure of the  circumstellar medium at the time of the explosion and we consider that GCR start being accelerated in the Sedov-Taylor (ST)  phase of the SN remnant (see e.g. Ptuskin et al. 2010).  By requiring  the resulting IMF averaged \\neo \\ ratio to equal the observed one  $R_{Obs}$=(\\neo)$_{GCR}$/(\\neo)$_{\\odot}$=5.3$\\pm$0.3 we are able to constrain the forward shock velocity to  values $>$1900 km/s for rotating stars (and to $>$2400 km/s for non rotating ones), i.e. we find that GCR are accelerated during the early ST phase, lasting for  a few 10$^ 3$ yr. Assuming that 10\\% of the SN kinetic energy is converted to GCR, we find that during the acceleration  period a few particles out of a million encountered by the forward shock are accelerated.  Finally, we reassess the superbubble paradigm for the origin of GCR, by evaluating consistently  the \\neo \\ ratio with the aforementioned  stellar yields. We find that it can not be as high as observed, unless some extremely favorable assumptions are made (only the early period of the superbubble lifetime considered, no gas left over from the formation of the OB association). We conclude that superbubbles cannot be at the origin of the bulk of GCR.  The plan of the paper is as follows. In Sec. 2 we present the general \"set-up\" of our model: the adopted stellar models (Sec. 2.2) and  their wind yields (Sec. 2.3), the description of the circumstellar environment (Sec. 2.4) and the evolution of the forward shock in the ST phase (Sec. 2.5). In Sec. 3 we present our results for the (time-dependent) composition of the accelerated particles, the limits imposed on the shock velocity by the observed \\neo \\ ratio and the efficiency of the particle acceleration. Finally, in Sec. 4 we explore the \\neo \\  ratio of GCR,  assumed to be accelerated inside a superbubble, and we show that it cannot match the oberved one (unless extreme assumptions are made). The results are summarized in Sec. 5. ",
        "conclusions": "In this work we explore some implications of the idea that cosmic rays are accelerated by the forward shock in supernova remnants, during their ST phase. We focus on the chemical composition resulting from such an acceleration and, in particular, on the \\neo \\ ratio, which is the most characteristic feature of the  observed GCR source composition and is unaffected by atomic effects. For that purpose, we adopt recent models of the nucleosynthesis and evolution of massive stars with mass loss: those of LC06 with no rotation and those of HMM05 with rotation.  In Sec. 2 we present a detailed summary of the properties of those models and, in particular, of the chemical composition of their winds, insisting on the fact that rotating models release more \\neb \\ in their winds than non-rotating ones. We also present the adopted model for the evolution of a SN remnant within a stellar wind, based on the ideas of Ptuskin and Zirakashvili (2005) and the equations summarized in Caprioli (2011).  In the framework of the simple model adopted here, we follow the time-dependent composition of  GCR, accelerated by the forward shock as it runs through either the ISM (in the case of stars with M$<$25-35 \\ms, depending on rotation) or through the stellar wind (in the case of more massive stars). In fact, during the largest part of the ST phase, the shock runs through ISM and encounters a $\\sim$solar composition. In order to reproduce the observed high value of (\\neo)$_{CR}$    ($R_{Obs}$=5.3$\\pm$0.3 in solar units)  after accounting for the stellar IMF, we have then to assume that acceleration is efficient only during a short early period in the ST phase. We chose to use the shock velocity as a criterion for efficient acceleration and, based on the aforementioned CR composition argument, we find that shock velocities larger than $\\sim$1900 km/s (for the rotating stellar models) or 2400 km/s ( for the non-rotating ones) are required. This result is obtained by assuming a step function for the efficiency $f$ of particle acceleration ($f$=0 before the ST phase and after $\\upsilon_{min}$, and $f$=1 between the two). For the - perhaps, more realistic - assumption of a velocity-dependent efficiency  $f\\propto \\upsilon^2$, we find slightly lower values for $\\upsilon_{min}$ (1600 km/s for HMM05 yields and 2150 km/s for CL06 yields, respectively). In the framework of the adopted models, this corresponds to a circumstellar mass of  several tens of \\ms \\ encountered by the forward shock. Assuming, furthermore, that 10\\% of the SN kinetic energy is used in acceleration of escaping cosmic rays with standard energy spectra, allows us to evaluate the efficiency of that acceleration: we find that a few particles out of a million encountered by the forward shock  are accelerated to CR energies. We also notice that this scheme of GCR acceleration does not suffer from problems related to the absence of unstable $^{59}$Ni in observed GCR composition: this heavy nucleus is well inside the SN ejecta and is not reached by the forward shock which accelerates only wind material and ISM.  The aforementioned scenario assumes that even the most massive stars, up to 120 \\ms, develop strong forward shocks and accelerate the particles of their WR winds. For non-rotating stars, this is a rather extreme assumption, since it has been argued (Heger et al. 2003) that non-rotating masssive stars of about solar metallicity collapse into black holes, if their mass is in the 30-60 \\ms \\ range (see their Fig. 1). Notice that stars in the 60-120 \\ms \\ range (the most important \\neb \\  producers) end as  black holes in their scheme. However,  for slightly higher metallicities - such as those encountered in the inner Galactic disk - they find that  only neutron stars are formed, because higher stellar mass losses result in a less massive star at explosion. It should be noticed that the details of massive star explosions remain poorly understood at present (see e.g. Hanke et al. 2011) and so is the fate of a massive star above 30 \\ms \\ (see Fryer et al. 2011 for a recent - but certainly not  definitive - assessment). The fate of the rotating mass losing stars considered here is even less well known.  In view of the  aformentioned uncertainties, we feel that the scenario proposed here can be considered as valid at present, although future refinements in our understanding of massive star explosions may change it quantitatively (and even qualitatively, if it turns out that most masive stars above, say, 50 \\ms, end up as black holes).    Finally, we explore the idea that CR are accelerated in superbubbles, in which case their composition results from the ejecta of both stellar winds and SN explosions. We first notice that simple nucleosynthesis arguments suggest that the resulting composition, averaged over the stellar IMF, should be very close to solar. We  demonstrate this quantitatively, with a simple model for the evolving composition of a superbubble, enriched first by the (\\neb \\ rich) winds of the most massive stars, then by the (\\nea \\ rich) SN ejecta of less massive stars. We find that, after a few Myr the superbubble (\\neo)$_{SB}$  \\ ratio tends to solar, and so does the average (\\neo)$_{CR}$ \\ ratio in accelerated particles. We conclude that superbubbles cannot provide the observed high  $R_{Obs}$ value of CR sources and, therefore, are not the main site of CR acceleration. On the contrary, SN remnants - including those expanding in the pre-explosion environment of a stellar wind - appear as suitable sites of GCR acceleration."
    },
    "1112/1112.3878_arXiv.txt": {
        "abstract": "We assess the viability of successful reconstruction of the evolution of the dark energy equation of state using varying fundamental couplings, such as the fine structure constant or the proton-to-electron mass ratio. We show that the same evolution of the dark energy equation of state parameter with cosmic time may be associated with arbitrary variations of the fundamental couplings. Various examples of models with the same (different) background evolution and different (the same) time variation of fundamental couplings are studied in the letter. Although we demonstrate that, for a broad family of models, it is possible to redefine the scalar field in such a way that its dynamics is that of a standard quintessence scalar field, in general such redefinition leads to the breakdown of the linear relation between the scalar field and the variation of fundamental couplings. This implies that the assumption of a linear coupling is not sufficient to guarantee a successful reconstruction of the dark energy dynamics and consequently additional model dependent assumptions about the scalar field responsible for the dark energy need to be made. ",
        "introduction": "Introduction}  More than one decade ago type Ia supernovae observations suggested, for the first time, that the expansion of the universe is accelerating \\cite{Perlmutter:1998np,Riess:1998cb}. Since then, increasingly precise cosmological observations \\cite{Percival:2009xn,Komatsu:2010fb,Amanullah:2010vv} led to a well tested cosmological model presently dominated by an exotic dark energy form, violating the strong energy condition. In fact, if General Relativity is valid on large cosmological scales then dark energy \\cite{Ratra:1987rm,Frieman:1995pm,Caldwell:1997ii,Ferreira:1997hj,ArmendarizPicon:2000dh,ArmendarizPicon:2000ah} provides the only convincing explanation for the observed acceleration of the universe. Understanding the nature of dark energy is therefore one of the most important challenges of modern cosmology with one of the primary goals being determining whether its energy density is constant or slowing varying (see \\cite{Copeland:2006wr,Frieman:2008sn,Caldwell:2009ix,Silvestri:2009hh,Li:2011sd} for recent dark energy reviews). A fundamental problem associated to a cosmological constant is that its magnitude is constrained to be much smaller than particle physics predictions. On the other hand, it is also not clear if there is a deep physical reason which explains why it became the dominant component of the Universe just around the present day \\cite{Barreira:2011qi}.  An arguably better motivated alternative to the cosmological constant is the possibility that dark energy might be described by a dynamical scalar field. One important parameter characterizing dynamical dark energy is its equation of state, the ratio $w$ between the dark energy pressure and energy density. Constant $w$ models are unrealistic unless $w=-1$, which corresponds to the cosmological constant case \\cite{Avelino:2009ze,Avelino:2011ey}. Hence, a measure of $w \\neq -1$ at any redshift or redshift band should be indicative of dynamical dark energy. Considerable efforts are being put forward to constrain the dynamics of $w$ at low redshifts (see \\cite{EditorialTeam:2011mu} for expected future developments with the Euclid mission) using type Ia supernova, galaxy clustering or weak lensing. These are indirect probes which rely on the impact of dark energy on the overall dynamics of the universe. However, dark energy is expected to become subdominant at early times and, consequently, it is not possible to strongly constrain its dynamics at high redshift using standard methods.  Still, in realistic models, dynamical scalar fields may couple to other fields, possibly leading to measurable variations of nature's fundamental \"constants\" \\cite{Carroll:1998zi}. The coupling between a quintessence field and fundamental couplings such as $\\alpha$ or $\\mu$ has been investigated by several authors \\cite{Carroll:1998zi,Chiba:2001er,Wetterich:2002ic,Nunes:2003ff,Anchordoqui:2003ij,Copeland:2003cv,Avelino:2004hu,Parkinson,Doran:2004ek,Marra:2005yt,Avelino:2006gc,Avelino:2008dc,Avelino:2009fd,Dent:2009,Calabrese:2011nf,Amendola:2011qp}. The dynamics of $\\alpha$ over the redshift range $z=0-10^{10}$  is severely constrained using using both cosmological and laboratory experiments (see \\cite{Uzan:2010pm} for a recent review). At low redshifts laboratory experiments \\cite{Peik:2006xy,Rosenband:2008} and the Oklo natural nuclear reactor \\cite{Gould:2006,Petrov:2006,Onegin:2010kq}  provide very stringent limits on the time-variation of $\\alpha$ and $\\mu$, while at high redshift cosmic microwave background temperature and polarization anisotropies \\cite{Avelino:2000ea,Avelino:2001nr,Martins:2003pe,Rocha:2003gc, Stefanescu:2007aa,Nakashima:2008cb} and light element abundances \\cite{Bergstrom:1999wm,Avelino:2001nr,Nollett:2002da} constrain the value of $\\alpha$ at $z \\sim 10^{10}$ and $z \\sim 10^3$ to be within a few percent of its present day value. Despite a few positive claims for a detection of a variation of the fine-structure constant $\\alpha$ \\cite{Webb:1998cq,Murphy:2006vs} or the proton-to-electron mass ratio $\\mu$ \\cite{Ivanchik:2005ws,Reinhold:2006zn} in the redshift range $z=1-4$, and the more recent claims for a significant spatial variation of $\\alpha$ \\cite{Webb:2011,King:2012}, there is presently no unambiguous  evidence for such variation (see, for example, \\cite{Chand2004,Chand2007,Murphy2008} for some strong negative results). Nevertheless, it has been shown that varying couplings may be used to determine the evolution of the dark energy equation of state over a larger redshift range than standard methods, if a number of conditions are verified \\cite{Avelino:2006gc,Avelino:2009fd}. These are: i) that the dark energy can be described by a standard quintessence field; ii) that the relation between the quintessence field and varying fundamental couplings  is linear; iii) that such variations are within reach of forthcoming experiments. In this letter we shall relax assumptions i) and ii) and consider more general k-essence models for dark energy, thus testing the robustness of the varying fundamental couplings method for the reconstruction of the evolution of the dark energy equation of state.  Throughout this letter we shall use units with $c=8\\pi G/3=H_0=1$ and a metric signature $(+,-,-,-)$. ",
        "conclusions": ""
    },
    "1112/1112.4175_arXiv.txt": {
        "abstract": "{We introduce a new technique for imaging the polarized radio sky using interferometric data. The new approach, which we call Faraday synthesis, combines aperture and rotation measure synthesis imaging and deconvolution into a single algorithm. This has several inherent advantages over the traditional two-step technique, including improved sky plane resolution, fidelity, and dynamic range. In addition, the direct visibility- to Faraday-space imaging approach is a more sound foundation on which to build more sophisticated deconvolution or inference algorithms. For testing purposes, we have implemented a basic Faraday synthesis imaging software package including a three-dimensional CLEAN deconvolution algorithm. We compare the results of this new technique to those of the traditional approach using mock data. We find many artifacts in the images made using the traditional approach that are not present in the Faraday synthesis results. In all, we achieve a higher spatial resolution, an improvement in dynamic range of about 20\\%, and a more accurate reconstruction of low signal to noise source fluxes when using the Faraday synthesis technique.} ",
        "introduction": "Rotation measure (RM) synthesis, introduced by \\citet{brentjens_faraday_2005}, is a technique that makes use of the Faraday effect to improve the sensitivity of polarimetric observations by combining data over wide ranges in frequency. RM synthesis allows for the separation of polarized sources along the line of sight (LOS) by decomposing the observed polarized emission into parts originating from different Faraday depths (in the simplest case, Faraday depth is equivalent to RM), allowing one to generate a 3D representation of the polarized sky. While the Faraday depth axis cannot be mapped to a physical depth, the Faraday depth information can be of significant scientific value since the Faraday depth traces the projected strength and orientation of magnetic fields along the LOS.  A more detailed introduction to RM synthesis is provided later.  The RM synthesis technique has only recently become viable, owing to the availability of broadband receivers in the next generation of radio telescopes such as the Expanded Very Large Array (EVLA), the upgraded Westerbork Synthesis Radio Telescope (WSRT), and the pathfinder projects leading to the Square Kilometer Array (SKA) such as the Low-Frequency Array (LOFAR). Recently, there have been many successful applications of RM synthesis, and interest is rapidly increasing as new radio telescopes are being commissioned. Applications have included studies of the diffuse polarized emission in the Perseus field \\citep{de_bruyn_diffuse_2005,brentjens_wide_2010} and the Abell 2255 field \\citep{pizzo_deep_2010}, analysis of the polarized emission in nearby galaxies in the WSRT-SINGS survey \\citep{heald_westerbork_2009}, and the detection of a shell of compressed magnetic fields surrounding a local HI bubble \\citep{wolleben_antisymmetry_2010}. The RM synthesis technique will play a critical role in several upcoming polarization surveys, e.g. POSSUM \\citep{gaensler_possum_2010}, GMIMS \\citep{wolleben_gmims_2009}, and future surveys with LOFAR. RM synthesis is very useful for studying magnetism.  For instance, \\citet{bell_caustics_2011} have shown that prominent asymmetric features in RM synthesis images known as Faraday caustics, which are related to LOS magnetic field reversals, can be used to study the structural and statistical properties of magnetic fields.  In addition to the considerable interest in applications, there has been recent interest in improving RM synthesis imaging techniques. The RMCLEAN deconvolution algorithm was introduced by \\citet{heald_westerbork_2009} and was quickly adopted, owing to its simplicity and similarity to techniques used in aperture synthesis imaging. \\citet{frick_wavelet-based_2010} have proposed a wavelet-based RM synthesis technique. With this approach, one obtains a decomposition of the size scale of structures in addition to their Faraday depth location. There has also been growing interest in applying compressed sensing to RM synthesis \\citep{li_compressed-sensing-rmsynth_2011,andrecut_sparse-rmsynth_2011}. Compressed sensing is a method of reconstructing signals that are sparse in some set of basis functions. If the signal is sufficiently sparse, it can be reconstructed using fewer measurements than indicated by the Nyquist-Shannon sampling theorem. Since compressed sensing techniques make use of wavelet bases, but are implemented in a noise-aware fashion, they can be expected to be superior to pure wavelet based methods.  These new imaging techniques have thus far focused on the problem of one-dimensional (1D) reconstruction. However, the product of RM synthesis is a function not only of Faraday depth, but also of position on the sky, making its reconstruction an inherently three-dimensional (3D) problem. In many cases, RM synthesis is performed on sky brightness images that have been produced from radio interferometric data. Observations performed with a radio interferometer sample the aperture plane rather than the image plane, and this is done at many different frequencies. This data space is the 3D Fourier space representation of the polarized sky brightness as a function of Faraday depth. Imaging algorithms should ideally make use of the entire data space to inform the reconstruction at each pixel, since the information about the sky brightness at each pixel is spread throughout Fourier space. However, the current approach is to perform the imaging in a piecewise fashion, first reconstructing the 2D sky plane images at each frequency before doing 1D RM synthesis imaging along each LOS. Therefore, each step of the traditional imaging approach is done with a limited subset of the data, which will reduce the overall sensitivity and degrade fidelity in the final image.  In this paper, we introduce a new technique for imaging the Faraday spectrum directly from radio interferometric data. We call this technique \\emph{Faraday synthesis}. Using this approach, one images the polarized emission as a function of sky position and Faraday depth from the visibility data itself, rather than using the traditional piecewise prescription. Faraday synthesis is a natural extension of the aperture synthesis plus RM synthesis techniques that provides improvements in image fidelity and sensitivity.  We note that a similar technique was briefly discussed in \\citet{pen_GMRT-eor_2009}, although it was not considered in any detail, nor was it compared to the traditional approach to RM synthesis imaging. Furthermore, deconvolution was not considered.  With the advent of RM synthesis the concept of rotation measure, defined to be the amount that the observed polarization angle changes as a function of frequency, has become somewhat outdated. With RM synthesis, one does not measure RMs, but instead reconstructs the polarized intensity as a function of Faraday depth. In the simplest case, where a single, discrete source of polarized emission is positioned behind a Faraday rotating medium, the RM is equal to the Faraday depth. In all other cases this is not true. In general, RM cannot be used as a proxy for Faraday depth, and the full distribution of polarized brightness as a function of Faraday depth is the most appropriate quantity to study. Therefore, we avoid use of the term RM to describe this new method, and instead call it Faraday synthesis. Throughout the remainder of this paper, RM synthesis will refer to the LOS imaging method developed by \\citet{brentjens_faraday_2005}. The traditional practical approach of first imaging individual frequencies using 2D aperture synthesis techniques and then reconstructing the LOS brightness distribution on a pixel-by-pixel basis will be referred to as 2+1D imaging, in contrast to Faraday synthesis, which we will often refer to as 3D imaging.  In Sec. \\ref{sec:Synthesis-imaging} we briefly review the theories of aperture and RM synthesis imaging. In Sec. \\ref{sec:Faraday-synthesis} we introduce the Faraday synthesis imaging technique. In Sec. \\ref{sec:Proof-of-concept} we describe the proof of concept software that we have implemented, and in Sec. \\ref{sec:Tests} we compare test results obtained by imaging mock data using both the 3D and 2+1D techniques. We conclude in Sec. \\ref{sec:Discussions-and-conclusions} with a summary and discussion of our results. ",
        "conclusions": "}  We have described a new approach to imaging linearly polarized visibility data that we call Faraday synthesis. With this approach, one directly reconstructs the Faraday spectrum, or the polarized intensity as a function of Faraday depth, from the visibility data. This takes the place of the usual approach of first performing aperture synthesis imaging on the visibility data at each frequency, then performing RM synthesis along each line of sight independently.  These two approaches would be equivalent if deconvolution were not required. With Faraday synthesis, the deconvolution is done in one step using the entirety of the broad-band data. In contrast, the traditional approach requires deconvolving the images individually at each frequency. The sensitivity in the narrow-band images is significantly limited, and residual artifacts remain in these images that limit the dynamic range and image fidelity achieved during RM synthesis. Moreover, another deconvolution procedure is necessary to account for the point spread function due to the incomplete sampling of the wavelength space. Artifacts introduced by the first deconvolution algorithm will be compounded during this procedure, further reducing image fidelity.  Indeed, we found in our proof-of-concept testing that artifacts were significantly higher in the traditional imaging method than with Faraday synthesis. The noise was roughly 20\\% lower when using the Faraday synthesis technique, and the strongest artifact was about half as bright.  We found that one is able to achieve a better resolution in the final image using the new approach. The main lobe of the 3D dirty beam was 30\\% smaller in the sky plane than that of the traditional method. With the 2+1D method, stacking of the images at each frequency requires tapering the visibility data such that the higher frequency images have the same resolution as those at the lower frequencies. We find that the inaccuracies inherent in the stacking process are a significant source of artifacts in the traditional RM synthesis technique.  Furthermore, this procedure requires severely down-weighting a large portion of the available data, again limiting sensitivity. With the Faraday synthesis approach, no such down-weighting is required.  The Faraday synthesis approach of working directly between visibility and Faraday space is a much better foundation on which to build new image reconstruction algorithms because with it one is able to accurately describe the instrument response. The effects of the intermediate, non-linear deconvolution procedure can not be easily understood or modeled. Many signal inference techniques, like those derived using Information Field Theory \\citep{ensslin_ift_2009,ensslin_gibbs_2010}, depend on a complete description of the instrument response and would need to be built on the framework of Faraday synthesis.  While the CLEAN algorithm has worked quite well in radio astronomy for decades, the implicit assumption of a sky sparsely populated by point sources is not well suited for the kinds of diffuse polarized signals that one finds, for example, in the polarized emission from the Milky Way. Novel image reconstruction algorithms built using more appropriate constriants or assumptions are likely justified. Such algorithms could make use of statistical correlations inferred from the data, similar to the extended critical filter algorithm developed by \\citet{oppermann_extcritfilt_2011}."
    },
    "1112/1112.1420_arXiv.txt": {
        "abstract": "We carry out a comprehensive smooth particle hydrodynamics simulation survey of double-degenerate white dwarf binary mergers of varying mass combinations in order to establish correspondence between initial conditions and remnant configurations. We find that all but one of our simulation remnants share general properties such as a cold, degenerate core surrounded by a hot disk, while our least massive pair of stars forms only a hot disk. We characterize our remnant configurations by the core mass, the rotational velocity of the core, and the half-mass radius of the disk. We also find that some of our simulations with very massive constituent stars exhibit helium detonations on the surface of the primary star before complete disruption of the secondary. However, these helium detonations are insufficiently energetic to ignite carbon, and so do not lead to prompt carbon detonations. ",
        "introduction": "Type Ia supernovae are commonly accepted to be the observed transient produced after a thermonuclear detonation inside a white dwarf star. While the preferred mechanism for producing SNeIa involves accretion from an evolved main sequence star onto a white dwarf (Whelan \\& Iben 1973; Nomoto 1982; Hillebrandt \\& Niemeyer 2000), the observed SNeIa rate is incompatible with the narrow range of helium accretion rates that initiate a carbon detonation as opposed to accretion induced collapse or classical novae (Nomoto \\& Kondo 1991; Hardin \\etal 2000; Pain \\etal 2002; Ruiter \\etal 2009). Moreover, many recent observations of abnormally luminous SNeIa have been interpreted as having derived from double-degenerate systems involving two white dwarfs.  For example, photometric observations of SN~2007if suggest 1.6$\\pm$0.1\\msol of \\nickel[56] was formed, implying a progenitor mass of 2.4$\\pm$0.2\\msol (Scalzo \\etal 2010), which is well above the Chandrasekhar limit (Chandrasekhar 1931). Spectroscopic observations of SN~2009dc suggest $\\apprge1.2$\\msol of \\nickel[56] (Tanaka \\etal 2010), depending on the assumed dust absorption. Since 0.92\\msol of \\nickel[56] is the greatest yield a Chandrasekhar mass can produce (Khokhlov \\etal 1993), this yield also implies a super-Chandrasekhar progenitor mass. And observations of SN~2003fg by Howell \\etal (2006) and of SN~2006gz by Hicken \\etal (2007) infer $\\sim1.3$\\msol of \\nickel[56] each.  Generally, for the purposes of cosmological measurements, obvious double-degenerate candidates are excluded from SNeIa surveys. The Phillips relation, or the width-luminosity relation (WLR), which established SNeIa as standard candles (Phillips 1993), relates the peak luminosity of a SNIa to the change in magnitude after 15 days. The WLR for standard SNeIa indicates that SNeIa with bright peak magnitudes also decay at a slower rate than dimmer SNeIa. This is often thought to be the result of a relationship between the \\nickel[56] yield and the opacity of the ejecta material, assuming a total mass not exceeding the Chandrasekhar mass. However, a double-degenerate system may have up to two times the Chandrasekhar mass, and so the relationship between the \\nickel[56] yield and the ejecta opacity need not be similar to single-degenerate scenarios, and the WLR may not be applicable to these SNeIa. In fact, it is more likely that for a given \\nickel[56] production and energy deposition history, an increased ejecta mass results in an increased opacity, reducing the peak magnitude and broadening the lightcurve in a fashion that is the inverse of the standard WLR (Pinto \\& Eastman 2000; Mazzali \\etal 2001; Mazzali \\& Podsiadlowski 2006; Kasen, R\\\"opke, \\& Woosley 2009).  Complicating matters is the possibility that SNeIa deriving from double-degenerate progenitors can have ordinary \\nickel[56] yields, depending on the progenitor mechanism (see \\eg collisional mechanisms in Raskin \\etal 2009, 2010 and Rosswog \\etal 2009) and the final, central densities of the degenerate material before ignition. Thus, double-degenerate SNeIa may be masquerading as typical SNeIa. If they do not conform to the WLR, they may introduce systematic errors into cosmological surveys. In order to reduce the scatter in the Hubble diagram, we must first establish whether double-degenerate SNeIa are standardizable, and if not, we must identify the tell-tale signatures of a double-degenerate progenitor mechanism.    The most probable double-degenerate progenitor scenario involves two white dwarfs in a tight binary, though other progenitor systems have been considered (Benz \\etal 1989a; Raskin \\etal 2009; Rosswog \\etal 2009). Binary white dwarf systems were first seriously explored as plausible SNeIa progenitors by Iben \\& Tutukov (1984) and Webbink (1984). In such a system, tidal dissipation and gravitational radiation drive the binary pair into an ever closer orbit. Eventually, the least massive white dwarf, being physically larger as $R\\propto M^{-1/3}$, overflows its Roche lobe and begins to accrete material onto the primary, or more massive companion star. For many mass combinations, this is a fundamentally unstable process in which the loss of mass from the secondary causes it to outgrow its Roche lobe faster than its orbit widens due to conservation of angular momentum (see \\eg Marsh \\etal 2004).  Benz \\etal (1990) performed one of the first simulations of double-degenerate mergers, examining a binary system consisting of a 1.2\\msol white dwarf primary and a 0.9\\msol white dwarf secondary. They used a smooth particle hydrodynamics code with 3000 particles per star and found that the merger remnant consisted of a 1.7\\msol core surrounded by a rotationally supported disk. This is more massive than the Chandrasekhar limit, but they concluded that the central object was not entirely degenerate, having been considerably heated, and thus, much of the object's support against gravitational collapse came simply from thermal pressure.  Since this pioneering work, others have revisited the white dwarf merger scenario with up-to-date simulation codes and higher resolutions than were possible in the past (see \\eg Yoon \\etal 2007; Lor\\'{e}n-Aguilar \\etal 2009). In this paper, we revisit white dwarf mergers simulations, examining a wide range of possible mass combinations with high resolution, more accurate initial conditions, and up-to-date physics for the equation of state and nuclear reaction network. The structure of this paper is as follows. In \\S2, we outline our methods and initial conditions. We discuss the results of our simulations in \\S3, and compare our work to previous studies of white dwarf mergers in \\S4. Finally, in \\S5, we summarize our results and conclusions and discuss possible avenues for advancing remnant evolution in future studies. ",
        "conclusions": "Double degenerate progenitor scenarios are drawing new interest among the supernova community. They have the potential to explain some of the confounding mysteries that remain about SNeIa, and to enhance their usefulness as cosmological probes. Here, we have conducted a large survey of white dwarf binary merger models in order to begin to understand how these systems evolve, and to piece together some of the observational signatures of double degenerate SNeIa. %  In each of our simulations, the merger remnant consisted of a cold, degenerate core surrounded by a hot accretion disk, with $h/r_d\\approx0.1$. The only exception to this remnant configuration was the simulation of two 0.64\\msol white dwarfs, where the cores of both stars merged at the center of mass, heating the remnant core considerably and lifting most of its degeneracy. Since it is unlikely that two identical mass white dwarfs would form in a binary in nature, this scenario might seem of trivial importance, but the merging of the cores was less the result of their masses being identical than it was of the stars being highly susceptible to tidal disruption. Our other equal mass simulations with more massive constituent stars featured merger scenarios wherein one of the stars was completely disrupted into an accretion disk around the core of the other, provided there was a non-trivial asymmetry between the two. Therefore, it is likely that merger scenarios involving slightly unequal but low-mass white dwarfs might also exhibit core merging, though more simulations are certainly needed to confirm this conjecture.  Evolving these meta-stable remnants further will require the implementation of new physics beyond hydrodynamics and nuclear burning. Since the disks are optically thick, radiative losses from the surface are the dominant evolutionary mechanism, and efforts are underway to incorporate this process. As the surface layers lose energy through radiation, the material will sink to lower latitudes, driving convective currents that will increase the disk viscosity. We expect this viscosity will mimic an $\\alpha$-disk prescription, and thus, the accretion times will be of order 100 orbital times. However, the three most massive combinations simulated have disk configurations that suggest the cores might grow to become super-Chandrasekhar masses on shorter timescales, after cooling has removed some of the thermal support for the inner portions of the disks.  While none of our simulations exhibited prompt carbon detonations, all of the simulations that included a 1.06\\msol primary did exhibit prompt helium detonations on the surface of the primary. These helium detonations were not sufficiently energetic to significantly burn much of the carbon or to unbind either of the constituent stars, but the detonation shocks they produced did alter the structure of the primary before complete merger. Moreover, at $\\sim10^{49}$ ergs, the energy these detonations released is likely sufficient to be observable by many of the upcoming transient surveys, such as LSST.  Whether white dwarf mergers produce SNeIa is still an open question. Our simulations and others' represent only the beginnings of our exploration of this progenitor mechanism. Preliminary results look very promising, and it is doubtless that the viability of mergers as SNeIa progenitors will be established in the near term. Meanwhile, it remains an exciting time to be exploring these dynamical scenarios that continue to surprise us."
    },
    "1112/1112.5035_arXiv.txt": {
        "abstract": "{The largest uncertainty for cosmological studies using clusters of galaxies is introduced by our limited knowledge of the statistics of galaxy cluster structure, and of the scaling relations between observables and cluster mass. A large effort is therefore undertaken to compile global galaxy cluster properties in particular obtained through X-ray observations and to study their scaling relations. However, the scaling schemes used in the literature differ.} {The present paper aims to clarify this situation by providing a thorough review of the scaling laws within the standard model of large-scale structure growth and to discus various steps of practical approximations.} {We derive the scaling laws for X-ray observables and cluster mass within the pure gravitational structure growth scenario. Using N-body simulations we test the recent formation approximation used in earlier analytic approaches which involves a redshift dependent overdensity parameter. We find this approximation less precise than the use of a fiducial radius based on a fixed overdensity with respect to critical density.} {Inspired by the comparison of the predicted scaling relations with observations we propose a first order modification of the scaling scheme to include the observed effects of hydrodynamics in structure formation. This modification involves a cluster mass dependent gas mass fraction. We also discuss the observational results of the reshift evolution of the most important scaling relations and find that also a redshift dependence of the gas mass to total mass relation has to be invoked within our modification scheme.} {We find that the current observational data are within their uncertainties consistent with the proposed modified scaling laws.} ",
        "introduction": "Galaxy clusters form from overdense regions in the large-scale matter distribution, which have small amplitudes at early epochs and only collapsed to objects very recently. In the standard cosmological model the large-scale matter distribution is described by a random Gaussian field characterized by a power spectrum with a smoothly changing power law index over relevant length scales. This implies that the structure evolution will feature a large degree of self-similarity in scale and time (e.g. Peebles 1980). Consequently galaxy clusters, which form an integral part of this large-scale structure, also show an imprint of this general self-similarity. This connection between the framework of the evolution of the gravitating matter on large scales and galaxy cluster formation and their observed structure was realized in early studies, e.g. by Gunn \\& Gott (1972), Fillmore \\& Goldreich (1984), Bertschinger (1985), Hoffman \\& Shaham (1985) and in simulations e.g. Frenk et al. (1985), Zureck, Quinn \\& Solomon 1988, Efstathiou et al. 1988, West, Dekel \\& Oemler 1987). It resulted in a comprehensive description of the structure of dark matter halos, of which galaxy clusters are the most massive representatives, in a series of papers by Navarro, Frenk \\& White (1995, 1996, 1997) and follow-up literature. In this picture of purely dark matter structure growth, dark matter halos form a nearly self-similar, two-parameter family, with the two parameters being mass and a concentration or time-of-formation parameter. This structural model of clusters describes an average behavior of the cluster population, where the different statistical realizations of mass distributions in the protoclusters produce a significant scatter in the observed structural parameters around this mean. Deviations from the equilibrium state after merger events further contribute to this scatter.  The addition of baryons to this model leads to a modification of this picture, which for clusters can be seen as a perturbation of the dark matter structure evolution. In this sense galaxy clusters mark the very interesting transition region, where a first order description involving only the large-scale structure gravitational physics provides a very effective guideline and the more complicated hydrodynamics, including radiative cooling and feedback from star formation as well as AGN activity, constitutes a perturbative refinement. At smaller scales, for galaxies the gaseous astrophysics acting on small scales becomes dominant for the appearance of the visible objects and the observed evolution of the large-scale structure on galaxy scales becomes very non-linear. It is therefore on galaxy cluster scales where we can still very successfully apply analytical descriptions as a useful guideline for the understanding of structure evolution. The paper builds on this property of galaxy clusters.  X-ray observations are currently providing the most detailed account of galaxy cluster structure and are consequently used extensively to test the predictions of the large-scale structure growth models. However, they do not directly provide a picture of the dark matter halo distribution, but the distribution of the hot intracluster medium (ICM) that fills the entire cluster volume and radiates in X-rays. Therefore the X-ray appearance of clusters includes aspects of the hydrodynamics how the gas reacts to the dark matter density distribution and how the ICM evolves in its thermodynamic properties (e.g. Voit 2005). One can use it as a tracer of the dark matter distribution, for example through the assumption that it is located in the dark matter potentials in hydrostatic equilibrium. We can thus expect that the X-ray appearance of galaxy clusters is featuring some hydrodynamic modification compared to the more readily described dark matter distribution.  Because the galaxy cluster formation is so tightly connected to the large-scale structure evolution and the fact that there is a well described self-similar structure statistics in the purely gravitational cluster formation model (Navarro et al. 1997) we can expect that there are simple analytically derivable scaling relations for the global X-ray observables as a function of cluster mass. These relation have been studied already early by e.g. Kaiser (1986), and Evrard \\& Henry (1991) and this theoretical work has been supported by simulations (e.g. Bryan \\& Norman 1998, Borgani 2004, Kravtsov et al. 2006, Evrard et al. 2008, Stanek et al. 2010, Short et al. 2010, Borgani \\& Kravtsov 2010).  With the event of detailed observational studies of cluster structure in X-rays by means of the advanced X-ray observatories {\\sl Chandra} and {\\sl XMM-Newton}, large sets of observational data on cluster structure and scaling relations have become available now and the detailed testing of the theoretical predictions for the scaling laws is in full swing (e.g. Markevitch et al. 1998, Arnaud \\& Evrard 1999, Mohr \\& Evrard 1997, Finoguenov et al. 2001, Ikebe et al. 2002, Reiprich \\& B\\\"ohringer 2002, Ponman et al. 2003, Ettori et al. 2004, Vikhlinin et al. 2005, Pointecouteau et al. 2005, Arnaud et al. 2005, Pratt et al. 2006, Kotov \\& Vikhlinin 2006, Zhang et al. 2006, Maughan et al. 2006, Maughan 2007, Arnaud et al. 2007, Pratt et al. 2009, Mantz et al. 2010, Arnaud et al. 2010, Sun et al. 2011, Reichert et al. 2011). An investigation of the relevant literature shows, however, that a number different methods are used for the scaling of the data at different redshifts. The aim of this paper is therefore to critically review these methods and to determine the best approach based on comparison with simulations and observations.  In the above mentioned literature mostly analytical formulations of the scaling relations have been used, based on general considerations of structure formation. In order to provide the ground for higher precision in the analysis of the evolution of cluster structure, the simplifications made in the analytical models should be replaced by tests and calibrations with N-body simulations. This situation can be compared to that of the theoretical prediction of the dark matter halo (galaxy cluster) mass function, where the analytical model by Press and Schechter (1974) has paved the way for the general formulation of the solution for the mass function, but the actual formulae now applied, are the result of careful calibration with N-body simulations (e.g. Jenkins et al. 2001, Evrard et al. 2002, Warren et al. 2006, Tinker et al. 2008). Here we adopt a similar approach. We first present the theoretical framework for the description of the evolution of the scaling relations in the classical form based on the assumption of the recent formation approximation and compare it to an alternative scheme used in the literature. We then resort to the results of N-body simulations to test the predictions of the analytical approaches and discuss which of the presently used methods in the literature is best.  In the second part of the paper we compare the theoretical scaling relations to observations, inspect the deviations of the observed scaling relations from the predictions based on dark matter structure evolution (often called ''gravitational scaling relations''), and discuss these deviations in the context of the influence of hydrodynamical processes. We then explore a simply empirical modification scheme of the scaling relation using a mass dependent depletion factor for the ICM gas to account for the hydrodynamical scaling effects and compare the so obtained set of scaling relations to observations. In a last step we consider how this modification should depend on redshift to be consistent with the observational data.  The paper is structured as follows. In section 2 we derive the ''gravitational scaling relations'' based on dark matter structure evolution with the assumption that the baryonic matter follows the dark matter. In section 3 we investigate the dependence of these scaling relations on the redshift dependent overdensity parameter in a $\\Lambda$CDM cosmology and test this model against the method using a fixed overdensity parameter in section 4 using numerical simulations. In section 5 we discuss the redshift dependence of the overdensity parameter, which defines the proper fiducial radius of the clusters, in the context of numerical studies of the redshift dependence of the concentration parameter of galaxy clusters. Section 6 then starts the second part of the paper where we discuss the modification of the scaling relations to include hydrodynamical effects. After a comparison of the description of the evolution of the scaling relations in terms of the parameters E(z) and (1+z) in section 7 and a brief comparison with some recent simulations in section 8, we provide a comprehensive comparison of scaling relation results in the literature with the model predictions in section 9. Finally section 10 contains a discussion and conclusions. ",
        "conclusions": "Studying in detail the evolution of the dark matter mass density profiles of simulated galaxy clusters, we have shown that the older model of self-similar scaling relations based on the recent formation approximation which uses a scenario where the fiducial overdensity radius, $r_{\\Delta(z)}$, is taken to be redshift dependent, is not accurate and a scaling with a fixed overdensity provides a better and currently sufficiently precise description of self-similar evolution. For a precise analysis of future cosmological surveys of the galaxy cluster population, we should improve this description further. We plan to do this with larger N-body/hydrodynamical simulations which are performed at present and therefore give no detailed recipes for the scaling corrections in this paper. The corrections as shown in Fig. 5 depend on the cosmological model used and therefore these corrections will have to be worked separately for each model case studied.  Studying the different types of scaling relations involving parameters derived from X-rays, we can distinguish two types of scaling behavior: the mass - temperature relation is mostly dependent on the scaling of the dark matter potentials and is therefore very close to the gravitational scaling prediction. Most other relations involving gas density or gas mass are affected by the non-constancy of the gas mass to total mass ratio. With the introduction of modified scaling relations to take this hydrodynamical effects into account, we can describe the currently available data sets within the given uncertainties.  Looking at the evolution of the scaling relations with redshift, we find an analogous situation: the $M-T$ relation corresponds within the current uncertainties to the prediction of the gravitational self-similar scenario. All other relations involving parameters which depend on the gas density show deviations, which implies that the gas mass fraction is not constant for given cluster mass with redshift. This is just the consequence of the following effects. At higher redshifts clusters of given mass are more compact and the ICM has to be squeezed into a deeper and narrower potential. Since in preheating models, which seem to explain the data best, the gas starts out with an elevated entropy before cluster formation is complete, the gas is less tightly squeezed into the earlier, narrower potentials than into the later wider potentials.  Our modified scaling relation model does not describe all observational effects. As shown in Pratt et al. (2010), the entropy scaling depends on the radius at which the entropy is measured. These results imply that the ICM depletion is larger in the center of groups and clusters than in the outer parts. More data are required that extend out to large cluster radii (to $r_{500}$ and beyond) to substantiate this result.  The redshift evolution of the $L-T$ scaling relation, which is predicted to be $\\propto E(z)$ in the simple self-similar scenario, is now found in recent studies to be much less positive or even negative (e.g. O'Hara et al. 2007, Reichert et al. 2011). This has the important consequence that one will find less high redshift galaxy clusters in future X-ray and SZE surveys, than predicted based on the simple scaling models (Reichert et al. 2011). More importantly, a precise measurement of this evolution effect is crucial for using the future X-ray survey data on galaxy clusters for the test of cosmological models.  As shown by the comparison of the evolution of the scaling relations compiled by Reichert et al. (2011) and the simulations by Short et al. (2010), the study of the evolution of the scaling relations also provides important insight into the astrophysics of the ICM. Currently the observational data strongly favour a model with ealry preheating of the ICM.  In the coming years both the observational data as well as the simulation results will experience further strong improvements. Therefore we see the importance of this paper more in elucidating the way how the scaling relations should be analysed and applied, rather than already providing the best parameterization of the results."
    },
    "1112/1112.0756_arXiv.txt": {
        "abstract": "Following the discovery of a new radio component right before the GeV $\\gamma$-ray detection since 2008 August by {\\it Fermi} Gamma-ray Space Telescope, we present a detailed study of the kinematics and lightcurve on the central sub-pc scale of 3C~84 using the archival VLBA 43-GHz data covering the period between 2002 January to 2008 November. We find that the new component ``C3'', previously reported by the observations with the VLBI Exploration of Radio Astrometry (VERA), was already formed in 2003. The flux density of C3 increases moderately until 2008, and then it becomes brighter rapidly after 2008. The radio core, C1, also shows a similar trend. The apparent speed of C3 with reference to the core C1 shows moderate acceleration from $0.10c$ to $0.47c$ between 2003 November to 2008 November, but is still sub-relativistic. We further try to fit the observed broadband spectrum by the one-zone synchrotron self-Compton (SSC) model using the measured apparent speed of C3. The fit can reproduce the observed $\\gamma$-ray emission, but does not agree with the observed radio spectral index between 22 and 43~GHz. ",
        "introduction": "\\label{sec:intro}  The radio source 3C~84 is associated with the giant elliptical galaxy NGC~1275 ($z=0.0176$). It is well known that 3C~84 has a pair of compact radio lobe structures on the central 10-pc scale \\citep{1994ApJ...430L..41V, 1994ApJ...430L..45W, 2000ApJ...530..233W, 2006PASJ...58..261A}. Because of its proximity, it allows us to study the region within the central sub-pc scale using a Very Long Baseline Interferometer (VLBI). It is an ideal laboratory to investigate the formation mechanism of relativistic jet ultimately powered by super-massive black holes and the interaction between the jets and ambient matter in the galactic central regions.  Thanks to the recent observations with {\\it Fermi}/LAT, we have a new opportunity to explore GeV $\\gamma$-ray production mechanism in misaligned radio-loud AGNs \\citep{2010ApJ...720..912A}. With CGRO/EGRET, a few extragalactic radio galaxies, such as Centaurus A \\citep{1999APh....11..221S}, 3C~111 \\citep{2008ApJ...688..852H}, and NGC~6251 \\citep{2002ApJ...574..693M} had been already detected. However, the detection of NGC~1275 \\citep{2009ApJ...699...31A} by {\\it Fermi}/LAT in 2008 August is particularly noteworthy because NGC~1275 was not detected by CGRO/EGRET \\citep{2003ApJ...588..155R}. The flux density detected by {\\it Fermi}/LAT is about 7 times higher than the upper limit of EGRET sensitivity. Intriguingly, the radio monitoring also shows the flux increase starting in 2005. The time variation of $\\gamma$-ray flux density shows a similar trend with the radio flux density on the timescale of decades, implying the possible connection between the $\\gamma$-ray emission and variable radio component.  Overall spectral energy distribution (SED) of NGC~1275 from radio to $\\gamma$-ray can be explained by the synchrotron-self Compton (SSC) model and the deceleration jet model \\citep{2009ApJ...699...31A} adopting sub-pc for the size of emitting region. The authors derived Lorentz factors of the emitting region  $\\Gamma=1.8$ ($\\theta=25$ deg, $\\delta=2.3$) for the case of SSC model and $\\Gamma$ varying from $10$ to $2$ ($\\theta=20$ deg, $\\delta$ varying 1.6 to 2.7) for the decelerating model, where $\\theta$ and $\\delta$ are the jet angle to the line of sight and the beaming factor of the emitting region, respectively. Here we should stress that the size of the $\\gamma$-ray emitting region adopted in the SED models is comparable to a VLBI component size. Therefore, it is essential to test the scenario of co-spatiality of GeV $\\gamma$-ray and radio emitting regions by VLBI observations.  In order to find the radio counterpart of the GeV $\\gamma$-ray emitting region, we have conducted VLBI observations using the VLBI Exploration of Radio Astrometry (VERA) at 22~GHz between 2006 June 14 and 2009 April 24 (\\citet{2010PASJ...62L..11N}, hereafter Paper I). Surprisingly, we found that the monotonic increase of radio flux density mainly originated in the newly born bright component C3  (Paper I). The measured projected speed of C3, however, was $(0.23 \\pm 0.01)c$ on average between 2007 October 24 and 2009 April 24 in the central sub-pc to pc scale. Given previously estimated jet angle to the line of sight ($11^{\\circ}$ - $55^{\\circ}$: \\citet{2006PASJ...58..261A, 2009AJ....138.1874L, 2009ApJ...699...31A}), the de-projected speed corresponds to 0.24$c$ - 0.55$c$, i.e. slower than the speed of jet derived from the SED modeling \\citep{2009ApJ...699...31A}. The result of VERA observation implies $\\delta < 2$, and there seems to be a discrepancy between the SED models and VLBI observations.  A possible idea to explain this discrepancy is that a jet component with relativistic speed other than C3, located much closer to the central black hole (within the central sub-pc), is responsible for the GeV $\\gamma $-ray emission. If such a relativistic motion in the jet of 3C~84 occurs only in vicinity of the black hole, the speed of C3 could have been relativistic soon after the ejection from the core, and then it underwent to a rapid deceleration. \\citet{2010A&A...516A...1L} also suggested that jets are relativistic in bright cluster galaxies but they decelerate very soon (sub-parsec scale) because of a strong interaction with the ISM. In Paper I, we argued possible detection of the relativistic speed at the early stage of emergence of C3. However, this was not conclusive because of (1) the lack of spatial resolution of VERA at 22~GHz and (2) possible screening by free-free absorption.  In order to overcome the above potential problems, higher resolution observations at higher frequency are required. We explore the possible relativistic flow of C3 at the early stage of its emergence from the core with the archival VLBA data at 43~GHz. We selected only data obtained from 2002 to 2008, i.e. considering earlier epochs than those discussed in Paper I. We also discuss the SED fit using one-zone SSC model involving the actual measured apparent motion of C3 and possible ideas for incorporating the VLBI observation with GeV  $\\gamma $-ray emission.  Throughout this paper, we adopt the following cosmological parameters; $H_{0}=71$~km sec$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\mathrm{M}}=0.27$, and $\\Omega_{\\mathrm{\\Lambda}}=0.73$ (1~mas = 0.353~pc, and 0.1~mas~yr$^{-1}$ = $0.113c$). ",
        "conclusions": "In this section, we try to reproduce the broadband flux from $\\gamma$-ray emitting region of 3C~84 based on the quantities of C3 obtained as in \\S \\ref{sec:results}. In Figure \\ref{fig:SED}, we show the observed SED of 3C~84 and one-zone SSC model fit to the observed SED using the size, flux, and $\\beta_{\\rm app}$ of C3 component measured by the VLBA observations. The flux data are the same as the ones in \\citet{2009ApJ...699...31A} but we add the 43-GHz flux of C3 measured on 2008 August 27. We assume that $\\gamma$-ray and VLBA 43-GHz flux are emitted from the same sub-pc region. Optical data (after the extraction of host galaxy contribution; c.f., \\citet{2009ApJ...699...31A}) is also used as flux upper limit of C3. The $\\beta _{\\rm app}$ at 2008 August is $0.44$ (in Figure \\ref{fig:gamma}). Assuming the jet viewing angle as $25^{\\circ}$ (\\citet{2009ApJ...699...31A}), the intrinsic jet speed $\\beta$ can be estimated as $0.54$, which corresponds to $\\Gamma=1.19$ ($\\delta=1.65$). The Gaussian-fitted model size of C3 is $\\sim 0.2$-$0.3~{\\rm mas}$. To fit the 43-GHz flux, the size of $\\gamma$-ray emitting region is adopted as $1.45 \\times 10^{17}~{\\rm cm}~{\\rm(0.133~mas)}$ which is slightly smaller than the C3 model size. The detailed numerical treatment of synchrotron and inverse Compton scattering processes is shown in \\citet{2002ApJ...564...97K} and references therein.  If we adopt $\\gamma_{\\rm min} \\sim 10^{3}$ as is the case in \\citet{2009ApJ...699...31A}, 43~GHz lies below the frequency of low energy cutoff and therefore the spectral index should be $\\alpha _{_K}^{_Q} = 1/3$ ($\\nu f_{\\nu} \\propto 4/3$). This is inconsistent with the observed spectral index between 22 and 43~GHz $\\alpha _{_K}^{_Q} = -0.87\\pm0.3$ as shown in Figure \\ref{fig:spectrum}. Then, we adopt $\\gamma_{\\rm min} = 1$ to raise the cutoff frequency. This fit is indicated by a solid line in Figure \\ref{fig:SED}. The model spectrum requires an extremely hard electron energy index $s=1.2$ occasionally seen in TeV blazars (e.g., \\citet{2000ApJ...528..243K}) if we try to reproduce the flux data at other wavelengths. The synchrotron opacity becomes unity at $\\sim 30$~GHz, and the resultant model spectrum shows $\\alpha _{_K}^{_Q} \\sim -0.05$, which is still inconsistent with the observed spectral index. Hence it seems difficult to attribute the broadband SED to the one-zone SSC emission from C3. This indicates that some other components C1 and/or C2 also should be taken in account for the candidates for $\\gamma$-ray emitting region.    Following the detection of 3C~84 by {\\it Fermi}/LAT in 2008, we discussed the kinematics of the newly formed jet component C3 using our VERA radio observations for the period 2006 June - 2009 April in Paper I. In this paper, we further explored the kinematics and lightcurves of C1, C2, and C3 using archival 43-GHz VLBA data spanning a time interval from 2002 January to 2008 November, i.e. before the VERA observations. Summary and discussions are as follows:  \\begin{enumerate} \\item In the multi-epoch high resolution images obtained by VLBA at 43~GHz (Figure \\ref{fig:3C84_Q_all}), we find that C3 has been visible since 2003 November. We succeed in obtaining a strong constraint on the time of the C3 birth before 2003 November.  \\item The lightcurve of VLBA at 43~GHz shows a trend similar to that of Mets\\\"{a}hovi at 37~GHz, indicating that the increase of radio flux density arises from the region within the central $\\sim1$~pc.  In particular, the rapid increase of flux density has been seen since the middle of 2008. Both C1 and C3 show the largest change of flux density for this flare.  This indicates that the radio flare in the central $\\sim1$~pc region originates in both C1 and C3.  The flux density increase of C1 was not prominent in Paper I, possibly because of large optical depth at 22~GHz.  \\item The apparent speed of C3 with respect to C1 changes from $(0.10 \\pm 0.05)c$ to $(0.47 \\pm 0.08)c$ between 2003 November 20 and 2008 November 27 with acceleration rate $\\dot{\\beta }_{\\rm app} = 0.08 \\pm 0.02 $ yr$^{-1}$. No superluminal motion is detected with VLBA observations at 43~GHz before 2008 November.  \\item The one-zone SSC model using measured apparent speed of C3 is fitted to the observed broadband spectrum, in particular, the 43-GHz and GeV $\\gamma $-ray fluxes. Our model fit is difficult to reproduce the optically-thin radio spectrum $\\alpha _{_K}^{_Q} \\sim -0.9$ that is measured by the VLBA observations. This indicates that it seems difficult to attribute the broadband SED to the one-zone SSC emission from C3. Therefore, the other components should not be excluded as candidates for $\\gamma$-ray emitting regions.  \\end{enumerate}"
    },
    "1112/1112.6084.txt": {
        "abstract": "We study the effect of primordial magnetic fields (PMFs) on the anisotropies of the cosmic microwave background (CMB). We assume the spectrum of PMFs is described by log-normal distribution which has a characteristic scale, rather than power-law spectrum. This scale is expected to reflect the generation mechanisms and our analysis is complementary to previous studies with power-law spectrum. We calculate power spectra of energy density and Lorentz force of the log-normal PMFs, and then calculate CMB temperature and polarization angular power spectra from scalar, vector, and tensor modes of perturbations generated from such PMFs. By comparing these spectra with WMAP7, QUaD, CBI, Boomerang, and ACBAR data sets, we find that the current CMB data set places the strongest constraint at $k\\simeq 10^{-2.5}$ Mpc$^{-1}$ with the upper limit $B\\lesssim 3$ nG. ",
        "introduction": "Introduction} %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ Many researchers have studied magnetic fields for wider ranges in the early Universe. Recently, magnetic fields with strength $\\sim 1 \\mu$ G have been observed in clusters of galaxies, and primordial magnetic fields (PMFs) have been studied by many authors to explain their origins. The strength of PMFs is constrained as $B\\lesssim 1$ nG from cosmological observations, such as temperature and polarization anisotropies of cosmic microwave background (CMB) and matter power spectra \\cite{Subramanian:1998fn, Mack:2001gc, Subramanian:2002nh, Lewis:2004ef, Yamazaki:2004vq, Kahniashvili:2005xe, Challinor:2005ye, Dolgov:2005ti, Gopal:2005sg, Yamazaki:2005yd, Kahniashvili:2006hy, Yamazaki:2006mi, Yamazaki:2006bq, Yamazaki:2006ah, Giovannini:2006kc, Yamazaki:2007oc, Paoletti:2008ck, Finelli:2008xh, Yamazaki:2008bb, 2008nuco.confE.239Y, Sethi:2008eq, Kojima:2008rf, 2008PhRvD..78f3012K, Giovannini:2008aa, Yamazaki:2009na, Shaw:2009nf, 2010PhRvD..81b3008Y, 2010PhRvD..81j3519Y, Yamazaki2010aa}.  Many theoretical models have been proposed for generating PMFs of cosmological scales. Generation models of PMFs from inflation can create strong fields, whose amplitude is of the order of $1$ nG at redshift $z\\sim 0$, depending on assumed hypothetical fields and couplings \\cite{Turner:1987bw,Ratra:1991bn,Bamba:2004cu,2004PhRvD..69d3507B, 2007JCAP...02..030B}. Furthermore, PMFs can be generated by primordial perturbations of density fields \\cite{Takahashi:2005nd,ichiki:2006sc}, the Beiermann battery mechanism in primordial supernovae remnants \\cite{Hanayama:2005hd} or the Weibel instabilities \\cite{2005MNRAS.364..247F}. These PMFs can evolve into the magnetic fields observed in galaxies and/or galaxy clusters directly or through the dynamo process.  While magnetogenesis during inflation can produce PMFs beyond the horizon scale, PMFs generated by the other mechanism are expected to have the characteristic scale because they are based on causal processes. Here it should be noted that the coherence length of PMFs can grow after generation due to inverse cascade \\cite{2001PhRvE..64e6405C,Banerjee:2004df}, although the efficiency is still under study.  In the previous studies to put constraints on PMFs from CMB anisotropies, the spectrum of PMFs was assumed to be power-law shape. Although power-law spectrum is natural for magnetogenesis during inflation, this is not the case for the other mechanisms based on causal processes. Furthermore, in previous studies with power-law PMFs, it was not so clear which scale of PMFs mainly contributes to the constraints. Since inflationary mechanisms tend to generate power-law magnetic fields, the magnetic fields which have a characteristic scale at cosmological scales are difficult to be produced and seem somewhat artificial. However, it is still possible that such fields can be constrained by observation, and if observed it would give us useful information about generation mechanism of cosmological magnetic fields.  Hence, in this article, as a toy example we use a log-normal distribution (LND) for the PMF spectrum $f_\\mathrm{LND}(k;k_\\mathrm{M},\\sigma_\\mathrm{M})$ which has a characteristic scale expressed by $k_\\mathrm{M}$, and aim to constrain PMFs scale by scale instead of the power-law PMF spectrum. These quantities reflect the generation mechanism but here we regard them just as parameters. We study features of angular power spectra of the CMB from the PMFs with log-normal distribution. Finally we constrain the strength of the PMFs for fixed sets of the parameters. %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ ",
        "conclusions": "%_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ It is natural to assume that the power spectrum of the PMF from inflation is the power law.  However, when we consider PMF generated by causal mechanisms, it would be appropriate to assume a spectrum with a characteristic scale. In this paper, we assumed the log-normal distribution as the PMF spectrum and studied the effects of the PMFs on the temperature and polarization anisotropies of the CMB.  We have revealed the features of the LND-PMF effects on the CMB as follows: (1) For larger $k_\\mathrm{M}$ and/or $\\sigma_\\mathrm{M}$, the CMB spectra of TT, TE and EE from the LND-PMF are dominated by the vector mode. Meanwhile, in the opposite case, these spectra are dominated by the scalar mode.  (2) The CMB spectrum of the BB mode for all scales and other modes for smaller scales is dominated by the vector modes. Because three parameters which characterize the LND-PMF affect the CMB power spectra differently at small and large scales we expect that tight constraints can be placed on these parameters without degeneracy. We found that the LND-PMF which generates CMB anisotropies among 1000 $<\\ell<$ 2000 is most effectively constrained by the current CMB data sets. For example, $B_\\mathrm{LND}$ for $10^{-3}\\mathrm{Mpc}^{-1}~\\le~k/h~\\le 10^{-2}\\mathrm{Mpc}^{-1}$ and $\\sigma_\\mathrm{M} = 1.0$ is limited most strongly as shown in Fig.\\ref{fig6}. In the near future, the tighter constraints on $B_\\mathrm{LND}$ at $k_\\mathrm{M}/h>10^{-2}$Mpc$^{-1}$ will be expected from the observations, such as the Planck, QUIET, and PolarBear missions. \\begin{figure} \\includegraphics[width=1.0\\textwidth]{fig2}% Here is how to import EPS art \\caption{\\label{fig2} CMB spectra from the LND-PMF at $\\sigma_\\mathrm{M}=1.0$ and $B_\\mathrm{LND} = $10nG. TT, EE, BB and TE in each panel indicate the temperature auto-correlation, the E-mode auto-correlation, the B-mode auto-correlation and the temperature E-mode cross correlation, respectively Curves in all panels are theoretical lines as indicated in the legend box at the left bottom. Curves of all TE modes are plotted in the absolute value. } \\end{figure}  \\begin{figure} \\includegraphics[width=1.0\\textwidth]{fig3}% Here is how to import EPS art \\caption{\\label{fig3} CMB spectra from the LND-PMF at $k_\\mathrm{M}=10^{-3}$ and $B_\\mathrm{LND} = $10nG. Curves in all panels are theoretical lines as indicated in the legend box at the left bottom. Curves of all TE modes are plotted in the absolute value. } \\end{figure}  \\begin{figure} \\includegraphics[width=0.8\\textwidth]{fig1}% Here is how to import EPS art \\caption{\\label{fig1} Power spectrum of the energy density of the PMF. Panels (a) and (c) on this figure are for the fixed $\\sigma_\\mathrm{M}$ value at $\\sigma_\\mathrm{M}=1.0$, and panels (b) and (d) are for the fixed $k_\\mathrm{M}$ value at $k_\\mathrm{M}=10^{-3}$. Curves in different colors in all panels correspond to the different PMF parameters as indicated in the legend box in the figure. Note that from Eqs.(\\ref{p:se})-(\\ref{p:t}) we see that the other spectra, such as magnetic tension, anisotropic stress and so on, have very similar shape to that of the energy density so we only show the spectrum of the energy density in the figure.} \\end{figure}  \\begin{figure} \\includegraphics[width=1.0\\textwidth]{fig4}% Here is how to import EPS art \\caption{\\label{fig4} Comparison of total CMB spectra from the LND-PMF at $\\sigma_\\mathrm{M}=1.0$ and $B_\\mathrm{LND} = $10nG with the best-fitting CMB power spectra without the LND-PMF (dashed line). Curves in all panels are the theoretical lines as indicated in the legend box on the figure. Curves of the TE mode are plotted in the absolute value. } \\end{figure}  \\begin{figure} \\includegraphics[width=1.0\\textwidth]{fig5}% Here is how to import EPS art \\caption{\\label{fig5} Comparison of total CMB spectra from the LND-PMF at $k_\\mathrm{M}=10^{-3}$ and $B_\\mathrm{LND} = $10nG with the best-fitting CMB power spectra without the LND-PMF (dashed line). Curves in all panels are the theoretical lines as indicated in the legend box on the figure. Curves of the TE mode are plotted in the absolute value. } \\end{figure}  \\begin{figure} \\includegraphics[width=1.0\\textwidth]{fig6}% Here is how to import EPS art \\caption{\\label{fig6} Constraint on the strengths of LND-PMF for $10^{-5}< k_\\mathrm{M} <10^{-1}$. The bold curve is the 2 $\\sigma$ upper limits of $B_\\mathrm{LND}$ [nG]. We fix the standard cosmological parameters and use the best-fitted value from WMAP 7th + tensor mode\\cite{2011ApJS..192...16L}. } \\end{figure}"
    },
    "1112/1112.3033_arXiv.txt": {
        "abstract": "Mapping one-dimensional stellar profiles onto multidimensional grids as initial conditions for hydrodynamics calculations can lead to numerical artifacts, one of the most severe of which is the violation of conservation laws for physical quantities such as energy and mass.  Here we introduce a numerical scheme for mapping one-dimensional spherically-symmetric data onto multidimensional meshes so that these physical quantities are conserved. We validate our scheme by porting a realistic 1D Lagrangian stellar profile to the new multidimensional Eulerian hydro code \\CASTRO.  Our results show that all important features in the profiles are reproduced on the new grid and that conservation laws are enforced at all resolutions after mapping. ",
        "introduction": "Multidimensional simulations shed light on how fluid instabilities arising in supernova explosions mix ejecta \\cite{herant1994,candace2009,candace2010,candace2011}. Unfortunately, computing the full self-consistent three-dimensional (3D) stellar evolution initial models for the explosion setup is still beyond the realm of contemporary computational power. One alternative is to first evolve the main sequence star in a 1D stellar evolution code in which the equations of momentum, energy and mass are solved on a spherically symmetric Lagrangian grid, such as \\KEPLER{} \\cite{kepler} or \\MESA{} \\cite{mesa}. Once the star reaches the pre-supernova phase, its 1D profiles can then be mapped into multidimensional hydro codes such as \\CASTRO{} \\cite{zhang2011,ann2010} or \\FLASH{} \\cite{flash} and continue to be evolved until the star explodes.  Differences between codes in dimensionality and coordinate mesh can lead to numerical issues such as violation of conservation of mass and energy when profiles are mapped from one code to another.  A first, simple approach could be to initialize multidimensional grids by linear interpolation from corresponding mesh points on the 1D profiles. However, linear interpolation becomes invalid when the new grid fails to resolve critical features in the original profile such as the inner core of a star.  This is especially true when porting profiles from 1D Lagrangian codes, which can easily resolve very small spatial features in mass coordinate, to a fixed or adaptive Eulerian grid.  Besides conservation laws, some physical processes such as nuclear burning are very sensitive to temperature, so slight errors in mapping can lead to very different outcomes for the simulation. Only a few studies have examined mapping 1D profiles to 2D or 3D meshes \\cite{zingale2002}, and none address the conservation of physical quantities by such procedures.  We investigate these issues and introduce a new scheme for mapping 1D data sets to multidimensional grids.  We first describe our mapping algorithm in \\Sect{method} and then present results of porting a massive star model from \\KEPLER{} to \\CASTRO{} in \\Sect{result}. Finally, we conclude in \\Sect{conclusion}. ",
        "conclusions": "\\lSect{conclusion} Multidimensional stellar evolution and supernova simulations are numerically challenging because multiple physical processes (hydrodynamics, gravity, burning) occur on many scales in space and time.  For computational efficiency, 1D stellar models are often used as initial conditions in 2D and 3D calculations.  Mapping 1D profiles onto multidimensional grids can introduce serious numerical artifacts, one of the most severe of which is the violation of conservation of physical quantities. We have developed a new mapping algorithm that guarantees that conserved quantities are preserved at any resolution and reproduces the most important features in the original profiles.  Our method is practical for 1D and 2D calculations, and we are now developing integral methods (an explicit integral approach instead of using volume subsampling) that are numerically tractable in 3D."
    },
    "1112/1112.5565_arXiv.txt": {
        "abstract": "We investigate baryogenesis in the $\\nu$MSM (neutrino Minimal Standard Model), which is the MSM extended by three right-handed neutrinos with masses below the electroweak scale.  The baryon asymmetry of the universe can be generated by the mechanism via flavor oscillation of right-handed (sterile) neutrinos which are responsible to masses of active neutrinos confirmed by various experiments.  We present the kinetic equations for the matrix of densities of leptons which describe the generation of asymmetries. Especially, the momentum dependence of the matrix of densities is taken into account.  By solving these equations numerically, it is found that the momentum distribution is significantly distorted from the equilibrium one, since the production for the modes with lower momenta $k \\ll T$ ($T$ is the temperature of the universe) is enhanced, while suppressed for higher modes.  As a result, the most important mode for the yields of sterile neutrinos as well as the baryon asymmetry is $k \\simeq 2 T$, which is smaller than $\\langle k \\rangle$ inferred from the thermal average.  The comparison with the previous works is also discussed. ",
        "introduction": "\\label{sec:Introduction} The baryon asymmetry of the universe (BAU) is one of the most puzzling problems in particle physics and cosmology.  It is certain that the Minimal Standard Model (MSM) cannot account for its origin. There have so far been proposed various scenarios of baryogenesis by considering physics beyond the MSM. (See, for example, a recent review \\cite{Riotto:1999yt}.)  Among them leptogenesis~\\cite{Fukugita:1986hr} by superheavy right-handed neutrinos is one of the most motivated scenarios.  This is because these new particles allow us to give non-zero masses to neutrinos which have been confirmed in many oscillation experiments.  Being singlet under the MSM gauge group they can obtain Majorana masses which are completely independent on the electroweak scale.  Right-handed neutrinos having superheavy masses then can generate the lepton asymmetry by their decays which can be a source of the BAU.  In addition, the observed smallness of active neutrino masses is naturally explained by such fermions through the seesaw mechanism~\\cite{Seesaw}.  The required masses are so heavy that it is impossible to directly test these new particles in near future experiments.  Akhmedov, Rubakov and Smirnov have proposed~\\cite{Akhmedov:1998qx} another attractive scenario of baryogenesis by using right-handed neutrinos (mechanism via neutrino oscillation).  One of the most important features is that the mechanism works when their masses are smaller than the electroweak scale, which are within the reach of current experiments.  The flavor oscillation among right-handed neutrinos in the early universe leads to the separation of the lepton number into right-handed neutrinos and left-handed leptons.  The sphaleron process then converts the asymmetry of the left-handed leptons into the baryon asymmetry for high temperatures~\\cite{Kuzmin:1985mm}.  This mechanism can be realized in a simple and attractive framework called as the neutrino MSM ($\\nu$MSM)~\\cite{Asaka:2005an,Asaka:2005pn} in which three right-handed neutrinos are introduced with Majorana masses below the electroweak scale.  Because of the very suppressed neutrino Yukawa interaction the seesaw mechanism is still effective, which leads to three active neutrinos $\\nu_i$ ($i=1,2,3$) and three sterile neutrinos $N_I$ ($I=1,2,3$) as mass eigenstates.  The former ones are responsible to the neutrino oscillations observed in experiments.  The latter ones give the solutions to the cosmological problems.  The lightest sterile neutrino $N_1$ with mass $\\sim 10$ keV can be a candidate of dark matter, while the rest two $N_2$ and $N_3$ with masses $\\sim 1$ GeV can generate the BAU through the mechanism above. (For a review see Ref.~\\cite{Boyarsky:2009ix}.)  Furthermore, the non-minimal coupling of Higgs field to gravity allows to realize the cosmic inflation~\\cite{Bezrukov:2007ep}.  Various aspects of baryogenesis in the $\\nu$MSM has been studied until now~\\cite{Asaka:2005pn,Shaposhnikov:2008pf,Asaka:2010kk,Canetti:2010aw}. The generation of the asymmetry is described by the matrix of densities~\\cite{Dolgov:1980cq,Barbieri:1990vx,Sigl:1992fn} of sterile (right-handed) neutrinos and active (left-handed) leptons.  The previous works were based on the kinetic equations for the matrix of densities in Ref.~\\cite{Asaka:2005pn}, where the terms describing the transfer of asymmetries between sterile neutrinos and active leptons are added to the original ones in \\cite{Akhmedov:1998qx}.  As shown in \\cite{Asaka:2005pn}, such terms are essential to generate enough amount of the BAU in the $\\nu$MSM.  In spite of their significance such terms are introduced by a heuristic approach.  The main motivation of this article is to improve the estimation of the BAU in the $\\nu$MSM.  One of the most unsatisfactory points in the previous works is that the evolution of the matrix of densities has been traced only by an approximate way.  Namely, the kinetic equations analyzed previously are for the typical, single mode with momentum $k \\sim T$ ($T$ is the temperature of the universe), which is expected to give a dominant contribution to the asymmetry.  By extrapolating the result of such a mode to other modes the asymmetry in the number density is estimated.  In this article, thus, we derive the kinetic equations in which the momentum dependence of the matrix of densities is fully taken into account.  For this purpose we shall re-evaluate the destruction and production rates of sterile neutrinos as well as active leptons paying attention to the two points.  One is the momentum dependence of the particle of interest.  The other is the corrections coming from the fact that the states in the destruction and production processes may differ from the thermal equilibrium states (\\ie, the deviations from the thermal equilibrium for sterile neutrinos and the chemical potentials for active leptons).  As we will show later, the kinetic equations are written as simultaneous integrodifferential equations for the matrices of densities of sterile neutrinos and active leptons.  Interestingly, our kinetic equations include the terms connecting between sterile and active sectors, which is crucial for the baryogenesis in the $\\nu$MSM as mentioned before.  It will be shown that such terms arise automatically as the corrections to the destruction and production rates owing to the deviations from the equilibrium states in the scattering processes.  It turns out that the coefficients of these terms are different from~\\cite{Asaka:2005pn} and, in addition, there appears a new type of terms which couples sterile neutrinos to active anti-leptons (rather than active leptons).  This paper is organized as follows: In Sec.~\\ref{sec:BGinNuMSM} we briefly review the framework of the analysis and the baryogenesis scenario in the $\\nu$MSM.  In the Sec.~\\ref{sec:KE}, we derive the kinetic equations which describe the generation of the baryon asymmetry paying attention to the momentum dependence in the matrices of densities.  Sec.~\\ref{sec:NS} is devoted to study the numerical solutions to the obtained equations.  We show the mode by mode evolution of the matrices of densities and their momentum distributions.  The comparison with the previous works is done in Sec.~\\ref{sec:Comparison}.  We perform the quantitative comparison of the yields of the BAU, but also clarify the differences between the kinetic equations in the literature.  We conclude in Sec.~\\ref{sec:Conc}. ",
        "conclusions": "\\label{sec:Conc} In this paper we have investigated baryogenesis induced by the flavor oscillation between sterile (right-handed) neutrinos $N_2$ and $N_3$ in the $\\nu$MSM.  We have presented the kinetic equations for the matrices of densities of sterile neutrinos $\\rho_N$ and $\\rho_{\\bar N}$ and chemical potentials of active (left-handed) leptons $\\mu_{\\nu_\\alpha}$ shown in Eqs.~(\\ref{eq:KE1}) and (\\ref{eq:KE3}). By using these equations the time evolution for the matrices of densities can be traced for each mode, and the momentum distributions can be found at any moments relevant for baryogenesis.  As a result they allow us to estimate the yield of the baryon asymmetry more precisely.  We have evaluated the destruction and production rates of sterile and active leptons for the kinetic equations.  Especially, we have paid a special attention to the following two respects.  The first is to include the momentum dependence of the rate correctly, which is mandatory to obtain the equations of $\\rho_{N, \\bar N}$ for a given momentum.  The second is to include the effect from the deviation of the initial or final state from the thermal equilibrium in the scattering process.  As explained in Sec.~\\ref{sec:KE}, the second point is crucially important in the construction of the kinetic equations for baryogenesis in the $\\nu$MSM. We have shown that the terms connecting between sterile and active sectors, which are required for the successful baryogenesis, arise in a consistent way from the second effect to the rate.  To be concrete, the terms connecting the matrices of densities $\\rho_N$ and $\\rho_L$, which introduced in Ref.~\\cite{Asaka:2005pn}, appear as the corrections to the production rates $\\delta \\Gamma_N^p$ and $\\delta \\Gamma_\\nu^p$ via the scattering processes (A) and (B). As shown in Sec.~\\ref{sec:Comparison} the prefactor is different by 2/3 to the previous result in Ref.~\\cite{Asaka:2010kk}.  Moreover, we have found the new terms which link $\\rho_N$ to $\\rho_{\\bar L}$ (rather than $\\rho_L$).  They are induced as the corrections $\\delta \\Gamma_N^d$ and $\\delta \\Gamma_\\nu^d$ to the destruction rates via the process (C).  Importantly, such a term is proportional to the product of $\\rho_{N , {\\bar N}}$ and $\\mu_{\\nu_{\\alpha}}$.  Thus, our kinetic equations are no longer linear in $\\rho_N , \\rho_{\\bar N}$ and $\\mu_{\\nu_{\\alpha}}$.  It should also be stressed that both types of these terms are required to ensure the conservation of the total lepton number.  The impacts of these terms on the generation of the baryon asymmetry have been described in detail.  We have then investigated the numerical solutions of the kinetic equations.  Because of the momentum dependence in the rates the production of sterile neutrinos is enhanced for the modes with $k/T \\lesssim 1$ while diminished for the higher momentum modes. Consequently, the momentum distributions of $\\rho_N$ and $\\rho_{\\bar N}$ are significantly distorted from the equilibrium one.  This clearly shows the importance to utilize the momentum-dependent kinetic equations.  It has also been found that the low modes get in the thermal equilibrium and the asymmetries carried by them receive the wash-out effect at higher temperatures.  In our choice of the parameters, the mode with $k/T \\simeq 2$ is the most significant for the occupation numbers as well as the asymmetries including the BAU.  The comparison with the previous works has also been discussed.  We have first explained what kind of approximations are needed to reproduce the previous kinetic equations staring from the present ones.  It is shown that the use of the destruction rate~\\cite{Akhmedov:1998qx} obtained by the thermal average is crucial to ensure the lepton number conservation when we treat the momentum dependence of $\\rho_N$ approximately.  We have then compare quantitatively the yields of the BAU.  Apart from the correct treatment of the momentum dependence, the present equations gives the smaller BAU by a factor of $\\simeq 2/3$ compared with Ref.~\\cite{Asaka:2010kk}.  Within the specific choice of parameters, the inclusion of the momentum dependence enhances the BAU by about 40\\% and 10\\% for $M_N=100$ MeV and 10 GeV, respectively.  This is because the enhancement in the generation of the asymmetry from the lower modes overcomes the suppression induced by the higher modes. On the other hand, such an amplification may be disturbed by the wash-out effect for large neutrino Yukawa coupling constants.  The present analysis shows that the momentum distribution of the matrices of densities differs from the equilibrium one and the inclusion of  their momentum dependence modifies the yield of the BAU in non-trivial manner depending on the parameters of neutrino Yukawa coupling constants and masses of sterile neutrinos.  Therefore, it is important to explore the full parameter space of the $\\nu$MSM accounting for the observed BAU by using the kinetic equations presented in this paper.  This issue will be performed elsewhere."
    },
    "1112/1112.2942_arXiv.txt": {
        "abstract": "We use the photoionisation and dust radiative transfer code MOCASSIN to create a model of the dwarf irregular galaxy NGC 4449. The best-matching model reproduces the global optical emission line fluxes and the observed spectral energy distribution (SED) spanning wavelengths from the UV to sub-mm, and requires the bolometric luminosity of 6.25$\\,\\times\\,$10$^9$ $L_\\odot$ for the underlying stellar component, $M_{\\mathrm{dust}}$/$M_{\\mathrm{gas}}$ of 1/680 and $M_{\\mathrm{dust}}$ of 2.2$\\,\\times\\,$10$^6$ $M_\\odot$. ",
        "introduction": "The spectral energy distribution (SED) of a galaxy emerges as a combination of the underlying stellar emission, the degree of reprocessing taking place in the interstellar medium (ISM) and thermal or non-thermal emission due to dynamical or evolutionary processes. As a consequence, galaxies emit across the entire range of the electromagnetic spectrum and the distribution of the energy output as a function of wavelength provide clues to the galaxy's past and present. Numerical models that reproduce the observed global properties of a galaxy over a wide range of wavelengths can be useful in understanding the physical and chemical nature of the underlying stellar and dust components. Here we present a MOCASSIN model of the dwarf irregular galaxy NGC~4449.  MOCASSIN \\cite[(Ercolano \\etal\\ 2003, 2005, 2008)]{Ercolano2003,Ercolano2005,Ercolano2008} is a fully 3-dimensional photoionisation and dust radiative transfer code. Recent applications of the code include modelling the planetary nebula NGC~6302 \\cite[(Wright \\etal\\ 2011)]{Wright2011} and studying dust formation by SN~2008S \\cite[(Wesson \\etal\\ 2010)]{Wesson2010}. Modelling entire galaxies is technically challenging due to their large physical sizes and the high level of complexity of the galactic environments. Dwarf galaxies are ideal candidates due to their intrinsically small sizes, relatively simple star formation histories and the low degree of processing of their interstellar material.  NGC~4449 (Fig.\\,\\ref{fig1}) is a metal-poor magellanic-type irregular galaxy. It is not a typical dwarf irregular: at the distance of 4.21 Mpc its H~{\\sc i} envelope extends to a radius of 14.2 kpc (11$.\\!\\!^{\\prime}$6), equal to approximately 6$R_{25}$ \\cite[(Swaters \\etal\\ 2002)]{Swaters2002}, and exhibits a counter-rotating core. It was selected for this study for its extensive coverage of photometric measurements for wavelengths ranging from the UV to sub-mm. Recently, new FIR observations were acquired by the {\\it Herschel Space Observatory} \\cite[(Pilbratt \\etal\\ 2010)]{Herschel} as part of the Guaranteed Time Key Project: Dwarf Galaxy Survey (PI: Suzanne Madden).  The emission line fluxes, used to constrain the models, were derived from global spectra, acquired with the drift scanning technique, and are representative of the entire galaxy \\cite[(Kobulnicky \\etal\\ 1999)]{Kobulnicky1999}.  \\begin{figure} \\vspace{2mm} \\begin{center} \\includegraphics[width=5.4cm]{./n4449_bw_sdssr.jpg} \\caption{SDSS \\textit{r}-band image of NGC~4449. North is up, east is to the left. The image is centred at 12$^{\\rm h}$28$^{\\rm m}$11$.\\!\\!^{\\rm s}$1, +44$^{\\circ}$05$^{\\prime}$37$^{\\prime\\prime}$ (J2000). The bar is 2$^{\\prime}$ in length.} \\label{fig1} \\end{center} \\end{figure}  \\vspace{-2mm} ",
        "conclusions": "The gas and dust masses derived from our MOCASSIN model of NGC~4449 are consistent with earlier estimates (\\cite[Hunter \\etal\\ 1999]{Hunter1999}; \\cite[Engelbracht \\etal\\ 2008]{Engelbracht2008}). Our best-matching star formation scenario broadly agrees with the presence of an old (3--5 Gyr) and a young population of stars, and continuous star formation over the last 1~Gyr (\\cite[Bothun 1986]{Bothun1986}; \\cite[Martin \\& Kennicutt 1997]{Martin1997}; \\cite[Annibali \\etal\\ 2008]{Annibali2008}). However, the predicted present-day star formation rate is only one-fourth of the estimate based on the H$\\alpha$ luminosity \\cite[(Hunter \\etal\\ 1999)]{Hunter1999}. The discrepancy between these two independent estimates highlights the differences in the adopted initial mass function (IMF) as well as the limitations of the two-episode model.  The presented technique assumes spherical symmetry for NGC~4449 and a two-episode star formation history, both of which are great simplifications. Also, at present PAHs are not included in the models and, as a consequence, the 4--20~$\\mu$m range is not correctly predicted. Nevertheless, most of the global parameters are predicted to within the observational uncertainties, and therefore this technique may prove useful in studying the characteristics of dust and the underlying stellar components in other dwarf galaxies.  \\begin{table} \\vspace{0.2cm} \\begin{center} \\begin{tabular}{l@{\\hspace{8mm}}cc} \\cline{1-3}\\noalign{\\medskip}  parameter & \\multicolumn{2}{c}{MOCASSIN model of NGC~4449}\\tabularnewline \\hline\\noalign{\\smallskip} $L_\\star$ & \\multicolumn{2}{c}{6.25$\\,\\pm\\,$0.25$\\,\\times\\,$10$^9$ $L_\\odot$} \\tabularnewline $M_\\star$ & \\multicolumn{2}{c}{1.2$\\,\\pm\\,$0.1$\\,\\times\\,$10$^9$ $M_\\odot$} \\tabularnewline $M_\\mathrm{gas}$ & \\multicolumn{2}{c}{1.5$\\,\\pm\\,$0.2$\\,\\times\\,$10$^9$ $M_\\odot$} \\tabularnewline $M_\\mathrm{dust}$ & \\multicolumn{2}{c}{2.2$\\,\\pm\\,$0.2$\\,\\times\\,$10$^6$ $M_\\odot$} \\tabularnewline $M_{\\mathrm{dust}}$/$M_{\\mathrm{gas}}$ & \\multicolumn{2}{c}{ 1/680 (1/850 to 1/540)} \\tabularnewline dust composition & \\multicolumn{2}{c}{100\\% amorphous carbon} \\tabularnewline dust grain sizes & \\multicolumn{2}{c}{0.005--1~$\\mu$m, $n \\propto a^{-3.5}$} \\tabularnewline \\noalign{\\medskip} SF episodes & 6$\\,\\pm\\,$2~Gyr to 120$\\,\\pm\\,$40~Myr & 120$\\,\\pm\\,$40~Myr  to present \\tabularnewline relative $M_\\star$ & 120 & 1 \\tabularnewline $<$SFR$>$ & 0.2 to 0.3 $M_\\odot$\\,yr$^{-1}$ & 0.06 to 0.12 $M_\\odot$\\,yr$^{-1}$ \\tabularnewline \\noalign{\\smallskip}\\cline{1-3} \\end{tabular} \\vspace{0.2cm} \\caption{Summary of global parameters derived from the best-matching MOCASSIN model of NGC~4449 assuming two continuous episodes of star formation. The global emission line fluxes are matched to within 20 per~cent or better, with the exception of the sulphur lines [S~{\\sc ii}] $\\lambda\\lambda$6717,31, which are overpredicted by a factor of four.} \\label{tab1} \\end{center} \\end{table}  \\vspace{-4mm}"
    },
    "1112/1112.0176_arXiv.txt": {
        "abstract": "{The Small Magellanic Cloud (SMC) hosts a large number of Be/X-ray binaries, however no Be/white dwarf system is known so far, although population synthesis calculations predict that they might be more frequent than Be/neutron star systems.} {\\xmms was found as a new faint super-soft X-ray source (SSS) with a likely Be star optical counterpart. We investigate the nature of this system and search for further high-absorbed candidates in the SMC.} {We analysed the \\xmm X-ray spectrum and light curve, optical photometry, and the $I$-band OGLE III light curve.} {The X-ray spectrum is well represented by black-body and white dwarf atmosphere models with highly model-dependent temperature between 20 and 100 eV. The likely optical counterpart AzV\\,281 showed low near infrared emission during X-ray activity, followed by a brightening in the $I$-band afterwards. We find further candidates for high-absorbed SSSs with a blue star as counterpart.} {We discuss \\xmms as the first candidate for a Be/white dwarf binary system in the SMC.} ",
        "introduction": "\\label{sec:introduction} Super-soft X-ray sources \\citep[SSSs, e.g. ][]{2006AdSpR..38.2836K} are a phenomenological class of X-ray sources, defined by a very soft X-ray spectrum without significant emission above 1 keV. They were discovered first in the Magellanic Clouds due to their close distance and low foreground absorption. The general scenario for a SSS is thermonuclear burning on the surface of an accreting white dwarf (WD), which can be stable for certain accretion rates \\citep{2007ApJ...663.1269N}. SSSs are associated with cataclysmic variables (CVs), symbiotic stars, and post-outburst optical novae. Super-soft X-ray emission at lower luminosity can e.g. originate from some CVs, cooling neutron stars (NSs), active galactic nuclei (AGN), PG 1159 stars (hot cooling isolated WDs), and planetary nebulae. Also Be stars \\citep[e.g. ][]{2003PASP..115.1153P} are believed to harbour accreting WDs. By a mechanism not well understood and in combination with high rotation, Be stars eject material in the equatorial plane causing the build up of a circumstellar decretion disc \\citep{2001PASJ...53..119O,2002MmSAI..73.1038Z}. This disc causes line emission and a near infrared (NIR) excess, both showing high variability pointing to the instability of these discs. How far the generation of the Be phenomenon is related to close binary evolution or another mechanism is still under debate \\citep{2000ASPC..214..668G}. In the case of close binary evolution, matter transfer causes a spin up of the gainer, becoming a Be star, whereas the donor turns into a He star, WD or NS. Hereafter, accretion from the decretion disc onto the compact object causes X-ray emission.   A high recent star formation \\citep{2010ApJ...716L.140A} and low metalicity \\citep{2006MNRAS.370.2079D} are probably responsible for the remarkably high number of $\\sim$90 known Be/X-ray binaries in the SMC. In these systems a NS accretes matter from the decretion disc of the Be star \\citep[for a recent review see ][]{2011Ap&SS.332....1R}. Binary system evolution models predict, that Be/WD systems might be more frequent than Be/NS systems \\citep{1991A&A...241..419P,2001A&A...367..848R}. \\citet{2001A&A...367..848R} obtain an abundance ratio for He star/NS/WD of 2/1/7, but no Be/WD system is known so far in the SMC in contrast to the large number of Be/NS systems. Similar to an accreting NS, WDs are expected to show hard X-ray emission powered by accretion, but at lower luminosity of $10^{29-33}$ erg s$^{-1}$, compared to $10^{34-38}$ erg s$^{-1}$ for NSs \\citep{1989A&A...220L...1W}. Suggested candidates are $\\gamma$\\,Cas like objects \\citep{1995A&A...296..685H,2006A&A...454..265L}. In addition, nuclear burning on the WD surface can produce super-soft X-ray emission at luminosities of $10^{35-38}$ erg s$^{-1}$, which might be absorbed by the decretion disc in most cases, complicating the discovery of such systems \\citep{1991A&A...248..139A}. Only one Be/WD SSS system has been proposed to date, \\wdbelmc \\citep{2006A&A...458..285K} in the Large Magellanic Cloud (LMC). To investigate the observational boundary conditions for Be/WD systems and to constrain the evolutionary models in general, it is of high importance to find more such systems.   The \\xmm \\citep{2001A&A...365L...1J} large programme survey of the SMC (Haberl et al. 2012, in preparation), in combination with archival observations, provides a flux limited sample of X-ray sources of the central field of the SMC. This enables comprehensive studies of SSSs \\citep[e.g.][]{2010A&A...519A..42M,2011A&A...529A.152S} as well as Be/X-ray binaries \\citep[e.g.][]{2011A&A...527A.131S,2011MNRAS.414.3281C}. In a search for new faint SSSs in our point source catalogue (Sturm et al. 2012, in preparation), we found one candidate, correlating with an early type SMC star. In this study, we present the results of the analysis of the \\xmm EPIC X-ray data, together with optical information. This led to the discovery of the first candidate Be/WD system in the SMC -- \\xmms\\ -- and illustrates, that \\wdbelmc in the LMC is not a unique case. ",
        "conclusions": "\\label{sec:discussion}  We analysed \\xmm X-ray data and OGLE III photometry of \\xmms. In the following we will discuss possible classifications of the source:  {\\it A Galactic star} can produce soft coronal X-ray emission. In this case the fainter object could be the true counterpart. But the emission of stars extends to higher energies than we see from our SSS causing a positive $HR_1$ in general \\citep{2004A&A...426...11P,2011A&A...534A..55S}. Also, the derived $\\chi_{\\rm red}^2$ of the apec model suggests formally that a description of the spectrum by a thermal plasma is less likely than by the other models.   {\\it A cooling isolated compact object}, i.e. a NS or PG 1159 star, can emit super-soft X-rays.  We would not expect to see an optical counterpart. In both cases, we cannot explain the variability, seen in the X-ray light curve.  For an {\\it AGN} with very soft X-ray emission and with a faint undetected optical counterpart, we expect to see the Galactic and total SMC absorption in the line of sight. The total SMC column density in the direction of the source is $N_{\\rm H, SMC} = 5.1 \\times 10^{21}$ cm$^{-2}$ \\citep{1999MNRAS.302..417S}. We derive a  $N_{\\rm H, SMC} < 1.5 \\times 10^{21}$ cm$^{-2}$ for the power-law model, which seems to contradict the AGN assumption.  {\\it A CV in the SMC} which by chance correlates with AzV\\,281 can account for the super-soft emission. However, there is no accretion powered system known with super-soft emission exceeding $10^{33}$ erg s$^{-1}$, which is  below the observed luminosity of \\xmms for the distance of the SMC. In the case of nuclear surface burning, higher luminosities can be reached (e.g. CAL 83 in the LMC). This would need a highly obscured WD, as discussed below. Another possibility for variable super-soft X-ray emission would be the rare case of a nova explosion. Here, super-soft emission is observed up to 2--5 years \\citep{2011A&A...533A..52H} after outburst. However, a nova explosion in this system is not known, which of course could have been missed. Due to the low nova rate in the SMC (only 3 novae were discovered during the last 10 years) this is an unlikely scenario.  {\\it A Galactic CV} cannot be excluded with the available data. E.g. the Galactic star could contain a WD companion. In this case, a soft intermediate polar \\citep{1995A&A...297L..37H} could account for the spectral properties, but only a handful of these systems is known in the Galaxy so far.  During the \\xmm SMC survey data analysis, we did not find any other SSS candidate correlating with an early-type emission line star. To estimate the chance for a random coincidence of a SSS or SSS candidate (14 in total) with a blue emission line star, we used a subset of early type SMC stars ($V<17$ mag, $-0.5$ mag $<B-V< 0.5$ mag, $-1.5$ mag $<U-B< -0.2$ mag) from the MCPS which correlate with an emission line object from MA93. This sample (1748 sources inside the observed \\xmm\\ field) was correlated with X-ray sources, which fulfill our criteria for super-soft X-ray emission and whose coordinates were shifted by multiples of the maximal correlation distance. We accepted correlations with an angular separation of $d\\leq3.439\\times (\\Delta r_{1}^2 + \\Delta r_{2}^2 )^{1/2}$, where $\\Delta r$ was estimated to 0.3\\arcsec and 2.0\\arcsec for the MCPS and MA93 sources, respectively. The \\xmm\\ sources have source-individual position uncertainties of $\\sim$1\\arcsec. In 440 runs, we obtained only 3 chance correlations. Analogous, using a set of USNO-B1.0 sources with proper motions $>$50 mas yr$^{-1}$ (134780 sources) results in 773 chance coincidences. This further affirms AzV\\,281 as the most likely optical counterpart. Also the coincidence of the X-ray emission with the NIR flux minimum further supports the identification of the super-soft X-ray source with AzV\\,281.  From a {\\it Be/NS system} we would expect outbursts with hard X-ray emission during accretion, which has not been observed for this system, although it was monitored regularly during the last 11 years. Also, observing only supersoft X-ray emission is atypical for a Be/NS system, as dominant hard emission is expected in additon to a soft excess \\citep{2004ApJ...614..881H}. Even in the extreme case of an outburst of RX\\,J0103.6-7201 \\citep{2008A&A...491..841E}, a hard X-ray component is present and the soft component extends to higher energies. From the bb+po model we obtain a limit for $L_{\\rm bb}/L_{\\rm X}$ of $>$0.5 in the (0.15-10.0) keV band, proofing the dominance of the thermal component. Therefore, we argue, that the system is less likely a classical Be/X-ray binary.   In contrast to the former possibilities, a {\\it WD/Be system} can account for both the super-soft X-ray emission and the variability. Black-body emission is a crude approximation to the spectra of SSSs \\citep{2010AN....331..146R}. Physically more meaningful are nlte models, but higher statistics and high resolution spectra would be necessary to constrain the parameters. To demonstrate the parameter dependence on elemental abundances we list two nlte models in Table~\\ref{tab:spec}. Solar abundances result in similar parameters as the black-body model with an emission radius of (73--106) km compared to (19--75) km from the bb model. Halo abundances result in a higher absorption, higher luminosity and an emission radius of (90000--190000) km. For an object in the SMC we expect abundances between solar and halo values \\citep{1992ApJ...384..508R,2010A&A...520A..85D}. Therefore the black-body emission radius is rather a lower limit. This further demonstrates, that luminosity and absorption can not be determined uniquely for \\xmms and that the spectra might be compatible with absorbed emission from a WD and luminosities sufficiently high for surface burning.  The presumable optical counterpart shows properties typical for a Be star. This makes the system the second candidate for a Be/WD binary after \\wdbelmc in the LMC \\citep{2006A&A...458..285K}.           \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0,clip=]{opt_lc.ps}}  \\caption{ {\\it Top:} $I$-band light curve from OGLE III. {\\it Bottom:} X-ray flux in the (0.2--1.0) keV band. } \\label{fig:lc} \\end{figure}   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[clip=]{opt_spec.ps}} \\caption{ Dereddened optical \\citep{2002AJ....123..855Z} and NIR \\citep{2007PASJ...59..615K} fluxes (squares) compared to a Kurucz model atmosphere for a B1V (upper line) and an O7V (lower line) star both normalised to the $U$-band flux. } \\label{fig:spec_opt} \\end{figure}    The X-ray turn-off around MJD 52500 can be caused by an exhaustion of the nuclear burning on the WD, as it occurs for post-nova SSSs, if the accretion rate is too low for stable nuclear burning. In the case of ongoing mass transfer, reignition of nuclear burning on the WD is possible in a nova explosion, followed by another super-soft X-ray emission state. More likely, the decrease in soft X-rays in the later observations might be explained by increasing photo-electric absorption due to a build-up of the decretion disc around the Be star. An additional absorbing column density of \\nh$\\gtrsim 10^{22}$ cm$^{-2}$ would be required for a non-detection of \\xmms\\ in the merged images. Be stars can loose and rebuild their discs on a time scale of some years including intervals without a disc of some hundred days \\citep{2011Ap&SS.332....1R,2010ApJ...709.1306W}. The decrease of the NIR flux before or during the X-ray active phase and the low photo-electric absorption during the three X-ray detections suggest this scenario. The X-ray detections have only a small overlap with the optical light curve, but since X-ray turn-off, the $I$-band shows an increase in flux, which can be caused by the recreation of a disc. We note, that the optical light curve of \\wdbelmc \\citep[cf. Fig. 4 in ][]{2006A&A...458..285K} has similar characteristics of both a long-term increase in the $I$-band and a periodicity around one thousand days.   The low X-ray absorption suggests a location of the system on the near side of the SMC bar. Also, the system intrinsic absorption along the line of sight must have been low during the X-ray detections. If the 1264 d period is caused by binarity, we get a relatively large separation of 5.7 AU ($\\sim$$140 R_{\\ast}$) for $M_{\\rm WD}=1\\msun$ and $M_{\\rm \\ast}=15\\msun$. But we note, that the realness and nature of the periodicity are uncertain, since the OGLE light curve covers only about two cycles. While \\citet{2001A&A...367..848R} argue that most Be/WD systems are not detectable in X-rays due to absorption by the decretion disc, truncation of the disc, as it is known from Be/NS systems with low excentricity \\citep{2001A&A...377..161O} might only allow accretion during disc instabilities. Also, the absorption by the SMC has to be low, to be able  to observe such systems in X-rays. We estimate the completeness in our SMC catalogue in the (0.2--1.0) keV band to a flux limit of 2.5\\ergcm{-15}, which would require a SMC absorption of less than 3 $\\times 10^{21}$ cm$^{-2}$ for the best-fit black-body model. Therefore, \\xmms might be the tip of the iceberg of the SMC Be/WD system population.       We report the discovery and analysis of a new SSS in the SMC, having a presumable early type star companion similar to \\wdbelmc in the LMC. This star has a variable NIR excess and emission lines, typical for a Be star. This makes the system the first candidate for a Be/WD system in the SMC. The X-ray turn off of the source around MJD 52400 can be caused by an increasing photo-electric absorption, possibly caused by the build up of a decretion disc around the Be star. We list further candidates for brighter but more absorbed Be/WD systems."
    },
    "1112/1112.0495_arXiv.txt": {
        "abstract": "Extending previous studies, we derive generic predictions for lower order cumulants and their correlators for individual tomographic bins as well as between two different bins. We derive the corresponding one- and two-point joint probability distribution function for the tomographic convergence maps from different bins as a function of angular smoothing scale. The modelling of weak lensing statistics is obtained by adopting a detailed prescription for the underlying density contrast. In this paper we concentrate on the convergence field $\\kappa$ and use top-hat filter; though the techniques presented can readily be extended to model the PDF of shear components or to include other windows such as the compensated filter. The functional form for the underlying PDF and bias is modelled in terms of the non-linear or the quasilinear form depending on the smoothing angular scale. Results from other semi-analytical models e.g. the lognormal distribution are also presented. Introducing a {\\em reduced} convergence for individual bins, we are able to show that the tomographic PDFs and bias for each bin sample the same functional form of the underlying PDF of density contrast but with varying variance. The joint probability distribution of the convergence maps that correspond to two different tomographic bins can  be constructed from individual tomographic PDF and {\\em bias}. We study their dependence on cosmological parameters for source distributions corresponding to the realistic surveys such as LSST and DES. We briefly outline how photometric  redshift information can be incorporated in our computation of cumulants, cumulant correlators and the PDFs. Various  approximate results for cumulants and their correlators are presented. Connection of our results to the full 3D calculations is elucidated. Analytical results for inclusion of realistic noise and finite survey size are presented in detail. ",
        "introduction": "Following the first weak lensing measurements \\citep{BRE00,Wittman00,KWL00,Waerbeke00} , the field of weak lensing has witnessed a tremendous progress in all fronts (see \\citet{MuPhysRep08} for a review). Currently, in terms of cosmological observations, weak lensing plays a role complementary to both Cosmic Microwave Background (CMB) studies and studies involving large scale structure (LSS) surveys. The ability of weak gravitational lensing to reveal cosmological information, particularly the dark energy equation of state is considerably enhanced by the inclusion of tomographic information. The impotance of weak lensing has spurred tremendous progress on the technical front in terms of specification and control of systematics. There are many ongoing and future weak lensing surveys such as the CFHT{\\footnote{http://www.cfht.hawai.edu/Sciences/CFHLS/}} legacy survey, the Pan-STARRS{\\footnote{http://pan-starrs.ifa.hawai.edu/}} and the Dark Energy survey{\\footnote{https://www.darkenergysurvey.org/}}, and further in the future, the Large Synoptic Survey Telescope{\\footnote{http://www.lsst.org/llst\\_home.shtml}}, Joint Dark Energy Mission or JDEM{\\footnote{http://jdem.gsfc.nasa.gov/}} that will map the dark matter and dark energy distribution of the entire sky in unprecedented details. In particular, owing to the large fraction of the sky coverage and tighter control on systematics as well as dense sampling of source galaxy populations it will be soon possible to study gravity induced non-Gaussianity with extreme accuracy. The gravity induced non-Gaussianity is typically probed using real space correlation functions as well as in the harmonic domain using their harmonic counterparts i.e. the multispectra (see e.g. \\cite{Pen03}). These correlation functions provide a set of tools to go beyond the usual power spectrum analysis. The higher-order correlation functions are important not only to break the parameter degeneracy  inherent in power spectrum analysis (e.g. between the amplitude of the matter power spectrum $\\sigma_8$ and the matter density parameter $\\Omega_{\\rm M}$) but also to understand error-estimates of lower-order correlations functions. Starting with the study of the three-point correlation function \\citep{Vil96,JainSeljak97} higher order statistics of weak lensing shear, convergence or flexions are now well understood from a theoretical point of view.  The power spectrum of density perturbations remains the most commonly used statistic in many cosmological studies. Weak lensing surveys probe the non-linear regime and are sensitive to non-Gaussianity which can not be probed using only the two-point correlation function or its harmonic analog the power spectrum. The statistics of shear or convergence probe the statistics of underlying mass distribution in an unbiased way \\citep{JSW00,MuJai01,Mu00,MuJai00,Valageas00, VaMuBa05,TakadaWhite03,TakadaJain04}, sensitive to nonlinear evolution due to gravitational clustering. Various analytical schemes from perturbative calculations to halo models have been employed to model the weak lensing statistics \\citet{Fry84,Schaeffer84, BerSch92,SzaSza93, SzaSza97, MuBaMeSch99, MuCoMe99a, MuCoMe99b, MuMeCo99, MuCo00, MuCo02, MuCo03, CooSeth02}). In addition to studying the statistics in projection they have also been studied in 3D using photometric redshifts. This approach can further tighten the constraints on e.g. the neutrino mass as well as the dark energy equation of state \\citep{Heav03,HRH00, HKT06, HKV07, Castro05, Kit08}. Tomographic techniques have also been employed as an intermediate strategy between projected surveys and 3D mapping \\citep{Hu99,TakadaJain04,TakadaJain03,Semboloni08}.  In this paper we extend previous results \\citep{JSW00,MuJai01,Mu00,MuJai00,Valageas00} on projected surveys by analysing the entire one-point PDF and the two-point PDF with tomographic information.The PDF contains information about the correlation hierarchy to an arbitrary order; the correlation hierarchy of the convergence field is directly related to that of the underlying mass distribution. We employ a generating function formalism that relies on {\\em hierarchical ansatz} on smaller angular smoothing scales and on perturbative results on larger scales. We define a reduced convergence for each bin and show that the different bins sample the same underlying PDF and bias functions (to be defined later) for the density contrast. The entire joint two-point PDFs for different pairs of redshift bins and individual PDF for each bins can be constructed from the PDF and the bias associated with individual bins because the joint PDF is factorisable in terms of the individual PDFs, bias and cross-correlations among various bins and different angular scales. We will show that individual redshift-resolved tomographic maps can be used to map out the PDF of the underlying mass distribution for a wide range of variance. This underlying PDF of the density contrast can be used to recover the tomographic PDF with the use of just two individual variables $\\kappa_{min}$ and the reduced variance for each bin; both of these variables are uniquely determined by the geometry and matter content of the Universe. The results are applicable not only to the PDFs as determined under hierarchical ansatz but also for other well motivated approximations for PDF such as the lognormal distribution.  Recent cosmological observations favour an accelerating Universe. This implies existence of energy of unknown nature (dark energy) which has negative pressure \\citep{Amen10,Wang10}. Current data continues to be consistent with dark energy being a non-zero cosmological constant. Though many other alternative dark energy candidates have been consider which are consistent  with data as well, e.g. quinessence, k-essence, spintessence. Different dark energy models can be classified according to the equation of state of of the dark energy component $w_{\\rm X}$. For quintessence model  $dw_{\\rm X}/dz>0$ while for k-essence models  $dw_{\\rm X}/dz<0$. There are many complimentary probes for dark energy, the distance-redshift relation of cosmological standard candles; Cosmic Microwave Background Anisotropy; volume redshift relations using galaxy counts; the evolution of galaxy clustering; weak lensing, etc. The different methods to probe dark energy are complementary to other and can provide important consistency check. Weak lensing surveys are particularly suitable for dark energy studies. All major weak lensing surveys has dark energy as their one of prime science driver. We will use the techniques developed in this paper to study two different dark energy model and compare the predictions against those of standard $\\Lambda$CDM model. The methods presented here are complementary to the usual Fisher matrix based approach that rely on two-point correlation functions or the power spectrum as it includes non-Gaussian information upto order .  This paper is organised as follows. In \\textsection2 we introduce our notation and present some standard results. In \\textsection3 we link the lower order statistics of weak lensing convergence to that of the underlying density distribution. In \\textsection4 we briefly review the hierarchical ansatz in the context of generating function formalism. In \\textsection6 we discuss the lognormal model in the context of weak lensing statistics. In \\textsection7 we derive the PDF and bias for various tomographic bins. The results are quite generic and can be used for arbitrary source redshift distribution. Finally the \\textsection8 is left for  discussion of our results. In an appendix we outline how in the context of tomographic binning the evolution topological estimators such as Minkowski Functionals can be studied using the lognormal distribution. ",
        "conclusions": "\\label{sec:conclu} Previous tomographic studies of weak lensing have typically worked with the lower order cumulants; we have generalized here these results to the case of the entire one- and two-point PDF, which contain information about the cumulants to an arbitrary order. The construction was performed using a generating function formalism based on hierarchical ansatz and a lognormal model. Our analysis generalizes previously obtained results derived for projected or 2D surveys. Though we have considered a top-hat filer convergence maps due to their simplicity, similar results can be obtained for related statistics such as the shear components or aperture mass $M_{ap}$ \\citep{BV00}.  The PDFs for the individual bins are constructed by generalization of the previously introduced global variable $\\kappa^{\\rm min}$, for individual bins i.e.$\\kappa^{\\rm min}$, that was used in the context of 2D projected maps. Next, using $\\kappa_{(i)}^{\\rm min}$,  reduced variable $\\eta^{(i)}$ is defined for each individual bins whose statistics can directly be linked to that of underlying density contrast $\\delta$. The convergence in individual bins can then be mapped to unique values of $\\eta=1+\\delta$ for a given smoothing angular scales $\\theta_0$.  For modelling the statistics of underlying density contrast $\\delta$ we have assumed two completely different model: the hierarchical ansatz along with its perturbative counterpart as well as the lognormal distribution. Both these approximations have been used successfully in various cosmological contexts. There are a wide class of models that are available under the general category of hierarchical ansatz. The main motivation behind our choice of a specific hierarchy is simplicity. In recent years more sophisticated models of hierarchical clustering have been proposed which rely more on inputs from numerical simulations. The generic results we have derived here indeed can be improved using such modelling though the fundamental characteristics will remain unchanged. \\begin{figure} \\begin{center} {\\epsfxsize=6. cm \\epsfysize=6. cm {\\epsfbox[27 426 315 709]{noise_diff.eps}}} \\end{center} \\caption{The difference of noisy $\\Lambda$CDM PDF and dark energy models $\\Delta_p(\\kappa)=p(\\kappa)-p_{\\Lambda \\rm CDM}(\\kappa)$ is plotted as a function of $\\kappa$. The smoothing angular scale, bin size and galaxy number density is as depicted. The scatter in estimation is smaller compared to the difference in the PDFs considered. The comsological parameters considered are the same as the ones in Figure-(\\ref{fig:dark}). The two survey configurations that we have considered both produces near identical results.} \\label{fig:10} \\end{figure} In our treatment we find, in agreement with \\cite{MuWa03}, the dynamical and geometrical contribution can be treated separately. The geometrical effects are completely encoded in a parameter $\\kappa^{min}$. The reduced convergence as defined is independent of the background geometry of the universe and essentially probe the evolution of gravitational clustering. We showed that a set of $\\kappa^{\\rm min}_{(i)}$ defined for a given set of redshift slices are adequate to characterize not only individual PDFs for each bin but it is also sufficient to study the joint two-point PDF among two different bins. The PDF of the reduced convergence $\\eta^{(i)}$ for individual bins or joint PDFs for a pair of bins generalizes the earlier studies where the projected or 2D maps were considered in a straight forward manner.  We also note that the construction of convergence maps is difficult compared to the direct evaluation of non-Gaussian statistics from shear maps. On the other hand convergence statistics can directly be modelled at arbitrary order whereas for shear field the computation is done mostly order by order manner. An independent analysis of convergence maps constructed from shear maps should therefore be useful in constraining various errors which might get introduced during various stages of data reduction.  In our analysis we have ignored the noise from intrinsic ellipticity distribution of galaxies as well as from shot noise resulting from finite number of galaxies that are used to trace the underlying mass distribution. These issues have been dealt with in great detail in \\cite{MuCo03,VaBaMu04}. Dividing the source population into bins reduced the number-density of sources. This in turn will increase the level of noise or the scatter in the estimator. In our analysis we have considered two different survey configurations, i.e. LSST and DES and found that for our choice of tomographic bins the one- and two-point PDFs are very similar in nature.  The lognormal distribution has already been used to model the statistics of weak lensing observables \\citep{Mu00,TTHF02} and the clustering of Lyman alpha absorption systems e.g. \\citep{BD97}. One-to-one mapping of initial density fields to evolved density fields using maps that are consistent with lognormal distribution function was not found to be very successful and the success of a lognormal distribution function in reproducing the statistics of gravitational clustering still remains somewhat unclear.  Tomographic weak lensing surveys can be cross-correlated with external data sets including frequency cleaned maps of secondaries from  ongoing CMB surveys; e.g. the thermal Sunyaev-Zeldovic (tSZ) maps or $y$-maps that will be available from CMB surveys such as Planck. The cross-correlation with tomographic information can help to understand the evolution of cosmological pressure fluctuations responsible for tSZ effect with redshift. The formalism presented here is perfectly suitable for such analysis. Detailed results of such analysis will be presented elsewhere. In addition to the weak lensing surveys the Supernova pencil beam surveys might also benefit for the results presented here.  To summarize, we have extended results derived in three different previous papers \\citep{MuCo03,VaMuBa05,VaBaMu04} to take into account tomographic bins within which the photometric redshift are available. The results obtained previously for one-point PDF are now extended to two-point PDF. These results can provide an alternative to usual Fisher-matrix analysis that is employed to optimize survey strategies. We have concentrated mainly on analytical results in this paper. The numerical results regarding optimization of survey strategy using these results will be considered elsewhere."
    },
    "1112/1112.5279_arXiv.txt": {
        "abstract": "{In the last few years Fermi and AGILE observations have indicated the existence of a possible population of transient MeV-GeV sources located on the Galactic plane and characterized by fast flares lasting only a very few days. Notably, no blazar-like counterparts are known within their error boxes so they could represent a completely new class of Galactic transient high energy emitters. The task of identifying their counterparts at lower energies remains very challenging. Despite this difficulty, INTEGRAL observations have provided intriguing hints that reliable candidate counterparts for these unidentified MeV-GeV transients could be found among the members of the recently discovered class of Supergiant Fast X-ray Transients (SFXTs). In this context, to date the best test case is represented by the association between the two sources IGR J17354$-$3255 and AGL J1734$-$3310. We will discuss their possible physical link and implications stemming from this association.}  \\FullConference{The Extreme and Variable High Energy Sky - extremesky2011,\\\\ September 19-23, 2011\\\\ Chia Laguna (Cagliari), Italy}   \\begin{document} ",
        "introduction": "The field of high energy astronomy is relatively young. Breakthrough results have been obtained only in the last twenty years thanks to satellites  carrying instruments whose survey capabilities  unveiled the extreme richness of objects at hard X-rays (E$>$20 keV, e.g. INTEGRAL/IBIS, Swift/BAT) as well as gamma-rays (E$>$100 MeV, e.g CGRO/EGRET, AGILE/GRID, Fermi/LAT). Interestingly, the great majority of such objects  are still unidentified, with no firmly established counterparts at lower energies. Their identification  is one of the great challenges of current high energy astronomy, it could leave some room for novel and unexpected discoveries. In this context, recent AGILE/GRID and Fermi/LAT observations have indicated the existence of a possible population of fast transient MeV-GeV sources located on the Galactic plane and characterized by flares lasting only a very few days (e.g. Hays et al. 2009, Bulgarelli et al. 2009, Chen et al. 2007). Notably, no blazar-like counterparts are known  within their error boxes so they could represent a completely new class of Galactic fast high energy transients. The task of identifying their counterparts at lower energies remains very challenging, mainly because of their often  large error circles (e.g radii typically from 10 arcmin  to 0.5 degrees). INTEGRAL/IBIS is particularly suited to search for reliable  best candidate counterparts thanks to i) a large field of view (FoV) which ensure a total coverage of the gamma-ray error box ii) a good angular resolution which is crucial to disentangle the hard X-ray emission of different sources in crowded fields such as those on the Galactic Plane iii) a good sensitivity above 20 keV. In particular, recent INTEGRAL results provided intriguing hints that reliable best candidate counterparts could be found among the members of the recently  discovered class of Supergiant Fast X-ray Transients (Sguera et al. 2009, Sguera et al 2011, Sguera 2009).  Supergiant Fast X-ray Transients (SFXTs) are a new class of High Mass X-ray Binaries (HMXBs) mainly unveiled thanks to INTEGRAL observations of the Galactic plane. They host a massive OB supergiant star as companion donor (Negueruela et al. 2006) and display X-ray flares lasting from a few hours to a few days (Sguera et al. 2005, 2006). It is generally assumed that all SFXTs host a neutron star as compact object  because their broad band X-ray spectra (0.2--100 keV) strongly resemble those of accreting X-ray pulsars in HMXBs, i.e. absorbed cutoff power law shape (e.g. Sidoli et al. 2009). This idea is supported by the detections of X-ray pulsations in some systems (e.g. Sguera et al. 2007). The typical  dynamic range of classical SFXT is 10$^{3}$--10$^{5}$, however some systems show a lower value of $\\sim$10$^{2}$ and so they have been named as intermediate SFXTs (Sidoli et al. 2011,  Clark et al.  2010, Walter \\& Zurita 2007). The physical mechanism driving the  peculiar X-ray behaviour of SFXTs  is still unclear and highly debated. Several different models  have been proposed in the literature (see Sidoli 2009  for a review).  We note that in principle SFXTs have all the \"ingredients\" to possibly be MeV-GeV emitters since they host a compact object (i.e. neutron star) as well as a bright and massive OB star which could act as source of seed photons (for the Inverse Compton emission) and target nuclei (for hadronic interactions).  In this respect it is important to point out that in the last few years a few classical HMXBs, having the same \"ingredients\" of SFXTs in terms of compact object and companion stellar donor, have been firmly detected at MeV-TeV energies as persistent and variable sources (e.g. LS 5039 and LS I +61 303, Paredes 2008, Hill et al. 2010) or as fast flaring transients (e.g. Cyg X-3 and Cyg X-1, Tavani et al. 2009, Sabatini et al. 2010), providing evidence that particles can be efficiently accelerated to very high energies in HMXBs. At odds with such gamma-ray binaries, the eventual MeV-GeV emission from SFXTs must be in the form of fast flares and  should be expected only for a very small fraction of time (i.e. from few hours to few days), making very difficult their detection. Despite this drawback, some observational evidences have been recently  reported in the literature on SFXTs as best candidate counterparts of unidentified transient MeV-GeV sources located on the Galactic Plane (Sguera et al. 2009, Sguera et al 2011, Sguera 2009). These evidences are merely based on intriguing hints such as a spatial correlation and a common transient behaviour on  similar, though as yet not simultaneous, short time scales. This scenario is also supported from an energetic  standpoint by a theoretical model based in the microquasar accretion/jet  framework (Sguera et al. 2009). The so far proposed associations represent an important first step towards obtaining reliable candidates on which to concentrate further efforts to obtain quantitative proofs for a real physical association.  In this respect, so far, the best test case is represented by the association between the two sources IGR J17354$-$3255 and AGL J1734$-$3310. ",
        "conclusions": ""
    },
    "1112/1112.0340_arXiv.txt": {
        "abstract": "Radio observations of galaxy clusters show that the intra cluster medium is permeated by $\\rm{\\mu G}$ magnetic fields. The origin and evolution of these cosmological magnetic fields is currently not well understood and so their impact on the dynamics of structure formation is not known. Numerical simulations are required to gain a greater understanding and produce predictions for the next generation of radio telescopes.  We present the galactic chemodynamics smoothed particle magnetohydrodynamic (SPMHD) code  (GCMHD+), which is an MHD implementation for the cosmological smoothed particle hydrodynamic code GCD+. The results of 1, 2 and 3 dimensional tests are presented and the performance of the code is shown relative to the ATHENA grid code. GCMHD+ shows good agreement with the reference solutions produced by ATHENA. The code is then used to simulate the formation of a galaxy cluster with a simple primordial magnetic field embedded in the gas. A homogeneous seed field of $10^{-11} \\rm{G}$ is amplified by a factor of $10^3$ during the formation of the cluster. The results show good agreement with the profiles found in other magnetic cluster simulations of similar resolution.  \\vspace{2mm}  \\noindent\\textbf{Key words:} magnetic fields - MHD - methods: numerical - galaxies: clusters: general \\vspace{0.1cm} ",
        "introduction": "Cosmological magnetic fields are thought to be ubiquitous through out the universe and have been detected by radio observations on scales as large as galaxy clusters. On the galactic cluster scale, magnetic fields have been detected via Faraday rotation measurements of distant quasars/AGN, by the observation of radio galaxies within the cluster and by the detection of diffuse synchrotron emission from radio haloes \\citep[e.g.][]{obs94,obs02,obs04,obs11,obs01}. The current observations find magnetic fields of $\\rm{\\mu G}$ strength permeates the intra-cluster medium (ICM) of most galaxy clusters. Beyond the cluster scale, measurements of the magnetic field are far less certain. Cosmological fields can be generated via Weibel's instability \\citep{sch03,med06}, Biermann's battery \\citep{bbm50,sub94,kul97,gne00}, structure formation \\citep{har70,ici06}, galactic winds \\citep{bec96,ber06,zfe09}, relativistic charged particles \\citep{min11} and various processes in the early universe \\citep{the02}. Current observations of the structural detail of these fields is limited, but the next generation of radio telescopes, such as the Square Kilometer Array (SKA) \\citep[e.g.][]{ska06,ska08}, will produce a wealth of observational data on the strength and structure of cosmological magnetic fields. To compare these observations with our knowledge of cosmological magnetic fields we require numerical simulations. The predictions made by simulations allow for a comparison of the theory with the observation. This will produce a more detailed understanding of cosmological magnetic fields.  The evolution of primordial magnetic fields in a cosmological setting has been studied using both smoothed particle magnetohydrodynamics \\citep[SPMHD, e.g.][]{cp09,niMHD11} and adaptive mesh refinement codes \\citep[e.g.][]{pmc08,pmc11}. The different techniques produce good agreement on the predicted strength and profile of a magnetic field in a galaxy cluster. They show that the initial structure of the field has little influence on the final field. The central cluster field strength predicted by these simulations agrees with the values inferred from observation. A few simulations have followed the creation of a cosmological magnetic field from either galactic wind pollution of the ICM \\citep{zfe09}, AGN pollution of the cluster \\citep{zfe10} or the Biermann battery mechanism during structure formation \\citep{zfe98} to generate a field from zero field initial conditions. The magnetic field strength can also be predicted a posteriori using the velocity fields from a purely hydrodynamical cosmological simulations \\citep{zfe08}. These simulations produce different field strengths than the values found from following the evolution of a primordial magnetic field, especially in filament structures. Further simulations are required to examine the different origins of a magnetic field in the ICM and to test the validity of the predictions made.  We present the implementation of magnetohydrodynamics (MHD) for the existing galactic chemodynamics code (GCD+) \\citep{gcd03,gcd09,gcd11}. GCD+ is a three-dimensional tree N-body/Smoothed Particle Hydrodynamics (SPH) code which incorporates self-gravity, hydrodynamics, radiative cooling, star formation, supernova feedback and metal enrichment.  The following MHD implementation is fully compatible with all of the original features of the code and allows for a complete cosmological simulation to be run. We discuss the choice of method to ensure tensile instability is avoided to an appropriate level. We also discuss the effect of an applied time dependant dissipation scheme \\citep{tdd97} for the magnetic field and the effect of the Balsara switch \\citep{cp95} on the test simulations, including its need in a cosmological simulation. For the tests shown in this paper we ignore the additional processes and concentrate on the effect of introducing MHD to the simulations and the accuracy of the non-radiative solutions produced. As such, radiative cooling, star formation, supernovae feedback and metal enrichment are switched off in all results presented.  In Section 2 we present the numerical implementation of the MHD equations used in GCMHD+. Section 3 shows the performance of the code in various test simulations. Section 4 shows the results of a cosmological simulation for the formation of a galaxy cluster with a simple primordial magnetic field embedded in the gas. Our conclusions are then given in Section 5. ",
        "conclusions": "We have introduced the MHD component to the N-body/SPH code GCD+. We discussed the addition of the equations of ideal MHD, the choice of instability correction and the addition of dissipative terms to treat discontinuities in the magnetic field. We implement schemes to remove the tensile instability, suggested by \\citet{ic01}, and for artificial magnetic dissipation, following \\citet{cp04a}. The code's ability to vary this dissipation in the simulation and allow each particle to evolve it own dissipation constant is presented. We put the code through a set of standard 1D shocktube tests, the fast rotor test, the Orszag-Tang vortex test and the MHD Point like explosion test. The numerical parameters were varied for all tests and the best compromise between noise reduction and minimised smoothing was found. The code with the best compromised parameter set performs very well in these tests and agrees with the reference solutions provided by the ATHENA mesh code, where they are available. The code shows no sign of the tensile instability and the magnetic dissipation scheme produces very little smoothing, while allowing the code to accurately capture the features. We then applied the code to a cosmological simulation for the formation of the Santa Barbara galaxy cluster. The code produces the expected hydrodynamic parameters of the cluster. The magnetohydrodynamic parameters are also well captured and show a similar level of magnetic field to the literatures with similar resolution \\citep{cp09}. We demonstrate that no minimum limit for the parameter of the dissipation of magnetic field, $\\alpha^B_{min}=0.0$, is necessary to minimise the artificial dissipation for the cluster scale magnetic field. This requirement is significantly lower than the previous SPMHD implementations, except for \\citet{niMHD11}. Our extensive test simulations in Section 3 demonstrate that $\\alpha^B_{min}=0.0$ still leads to satisfactory results. Encouraged by the success of our new MHD code, GCMHD+, we will apply it to higher resolution cosmological simulations, and study how magnetic fields developed in the evolving universe."
    },
    "1112/1112.4506_arXiv.txt": {
        "abstract": "We examine the {\\it Kepler} light curves of V1504 Cyg and V344 Lyr, encompassing $\\sim$736 d  at 1 min cadence. During this span each system exhibited $\\sim$$64-65$ outbursts, including six superoutbursts. We find that, in both systems, the normal outbursts lying between two superoutbursts increase in duration over time by a factor $\\sim$1.2$-$1.9, and then reset to a small value after the following  superoutburst. In both systems the trend of quiescent intervals between normal outbursts is to increase  to a local maximum about half way through the supercycle $-$ the interval from one superoutburst to the next $-$ and then to decrease back to a small value by the time of the next superoutburst. This is inconsistent with Osaki's thermal-tidal model, which predicts a monotonic increase in the quiescent intervals between normal outbursts during a supercycle. Also, most of the  normal outbursts have an asymmetric, fast-rise/slower-decline shape, which would be consistent with outbursts triggered at large radii. The exponential rate of decay of the plateau phase of the superoutbursts is 8 d mag$^{-1}$ for V1504 Cyg and 12 d mag$^{-1}$ for V344 Lyr. This time scale gives a direct measure of the viscous time scale in the outer accretion disk given the expectation that the entire disk is in the hot, viscous state during superoutburst. The resulting constraint on the Shakura-Sunyaev parameter, $\\alpha_{\\rm hot}\\simeq0.1$, is consistent with the value inferred from the fast dwarf nova decays. By looking at the slow decay rate for superoutbursts, which occur in systems below the period gap, in combination with the slow decay rate  in one long outburst above the period gap (in U Gem), we infer a steep dependence of the decay rate on orbital period for long outbursts. We argue that this relation implies a steep dependence of $\\alpha_{\\rm cold}$ on orbital period, which may be consistent with recent findings of Patterson, and is consistent with tidal torquing as being the dominant angular  momentum transport mechanism in quiescent disks in interacting binary systems. ",
        "introduction": "Cataclysmic variables (CVs) are semi-detached interacting binaries in which a Roche lobe filling K or M secondary transfers matter to a white dwarf (WD). CVs evolve to shorter orbital periods and show a ``gap'' between $P_h=2$ and 3 (where $P_h = P_{\\rm orbital}/1$ hr) during which time the secondary star loses contact with its Roche lobe and mass transfer ceases. Thus the binary becomes fully detached. At $P_h=2$ the secondary  comes back into contact with its Roche lobe and mass transfer resumes. For $P_h < 2$ angular momentum loss from the binary is thought to be due solely to gravitational radiation. The CV subclass of dwarf novae (DNe) are characterized by their semi-periodic outbursts. SU UMa stars are DNe lying below the period gap that exhibit  short, normal outbursts (NOs) and superoutbursts (SOs). We refer the time from one SO to the next as the supercycle. SOs show superhumps which are modulations in the light curve at  periods slightly exceeding the orbital  period. There are two further subdivisions within the SU UMa grouping: (i)  the VW Hyi stars at long orbital periods, near $P_h=2$,  for which the decay rate is fairly constant during a SO, and (ii) the WZ Sge stars at short orbital periods, a little greater than $P_h=1$, which have less frequent, larger amplitude SOs, for which the decay rate decreases during a SO. DNe outbursts are thought to be due to a limit cycle accretion disk instability (Lasota 2001) in which material is accumulated in quiescence and then dumped onto the WD during outburst. During short outbursts in longer period DNe, a  few percent of the stored mass is accreted, and during long outbursts a significant fraction $\\sim$0.2 of the stored mass is accreted. For the SU UMa stars, a SO  is thought to accrete   $\\ga0.7-0.8$ of the stored mass. Although the accretion disk is never in steady state during the limit cycle, it is close to steady state during SO, with the local rate of accretion inside the disk $|{\\dot M}(r)|$ decreasing linearly from a maximum value at the disk inner edge to zero at the outer edge. The accretion disk modeling has traditionally been done within the Shakura \\& Sunyaev (1973, hereafter SS) formalism, using two values for the $\\alpha$ viscosity parameter, $\\alpha_{\\rm hot}$ for gas in the hot, ionized disk, and $\\alpha_{\\rm cold}$ for gas in the quiescent disk.   \\begin{figure*} \\centering \\epsscale{1.0} \\includegraphics[scale=0.75]{figure1.jpg} \\vskip .35cm \\figcaption{ {\\it Kepler} light curve for V1504 Cyg at 1 min cadence, spanning $\\sim$736 d. The light curve covers six  superoutbursts and $59$ normal outbursts. The vertical red lines indicate the local maxima for the normal outbursts in which coverage permits a reliable determination, and with well-sampled decays down to $\\sim$2 mag below maximum. \\label{fig1}} \\smallskip \\end{figure*}     There are two bright SU UMa stars in the {\\it Kepler} field exhibiting a variety of temporal behavior that make them worthy of a detailed statistical study of their outbursts, V1504 Cyg ({\\it Kepler} ID 7446357; $P_h=1.67$) and V344 Lyr ({\\it Kepler} ID 7659570; $P_h=2.1$). These are members of the VW Hyi subdivision. To date the two light curves have amassed 736.4 d at 1$-$min cadence. Excluding gaps and bad data points, the light curves contain 1000431 and 1000345 data entries, respectively. Previous studies  of the  {\\it Kepler} data  on SU UMa stars have  found quiescent superhumps in V344 Lyr (Still et al. 2010), presented numerical models of the long term light curve of  V344 Lyr (Cannizzo et al. 2010; hereafter C10), and studied superhumps, both positive and negative,  in  the long term V344 Lyr light curve  (Wood et al.  2011).   \\begin{figure*} \\centering \\epsscale{1.0} \\includegraphics[scale=0.75]{figure2.jpg} \\vskip .35cm \\figcaption{ {\\it Kepler} light curve for  V344 Lyr at 1 min cadence, spanning $\\sim$736 d. The light curve encompasses six  superoutbursts and $58$ normal outbursts. The conventions are the same as in Figure 1. \\label{fig2}} \\smallskip \\end{figure*}   Statistical studies of DNe have been useful in delineating the   long-term properties of outbursts, e.g.., the outburst duration, recurrence times, and quiescent intervals, and placing physical constraints on models of outbursts (e.g., Campbell  \\& Shapley 1940, Sterne,  Campbell  \\& Shapley 1940, Bath \\& van Paradijs 1983, van Paradijs 1983, Szkody  \\& Mattei 1984, Cannizzo \\& Mattei 1992, 1998, Ak, et al. 2002, Simon 2004). Several interesting studies of the SU UMa stars have been carried out. For instance, van Paradijs (1983) studied the variation in outburst duration $t_b$ with orbital period for 15 DNe spanning the $2-3$ hr CV period gap. He  found that short outburst durations increase with orbital period, whereas long outburst durations are relatively constant with orbital period. Therefore for the SU UMa systems, which lie at short orbital period, the ratio $t_b({\\rm SO})/t_b({\\rm NO})$ is large. The relation of superoutbursts to normal outbursts for DNe below the period gap and the relation of long outbursts to normal outbursts for DNe above the period gap are part of a general trend; superoutbursts are just long outbursts in short orbital period DNe. This finding was amplified by Ak et al. (2002) using a larger sample. In this work we seek to extend previous studies by analyzing the high fidelity {\\it Kepler} light curves of two SU UMa stars to investigate the properties of the outbursts. In Section 2 we examine the outburst properties of the  NOs and SOs, in  Section 3 we  discuss the results, in particular the scaling of the SO decay rate with orbital period, and in Section 4 we summarize our results. ",
        "conclusions": "We have analyzed  long term {\\it Kepler} light curves of V1504 Cyg, containing 65 outbursts, and V344 Lyr, containing 64 outbursts.  Each system showed six superoutbursts. The findings are:  \\smallskip  (1) The NO decays are faster-than-exponential, with most of the deviation occurring near maximum light. Near quiescence the decays are close to exponential, $\\sim$0.7 d mag$^{-1}$ for V1504 Cyg and $\\sim$0.6 d mag$^{-1}$ for V344 Lyr. These rates are in line with that expected from the Bailey relation. Our values are based on frequency histogram distributions of only $\\sim$50 NOs and should improve as more outbursts are accumulated.  \\smallskip   (2)   The NO outburst durations $t_b$ increase by a factor $\\sim$1.2$-$1.9 over the time interval between consecutive superoutbursts, resetting to a small value after a  superoutburst. This is consistent with a monotonic increase of the accretion disk mass with each successive NO.  \\smallskip     (3)    The quiescent intervals $t_q$ between  normal outbursts show a general trend in both systems of increasing to a local maximum about mid-way between superoutbursts. (There was also an anomalously long quiescence interval at the start of   the V344 Lyr light curve, potentially the result of a tilted disk.) This behavior for $t_q$ is inconsistent with Osaki's thermal tidal instability model for SOs, in which one expects a monotonic increase in $t_q$ values accompanying the monotonic build-up in disk mass and angular momentum leading  to a SO. If $t_q$ is correlated with the NO triggering radius $r_{\\rm trig}$, then $r_{\\rm trig}$ moves outward with each successive NO only through roughly the first half of a supercycle, and then recedes.  \\smallskip  (4) The NOs in both systems are asymmetric, i.e., $t_{\\rm rise}/t_b < 0.5$, consistent with ``Type A'' outbursts (Smak 1984).  \\smallskip  (5) The inference of the steep relation  $\\alpha_{\\rm cold} \\simpropto P_h^{4}$ gives strength to the notion of tidal torquing as the dominant angular momentum transport mechanism operating in quiescent accretion disks in interacting binaries.  \\smallskip  If one considers the slow decay rate of the long, viscous outbursts spanning DNe from WZ Sge  to U Gem, one infers an overall variation $\\simpropto P_h^{1.6}$. It is unfortunate that only one long outburst in one DN longward of the $2-3$ hr period gap exhibited enough dynamic variation in flux during its plateau phase that a decay rate can be reliably extracted for the slow decay portion. It is fortuitous, given the general scatter in DN outburst properties at a given orbital period, that the decay rate between V1504 Cyg and V344 Lyr shows a similar law $\\simpropto P_h^{1.8}$. These dependencies are much steeper than the $P_h^{0.25}$ expected from the previous theoretical expression, due to Warner (1995a, 1995b). By starting with the definition for the viscous time scale in the outer disk during superoutburst, and taking into account the scaling for material accumulated during quiescence, we derive a more physically motivated formula, and find that a steep dependence of the quiescent state value of $\\alpha$ is required, $\\alpha_{\\rm cold}\\simpropto P_h^{4.2}$.  In simple terms, the dramatic increase in stored quiescent mass with decreasing orbital period brings about a decrease in the viscous time scale of matter in the superoutburst disk, which is formed out of redistributed gas from the quiescent state. Recent work by Patterson (2011) on the recurrence time scale in SU UMa stars for SOs also supports a steep dependence of $\\alpha_{\\rm cold}$ on $P_h$. The long outbursts in systems above the period gap and superoutbursts in systems below the gap both represent viscous decays, with the absence of transition waves in the disk. It is somewhat counterintuitive that the recurrence time for these long, viscous outbursts varies as $P_h^{-2.6}$, and yet their decay time varies as $P_h^{1.8}$. The steepness of the  $\\alpha_{\\rm cold}(P_h)$  relation explains why  accretion disk modelers have had to adopt such small $\\alpha_{\\rm cold}$ values for the SU UMa stars, compared to DNe above the period gap, and also probably accounts for the fact that the extreme mass ratio black hole transient systems undergoing limit cycle accretion disk outbursts,  like A0620-00,  have such long recurrence times for outbursts.   The {\\it Kepler} data have already  provided important constraints on the physics of accretion disks, in particular the Shakura-Sunyaev $\\alpha$ parameter: C10 showed that an interpolation between $\\alpha_{\\rm cold}$ and $\\alpha_{\\rm hot}$ based on the degree of partial ionization of gas was needed to account for the shoulders in the SOs of V344 Lyr, and in this work we have shown that the scaling of SO decay rate with orbital period mandates a steep  dependence of $\\alpha_{\\rm cold}$  on orbital period, enforcing the notion of tidal torquing as being the dominant viscosity mechanism in quiescent accretion disks.   \\smallskip    We acknowledge the contributions of the entire {\\it Kepler}  team.     \\def\\mnras{MNRAS} \\def\\apj{ApJ} \\def\\apjs{ApJS} \\def\\apjl{ApJL} \\def\\aj{AJ} \\def\\araa{ARA\\&A} \\def\\aap{A\\&A} \\def\\aapl{A\\&AL} \\def\\pasj{PASJ}"
    },
    "1112/1112.3260_arXiv.txt": {
        "abstract": "We present a multifrequency approach which optimizes the constraints on cosmological parameters with respect to extragalactic sources and secondary anisotropies contamination on small scales. We model with a minimal number of parameters the expected dominant contaminations in  intensity, such as unresolved point sources and the thermal Sunyaev-Zeldovich effect. The model for unresolved point sources, either Poisson distributed or clustered, uses data from {\\it Planck} early results. The overall amplitude of these contributions are included in a Markov Chain Monte Carlo analysis for the estimate of cosmological parameters. We show that our method is robust: as long as the main contaminants are taken into account the constraints on the cosmological parameters are unbiased regardless of the realistic uncertainties on the contaminants. We show also that the two parameters modelling unresolved points sources are not prior dominated. ",
        "introduction": "The CMB anisotropy data constitute a fundamental tool to test the standard cosmological model and its extensions. In particular small scale CMB anisotropies have a great importance in cosmology and are one of the current frontiers in CMB observations. Many experiments have been dedicated to the observations of small scale anisotropies. The {\\it Planck} satellite \\citep{Bluebook} will cover all the scales up to $\\ell\\sim 2500$ but there are also ground based experiments observing smaller regions of the sky but with higher angular resolutions, such as: ACBAR \\citep{ACBAR}, CBI \\citep{CBI}, QUaD \\citep{QUAD}, SPT \\citep{SPT}, ACT \\citep{ACT}. At these scales, the Silk damping suppresses the primordial CMB contribution with respect to extragalactic contamination and secondary anisotropies. These contaminations affect the constraints on cosmological parameters derived from small scale data introducing biases and therefore it is necessary to properly account for them in the data analysis. We present an approach which has been developed to optimize, with respect to the contamination by point source contributions, the constraints on cosmological parameters which could be obtained from high resolution multi-frequency data.\\\\ The problem of extragalactic sources and secondary anisotropies contamination of the angular power spectrum and consequentely of the cosmological parameters has already been addressed in several recent articles \\citep{Douspis2006,TaburetBias,TaburetSZ,EfstathiouGratton,Millea,IRCTemplate,SZGeorge}, together with the specific tools developed by the data analysis groups of the most recent small scale experiments \\citep{SPT2011,ACT2010}. \\\\  The approach we present in this article consists in modelling the astrophysical contributions to the angular power spectrum with minimal parametrizations, where minimal relates to the number of parameters necessary to describe the signals, and considering the frequency dependences of the different contributions. This is the most important difference with respect to previous approaches \\citep{EfstathiouGratton,Millea,IRCTemplate,SPT2011,ACT2010}. These parametrizations are included as additional contributions to the primary CMB in the Markov Chain MonteCarlo analysis to constrain cosmological parameters. The parameters characterizing the astrophysical contributions are included in the analysis together with the cosmological ones. As described in the following, we focus on the frequencies: $70,100,143,217,353$ GHz. We use the available data from {\\it Planck} early results \\citep{PlanckRadio,PlanckCIB} to derive the parametrizations to describe the astrophysical signals. When no data are available, we rely on predictions from theoretical models and empirical simulations.\\\\  In the following, we use a fiducial cosmological model with the parameters from \\cite{WMAP7} which are reported in the third column of  Table 4. The polarization from point sources and SZ effect is negligible we thus do not take it into account. We also do not account for residual signal, in intensity nor in polarization, from our Galaxy. \\\\  The article is organized as follows.  In Section 2 we describe the astrophysical signals and we derive the parametrizations of their contributions to the angular power spectrum. In Sect 3 we describe the analysis method and the frequency channel combinations. In Section 4 we present the results of the Markov Chain MonteCarlo (MCMC) analysis for the cosmological and astrophysical parameters. In section 5 we present the discussion and conclusions. ",
        "conclusions": "In the present study, we developed a multi-frequency approach of the analysis of the CMB power spectrum aiming at obtaining unbiased cosmological parameters in the presence of unresolved extragalactic point sources and TSZ signal.  We have considered the three main extra-galactic contributions to the CMB signal at small scales, such as radio sources, infrared sources and SZ clusters. We have modelled these three signals in a minimal way using the observational constraints from {\\it Planck}'s early results \\citep{PlanckRadio,PlanckCIB}. In particular, the radio and IR sources were modelled through a unique Poissonian term and a unique clustering term. In total, three amplitude terms suffice to account for and characterize the extra-galactic contributions. This minimal number of parameters was obtained by considering the frequency dependences within our parametrizations of the TSZ signal, and IR and radio point sources.  The use of a data-based physical model including the frequency dependence of the astrophysical signals under consideration to build the parametrizations of the small-scale contribution to the CMB is one of the original points of the present work. Previous studies \\citep{EfstathiouGratton,Millea} chose blind parametrizations of the signals which consider generic shapes and parametrize each frequency with a different set of parameters. Such an approach applied for a frequency channels, leads to either a rapid increase of the number of parameters necessary to describe the astrophysical signals, like in \\cite{Millea}, or it requires the use of a limited number of frequencies, like in \\cite{EfstathiouGratton}. Other studies focused only on one single contribution using more detailed and physical based parametrizations which involve a higher number of parameters with respect to our approach (see for example \\cite{SZGeorge} for the TSZ) or considered the frequency dependence on a much simpler level like for the clustering contribution in \\cite{IRCTemplate}. A different approach, relying directly on the data, has been used by the SPT team \\citep{SPT2011,SPT}. They took advantage of the high resolution to acquire data on scales where the CMB is suppressed. They used the data on smallest scales to measure the amplitudes of their blind templates for the astrophysical signals and consequently to remove the associated contamination from the CMB data.  We have applied our multi-frequency minimal parametrization of the extra-galactic contamination to estimate the cosmological parameters for a standard six-parameter $\\Lambda$CDM model and for a model with tensor perturbations and a running of the spectral index, without including residuals from our Galaxy and without using the $B$-mode information. We show that the best combination of frequency is the one which considers the frequencies from 70 to 353 GHz. It allows to estimate the cosmological parameters and at the same to put the tightest constraints on the amplitude of the IR clustering term. We show, both in the standard $\\Lambda$CDM model and in the \"extended-parameter\" model, that the input cosmological parameters together with the astrophysical ones (amplitudes of the TSZ, Poisson and clustering terms) are recovered as unbiased and with slighlty enlarged error bars. We show that the Poisson and clustering terms are degenerate.  Our approach instead of blind parametrizations prefers to use the whole available information on point sources and theoretical knowledge of the TSZ to have models of the astrophysical signals which are the most similar as possible to the expected signal. The use of the frequency dependence and the technique used to fit the data allowed a very restricted number of parameters to characterize the signals, in particular only a total of three parameters for all the signals.  We have investigated for both cosmological models, standard and extended, the dependence of our approach and of the derived parameters on the accuracy of our model of extra-galactic contributions. In this way we tested its robustness against our theoretical priors on frequency dependence. We tested different Poissonian and clustering contributions by varying the associated terms both in amplitude and in frequency dependence. We showed that for both cosmological models the cosmological parameters are recovered unbiased even in situations where the mock data differ by several sigmas from the model. The astrophysical parameters, amplitudes $A_{\\rm PS}$ and $A_{\\rm CL}$, are recovered well in the case of a ``real'' Poissonian term different from that of the parametrisation. In the case of a different clustering term, the degeneracy with the Poissonian term induces a shift also in the Poissonian term.  We have provided a minimal parametrisation to take into account the extragalactic contaminations on small scales. Our minimal analysis can be extended to include additional contributions like residuals from the Galaxy contamination. However, it is ideal in the sense that systematic effects, beam uncertainties etc are not accounted for.   \\begin{table} \\centering \\small\\addtolength{\\tabcolsep}{-4pt} \\begin{tabular}{|l|c|c|c|c|c|c|} \\hline Channel & $70$ & $100$  & $143$  & $217$ & $353$\\\\ (GHz) &  &   &   &  & \\\\ \\hline $A^{\\mathrm Low}$& 8.06 &7.76  & 7.47   & 7.16 & 6.81 \\\\ $B^{\\mathrm Low}$& $10^{14}$ & $10^{14}$  & $10^{14}$   & $10^{14}$  & $ 10^{14}$ \\\\ $C^{\\mathrm Low}$& 1 & 1  & 1   & 1 &1 \\\\ $\\alpha^{\\mathrm{Low}}$& 0.45 & 0.45 &   0.45   & 0.45 & 0.45 \\\\ $\\beta^{\\mathrm{Low}}$& 2.5 & 2.5 &  2.5   & 2.5 & 2.5 \\\\ $w_{\\nu}^{\\mathrm Low}$& 4.4 & 4.0 &  5.0   &  7.5 &  4.5 \\\\ \\hline $A^{\\mathrm{Med}}$& 306.98 &224.99  & 241.08   & 203.69 & 165.85 \\\\ $B^{\\mathrm{Med}}$& 20408 & 40000 & 111111  &   308642 &  501187 \\\\ $C^{\\mathrm Med}$& 1 & 1  & 1   & 1 &1 \\\\ $\\alpha^{\\mathrm{Med}}$& 0.75 & 0.74 & 0.75   & 0.73 & 0.7 \\\\ $\\beta^{\\mathrm{Med}}$& 2. & 2. & 2.  &  2 & 1.9 \\\\ $w_{\\nu}^{\\mathrm Med}$&2.3 & 2.3 &  2.7   &  2.7 &  3.8 \\\\ \\hline $A^{\\mathrm High}$& 58.41 & 22.66 & 70.89   & 23.12 &51.17 \\\\ $B^{\\mathrm High}$& 1. & 0.91 &  3.52  &   1.48 &  0.93 \\\\ $C^{\\mathrm High}$& 1 & 0.25  & 1   & 0.29 &1 \\\\ $\\alpha^{\\mathrm High}$& 0.75 & 0.85 & 0.8   & 0.85 & 0.75 \\\\ $\\beta^{\\mathrm High}$& 1.45 & 1.0 & 1.2   & 1.1 &  1.5 \\\\ $w_{\\nu}^{\\mathrm High}$&1.0 & 1.0 & 1.04   &  1.0 & 1.03 \\\\ \\hline \\end{tabular} \\label{FPOI} \\caption{The exponents $\\alpha$, $\\beta$ and the coefficients $A$, $B$, $C$ and $w_{\\nu}$ for each term of the fits of the radio galaxies composite number counts. } \\end{table}   \\begin{table} \\centering \\small\\addtolength{\\tabcolsep}{-4pt} \\begin{tabular}{|l|c|c|c|c|c|c|} \\hline Channel & $70$ & $100$  & $143$  & $217$ & $353$\\\\ (GHz) &  &   &   &  & \\\\ \\hline $P_1$& $1.1 10^{-4}$ & $3.6\\times 10^{-5}$ & $1.1\\times 10^{-5}$ & $4.3\\times 10^{-6}$ & $1.8\\times 10^{-5}$ \\\\ $\\chi_1$& 0.95  & 0.95  &  0.95  & 0.95  &0.95 \\\\ $a_1$& 0.34 &  0.34 &  0.34  &  & \\\\ $b_1$& 1 & 1 & 1     & 1 & 0.34 \\\\ $c_1$& 1.334 & 1.34 & 1.34    &  1.34 & 1\\\\ $d_1$&$-10^{14}$ &$ -10^{14}$ & $ -10^{14}$  & $-10^{14}$  & $ -10^{14}$ \\\\ $e_1$& 2.5 & 2.8 &  2.8   & 2.8 &2.8  \\\\ \\hline $P_2$&$ 5.9\\times 10^{-3}$ &  $1.4 \\times 10^{-3}$ & $ 3.6 \\times 10^{-4}$ & $1.5 \\times 10^{-4}$ & $4.1 \\times 10^{-4}$\\\\ $\\chi_2$&1.28  & 1.24 & 1.2  & 1.15   &1.2 \\\\ $a_2$& 0.64 & 0.62  & 0.6   &0.58  & 0.63\\\\ $b_2$& 1 & 1. &   1.  &  1. & 1.\\\\ $c_2$& 1.64 &  1.62 &  1.6   & 1.58 &1.63 \\\\ $d_2$& -20408.2 & -40000  & -111111.    &-308642. & -501187  \\\\ $e_2$& 2 & 2 &  2   &2 &  1.9 \\\\ \\hline $P_3$&$2.6\\times 10^{-3}$  &$1.2\\times 10^{-3}$ & $3.6 \\times10^{-4}$ & $2.5 \\times 10^{-4}$  & $4.5 \\times 10^{-4}$ \\\\ $\\chi_3$&1.25 & 1.35  & 1.3   & 1.35 &1.25 \\\\ $a_3$& 0.86 & 1.  & 1   & 1. & 0.83\\\\ $b_3$& 1. & 1.35 &  1.08   & 1.22727 & 1.\\\\ $c_3$&1.86  & 2.35 & 2.08    & 2.23 & 1.83\\\\ $d_3$&  -1 &  -3.64 &  -3.52   &  -5.10 & -0.93\\\\ $e_3$& 1.45 & 1 & 1.2    & 1.1 & 1.47\\\\ \\hline \\end{tabular} \\label{ClPTerms} \\caption{ The coefficients and exponents of the radio galaxy Poissonian contribution to the angular power spectrum Eq.\\ref{RadioP}. } \\end{table}  \\begin{table} \\centering \\small\\addtolength{\\tabcolsep}{-5pt} \\begin{tabular}{|l|c|c|c|c|c|} \\hline $a_i^j$&j & 0 & 1 & 2  & 3     \\\\ \\hline i    &  &     &   &    &     \\\\ \\hline 0    &  & 14541.6 & -298.31 & 1.84   &$ - 3.26\\times 10^{-3}$ \\\\ \\hline 1    &  &   -718.25     &15.14  &  $- 9.8\\times 10^{-2}$  & $1.9\\times 10^{-4}$ \\\\ \\hline 2    &  &     19.52   & - 0.42  & $2.8\\times 10^{-3}$   &$- 5.99 \\times10^-6$  \\\\ \\hline 3    &  &    -0.19    & $4.2\\times 10^{-3}$  & $ - 2.8\\times 10^{-5}$  & $6.1\\times 10^{-8} $ \\\\ \\hline 4    &  &  $ -2.\\times 10^{-4} $    & $4.4\\times 10^{-6}$  &  $- 3.03\\times 10^{-8}$  &   $6.72\\times 10^{-11}$ \\\\ \\hline 5    &  &  $ 1.83\\times 10^{-8}$    & $- 3.99\\times 10^{-10} $ & $ 2.78\\times 10^{-12}$  & $6.18\\times 10^{-15}$ \\\\ \\hline 6   &   &   $ 4.58\\times 10^{-13} $  & $- 9.72\\times 10^{-15}$  &   $6.50\\times 10^{-17}$ & $-1.36\\times 10^{-19}$ \\\\ \\hline 7   &   &    -108.25   & 2.40  & $ - 1.67\\times 10^{-2}$  & $ 3.66\\times 10^{-5}$\\\\ \\hline 8   &   &      3.81 & $ - 6.03\\times 10^{-2}$ & $ 3.87\\times 10^{-4} $  & $-6.53\\times 10^{-7}$ \\\\ \\hline 9   &   &    2.76   & $- 2.08\\times 10^{-3}$  & $- 1.13\\times 10^{-5}  $ & $ 4.07\\times 10^{-8}$ \\\\ \\hline \\end{tabular} \\label{ClusteringTable} \\caption{ The coefficients and exponents of the fits for the clustering term Eq.\\ref{clfit}. } \\end{table}   \\begin{figure} \\includegraphics[{width=85mm,height=90mm}]{CMBPSSZCL_70353_RUNGW.eps} \\caption{Results of the analysis which considers together with cosmological and astrophysical parameters also the variation of the tensor to scalar ratio and the running of the spectral index. Comparison between the results varying or not the astrophysical contributions: the black (solid) curve considers the astrophysical signals whereas the red (dashed) curve represent the results with astrophysical contributions fixed to the fiducial value. Vertical bars are the input parameters.} \\label{GWRUNFG} \\end{figure} \\begin{figure} \\includegraphics[scale=0.3]{CMBPSSZCL_70353_RUNGW_tri.eps} \\caption{Results of the analysis which considers together with cosmological and astrophysical parameters also the variation of the tensor to scalar ratio and the running of the spectral index. Comparison between the results varying or not the astrophysical contributions: the black (solid) curve considers the astrophysical signals whereas the red (dashed) curve represent the results with astrophysical contributions fixed to the fiducial value. We show the comparison of the two analysis for the running of the spectral index and the scalar spectral index. Curves are the $68\\%$ and $95\\%$ confidence level.} \\label{GWRUNFG_2D} \\end{figure} Acknowledgements: We wish to thank Nicolas Taburet for useful discussions and for providing the TSZ power spectrum template and Marco Tucci for provinding the number counts of \\cite{Tucci}.This work is partially supported by ASI contract Planck-LFI activity of Phase E2, by LDAP. NA wishes to thank IASF Bologna for hosting and DP thanks IAS for hosting. The simulations for this work have been carried on IASF Bologna cluster. We wish thank the financial support by INFN IS PD51 for NA visit in Bologna.   \\begin{figure} \\includegraphics[{width=85mm,height=100mm}]{CMBPSSZCL_70353_FGVariation_RUNGW.ps}\\\\ \\caption{Robustness test part 4. Cosmological and astrophysical parameters with different astrophysical signals with respect to the fiducial model considering also the tensor perturbations and the running of the spectral index. The black (solid) curve are the results of the analysis which considers a Poissonian contribution double its fiducial value. The red (dashed) curve are the results of the analysis which considers the 100, 143 and 217 GHz clustering terms double their fiducial value. The blue (dot-dashed) curve represent the analysis with a clustering term for the 217 GHz half its fiducial value whereas double its fiducial value for 100 and 143 GHz.} \\label{FGV_RG} \\end{figure}"
    },
    "1112/1112.1786_arXiv.txt": {
        "abstract": "In this paper we investigate gravitationally bound, spherically symmetric equilibrium configurations consisting of ordinary (polytropic) matter nonminimally coupled to an external chameleon scalar field. We show that this system has static, regular, asymptotically flat general relativistic solutions. The properties of these spherical configurations, such as total mass, distribution of matter, and size, depend strongly on the surrounding scalar field. The mass is found in terms of the parameter $\\sigma$, which is proportional to the central mass density of the ordinary matter. We perform a stability analysis of these spherical solutions and find an upper bound for $\\sigma$ where dynamical instability occurs. ",
        "introduction": "Scalar fields have been applied to a wide range of model building in cosmology and astrophysics. At the cosmological scale scalar fields have been used to model inflation and the late time acceleration due to dark energy. At the galactic scale scalar fields have been tried as models for dark matter (for a review with extensive references to scalar field models of dark energy and dark matter, see \\cite{sahni:2004}). At the astrophysical scale there are boson stars (see the reviews \\cite{Jetzer:1991jr,Schunck:2003kk}). Boson stars are stellar size or smaller, self-gravitating configurations of scalar fields. The scalar fields may be self-interacting or not, and one can include a fluid which interacts with the scalar field either only gravitationally \\cite{Henriques:1989ar,Henriques:1989ez} or through  direct coupling \\cite{Lee:1986tr}. The properties of these boson stars strongly depend on the type of scalar fields of which they are composed as well as how these fields interact with each other and the surrounding matter/fluid.  In scalar-tensor gravitational theories there is also the possibility of constructing mixed configurations of a scalar field plus ordinary matter. In such theories, a conformal transformation from the original Jordan frame to the Einstein frame is always accompanied by the appearance of a nonminimal coupling between ordinary matter and a scalar field, which leads to new features not present in general relativity. In particular, this nonminimal coupling between the matter of a neutron star and the scalar field can lead to the effect of spontaneous scalarization of the neutron star \\cite{Damour:1993hw,Damour:1996ke}. In turn, this leads to a redistribution of the scalar field inside the neutron star, which   has a significant influence on the process of the neutron star's collapse  \\cite{Novak:1997hw}.  In cosmology, this ability of scalar fields to interact nongravitationally with other fields has been used within the framework  of chameleon cosmology  \\cite{Khoury:2003aq,Khoury:2003rn,Brax:2004qh}, where a chameleon scalar field $\\phi$ interacts directly with ordinary matter through a conformal coupling of the form $e^{\\alpha \\phi}$.  The name ``chameleon scalar field'' was suggested in \\cite{Khoury:2003aq} since the properties of the scalar field, such as its  mass, depend sensitively on the environment in which the scalar field finds itself. The idea that the properties of a scalar field could be influenced by the environment/matter surrounding the scalar field was studied earlier, in particular, in the papers \\cite{Ellis:1989as,Mota:2003tc} where the interaction between matter and the scalar field was used to model a dependence of fundamental  coupling  constants on the local environment. The possibility of having scalar fields interact directly with ordinary matter has been used in~\\cite{Brax:2004qh} to describe the present accelerated expansion of the Universe. These cosmological models were developed further in \\cite{Farajollahi:2010pk}, where the authors choose a generalized expression for the nonminimal coupling between the scalar field and the matter having the form $f(\\phi) L_m$, with $f(\\phi)$ being an arbitrary coupling function and $L_m$ being the Lagrangian of ordinary matter. In this case, for certain parameters of the model, the numerical calculations of \\cite{Farajollahi:2010pk} showed that it was possible to describe the evolving Universe with a contraction phase, a subsequent bounce phase, and finally a phase of accelerated expansion. One can also find analytical solutions \\cite{Cannata:2010qd} within the framework of the cosmological model of \\cite{Farajollahi:2010pk}. These analytical solutions describe the evolution in time of the chameleon scalar field.  It is also possible to consider the effect of a chameleon scalar field at astrophysical scales rather than cosmological scales. In \\cite{Dzhunushaliev:2011ma} the  influence of a chameleon field on the structure of a star composed of a polytropic fluid nonminimally coupled to a chameleon scalar field was studied. It was found that for special choices of the scalar field potential energy and the coupling function $f(\\phi)$, one could obtain static, regular, asymptotically flat solutions in both relativistic and nonrelativistic limits. However, a preliminary stability analysis of the solutions found in \\cite{Dzhunushaliev:2011ma} indicated that they were not stable in that the matter of the solutions would either collapse or disperse. In \\cite{Folomeev:2011uj} it was shown that taking an isothermal fluid coupled to a chameleon scalar field led to both stable and unstable static solutions, and for the  unstable solutions,  the problem of the gravitational collapse of the nonrelativistic configurations has been considered.  In this paper we continue the study of the influence of a chameleon scalar field on the structure of polytropic spheres begun in \\cite{Dzhunushaliev:2011ma}.  We consider a relativistic, polytropic, spherically symmetric fluid embedded in an external  chameleon scalar field, which is distributed homogeneously and isotropically over the Universe.  When polytropic matter is placed in such a homogeneous background, it begins to interact with the background scalar field both gravitationally and through the nonminimal coupling. This results in a change of the inner structure of the polytropic star depending on the properties of the surrounding chameleon scalar field. In general, the scalar field will evolve over cosmological time, and this will lead to a time dependence for the term $f(\\phi) L_m$, which in turn can have a strong influence on the structure of the compact fluid configuration which is embedded in the scalar field. In this paper we restrict ourselves to only a simplified version of the problem, when we just study changes in the structure of a {\\it static} polytropic, fluid sphere embedded in an external {\\it dynamic} chameleon scalar field. For the time variation of the scalar field we take the time dependence suggested in \\cite{Cannata:2010qd} (see below in Sec.~\\ref{num_res}).  The paper is organized as follows: In Sec.~\\ref{gen_equations_cham_star} the general equations describing an equilibrium configuration consisting of a real scalar field coupled to a perfect polytropic fluid are derived. In Sec.~\\ref{static_sol_exp} these equations are written for a particular case when the coupling function $f(\\phi)$ is chosen to be the exponential one. In Sec.~\\ref{num_res} we give the results of the numerical calculations for  this choice of nonminimal coupling function. In Sec.~\\ref{stab_gen_exp} we address the issue of the stability of the obtained equilibrium configurations against small radial perturbations. Finally, in Sec.~\\ref{concl_exp} we summarize the main results. ",
        "conclusions": "\\label{concl_exp}  Based on the assumption that   time evolving chameleon  scalar fields may exist in the Universe, this paper studies the possible influence that such fields may have on the inner structure and stability of compact gravitating configurations consisting of ordinary matter and embedded in an external, homogeneous chameleon  scalar field. The ordinary matter is taken to have a polytropic equation of state which is parametrically given by expressions \\eqref{eqs_cham_star}-\\eqref{pressure_fluid_theta}. We choose a specific form for the nonminimal coupling between the ordinary matter and the scalar field given  by \\eqref{fun_f_exp}. This form was used in \\cite{Farajollahi:2010pk} within the framework of chameleon cosmology. Here we study the consequences  of the presence of a chameleon scalar field on the inner structure of compact astrophysical objects (i.e. ``stars\") rather than the cosmological consequences of chameleon fields. In Ref.~\\cite{Dzhunushaliev:2011ma} this type of gravitating configuration of a polytropic fluid nonminimally coupled to a scalar field was called a ``chameleon star.''  One of the main conclusions of the present work is that the characteristics of the chameleon stars  -- their masses, sizes, matter distribution -- depend strongly on  the parameters of the surrounding scalar field $\\phi$. Assuming that  this background scalar field changes with cosmological times, we predict that the characteristics of the stars embedded in the scalar field should also change. For example, we considered three different values of the background scalar field, $\\phi_{\\text{ext}}=-2, 0, 2$, which could represent the time evolved value of the field at different times according to the law \\eqref{cham_field_evol}. These three values could correspond to three different periods of time in the evolution of the Universe -- from the relatively early Universe to the future Universe, for example, already  near the big rip singularity. The results of the numerical calculations presented in Figs.~\\ref{energ_fig}-\\ref{fig_bin_energy_exp} (which cover the values of the polytropic index, $n=1.5$ and $n=3.0$) indicate that the characteristics of the configurations have a strong dependence on the value of the scalar field $\\phi_{\\text{ext}}$. In particular, from Fig.~\\ref{energ_fig} one finds that, at the center of the configurations considered, there can be either a larger or smaller concentration of the matter density compared to the case of a star without a scalar field. This would have an impact  on the rate of fusion reactions inside the stars, and thus on their lifetimes so that the life span of the star could be altered by the properties of the background chameleon scalar field surrounding the star. Some other speculations about possible physical applications of the chameleon star model can be found in Ref.~\\cite{Dzhunushaliev:2011ma}.  Another important consequence of the presence of the chameleon scalar field is its influence on the stability of the configurations studied here. From previous investigations of the stability of polytropic gravitating fluids it is known that for equilibrium configurations instabilities occur once one moves to the right of the first peak in the binding energy (or the first peak in the mass) \\cite{Tooper2}. A similar situation also occurs for purely boson stars (see the erratum from Ref.~\\cite{Gleiser:1988rq}). However, the dynamical stability analysis performed in Sec.~\\ref{stab_gen_exp} shows that the configurations considered here have their stable points shifted to the right of the first peak in the binding energy. These shifts are shown by the bold dots in Figs.~\\ref{fig_mass_exp} and \\ref{fig_bin_energy_exp}. Thus, in contrast to the usual polytropic and boson stars, the  expansion  of  the region of stability is in the direction where $\\sigma$ increases. It is obvious that this new effect is caused by the presence in the system of the nonminimal coupling between the ordinary matter and the scalar field.  The study of mixed configurations with scalar fields and ordinary matter sources is of interest since various types of scalar fields are believed to exist in the Universe. Most often, such scalar fields play a cosmological role~-- as a mechanism to drive inflation (the inflaton field), as dark matter, or as dark energy. Here we consider the possibility that one type of scalar field (a chameleon scalar field) might also have some influence on astrophysical objects such as stars."
    },
    "1112/1112.1265_arXiv.txt": {
        "abstract": "In order to mimic the phase changes in  the primordial Big Bang, several {\\it cosmological} solid-state experiments have been conceived, during the last decade, to investigate the spontaneous symmetry breaking in superconductors and superfluids cooled through their transition temperature. In one of such experiments the number of magnetic flux quanta spontaneously trapped in a superconducting loop was measured by means of a long Josephson tunnel junction built on top of the loop itself. We have analyzed this system and found a number of interesting features not occurring in the conventional case with simply connected electrodes. In particular, the fluxoid quantization results in a frustration of the Josephson phase, which, in turn, reduces the junction critical current. Further, the possible stable states of the system are obtained  by a self-consistent application of the principle of minimum energy.  The theoretical findings are supported by measurements on a number of  samples having different geometrical configuration. The experiments demonstrate that a very large signal-to-noise ratio can be achieved in the flux quanta detection. ",
        "introduction": "Long Josephson tunnel junctions (LJTJs) were traditionally used to investigate the physics of non-linear phenomena\\cite{barone}. In the last decade they have been employed to shed light on other fundamental concepts in physics such as the symmetry principles and how they are broken\\cite{PRLS,PRB08,gordeeva}. A recent experiment\\cite{PRB09} has demonstrated  spontaneous symmetry breaking during the superconducting phase transition of a metal ring and both fluxoids or antifluxoids can be trapped in the ring while it is cooled rapidly through the superconducting critical temperature. The basic phenomenon of quantization of magnetic flux in a multiply connected superconductor was suggested long time ago as one among several possible condensed matter \\textit{cosmological} experiments\\cite{zurek2} suitable to check the validity of the \\textit{causality} principle in the early Universe\\cite{kibble1}. In the experiment of Ref.\\cite{PRB09} the magnetic flux quanta are spontaneously trapped in the ring during its cooling through the transition temperature.  Much later at lower temperature when superconductivity is fully established, the number $0,\\pm 1, \\pm 2$ ... of flux quanta is registered as a function of the quench rate.  This can be done in a variety of ways, one of which is the detection of the induced persistent currents by the magnetic field modulation of the critical current of a planar LJTJ \\textit{built} on top of the ring. In these experiments, the quench rate can be varied over four decades. This allows for an accurate check of the theoretical predictions of the involved second-order phase transitions. This is of  interest within cosmology and of major importance for the physical understanding of many order-disorder processes. However, the working principles of that experiment had not yet been reported. The general task of this work is to study the static properties of a planar LJTJ for which at least one of the superconducting electrodes is multiply-connected, i.e., not every closed path can be transformed into a point. In the simplest case, one of the superconducting thin-film  stripes forming the LJTJ is shaped as a doubly-connected loop. This configuration is illustrated in Fig.\\ref{geometry}(a) in which the ring-shaped base electrode is in black, while the top electrode is in gray and the junction area is in white. The geometry of the loop is not critical to our discussion; however, a ring-shaped bottom electrode simplifies the analysis.  \\begin{figure}[tb] \\centering \\subfigure[ ]{\\includegraphics[width=7.0cm]{geometrya.png}} \\subfigure[ ]{\\includegraphics[width=7.0cm]{geometryb.png}}  \\caption{ a) Sketch of an in-line Josephson tunnel junction with a single doubly connected ring-shaped base electrode. $L_{loop}=L_1+L_2+L_b$. b) Sketch of an in-line Josephson tunnel junction with two doubly connected electrodes. The base electrode is in black, the top electrode is in gray, and the tunneling insulating layer is white. The dotted withe arrows indicate the direction of the circulating currents.\\label{geometry}}  \\end{figure}  \\noindent For the sake of generality, in our analysis, we will include an external flux $\\Phi_e$ linked to the loop by some externally applied field and the presence of an integer number $n$ of flux quanta trapped in the loop; altogether they induce a current $I_{cir}=(n\\Phi_0-\\Phi_e)/L_{loop}$ circulating clockwise in the loop and inversely proportional to its inductance, $L_{loop}$; in turns, the circulating current produces at the loop surface a radial magnetic field, $H_{rad} \\propto I_{cir}$, that adds to any external  field, $H_{app}$, applied in the loop plane. With no loss of generality, we will assume that the width, $W_t$, of the top film does not exceed the width, $W_b$, of the bottom one, $W_t\\leq W_b$ and, to simplify the analysis, we will also assume that both widths are much smaller than the mean radius, $R$, of the ring; in this narrow ring approximation, the current distribution in the ring and the surface radial field become radially independent\\cite{brandt}. A dc current $I$ is injected into the loop at an arbitrary point $O$ along the ring and is inductively split in the two loop arms before going through the LJTJ; let $\\alpha$ (1-$\\alpha$) be the fraction of the bias current $I$ diverted in the left (right) side of the loop. In principle, $\\alpha$ values outside the $[0,1]$ range are possible if the current $I$ would include also a persistent current $I_{cir}$ circulating in the loop: however, since the two currents are independent, we will treat them separately. Independently of the $\\alpha$ value, the bias current $I$ is extracted at one end of the junction via the top electrode. With the current entering and exiting at the junction extremities we have the well-known case of the so-called \\textit{in-line} configuration treated in the pioneering works on LJTJs soon after the discovery of the Josephson effect\\cite{ferrel,OS,stuehm,basa,radparvar85}.  \\noindent Throughout the paper we will limit our interest to LJTJs in the zero-voltage time-independent state; this can be achieved as far as the applied current $I$ is smaller than the junction critical current $I_c$. To further simplify the analysis, we assume that the Josephson current density $J_c$ is uniform over the barrier area and that the junction width $W_{}$ is smaller than the Josephson penetration depth $\\lambda _J\\equiv \\sqrt{\\Phi_0/ 2\\pi \\mu _{0}d_{e} J_{c}}$ setting the length unit of the physical processes occurring in the \\Jos \\jun  (here $\\Phi_0$ is the magnetic flux quantum, $\\mu_0$ the vacuum permeability and $d_e$ the junction magnetic thickness). The gauge-invariant phase difference $\\phi$ of the order parameters of the superconductors on each side of the tunnel barrier obeys the Josephson equations\\cite{joseph}:  \\begin{equation} \\label{jos} J_Z(X)=J_c \\sin \\phi(X), \\end{equation}  \\noindent and  \\begin{equation} \\label{gra} \\kappa {\\bf \\nabla} \\phi(X) = {\\bf H}\\times {\\bf \\hat{n}}, \\end{equation}  \\noindent in which $-{\\rm{L}}/2\\leq X \\leq {\\rm{L}}/2$ is a curvilinear coordinate and ${\\rm{L}}$ is the long dimension of the junction. The net current crossing the tunnel barrier is $I \\equiv  W_{} \\int_{-{\\rm{L}}/2}^{{\\rm{L}}/2} J_Z(X)dX$. The last equation states that the phase gradient is everywhere proportional to the local magnetic field ${\\bf H}$ and parallel to the barrier plane. Therefore, in the case of a curvilinear one-dimensional junction, a uniform external field applied in the junction plane has to be replaced by its radial component\\cite{PRB96}. $\\kappa \\equiv {\\Phi_0}/{2\\pi d_e \\mu _0}=J_c \\lambda_J^2$ has the dimension of a current ($\\kappa\\approx 2.5\\,$mA when $d_e \\approx 100\\,$nm, which is typical of all-Niobium \\Jos junctions) and ${\\bf \\hat{n}}$ is the versor normal to the insulating barrier separating the two superconducting electrodes. It is well known\\cite{ferrel,OS} that combining Eqs.(\\ref{jos}) and (\\ref{gra}) with the static Maxwell's equations, a static sine-Gordon equation is obtained that describe the behavior of a one-dimensional in-line LJTJ:  \\begin{equation} \\lambda_J^2 \\frac{d^2 \\phi}{d X^2} = \\sin \\phi(X). \\label{sG} \\end{equation}  \\noindent  Equation(\\ref{sG}) was first introduced in the analysis of \\textit{asymmetric} in-line LJTJs by Ferrel and Prange\\cite{ferrel} in 1963; few years later, Owen and Scalapino\\cite{OS} reported an extensive study of its analytical solutions for \\textit{symmetric} in-line \\Jos junctions (provided that ${\\rm{L}}\\geq\\pi \\lambda_J /2$). Ampere's law applied along the barrier perimeter requires that the magnetic fields at the two ends of the junctions differ by the amount of the enclosed current: $I=W_{} [H_Y ({\\rm{L}}/2) -H_Y (-{\\rm{L}}/2)]$. We remark that Eqs.(\\ref{jos}), (\\ref{gra}) and (\\ref{sG}) automatically satisfy the Ampere's law.  \\noindent As it is usually done in the modeling of Josephson interferometers, it is useful to divide the loop inductance $L_{loop}$ in three inductive paths characterized by positive coefficients $L_1$, $L_2$ and $L_b$ having units of inductance such that $L_{loop}=L_1+L_2+L_b$ so that any current $I_{cir}$ circulating around the loop \\textit{sees} them in series; furthermore, $\\alpha=1$ ($\\alpha=0$) in the limit of $L_2>>L_1+L_b$ ($L_2<<L_1+L_b$). More specifically, $L_{1,2}$ is the angular fraction $\\vartheta_{1,2}/2\\pi$ of the inductance for an isolated superconducting narrow ($W_b<<R$) ring\\cite{brandt} $L_{ring}=\\mu_0 R (\\ln 8R/W_b - 2 + \\ln 4)$ and $L_b$, that should not be mistaken as the inductance of the LJTJ, is given by the junction physical length  ${\\rm{L}}$ times the inductance per unit length of the junction bottom strip-line $\\mathcal{L}_b$, related to the magnetic energy stored within a London penetration distance of its surface. If the bottom and top superconducting films have, respectively, thickness $d_{b}$ and $d_{t}$ and bulk London penetration depths $\\lambda_{Lb}$ and $\\lambda_{Lt}$, then, in presence of a quasi-static magnetic field, their effective penetration depth\\cite{wei} is $\\lambda_{b,t} =\\lambda_{Lb,t} \\tanh (d_{b,t}/2 \\lambda_{Lb,t})$ which reduces to $\\lambda_{Lb,t}$ in the case of thick superconducting films ($d_{b,t} >5 \\lambda_{Lb,t}$).  The magnetic penetration $d_e$ of a tunnel barrier with negligible height $t_{ox}<< d_{b,t}$ is\\cite{wei} $d_e \\simeq \\lambda_{b} + \\lambda_{t}$. Insofar as the width $W_b$ is much larger\\cite{swihart,orlando} than the strip-line magnetic thickness $d_e$, then $\\mathcal{L}_b \\simeq \\mu_0 \\lambda_{b}/W_b$; this expression also takes into account the kinetic inductance due to the motion of the superelectrons. In the wide-strip approximation, most of the magnetic energy is confined in the region between the plates and the fringing field can be ignored; as the strip width becomes narrower, the fringe field effects become more important and may dominate if $W_b$ and $d_e$ are comparable\\cite{chang}. It is worth pointing out that, since, in all practical cases, the width of the loop is much larger than the London penetration depth, then $\\mathcal{L}_b$ is considerably smaller than the inductance per unit length along the ring $L_{ring}/2\\pi R$; this is due to the presence of a counter electrode acting as a superconducting ground plane\\cite{chang,vanDuzer}. Similarly, we introduce $\\mathcal{L}_t \\simeq \\mu_0 \\lambda_{t}/W_t$, the inductance per unit length along the current direction of the top plate. Since the electrodes have different widths and penetration depths, in general, $\\mathcal{L}_t \\neq \\mathcal{L}_b$. In Section III we will show that for our high quality all-Niobium LJTJs, having the base electrode thinner and wider than the top one, we found $\\mathcal{L}_t \\approx 3 \\mathcal{L}_b$. According to the theory of the two-conductor transmission lines\\cite{kraus}, the inductance per unit length, $\\mathcal{L}_J$, of a LJTJ, seen as a transmission line structure, is simply obtained as the sum of the inductances/unit lengths of the bottom and top stripes, i.e.,  \\begin{equation} \\mathcal{L}_J=\\mathcal{L}_b + \\mathcal{L}_t=\\mu_0 \\left( \\frac{\\lambda_{b}}{W_b}+ \\frac{\\lambda_{t}}{W_t} \\right). \\label{indu} \\end{equation}  \\noindent Historically, the boundary conditions for Eq.(\\ref{sG}) were derived under the implicit assumption that $W_b = W_t=W_{}$, so that\\cite{scott76,vanDuzer} $\\mathcal{L}_J=\\mu_0 d_e/W_{}$. However these conditions are not fulfilled in real samples, especially for window-type LJTJs used nowadays whose electrodes have quite different widths; typically, $W_b>W_t>W_{}$.  \\noindent The paper is organized in the following way. In Section II we will overcome this limitation by extending the existing theoretical model\\cite{OS} to LJTJs having different electrodes widths, strictly speaking, different inductances per unit length. At the same time we will derive the most general boundary conditions for Eq.(\\ref{sG}) needed to correctly describe any self-field effect in LJTJs. Next we will focus on the specific case of LJTJs with doubly connected electrode(s). Later on we will consider the consequences of the fluxoid quantization and energy minimization principles. In the next section we will describe our experimental setup and our samples; in addition we will present their magnetic diffraction patterns and discuss how the experimental findings can be unambiguously interpreted in term of our  modeling.  Finally, the conclusions will be drawn in Section IV. ",
        "conclusions": "\\noindent In this paper we have revisited the theory of the self-field effects that characterize the long Josephson tunnel junctions and made them interesting for the investigation of non-linear phenomena. Our analysis goes beyond the previous works in  two ways: (i) it takes into account the different inductances per unit length of the electrodes forming the junction and (ii) it provides the boundary conditions for the most general junction biasing configuration. We applied the theory to the specific case of long Josephson tunnel junctions with not simply connected electrodes. Apart from their intriguing physical properties, the interest for LJTJs built on a superconducting loop stems from the fact that they were successfully used to detect trapped fluxoids in a {\\it cosmological} experiment aimed to study the spontaneous defect production during the fast quenching of a superconducting loop through its normal-to superconducting transition temperature\\cite{PRB09}. We found that the single-valuedness of the phase $\\phi_1$ of the order parameter of the bottom (or top) superconducting electrode (fluxoid quantization) gives raise to a variety of unexpected non-linear phenomena when coupled to the sine-Gordon equation for the phase difference $\\phi_2-\\phi_1$ of the order parameters in the superconductors on each side of the tunnel barrier. The principle of energy minimization was also invoked to determine the possible states of the system. We have focused on static phenomena such as the reduction of the junction critical current and its dependence on magnetic fields applied in and out of the loop plane. Nevertheless, also the dynamic properties, such as the propagation of non-linear waves, are expected to be drastically affected by the doubly connected electrode(s). Our experiments unambiguously corroborate the analytical findings and provide hints to implement the modeling. Future work should go in the direction of investigating the consequences of the fluxoid quantization for small and intermediate length Josephson tunnel junctions and on the dynamic properties of long junctions."
    },
    "1112/1112.1862_arXiv.txt": {
        "abstract": "The relation between redshift and the CMB temperature, $T_{\\rm CMB}(z)=T_0(1+z)$ is a key prediction of standard cosmology, but is violated in many non-standard models. Constraining possible deviations to this law is an effective way to test the $\\Lambda$CDM paradigm and search for hints of new physics. We present state-of-the-art constraints, using both direct and indirect measurements. In particular, we point out that in models where photons can be created or destroyed, not only does the temperature-redshift relation change, but so does the distance duality relation, and these departures from the standard behaviour are related, providing us with an opportunity to improve constraints. We show that current datasets limit possible deviations of the form $T_{\\rm CMB}(z)=T_0(1+z)^{1-\\beta}$ to be $\\beta=0.004\\pm0.016$ up to a redshift $z\\sim 3$. We also discuss how, with the next generation of space and ground-based experiments, these constraints can be improved by more than one order of magnitude. ",
        "introduction": "Introduction}  Cosmology and particle physics are presently experiencing a truly exciting period. On the one hand, both have remarkably successful standard models, which are in agreement with a plethora of experimental and observational data. On the other hand, there are also strong hints that neither of these models is complete, and that new physics may be there, within the reach of the next generation of probes.  There are three compelling and firmly established observational facts that the standard model of particle physics fails to account for: neutrino masses, the existence of dark matter, and the size of the baryon asymmetry of the Universe. For each of these, the model makes very specific statements, failing however to reproduce the experimental evidence. It is precisely our confidence in the model and our ability to calculate its consequences that lead us to the conclusion that it is incomplete, and new phenomena must be anticipated. This is, of course, the reason for the LHC project.  Similarly, the last decade saw the emergence of the so-called concordance model of cosmology. This can reproduce all the available observations with only a small number of parameters, but also requires that about $96\\%$ of the content of the universe is in a form that has never been seen in the laboratory (and is only known indirectly from its gravitational properties). It is thought that dark matter is a subdominant part of this, while the dominant one is an even more mysterious component usually called dark energy.  In this context, it is important to identify laboratory or astrophysical probes that can give us more information about the nature and properties of this still unknown physics. In this work we will discuss one such probe---the temperature-redshift relation---, and lay the foundations for exploring its cosmological implications.  One of the most precise measurements in cosmology is the intensity spectrum of the cosmic microwave background radiation: the COBE-FIRAS experiment revealed a very precise black-body spectrum \\cite{Mather}. However, this measurement tells us nothing about the behaviour of the Cosmic Microwave Background (CMB) at non-zero redshift. If the expansion of the Universe is adiabatic and the CMB spectrum was a black-body at the time it originated, this shape will be preserved with its temperature evolving as $T(z)=T_0(1+z)$. This is a robust prediction of standard cosmology, but it is violated in many non-standard models, including string theory motivated scenarios where photons mix with other particles such as axions (see~\\cite{JaeckRing} for a recent review), and those where dimensionless couplings like the fine-structure constant vary \\cite{RoySoc}.  A few measurements of $T(z)$ already exist, but the currently large uncertainties do not allow for strong constraints on the underlying models to be set. However, with future datasets this will become a competitive probe. It is therefore timely to discuss what these measurements can tell us about the underlying cosmological paradigms.  At low redshifts, say $z<1$, the $T(z)$ relation can be measured via the Sunyaev-Zeldovich (SZ) effect towards galaxy clusters. This method was applied to ground-based CMB observations \\cite{Battistelli,Luzzi}, which demonstrated its potential. With a new generation of ground experiments becoming operational and a forthcoming all-sky survey of SZ clusters to be carried out by Planck \\cite{Horellou}, the potential of this method will come to fruition. At higher redshifts, $z>1$, $T(z)$ can be evaluated from the analysis of quasar absorption line spectra which show atomic and/or ionic fine structure levels excited by the photon absorption of the CMB radiation \\cite{Srianand}. (The CMB is an important source of excitation for species with transitions in the sub-millimeter range.) Although the suggestion is more than four decades old, measurements (as opposed to upper bounds) were only obtained in the last decade, and the best ones so far still have errors at the ten percent level \\cite{Noterdaeme}.  Here we will study these issues in detail, but we will also place them in a wider context. For example, in models where photons can be created or destroyed, not only does the temperature-redshift relation vary, but so does the distance duality relation (also known as the Etherington relation \\cite{Etherington1}), and these two different departures from the standard behaviour are quantitatively related. One issue that has been overlooked so far is that in such models, where photon number is not conserved, this relation between $T(z)$ and distance duality provides us with an opportunity to improve constraints. By combining data from different observations one not only reduces the statistical uncertainties on underlying phenomenological parameters but, given the different nature of both observational datasets, one also has a much better control over possible systematics.  We therefore discuss in detail the origin of the above relation, as it can be a unique consistency test for the standard paradigm and, at the same time, a valuable tool for probing new physics beyond the standard model. We also study further imprints of these models in the CMB, and present forecasts for improvements that Planck, as well as planned Baryon Acoustic Oscillations (BAO) missions and spectrographs planned for the VLT and the E-ELT, will soon make possible. Last but not least, we derive the strongest constraints to date on deviations of these relations from their standard behaviour, and quantify the improvements to be expected from the aforementioned forthcoming experiments. ",
        "conclusions": ""
    },
    "1112/1112.4992_arXiv.txt": {
        "abstract": "{Through an optical campaign performed at 4 telescopes located in the northern and the southern hemispheres, plus archival data from two on-line sky surveys, we have obtained optical spectroscopy for 28 counterparts of unclassified or poorly studied hard X-ray emitting objects detected with Swift/BAT and listed in the 39 months Palermo Swift/BAT hard X-ray catalogue. We have been able to pinpoint the optical counterpart of these high energy sources by means of X-ray observations taken with Swift/XRT or XMM which allowed us to restrict the positional uncertainty from few arcmin to few arcsec; satellite data also provided information on the X-ray spectra of these objects. We find that 7 sources in our sample are Type 1 AGN while 20 are Type 2 AGN, with their redshifts lying between 0.009 and 0.075; the remaining object is a Galactic cataclysmic variable (CV). In this work we provide optical information for all  28 sources and the results of the soft X-ray analysis of 3 out of 5 AGN observed with XMM/Newton.}  \\FullConference{The Extreme and Variable High Energy Sky - extremesky2011,\\\\ September 19-23, 2011\\\\ Chia Laguna (Cagliari), Italy}   \\begin{document} ",
        "introduction": "The {\\it Swift} mission was designed to study cosmic gamma-ray bursts (GRBs) in a multiwavelength context ([7]), but it is also able to perform dedicated X-ray and UV-optical observations as well as surveys of the entire sky. {\\it Swift} carries three instruments, i.e. the burst alert telescope (BAT; [1]), the X-ray telescope (XRT; [2]) and the ultraviolet/optical telescope (UVOT; [13]) and therefore can detect and follow up X-ray emitting objects over a wide range of wavelengths.  In particular, BAT, the high energy instrument, is a coded mask detector operating with good sensitivity in the energy range 14--195 keV over a field of view of 1.4 sr with a point source location accuracy of $1^{\\prime}-4^{\\prime}$ ([7]) depending on the source intensity. Its sensitivity is estimated to be  $\\sim$1 mCrab at high Galactic latitudes and $\\sim$3 mCrab over the Galactic plane.  This instrument is not only able to detect GRBs, but also to perform  highly sensitive hard X-ray surveys (e.g. [4], [5], [14]). In particular, the BAT surveys allow the study of the extragalactic X-ray sky, and the observation of many absorbed AGNs which are often missed by lower energy instruments. Quantifying the number of such absorbed objects, especially at low redshifts, is very important if one wants to understand the accretion \\mbox{mechanisms} at work in AGNs and to estimate the contribution of all AGN  to the cosmic X--ray background ([3]). However, many of the objects listed in the BAT surveys are still unclassified or poorly studied and hence they need optical follow up work to be fully characterized.  For this work we have selected from the 39 months Palermo {\\it Swift}/BAT AGN survey ([4]), a group of objects (28 in total) either without optical identification, or not well studied or without published optical spectra. Following the method applied by [8] and references therein or [11] for the optical spectroscopic follow up work of unidentified {\\it INTEGRAL} and/or BAT sources, we determine the nature of these 28 selected objects by means of X-ray observations (to pinpoint the likely X-ray counterpart) and optical measurements (to provide the source classification).  In the following sections we show the results obtained with our optical spectroscopic campaigns and we discuss in detail the results of the X-ray analysis of 3 out of 5 objects observed by XMM/Newton. ",
        "conclusions": "With this work we have been able to either give or confirm or correct  the optical classification of 28 {\\it Swift} sources belonging to the Palermo 39 months {\\it Swift} catalogue (see also [12] for details). This was achieved through a multisite observational campaign in Europe, South Africa, Central and South America.  We found that our sample is composed of 27 AGNs (7 of Type 1 and 20 of Type 2), with redshifts between 0.009 and 0.075, and 1 CV. Among these sources we found some peculiar objects, such as 3 likely LINERs and 1 narrow line seyfert 1.  The X-ray spectral analysis of 3 out of 5 sources observed with XMM-Newton shows a complex best-fit model with an absorbed power law component as a primary emission model; all 3 require an extra component at low energies to fit an emission excess below few keV, while only two display iron line emission features.  This work shows the importance of optical spectroscopic follow up observations for sources discovered by hard X-ray  surveys and either unclassified or poorly studied. By increasing the number of the identifications in hard X-ray catalogues, it is possible to perform more reliable statistical studies as multiwavelength characterization of the sources, thus allowing a better understanding of  the physical processes that drive the powerful AGN detected."
    },
    "1112/1112.1053_arXiv.txt": {
        "abstract": "The IceCube Neutrino Observatory is the world's largest high energy neutrino telescope, using the Antarctic ice cap as a Cherenkov detector medium. DeepCore, the low energy extension to IceCube, is an infill array with a fiducial volume of around 30 MTon in the deepest, clearest ice, aiming for an energy threshold as low as 10 GeV and extending IceCube's sensitivity to indirect dark matter searches and atmospheric neutrino oscillation physics. We will discuss the analysis of the first year of DeepCore data, as well as ideas for a further extension of the particle physics program in the ice with a future PINGU detector. ",
        "introduction": "\\label{sec:intro}  The IceCube neutrino telescope, now fully operational at depths of 1450-2450 m below the surface of the Antarctic ice cap, was designed to detect high energy neutrinos from astrophysical accelerators of cosmic rays.  Although the energy threshold of a large volume neutrino detector is not a sharp function, the original IceCube design focused on efficiency for neutrinos at TeV energies and above.  Recently, the IceCube collaboration decided to augment the response of the detector at lower energies with the addition of DeepCore, a fully contained subarray aimed at improving the sensitivity of IceCube to neutrinos with energies in the range of 10's of GeV to a few hundred GeV.  This energy range is of interest for several topics related to particle physics, including measurements of neutrino oscillations and searches for neutrinos produced in the annihilation or decay of dark matter.  DeepCore consists of an additional eight strings of photosensors (Digital Optical Modules, or DOMs) comprising 10'' Hamamatsu photomultiplier tubes and associated data acquisition electronics housed in standard IceCube glass pressure vessels.  For most of the DeepCore DOMs, the standard IceCube R7081 PMTs were replaced with 7081MOD PMTs with Hamamatsu's new super-bialkali photocathode.  These PMTs provide approximately 35\\% higher quantum efficiency (averaged over the detected Cherenkov spectrum) than the standard bialkali PMTs.  Sited at the bottom center of the IceCube array, DeepCore benefits from the high optical quality of the ice at depths of 2100-2450 m, with an attenuation length of approximately 50 m in the blue wavelengths at which most Cherenkov photons are detected in ice. DeepCore also benefits from the ability of the standard IceCube sensors to detect atmospheric muons penetrating the ice from cosmic ray air showers above the detector, allowing substantial reduction in the background rate by vetoing events where traces of penetrating muons are seen.  Each DeepCore string bears 50 DOMs in the fiducial region, with an additional 10 DOMs deployed at shallower depths to improve the vetoing efficiency for steeply vertical muons.  In addition to the new DeepCore strings, the DeepCore fiducial volume for analysis includes 12 standard IceCube strings, chosen so that the fiducial region is shielded on all sides by a veto region consisting of three rows of standard IceCube strings, as shown in Fig.~\\ref{fig:layout}.  The random noise rate of IceCube DOMs is quite low (around 500 Hz, on average) due to the low temperatures and radiopurity of the ice cap. This permits DeepCore to be operated with a very low trigger threshold, demanding that 3 DOMs within the DeepCore fiducial region detect light in ``local coincidence'' within a period of no more than 2500 ns.  The local coincidence criterion counts DOMs as being hit (i.e., having detected light) only if one of the four neighboring DOMs on a string (two above and two below) also registers a hit within $\\pm1 \\;\\mu$s. Most of the resulting 185 Hz of triggers are due to stray light from muons which simultaneously satisfy the main IceCube trigger condition of 8 DOMs hit in local coincidence within 5 $\\mu$s, but the DeepCore trigger contributes an additional (exclusive) rate of around 10 Hz.  \\begin{figure} \\begin{center} \\includegraphics[width=0.8\\columnwidth]{icecube_deepcore_pingu_koskinen_new.pdf} \\caption{Schematic layout of DeepCore within IceCube.  The shaded region indicates the fiducal volume of DeepCore, at the bottom center of IceCube, plus the extra veto cap of DOMs deployed at shallower depths to reinforce the veto against vertically-downgoing atmospheric muons.  This schematic depicts both the DeepCore configuration used in 2010, when 79 IceCube strings were operational, and the final DeepCore layout and fiducial region used in the 2011 run. \\label{fig:layout}} \\end{center} \\end{figure}  The vast majority of the events which trigger DeepCore, irrespective of whether they also trigger IceCube, are due to either penetrating atmospheric muons or random coincidences of dark noise.  Immediately after data acquisition, events triggering DeepCore are subjected to an online data rejection algorithm which calculates a characteristic time and location for the activity observed in the DeepCore fiducial region, as an initial estimate of the putative neutrino vertex.  The estimated location is the average position of the hit DOMs, and the time is determined by subtracting the time of flight $dn/c$ of an unscattered photon emitted from that location from the observed arrival time of the first photon to hit each DOM.  After outliers due to dark noise or scattered light are removed, the average inferred emission time is used as the estimated time of the underlying physics event.  Based on this estimated time and location, every locally coincident hit recorded in the veto region prior to the vertex time is examined to determine whether it lies on the light cone connecting it with the estimated event vertex.  The distributions of the inferred speed required to connect hits in the veto region to the DeepCore vertex, for both simulated atmospheric muons and simulated neutrinos, is shown in Fig.~\\ref{fig:veto}; positive speeds indicate hits occuring in the veto region prior to the DeepCore vertex time.  If any hits are found with inferred speeds between +0.25 and +0.4 m/ns, the event is rejected as being most likely due to an atmospheric muon.  This algorithm reduces the event rate by more than two orders of magnitude, to 18 Hz, while retaining over 99\\% of simulated triggered events due to neutrinos interacting within the fiducial volume.  Additional background rejection criteria are applied offline, depending on the goals of each physics analysis making use of these data.  \\begin{figure} \\includegraphics[width=\\columnwidth]{ParticleSpeedProbabilities_v22.pdf} \\caption{Distribution of probabilities of observing hits leading to a given inferred particle speed, for simulated atmospheric muons (dashed line) and atmospheric neutrinos (solid). Positive speeds indicate activity in the veto region prior to that in the DeepCore volume, and a peak around $c =$ 0.3 m/ns is visible for penetrating muons.  The integral of each distribution corresponds to the mean number of hits observed in the veto region for the given class of events. \\label{fig:veto}} \\end{figure}  The effective volume of the DeepCore detector for detection low energy muon neutrinos, accounting for this online data filter, is shown in Fig.~\\ref{fig:nuMuVolume}.  It should be stressed that this effective volume curve does \\emph{not} include losses due to later background rejection or event quality criteria.  The contribution of DeepCore to low energy analysis is evident in the fact that despite its relatively small geometric volume, around 3\\% that of IceCube, the overall sample of neutrino events below 100 GeV consists primarily of those detected by DeepCore.  This energy range is of considerable interest for several topics in particle physics, including searches for dark matter and measurements of neutrino oscillations.  While DeepCore does not have a sharp energy threshold, it retains around 7 megatons of effective volume at energies as low as 10 GeV.  Further details regarding DeepCore's instrumentation and performance are available in Ref.~\\cite{Collaboration:2011ym}.  \\begin{figure} \\includegraphics[width=\\columnwidth]{effectiveVolume_IC86_numu_GENIE_effVolumes_logScale_prelim.pdf} \\caption{Effective volume of DeepCore for muon neutrinos at trigger level (solid) and after application of the online veto algorithm described in the text (dot-dashed line).  The effective volume of IceCube as originally proposed is shown for comparison. \\label{fig:nuMuVolume}} \\end{figure} ",
        "conclusions": "The effectiveness of IceCube at energies below 100 GeV has been significantly enhanced by the addition of DeepCore, which extends IceCube's reach to energies of 10's of GeV.  This range is of interest for observations of neutrino oscillations, as well as searches for dark matter.  As a first step toward these studies, we have observed a significant sample of atmospheric neutrino-induced cascades, enabled by the ability of the IceCube detector to identify and veto atmospheric muons penetrating to the DeepCore volume.  We are also investigating the potential for a further reduction in the energy threshold of IceCube with an additional extension known as PINGU, which could extend IceCube's reach to energies as low as a few GeV"
    },
    "1112/1112.5994_arXiv.txt": {
        "abstract": "Spectroscopic monitoring with Mercator-HERMES over the past two years reveals that MWC~314 is a massive binary system composed of an early B-type primary LBV star and a less-luminous supergiant companion. We determine an orbital period $\\rm P_{\\rm orb}$ of 60.85 d from optical S~{\\sc ii} and Ne~{\\sc i} absorption lines observed in this single-lined spectroscopic binary. We find an orbital eccentricity of $e$=0.26, and a large amplitude of the radial velocity curve of 80.6 $\\rm km\\,s^{-1}$. The ASAS $V$ light-curve during our spectroscopic monitoring reveals two brightness minima ($\\Delta$$V$$\\simeq$$0^{\\rm m}$.1) over the orbital period due to partial eclipses at an orbital inclination angle of $\\sim$70$\\deg$. We find a clear correlation between the orbital phases and the detailed shapes of optical and near-IR P Cygni-type line profiles of He~{\\sc i}, Si~{\\sc ii}, and double- or triple-peaked stationary cores of prominent Fe~{\\sc ii} emission lines. A preliminary 3-D radiative transfer model computed with {\\sc Wind3D} shows that the periodic P Cygni line profile variability results from an asymmetric common-envelope wind with enhanced density (or line opacity) in the vicinity of the LBV primary. The variable orientation of the inner LBV wind region due to the orbital motion produces variable P Cygni line profiles (with wind velocities of $\\sim$200~$\\rm km\\,s^{-1}$) between orbital phases $\\phi$ = 0.65 to 0.85, while weak inverse P Cygni profiles are observed half an orbital period later around $\\phi$ = 0.15 to 0.35. We do not observe optical or near-IR He~{\\sc ii}, C~{\\sc iii}, and Si~{\\sc iii} lines, signaling that the LBV's spectral type is later than B0. Detailed modeling of the asymmetrical wind properties of massive binary MWC~314 provides important new physical information about the most luminous hot (binary) stars such as $\\eta$~Carinae. ",
        "introduction": "MWC~314 (V1492~Aql; BD$+$14$\\deg$3887; $V$=$9^{\\rm m}$.9) is a candidate Luminous Blue Variable (LBV) previously proposed to be one of the most luminous stars of the Galaxy by \\citet{miro1} with log(${L}_{\\star}/{L}_{\\odot}$)$\\simeq$6.1$\\pm$0.3, $T_{\\rm eff}$$\\simeq$25 to 30 kK, and $\\dot{M}$$\\sim$ $\\rm 3\\,10^{-5}$ $\\rm M_{\\odot}\\,{yr}^{-1}$. More recently, \\citet{mura1} found that the star is a binary system with an orbital period of $\\sim$30 d using optical spectra, however without determining other orbital parameters. We therefore observed 12 high-resolution spectra over the past two years with Mercator-HERMES ($R$=80,000) at La Palma (Spain). HERMES is a high-efficiency {\\'e}chelle spectrograph covering 420 nm to 900 nm \\citep{rask1}. We observed the spectra of MWC~314 with large SNR$\\sim$100 for accurate radial velocity (RV) measurements and detailed line profile studies. On 5 \\& 9 Sep 2009, and on 17 \\& 20 Mar 2011 we also observed two spectra within 5 d to investigate possible short-time spectroscopic variability in MWC~314 \\citep[see also][]{lobe1}. ",
        "conclusions": "Based on long-term spectroscopic monitoring with Mercator-HERMES we confirm the binarity of candidate LBV MWC 314, first conjectured by Muratorio et al. in 2008. However, we determine an orbital period of 60.85 d (twice longer than $P_{\\rm orb}$$\\sim$1 m they proposed) from an accurate solution of the RV-curve. We also compute an orbital eccentricity of 0.26 with LBV periastron passage oriented almost towards the observer. The visual brightness curve reveals two unequal minima signaling partial eclipses at an orbital inclination angle of $\\sim$70$\\deg$ in the plane of the sky. We also confirm the LBV character of MWC 314 with strong P Cygni-type line profiles observed during the orbital phases of fastest approach around periastron passage. A 3-D radiative transfer model we compute for the wind of MWC~314 shows that the P Cyg profiles result from enhanced LBV wind density inside a circumbinary expanding wind envelope."
    },
    "1112/1112.2785_arXiv.txt": {
        "abstract": "In this paper we study the non-Gaussian features of the primordial fluctuations in loop quantum cosmology with the inverse volume corrections. The detailed analysis is performed in the single field slow-roll inflationary models. However, our results reflect the universal characteristics of bispectrum in loop quantum cosmology. The main corrections to the scalar bispectrum come from two aspects: one is the modifications to the standard Bunch-Davies vacuum, the other is the corrections to the background dependent variables, such as slow-roll parameters. Our calculations show that the loop quantum corrections make $f_{{\\rm NL}}$ of the inflationary models increase $0.1\\%$. Moreover, we find that two new shapes of non-Gaussian signal arise, which we name $\\mathcal F_{1}$ and $\\mathcal F_{2}$. The former gives a unique loop quantum feature which is less correlated with the local, equilateral and single types, while the latter is highly correlated with the local one. ",
        "introduction": "As a non-perturbative and background-independent theory, Loop Quantum Gravity (LQG) \\cite{rovelli04,ashtekar04,thiemann07} has achieved great successes in past years: derivations of the quantized area and volume operators \\cite{rovelli95,ashtekar97,ashtekar98,thiemann98a}, calculations of black holes entropy \\cite{rovelli96} and Loop Quantum Cosmology(LQC) \\cite{bojowald08a}, {\\it etc.} And the nonperturbative quantization procedure of LQG is also valid for a more general class of four-dimensional metric theories of gravity~\\cite{zhang11a,zhang11b,zhang11c}. As an example of LQG, LQC gives a quantization scheme of LQG for a symmetry-reduced model in the homogeneous and isotropic Friedmann-Lema\\^ itre-Robertson-Walker universe. The discrete spacetime geometry in LQC scenarios predicts a non-singular bouncing universe in some simplified models, which satisfies most of the astronomical and cosmological observational constraints. Although the quantum correction effects are being diluted with the expansion of our universe, it remains present in a weaker form, especially on/near the super-horizon scales.  Recently, the gauge invariant cosmological perturbation theory has been systematically constructed in \\cite{bojowald06,bojowald08,bojowald09} for inverse volume corrections and in \\cite{wu10,Cailleteau11} for holonomy corrections. Some relevant applications have been considered in~\\cite{bojowald11a,bojowald11b,bojowald11c,zhu11}. The inverse volume corrections and the holonomy corrections are two main quantum corrections in LQC. The inverse volume corrections come from the quantization of the inverse of the volume operator in LQG. The inverse volumes exist in the Hamiltonian constraint of gravity and the usual matter Hamiltonian, especially in the kinematic terms. Since the volume can be taken the value zero, there does not exist well defined inverses of the volume operator. Fortunately, with the Thiemann trick~\\cite{thiemann98}, we can construct well defined inverse volume operators, which bring the quantum corrections. While the holonomy corrections arise from the loop quantization based on the holonomies instead of the direct connection. The holonomy corrections become important when the energy scale of our Universe approaches the Planck one. Both the modifications to the scalar and tensor primordial power spectra from the inverse volume corrections are carefully investigated by the authors of \\cite{bojowald11a}. Their results show that the inverse volume corrections could give rise to the enhancement of the power spectra on the large scales, i.e., a red-tilt one. However, some other mechanisms such as the non-commutative geometry or string theory~\\cite{tsujikawa03,piao04,calcagni04}, could also lead to similar features. Therefore to seek for signature of loop quantum cosmology, the study for the non-Gaussianity features in loop quantum cosmology is necessary.   Because the primordial non-Gaussianities are quite helpful to distinguish inflationary models, so far a lot of papers have been devoted to studying the non-Gaussianities in different inflation models, see the relevant references in the nice reviews \\cite{Bartolo:2004if,Chen:2010xk}. Inspired by the studies of \\cite{bojowald11a}, in this paper we mainly consider the non-Gaussianities from the inverse volume corrections in LQC. The reason for consider inverse volume corrections only is as following. We denote $\\delta_{\\rm inv}$ as the correction term coming from the inverse volume operator and $\\delta_{{\\rm hol}}$ as the correction term from the holonomy corrections. We can estimate the inverse volume correction as\\cite{bojowald11c} \\begin{equation}\\label{compare} \\delta_{{\\rm inv}}\\sim\\left(\\frac{8\\pi}{3}\\frac{\\rho}{\\rho_{{\\rm Pl}}}\\delta_{{\\rm hol}}^{-1}\\right)^2\\;, \\end{equation} where the Planck density $\\rho_{{\\rm Pl}}$ is assumed as the quantum gravity scale. From the above expression, we can see that the inverse volume corrections behave very differently from what is normally expected for quantum gravity. For low densities, the holonomy corrections is small, but the inverse volume one may still be large because they are magnified by the inverse of $\\delta_{{\\rm hol}}$. For an example, the small holonomy corrections of size $\\delta_{{\\rm hol}}<10^{-6}$ then requires the inverse volume corrections larger than $\\delta_{{\\rm inv}}>10^{-6}$ even at scale $\\rho\\approx10^{-9}\\rho_{\\rm Pl}$. These novel features make the investigations on the inverse volume corrections more interesting than the former at sub-Planckian inflationary scales. So we only consider the inverse volume correction in this work.  Explicitly, in the perturbation theory in LQC, the inverse volume operator can be captured by a correction function such as $\\bar{\\alpha}\\simeq1+\\alpha_0\\delta_{\\rm inv}\\simeq 1+\\alpha_0(a_{\\rm inv}/a)^{\\sigma}$, where $a$ is the scale factor of the FLRW universe and $a_{\\rm inv}$ is introduced to describe the characteristic scale of the inverse volume correction, which is not the Planck one in general. When $a_{\\rm inv}/{a}\\ll 1$, we can ignore the correction term. However, if $a_{inv}/a \\lesssim 1$ during inflation, one cannot neglect inverse volume corrections. In this case, the correction approximates $\\alpha_0\\delta_{\\rm inv}(k)\\approx\\delta(k_0)(k_0/k)^{\\sigma}$, where $k$ and $k_0$ are, respectively, the considered perturbation wave number and some characteristic number involved in the inverse volume correction. In addition, many works about LQC imply that $\\sigma\\in[0,6]$ \\cite{bojowald11b}. From the form of the inverse volume correction, we can see that a small sigma corresponds to a small the inverse volume correct, vice versa. Therefore, as an example, following \\cite{bojowald11b}, we take $\\sigma=2$ in this paper. For other $\\sigma$ values the behavior will be similar. Furthermore, in terms of spherical multiples the wave number could be expressed as $k\\approx10^{-4}hl$, with $h$ representing for the reduced Hubble parameter $h=0.7$ and $l$ for the spherical multiples. In the typical linear regime of Cosmic Microwave Background (CMB), the multiples $l$ range in $2<l<1000$ or more. Given all the mentioned observations, we expect some new features in the non-Gaussianities will arise. The purpose of this work is to investigate the characterized sizes and shapes of bispectra in LQC scenarios.   Note that the bispectra for the single field slow-roll inflationary model have been calculated in the papers~\\cite{arXiv:0209156,maldacena03,Seery:2005wm,Chen:2006nt,Chen:2006xjb}. For simplicity and comparison, we study the simplest single field slow-roll inflationary model in LQC. Our results show that the quantum corrections mainly come from the third order interaction Hamiltonian, the corrected vacuum state and the corrections to the slow-roll parameters. Our paper is organized as follows. In Sec.~\\ref{sec:review}, the canonical formulism and the slow-roll inflationary model in LQC scenarios are briefly reviewed. In Sec.~\\ref{sec:power}, we study the power spectrum in LQC and recover the previous results. The effect of the inverse volume corrections on the non-Gaussianity in LQC is investigated in Sec.~\\ref{sec:bispectrum}. The detailed analysis for the sizes, shapes and shape correlations is presented in Sec.~\\ref{sec:bispectrum}and Sec.~\\ref{sec:shape}, respectively. Throughout this paper we set $8\\pi \\gamma G=1$ and Einstein's summing convention is always adopted. ",
        "conclusions": "\\label{sec:discussion} In this paper we investigated the contributions to the cosmic primordial scalar bispectrum from the inverse volume corrections in LQC scenarios. We derived the interaction Hamiltonian, however, we found that the new interactions contribute greatly to the modes with $k_1+k_2+k_3\\ll \\bar{k}$. Because the scales corresponding to $\\bar{k}$ is very large, here we take $\\bar{k}\\approx0.00014{\\rm Mpc}^{-1}$ corresponding to quadruple mode $\\bar{l}=2$, it means that the three wave lengthes in the bispectrum are all on the super-horizon scales. On so large scales the cosmic variance usually dominates over the signals, we hence neglected these new interactions in our calculations. That is to say, the interaction Hamiltonian we used is the same form as the usual one for the single field inflation models in Einstein gravity. This greatly simplifies our calculations, and more importantly, makes our results robust, i.e., our results should hold for other inflationary models in LQC scenarios.  Although the Hamiltonian shares the same forms as the one in Einstein gravity, the inverse volume corrections $\\delta_{{\\rm inv}}\\propto(k_0/k)^{\\sigma}$ still contribute to the bispectrum. Roughly speaking, there are two aspects, one comes from the deviations from the standard Bunch-Davies vacuum; the other attributes to the non-trivial gauge transformations $\\zeta(k)=u(k)/z(k)$ from the spatially flat gauge to the observable curvature perturbations. Consequently, we obtained the three-point functions of the gauge invariant curvature perturbations. We found that, except for the usual single component in slow-roll inflation models, two new shapes arise due to the corrections, namely $\\mathcal F_{1,2}$. Furthermore, we performed a careful analysis on the new shapes. We found that, the whole profiles for all the three shapes (single, $\\mathcal F_{1,2}$) are visually similar, i.e. they peak at the squeezed limit. However, the substructures among them are different. Compared to the single shape, $\\mathcal F_1$ shape upraises at another corner (See. Fig. \\ref{fig_bisp1} and \\ref{fig_bisp2}), while $\\mathcal F_2$ flattens at the same point, and $\\mathcal F_2$ peaks more dramatically in the squeezed corner than $\\mathcal F_1$. In addition, we investigated the correlations among five shapes, including local, equilateral, single and $\\mathcal F_{1,2}$. The results show that $\\mathcal F_{2}$ is highly correlated with the local type, while $\\mathcal F_{1}$ is less. It means that the latter can provide a new window for probing the loop quantum mechanisms using cosmic primordial bispectrum information. Finally, we estimated the order of observable parameter $\\Delta f_{{\\rm NL}}^{{\\rm local}}\\sim \\mathcal O(10^{-3})\\times f_{{\\rm NL}}^{{\\rm local}}({\\rm inflation})$ from the inverse volume corrections.  The non-Gaussianity from the inverse volume corrections in LQC scenarios is tiny and still undetectable currently, however, considering they are generated by the quantum effect, which is naively expected of the order $\\mathcal O({\\rm GUT}/{\\rm Planck})^{\\sigma}\\sim \\mathcal O(10^{-5})^{\\sigma}$, our finding becomes non-trivial. Especially, the results obtained in this work should be generalized directly to other inflationary models with large non-Gaussianities, in which the inverse volume corrections also becomes large therein. In addition, in this paper we only investigated the non-Gaussianities from the inverse volume corrections, while ignored those from the holonomy corrections, in which the sound speed of scalar perturbations are typically changed. Besides, we only investigated in this work the bispectrum from scalar modes, left those from tensor mode unexplored. These topics are worth investigating further."
    },
    "1112/1112.0439_arXiv.txt": {
        "abstract": "The light curve parameters of the normal type Ia SN\\,2011fe are derived from the rich archive of the AAVSO. This leads, together with the TRGB distance modulus of $(m-M)=29.39\\pm0.05$ of the parent galaxy M101, to maximum magnitudes of the unreddened SN of $M_{B} =-19.45\\pm0.08$, $M_{V} =-19.46 \\pm0.08$, and $M_{I}=-19.25\\pm0.06$ (for the standard decline rate of $\\Delta m_{15}=1.1$). When these values are inserted into the Hubble line defined by 62 SNe\\,Ia with $3000 < v < 20,000\\kms$ --- and considering also four other SNe\\,Ia with TRGB distances --- one obtains a large-scale value of the Hubble constant of $H_{0}=64.3\\pm1.9\\pm 3.2$. This value can be much improved in the future by using only TRGB distances of SNe\\,Ia. ",
        "introduction": "} \\label{sec:1} A single SN\\,Ia with a reliable distance can be important for the luminosity calibration of SNe\\,Ia as a class and hence for the determination of the cosmic value of $H_{0}$. The Hubble diagram of SNe\\,Ia out to $20,000\\kms$ \\citep[][in the following RTS\\,05]{RTS:05} shows that their intrinsic luminosity dispersion is $0.14\\mag$ --- after corrections for internal absorption, the decline rate, and if necessary for the $z$-dependent $K$-effect. From this follows that the luminosity of only one SN\\,Ia determines the distances of all SNe\\,Ia --- and hence the value of $H_{0}$ --- to within a statistical error of $0.14\\mag$ or 7\\% in linear distance.  The SN\\,Ia 2011fe occurred in the outer region of M101. Since the Cepheid distance of this galaxy is rather ambiguous (see Section~\\ref{sec:3} below), we rely on the tip of the red-giant branch (TRGB) as the most secure distance indicator for this galaxy. Attempts to calibrate other SNe\\,Ia with TRGBs have been made before (\\citealt{TSR:08a}; see also \\citealt{Mould:Sakai:09b}). SN\\,2011fe is particularly favorable for the purpose because its light curve parameters are well determined and its internal absorption is close to zero.  The method of the TRGB seems problem-free if applied to old halo populations. The absolute magnitude $M^{*}_{I}$ of the TRGB is well determined through 22 galaxies with RR~Lyr star distances. The result of $M^{*}_{I} = -4.05\\pm0.02, \\sigma = 0.08$ \\citep{TSR:08a}, agrees exactly with two earlier, independent determinations by \\citet{Sakai:etal:04} and \\citet{Rizzi:etal:07}. The color/ metallicity dependence of the TRGB has been discussed several times with contradictory results. But in the relevant color range around  $(V\\!-\\!I)^{*}=1.6$ the variation is not larger than $0.05\\mag$ \\citep{TSR:08a} and is neglected in the following. Also internal absorption poses a minimum problem in the halo fields. The most severe problem of the TRGB are very red stars like young AGBs, that can become brighter than the TRGB in dense fields of mixed populations. In that case the TRGB may get drowned by contaminating stars, as further discussed in Section~\\ref{sec:3}. This always leads to an underestimate of the distance. TRGB distances are so far restricted to $\\sim\\!22\\;$Mpc \\citep{Schweizer:etal:08}, but M101 lies well within the reach. ",
        "conclusions": "} \\label{sec:5} The importance of the RR~Lyr star-calibrated TRGB is that it uses exclusively Population~II objects and provides an independent check on the Population~I Cepheid distance scale. The determination of Cepheid distances is time-consuming and complex. The metallicity dependence of their colors makes the determination of internal absorption uncertain and implies also that their period-luminosity relation changes with metallicity \\citep{ST:08}. Notwithstanding the general agreement, already cited, of Cepheid and TRGB distances, the discrepancies of the Cepheid distances of the two galaxies, where these variables were observed in two separate fields, viz. M101 (see Section~\\ref{sec:3}) and even more pronounced in case of NGC\\,4258 \\citep{TR:11}, are unexplained. Alarming are also the Cepheids in NGC\\,1309 by \\citet{Riess:etal:09b} whose observed $(V\\!-\\!I)$ color --- uncorrected for internal reddening --- is \\textit{bluer} by $0.15\\mag$ on average than that of any other long-period Cepheids \\citep{TR:11}. Also the Cepheids of NGC\\,3021 \\citep{Riess:etal:09b} are very blue. Other blue Cepheids may be concealed by internal reddening. The properties of blue Cepheids must be influenced by one or more hidden parameters, as for instance the Helium content. Since the luminosity of these Cepheids is unknown, they cannot be used for a distance determination.  The difficulties with some Cepheids illustrate the need for an independent confirmation of the large-scale distance scale. The test can be provided by TRGB distances which carry the distance scale --- if used for the luminosity calibration of SNe\\,Ia --- to the limit where SNe\\,Ia can be observed.  The present result will be much improved by future TRGB distances of galaxies that have produced a SNe\\,Ia. It is desirable that the search for the TRGB in NGC\\,3368 and 3627 in the Leo~I group will be repeated in some uncontaminated halo fields. The next easiest targets at present are the Virgo cluster galaxies NGC\\,4526 with SN\\,1994D and NGC\\,4639 with SN\\,1990N. Slightly more difficult are the Fornax cluster galaxies NGC\\,1316 with three(!) SNe\\,Ia (1980N, 1981D, and 2006dd) and NGC\\,1380 with SN\\,1992A. These SNe, except SN\\,1981D, are blue and hence little affected by internal absorption. The Virgo and Fornax galaxies will be rather nearer than NGC\\,4038 where \\citet{Schweizer:etal:08} have demonstrated that the TRGB can be reached.  It is foreseeable that the route to a precision value of $H_{0}$ through the TRGB and SNe\\,Ia will become highly competitive with any Cepheid-based distance scale."
    },
    "1112/1112.3676_arXiv.txt": {
        "abstract": "Recently, a supersymmetric model of dark energy coupled to cold dark matter, the supersymmetron, has been  proposed. In the absence of cold dark matter, the supersymmetron field converges to a supersymmetric minimum with a vanishing cosmological constant. When cold dark matter is present, the supersymmetron evolves to a matter dependent minimum where its energy density does not vanish and could lead to the  present acceleration of the Universe. The supersymmetron generates a short ranged fifth force which evades gravitational tests. It could lead to observable signatures on structure formation due to a very strong coupling to dark matter. We investigate the cosmological evolution of the field, focusing on the linear perturbations and the spherical collapse and find that observable modifications in structure formation can indeed exist. Unfortunately, we find that when the growth-rate of perturbations is in agreement with observations, an additional  cosmological constant is required to account for dark energy. In this case, effects on large scale structures are still present at the non-linear level which are investigated using the spherical collapse approach. ",
        "introduction": "Dark energy, the component responsible for the late time acceleration of our Universe, is currently well described by a cosmological constant in the $\\Lambda$CDM concordance model. $\\Lambda$CDM has been very successful in explaining a large range of observations probing a vast range in length scales, but from a theoretical point of view the model suffers from  the fine-tuning problem and the coincidence problem \\cite{Copeland:2006wr}. This has led  to  more general models for dark energy. Scalar field models have been particularly popular over the last decade, and are predicted to exist in many theories of high energy physics, like string theory and supergravity (see e.g. \\cite{ss_book,Linde:2007fr} and references therein).  However, many of the dark energy models that have been constructed so far suffer from  problems akin to the ones they are trying to solve or introduce new issues themselves. At best they can be treated as low energy field theories valid  well below the electron mass, corresponding to the very late phase of the Universe. Hence these models need to be embedded in a better defined theory whose ultra violet behaviour is under control. So far, no such complete scenario has been constructed. Dark energy models also seem to require the existence of a very light scalar field whose coupling to matter leads to a long ranged fifth force whose presence is at odds with current gravitational tests. Screening mechanisms \\cite{Vainshtein:1972sx,Deffayet:2001uk,ArkaniHamed:2002sp,Khoury:2003rn,Khoury:2003aq,Hinterbichler:2010es,Khoury:2010xi} have been invoked in order to alleviate this problem. Axion-like particles with derivative couplings to matter are also possible candidates \\cite{Choi:1999xn}.  On the other hand, it could well be that the dark sector of the Universe, composed of the still undiscovered Cold Dark Matter (CDM) and Dark Energy (DE), could be  described by a globally supersymmetric theory \\cite{Copeland:2000vh,Brax:1999gp,Brax:1999yv}. In such a case the vanishingly small amount of dark energy which is necessary to generate the acceleration of the Universe could result from a small cosmological breaking of supersymmetry due to the non-zero CDM energy density. Such a scenario would naturally lead to a close relationship between the dark energy and the CDM energy densities. Of course, one should also ensure that corrections to the globally supersymmetric scalar potential coming from the soft supersymmetry breaking in the MSSM sector do not spoil the CDM-DE correspondence and the properties of the scalar potential in the late time Universe.  Recently \\cite{Brax:2011qs}, such a supersymmetric model of dark energy coupled to cold dark matter, the supersymmetron, was proposed by two of us. In the absence of cold dark matter, the supersymmetron converges to a supersymmetric minimum with a vanishing cosmological constant. When cold dark matter is present, the supersymmetron evolves to a matter dependent minimum where its energy density does not vanish and can contribute to the dark energy budget of the Universe.  The supersymmetron generates a short ranged fifth force between the CDM and the DE which evades gravitational tests, but could lead to observable signatures on structure formation as found in similar modified gravity theories \\cite{Davis:2011pj}-\\cite{Oyaizu:2008tb}.  In this paper we  analyse the cosmological evolution of the supersymmetron at the background level and the evolution of dark matter perturbations in the linear and non-linear regime. The non-linear regime is studied by using spherical collapse. Due to the highly non-linear behaviour of the field during the spherical collapse we are able to extract constraints on the model parameters, which are then used to constrain the background cosmology. The spherical collapse model has been previously used in models with a simple Yukawa-type modification of gravity, in the so-called $f(R)$ / chameleon models \\cite{Schmidt:2008tn,Brax:2010tj,Brax:2005ew}, in brane-world cosmologies \\cite{Schmidt:2009yj} and in models which allow for dark energy fluctuations \\cite{Mota:2004pa}-\\cite{Valageas:2010ki}. We find that a cosmological constant (CC) must be included in the model and that linear perturbations do not deviate from their $\\Lambda$CDM counterparts. On the other hand, non-linear effects are significant on astrophysical scales.  The outline  of this paper is as follows. In Sec.~(\\ref{sec1}) we present the supersymmetric formulation of the model, in Sec.~(\\ref{sec2}) we derive static solutions to the field equation and then in Sec.~(\\ref{sec3}) we study the cosmological evolution of the supersymmetron including the cosmological background evolution, linear perturbations and the  spherical collapse. In Sec.~(\\ref{sec4}) we revisit the original mass scales of the model before summarizing and concluding in Sec.~(\\ref{sec5}). ",
        "conclusions": "\\label{sec5}  We have studied the cosmological evolution of the supersymmetron and its possible effects on structure formation. Requiring that  linear perturbations are in agreement with $\\Lambda$CDM on  large scales, we find that the energy density in the supersymmetron is negligible compared to the dark matter density and a pure cosmological constant must be introduced to play  the role of  dark energy.  The non-linear evolution of the model was also investigated by using the spherical collapse model. Spherically symmetric solutions to the field equation have been  derived and used to predict the fifth-force effects on a collapsing halo. The effective gravitational constant at the edge of a spherical overdensity has been  found to be much larger than the linear prediction due to the highly non-linear properties of the model. The model parameter must be tuned such that the spherical collapse is under control on large scales. This implies that linear perturbations reduce to that of $\\Lambda$CDM. On non-linear scales the model then predicts a faster collapse than that of $\\Lambda$CDM. In particular we find that the supersymmetron predicts an excess of small mass halos compared to $\\Lambda$CDM. However for this to be the case, one or more of the mass-scales in the theory must be fine-tuned.  On very small scales, i.e. galaxy-scales, the matter density is large enough to effectively screen the fifth-force via the chameleon mechanism. This non-linear regime could in principle be probed using N-body simulations. However, due to the enormous mass of the field this poses a severe challenge for existing methods."
    },
    "1112/1112.1997_arXiv.txt": {
        "abstract": "In previous warm-up chemical models of the low-mass star-forming region L1527, we investigated the evolution of carbon-chain unsaturated hydrocarbon species when the envelope temperature is slightly elevated to $T\\approx 30$~K.  These models demonstrated that enhanced abundances of such species can be explained by gas-phase ion-molecule chemistry following the partial sublimation of methane from grain surfaces.  We also concluded that the abundances of  hydrocarbon radicals such as the C$_{\\rm n}$H family should be further enhanced as the temperatures increase to higher values, but this conclusion stood in contrast with the lack of unambiguous detection of these species toward hot core and corino sources.  Meanwhile, observational surveys have identified C$_2$H, C$_4$H, CH$_3$CCH, and CH$_3$OH toward hot corinos (especially IRAS 16293-2422) as well as towards L1527, with lower abundances for the carbon chain radicals and higher abundances for the other two species toward the hot corinos.  In addition, the {\\it Herschel Space Telescope} has detected the bare linear chain C$_3$ in 50 K material surrounding young high-mass stellar objects.  To understand these new results, we revisit previous warm-up models with an augmented gas-grain network that incorporated reactions from a gas-phase network that was constructed for use with increased temperature up to 800 K.  Some of the newly adopted reactions between carbon-chain species and abundant H$_2$ possess chemical activation energy barriers.  The revised model results now better reproduce the observed abundances of unsaturated carbon chains under hot-corino (100~K) conditions and make predictions for the abundances of bare carbon chains in the 50 K regions observed by Herschel HIFI. ",
        "introduction": "The synthesis and evolution of carbon-chain species have once again become topics of interest, specifically pertaining to early stages of low-mass star formation in which the envelope temperature is warmer than the parent dark cloud temperature of 10 K.  Recent observations indicate that two specific envelope stages for low-mass protostars, hot and lukewarm corinos, show differences in their abundances of carbon chain species.  Hot corinos are relatively condensed ($n_{\\rm H}\\sim 10^{7}$ cm$^{-3}$) and hot ($T\\gtrsim 100~$K) regions, and feature a molecular composition rich in hydrogen-rich organic species as well as oxygen-bearing species.  They can be considered  low-mass analogs to the hot cores that surround high-mass star-forming regions.  Some of the first hot corinos to be categorized were IRAS 16293-2422, NGC 1333 IRAS 4A and 4B \\citep{bclwcccmpt04,jsv04,mea04,cea03}. As in hot cores, the complex composition is the combined result of grain surface and gas phase processes \\citep{gh06}.  Observations of L1527, an envelope surrounding the transitional Class 0/I protostar IRAS 04368+2557 in Taurus, indicated that enhanced abundances of unsaturated (hydrogen-poor) hydrocarbons are likely related to a slightly elevated envelope temperature of $T \\approx 30$~K via a chemistry labeled ``warm carbon-chain chemistry''  (WCCC) \\citep{sea07}.  Although the chemical composition resembles that of a dark cloud core,  the composition of WCCC regions can also be attributed to a process beginning with grain surface chemistry.    The elevated temperature allows methane to evaporate from the ice mantles surrounding dust particles and act as a precursor for a carbon-chain rich ion-molecule chemistry. The resulting chemical composition is a defining characteristic of the source, which we designated a ``lukewarm corino\".  The idea of WCCC has proved somewhat controversial, however, and alternative explanations, such as standard gas-phase ion-molecule chemistry in a collapsing source, have been advanced \\citep{charn10}.   A recent interferometric study provides  evidence for WCCC chemistry within the central protostar and normal ion-molecule chemistry in cold cloud remnants  \\citep{sea11}. It has not been firmly established if the lukewarm corino phase is one step on an evolutionary track leading to the hot corino phase, or if it represents a terminal stage.  Since the observation of L1527, additional observations have identified a similar kinetic temperature and similar enhancements for hydrocarbon chains towards the  source IRAS 15398-3359 in Lupus, alternatively identified as B228 \\citep{sea09}, so that B228 is the second lukewarm corino to be studied.  Also, \\citet{sea09} reported detections of some carbon chains, especially C$_{4}$H, toward a number of hot corinos at temperatures near 100 K, but with abundances significantly lower than those in the two lukewarm corino sources,  suggesting a diminution in abundance with increasing temperature.  More recently, linear C$_3$ has been detected in absorption in regions at a temperature of $\\approx$ 50~K  against continua in the THz region from the sources W31C and W49N \\citep[hereafter referred to as the HIFI sources]{mea10c3}.  We previously reported the use of a gas-grain chemistry network with a simple warm-up following a cold stage to simulate the evolution of molecular abundances in L1527 as a result of increasing temperature \\citep{hhg08}(hereafter HHG). A comparison of model abundances with the observations of \\citet{jsv04} and \\citet{sea08,sea07} showed that the best overall agreement occurs at times when the temperature approaches $T=30$ K, the indicated temperature for the envelope.  Reasonable agreement with observation in L1527 was found near this temperature whether or not the warm-up was terminated at 30 K or continued to higher temperatures.  Similar chemistry was also found in the hydrodynamic-chemical model of prestellar to protostellar core formation of \\cite{aik08}.  Some of the warm-up models were extended to maximum temperatures of 100 and 200 K, in a similar fashion to earlier hot core models, in which it was shown that the standard oxygen-containing complex molecules in such sources could be partially produced by surface chemistry followed by desorption \\citep{gh06,gwh08}.  These prior studies emphasized the typical hot-core species instead of the calculated abundances of carbon-chain species.  It was later noted  that in all of our warm-up models, carbon chain species are quite abundant at the higher temperatures for long periods of time.  This result contrasts with the lack of strong detections of such species toward hot core and hot corino sources, with the general exception of HC$_{3}$N \\citep{hvd09}.  In the survey of \\citet{sea09}, the fractional abundance values obtained for C$_{4}$H vary, but are generally smaller than $X_{\\rm C_{4}H} = N_{\\rm C_{4}H}/N_{\\rm H} =10^{-10}$ in hot corinos, as compared with $X_{\\rm C_{4}H} \\approx$ 2--3 $\\times 10^{-9}$ for the two lukewarm corino sources.  The observed hot corino result is smaller than the peak fractional abundance of $X_{\\rm C_{4}H} > 10^{-8}$ found in the extended warm-up models of HHG.  Some other observational studies have been directed at the question of unsaturated molecules in star-forming regions. In their review of the chemistry of complex molecules, \\citet{hvd09} mentioned a study of the C$_{4}$H/CH$_{3}$OH and HC$_{5}$N/CH$_{3}$OH abundance ratios in seven hot cores that resulted in upper limits of 0.01--0.1 for the ratios (S. and J. J{\\o}rgensen-Bisschop, private communication).  In addition, an average abundance of $N_{\\rm CH_{3}OH}/ N_{\\rm H} = 9.5\\times 10^{-7}$ was reported for the seven hot cores \\citep{BJDW07}, leading to upper limits for C$_{4}$H and HC$_{5}$N high enough to overlap with the $T_{\\rm max} = 100$ \\& 200 K peak results of HHG for both species.   It would appear at present that C$_{4}$H is detected in largest abundance in the two known lukewarm corinos, that it is detectable in hot corino sources although to much lower abundance than predicted, and that it has a high upper limit with respect to methanol in hot cores.  The survey of \\citet{sea09b} also included the detection of CH$_3$OH toward L1527 with a fractional abundance of $X_{\\rm CH_{3}OH}\\approx 10^{-9}$, similar to what is observed in cold cores such as TMC-1.  Unlike C$_{4}$H, and possibly other radicals of this class, methanol is a molecule typically associated with high gas-phase abundances only in hot cores and corinos, where it is desorbed from grains.  Although the picture is clearly incomplete at the present time, comparison of carbon chain species and species known to prefer high temperatures  provides an initial basis for chemical distinction of lukewarm and hot corino sources.  Although one can think of a lukewarm corino as a prior stage of evolution than a hot corino,   \\citet{sea09b} have also suggested that the difference between lukewarm and hot corinos has to do at least partially with the rate of collapse of the protostar; the slower the rate, the more chance for CO to form in the gas, accrete onto grains, and lead to oxygen-containing species; the faster the rate, the greater the chance that C atoms land on grains and form methane.  In a recent study, \\citet{har10} reported the extension of the temperature range of the OSU gas-phase chemical model  to 800 K by the inclusion of many new reactions, including those with non-zero activation barriers.  Select reactions might also be relevant at hot core and corino temperatures.  In particular, reactions between the carbon chain radicals C$_{\\rm n}$H and C$_{\\rm n}$ and H$_2$ were omitted from previous networks because they possess barriers of $E_{\\rm{A}} \\gtrsim 10^3$ K, and are inefficient in cold models.  The effect of adding these gas-phase reactions to our warm-up gas-grain model network would likely be to decrease the calculated abundances of these radicals at higher temperatures, so as to make the results more in line with hot corino observations of C$_{4}$H.  On the other hand, activation energy barriers for reactions between normal molecules in ground electronic singlet states, such as cyanopolyynes, and H$_{2}$ should be much larger so that these reactions are still not likely to be important in hot corinos.  The updated reaction network should not affect the abundances at lukewarm corino conditions.  In this paper, we present revised gas-grain warm-up models including addition of gas-phase radical-H$_{2}$ reactions as well as many other reactions added by \\citet{har10}, and use the results to determine a revised temperature dependence of the abundances of carbon-chain species in the range from 30~K (lukewarm corinos) to 100~K (hot corinos). ",
        "conclusions": "We have calculated gas-grain chemical models using an expanded reaction network that includes non-zero barrier reactions \\citep{har10} in order to better understand the abundances of carbon chains and other species in star-forming regions from lukewarm to hot corinos and in dense, cool regions detected by Herschel surrounding young high-mass stellar objects.  Our models include a cold core period at 10 K, followed by a warm-up period to a maximum asymptotic temperature ranging from 30 K to 200 K.  The specific temperature of a source can either be associated with an asymptotic temperature or a temperature in the range from 10 K to the asymptotic temperature of a model. It is not presently known if a lukewarm corino such as L1527 ($T \\approx 30$~K) represents a terminal stage of chemical evolution or a transient phase on an evolutionary track toward the hot corino stage (100-200~K), but the present models allow either scenario, because,  generally speaking, the optimum results we get are not strongly dependent on whether a source temperature is considered to be an intermediate temperature or an asymptotic temperature.  Recent observations show that a range of physical conditions and chemical compositions can exist not only toward complex hot cores \\citep{mea10c3} but also towards lukewarm corinos \\citep{sea11}, and the results provide a basis for future modeling investigations.   The present models are better suited to address the more specific question of whether large abundances of carbon chain species at $T\\gtrsim 100$~K exist.   Our four main conclusions are as follows:  \\begin{enumerate}  \\item{\\it The inclusion of new reactions into our gas-grain network does not significantly affect the results and conclusions obtained from previous warm-up model fits to lukewarm corino observations.} As a benchmark calculation, we revisited the simulations of the lukewarm corino L1527 and evaluated fits  using the same criteria as \\citet{hhg08}.  The times of peak agreement, $t_{\\rm opt}$, and associated temperatures  (25 - 30 K) do not change as a result of using an expanded reaction network with reactions containing barriers.  The values of the optimum mean confidence level, $\\kappa_{\\rm max}$, are only slightly diminished by comparison with the previously reported values, partly because of the addition of more species to the observation set.   Thus, the idea  that the chemistry of a lukewarm corino involves a warm regeneration of carbon chains resulting from CH$_{4}$ desorption at 25 - 30 K \\citep{sea07}, incorporated into both our old and updated models, is supported.  \\item{\\it The inclusion of the new reactions tends to lower carbon-chain radical abundances at higher temperatures.}  Prior models showed significantly increased abundances over significant periods of time for several carbon-chain species including the C$_{\\rm n}$H and HC$_{\\rm n}$N families above the L1527 observed values as the temperature increased from lukewarm corino ($T\\sim 30~$K) toward hot corino/hot core levels ($T\\gtrsim 100$~K).  In the updated networks, reactions with activation energy barriers between molecular hydrogen and polyyne radicals become the dominant destruction mechanism for these species in star-forming regions as temperatures exceed $T\\gtrsim 80$~K.  The inclusion of these reactions reduces the abundances of most carbon chain radical species at higher temperatures.  \\item{\\it Most importantly, the new observations and models begin to chemically distinguish lukewarm and hot corinos.} The species C$_2$H, C$_4$H, CH$_3$CCH, and CH$_3$OH have been observed with abundance differences of approximately two orders of magnitude between these types of sources, and previous models could not account for the depletion of C$_2$H and C$_4$H at larger temperatures.  Updated results at increased temperatures of $T\\gtrsim 80$~K show agreement with the diminished abundances of C$_2$H and C$_4$H observed toward hot corinos compared with lukewarm ones.  The reactions most responsible for diminishing the abundances of the carbon chain (polyyne) radicals C$_2$H and C$_4$H and distinguishing lukewarm from hot corinos in their carbon chain inventories are those with molecular hydrogen, listed in Table~\\ref{tbl-newrxns}.  Even though they possess small barriers, these reactions are still competitive at temperatures above 80 K.  The model results indicate that the ``excess\" polyynes are primarily converted to unobservable polyacetylene forms of C$_{\\rm n}$H$_2$, although it is not currently possible to directly verify this.  Unlike the case for the unsaturated polyyne radicals, the model results, in agreement with observation, show that the  the abundance of CH$_3$CCH has  the opposite trend of lower abundances in lukewarm sources and distinctively larger abundances in hot corino sources.  The observed abundances of CH$_3$OH show a similar pattern of smaller values toward a lukewarm source and larger values toward hotter sources, which the model results also accurately replicate. The trends observed in these four species provide the framework for a molecular distinction between the lukewarm corino and hot corino classes of protostellar envelopes.  A recent observation of L1527 indicates depletions toward the innermost region, suggesting that such differentiation occurs within an individual source \\citep{sea11}.  This observation may be consistent with that of \\citet{Oberg10} where complex oxygen-containing molecules appear to be associated with a cold portion of the low-mass protostar B1 -b.  \\item {\\it The high abundance of C$_{3}$ observed in absorption in dense regions with $T\\sim 50$~K towards W31C and W49N is explained by warm-up models using both our previous and updated gas-grain network. }  Further, this agreement is reached without requiring the photo-destruction of PAH molecules, another suggested source of this species \\citep{mea10c3}.  The new models presented here provide evidence of only a slight diminution of abundance as a function of increasing chain length for more complex bare carbon chains.  The formation of C$_3$ from acetylene during the warm-up implies the presence of lukewarm corino-like regions within the complex sources.  \\end{enumerate}  Each of these conclusions could be further tested and improved with more observational data.  The survey of \\citet{sea09} has indicated that some additional sources such as B1 and L1448N may also classify as lukewarm corinos on the basis of enhanced C$_4$H abundance, while \\citet{ccbwm11} indicated that the enhanced abundance of several carbon chain species toward L1251A (IRS3) could possibly be attributed to a similar effect.  We encourage follow-up observations of both lukewarm and hot corinos to investigate other carbon-chains including CH$_3$CCH, as well as CH$_3$OH.  The depleted abundances of C$_2$H and C$_4$H toward hot corinos may make their detection more difficult, but the CH$_3$C$_{\\rm n}$H chains (n=2, 4, 6, $\\dots$) should be present in enhanced abundance toward these sources.  The model results for the HIFI sources would benefit from additional detections of longer bare carbon chains or an expanded search for other  carbon chains in the inventory  of lukewarm or hot corinos.  The authors gratefully acknowledge contributions useful discussions of observational data with Nami Sakai, Maryvonne Gerin, and Bhaswati Mookerjea.  G.H. thanks Rose Finn for computational time.  E.H. would like to thank the National Science Foundation for supporting his research program in astrochemistry and NASA for support to study the formation of pre-planetary matter and the analysis of Herschel data."
    },
    "1112/1112.2393_arXiv.txt": {
        "abstract": "We report on the first 180 days of RXTE observations of the outburst of the black hole candidate IGR~J17091--3624. This source exhibits a broad variety of complex light curve patterns including periods of strong flares alternating with quiet intervals. Similar patterns in the X-ray light curves have been seen in the (up to now) unique black hole system GRS~1915+105. In the context of the variability classes defined by \\citet{Belloni00} for GRS~1915+105, we find that IGR~J17091--3624 shows the $\\nu$, $\\rho$, $\\alpha$, $\\lambda$, $\\beta$ and $\\mu$ classes as well as quiet periods which resemble the $\\chi$ class, all occurring at 2-60 keV count rate levels which can be 10-50 times lower than observed in GRS~1915+105. The so-called $\\rho$ class ``heartbeats'' occur as fast as every few seconds and as slow as $\\sim$100 seconds, tracing a loop in the hardness-intensity diagram which resembles that previously seen in GRS~1915+105. However, while GRS~1915+105 traverses this loop clockwise, IGR~J17091--3624 does so in the opposite sense. We briefly discuss our findings in the context of the models proposed for GRS~1915+105 and find that either all models requiring near Eddington luminosities for GRS~1915+105-like variability fail, or IGR~J17091--3624 lies at a distance well in excess of 20 kpc or, it harbors one of the least massive black holes known ($< 3 M_\\odot$). ",
        "introduction": "\\label{sec:intro}  Observations with the \\textit{Rossi X-ray Timing Explorer} (RXTE) have led to extraordinary progress in the knowledge of the variability properties of many different types of sources, particularly of black hole candidates (BHCs) and neutron stars (NSs) in low-mass X-ray binaries \\citep[e.g,][]{Vanderklis06}. Both BHCs and NS are known to exhibit distinct ``accretion states'' (usually defined in terms of their X-ray spectral shape and variability), whose characteristics are thought to be intimately related to the physics of the accretion flow and its interaction with the central compact object \\citep[e.g.,][]{Belloni10b, Remillard06, Vanderklis06,Belloni11}.  Of all Galactic BHC X-ray binaries known today, one of the most prolific in terms of state transitions is GRS~1915+105. In outburst since its discovery in 1992 with WATCH \\citep{Castro-Tirado92}, it is a 33 days orbital period binary system \\citep{Greiner01} harboring a $14\\pm4.4 \\ M_{\\odot}$ black hole \\citep{Greiner01,Harlaftis04}. At a distance of $\\sim12.5$ kpc \\citep{Mirabel94}, GRS~1915+105 is very often at Eddington or super-Eddington luminosity \\citep[e.g., ][]{Done04}.  \\begin{figure*} \\centering \\resizebox{2\\columnwidth}{!}{\\rotatebox{-90}{\\includegraphics[clip]{f1.eps}}} \\caption{\\textit{Upper Panel}: \\textit{Swift/BAT} light curve. \\textit{Lower Panel}: Circles: RXTE Crab normalized 2-20 keV intensity.  RXTE values are averaged per observation, background subtracted and dead-time corrected.  Triangles: RXTE 2-50 keV flux RXTE observations between day 55601 and 55615 are affected by the nearby source GX~349+2 (Sco X-2); the intensity during this period was arbitrarily fixed to 65 mCrab and no X-ray flux is reported. Shaded areas mark quiet periods (see text). } \\label{fig:outburstlc} \\end{figure*}   GRS~1915+105 is known to show quasi-periodic oscillations (QPOs) in the 0.001--70 Hz range, some of which are the same as those generally seen in other BHCs \\citep[]{Morgan97, Markwardt99, Reig00a, Strohmayer01b, Soleri08}. However, so far GRS~1915+105 has been unique in that its X-ray light curves exhibit more than a dozen different patterns of variability usually called ``classes'' (which are referred to with Greek letters), most of which are high amplitude and highly-structured \\citep[e.g.,][]{Belloni00}. Most of this structured variability is thought to be due to limit cycles of accretion and ejection in an unstable disk \\citep[see, e.g., ][and references within]{Belloni97, Mirabel98, Tagger04, Neilsen11}. Much effort has been expended in order to understand why GRS~1915+105 is so unusual among BHCs \\citep[e.g.,][]{Fender04}. It has been proposed that the high accretion rate estimated for GRS~1915+105 might be the determining factor \\citep[e.g.,][]{Done04}; however, the lack of any other source showing similar characteristics has prevented definite conclusions.     \\begin{figure} \\centering \\resizebox{1\\columnwidth}{!}{\\rotatebox{0}{\\includegraphics[clip]{./f2.eps}}} \\caption{The upper and lower panels of each frame show a light curve of IGR~J17091--3624 and GRS~1915+105, respectively. Count rates are given in 1 s bins, per PCU, 2-60 keV and background subtracted. The IGR~J17091--3624 A-C intervals last $T1$ seconds and come from observations 96420-01-05-00, -06-00 and -07-01, respectively.  For GRS~1915+105 they last $T2$ seconds and come from observations 10408-01-40-00, 20402-01-34-00 and 93701-01-02-01, respectively. $R \\equiv MIN_{GRS~1915+105} / MIN_{IGR~J17091-3624}$, where $MIN$ is an average count rate during minima in the above light curve and is given for each frame as an order of magnitude estimate of the flux ratio. Greek letters on the side indicate variability the class. } \\label{fig:complc1} \\end{figure}   IGR~J17091--3624 was discovered with INTEGRAL/IBIS during a Galactic Center observation on April 2003 \\citep{Kuulkers03b} and later again in 2007 \\citep{Capitanio09}. Reexamination of archival data from different missions showed that IGR~J17091--3624 was also active in 1994, 1996, 2001 \\citep{Revnivtsev03a, Intzand03, Capitanio06}. Although a neutron star could not be excluded based on spectral characteristics and the outburst radio/X-ray flux ratio, \\citet{Capitanio06} concluded that IGR~J17091--3624 is most probably a black hole. A new outburst was detected with \\textit{Swift/BAT} in February 2011 \\citep{Krimm11}; the radio/X-ray characteristics during the first $\\sim$40 days, combined with the discovery of QPOs \\citep{Rodriguez11a}, further suggested that IGR~J17091--3624 is a black hole.   Recently we reported the discovery of 10 mHz QPOs in RXTE observations of IGR~J17091--3624 \\citep{Altamirano11a} very similar to those in GRS~1915+105. The suggested link between these sources was strengthened by our discovery of (i) a continuous progression of regular, quasi-periodic flares occurring at a rate of 25--30 mHz \\citep{Altamirano11b}, and (ii) a broad variety of complex patterns alternating with quiet intervals \\citep{Altamirano11b,Altamirano11c}, resembling respectively, the so called $\\rho$ class (``heartbeat'') oscillations and the complex $\\beta$ class patterns, both so far seen only in GRS~1915+105.   The existence of a source showing similar X-ray variability to that seen in GRS~1915+105 opens a new window of opportunities to understand the physical mechanism that produces the highly structured X-ray variability. Multi-wavelength and high spectral resolution studies of IGR~J17091--3624 as compared with GRS~1915+105 can help gaining further insights into the role of disc-jet coupling \\citep[e.g., review by][]{Fender04} and accretion disk winds \\citep[e.g, ][]{Neilsen11} in accreting BHCs. Furthermore, it allows the possibility to test the role of the accretion disk size, evolutionary state of the companion star, and the evolution of the disc structure as an explanation of the long- and short-term X-ray variability seen in GRS~1915+105 \\citep[e.g.,][]{Done04}.  In this Letter we use RXTE data to describe the phenomenological similarities between IGR~J17091--3624 and GRS~1915+105. We demonstrate that IGR~J17091--3624 shows the same type of high amplitude and highly-structured variability previously observed only in GRS~1915+105, although at much lower count rates. The analysis presented here is the first step of a larger program to compare IGR~J17091--3624 and GRS~1915+105 in detail using data obtained with \\textit{RXTE}, \\textit{Swift}, \\textit{XMM-Newton} as well as optical facilities.      \\begin{figure} \\centering \\resizebox{1\\columnwidth}{!}{\\rotatebox{0}{\\includegraphics[clip]{f3.eps}}} \\caption{Similar to Figure~\\ref{fig:complc1}. The IGR~J17091--3624 D-F intervals come from observations 96420-04-03, -08-03 and -09-06, respectively.  For GRS~1915+105 they come from observations 20187-02-01-00, 95701-01-31-00 and 10258-01-10-00, respectively. } \\label{fig:complc2} \\end{figure}   \\begin{figure*} \\centering \\resizebox{2\\columnwidth}{!}{\\rotatebox{0}{\\includegraphics[clip]{f4.eps}}} \\caption{\\textit{Left panels} show the hardness-intensity diagram for flares observed during the $\\rho$ variability class in IGR~J17091--3624 (top; ObsID 96420-01-04-03) and GRS 1915+105 (bottom; ObsID 96378-01-01-00) occurring at an average period of T1$=$70.96 seconds and T2$=$63.72 seconds, respectively. Arrows mark the time evolution. Inset shows representative flares.  Light curves and colors are estimated from 1 sec averages. Intensity is the count rate in the 2-60 keV range (absolute channels 0-240) and hard color is the 6.5--15.0~keV / 2--6.5~keV count rate ratio (channels 15-35 and 0-14, respectively).  \\textit{Right panels} show representative power spectra from averages of 512 sec segments during the $\\rho$ variability class for IGR~J17091--364 (top, ObsID:96420-01-05-000, MJD 55647.9) and for GRS~1915+105 (bottom, ObsID:40703-01-07-00, MJD 51235.3).} \\label{fig:loopPDS} \\end{figure*} ",
        "conclusions": "An extensive literature exists attempting to understand the complex variability observed in GRS~1915+105.  Some authors propose that the high luminosity (close to, or super-Eddington ) of GRS~1915+105 is the determining factor \\citep[e.g.,][and references therein]{Belloni97, Vilhu98, Belloni00, Nayakshin00, Janiuk02, Done04, Neilsen11}. Although it is generally accepted that the complex X-ray variability in GRS~1915+105 results from disk instabilities, the exact nature of the instability remains unknown.  The lack of at least a second source showing similar characteristics has prevented definite conclusions.  In this Letter we show for the first time that another source, the BHC IGR~J17091-3624, can show the same broad variety of complex light curves as GRS~1915+105 (at least 7 of the 12 variability classes observed in GRS~1915+105, in addition to unclassified ones; note that at the time of submission of this Letter IGR~J17091-3624 is still active). Although the comparisons presented in this paper constitute only a first step, the observed similarities suggest that the complex light curves of the two sources are produced by the same physical mechanisms. If true, the low flux of IGR~J17091-3624 compared with GRS~1915+105 combined with the circumstance that currently neither the distance to IGR~J17091-3624 nor the mass of its compact object are known, raises the fundamental question: is IGR~J17091-3624 close to Eddington or not at times when showing the same characteristic X-ray variability?  A scenario in which IGR~J17091-3624 is not emitting at close to the Eddington but only at a few percent, is at variance with models where the variability is explained as due to disk instabilities that \\textit{only} occur at high luminosity \\citep[e.g.,][and references therein]{Vilhu98, Belloni00, Nayakshin00, Janiuk02, Done04, Neilsen11}. However, in this scenario IGR~J17091-3624 would follow radio/X-ray correlation of some or most BHCs \\citep[depending if the distance is closer to $\\sim11$ kpc or $\\sim17$ kpc, respectively;][]{Rodriguez11}.     \\begin{figure} \\centering \\resizebox{1\\columnwidth}{!}{\\rotatebox{-90}{\\includegraphics[clip]{f5.ps}}} \\caption{Mass of the compact object versus distance to the binary system assuming IGR~J17091-3624 is emitting at Eddington. The three curves correspond to different (1, 2 and 3) correction factors between the 2-50 keV and the bolometric flux. } \\label{fig:edd} \\end{figure}   A scenario in which IGR~J17091-3624 is emitting at close to the Eddington limit puts constraints on its mass and distance. Although this scenario would imply that, independently of the distance, IGR~J17091-3624 does not follow the standard Radio/X-ray luminosity plane for BHCs \\citep{Rodriguez11}, this would be similar to GRS~1915+105, as it does not follow the radio/X-ray flux plane either \\citep[see, e.g., figure 4 in ][ and references therein]{Rodriguez11}.   We therefore made a rough estimate of the source flux in each observation by calculating the average-per-observation energy spectrum. In the upper panel of Figure~\\ref{fig:outburstlc} we plot the 2-50 keV unabsorbed flux. In IGR~J17091-3624 we observe the $\\alpha$, $\\beta$, $\\nu$, $\\rho$ and $\\mu$ classes when the flux is between $\\sim$2 and $\\sim$3.7 $10^{-9}$ erg~cm$^{-2}$~s$^{-1}$ (the {\\tt diskbb} temperature {\\tt kT} vary between 1 and 2 keV while the power law index between 2 and 3). Assuming that IGR~J17091-3624 is emitting at Eddington rates, we derived the mass of the compact object as a function of distance (Figure~\\ref{fig:edd}) for a 2-50 keV flux of $4 \\cdot 10^{-9}$~erg~cm$^{-2}$~s$^{-1}$. Given that the correction factor linking the 2-50 keV and the bolometric flux is not exactly known, we plot curves for factors of 1, 2 and 3. Figure~\\ref{fig:edd} implies that if IGR~J17091-3624 emits at Eddington, then either it harbors the lowest mass black hole known today ($<3 M_\\odot$ for distances lower than 17 kpc), or, it is very distant. Such a large distance, together with its $b\\simeq$2.2$^\\circ$ Galactic latitude, would imply a significant, but not necessarily implausible, altitude above the disk \\citep[e.g., $b\\simeq$-2.8$^\\circ$ and $d>25$kpc for the BHC GS~1354--64,][]{Casares09}.    Clearly, constraining the distance to IGR~J17091-3624 is fundamental for understanding the physical processes that govern the X-ray variability in IGR~J17091-3624, and hence those in GRS~1915+105. The interstellar absorption in the direction of IGR~J17091-3624 is high \\citep[$\\sim1 \\times 10^{22}$ cm$^{-2}$, see][]{Rodriguez11}, implying that optical studies when the binary system is in quiescence \\citep[e.g., ][]{Casares04} might be challenging. However, accurate parallax distance estimates like those reported for the black hole X-ray binary V404 Cyg \\citep{Miller-Jones09} could still be possible.  \\textbf{Acknowledgments:} We thank Dave Russell and JamesMiller-Jones for insightful discussions. M.L. acknowledges support from an NWO Rubicon fellowship. TB has received funding from the European Community\u2019s Seventh Framework Programme (FP7/2007-2013) under grant agreement number ITN 215212 \u201cBlack Hole Universe,\u201d and TMD from the Spanish MEC under theConsolider-Ingenio 2010 Programme grant CSD2006- 00070: \u201cFirst Science with the GTC.\u201d"
    },
    "1112/1112.2670_arXiv.txt": {
        "abstract": "We report the results from spectroscopic observations of 113 ultra-wide, low-mass binary systems, composed largely of M0--M3 dwarfs, from the {\\slowpokes} catalog of common proper motion pairs identified in the Sloan Digital Sky Survey. Radial velocities of each binary member were used to confirm that they are co-moving and, consequently, to further validate the high fidelity of the {\\slowpokes} catalog. Ten stars appear to be spectroscopic binaries based on broad or split spectral features, supporting previous findings that wide binaries are likely to be hierarchical systems. We measured the {\\ha} equivalent width of the stars in our sample and found that components of 81\\% of the observed pairs has similar {\\ha} levels. The difference in {\\ha} equivalent width amongst components with similar masses was smaller than the range of {\\ha} variability for individual objects. We confirm that the Lepine et al. $\\zeta$-index traces iso-metallicity loci for most of our sample of M dwarfs. However, we find a small systematic bias in $\\zeta$, especially in the early-type M dwarfs. We use our sample to recalibrate the definition of $\\zeta$. While representing a small change in the definition, the new $\\zeta$ is a significantly better predictor of iso-metallicity for the higher mass M dwarfs. ",
        "introduction": "\\label{Sec: intro} Low-mass stars, generally defined as the regime bracketed by the hydrogen-burning limit ($\\sim$0.08~{\\Msun}) and the onset of molecular lines in the photosphere ($\\sim$0.8~{\\Msun}), make up $\\sim$ 70\\% of the Milky Way's stars \\citep{Bochanski2010} and are, perhaps, the best tracers of the structure, dynamics, and evolutionary history of the Galaxy. However, their intrinsic faintness has historically limited the construction of large samples. In addition, the ubiquitous molecular features in their photospheres and the resulting incomplete line lists has restricted the accuracy and usefulness of theoretical atmospheric models. Large surveys, such as the Sloan Digital Sky Survey \\citep[SDSS;][]{York2000} and the Two Micron All Sky Survey \\citep[2MASS;][]{Skrutskie2006}, have played a large role in advancing our understanding of low-mass stars. With a photometric catalog of more than 33~million \\citep{Bochanski2010} and a spectroscopic catalog of more than 70,000 \\citep{West2011} M dwarfs, SDSS has enabled studies of the spatial \\citep{Bochanski2010} and kinematic distributions \\citep{Bochanski2007b, Fuchs2009} in the Milky Way; the mass and luminosity functions \\citep{Covey2008, Bochanski2010}; and magnetic activity \\citep[e.g.,][]{West2008a, West2011, Kruse2010, Kowalski2009, Hilton2010} of low-mass stars.  The metallicity of low-mass stars remains an elusive parameter to measure. Given the large number of M dwarfs in the Milky Way, an absolute metallicity scale tied to an easily observable spectral index would allow for the tracing of the formation history and the chemical evolution of the Galaxy \\citep[e.g.,][]{West2008a}, the dependence of the fundamental mass--radius relation on metallicity at the bottom of the main sequence \\citep[e.g.,][]{Lopez-Morales2007}, and the relationship between metallicity and the presence of planets \\citep[e.g.,][]{Laws2003, Valenti2008}. While spectral modeling has allowed for metallicity determinations and well-defined metallicity indices for warmer stars, such efforts in the late-K and M spectral type regimes (e.g., \\citealt*{Hauschildt1999a}; \\citealt{Witte2011}) have met with notable problems due to the onset of broad molecular lines at $\\lesssim$4300~K and due to incomplete molecular line lists. Some authors have tried to use photometric indices to infer the metallicity \\citep{Bonfils2005, Johnson2009, Schlaufman2010}, but these techniques rely on trigonometric parallax measurements which are uncommon for M dwarfs.  Some useful spectral features that correlate with metallicity have been identified. In the near-infrared, \\citet{Rojas-Ayala2010} developed a metallicity indicator based on the strength of the {\\Nai} doublet, the {\\Cai} triplet, and a temperature-sensitive water index. This technique has so far only been calibrated over a limited range but delivers the greatest precision among current techniques. Meanwhile, much effort has gone into optical spectra. As the TiO band in the optical spectrum becomes weaker with decreasing metallicity \\citep{Bessell1982}, the ratio of CaH and TiO molecular bands has been used to distinguish M dwarfs from M subdwarfs \\citep*{Kirkpatrick1991, Reid1995, Gizis1997, Lepine2003, Burgasser2006a}. Building on these studies,  \\citet*[][hereafter \\citetalias{Lepine2007b}]{Lepine2007b} defined the  metallicity-dependent quantity $\\zeta$ using the \\citet{Reid1995} CaH2, CaH3, and TiO5 molecular band heads; this allowed for the segregation of low-mass dwarfs into four classes: dwarfs (dMs), subdwarfs (sdMs), extreme subdwarfs (esdMs), and ultra subdwarfs (usdMs). These classes may also trace the Galactic populations to which these stars belong: dMs were formed in the thin disk, sdMs in the thick disk, and esdMs/usdMs in the halo. \\citetalias{Lepine2007b} calibrated the definition of $\\zeta$ using the visual binary pairs known at the time, including four sdM and two esdM pairs. \\citet*{Woolf2009} mapped the $\\zeta$ index to an absolute metallicity scale using dM binaries with a FGK companion of measurable absolute metallicity; but it suffers from significant scatter ($\\sim$0.3~dex).  Wide binary (or multiple) systems are ideal, coeval laboratories to constrain and calibrate the observable properties of stars as the components were presumably formed at the same time and from the same primordial material but have evolved independently. In \\citet[][hereafter \\citetalias{Dhital2010}]{Dhital2010} we identified the Sloan Low-mass Wide Pairs of Kinematically Equivalent Stars ({\\slowpokes}) catalog  consisting of 1342 ultra-wide, low-mass common proper motion (CPM) binary systems from the SDSS Data Release 7 \\citep[DR7;][]{Abazajian2009} by matching angular separations, photometric distances, and proper motions.  The binary systems in the catalog have at least one low-mass (spectral subtype K5 or later) component, projected physical separations of $\\sim$10$^3$--10$^5$~AU, and distances of $\\sim$50--800~pc. While most {\\slowpokes} pairs are disk dwarfs, 70 low-metallicity sdM and 21 white dwarf--dM pairs were identified based on their reduced proper motions. A Galactic model---based on empirical constraints on the stellar number density \\citep{Bochanski2010, Juric2008} and velocity \\citep{Bochanski2007a} distributions in the Milky Way---was used to assess the probability that the candidates were a chance alignment of random stars; only pairs with such probabilities $\\leq$ 5\\% were published in the {\\slowpokes} catalog.  The overall fidelity of the catalog is expected to be $\\sim$98\\%. Hence, the {\\slowpokes} catalog is a very clean and diverse source of CPM binary systems to be used in follow-up studies. As the {\\slowpokes} catalog spans a wide range in mass  and a smaller, but still considerable, range in metallicity, it is an ideal sample to constrain the $\\zeta$ index as well as to eventually map it to an absolute metallicity space.  Magnetic activity has been shown to decline with age, with activity lifetimes of $\\sim$1--2~Gyr for M0--M3 and $\\sim$7--8~Gyr for M5--M7 dwarfs \\citep{West2006, West2008a, West2011}. This monotonic decline of activity with age is a signature of stellar spin-down and suggestive of a gyrochronology-like age--rotation--activity relationship in M dwarfs \\citep{Skumanich1972, Barnes2003, Barnes2007, Delorme2010}. Leveraging the coevality of components of the {\\slowpokes} pairs is a good way of testing this relationship.  We have carried out a spectroscopic follow-up study of 113 CPM pairs from the {\\slowpokes} catalog. Section~\\ref{Sec: observation} details our observations and the data reduction procedures. In Section~\\ref{Sec: results} we use our radial velocities to assess the fidelity of the observed {\\slowpokes} pairs, use them to redefine the \\citetalias{Lepine2007b} $\\zeta$ index, and examine the magnetic activity properties of the {\\slowpokes} pairs. The conclusions are presented in Section~\\ref{Sec: conclusions}. ",
        "conclusions": "\\label{Sec: conclusions} We have carried out a spectroscopic follow-up study of 113 ultra-wide, low-mass CPM binary systems from the {\\slowpokes} catalog \\citepalias{Dhital2010} using the GOLDCAM spectrograph on the KPNO 2.1~m telescope. We measured the radial velocities of each component by cross-correlating them with appropriate standards and used them to assess the fidelity of pairs in the {\\slowpokes} catalog. 95 of the 113 (84\\%) of the pairs have the same radial velocity within 3~$\\sigma$. At least five of the pairs with discrepant radial velocities are candidate SBs, which would explain the difference. There may be additional spectroscopic companions undetected in our low-resolution spectra. \\citet{Law2010} found that 45\\% of the {\\slowpokes} systems are either hierarchical triples or quadruples. Either high-resolution spectroscopy or imaging would be needed to identify the close companions and to further quantify the incidence of higher-order systems in wide binaries.  We examined the {\\ha} activity in our observed sample. The components of binary pairs exhibited overwhelmingly comparable levels of {\\ha} activity. Moreover, the {$\\Delta$\\ha} of the pairs with similar {\\rz} colors and two active components, while large, was several times smaller than the variation seen in single M dwarfs. Our results corroborate that low-mass stars of the same mass should spin-down at the similar rates over time. However, larger samples of active are needed to confirm this finding and to constrain the rate of this spin-down.  We tested the \\citetalias{Lepine2007b} $\\zeta$-index and found that, to first-order, it is a robust measure of relative metallicity. The value of $\\zeta$ for the two components in each binary system match within the error bars for most pairs, indicating a common metallicity as expected. However, we find a systematic bias for the higher-mass M dwarfs such that $\\zeta$ overestimates the metallicity. Assuming all of the pairs are physically associated systems and have the same metallicity, we have redefined $\\zeta$. While the shift is small, it better represents iso-metallicity lines in the high-metallicity regime and represents an incremental step towards defining an absolute metallicity scale for low-mass dwarfs. Planned further observations should extend the calibration of $\\zeta$ as well as map it to an absolute metallicity scale in the near future."
    },
    "1112/1112.4481.txt": {
        "abstract": "Modified gravity is one of the most promising candidates for explaining the current accelerating expansion of the Universe, and even its unification with the inflationary epoch. Nevertheless, the wide range of models capable to explain the phenomena of dark energy, imposes that current research focuses on a more precise study of the possible effects of modified gravity may have on both cosmological and local levels. In this paper, we focus on the analysis of a type of modified gravity, the so-called $f(R,G)$ gravity and we perform a deep analysis on the stability of important cosmological solutions. This not only can help to constrain the form of the gravitational action, but also facilitate a better understanding of the behavior of the perturbations in this class of higher order theories of gravity, which will lead to a more precise analysis of the full spectrum of cosmological perturbations in future. ",
        "introduction": "In the recent years, {\\it modified gravity theories} have become one of the most popular candidates for explaining the current accelerated expansion of the Universe. As it is very well known,  general relativity (GR) in its standard form can not explain such behavior of the Universe expansion without either extra terms in the gravitational Lagrangian (for  reviews on modified theories of gravity, see Refs.~\\cite{review,book}) or exotic fluid components (see Refs.~\\cite{Ref1}-\\cite{Ref3}). % Modified gravity theories have been widely studied, and it has been shown that they are not only capable to mimic the dark energy epoch, but also the inflationary era \\cite{unification}. % Therefore, by the only use of large scale observations (Ia type supernova, BAO, or the cosmic microwave background) which depend uniquely on the evolution history of the Universe,  the nature and the origin of DE cannot be determined due to the fact that identical evolutions for the cosmological background can be explained by a diverse number of theories. This is the so called {\\it degeneracy problem}. It is thus required, in order to confirm or discard the validity of these theories, to obtaining solutions that can also describe correctly, e.g., the growth factor of scalar perturbations (see Refs.~\\cite{Perturbations}), the stability of cosmological solutions against small perturbations and the existence of GR-predicted astrophysical objects such as black holes \\cite{BH}.  \\\\ %In this sense research of the evolution of small perturbations around a given cosmological background may help shedding some light about the stability of proposed models and the viability of alternative gravity theories depending on the values of the parameters of the model under consideration.   In this sense the simplest, and in fact the most studied, modification of GR is the one where the Hilbert-Einstein action is generalized to a general function of the Ricci scalar $R$, the so-called $f(R)$ gravities  \\cite{Ref5}-\\cite{Ref6}. These theories are able to mimic the behavior of the cosmological constant (see for instance Ref.~\\cite{gr-qc/0607118}), but also can reproduce the entire cosmological history (see Ref.~\\cite{unification}).  In addition, they seem to behave quite well at local scales, where the GR limit must be recovered \\cite{Viable}, and the existence of GR-predicted astrophysical objects such as black holes can be achieved \\cite{BH_Dombriz}. Nonetheless, these candidates have their own shortcomings \\cite{Ref6} and have to pass rigorous theoretical and observational scrutiny before they can be accepted as viable theories \\cite{Viable}. % Another possible modification of the standard gravitational Lagrangian includes a wider number of curvature invariants ($R$, $R_{\\mu\\nu}R^{\\mu\\nu}, R_{\\mu\\lambda\\nu\\sigma}R^{\\mu\\lambda\\nu\\sigma}$ among others). Within these modifications, the so-called Gauss-Bonnet gravity can be included. In these theories the gravitational Lagrangian consists of a function $f(R,G)$ where $G$ holds for the usual Gauss-Bonnet invariant. % This modified gravity has been also widely studied and it is known that can also reproduce any kind of cosmological solution (see Refs.~\\cite{Cog1}-\\cite{N2}), where special attention has been already paid to models able to mimic the $\\Lambda$CDM model, as well as other important cosmological solutions (see Refs.~\\cite{arXiv:1001.3636, arXiv:1009.0902}). Finally, the cosmological perturbations have been explored within different standard scenarios for this class of theories \\cite{Felice_malo}. \\\\  In this investigation, we are interested in studying the stability behavior of several kind of cosmological solutions in the framework of Gauss-Bonnet gravities when subjected to homogeneous perturbations. % It is well known that any cosmological solution can be reconstructed in this class of theories,  but its behavior and stability is not well understood yet. % %%%%%%%%%%%%%% Pedazo a\u0096adido para contentar al referee, October 2012 %%%%%%%%%%%%%%%%%%%% Our analysis will therefore exclude anisotropic, i.e. cosmological scalar, perturbations.  Homogeneous and isotropic perturbations have been historically considered as the mean to determine the stability of different modified gravity theories (see for instance Refs. \\cite{Ref1,Barrow, arXiv:1011.2090}). % The usual approach in these references consisted in perturbing both the Hubble parameter and the matter density in order to check the background stability in the time evolution. This way, one can determine the stability of the modified Einstein equations when treated as differential equations. With respect of the cosmological perturbations, the full anisotropic analysis of cosmological perturbations in modified gravity theories is out of the scope of this investigation whereas some significant advances have been made in the last years, in particular for $f(R)$ theories \\cite{Perturbations}.\\\\  % [3] J. D. Barrow and A. C. Ottewill J. Phys. A 16 2757 (1983). % [4] A. Dobado and A. L. Maroto Phys. Rev. D 52 1895 (1995). % [5] D. Saez-Gomez, Phys. Rev. D 83 064040 (2011).  % 3 \\cite{Barrow} % 4 \\cite{Ref1} % 5 \\cite{arXiv:1011.2090}   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% With regard to Gauss-Bonnet gravity theories, %this kind of actions a particular subclass has been studied in Refs.~\\cite{arXiv:0810.5712,arXiv:0911.1811}, where the Hilbert-Einstein action plus a function $f(G)$ is considered. However,  the extension to a more general form for the $f(R,G)$ gravitational theory is a mandatory task that may help to understanding the viability and features of more general Lagrangians. %their cosmic evolution. % These theories may present ghost degrees of freedom in an empty anisotropic universe, i.e. the Kasner-type background \\cite{Felice_peor}. Nonetheless, these degrees of freedom are absent on Friedmann-Lema\\^itre-Robertson-Walker (FLRW) backgrounds (see also \\cite{Felice_peor}). This is precisely the type of cosmological background that will be considered in our investigation. % Hence, in the present paper we analyze some important cosmological solutions in $f(R,G)$ gravity, both in a vacuum scenario and with the presence of standard perfect fluids. The stability of cosmological solutions has been studied for $f(R)$ gravity in Ref.~\\cite{Barrow}, as well as for other curvature invariants in Ref.~\\cite{Ref1}. For some particular cases of $f(R,G)$, the perturbations on the solutions have been analyzed in Ref.~\\cite{Saltas:2010tt}, whereas for Ho\\v{r}ava-Lifshitz gravity, the analysis has been performed in Ref.~\\cite{arXiv:1011.2090}. Here we extend the analysis to more general actions, where we shall find the stability conditions for different cosmological evolutions in FLRW universes such as inflationary epoch and late-time accelerated era as described by the $\\Lambda$CDM model. % Therefore, these analyses can help understanding the viability of cosmological evolution provided by this kind of modified gravity, and constraining the viable candidates for the underlying gravitational action. \\\\  %We will show that inflation and dark energy epochs can be reproduced by these theories, paying special attention to LCDM model. % but  the exact $\\Lambda$CDM model presents some instability issues, which  restrict  this kind of actions to other dark energy solutions, and consequently to other type of cosmological evolution. \\\\  The paper is organized as follows: in Section II we present the general features of the $f(R,G)$ gravity theories by writing the corresponding modified Einstein equations. In Section III we introduce the evolution equations of perturbations appearing in these scenarios once that a FLRW background is assumed. Sections IV  and V are then devoted to the study of stability around the de Sitter and power-law solutions respectively. In the last case, we pay special attention to configurations including perfect fluids such as radiation and dust. Section VI is finally devoted to study the stability of the $f(R,G)$ model able to mimic the $\\Lambda$CDM cosmological evolution without any cosmological constant. % We conclude the paper by giving our conclusions in Section VII. An appendix is included at the end of the communication to provide explicitly the coefficients of perturbations equations introduced in Section III. ",
        "conclusions": "In this work we have extended the study of $f(R,G)$ modified gravity theories to provide an accurate description of the instabilities that these theories may present. % For such scenarios, the linearized perturbed equations have been derived. They were obtained once the modified Einstein equations are implemented with perturbations in both the Hubble parameter and matter density. The resulting coefficients have been explicitly presented  for the first time in the existing literature and a strong dependence on the chosen $f(R,G)$ model was observed.  We have studied three of the most important cosmological solutions in the standard cosmological concordance model around a spatially flat FLRW background: de Sitter expansion, power laws and the scale factor solution as provided for $\\Lambda$CDM model. The required $f(R,G)$ gravitational Lagrangians to provide de Sitter and power-law solutions are explicitly determined whereas the model mimicking $\\Lambda$CDM is revised . Concerning the perturbations, in the first case, we have found that a certain combination in the sign of some derivatives of the gravitational Lagrangian evaluated in the cosmological background is required to ensure the stability of the solution. This may provide a way to understand  the end of an inflationary era produced by one of these de Sitter points (which appear as natural solutions of $f(R,G)$ gravity), as well as for understanding the evolution of dark energy epoch. \\\\   With respect to the power-law solutions, the complexity of the coefficients for the perturbed equations led us to study two representative models:  either a sum of functions of scalar curvature and Gauss-Bonnet term or powered products of these two scalars. These models  may encapsulate the main features for cosmologically viable $f(R,G)$ theories in some asymptotic regimes. Their rich phenomenology was summarized in Table \\ref{Table_II}.  Besides, the usual power-law behavior for the scale factor in general relativity when either dust or radiation are the dominant components in the fluid sector, was found to be mimicked, even in the absence of such fluids, by appropriate choices in the parameter space of these models. % The radiation dominated evolution, i.e. $a(t)\\propto t^{1/2}$, deserves special attention in the case of actions of the type $f(R,G)=f(G)+f(R)$. For this case, we have shown that the background evolution requires the action reduces to Hilbert-Einstein action, a result that was also naturally extended to $f(R)$ theories. Consequently, the radiation dominated era can not be described in principle for this kind of Lagrangians, or at least,  effects of such terms should be negligible during that epoch. % % Concerning the evolution of the perturbations, we have shown that for models of the form $f(R,G)=f_1(G)+f_2(R)$, the perturbations oscillate tending to zero at $z=0$ provided that the initial condition $\\delta_{0}'$ is assumed to be null, and the background evolution ($\\alpha$) coincides with the EoS parameter of the perfect fluid. % Also note that independently of the model, the final values has a direct linear relation with the initial ones imposed. Note also that in the case $\\alpha=2$ (accelerating expansion), the perturbed equation presents a pole at a particular redshift, while the perturbations remain small up to the divergence.  \\\\   Finally, the case of the $f(R,G)$ model able to mimic the $\\Lambda$CDM scale factor led us as well to relevant conclusions. According to previous literature, this model presents two branches. Our analysis showed the decreasing and oscillatory character of perturbations in the positive branch regardless the initial conditions. Therefore, this solution may be considered as stable with respect to small perturbations. On the other hand, the negative branch showed that matter perturbations are not  constrained in amplitude even if the perturbations for the Hubble parameter approach zero for a wide choice of initial conditions.  Hence, the wide study on both cosmological solutions and stability of homogeneous perturbations for in $f(R,G)$ gravity carried out in this paper may help to a better understanding for such higher order theories of gravity. We have shown that gravitational action plays a very important role in the stability of the solutions depending both upon the form of the  $f(R,G)$ theory and the parameters of the model, which means a great  difference with respect to general relativity. The search for stability in widely accepted cosmological solutions helps to constrain the parameters space of the $f(R,G)$ models that may be viable and therefore may deserve further study by future analyses of the full spectrum of cosmological perturbations in modified gravity.   \\begin{table}[h] \\begin{center} \\begin{tabular}{|c|c|c|} \\hline \\hline \\multicolumn{3}{|c|}{Model $f(R,G)=f_1(G)+f_2(R)$} \\\\ \\hline \\hline $a(t)\\propto t^{\\alpha}$ & EoS & Stability \\\\ \\hline % \\multirow{2}{*}{$\\frac{2}{3}$}  & $w_m=0$ &  Stable with $\\dot{\\delta}_0=0$. $\\delta_m$ grows very fast when  $\\dot{\\delta}_0\\neq0$. \\\\ % % & $w_m=1/3$ & Unstable: $\\delta$ and $\\delta_m$ grow with $z$.  \\\\ % \\hline % % $2$  & $w_m=0$  & Existence of a pole $z_{pole}$ fixed at large redshift. Stable close to $z=0$. \\\\ \\hline \\end{tabular} \\end{center} \\begin{center} \\begin{tabular}{|c|c|c|c|} \\hline \\hline \\multicolumn{4}{|c|}{Model $f(R,G)=\\mu R^{\\beta}G^{\\gamma}$} \\\\ \\hline \\hline Configuration& $a(t)\\propto t^{\\alpha}$ & EoS & Stability \\\\ \\hline % \\multirow{2}{*}{Vacuum} & $\\alpha=\\frac{1}{2}$ & - & $\\beta>1$,   $\\gamma$ one or even number and $\\gamma<\\frac{1}{8}\\left(5-4\\beta\\right)$ \\\\ % % & $\\alpha=\\frac{2}{3}$   & -  & Any $\\beta$ $\\gamma$ one or even number and $\\gamma=\\gamma_{-}$ given by eqn. (\\ref{gammas}) \\\\ % \\hline % \\multirow{5}{*}{Fluid} & $\\alpha=\\frac{1}{2}$  & $\\omega=\\frac{1}{3}$ & Impossible to satisfy. \\\\ % % & $\\alpha=\\frac{2}{3}$  & $\\omega=0$  & Eqns. (47) and (48) must be satisfied,  \\\\ % % & &  & $\\delta_m$ decreases and $\\delta$ increases. \\\\ % % & $\\alpha\\neq\\frac{1}{2}$  & $\\omega= \\frac{1}{3}$  & Impossible to satisfy. \\\\ % % & $\\alpha=\\frac{4}{3}$  & $\\omega= 0$  & $\\beta=-1$, $\\gamma=\\frac{3}{2}$, $\\delta$ tends to zero and $\\delta_m$ increases. \\\\ % \\hline % \\end{tabular} \\end{center}  \\caption{ \\footnotesize{Summarization of the stability of power-law solutions for $f(R,G)$ models given by expressions (\\ref{2.9}) and (\\ref{product_power_laws}). In the cases where no fluids are considered (Vacuum) it can be seen that stability of the cosmological solutions can be achieved for appropriate choice in the parameters space. Once either dust of radiation fluids are considered, one sees that usual power-law evolution is on the one hand not feasible for radiation and, on the other hand, the stability depends on the initial conditions for dust. By the a particular action for this kind of cosmological solutions, modified gravity may contribute during the matter/radiation dominated eras, but also reproduce dark energy epoch($\\alpha=2$)}.} \\label{Table_II} \\end{table}   %"
    },
    "1112/1112.0505_arXiv.txt": {
        "abstract": "We combine a convectively driven dynamo in a spherical shell with a nearly isothermal density-stratified cooling layer that mimics some aspects of a stellar corona to study the emergence and ejections of magnetic field structures. This approach is an extension of earlier models, where forced turbulence simulations were employed to generate magnetic fields. A spherical wedge is used which consists of a convection zone and an extended coronal region to $\\approx1.5$ times the radius of the sphere.  The wedge contains a quarter of the azimuthal extent of the sphere and $150\\degr$ in latitude.  The magnetic field is self-consistently generated by the turbulent motions due to convection beneath the surface.  Magnetic fields are found to emerge at the surface and are ejected to the coronal part of the domain. These ejections occur at irregular intervals and are weaker than in earlier work. We tentatively associate these events with coronal mass ejections on the Sun, even though our model of the solar atmosphere is rather simplistic. ",
        "introduction": "Recent observations of the {\\it Solar Dynamics Observatory} \\citep[SDO:][]{PTC12} have provided us with a record of impressive solar eruptions. These eruptions are mostly associated with coronal mass ejections (CMEs). These are events through which the Sun sheds hot plasma and magnetic fields from the corona into the interplanetary space. The energy causing such huge eruptions is stored in the magnetic field and can be released {\\it via} reconnection of field lines \\citep{St80,ADK99}. Some of the CMEs are directed towards the Earth, hitting its magnetosphere and causing phenomena like aurorae. Furthermore, encounters with CMEs can cause sudden outages of GPS signals due to ionospheric scintillation. The resulting radiation dose from such events poses risks to astronauts. This is now also of concern to airlines, because the radiation load during polar flights can reach annual limits, especially for pregnant women. This leads to great interest of scientists in many fields of physics. However, there is an additional motivation which comes along with space weather effects. The solar dynamo, which is broadly believed to be responsible for the generation of the solar magnetic field, needs to be sustained by shedding magnetic helicity from the Sun's interior \\citep{BB03}. Mean-field and direct numerical simulations have shown that the magnetic field generation is catastrophically quenched at high magnetic Reynolds numbers in closed systems \\citep{VC92} that do not allow magnetic helicity fluxes out of the domain \\citep{BF00a,BF00b,BS04}, or between different parts of it \\citep{BCC09,MCCTB10,HB10}. The magnetic Reynolds number, which quantifies the relative importance of advective to diffusive terms in the induction equation, is known to be very large in the Sun, therefore implying the possibility of catastrophic quenching in models of the solar dynamo, unless efficient magnetic helicity fluxes occur, for example through CMEs \\citep{BB03}. Indeed, CMEs are well known to be closely associated with magnetic helicity \\citep{LO01}. In particular observations \\citep{PVSK00,RAK02} and a recent study by \\cite{TKT11}, where the observations are compared with numerical models, suggest that CMEs have a twisted magnetic structure, implying that CMEs transport helicity outwards.  There has been significant progress in the study of CMEs in recent years. In addition to improved observations from spacecrafts, \\textit{e.g.} SDO or the {\\it Solar TErestical RElation Observatory} \\citep[STEREO:][]{KKTD08}), there have also been major advances in the field of numerical modeling of CME events \\citep{RFG03,AHSGD09}. However, the formation and the origin of eruptive events like CMEs is not yet completely understood. Simulating CMEs and their formation is challenging. Leaving the difficulties of modeling the interplanetary space aside, a CME, after being ejected into the chromosphere or the lower corona, travels over an extended radial distance to the upper corona. In this environment, density and temperature vary by several orders of magnitude, which is not easy to handle in numerical models. Additionally, the origin of the CMEs is assumed to relate to the magnetic fields and the velocity pattern at the surface. However, the surface magnetic and velocity fields are rooted in the solar convection zone, where convective motions, in interplay with differential rotation, generate the magnetic field and the velocity patterns that are observed at the surface. The majority of researchers modeling CMEs do not include the convection zone in their setup, and thereby neglect the effect of the magnetic and velocity fields being rooted to this layer. Often the initial conditions for the magnetic and velocity fields are prescribed or taken from 2D observations; see for example \\cite{ADK99} and \\cite{ALML99} as well as \\cite{TK03}.  Another approach is to study the emergence of flux ropes from the lower convection zone into the corona. In the presence of strong shear, convection simulations have been showing the formation of flux tubes \\citep{GK11,Nelson11}, but such structures are similar to vortex tubes whose diameter is known to relate with the visco-resistive scale \\citep{BPS95}. In other approaches flux ropes are inserted in a self-consistent model, but their origin is left unexplained. In several recent papers \\citep{MS08,JB09,FMAH10}, the focus lies on the emergence of magnetic flux and the resulting features in the solar atmosphere. However, eruptive events have not been investigated with this setup. In earlier work \\citep[][hereafter WB]{WB10} a different approach was developed. The solar convection zone was combined with a simple model of the solar corona. The magnetic field, which was here generated by dynamo action beneath the solar surface, emerged through the surface and was ejected out of the domain. The focus was on the connection of the dynamo-generated field and eruptive events like CMEs through the dynamo-generated twist. WB used a simplified coronal model and drove the dynamo with forced turbulence. These simplifications allowed them to study the emergence and a new mechanism to drive ejections in great detail. In subsequent work \\citep[][hereafter WBM]{WBM11}, the setup of WB was improved by using a spherical coordinate system and helical forcing with opposite signs in each hemisphere to mimic the effects of rotation on inhomogeneous turbulence. In addition, WBM included the stratification resulting from radial gravity for an isothermal fluid. To improve this model, we now employ convection to generate the velocity field. In a related approach, \\cite{PB11} considered convective overshoot into the chromosphere and the excitation of gravity waves therein, but dynamo-generated twist seemed to be unimportant in their work. The turbulent motions driving the generation of magnetic field are now self-consistently generated by convective cells operating beneath the surface. The setup of the convection zone follows ideas of \\cite{KKB08}, \\cite{KKBMT10,KKGBC11} and \\cite{KMB11,KMB12}. There are other approaches simulating convection in hot massive stars, which have thin subsurface convection zones \\citep{CBBDKL11a}. But we now use an extended cooling layer to describe some properties of a solar corona. The results of this work complement those of earlier work and can be compared with observations. The model of the solar atmosphere is still a very simplified one, but can be regarded as a preliminary step, which will provide a reference point for improved work in that direction. ",
        "conclusions": "\\label{concl}  In the present paper we have presented an extension of the two-layer approach of WB and WBM by including a self-consistent rotating convection zone into the model. We find a large-scale magnetic field generated by the convective turbulent motion in the convection zone. At moderate rotation rates, for a Coriolis number larger than 3, we obtain a differential rotation pattern showing super-rotation, i.e., an equator rotating faster than the poles. The dynamo solutions we find are different and some of them have a periodic oscillatory behavior, where the large-scale magnetic field does not change sign; only the strength is varying. At the maxima, the velocity is suppressed due to the backreaction via the Lorentz force. Small-scale magnetic structures seem to show an equatorward migration near the equator and a poleward one near the poles.  Using a convectively driven dynamo complicates the generation of ejections into a coronal part due to lower relative kinetic helicity. However, it was possible to produce ejections in two of the runs. The shape and the bipolar helicity structure are comparable with those of WBM. Due to the relatively high plasma $\\beta$ in the outer parts of our model (compared with the solar corona), the ejections produce local minima of density which are carried along and ejected out to the top of the domain. The ejections occur recurrently, but not clearly periodically, which is similar to the Sun.  Note that our results have to be interpreted cautiously, given the use of a simplistic solar atmosphere. We neglect the effects of high temperature and low plasma $\\beta$. However, we feel that the mechanism of emergence of magnetic structures driven by dynamo action from self-consistent convection may not strongly depend on these two conditions. This suggestion has to be proven in more detail in forthcoming work.  An extension of the present work would require a detailed parameter study of cause and properties of the ejections. This also includes an advanced model for the solar corona with a lower plasma $\\beta$ and more efficient convection, which has a stronger stratification and is cooled by radiation. Another important aspect would be the generation of a self-consistent solar wind which supports and interacts with the ejections.  \\begin{acks} The authors thank Hardi Peter for discussion of the dynamics and rotation behavior of the solar corona. We also thank the anonymous referee for many useful suggestions. We acknowledge the allocation of computing resources provided by the Swedish National Allocations Committee at the Center for Parallel Computers at the Royal Institute of Technology in Stockholm, the National Supercomputer Centers in Link\\\"oping and the High Performance Computing Center North in Ume\\aa. Part of the computations have been carried out in the facilities hosted by the CSC  -- IT Center for Science in Espoo, Finland, which are financed by the Finnish ministry of education. This work was supported in part by the European Research Council under the AstroDyn Research Project No.\\ 227952, the Swedish Research Council Grant No.\\ 621-2007-4064, and the Academy of Finland grants 136189, 140970 (PJK) and 218159, 141017 (MJM) as well as the HPC-Europa2 project, funded by the European Commission - DG Research in the Seventh Framework Programme under grant agreement No.\\ 228398.  \\end{acks}"
    },
    "1112/1112.5044_arXiv.txt": {
        "abstract": "{Presented in this work is the most comprehensive INTEGRAL AGN sample: it includes 272 active galaxies for which optical identifications and redshifts are available.  2-10 keV fluxes and  column density measurements are also collected for all the sample sources. For 33 new hard X-ray discovered AGN with no previous X-ray coverage, XRT and XMM data analysis is reported for the first time in this work. Examples of future studies which are being developed using  this large X-ray selected sample are presented together with some early results. }  \\FullConference{The Extreme and Variable High Energy Sky - extremesky2011,\\\\ September 19-23, 2011\\\\ Chia Laguna (Cagliari), Italy}   \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.4568_arXiv.txt": {
        "abstract": " ",
        "introduction": "The difference between the mass-squares of different-generation neutrinos has been measured by atmospheric neutrino oscillation experiments with Super-Kamiokande \\cite{SK1,SK2},  neutrino beam experiments with K2K \\cite{K2K} and MINOS \\cite{MINOS1,MINOS2}, solar neutrino oscillation experiments with Super-Kamiokande \\cite{solar-SK} and SNO \\cite{SNO1,SNO2} and a nuclear reactor oscillation experiment with KamLAND \\cite{KamLAND1,KamLAND2}, but the neutrino mass itself has not been measured yet.  Detection of neutrino radiative decay enables us to measure a quantity independent of the difference between the mass-squares of different-generation neutrinos. Thus we can determine the neutrino mass itself from these two independent measurements, the neutrino oscillation and the neutrino radiative decay.  As the neutrino lifetime is so long as to be much larger than the age of the universe, the most promising method is to  observe the decay of the cosmic background neutrino. Detection of this radiative decay  means a discovery of the cosmic background neutrino predicted by standard cosmology.  The cosmic background neutrino has a temperature of 1.9K and a particle density $\\rho$ of 110 cm$^{-3}$ per generation.  Atmospheric neutrino data with Super-Kamiokande and neutrino beam data with K2K and MINOS gives the mass-square difference between $\\nu_3$ and $\\nu_2$, $\\Delta m_{32}^2$ of $( 2.43 \\pm 0.13 ) \\times 10^{-3}eV^2$  \\cite{pdg} and the mixing angle of sin$^{2}2\\theta _{23} > $ 0.92 at 90\\% confidence level \\cite{pdg}.  Solar neutrino data with Super-Kamiokande, SNO and the reactor neutrino data with KamLAND give the mass-square difference between $\\nu_2$ and $\\nu_1$ of $\\Delta m_{12}^2 = ( 7.59 +0.19/ -0.21 ) \\times 10^{-5} eV^2$ \\cite{pdg} and the mixing angle of sin$^{2}2\\theta_{12} = 0.87 \\pm 0.03$ \\cite{pdg}.  We search for the neutrino radiative decay under these constraints on the mass-square differences and the mixing angles between the neutrino generations. For this search, we  make a fit of the photon energy spectrum of Cosmic Infrared Background (CIB) to a sum of a neutrino decay photon energy spectrum and a continuum spectrum.  In the left-right symmetric model, we make a feasibility study of the neutrino decay search in the CIB energy spectrum measurement. ",
        "conclusions": "We propose to search for the neutrino radiative decay by fitting a photon energy spectrum of the cosmic infrared background to a sum of the photon energy spectrum from the neutrino radiative decay and a continuum. By comparing the present cosmic infrared background energy spectrum observed by AKARI and Spitzer to the photon energy spectrum expected from neutrino radiative decay with a maximum likelihood method, we obatined a lower limit of netrino lifetime of 3.1 $\\times 10^{12}$  to 3.8  $\\times 10^{12}$ years at 95 \\% confidence level for $\\nu_3$  in the mass range between 50 \\mmev\\ and 150 \\mmev\\ under the present constraints by neutrino oscillation measurements.  In the left-right symmetric model, the lifetime is predicted to be much shorter than in the standard model. The present mass limit  of the right-handed weak boson $W_R$ is $m_R >$ 0.715 TeV and the mixing angle between $W_L$ and $W_R$ is sin$\\zeta <$ 0.013 .  In this model, the minimum lifetime of $\\nu_3$ is predicted to be $1.5 \\times 10^{17}$ years  for $m_3$ of 50 \\mmev.  The present lifetime limit is shorter than this prediction by  a factor of 2 $\\times$ $10^{-5}$. By measuring the continuous energy spectrum of the cosmic infrared background  in the photon energy region between 5 meV ($\\lambda$ = 250 $\\mu$m) and 35 meV ($\\lambda$ = 35 $\\mu$m), we  can expect  5$\\sigma$ observation of the radiative decay of neutrino with a lifetime of $1.5 \\times 10^{17}$ years.\\\\   \\begin{Large} \\bf{Acknowlegements}\\\\ \\end{Large}  We thank M. Yamauchi and Y. Okada at KEK for discussions regarding the neutrino decay and its lifetime calculation. This work was supported by the Ministry of Education, Science, Sports and Culture of Japan. Part of this work was supported by KAKENHI(19540250 and 21111004).   {\\normalsize"
    },
    "1112/1112.3348_arXiv.txt": {
        "abstract": "{We have obtained Herschel images at five wavelengths from 100 to 500 \\mums of a $\\sim 5.5\\times2.5$ degree area centred on the local galaxy M31 (Andromeda), our nearest neighbour spiral galaxy, as part of the Herschel guaranteed time project  ``HELGA''. The main goals of HELGA are to study the characteristics of the extended dust emission, focusing on larger scales than studied in previous observations of Andromeda at an increased spatial resolution, and the obscured star formation.} {In this paper we present data reduction and Herschel maps, and provide a description of the far-infrared morphology, comparing it with features seen at other wavelengths.} {We use high--resolution maps of the atomic hydrogen, fully covering our fields, to identify dust emission features that can be associated to M31 with confidence, distinguishing them from emission coming from the foreground Galactic cirrus.} {Thanks to the very large extension of our maps we detect, for the first time at far-infrared wavelengths, three arc-like structures extending out to $\\sim 21$, $\\sim 26$ and $\\sim 31$ kpc respectively, in the south-western part of M31. The presence of these features, hosting $\\sim 2.2\\times 10^6$ M$_\\odot$ of dust, is safely confirmed by their detection in HI maps. Overall, we estimate a total dust mass of $\\sim5.8\\times 10^7$ M$_\\odot$, about 78\\% of which is contained in the two main ring-like structures at 10 and 15 kpc, at an average temperature of 16.5 K. We find that the gas-to-dust ratio declines  exponentially as a function of the galacto-centric distance, in agreement with the known metallicity gradient, with values ranging from 66 in the nucleus to $\\sim 275$ in the outermost region.} {Dust in M31 extends significantly beyond its optical radius ($\\sim 21$ kpc) and what was previously mapped in the far-infrared. An annular--like segment, located approximately at R$_{25}$, is clearly detected on both sides of the galaxy, and two other similar annular structures are undoubtedly detected on the south-west side even further out.} ",
        "introduction": "An overall view of the interstellar medium of the Milky Way is hampered by the huge column densities that any line of sight has to go through. Andromeda is the closest galaxy offering broadly similar characteristics and an unencumbered view, thus probably being the best target to test our knowledge of the physical processes that govern the formation and evolution of massive spiral galaxies. In fact, high quality data of M31 are available at a wealth of different spectral ranges: X-rays \\citep{pietsch05}, ultraviolet \\citep{thilker05}, optical (e.g. SDSS, H$\\alpha$ by \\citealt{devereux94} or the PAndAS survey, see e.g. \\citealt{mcconnachie08}), near and far-infrared (FIR) \\citep{haas98,barmby06,gordon06,beaton07,krause11}, millimeter \\citep{nieten06} and radio \\citep[e.g.][]{thilker04,braun09}. Despite this, the observational coverage in the sub-millimeter domain was until now limited by the absence of data longward of  170 \\mum, except for the extremely low resolution (with a beam size larger than $40'$) data from the Diffuse Infrared Background Experiment (DIRBE) at 240 \\mums \\citep{odenwald98}.  This situation has left M31 unexplored at far-infrared and sub-mm wavelengths in adequate spatial resolution. The advent of the {\\it Herschel} Space Observatory \\citep{pilbratt10}, with its two instruments PACS (Photodetector Array Camera and Spectrometer) \\citep{poglitsch10} and SPIRE (Spectral and Photometric Imaging Receiver) \\citep{griffin10}, allows us to fill this observational gap and complete the view of our neighbour galaxy down to sub-mm wavelengths. The other important advantage of {\\it Herschel} is the gain in terms of spatial resolution which, at 160 \\mum, is about 2.5 times higher than that of the Multiband Imaging Photometer for {\\it Spitzer} (MIPS) at similar wavelengths, making it possible to resolve physical scales of about 50 pc at the distance of Andromeda.  On the one hand, M31 is considered in many respects to be very similar to the MW, sharing some characteristics  and similarities such as the Hubble type, the luminosity, and gas content in the respective disks. The fact that it is so close makes it not only the ideal candidate for a comparison with similar studies and ongoing surveys taking place with {\\it Herschel} on our own Galaxy \\citep[such as Hi-Gal,][]{molinari10}, but also the ideal place for studying the interstellar medium of a spiral galaxy. Hinting at some common aspects between the MW and M31, \\cite{yin09} find very similar gas mass fraction profiles (similar scale lengths but different normalizations) and abundance gradients in the two disks.  On the other hand, some differences are evident between the two galaxies: not only does M31 have about twice the baryonic mass of the MW, it has also a larger disk, by a factor of about 2.4. The mass of gas contained in the respective disks, however, is remarkably similar. An overview of the main properties of the two galaxies can be found in \\cite{yin09} (see their Table 1).  So far, all the studies of M31's far-infrared Spectral Energy Distribution (SED), mainly relying on data shortwards of $\\sim160$ \\mum, have speculated about the possible presence of a very cold component ($\\sim 10$ K). Such a component would have a SED that will peak at wavelengths longer than 200 \\mums \\citep{haas98,montalto09}. Due to the wavelength limitations of the data, no constraints could be put on its presence and, even less, its characteristics. The estimates of the dust mass were therefore believed to be only lower limits to the real value. With {\\it Herschel} data we are in the position of setting more stringent constraints both to the presence and spatial distribution of a very cold dust component, and its emissivity.  The data we exploit in our survey allow us to also study the dust morphology in the very outer regions of the galaxy. The far-infrared morphology of Andromeda has been the topic of several studies exploiting, over the years, the available infrared facilities such as the Infra Red Astronomical Satellite \\citep[IRAS,][]{habing84,xu96}, the Infrared Spatial Observatory \\citep[ISO,][]{haas98} and {\\it Spitzer} \\citep{barmby06,gordon06} each facility improving in both sensitivity and spatial resolution. Evidence for a complete dusty ring at a radius of $\\sim 10$ kpc came from the very first  IRAS observations, while following studies exploiting higher resolution were able to detect details such as a discontinuity, most likely caused by an interaction with the satellite galaxy M32 \\citep[see, e.g.][]{gordon06}. A fainter ring-like structure at $\\sim 15$ kpc was already noted by \\cite{haas98} and confirmed by \\cite{gordon06}, but nothing more could be said on larger radii because of the size limitations of the maps.  While dust located beyond the regions sampled by such observations was indeed detected indirectly by studying the characteristics of the stellar populations in the extreme outskirts of M31 \\citep[e.g.][]{cuillandre01,bernard12}, it had never been possible to study its properties and quantify its presence before.  In this work we describe {\\it Herschel} PACS and SPIRE observations of M31, at 100, 160, 250, 350 and 500 \\mum, extending out to about 34 kpc ($\\sim2.5^\\circ$) from its centre, presenting the most complete far-infrared/submillimeter survey of Andromeda. We describe the observations and the data reduction steps in \\S \\ \\ref{sec:datared}, while in \\S \\ \\ref{sec:ancillary} we give a general overview of the ancillary dataset that we exploit in our project, together with a comparison to {\\it Herschel} maps. In \\S \\ \\ref{sec:results} we use high and low-resolution {\\sc Hi} maps, entirely covering our fields, to disentangle the dust dynamically connected to Andromeda from the foreground Galactic cirrus, and in \\S \\ \\ref{sec:dust} we use these foreground--corrected maps to study the dust properties. Conclusions and remarks are given in \\S \\ \\ref{sec:conclusions}.  Throughout this paper and in all works within the HELGA project, we will assume a distance to Andromeda of $785$ kpc \\citep{mcconnachie05}, which turns into a linear scale of 3.8 pc/$''$. As for the inclination and position angles, we are using $i=77^\\circ$ and PA$=38^\\circ$, respectively.  \\begin{figure*} \\centering \\includegraphics{fritz12_01.pdf} \\caption{Images in the five {\\it Herschel} bands (from top to bottom: PACS 100 and 160 \\mum, SPIRE 250, 350 and 500 \\mum). These are the original images, where no correction for the foreground cirrus was applied (see Sec. \\ref{sec:identification}). The green rectangles in the 250 \\mums map represent tentative detections of dust in the outskirts (see discussion in Sec. \\ref{sec:fir_struct}). The directional axes length corresponds to $15'$ ($\\sim 3.42$ kpc).} \\label{fig:herschelimages} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusions} We have presented and described {\\it Herschel} PACS (100 and 160 \\mum) and SPIRE (250, 350 and 500 \\mum) observations of a $\\sim 5.5^\\circ \\times 2.5^\\circ$ field centered on the local galaxy M31. Using high spatial and velocity resolution {\\sc Hi} maps, roughly covering the same area, we have identified the dusty structures that surround the galaxy outside the already well-observed $\\sim 3^\\circ \\times 1^\\circ$ region \\citep{haas98,gordon06,barmby06,krause11}.  We have performed a broad analysis of the characteristics of the dust as a function of the radius. We have considered 5 main regions, described by projected circular apertures (plus an elliptical one), sampling the most significant morphological structures: the inner 1.5 kpc ring, the intra-arm region enclosed within the 10--kpc ring, the 10--kpc ring itself, the outer ($\\sim15$ kpc) ring, and a fifth one, containing the outermost regions. A simple one--temperature modified black body model is able to successfully reproduce the observed PACS/SPIRE SED. Average quantities, such as the dust temperature, and the gas--to--dust ratio, are derived, and they are found to strongly depend on the radius. The gas--to--dust ratio, in particular, is changing by a factor of $\\sim 2$, following an exponential increase, from the inner regions to the outskirts, as expected from the presence of a metallicity gradient with the opposite sign, across the galaxy.  According to our simple emission model the two main rings, at 10 and 15 kpc from the centre, contain the bulk (about 78\\%) of the dust in all the galaxy. The dust surface density is the highest in within the 10 kpc one, 0.41 M$_\\odot$/pc$^2$, being higher by a factor of $\\sim 4$ with respect to the average value, and also higher than in the 15 kpc ring, where it is about $0.25$ M$_\\odot$/pc$^2$.  The dust temperature is highest in the innermost aperture, but the value we derive from our SED fitting, 22 K, is considerably lower than previous studies  \\citep[e.g.][derive temperatures of $\\sim 28$ and $\\sim 30$ K, respectively, for the nuclear dust]{haas98,tabatabaei10}. This is most likely due to the fact that, ignoring data shortwards 100 \\mum, we are not sensitive to the presence of a warmer component \\citep[e.g.][]{bendo10}: this would, anyway, constitute a minor fraction of the total dust content.  We will give a more detailed description of the IR morphology in a forthcoming paper \\citep{kirk12}, where an analysis of the structures on deprojected {\\it Herschel} maps will be performed.  The maps we describe in this work present two distinct situations: in the north--eastern part, the combination of Andromeda's systemic and rotation velocities, yielding values that are similar to those expected for the interstellar medium in our galaxy, together with the presence of a substantial amount of gas and dust, makes the identification of emission regions that actually belong to Andromeda very difficult, and sometimes ambiguous. On the other hand, at the southern side, not only is the velocity of gas and dust that belong to Andromeda expected to have a clearly different pattern with respect to that of our Galaxy, but is also much less affected by the presence of interstellar medium (see e.g. Fig. \\ref{fig:ident_paper} for a panoramic view of the M31 neighbourhood).  In this region, we identify in our maps three different ring-shaped, concentric structures -- which we label E, F and G --  that can be safely attributed to M31, based on the shape and velocity of their {\\sc Hi} counterparts, extending out to a radius of $\\sim 31$ kpc. Dust in the outskirt of Andromeda had been already indirectly detected, through extinction of both stars and background galaxies, by \\cite{cuillandre01}, who exploited deep optical photometry in a $\\sim 28' \\times 28'$ field roughly centred on our ``E'' region. This result, recently confirmed on a smaller scale by \\cite{bernard12} (their ``Outer Disc field''), implies an atomic gas--to--dust ratio about three times higher with respect to the average Galactic value, and is considered as an upper limit since it only takes into account the dust in front of the observed stars. Instead, the {\\sc Hi}--to--dust ratio we estimate for this particular region, $\\sim 204$, is a factor of $\\sim 1.5$ larger than the 137 value estimated for the MW  \\citep[see Table 2 in][]{draine07}, that one would expect given the roughly solar metallicity of stars observed in a nearby field. This value can provide a more stringent constraint to the amount of dust in this region, and to the chemical enrichment of the interstellar medium in Andromeda's outskirts."
    },
    "1112/1112.4797_arXiv.txt": {
        "abstract": " ",
        "introduction": "The existence of magnetic fields in the universe has long been established by observations on small upto very large scales, that is on the scales of galaxies, galaxy cluster and superclusters \\cite{obsgal}. Recent observations even indicate evidence for magnetic fields at truly cosmological scales \\cite{obscosmag}. Over recent years there has been an increasing interest in cosmological magnetic fields, from generation mechanisms to observational imprints on the cosmic microwave background (CMB) (for recent reviews, e.g., \\cite{reviews}).  For a primordial magnetic field to serve as an initial seed field to explain the galactic magnetic field, it is assumed to be generated either in the very early universe, such as due to the amplification of perturbations in the electromagnetic field during inflation (e.g. \\cite{tw}), or after inflation has ended, e.g., during a phase transition such as the electroweak or QCD phase transitions in which case it has been shown that the resulting magnetic fields are helical \\cite{pt}.  The generation of helical magnetic fields during inflation has been studied in \\cite{infhel}. However, in any case it is natural to assume that the magnetic field existed long before initial conditions are set for the evolution of the perturbations determining the temperature anisotropies and polarization of the CMB. Thus observations of the CMB could in principle be used to put limits on the parameters of a primordial magnetic field. The aim here is to calculate the angular power spectra of the temperature anisotropies and polarization for a Gaussian stochastic helical magnetic field.  In this case the magnetic field two-point correlation contains a symmetric and an asymmetric part. Each of which are characterized by one amplitude and one spectral index. The helical part results in non-vanishing correlations of the temperature anisotropies and the polarization B-mode on the one hand and of the polarization E-mode and B-mode of the CMB on the other hand. These have been first considered in \\cite{pvw}, where the vector modes were studied. The spectra of the CMB anisotropies due to the tensor mode have been treated in \\cite{cdk} and vector and tensor modes in \\cite{kr}.  Moreover, the helical part also contributes to all other correlation functions, that is the temperature and polarization E-mode and B-mode autocorrelation functions as well as the temperature polarization E-mode cross correlation of the CMB. Here  a consistent numerical treatment is presented taking into account the magnetic field in the initial conditions and the evolution equations as well as calculating numerically the correlation functions encoding the contribution of the magnetic field.  Using a Gaussian window function the magnetic field spectrum is effectively cut-off at a damping scale. Moreover, the angular power spectra of the temperature anisotropies and polarization of the CMB are calculated for scalar, vector and tensor modes.  The calculation of the CMB temperature anisotropies and polarization requires to solve the evolution equations of the perturbation equation of the geometry and matter components together with the Boltzmann equations for photons and neutrinos. Since the first numerical program COSMICS \\cite{cosmics} to do this and with a significant improvement in speed using the line-of-sight integration in CMBFAST \\cite{cmbfast} there has been a steady evolution of numerical codes such as CAMB \\cite{camb}, CMBEASY \\cite{cmbeasy}  and CLASS \\cite{class}  being the latest addition. The effect of a primordial nonhelical magnetic field on the CMB anisotropies has been calculated using different approaches: synchronous gauge and thus  a modified version of CMBFAST in \\cite{grant,GK}  or the covariant formalism and a modified version of CAMB in \\cite{fin,le}. Finally using the gauge-invariant formalism for the scalar perturbations the CMB anisotropies have been calculated in the presence of a stochastic nonhelical magnetic field using CMBEASY in \\cite{kek}.  Here the modified version of CMBEASY \\cite{kek} has been  expanded to include firstly the numerical calculation of the correlation functions of a helical stochastic Gaussian magnetic field and secondly a new part to solve the Boltzmann equation and calculate the CMB temperature anisotropies and polarization for vector and tensor modes. These are calculated using the total angular momentum approach of Hu and White \\cite{hw}. In section \\ref{sect2} the decomposition of the magnetic field contribution for scalar, vector and tensor modes is described. Section \\ref{sect3} is devoted to the calculation of the relevant correlation functions of  the contribution of a Gaussian stochastic helical magnetic field. As it will be shown there are contributions due to the helical part of the spectrum to the correlation functions for all three modes, namely, scalar, vector and tensor modes. In section \\ref{sect4} the perturbation equations for the vector and tensor modes in the presence of a magnetic field are presented in the gauge-invariant formalism \\cite{ks}. The corresponding equations for the scalar mode can be found in \\cite{kek}. Moreover, the initial conditions are presented. In section \\ref{sect5} results of the numerical calculation of the angular power spectra determining the autocorrelation and cross correlation functions of the temperature (T) and polarization, that is the E-mode (E) and B-mode (B),  of the CMB are presented. Due to the helical nature of the magnetic field apart from the autocorrelation functions of T, E and B and the temperature polarization E-mode cross correlation TE there are also non vanishing correlations between E and B as well as  T and B. The latter one, in particular, is compared for a choice of parameters with the WMAP7 data \\cite{wmap7}. Section \\ref{sect6} contains the conclusions. ",
        "conclusions": "\\label{sect6}  The CMB anisotropies and polarization in the presence of a stochastic helical magnetic field have been calculated for scalar, vector and tensor modes. For this purpose the modified version of CMBEASY  \\cite{kek} calculating  the CMB anisotropies due to the scalar perturbations in the presence of a non helical magnetic field has been expanded. Firstly the numerical solution for the corresponding correlation functions has been included for scalar, vector and tensor modes. Secondly a new part has been added to include the calculation of the CMB anisotropies and polarization due to vector and tensor modes using the total angular momentum approach of Hu and White \\cite{hw}. A Gaussian window function is used to effectively cut-off the magnetic field spectrum at a wave number corresponding to the magnetic damping scale. In the case of the scalar and tensor perturbations the initial conditions for the numerical solution which are set long after neutrino decoupling include a contribution encoding the evolution of relevant quantities due to the presence of the magnetic field before neutrino decoupling. In the case of the scalar perturbations the comoving curvature perturbation grows upto the time of neutrino decoupling due the magnetic anisotropic stress. After neutrino decoupling the neutrino anisotropic stress compensates the magnetic anisotropic stress and the comoving curvature perturbation becomes a constant on superhorizon scales \\cite{kkm,le,BC}. In the case of the tensor modes a similar behaviour is found for the amplitude of the tensor mode which has been shown here explicitly. The presence of a magnetic field prior to neutrino decoupling does not affect significantly the evolution of the vector modes. Using a standard Boltzmann solver the contribution due to presence of the magnetic field before neutrino decoupling has been included in the initial conditions for  the numerical solution which are set long after neutrino decoupling.  However, since the presence of a magnetic field affects the evolution of the scalar and tensor perturbations before neutrino decoupling in order to be more precise one would have to start the numerical evolution and thus set the initial conditions for the Boltzmann code before neutrino decoupling. A problem we hope to address in the future.  In the case of a helical magnetic field in addition to the temperature (T) and polarization E- and B- mode autocorrelation spectra and cross correlation TE angular power spectrum there are also the cross correlation EB and TB angular power spectra. The latter one has been used to illustrate the use of the current WMAP7 data to constrain the magnetic field parameters.  In comparison to earlier work on the effects of a helical magnetic field \\cite{pvw,cdk,kr} on the CMB here a full numerical treatment has been provided including the correct initial conditions and evolution equations in the presence of a magnetic field as well as including numerical solutions for the different correlation functions of the magnetic field contributions. Moreover, the magnetic field spectrum  is effectively cut-off using a Gaussian window function."
    },
    "1112/1112.6101_arXiv.txt": {
        "abstract": "It is well known that a large fraction of galaxies have cuspy luminosity profiles in their central regions, at least within the observational resolution. In such cases, the often used, simplified, local approximation for the dynamical friction braking classical term fails when the massive satellite moves through the inner parts of the galaxy, although the scattering integral still converges for phase space  distribution singularities that are not too sharp. Here we present preliminary results of our work aiming at finding better and more reliable results from the integration of motion of massive objects (globular clusters) in galaxies where the density diverges to the center in a power law form, with exponent greater than $-2$. ",
        "introduction": "The frictional deceleration of the motion of a $test$ particle of mass $M$ placed, at time $t$, at ${\\mathbf r}$ is given by the cumulative effect of a background of $target$ stars of mass $m$  at varying the impact vector ${\\mathbf b}$, which can be expressed as  \\begin{equation} \\frac{d{\\mathbf v}_M}{dt} = -{{2m}\\over {M+m}} \\int_{{\\mathbf b}}\\int_{{\\mathbf v}_m} f(\\mathbf{r+b,v}_m) \\frac{{\\mathbf v}_M-{\\mathbf v}_m}{1+\\frac{b^2 |{\\mathbf v}_m-{\\mathbf v}_M|^4}{G^2(M+m)^2}}{|{\\mathbf v}_M-{\\mathbf v}_m| \\over b} d^3\\mathbf{v_m}d^3\\mathbf{b}, \\label{dfcorrect1} \\end{equation} where $f(\\mathbf{r+b,v}_m)$ is the stellar background distribution function (DF) at position $\\mathbf{r}'=\\mathbf{r+b}$. Even for a DF inspace and velocity spherical simmetry, i.e. $f(r',v_m)$, the integral in Eq.(\\ref{dfcorrect1}) is very complicated whenever the test particle is out of the origin ($r\\neq 0$).  \\begin{figure} {\\includegraphics[width=6.5cm]{./ArcaSedda_M_Fig1.eps}} \\hspace{0.5cm} {\\includegraphics[width=6.5cm]{./ArcaSedda_M_Fig2.eps}} \\caption{The time evolution of the distance to the galactic center of a massive test object in the case of initially circular orbits (regular curves) and radial orbits (oscillating curves) of same initial orbital energy, for two values of $\\gamma$, as labeled, and two values of the object mass (upper curves refer to $M=10^6$ M$_\\odot$. lower curves to $M=4 \\times10^6$ M$_\\odot$).} \\label{fig1} \\end{figure}   To overcome, partially, this problem, the common approach consists in taking out of the integral in Eq. (\\ref{dfcorrect1}) the DF, assumed as independent of ${\\mathbf b}$ and isotropical in the velocity space, thus yielding to what is often referred as {\\it local} approximation formula \\cite{bt}  \\begin{equation} \\left({ \\frac{d{\\mathbf v}_M}{dt} }\\right)_{loc}= -4 \\pi G^2 M \\rho(r) F(r,v_M) \\ln\\Lambda {{{\\bf v}_M} \\over {v_M^3}}, \\label{ads} \\end{equation} the function $0\\leq F(r,v_M)\\leq 1$ being the local fraction of target stars slower than the test object ($v_m \\leq v_M$). In the local approximation, the mass density, $\\rho$, of field stars in Eq. (\\ref{ads}) is assumed at its {\\it local} value, i.e. that evaluated at the position (${\\mathbf r}$) of the test particle; $\\ln\\Lambda$ is the usual Coulomb's logarithm.  Now, whenever the test particle is significantly off center with respect to the stellar system (star cluster, galaxy, etc.) wherein it moves,  the local expression gives (for symmetry reasons) a quite acceptable approximation; on the contrary, it loses its validity in the neighbourhood of the centre. In this case the local approximation is clearly an overestimate of the actual dynamical friction, because it corresponds to weighting the contribution of the gravitational encounters with objects at a given distance from the test particle not with the, correct, density of target stars at that distance but, rather, with the density of targets evaluated at the location of the test particle itself, that is maximum at the origin of any self-gravitating system. This overestimate is particularly relevant when dealing with {\\it cuspy} galaxies, where the spatial density of stars diverges at the galactic center, making the local approximation unviable. \\\\ To solve this problem, we follow the approach presented in \\citet{vicetal} and deeply discussed in \\citet{cdas11}), which is based on taking into account that at the origin ($r=0$) and in the hypothesis of isotropic DF the integral in Eq. (\\ref{dfcorrect1}) simplifies into the form  \\begin{equation} \\left({ \\frac{d{\\mathbf v}_M}{dt} }\\right)_{cen} = -{{8\\pi m}\\over {M+m}} \\int_{b_{min}}^{b_{max}}\\int_{{\\mathbf v}_m} f(b,\\mathbf{v}_m) \\frac{{\\mathbf v}_M-{\\mathbf v}_m}{1+\\frac{b^2 |{\\mathbf v}_m-{\\mathbf v}_M|^4}{G^2(M+m)^2}}|{\\mathbf v}_M-{\\mathbf v}_m|b d^3\\mathbf{v_m} db. \\label{dfcorrect2} \\end{equation} which can be performed numerically, although with the care due to the double singularity in phase space. Once that the integral in Eq. (\\ref{dfcorrect2}) is evaluated, an estimate of the dynamical friction deceleration along the test object trajectory is given by an interpolation between the local approximation as given by Eq. \\ref{ads} and the central value given by Eq. (\\ref{dfcorrect2}), in the form  \\begin{equation} \\frac{d{\\mathbf v}_M}{dt} = p(r)\\left({ \\frac{d{\\mathbf v}_M}{dt} }\\right)_{cen}+ [1-p(r)] \\left({ \\frac{d{\\mathbf v}_M}{dt} }\\right)_{loc} , \\label{dfinterp} \\end{equation} where the interpolation function $0\\leq p(r)\\leq 1$ is such that $p(0)=1$. We found $p(r)=e^{-r/R}$ (where $R$ is a proper length scale) as a good interpolation formula. We have applied this to the orbital evolution of massive objects in a spherical potential resembling that of the Milky Way (disk excluded), as given by the sum of three components: i) a cuspy inner bulge, represented as  a Dehnen's $\\gamma$-model \\citep{deh}, ii) an extended bulge-halo represented as a Plummer sphere \\citep{plu} and, iii) a dark matter halo in the Navarro, Frenk and White form \\citep{nfw}. ",
        "conclusions": ""
    },
    "1112/1112.3809_arXiv.txt": {
        "abstract": "We have embarked upon a project to model the UV spectra of BALQSOs using a Monte Carlo radiative transfer code previously validated through modelling of the winds of cataclysmic variable stars \\citep[e.g.][]{2010ApJ...719.1932N}. We intend to use the simulations to investigate the plausibility of geometric unification \\citep[e.g.][]{2000ApJ...545...63E} of the different classes of QSO. Here we introduce the code, and present some initial results. These demonstrate that for reasonable geometries and mass loss rates we are able to produce synthetic spectra which reproduce the important features of observed BALQSO spectra. ",
        "introduction": "Broad absorption line quasars (BALQSOs) are a class of QSO exhibiting broad, blue shifted absorption lines of highly ionized species such as N~{\\sc v}, C~{\\sc iv}, O~{\\sc vi} and Si~{\\sc iv}. These objects make up perhaps 17\\% of the population of QSOs \\citep[e.g.][]{2008MNRAS.386.1426K, 2003AJ....125.1784H}. Their unusual UV spectra are believed to be the result of fast (0.01-0.2c) winds originating from the accretion disk \\citep[e.g.][]{1992ApJ...401..529K}, although the geometry and kinematics of these winds are not well understood, and even the driving mechanism is unclear. It is a source of discussion as to whether this observation demonstrates that all QSOs spend 17\\% of their lifetime as a BALQSO (`evolutionary unification'; e.g. \\citealt{2000ApJ...538...72B}), or whether all QSOs have these observational signatures when observed from 17\\% of possible directions (`geometrical unification';  e.g. \\citealt{2000ApJ...545...63E}). Our aim in developing a code to simulate UV spectra for quasar disk winds is to permit quantitative evaluation of prospective geometries for geometrical unification. ",
        "conclusions": "We have begun a project to use Python, a Monte Carlo radiative transfer code, to simulate the UV spectra of QSO winds. We present initial results for a geometry suggested by \\cite{2000ApJ...545...63E}, illuminated by a disk only, and observe that the simulated spectra exhibit features reminiscent of those seen in BALQSOs.  At present, for the model parameters we have tried, the wind becomes over-ionized whenever a strong central X-ray source is included, and none of the expected UV lines are observed.  We are incorporating a dense shielding region into the geometry to see if this produces BAL-like spectra even with a strong central X-ray source.  We will produce synthetic spectra for a range of geometries, viewed from different sight-lines to and so inform the discussion of the evolutionary vs orientation interpretation of the BALQSO phenomenon."
    },
    "1112/1112.3038_arXiv.txt": {
        "abstract": "We perform a global fit to $\\sim\\! 5,000$ radial velocity and $\\sim\\!177,000$ magnitude measurements in $29$ photometric bands covering $0.3\\,\\mu$m to $8.0\\,\\mu$m distributed among $287$ Galactic, LMC, and SMC Cepheids with $P>10$\\,days. We assume that the Cepheid light curves and radial velocities are fully characterized by distance, reddening, and time-dependent radius and temperature variations. We construct phase curves of radius and temperature for periods between $10$ and $100$\\,days, which yield light curve templates for all our photometric bands and can be easily generalized to any additional band. With only $4$ to $6$ parameters per Cepheid, depending on the existence of velocity data and the amount of freedom in the distance, the models have typical rms light and velocity curve residuals of $0.05$\\,mag and $3.5$\\,km\\,s$^{-1}$. The model derives the mean Cepheid spectral energy distribution and its derivative with respect to temperature, which deviate from a black body in agreement with metal-line and molecular opacity effects. We determine a mean reddening law towards the Cepheids in our sample, which is not consistent with standard assumptions in either the optical or near-IR. Based on stellar atmosphere models we predict the biases in distance, reddening, and temperature determinations due to the metallicity and we quantify the metallicity signature expected for our fit residuals. The observed residuals as a function of wavelength show clear differences between the individual galaxies, which are compatible with these predictions. In particular, we find that metal-poor Cepheids appear hotter. Finally, we provide a framework for optimally selecting filters that yield the smallest overall errors in Cepheid parameter determination, or filter combinations for suppressing or enhancing the metallicity effects on distance determinations. We make our templates publicly available. ",
        "introduction": "The classical Cepheids play a key role in the extragalactic distance determination, both as a direct way of obtaining distances of nearby galaxies and as a calibrator for other methods (e.g.\\ the review by \\citealp{freedman10a}, and \\citealp{madore91,feast97,freedman01,benedict07,riess09a,riess09b,riess11a,riess11b}). This is enabled by the relatively tight relation between the pulsational period of a Cepheid and its luminosity  (the $PL$ relation). Accurate distance estimates using Cepheids require increasingly better understandings of systematic effects such as extinction, composition (metallicity) and blending.  For example, the light from Cepheids is typically extinguished by dust, which can have different properties in different galaxies \\citep[e.g.][]{cardelli89,laney93,falco99,misselt99,motta02,gordon03,indebetouw05,marshall06,laney07}. Cepheid observations must be corrected for extinction before any subsequent analysis. This usually makes use of an assumption that Cepheids with a given period have a specific color and that there is a known, universal extinction law. This holds for the Baade-Wesselink and related methods as well as purely photometric analyses. To alleviate these shortcomings, recent efforts have shifted from the optical (Hubble Key Project) to the infrared wavelengths ($H$, $I$ and Spitzer IRAC), where the effects of reddening are reduced \\citep[e.g.][]{mcgonegal82,gieren05,ref167,freedman08,madore09,ngeow10,storm11a,storm11b}. There are, however, differences in the near/mid-infrared extinction law for different sightlines in the Galaxy \\citep[e.g.][]{roman07,nishiyama09}.  Cepheid properties must also change as a function of chemical composition, but the associated observational signatures have been a matter of considerable debate. While the general conclusion of observational studies is that metal-rich Cepheids are brighter than their metal-poor counterparts \\citep[e.g.][]{stothers88,kochanek97,sakai04,macri06,valle09,shappee11}, results from stellar pulsation models and some observations yield the opposite dependence \\citep[e.g.][]{fiorentino02,fiorentino07,marconi05,bono08,romaniello08,freedman11}. Estimates of the metallicity correction vary by almost an order of magnitude, as can be seen in Figure~14 of \\citet{gerke11}. It is also now clear, at least in the extragalactic context, that uncertainties in the appropriate metallicities and metallicity gradients are nearly as important as the actual metallicity dependence of Cepheid parameters \\citep[see][and references therein]{bresolin11,gerke11,shappee11}. Furthermore, studies find different trends in different filters \\citep{bono08,bono10,ngeow11,storm11b}. Many of these problems arise because the effects of metallicity have filter-dependent degeneracies with reddening and distance, making it crucial to fully recover all covariances in the final results, as emphasized by \\citet{gould94}, \\citet{kochanek97}, \\citet{sasselov97}, and \\citet{riess11a}. As a result, the question of metallicity effects on the Cepheid distance scale remains a controversial issue. The problem of blending then adds further complications for more distant galaxies \\citep{mochejska00}.  In many circumstances, light curve templates for the individual observational bands are needed to accurately determine the mean magnitudes of extragalactic Cepheids because of their sparse light curves. \\citet{stetson96} constructed $V$- and $I_{\\rm C}$-band templates for the LMC, SMC, and Galactic Cepheids with periods between $7$ and $100$\\,days by fitting Fourier series to the light curves. He assumed that the amplitudes and phases of individual modes vary continuously with period, thus forming the well-known Hertzsprung sequence \\citep{hertzsprung26}. \\citet{hendry99}, \\citet{ngeow03}, \\citet{tanvir05} and \\citet{yoachim09} expanded on these models using Principal Component Analysis (PCA) techniques to build new templates for analysis of HST Cepheid data. All of these templates are limited in the sense that they are defined only for a small set of filters (typically $V$ and $I_{\\rm C}$), and there is no well-defined means of shifting them to other, similar filters, let alone to very different wavelengths. \\citet{freedman10b} show that all light curves are related by a linear transformation which provides a path towards addressing this problem.  In this work, we address these issues by constructing a global model of Cepheid light curves and radial velocity curves in the Galaxy, LMC, and SMC. We self-consistently determine distances, reddenings, radii, and temperatures of individual Cepheids along with a global reddening law, the mean SED of a Cepheid and its dependence on temperature, and the phase variations of the radii and temperatures as a function of period. We self-consistently determine the uncertainties in {\\it all} parameters including {\\it all} their covariances. Our model enables us to construct a light curve template for an arbitrary filter given just a single calculable parameter for the relative contribution of radius and temperature variations to that filter. We evaluate metallicity effects on Cepheid observations first from a theoretical point of view, and then we search for metallicity effects in the estimated parameters and residuals of the fit.  The approach to the problem is closely related to the Baade-Wesselink method \\citep{baade26,wesselink46} and assumes that the magnitude of a Cepheid at a given time depends only its distance, reddening, and instantaneous radius and temperature. Rather than carrying out a study of individual Cepheids, we simply fit a global model to all the available photometric and velocity data. A simpler variation on this idea was recently proposed by \\citet{freedman10b}. This paper is organized as follows. In Section~\\ref{sec:model_data}, we outline our physical model, priors, dataset and the method of fitting. In Section~\\ref{sec:results}, we describe results of the fit, discussing in turn the fit residuals, global quantities and individual properties of the Cepheids. In Section~\\ref{sec:discussion}, we discuss our results and search for signs of additional physics, principally metallicity effects. We review our findings and outline future directions in Section~\\ref{sec:conclusions}. An Appendix outlines a general quantitative approach to selecting filters for Cepheid studies. ",
        "conclusions": "\\label{sec:conclusions}  In this work, we present a method for self-consistently determining distances, reddenings, radii, and temperatures of individual Cepheids along with the mean reddening law and phase-dependent variations of radius and temperature as a function of period. Our approach is in some senses a global version of the Baade-Wesselink method and it is in some senses an implementation of the ideas of \\citet{freedman10b}. Our method provides statistically consistent error estimates for all parameters of the model (Section~\\ref{sec:model_data}). We fit our physical model to $\\sim\\! 5,000$ radial velocity and $\\sim\\! 177,000$ magnitude measurements in 29 bandpasses covering wavelength range from $0.3\\,\\mu$m to $8\\,\\mu$m for $287$ Cepheids from the Galaxy, LMC, and SMC. With only four to six variables per Cepheid, depending on the availability of velocity data and the amount of freedom in the distance estimates, we fit the complete phase dependent data set with a magnitude scatter of $0.051$\\,mag and radial velocity scatter of $3.5$\\,km\\,s$^{-1}$, as shown in Figures~\\ref{fig:residuals} and \\ref{fig:residuals_irac}.  Our templates are built to span the period range $10 \\le P \\le 100$\\,days and can be used to produce model light curves or $PL$ relations for any of the $29$ bands. Predicting the light curve for an uncalibrated band is straightforward provided that the value of a single parameter $\\beta_i$, the logarithmic flux derivative with respect to the temperature (Eq.~[\\ref{eq:beta_general}], Figure~\\ref{fig:beta}), can be calculated or estimated for the new filter. The templates do not provide perfect fits, particularly for periods between $10$ and $13$\\,days where strong resonance effects are present. The mean phase radius and temperature variations we obtain match the Hertzsprung progression in the longer period stars.  We obtain a spectral energy distribution (Figure~\\ref{fig:mbar}) and the logarithmic flux derivative with respect to the temperature $\\beta_i$ (Figure~\\ref{fig:beta}) for a ``mean'' Cepheid. When we compare our results with estimates based on the black body law or theoretical stellar atmosphere models, we completely rule out black body models. We find better agreement with the \\citet{castelli04} model atmospheres, but also strong evidence for problems in the atmosphere models \\citep[see][]{fremaux06,hauschildt99a,hauschildt99b}. The coefficients $\\beta_i$, which correspond to the logarithmic derivative of the spectrum with respect to temperature, deviate from the black body law due to metal line blanketing and molecular opacities again in general agreement with theoretical atmosphere models. We are even able to resolve several spectral features. However, the theoretical atmosphere models are inadequate in describing the infrared spectrum, probably indicating problems with the molecular opacities in the models. There are no difficulties separating temperature and reddening in our models (Figures~\\ref{fig:beta_r} and \\ref{fig:tem_ext}) up to degeneracies created by the effect of composition.  We obtain a very high precision measurement of the mean extinction law for $29$ filters from $0.3\\,\\mu$m to $8\\,\\mu$m. The mean reddening law of Cepheids in our sample (Figure~\\ref{fig:extinction}) shows departures from the $\\mathscr{R}_V = 3.3$ \\citet{cardelli89} law usually assumed for Cepheids. In part this is due to averaging over the three galaxies, but it clearly illustrates that extinction laws need to be considered in efforts to increase the precision of Cepheid distances. This is true despite the trend towards infrared wavelengths, because the near-IR $\\mathscr{R}$ values have larger fractional uncertainties. Our results for the extinction curve also let us directly evaluate the Wessenheit factor $\\mathcal{R}_{VI_{\\rm C}} = \\mathscr{R}_V/(\\mathscr{R}_V-\\mathscr{R}_{I_{\\rm C}})$ commonly used in Cepheid studies \\citep[e.g.][]{madore82,madore91,madore_freedman09,ngeow05}. For $\\rv\\equiv 3.3$ $(3.1)$ we find that $\\mathcal{R}_{VI_{\\rm C}}= 2.521 \\pm 0.005$ ($2.475\\pm 0.005$), where the error comes only from the determination of $\\mathscr{R}_{I_{\\rm C}}$, because we hold $\\mathscr{R}_V$ fixed. The result for $\\rv\\equiv 3.3$ is significantly different from the usually assumed $2.45$ \\citep{wozniak96,freedman01} and introduces a distance shift of about $\\Delta\\mu \\simeq 0.01$\\,mag for a Cepheid with $E(V-I_{\\rm C})=0.10$\\,mag. More generally, we can calculate the Wessenheit factors for any combination of filters and we can correctly include the mutual covariances between the filters in the error estimate. For example, \\citet{riess11a} assumed the optical/near-IR Wessenheit factor of  $\\mathcal{R}_{HVI_{\\rm C}} = \\mathscr{R}_H/(\\rv-\\mathscr{R}_{I_{\\rm C}}) = 0.410$ to reduce the uncertainty on the measurement of the Hubble constant. From our analysis, we find $\\mathcal{R}_{HVI_{\\rm C}} = 0.322\\pm 0.006$ ($0.290 \\pm 0.006$) again at fixed $\\rv\\equiv 3.3$ ($3.1$), inconsistent with the assumptions by \\citet{riess11a} in either case. This illustrates that extinction laws need to be better understood to correctly measure precise distances in the universe, particularly since these estimates still all assume $\\rv\\equiv 3.3$ (or $3.1$) and so underestimate the uncertainties.  We find weak evidence of metallicity effects in the fit residuals. They are weak because much of the metallicity effect can be mimicked by shifts in the distance, reddening, and temperature of a Cepheid (Fig.~\\ref{fig:msmc-mgal}). The effect of metallicity on stellar atmospheres  that is orthogonal to these parameters only leads to residuals $\\lesssim 0.05$\\,mag for a metallicity difference of $0.5$\\,dex (Fig.~\\ref{fig:metal_residuals}) which is approaching the regime where we may also need to consider absolute calibration differences arising from the heterogeneous data sets. The clearest effect of metallicity is seen in the temperature distributions of the Cepheids, because with $29$ bands the temperature is well constrained independently of the extinction and distance/radius. The SMC Cepheid temperatures are typically higher than those of the LMC, which are in turn higher than for the Galaxy, although the distributions overlap. This can be explained as a combination of projecting the metallicity effects on stellar atmospheres into changes in temperature as well as shifts in the period--temperature relation. The shifts we observe are roughly consistent with expectations. There may also be small shifts in radius (Fig.~\\ref{fig:deltar_deltat}), but for our sample this is difficult to disentangle from the effects of distance priors and would be better addressed given samples of Cepheids truly at a common distance and with a range of metallicities. Similarly, we find small differences in the zero-points of our $PL$ relations, which can also be attributed to the projection of metallicity effects, but this has many of the same limitations as the search for shifts in radius. We do find shifts in the $PL$ relation slopes, but our limited period range ($P \\ge 10$\\,days) is not well suited for a robust investigation of slopes.  In Appendix~\\ref{app:opti} we outline a procedure to select optimal filter sets for disentangling distance, temperature, reddening and metallicity or to maximize or minimize metallicity effects when one is concerned solely with distance determination. We find that sets of filters spanning the longest wavelength range possible, for example $UV$[4.5], yield the best fit results for parameter estimation. Such filter combinations produce parameter uncertainties $0.5$ to $1$ order of magnitude smaller than the $BVI_{\\rm C}$ combination originally planned for the HST Key Project. We find that the maximum effect on distance determination due to metallicity can be expected for filter combinations like $UBV$ and $BVR_{\\rm C}$. On the other hand, the slope of the dependence of distance on metallicity can be reduced to less than $10^{-3}$\\,mag\\,dex$^{-1}$ for filter combinations like $VHK$ or $I_{\\rm C}$[3.6][4.5]. However, some filter combinations with infrared filters are quite sensitive to metallicity (e.g.\\ $UVJ$). For these metrics, there are no particular benefits to the mid-IR over the near-IR, so Cepheid studies in the JWST era might better focus on the near-IR where Cepheids are brighter and the PSF will be more compact (to reduce systematic errors from crowding and blending). Also, mid-IR Cepheid measurements might be affected by circumstellar dust \\citep[see, e.g.][]{kervella06,gallenne11}. These statements assume the validity of the \\citet{castelli04} model atmospheres, which we did find are partly inconsistent with the Cepheid data, so filter choices should be evaluated for a broader range of atmosphere models.  Our goal in this paper was simply to carry out a complete analysis with the simplest possible model. That the model does so well helps to explain why efforts to identify higher-order corrections such as metallicity are so challenging. At least based on the \\citet{castelli04} atmosphere models, the next most important terms are the second derivative of the zero point with respect to the temperature followed by the first-order metallicity correction and the metallicity correction to the logarithmic flux derivative with respect to the temperature. The second derivative of the zero point with respect to metallicity, the first order surface gravity derivative  and the dependence of the extinction correction $\\ri$ on temperature are relatively unimportant unless working in the UV. There is also a clear metallicity dependence to the mean temperature, which also indicates the need for higher-order terms in the period-radius and period-temperature relations. These estimates of higher-order effects are again based on the theoretical atmosphere models of \\citet{castelli04}.  There is clearly a broad range of possible future extensions. Adding additional data or galaxies is trivial, as is extending to shorter periods. It is also possible to fit for signatures of a companion to the Cepheid which will distort the SED. The virtue of our approach, particularly in a cosmological setting, is that it forces a correct use of prior assumptions in the parameter estimation and uncertainties, allows the data to overrule those assumptions, if necessary, and produces error estimates incorporating the full uncertainties of the entire model. Improved stellar atmosphere models would lead to better priors, although this is less of a concern given the demonstrated ability of the data to constrain the SED of Cepheids and their dependence on temperature. Particularly with better atmosphere models it should be possible to relax the assumptions about extinction ($\\rv\\equiv 3.3$ and mean $\\langle E(B-V)\\rangle =0.147$\\,mag for the LMC based on \\citealt{udalski99}) and put the Cepheids on their own absolute extinction scale. This would particularly help in finding the expected environmental dependencies of the extinction law and examining their effects on distance uncertainties.  The final important lesson from this study, as also previously stated by \\citet{freedman10b}, is that there is no need to obtain a light curve in order to add almost all the information available from observing at a new wavelength. Given one band with enough data to determine the period and phase of the Cepheid at the time of any observation at another band, a single observation in this new band suffices because the model correctly includes all the phase correlations between the bands. In practice, it would be wise to obtain at least two observations so as to have an internal check on the results. This removes any major barrier to systematically using the short wavelength filters that provide the greatest leverage for the control and measurement of systematic errors. Much of the drive towards longer wavelengths has been that with lower amplitudes the light curves can be more poorly sampled to still yield an accurate mean magnitude.  With our models, {\\it any} band can be poorly sampled provided its phase is well-determined. In some sense, this is a multi-band version of the method used by \\citet{gerke11}, where Cepheids in M81 were identified and phased with ground based images, and then calibrated using single epochs of HST data combined with the \\citet{stetson96} template models. Furthermore, when fitting an individual Cepheid, all the bands can be simultaneously co-phased by using our models and measuring the goodness of fit to all the data as a function of period."
    },
    "1112/1112.0936_arXiv.txt": {
        "abstract": "Some studies have claimed the existence of a stellar upper-mass limit of 150$M_{\\odot}$. A factor that is often overlooked concerns the issue that there might be a significant difference between the present-day and the initial mass of the most massive stars -- as a result of mass loss. The upper-mass limit may be substantially higher, possibly exceeding 200$M_{\\odot}$. The issue of the upper mass-limit will however remain uncertain as long as there is only limited quantitative knowledge of mass loss in close proximity to the Eddington (=$\\Gamma$) limit. For this reason, we present mass-loss predictions from Monte Carlo radiative transfer models for very massive stars up to 300$M_{\\odot}$. Using our new dynamical approach, we find an upturn or ``kink'' in the mass-loss versus $\\Gamma$ dependence, at the point where our model winds become optically thick. These are the first mass-loss predictions where the transition from optically {\\it thin} O-star winds to optically {\\it thick} Wolf-Rayet winds has been resolved. ",
        "introduction": "Stellar winds from massive O and Wolf-Rayet (WR) stars are ubiquitous and form an important aspect of both stellar and galaxy evolution. Despite recent reductions in O-star mass-loss rates by factors of $\\sim$3 in comparison to unclumped empirical rates, population synthesis models by Voss et al. (2009) -- utilizing rotating stellar models such as those of Meynet \\& Maeder (2003) that employ theoretical Vink et al. (2000) rates that match the factors $\\sim$3 reduced rates -- show that the role of stellar winds in massive-star feedback has {\\it increased} over the last decade. The reason for this possibly counter-intuitive finding is that stellar rotation has increased luminosities, whilst the objects also enter the WR-phase sooner. Strong stellar winds may prevent the most massive stars from entering the eruptive Luminous Blue Variable (LBV) phase, at least at Galactic metallicity. It is these very massive stars (VMS) with their strong stellar winds that are thought to dominate the radiative and kinetic energy input in young ($<$10 Myr) clusters, as well as the radio-active 26-Al input that gives rise to gamma-ray lines from Galactic star-forming regions like Orion and Carina (Voss et al. 2010, 2012).  Our understanding of stellar mass loss from the most massive stars is also crucial for the determination of their final fates, with respect to the formation of intermediate mass-black holes (Belkus et al. 2007; Yungelsen et al. 2008), the occurrence of pair-instability supernovae (Langer et al. 2007), and the mass distribution of neutron stars and black holes (Heger et al. 2003; Eldridge \\& Vink 2006)  A most relevant issue concerns the stellar upper mass limit. Until recently, many studies accepted the existence of an upper limit of 120-150$M_{\\odot}$, but a recent study by Crowther et al. (2010) claimed much higher luminosities and masses for the central WN5h objects of the R136 cluster than previous estimates (de Koter et al. 1997). One of the potential problems with the Crowther et al. luminosities is that the stars are clustered, and whilst Crowther et al. took great care in the assessment of a potential binary nature of the objects, there is a non-negligible chance that the stars are ``polluted'' with the light from line-of-sight objects.  In the context of the VLT Flames Tarantula Survey, we have recently discovered and analysed a new WN5h object, VFTS 682, a near-identical twin of R136a3 (Evans et al. 2011, Bestenlehner et al. 2011). This object is in apparent isolation from the R136 cluster, some 30\\,pc away, which raises interesting questions about its origin (see Bestenlehner et al. for a discussion on the object having formed in isolation or being a slow runaway). It also allowed us to assess the reliability of the Crowther et al. luminosity claims for the WN5h stars in the central cluster. Our finding of $\\log L/L_{\\odot} = 6.5 \\pm 0.2$ for the ``isolated'' object VFTS\\,682 seems to support the claim for high luminosities of WN5h stars in the Tarantula region. As VFTS\\,682 is in relative isolation from the dense cluster, we argue that the chance of a line-of-sight coincidence is much lower than for the WN5h stars in the core of R136, and we thus have independent evidence that the ``traditional'' 120 or 150$M_{\\odot}$ stellar upper-mass limit has been broken. The upshot is that mass-loss rates for stars with masses above 150$M_{\\odot}$ are needed to establish the evolution and fate of the most massive stars. Here, we present predictions from 60 up to 300$M_{\\odot}$. In the VMS regime, stars are identified as late-type nitrogen-rich WNh stars. These objects are extremely close to the Eddington limit $\\Gamma~=~g_{\\rm rad}/g_{\\rm grav}~=~\\kappa L/(4 \\pi c G M)$. ",
        "conclusions": "We presented mass-loss predictions from Monte Carlo radiative transfer models for VMS in the mass range 40-300$M_{\\odot}$\\ and with Eddington factors $\\Gamma$\\ in the range 0.4--1.0. Our models suggest that the spectral transition from Of to WN corresponds to a transition from relatively low $\\Gamma$\\ to high $\\Gamma$ values (and larger velocity law parameter $\\beta$) for WN stars. This assertion is not only based on the numerically larger mass-loss values, but more specifically on the finding of a ``kink'' at $\\Gamma$ $=$ 0.7 where the mass-loss slope in function of $\\Gamma$ shows an upturn."
    },
    "1112/1112.0107_arXiv.txt": {
        "abstract": "Silicon Photomultipliers (SiPM) represent a promising alternative to classical photomultipliers for the detection of photons in high energy physics and medical physics, for instance.~In the present work, electrical characterizations of test devices - manufactured by STMicroelectronics - are presented.~SiPMs with an area of $3.5 \\times 3.5\\,\\textrm{mm}^2$ and a cell pitch of $54\\,\\mu$m were manufactured as arrays of $64 \\times 64$ cells and exhibiting a fill factor of 31\\%.~The capacitance of SiPMs was measured as a function of reverse bias voltage at frequencies ranging from about 20\\,Hz up to 1\\,MHz and temperatures from 310\\,K down to 100\\,K.~Leakage currents were measured at temperatures from 410\\,K down to 100\\,K.~Thus, the threshold voltage - i.e., the voltage above a SiPM begins to operate in Geiger mode - could be determined as a function of temperature.~Finally, an electrical model capable of reproducing the frequency dependence of the device admittance is presented. ",
        "introduction": "In recent years Silicon Photomultipliers (SiPM) has been developed for usage in photon detection.~The high gain, insensitivity to magnetic field and low reverse bias voltage of operation make SiPM a promising candidates as replacement of classical photomultipliers\\cite{intro} in several of their applications.~The SiPM is an array of parallel-connected single photon avalanche diodes (SPAD)\\cite{Maz1}\\,.~Every SPAD operates in Geiger mode with a quenching resistor in series to prevent an avalanche multiplication process from taking place.~The overall output of a SiPM depends on how many SPADs are simultaneously ignited\\cite{intro1}$^{,}$\\cite{Maz2}\\,.~ \\par The current SiPM device was manufactured by STMicroelectronics (details of the technology can be found in Refs.\\cite{Maz1,Condorelli}).~It consisted of an array of $64 \\times 64$ SPAD cells with $3.5 \\times 3.5\\,\\textrm{mm}^2$ effective area, a cell pitch of 54$\\,\\mu$m and a fill factor of $\\approx 31$\\%.~ \\par In the present article, the electrical characteristics of these devices are shown as function of test frequency from 1\\,MHz down to about 20\\,Hz and temperature from 410\\,K down to 100\\,K (Sects.~\\ref{Elec_Cha_SiPM}--\\ref{Cur_V_SiPM}).~In addition, in Sects.~\\ref{Elec_Mod} and~\\ref{Elec_Mod_SiPM} an electrical model is discussed and compared with data obtained from dependencies of capacitance and resistance on the test frequency for both a photodiode and a SiPM devices. \\begin{figure}[b] \\begin{center} \\vskip -.3cm \\includegraphics[width=0.8\\textwidth]{SiPM32bis_CV.eps} \\vskip -.3cm \\caption{Capacitance (in nF) as a function of reversed bias voltage (in V) using the LCZ meter with test frequencies of 100\\,Hz ($\\bullet$), 1\\,kHz ($\\blacksquare$), 10\\,kHz ($\\circ$) and BC with 1\\,MHz ($\\square$).}\\label{cv1} \\end{center} \\end{figure} \\begin{figure}[t] \\begin{center} \\includegraphics[width=0.8\\textwidth]{SiPM32bis_CF.eps} \\vskip -.3cm \\caption{Capacitance (in nF) as a function of the frequency (in Hz) for a sample operated with a reverse bias of 5\\,V: data points indicated with $\\bullet$ ($\\circ$) were obtained using the LCR (LCZ) meter, the continuous line is obtained from Eq.~(\\ref{susceptance_array}) (see discussion in Sect.~\\ref{Elec_Mod_SiPM}).}\\label{cv2} \\end{center} \\end{figure} \\begin{figure}[t] \\begin{center} \\vskip -.3cm \\includegraphics[width=1.0\\textwidth]{CVT.eps} \\vskip -.3cm \\caption{Capacitance response (in nF or pF) obtained using the LCZ meter as a function of reverse bias applied (in V) and temperature from 310\\,K down to 100\\,K: 100\\,Hz (upper left), 1\\,kHz (upper right), 10\\,kHz (bottom left) and - using the BC - 1\\,MHz (bottom right).}\\label{cvt} \\end{center} \\end{figure} ",
        "conclusions": "\\label{SiPM_Conclusion} The electrical characteristics of SiPM devices were investigated as a function of applied reverse bias voltage ($V_{\\rm r}$) and temperature from 410 down to 100\\,K.~The capacitance response was also studied as a function of test frequency.~One finds that the measured capacitance decreases i) with increasing the test frequency above 10\\,kHz at 300\\,K and ii) with the decreasing of the temperature.~Furthermore, the measurement of the leakage current allowed one to determine the value (and its temperature dependence) of the threshold voltage ($V_{\\rm th}$) above which a SiPM begins to operate in Geiger mode: $V_{\\rm th}$ decreases with lowering temperature at the rate of $\\approx - 29 \\,$mV/K from 300 down to 100\\,K. \\par It was developed an electrical model to treat the frequency dependence of a SiPM device operated under reverse bias voltage.~This model is based on the so-called small-signal ac impedance of junction diode (SIJD) used for radiation detectors and photodiodes.~The model was found to be well suited to account for the SiPM dependencies of capacitance and resistance on frequency. \\vspace{-.3cm}"
    },
    "1112/1112.5664_arXiv.txt": {
        "abstract": "We investigate the expected neutrino emissivity from nuclear collisions in a magnetically dominated model of gamma-ray bursts motivated by recent observational and theoretical developments. The results indicate that small multi-GeV neutrino fluxes are expected for model parameter values which are typical of electromagnetically detected bursts. We show that for detecting at least one muon event in Icecube and its Deep Core sub-array, a single burst must be near the high end of the luminosity function, and at low redshifts $z\\lesssim 0.1$, where the burst rate is very low. We also calculate the luminosity and distance ranges that can generate $0.01-1$ muon events per GRB in the same detectors, which may be of interest if simultaneously detected electromagnetically, or if measured with future extensions of Icecube or other neutrino detectors with larger effective volume and better sensitivity. ",
        "introduction": "\\label{sec:intro}  Observations of multi-GeV photons from gamma-ray bursts (GRB) recently accumulated by the {\\it Fermi} satellite (e.g. \\cite{Abdo+09grb080916c, Ackermann+10grb090510}) have pointed out the need to re-evaluate the type of models used to explain the prompt photon emission mechanisms and the location of the emission regions in these objects (e.g. \\cite{Peer11-fermigrb}).  In particular, concerns about the radiative efficiency of usual internal shock models, and the larger radii required to avoid two-photon degradation of the spectra have spurred the investigation of baryonic (non-MHD) jet models where the radiation arises in a jet photosphere \\citep{Meszaros+00phot, Rees+05photdis,Beloborodov10phot}. In such baryon-loaded jet models, at a certain radius the timescale of the nuclear elastic collisions (which couple the proton $p$ and neutron $n$ components) becomes longer than the expansion timescale, i.e. the collision optical depth falls below $o(1)$, and the two components decouple from each other \\citep{Derishev+99,Bahcall+00pn}. The protons can continue to be accelerated by the radiation, while the neutrons, which have zero electric charge, start to coast with a constant Lorentz factor. Starting at this decoupling radius and for some distance beyond, the longitudinal drift velocity between the $n$ and $p$ becomes $\\Delta v \\gtrsim 0.5 c$, and they start to collide inelastically.  Such large relative velocities between $n$ and $p$ components also can arise in realistic jets where the bulk Lorentz factor $\\Gamma$ depends on the polar angle $\\theta$, which also leads to inelastic collisions \\citep{Meszaros+00gevnu, Beloborodov10phot}, as neutrons from the outer parts (sheath) of the jet thermally drift into the jet core. In both pictures, pions are created, which in the case of the dynamics being dominated by the baryons results in multi-GeV photons and neutrino production \\citep{Bahcall+00pn,Beloborodov10phot}.  A different approach towards resolving the radiative efficiency of GRBs involves consideration of magnetically dominated jets \\citep{Giannios+07photspec, Tchekhovskoy+10grb,Meszaros+10pop3}. Some of these magnetic models assume an almost baryon-free outflow \\citep{Usov94,Meszaros+97poynting,Meszaros+10pop3, Toma+11pop3,Gao+11pop3nu}, while in other cases a  substantial but dynamically sub-dominant baryon load is assumed \\citep{Thompson94,Drenkhahn+02,Lyutikov+03grbmag, Giannios+07photspec,Tchekhovskoy+10grb,McKinney+11magphot,Metzger+11grbmag,Giannios07pn}. The baryons in such jet models are expected to accelerate at a different rate than in baryonic (non-MHD) jet models, and the different dynamics leads to quantitatively different predictions for the photon spectrum \\citep{Meszaros+11gevmag}. Similarly, it should lead to quantitatively different neutrino spectra, which we investigate in this paper.  Unlike in the previous investigations cited above, here we consider the neutrino spectra arising from nuclear collision effects in magnetically dominated GRB outflows.  In this case, both the radial $np$ decoupling  radius as well as the photosphere occur at larger radii from the central engine  than in the non-magnetic case, and also the transverse drift of neutrons from the periphery of the jet into the jet core becomes important at different radii, where the physical conditions differ from those previously considered. As a consequence, multi-GeV photons are produced at somewhat softer energies and with appreciable time delays \\citep{Meszaros+11gevmag} respect to the MeV photospheric photons, but the detailed neutrino spectrum for such magnetically dominated jets has not been considered so far. In this paper we investigate numerically the neutrino spectrum expected in magnetically dominated baryon-collisional GRB models, taking into account both longitudinal and transverse $n,p$ decoupling and inelastic collisions.  These neutrinos are in the energy sensitivity range of Icecube and its DeepCore \\citep{DeYoung+11RICAP,IcecubeDC+09,IcecubeICRC+11} sub-array.  In \\S \\ref{sec:col} we briefly introduce the astrophysical model and present the method of neutrino emission calculation. In \\S \\ref{sec:nuflux} we present the results for the expected neutrino fluxes and muon events, as well as the detection prospects with Deep Core and IceCube. A discussion and summary of the results is given in \\S \\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  We have calculated the neutrino emission in the range of a few GeV to a few hundred GeV arising in magnetized collisional GRB models, such as have been recently used for interpreting the electromagnetic properties of these objects. The neutrino emission considered here arises partly from longitudinal proton and neutron collisions following `their decoupling in the jet, and in larger part from collisions caused by transverse thermal drift of neutrons from an outer jet sheath into the jet core, having different bulk Lorentz factors. This neutrino production model differs from the commonly considered photohadronic $p\\gamma$ models, and it also differs from previous $pn$ model calculations in incorporating explicitly the (dominant) transverse drift collision effects. Also, the emission region characteristics are here determined by the magnetically dominated dynamics, differing from those in previous neutrino calculations, which mostly use baryonic dynamics.  Furthermore, the $pn$ neutrino spectra are calculated numerically, using two different codes which are suited for the GeV range energies considered.  Our present results indicate that for the burst parameters suitable to explain the photon spectra and the MeV-GeV photon lags \\cite{Meszaros+11gevmag} indicated by the recent {\\it Fermi} satellite observations, a low level of neutrino emission is expected at $\\sim 5-100$ GeV energies. This is in the sensitivity range of the Deep Core sub-array of IceCube, and extends into the lower range of the main IceCube array.  For a neutrino detection of an individual burst, unaided by a coincident electromagnetic detection, one would require $\\gtrsim 1$ muon event. For neutrino-induced muon track events in IceCube and Deep Core, considering the angular resolution (e.g. $\\Omega \\lesssim 10^2~{\\rm deg}^{2}$) and time search bins $\\sim o(1)\\times20s$ per burst, the atmospheric neutrino background is low.  During the total sum of the search bins (e.g. assuming 300 electromagnetically observed long GRBs in a year, from the atmospheric neutrino spectrum the expected number of background muon events is $\\lesssim o(1)\\times0.09$.  However, any GRBs which might be expected to yield $\\gtrsim1$ muon event need to be at the high end of the luminosity function, and located at very low redshifts, $L_{\\gamma,iso} \\gtrsim 10^{53}$ erg/s and $z\\lesssim 0.1$.  From the discussion of \\S \\ref{sec:nudiff}, the occurrence of such GRB is estimated to be very rare, $\\lesssim 1/17$ per year.  This estimate is uncertain, because of the poor statistics in the determination of the redshift distribution in this range.  Such bursts would in general also be detectable by photon detectors such as {\\it Swift} and {\\it Fermi}, except for Earth occultations or possible outages.  For the more frequent weaker or more distant bursts, taking into account fluctuations in the average quantities, a neutrino observation correlated with a photon detection can narrow the time bin search, increasing the effective sensitivity of the detection.  The rate of occurrence of such lower fluence GRBs can be calculated from the luminosity function and redshift distribution, and is shown in Fig.  [\\ref{contour}]).   Because of the low occurrence rate of bursts which can be expected to be detected individually, it is useful to consider also the diffuse neutrino fluxes. The same luminosity function and redshift distribution as above were used for this, as discussed in \\S \\ref{sec:nudiff}, extending the integration to all bursts with $z\\ge 0.01$.  In Fig.[\\ref{diff}] we show the diffuse neutrino fluences over the period of a year, for a set of burst parameters $\\eta_{core}=10\\eta_{out}=300,700,1000$ and $\\epsilon_{b}=10\\epsilon_{e}\\approx1$. For these paramaters, the corresponding number of average expected muon events in Icecube and its DeepCore array are estimated to be $0.03$, $0.4$ and $0.5$ per year over the whole sky, respectively. Due to the poor statistics in the rate of low redshift GRBs, the uncertainty in these numbers could be a factor $\\sim 2$. The Icecube IC40+IC59 muon event rates per year being used to set upper limits on the TeV-PeV neutrino flux from putative Waxman-Bahcall GRB models (different from the present model) are larger than the the event rates discussed here, but within one order of magnitude of those for the highest $\\eta$ values.  Thus, imposing weak limits may perhaps be possible in the long term.  \\medskip  In conclusion, both the individual burst fluences and the expected diffuse flux in the 5-100 GeV range are significantly low. Fortuitous fluctuations above the mean values could increase this somewhat, but any conclusions based on the current IceCube and Deep Core arrays are likely to require years of data accumulation with the full array. The proposed upgrades to these installations would help, but a next generation of larger effective volume neutrino detectors could be the best way to accelerate the detection or non-detection of GRB neutrinos in this energy range, and to test GRB models such as discussed here.  \\medskip  We thank P. Veres, M. Strikman, M. Smith, D. Cowen and especially T. DeYoung for useful discussions, the referee for valuable critical comments, and NASA NNX09AL40G, NSF PHY-0757155 for partial support.  \\newpage"
    },
    "1112/1112.2272_arXiv.txt": {
        "abstract": "We observe that the dominant one loop contribution to the graviton propagator in the theory of $N$ ($N\\gg 1$) light scalar fields $\\phi_a$ (with masses smaller than $\\mpl/\\sqrt{N}$) minimally coupled to Einstein gravity is proportional to $N$ while that of  graviton-scalar-scalar interaction vertex is $N$ independent. We use this to argue that  the coefficient of the $R\\phi_a^2$ term appearing at one loop level is $1/N$ suppressed. This observation provides a resolution to the $\\eta$-problem, that the slow-roll parameter $\\eta$ receives order one quantum loop corrections for inflationary models built within the framework of scalar fields minimally coupled to Einstein gravity,  for  models involving large number of fields. As particular examples, we employ this to argue in favor of the absence of $\\eta$-problem in M-flation and N-flation scenarios. ",
        "introduction": "Recent CMB observations \\cite{arXiv:1001.4538} indicate that the early Universe has passed through an inflationary period with Hubble parameter $H\\lesssim10^{-5}M_{\\mathrm{pl}}$. The standard theoretical setup for inflationary models generically involve some scalar inflaton fields slowly rolling down their potential. The slow-roll that is needed to ensure a resolution to the problems of standard big bang cosmology, and consistency with the CMB results, demands $\\epsilon=-\\frac{\\dot{H}}{H^{2}}$ and $\\eta=\\frac{\\ddot{H}}{H\\dot{H}}$ to be smaller or of order $10^{-2}$ \\cite{arXiv:1001.4538}. In the context of simple single scalar models with potential $V(\\phi)$, $\\epsilon$ and $\\eta$ are a measure of flatness of the potential and, specifically, $\\eta\\sim m^{2}/H^{2}$, where $m$ is the effective mass of the inflaton field. Therefore, the theoretical framework invoked for inflationary model building should also provide mechanisms to protect the potential and its flatness against quantum and/or quantum gravity corrections. In physics models we generically associate smallness and protection of a quantity like $\\eta$ with an approximate symmetry, such that when the symmetry is exact this parameter (here the effective mass of inflaton) is zero. Supersymmetry, for example, can be such a symmetry. It turns out that in the presence of gravity, as in our case where we are dealing with inflation models, symmetries protecting $\\eta$ are all broken, inducing an inflaton mass term of order the Hubble scale $H$ of the background. We hence end up with an order one $\\eta$, i.e., the $\\eta$-problem.  From the above discussion it is seen that the $\\eta$-problem may appear in two ways: In a top-down approach where we invoke a theory of quantum gravity like string theory for inflationary model building. Or in a bottom-up approach where we take the usual field theory setup of Einstein gravity plus scalar inflaton fields, assuming that this framework is valid up to Planck scale $M_{\\mathrm{pl}}$. In the top-down approach the $\\eta$-problem appears as a classical (not loop) effect, usually due to the interaction of the inflationary sector with the \\textquotedblleft moduli stabilization sector\\textquotedblright, see, e.g., \\cite{Stringy-eta}. In these models it turns out to be easy to have small $\\epsilon$ with controlled back-reactions on it, but $\\eta$ receives order one corrections. Intuitively, in these top-down setups, the $\\eta$-problem can be understood as follows. With a vacuum energy of order $V$ all scalars, including the inflaton, will be endowed with soft masses of order $V/M_{\\mathrm{pl}}^{2}=3H^{2}$ since the supersymmetry breaking scale is not below $H$. In this work, however, we will focus on the $\\eta$-problem in the bottom-up approach.  In the bottom-up approach the $\\eta$-problem arises from \\textit{quantum} loop corrections to the tree-level graviton-scalar-scalar vertex. Despite being non-renormalizable, one can still apply the (Wilsonian) effective field theory techniques to the Einstein gravity theory and consider loop corrections, e.g. see \\cite{Gravity-loops,Duff:1974hb,'tHooft:1974bx} . In the presence of a scalar field $\\phi$ minimally coupled to Einstein gravity, as in generic inflationary models, these corrections at one loop level generically involve a $R\\phi^{2}$ term, a term whose presence was noted long ago \\cite{Sakharov-Adler}. As we will review below, such a term is quadratically divergent and in the one loop effective action appears as \\begin{equation} \\zeta\\frac{\\Lambda^{2}}{M_{\\mathrm{pl}}^{2}}R\\phi^{2}, \\label{conf-mass-zeta}% \\end{equation} where $\\Lambda$ is the UV cutoff of the theory and $\\zeta$ is an order one coefficient. Assuming a Planckian cutoff scale, $\\Lambda\\sim M_{\\mathrm{pl}}$, in an inflationary background where $R\\sim H^{2}$, this term yields a correction of order $H^{2}$ to the inflaton mass, causing the $\\eta$-problem.    The $\\eta$-problem, or the Hubble scale mass term for the effective inflaton field, seems quite generic and one may put forward the idea of \\emph{kinematically} reducing the coefficient $\\zeta\\Lambda^{2}/M_{\\mathrm{pl}% }^{2}$. In this Letter we explore this possibility. One obvious possibility is to choose the cutoff $\\Lambda$, the scale where quantum gravity effects become important, to be one or two orders of magnitude smaller than $M_{\\mathrm{pl}}$ \\cite{hep-ph/9803315}. In this case the coupling constant of gravitons will be reduced like the momentum UV cutoff and the $\\eta$ problem persists. Alternatively, one may explore the idea that $\\zeta$ is a kinematical factor that for some reason is not of order one, while the bare cutoff $\\Lambda$ is $M_{\\mathrm{pl}}$. In fact, similar suppressions are already very well known in the context of large $N$ gauge theories \\cite{'tHooft:1973jz}: the nonplanar part of a given Feynman diagram comes with powers of $1/N$ suppression compared to the planar part of the same diagram. As we will show similar analysis can be repeated for the theories involving large number of scalar fields minimally coupled to gravity. In particular, if we have $N$ number of \\textquotedblleft light fields\\textquotedblright, lighter than $M_{\\mathrm{pl}}/\\sqrt{N}$, $\\zeta$ turns out to have a $1/N$ suppression factor. In a sense, as if, the diagram leading to \\eqref{conf-mass-zeta} is a nonplanar diagram. This observation is closely related to the \\textquotedblleft species dressed gravity cutoff scale\\textquotedblright% \\ ideas discussed in \\cite{Dvali:2007hz,Dvali:2008jb}, in light of which the $\\zeta\\sim1/N$ result may be interpreted as dealing with a \\textquotedblleft dressed cutoff\\textquotedblright\\ $\\Lambda/\\sqrt{N}$ while $\\Lambda\\sim M_{\\mathrm{pl}}$.    Inflationary models with many scalar fields have recently got attention in view of their success in providing a natural explanation for the smallness of the inflaton self-couplings (the issue of steepness of the potential) and for the super-Planckian excursion of the effective inflaton in the field space \\cite{Liddle:1998jc}. This idea is not exotic to string theory motivated inflationary settings where it is quite common to have an abundant number of fields/degrees of freedom with masses below the dressed cutoff $M_{\\mathrm{pl}}/\\sqrt{N}$, see, e.g, \\cite{Nflation, Becker:2005sg,Mflation}. Even though in some of these setups, like N-flation, the individual field excursion is greater than the dressed UV cutoff, some, like Gauged M-flation \\cite{gauged-M-flation} or multiple M5-branes Inflation \\cite{Becker:2005sg}, remain immune to this\\textquotedblleft beyond-the-cutoff\\textquotedblright\\ problem.  In this work we examine the above proposed $1/N$ resolution to the $\\eta $-problem. We assume that there is a hierarchy of scales between $H$, the dressed gravity cutoff $\\Lambda_{\\mathrm{dressed}}\\sim M_{\\mathrm{pl}}% /\\sqrt{N}$ and $M_{\\mathrm{pl}}$: $H\\ll\\Lambda_{\\mathrm{dressed}}\\ll M_{\\mathrm{pl}}$ which is easily achieved by e.g. $N\\gtrsim10^{2}$. This provides a window where one can safely use the standard techniques of quantum field theory and effectively deal with a system that could be described by Einstein-Hilbert gravity, the inflaton sector, and other heavy remnants of the theory of quantum gravity whose masses $M_{a}$ falls below the new gravity cutoff, i.e. $M_{a}\\lesssim\\Lambda_{\\mathrm{dressed}}$.  The outline of this work is as follows. We consider a system of $N$ light scalars minimally coupled to Einstein gravity and work out basic Feynman rules of the theory and compute the quadratically-divergent part of the one loop contributions to the graviton propagator and graviton-scalar-scalar vertices. We show that one loop graviton two-point function has a linear $N$ parametric dependence while the graviton-scalar-scalar vertex has no $N$ dependence. Therefore, if we (re)normalize the graviton two-point function, the vertex will have a factor of $1/\\sqrt{N}$. This latter leads to $\\zeta\\sim1/N$ (\\emph{cf.} \\eqref{conf-mass-zeta}). We discuss how this can resolve the $\\eta$-problem in the context of many-field models like N-flation \\cite{Nflation} or M-flation \\cite{Mflation,{gauged-M-flation}}. ",
        "conclusions": "The bottom-up $\\eta$-problem seems quite generic to all models of inflation that involve a scalar field minimally coupled to gravity. Even if the parameters of the inflaton potential are chosen meticulously at tree-level, the loop corrections that arise from interactions of the graviton with the scalar field create the quadratically divergent conformal mass term which leads to the $\\eta$-problem, if the UV cutoff of the theory is of order $M_{\\mathrm{pl}}$. This kind of $\\eta$-problem is of course different from the ``top-down'' $\\eta$-problem  arising within the string theory setups in which the volume modulus stabilization often resurrects the $\\eta$-problem. The precision that should be enforced upon the tree-level parameters are often not needed to sustain inflation, but to match the observed density of perturbation \\cite{arXiv:1001.4538}. In this letter we examined the possibility of circumventing the $\\eta$-problem, in the former sense discussed above, in many-field models of inflation that are minimally coupled to gravity. These many-fields whose masses are assumed to be smaller than $\\mpl/\\sqrt{N}$, have a natural appearance in effective low energy field theory description of quantum gravity models. As we argued $N\\gtrsim 100$ will resolve the $\\eta$-problem. Even though it is not necessary for our argument, the scalar fields should be non-interacting to realize inflation.  One example of such many-field models is N-flation \\cite{Nflation} which has the $O(N)$ symmetric potential $V=\\sum_{i=1}^{N}\\frac{1}{2}m^2\\phi_i^2$. As stated above, a few hundred scalar fields will be enough to circumvent the $\\eta$-problem. Like its chaotic counterpart, the mass parameter $m$ has to be $\\sim 10^{-6}\\mpl$ which is smaller than $\\Ld$ unless one resorts to an unnaturally large number of scalar fields, \\textit{i.e.} $N\\sim 10^{12}$. This of course comes at the price of  exposing the model to quantum instability of the type discussed in \\cite{gauged-M-flation}. Namely, the quantum fluctuations of these light fields may dominate over the classical evolution of the inflaton. Assisted model with quartic potential $V=\\sum_{i=1}^{N}\\frac{1}{2}\\lambda\\phi_i^4$ is another possibility. To have an observationally viable  model, the effective coupling must be around $10^{-14}$ and thus with $\\lambda\\simeq 1$, one needs around $10^{14}$ scalar fields. This scenario also suffers from the above quantum instability with such a large number of massless scalar fields. Another disadvantage of both these two models is that the physical excursion of the fields is larger than $\\Ld$ during the required $60$ e-folds of inflation.   Another explicit example is M-flation \\cite{Mflation} or its gauged  version \\cite{gauged-M-flation} where the inflaton potential is constructed by three $N\\times N$ non-commutative hermitian matrices whose action is invariant under $U(N)$. The classical dynamics is simplified considerably in the $SU(2)$ sector where these three scalar fields are proportional to the generators of the $SU(2)$ algebra. Gauged M-flation, in addition to the above ingredients, has an extra $U(1)$ field, associated with ``center of mass '' $U(1)\\in U(N)$. M-flation in the $SU(2)$ sector besides the inflaton field contains some number of ``spectator fields'' which do not contribute to classical inflationary trajectory while can be excited quantum mechanically and appear in the loops. For the gauged M-flation there are $2N^2+1$ such scalar modes and $3N^2-1$ massive vector modes. These modes have a hierarchical spectrum, {\\it i.e.} they can be lighter or heavier than the Hubble scale $H$, for the explicit masses see \\cite{gauged-M-flation}. Not all these modes are light enough to be counted in the dressed cutoff. As discussed in \\cite{gauged-M-flation}, the number of ``contributing species'' $N_s$ varies between $3\\times10^{5}$ and $10^{6}$, depending on the region of the potential inflation happens. Hence, the species dressed UV cutoff is $10^{-3}M_{\\mathrm{pl}% }\\lesssim\\Ld\\lesssim5\\times10^{-3}M_{\\mathrm{pl}}$. Consequently the $R\\phi^{2}$ term is suppressed by a factor of $\\lesssim10^{-5}$ and could be safely ignored. Gauged M-flation could be motivated from the branes dynamics in an appropriate flux where the above matrices correspond to three of the perpendicular directions of a stack of $N$ D3-branes which are scalars in the adjoint representation of the $U(N)$. As such, although the ``quantum\" $\\eta$-problem is resolved for M-flation, embedded within string theory, one should still deal with the ``stringy $\\eta$-problem'' (cf. introduction for further references and discussions).  Finally, the term leading to the $\\eta$-problem, $R\\phi_a^2$, is a one-loop but marginal operator. One may naturally worry about other loop corrections, that enhancement factors of $N$ will dominate over ${\\Ld}/{\\mpl}$ suppressions. It is straightforward to show that the largest such $N$ enhancement factor appears in the graviton two point function (at higher loops) for which this factor is $\\left(N{\\Ld^2}/{\\mpl^2}\\right)^l$, where $l$ is number of loops. All the other diagrams will be of the form $N^k(\\frac{\\Ld^2}{\\mpl^2})^{l}$ with $k<l$. Hence our one loop result seems to be also valid to all orders."
    },
    "1112/1112.5949_arXiv.txt": {
        "abstract": "{ According the Collapsar model long gamma-ray bursts (LGRBs) involve relativistic jets that puncture the envelope of a collapsing star, and produced the $\\gamma$-rays after they break out. This model provides a theoretical framework for the well known association between LGRBs and massive stars. However although this association is supported by a wealth of observations, to this date there is no direct observational evidence for the emergence of the jet from the star. In other words there is no direct evidence for the Collapsar model. Here we show that a distinct signature of the Collapsar model is the appearance  of a plateau in the duration distribution of the prompt GRB emission at times much shorter than the typical breakout time of the jet. This plateau is evident in the data of all major GRB satellites, and provides a direct evidence supporting the Collapsar model. It also enables us to place limits on the sizes and masses of LGRB progenitors; suggests the existence of a large population of choked (failed) GRBs; and indicates that the 2 s duration commonly used to separate Collapsars and non-Collapsars holds for BATSE and possibly Fermi GBM GRBs, but it is inconsistent with the duration distributions of {\\it Swift} GRBs.}  \\FullConference{The Extreme and Variable High Energy Sky - extremesky2011,\\\\ September 19-23, 2011\\\\ Chia Laguna (Cagliari), Italy}   \\begin{document} ",
        "introduction": "The connection of long GRBs (LGRBs) to collapsing massive stars is strongly supported by a wealth of observational evidence. The most notable evidence is the association of half a dozens GRBs with spectroscopically confirmed broad-line Ic supernovae (SNe), and the identification of \"red bumps\" in the afterglows of about two dozens more, which shows a photometric evidence of underlying SNe. The model that provides the theoretical framework of this association is known as the Collapsar model \\citep{MacWoos99,MacFadyen01}. According to this model, following the core collapse of a massive star, a bipolar jet is launched at the center of the star. The jet drills through the stellar envelope and breaks out of the surface before producing the observed $\\gamma$-rays. However, although this model is supported indirectly by the LGRB-SN association, to this date we could not identify a clear direct observational imprint of the jet-envelope interaction, thus there is no direct confirmation yet of the Collapsar model.  Here we show that under very general conditions the time that the jet spends drilling through the star leads to a plateau in the duration distribution of the GRBs at times much shorter than the breakout time of the jet. This plateau exists in the duration distribution of all major GRB satellites, and provides a strong observational support for the Collapsar model. An analysis of this plateau also (i) supports the hypothesis of compact stellar progenitors, (ii) implies the existence of  a large population of chocked jets that fail to break out of their progenitor stars and (iii) enables us to determine statistically the fraction of non-Collapsars from the total GRB sample and the threshold duration that separates Collapsars from non-Collapsars. Specifically it shows that this time differs from one satellite to the other. ",
        "conclusions": "The observed plateaus in all three duration distributions, and most notably in the distribution of soft BATSE bursts, provide a direct support for the Collapsars model for LGRBs. An inspection of different regions of the observed temporal distribution (Fig. 1), under the interpretation of the plateau as an imprint of the time it takes the jet to break out of the envelope, provides further important information. \\begin{enumerate}[itemsep=2pt,parsep=2pt,leftmargin=0.5cm] \\item{ The end of the plateau and the decrease in the number of GRBs at long durations, allow us to estimate the typical time it takes a jet to breakout of the progenitor's envelope. All three distributions are flat below $\\sim 10$ s in the GRB frame, implying that $\\hat t_b \\sim$ a few dozen seconds. This fits nicely with the canonical GRB parameters used in Eq. \\ref{eq:tB_GRB}, and provides another indication that the stellar progenitors of Collapsars must be compact \\citep{Matzner03}.}  \\item{ The distribution at long durations can be used to set a lower limit on the number of chocked jets, by extrapolating the slope at $T_{90}\\gg \\tbh$ to durations shorter than $\\tbh$ (see section \\ref{sec theory}). At durations $T_{90}\\gtrsim 100$ s $p_\\gamma$ can be fitted by a power law, $p_\\gamma \\propto T_{90}^{\\alpha}$ with $-4<\\alpha<-3$. Extrapolating to $T_{90} < \\hat t_b$ we find that if $p_e$ continues with this power law to $t_e < \\tbh$ there are significantly more chocked GRBs than long ones. For example, even if we extrapolate  this distribution only down to $t_e = \\hat t_b/2$ there are $\\sim10$ times more chocked GRBs than long GRBs. This prediction is consistent with the suggestion that shock breakout from these choked GRBs produces low luminosity smooth and soft GRBs \\citep[][and references there in]{Nakar11}. Indeed the rate of such low luminosity GRBs is far larger than that of regular long GRBs \\citep{Soderberg06}.}   \\item{At short durations, the duration distribution is dominated by non-Collapsars SGRBs. These bursts are hard to classify and are commonly defined by their shorter duration. Our analysis shows that putting the dividing line between Collapsars and non-Collapsars at 2 s is statistically reasonable for BATSE, and possibly also for Fermi bursts. However, it is clearly wrong to do so for {\\it Swift}. We also calculate the probability as a function of $T_{90}$ that a burst with a given $T_{90}$ is not produced by a Collapsar.}  \\item{While the difference in the lower limit of the plateaus is understood qualitatively as it depends on the spectral range of the detector, the variance in the upper limit is less obvious. It may reflect various selection effects in triggering algorithms. A more interesting possibility is that it reflects a physical origin, e.g. different satellites probing populations with different $\\hat t_b$.  This could be explored when the statistical sample of Fermi GBM becomes sufficiently large. }   \\end{enumerate}"
    },
    "1112/1112.5178_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec1}  Initial singularity is one of the most serious problems in modern cosmology. It is inevitable under several physically reasonable conditions in general relativity as a consequence of the singularity theorems~\\cite{singularity,singularity2}. Since the quantum effect of gravity dominates close to curvature singularities, the completion of the quantum theory of gravity is necessary to answer to this problem. Although the full quantum theory of gravity is not available at present, (Super)string/M-theory is one of the most promising candidates. Given the present circumstances, the results in the low-energy effective theory must be useful to gain insights and suggestions for the problem. As a possibility to avoid the initial singularity, the bouncing universe is one intriguing scenario~\\cite{gv1993}. Because the effective energy-momentum tensor given from the semiclassical correction may violate the energy conditions, the quantum-gravity effect could cause the big bounce of the universe, a transition from the contracting phase to the expanding phase. (See~\\cite{bounce} for a review.)  Higher-dimensional cosmology based on string-generated gravity models is an ambitious way  to explain problems in modern cosmology such as dark energy or initial singularity in a harmonic way by the effect of extra dimensions. After the publication of the papers by Randall and Sundrum~\\cite{Randall}, string-inspired braneworld cosmology, in which the observable universe is a (3+1)-dimensional timelike hypersurface (three-brane) embedded in a higher-dimensional anti-de~Sitter (AdS) bulk space-time, has been investigated enthusiastically. (See~\\cite{braneworld} for a review.) The dynamics of the Friedmann-Robertson-Walker (FRW) brane universe has been fully investigated to explain cosmic inflation in the early universe or dark energy. In the present paper, we focus on the possibility of the bouncing brane universe.  Among two Randall-Sundrum models, we consider the single-brane model. Assuming a perfect fluid obeying an equation of state $p=(\\gamma-1)\\rho~(0<\\gamma\\le 2)$ on the FRW brane in the simplest braneworld with $Z_2$ symmetry~\\cite{krausida}, we can show that the bouncing universe is not realized for $2/3<\\gamma \\le 2$. This implies that, although the condition is milder than the standard four-dimensional FRW cosmology, a matter field with negative pressure violating the strong energy condition is still necessary for the big-bounce. Moreover, the bouncing universe requires the negative mass parameter in the Schwarzschild-Tangherlini-type bulk solution which causes a naked singularity.  Nevertheless, there are models providing the bouncing universe on the brane with ordinary matter. One possible way is to introduce a matter field in the bulk spacetime. The low-energy effective theory of (Super)string/M-theory predicts several fields in the bulk spacetime such as a dilaton field or gauge fields. For example, by the dimensional reduction from the 11-dimensional Ho{\\v r}ava-Witten model to the five-dimensional bulk spacetime, a dilaton field and a U(1) gauge field appear~\\cite{hw1996}. In the presence of a U(1) gauge field in the bulk, the bouncing universe is generically realized on the brane~\\cite{bounce-U(1)}. However, it is expected that the bulk spacetime suffers from the mass inflation instability~\\cite{pi1990,massinflation} when the brane approaches the inner Cauchy horizon in the bulk~\\cite{hm2003}. In this research direction, a model of the bouncing brane universe without a bulk inner horizon was constructed in the presence of a SU(2) Yang-Mills field, in which the bulk spacetime is regular and free-from singularities~\\cite{om2003,om2004}.  In the present paper, we reconsider the possibility of the bouncing brane universe with ordinary matter on the brane. As mentioned above, this is impossible in the standard Randall-Sundrum scenario based on general relativity unless the extra dimension is spacelike. Actually the bouncing universe is generically realized if the extra dimension is timelike~\\cite{ss2003}, but we do not consider this radical possibility here. Instead, we consider the braneworld in the framework of Einstein-Gauss-Bonnet gravity. While the bouncing universe in the Gauss-Bonnet braneworld was discussed also in~\\cite{mp2010} with a U(1) gauge field in the bulk, we consider the vacuum bulk here.   Einstein-Gauss-Bonnet gravity is a natural higher-dimensional generalization of general relativity as a quasilinear second-order theory. Also, the quadratic Lanczos (Gauss-Bonnet) Lagrangian appears in the low-energy limit of heterotic string theory together with a dilaton~\\cite{Gross}. The early stage of the research history of the Gauss-Bonnet braneworld was somehow winding. There was a disagreement among the authors on the junction condition for the bulk spacetime and hence the resulting Friedmann equation on the brane. This confusion was settled finally by the establishment of the generalized Isreal junction condition by Davis~\\cite{davis2003} and independently by Gravanis and Willison~\\cite{GB-junction}. (See also~\\cite{onGB-juncton} and the comment in~\\cite{davis2003}.) The correct Friedmann equation in the Gauss-Bonnet braneworld was first derived by Charmousis and Dufaux in~\\cite{GBbrane-FRWequation} and the dynamics of the Friedmann brane has been investigated by many authors up to now~\\cite{GBbrane-dynamics,brownPhD}. (See~\\cite{GBbrane-equations} for the covariant gravitational equations on the brane.) Also, its viability~\\cite{GBbrane-viability}, creation of the universe~\\cite{GBbrane-creation}, cosmological perturbations~\\cite{GBbrane-perturbation} have been investigated. The maximally symmetric brane was studied in~\\cite{GBbrane-(A)dSbrane}. Other aspects of the Gauss-Bonnet braneworld have also been studied~\\cite{GBbrane-stringy}. (See Sec. 5.7 in~\\cite{Review} for a review.) The purpose of the present paper is to point out a new and intriguing dynamical property of the Gauss-Bonnet brane universe.  In Einstein-Gauss-Bonnet gravity, there is an exact  symmetric vacuum solution which is a generalization of the Schwarzschild-Tangherlini solution in general relativity~\\cite{bd1985,GBBH}. One of the most characteristic properties of this solution is the existence of a curvature singularity at a nonzero physical radius for negative mass parameter, which is called the {\\it branch singularity}. The qualitative behavior of the Gauss-Bonnet brane universe has been analyzed by many authors; however, the role of the branch singularity in this context has not been clarified yet. In the general relativistic case, the moment when the brane reaches the central singularity in the bulk corresponds to the big-bang (or big-crunch) time on the brane. Therefore, one might naively think that the brane universe ends in (or starts from) some curvature singularity on the brane when it hits the branch singularity in the bulk. However, we will show that it is not the case.  In the present paper, we will show that the branch singularity is harmless in the sense that a finite body is able to reach there without being crushed to a point or ripped apart. We will also show that the vacuum bulk spacetime can be extended beyond the branch singularity in the  $C^0$ sense and then the branch singularity is considered as a massive thin shell. As a consequence, it is concluded that the brane may reach and pass through the branch singularity into the extended region of the spacetime. This process provides a new scenario for the bouncing universe on the brane.    The rest of the present paper is constituted as follows. In the following section, we present the bulk spacetime and the dynamical equation for the FRW brane in Einstein-Gauss-Bonnet gravity. In Sec.~III, we explain our new scenario for the bouncing universe. In Sec.~IV, the weakness of the branch singularity is shown. Concluding remarks and discussions including future prospects are summarized in Sec.~V. In Appendix A, we present the geometric quantities of the bulk spacetime. In Appendix B, the asymptotic behavior of the brane universe for $a\\to \\infty$ is presented.  Our basic notation follows~\\cite{wald}. The convention for the Riemann curvature tensor is $[\\nabla _\\rho ,\\nabla_\\sigma]V^\\mu ={R^\\mu }_{\\nu\\rho\\sigma}V^\\nu$ and $R_{\\mu \\nu }={R^\\rho }_{\\mu \\rho \\nu }$. The Minkowski metric is taken as diag$(-,+,+,+,+)$, and Greek indices run over all spacetime indices. We adopt the units in which only the five-dimensional gravitational constant $G_5$ is retained. ",
        "conclusions": "In this paper, we have presented a novel scenario for the bouncing universe in the five-dimensional braneworld in the framework of Einstein-Gauss-Bonnet gravity. In this scenario, the branch singularity located at the finite physical radius in the bulk, which appears for the negative mass parameter,  plays an essential role. We have shown that the branch singularity can be considered as a massive thin shell and a three-brane may pass through it. As a result, the bouncing universe is realized on the brane.  The bulk spacetime is extended beyond the branch singularity in the $C^0$ sense and then the branch singularity is identified as a massive thin shell. From the bulk point of view, the three-brane collides with another shell of branch singularity and continues to evolve into the extended spacetime. From the brane point of view, the moment of the collision is the moment of the bounce. This is not the conventional bounce because the derivative of the scale factor is discontinuous at the bounce moment. Therefore, there appears a shell-type instantaneous singularity on the brane but the curvature invariants never blow up before or after then. This claim is strongly supported by the fact that the branch singularity is radially deformationally weak, which implies that the singularity is harmless for a finite body moving radially. The present result opens a completely new possibility to achieve the bouncing brane universe as an effect of the higher-curvature terms. Our scenario is not sensitive about the equation of state for the matter on the brane and does work even with ordinary matter.      Here we should make a brief comment on the anisotropy of the universe. Although the assumption of the isotropy of the universe is justified at the recombination era by the distribution of the cosmic microwave background in the WMAP-Planck data, its validity is a nontrivial problem in the very early universe before that. In the standard FRW cosmology, it is expected by the BKL (Belinsky-Khalatnikov-Lifshitz) conjecture that such an assumption is no more valid and the behavior of the generic spacetime close to the initial singularity is quite anisotropic~\\cite{we1997}. In contrast, this claim is not necessarily true in the Randall-Sundrum braneworld. It has been shown within the Bianchi class anisotropic cosmological models that the brane universe is generically isotropic close to the initial singularity. (See~\\cite{BianchiBrane} for review.) However, the similar analysis has not been done in the Gauss-Bonnet braneworld so far. This is one of the important future subjects to clarify the validity of our bouncing scenario.    Our result also gives a suggestion on what would happen if the brane approaches the inner Cauchy horizon in the generalized Reissner-Nordstr\\\"om-AdS bulk spacetime with a U(1) gauge field in the Randall-Sundrum braneworld~\\cite{bounce-U(1),hm2003}. The rigorous analysis in the four-dimensional asymptotically flat case suggests that the mass inflation instability transforms the inner horizon into a weak singularity. More precisely, the resulting singularity appears at a finite physical radius where the metric and its inverse are finite~\\cite{massinflation}. Therefore, the $C^0$ extension is possible beyond this singularity. Then a natural question is whether the brane can reach and pass through it or not. Of course, the appearance of the singularity means that the theory is no more reliable around there. However, the result in the present paper suggests that the three-brane could safely pass through the inner horizon in the bulk.   Since the existence of the branch singularity stems from the quadratic nature of the theory, such a singularity is characteristic and must be quite generic in higher-curvature theories. In Lovelock higher-curvature gravity~\\cite{lovelock}, which contains general relativity and Einstein-Gauss-Bonnet gravity as special cases, such a singularity appears in this class of vacuum solutions rather generically~\\cite{mwr2011}. Interestingly, the central singularity is totally absent and the branch singularity is generic independent of the mass paraeter $\\mu$ if there is a U(1) gauge field in the bulk spacetime in Einstein-Gauss-Bonnet gravity~\\cite{GBBH-charge,tm2005b}. Undoubtedly, the effect of such singularities in cosmology is an interesting problem and should be investigated further.            \\subsection*{Acknowledgements} The author thanks S.~Willison, Y.~Shtanov, and V.~Sahni for useful comments and discussions. He also thanks F.~Canfora for comments on the asymptotic analysis. This work has been funded by the Fondecyt Grants No. 1100328 and No. 1100755 and by the Conicyt grant \"Southern Theoretical Physics Laboratory\" ACT-91. This work was also partly supported by the JSPS Grant-in-Aid for Scientific Research (A) (22244030). The Centro de Estudios Cient\\'{\\i}ficos (CECs) is funded by the Chilean Government through the Centers of Excellence Base Financing Program of Conicyt.    \\appendix"
    },
    "1112/1112.1792_arXiv.txt": {
        "abstract": "{Nearby, Galactic gamma-ray bursts (GRBs) may affect the terrestrial biota if their radiation is beamed towards the Earth. Compact stellar binary mergers are possible central engines of short GRBs and their rate could be boosted in globular clusters.} {The occurrence, impact on the terrestrial biota and the point in time of the explosion of the closest GRB launched in a globular cluster are explored.} {Globular cluster typically follow well defined orbits around the galactic center. Therefore their position relative to the solar system can be calculated back in time. This fact is used to demonstrate that globular cluster - solar system encounters define possible points in time when a nearby GRB could have exploded. Additionally, potential terrestrial signatures in the geological record connected to such an event are discussed.} {Assuming rates of GRBs launched in globular cluster found from the redshift distribution of short burst and adopting the current globular cluster space-density around the solar system it is found that the expected minimal distance $d_\\mathrm{min}$ for such a GRB in the last Gyr is in the range $d_\\mathrm{min} \\sim 1 - 3.5$~kpc. From the average gamma-ray luminosity of a short GRB significant depletion of the terrestrial ozone-layer is expected if such an event explodes at a distance of $\\sim1$~kpc. In the last Gyr a few globular cluster passages are expected within a distance of $d_\\mathrm{min}$ from the solar system and a GRB should have exploded during one of these passages.} {Globular cluster - solar system encounters and events of mass extinction in the history of life can be correlated to investigate the impact of a nearby GRB on the terrestrial biota. To explore such a correlation reliable globular cluster positions relative to the solar system have to be calculated for the time span of the fossil record of the last 600~Myrs. The upcoming GAIA mission will be crucial to determine the possible time intervals of the occurrence of nearby GRBs launched in globular clusters.} ",
        "introduction": "Gamma-ray bursts (GRBs) are powerful cosmological explosions which irradiate their surroundings in two cones with high energy photons \\citep[see][for a review]{gehrels2009}. It has been recognized that due to this radiation, galactic GRBs can have a dramatic effect on the earth atmosphere and on the biota \\citep{thorsett1995}. Furthermore, it has been pointed out that beside the extinction of species cosmic rays connected to GRBs can cause biological mutations leading to fast appearance of new species after these mass extinctions events \\citep{dar1998}. Indeed the evolution of live on earth was affected by several events of mass-extinction \\citep{raup1982} and show also periods of rapid development of new species \\citep[e.g. Cambrian explosion][]{marshall2006}. From the observed rate of GRBs it is evident that a nearby ($d \\sim 1$~kpc) burst is likely during the geological record and such an event has been linked to the late Ordovician mass extinction event that happened 440~Myr in the past \\citep{melott2004}. Additionally,  \\citet{horvath2003} has suggested that the Cambrian explosion 543~Myr ago could have been triggered by a nearby GRB. In general, the exact point in time when a nearby GRB has happened is difficult to identify. This challenges the correlation between a certain event in the history of the Earth and the occurrence of a nearby GRB.  GRBs appear to be launched by two distinct populations of progenitors. Long bursts (duration $> 2$~s) appear to be connected to the death of massive stars whereas short burst (duration~$<~2$~s) seem to result from the merger of two compact objects \\citep[see][for a review]{gehrels2009}.  For short bursts it has been shown that their rate can be boosted in globular clusters due to the efficient production of close stellar binaries in their cores \\citep{grindlay2006,lee2010}. Several observational facts support the association between short GRBs and globular clusters. Firstly, in the globular cluster M~15 a close binary consisting of two neutron stars has already been discovered \\citep{anderson1990}. Secondly, from the off-set distribution of short bursts from their host galaxies it has been inferred that a subset of bursts can originate in globular clusters \\citep{berger2010,salvaterra2010,church2011}. Thirdly, from the redshift distribution of short burst it has been concluded that in the local universe the short GRB rate is dominated by dynamically formed compact binaries in globular clusters \\citep{salvaterra2008,guetta2009}. Finally the occurrence of short GRBs in globular clusters in the Milky Way can also be probed by the presence of corresponding remnants \\citep{domainko2005,domainko2008}. Recently, indeed a source of very-high-energy gamma-rays, HESS~J1747-248, in the direction of the globular cluster Terzan~5 has been discovered \\citep{abramowski2011}. This source is potentially accompanied by diffuse X-ray emission \\citep{eger2010} and diffuse radio emission \\citep{clapson2011} and has been interpreted in the framework of a GRB remnant \\citep{domainko2011}. To conclude there seems to be evidence that a considerable rate of GRBs are launched in globular clusters.  Globular clusters typically follow well defined orbits around the center of the Milky Way. Proper motions are available for a number of globular clusters \\citep{dinescu1997,dinescu1999a,dinescu1999b} and this allows to determine their orbital motions \\citep{allen2006}. Thus, their position relative to the Earth can be traced back in time \\citep{vandeputte2009}. Encounters of globular clusters and the Earth are potential time intervals during which to the biota hazardous, nearby bursts can occur. Consequently it is possible to identify such periods by tracing back the distance of globular cluster to the Earth.  The observed values for the proper motion of globular clusters are still affected by considerable errors. \\citet{vandeputte2009} found that the relative position between the Earth and various globular clusters can be traced back in time reliably only for about 50~Myr. Since this time span is considerably shorter than the geological record, in this paper the possibility of a correlation between globular cluster passages and events in the history of the Earth are investigated on a statistical basis.  This paper is organized as follows: firstly the expected minimum distance to a GRB launched in a globular cluster will be determined. Secondly, potential terrestrial signatures connected to a GRB happening at this minimal distance are investigated. Thirdly, it will be shown that with determining the time intervals of globular cluster - solar system encounters points in time can be identified, where in principal nearby GRBs can happen. And finally the prospects for correlating such globular cluster passages and events of mass extinction are discussed. ",
        "conclusions": ""
    },
    "1112/1112.4038_arXiv.txt": {
        "abstract": "Swift J1644+57 is an unusual transient event, likely powered by the tidal disruption of a star by a massive black hole. There are multiple short timescales X-ray flares were seen over a span of several days. We propose that these flares could be produced by internal shocks. In the internal shock model, the forward and reverse shocks are produced by collisions between relativistic shells ejected from central engine. The synchrotron emission from the forward and reverse shocks could dominate at two quite different energy bands under some conditions, the relativistic reverse shock dominates the X-ray emission and the Newtonian forward shock dominates the infrared and optical emission. We show that the spectral energy distribution of Swift J1644+57 could be explained by internal shock model. ",
        "introduction": "Swift J1644+57 was triggered by the Swift/BAT on 28 March 2011 (Cummings et al. 2011). Swift J1644+57 was initially discovered as a long-duration gamma-ray burst (GRB 110328A) by the Swift satellite, but the light curve soon showed that it was quite different. It remained bright and highly variable for a long period, and re-trigger the BAT three times over the next 48 hours (Sakamoto et al. 2011). The isotropic X-ray luminosity of Swift J1644+57 ranges from $10^{45}-4\\times 10^{48}~$erg~s$^{-1}$, and the total isotropic energy is about $3\\times 10^{53}$ erg during the first 30 days after the BAT trigger (Burrows et al. 2011). From the strong emission lines of hydrogen and oxygen, Levan et al. (2011) estimate the redshift of Swift J1644+57 is $z=0.35$. From the astrometric observation of the X-ray, optical, infrared, and radio transient with the light-centroid of the host galaxy, it is found that the position of this source is consistent with arising in the nucleus of the host galaxy (Bloom et al. 2011; Zauderer et al. 2011).  The X-ray light curve of Swift J1644+57 exhibits repeated extremely short timescale flares. The flares have rise-times as short as $100$ s (Burrows et al. 2011). These flares are similar as the flares discovered in the GRB afterglow (Burrows et al. 2005), which may indicate the same origin of them. The internal shock model can produce the X-ray flares observed in GRB afterglows (Burrows et al. 2005; Fan \\& Wei 2005; Zhang et al. 2006; Yu \\& Dai 2009).  After the Swift J1644+57 was discovered, several models were proposed to explain it, most concentrating on a picture that a main sequence star was tidally disrupted by passing too close to a $10^6-10^7M_{\\odot}$ black hole (Bloom et al. 2011; Burrows et al. 2011; Cannizzo et al. 2011; Socrates 2011; Shao et al. 2011). Krolik \\& Piran (2011) suggest that this event may be produced by a white dwarf tidally disrupted by a $10^4M_\\odot$ black hole. The process is as follow: a star is disrupted as it passes near a supermassive black hole, and much of its mass is distributed into an accretion disk around the black hole. A powerful jet is then launched. In these models, the X-ray emission is thought to be produced by external inverse Compton (EIC) (Bloom et al. 2011) or synchrotron emission (Burrows et al. 2011). But on the high frequency side, the Fermi LAT (Campana et al. 2011) and VERITAS upper limits (Aliu et al. 2011) require that the synchrotron self-Compton (SSC) component is suppressed by $\\gamma-\\gamma$ pair production. The soft photons of $\\gamma-\\gamma$ pair production are thought to be generated from the thermal emission of the accretion disk or the disk outflow. In the SSC model, soft photons originated from the thermal emission of the accretion disk may not provide an efficient source for the $\\gamma-\\gamma$ production. Because the condition of $\\gamma-\\gamma$ production is $E_XE_\\gamma(1-\\cos\\theta)\\geq 2(m_ec^2)^2$, where $\\theta$ is the angle between the directions of soft seed photon and high-energy photon. Only a fraction of high-energy emission can be absorbed by soft photons. So the soft photons from the disk outflow may be a better candidate (Strubbe \\& Quataert 2009). In the synchrotron emission model, the jet must have a strong magnetic field (Poynting-flux-dominated) and has ongoing in situ acceleration of electrons (Aliu et al. 2011; Burrows et al. 2011).  In this letter, we use the internal shock model to explain the X-ray flares of Swift J1644+57. The internal shock produces the prompt emission of GRB in the standard fireball model (Rees \\& M\\'{e}sz\\'{a}ros 1994; Paczy\\'{n}ski \\& Xu 1994). The internal shock model also the leading model of X-ray flares in GRBs, the external shock model is very hard to account for the X-ray flares (Burrows et al. 2005; Fan \\& Wei 2005). The central engine of this event may be formed as follow. When a supermassive black hole tidal disrupts a star, a disk is formed. The magnetic field could be produced by disk instability. The disk can then anchor and amplify the seed magnetic field to a strong ordered poloidal field , which in turn threads the black hole with debris material in the inner region of the disk. A large amount of the rotational energy of the black hole can be extracted via the Blandford-Znajek (BZ) process, which creates a jet along the rotation axis (Blandford \\& Znajek 1977). The magnetic field lines will break the disk into blobs, so many shells could be ejected (Cheng \\& Lu 2001). When fast shell catches up with early slow shell, internal shock is generated. Other models of central engine also discussed, such as the episodic accretion onto a central object due to a chopped accretion disk (Perna et al. 2006), or episodic accretion due to a modulation of the accretion flow by a magnetic barrier (Proga \\& Zhang 2006).   The structure of this letter is as follow. In next section, we describe the dynamics of internal shock arising from a collision between two shells and the synchrotron radiation of the shocked electrons. In Section 3, we apply the model to the Swift J1644+57. Finally, a summary is given in Section 4. ",
        "conclusions": "In this letter, we propose the internal shock model to explain the X-ray flares of Swift J1644+57. In the internal shock model, collisions between a series of relativistic shells generate many pairs of forward and reverse shocks. The synchrotron emission produced by the forward and reverse shocks could dominate at two quite different energy bands if the Lorentz factors of these two types of shocks are significantly different from each other. We show that the spectral energy distribution of Swift J1644+57 could be fitted in internal shock model, in which the reverse shock is relativistic and the forward shock is Newtonian. A moderate extinction ($A_V=3-5$) is required, this value is consistent with that used in Bloom et al. (2011) and Burrows et al. (2011). Burrows et al. (2011) showed that the high frequency spectrum is produced by the synchrotron and SSC mechanisms, similar to poynting flux-dominated blazar jet model. The radio fluxes come from a larger region of the other jet. This model requires continuous in situ re-acceleration of electrons to maintain a low energy cut-off in the electron distribution (Aliu et al. 2011). Bloom et al. (2011) presented two models for the spectrum: one is two-component blazar emission model, the other is forward shock emission from jet-ISM interaction plus EIC emission model. But on the high frequency side, the LAT and VERITAS upper limits require that the SSC component is suppressed by $\\gamma-\\gamma$ pair production. The soft photons from the disk outflow may provide sources for the $\\gamma-\\gamma$ production.  The rapid rise and decline of the light curve may indicate the internal shock origin of these flares. The external shock is very hard to account for the X-ray flares (Burrows et al. 2005; Fan \\& Wei 2005; Zhang et al. 2006). During the high state, the emission of internal shock will dominate at X-ray band and a broken power-law spectrum is shown. During the low state, there is no internal shock and the emission is from external shock. The spectrum at X-ray band will be shown as a single power-law if no energy injection happens. In both states, the radio emission is from the external shock."
    },
    "1112/1112.3104_arXiv.txt": {
        "abstract": "Hot (OB) post-AGB stars are immediate progenitors of planetary nebulae (PNe). Very few hot post-AGB stars are known. Detecting new hot post-AGB candidates and follow-up multiwavelength studies will enable us to further understand the processes during the post-AGB evolution that lead to the formation of PNe. Case-Hamburg OB star surveys and their extension (LS, LSS, and LSE catalogues) and IRAS (point source) catalogues are good sources for detecting new hot post-AGB candidates from low resolution spectroscopy. Spectral types are determined from low resolution optical spectra of 44 stars selected from the LS, LSS, and LSE catalogues. Unlike the stars in the first paper, the stars in this paper were selected using criteria other than positional coincidence with an IRAS source with far IR (IRAS) colours similar to post-AGB supergiants and planetary nebulae. These included high galactic latitude, spectral types of O, B, A  supergiants, emission lines in the spectrum and known spectral peculiarity. From the present study we find that LSS 1179, LSS 1222, LSS 1256, LSS 1276, LSS 1341, LSS 1394, LSS 2241, LSS 2429, LSS 4560, LSE 16, LSE 31, LSE 42, and LSE 67 to be new hot post-AGB candidates. Further study of these candidates is needed. ",
        "introduction": "From an analysis of the IRAS point source catalogue data, several post-asymptotic giant branch (post-AGB) candidates were detected \\citep{partha86,partha89,pot88a,partha93b,kwok93,vanwin03}. These stars have far-infrared (IRAS) colours similar to proto-planetary nebulae (PPNe) and planetary nebulae (PNe) \\citep{ vander88,pot88b}. Several of them are at high galactic latitude, like HD 161796 (F3Ib) \\citep{partha86} and LS II + 34 26 (B1Iae) \\citep{partha93b}. Multiwavelength studies of these objects by several investigators have confirmed that these are indeed in post-AGB stage of evolution \\citep{kwok93, vanwin03}. The post-AGB stars seem to form an evolutionary sequence ( K, G, F, A to OB supergiant types) in the transition region from the tip of the AGB to the early stages of planetary nebulae \\citep{partha93b,partha93a}. In order to understand the evolution from cooler to hotter post-AGB types and then into the young PN stage, it is important to detect and study several hot post-AGB candidates. \\citet{vij97, vij98} and \\citet{partha00} presented the results of low resolution blue spectra of 40 IRAS sources with far-IR colours similar to PNe and PPNe \\citep{vander88, pot88b}. A few of these objects have shown rapid changes in MK spectral type \\citep{partha93a, partha95}.  The importance of the detection of new hot post-AGB candidates relies in objects such as SAO 244567 (Hen 3-1357) and LS II +34 26. SAO 244567 evolved rapidly from  a B-type post-AGB star into a young PN within a period of 20 years \\citep{partha93a, partha95, bob98}. The variations in the spectrum of SAO 244567 were dramatic. Another star LS II +34 26 initially classified as a massive B supergiant \\citep{turn84} turned out to be a rapidly evolving B-type post-AGB star \\citep{partha93a, smith94, garcia97}. Other such object are  LSE 162 (SAO 85766) \\citep{volk89, arkh07} (references therein) and LSIV -12 111 \\citep{conlon93}. Another reason to detect more hot post-AGB candidates is their abundance peculiarities \\citep{mccau92} which are different when compared with the chemical composition of cooler post-AGB stars \\citep{vanwin03}. Very few hot post-AGB stars are known and most of them are at high galactic latitudes and several of them are not IRAS sources \\citep{mccau92, moeh98}. We started a program to identify new hot post-AGB candidates by obtaining spectra of selected stars from LS, LSS and LSE catalogues \\citep{hard59,step71,dril94,dril95}, in the hope that some of them would reveal their post-AGB nature by variations in the spectrum. And, if we have a siginificant sample of hot post-AGB candidates subsequent chemical composition study of these candidates may enable us to understand their abundance peculiarities. LS, LSS and LSE catalogues are best sources for selecting high latitude luminous hot stars. In this paper we present new spectral types for 44 of these stars based on observations made with the 1m telescope of the Cerro Tololo Inter-American Observatory (CTIO), Chile. Included are several IRAS point sources which were not included in the first paper \\citep{partha00}. ",
        "conclusions": "The presence of circumstellar dust with far-IR colours similar to PNe, high galactic latitude, OB supergiant type spectrum, and emission in the Balmer lines are some of the characteristics of hot post-AGB stars. In the sample of observed stars we found a few cases of B type dwarfs with H${\\beta}$ emission. The presence of circumstellar material around Be dwarfs indicates that they may be related to other Be stars, shell stars or Herbig Ae/Be stars. In Table 1 we have a few other objects which are not known to be post-AGB stars. They were included in the observing program as some of them have IRAS colours overlapping with the IRAS colours of some known post-AGB stars.  The comments given in Table 1 for some of the stars are from the OB star catalogues and SIMBAD data base. For more information on previous spectral classification, notes and comments etc. refer to OB star catalogues and SIMBAD data base. The m$_{V}$ given in Table 1 is from the SIMBAD data base.  New spectral types for 44 O, B and A stars in the LS, LSS, and LSE catalogs have revealed several new hot (OBA supergiant) post-AGB candidates on the basis of either coincidence with IRAS point sources or high galactic latitude. None of these objects shows any evidence of a changing spectral type, but we plan to continue our monitoring program. Some of the post-AGB candidates we detected in our sample are not associated with an IRAS source, indicating they do not have dust shells. These are low mass objects and their evolution in the HR diagram from the tip of the AGB is rather very slow and by the time they evolve to A, B, O post-AGB spectral type, the dust shells seems to have disappeared and they may never appear as planetary nebulae, hence they can be called naked post-AGB stars (e.g. BD +39 4926). Mutliwavelength study of post-AGB candidates discussed in this paper is needed to further understand their chemical composition and evolutionary stage. Some of the post-AGB candidates given in Table 1 may show light variations similar to the high galactic latitude hot post-AGB stars LS II +34 26 and SAO 85766 \\citep{arkh07}.      All the spectral data used in this study are given in the digital form in electronic table E, which is available at the PASJ web site.    \\bigskip \\bigskip"
    },
    "1112/1112.3935_arXiv.txt": {
        "abstract": "Blue Supergiants (BSGs) are the brightest stars in the universe at visual light with absolute magnitudes up to M$_{V}$ = -10 mag. They are ideal stellar objects for the determination of extragalactic distances, in particular, because the perennial uncertainties troubling most of the other stellar distance indicators, interstellar extinction and metallicity, do not affect them. The quantitative spectral analysis of low resolution spectra of individual BSGs provides accurate stellar parameters and chemical composition, which are then used to determine accurate reddening and extinction from photometry for each individual object. Accurate distances can be determined from stellar gravities and effective temperatures using the \"Flux Weighted Gravity - Luminosity Relationship (FGLR)\".  Most recent results of the quantitative spectral analysis of BSGs in galaxies within and beyond the Local Group based on medium and low resolution spectra obtained with the ESO VLT and the Keck telescopes on Mauna Kea are presented and distances obtained with the FGLR-method are discussed together with the effects of patchy extinction and abundance gradients in galaxies. BSG metallicities and metallicity gradients are compared with results from strong-line \\hii~region studies and the consequences for the empirical calibration of the metallicity dependence of the Cepheid period - luminosity relationship are pointed out. The perspectives of future work are discussed, the use of the giant ground-based telescopes of the next generation such as the TMT on Mauna Kea and the E-ELT and the tremendous value of the GAIA mission to allow for the ultimate calibration of the FGLR using galactic BSGs. ",
        "introduction": "After the detection of the accelerated expansion of the universe the physical explanation of dark energy has become the major challenge of astronomy and physics. One way to constrain the physics behind dark energy is to measure the equation-of-state parameter w = p/($\\rho$c$^{2}$). This requires an extremely accurate determination of the extragalactic distance scale and the Hubble Constant H$_{0}$. As is well known \\citep{macri06}, the determination of cosmological parameters from the cosmic microwave background is affected by degeneracies in parameter space and cannot provide strong constraints on the value of H$_{0}$ \\citep{spergel06, tegmark04}. Only if additional assumptions are made, for instance that the universe is flat, H$_{0}$ can be predicted with high precision (i.e. 2\\%) from the observations of the cosmic microwave background, baryonic acoustic oscillations and type I high redshift supernovae. If these assumptions are relaxed, then much larger uncertainties are introduced \\citep{spergel07, komatsu09}. The uncertainty of the determination of w is related to the uncertainty of H$_{0}$ through $\\Delta$w/w $\\approx$ 2$\\Delta$H$_{0}$/H$_{0}$. Thus, an independent determination of H$_{0}$ with an accuracy of 5\\% will allow the uncertainty of w to be reduced to $\\pm$0.1 or, even more ambitious, H$_{0}$ accurate to 1\\% will yield w accurate to 0.02.  Extremely promising steps towards this goal have been made recently by \\citet{macri06} and by \\citet{riess09a,riess09b,riess11}, who in a most ambitious approach use the maser galaxy NGC 4258 and its Cepheids as a new anchor point for the extragalactic distance scale. HST detections of Cepheids out to 30 Mpc in galaxies with SNIa, which were detected recently and have accurate and well understood light cuves, are then used to consistently calibrate SNIa as far reaching standard candles using HST H-band photometry of the host galaxy Cepheids. In this way, \\citet{riess11} have been able to determine H$_{0}$ with unprecedented accuracy. They estimate an error of only 4\\% including both systematic and random uncertainties (see papers by Lucas Macri and Adam Riess in this volume).  However, it is clear that the complexity of this bold and courageous approach requires additional and independent tests. There are two long-standing issues with the use of Cepheids as stellar distance indicators, the perennial problem of interstellar extinction and the uncertainty about a potential metallicity dependence of the period-luminosity relationship. While both problems can be mitigated to some extend by near IR observations and a strictly differential approach, as applied by \\citet{riess11}, at a few percent accuracy level of H$_{0}$ they continue to be a major concern.  In order to address these concerns, an independent and complementary method is desirable, which can overcome the problems of interstellar extinction and variations of chemical composition. In this paper, we introduce such a method, the flux weighted gravity - luminosity relationship (FGLR), which is based on the quantitative spectroscopy of blue supergiant stars (BSGs).  In the following we briefly illustrate the still acute problems caused by extinction and metallicity. We then introduce the spectroscopy of BSGs as an ideal tool to obtain accurate information about reddening, metallicity and distance. We will discuss most recent applications of the BSG-method and future work. ",
        "conclusions": ""
    },
    "1112/1112.5912_arXiv.txt": {
        "abstract": "We report a detailed examination of the ``red asymmetry'' of H$\\alpha$ emission line seen during the 2001 April 10 solar flare by using a narrowband filtergram. We investigated the temporal evolution and the spatial distribution of the red asymmetry by using the H$\\alpha$ data taken with the 60cm Domeless Solar Telescope at Hida Observatory, Kyoto University. We confirmed that the red asymmetry clearly appeared all over the flare ribbons, and the strong red asymmetry is located on the outer narrow edges of the flare ribbons, with the width of about $1.^{\\prime\\prime}5$ -- $3.^{\\prime\\prime}0$ (1000 -- 2000~km), where the strong energy releases occur. Moreover, we found that the red asymmetry, which also gives a measure of the Doppler shift of the H$\\alpha$ emission line concentrates on a certain value, not depending on the intensity of the H$\\alpha$ kernels. This implies not only that the temporal evolutions of the red asymmetry and those of the intensity are not in synchronous in each flare kernel, but also that the peak asymmetry (or velocity of the chromospheric condensation) of individual kernel is not a strong function of their peak intensity. ",
        "introduction": "Solar flares have been extensively observed in H$\\alpha$ line, since they show spectacular phenomena in this wavelength, such as filament/prominence eruptions, flare ribbons, post flare loops (or coronal rains), and so on (e.g. \\cite{Tan88}). H$\\alpha$ kernels, very bright and compact H$\\alpha$ emission sources seen in flare ribbons, are also one of the prominent features that appear during a flare. Such flare ribbons and kernels have been thought to be caused by the energy deposition from the corona into the chromosphere.  To explain these various phenomena successfully, magnetic reconnection models have been widely discussed for last several decades. Especially, the so-called CSHKP magnetic reconnection model \\citep{Car64,Stur66,Hira74,Kopp76} is often accepted as the ``standard'' model. In the scenario, the connectivity of magnetic field lines is topologically changed (i.e. re-connected) in the corona, and then, vast magnetic energy stored in the corona is released. The released energy is converted to various kinds of energy; kinetic energy spent for ejections of plasma, thermal energy to generate hot plasma of about 10~MK or more, nonthermal energy to accelerate energetic particles, and so on. The thermal conduction and/or the nonthermal particles rapidly travel along newly connected magnetic field lines. Then, they reach the chromosphere, followed by the bombardment to the chromospheric plasma. H$\\alpha$ flare ribbons and kernels are caused by such bombardment. The CSHKP model further suggests that magnetic field lines successively reconnect in the corona. During the successive magnetic reconnection, the reconnection points (X-points) move upward. As a result, newly reconnected field lines have their footpoints farther out from the neutral line than the footpoints of the field lines that have already reconnected. Therefore, the separation of the flare ribbons is not a real motion, but an apparent one (see also the reviews of the CSHKP model by \\cite{Stur92} and \\cite{Sve92}).  The energy deposition into the chromosphere also causes strong distortions of H$\\alpha$ line (line width and line shift) besides the emission. The enhancements of the intensity at red wings are particular, and have been reported as ``red asymmetry'' of the lines \\citep{Sve76}, while ``blue asymmetry'' is also reported to be sometimes seen at very early phase of a flare. Comparing between the red and the blue wing images, it has known that such red asymmetry is seen overall the flare ribbons \\citep{Jan70,Sve76,Tang83}. \\citet{Hana03} examined the H$\\alpha$ impact polarization (liner polarization) due to collisions of high-energy particles at flare kernels, and reported that red asymmetry also appears at such kernels during the rise phase of kernel brightenings. Observations on the red asymmetry with filtergraphs, however, have not been progressed so much since then, and the further characteristics on the spatial distribution of the red asymmetry remains to be revealed.  \\citet{Ichi84} reported that the red-shift corresponds to the downward motion with speed of several tens of km~s$^{-1}$, based on their excellent spectroscopic observations. \\citet{Fal97} further reported that strong red asymmetry, which corresponds to the downward velocity of the order of tens of km~s$^{-1}$, is observed at the outer edge of the flare ribbons. The width of the edges is about 2 -- 3$^{\\prime\\prime}$ ($\\sim$ 1500 -- 2000 km). This is reasonable in the context of magnetic reconnection theories suggesting that strong energy depositions occur at the outer edges of the flare ribbons connected by the newly reconnected loops. \\citet{Ichi84} also quantitatively examined the red asymmetry. The strong enhancement of the chromospheric plasma due to the precipitation of nonthermal plasma and/or thermal conduction causes not only red asymmetry of the chromospheric lines but explosive upflow of the heated plasma along the flare loops. The latter is called ``chromospheric evaporation'', and has been observed as blue-shift of coronal emission lines \\citep{Hira74,Anto84}. \\citet{Ichi84} found a momentum balance between the downward motion of the chromospheric condensation and the chromospheric evaporation, and confirmed that the red asymmetry is a counteraction of the chromospheric evaporation.  Red asymmetry, in this way, tells us the quantitative information of the precipitation (nonthermal particles / thermal conduction). Understanding the features of the chromospheric plasma at the footpoints of flare loops is strongly related with understanding energy release/deposit mechanisms and is crucially important for solar flare studies. It has been, however, a challenge theoretically and numerically to explain exactly those observational facts (e.g. \\cite{Fish85a,Fish85b}), since the distortion of chromospheric lines is too complicated. Observations on the red asymmetry, have not been progressed so much since \\citet{Ichi84}. Especially, spatial distribution of the red asymmetry using narrowband filtergram, remains to be revealed, although \\citet{Tang83} reported that red asymmetry features appear all over flare ribbons and continue from onset until late in the decay phase by using narrowband filtergram of H$\\alpha$ line. In this paper we examined the spatial distribution of the red asymmetry at H$\\alpha$ kernels and the temporal evolution, by using H$\\alpha$ wing images. In \\S~2 we summarize the observational data and the features of the flare. In \\S~3 we show the features of the red asymmetry. In \\S~4 a summary and discussions are given. ",
        "conclusions": "We examined the spatially resolved red asymmetry distributions seen in the 2001 April 10 solar flare and their temporal evolution, by using the H$\\alpha$ wing images taken with DST at Hida Observatory. We found that red asymmetry appears all over the flare ribbons, both for bright kernels generated by the bombardment of nonthermal particles and for other fainter regions generated by the (thermal) conduction. We confirmed that the strong $RA$ appears at the outer edge of the flare ribbons that are the footpoints of newly reconnected flare loops. The width of about 1.5 -- 3.0$^{\\prime\\prime}$ ($\\sim$ 1000 -- 2000~km). This supports that the strong energy release occurs on the magnetic loops that connect the outermost edges of the flare ribbons. We found a clear tendency that the larger difference between the H$\\alpha$ red and the blue wing intensities $DF$ is associated with the brighter H$\\alpha$ kernels $SM$. It turns out to be that the red asymmetry ($RA = (I_{red} - I_{blue}) (I_{red} + I_{blue})^{-1} =DF/SM$) does not depend on the intensity of the H$\\alpha$ kernels, and scatters around a mean value of about 0.37.  In Figure~3 we showed that the temporal evolution of $RA$ does not synchronize with that of the intensity, and the peak times of $RA$ time profiles precedes those of the intensities. Moreover, the enhancements of $RA$ time profiles procede to those of the nonthermal emissions observed in hard X-rays and/or in microwaves. \\citet{Zarro89} also mentioned that redshift in the H$\\alpha$ wings starts to appear before the enhancement of the hard X-ray emission, that is, strong energy release occurs. At flare kernels, not only nonthermal particles but thermal conduction precipitate from the energy release site in the corona into the chromosphere. The preceding enhancement of $RA$ to the hard X-ray emission could be caused by such other mechanisms rather than by nonthermal particles.  To make clear the temporal evolutions of $RA$ and intensity, we overlaid the temporal trajectories of the H$\\alpha$ kernels, shown in Figure~3 with the $\\times$ and $\\triangle$ sings, on the scatter plot between $RA$ and the summed intensity $SM$, which is presented as Figure~6. The $RA$ plots move counter-clockwise on the scatter plot, and after the $RA$ enhancements, the intensities reaches the peak. The absence of a clear correlation with the intensity is because the time profile of $RA$ for individual kernels evolves independently with a manner that the peak downward velocity is not a strong function of their peak intensity. \\citet{Can90b} also reported that there is no one-to-one relation between brightness of H$\\alpha$ line center and redshift, while they found many redshift features at bright kernels at H$\\alpha$ line center.  If we assume a constant width of the H$\\alpha$ emission line, we can regard $RA$ as an indicator of the velocity. For example, assuming a gaussian function with the full-width-half-maximum $FWHM$ in {\\AA}, $RA$ can be related to the Doppler (redward) velocity as $2.75 \\times \\ln{(1+RA)(1-RA)^{-1}} (FWHM ({\\rm{\\AA}}))^2$~km~s$^{-1}$. The estimation of the width of the H$\\alpha$ emission line is, however, very difficult, because the shape of the line is drastically changed (it cannot be fitted with a gaussian function) and especially the Stark effect broadens the line at the wings \\citep{Sve76}. From the results of \\citet{Ichi84} (in the Figs. 3b and 4b), we estimated $RA$ for their flares, and they are about 0.2 -- 0.5, which is comparable to the 2001 April 10 flare. They also reported that the Doppler velocities are 10 -- 100~km~s$^{-1}$, the FWHM of about 4.5 -- 6.0 is expected. If we assume that the FWHM is about 5.0, here, $RA$ of 0.37 turns out the Doppler (downward) velocity of about 53~km~s$^{-1}$. Scattering of $RA$ around a certain value of about 0.37 is also because the time profiles of $RA$ keep high values for several minutes even after the intensity peaks. This means that the downward motion at flare ribbons continues with the velocity of about several tens of km~s$^{-1}$ independently from the intensity.   \\begin{figure} \\begin{center} \\FigureFile(150mm,75mm){f6.eps} \\end{center} \\caption{ Temporal trajectories of H$\\alpha$ kernels in the $RA$ - $SM$ scatter plot. The temporal evolutions of $RA$ are shown in the right panels of Figure~3, and the marks of the H$\\alpha$ kernels (cross $\\times$ and triangle $\\triangle$ signs) are also the same as those shown in Figure~3. The arrows guides the temporal evolutions. The background $RA$ -- $SM$ scatter plots are the same as the right panel of Figure~5. }\\label{fig:scat} \\end{figure}   \\bigskip %  Acknowledgements  We first acknowledge an anonymous referee for her/his useful comments and criticisms. 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."
    },
    "1112/1112.0530_arXiv.txt": {
        "abstract": "Galactic X-ray emission is a manifestation of various high-energy phenomena and processes. The brightest X-ray sources are typically accretion-powered objects: active galactic nuclei and low- or high-mass X-ray binaries. Such objects with X-ray luminosities of $\\gtrsim 10^{37} {\\rm~ergs~s^{-1}}$ can now be detected individually in nearby galaxies. The contributions from fainter discrete sources (including cataclysmic variables, active binaries, young stellar objects, and supernova remnants) are well correlated with the star formation rate or stellar mass of galaxies. The study of discrete X-ray sources is essential to our understanding of stellar evolution, dynamics, and end-products as well as accretion physics. With the subtraction of the discrete source contributions, one can further map out truly diffuse X-ray emission, which can be used to trace the feedback from active galactic nuclei, as well as from stars, both young and old, in the form of stellar winds and supernovae. The X-ray emission efficiency, however, is only about 1\\% of the energy input rate of the stellar feedback alone. The bulk of the feedback energy is most likely gone with outflows into large-scale galactic halos. Much is yet to be investigated to comprehend the role of such outflows in regulating the ecosystem, hence the evolution of galaxies. Even the mechanism of the diffuse X-ray emission remains quite uncertain. A substantial fraction of the emission cannot arise directly from optically-thin thermal plasma, as commonly assumed, and most likely originates in its charge exchange with neutral gas. These uncertainties underscore our poor understanding of the feedback and its interplay with the galaxy evolution. ",
        "introduction": "Why are X-rays from galaxies interesting? X-ray emission from a galaxy typically arises from high-energy (high temperature or relativistic) processes in extreme conditions. % Such processes are important to understand in their own right, but are also believed to play a major role in shaping the structure and evolution of galaxies via various types of feedback (e.g., Tang et al. 2009; Oppenheimer et al. 2010). In particular, the X-ray emission effectively traces diffuse hot plasma, which may dominate the volume and account for the bulk of the baryon mass in and around galaxies, especially massive ones (e.g,, Crain et al. 2010). The X-ray wavelength range also contains major atomic transitions (e.g., all K-shell transitions of carbon through iron), providing key thermal, chemical, and kinematic diagnostics of the interstellar medium (ISM) in all phases (cold, warm, and hot) and all forms (atomic, molecular, and dust). X-rays are also energetic and penetrating, which may heat and evaporate dust grains/molecules and even companion stars in X-ray binaries, significantly affecting stellar and interstellar radiation. X-rays can often be observed through dense clouds with column densities up to $\\gtrsim 10^{24} {\\rm~cm^{-2}}$. Therefore, X-ray observations are a powerful tool in examining the stellar and interstellar properties of galaxies.  Substantial progress in understanding the X-ray emission from nearby galaxies has been made over the past decade or so. \\Chandra\\ and \\xmm\\ X-ray Observatories, in particular, have provided a combination of  superb spatial and spectral resolutions (reaching $\\sim 1^{\\prime\\prime}$ and $500 {\\rm~km~s^{-1}}$ FWHM) as well as large collecting areas (up to $\\sim 3000$ cm$^2$) over a broad energy range of 0.3-10 keV. These capabilities have enabled extensive studies of various classes of X-ray sources: their luminosity functions, spectral characteristics, and relationships to other galactic properties such as the star formation rate (SFR) and total stellar mass. These are the topics that I am going to review, focusing on recent results. ",
        "conclusions": "I have reviewed some basic properties of various types of X-ray sources, particularly in relation to the spectral energy distribution (SED) of galaxies. The key points are as follows: \\begin{itemize} \\item AGNs are typically the brightest X-ray sources in galaxies and can be detected relatively easily even at high redshifts with little confusion. Such detection is important in modeling the SED of galaxies. Even low-$L_x$ AGNs, as typically seen in nearby galaxies, can be important in regulating galactic nuclear environments.  \\item In nearby galaxies, the bulk of HMXBs and LMXBs can readily be detected. Their overall galactic populations are well correlated with the SFR and stellar mass and can also be independently estimated (e.g., with infrared observations). The comparisons of these independent estimates can then be used to constrain the presence of AGNs in distant galaxies, even if they are not well resolved.  \\item The X-ray source population  of a stellar system is particularly sensitive to its binary fraction. Thus X-ray observations provide a powerful tool to probe the initial binary fraction and dynamical evolution of stellar clusters and even galactic spheroids.  \\item Diffuse X-ray emission traces the galactic feedback, primarily due to massive SF (via fast stellar winds and core-collapsed SNe) in galactic disks, to old stars (mass-loss and Ia SNe), and to possible AGNs. However, the  luminosity of the X-ray emission from a galaxy typically accounts for $\\lesssim 1\\%$ of the feedback energy. The bulk of the energy is released in outflows of hot plasma and cosmic-rays/magnetic fields from stellar spheroids, as well as from SF regions.  \\item A substantial fraction of the diffuse X-ray emission may arise from the CX at the interface between hot plasma and cool gas. The CX, in principle, can be used as a new tool to probe the thermal, chemical, and kinematic properties of hot plasma tracing various high-energy feedback processes in the galaxies. \\end{itemize}  While much can still be done with the existing X-ray observatories, we will also have new tools in the near future. {\\sl NuSTAR} (scheduled to be launched in late 2012) will offer the first imaging capability in the 8-80 keV range, while {\\sl Astro-H} (2013) will provide the calorimeter-based imaging spectroscopic capability in the 0.5-10 keV range. Furthermore, {\\sl eROSITA} (2013) will conduct the first all-sky X-ray CCD imaging survey over the 0.3-10 keV range, while {\\sl GEMS} (2014) will enable polarization measurements in the 2-10 keV range. With these new capabilities, we can expect a substantial advance in our understanding of X-ray sources in galaxies.  I am grateful to the organizers of the IAU Symposium for inviting me to give a review talk on this topic, which the present article is based on, and to my students and collaborators for their contributions to some of the work described above. My research is largely supported by NASA via various grants."
    },
    "1112/1112.5315_arXiv.txt": {
        "abstract": "The low mass X-ray binary (LMXB) \\source in the globular cluster Terzan 5 harbors an 11 Hz accreting pulsar. This is the first object discovered in a globular cluster with a pulsar spinning at such low rate. The accreting pulsar is anomalous because its characteristics are very different from the other five known slow accreting pulsars in galactic LMXBs. Many features of the 11 Hz pulsar are instead very similar to those of accreting millisecond pulsars, spinning at frequencies $>\\,100$Hz.  Understanding this anomaly is valuable because \\source can be the only accreting pulsar discovered so far which is in the process of becoming an accreting millisecond pulsar. We first verify that the neutron star (NS) in \\source is indeed spinning up by carefully analyzing X-ray data with coherent timing techniques that account for the presence of timing noise.  We then study the present Roche lobe overflow epoch and the two previous spin-down epochs dominated by magneto dipole radiation and stellar wind accretion. We find that \\source is very likely a mildly recycled pulsar and suggest that it has started a spin-up phase in an exceptionally recent time, that has lasted less than a few $10^7$ yr. We also find that the total age of the binary is surprisingly low ($\\simless10^8$ yr) when considering typical parameters for the newborn NS and propose different scenarios to explain this anomaly. ",
        "introduction": "According to the recycling scenario (\\citealt{alp82},~\\citealt{rad82}), millisecond radio pulsars are produced by spin-up of the neutron stars (NSs) via transfer of angular momentum through accretion in low mass X-ray binaries (LMXBs). During this phase, channeled accretion of plasma along the NS magnetic field lines can produce X-ray pulsations that reveal the spin frequency of the accreting NS.  A confirmation of this scenario came with the discovery of the first accreting millisecond X-ray pulsar SAX J1808.4-3658, spinning at a frequency of about 401 Hz \\citep{wij98}.  Thirteen more accreting millisecond pulsars have been discovered so far with spin frequencies ranging from about 180 up to 600 Hz.  In 2010 October the new accreting pulsar IGR J17480-2446 has been discovered in the globular cluster Terzan 5 \\citep{bor10, str10}. \\citet{poo10} localized the source in outburst with \\textit{Chandra} observations and identified it with a previously known candidate quiescent LMXB, named CX25 in the X-ray source catalog of \\citet{hei06}. The source is in core of Terzan 5, one of the densest among globular clusters \\citep{fer09}.  The accreting pulsar spins with a frequency of 11 Hz and the LMXB has a period of 21.3 hr around a companion with mass $M>0.4\\,M_{\\odot}$ \\citep{pap11}. Although this pulsar is not a millisecond one, its location in a globular cluster makes it a unique system whose properties are of particular importance to study the recycling mechanism. \\source is different from the other five known slow accreting pulsars in LMXBs (4U 1626$-$67, 2A 1822$-$371, GRO 1744-28, GX 1+4 and Her X-1) with properties which are in between those systems and the accreting millisecond pulsars \\citep{lin12}.  Slow accreting pulsars are a very heterogeneous group of LMXBs, with white dwarf, main sequence and giant companions and orbital periods ranging from less than 1 hr to more than 11 days. Despite this, the five slow pulsars have a high magnetic field $B\\sim10^{11}-10^{13}\\rm\\,G$ and do not show thermonuclear bursts. This high $B$ field is difficult to reconcile with the long lifetime of these slow systems, and several scenarios have been proposed (see, for example, \\citealt{ver90} for a discussion).  \\source instead shows thermonuclear bursts (\\citealt{che10}, \\citealt{lin11}) with burst oscillations phase locked with the accretion powered pulsations \\citep{cav11}, and the $B$ field has been constrained to be in the range $10^9-10^{10}$G (\\citealt{mil11}, \\citealt{cav11}, \\citealt{pap11}).  All these field estimates point toward a mildly recycled pulsar, since accreting millisecond pulsars have shown so far field strengths of the order of $10^8$G (in agreement with the $B$ fields of radio millisecond pulsars). Phase locking between burst oscillations and accretion powered pulsations has also been observed in at least one accreting millisecond pulsar (\\citealt{str03}, \\citealt{wat08}). The 11 Hz rotation rate is, however ,still unusually slow for an old NS, the typical rotation periods for accreting or rotation powered pulsars in globular clusters and old stellar populations being in the millisecond range.  The reason for old NSs possessing dipole magnetic fields smaller by about three orders of magnitude from the typical $10^{12} \\rm\\,G$ fields of young NSs has elicited several explanations involving binary evolution. The burial of the field under accreted material has been proposed as a viable mechanism to dissipate the field via ohmic decay in the heated NS crust (see for example \\citealt{rom90}). This mechanism, however, has a dramatic consequence for the final magnetic field: as the accreted matter keeps falling on the NS surface, the magnetic field lines will be pushed to deeper (and colder) regions of the crust where the conductivity is much larger, thus freezing at the crust-core interface. This will leave a residual field that is compatible with the observed weak fields of millisecond pulsars.  Another particularly simple and elegant explanation relies on the superfluid-superconducting nature of the NS interior, requiring that the rotation is carried by quantized vortex lines in the neutron superfluid and the magnetic flux by quantized flux lines in the proton superconductor \\citep{sri90}. Spin down of the star is achieved by outward motion of the vortex lines, which inevitably entangle and drag flux lines with them. Srinivasan et al.  noted that many NS binary evolution scenarios have an early first mass transfer epoch when the companion, not yet filling its Roche lobe, is losing mass by a stellar wind, some of which is captured by the NS in a weak quasi spherical accretion with low specific angular momentum. The NS will be spinning down in this epoch, thus reducing its dipole field in proportion to its rotation rate. The wind spin-down epoch ends when the companion fills its Roche lobe. The ensuing disk accretion now starts to spin-up the NS.  If IGR J17480-2446 is a primordial binary, then its spin frequency is surprising low (the globular cluster Terzan 5 is about 12 Gyr old, see \\citealt{fer09}) and requires investigation to understand whether this is the consequence of a peculiar evolutionary history.  At present \\source is likely to be in its spin-up epoch \\citep{pap11}. The results reported by \\citet{pap11} are however affected by the usual problem of X-ray timing noise in the pulse time of arrivals (TOAs) that when not taken properly into account might strongly affect the significance of the spin frequency derivatives (see for example the case of the other three accreting millisecond X-ray pulsars in \\citealt{har08,har09}, \\citealt{pat10} \\citealt{has11}).  In this work we will re-investigate the existence of a spin-up phase by taking into account the effect of timing noise. If the spin-up is confirmed, then the comparatively long spin period indicates either that \\source is at an exceptionally early phase of the spin-up epoch or that it is close to spin equilibrium, with an unusually long equilibrium period reflecting a magnetic field stronger than in most NSs in LMXB.  The aperiodic variability can also help to understand the peculiar behavior of the source. \\citet{alt10} reported several quasi-periodic oscillations (QPOs) at frequencies ranging from about 48 up to 815 Hz.  This sequence of frequencies make possible several alternative estimations of the inner disk radius at present. Each such model gives an estimate of the dipole magnetic field, equilibrium rotation rate and expected spin-up rate.  The location of \\source in the globular cluster Terzan 5 allows also to place stronger constraints on the evolutionary history of the binary since only a narrow range of donor masses is allowed. The distance of the system is also constrained by the globular cluster and leads to a precise determination of the X-ray luminosity. In Section~\\ref{xray} we describe the X-ray observations carried during the 2010 outburst, and in Section~\\ref{results} we verify the existence of a spin-up phase.  In Section~\\ref{evolution} we present the evolutionary history of \\source. We discuss three distinct epochs: in the initial two the NS has been spun down first via magneto dipole radiation and then by wind accretion, whereas in the third epoch the NS is spun-up via accretion of mass provided by the donor in Roche lobe overflow (RLOF).  The evolution of the magnetic moment of the NS through the earlier dipole and wind spin-down epochs is also discussed.  In Section~\\ref{constraints} we place further constraints on the binary parameters by using the available information on the globular cluster Terzan 5 and the results of the evolutionary history. We conclude with a discussion of the results in Section~\\ref{discussion} and infer that \\source is in an exceptionally early RLOF phase that explains why its spin frequency is still so small when compared to accreting millisecond pulsars. ",
        "conclusions": "\\label{discussion}  We have investigated the evolutionary history of \\source discussing three distinct evolutionary epochs: a dipole dominated spin-down epoch, a wind epoch and the currently observed RLOF phase. By taking into account the effect of timing noise in the X-ray pulsations we are able to confirm the long-term average spin-up of $1.4\\times10^{-12}\\rm\\,Hz\\,s^{-1}$ observed during the 2010 outburst (\\citealt{pap11,cav11}). If the X-ray flux represents most of the bolometric flux, the short term spin frequency derivatives $\\dot{\\nu}$ do not scale with the X-ray flux in the expected way. The long-term $\\dot{\\nu}$, however, leads to exclude a scenario in which \\source has already reached spin equilibrium. A long-term spin equilibrium with magnetic dipole spin down in quiescence balancing the accretion torques in outburst is also excluded based on binary evolution considerations.  The spin-up timescale suggests that \\source has been spun-up in the RLOF phase for a few $10^{7}$ yr. An interpretation of the observed QPOs as the equilibrium frequencies of \\source indicates a similar timescale for the spin-up epoch up to the present.  The spin-down epochs suggest that the NS in \\source had reached a spin frequency of the order of $\\sim1$Hz at the onset of the RLOF phase, if we assume that the newborn NSs had typical initial values for $\\Omega_0$ and $\\mu_0$.  This reinforces the idea that the time elapsed in RLOF has been very short, since the time required to spin-up a 1 Hz NS to 11 Hz with a spin frequency derivative of $10^{-12}\\rm\\,Hz\\,s^{-1}$ and a duty cycle of 0.01 is a few $10^7$yr.  Based on these findings we conclude that \\source is in an exceptionally early RLOF phase. The total spin-up timescale to transform \\source from a slow pulsar ($\\sim\\,1\\rm\\,Hz$) into a millisecond one ($\\nu>100\\rm\\,Hz$) is $\\simgreat10^8-10^{9}$yr.  Since the current RLOF phase will last for $\\sim10^9$yr, we would expect to observe today $\\sim1-10$ accreting millisecond pulsars that have followed a similar evolutionary history as \\source. The accreting millisecond pulsar SAX J1748.9-2021 ($\\nu\\simeq442\\rm\\,Hz$) in the globular cluster NGC 6440 has a companion with mass and orbital period compatible with being in a slightly evolved post main sequence phase \\citep{alt08}. It is likely that the companion of SAX J1748.9-2021 has followed an evolutionary history similar to \\source with the difference that it spent a longer time in the RLOF phase thus turning its NS into a millisecond pulsar. The scenario outlined above is compatible with the observed population of accreting pulsars, although it is difficult to assess the problem in a robust way given the small number statistics. It is also possible that dynamical interactions have played a role by decreasing the number of similar systems in globular clusters via ionization interactions. This however requires further investigations along with detailed binary evolution calculations.  The prior spin-down history gives results which are not compatible with a binary having an age of several $10^9$ yr as expected if \\source is primordial or if it has not suffered exchange interactions in the last few $10^9$ yr. Different formation scenarios are available to explain the apparent discrepancy between the age of the cluster and the binary.  A first possibility is that a recent dynamical encounter has played a role in forming the binary or in accelerating the onset of the RLOF phase. If an exchange interaction has taken place in the last few $10^7-10^8$ yr then this would explain why the apparent age of \\source is so short compared with the age of the cluster. The location of \\source in the high mass and high central density cluster Terzan 5 suggests that the interaction rate might be particularly high in this environment. Dynamical simulations to calculate the typical interaction rates in the core of Terzan 5 are required to assess this question.  A second possibility for the origin of \\source is formation of the accreting pulsar via accretion induced collapse (AIC) of a massive white dwarf (\\citealt{miy80}, \\citealt{nom87}). In this scenario a ONeMg white dwarf accretes matter until its mass exceeds the Chandrasekhar limit and it collapses to form an NS via electron captures on Mg and Ne nuclei. In this case the binary must have gone through a preliminary contact phase during which the white dwarf was accreting from a donor star in RLOF. During the collapse approximately 0.2$\\msun$ are ejected from the binary (see for example \\citealt{van11} for a discussion) which causes a sudden expansion of the orbit turning the system into a detached binary. At this point the binary follows the three evolutionary epochs described in Section~\\ref{evolution} with the NS age set by the onset of AIC.  The formation of NSs from AIC has been discussed by \\citet{lyn96} to explain the young radio pulsar population in globular clusters. Lyne et al. note that the formation rate of young pulsars in globular clusters is larger than that of millisecond pulsars in globular clusters. The young pulsars are found in metal rich globular clusters with high mass and high central densities, like Terzan 5 (\\citealt{lyn96,boy11}). Further evidence that globular clusters might contain young NSs was also presented by \\citet{fre11}, who reported the discovery of the youngest millisecond pulsar ever found in the globular cluster NGC 6624, with a characteristic age $t<24$ Myr (see also \\citealt{lyn87} and \\citealt{cog96} on the ``young'' millisecond pulsar PSR 1821-24A in the globular cluster M28).  We can therefore speculate that a similar evolutionary process has taken place for \\source thus explaining the relatively short lifetime of the NS through the dipole and wind spin-down and present RLOF spin-up epochs.  A possible test for the AIC scenario can be performed by measuring the gravitational mass of the NS in \\source. Since the RLOF phase has lasted for a very short time, the NS mass is currently almost identical to the value expected from the AIC scenario, which is $M\\simeq1.25\\msun$. A gravitational NS mass larger than this value would immediately rule out the AIC scenario."
    },
    "1112/1112.0706_arXiv.txt": {
        "abstract": "I present a review of the research and analysis paths that converged to make Type Ia SNe the most mature cosmological distance estimator of the present time. The narrative starts with the first works in the early decades of the 20th century and finishes with the more recent results, covering the surprising discovery of Cosmic Acceleration in 1998.  The review was written thinking of physicists with a strong interest in Cosmology, who might have pondered why was that, after decades of not being able to agree upon the rate of cosmic expansion, astronomers were so quick to concur on cosmic acceleration. ",
        "introduction": "\\label{intro} Supernovae (SNe) have been with us from the beginning, since before we were ``us''. Occasionally superimposed on the background of a familiar, regular and repetitive sky, they were among the variable phenomena that puzzled, marveled and scared our human ancestors. The modern view associates them with the final stages of stellar evolution in scenarios where the shifting balance between pressure sources and gravitation reaches unstable regimes. When the instabilities lead to an explosion powerful enough to disrupt the star we have SNe.  The current paradigms purport SNe as critical ingredients in the micro-physics of Cosmology. They create the chemical history of the Universe, stir the interstellar medium of galaxies enriching primordial matter with novel chemical elements, and, if exploding in dense environments, generate shock waves that promote the birth of newer generations of stars. In addition, from the practical point of view of today$'$s astronomers, they provide outstanding estimators for cosmologically relevant distances.  This paper is a review of the research paths that converged to make Type Ia SNe, at the time, the most mature distance estimator for use in Cosmology, the surprising discovery of Cosmic Acceleration in 1998, and touches the current work trying to use SNe to obtain even finer cosmological inferences. Opposite to what the organizers of the meeting expected, the paper was written not before, but {\\em after} the presentation at the workshop. The positive side of having done it this way, is that I could tune it to answer the questions I received after the talk. It was the first time I presented the subject to an audience made up, primarily, of physicists, and the questions were different to what I was typically used to answer. Hence, I prepared this work thinking mainly of colleagues who are physicist that work in Cosmology. I will touch some topics considered ``basic'' in SNe studies, which are available in other reviews or papers, but are scattered around mainly in astronomical journals. My hope is to present basic concepts, give an overview of the advances as they proceeded in time, and organize some relevant bibliographic sources to facilitate further study from scholars who are not specialist in the field.  I must warn the reader that the text, as a review of SNe, is not complete. It cannot be. It is biased towards the developments that lead Type Ia SNe to become the outstanding distance estimators that they are today. The rest of the SNe appear and are mentioned just as the background out of which Type Ia SNe appeared and were recognized.  I shall also warn the reader that I have made an attempt to merge in this manuscript the different styles of referencing literature used by physicists, like in Phillips \\cite{p1993}, and astrophysicists, like in Phillips (1993). Being myself an astrophysicist I have come to appreciate the importance of having the date of the papers inserted in the text, to figure out the pace of developments and/or possible historical relation between different pieces of work. The result, Phillips (1993 \\cite{p1993}) is slightly cumbersome, but I hope that will suite the prejudices and slants of both communities.  The chart of this paper is as follows. Section 2 is a historical narrative. In it I describe the evolution of the area that I loosely call ``SN Studies,'' from the times of the pioneers, in the early decades of the 20th Century until all the major advances in instrumentation, observation, analysis, and theoretical interpretation had taken place, leaving the field poised for a major breakthrough by 1993. In Section 3, I describe the fast developments between 1994 and 1998, when the Cosmic Acceleration was discovered. In Section 4, I present the results of 1998, using as an illustration the complete set of Type Ia SNe published by the High Z SN Search Team, augmented by selected later discoveries. In Section 5, I describe what could be called the ``Second Generation'' SN surveys started soon after the turn of the century, with the goal of constraining the nature of the now called ``Dark Energy''. ",
        "conclusions": ""
    },
    "1112/1112.5253_arXiv.txt": {
        "abstract": "We construct generic spherically symmetric thin-shell gravastars by using the cut-and-paste procedure. We take considerable effort to make the analysis as general and unified as practicable; investigating both the internal physics of the transition layer and its interaction with ``external forces'' arising due to interactions between the transition layer and the bulk spacetime.  Furthermore, we discuss both the dynamic and static situations. In particular, we consider ``bounded excursion'' dynamical configurations, and probe the stability of static configurations. For gravastars there is always a particularly compelling configuration in which the surface energy density is zero, while surface tension is nonzero.  \\bigskip \\noindent PACS: 04.20.Cv, 04.20.Gz, 04.70.Bw  \\bigskip \\noindent \\TODAY;  \\LaTeX-ed \\today.  \\bigskip \\noindent Published: JCAP {\\bf03} (2012) 034.   doi:10.1088/1475-7516/2012/03/034  \\bigskip \\noindent ArXiv eprint: 1112.5253 ",
        "introduction": "Gravastars (gravitational-vacuum stars) are hypothetical objects mooted as alternatives to standard Schwarzschild black holes~\\cite{gr-qc/0012094, gr-qc/0109035}.  Typically the interior is some simple nonsingular spacetime geometry such as de~Sitter space, the exterior is some close approximation to the Schwarzschild geometry, and there is some complicated transition layer near the location where the event horizon would otherwise have been expected to form. Thus, in the traditional gravastar picture, the transition layer replaces both the de Sitter and the Schwarzschild horizons, and consequently the gravastar model has no singularity at the origin and no event horizon, as its surface is located at a radius slightly greater than the Schwarzschild radius. In this model, the quantum vacuum undergoes a phase transition at or near the location where the event horizon is expected to form. Considerable attention has been devoted to these objects, and to closely related ``dark stars'', ``quasi black holes'', ``monsters'', ``black stars'', and the like~\\cite{darkstar, quasi-black-hole, monster, black-star}. Related models, analyzed in a different context, have also been considered by Dymnikova \\cite{Dymnikova}. Some models use a continuous distribution of stress-energy, which on rather general grounds must be anisotropic in the transition layer~\\cite{anisotropic, anisotropic2}. Other models idealize the transition layer to being a thin shell~\\cite{gr-qc/0310107} and apply a version of the Sen--Lanczos--Israel junction condition formalism~\\cite{junction_formalism}. In fact, the latter approaches have been extensively analysed in the literature, and applied to a wide variety of scenarios \\cite{gravastar2,thinshell2}.  Several  criteria related to potential observability have been explored \\cite{observations}.  The key point of the present paper is to develop an extremely general and robust framework that can quickly be adapted to wide classes of generic thin-shell gravastars. We shall consider standard general relativity, with gravastars that are spherically symmetric, with the transition layer confined to a thin shell. The bulk spacetimes (interior and exterior) on either side of the transition layer will be spherically symmetric and static but otherwise arbitrary.  (So the formalism is simultaneously capable of dealing with gravastars embedded in Schwarzschild, Reissner--Nordstr\\\"om, Kottler, or de~Sitter spacetimes, or even ``stringy'' black hole spacetimes.  Similarly the gravastar interior will be kept as general as possible for as long as possible.)  The thin shell (transition layer) will be permitted to move freely in the bulk spacetimes, permitting a fully dynamic analysis. This will then allow us to perform a general stability analysis, where gravastar stability is related to the properties of the  matter residing in the thin-shell transition layer.  Many of the purely technical aspects of the analysis are very similar in spirit to that encountered in a companion paper analyzing thin-shell traversable wormholes~\\cite{GLV} --- mathematically there are a few strategic sign flips --- but physically the current framework is significantly  different. Consequently we shall very rapidly find our analysis diverging from the traversable wormhole case. Further afield, we expect that the mathematical formalism developed herein will also prove useful when considering spacetime ``voids'' (manifolds with boundary)~\\cite{voids}.  This paper is organized in the following manner: In Section \\ref{secII} we outline in great detail the general formalism of generic dynamic spherically symmetric thin shells, and provide a novel approach to the linearized stability analysis around a static solution. In Section \\ref{secIII}, we provide specific examples and consider a stability analysis by applying the generic linearized stability formalism outlined in section \\ref{secII}. Finally, in Section \\ref{conclusion}, we shall draw some general  conclusions. Throughout this work, we adopt the sign conventions of Misner--Thorne--Wheeler \\cite{MTWbook}, with $c=G=1$. ",
        "conclusions": "\\label{conclusion}  In this work, we have developed an extremely general and robust framework leading to the linearized stability analysis of dynamical spherically symmetric thin-shell gravastars, a framework that can quickly be adapted to wide classes of generic thin-shells. We have built on a companion paper, which analyzed thin-shell traversable wormholes~\\cite{GLV}. We emphasize that mathematically there are a few strategic sign flips, which physically implies significant changes in the analysis. An important difference is the possibility of a vanishing surface energy density at certain specific shell radii, for the generic thin-shell gravastars considered in this work. Due to the key sign flips, ultimately arising from the definition of the normals on the junction interface, the surface energy density is always negative for the thin-shell traversable wormholes considered in \\cite{GLV}. We have also explored static gravastar configurations to some extent, and considered the generic qualitative behavior of the surface stresses. Relative to the conservation law of the surface stresses, we have analysed in great detail the most general case, widely ignored in the literature, namely, the presence of a flux term, corresponding to the net discontinuity in the (bulk) momentum flux which impinges on the shell. Physically, the latter flux term can be interpreted as the work done by external forces on the thin shell.  In the context of the linearized stability analysis we have reversed the logic flow typically considered in the literature and introduced a novel approach. More specifically, we have considered the surface mass as a function of the potential, so that specifying the latter tells us how much surface mass we need to put on the transition layer. This procedure implicitly makes demands on the equation of state of the matter residing on the transition layer and demonstrates in full generality that the stability of the gravastar is equivalent to choosing suitable properties for the material residing on the thin shell. We have applied the latter stability formalism to a number of specific cases, namely, to the traditional gravastar picture, where the transition layer separates an interior de Sitter space and an exterior Schwarschild geometry; bounded excursion; close-to-horizon models, in both the extremal and non-extremal regimes; and finally, to specific example of an interior de Sitter spacetime matched to a charged dilatonic exterior. This latter case is particularly interesting, as it involves the flux term.  In conclusion, by considering the matching of two generic static spherically symmetric spacetimes using the cut-and-paste procedure, we have analyzed the stability of thin-shell gravastars. The analysis provides a general and unified framework for simultaneously addressing a large number of gravastar models scattered throughout the literature.  As such we hope it will serve to bring some cohesion and focus to what is otherwise a rather disorganized and disparate collection of results.  \\ack  PMM acknowledges financial support from a FECYT postdoctoral mobility contract of the Spanish Ministry of Education through National Programme No.~2008--2011. NMG acknowledges financial support from CONACYT-Mexico. FSNL acknowledges financial support of the Funda\\c{c}\\~{a}o para a Ci\\^{e}ncia e Tecnologia through Grants PTDC/\\-FIS/\\-102742/2008 and CERN/FP/116398/2010. MV acknowledges financial support from the Marsden Fund, administered by the Royal Society of New Zealand."
    },
    "1112/1112.1470_arXiv.txt": {
        "abstract": "We present the ultraviolet (UV) color--color relation of early-type galaxies (ETGs) in the nearby universe ($0.05<z<0.12$) to investigate the properties of hot stellar populations responsible for the UV excess (UVX). The initial sample of ETGs is selected by the spectroscopic redshift and the morphology parameter from the Sloan Digital Sky Survey (SDSS) DR7, and then cross-matched with the Galaxy Evolution Explorer (GALEX) Far-UV (FUV) and Near-UV (NUV) GR6 data. The cross-matched ETG sample is further classified by their emission line characteristics in the optical spectra into quiescent, star-forming, and AGN categories. Contaminations from early-type spiral galaxies, mergers, and morphologically disturbed galaxies are removed by visual inspection. By drawing the $FUV-NUV$ (as a measure of UV spectral shape) $vs.$ $FUV-r$ (as a measure of UVX strength) diagram for the final sample of $\\sim$3700 quiescent ETGs, we find that the ``old and dead'' ETGs consist of a well-defined sequence in UV colors, the ``UV red sequence'', so that the stronger UVX galaxies should have a harder UV spectral shape systematically. However, the observed UV spectral slope is too steep to be reproduced by the canonical stellar population models in which the UV flux is mainly controlled by age or metallicity parameters. Moreover, 2~mag of color spreads both in $FUV-NUV$ and $FUV-r$ appear to be ubiquitous among any subsets in distance or luminosity. This implies that the UVX in ETGs could be driven by yet another parameter which might be even more influential than age or metallicity. ",
        "introduction": "The ultraviolet excess (UVX) refers to the phenomenon that the early-type galaxies (ETGs) exhibit a rising flux with a decreasing wavelength from 2500{\\AA} to the Lyman limit. Theoretical and observational evidence shows that the phenomenon originates from the hot helium-burning horizontal-branch (HB) stars and their progeny \\citep{oco99,bro00}. One of the most important observational constraints to understanding this phenomenon was given by \\citet{bur88}. They first explored the systematics of the UV radiation from the 31 local ETGs and the bulge of M31, and found that the UVX strength correlates with the nuclear spectral line index Mg$_2$ (and, weakly, with the central velocity dispersion and luminosity) in $quiescent$ galaxies. With the advent of the Galaxy Evolution Explorer (GALEX) satellite \\citep{mart05}, a large area, deep UV imaging survey gathers high-quality data for an unprecedentedly large sample of ETGs in the nearby universe. With the GALEX UV data cross-matched with the Sloan Digital Sky Survey (SDSS), \\citet{ric05} investigated the systematics of the ETGs at $z<0.2$, but found no correlation of the UVX strength against the Mg$_2$, D4000 indices or stellar velocity dispersion. On the other hand, \\citet{don07} found a tight correlation between the GALEX UV color and the Mg$_2$ index among the morphologically sorted ETGs in the RC3 catalog \\citep{dev91}. Recently, \\citet{bur11} recovered the correlation of Burstein et~al. from the GALEX UV imagings combined with the high-quality data of the SAURON sample, and also with the SDSS data for a more distant ETG sample. They argued that such correlations could be seen among the $quiescent$ ETGs only when the contaminations of star-forming galaxies and low-quality data were much more stringently removed than was done in \\citet{ric05}.  The UV color--color diagram is another powerful tool to investigate the properties of hot stellar populations responsible for the UVX. Particularly, the capability of the GALEX to observe the far-UV (FUV; 1528{\\AA}) and near-UV (NUV; 2271{\\AA}) bandpasses simultaneously \\citep{mor05} let us securely estimate the UV spectral shape without instrumental biases. As demonstrated by \\citet{dor95}, the UV color--color diagram, $FUV-NUV$ (as a measure of UV spectral shape) $vs.$ $FUV-r$ (or $FUV-V$; as a measure of UVX strength) can trace the UV spectral shape and amplitude. It is also useful to segregate different stellar population systems. \\citet{ric05} employed this tool but found no systematic variation among their ETGs at $z<0.2$. \\citet{ree07} presented another example to show that the brightest cluster galaxies (BCGs) are homogenous populations following an old stellar sequence with considerable amounts of color spread.  In this $Letter$, we expand significantly the sample size of previous studies from the latest data release of the GALEX and SDSS to construct the UV color--color relation of the ETGs in the nearby universe, and reexamine our current understanding of UVX. ",
        "conclusions": "In summary, we made two main findings with the latest GALEX UV data of the ETGs in the nearby universe. 1) The ``old and dead'' ETGs consist of a well-defined sequence in UV colors in which the stronger UVX galaxies should have a harder UV spectral shape systematically. The UV color distributions of ETGs at $0.05<z<0.12$ are very similar to those observed for the local RC3 galaxies. The comparison with the population synthesis models indicates that the hot HB stars are responsible for the observed ``UV red sequence'', while the contamination of recently-formed young stars can be effectively discriminated in the UV color--color diagram. 2) The observed slope of the UV color-color relation of quiescent ETGs is too steep to be reproduced by the canonical stellar population models in which the UV flux is mainly controlled by age or metallicity parameters. Moreover, 2~mag of color spreads both in $FUV-NUV$ and $FUV-r$ appear to be ubiquitous among any subsets in distance or luminosity, which can hardly be explained by the canonical models with even large spreads in age ($>$ $\\pm$2~Gyrs) and mean metallicity ($>$ $\\pm$0.3~dex) assumed. This implies that the UVX in ETGs could be driven by yet another parameter which might be even more influential than age or metallicity.  The recent debate on the helium-enhanced populations has an interesting implication for this issue. Observations suggest that there are some helium-enhanced subpopulations in the extended HB globular clusters \\citep{dan04,nor04,lee05,pio05,rec06,lee07} and in the massive globular clusters in M87 \\citep{soh06}. If present, such subpopulations would steepen the UV spectral slope systematically by increasing the mean temperature of helium burning stars, as demonstrated in Fig.~3. They would also enhance the UV color spreads among the ETG stellar populations with even small variations in the number fractions or in the amount of helium enhancement \\citep[e.g.,][]{kav07c,chu11}. Whether the postulated helium enhancement in old stellar systems would occur also in galactic scales or in the intracluster medium of galaxy clusters \\citep{pen09} is in question. Observational tests by \\citet{lou11} and \\citet{yi11} do not support the helium sedimentation hypothesis, while \\citet{car11} argue that the helium abundance is a plausible candidate to explain the UV radial gradient. By the comparison between the UV color distribution and some toy models presented in this $Letter$, it is suggested that the observed steep slope of the UV color--color relation of ETGs would be better explained with the helium-enhanced subpopulations. The excessive spreads in the UV colors observed among the ETGs, besides the variations expected by the age/metallicity spreads, would be attributed to small variations in helium abundance or to the low-level star formation that is hard to be identified with the current survey depths and resolutions of the GALEX and SDSS data."
    },
    "1112/1112.4639_arXiv.txt": {
        "abstract": "\\noindent The most intense gamma-ray line observable from novae is likely to be from positron annihilation associated with the decay of $^{18}$F. The uncertainty in the destruction rate of this nucleus through the $^{18}$F(p,$\\alpha$)$^{15}$O reaction presents a limit to interpretation of any future observed gamma-ray flux. Direct measurements of the cross section of both this reaction and the $^{18}$F(p,p)$^{18}$F reaction have been performed between center of mass energies of 0.5 and 1.9~MeV. Simultaneous fits to both data sets with the R-Matrix formalism reveal several resonances, with the inferred parameters of populated states in $^{19}$Ne in general agreement with previous measurements. Of particular interest, extra strength has been observed above~$E_{CM}\\sim$1.3~MeV in the $^{18}$F(p,p)$^{18}$F reaction and between 1.3--1.7~MeV in the $^{18}$F(p,$\\alpha$)$^{15}$O reaction. This is well described by a broad 1/2$^{+}$ state, consistent with both a recent theoretical prediction and an inelastic scattering measurement. The astrophysical implications of a broad sub-threshold partner to this state are discussed. ",
        "introduction": "Introduction}  \\noindent Novae are the most common astrophysical explosion sites in the Universe and provide an excellent opportunity for the study of the nucleosynthesis of radioisotopes through detection of their gamma-ray emission. The current INTEGRAL satellite includes such an objective within its mission goals. A robust prediction of simulations of CO- and ONe-type novae events is that their gamma-ray emission will be dominated by the 511~keV gamma-rays produced by positron annihilation following the $\\beta^{+}$ decay of $^{18}$F~\\cite{Hernanz1999}. Key reasons for the high flux are that this isotope is produced relatively abundantly, and its lifetime of $\\sim$158~minutes is well matched to the timescale for nova ejecta to become transparent to gamma-ray emission. There remains however a large uncertainty on the absolute flux, and by implication the detectability distance for novae, because of uncertainty in the rates of the nuclear reactions producing and destroying $^{18}$F. Of these, the reaction which contributes the most uncertainty to the final abundance of $^{18}$F is that of $^{18}$F(p,$\\alpha$)$^{15}$O. It is this reaction which is the focus of the present work.  Despite significant recent effort expended on determination of this reaction rate, the situation remains unclear. In the energy region of interest, several states in $^{19}$Ne may make resonant contributions to the $^{18}$F(p,$\\alpha$)$^{15}$O reaction rate (see~\\cite{Nesaraja2007} for a recent summary). $S$- and $p$-wave resonances just above and below the $^{18}$F+p threshold at 6411~keV have proven difficult to measure, although some recent progress has been made~\\cite{Adekola2011}. The impact of interference between states of the same spin-parity is still unclear, although again, some recent progress has been made~\\cite{Beer2011}. It is also expected that there are additional resonances in the region of interest that are yet to be seen. One such prediction is for a broad $J^{\\pi}=1/2^{+}$ state, below threshold, but which has sufficient width to play an important role in the nova burning energy range~\\cite{Dufour2007}.  The Generator Coordinate Method model that made this prediction also suggests the existence of another broad 1/2$^{+}$ state ($\\Gamma_{p} =$ 157~keV, $\\Gamma_{\\alpha} =$ 139~keV) at a resonance energy around 1.49~MeV (E$_{x}$($^{19}$Ne) $=$ 7894~keV). Data from a recent $^{1}$H($^{19}$Ne,p)$^{19}$Ne$^{*}$(p)$^{18}$F inelastic scattering measurement revealed a broad feature beneath several well defined peaks that is consistent with the presence of such a state~\\cite{Dalouzy2009}. However, a direct measurement using the thick target method~\\cite{Murphy2009} revealed no such feature, though a narrow state of different spin was observed in the vicinity. This measurement only took data up to $E_{CM}\\sim1.6~MeV$, hence the assignment of parameters to this state is hampered by its vicinity to this high energy cut off. ",
        "conclusions": ""
    },
    "1112/1112.6316_arXiv.txt": {
        "abstract": "Stochastic acceleration of charged particles due to their interactions with plasma waves may be responsible for producing superthermal particles in a variety of astrophysical systems. This process can be described as a diffusion process in the energy space with the Fokker-Planck equation. In this paper, a time-dependent numerical code is used to solve the reduced Fokker-Planck equation involving only time and energy variables with general forms of the diffusion coefficients. We also propose a self-similar model for particle acceleration in Sedov explosions and use the TeV SNR RX J1713.7-3946 as an example to demonstrate the model characteristics. Markov Chain Monte Carlo method is utilized to constrain model parameters with observations. ",
        "introduction": "\\label{intro}  Observations of high-energy emission from many astrophysical sources imply the presence of significant populations of relativistic particles. Relativistic cosmic rays (CRs) with energies lower than $\\sim10^{15}$eV are commonly believed to be accelerated at the shock of supernova remnants (SNRs) in our Galaxy \\cite{h05}. Evidence of particle acceleration by shocks of SNRs first comes from radio observations of the remnants, where the radio emission is produced by relativistic electrons via the synchrotron process. The discovery of synchrotron X-ray emission from SNR 1006 reveals acceleration of TeV electrons by these shocks \\cite{k95}. However, direct evidence of proton and ion acceleration up to the spectral ``knee'' by these shocks remains elusive.  Nonthermal distributions are naturally produced via the Fermi process. The first-order Fermi mechanism corresponds to a convection term in the energy space. The second-order process, as it was originally conceived by Fermi, involved the stochastic acceleration (SA) of particles scattering with randomly moving magnetized clouds \\cite{f49}. Later refinements of this idea replaced the magnetized clouds with magnetohydrodynamic (MHD) waves \\cite{m80}. The second-order stochastic Fermi process recently finds applications in a wide range of astrophysical environments (See Refs.~\\refcite{pb08}--\\refcite{b06} for more details).  One commonly assumed agent of SA is plasma wave turbulence, which is expected to be present in nonequilibrium conditions of highly magnetized plasmas. The wave-particle interactions in magnetized plasmas have influences on the dynamics of energetic charged particles and can play a significant role in the propagation and acceleration of these particles. Charged particles which spiral along magnetic field lines are accelerated through resonant interactions with plasma waves. This process can be described as a diffusion process in the energy space through the Fokker-Planck (F-P) equation with a diffusion coefficient whose magnitude and form depend on the power spectrum and other characteristics of the plasma turbulence.  The resultant equation in general is complicated and one has to resort to some simplifying approximations. If we are mainly interested in the spatially integrated results, one may obtain a reduced F-P equation with time and energy as the variables. There have been a few studies for analytic solutions of this reduced F-P equation in simple cases \\cite{b06}\\cdash\\cite{m11}. However, numerical methods need to be developed to solve more general forms of the reduced equation. Park \\& Petrosion (1996) \\cite{pp96} studied in detail the numerical methods for the reduced F-P equation. In this paper, we use their fully implicit  Chang-Cooper method, which is proposed by Chang \\& Cooper (1970) \\cite{cc70}. Our primary goal in this paper is to present a detailed derivation of the exact time-dependent numerical solution for the reduced F-P diffusion equation. We utilize the Markov Chain Monte Carlo (MCMC) \\cite{l11} method to constrain parameters of a model for multi-wavelength observations of TeV SNR RX J1713.7-3946. ",
        "conclusions": "\\label{dis}  The F-P equation in general does not have analytical solutions for the diffusion coefficient in the SA theory. To obtain the electron distribution from the F-P equation, one may resort to some simplifying approximations involving only the time and energy variables. In the paper, we develop a numerical code to solve this reduced F-P equation, which is applicable for arbitrary forms of the diffusion coefficient.  Recent observations of SNR RX J1713.7-3946 made with the {\\it Fermi} space telescope show that the spectrum of this source in the GeV band is very hard with a power-law photon index of $\\Gamma=1.5\\pm0.1$, which is difficult to accommodate in a hadronic scenario, but agrees well with the IC origin of $\\gamma$-rays in the leptonic scenario \\cite{a11}. On the other hand, by considering the potentially high inhomogeneity of the shocked interstellar medium (ISM), \\cite{i11} argue that a hadronic model may still explain the Fermi observation. The nature of the TeV emission from SNR is an issue still open to investigation. In the work, we show that the SA of electrons by turbulent plasma waves in the shock downstream might naturally explain these observations \\cite{li11}. Future observations with the {{\\it HXMT}} and {{\\it NuSTAR}} can test the model."
    },
    "1112/1112.1700_arXiv.txt": {
        "abstract": "We present a carefully selected sub-sample of {\\it Swift} Long Gamma-ray Bursts (GRBs), that is complete in redshift. The sample is constructed by considering only bursts with favorable observing conditions for ground-based follow-up searches, that are bright in the 15-150 keV {\\it Swift}/BAT band, i.e. with 1-s peak photon fluxes in excess to 2.6 ph s$^{-1}$ cm$^{-2}$. The sample is composed by 58 bursts, 52 of them with redshift for  a completeness level of $90$\\%, while another two have a redshift constraint, reaching  a completeness level of 95\\%. For only three bursts we have no constraint on the redshift. The high level of redshift completeness allows us for the first time to constrain the GRB luminosity function and its evolution with cosmic times in a unbiased way. We find that strong evolution in luminosity ($\\delta_l=2.3\\pm 0.6$) or in density ($\\delta_d=1.7\\pm 0.5$) is required in order to account for the observations. The derived redshift distribution in the two scenarios are consistent with each other, in spite of their different intrinsic redshift distribution. This calls for other indicators to distinguish among different evolution models. Complete samples are at the base of any population studies. In future works we will use this unique sample of {\\it Swift} bright GRBs to study the properties of the  population of long GRBs. ",
        "introduction": "Gamma-ray bursts are powerful flashes of high--energy photons occurring at an average rate of a few per day throughout the Universe. They are detected at all redshifts, from the local Universe up to the extreme high redshifts (Salvaterra et al. 2009; Tanvir et al. 2009; Cucchiara et al. 2011a). Our knowledge of the distribution of long GRBs through cosmic times is still hampered by the fact that most of the observed {\\it Swift} GRBs are without redshift. Indeed, the measure of the distance has been secured for only $\\sim 1/3$ of the cases. Given the low completeness level in redshift determination, the effect of possible observational biases could be important in shaping their redshift distribution (Fiore et al. 2007). This fact strongly limit the possibility of well grounded statistical studies of the rest-frame properties of long GRBs and their evolution with cosmic time. Therefore, it is of paramount importance to obtain an unbiased complete sample of GRBs, capable to fully represent this class of object.  To this end, we present in this paper a well selected sub-sample of the full {\\it Swift} database. We select bursts that have favorable observing conditions for redshift determination from ground  and that are bright in the 15-150 keV {\\it Swift}/BAT band. We find 58 bursts matching our selection criteria with a completeness level in redshift determination of $90$\\%. The completeness level increases up to $\\sim 95$\\% by considering the redshift constraints imposed by the detection of the afterglow or host galaxy in some optical filters. Therefore, our selection criteria allow us to construct a sizable sample of long bursts that is (almost) complete in redshift, providing the solid basis for the study of the long GRB population in an unbiased way. In particular, since our selection is based on the brightness in the {\\it Swift}/BAT band, our sample is not biased against the detection of dark bursts, thus providing a complete description of the whole long GRB population.  In the present paper, we will take advantage of the high completeness level of our sample to constrain the GRB luminosity function (LF) and its evolution with cosmic time. In the last few years, this problem has been faced by many different authors ((e.g. Porciani \\& Madau 2001, Firmani et al. 2004, Guetta et al. 2005, Natarajan et al. 2005, Daigne et al. 2006, Salvaterra \\& Chincarini 2007, Salvaterra et al. 2009b, Butler et al. 2010, Wanderman \\& Piran 2010, Campisi et al. 2010, Qin et al. 2010, Virgili et al. 2011, Robertson \\& Ellis 2012). There is a general agreement about the fact that GRBs must have experienced some sort of evolution through cosmic time, whereas the nature and the level of such evolution is still matter of debate. Most of the previous works relied on the assumption that bursts lacking of redshift measurements follow closely the redshift distribution of bursts with known $z$. In the past, we tried to overcome this assumption by deriving conservative lower limit on the level of evolution on the basis of the number of bursts detected at $z>2.5$ (Salvaterra \\& Chincarini 2007) and of bursts with peak luminosity $L>10^{53}$ erg s$^{-1}$ (Salvaterra et al. 2009b). For the first time, thanks to our well selected, complete sub-sample of {\\it Swift} GRBs, we can tackle this issue in a unbiased way. In future works we will use this sample to study the correlation between physical parameters of the bursts, the properties of the burst light curves and of the environment in which they explode.  This paper is organized as follow. In Sect.~2 we describe our selection criteria and present our sample. We present our models of the GRB LF and redshift distribution in Sect.~3, while the results are given in Sect.~4. We extrapolate our findings to the detection limit of {\\it Swift} in Sect.~5. Finally, in Sect.~6 we draw our conclusions.  \\begin{table*} \\begin{center} \\begin{tabular}{lccclccclccclcc} \\hline \\hline GRB & redshift & ref. & & GRB & redshift & ref & & GRB & redshift & ref & & GRB & redshift & ref \\\\ \\hline  050318 & 1.44 & 1 &  $|$ & 061007 & 1.26 & 2  & $|$  &080603B & 2.69 & 2   & $|$ & 090709A & $<3.5$ & 15\t     \\\\ 050401 & 2.90 & 2  & $|$ &  061021 & 0.35 & 2 &  $|$  &  080605 & 1.64 & 2 & $|$ & 090715B & 3.00 & 16 \\\\ 050416A & 0.65 & 3 & $|$ & 061121 & 1.31 & 2 &  $|$ &  080607 & 3.04 & 2   & $|$ & 090812 & 2.45 & 17 \\\\ 050525A & 0.61 & 4 & $|$ & 061222A & 2.09 & 3 & $|$ &   080613B & - & -   & $|$ & 090926B & 1.24 & 18 \\\\ 050802 & 1.71 & 2 &  $|$ & 070306 & 1.50 & 2 &  $|$  &  080721 & 2.59 & 2 & $|$ & 091018 & 0.97 & 19 \\\\ 050922C & 2.20 & 2 & $|$ & 070328 & $<4^{(c)}$ &  &  $|$  & 080804 & 2.20 & 2 & $|$ & 091020 & 1.71 & 20 \\\\ 060206 & 4.05 & 2 &  $|$ & 070521 & 1.35 & 3 &  $|$  &  080916A & 0.69 & 2 & $|$ & 091127 & 0.49 & 21 \\\\ 060210 & 3.91 & 2 &  $|$ & 071020 & 2.15 & 2 &  $|$ &   081007 & 0.53 & 8  & $|$ & 091208B & 1.06 & 22 \\\\ 060306 & 3.5$^{(a)}$  &  & $|$ & 071112C & 0.82 & 2 & $|$  &  081121 & 2.51 & 9 & $|$ & 100615A & - & - \\\\ 060614 & 0.13 & 2 &  $|$ & 071117 & 1.33 & 2 &  $|$  &  081203A & 2.10 & 10  & $|$ &   100621A & 0.54 & 23 \\\\ 060814 & 1.92$^{(b)}$ &  &  $|$ &080319B & 0.94 & 2 & $|$  & 081221 & 2.26$^{(a)}$ &  & $|$ &       100728B & 2.106 & 24 \\\\ 060904A & - & - &    $|$ & 080319C & 1.95 & 2 & $|$  & 081222 & 2.77 & 11  & $|$ &   110205A & 2.22 & 25 \\\\ 060908 & 1.88 & 2 &  $|$ & 080413B & 1.1 & 2 & $|$  &  090102 & 1.55 & 12  & $|$ &  110503A & 1.613 & 26,27 \\\\ 060912A & 0.94 & 5 & $|$  & 080430 & 0.77 & 6 & $|$  &  090201 & $<4$ & 13 & $|$ & & &  \\\\ 060927 & 5.47 & 2 &  $|$ & 080602 & $\\sim 1.4^{(d)}$ & 7 & $|$ &090424 & 0.54 & 14   & $|$ & & &  \\\\ \\hline \\hline \\end{tabular} \\end{center} \\caption{List of the bursts matching our selection criteria.  Redshifts or limits are provided in the following references: [1] Berger et al. 2005, [2] Fynbo et al. 2009a and references therein, [3] Perley et al. 2009, [4] Foley et al. 2005, [5] Levan et al. 2007, [6] Cucchiara \\& Fox 2008, [7] Rossi et al. 2012, [8] Berger et al. 2008, [9] Berger \\& Rauch 2008, [10] Kuin et al. 2009, [11] Cucchiara  et al. 2008, [12] de Ugarte Postigo  et al. 2009a, [13] D'Avanzo et al. 2009, [14] Chornock  et al. 2009, [15] Perley et al. in prep., [16] Wiersema  et al. 2009a, [17] de Ugarte Postigo  et al. 2009b, [18] Fynbo  et al. 2009b, [19] Chen et al. 2009, [20] Xu et al. 2009, [21] Cucchiara et al. 2009, [22] Wiersema et al. 2009b, [23] Milvang-Jensen et al. 2010, [24] Flores  et al. 2010, [25] Cucchiara  et al. 2011b, [26] de Ugarte Postigo  et al. 2011, [27] D'Avanzo et al. 2011. } \\begin{flushleft} \\begin{small}  (a) based on VLT/X-shooter spectra of the host galaxies obtained within the program 087.A-0451 (PI: H. Flores). The spectra were reduced using the X-shooter data reduction pipeline version 1.3.7 (see Goldoni et al. 2006). In the NIR arm spectrum of GRB\\,060306 we identified at the afterglow position the [OII] doublet emission at $z=3.5$, while in the NIR arm spectrum of GRB\\,081221 the [OIII] and H$\\alpha$ emission lines are present at $z=2.26$ at the afterglow position.  (b) for this GRB a redshift of $z=0.84$ was reported by Thoene et al, (2007). Images and spectra at the afterglow position have been taken also using the VLT/FORS within the program 177.A-0591 (PI: J. Hjorth). We downloaded these data from the ESO Archive and reduced them with standard procedures using the ESO-MIDAS package. At the afterglow position we could identify an object, showing a continuum signature in the spectra, that we therefore consider to be the host galaxy of GRB\\,060814. We can also identify the galaxy reported by Thoene et al. (2007), at $z=0.84$, but this object is offset from the afterglow position. VLT/X-shooter spectroscopy of the host galaxy has been performed within the program 084.A-0303 (PI: J. Fynbo). We reduced these spectra using the X-shooter data reduction pipeline version 1.2.0 (see Goldoni et al. 2006). Thanks to the identification in the NIR arm of the [OII] doublet, [OIII]\\,$\\lambda$5007 and H$\\alpha$ emission lines associated with the host galaxy, we can establish a redshift of $z=1.92$ for GRB\\,060814.  (c) based on the $R-$band host galaxy detection in ESO-VLT/FORS2 imaging data obtained  within the program 177.A-0591 (PI: J. Hjorth), taken from the ESO Archive.  (d) photometric redshift on the bases of the most probable host galaxy association in the XRT error circle (Rossi et al. 2012). \\end{small} \\end{flushleft} \\end{table*} ",
        "conclusions": "We select a sub-sample of {\\it Swift} long GRBs that is complete in redshift. The sample is composed by bursts with favourable observing conditions and with  1-s peak photon fluxes $P\\ge 2.6$ ph s$^{-1}$ cm$^{-2}$. It contains  58 bursts with a completeness level of $\\sim 90$\\% and provides the basis for statistical studies of the properties of long GRBs and their evolution with redshift in a unbiased way. GRBs can be used to study fundamental issues in astronomy and astrophysics, such as the star formation rate and the stellar and metal abundances evolution. They can be use as tracers of the galaxy evolution, of the ISM composition and to investigate the early universe. Complete and fully representative samples of GRBs are therefore unique tools to perform these investigations.  Here, we use the observed burst redshift distribution of our complete sample to probe and constrain the evolution of the long GRB population in redshift. We confirm that GRBs must have experienced some sort of evolution being more luminous or more numerous in the past than observed today. We found that in order to match the observed distribution, the typical burst luminosity should increase as $(1+z)^{2.3\\pm 0.6}$ or the GRB rate density as $(1+z)^{1.7\\pm 0.5}$ on the top of the known cosmic evolution of the SFR. This result does not depend on the assumed expression of the GRB luminosity function. We also explore models in which GRBs form preferentially in low-metallicity environments. We find that the metallicity threshold for GRB formation should be lower than 0.3 $\\Zsun$ in order to account for the observations assuming no evolution of the GRB LF. This value, while consistent with the expectations of collapsar models, seems to be at odd with the observed properties of $z<1$ GRB hosts  (Levesque et al. 2010; Mannucci et al. 2011; Campisi et al 2011a).  Extrapolating our results to $P=0.4$ ph s$^{-1}$ cm$^{-2}$, the predictions of evolution models are consistent with available observational constrain without the need of any adjustment of the model free parameters. We also note that the predicted redshift distributions at the sensitivity of {\\it Swift} for different evolution models are very similar, implying that it would be extremely difficult to distinguish among different kind of evolutions only on the basis of the observed redshift distribution.  On the bases of our models we predict that $3-5$\\% of the bursts detected by {\\it Swift} lie at $z>5$ consistently with the most recent observational estimate of $5.5\\pm 2.8$\\% (Greiner et al. 2011). This indicates that high-$z$ bursts can contribute only marginally to the observed fraction of dark bursts. Finally, we expect 30-50 bursts per year over the entire sky with luminosities in excess to $10^{53}$ erg s$^{-1}$ and 2-5 with luminosities exceeding $10^{54}$ erg s$^{-1}$. Extreme luminous GRBs, such as the current {\\it Swift} record holder GRB~080607 with $L_{iso}= 2.26\\times 10^{54}$ erg s$^{-1}$ (Perley et al. 2011), should be rare, exploding once every 0.5-1 years in the Universe. This corresponds to one detection every $\\sim 0.7-1.5$ years in the field of view of {\\it Fermi}/GBM (and one every 5-12 years for {\\it Swift})."
    },
    "1112/1112.0310_arXiv.txt": {
        "abstract": "The differential equations governing the late-time ring-down of the perturbations of the Kerr metric, the Teukolsky Angular Equation and the Teukolsky Radial Equation, can be solved analytically in terms of confluent Heun functions.  In this article, for the first time, we use those exact solutions to obtain the electromagnetic (EM) quasinormal spectra of the Kerr black hole . This is done by imposing the appropriate boundary conditions on the solutions and solving numerically the so-obtained two-dimensional transcendental system.  The EM QNM spectra are compared with already published results, evaluated through the continued fraction method. The comparison shows that the modes with lower $n$ coincide for both methods, while those with higher $n$ may demonstrate significant deviations. To study those deviations, we employ the $\\epsilon$-method which introduces small variations in the argument of the complex radial variable. Using the $\\epsilon$-method one can move in the complex $r-$plane the branch cuts of the solutions of the radial equation and examine the dependence of the spectrum on their position.  For different values of $\\epsilon$, one can obtain both the frequencies evaluated through the well-established continued fraction method or somewhat different spectra calculated here for the first time. This result raises the question what spectrum should be compared with the observational data and why. This choice should come from better understanding of the physics of the problem and it may become particularly important  considering the recent interest in the spectra of the electromagnetic counterparts of events producing gravitational waves. ",
        "introduction": "During the long history of the study of the quasi-normal modes (QNMs) of a black hole (BH) (\\cite{RWE,ZRE,vish,Teukolsky0,Teukolsky_0, Teukolsky2, Teukolsky21, teukolsky_, chan0, QNM, chan1,det1,chan2,QNM2, QNM0, Leaver, Leaver0, Q_N_M, high, special1,extr1, Fiziev1, QNM1,extr2, special3, Fiziev2, Fiziev3, BHB1, QNM21, GW7}), the case of electromagnetic (EM) perturbations has been often ignored because of their perceived irrelevance to the problem of finding gravitational waves. The reasons behind this are that:  1. The expected luminosity of the gravitational output in the most often studied process of BH binary merger is much bigger than the electromagnetic one (\\cite{GW3}).  2. The EM waves strongly interact with the surrounding medium. This may lead to essential deviation of the observed spectrum from the one expected from the no-hair theorem and makes the object's fingerprint harder to detect.  3. Most importantly, because of this interaction, the electromagnetic perturbations are strongly absorbed by the interstellar medium, thus making the detection of the signal almost impossible at the predicted low frequencies for the electromagnetic QNMs.  On the other hand, the gravitational waves (GW) interact very weakly with matter and thus they can be detected at big distances, without getting absorbed or scattered, i.e. without obscuring the signature of the body that emitted them. It is, therefore, reasonable to expect that the GW should be much better suited for studying the central engine of astrophysical events, such as gamma-ray bursts (GRBs), while the EM waves should be seriously influenced by its environment.  For now, however, there are no gravitational waves detected. Although both LIGO and VIRGO detectors already work at design sensitivity, both detectors still fail to ``see`` gravitational waves (\\cite{LIGO1,LIGO2,LIGO3,LIGO4,LIGO6,LIGO7,LIGO8}). Particularly puzzling is the lack of GW detection from short GRBs (\t\\cite{LIGO5,LIGO9}) whose progenitors are expected to emit GWs in the range of sensitivity of the detectors.  The simplest explanation of those negative results may be a new mechanism of generation  of short GRBs which in  good approximation preserves the spherical symmetry of the process in the central engine and admits only significant dipole radiation. As a result, no significant gravitational waves will be generated during the short GRBs, since the gravitational waves have a quadrupole character. A similar situation is observed in the long GRBs. Such a new hypothesis for short GRBs is supported by the strong observational indications that both types of GRBs may have the same nature and differ only in their time scales (\\cite{GRB1,GRB2}). If so, we may expect that most of the energy release from GRBs is in the form of electromagnetic radiation.  A more traditional point of view is a physical process which yields both electromagnetic and gravitational radiation from GRBs. Actually, the ratio of the energy release in the form of electromagnetic waves and in the form of gravitational waves is still an open problem. Its solution strongly depends on the details of the hypothetical mechanism of GRBs which is still far from being well established. For example, the expected (but not yet observed) energy output in the case of GWs from a BH merger is $\\sim10^{53}erg$, which coincides (up to a factor due to collimation) with the {\\em observed} electromagnetic energy output of GRBs.   While hopes are laid on the Advanced LIGO and Advanced VIRGO which should start operating in the next years, this situation offers a good motivation for optimizing the GW search strategy and understanding better the physics of the GW sources. Particularly, this points to the advantages of studying the EM counterpart of the GW emission, which can help the localization of the source (improving on the big error box of the GW detector) and also it may give additional clues to the physics of the event (\\cite{GW1,GW3}). Numerical simulations already explore the detectability of the EM counterpart in different cases (for the case of supermassive black holes mergers, for example, see \\cite{GW4, GW5}, for neutron stars mergers \\cite{GW1,GW_}). The first results of LIGO and VIRGO searches using such a multimessenger approach were also published \\cite{LIGO10}. The idea behind those searches (for details, see \\cite{GW8,LIGO10}) is to use the GRBs as the EM counterparts of the GW, since the suspected GRB progenitors (collapsars for long GRBs and binary system mergers of neutron stars and/or black holes for the short GRBs) should emit GWs as well as EM and there is already a well working mechanism for observing the extremely EM luminous GRBs. Although the theoretical results from the multimessenger approach are still being analyzed, the intensive activity in this field shows that the EM counterpart of the GW emission can both facilitate and improve the information obtained from the GW observations.  The discrete spectrum of complex frequencies called QNMs describes only the linearized perturbations of the metric. Hence, they cannot describe completely the dynamics of the process during the early, highly intensive period of those events when the linearized theory is not applicable. On the other hand, it is known from full numerical simulations that it is the QNMs which dominate the late-time evolution of the object response to perturbations \\cite{high}. Thus, from an observational point of view, the QNM are important, since we may observe only the tails from the corresponding events, being far from them. This conclusion is supported by the recent numerical observation of two lowest gravitational modes of QNMs in the spectrum of the signal obtained from the full 3d general relativistic head-on collision of non-spinning BH  (see \\cite{headon}, and for further information \\cite{headon0}). This result is not isolated -- there are number of works in which the QNMs approximate well the signal of full 3d general relativistic simulations of mergers (for example \\cite{bin1,bin2,bin3,bin4,bin5} and also the pioneer works discussed in \\cite{special3}). This clearly implies that studying QNMs can bring new insights to the physics in the processes which include strong-field regime.  Those numerical results also point to another possible use of the QNMs in astrophysical observations. The QNMs correspond to particular boundary conditions characteristic for the object in question, and since in the case of BH the no-hair theorem states that they should depend only on the parameters of the metric (the mass $M$ and the rotation $a$ for the case of Kerr BH), measuring those frequencies observationally can be used to test the nature of the object -- a black hole or other compact massive objects like super-spinars (naked singularities), neutron stars, black hole mimickers  etc. \\cite{special31,NB1,NB2,NB3,NB4,spectra}. It also can constrain additionally the no-hair theorem which was recently put into question in the case of black holes formed as a result from the collapse of rotating neutron stars \\cite{GW6}. An interesting possibility is to find a way to use the damping times of the EM quasi-normal modes for comparison with observations. While the frequencies are subject to interaction with the surrounding matter which can significantly change the spectra, their damping times should be much less prone to deviation. A suggestion for such use can be found in simulations of jet propagation, which imply that the short-scale variability of the light-curve should be due to the central engine and not to the interaction of the jet with the surrounding medium (see \\cite{time-scale} and reference therein).  One more important application of the QNM spectrum can be found in the study of the central engines of the GRBs, whose extreme luminosity($\\sim10^{51}-10^{53}erg/s$) and peculiar time-variability cannot yet be fully explained in the frames of current models.  Even though numerical simulations proved to be capable of describing {\\em some} of the features of the GRBs light curves (for a recent review on GRBs, see \\cite{GRB}), the biggest stumbling stone seems to be the lack of proper understanding of the central engine of the GRBs.  Common ingredients of the existing GRB models include a compact massive object (black hole or a milli-second magnetar) and extreme magnetic field ($\\sim10^{15}G$ ) which accelerate and collimate the matter via different processes. Although those processes are still an open question for both theory and numerical simulations, the very central engine can be studied approximately by the linearized EM (and also GW if data is available) perturbations of the Kerr metric. When finding the electromagnetic QNMs, one does not care for the origin of the perturbation, but only for its spin and the parameters of the compact massive object. In the idealized EM case, the perturbation is described by free EM waves in vacuum. While the astrophysical black holes are thought to be not charged, they are immersed into EM waves with different energy and origin. The black hole response to such EM perturbations in linear approximation will be then the QNM spectrum defined by the appropriate boundary conditions \\footnote{Other   conditions more suitable for describing a primary jet were studied in \\cite{spectra}.}.  Studying the so obtained electromagnetic spectrum can give important insights into the key parameters of the physics occurring during  high-energy events as GRBs. In particular, the electromagnetic QNMs are subject to resonant amplification (the idea of the black hole bomb, \\cite{BHB1,GW7,teukolsky_}) and additionally, it is known from previous evaluations of the spectrum, that they exhibit very low damping in the limit $a\\to M$. For the moment, there are no observations of the rotations of the GRB progenitors, but the theoretical expectations are that they should be highly rotating in order to produce jets with such luminosity and collimation. Available observationally measured rotations of astrophysical compact massive objects show that there are many cases of near extremal values thus studying the extremal limit could be relevant to such objects. For example, recent evaluations of the spin parameter of astrophysical black holes give for the spin parameter $a=0.63,0.90$ and $a=0.89,0.99$ for $M=1,0.1M_{\\odot}$ for Sw J1644+57 and Sw J2058+05 respectively (most probable values, see \\cite{rot}), and also $a>0.98$ for GRS 1915+105 (\\cite{rot1}) and $a=0.989$ for MCG-6-30-15 (\\cite{rot2}).  Moreover, because of the relatively good coverage of the GRBs observations, there is a great amount of data, in a wide energy range (from optical to GeV energies) and from different epochs of the bursts which can be used to test the eventual applicability of the QNM spectrum in the late-time epoch of the burst. It may be hard to extract EM QNM spectra from the existing crude GRB spectra since the basic EM QNM frequencies are very low (from a small part of Hz --- for supermassive BH to several kHz -- for stellar mass BH) and the intensity of the higher EM QNM may be very low. To the best of our knowledge such attempts haven't been made. A new space mission, which will additionally help the EM observations in the radio range -- RadioAstron -- will offer unprecedented resolution (up to 1$\\mu$arcsec) in a  wide range of high frequencies (from 0.3GHz to 18-25GHz) accompanied by continuum, polarized and spectral imaging (for details see RadioAstron website \\footnote{\\url{http://www.asc.rssi.ru/radioastron}}). One may hope to use this new mission for a more detailed study of the spectra of EM radiation from the compact objects but its sensitivity is also far from the area of the basic EM QNM.  Theoretical calculations of the QNMs, however, are not simple. The linear perturbations of the rotating BHs are described by two second-order linear differential equations: the Teukolsky radial equation (TRE) and  the Teukolsky angular equation (TAE) on which specific boundary conditions should be imposed (\\cite{QNM,QNM0}). Until recently, solving those equations analytically was considered impossible in terms of known functions, so approximations with  simpler wave functions were used instead.  The resulting system of spectral equations -- a connected problem with two complex spectral parameters: the frequency $\\omega$ and the separation constant $E$ -- has been solved using different methods (\\cite{special3,special1,QNM1,QNM21}) with notably the most often used being the method of the continued fraction. This method was adapted by Leaver from the problem of the hydrogen molecule ion in quantum mechanics \\cite{Leaver,Leaver0}. While being successful in obtaining the QNMs spectra, Leaver's method has the disadvantage of not being directly connected with the physics of the problem, thus making it harder to further explore the spectra -- for example studying its dependence on the choice of the branch cuts of the exact solutions of the radial equation. In addition, one has some specific numerical problems in calculation of particular modes, for example, in calculation of the  $9^{th}$ one in the gravitational case \\cite{special3, special1,Q_N_M}.  The analytical solutions of the TRE and the TAE can be written in terms of the confluent Heun function (for $a\\neq M$) as done for the first time in \\cite{Fiziev1,Fiziev3, Fiziev4, Fiziev2}. Those functions are the unique local Frobenius solutions of the second-order linear ordinary differential equation of the Fuchsian type \\cite{heun3_,heun,heun1_,heun2_} with 2 regular singularities ($z=0,1$) and one irregular ($z=\\infty$) (for details see \\cite{Fiziev3}) and in the \\textsc{maple} notation, they are denoted as: $\\text{HeunC}(\\alpha,\\beta,\\gamma,\\delta,\\eta,z)$ (normalized to $\\text{HeunC}(\\alpha,\\beta,\\gamma,\\delta,\\eta,0)=1$). While the theory of the Heun functions is still far from being complete, they are implemented in the software package \\textsc{maple} and despite the problems in that numerical realization (see the discussion in \\cite{arxiv1}), the confluent Heun function was used successfully in our previous works \\cite{Fiziev1,spectra,arxiv1,arxiv3}. The advantage of using the analytical solutions is that one can impose the boundary conditions on them {\\em directly} (see \\cite{Fiziev1,spectra}), and thus be able to control all the details of the physics of the problem.  In a series of articles, we developed a method for solving numerically two-dimensional systems featuring the Heun functions (the two-dimensional generalization of the M\\\"uller method described in \\cite{arxiv}) and we used it successfully in the case of gravitational perturbation $s=-2$ of the Schwarzschild metric \\cite{arxiv1}. The so obtained frequencies repeat with high precision the results already published by other authors. Additionally, we used the epsilon-method (see below) to study the branch cuts of the solutions, which are particularly important in the case of the $9^{th}$ mode for $s=-2$. The latter, because of its very small real part, was often wrongly considered to represent the purely imaginary algebraically special mode. While the analysis of the potentials of the Regge-Wheeler equation (RWE) and the Zerilli equation (ZE) showed that there is a branch cut on the imaginary axis for this mode \\cite{special2,AS} which leads to its interesting properties, this result is directly obtainable from the actual solutions of the RWE and ZE in terms of the confluent Heun functions.  Furthermore, the stability of the solutions with respect to movement of the position of those branch cuts was studied in the whole interval of applicability of the method. Such a study is impossible with the continued fraction method, where the radial variable does not explicitly enter the equations being solved and which cannot be used for purely imaginary frequencies (\\cite{Leaver0}, p.8). If one looks at the equations used to derive this method in detail, it turns out that the angular equation \\cite{Leaver} in the continued fraction method coincides with the the three-term recurrence defining the confluent Heun function, solution of the TAE, in the neighborhood of the two regular singular points, $u=-1,1$, where $u=\\cos(\\theta)$ (\\cite{heun1_} Eq. (1.9-1.10). The radial equation in the continued fraction method, however, differs from the solution of the TRE in terms of the confluent Heun functions.  This is because in  Leaver's paper, the series from which the continued fraction are obtained are developed for the powers of $\\frac{r-r_+}{r-r_-}$ (due to switching the places of the singular points, see \\cite{Leaver0}, p.7), while the asymptotic three-term recurrence of the confluent Heun function at infinity is developed for $\\frac{1}{r-r_-}$. Note that in \\textsc{maple}, for $r>r_+$, the evaluation of the confluent Heun functions at infinity is obtained by numerical integration from the second singularity $r=r_+$.  In this article, we continue the exploration of the application of the confluent Heun functions by studying the QNMs of the Kerr BH. Our results show that using the confluent Heun function, one can obtain the QNMs for a wide range of modes and rotational parameters, and that there is very good agreement between our results and those obtained within other methods. Using the $\\epsilon$-method made it possible, for the first time to study the dependence of the so obtained frequencies on small deviations in the phase condition and it is shown how this nontrivial dependence evolves with $n$ and $a$. In this case, some of the modes are independent of $\\epsilon$ which should be expected since the frequencies should not depend on the radial variable. Other modes, however, depend critically on the value of $\\epsilon$ and they can differ seriously from the already published results. Additionally, details how the modes change in the interval of validity of the steepest descent method are presented. ",
        "conclusions": "From the recent developments in the field of gravitational waves detection it is clear that finding the EM counterpart to those events can prove to be very useful. In this case, it is needed to better understand the fundamental physics of quasi-normal ringing. In this paper, our team offered a new approach to finding the QNMs for the KBH based on directly solving the system obtained by the analytical solutions of the TRE and TAE in terms of the confluent Heun function. This approach has the advantage of being more traditional  (i.e. imposing directly the corresponding boundary conditions on the exact analytical solutions of the problem) and hence it should allow better understanding of the peculiar properties of the EM QNMs and the physics they imply.  It was shown that using this approach, one can reproduce the frequencies already obtained by other authors, but without relying on approximate methods. Particularly important is the ability to impose the boundary condition {\\em directly} on the solutions of the differential equations. We require the standard regularity condition on the TAE and we explore in detail the radial boundary condition (the BHBC). Critical in it is the use of the direction of steepest descent, which secures the purely outgoing wave at infinity. By using small deviations from this direction (and the phase-condition it defines), we were able to move around the branch cut in the solutions of the radial equation and thus to study its effect on the so obtained spectra. While this movement had no significant effect for the lower modes $n<3$, for the higher modes it led to significant deviations from the already published results. This behavior is persistent for the modes with $m=0,1,2$ and $l=1,2$. This observation raises the important question: What are the electromagnetic QNMs for which one has to look in astrophysical data.  Also interesting is that while the $\\epsilon$-method leads to significant changes of the  frequencies $\\omega_{m,n}$, it affects much less the separation parameter $E_{m,n}$ which here for the first time was obtained directly as a solution of the two-dimensional system without any prior approximations for it.  Another general result is that the confluent Heun function proved to be an effective tool for physical problems. Even though its \\textsc{maple} realization still has many flaws, its precision proved to be good enough to repeat  the already published results, and also studying the solutions, we were able to reveal new properties of the numerical stability of the EM QNMs with respect to changes in the phase-condition.  An interesting question is the results obtained using this method for $a>M$ or the so called naked singularity regime. Preliminary results show that the method is applicable in this case as well and the results will be published elsewhere."
    },
    "1112/1112.5473_arXiv.txt": {
        "abstract": "{The Large Observatory For X-ray Timing, LOFT, was selected by the European Space Agency as one of the four Cosmic Vision M3 candidate missions to compete for a launch opportunity at the start of the 2020s. Thanks to an innovative design and the development of large-area monolithic silicon drift detectors, the Large Area Detector (LAD) on board LOFT will operate in the 2-30 keV range (up to 50 keV in expanded mode), and achieve an effective area of $\\sim$10~m$^2$ at 8~keV, a time resolution of $\\sim$10~$\\mu$s, and a spectral resolution of $\\sim$260 eV (FWHM at 6~keV). These characteristics make LOFT a perfectly suited instrument to perform high-time-resolution X-ray observations of collapsed objects in our galaxy and brightest supermassive black holes in active galactic nuclei. LOFT will yield unprecedented information on strongly curved spacetimes and matter under extreme conditions of pressure and magnetic field strength, thus addressing two of the fundamental questions of the Cosmic Vision Theme ``Matter under extreme conditions'': does matter orbiting close to the event horizon follow the predictions of general relativity? What is the equation of state of matter in neutron stars? }  \\FullConference{The Extreme and Variable High Energy Sky - extremesky2011,\\\\ September 19-23, 2011\\\\ Chia Laguna (Cagliari), Italy}   \\begin{document} ",
        "introduction": "Collapsed objects, like neutron stars (NS) and black holes (BH), are characterized by exceptionally intense gravitational and magnetic fields. They provide a unique opportunity to reveal a variety of general relativistic effects and probe the conditions of matter in presence of supercritical magnetic fields and bulk densities exceeding those of atomic nuclei \\citep{psaltis08}. Carrying out these investigations requires the exploitation of diagnostics that are capable of probing {\\it in situ} matter motion in close vicinity of collapsed objects and/or measuring the rotation of the collapsed object itself. In this respect, both techniques based on X-ray spectroscopy and fast X-ray timing measurements are crucial. The former permit to reveal very broad profiles of X-ray emission lines (Fe K-shell in particular) from accretion disks extending down to the radius of the NSs and BHs marginally stable orbit. In these regions, velocities are comparable to the speed of light and well known relativistic effects (such as gravitational redshift, light-bending, frame-dragging, and Doppler shifts) distort the iron line shape into a characteristic broad shape. Measurements of the line profile then gives a detailed overview of the motion of matter very close to the central object. Among the available timing features, fast quasi periodic oscillations (QPOs) from BHs and NSs and coherent pulsations from the rotation of NSs and white dwarfs (WDs) are widely used to probe the intimate physical structure of these objects. QPOs in X-ray binaries and very massive BH in Active Galactic Nuclei (AGN), appear at the characteristic dynamical time-scales of the innermost disk regions, and are thus associated to the matter motion in strong field gravity. In the case of NSs, the detection and accurate measurement of coherent pulsations has been demonstrated effective in constraining its mass and radius, thus providing information on the NS equation of state \\citep{stella09,klis06,strohmayer06}.  Many decisive progresses in all the above fields of research were made possible in particular through the exploitation of the time resolution and effective area of the proportional counters array (PCA, $\\sim$0.67~m$^2$) on-board RXTE \\citep{jahoda}. LOFT, the Large Observatory For X-ray Timing, will constitute the next giant leap forward with respect to any existing X-ray timing devoted experiment, and it is being proposed as the natural follow-up to the more than 15 years of operations of RXTE. The instruments on-board LOFT will be endowed with an unprecedentedly large effective area ($\\sim$10~m$^2$ at 8~keV) and a spectral resolution of $\\sim$260~eV (FWHM at 6~keV). This will revolutionise the research field of compact objects and yield a great opportunity to address many of the fundamental questions on the matter under extreme conditions that were left open by the previous generation of X-ray instruments. \\begin{figure} \\centering \\includegraphics[scale=0.40]{sketch.ps} \\includegraphics[scale=0.32]{area.ps} \\caption{\\scriptsize {\\it Left panel}: Sketch of the LOFT satellite. The 6 deployable panels of the LAD, the WFM and the solar array are indicated. {\\it Right panel}: effective area of the LOFT/LAD as a function of energy.} \\label{fig:area} \\end{figure} ",
        "conclusions": "\\label{sec:conclusions} \\begin{figure} \\centering \\includegraphics[scale=0.22]{ns.ps} \\includegraphics[scale=0.30]{agn.ps} \\caption{\\scriptsize {\\it Left panel}: Mass-radius curves of representative EOSs for nucleonic (blue), nucleonic plus exotic (pink), and strange quark (green) matter. Horizontal bars represent mass determinations of binary radio pulsars. Regions to the upper left are disallowed by General Relativity and causality. The lower right region is excluded by the highest-frequency pulsar known (716 Hz). The pulse profile rising phase of a Type I X-ray burst was simulated by assuming the LAD capabilities in a way similar to \\cite{strohmayer04}. The derived 90\\% confidence level limits on the NS M and R from observations with LOFT are shown by the green ellipse. The uncertainty is $<$5\\% on both M and R. The M-R constraint derived from fitting simulated pulse profiles of SAX\\,J1808.4-3658, the first millisecond X-ray pulsar discovered (401 Hz), are shown by a gray ellipse. The same technique applied to simulations of the signal of the first eclipsing millisecond X-ray pulsars, Swift\\,J1749.4-2807 (518 Hz), gives in a much smaller uncertainty region (red); this is mainly due to the higher system inclination. {\\it Right panel}: Simulated LOFT/LAD spectra of a 3 mCrab (2-10 keV) AGN. The lower plot shows the broad relativistic Fe line produced in the innermost region of the accretion disc extending down to the marginally stable for an almost-maximally rotating  BH ($a$=0.99) with mass of 3.6$\\times$10$^6$~M${\\odot}$, viewed at an inclination of 45$^{\\circ}$. The plot shows the variable Fe K feature produced by an orbiting spot at $r$=10~$GM$/$c^2$ on the disk surface that is illuminated by a flare. The orbital period is 4~ks and the total exposure is 16~ks. LOFT/LAD will track the line variation on 1~ks time scale, thus allowing a determination of the orbital period through measurements at four different orbital phases over four cycles. The inset shows two 10~ks spectra from simulated high (6~mCrab in 2-10 keV) and low (2~mCrab) flux states LOFT/LAD observations of MCG-6-30-15. Thanks to the wide energy band and the large effective area above 10~keV, the variation of the reflection hump can be measured with 2\\% accuracy.} \\label{fig:science} \\end{figure}  LOFT was primarily designed to address the fundamental questions raised by the Cosmic Vision under the ``matter under extreme conditions'' theme, through the study of the spectral and fast flux variability of compact X-ray sources. The high throughput required by this type of investigation is achieved by an unprecedented large collecting area ($\\sim$10~m$^2$ at 8~keV), coupled with a spectral resolution of $\\sim$260~eV (FWHM at 6~keV) . A simulation of the LOFT capabilities is shown in Fig.~\\ref{fig:science} for two of the main science drivers of the mission: the measurement of the NS mass and radius in Galactic X-ray binaries and the detection of broad iron lines in bright AGNs. The science investigations that will be accessible with LOFT are far too broad to be efficiently summarized here, and we remind the reader to the documentation available on the official LOFT website\\footnote{http://www.isdc.unige.ch/loft}.  Following the selection of LOFT as one of the four candidate M3 missions for the Cosmic Vision programme in 2011, a further down-selection is expected in 2013. This will bring only one of the candidate missions to a launch in the 2020s."
    },
    "1112/1112.2253_arXiv.txt": {
        "abstract": "Recently,  Fermi-LAT detected GeV emission during the X-ray flaring activity in GRB 100728A. We study various scenarios for its origin. The hard spectrum of the GeV emission favors  the external inverse-Compton origin  in which X-ray flare photons are up-scattered by  relativistic electrons in the external forward shock. This external inverse-Compton scenario, with anisotropic scattering effect taken into account, can reproduce the temporal and spectral properties of the GeV emission in GRB 100728A. ",
        "introduction": "One of the key discoveries by the {\\it Swift} satellite is the presence of X-ray flares during the early afterglow phase in a large fraction of gamma-ray bursts (GRBs) (see, e.g., Burrows et al. 2005; Zhang et al. 2006; Nousek et al. 2006). The rapid rise and decay behavior of X-ray flares suggests that they are caused by internal dissipation of energy due to late central-engine activity. It was predicted that when the inner flare photons pass through the forward shocks, they will be up-scattered by forward shock electrons, producing a GeV flare (Wang et al. 2006; Galli \\& Piro 2007; Fan et al. 2008). The angular dispersion effect at the forward shock front will wash out any shorter temporal structure (Beloborodov 2005), so the high-energy inverse-Compton (IC) emission has a much smoother temporal structure determined by the forward shock dynamical time. The incoming X-ray flare photons are anisotropic as seen by the isotropically distributed electrons in the forward shock, so the scatterings between the flare photons and electrons in the forward shock are anisotropic. This effect may decrease the IC emission in the $1/\\Gamma$ cone along the direction of the photon beam, where $\\Gamma$ is the bulk Lorentz factor of the forward shock, but enhance the emission at larger angles, which leads to a delayed arriving time of the IC photons relative to the flare photons (Wang et al. 2006; Fan et al. 2008). On the other hand, the self inverse-Compton emission of X-ray flares may give rise to a high-energy component which should have similar light curves as the flares (Wang et al. 2006; Fan et al. 2008; Yu \\& Dai 2009). During the early afterglows, the afterglow synchrotron emission can also produce an extended high-energy component, which is thought to be responsible for the decaying long-lived high-energy emission detected from several GRBs by {\\em Fermi}/LAT (e.g. Kumar \\& Barniol Duran 2009, 2010; Ghisellini et al. 2010; Wang et al. 2010; De Pasquale et al. 2010; Corsi et al. 2010; He et al. 2011, Liu \\& Wang 2011).   GRB 100728A is the second case (after GRB 090510) with simultaneous detections by {\\em Swift} and {\\em Fermi}/LAT (Abdo et al. 2011). In this paper, we first analyze the {\\em Swift} XRT data and the {\\em Fermi}/LAT data of GRB 100728A, and obtain the X-ray and high-energy ($>$100 MeV) light curves (\\S 2). The early (167-854 seconds) X-ray afterglow exhibits intense and long-lasting flaring activity, with a total of 8 flares (Abdo et al. 2011). Two X-ray flares are also seen at the late stage (80-167 seconds) of the prompt phase (see Fig.1), indicating that the X-ray activity continues from the prompt phase to the afterglow phase. Interestingly {\\em Fermi}/LAT detected high-energy gamma-ray emission during the X-ray flaring activities (Abdo et al. 2011). These simultaneous X-ray and GeV observations offer a good case to study the nature of high energy emission. In \\S3, we confront the observations with the theoretical models and find that the external inverse-Compton (EIC) scattering of X-ray flare photons by electrons in the forward shock provides the best explanation for the GeV emission. We calculate the light curve of the EIC emission by taking into account the anisotropic scattering effect, and compare it with that of the observed GeV emission by {\\em Fermi}/LAT in \\S4. Finally, we give our conclusions in \\S5. ",
        "conclusions": "For the first time, GeV emission from a GRB during the X-ray flaring activity has been detected by the {\\em Fermi}/LAT, in GRB100728A. The temporal coincidence between the GeV emission and the X-ray flares suggests that the GeV emission should be related to the flares in some way. Here we have shown that an EIC scenario, where X-ray flare photons are up-scattered by electrons  in the external forward shocks, provides the best explanation for the GeV emission, supporting the earlier prediction of GeV emission from GRB X-ray flares (Wang et al. 2006). The hard spectrum of the GeV emission can be readily explained by the EIC emission. The delayed behavior of the GeV emission relative to the X-ray flares  also naturally arises in the EIC scenario. Synergistic observations between {\\em Swift} and {\\em Fermi} should be able in the future to find other similar cases, allowing further tests of this EIC scenario.  We are grateful to Bing Zhang and Zigao Dai for valuable discussions. This work is supported by the NSFC under grants 10973008 and 11033002, the 973 program under grant 2009CB824800, the program of NCET, the foundation for the Authors of National Excellent Doctoral Dissertations of China, Jiangsu Province Innovation for PhD candidate CXZZ11$\\_$0031, the Fok Ying Tung Education Foundation, NASA NNX08AL40G and NSF PHY 0757155. It is also supported in part by NASA via the Smithsonian Institution grant SAO SV4-74018,A035 [BBZ] at PSU."
    },
    "1112/1112.4475_arXiv.txt": {
        "abstract": "{One formation channel discussed for ultra-compact dwarf galaxies (UCDs) is that of massive star clusters, and the other main scenario is that of tidally transformed dwarf galaxies.} {We aim at quantifying the specific frequency of UCDs in a range of environments and at relating this to the frequency of star clusters and potential progenitor dwarf galaxies. Are the frequencies of UCDs consistent with being the bright tail of the globular cluster luminosity function (GCLF)?} {We propose a definition for the specific frequency of UCDs, $S_{N,UCD} = N_{UCD} 10^{0.4 (M_{V,host}-M_{V,0})} c_{w}$. The parameter $M_{V,0}$ is the zeropoint of the definition, chosen such that the specific frequency of UCDs is the same as those of globular clusters, $S_{N,GC}$, if UCDs follow a simple extrapolation of the GCLF. Considering UCDs as compact stellar systems with $M_V<-10.25$ mag (mass above $\\sim 2 \\times 10^6$ M$_{\\odot}$), it is $M_{V,0}=-20$ mag. The parameter $c_{w}$ is a correction term to take the dependence of the GCLF width $\\sigma$ on the host galaxy luminosity into account. We apply our definition of $S_{N,UCD}$ to results of spectroscopic UCD searches in the Fornax, Hydra and Centaurus galaxy clusters, two Hickson Compact Groups, and the Local Group. This includes a large database of 180 confirmed UCDs in Fornax. } {We find that the specific frequencies derived for UCDs match those of GCs very well, to within 10-50\\%. The ratio $\\frac{S_{N,UCD}}{S_{N,GC}}$ is $1.00 \\pm 0.44$ for the four environments Fornax, Hydra, Centaurus, and Local Group, which have $S_{N,GC}$ values. This good match also holds for individual giant galaxies in Fornax and in the Fornax intracluster-space. The error ranges of the derived UCD specific frequencies in the various environments then imply that not more than $\\sim$50\\% of UCDs were formed from dwarf galaxies. We show that such a scenario would require $\\gtrsim$90\\% of primordial dwarfs in galaxy cluster centers ($<$100 kpc) to have been stripped of their stars.} {We conclude that the number counts of UCDs are fully consistent with them being the bright tail of the GC population. From a statistical point of view there is no need to invoke an additional formation channel.}  \\titlerunning{Specific frequencies of UCDs} ",
        "introduction": "Ultra-compact dwarf galaxies (UCDs) were recognized as a populous and potentially distinct class of objects about a decade ago (Drinkwater et al.~\\cite{Drinkw03}), following the results of various spectroscopic surveys in the Fornax cluster (Minniti et al.~\\cite{Minnit98}, Hilker et al.~\\cite{Hilker99}, Drinkwater et al.~\\cite{Drinkw00}, Phillipps et al.~\\cite{Philli01}). UCDs are generally considered as compact stellar systems with masses above $\\simeq$ 2$ \\times 10^6$ M$_{\\odot}$ and sizes below $\\sim$ 100 pc (e.g. Ha\\c{s}egan et al.~\\cite{Hasega05}, Mieske et al.~\\cite{Mieske08}). Since their discovery in Fornax, more UCDs have been detected in a range of environments, from loose and compact galaxy groups (e.g. Evstigneeva et al.~\\cite{Evstig07}, da Rocha et al.~\\cite{Daroch11}) to dense galaxy clusters like Virgo, Centaurus, Hydra, Coma and Fornax itself (e.g. Jones et al.~\\cite{Jones06}, Ha\\c{s}egan et al.~\\cite{Hasega05}, Misgeld et al.~\\cite{Misgel08} \\& ~\\cite{Misgel11}, Mieske et al.~\\cite{Mieske04} \\& ~\\cite{Mieske07}, Gregg et al.~\\cite{Gregg09}, Chiboucas et al.~\\cite{Chibou10a} \\& ~\\cite{Chibou10b}).  The luminosity and size distribution of UCDs shows a smooth transition to the regime of globular clusters (GCs) (e.g. Ha\\c{s}egan et al.~\\cite{Hasega05}, Mieske et al.~\\cite{Mieske04} \\& ~\\cite{Mieske06}). It is thus not surprising that one of the two discussed formation channels of UCDs is that of massive star clusters created in the same or a similar way to the main star cluster population of early-type galaxies (e.g. Fellhauer \\& Kroupa ~\\cite{Fellha02} \\& ~\\cite{Fellha05}, Murray ~\\cite{Murray09}, Gieles et al.~\\cite{Gieles10}).  The other proposed formation channel is that of tidally stripped dwarf galaxies (e.g. Zinnecker et al.~\\cite{Zinnec88}; Drinkwater et al. ~\\cite{Drinkw03}, Bekki et al.~\\cite{Bekki03}, Goerdt et al.~\\cite{Goerdt08}), as has also been frequently suggested for the Local Group object $\\omega$ Cen (e.g. Majewski et al.~\\cite{Majews00}, Carraro \\& Lia~\\cite{Carrar00}, Hilker \\& Richtler~\\cite{Hilker00}, Bekki \\& Freeman~\\cite{Bekki03b}, Noyola et al.~\\cite{Noyola08}, da Costa \\& Coleman~\\cite{daCost08}). Using the mass limit above for the definition of a UCD, indeed $\\omega$ Cen is the only UCD associated with the Milky Way (see e.g. the Milky Way GC catalog of Harris et al.~\\cite{Harris96}).  In this paper we test whether the numbers of UCDs are consistent with being merely the bright ($\\sim 1\\%$) tail of the globular cluster luminosity function (GCLF). A clear excess of UCDs above the canonical GCLF would suggest that a separate process, distinct from the GC formation process, is responsible for the overabundance of UCDs. To this end we define the specific frequency of UCDs, by relating the number of UCDs to the luminosity $M_{V,host}$ of their host galaxy. The zero-point of the definition is chosen such that the UCD specific frequency is equal to that of GCs for the case that the numbers of UCDs are equal to a simple extrapolation of the GCLF. We calculate this specific frequency in a range of environments, and compare it to that of star clusters and potential progenitor dwarf galaxies. Special attention is given to the Fornax galaxy cluster, where a large database of confirmed UCDs and GCs exists. ",
        "conclusions": "\\label{conclusions}  In this paper we have proposed a definition for the specific frequency of UCDs, which we denote as $S_{N,UCD}$. We adopt the following functional form:\\\\  $S_{N,UCD} = N_{UCD}   10^{0.4    (M_{V,host}-M_{V,0})}   c_{w}$.\\\\  This definition normalizes the number of UCDs in a given environment to a unit host galaxy luminosity $M_{V,host}$, analogous to the definition of the specific frequency $S_{N,GC}$ for GCs. The premise of our definition is that if UCDs follow the extrapolation of the GCLF to bright magnitudes, then $S_{N,UCD} = S_{N,GC}$. This premise defines the value of the zero point $M_{V,0}$. Considering UCDs as compact stellar systems with $M_V<-10.25$ mag, we find that the value of $M_{V,0}=-20$ mag fulfills the above premise. For the case of extragalactic surveys with a brighter completeness limit $M_V<-11.0$ mag, we need to adopt $M_{V,0}=-22$ mag. The term $c_{w}$ is introduced to correct the specific frequency for the well known systematic variation of the GCLF width $\\sigma$ with host galaxy magnitude. Details of this correction are outlined in Sect.~\\ref{sndefsec}. In Sect.~\\ref{envsn} we apply our proposed definition of $S_{N,UCD}$ to results of spectroscopic UCD searches performed by our group in the Fornax, Hydra and Centaurus galaxy clusters, and two Hickson Compact Groups. We also include the Local Group.  Our main finding is that the specific frequencies derived for UCDs match those of GCs very well. For four of the six investigated environments, there are GC specific frequency measurements available, allowing a direct comparison. For those four environments we find a mean $S_{N,UCD}=6.7 \\pm$ 2.0, vs. $S_{N,GC}=6.7 \\pm 2.2$. The ratio of UCD-to-GC specific frequency is therefore $\\frac{6.7 \\pm 2.0}{6.7 \\pm 2.2}=1.00 \\pm 0.44$. The mean $S_{N,UCD}$ of all six investigated environments is 5.3 $\\pm$ 1.7, in agreement with the average GC specific frequency of 5.0 $\\pm$ 0.7 of the available literature data for the corresponding host magnitude range $M_V<-21$ mag. Our findings are consistent with the hypothesis that most UCDs are formed by the same process as the overall GC population, and with a similar formation efficiency to GCs.  We also present an extension of our analysis in the Fornax cluster by using the large available data set of spectroscopically confirmed UCDs and GCs. This literature data set comprises about 180 confirmed UCDs with $M_V<-10.25$ mag. We find that the specific frequencies of UCDs around giant galaxies and in the intracluster space are consistent with their being drawn from the bright tail of the GCLF. There is still room for a possible UCD overabundance in the intracluster space by a factor of $\\sim 2$, given the incompleteness in spectropscopic coverage in those regions. We do not find significant evidence of a different spatial distribution between UCDs and GCs.  It has been proposed that the present-day population of UCDs is indeed {\\it a superposition} of sources formed via tidal stripping of dwarf galaxies and sources formed in the same process as the main GC population (e.g. Ha\\c{s}egan et al.~\\cite{Hasega05}, Mieske et al.~\\cite{Mieske06}, Chilingarian et al.~\\cite{Chilin11}, Norris et al.~\\cite{Norris11}, da Rocha et al.~\\cite{Daroch11}). The error bars of the specific frequencies derived for UCDs suggest $\\sim$50 \\% as an upper limit for the importance of the dwarf galaxy channel. We show in Sect.~\\ref{progen} that this would require at least $\\gtrsim$ 90\\% of primordial dwarf galaxies in the central $\\sim$50-70 kpc of the considered galaxy clusters to have already been disrupted. If indeed a significant amount of UCDs were formed from tidal processes, then the tidal stripping of stars from low-mass dark matter halos in the central Fornax cluster has been extremely efficient.\\\\  \\noindent We conclude that the number counts of UCDs are fully consistent with them being the bright tail of the GC population. From a statistical point of view there is no need to invoke an additional formation channel. The statistical error bars constrain the fraction of tidally stripped dwarfs to not more than 50\\% of UCDs."
    },
    "1112/1112.1066_arXiv.txt": {
        "abstract": "Continuing our study of the role and evolution of accretion disc winds in the microquasar GRS 1915+105, we present high-resolution spectral variability analysis of the $\\beta$ and $\\gamma$ states with the \\textit{Chandra} High Energy Transmission Grating Spectrometer. By tracking changes in the absorption lines from the accretion disc wind, we find new evidence that radiation links the inner and outer accretion discs on a range of time-scales. As the central X-ray flux rises during the high-luminosity $\\gamma$ state, we observe the progressive over-ionization of the wind. In the $\\beta$ state, we argue that changes in the inner disc leading to the ejection of a transient `baby jet' also quench the highly-ionized wind from the outer disc. Our analysis reveals how the state, structure, and X-ray luminosity of the inner accretion disc all conspire to drive the formation and variability of highly-ionized accretion disc winds. ",
        "introduction": "\\label{sec:intro} In the last two decades, \\textit{ASCA, Chandra,~XMM-Newton,} and \\textit{Suzaku} have discovered a multitude of highly ionized absorbers in the X-ray spectra of black hole and neutron star X-ray binaries. In some cases, without clear evidence for blueshifts (e.g.\\ \\citealt{Ebisawa97a,Kotani97a,K00,Kotani2000a,L02,M04a,Xiang09}), these absorbers have been attributed to hot or extended disc atmospheres. Many observations, however, provide unambiguous evidence for outflows from the accretion disc (e.g.\\ \\citealt{BS00,Schulz02,U04,M06a,M06b,M08}; \\citealt*{U09}; \\citealt{NL09,Blum10,ReynoldsM10,Miller11}; \\citealt*{N11a}, hereafter Paper II). These results suggest that accretion disc winds may be as common as relativistic jets in accreting systems. \\defcitealias{NL09}{Paper I} \\defcitealias{N11a}{Paper II} \\defcitealias{B00}{B00}  Of particular interest, then, is the variability of these disc winds and their role in the evolution of the accretion flow. Ionized absorbers in X-ray binaries are now known to vary on time-scales from 5 seconds (GRS 1915+105, \\citetalias{N11a}) to 300 s (H 1743-322, \\citealt{M06b}) to ks or more (Cir X-1, \\citealt{Schulz02}; GRS 1915+105, \\citealt{L02,U10}). Furthermore, there is mounting evidence that this variability cannot be due solely to changes in the ionizing luminosity (\\citealt{L02,Schulz02,Blum10}; \\citetalias{N11a}), which implies that high-spectral-resolution X-ray observations of accreting X-ray binaries may be able to probe either the physical processes that link outflows to the behavior of the inner accretion flow, or inhomogeneities in the structure and density of accretion disc winds.  In a study of 10 years of \\textit{Chandra} High-Energy Transmission Grating Spectrometer (HETGS; \\citealt{C05}) observations of GRS 1915+105, we demonstrated that the strength of absorption lines in the accretion disc wind is anticorrelated with the fractional hard X-ray flux (and the strength of the jet; \\citealt{NL09}, hereafter \\citetalias{NL09}; for a preliminary discussion of the absorption lines, see \\citealt{M08}). Based on this anticorrelation and estimates of the mass loss rates in the wind and the jet, we argued that accretion disc winds may be able to suppress or quench jets by directly draining their matter supply on long time-scales. The fractional hard X-ray flux is therefore a useful indicator of both the accretion state and the physics of outflows around black holes. \\begin{figure*} \\centerline{\\includegraphics[width=0.75\\textwidth]{f1.eps}} \\caption{The \\textit{Chandra} (black) and \\textit{RXTE} (red) lightcurves of the $\\beta$ (top, middle) and $\\gamma$ (bottom) states. The \\textit{RXTE} lightcurves have been renormalized to match the mean \\textit{HETGS} count rate, with excellent agreement between the two missions. In the top and middle panels, we use horizontal lines to indicate the locations of the X-ray dips analysed in Section \\ref{sec:betahetgs}.} \\label{fig:lc} \\end{figure*} Yet much more can be said for GRS 1915+105, whose X-ray lightcurve exhibits at least 14 different patterns of variability (\\citealt{B00}, hereafter B00; \\citealt{K02,Hannikainen05}). Many of these phenomenological patterns, which are highly structured and high-amplitude, are interpreted as limit cycles of accretion and ejection in an unstable disc \\citep{B97b,M98,T04,FB04}. In order to understand the physics linking disc winds, jets, and this fast X-ray variability, we presented the first detailed phase-resolved spectral analysis of a variability class in GRS 1915+105, using a joint \\textit{RXTE/Chandra} HETGS observation of the $\\rho$ state \\citepalias{N11a}. We showed for the first time that changes in the broadband X-ray spectrum produce changes in the \\textit{structure} and mass loss rate of the accretion disc wind on time-scales as short as 5 seconds. Furthermore, our analysis indicated that the wind in the $\\rho$ state may be sufficiently massive to excite long-term oscillations in the mass of the disc (\\citealt{Shields86,Luketic10}). By establishing clear causal links between the X-ray oscillations and the disc wind, \\citetalias{N11a} provides a physical mechanism that could be responsible for both the suppression of jets and (possibly) state transitions in black hole binaries.  In this paper, we continue our follow-up wind variability studies with spectral timing analysis of three additional archival \\textit{Chandra} HETGS observations of GRS 1915+105. Two of these observations (ObsIDs 6579 and 6580) fall into the $\\beta$ class, a wild $\\sim30$-min cycle that is associated with the production of impulsive/discrete (``baby'') radio jets (\\citealt{PF97,Fender97,M98,E98a,Mirabel99,B00,K02,FB04}). Each $\\beta$ cycle begins with a 10--15 minute dip in the lightcurve, which is  spectrally hard and exhibits a QPO whose frequency tracks the X-ray flux (e.g.\\ \\citealt{Markwardt99,Mikles06,Rodriguez08b}). The dip ends with a bright spike, followed by 10--15 minutes of soft, bright X-ray flares. These flares are coincident with strong optically-thin radio/infrared oscillations (baby jets), but there is an ongoing debate as to whether these discrete ejections are produced during the dip itself or at the time of the X-ray spike \\citep{M98,E98a,Migliari03,Rothstein05,Rodriguez08a}. \\setcounter{table}{0} \\begin{table} \\caption{\\textit{Chandra} HETGS observations of GRS 1915+105} \\label{tbl:obs} \\begin{tabular}{@{}ccccc} \\hline Obs.       & X-ray  &               & Elapsed    & T$_{\\rm exp}$ \\\\ ID         & State  & Date          & Time (ks)  & (ks)\\\\ \\hline 6579       & $\\beta$ & 2005 Dec.\\ 1 & 20.65 & 12.30 \\\\ 6580       & $\\beta$ & 2005 Dec.\\ 1 & 21.62 & 12.14 \\\\ 6581       & $\\gamma$ & 2005 Dec.\\ 3 & 20.65 & ~9.73 \\\\ \\hline \\end{tabular} \\end{table}  During the third observation (ObsID 6581), GRS 1915+105 exhibited irregular variability with very high X-ray flux (\\textit{RXTE} count rate $\\ga6000$ counts s$^{-1}$ PCU$^{-1}$) and extremely strong red noise in the power density spectrum ($P_{\\nu}\\sim\\nu^{-2}$). Although we identified this as the $\\gamma$ state in \\citetalias{NL09} and continue to refer to it thus here, this particular observation is rather unusual since it also displays significant similarities to typical $\\delta$ and $\\mu$ states. Relatively little has been said about the $\\mu,~\\delta,$ and $\\gamma$ states in the literature, as none have been associated with strong radio jets \\citep{K02}. However, in a study of the $\\mu$ state, \\citet*{Soleri08} reported the discovery of a transient low-frequency QPO, which they argue may be linked to state transitions or transient jets. This leaves open the possibility that the $\\mu/\\delta/\\gamma$ states exhibit physically-interesting outflow behavior, despite their weak radio flux. In this respect, these states could be similar to the $\\rho$ state, which has very low radio brightness \\citep{K02} but exhibits extreme disc wind variability and may produce short-lived jets \\citepalias{N11a}.  Here we address the variability of the known accretion disc wind in the \\textit{Chandra} HETGS observations of the $\\beta$ and $\\gamma$ states, whose lightcurves are shown in Figure \\ref{fig:lc}. By comparing changes in the ionized absorption to changes in the X-ray flux, we show how the wind is linked to the oscillations in the inner disc. The paper is organized as follows: in Section \\ref{sec:obs} we describe our observations and data reduction. In Section \\ref{sec:beta} we perform spectral analysis of the wind in the $\\beta$ state with the \\textit{Chandra} HETGS and \\textit{RXTE} PCA; in \\ref{sec:gamma} we use the HETGS to study the wind in the $\\gamma$ state. We discuss the implications of our results in Section \\ref{sec:discuss}.\\vspace{-7mm} ",
        "conclusions": "\\label{sec:conc} In this paper, we have presented a time- and flux-resolved high-resolution spectral analysis of GRS 1915+105 in the $\\beta$ and $\\gamma$ states. Although only the $\\beta$ state is associated with transient jet formation, both states exhibit strong erratic X-ray variability, which we have used to probe the relationship between the central X-ray luminosity and the properties of the accretion disc wind. In particular, our new results confirm our previous finding that radiation mediates rapid interactions between the inner accretion disc and the outer accretion disc. This interaction is most evident in the $\\beta$ state, where the wind is actually quenched during the hard X-ray dip, possibly because the irradiation of the outer disc is affected by changes in the structure of the inner disc leading to the formation of a transient jet. To the best of our knowledge, this is the first time that a wind has been observed to turn off within a single observation, and the mechanism ($\\dot{M}$ and/or structural changes in the inner disc) requires intricate links between radiation, winds, and jets.  This result is therefore an exciting step towards a complete understanding of wind-jet interactions on all time-scales. With our sensitivity and limited time resolution, our best estimates place the wind in the $\\beta$ state $\\sim10^{11}$ cm from the black hole, leaving it effectively unable to affect the transient jet on such short time-scales. However, if future multiwavelength studies unaffected by telemetry saturation are able to place better constraints on the location of the wind, and require it to be very close to the black hole, then the $\\beta$ state might produce the first evidence for direct wind-jet interactions on any time scale. Such a result would significantly enhance our understanding of accretion and ejection as profoundly inter-related processes around black holes.\\vspace{-7mm}"
    },
    "1112/1112.3125_arXiv.txt": {
        "abstract": "We report the results of quasi-simultaneous two filter optical monitoring of two high-energy peaked blazars, 1ES 1959+650 and 1ES 2344+514, to search for microvariability and short-term variability (STV). We carried out optical photometric monitoring of these sources in an alternating sequence of B and R pass-bands, and have 24 and 19 nights of new data for these two sources, respectively. No genuine microvariability (intra-night variability) was detected in either of these sources.  This non-detection of intra-night variations is in agreement with the conclusions of previous studies that high-energy peaked BL Lacs are intrinsically less variable than low-energy peaked BL Lacs in the optical bands. We also report the results of STV studies for these two sources between July 2009 and August 2010. Genuine STV is found for the source 1ES 1959+650 but not for 1ES 2344+514.  We briefly discuss possible reasons for the difference between the intra-night variability behaviour of high- and low-energy peaked blazars. ",
        "introduction": "Blazars constitute an enigmatic subclass of radio-loud active galactic nuclei (AGNs). Blazars include both BL Lacertae (BL Lac) objects and flat spectrum radio quasars (FSRQs), as both are characterized by strong and rapid flux variability across the entire electromagnetic spectrum. In addition, BL Lacs show largely featureless optical continua.   Blazars exhibit strong polarization from radio to optical wavelengths and usually have core-dominated radio structures. According to the orientation based unified model of radio-loud AGNs, blazar jets usually make an angle $\\leq 10^{\\circ}$ to our line-of-sight (e.g., Urry \\& Padovani 1995).   These extreme AGNs provide a natural laboratory to study the mechanisms of energy extraction from the central super-massive black holes and the physical properties of astrophysical jets. \\par The electromagnetic (EM) radiation from blazars is predominantly non-thermal. At lower frequencies (through the UV or X-ray bands) the mechanism is almost certainly synchrotron emission, while at higher frequencies the emission mechanism is probably due to Inverse-Compton (IC) emission (e.g., Sikora \\& Madejski 2001; Krawczynski 2004). The spectral energy distributions (SEDs) of blazars have a double-peaked structure (e.g., Giommi, Ansari \\& Micol 1995; Ghisellini et al.\\ 1997; Fossati et al.\\ 1998). Based on the location of the first peak of their SEDs, $\\nu_{\\rm peak}$, and blazars are often sub-classified into low energy peaked blazars (LBLs) and high energy peaked blazars (HBLs)  (Padovani \\& Giommi 1995). The first component peaks in the near-infrared/optical for LBLs and in the UV or X-rays for HBLs, while the second component usually peaks at GeV energies for LBLs and at TeV energies for HBLs.  Observations of blazars reveal that they are variable at all accessible timescales, from a few tens of minutes to years, and even decades, at many frequencies (e.g., Carini \\& Miller 1992; Wagner \\& Witzel 1995; Gupta et al.\\ 2004; Villata et al.\\ 2007). These different timescales allow the variability of blazars to be broadly divided into three classes, e.g., intra-day variability (IDV), short-term variability (STV), and long-term variability (LTV). Variations in the flux of a source of a couple of hundredths of a magnitude up to a few tenths of magnitude over a time scale of a day or less is known as IDV (Wagner \\& Witzel 1995) or microvariability or intra-night optical variability. Flux changes observed over days to a few months are often considered to be STV, while those taking from several months to many years are usually called LTV (e.g., Gupta et al.\\ 2004).  For a better understanding of how the properties of HBLs relate to the previously known LBLs, we need to discriminate among the theoretical models that try to explain them. Optical observations offer a wealth of information on the variability of blazars. Over the past many years, a large amount of  optical variability data for LBLs have been reported and rapid optical variability has been shown to be common (Rani et al. 2010, 2011 and references therein). However, because most of the HBLs were more recently found by X-ray and gamma-ray sky surveys (Perlman et al. 2006, Abdo et al. 2010 and refernces therein) they have been less extensively monitored, and the properties of the optical variability of HBLs are not yet clear. Preliminary observations of  HBLs  suggested that they show statistically lesser amounts of optical variability  (Heidt \\& Wagner 1996, 1998) and polarized light than that of LBLs (Stoke et al. 1985, 1989, Jannuzi et al. 1993, 1994, Villata et al. 2000). Heidt \\& Wagner (1998) noticed these objects likely display different duty cycles and variability amplitudes from those of  the LBLs. Romero, Cellone \\& Combi (1999) found such differences in the case of  microvariability, as just one out of three HBLs in their sample showed micro-variations, while eight LBLs out of 12 did so.  These differences were attributed to the possible presence of stronger magnetic fields in the high-energy peaked BL Lacs which could prevent the formation of small-scale jet inhomogeneities in these objects. However, due to the relative paucity of optical monitoring for HBLs, we can not yet definitively say whether and how HBLs are different from LBLs in optical variability and therefore any additional high quality observations of HBLs are particularly useful.  Here we report quasi-simultaneous observations of two HBLs, 1ES 1959$+$650 and 1ES 2344$+$514 in two optical bands (B and R) for the first time. In addition to looking for any microvariability in these two optical bands, our observations also give additional information on any microvariations in (B$-$R) colour. We also have searched for  short-term variability timescales for these HBLs in the optical bands. We used five telescopes located in India, Bulgaria and Greece to monitor the optical variability of 1ES 1959$+$650 and 1ES 2344$+$514  between July 2009 and November 2010.  The paper is structured as follows.  In Section 2, we discuss the key features of these two HBLs.  Section 3 describes the observations and data reduction procedure.  Section 4 reports our results and we present our discussion and conclusions in Section 5. ",
        "conclusions": "We have presented our results  of quasi-simultaneous and relatively dense temporal observations of two HBLs, 1ES 1959$+$650 and 1ES 2344$+$514, in two optical bands (B and R) during 2009--2010 . These are the first extensive observations of this type of these two blazars and thus adds to the rather small amount of  data on HBL microvariability and STV. We have observed the sources 1ES 1959$+$650 and 1ES 2344$+$514 on 24 and 19 nights respectively, but no significant intra-day variability was observed during any night for those two targets. In addition, we can report the results of short-term variability studies of these sources. The blazar 1ES 1959$+$650 has shown genuine short-term variability and it can be characterized by a preliminary fading trend of $\\sim$0.6 mag (from JD 2455036 to 2455161) and then a brightening trend of $\\sim$0.74 (from JD 2455326 to 2455506 in V band). However, despite nominally substantial magnitude changes, genuine short-term variability was not found for the source 1ES 2344+514. We found that most of the apparent variations by the blazar could have arisen from slow variations in the ``standard'' stars in the field. The results presented here make an important contribution  to the study of these target objects, as the such extensive studies have not yet been performed on their intra-night timescales and the optical behaviour of this class of BL Lacs is still largely unknown. Our measurements provide quasi-simultaneous measurements in two bands, unlike most earlier microvariability studies.\\par  There are several theoretical models that might be able to explain the observed variability over wide time-scales for all bands, with the leading contenders based upon shocks propagating down relativistic jets (e.g., Marscher \\& Gear 1985; Qian et al.\\ 1991; Hughes, Aller \\& Aller 1991; Marscher, Gear \\& Travis 1992; Wagner \\& Witzel 1995; Marscher 1996). In this class of models much of the variability arises when a shock strikes a feature in the jet. Some of the variability may arise from helical structures, precession, or other geometrical effects occurring within the jets (e.g., Camenzind \\& Krockenberger 1992; Gopal-Krishna \\& Wiita 1992). These models seem to provide reasonable, albeit generic, representations of the variations observed on timescales from about a day up through months and probably, years. Hot spots or disturbances in or above accretion disks surrounding the black holes at the centers of AGNs (e.g., Mangalam \\& Wiita 1993; Chakrabarti \\& Wiita 1993) are likely to play a key role in the variability of non-blazar AGNs and might provide seed fluctuations that could be advected into the blazar jet and then Doppler amplified in the relativistic jets. \\par  Previous observational results have indicated that HBLs show a statistically significantly lesser amount of optical variability then that do LBLs (Stocke et al.\\ 1985, 1989, Heidt \\& Wagner 1996), and this conclusion is supported by our results, in that neither of these two HBLs ever showed microvariability. Sambruna et al.\\ (1996) investigated the multifrequency spectral properties of LBLs and HBLs and mentioned this spectral properties requires a systematic change of intrinsic physical parameters, such as magnetic field, jet size, and maximum electron energy and this change is in the sense that HBLs have higher magnetic fields/electron energies and smaller sizes than LBLs. The difference in the intra-day variability behaviour of HBLs could be due to the effect of stronger magnetic fields and/or electron energies in these objects (Sambruna et al.\\ 1996) that could prevent or delay the formation of features like density inhomogeneities and bends in the bases of the jets by Kelvin-Helmholtz instabilities (Romero et al.\\ 1999). For instance, in the model proposed by Sol et al.\\ (1989), extragalactic jets are composed of two different fluids, a highly relativistic spine and a slower sheath. Kelvin-Helmholtz instabilities must arise at the interface between these two fluids with different bulk velocities. Such instabilities might be responsible for many of the knots and bends observed in the large scale structures of jets. Romero (1995) showed that axial magnetic fields prevent the development of Kelvin-Helmholtz instabilities in sub-parsec to parsec scale jets if their values exceed the critical value B$_{c}$ given by \\begin{equation} B_{c}=[4 \\pi n m_{e}c^{2}(\\gamma^{2}-1)]^{1/2} \\gamma^{-1}, \\end{equation} where $n$ is the local electron density, $m_{e}$ is the electron rest mass, and $\\gamma$ is the central flows bulk Lorentz factor. So, if $B > B_{c}$ this class of instabilities will be suppressed.  However, if B $< B_{c}$ the Kelvin-Helmholtz instabilities can produce significant changes in the jet morphology, and these features could be responsible for rapid variability when the relativistic shock waves interact with them.  So if HBLs typically possess stronger magnetic fields for given electron densities or lower electron densities for given field strengths, than we suggest that such a reduction in Kelvin-Helmholtz instabilities would reduce the incidence of microvariability in the optical light curves.  Another possible explanation for the difference between the optical variability behaviour of the LBL and HBL classes of blazars can dispense with the hypothesis of suppressed Kelvin-Helmholtz instabilities. We recall that the optical band is near or above the peak of the first hump in the SED for LBLs but below that peak for HBLs. Therefore changes in the efficiency of acceleration of, and/or in the rates at which energy is radiated by, the highest energy electrons available for synchrotron emission would have a large and rapid effect on the emitted flux in those bands for LBLs.  However, those variations would have a more modest and more retarded effect on optical variability in HBLs, since their optical emission is below the peak of the synchrotron emission. Such variations in acceleration efficiency could arise from changes in the local number density of the most energetic electrons or the strengths of the maximum localized magnetic fields.  While magneto-hydrodynamic instabilities might be the origin of such variations, so might the presence of turbulence in the vicinity of the shock (e.g., Marscher, Gear \\& Travis 1992). If this is the case, then X-ray variability should be more pronounced for HBLs than for LBLs, in that the peak of the synchrotron emission lies near the X-ray band for the former class."
    },
    "1112/1112.4527_arXiv.txt": {
        "abstract": "With high resolution ($0.46h^{-1}$kpc), large-scale, adaptive mesh-refinement Eulerian cosmological hydrodynamic simulations we compute properties of O~VI and O~VII absorbers from the warm-hot intergalactic medium (WHIM) at $z=0$. Our new simulations are in broad agreement with previous simulations with $\\sim 40\\%$ of the intergalactic medium being in the WHIM. Our simulations are in agreement with observed properties of O~VI absorbers with respect to the line incidence rate and Doppler width-column density relation. It is found that the amount of gas in the WHIM below and above $10^6$K is roughly equal. Strong O~VI absorbers are found to be predominantly collisionally ionized. It is found that $(61\\%, 57\\%, 39\\%)$ of O~VI absorbers of $\\log{\\rm N(OVI)~cm^2}=(12.5-13,13-14,>14)$ have $T<10^5$K. Cross correlations between galaxies and strong [${\\rm N(OVI)}>10^{14}$cm$^{-2}$] O~VI absorbers on $\\sim 100-300$kpc scales are suggested as a potential differentiator between collisional ionization and photoionization models. Quantitative prediction is made for the presence of broad and shallow O~VI lines that are largely missed by current observations but will be detectable by COS observations. The reported $3\\sigma$ upper limit on the mean column density of coincidental O~VII lines at the location of detected O~VI lines by Yao \\etal is above our predicted value by a factor of $2.5-4$. The claimed observational detection of O~VII lines by Nicastro et al, if true, is $2\\sigma$ above what our simulations predict. ",
        "introduction": "Physical understanding of the thermodynamic evolution of the intergalactic medium (IGM) has been substantially improved  with the aid of {\\it ab initio} cosmological hydrodynamic simulations. One of the most robust predictions is that $40-50\\%$ of all baryons in the present universe is in the WHIM of temperature $10^5-10^7$K and overdensity $10-300$ \\citep[e.g.,][]{1999Cen, 2001Dave}. The predicted WHIM provides an attractive solution to the long standing missing baryons problem \\citep[][]{1992Persic, 1998Fukugita}. Let us first clarify the nomenclature of several related gas phases. The intra-group and intra-cluster medium (ICM) is defined to be gas within these virialized regions (i.e., overdensity $>100$). The high density portion (overdensity $\\ge 500$) of the ICM has traditionally been detected in X-ray emission; thermal Sunyaev-Zeldovich effect and more sensitive X-ray measurements can now probe ICM to about the virial radius. The circumgalactic medium (CGM) is usually defined to be gas that embeds the stellar components in galactic halos and may be made up of gases of a wide range of temperatures ($10^4-10^7$K) and densities. It is likely, at least for large galaxies, that a significant fraction of the CGM falls into the same temperature range of the WHIM. Of particular interest is some of the CGM that has been heated up by star formation feedback shocks to the WHIM temperature range \\citep[e.g.,][]{2006Cen, 2011Cen}. In the present analysis we define WHIM as gas of temperature $10^5-10^7$K with no density limits. Most of the WHIM gas is truly intergalactic with overdensity $<100$ (see Figure 7) and mostly easily probed in absorption.     The reality of the WHIM, at least its low temperature ($T\\le 10^6$K) portion, has now been fairly convincingly confirmed by a number of observations in the FUV portion of QSO spectra from HST and FUSE, through the O~VI $\\lambda \\lambda$1032, 1038 absorption lines that peak at $T\\sim 3\\times 10^5$K when collisionally ionized \\citep[e.g.,][]{2000Tripp, 2000bTripp, 2000Oegerle, 2002Savage, 2004Prochaska, 2004Sembach, 2005Danforth, 2006Danforth, 2008Danforth, 2008Tripp, 2008bThom, 2008Thom, 2008Cooksey} and Ne VIII $\\lambda \\lambda$770, 780 absorption lines that peak at $T\\sim 7\\times 10^5$K in collisional ionization equilibrium \\citep[][]{2005Savage, 2006Savage, 2009Narayanan, 2011Narayanan, 2011Tripp} as well as broad $\\lya$ absorption lines (BLAs) \\citep[][]{2010Danforth, 2011bSavage, 2011Savage}. In agreement with simulations, the part of WHIM detected in O~VI absorption is estimated to constitute about 20-30\\% of total WHIM. The detection of Ne VIII lines along at least some of the sight lines with O~VI detection provides unambiguous evidence for the WHIM origin, instead of lower temperature, photoionized gas, under physically plausible and observationally constrained situations. %  X-ray observations performed to search for X-ray absorption of the higher temperature portion ($T\\ge 10^6$K) of the WHIM associated with known massive clusters have also been successful. An XMM-Newton RGS spectrum of quasar LBQS 1228+1116 revealed a feature at the Virgo redshifted position of O~VIII Ly$\\alpha$ at the $95\\%$ confidence level \\citep[][]{2004Fujimoto}. Using XMM-Newton RGS observations of an AGN behind the Coma Cluster, the Seyfert 1 X Comae, \\citet[][]{2007Takei} claimed to have detected WHIM associated with the Coma cluster. Through the Sculptor Wall \\citet[][]{2009Buote} and \\citet[][]{2010Fang} have detected WHIM O~VII absorption at a column greater than $10^{16}$cm$^{-2}$. There is evidence of detection in soft X-ray emission along the filament connecting clusters A222 and A223 at $z=0.21$ that may be associated with the dense and hot portion of the WHIM \\citep[][]{2008Werner}.  However, the search for X-ray absorption of WHIM along random lines of sight turns out to be elusive. Early pioneering observations \\citep[][S5 0836+710, PKS 2149-306, PKS 2155-304]{2001Fang, 2002Fang} gave the first O~VII detection (O~VIII for PKS 2155-304), which has not been convincingly confirmed subsequently \\citep[][]{2004Cagnoni, 2007bWilliams, 2007Fang}. \\citet[][]{2003Mathur} performed a dedicated deep observation (470 ks) with the Chandra LETGS of the quasar H 1821+643, which has several confirmed intervening O V I absorbers, but found no significant ($>>2\\sigma$) X-ray absorption lines at the redshifts of the O V I systems. \\citet[][]{2005Nicastro, 2005bNicastro} embarked on a campaign to observe Mrk 421 during its periodic X-ray outbursts with the Chandra LETGS with a total of more than 7 million continuum counts and presented evidence for the detection of two intervening absorption systems at $z=0.011$ and $z=0.027$. But the spectrum of the same source observed with the XMM-Newton RGS did not show these absorption lines \\citep[][]{2007Rasmussen}, despite higher signal-to-noise and comparable spectral resolution. \\citet[][]{2006Kaastra} and \\citet[][]{2012Yao} reanalyzed the Chandra LETGS data and were in agreement with \\citet[][]{2007Rasmussen}. The detection of O~VII lines may also be at odds with recent BLA measurements \\citep[][]{2011Danforth}, under simplistic assumptions about the nature of the absorbing medium. However, it has been argued that the reported XMM-Newton upper limits and the Chandra measurements may be consistent with one another, when taking into consideration certain instrumental characteristics of the XMM-Newton GRS \\citep[][]{2006Williams}. Moreover, an analysis of the two candidate X-ray absorbers at $z=0.011$ and $z=0.027$ yields intriguing evidence of two large-scale filaments at the respective redshifts, one of which has only $5-10\\%$ probability of occurring by chance \\citep[][]{2010Williams}. Observations of 1ES 1028+511 at $z=0.361$ by \\citet[][]{2006Steenbrugge} yield no convincing evidence for X-ray WHIM absorption.  What is perceived to be more disconcerting is the lack of detection of O~VII absorbers at the redshifts of detected O~VI absorbers along some random lines of sight. This is because, overall, the O~VII line is predicted to be the most abundant and anecdotal evidence suggests substantial coincidence between O~VI and O~VII \\citep[e.g.,][]{2006bCen}. A statistically significant upper limit placed on the mean column density of O~VII absorbers at the locations of a sizeable set of detected O~VI absorbers using stacking techniques by \\citet[][]{2009Yao} prompts them to call into question the very existence of the high temperature ($T\\ge 10^6$K) portion of the WHIM, although the limited sensitivity and spectral resolution of the current X-ray observations may render any such conclusions less than definite.   Therefore, at this juncture, it is pressing to statistically address this lack of significant coincidence between O~VI and O~VII absorbers and other issues theoretically, through higher resolution simulations that are necessary in order to well resolve the interfaces of multi-phase media. This is the primary purpose of this paper. We use two simulations of high resolution of $0.46h^{-1}$kpc and box size of $20-30h^{-1}$Mpc to perform much more detailed characterization of O~VI and O~VII lines to properly compare to extant observations. This high resolution is to be compared with $83h^{-1}$kpc resolution in our previous simulations \\citep[][]{2006Cen, 2006bCen}, $25-49h^{-1}$kpc resolution in \\citet[][]{2011Smith} and \\citet[][]{2011Shull}, $1.25-2.5h^{-1}$kpc in \\citet[][]{2012Oppenheimer} % and $1.25-2.5h^{-1}$kpc in \\citet[][]{2011TepperGarcia}, % resolves the Jeans scale of WHIM by 2-3 orders of magnitude and interfaces between gas phases of different temperatures in a multi-phase medium. It is useful to distinguish, in the case of SPH simulations, between the gravity force resolution and the resolution of the hydrodynamics solver, with the latter being worse than the former by a factor of order a few. It is also useful to keep in mind the initial cell size or interparticle separation, because in both SPH and adaptive mesh refinement (AMR) simulations not all regions are resolved by the maximum resolution. Calling this ``mean region resolution\" $\\Delta_{\\rm root}$, $\\Delta_{\\rm root}=(117, 25-49, 125, 195)h^{-1}$kpc for [this paper, \\citet[][]{2011Smith}, \\citet[][]{2012Oppenheimer}, \\citet[][]{2011TepperGarcia}]. We note that a region of overdensity $\\delta$ is approximately resolved at a resolution of ${\\rm C}\\Delta_{\\rm root}\\delta^{-1/3}$ (up to a pre-specified highest resolution), where the pre-factor $C$ is about unity for AMR simulations and $\\sim 2$ for SPH simulations. Using Lagrangian SPH or AMR approaches becomes necessary for regions $\\delta\\ge 300$, because simulations of a similar resolution with the uni-grid method become increasingly impractical (largely due to limitations of computer memory). A more important advantage with very high resolution simulations has to do with the need to resolve galaxies, which in turn allows for a more self-consistent treatment of the feedback processes from star formation, namely, the temporal and spatial distribution of metals and energy deposition rates to the CGM and IGM and their effects on subsequent star formation.  Our new simulations, in agreement with earlier findings, reaffirm quantitatively the existence of WHIM and furthermore show that the properties of the WHIM with respect to O~VI line and O~VI-O~VII relations are fully consistent with observations. In particular, the observed upper limit of the mean coincidental O~VII column density of detected O~VI absorbers is higher than what is predicted by the simulations by a factor of $\\sim 2.5-4$. Higher sensitivity X-ray observations or a larger sample by a factor of $\\sim 10$ should test this prediction definitively. The outline of this paper is as follows. In \\S 2.1 we detail simulation parameters and hydrodynamics code, followed by a description of our method of making synthetic O~VI and O~VII spectra in \\S 2.2, which is followed by a description of how we average the two separate simulations C (cluster) and V (void) run in \\S 2.3. Results are presented in \\S 3. In \\S 3.1 we present some observables for O~VI to compare to observations to provide additional validation of the simulations. In \\S 3.2 we dissect the simulations to provide a physical analysis of the O~VI and O~VII absorbers. In \\S 3.3 results on the coincidences between O~VI and O~VII lines are given. Conclusions are summarized in \\S 4. ",
        "conclusions": "Utilizing high resolution ($0.46h^{-1}$kpc), adaptive mesh-refinement Eulerian cosmological hydrodynamic simulations we examine properties of O~VI and O~VII absorbers in the warm-hot intergalactic medium (WHIM) at $z=0$, along with a physical examination. We find that our new high resolution simulations are in broad agreement with all other simulations with respect to the thermal distribution of baryons in the present universe. In particular, we find that about $40\\%$ of the intergalactic medium is in the WHIM. We find that our simulations are in excellent agreement with observed properties of O~VI absorbers, including line incidence rate, Doppler width-column density relation, and consistent with observed Doppler width-temperature relation. Physical properties of O~VI and O~VII absorbers are given, including inter-relations between metallicity, temperature, density, Doppler width, to facilitate a coherent understanding. We highlight some of the important or new findings.  (1) We find that strong O~VI absorbers are predominantly collisionally ionized, whereas for weaker absorbers the contributions from  photoionization become progressively more important. We find that $(39\\%, 43\\%, 61\\%)$ of O~VI absorbers in the column density ranges of $\\log{\\rm N(OVI)cm^2}=(12.5-13,13-14,>14)$ have temperature greater than $10^5$K. This may be contrasted with the results of \\citet[][]{2009Oppenheimer} where low temperature ($\\sim 10^4$K), high metallicity, photoionized O~VI absorbers dominate even at high column densities ($\\log{\\rm N(OVI)cm^2}>14$). We concur that lack of metal mixing in SPH simulations, which in turn causes overcooling of high-metallicity feedback SPH particles, most severely in high density regions, may be able to account for the discrepancy. We suggest that cross correlations between strong [${\\rm N(OVI)}\\ge 10^{14}$cm$^{-2}$] O~VI absorbers and galaxies on $\\sim 100$kpc scales may be able to differentiate between the models.  (2) Velocity structures within absorbing regions are a significant, and for large Doppler width clouds, a dominant contributor to the Doppler widths of both O~VI and O~VII absorbers. Doppler width is thus a poor indicator of temperature.  (3) Quantitative prediction is made for the presence of broad and shallow O~VI lines, which current observations have largely failed to detect. Upcoming observations by COS may be able to provide a test.  (4) The coincidence rates between O~VI and O~VII lines are found to be small, for which physical explanations are given. We find that the reported $3\\sigma$ upper limit on the mean column density of coincidental O~VII lines at the location of detected O~VI lines by \\citet[][]{2009Yao} is above the predicted value by a factor of $2.5-4$, implying that a factor of $\\sim 10$ increase in sample size or sensitivity will be able to yield a definitive detection of O~VII column density using the stacking technique even without detection of individual O~VII absorbers.  (5) We show that, if the previously claimed observational detection of O~VII lines by \\citet[][]{2005Nicastro} is true, our predicted O~VII line density is $2\\sigma$ below that. This shows that higher sensitivity X-ray observations of this or other viable targets will be very useful to potentially place strong constraints on the model.      \\vskip 1cm  I would like to thank Dr. M.K.R. Joung for help on generating initial conditions for the simulations and running a portion of the simulations and Greg Bryan for help with Enzo code. I would like to thank the referee Mike Shull for critical and constructive reports. I would like to thank Dr. Edward Jenkins for a careful reading of the manuscript and helpful discussion, Dr. Charles Danforth for kindly providing the observational data and useful discussion, Dr. Jeremiah P. Ostriker for useful discussion and Drs. John Wise, Matthew Turk and Cameron Hummels for help with visualization program yt \\citep[][]{2011Turk}. Computing resources were in part provided by the NASA High- End Computing (HEC) Program through the NASA Advanced Supercomputing (NAS) Division at Ames Research Center. This work is supported in part by grants NNX08AH31G and NAS8-03060. The simulation data are available from the author upon request."
    },
    "1112/1112.1072_arXiv.txt": {
        "abstract": "We show that the factor $\\sim$2 discrepancy between spectroscopic and photometric age determinations of the Galactic bulge main-sequence turnoff can be naturally explained by positing an elevated helium enrichment for the bulge relative to that assumed by standard isochrones. Helium-enhancement relative to standard isochrones is confirmed at the 2.3$\\sigma$ level. We obtain an upper bound on the helium enrichment for the metal-rich ([Fe/H]$\\approx+0.30$) stars of ${\\Delta}Y\\approx+0.11$ relative to canonical expectations, given the requirement that the spectroscopic and photometric ages be consistent and the limiting condition of instantaneous star formation. We discuss phenomenological evidence that the bulge may have had a chemical evolution that is distinct from the solar neighborhood in this manner, and we make several testable predictions. Should this emerging picture of the bulge as helium-enhanced hold, it will require the development of new isochrones, new model atmospheres, and modified analysis and cosmological interpretation of the integrated light of other bulges and elliptical galaxies. ",
        "introduction": "\\label{sec:Introduction} The mean age and age-spread of the Galactic bulge, $t_{\\rm{Bulge}}$ and ${\\Delta}t_{\\rm{Bulge}}$, are fundamental parameters of Galactic evolution, with hierarchical formation theories predicting an older bulge with more rapid formation timescale than models where bulge formation was linked to early disk \\citep{2004ARA&A..42..603K,2011A&A...525A.126C,2011arXiv1109.2898I}.  However, age determinations of bulge stars are currently in a state of dissonance, with spectroscopic ages extending to much lower ages than photometric determinations.  A seemingly secure picture had emerged from \\textit{Hubble Space Telescope (HST)} photometry of the Galactic bulge main-sequence turnoff (MSTO). \\citet{2002AJ....124.2054K} and \\citet{2003A&A...399..931Z} both showed that the bulge MSTO is dim relative to the bulge horizontal branch (HB), with a brightness offset indistinguishable from that found in Galactic globular clusters (GCs), suggesting an old age. \\citet{2011ApJ...735...37C} refined this approach by searching for eclipsing binary blue-straggler stars near the MSTO. They found that whatever sparse population of bright stars there was could be entirely explained by bulge blue stragglers, and concluded that no more than 3\\% of bulge stars are younger than 5 Gyr. Thus, the statement that the bulge stellar population is old appeared robust: bulge MSTO stars are dim.  However, this picture has been challenged by  \\citet{2010A&A...512A..41B,2011A&A...533A.134B}, who used gravitational microlensing events to obtain high-resolution spectra of bulge MSTO and subgiant (SGB) stars, which yielded detailed measurements of $\\log{g}$, $T_{\\rm{eff}}$, [Fe/H] and 12 additional chemical abundances. Their sample, microlensed source stars (only $\\sim$3\\% of microlensed sources are expected to be disk stars) with coordinates $(|l| \\lesssim 5, 2 \\lesssim |b| \\lesssim 5)$ and radial velocity dispersion of $\\sim$100 km/s, is  expected to be a predominantly bulge sample. Comparison of the measurements to isochrones led to a big surprise: a vast population of young stars. Seven of their 26 spectra (27\\%) yield ages, $t_{\\rm{Inferred}}\\leq5$ Gyr, an outcome with probability $P < 10^{-5}$ if one assumes the 3\\% constraint. This method may seem complex at first glance, but upon close inspection it also appears robust. Consider the MSTO star MOA-2008-BLG-311S. At its high metallicity, [Fe/H]$=+0.36\\pm0.07$, standard isochrones \\citep{2008ApJS..178...89D} predict it should be no hotter than $T_{\\rm{eff}} \\approx 5450 K$ if one assumes an old age. However, its measured temperature is $T_{\\rm{eff}} = 5944 \\pm 68 K$, yielding the spectroscopic fit $t_{\\rm{Inferred}} \\sim 2.9$ Gyr. Next, consider the metal-rich ([Fe/H]$=+0.37\\pm0.05$) SGB star MOA-2009-BLG-259S. At its temperature, $T_{\\rm{eff}} = 4953 \\pm 93 K$, the expected surface gravity is $\\log{g} \\approx 3.90$. However, the measured value is $\\log{g}=3.40 \\pm 0.24$, for $t_{\\rm{Inferred}} \\sim 3.0$ Gyr. Therefore, the paradigm-challenging data reported by \\citet{2010A&A...512A..41B,2011A&A...533A.134B} pass the consistency test of being independently demonstrated in two distinct phases of stellar evolution. In summary: \\begin{itemize} \\item \\textit{HST} photometry: a faint MSTO, indicating an old stellar population \\citep{2002AJ....124.2054K,2003A&A...399..931Z,2010ApJ...725L..19B,2011ApJ...735...37C}. The similarity with Galactic GCs suggests $t_{\\rm{Bulge}} \\approx$ 12.8 Gyr \\citep{2009ApJ...694.1498M}. \\item High-resolution spectroscopy of [Fe/H]$>$0 stars: high  $T_{\\rm{eff}}$ for the MSTO and low $\\log{g}$ for the SGB, indicating a young stellar population \\citep{2010A&A...512A..41B,2011A&A...533A.134B}. They find a weighted mean age of 5.9 Gyrs for the super-solar metallicity stars. \\end{itemize} In this \\textit{Letter} we suggest that the origin of this discrepancy is with the isochrones that are used to interpret the data and estimate the ages.  Specifically, we question whether the adopted assumption of scaled-solar helium abundance is a valid approximation for the chemical evolution of the Galactic bulge at the metal-rich end. We demonstrate that the resulting impact on stellar parameter determination leads to a bias in inferred ages: photometric determinations are too old and spectroscopic determinations are too young.      \\begin{figure}[H] \\begin{center} \\includegraphics[totalheight=0.58\\textheight]{DartmouthIsochrone4} \\end{center} \\caption{Applying standard isochrones to helium-enhanced populations overestimates photometric ages and underestimates spectroscopic ages. TOP: [Fe/H]$=$0, Y$=$0.274 (solar) isochrones with ages 6-13 Gyr in black (thicker for t$=$10 Gyr), and an [Fe/H]$=$0, Y$=$0.400 (enhanced), t$=$10 Gyr isochrone plotted in red. On the spectroscopic plan, helium-enhanced populations mimic younger, standard populations with the effect minimized near the MSTO. BOTTOM: Same color scheme as above, but with tracks shown on the photometric $(V-I,I)$ plane. The 10 Gyr helium-enhanced track has a much fainter MSTO, equivalent to a t$\\sim$13.5 Gyr track of the same metallicity and standard helium. A large difference in helium, ${\\Delta}Y=$0.126,  has been selected for demonstrative purposes of clarity. Isochrones are from the Dartmouth stellar evolution database \\citep{2008ApJS..178...89D}. } \\label{Fig:Dartmouth4} \\end{figure} ",
        "conclusions": "\\label{sec:Discussion}  We have shown in this \\textit{Letter} that the discrepancy in Galactic bulge turnoff age estimates can be well-explained by the hypothesis that the chemical evolution of the Galactic bulge is helium-enhanced relative to that assumed by standard isochrones. We first found evidence for helium-enhancement at the 2.3$\\sigma$ level, suggesting ${\\Delta}Y=+0.098$ for the metal-rich stars. We also derived an upper bound on the helium enrichment of the bulge, ${\\Delta}Y \\lesssim 0.11$ for the metal-rich ([Fe/H]$\\approx+0.30$) stars. We deliver two testable predictions, that the absolute magnitudes will be systematically fainter than the best-fit values from standard isochrones, and that inferred ages will be highest near the MSTO. Both diagnostics suggest helium-enrichment for the metal-rich stars, with no such evidence for the metal-poor stars.  There are multiple roads forward. On the theory side, chemical evolution models will be required to ascertain what conditions are required in the primeval Milky Way for a high value of $({\\Delta}Y/{\\Delta}Z)_{\\rm{Bulge}}$. A broader range of isochrones than currently available will be required to interpret the available photometric and spectroscopic data.  Finally, it is possible that the stellar atmosphere models used by spectroscopists fail in the case of atmospheres that are extremely helium-enriched.  The [Fe/H] distribution of MSTO+SGB stars \\citep{2011A&A...533A.134B} is inconsistent with that derived from giants \\citep{2008A&A...486..177Z,2011A&A...534A..80H,2011ApJ...732..108J}, as it has a trough near [Fe/H]$=$0, where the giant metallicities show a peak. Could this be the reason? Observationally, an uncertainty in the value of $({\\Delta}Y/{\\Delta}Z)_{\\rm{Bulge}}$ should be incorporated in analyses of bulge stars. The age discrepancy discussed in this paper constitutes the fourth, independent line of evidence for a helium-enriched bulge, following previous investigations of star counts on the upper RG branch, on the RGBB, and chemical evolution arguments based on the high values of [$\\alpha$/Fe].  The helium abundances suggested here are difficult to attain with chemical evolution models, as ${\\Delta}Y{\\Delta}Z \\lesssim 3.0$ for general assumptions \\citep{2002A&A...395..789C,2008RMxAA..44..341C}. Should the inferred helium-enhancement for the bulge hold up under further scrutiny it may necessitate rethinking the underlying assumptions of chemical evolution models. Additionally, even in the disk, there may be surprises with respect to helium-enrichment: \\citet{2010A&A...518A..13G} find ${\\Delta}Y/{\\Delta}Z=5.3\\pm1.4$ in a sample of local K-dwarfs.  The preliminary estimates of this Letter result from the assumption of instantaneous star formation. How much age spread is reasonable depends on assumptions of the initial conditions of the Galaxy. The assumption of rapid gravitational infall yields ${\\Delta}t_{\\rm{Bulge}}\\approx0.5$ Gyr \\citep{2011A&A...525A.126C}, whereas the ``clump-origin bulge'' model, in which the bulge is formed by stellar clumps forming in a gas-rich disk and accreting to its center by dynamical friction, yields ${\\Delta}t_{\\rm{Bulge}}\\approx2.0$ Gyr \\citep{2011arXiv1109.2898I}. However, recent investigations of star counts and radial velocities find that bulge kinematics are consistent with a pure N-body bar that evolved from secular instabilities, which would imply that bulge stars are just disk stars on bar orbits \\citep{2010ApJ...720L..72S,2011ApJ...734L..20M}. This would suggest a duration of star formation as extended as that of the inner disk.     The results of these endeavors may have significant cosmological implications. Population synthesis studies of other galaxies assume a solar value of ${\\Delta}Y/{\\Delta}$Z. If this parameter is demonstrated to be non-universal for field stars, it will affect age and mass determinations of field galaxies."
    },
    "1112/1112.3846_arXiv.txt": {
        "abstract": "The termination structures of the jets of Fanaroff \\& Riley (FR) galaxies are observed to produce extended non-thermal emission in a wide frequency range. The study of these structures can provide valuable insights on the conditions for particle acceleration and radiation at the shock fronts. We have studied the thermal and non-thermal emission that can be expected from the jet termination regions of Fanaroff \\& Riley type I sources. The broadband emssion from these galaxies has been recently extended to include the high-energy gamma-ray domain, owing to the \\textit{Fermi} detection of Cen~A lobes. Exploring the physics behind the jet/medium interactions in FRI can provide valuable insights on the conditions for particle acceleration and radiation in the jet termination shocks. Making use of the results of a fully relativistic numerical simulation code of the evolution of a FRI jet we model the expected radiative output and predict spectra and lightcurves of both thermal and non-thermal emission at different source ages. ",
        "introduction": "Fanaroff-Riley sources of type I and II (FRI-II) are radio loud Active Galactic Nuclei (AGN) displaying large-scale jets that interact with the surrounding ISM/ICM medium.\\cite{fr74}. The study of the extended emission in FRI/II sources can be used to derive the properties of the flows and the surrounding medium, and to determine the conditions for particle acceleration. In this regard a simple evolutionary model can be used to conclude that thermal X-ray emission produced in the shell could be detected only for dense enough media even for powerful FRII jets,\\cite{hrb98} whilst thermal MeV radiation from the cocoon region would be significant only for sources with ages $\\ll 10^7\\,\\rm{yrs}$.\\cite{kki07}. Non-thermal radiation can also be produced in the jet/medium interaction regions. X-ray synchrotron and/or IC emission has been observed in several sources (see e.g. \\refcite{ka03,cr09}), and the {\\it Fermi} Collaboration recently reported the detection of extended GeV emission in the radio lobes of Cen~A, demonstrating that acceleration up to very high energies is indeed taking place in the disrupted jet region.\\cite{ab10a}  Numerical simulations, on the other hand, have demonstrated to be a powerful tool to study the jet/medium interactions. Perucho \\& Mart\\'i (2007) performed a simulation aimed to test the FRI jet evolution paradigm.\\cite{pm07,b84} We computed the thermal and non-thermal emission from FRI sources obtained with a radiative model adapted from Bordas et al. (2009) that takes as input the results of those simulations as input and the radiation model used to characterize the conditions of the shocked regions.\\cite{br09} Below, we briefly introduce the numerical simulations and the radiation model used. We then discuss the obtained results and the relevance of the thermal and non-thermal radiation at different energy bands. For a more extended discussion, the reader is referred to Perucho \\& Mart\\'i (2007) and Bordas et al. (2011).\\cite{pm07, br11} ",
        "conclusions": "Our results can be discussed on the ground of the thermal and non-thermal fluxes predicted at different energy bands. In radio, the cocoon dominates with fluxes up to $\\sim few \\times 10^{-12}\\,(d/100~$Mpc$)^{-2}$~erg~cm$^{-2}$~s$^{-1}\\sim 10$~Jy at 5~GHz from a region of few times $10'\\,(d/100~$Mpc$)^{-1}$ of angular size. This radio fluxes are similar to those observed for instance in 3C~31 and 3C~15,\\cite{a92,ka03}.The shell radio flux is sligthly below the cocoon one, although limb brightening may enhance its detectability. The radio lightcurve is quite steady, since the accumulation of electrons compensates the weakening of the magnetic field with $t_{\\rm src}$.   At X-rays, the different thermal emitting regions in the shell would lead to a bolometric flux from $\\sim 10^{-15}\\,(d/100~$Mpc$)^{-2}$ ($10^5$~yr) to a few times $10^{-13}\\,(d/100~$Mpc$)^{-2}$~erg~cm$^{-2}$~s$^{-1}$ ($10^8$~yr) in 3C 15. The hard X-rays would be located at the apex of the bow shock, with a typical angular size of a few $1'\\,(d/100~$Mpc$)^{-1}$ (for $t_{\\rm src}\\sim 10^8$~yr), whereas lower energy X-rays would come from the whole shell, with of a few $10'\\,(d/100~$Mpc$)^{-1}$. Regarding non-thermal X-rays, the cocoon dominates the total output, with fluxes $\\sim 10^{-13}\\,(d/100~$Mpc$)^{-2}$~erg~cm$^{-2}$~s$^{-1}$, although limb brightening effects may again increase the shell brightness. In fact, in the case of Cen~A, the shell seems to be the dominant source of non-thermal X-rays.\\cite{cr09} This difference could be explained by a higher $E_{\\rm max}$ in the shell of that source. The non-thermal X-ray fluxes $\\sim 10^{-14}$~erg~cm$^{-2}$~s$^{-1}$ of 3C 15 (see e.g. \\refcite{ka03}) imply a non-thermal luminosity of $\\approx 1.3 \\times 10^{41}$~erg~s$^{-1}$ at 300~Mpc, in good agreement with the values predicted here. Furthermore, synchrotron emission concentrated around the recollimation shock is compatible with the large-scale jet X-ray emission found in 3C~31.\\cite{hr02} If a strong recollimation shock is indeed the origin of these large-scale jet X-rays, then the hypothesis that jet disruption in 3C~31 is caused by shock triggered instabilities is favored against stellar wind mass-load.\\cite{lb02} Finally, we note that for sources older than those considered here and/or lower $B$-values, the synchrotron emission would be less relevant and IC would dominate the X-ray output.      The moderate velocities of the bow shock found in the numerical simulations prevents thermal photons from reaching energies as high as those discussed in \\refcite{kki07}. In the gamma-ray domain, the SED is close to flat at HE, and becomes steeper at VHE. We have not accounted for the gamma-ray absorption induced by the extragalactic background light, which would become significant at distances larger than 100~Mpc. The overall emission increases with time mainly due to the increasing efficiency of the IC channel (considering the cosmic microwave background only) as compared with synchrotron and adiabatic losses. Gamma-ray fluxes for a source with $t_{\\rm src}\\sim 10^8$~yr are around $\\sim 10^{-12}\\,(d/100\\,{\\rm Mpc})$~erg~cm$^{-2}$~s$^{-1}$. At HE, such a source may require very long exposures to be detected by, e.g., {\\it Fermi}, unless it is very nearby (as it is the case of Cen~A, see \\refcite{ab10a}) or shows a higher non-thermal efficiencies and/or jet power. At VHE, the fluxes could be observed by the current Cherenkov telescopes, although the extension of the source, of tens of arcminutes at 100~Mpc, and the steep spectrum above $\\sim 100$~GeV, may make a detection possible only after a long exposure. The forthcoming CTA observatory may however allow the detection of VHE emission from FRI jet lobes, and possibly carry out energy-dependent morphological studies.    From our study we conclude that, for moderate non-thermal luminosities, radio lobes of FRI radio galaxies are good candidates to be detected in the whole spectral range, with the radiation appearing extended in most of the energy bands. Soft X-rays will be likely dominated by synchrotron emission up to ages $\\sim 10^8$~yr, with IC tending to be dominant for older sources. Thermal X-rays seem unavoidable and may dominate in hard X-rays in old sources even if a non-thermal component is present. The low surface brightness may require long observation times for the detection in X- and gamma-rays, although the steady nature of these sources can help in this regard."
    },
    "1112/1112.6187_arXiv.txt": {
        "abstract": "Measurements of the power spectrum of the seeing in the range 0.001-1 Hz have been performed in order to understand the criticity of the transits' method for solar diameter monitoring. ",
        "introduction": "We first measured the daytime seeing by projection of the solar image on a regular grid during a drift-scan observation.[\\cite{Sigismondi}] The transit of the solar limbs above the grid are recorded, and the time intervals required to cover the evenly spaced intervals of the grid are measured by a frame by frame inspection of the video. The standard deviation of these time intervals $\\sigma$ [s] is related to the seeing $\\rho$ [arcsec] by the approximate formula: $\\rho = \\sigma \\cdot 15\\cdot \\cos(\\delta_{\\odot})$ where $\\delta_{\\odot}$ is the declination of the Sun at the moment of the observation. The frequency of the seeing depends on duration of these time intervals: we used 1, 2, 4 and 8 times the single space, to sample the first 4 frequencies at the 6.3 cm Lucernaria lenses in Santa Maria degli Angeli in Rome:[\\cite{Cuevas}] they are fixed in the dome and they have a focal lenght of 20 m.  These measurements have been extended to more frequencies at the 45 cm telescope of IRSOL and at the 10 cm Huariou solar station's SMAT telescope.  \\begin{figure} \\centerline{\\includegraphics[width=0.8\\textwidth,clip=]{2.EPS}} \\caption{Seeing amplitude at different frequencies with Lucernaria lenses.} \\label{Fig. 1} \\end{figure} ",
        "conclusions": "The difference between two following measurements of the solar diameter always experienced with drift-scan method,[\\cite{Wittmann}] can be explained by the identification of this slow image motion of the whole solar disk with 1 arcsec amplitude. This can also explain the need of averaging over several measurements to give a reference value for the solar diameter in a given day. But different meteorological conditions, even in the same day, occur during the measurements and they cannot be considered as statistically independent, as the Gaussian hypothesis requires. That is one of the reasons to explain the great scatter of the yearly averages published from Greenwich and Capitol Observatory.[\\cite{Gething}]  \\thanks{The authors C. S., M. M. and L. N. {\\bf acknowledge} a grant of FRAEN Corporation to the Clavius projet for the accurate measurement of the solar diameter.}"
    },
    "1112/1112.3077_arXiv.txt": {
        "abstract": "A perturbed black hole emits gravitational radiation, usually termed the ringdown signal, whose frequency and time-constant depends on the mass and spin of the black hole. I investigate the case of a binary black hole merger resulting from two initially non-spinning black holes of various mass ratios, in quasi-circular orbits. The observed ringdown signal will be determined, among other things, by the black hole's spin-axis orientation with respect to Earth, its sky position and polarization angle - parameters which can take any values in a particular observation. I have carried out a statistical analysis of the effect of these variables, focusing on detection and measurement of the multimode ringdown signals using the reformulated European LISA mission, Next Gravitational-Wave Observatory, NGO, the third generation ground-based observatory, Einstein Telescope and the advanced era detector, aLIGO. To the extent possible I have discussed the effect of these results on plausible event rates, as well as astrophysical implications concerning the formation and growth of supermassive and intermediate mass black holes. ",
        "introduction": "Astrophysical observations to date have provided sturdy evidence that black holes (BHs) may exist and play an important role in many physical processes \\cite{Nandra:2009br, Somerville:2008ch, AmaroSeoane:2009ui}. With direct evidence still lacking, it is expected that  observation of gravitational waves (GWs) from merged BHs will not only provide indisputable evidence for the existence of BHs, but also the ability to extract accurate information about the progenitor system and the BH.  The ringdown radiation consists of a superposition of, in principle, an infinite number of essentially damped sinusoids, termed quasi-normal modes. Their frequencies and time-constants depend only on the mass and spin of the BH -- a consequence of the {\\em no-hair} theorem \\cite{BCW05}. In a recent work \\cite{PhysRevD.85.024018} we have argued that the amplitude terms of the various quasi-normal modes encode important information about the origin of the perturbation that caused them, such as the component masses of the progenitor binary. This allows performing parameter estimation on the system from the strong-field regime, as opposed to using the inspiral phase. However, in that study, as well as in previous studies of parameter estimation from ringdown signals \\cite{BCW05,Berti:2007a}, only a small region of the parameter space was explored. In a realistic scenario, a BH ringdown signal can have any sky location and polarization. All the while, it could have originated from a BH with any spin-axis orientation with respect to Earth. These variables have a significant effect on the observed ringdown signal and a direct impact on the science we can achieve by observing GWs from merged BHs.  In the present study, I have investigated in detail how these angular parameters affect the detection and measurement of the ringdown signals. To this end, I have varied the angular parameters over their full range, thereby considering a large population of BH mergers. I have considered supermassive BHs (SMBHs of mass $\\geq 10^6M_\\odot$) visible in NGO and intermediate mass BHs (IMBHs of mass $\\sim10^3M_\\odot$) observable in ET and aLIGO. I have computed the probability distribution functions of signal to noise ratios (SNR), as well as measurement errors of a chosen set of parameters, for a wide range of the BH mass and for mass ratios between 2 and 20. Finally, I translate these probabilities to proportions of observed events in NGO and ET that will yield parameter errors below certain thresholds and discuss how observations of ringdown signals could help in dealing with open questions on the existence and history of SMBHs and IMBHs. \\vspace{-0.19in}  \\begin{figure}[ht] \\begin{minipage}[b]{0.5\\linewidth} \\includegraphics[scale=0.3]{PSDs-eps-converted-to.pdf} \\caption{Noise power spectral densities for various proposed configurations of LISA-like space detectors. The dashed red line corresponds to the original LISA mission, while the rest of them refer to European designs of LISA. In all cases, the galactic binary white dwarf confusion noise \\cite{Nelemans:2003ha, VanDenBroeck:2010fp} is included, which has a negligible visible effect on the newer LISA curves though, due to their, almost two orders of magnitude worse sensitivity. Additionally, a low frequency cut-off - not shown - was induced at $5\\times10^{-5}Hz$. In this study, I am using the latest arrangement for the European mission of LISA, labeled eLISA-NGO in this graph.} \\label{fig:PSDs} \\end{minipage} \\hspace{5pt} \\begin{minipage}[b]{0.5\\linewidth} \\includegraphics[scale=0.3]{RelativeAmplitudes-eps-converted-to.pdf} \\caption{This plot shows the relative amplitudes of modes $(3,3)$, $(2,1)$ and $(4,4)$ over $(2,2)$ as a function of the mass ratio. The diamond points show the values that correspond to a time 10$M$ after the peak luminosity of the $22$ mode in the equal mass case. (see also Fig.3 in \\cite{PhysRevD.85.024018}) Fits to these points were used in this previous work. On the other hand, the circle points, which are used in the present study, were computed by taking into account all the points in the waveform, in a time region starting at the peak luminosity of $22$ and ending 30$M$ later. The solid lines shown, constitute fits to these circle points, given by expressions (\\ref{eq:a33})-(\\ref{eq:a44}).} \\label{fig:relampl} \\end{minipage} \\end{figure} \\vspace{-0.05in} ",
        "conclusions": "The present study constitutes a sensible and realistic approach to the subject of parameter estimation from a multimodal ringdown signal, inasmuch as it is supposed to be emitted from a merged binary in a generic configuration. Parameters such as the inclination angle $\\iota$, along with the sky location $\\theta$ and $\\varphi$ and the signal polarization $\\psi$ have an effect on the observed quasi-normal mode spectrum. Their impact on the detectability and parameter estimation has been assessed, by performing a large sample of Fisher-matrix analysis simulations, allowing for a simple statistical analysis of the results.  I am quoting frequency distributions for the errors at the representative distances of $z\\simeq1$ (6.73 Gpc) for supermassive BHs, in NGO, and of $z\\simeq0.2$ (1 Gpc) for intermediate-mass BHs, in the Einstein Telescope and advanced LIGO. An additional simulation was performed for a $5\\times10^6M_\\odot$ BH at $z\\simeq3$. The results are quite satisfactory in determining the mass and spin, especially for the low mass ratios from 1 to $\\sim5$, where in typically $90\\%$ of the cases their errors fall below $10\\%$. The effectiveness in measuring the luminosity distance, mass ratio and inclination angle is almost an order of magnitude worse.  The effects of these results on, as much as possible, realistic event rates in NGO and ET are discussed. The likelihood to have a waveform parameter measured to an accuracy of a certain threshold translates to the same proportion of observed events featuring error values below that threshold. As an example, if supermassive BHs within a luminosity distance of $z=1$ coalesce at a rate of $\\simeq10$, then NGO could act as a SMBH dynamics probe, as almost all of these events will yield very low errors $1\\%-10\\%$ in the BH mass and spin, while in half of the events the errors in the luminosity distance, mass ratio and inclination angle will be of the order of $20\\%$."
    },
    "1112/1112.3594_arXiv.txt": {
        "abstract": "Mestel determined the surface mass distribution of the finite disk for which the circular velocity is constant in the disk and found the gravitational field for points in the $z=0$ plane. Here we find the exact closed form solutions for the potential and the gravitational field of this disk in cylindrical coordinates over all the space.  The Finite Mestel Disk (FMD) is characterized by a cuspy mass distribution in the inner disk region and by an exponential distribution in the outer region of the disk.  The FMD is quite different from the better known exponential disk or the untruncated Mestel disk which, being infinite in extent, are not realistic models of real spiral galaxies.  In particular, the FMD requires significantly less mass to explain a measured velocity curve. ",
        "introduction": "\\cite{bin08} give a summary of some of the explicit potential-density pairs available for the study of spiral galaxies and the mathematical methods used to discover these relationships. At the best of our knowledge, no finite thin-disk solutions are available, and this lack has hampered the study of the dynamics of spiral galaxies. Solutions on the $z=0$ plane exist for the exponential disk and for the untruncated Mestel disk \\citep{bin08} and these have been most frequently used to model spiral galaxies. A major problem is that these disks of infinite extent can over-predict the total mass and mass surface density associated with a measured velocity profile.  \\cite{sch09} recently presented the complete closed form solution of the Maclaurin disk and a few related disks over all the space.  The Maclaurin disk is applicable to the study of spiral galaxies for which the rotational velocity rises to the edge of the disk as is seen in irregular galaxies.  Here we present the explicit expressions for the potential and gravitational field  of the finite Mestel disk (hereafter FMD) which was first introduced by \\cite{mes63}.  Until now the force vectors of the FMD were only known for points in the plane of the disk, and there was no analytical solution of the potential.  Throughout this work we use the notation that $M$ is the total disk mass, $\\alpha$ is the disk radius, $G$ is the gravitational constant, and $(R,z)$ are the cylindrical co-ordinates in each meridional plane. ",
        "conclusions": "We have presented the complete closed form solution in cylindrical coordinates for the gravitational potential and force field of the FMD. The mass distribution and the rotation curve of the FMD match those of typical spiral galaxies and so this work provide a useful model for the study of the properties of disk galaxies.  The FMD is defined by two parameters, disk mass and disk diameter, in contrast to the commonly used  maximum disc solution which permits many free parameters.  This simpler description will allow us to standardize the characterization of populations of spiral galaxies which will aid in the study of, e.g., relations such as the baryonic Tully-Fisher relation or the Magnitude-Diameter relation."
    },
    "1112/1112.6378_arXiv.txt": {
        "abstract": "Model-independent parametrisations for examining departures from General Relativity have been increasingly studied over the past few years. Various observables have been used to constrain the parameters and forecasts for future surveys have been carried out. In one such forecast \\cite{paper2}, galaxy cluster counts were used to constrain the parameters. Here, we carry out a limited set of $N$-body simulations, with a modified Poisson equation, to examine the accuracy of existing mass functions for modified gravity cosmologies. As well as altering the gravitational calculation, we include the effect of a screening scale to ensure consistency of the theory with solar system tests. Our results suggest that if a screening scale exists its effect can be taken into account in the cluster count calculation through its effect on the linear matter power spectrum. If this is done, the accuracy of the standard mass function formalism in modified gravity theories with reasonably small departures from General Relativity, as tested in this work, is comparable to the standard case. ",
        "introduction": "Several explanations have been put forward to explain the apparent accelerating cosmological expansion. The most popular idea is that the acceleration is due to a distinct form of energy in the universe, generically known as dark energy. The simplest example of what could constitute dark energy is a cosmological constant $\\Lambda$, a non-dynamical form of energy that may be interpreted as arising as an unconstrained integration constant in the equations of motion or from the vacuum energy generated by the quantised fields that make up the matter content of the universe. Dynamical forms of dark energy involving scalar fields that are evolving with time can also lead to late-time acceleration. Some have also argued that the observation that the expansion is accelerating results from the misuse of the Cosmological Principle in interpreting the expansion history from the inhomogeneous distribution of matter around us. An alternative perspective to the above both lines of thought above is to explain the departure from the expected late-time behaviour by introducing modifications to the theory of General Relativity (GR) itself.  There has been an increased interest in modified gravity theories in cosmology over the past decade or so. The large number of proposed alternatives to General Relativity (GR) has led to interest in examining possible deviations from GR in a more model independent way. Much work has gone into creating these parametrisations, examining their consistency, constraining them with current data and forecasting future experiments (see \\cite{mgreview} for a comprehensive review). Many different observations and combinations of observations have been considered for constraining these parametrisations. The key point with the combinations is the issue of distinct degeneracies in different sets of observations. In particular, observations such as the integrated Sachs-Wolfe effect in the Cosmic Microwave Background (CMB) and cosmic shear surveys, both being effects on the energies and trajectories of photons that are relativistic, constrain the sum of the two scalar potentials in the metric $\\Psi + \\Phi$ and thus constrain a particular combination of the Modified Gravity Parameters (MGPs) as will be introduced later. This degeneracy is best broken with an observation that depends on a different linear combination of the scalar potentials. Observations relating to the growth of structure in non-relativistic matter, such as redshift space distortions or galaxy surveys, are a natural choice for breaking the degeneracy. Recently \\cite{paper2}, we argued that combining future cluster counts with CMB and cosmic shear observations could yield high precision constraints on the simplest form of MGPs particularly in the case where the modifications are constant at low redshift.  Galaxy cluster counts can be predicted for a particular combination of cosmological parameters by starting from the linear matter power spectrum and using mass functions to calculate how many bound structure will form at different mass scales. The mass functions are semi analytic formulas whose form derives from the original Press-Schechter ideas \\cite{ps} and its extensions \\cite{bond91,st99a,st99b,st02}. The functions are now calibrated using detailed and extensive $N$-body simulations. In this \\textit{paper}, we present the results of a limited set of of $N$-body simulations aimed at investigating the accuracy of existing mass functions for the MGP formalism. Due to the computational demand of running the required $N$-body simulations we have restricted our investigation to a single point in the phase space of MGPs i.e. a single choice in the modification of the strength of gravity that is constant at low redshift. As such we are not attempting to accurately calibrate an extension of the mass function formalism in the phase space of MGPs. We are instead addressing the question of whether the existing mass function formalism can be used {\\sl as is} in forecasts involving cluster counts such as that carried out in \\cite{paper2}.  This {\\sl paper} is organised as follows. In section~\\ref{sec:massf} we discuss the mass function formalism and some of the results in the literature regarding modified gravity models. The modifications required to run $N$-body simulation in such models of parametrised modified gravity are summarised in section~\\ref{sec:nbody} and the results of running our suite of simulations are described in section~\\ref{sec:res}. We conclude with a discussion of the results in section~\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  We have shown that the mass function formalism can work for simple phenomenological models of modified gravity without a screening scale. To model modified gravity with a screening scale, its effect needs to be included in the calculation of the linear matter power spectrum in order to retain the accuracy of the mass function formalism. Of course, the mass function formalism is likely to break down for more extreme departures from GR, however it isn't obvious that large deviations are allowed by current data. In addition, more complicated modelling of the $\\mu$ parameter, such as including scale dependence may reduce the accuracy of the mass function predictions. However, as long as the scale dependence manifests in the linear matter power spectrum and the deviations from GR are relatively small then the mass function prediction may still be accurate. This will require further work. The work presented here needs to be extended in other ways as well, to more precisely determine the range of validity of the mass function formalism. Varying the cosmological parameters, linking length and parametrisation of $\\mu$ will further test the validity of the mass functions. In addition, running bigger boxes and boxes with improved resolution will test the mass function over a larger range of cluster masses. The mass function should also be tested over a range of redshifts and with baryons included, rather than solely dark matter particles. In addition, an examination of spherical collapse in model independent modified gravity and/or an examination of virialisation in these theories may shed further light on finding halos and how mass functions work in modified gravity simulations."
    },
    "1112/1112.0418_arXiv.txt": {
        "abstract": "We show that the cosmological moduli problem is solved, without relying on huge late-time entropy production, if the universal cutoff scale of the theory is a few orders of magnitude smaller than the Planck scale. We obtain a general estimate of the modulus abundance in terms of the inflationary scale and the reheating temperature, and find in particular that the reheating temperature can be high enough for the  non-thermal leptogenesis to work. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.1839_arXiv.txt": {
        "abstract": "We report the detection of an orbital period decay of $\\dot P=-1.83\\pm0.66$~ms~yr$^{-1}$ in the black hole X-ray binary \\mbox{XTE J1118+480}. This corresponds to a period change of $-0.85\\pm0.30$~$\\mu $s per orbital cycle, which is $\\sim 150$~times larger than expected from the emission of gravitational waves. These observations cannot be reproduced by conventional models of magnetic braking even when including significant mass loss from the system. The spiral-in of the star is either driven by magnetic braking under extremely high magnetic fields in the secondary star or by a currently unknown process, which will have an impact on the evolution and lifetime of black hole X-ray binaries. ",
        "introduction": "Angular momentum loss (AML) in short-period black hole X-ray binaries are driven essentially by magnetic braking~\\citep{ver81}, gravitational radiation~\\citep{lan62,tay82}, and mass loss~\\citep{rap82}, with possible additional contributions from jets~\\citep{kin99}. Measurements of orbital period variations with time constrain the role of these processes in the evolution of such binaries~\\citep{ver93}.  Magnetic braking is proposed as the main mechanism for AML in compact binaries~\\citep{rap83}. It produces the shrinking of the binary orbit and maintains the donor star in contact with its Roche lobe, therefore, sustaining stable mass transfer. Although this is an essential mechanism to understand compact binary evolution its actual prescription is currently debated and not well established~\\citep{iva06}.  \\mbox{XTE J1118+480} is one of only 18 galactic X-ray binaries which contain a dynamical black hole~\\citep{cas07}. With an orbital period of 4.1 hours it is the most compact black hole X-ray binary known. Determination of the times at the inferior conjunction of the $\\sim 0.2$~\\Msun secondary star with respect to the $\\sim 8$~\\Msun black hole have been obtained at different epochs since its discovery on UT 2000 March 29 by the Rossi X-ray Timing Explorer~\\citep{rem00}. Previous attempts to estimate the orbital period derivative failed due to the small baseline of the observations~\\citep{joh09a}. ",
        "conclusions": "}  General Relativity~\\citep{mis73} predicts a secular periastron precesion period of $\\sim 23$~yr for \\mbox{XTE J1118+480}. In the case of non-zero orbital eccentricity, this would affect the period inferred from succesive inferior conjunctions ~\\citep[see e.g. eq.~20 in][]{pal08}. Using our data we set an upper-limit to the eccentricity of the black hole binary orbit of $e < 0.0067$ (95\\% confidence level), and conclude that relativistic periastron precesion could only explain up to 50\\% of the measured orbital period decay.  \\begin{deluxetable}{lccc} \\tabletypesize{\\scriptsize} \\tablecaption{Time at inferior conjunction of the secondary star in \\mbox{XTE J1118+480}.\\label{ttn}} \\tablewidth{0pt} \\tablehead{\\colhead{$N$} & \\colhead{$T_n-2450000$}\\tablenotemark{a} & \\colhead{$\\delta T_n$}& \\colhead{Reference}} \\startdata 0     & 1868.8916  & 0.0004 & [1] \\\\ 66    & 1880.1086  & 0.0004 & [2] \\\\ 904   & 2022.5122\\tablenotemark{b}  & 0.0004 & [3] \\\\ 6950  & 3049.93347 & 0.00007 & [4] \\\\ 21772 & 5568.6936  & 0.0003 & [5] \\\\ 21960 & 5600.6413  & 0.0002 & [5] \\\\ 22407 & 5676.6017  & 0.0002 & [5] \\\\ \\enddata \\tablenotetext{a}{Time at inferior conjunction, $T_n$, of the secondary star in the low-mass black hole X-ray binary \\mbox{XTE J1118+480}, and uncertainties, $\\delta T_n$, corresponding to a given orbital cycle number, $n$.} \\tablenotetext{b}{From photometric measurements.} \\tablecomments{References: [1]~\\citet{wag01}; [2]~\\citet{tor04}; [3] ~\\citet{zur03}; [4]~\\citet{gon08}; [5]~This work.} \\end{deluxetable}  The loss of energy through emission of quadrupole gravitational radiation, as predicted by general relativity, provides an orbital period derivative for this binary system of $\\dot P_{\\rm GR} \\sim -0.01$~ms~yr$^{-1}$ \\citep[from eq.~8 in][]{tay82} which is far below the observed rate of variation of the orbital period. On the other hand, the presence of jets in \\mbox{XTE J1118+480} has also been postulated based on X-ray and infrared observations~\\citep[see e.g.][]{mil02}. However, the consequential AML due to jets even for relativistic jets \\citep{kin99} is $\\leq 1000$ times lower than the AML due to magnetic braking. The presence of a circumbinary disk is also proposed as possible AML mechanism in cataclysmic variables \\citep[see e.g.][]{taa03}, and the detection of mid-IR 4.5--8~$\\mu$m excess emission~\\citep{mun06} do support its existence in \\mbox{XTE J1118+480}. However, the estimated total mass of this disk is too small to affect the orbital angular momentum of this X-ray binary. Therefore, these effects cannot account for any significant additional contribution to the apparently too large orbital period decay measured in this black hole binary.  Most conventional magnetic braking and mass loss models for the binary parameters of \\mbox{XTE J1118+480} predict a much lower orbital period decay than we measure (see Fig.~\\ref{fmbml}). Our observations appear inconsistent with these models. Only a very restrictive, rather unplausible subset of models, where nearly all the mass transferred by the secondary star is also lost by the system ($\\beta=-\\dot M_{\\rm BH}/ \\dot M_2\\sim0$), may lead to values of the orbital period derivative close to our observations (see Fig.~\\ref{fmbml}).  \\begin{figure}[ht!] \\centering \\includegraphics[width=8.5cm,angle=0]{apjlf3.eps} \\caption{\\scriptsize{Theoretical estimate of the orbital period derivative in the low-mass X-ray black hole binary \\mbox{XTE J1118+480} due to magnetic braking and mass loss, $\\dot P_{\\rm MB,ML}$ given in ms/yr, derived from the second term of eq.~28 in~\\citet{joh09b}}, as a function of two parameters: $\\beta=-\\dot M_{\\rm BH}/ \\dot M_2$, the fraction of mass lost by the secondary star that is captured by the black hole; and $\\gamma$, an index that characterizes the strength of the magnetic braking. The specific angular momentum, $j_w$, carried away by the mass lost from the system has been set to 1 which indicate that all the mass lost from the system is lost from the neighborhood of the black hole \\citep[or its accretion disk, see][]{pod02}. Dashed lines represent the value derived from observations and the dotted lines, the values at 1-$\\sigma$ ver the uncertainty on the orbital period derivative, $\\Delta\\dot P$.} \\label{fmbml} \\end{figure}  Magnetic braking could be enhanced by anomalously high magnetic fields in the secondary star. We have estimated, from eq.~5 in~\\citet{jus06} and assuming that the mass lost by wind is equal to the mass transfer rate (derived from eq.~9 in~\\citet{kin96}), that a magnetic field at the surface of the companion star $B_s\\ge 10-20$~kG is needed to explain the observations. The angular momentum loss due to magnetic braking is determined from eq.~9 in~\\citet{joh09b}. This magnetic field is 1--2 orders of magnitude larger than typical magnetic field strengths in highly-rotating low-mass stars~\\citep{pha09}. However, we cannot discard that the donor star in \\mbox{XTE J1118+480} descends from a magnetically peculiar Ap/Bp star that has retained most of its primordial magnetic field. In such stars very high magnetic fields of $\\sim 20$~kG have been observed~\\citep{elk10}. Binaries with intermediate-mass Ap/Bp stars have indeed been postulated as the progenitors of compact black hole X-ray binaries~\\citep{jus06}. In particular, magnetically coupled, irradiation driven stellar winds can lead to substantial loss of systemic angular momentum~\\citep{jus06} required to form low-mass X-ray binaries out from intermediate-mass binaries~\\citep{pod02}. The detection of CNO-processed material in \\mbox{XTE J1118+480} strongly suggests that the donor descends from an intermediate-mass star~\\citep{has02}. Our discovery of a large period decay in XTE J1118+480 is also consistent with this scenario although we cannot rule out the possibility that other mechanism, not yet identified, plays a major role in the loss of angular momentum and hence the evolution of black hole X-ray binaries. In any case, our observations suggests a faster evolution and shorter lifetimes than previously assumed. This would help to reconcile the population number of low-mass X-ray binaries and millisecond pulsars, a long-standing problem in galactic astronomy~\\citep{pod02}. Follow-up spectroscopy in the coming years may provide a determination of the second derivative of the orbital period with strong implications on our knowledge on the formation and evolution of black hole X-ray binaries.  The black hole in \\mbox{XTE J1118+480} was possibly formed in a violent supernova explosion that launched the system via an asymmetric natal kick~\\citep{gua05} from its formation region in the Galactic thin disk~\\citep{gon06} to its present location in the Galactic halo~\\citep{mir01}. A lower limit to the age of the system of $\\geq 11$~Myr was derived from its peculiar location and kinematics. This limit has been used to set constraints on the rate at which black holes  can evaporate in the Anti-de Sitter (AdS) braneworld Randall-Sundrum gravity model~\\citep[for details see the introduction in][]{joh09b,joh09a} via the emission of a large number of conformal field theory (CFT) modes~\\citep{emp03}.  An upper-limit to the asymptotic AdS curvature radius in the extra dimensions was established at $L\\leq 80$~$\\mu m$~\\citep{psa07}, restricting deviations from the gravitational inverse square law to manifest only at distances smaller than L.  The orbital period evolution in this system was studied before but with a less number of measurements of $T_n$ spread in a shorter interval of time, thus providing a period derivative consistent with zero at 1-$\\sigma$ \\citep{joh09a}, and an upper-limit of $L\\leq 97$~$\\mu m$. In their eq.~2, these authors assumed no angular momentum loss due to mass lost from the system ($j_w=0$), no accretion onto the black hole ($\\beta=0$), and the parameter $\\gamma$, which governs the strength of magnetic braking, equal to zero \\citep[see also Fig.~1 in][]{joh09a}. If we adopt the black hole and secondary masses given in Table~\\ref{tpar}, our determination of the orbital period decay provides a much tighter constraint on the asymptotic AdS curvature radius of $L\\leq 35$~$\\mu m$ at 2-$\\sigma$. This limit is more restricted than the best current table-top experiment upper-limit of $L\\leq 44$~$\\mu m$ ~\\citep{kap07}. We note here that the size of the extra dimensions has been also recently constrained from the age of a black hole in a extragalactic globular cluster, placing an upper-limit as low as $L\\leq 3$~$\\mu m$ ~\\citep[see][]{gne09}."
    },
    "1112/1112.6244_arXiv.txt": {
        "abstract": "The past 15 years have seen an explosion in the number of redshifts recovered via wide area spectroscopic surveys. At the current time there are approximately 2million spectroscopic galaxy redshifts known (and rising) which represents an extraordinary growth since the pioneering work of Marc Davis and John Huchra. Similarly there has been a parallel explosion in wavelength coverage with imaging surveys progressing from single band, to multi-band, to truly multiwavelength or pan-chromatic involving the coordination of multiple facilities. With these empirically motivated studies has come a wealth of new discoveries impacting almost all areas of astrophysics. Today individual surveys, as best demonstrated by the Sloan Digital Sky Survey, now rank shoulder-to-shoulder alongside major facilities. In the coming years this trend is set to continue as we being the process of designing and conducting the next generation of spectroscopic surveys supported by multi-facility wavelength coverage. ",
        "introduction": "This article briefly summarises the development of wide field spectroscopic survey programs in extragalactic astronomy (see Fig.~1, Table.~1), and their impact as measured in publications and citations (see Table.~2). The statistics shown in Table.~2 are necessarily crude but indicative of the major revolution which is overtaking our subject --- the rise of surveys over facilities. Whereas once the key question senior astronomers might ask of each other seemed to be ``so what are you building?'', in this age of internationally funded mega-facilities, the more pertinent question might be ``so what survey are you designing?''. The likely answer will be one which combines both spectroscopic information with panchromatic imaging. Traditionally imaging surveys, as pioneered by the various Schmidt surveys, were single bandpass (${b_J}$) and then later double bandpass ($b_J,r_F$). These were followed by a move to multi-band programs (e.g., $BVRI$, $ugriz$, $YJHK$) which used multiple filters to fill in the wavelength gaps but still remain single facility imaging surveys. Truly multiwavelength surveys span more than one imaging facility, and panchromatic surveys span multiple facilities covering a significant fraction of the em-spectrum (e.g., GOODS; Giavalisco et al.,~2004; GAMA; Driver et al.,~2011). This trend, towards broader spectral coverage, is driven not only by technological advancement but also by an appreciation that galaxies in particular emit significant levels of radiation at almost all wavelengths and to truly understand their nature will require robust distances combined with total energy measurements.  \\begin{figure}[h]  \\vspace*{-3.5 cm}  \\includegraphics[width=5.0in]{plot.ps}  \\caption{The rise in galaxy redshifts over the past 30 years.} \\label{fig1}  \\end{figure} ",
        "conclusions": ""
    },
    "1112/1112.0979_arXiv.txt": {
        "abstract": "About half of the hot subdwarfs are found to have spectra of composite types, indicating a main sequence companion of spectral type F--K, and the pulsators are no exception to this rule. The spectroscopic contamination from the main sequence stars makes it hard to reliably establish physical parameters for the hot component, and also makes pulsations harder to detect as the amplitudes are depressed. The binary fraction of the observed sample of hot subdwarf pulsators is discussed, as are the biases that are affecting it. Spectroscopic evidence is presented that clearly demonstrates that the well known sdB pulsator, EO\\,Ceti, is misclassified, and is actually an sdOV star. ",
        "introduction": "Already when the discovery of pulsations in the hot subdwarf B (sdB) stars were announced in the first papers by \\citet{kilkenny97}, \\citet{koen97} and \\citet{stobie97}, a strong binary fraction was evident, since all these first pulsators are in composite systems. At the time it was even speculated that the binarity is somehow involved in inducing the pulsations in these stars. Much has happened since then. It is now clear that about 50\\%\\ of sdB stars are found to be in close binaries with periods on the order of hours or days \\citep{copperwheat11}. The composite spectrum binaries are not included among those. The sdB stars with main sequence F-, G- or K-type companions are known to form through stable Roche-lobe overflow rather than common envelope ejection, and should end up with rather long periods. As recently demonstrated from {\\sc hermes} observations (\\O stensen \\& Van Winckel, these proceedings), the periods of such binaries are indeed exceedingly long, between one and three years. It has been demonstrated by \\citet{reed04}, by computing the IR excess of the known subdwarfs in the field, that about half of sdB stars are such composite spectrum binaries. Taken together, these results would seem to account for $\\sim$100\\%\\ of the sdBs.  When the South African group (whose most prominent member, Dave Kilkenny, we are celebratinig the 65th birthday of with this conference) initiated the rush of discoveries of pulsating sdB stars, the sample they observed was fairly unbiased with respect to binarity. The only stars excluded would have been those with main sequence companions of type A, and possibly the occasional (sub-)giant K or M stars, where the sdB star would be practically invisible in the optical. But when the collaboration that I have participated in joined the fray with observations from the Nordic Optical Telescope ({\\sc not}), we cheated a bit, using Uli Heber's extensive library of spectra and model grids to select targets most likely to be pulsators. Thus, we could increase our detection efficiency from 1.7\\%\\ to $\\sim$10\\% \\citep{sdbnot}. However, in the process we automatically selected against composite spectrum binaries, as the contamination makes the temperature determination very uncertain, especially in low-res spectra with modest S/N. This fact was somewhat alleviated by the fact that we pretty soon ran out of targets with well-determined physical parameters, and by the time we started observing the SDSS sample we were less concerned about the contamination. Still, the consequence of these selection effects is that any attempt to check if the fraction of binaries in the observed sample of sdB pulsators is the same as in the population as a whole will be biased towards non-composite systems. Revisiting the 49 stars compiled in Table~9 of \\citet{sdbnot} reveals 18 stars that are composite, i.e.~a fraction of 37\\%, which is close enough to the binary fraction of the field population to be consistent within 2-$\\sigma$ given the sample size.  We list the 18 composites in Table~1.  \\begin{table}[t] \\caption{Stars classified as sdB with a main sequence companion. The column labeled 'Disc.' identifies the target roughly in order of discovery (line number in Table~9 of \\citealt{sdbnot}), and the group responsible for the discovery. }\\smallskip \\begin{center}\\small \\begin{tabular}{llllll}\\tableline\\noalign{\\smallskip} Star name & Survey name & Class & $m_V$ & Disc. & References \\\\ \\noalign{\\smallskip}\\tableline\\noalign{\\smallskip} V361\\,Hya & EC\\,14026--2647 & sdB+G2 & 15.3 & 1/SA    & \\citet{kilkenny97} \\\\ EO\\,Ceti  & PB\\,8783        & sdB+F0 & 12.3 & 2/SA    & \\citet{koen97}, \\\\ &                 &        &      &         & \\citet{odonoghue98b} \\\\ UX\\,Sex   & EC\\,10228--0905 & sdB+G2 & 15.9 & 3/SA    & \\citet{stobie97} \\\\ LM\\,Dra   & PG\\,1618+563B   & sdB+F2$^1$ & 13.5 & 15/\\not & \\citet{silvotti00} \\\\ V387\\,Peg & HS\\,2151+0857   & sdOB+IR$^2$ & 16.5 & 20/\\not & \\citet{ostensen01b} \\\\ EP\\,Psc   & PG\\,2303+019    & sdB+IR$^2$ & 16.2 & 24/\\not & \\citet{silvotti02a} \\\\ & PG\\,0048+091    & sdB+F5 & 14.3 & 28/SA & \\citet{koen04} \\\\ & PG\\,0154+192    & sdB+G/K & 15.3 & 29/SA & \\citet{koen04} \\\\ & J1717+5805      & sdOB+G9 & 17.4 & 30/\\not & \\citet{solheim04} \\\\ & PG\\,1419+081    & sdOB+F/G & 14.9 & 32/\\not & \\citet{solheim06} \\\\ & J1445+0002      & sdOB+G0 & 14.9 & 33/\\not & \\citet{solheim06} \\\\ & J1642+4252      & sdOB+G0 & 14.9 & 34/\\not & \\citet{solheim06} \\\\ & EC\\,11583--2708 & sdB+F/G & 14.9 & 36/SA & \\citet{kilkenny06} \\\\ & PG\\,1657+416    & sdOB+G5 & 15.9 & 39/\\not & \\citet{oreiro07} \\\\ & JL\\,166         & sdOB+F/G & 15.0 & 43/BB & \\citet{barlow09} \\\\ & HE\\,2151--1001  & sdB+F/G & 15.6 & 46/\\not & \\citet{sdbnot} \\\\ & PG\\,1033+201    & sdB+F9 & 15.4 & 48/\\not & \\citet{sdbnot} \\\\ & HE\\,1450-0957   & sdB+F9 & 15.3 & 49/\\not & \\citet{sdbnot} \\\\ \\noalign{\\smallskip}\\tableline \\end{tabular} \\end{center} \\small {$^1$ The companion is a visual binary, not an actual close companion.}\\\\ {$^2$ The companion is identified based on IR-excess, not on spectroscopic features.} \\end{table}  The sdB+MS binaries are in general much more poorly studied than the rest of the sdBV sample. The complication of having a strong spectroscopic companion means several things. First of all the spectroscopic parameter determination becomes much more complicated and regardless of the care taken in correcting for the contamination, the resulting spectroscopic parameters will always be more uncertain. For asteroseismology there can be a significant suppression effect in the photometric amplitudes, an effect that is also colour dependent, since the cool companion suppresses the amplitudes in the red more than in the blue. Note also that the companion may itself be pulsating; \\cite{ostensen11b} found that several of the composite systems observed with the {\\em Kepler} spacecraft shows photometric variability at the 1\\%\\ level with periods on order of days or longer, consistent with gamma Dor pulsations or rotational effects associated with spots.  I have long wanted to rectify the bias towards non-composite spectroscopic binaries among the sdBVs in order to explore the possible difference in internal structure in sdBs that form through different evolutionary channels, i.e.~common envelope ejection, stable mass transfer or mergers. This was also the overarching topic of the Ph.D.~thesis of Maja \\citet{MajaPhD}. During the sample selection process for the long period binary study with {\\sc hermes} (\\O stensen \\&\\ Van Winckel, these proceedings), I was hoping to include some sdBVs in the sample, but as can be seen from Table~1, only EO\\,Ceti is bright enough to be observed with {\\sc hermes}, and even it at $V$\\,=\\,12.3 is on the limit of what can be done. However, as I explored low-resolution spectra of EO\\,Ceti and other sdB+F/G binaries kindly made available to be by Betsy Green and Uli Heber, I found an unexpected surprise, and the rest of this talk will be dedicated to this discovery.  \\begin{figure*}[!ht] \\centering \\includegraphics[height=11.4cm,angle=-90]{EOCeti_WHTfit.eps} \\includegraphics[height=11.4cm,angle=-90]{EOCeti_WHTcut.eps} \\caption{ WHT/ISIS spectrum of EO\\,Ceti. Also shown are model spectra with parameters as stated on the plot. The spectrum was response corrected by calibrating the instrumental response function with a single sdB star. The models have been scaled to fit the spectrum at 4050 and 5060\\,\\AA. } \\end{figure*} ",
        "conclusions": "After re-examination of new high quality spectroscopy it is clear that EO Ceti is not an sdBV but an sdOV star. Thus, as is the case with most other important discoveries in the field of hot subdwarf stars, the first sdO pulsator was actually discovered by Dave Kilkenny and his team, which I think is a very appropriate discovery to announce at this conference held in his honour.  The clear detection of the He\\,{\\sc ii}, together with weak or absent He\\,{\\sc i} effectively forces EO\\,Ceti to be classified as an sdO star, regardless of its actual temperature. Using the slope of the spectrum is no longer a viable option at temperatures in the sdO range, unless one has access to UV spectroscopy. So far no attempt has been made to reliably constrain the surface gravity of the component, and the analysis does not force the two stars to be at the same distance. However, since no orbit has been established, that may not actually be the case. We will continue monitoring EO\\,Ceti with high resolution spectroscopy in the hope of detecting the orbital period of the system. That may allow us to use spectral disentangling to decompose the spectra into its two components, which will greatly simplify the parameter estimation. However, the broadening from the pulsational velocity at the surface of the sdO star and the broadening due to the rotational velocity of the F star are both higher than the expected orbital velocity, which makes this a very challenging project.  At B=12.3 EO\\,Ceti is a much more encouraging target to study than J16007+0748, since high-precision photometry can be obtained with modest telescopes, and time-resolved spectroscopy is a lot more feasible. In fact, thanks to its early discovery, EO\\,Ceti is one of the best studied hot subdwarf pulsators with extensive campaign observations by \\citet{odonoghue98b}, multicolor photometry by \\citet{vuckovic10} and an early attempt at time-resolved photometry by \\citet{jeffery00}.  Can there be other sdO pulsators in the sample of known sdBVs (Table 9 of \\citealt{sdbnot})? Perhaps. I have checked the ones that I have spectroscopy on, but no clear candidates have emerged. But there are quite a few in that table that have very poor spectroscopy, and many have no established helium abundance. The spectroscopy also suffer from having been obtained with many different instruments and various resolutions and signal levels. In many cases only the original classification spectrum exists, so clearly, much more work has yet to be done."
    },
    "1112/1112.2717_arXiv.txt": {
        "abstract": "Tidal interaction between a gaseous disk and a massive orbiting perturber is known to result in angular momentum exchange between them. Understanding astrophysical manifestations of this coupling such as gap opening by planets in protoplanetary disks or clearing of gas by binary supermassive black holes (SMBHs) embedded in accretion disks requires knowledge of the spatial distribution of the torque exerted on the disk by a perturber. Recent hydrodynamical simulations by Dong et al (2011) have shown evidence for the tidal torque density produced in a uniform disk to change sign at the radial separation of $\\approx 3.2$ scale heights from the perturber's orbit, in clear conflict with the previous studies. To clarify this issue we carry out a linear calculation of the disk-satellite interaction putting special emphasis on understanding the behavior of the perturbed fluid variables in physical space. Using analytical as well as numerical methods we confirm the reality of the negative torque density phenomenon and trace its origin to the overlap of Lindblad resonances in the vicinity of the perturber's orbit --- an effect not accounted for in previous studies. These results suggest that calculations of the gap and cavity opening in disks by planets and binary SMBHs should rely on more realistic torque density prescriptions than the ones used at present. ",
        "introduction": "\\label{sect:intro}   Understanding the response of a gaseous disk to a periodic non-axisymmetric gravitational perturbation is an important astrophysical problem. It arises in a variety of contexts such as the disk-planet interaction in protoplanetary disks, orbital evolution of supermassive black hole (SMBH) binaries embedded in circumbinary disks, dynamics of accretion disks in cataclismic variables, neutron star and black hole binaries, and so on.  Gravitational coupling between a gaseous disk and a massive orbiting perturber inevitably leads to the angular momentum exchange between the two. This has two important consequences for the evolution of the system. First, the orbit of the perturber can change leading to planet migration in protoplanetary disks (Ward 1986) and to the black hole inspiral in the case of SMBH binaries embedded in gaseous disks (Ivanov \\etal 1999; Gould \\& Rix 2000). Second, the angular momentum lost by the perturber gets absorbed by the disk, which causes redistribution of the disk surface density. In the most dramatic cases, when the perturber is massive enough, gas may be completely depleted in some parts of the disk, resulting in gap opening in protoplanetary disks (Ward 1997) and cavity formation around the SMBH binaries (Armitage \\& Natarajan 2002).  It is important to note that the latter process directly affects the former since the strength of tidal coupling between the disk and the perturber leading to migration depends on the spatial distribution of the disk surface density (Ward 1997). Thus, the change of the disk properties caused by the gravitational influence of a perturber has an immediate and important effect on the orbital evolution of the perturber itself.  Disk-planet\\footnote{From now on we will refer as ``disk-planet'' or ``disk-satellite'' to any other type of system where the tidal coupling is present, e.g. a SMBH binary surrounded by the disk, an accreting white dwarf-main sequence star binary, etc.} interaction leads to the excitation of nonaxisymmetric density waves, which carry the angular momentum lost by the perturber away from their excitation site. Subsequent dissipation of the waves due to some linear or nonlinear process transfers this angular momentum to the disk material, ultimately driving evolution of the disk. The torque density (torque per unit radial distance) exerted on the disk material by damping waves $dT/dr|_d$ (i.e. the amount of angular momentum transferred by dissipating waves to the disk fluid per unit time and per unit radial distance) can thus be represented as \\ba \\frac{dT}{dr}\\Bigg|_d={\\cal L}*\\left(\\frac{dT}{dr}\\right), \\label{eq:dTdr_d} \\ea where $dT/dr$ is the {\\it excitation torque density} --- the amount of angular momentum {\\it added} to the density wave per unit time and per unit radial distance by planetary tide. Operator ${\\cal L}$ describes the wave damping --- transfer of the wave angular momentum to the disk material by some dissipation mechanism. This operator is intrinsically nonlocal since the amount of angular momentum transferred by the wave to the disk at some location depends on the amount of angular momentum accumulated by the wave prior to that point.  Thus, to know how the tidal coupling affects the disk one must understand both the {\\it excitation} of the density waves (i.e. the radial dependence of $dT/dr$) and the details of their {\\it damping} (i.e. the explicit form of operator ${\\cal L}$ in equation (\\ref{eq:dTdr_d})). This important point was previously highlighted by Lunine \\& Stevenson (1983), Greenberg (1983), Goldreich \\& Nicholson (1989). The behavior of ${\\cal L}$ has been previously explored for linear viscous dissipation by Takeuchi \\etal (1996) and for nonlinear damping by Goodman \\& Rafikov (2001) and Rafikov (2002). The results of these studies suggest that for low-mass perturbers wave dissipation is a nonlocal process and most of the wave damping occurs far from its driving region.  In this work we concentrate on the wave excitation and do not consider its propagation and dissipation. Driving of the density wave is characterized by the dependence of torque density $dT/dr$ on $r$, or the radial distance from the planet. Spatial distribution of $dT/dr$ has been previously derived from the results of direct numerical simulations of the disk-planet interaction (Bate \\etal 2003; D'Angelo \\& Lubow 2008). At the same time there have been few analytical studies of the spatial behavior of the torque density. Linear theory of density wave excitation has predominantly concentrated, with few exceptions (Korycansky \\& Pollack 1993, hereafter KP93; Goodman \\& Rafikov 2001; Muto \\& Inutsuka 2009), on the behavior of fluid variables in Fourier rather than physical space.  Nevertheless, some asymptotic results regarding the behavior of $dT/dr$ have been known since the pioneering study of the disk-satellite interaction by Goldreich \\& Tremaine (1980; hereafter GT80). According to this work, far from the perturber, at radial separations from its orbit $|r-r_p|$ ($r_p$ is the semi-major axis of the circular orbit of the planet) exceeding the disk scale height $h$, the torque density\\footnote{We assume that GT80 calculation refers to $dT/dr$ and not to $dT/dr|_d$ since no explicit dissipation mechanism was mentioned in their work.} is given by \\ba \\frac{dT}{dr} & \\to & \\mbox{sgn}(r-r_p)C_{\\rm GT80} \\frac{(GM_p)^2\\Sigma_0}{\\Omega^2|r-r_p|^4}, \\nonumber\\\\ C_{\\rm GT80} &=& \\frac{32}{81}\\left[2K_0\\left(\\frac{2}{3}\\right)+ K_1\\left(\\frac{2}{3}\\right)\\right]^2\\approx 2.50783. \\label{eq:dTdx_GT} \\ea Here $M_p$ is the planetary mass, $\\Sigma_0$ is the disk surface density assumed to be uniform on scales $\\sim |r-r_p|$, $\\Omega$ is the angular frequency of the disk at $r_p$, and $K_n$ is the modified Bessel function of order $n$.  To obtain this expression GT80 have computed the amplitudes of the torque produced by the individual azimuthal Fourier harmonics of the planetary potential and then assigned their action in physical space to the positions of the corresponding Lindblad resonances. With this prescription the one-sided (i.e. considering only $r>r_p$ or $r<r_p$) torque density maintains the same sign independent of the radial separation from the planetary orbit. Torque density behavior similar to that given by equation (\\ref{eq:dTdx_GT}) was also obtained by Lin \\& Papaloizou (1979) using impulse approximation.  The torque density prescription (\\ref{eq:dTdx_GT}) has been extensively used in studies of gap opening by planets (Lin \\& Papaloizou 1986; Trilling \\etal 1998; Bryden \\etal 1999; Armitage \\etal 2002; Varnier\\'e \\etal 2004; Crida \\etal 2006) and orbital evolution of the SMBH binaries surrounded by gaseous disks (Gould \\& Rix 2000; Armitage \\& Natarajan 2002; Lodato \\etal 2009; Chang \\etal 2010; Alexander \\etal 2011). Modifications of the prescription (\\ref{eq:dTdx_GT}) have sometimes been adopted e.g. with the proportionality coefficient different from $C_{\\rm GT80}$ (e.g. Papaloizou \\& Lin 1984; Armitage \\& Natarajan 2002). However, the key features of the equation (\\ref{eq:dTdx_GT}) namely the $|r-r_p|^{-4}$ dependence and the {\\it positive} sign of the asymptotic torque density for $r>r_p$ (and the negative sign for $r<r_p$) have rarely been questioned (cf. Crida \\etal 2006).  Recently Dong \\etal (2011) investigated the spatial behavior of the torque density using high-fidelity numerical simulations of the disk-planet interaction in the two-dimensional shearing sheet geometry. Quite unexpectedly, they found that the one-sided torque density does not maintain a constant sign but rather changes from positive to negative (for $r>r_p$) at $|r-r_p|\\approx 3.2 h$, in disagreement with the conclusions of GT80. Dong \\etal (2011) have shown this result to be quite robust and independent of the numerical issues (e.g. resolution of the simulations, simulation box size, etc.).  In this work we demonstrate that some of the linear results of GT80 need to be revised at the qualitative level even within the framework of the linear theory. In particular, we show that the torque density does indeed {\\it change sign} and becomes negative for $r>r_p$ (positive for $r<r_p$) beyond several scale heights in radial separation $|r-r_p|$ from the planet, in full agreement with the numerical results of Dong \\etal (2011). This has important implications for understanding the issues of the gap and cavity opening by planets in protoplanetary disks and SMBH binaries in circumbinary disks correspondingly.  Our paper is structured as follows. In \\S \\ref{sect:eqs} we summarize the linear equations for the perturbed fluid variables, which are subsequently solved numerically in \\S \\ref{sect:harmonics}. We demonstrate good agreement of the results with the analytical calculations based on the global Airy representation developed in \\S \\ref{subsect:Airy_rep}. We compute the spatial behavior of the torque density and the angular momentum flux in \\S \\ref{sect:torque} and \\S \\ref{sect:flux} correspondingly . The main result of this paper --- confirmation of the negative torque density phenomenon --- is described in \\S \\ref{subsect:neg_torque}. We discuss the origin of the negative torque density, performance of the Airy representation, and compare our results with existing studies in \\S \\ref{sect:disc}. Astrophysical implications of our results are outlined in \\S \\ref{sect:implications}. ",
        "conclusions": "\\label{sect:disc}    Most of the existing linear calculations of the disk-satellite coupling have been limited to the Fourier space (Ward 1986; Artymowicz 1993; Tanaka \\etal 2002). Notable exceptions include Meyer-Vernet \\& Sicardy (1987) who discussed the radial structure of the perturbed velocity eigenfunctions for a fixed azimuthal wavenumber under a variety of assumptions about the physics operating in the disk. Later on KP93 have done similar calculation for a disk with non-zero pressure (the same setup as used in our current work) and computed the run of $F_{H,k}$ as a function of $x$. But these studies have only looked at the spatial behavior of {\\it individual} modes driven by the planetary potential.  More recently Goodman \\& Rafikov (2001) have integrated the contributions of separate modes over the azimuthal wavenumber to obtain the full spatial distribution of the perturbed surface density. Our current work extends the study Goodman \\& Rafikov (2001) by computing the spatial behavior of other important fluid variables such as the torque density and the angular momentum flux.  The key result of our present study is the confirmation of the negative torque density phenomenon (discovered in the numerical study of Dong \\etal (2011)) in the framework of the linear theory of disk-satellite interaction, and the determination of the asymptotic behavior of $dT/dx$ far from the planet. The discrepancy between our result (\\ref{eq:dTdx_full}) and the conventional expression (\\ref{eq:dTdx_GT}) not just in amplitude of the effect but also in its sign quite naturally demands an explanation, which we provide below.    \\subsection{Origin of the negative torque density.} \\label{sect:origin}   All calculations of the torque and angular momentum flux excited by the planet in GT80 have been carried out in Fourier space for individual harmonics of the planetary potential. These results fully agree with our calculations in \\S\\S \\ref{sect:torque} \\& \\ref{sect:flux}, in particular we successfully reproduce their expression for $F_{H,k}$ in \\S \\ref{sect:flux}. However, GT80 also tried to convert their results on torque exerted by the planet to physical (coordinate) space, and in doing this they had to resort to certain assumptions. In particular, they postulated that the action of each potential harmonic, including the deposition of a corresponding torque contribution, is localized in the vicinity of a respective Lindblad resonance. Under this assumption they computed the torque density at each location in physical space as the simple product of the spatial density of Lindblad resonances $|dk_y/dx|$ and the amplitude of the torque component in Fourier space $T_k\\equiv\\int_{-\\infty}^\\infty (dT/dx)_k dx$, corresponding to $k_y$ with Lindblad resonance at this particular location (i.e. at $k_y=2/(3x)$). Equation (\\ref{eq:dTdx_GT}) represents the result of such a calculation in the limit $x/h\\gg 1$. Not surprisingly, one finds the sign of $(dT/dx)^{\\rm LR}$ to be the same for any $x$ simply because the Fourier amplitudes of the torque have the same sign independent of $m$ --- there is no disagreement\\footnote{Analytical calculation of $T_k$ in the limit $k_y h\\lesssim 1$ using equation (\\ref{eq:T_Hk}) and Airy representation (\\ref{eq:Imv}), (\\ref{eq:I}) for $\\Im(v)$ results in the expression analogous to equation (13) of GT80.} on this issue between us and GT80.  What we find in this work is that the Fourier contribution to the torque density of each azimuthal harmonic of the potential $(dT/dx)_k$ is in fact {\\it spread out} over a finite portion of the disk, rather than localized in a narrow region around each Lindblad resonance. In this situation, the overlap of many Fourier modes at a given point in space is inevitable and the calculation of the full torque density in real space represents a rather challenging task, requiring direct integration of the equation (\\ref{eq:T_Hk}). As a result of carrying out this exercise, one finds the distribution of the torque density different from (\\ref{eq:dTdx_GT}) at large values of $x$, as clearly demonstrated in Figure \\ref{fig:neg_torque}a.  This reasoning is best illustrated if we temporarily abandon the shearing sheet approximation and consider the disk-satellite interaction in the full cylindrical geometry. Then, as pointed out by GT80, coupling between the planetary potential and the disk occurs at discrete Lindblad resonances radially separated from the planetary orbit by $x_{L,m}\\approx 2r_p/(3m)$, where $m=k_yr_p$ in the azimuthal wavenumber of the potential harmonic and an assumption $m\\gtrsim 1$ has been made. The distance between the adjacent $m$-th and $(m+1)$-th resonances is $\\Delta x_{L,m}=x_{L,m}/m=x_{L,m}(h/r_p)(k_y h)^{-1}$.  We need to compare this distance to the width $\\delta x_m$ of the $m$-th resonance. According to Airy representation (see \\S \\ref{subsect:Airy_rep} and Appendix \\ref{sect:Airy}) the latter is set by the condition $t\\sim 1$ or $2a^{2/3}(\\xi-1)\\sim 1$ (valid near the Lindblad resonance), see equations (\\ref{eq:defs1}) and (\\ref{eq:near_peak}). One then finds that the resonance width $\\delta x_m$ is given by $\\delta x_m\\sim x_{L,m}a^{-2/3}\\sim x_{L,m}(k_y h)^{2/3}$ for $k_y h\\lesssim 1$. Note that $\\delta x_m\\ll x_{L,m}$.  A somewhat more illuminating way of deriving $\\delta x_m$ is based on the dispersion relation for the tightly wound density waves (Binney \\& Tremaine 2008), which is valid for $m\\gtrsim 1$ or $k_yh\\gtrsim (h/r)$: \\ba \\kappa^2+k_r^2c_s^2=m^2(\\Omega-\\Omega_p)^2\\approx \\frac{4m^2A^2}{r_p^2}x^2, \\label{eq:disp_rel} \\ea where $k_r$ is the radial wavenumber of the wave. The exact resonance location $x_{L,m}=r_pm^{-1}(\\kappa/2|A|)$ corresponds to the distance where $k_r=0$ and the wave couples best to the external potential. As the radial separation from the resonance increases $k_r$ grows until at $|x-x_{L,m}|\\sim\\delta x_m$ one gets $k_r\\delta x_m\\sim 1$. Beyond this point the wave eigenfunction starts oscillating rapidly, and coupling to the external potential becomes less effective. Plugging $k_r\\sim \\delta x_m^{-1}$ and $x=x_{L,m}+\\delta x_m$ into the dispersion relation (\\ref{eq:disp_rel}) and expanding it in terms of $\\delta x_m\\ll x_{L,m}$ one finds $\\delta x_m\\sim h(k_yh)^{-1/3}$, the same expression as that as resulting from the Airy representation.  Using these estimates of $\\Delta x_{L,m}$ and $\\delta x_m$ one can easily see that $\\delta x_m\\gtrsim \\Delta x_{L,m}$ and resonances located at distance $x$ from the perturber overlap as long as \\ba x\\lesssim h\\left(\\frac{r_p}{h}\\right)^{3/5}= r_p\\left(\\frac{h}{r_p}\\right)^{2/5}. \\label{eq:overlap} \\ea Equivalently, overlap occurs for modes having \\ba k_yh\\gtrsim \\left(\\frac{h}{r_p}\\right)^{3/5} ~~~~\\mbox{or} ~~~~m\\gtrsim \\left(\\frac{r_p}{h}\\right)^{2/5}. \\label{eq:overlap_ky} \\ea Interference of different modes takes place even though the width of each resonance $\\delta x_m$ is small compared to the resonance separation $x_{L,m}$ from the perturber's orbit.  Thus, whenever $r_p/h\\gg 1$, which is always the case for protoplanetary and many other kinds of accretion disks, resonances overlap in the vicinity of the perturber's orbit (for $x$ satisfying equation (\\ref{eq:overlap})) and the assumption of well-isolated resonances adopted by GT80 for their torque density calculation fails. Given that in our calculation the negative values of $dT/dx$ appear only for $x>x_-\\approx 3.2h$ (see \\S \\ref{subsect:neg_torque}) one needs the inverse of the disk aspect ratio to be $r_p/h\\gtrsim (x_-/h)^{5/3}\\approx 7$ for the negative $dT/dx$ to appear. This requirement is easily fulfilled in a large variety of astrophysical accretion disks, including protoplanetary disks.  In the framework of the shearing sheet approximation, which corresponds to the limit $r_p/h\\to \\infty$, it is not surprising, according to equation (\\ref{eq:overlap}), that resonances overlap for {\\it all} $x$ and the isolated resonance approximation adopted in GT80 never holds. In this limit our asymptotic expression (\\ref{eq:dTdx_full}) works fine for arbitrary $x$.  At the same time, far from the perturber's orbit, at separations not obeying the condition (\\ref{eq:overlap}) resonances become sufficiently spread apart to act in isolation. In this case one would expect the GT80's expression (\\ref{eq:dTdx_GT}) to be valid. Thus, when the full cylindrical geometry of the disk-planet interaction is accounted for $dT/dx$ at fist becomes negative close to the planet, for $x>x_-$, but then should change sign and become positive again (outside of the region defined by the inequality (\\ref{eq:overlap})) in accordance with the GT80 calculation. The details of this transition can only be elucidated via linear calculation of the disk-planet coupling in the cylindrical geometry, which is beyond the scope of this work.    \\subsection{Comparison with previous work.} \\label{sect:previous}   Quite interestingly, the existence of the negative torque density phenomenon could have been inferred from calculations done even prior to the work of Dong \\etal (2011). In particular, Bate \\etal (2003) performed global three-dimensional simulations of the disk-planet interaction in cylindrical geometry. As a part of this calculation they have computed torque density produced by planets of different masses. Close inspection of their Figure 12 shows that at small masses, $M_p\\lesssim 0.03$ M$_J$, the torque density changes sign at a finite separation from the planet. This effect does not disappear as $M_p$ is reduced strongly suggesting that it is a linear phenomenon.  More recently D'Angelo \\& Lubow (2008) have computed radial run of $dT/dx$ as one of the by-products of their three-dimensional global simulations. In their Figure 7 one can again see the torque density changing sign for low enough $M_p\\lesssim 10$ M$_\\oplus$, which is in good agreement with the torque density behavior that could have been inferred from the work of Bate \\etal (2003). Similar torque density pattern can also be seen in D'Angelo \\& Lubow (2010).  Finally, Muto \\& Inutsuka (2009) have carried out a local linear analysis of interaction between a planet and the surrounding viscous disk. Even though their results are affected to some extent by the assumed non-zero value of viscosity (our study assumes viscosity to be zero), their torque density still exhibits negative values beyond $x\\approx 3h$ (see their Figure 5) even at the lowest values of viscosity, in agreement with our work.  The most likely reason for why the negative torque density phenomenon went unnoticed in previous studies lies in the small amplitude of this effect. Indeed, the total negative torque contribution obtained by integrating $dT/dx$ from $x_-\\approx 3.2h$ to $\\infty$ amounts to $\\lesssim 4\\%$ of the total integrated torque $T(\\infty)=\\int_0^\\infty(dT/dx)dx$. Thus, it is not surprising that this real effect was previously ignored, most likely on the basis of it being due to some numerical issues (poor resolution, small simulation box size, etc.).  Nevertheless, the apparent presence of the negative torque phenomenon in the aforementioned earlier works provides additional support for the robustness of the results presented in this paper. In particular, one may wonder whether our conclusions based on a linear calculation performed in two-dimensional shearing sheet geometry would hold when some of our assumptions --- local approximation, linearity of the perturbation, two-dimensional geometry --- are relaxed. In this regard we simply point out here that both Bate \\etal (2003) and D'Angelo \\& Lubow (2008) ran three-dimensional, global simulations for a variety of planetary masses and one can still discern the presence of the negative torque phenomenon in their results. This gives us reassurance in the ubiquity of this rather subtle but important (see \\S \\ref{sect:implications}) feature of the disk-planet interaction even in more sophisticated and realistic geometries.    \\subsection{Accuracy of the Airy representation.} \\label{sect:acc_Airy}   Another interesting result of our study concerns the accuracy with which the analytical Airy approximation described in \\S \\ref{subsect:Airy_rep} and Appendix \\ref{sect:Airy} characterizes the details of the disk-planet interaction. Approximation of the behavior of the perturbed fluid variables in terms of Airy functions has been used by many authors before (Artymowicz 1993; Ward 1986) but only to represent the behavior in the vicinity of Lindblad resonances. As such the Airy representation served only as a {\\it local} analytical tool for computing the azimuthal Fourier components of the density wave properties such as the angular momentum flux.  In our study we use Airy representation in a {\\it global} sense, to describe the spatial variation of the fluid variables for the values of $x$ not limited to the locations of the Lindblad resonances. This procedure works well for describing the behavior of the Fourier harmonics of the azimuthal velocity perturbation $v$ and of the torque density $dT/dx$ at {\\it any} $x$ as long as the wavenumber of a particular mode satisfies the condition $a\\gtrsim 1$, which according to equation (\\ref{eq:C2}) is true whenever $k_yh\\lesssim 1$ or $k_yh\\gtrsim 1$. Moreover, for $x\\gtrsim h$ this approximation is quite reliable even for $k_yh\\sim 1$. As a result the Airy representation successfully reproduces the asymptotic behavior of $dT/dx$ and $F_{H,k}$ as a function of $x$ in the limit $x/h\\gg 1$ as we have shown in \\S \\ref{subsect:neg_torque} and \\S \\ref{sect:flux}. And as Figures \\ref{fig:dsigma} and \\ref{fig:dTdx} demonstrate one can also use Airy representation to reproduce with decent accuracy the density wave properties in physical space.  These conclusions are in agreement with the results of Heinemann \\& Papaloizou (2011) who found that the WKBJ approximation valid whenever $a\\gtrsim 1$ works very well for describing the evolution of the spiral density waves excited by the turbulence in the disk.     \\label{sect:summary}   We have carried out linear calculation of the gravitational coupling between the uniform gaseous disk and an embedded massive perturber. Our work goes beyond similar earlier studies by putting emphasis on understanding the behavior of the perturbed fluid variables in real space as opposed to Fourier space. This allows us to address the nature of the recently discovered negative torque phenomenon and explain its origin fully in the framework of linear theory.  We come up with the global analytical representation of the spatial variation of the perturbed fluid variables in terms of the Airy functions, which is shown to work quite well for describing the properties of the planet-generated density waves both in Fourier and in real space. Using this approximation we derive the asymptotic behavior of the torque density far from the perturber and show that it matches the results of numerical calculations. Our linear calculations confirm that the torque density changes sign at a finite separation from the perturber in full agreement with the results of the direct hydrodynamical simulations. This provides an important correction to the calculations of GT80.  These results have broad implications for understanding tidal interaction of planets with circumstellar disks and of binary SMBHs with the circumbinary accretion disks. In particular, our calculations provide a pathway for deriving the spatial torque distributions in such disks, which is important for understanding the issues of gap opening by planets and gas clearing by binary SMBHs."
    },
    "1112/1112.0524_arXiv.txt": {
        "abstract": "Direct detection experiments are poised to detect dark matter in the form of weakly interacting massive particles (WIMPs). The signals expected in these experiments depend on the ultra-local WIMP density and velocity distribution. Firstly we review methods for modelling the dark matter distribution.  We then discuss observational determinations of the local dark matter density, circular speed and escape speed and the results of numerical simulations of Milky Way-like dark matter halos. In each case we highlight the uncertainties and assumptions made. We then overview the resulting uncertainties in the signals expected in direct detection experiments, specifically the energy, time and direction dependence of the event rate. Finally we conclude by discussing techniques for handling the astrophysical uncertainties when interpreting data from direct detection experiments. ",
        "introduction": "Weakly Interacting Massive Particles (WIMPs) are a promising dark matter candidate as they are generically produced in the early Universe with roughly the right density. Furthermore supersymmetry (SUSY) provides a well-motivated concrete WIMP candidate in the form of the lightest neutralino (for reviews see e.g. Refs.~\\refcite{jkg,bhs,dmbook}). WIMPs can be directly detected in the lab via their elastic scattering off target nuclei~\\cite{dd}. Numerous direct detection experiments are underway and they are probing the regions of WIMP mass-cross-section parameter space populated by SUSY models (see e.g. Ref.~\\refcite{theory}).   The realisation that uncertainties in the velocity distribution, $f({\\bf v})$, will affect the direct detection signals dates back to some of the earliest direct detection papers written in the 1980s (e.g. Ref.~\\refcite{dfs}). We first discuss the standard halo model and other approaches to modelling the Milky Way halo (Sec.~\\ref{modelling}). We then discuss what observations and simulations tell us about the dark matter distribution (Secs.~\\ref{astro1} and \\ref{astro2} respectively). We then examine the direct detection signals (Sec.~\\ref{signals}) and how the astrophysical uncertainties affect these signals (Sec.~\\ref{affect}). We conclude in Sec.~\\ref{strategies} by discussing strategies for handling the uncertainties, including both `astrophysics independent' approaches and parameterizing the WIMP speed distribution. ",
        "conclusions": "Direct detection event rate calculations often assume the standard halo model, with an isotropic Maxwellian velocity distribution with dispersion $\\sigma = \\sqrt{3/2} \\, v_{\\rm c} = 270 \\, {\\rm km \\, s}^{-1}$ and a local WIMP density $\\rho_{0}= 0.3 \\, {\\rm GeV} \\, {\\rm cm}^{-3}$. However it has long been realised~\\cite{dfs} that uncertainties in the speed distribution will affect the signals expected in experiments.  We have discussed the standard halo model and other approaches to halo modelling, and the assumptions behind them. We then reviewed observational determinations of quantities that are relevant to direct detection experiments, namely the local dark matter density, the local circular speed (which is related to the velocity dispersion by the Jeans equations) and the local escape speed. While the statistical errors on these quantities are often small the systematic errors, from uncertainties in the modelling of the Milky Way, can be significantly larger.  Next we turned our attention to high resolution, dark matter only simulations of the formation of Milky Way-like halos in a cosmological context. They find velocity distributions that deviate significantly from the standard Maxwellian, and have features at high speeds.  The effect of baryonic physics on the dark matter distribution is not yet well understood. Some recent simulations have found that a dark disc may be formed, however its properties (and even existence) are highly uncertain. Simulations resolve the dark matter distribution on scales many orders of magnitude larger than those probed by direct detection experiments. The latest results suggest that the ultra-local dark matter distribution is largely smooth, but some features may exist.  We then reviewed the resulting uncertainties in the energy, time and direction dependence of the energy spectrum, and hence constraints on, or measurements of, the WIMP parameters. The uncertainty in the local circular speed has a significant effect on the event rate and hence exclusion limits and determinations of the WIMP mass. The uncertainty in the local density leads directly to an uncertainty in measurements of, or constraints on, the cross-section. Only the high energy tail of the energy spectrum is particularly sensitive to the exact shape of the speed distribution. The annual modulation is far more sensitive to the shape of the velocity distribution; its amplitude and phase can change significantly. The anisotropy of the recoil rate is robust to changes in the velocity distribution, however high speed features can change the peak direction of high energy recoils, and hence provide a way of probing the formation of the Milky Way.    Finally we discussed techniques for handling the astrophysical uncertainties when analysing direct detection data. Astrophysics independent comparisons between different experiments can be made using the integral of the velocity distribution, $g(v_{\\rm min})$. This approach is extremely useful for assessing the compatibility of various experiments, but requires the WIMP mass as input. Parameterising the WIMP speed distribution, and using data from multiple experiments to jointly constrain the WIMP mass and cross-section and speed distribution is a promising approach, but the optimal form for the parameterisation is not yet known."
    },
    "1112/1112.5816_arXiv.txt": {
        "abstract": "We have investigated the pulsar \\psr\\ and its X-ray jet with XMM-Newton. Apart from the long X-ray jet which is almost perpendicular to the direction of proper motion, a putative extended feature at the pulsar position, which oriented in the opposite direction of the proper motion, is also suggested by this deep X-ray imaging. Non-detection of any coherent X-ray pulsation disfavors the magnetospheric origin of the X-rays observed from the position of \\psr\\ and hence suggest that the interpretation of pulsar wind nebula is more viable. We have also probed the origin of \\psr\\ and identified a runaway star, which possibly originated from the Cygnus OB9 assoication, as a candidate for the former binary companion of the neutron star's progenitor. ",
        "introduction": "While the signature of rotation-powered pulsars is their pulsed emission, bulk of the rotational energy is going to power a relativistic outflow of magnetohydrodynamic wind particles with Lorentz factors up to $\\gamma\\sim10^8$. Although pulsar winds cannot be observed directly, their presence can be inferred through their interaction with the interstellar surroundings. If the wind reaches a termination shock, a pulsar wind nebula (PWN) is formed. The term PWN had been originally referred to the shocked emission associated with young and powerful pulsars (e.g. Crab) where the wind is effectively confined by the magnetic fields of the surrounding supernova remnants and relativistic leptons reach high energy densities which make these nebulae usually luminous synchrotron sources in the radio and X-ray regimes.  A different class of PWNe which is associated with older and less powerful pulsars is bow-shock nebulae. These nebulae are not embedded in supernova remnants but interact directly with the interstellar medium (ISM). If the pulsar has a high velocity, the nebula is confined by the ram pressure of the ISM and forms the typical bow-shock structure. The bow-shock size is determined by the balance between the wind and the ISM ram pressure. Even in this old and less powerful pulsars, there is a significant fraction of the whole pulsar spin-down power which is going to feed the pulsar wind (see Bucciantini \\& Bandiera 2001).  Among all the pulsars with PWNe detected, \\psr\\ is one of the most interesting system. This pulsar has a rotational period and a period derivative of $P\\sim0.68$~s and $\\dot{P}\\sim9.6\\times10^{-15}$~s~s$^{-1}$, respectively. This suggests \\psr\\ has a spin-down age of $1.13\\times10^{6}$ years. The spin parameters imply a spin-down luminosity and a surface dipole magnetic field strength of $\\dot{E}\\sim1.2\\times10^{33}$~erg~s$^{-1}$ and $B\\sim2.6\\times10^{12}$~G respectively. Concerning its distance, while the radio dispersion measure implies a value of 1.86 kpc (cf. Cordes \\& Lazio 2002), Chatterjee \\& Cordes (2004) found that by modeling the head of the shocked emission in H$\\alpha$ the pulsar distance could be as close as 1 kpc.  Hui \\& Becker (2007) have examined its X-ray properties with the Chandra data taken in October 2000. Brief results of the same observation have also been reported by Wong et al. (2003) and Zavlin \\& Pavlov (2004). Despite the photon statistics is limited in this data set, the spectral analysis suggests that the bulk of the X-ray emission from \\psr\\ is non-thermal. An independent analysis that includes also the second epoch Chandra observations taken in August-October 2006 has confirmed the non-thermal nature of the X-rays detected at the pulsar position (Johnson \\& Wang 2010). Nevertheless, the large frame time (i.e. 3.2 sec) of these Chandra observations do not allow us to constrain the temporal properties of the pulsar. Therefore, with the existing data, we cannot determine whether the bulk of these observed X-rays are originated from the pulsar magnetosphere or from the PWN.  The proper motion of \\psr\\ is among the largest in the currently known pulsar population with a transverse velocity of $v=865(d/1~{\\rm kpc})$~km~s$^{-1}$ (Harrison, Lyne \\& Anderson 1993). Its shock interaction with the ISM produces the most well-known \\emph{Guitar Nebula} as observed in H$\\alpha$ (Cordes et al. 1993). Although no evidence for the X-ray emission can be found in the region of the H$\\alpha$ nebula, an X-ray jet apparently associated with \\psr\\ has been detected by Chandra (Hui \\& Becker 2007; Wong et al. 2003). The most peculiar fact of this jet is that its orientation deviates by $\\sim118^{\\circ}$ from the pulsar's proper motion direction. Its X-ray spectrum is also found to be non-thermal.  Kargaltsev \\& Pavlov (2008) have argued that this jet-like feature may be associated with a nearby point source instead of with \\psr. However, this possibility has been ruled out by a recent analysis of the multi-epoch Chandra observations. Johnson \\& Wang (2010) have found that the sharp leading edge have a proper motion consistent with that of the pulsar, which secures the association between the jet and \\psr. Different pictures have been proposed to explain this X-ray jet (e.g. Bandiera 2008; Johnson \\& Wang 2010). However, there is still no consensus on its nature.   To further investigate the nature of the X-rays from the position of \\psr\\ and its enigmatic jet, a dedicated follow-up observation with improved temporal resolution and photon statistic is needed. In this paper, we present a deep XMM-Newton observation in order to study this extreme system in details. In \\S2, we describe the details of the XMM-Newton observation and the data reduction. In \\S3 and \\S4, we present the results of the X-ray analysis and the search for the $\\gamma-$ray emission from this system respectively. Finally, we discuss the physical implications of the observed results and search for the birth place of this fast-moving neutron star in \\S5. ",
        "conclusions": "\\subsection{Nature of the X-ray emission from \\psr\\ \\& its jet}  We have presented a detailed investigation of the X-ray emission from \\psr\\ and its jet with XMM-Newton. The non-detection of any coherent X-ray pulsation suggests that the magnetospheric origin is questionable. In the context of outer-gap model, the non-thermal pulsed X-rays are resulted from the back-flowing charge particles from the outer-gap (e.g. Cheng \\& Zhang 1999). Based on this model, the electromagnetic cascade developed in this flow should give rise to an X-ray photon index $\\leq2$. The photon index of the X-rays from \\psr\\ is $2.2^{+0.2}_{-0.3}$. The best-fit value is found to larger than that expected from the magnetospheric model, though they can be reconciled within $1\\sigma$ uncertainties.  On the other hand, with the deep X-ray imaging of \\psr, we have revealed an interesting feature at the pulsar position which is elongated along direction of proper motion as observed in different cameras (cf. insets in Fig.~\\ref{xmm_img}). Although we cannot confirm its extent unambiguously with the existing data, its orientation is interesting and worth discussing. X-ray tails have been found to be associated with many pulsars (see Kargaltsev \\& Pavlov 2010 for a recent review). The tails are usually oriented in the opposite direction of pulsar motion and behind the bow-shock, therefore it is interpreted as the flow of particles coming out of the ram-pressure dominated PWN. In this scenario, the X-ray emission from \\psr\\ is originated from the synchrotron radiation from the wind particles accelerated at the termination shock. This is consistent with the unpulsed non-thermal emission inferred from this observation. For the synchrotron emitting leptons of Lorentz factor $\\gamma$ distribute as $N\\left(\\gamma\\right)\\propto\\gamma^{-p}$, the observed photon index of $\\Gamma\\sim2.2$ from the pulsar location is consistent with the value expected from the fast-cooling scenario, $\\Gamma=\\left(p+2\\right)/2$, for the electron index lies within a typical range of $p\\sim2-3$ (Cheng, Taam, \\& Wang 2006).  This putative feature, which has an extent of $\\sim8''$ in EPIC images, should be resolved by Chandra. Nevertheless, in reexamining the archival Chandra data, its extent cannot confirm. On the other hand, the photon distribution of \\psr\\ in the merged Chandra image is slightly deviated from a point source with the deformation also toward the direction behind the pulsar proper motion. Although the evidence for the extended source cannot be conclusive with the existing data, the compact feature suggested by the Chandra data can possibly be the portion of the highest surface brightness. Complex structures of bow-shock PWNe have been observed in other systems. For example, XMM-Newton observation of a nearby pulsar PSR~B1929+10 has revealed a long X-ray tail of several arc-minutes (Becker et al. 2006). Apart from confirming the detection of this long trail, a follow-up Chandra observation has further resolved a brighter but more compact feature with an extent of $\\sim10''$ around the pulsar (Hui \\& Becker 2008). For a further investigation of \\psr, a dedicated deep Chandra observation, which simultaneously provides a sub-arcsecond spatial resolution and a photon statistic comparable with this XMM-Newton observation, is the only way to confirm or refute this suggested feature.  While this compact feature is still questionable, the evidence for the $\\sim2'$ long jet is unambiguous. Although all the X-ray studies suggest its non-thermal nature, its physical origin remains to be obscure. Very recently, a collimated $\\sim9'$ long X-ray feature has been discovered from a radio-quiet $\\gamma-$ray pulsar PSR~J0357+3205 (De Luca et al. 2011) which is very similar to the jet of \\psr. It is interesting to compare the observed properties between these two systems. The spin-down luminosity of PSR~J0357+3205, $\\dot{E}\\sim6\\times10^{33}$~erg/s, is $\\sim6$ times higher than that of \\psr, but it is the lowest among all known non-recycled $\\gamma-$ray pulsars. The low column absorption suggests it is a nearby pulsar with a distance of only a few hundred parsec. Assuming there is no inclination with respect to the sky plane, this suggests the physical extent of the trail associated with PSR~J0357+3205 at the order of $\\sim1$~pc which is similar to the case of \\psr. In both cases, there is no evidence of spectral variation along the feature. For a distance of 500 pc, the X-ray conversion efficiency $L_{x}/\\dot{E}$ of its trail in $0.5-10$~keV is $\\sim10^{-3}$. Adopting a distance of 1~kpc for \\psr, $L_{x}/\\dot{E}$ of its jet in the same energy band is found to be $\\sim10^{-2}$ which is an order larger than that for PSR~J0357+3205. This might be why the jet of \\psr\\ can still be detected in spite of its larger distance.  Following De Luca et al. (2011), we estimate the highest achievable energy of synchrotron-radiating particles injected by \\psr. First, we assume the highest energy of the electrons/positrons can be gained is comparable with maximum potential drop in the magnetosphere, which is $\\Delta\\Phi=\\left(3\\dot{E}/2c\\right)^{1/2}$ for an aligned pulsar (Goldreich \\& Julian 1969). Similar to the case of PSR~J0357+3205, this results in a maximum Lorentz factor of $\\gamma_{\\rm max}\\sim10^{8}$. In the presence of a magnetic field, these particles will radiate synchrotron emission with characteristic frequency of $\\nu_x=\\gamma^2eB/m_ec$, where $B_{\\mu G}$ is the local magnetic field strength in the emission region in units of microgauss. Adopted a typical magnetic field at an order of $\\sim5~\\mu$G, $\\gamma_{\\rm max}$ implies a synchrotron photon energy at the order of $h\\nu_x\\sim5$~keV.  One difficulty for the synchrotron scenario is that it requires the Lorentz factor of the emitting particles to achieve the same order of $\\gamma_{\\rm max}$ so as to produce X-rays at a few keV. On the other hand, inverse Compton (IC) scattering between the relativistic particles and the relic photons appears to be more viable in producing the observed X-rays. We note that the energy density of the relic photons is at the order of $\\sim10^{-12}$~erg~cm$^{-3}$, which is comparable with the magnetic energy density for a field strength of $\\sim5\\mu$G. Therefore the IC power cannot be ignored. In the IC scenario, the required Lorentz factor to produce $\\sim10$~keV photons is only at the order of $\\gamma\\sim10^{4}$ which is much easier to achieve.  The major difficulty for explaining the X-ray jet of \\psr\\ in the context of PWN comes from its orientation with respect to the pulsar's proper motion. The jet protrudes from the pulsar and extends toward northwest, which makes an angle of $\\sim118^{\\circ}$ from the proper motion direction of the pulsar (cf. Fig.~\\ref{xmm_img}). This peculiar orientation makes the bow-shock interpretation questionable as this scenario requires the nebular emission lies in an opposite direction of the pulsar's velocity vector. In view of this difficulty, different explanations for the nature of the jet association with \\psr\\ have been proposed, ranging from the magnetically confinement of relativistic particles leaking from the bow-shock (Bandiera 2008) to a picture similar to the AGN jet outflow (Johnson \\& Wang 2010). Nevertheless, there is no consensus on its physical origin yet.  Since the extended feature of PSR~J0357+3205 is akin to that of \\psr, a comparative analysis between these two systems might provide us a deeper insight on the nature of their jets. While the distance and the proper motion direction of \\psr\\ are well-constrained, these are the most important parameters to be determined in further investigation of PSR~J0357+3205. These would enable a more reliable estimations of its energetic and most importantly the orientation of the jet with respect to the pulsar's space velocity, which will allow us to ascertain if these two systems have a similar physical origin.  \\subsection{The origin of \\psr} \\psr\\ is certainly one of the most spectacular neutron stars regarding its motion through space as its transverse velocity, $v=865(d/1~{\\rm kpc})$~km~s$^{-1}$, is among the highest in the currently known pulsar population. On the other hand, the well measured bow shock with its inclination in the plane of the sky suggests a negligible radial velocity component (Chatterjee \\& Cordes 2004) which has been confirmed recently by Tetzlaff et al. (2009) who suggest the origin of \\psr\\ to lie within the Cygnus OB3 association or a small massive cluster nearby. This origin indicates a small radial velocity of $v_{r}=-21^{+81}_{-70}$~km~s$^{-1}$. The proposed kinematic age of $\\tau_{\\rm kin}\\approx 0.8$ Myr is consistent with a characteristic pulsar age of $\\tau_{\\rm char}=1.1$~Myr (Hobbs et al. 2004; Manchester et al. 2005).  To advance the analysis we presented earlier (Tetzlaff et al. 2009), we attempt to identify a runaway star which might have been the former companion of this pulsar's progenitor star (we refer to such runaway stars as binary supernova scenario (BSS) runaways; Blaauw 1961). Therefore, we take all runaway star candidates with full kinematics from the runaway star catalogue of Tetzlaff et al. (2011) and compare the past flight paths of \\psr\\ and each runaway star varying the observables by performing a Monte-Carlo simulation to find close encounters between the pulsar and the runaway star. Simultaneously, we calculate the distance to possible parent associations/clusters listed in Tetzlaff et al. (2010) (for a full description of the procedure we refer to Tetzlaff et al. 2009 and Tetzlaff et al. 2010).\\footnote{\\psr\\ reaches a distance to the Sun larger than 3 kpc (the catalogue limit of the current runaway positions) if it is older than 1.3 Myr (assuming $v_{r}=0$~km/s). Since the characteristic age ($\\sim1$ Myr) provides an upper limit on the true age of young (few Myr) pulsars in most cases, it is well possible to identify the associated runaway star (the former companion, if it exists) among the catalogue runaway star candidates.}   For \\psr\\ we take the following data for the equatorial coordinates, right ascension $\\alpha$ and declination $\\delta$, the distance to the Sun $d$ and the proper motion components $\\mu_{\\alpha}^*$ ($\\mu_{\\alpha}^{*}=\\mu_{\\alpha}\\cos\\delta$) and $\\mu_{\\delta}$ (Harrison et al. 1993; Taylor \\& Cordes 1993; Yuan et al. 2010);  \\begin{equation} \\begin{array}{l c l} \\alpha &=& 336^\\circ\\hspace{-0.8ex}.469671,\\ \\delta\\ =\\ 65^\\circ\\hspace{-0.8ex}.593217,\\\\ d &=& 2\\pm1~{\\rm kpc},\\\\ \\mu_{\\alpha}^* &=& 144\\pm3~{\\rm mas~yr}^{-1},\\ \\mu_{\\delta} \\ =\\ 112\\pm3~{\\rm mas~yr}^{-1}. \\end{array} \\end{equation}  Here we adopt $d=2$~kpc for the calculation as published by Taylor \\& Cordes (1993) and adopt a reasonable error of 1~kpc that covers the distance inferred by Cordes \\& Lazio (2002) and Chatterjee \\& Cordes (2004). Among 1703 runaway star candidates in the catalogue with full kinematic data, only nine appear to show possible close encounters with \\psr\\ (HIP 97359, 98773, 99283, 99580, 100308, 100409, 101186, 107418, 109393). For two of them, HIP 107418 and HIP 109393, separations between the runaway star and the pulsar are small at $\\tau_{\\rm kin}\\approx0$, i.e. at present; thus we exclude these two stars from our analysis.  Tracing back the seven remaining runaway stars and OB associations/ clusters listed in Tetzlaff et al. (2010), we find that four stars could have been ejected from their parent cluster within the last 3 Myr ($\\approx3\\cdot\\tau_{\\rm char}$) (see Table~\\ref{runaway}). For them, we simultaneously trace back \\psr, the runaway star and the potential parent association varying the observables within their confidence intervals (here we adopt $v_r=0\\pm100$~km s$^{-1}$ for \\psr) to find the closest approach between the pulsar and the runaway star (i.e. the minimum separation $d_{\\rm min}\\left(\\tau_{\\rm kin}\\right)$) as well as the associated time ($\\tau_{\\rm kin}$) and the distance to the respective association centre at that time.  In only one case -- HIP 99580 (=V2011 Cyg) with Cyg OB9 being the parent association -- the pulsar and the runaway star can be located within the association boundaries (radius of Cyg OB9 $\\approx40$~pc) at the same time in the past. The smallest separation $d_{\\rm min}$ between \\psr\\ and HIP 99580 that is found after 15 million Monte-Carlo runs is $0.4$~pc. Between 25~pc and 35~pc distance from the centre of Cyg OB9 \\footnote{We choose a distance range of 10~pc to make it comparable to the previously suggested neutron star/ runaway pair PSR B1929+10/$\\zeta$ Ophiuchi, see Hoogerwerf et al. (2001) and Tetzlaff et al. (2010)}, we find a peak of small separations $d_{\\rm min}$ between the pulsar and the runaway star. In 2540 simulations the distance of the pulsar as well as the distance of the runaway star to the centre of Cyg OB9 falls within this range. The first part of the distribution of the separations $d_{\\rm min}$ between \\psr\\ and HIP 99580 of these runs is consistent with the hypothesis that both stars once were at the same position $0.65\\pm0.05$~Myr in the past (see Figure~\\ref{histo}).  Although only a small fraction of runs yield small separations $d_{min}$ between the stars as well as the centre of Cyg OB9 the scenario seems likely considering the errors on the parameters involved (distance $d$, proper motion components $\\mu_\\alpha^\\ast=\\mu_\\alpha\\cos\\delta$ and $\\mu_\\delta$, radial velocity $v_r$ of both stars\\footnote{The errors on the motion of Cyg OB9 are also taken into account but are negligible compared to the stars.}). The fraction of successful runs is much smaller than in the case of the known neutron star/ runaway pair PSR B1929+10/ $\\zeta$ Ophiuchi (about one order of magnitude; for PSR B1929+10/ $\\zeta$ Oph three million Monte-Carlo runs were performed approximately twice the number of successful runs compared to 15 million here; see Hoogerwerf et al. 2001 and Tetzlaff et al. 2010); however, since \\psr\\ and HIP 99580 are far more distant from the Sun than PSR~B1929+10 and $\\zeta$ Oph (the error on the distance is larger by about one to two orders of magnitude), this is to be expected.  In Figure~\\ref{correlation}, we show the correlations between the astrometric parameters which were used as input for the 2540 runs for which both, \\psr\\ and the runaway star HIP 99580, were situated between 25~pc and 35~pc from the centre of Cyg OB9. For the runaway star only little correlation between $\\pi$ and $v_r$ can be seen and the most probable values of all parameters fall clearly within the error bars. Whereas for the pulsar there is a strong correlation between $\\pi$ and $v_r$. This is expected since an increase in distance (decrease in $\\pi$) can be compensated for by a larger radial velocity $v_r$ and vice versa. If \\psr\\ and the runaway star HIP 99580 were at the same place 0.65~Myr ago in Cyg OB9, i.e. once were in the same multiple system, the present distance of \\psr\\ needs to be $\\sim0.8$~kpc. Hence, we can confirm previous findings by Chatterjee \\& Codes (2004) which concludes the pulsar have a distance to the Sun of $\\sim1$~kpc.  If this supernova occurred $0.65$~Myr in the past, it was located at Galactic coordinates $l=77.4\\pm0.2~{\\rm deg}, b=2.5^{+0.4}_{-0.2}~{\\rm deg}$\\footnote{These coordinates are as viewed at the present time. (i.e. Without the correction for the Sun's motion through the Galaxy over the last $0.65$~Myr.)}, and a distance to Earth of $d_{\\rm Earth}=968^{+32}_{-28}$~pc.  If HIP 99580 is a BSS runaway star, it should show signs of stellar evolution in the former system such as a large rotational velocity and an enhanced helium abundance caused by momentum and mass transfer during evolution. It was claimed by Barannikov (1993) that HIP 99580 is a spectroscopic binary, however De Becker \\& Rauw (2004)\\footnote{We note that these authors quote a different radial velocity of $v_r=-23.1\\pm8.0$~km~s$^{-1}$ for HIP 99580 compared to $-44\\pm4$ km/s (Kharchenko et al. 2007) that we used here. In this case, the star would still be a runaway star as it was identified also from its tangential velocity (Tetzlaff et al. 2011). However, only a tiny fraction of runs would then yield small separations between HIP 99580 and \\psr\\ making the BSS scenario unlikely. But note that both radial velocities can be correct if the star is a spectroscopic binary, then it would be necessary to determine the systematic radial velocity.} and Turner et al. (2008) could not confirm this. According to Conti \\& Ebbets (1977), the O5 type star HIP 99580 is a rapid rotator with $v\\sin{i}=270$~km~s$^{-1}$. Penny (1996) investigated $v\\sin{i}$ of O-type stars and found that only a small number of very massive stars show such a large $v\\sin{i}$. Hence, it is reasonable to assume that it gained its large $v\\sin{i}$ during the earlier evolution in a binary system. Furthermore, BSS runaway stars should be blue stragglers, i.e. they appear younger, hence bluer, than their parent association (see Hoogerwerf et al. 2001 for other examples). In Figure~\\ref{bluestraggler} we show the colour-magnitude diagram of the potential parent association Cyg OB9 with the data adopted from Garmany \\& Stencel (1992). Apparently, the association consists of two major groups of different ages, the younger one being $\\sim4$ Myr old.\\footnote{\\bf Note that Uyan{\\i}ker et al. (2001) quote an age for Cyg OB9 of 8~Myr.} If HIP 99580 originated from Cyg OB9, it is clearly a blue straggler.  Assuming an association age of 4~Myr, the progenitor star of the neutron star must have been 3.4~Myr old as it exploded. Thus, its mass was $(67\\pm15)~M_{\\odot}$ (Tinsley 1980; Maeder \\& Meynet 1989; Kodama 1997). For such large masses, formation of black holes is expected rather than neutron stars (upper limit $\\approx30~M_{\\odot}$, Heger et al. 2003). However, due to mass transfer during binary evolution, even more massive stars ($\\gtrsim50-60~M_{\\odot}$) may eventually produce neutron stars instead of black holes (Heger et al. 2003; Belczynski \\& Taam 2008). Also, more massive progenitors may also produce larger speeds of the neutron star (Burrows \\& Young 2000) which is consistent with the extraordinary speed of \\psr."
    },
    "1112/1112.0744_arXiv.txt": {
        "abstract": "{Multi-wavelength observations of gamma-ray burst (GRB) afterglows provide important information about the activity of their central engines and their environments. In particular, the short timescale variability, such as bumps and/or rebrightening features visible in the multi-wavelength light curves, is still poorly understood.} {We analyze the multi-wavelength observations of the GRB\\,100219A afterglow at redshift 4.7. In particular, we attempt to identify the physical origin of the late achromatic flares/bumps detected in the  X-ray and optical bands.} {We present ground-based optical photometric data and \\textit{Swift} X-ray observations on GRB\\,100219A. We analyzed the temporal behavior of the X-ray and optical light curves, as well as the X-ray spectra.} {The early flares in the X-ray and optical light curves peak simultaneously at about 1000 s after the burst trigger, while late achromatic bumps in the X-ray and optical bands appear at about $2\\times 10^4$~s after the burst trigger. These are uncommon features in the afterglow phenomenology. Considering the temporal and spectral properties, we argue that both optical and X-ray emissions come from the same mechanism. The late flares/bumps may be produced by late internal shocks from long-lasting activity of the central engine. An off-axis origin for a structured jet model is also discussed to interpret the bump shapes. The early optical bump can be interpreted as the afterglow onset, while the early X-ray flare could be caused by the internal activity. GRB\\,100219A exploded in a dense environment as revealed by the strong attenuation of X-ray emission and the optical-to-X-ray spectral energy distribution.} {} ",
        "introduction": "Gamma-ray bursts (GRBs) are detected by high-energy observational satellites. Ground-based telescopes can be subsequently alerted and perform follow-up observations. The accurate positions delivered by the \\textit{Swift} satellite provide an excellent opportunity for multi-wavelength observations. The analysis of GRB light curves can provide plenty of information about the central engine and the surrounding environment. A canonical shape has been identified for the X-ray light curves (Nousek et al. 2006; Zhang et al. 2006), as obtained by the \\textit{Swift} X-ray Telescope (XRT). In the optical band, following the alert by the \\textit{Swift} Burst Alert Telescope (BAT), ground-based follow-up observations carried out by robotic telescopes increased the number of well-sampled optical light curves significantly (see, e.g., Melandri et al. 2008; Klotz et al. 2009; Rykoff et al. 2009; Cenko et al. 2009). Apparently, these light curves present different temporal behaviors. It is hard to identify a uniform characterization for them. An attempt to classify the temporal properties has been made by Melandri et al. (2008). Rykoff et al. (2009) attempted to find commonalities within a sample of optical light curves, finding that the external forward shock is a possible origin for the overall optical emission. Panaitescu et al. (2006) have studied in detail the decay of light curves both in the X-ray and optical bands from a theoretical point of view. Chromatic breaks identified by comparing optical and X-ray light curves indicate that most likely the optical and X-ray emissions arise from a different origin.  While light curves generically decay in time, in several cases rebrightenings or bumps are observed in the X-ray or optical bands. These features call for a more detailed investigation of the physics of GRB and afterglow. However, the situation is quite complicated. Most rebrightenings are only observed in the X-ray band. Some early X-ray bumps, usually called X-ray flares, have no corresponding optical features (see the statistics by Melandri et al. 2008 and Rykoff et al. 2009, as well as Uehara et al. 2010 for the individual cases of GRB\\,071112C and GRB\\,080506). In contrast, but less frequently, a rebrightening feature may be seen in the optical but not in the X-ray band (e.g. GRB\\,050721; Antonelli et al. 2006). Some GRBs with both X-ray and optical bumps shown before 1000 s in the observer's frame have been identified: GRB\\,060418 and GRB\\,060607A (Molinari et al. 2007), GRB\\,060904B (Klotz et al. 2008) and GRB\\,071031 (Kr\\\"{u}hler et al. 2009). The X-ray flare and optical bump of GRB\\,060418 have the same peak time. However, it is likely that the X-ray flare has an internal origin while the optical bump is the result of external shock onset (see the optical statistics from Oates et al. 2008). The optical bump and giant X-ray flare of GRB\\,060904B are clearly chromatic. The optical rising and X-ray flare of GRB\\,071031 are not exactly simultaneous, but their observed correlation suggests a common origin caused by late central engine activity.  More importantly, it is worth noting that GRBs with late bumps/rebrightenings shown after $10^3$--$10^4$~s in the observer's frame are very rare. We mention that the X-ray light curves of GRB\\,050502B, GRB\\,050724, GRB\\,050904, GRB\\,060413, GRB\\,060906, and GRB\\,070311 have similar late bump/flare/wiggle features. Some of them have been identified as late X-ray flares (Curran et al. 2008; Bernardini et al. 2011). It is rare, however, to have well-sampled, complete optical data at the time corresponding to the late X-ray bumps (see, e.g., Afonso et al. 2011). A broad, late optical rebrightening was found by Mundell et al. (2007) for GRB\\,061007, but no corresponding bump was seen in the X-ray band. The late-observed optical emission of GRB\\,050724 is likely related to the X-ray flare peaking at 41.8~ks, but the evidence for an optical rebrightening is not conclusive (Malesani et al. 2007). GRB\\,070311 is one case with relatively clear evidence of late bumps in both the X-ray and optical bands (Guidorzi et al. 2007), with comparable, although not simultaneous, peak times. The most interesting case is GRB\\,071010A, for which a late rebrightening feature was identified at 0.6 day after the trigger, visible simultaneously in both the optical/NIR and X-ray bands (Covino et al. 2008).  Several mechanisms responsible for the X-ray and optical bumps presented in GRB light curves have been proposed. Kumar et al. (2008a, b) suggested a model for the early flares powered by accretion of fall-back matter from the progenitor star. In general, early and late X-ray flares may have an internal origin although the detailed physics is unclear (Chincarini et al. 2010; Bernardini et al. 2011). An early rising of the optical light curve can be sometimes caused by the onset of forward shock (Oates et al. 2009), another discussion was given by Melandri et al. (2010). Some late optical bumps are interpreted by the external shock model. For example, the optical rebrightening feature of GRB\\,021004 seen at 0.1 days after the burst (Holland et al. 2003) can be explained by the interaction with a clumpy ambient (Lazzati et al. 2002; but see also Nakar et al. 2003). An early rising feature around 1000 s and a later rebrightening behavior after $10^4-10^5$ s can be described by the standard reverse-forward shock model as well (Zhang et al. 2003, Fan et al. 2005). However, for the short burst GRB\\,060313, the wiggles shown at about $10^4$ s have been proposed to be the result of late-time central engine activity (Roming et al. 2006). To explain the late rebrightening of GRB\\,070311 (Guidorzi et al. 2007), refreshed shocks were proposed (Rees \\& M\\'{e}sz\\'{a}ros 1998). The achromatic rebrightening feature of GRB\\,071010A may be produced by an episode of discrete energy injection (Covino et al. 2008).  In this paper, we analyze the X-ray and optical light curves of GRB\\,100219A. We report data taken from the 2.4-m telescope at Gao-Mei-Gu (GMG) and  from the Nordic Optical Telescope (NOT), with the former providing most of the data. The optical light curve shows a clear rebrightening around $2\\times10^4$~s after the GRB trigger. A simultaneous X-ray temporal wiggle/small bump is also detected. Thus, the late achromatic bumps of GRB\\,100219A provide us with an excellent opportunity to investigate the origin of the bumps/flares in the different observational bands in more detail.  We describe all observations of GRB\\,100219A in Section 2. The GMG and NOT telescopes have similar mirror size and detectors. Being located at roughly the same latitude in the Northern hemisphere, paired together they provide an excellent chance for continuing observations of GRB afterglows. In Section 3 we compare the optical light curves with the X-ray data, and we identify achromatic bumps in the X-ray and optical light curves. We also use the optical-to-X-ray spectral energy distribution (SED), aiming at investigating whether the late optical and X-ray bumps come from the same mechanism. The discussion about the possible physical origin is presented in Section 4. In Section 5, our conclusions are summarized.  Throughout the paper, we adopt the convention $F_\\nu(t,\\nu) \\propto t^{-\\alpha}\\nu ^{-\\beta}$ and cosmological parameters $H_0 = 70$ km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_\\mathrm{m} = 0.3$ and $\\Omega_\\Lambda = 0.7$. ",
        "conclusions": "We have presented the multi-wavelength observations of GRB\\,100219A using the data from \\textit{Swift} and ground-based telescopes. The early (1000 s) and late ($2\\times 10^4$ s) achromatic bumps visible in the X-ray and optical light curves are comprehensively discussed. In our analysis, we speculate that the early optical $I_{\\rm C}$-band bump is probably caused by the afterglow onset while the X-ray one might be caused by the internal activity. The late X-ray and optical bumps may be produced by the internal shell ejection from the long-active central engine. The jet structure could also play an important role for the achromatic property of the bumps. Moreover, the medium surrounding this GRB might be dense.  At present, many theoretical models have been proposed successfully to solve the problems of GRB energy release and to explain some major observational phenomena (such as prompt emission, early X-ray light curve and X-ray flares). However, it is still difficult to simply apply these models to specify the late achromatic bumps discussed in this paper. From this point of view, the theoretical explanations have to concentrate on more subtle observational features in the future."
    },
    "1112/1112.5357_arXiv.txt": {
        "abstract": "{} {We propose and explore the potential of a method to extract high signal-to-noise (S/N) integrated spectra related to physical and/or morphological regions on a 2-dimensional field using Integral Field Spectroscopy (IFS) observations by employing an optimization procedure based on either continuum (stellar) or line (nebular) emission features.} {The optimization method is applied to a set of IFS VLT-VIMOS observations of (U)LIRG galaxies, describing the advantages of the optimization by comparing the results with a fixed-aperture, single spectrum case, and by implementing some statistical tests.} {We demonstrate that the S/N of the IFS optimized integrated spectra is significantly enhanced when compared with the single aperture unprocessed case. In some cases, the optimization based on the emission-line allows to characterize some of its properties better than with standard integration methods. It also allows to retrieve well the weak continuum features, and therefore to constrain better the properties of the unresolved stellar population. The preferred method to integrate spectra over (part of) the Field-of-View is ultimately dependent on the science case, and it may require a complex the trade off among different variables (e.g. S/N, probe area, spatial resolution, etc). Therefore, we provide an iterative user-friendly and versatile IDL algorithm that, in addition to the above mentioned method, allows the user to spatially integrate spectra following more standard procedures. This is made available to the community as part of the PINGSoft IFS software package.} {} ",
        "introduction": "Large surveys of spatially integrated spectrophotometry provide a powerful means of investigating the physical properties of galaxies at different epochs in the history of the universe (e.g. SDSS, \\citealt{York:2000p2677}; 2dFGRS, \\citealt{Colless:2001p2675}; GEMS, \\citealt{Rix:2004p3993}; etc.). Spectral diagnostics based on integrated optical spectroscopy can be used to constrain the star formation rate, star formation history, stellar mass, chemical abundance, dust content, and other main drivers of galaxy evolution \\citep[e.g.][]{Tremonti:2004p1138}. The analysis of the integrated spectra in nearby objects allow us to compare local and distant samples, as well as to asses the limitations in high-redshift galaxy observations. In fact, in addition to signal-to-noise (S/N) limitations, an incomplete spatial coverage (or aperture bias) may be important given that many physical properties of galaxies vary depending on the geometry and position (e.g. stellar populations, metallicity, extinction, etc.).   The study of integrated spectral properties of galaxies poses several observational challenges with classical spectroscopic techniques. The main difficulty resides in obtaining an integrated spectrum of an extended object that typically is larger than the area covered by a long slit \\citep[e.g.][]{Moustakas:2006p307}. Usually the \"integrated spectrum\" refers to a single (nuclear) spectrum obtained within a fixed aperture, which corresponds to different linear scales at different redshifts, and generally there are limited possibilities to correct for these aperture effects.   Integral Field Spectroscopy (IFS) has become a popular observational technique to obtain spatially resolved spectroscopy of extended objects. One by-product of IFS data sets which was early recognised is the intrinsic capability of integrating (by adding up some or all) the individual spectra within an observed field or mosaic into a single spectrum, i.e. using the Integral Field Unit (IFU) as a large-aperture spectrograph to obtain the integrated spectra of a given Field-of-View (FoV), increasing significantly the S/N of the final integrated spectrum. For instance, \\citet{Mediavilla:1992p3969} averaged radial profiles from the IFS spectra to study the outwards variation of the physical conditions and kinematics of the circumnuclear region of NGC\\,4151. In a recent example, \\citet{Sanchez:2011p3844} obtained and analysed the integrated spectrum of NGC\\,628 based on the the largest IFS mosaic on a single nearby galaxy, by adding up more than 10 thousand spectra.   Obviously, the increase of S/N in the binned or integrated spectrum using IFS techniques has the price of a loss in angular resolution. However, a spectrum obtained from a full 2D FoV can be better used to study the real average spectroscopic properties of a given galaxy than those spectra taken from different regions, or with a limited extraction aperture which recovers only a fraction of the total optical light. This application of IFS has become increasingly important when dealing with objects that are in a photon starving regime, as it is usually in high-z studies, but it is also relevant when analysing faint features in local bright objects, and in cases where the analysis is very sensitive to S/N (e.g. stellar populations).   In principle there are several ways to combine the individual spectral of an IFS dataset to obtain an integrated spectrum. The preferred method obviously depends on the specific science application. For instance, if one wants to simulate a loss of angular resolution (due, for instance, to distant or seeing effects), a simple (weighted) integration may suffice. However, in other cases a more sophisticated methods may be preferred to obtain a trade between spatial information and S/N, for example the adaptive smoothing technique \\citep{Silverman:1986p4047}, which consists of correlating 2D neighbouring information; or the Voronoi binning technique \\citep[e.g.][]{Cappellari:2003p4046,Cappellari:2009p3997} which aims to preserve the maximum spatial resolution given a constraint on a minimum S/N by partitioning the data using a Voronoi tessellation (see \\citealt{Diehl:2006p3996} for a comparison between Voronoi binning and adaptive smoothing).   In this paper we describe the basis of a simple algorithm, part of the {\\sc PINGSoft} package for IFS visualization and analysis \\citep{RosalesOrtega:2011p3845}, that allows to combine IFS spectra following different prescriptions. We also propose a method to optimize (in terms of S/N) the extraction of an integrated spectrum from an IFS data set based on either continuum (stellar) or line (nebular) emission features. We show the advantages of this IFS S/N optimization method with respect to normal fixed-aperture extractions for the integrated spectra of galaxies by implementing the optimization to a IFS data set of nearby (Ultra)Luminous Infrared Galaxies\\footnote{LIRGs: $L_{IR} \\equiv L[8-1000 \\mu m]= 10^{11-12}L_{\\odot}$; ULIRGs: $L_{IR} > 10^{12}L_{\\odot}$} [(U)LIRGs], part of the IFS low-z (U)LIRG survey \\citep[][hereafter A08]{Arribas:2008p3550}. This paper is structured as follows: in \\S \\ref{sec:2} we describe the S/N spectra extraction optimization method proposed in this work,  which can be applied to any 2D field containing spectra with continuum and/or line emission; in \\S \\ref{sec:3} we apply the method to a set of IFS VLT-VIMOS observations of (U)LIRG galaxies, describing the advantages of the optimization by comparing the results with a fixed-aperture, single spectrum case, and by implementing some statistical tests; the conclusions of this work are presented in \\S \\ref{sec:4}. In Appen. \\ref{appen:1} and \\ref{appen:2} we describe publicly available scripts in order to apply the S/N optimization with any IFS data set. ",
        "conclusions": "\\label{sec:4}   In this paper we propose a method to extract integrated spectra from a 2D IFS data cube based on a S/N optimization by employing a procedure relying on either continuum (stellar) or line (nebular) emission features from the final-reduced observed spectra. The methodology of the spectra selection is based on a {\\em statistical} S/N defined in a different manner for the continuum and line-emission cases. The actual implementation of the extraction method relies on stacking the S/N-sorted spectra and calculating the {\\em cumulative} S/N of the integrated spectrum. We have shown several examples to illustrate qualitatively and quantitatively the potential of this optimization method for studies based on both, continuum and emission features. As for the nebular emission, we have shown a significant better definition of the kinematic profiles for the optimized spectrum. Regarding the weak stellar absorption features, we have found a significant better constrain of the output solution when applying the SSP fitting technique. In particular, Monte Carlo simulations in a particular case studied show that the uncertainties in the output Z, and $A_V$ are reduced by factors of 5 and 2, respectively, when compared with the unprocessed case.   We provide series of IDL scripts that can be applied to basically any IFS datacube in order to implement the S/N optimization here described. This tools are made available to the community as part of the {\\sc PINGSoft} IFS software package. In addition to the proposed method to optimize the S/N of the integrated spectrum, more standard procedures to obtain integrated spectra are possible with mentioned tools. In particular the following options are implemented:  \\begin{enumerate} \\item[i)] Integration of the spectra over a specific geometrical region of the galaxy/FoV \\item[ii)] Integration of the spectra above a given S/N threshold \\item[iii)] A combination of i) and ii), or \\item[iv)] Integration of the spectra that optimize the S/N over a given region, following the method here described. \\end{enumerate}  \\noindent These possibilities can be done on the basis of a (S/N)$_c$ or a (S/N)$_{em}$, with and without the velocity field correction (see Appen. \\ref{appen:1}).  The process for obtaining an integrated spectrum for an optimal science use is complex and it requires in practice a trade-off among several variables (e.g. S/N, geometrical region probed, spatial resolution, etc). Interactive user-friendly versatile tools, as the one presented here, are essential to carry out this process in an efficient manner. The application of adapting techniques in order to extract integrated information in future 2D spectroscopic surveys might be relevant when dealing with objects that are in a photon starving regime, as it is usually in high-z studies, but it is also relevant when analysing faint features in local bright objects, and in cases where the analysis is very sensitive to S/N. \\autoref{tab1} summarizes some general science scenarios in which the S/N optimization technique can be applied, together with the current main applications of other techniques. The S/N optimization based on emission lines could be advantageous when studying emission lines in poorly resolved high-z star forming galaxies or objects with weak emission lines, and when the systemic velocity field is removed, the method can be applied to objects with evidence of shocks, turbulence, AGNs, etc. to decouple the kinematic components of the integrated spectrum of the galaxy. The S/N optimization based on the continuum may help to retrieve weak continuum features in the integrated spectra of galaxies allowing to constrain in a more efficient way the derived physical parameters of the unresolved stellar populations using techniques such as spectral synthesis or age-metallicity sensitive absorption features. In all cases, the S/N optimization allows to select a subsample of spaxels to perform a spatially resolved studies and/or to use the outputs as a proxy for more sophisticated adaptive binning techniques. Note that the S/N optimization here proposed might be also applicable to IFS observations of nearby objects or any other 2D observation where S/N-sorted spectra may be relevant. The ultimate application of this method depends on the nature of the IFS data, the scientific case and/or the practical needs of a potential user."
    },
    "1112/1112.2788_arXiv.txt": {
        "abstract": "Despite the lack of important functions supported only by legacy applications, the current VO-compatible tools have enough capabilities to allow powerful analysis of stellar spectra using both public archives and local proprietary data. We give examples of the possible multi-wavelength analysis of changes of profiles in  different spectral lines  using the current SSAP-compatible VO tools. The potential of future VO applications supporting more elaborate methods for stellar analysis is discussed as well. ",
        "introduction": "Astronomical spectroscopy uses a wide range of techniques with different level of complexity to achieve its final goal --- to  estimate the most precise and reliable  information about celestial objects.  The  scientific analysis of such  spectra requires further processing by the variety of different methods. In certain studies a huge number of spectra has to be collected from different servers (in different spectral regions) and transformed into  common units.  Accomplishing the multi-spectral analysis in VO environment may benefit from automatic aggregation of distributed archive resources, seamless on-the-fly data conversion, common interoperability of all tools  (using PLASTIC or SAMP protocol) and powerful graphical visualisation of measured and derived quantities ",
        "conclusions": "The  large part of spectroscopic analysis today  has been accomplished by several independent non VO-compatible legacy packages as the currently available VO clients supporting the SSA protocol can provide with only simple, mostly interactive capabilities on individual spectra.  The processing of large amount of spectra in a consistent way requires the simple post-processing (e.g. cutout or rebinning) services to be implemented on the server side  A lot of interesting physical information about astronomical objected can be obtained from only several short spectral ranges covering selected spectral lines. Dynamic behaviour of interesting objects can be estimated from the changes of line profiles as well as  from period analysis of their moments (including equivalent width).  By introduction of  modern  VO-aware tools into the  astronomical spectral analysis  a remarkable increase of effectiveness of astronomical research can be achieved."
    },
    "1112/1112.3092_arXiv.txt": {
        "abstract": "Gravitational wave is a propagation of space-time distortion, which is predicted by Einstein in general relativity. Strong gravitational waves will come from some drastic astronomical objects, e.g. coalescence of neutron star binaries, black holes, supernovae, rotating pulsars and pulsar glitches. Detection of the gravitational waves from these objects will open a new door of \\textit{`gravitational wave astronomy'}. Gravitational wave will be a probe to study the physics and astrophysics.  To search these gravitational waves, large-scale laser interferometers will compose a global network of detectors. Advanced LIGO and advanced Virgo are upgrading from currents detectors. One of LIGO detector is considering to move Australia Site. IndIGO or Einstein Telescope are future plans.  LCGT (Large-scale Cryogenic Gravitational wave Telescope) is now constructing in Japan with distinctive characters: cryogenic cooling mirror and underground site. We will present a design and a construction status of LCGT, and brief status of current gravitational wave detectors in the world.  Network of these gravitational wave detectors will start in late 2016 or 2017, and may discover the gravitational waves. For example, these detectors will reach its search range for coalescence of neutron star binary is over 200 Mpc, and several or more events per year will be expected. Since most of gravitational wave events are from high-energy phenomenon of the astronomical objects, these might have counterpart evidences in electromagnetic radiation (visible light, X/gamma ray), neutrino, high energy particles or others. Thus, the mutual follow-up observations will give us more information of these objects. ",
        "introduction": "Gravitational waves are predicted by Einstein's general relativistic theory in 1916. With the perturbation of space-time metric in Einsteins's equation, a wave equation is derived as \\begin{equation} \\left({{\\nabla }^{2}\\ -\\ {\\frac{1}{c}}{\\frac{{\\partial }^{2}\\,}{\\partial \\mit \\,{t}^{2}}}}\\right)\\ {h}_{\\mu \\nu }\\ =\\ 0 \\label{eq:GWEq} \\end{equation} where, $h_{\\mu \\nu}$ respects the metric perturbation away from Minkowski metric as \\begin{equation} {g}_{\\mu \\nu }\\ ={\\eta }_{\\mu \\nu }\\ +\\ {h}_{\\mu \\nu }\\ . \\label{eq:hDefinition} \\end{equation} This equation suggest that small distortion of space-time $h_{\\mu \\nu}$ will be able to propagate as wave, which is called as {\\it ``gravitational wave''}.  The gravitational wave will emitted from a changing of mass quadrupole moment, just like as a non-spherical motion of mass distribution. The lowest order of a gravitational wave is quadrupole radiation. With transverse-traceless gauge, choosing $z$ axis as a propagation, the wave can be represented the linear combination of two polarization : \\begin{equation} {h}_{+}\\ =\\ a \\left({\\matrix{0&0&0&0\\cr 0&1&0&0\\cr 0&0&-1&0\\cr 0&0&0&0\\cr}}\\right)\\ ,\\ \\ {h}_{\\rm \\times }\\ =\\ b \\left({\\matrix{0&0&0&0\\cr 0&0&1&0\\cr 0&1&0&0\\cr 0&0&0&0\\cr}}\\right)\\ . \\label{eq:hCross} \\end{equation} It is a transverse wave, and it propagate with light speed. A gravitational wave will interact with masses. When it incident, a tidal force will be induced on free falling masses. The $h_{+}$ component will induce the distortion along $x$- and $y$-axis differentially, and $h_{\\times}$ will induce same but rotating 45 degree.  Unit of gravitational wave amplitude is dimensionless metric. It means that space-time distortion relativity. The gravitational interaction is very weak comparing with other fundamental inter, and also the gravitational wave amplitude is extremely small as \\begin{equation} h_{\\mu\\nu} \\ =\\ \\frac{2G}{Rc^4} \\ddot{I}_{\\mu\\nu} \\end{equation} , where $I_{\\mu\\nu}$ is a quadrupole moment of mass distribution. For example, two 100 kg masses connected with 1m beam, rotating 1000 cycles per second will emit 2kHz gravitational wave of amplitude only $\\sim 10^{-40}$ at a distance of 150km where is wavelength.  \\subsection{Gravitational Waves from Astronomical Objects}  Since the gravitational wave is weak, we expect sources consists of huge mass motions - astronomical objects. Also the sources must have large acceleration to make a large derivative of quadrupole motion. Typical sources of gravitational waves are coalescence of compact star (neutron star, black hole) binaries, stellar-core collapse of supernovae. These are occasionally gravitational wave sources. On the other hand, rotating pulsars or binaries might be a sources of continuous gravitational waves.  Neutron star binary is a most promising gravitational wave sources. PSR1913+16 is a binary pulsar is known as that their Kepler orbit period changing is well consistent with a transfer of energy and momentum from binary system by gravitational wave radiation\\cite{TaylorHulse}. There are some observations of such binaries\\cite{lrr-2008}. A Binary system will loss the energy emitting gravitational waves, will shrink its orbit, and become rotate faster. Finally, the binary coalesce and emit large gravitational wave. It is called as {\\it `the last three minutes'}\\cite{TheLastThreeMin}. The amplitude $h$ of gravitational wave neutron star binary from 200 Mpc away will be $10^{-24}\\sim 10^{-23}$ in frequency band of 100Hz -- 1kHz, and its frequency spectral density will be $h(f) \\sim 10^{-22} \\sim 10^{-23} {\\rm [1/\\sqrt{Hz}]}$. Gravitational waveforms of compact binary coalescence can be well predicted by post-newtonian \\cite{CBCwaveform}. The frequency and amplitude development along the time ate determined by the masses of stars. Thus, once we measure the waveform, it guess the absolute amplitude of gravitational wave from the source. Comparing observed amplitude, we can estimate the luminosity distance of sources only by gravitational wave measurement. In this mean, compact binary coalescence are called as {\\it `standard candle (siren)'}. The rate of coalescence of neutron star binary is estimated as $118^{+174}_{-79}$ events/Myr/Galaxy \\cite{Kim_GMR}. Even this expectation contains large statistical error and dependency of some models, it is encouraging us to try to detect the gravitational waves from compact star coalescence.  On the other hand, the weakness, i.e. small coupling constant of gravitational wave make possible that its come from an inside the star. Electromagnetic radiations are mostly emitted at the surface of stars. Neutrino may come from inside as a core of supernovae. Gravitational wave come from mass motion itself inside the stars, and pass through the materials. This fact suggest that gravitational wave will bring many information of deep inside of high-energetic astronomical object. In the case of black hole, only a gravitational wave is radiation from themselves. We expect a new window of {\\it `gravitational wave astronomy'} will be opened with current constructing/upgrading gravitational wave detectors.  \\begin{figure}[bt] \\begin{center} \\includegraphics[width=80mm,trim=0 10 0 0, clip]{IFOc.eps} \\end{center} \\vspace{-5mm} \\caption{\\label{IFOske}Schematic view of laser interferometric gravitational wave detector} \\end{figure}   \\subsection{Detection Principle of Gravitational Waves}  Gravitational wave induce the tidal force between free masses. Laser interferometric gravitational wave detectors are Michelson interferometer that consists of suspended mirrors, beam splitter and laser as light source. Mirrors and beamsplitter are suspended as pendulum that is a mechanical low pass filter which make mirror as a free mass. Laser light is divided to two perpendicular paths, round trip from mirrors and re-combine at the beam splitter. We can measure the metric amplitude of gravitational wave by small change of intensity of the combined light. To enhance a gravitational wave signal, long base-line of the Michelson beam path and Fabry-Perot cavity are employed. Figure \\ref{IFOske} shows the schematic view of the detector. ",
        "conclusions": ""
    },
    "1112/1112.4810_arXiv.txt": {
        "abstract": "{We present forecast errors on a wide range of cosmological parameters obtained from a photometric cluster catalogue of a future wide-field \\emph{Euclid}-like survey. We focus in particular on the total neutrino mass as constrained by a combination of the galaxy cluster number counts and correlation function. For the latter we consider only the shape information and the Baryon Acoustic Oscillations (BAO), while marginalising over the spectral amplitude and the redshift space distortions. In addition to the cosmological parameters of the standard $\\Lambda$CDM$+\\nu$ model we also consider a non-vanishing curvature, and two parameters describing a redshift evolution for the dark energy equation of state. For completeness, we also marginalise over a set of ``nuisance'' parameters, representing the uncertainties on the cluster mass determination. We find that combining cluster counts with power spectrum information greatly improves the constraining power of each probe taken individually, with errors on cosmological parameters being reduced by up to an order of magnitude. In particular, the best improvements are for the parameters defining the dynamical evolution of dark energy, where cluster counts break degeneracies. Moreover, the resulting error on neutrino mass is at the level of $\\sigma(M_\\nu)\\sim 0.9$ eV, comparable with that derived from present Ly$\\alpha$ forest measurements and Cosmic Microwave background (CMB) data in the framework of a non-flat Universe. Further adopting \\emph{Planck} priors and reducing the number of free parameters to a $\\Lambda$CDM+$\\nu$ cosmology allows to place constraints on the total neutrino mass of $\\sigma(M_\\nu) \\sim 0.08$ eV, close to the lower bound enforced by neutrino oscillation experiments. Finally, in the optimistic case where uncertainties in the calibration of the mass-observable relation were so small to be neglected, the combination of \\emph{Planck} priors with cluster counts and power spectrum would constrain the total neutrino mass down to $\\sigma(M_\\nu) \\sim 0.034$ eV, i.e. the minimum neutrino mass predicted by oscillation experiments would be detected in a $\\Lambda$CDM framework. We thus show that galaxy clusters from future wide galaxy surveys will be an excellent tool for studying cosmology and fundamental physics.}  \\begin{document} ",
        "introduction": "\\label{Intro} It is now established from solar, atmospheric, reactor and accelerator experiments that neutrinos have non-zero mass, and that a lower limit on the total neutrino mass is given by $M_\\nu\\equiv \\sum m_{\\nu}\\sim 0.05$ eV \\cite{lesgourgues2006, Wong2011}, where $m_\\nu$ is the mass of a single neutrino species. On the other hand the individual masses are still unknown. Since neutrino mass affects the evolution of the Universe in several observable ways, its measurements can be obtained from different cosmological probes, such as observations of the CMB, galaxy clustering, Ly$\\alpha$ forest, and weak lensing data \\cite{abazajian2011,Joudaki&Kaplinghat2011,Joudaki2012,Signe_etal2011}.  In particular, a thermal neutrino relic component in the Universe impacts both the expansion history and the growth of cosmic structures. Neutrinos with mass $\\lesssim 0.6$ eV become non-relativistic after the epoch of recombination probed by the CMB, and this mechanism allows massive neutrinos to alter the matter-radiation equality for a fixed $\\Omega_m h^2$. Massive neutrinos act as non-relativistic particles on scales $k>k_{\\rm nr}=0.018(m_\\nu/1{\\rm eV})^{1/2}\\Omega_m^{1/2} h$/Mpc, where $k_{\\rm nr}$ is the wave-number corresponding to the Hubble horizon size at the epoch $z_{\\rm nr}$ when the given neutrino species becomes non-relativistic, $\\Omega_m$ is the matter energy density and $h=H_0/100\\, {\\rm km\\, s^{-1} Mpc^{-1}}$. The large velocity dispersion of non-relativistic neutrinos suppresses the formation of neutrino perturbations in a way that depends on $m_\\nu$ and redshift $z$, leaving an imprint on the matter power spectrum for scales $k>k_{\\rm fs}(z)=0.82 H(z)/H_0/(1+z)^2 (m_\\nu/1{\\rm eV}) h$/Mpc \\cite{lesgourgues2006,takada2006}, where neutrinos cannot cluster and do not contribute to the gravitational potential wells produced by cold dark matter and baryons.  This modifies the shape of the matter power spectrum and the correlation function on these scales \\cite{doroshkevich1981,hu1998,abazajian2005,kiakotou2008,brandbyge2010,Matteo2,Takada2011}.  Massive neutrinos affect also the CMB statistics. WMAP7 alone constrains $M_\\nu < 1.3$ eV \\cite{Komatsuetal2010} and data from the ACT\\footnote{http://wwwphy.princeton.edu/act/} and SPT\\footnote{http://pole.uchicago.edu/} experiments constrain $M_\\nu < 0.948$ eV in the framework of a non-flat cosmology \\cite{Smith_etal2011}. Furthermore, thanks to the improved sensitivity to polarisation and to the angular power spectrum damping tail, forecasts for the \\emph{Planck} satellite alone give a 1--$\\sigma$ error on the total neutrino mass of $\\sim 0.2-0.4$ eV, depending on the assumed cosmological model and fiducial neutrino mass \\cite{perotto2006, kitching2008}. Moreover, the combination of present data-sets from CMB and large-scale structure (LSS) yields an upper limit of $M_\\nu<0.3$ eV \\cite{wang2005,vikhlinin2009,thomas2010,gonzalez2010,reid2010}. A further robust constraint on neutrino masses has been obtained using the Sloan Digital Sky Survey flux power spectrum alone, finding an upper limit of $M_\\nu<0.9$ eV ($2\\sigma$ C. L.) \\cite{Matteo1}. However, the tightest constraints to date in terms of a 2$\\sigma$ upper limit on the neutrino masses have been obtained by combining the Sloan Digital Sky Survey flux power from the Ly$\\alpha$ forest with CMB and galaxy clustering data. These constraints are partly driven by a discrepancy in the measured $\\sigma_8$ between Ly$\\alpha$ forest and CMB, and result in $\\Sigma m_\\nu<0.17$ eV \\cite{seljak2006}. Somewhat less constraining bounds have been obtained by \\cite{goobar2006}, while for forecasts on future CMB and Ly$\\alpha$ forest joint constraints we refer to \\cite{gratton2008}.  Moreover, the forecast sensitivity of future LSS experiments, when combined with \\emph{Planck} CMB priors, indicates that observations should soon be able to detect signatures of the cosmic neutrino background and measure the neutrino mass even in the case of the minimum mass $M_\\nu=0.05$ eV \\cite{kitching2008,hannestad2007,lsst2009,hannestad2010,lahav2010, Melita11}. Furthermore, these future surveys have been planned to measure with high accuracy the so-called ``dark energy'' equation of state. In fact, there is now strong evidence that the current energy density of the Universe is dominated by dark energy with an equation of state $w \\sim -1$, which is causing accelerated expansion. Accordingly, in this work we consider, in addition to massive neutrinos, the effect on LSS of a homogeneous dark energy component with a general time-varying equation of state $w(z)$.  In this respect, Ref.~\\cite{Melita11} has shown that future spectroscopic galaxy surveys, such as the ESA selected mission \\emph{Euclid}\\footnote{http://www.euclid-ec.org/} \\cite{RedBook}, will be capable of estimating the neutrino mass scale independently of flatness assumptions and dark energy parametrisation, if the total neutrino mass $M_\\nu$ is $>0.1$ eV. On the other hand, if $M_\\nu$ is $<0.1$ eV, the sum of neutrino masses, and in particular the minimum neutrino mass required by neutrino oscillations, can be measured in the context of a $\\Lambda$CDM model. For further discussion on neutrino mass constraints from different probes see e.g. \\cite{Wong2011,abazajian2011} and references therein.  In this picture, little attention has been dedicated to the constraining potential of galaxy clusters. Historically these systems have played a fundamental role in cosmology. For instance, the correlation function of galaxy clusters provided the very first hint toward the low-density Universe that is today commonly accepted as the standard cosmological model \\cite{Bahcall1,Bahcall2}. Galaxy cluster number counts at high-redshift later provided yet another early evidence for the matter density parameter being substantially smaller than unity \\cite{Carlberg,Donahue}. Galaxy clusters trace the large-scale matter distribution much as galaxies, with the difference that the former have a substantially larger bias than the latter, and hence their correlation is a factor of a few higher. At the same time, clusters are relatively rare objects, therefore their abundance is highly sensitive to features in the primordial matter power spectrum. These occurrences make clusters in principle valuable tools for constraining the details of the cosmological model, and in particular the neutrino masses. However their low spatial number density is also a drawback, in that it increases substantially the shot noise and the Poisson noise as compared to galaxies. Their effective cosmological power hence depends significantly on the chosen selection function and on the precision with which cluster masses can be estimated. In this work we study how the combination of cluster number counts with the shape of the cluster power spectrum, including information from the BAO, can constrain the total neutrino mass. We focus on the cluster catalogue that will be produced by the photometric part of \\emph{Euclid}, and provide constraints on the parameter set of an extended cosmological model which includes curvature, massive neutrinos, and a time-varying dark energy component. Since deviations from the standard model in the form of extra neutrino species are still uncertain \\cite{giusarma2011,gonzalez-morales2011}, in this paper we focus on standard neutrino families only and analyse the constraining power that future galaxy cluster catalogues, photometrically selected from galaxy surveys of \\emph{Euclid}-type, have on the total neutrino mass.  The rest of the paper is organised as follows. In \\S~\\ref{Clusters} we present the galaxy cluster modelling adopted in this work and the modifications due to uncertainties in the calibration of the mass-observable relation. In \\S~\\ref{Fiducial cosmologies} we describe the adopted fiducial model and the effect of massive neutrinos on the matter power spectrum. In \\S~\\ref{Fishers} we review our forecasting approach as applied to cluster counts and power spectrum, and specify characteristics of the galaxy cluster survey analysed in this work. In \\S~\\ref{Results} we present our results on the forecast neutrino mass errors, and finally in \\S~\\ref{Conclu} we draw our conclusions. ",
        "conclusions": "\\label{Conclu}  In this work we adopted a Fisher matrix approach to forecast constraints on a wide array of cosmological parameters from the photometric galaxy cluster catalog produced by future \\emph{Euclid}-like surveys. Specifically,  we focused attention on the total mass of neutrinos as derived by a combination of cluster number counts and the shape of the cluster power spectrum, including BAO information. In addition to the cosmological parameters of the standard $\\Lambda$CDM model and the total neutrino mass, we also considered non-vanishing curvature, as well as a dynamical evolution of the dark energy equation of state. We marginalised our constraints over the growth function, the amplitude of the scalar curvature power spectrum, and a set of nuisance parameters intended to represent the uncertainties in assigning true masses to galaxy clusters. In order to make our analysis more complete, we also added priors on cosmological parameters coming from the ongoing CMB experiment \\emph{Planck}. Our main results can be summarised as follows. \\begin{itemize} \\item Constraints coming from cluster counts and power spectrum shape+BAO separately are in general relatively weak. Errors on parameter estimated using the cluster correlation function alone are up to two orders of magnitude larger than what is obtained with the galaxy correlation function, due to the lower spatial number density of clusters. Number counts alone do better than this only for the matter density parameter and the present-day dark energy equation of state, while they perform worst in all other cases. Specifically, the error on the total neutrino mass is $\\sigma(M_\\nu) = 2.6$ eV for cluster correlation function alone and $\\sigma(M_\\nu) = 28.5$ eV for cluster counts alone in the context of a general cosmological model including curvature and evolving dark energy. \\item Combining the cluster power spectrum shape+BAO with cluster counts exploits the advantages of both approaches, and hence greatly improves the constraints of individual probes. The errors on cosmological parameters estimated by the combination of the two tests are reduced in general very substantially for all parameters, and by up to an order of magnitude for some of them. For a general cosmology, the error on the total neutrino mass is brought down to $\\sigma(M_\\nu) = 0.9$ eV, comparable with present constraints from the Ly$\\alpha$ forest \\cite{Matteo1} and CMB in non-flat models \\cite{Smith_etal2011}. \\item Constraints on cosmological parameters are also greatly tightened by either $i)$ adopting \\emph{Planck} priors, and $ii)$ reducing the number of free parameters by shrinking the cosmological model to a standard $\\Lambda$CDM one. For a general cosmology, adding \\emph{Planck} priors to the cluster power spectrum shape+BAO reduces the parameter errors by $1-2$ orders of magnitude, except for the dark energy parameters, while performing the same operation on the cluster counts also improves the constraining power by more than one order of magnitude. Moreover, these priors help to remove many of the degeneracies between cosmological parameters, thus reducing the error on the total neutrino mass from either probes down to $\\sigma(M_\\nu)\\sim 0.25$ eV, which goes down to $\\sigma(M_\\nu) \\sim 0.08-0.17$ if only a $\\Lambda$CDM background is considered. \\item The combination of cluster counts, power spectrum shape+BAO, and \\emph{Planck} priors produces very competitive constraints. In a $\\Lambda$CDM context, the errors on all cosmological parameters are compatible with those derived for a general cosmology from the galaxy power spectrum shape+BAO performed for the spectroscopic catalog of a future \\emph{Euclid}-like survey in combination with \\emph{Planck}. Errors on the total neutrino mass are now down to $\\sigma(M_\\nu) = 0.18$ eV, and $\\sigma(M_\\nu) = 0.076$ eV if the cosmological model is enforced to be standard $\\Lambda$CDM. The latter value is only $1.5$ times the minimum neutrino mass admitted by neutrino oscillation experiments. \\item Fixing the value of the nuisance parameters, which equals assuming a perfect knowledge of cluster true masses, also greatly improves parameter constraints. It is quite unlikely that this level of precision can be reached for all clusters in the \\emph{Euclid} photometric catalog, while it might be a reasonable assumption for a subsample of objects having an extensive multi-wavelength follow-up. Nonetheless, if this will indeed be the case, the errors on all estimated parameters would be reduced by up to a factor of a few. Specifically, the error on the total neutrino mass would go down to $\\sigma(M_\\nu) = 0.1$ eV for the general cosmology and to only $\\sigma(M_\\nu) = 0.034$ eV for the $\\Lambda$CDM one, which means that the minimum neutrino mass could be detected. \\end{itemize} \\noindent We emphasize that the analysis presented here is quite conservative from the point of view of the cosmological information, in that it does not include any guidance from the growth function in the observed galaxy cluster power spectrum, the amplitude of the scalar curvature power spectrum, and the redshift space distortions. At the same time, we assumed throughout this work a Gaussian distribution of initial density and curvature perturbations. It has been repeatedly shown that the primordial non-Gaussianity has a substantial impact on both the number counts of massive galaxy systems and their correlation function \\cite{Slosar_etal2008,Grossi1,Grossi2,Grossi3,Pillepich,Desjacques,Dalal,MaggioreRiotto,MV08,Sheth,Damico,McDonald}. Inclusion of this primordial non-Gaussianity would have the effect of adding at least another parameter to the study, hence loosening somewhat the constraints that have been found here \\cite{olga}.  The present investigation shows in a neat way that, whilst taken at face value the galaxy correlation function is superior to the cluster correlation function, the combination of the two, and of the latter with cluster number counts, can be of invaluable help in breaking the degeneracies between, and hence effectively improving the constraints on, cosmological parameters. Therefore galaxy clusters are expected to play a key role for precision cosmology with future large galaxy surveys, addressing even questions related to fundamental physics such as the value of the neutrino mass."
    },
    "1112/1112.4217_arXiv.txt": {
        "abstract": "Absolute spectrophotometric measurements of diffuse radiation at 1 $\\mu$m to 2 $\\mu$m are crucial to our understanding of the radiative content of the Universe from nucleosynthesis since the epoch of reionization, the composition and structure of the Zodiacal dust cloud in our solar system, and the diffuse galactic light arising from starlight scattered by interstellar dust.  The Low Resolution Spectrometer (LRS) on the rocket-borne Cosmic Infrared Background Experiment (CIBER) is a $\\lambda / \\Delta \\lambda \\sim 15-30$ absolute spectrophotometer designed to make precision measurements of the absolute near-infrared sky brightness between 0.75 $\\mu$m $< \\lambda <$ 2.1 $\\mu$m.  This paper presents the optical, mechanical and electronic design of the LRS, as well as the ground testing, characterization and calibration measurements undertaken before flight to verify its performance.  The LRS is shown to work to specifications, achieving the necessary optical and sensitivity performance.  We describe our understanding and control of sources of systematic error for absolute photometry of the near-infrared extragalactic background light. ",
        "introduction": "\\label{S:intro}  A measurement of the extragalactic background light (EBL) determines the integrated emission of all photon sources since the early universe. At near-infrared (NIR) wavelengths, the dominant physical process thought to be responsible for the generation of photons is nucleosynthesis in stars and gravitational energy released in active galactic nuclei (AGN). As most stars reside in galaxies, the brightness of the extragalactic cosmic NIR background (CNIRB) may be constrained by integrating the light resolved into discrete sources along a line of sight through the cosmos (e.g.~\\citet{Mad00, Tot01, Faz04, Keenan10}). However, since unresolved or faint sources of emission may also contribute to the CNIRB, such source counts are necessarily lower limits to the total emission.  The most complete way to determine the total CNIRB brightness is direct absolute photometry measurements  using suitably designed instruments, and the difference between absolute photometry measurements and source counts could reveal a diffuse background from the epoch of reionization. Because atmospheric emission at NIR is 100 times more than the total sky brightness, space-borne observations are required for absolute photometry.  Measurements from the Diffuse Infrared Background Explorer (DIRBE) on the Cosmic Background Explorer ({\\sc COBE}) \\citep{Hauser98, Cam01} and the Near-Infrared Spectrometer (NIRS) on the Infrared Telescope in Space ({\\sc IRTS}) \\citep{Matsumoto05} indicate that the total CNIRB brightness significantly exceeds the brightness determined from deep galaxy number counts \\citep{Mad00, Tot01, Faz04, Keenan10} after subtraction of the zodiacal light (ZL) from sunlight scattered by interplanetary dust grains. For example, the best galaxy counts give $14.7\\pm 2.4$ \\nw\\ at J-band ($1.25 \\, \\mu$m) \\citep{Keenan10} while DIRBE measured $54.0\\pm 16.8$ \\nw\\ at J-band \\citep{Cam01} and IRTS measured $70.1\\pm 13.2$ \\nw\\ at $1.43 \\, \\mu$m \\citep{Matsumoto05} with the DIRBE ZL model \\citep{Kelsall98}. The CNIRB derived from absolute photometry measurements depend critically on the choice of ZL model, and the ``strong no-Zodi'' foreground dust model \\citep{Wright98} produces significantly lower EBL results, $21\\pm 15$ \\nw\\ at $1.25 \\, \\mu$m \\citep{Wright2001, Levenson07}. An excess was also measured in optical bands with a combined Hubble Space Telescope ({\\sc HST}) and ground-based measurement \\citep{Bernstein02, Bernstein05}. However, \\citet{Mattila03} pointed out that their resulting EBL values should be corrected upwards and, because of large systematic errors, should be understood as upper limits only. This was confirmed  by \\citet{Bernstein07} whose reanalysis gave a formal new EBL value of $57\\pm 33$ \\nw\\ at 0.8 $\\mu$m. The cause of this discrepancy is unclear; possibilities range from the prosaic, for example residual ZL \\citep{Dwe05}, to the profound, such as Lyman-$\\alpha $ emission from the first stars \\citep{San02, Sal03, Cooray04, Madau05, Dwe05, Fernandez06}.  On the other hand, a recent re-analysis of Pioneer 10/11 data from outside of the zodiacal cloud gives $7.9\\pm 4.0$ \\nw\\ at 0.44 $\\mu$m and $7.7\\pm 5.8$ \\nw\\ at 0.64 $\\mu$m \\citep{Matsuoka11}, and a recent result from the dark cloud shadow method gives $7.2^{+4}_{-2}$ \\nw\\ at 0.40 $\\mu$m and $< 12$ \\nw\\ at 0.52 $\\mu$m \\citep{Mattila11}, which are more consistent with the source counts at optical wavelengths. In addition, indirect measurements of the CNIRB from TeV-energy $\\gamma$-ray attenuation via pair production are consistent with the source counts \\citep{Dwek05, Aha06, Aha07, Mazin07, Raue09}, and these results dispute the theory that a significant fraction of CNIRB comes from EBL. In these analysis, the authors calculated attenuation spectra of the blazars in various CNIRB levels by assuming their intrinsic spectra as $dN/dE \\propto E^{-\\Gamma }$, and compared with observed spectra. The general result is that a CNIRB of $>20$ \\nw\\ requires  $\\Gamma <1.5$, which is not consistent with estimates of intrinsic blazar spectra either from theory or from observations of nearby blazars \\citep{Aha06}. These are all viable methods of constraining the brightness of the CNIRB and thus the amplitude of the CNIRB and its origin are controversial. However, there is no substitute for direct photometry if the instrumental and astrophysical systematic effects can be mitigated with improved knowledge of local foregrounds. Spectroscopy at wavelengths 1 $\\mu$m to 2 $\\mu$m is especially important because the CNIRB from reionization is predicted to peak near 1 $\\mu$m and fall in brightness at shorter wavelengths. It is clear that our knowledge of the NIR EBL must improve in order to better constrain the Universe's total emission and thereby models of structure formation and galaxy evolution.  In addition to the EBL science, low resolution absolute spectroscopic measurements of the background from above the Earth's atmosphere allow determinations of the ZL and the diffuse galactic light (DGL) from starlight scattered by interstellar dust grains, neither of which are well measured near $1 \\, \\mu$m.  Both of these sources of emission are truly diffuse in nature and are astrophysically interesting in their own right.  For example, since the Zodiacal cloud is the nearest analog to extrasolar dust clouds and debris disks, understanding the local Zodiacal dust environment allows inferences to be made about the nature of distant systems.  The Low Resolution Spectrometer (LRS) on the Cosmic Infrared Background ExpeRiment (CIBER) \\citep{Bock06} is specifically designed to perform absolute spectrophotometric measurement at NIR wavelength (0.75 $\\mu$m $< \\lambda <$ 2.1 $\\mu$m). The LRS has a field of view (FoV) $5.5^{\\circ}$ along a slit and spectral resolution $\\lambda / \\Delta \\lambda \\sim$15-30.  The LRS has successfully flown on CIBER twice, and has yielded good results from both flights: for example, the LRS previously detected an unknown silicate feature in the ZL \\citep{Tsumura10}. Two additional flights of essentially the same configuration of LRS are planned. In this paper we present the LRS instrument in detail, concentrating on its optical, mechanical and electrical characteristics\\footnote{All uncertainties in the paper are standard uncertainties unless otherwise noted.}. In Section \\ref{S:instrument} we review the physical properties of the LRS, and in Section \\ref{S:lab} the laboratory testing and characterization of the LRS is presented.  Finally, we summarize the LRS instrument and its characterization in Section \\ref{S:summary}. ",
        "conclusions": "\\label{S:summary}  The LRS is one of three CIBER instruments. The primary scientific motivation for the LRS measurement is to determine the absolute spectrum of the diffuse sky brightness. The LRS is a 5-cm refracting telescope with a prism, operating at 0.75 $\\mu$m $< \\lambda <$ 2.1 $\\mu$m with a spectral resolution of $\\lambda /\\Delta \\lambda \\sim$15-30. We evaluated the performance of LRS in the laboratory as described in this paper and found the following:  \\begin{enumerate} \\setlength{\\itemsep}{-1pt} \\item The noise performance was consistent with design predictions. \\item A linearity correction up to 1.5\\% is applied to compare the flight data with the laboratory calibration. \\item Image persistence was confirmed and is detectable following observations of bright stars at the level of $\\sim$ 0.2 \\%. \\item A small electrical cross-talk of the incident signal injects $\\sim$ 0.35 \\% into other quadrants. \\item The best fit PSF size is consistent with one pixel and the encircled energy in the main pixel was $>$80\\%. \\item The modified baffling scheme for stray thermal emission removal provided a dramatic improvement for angles $> 20^{\\circ}$, and the stray light signal from the Earth is negligible for the range of angles observed in flight. \\item The calibration uncertainty is approximately $\\pm$5\\%. \\end{enumerate}  Combining all instrumental uncertainty described in this paper and the read noise obtained in flight gives a demonstrated sensitivity of $< 10$ \\nw\\ for a 25 second integration, 3$\\sigma $, 1000 pixels as shown in Figure~\\ref{limit}. The dominant component of uncertainty, $\\pm$5\\%, is from the instrument calibration. In the actual astronomical data analysis, there are other astronomical error elements rather than the instrumental systematics. The dominant error elements from astronomy come from the estimation of the ZL and DGL brightness to subtract. Details of such astronomical error evaluation will be discussed in future science papers.  The first result of LRS from CIBER 1st flight was summarized and published in \\citet{Tsumura10} based on the characterizations described in this paper, although some of them have been modified from the first flight to improve its performance. The 2nd flight data is under analysis now and the result will be reported in forthcoming papers. Owing to the modifications especially to combat the thermal stray light, data with good quality enough to obtain the CNIRB spectrum was acquired in the 2nd flight."
    },
    "1112/1112.2211_arXiv.txt": {
        "abstract": "We have conducted a long-term, wide-field, high-cadence photometric monitoring survey of $\\sim$50,000 stars in the Lagoon Nebula \\ion{H}{2} region. This first paper presents rotation periods for 290 low-mass stars in NGC 6530, the young cluster illuminating the nebula, and for which we assemble a catalog of infrared and spectroscopic disk indicators, estimated masses and ages, and X-ray luminosities. The distribution of rotation periods we measure is broadly uniform for $0.5 < P < 10$~d; the short-period cutoff corresponds to breakup. We observe no obvious bimodality in the period distribution, but we do find that stars with disk signatures rotate more slowly on average. The stars' X-ray luminosities are roughly flat with rotation period, at the saturation level ($\\log L_X / L_{\\rm bol} \\approx -3.3$). However, we find a significant positive correlation between $L_X / L_{\\rm bol}$ and co-rotation radius, suggesting that the observed X-ray luminosities are regulated by centrifugal stripping of the stellar coronae. The period--mass relationship in NGC 6530 is broadly similar to that of the Orion Nebula Cluster (ONC), but the slope of the relationship among the slowest rotators differs from that in the ONC and other young clusters. We show that the slope of the period--mass relationship for the slowest rotators can be used as a proxy for the age of a young cluster, and we argue that NGC~6530 may be slightly younger than the ONC, making it a particularly important touchstone for models of angular momentum evolution in young, low-mass stars. ",
        "introduction": "\\label{sec:intro}  Time-domain photometric monitoring surveys of young stars have been crucial to our understanding of a variety of fundamental questions related to low-mass stars in the pre--main-sequence (PMS) phase of evolution. Indeed, our empirical understanding of time-variable accretion processes \\citep[e.g.][]{gullbring96,bouvier99,stassun99b,bouvier04}, of magnetic activity \\citep[e.g.][]{walter87,feigelson02,feigelson05,feigelson07,walter04,stassun04b,stassun06,stassun07b}, of the early evolution of stellar angular momentum \\citep[e.g.][]{attridge92,bouvier93,choi96,stassun99a,rebull01,herbst02,makidon04,lamm05,irwin08a,irwin08b}, of the inner architectures of protoplanetary disks \\citep[e.g.][]{bouvier03,winn06,bouvier07,herbst08,herbst10}, of the formation of binary stars \\citep[e.g.][]{mathieu97,basri97,jensen07,stassun08}, of outbursts \\citep[e.g.][]{walter04,briceno04,kastner04,grosso05,aspin06,aspin11,covey11,bastien11}, of pulsations \\citep[e.g.][]{zwintz06,guenther07}, and of the fundamental masses and radii of PMS stars \\citep[e.g.][]{casey98,covino00,stassun04a,stassun06a,stassun07a,irwin07,stempels08,stassun08,hebb10}, has relied upon detailed light curves of T~Tauri stars (TTSs) in a variety of young clusters and star-forming regions spanning a range of ages and star-forming environments.  The most extensively monitored regions include the Taurus-Auriga association (age $\\sim$1--3 Myr), the Orion Nebula Cluster ($\\sim$1--2 Myr), NGC 2264 ($\\sim$3 Myr), NGC 2362 ($\\sim$3--4 Myr), IC 348 ($\\sim$4--5 Myr), and NGC 2547 ($\\sim$40 Myr). At a distance of $\\sim$400 pc \\citep[e.g.][]{menten07}, the Orion Nebula Cluster (ONC) and the larger star-forming region surrounding it is the nearest and perhaps the single best-studied massive star-forming region. This region contributes photometrically-determined rotation periods for hundreds of TTSs \\citep[see][and references therein]{herbst07}, nearly all of the known PMS eclipsing binary stars \\citep[see][and references therein]{stassun09}, including the only brown-dwarf eclipsing binary system yet discovered \\citep{stassun06a,stassun07a,gomez09,mohanty09,mohanty10}, and the most well-studied FUor/EXor type eruptive star discovered in recent times \\citep[V1647 Ori; see, e.g.,][and references therein]{bastien11}. Thanks in large part to the broad array of discoveries enabled by the extensive photometric monitoring surveys of this region, the ONC has become a crucial testbed for star formation theory, from PMS angular momentum evolution \\citep[see][]{stassun03,herbst07} to the nature of the initial mass function \\citep[see][]{hillenbrand97,dario10}. As wide-area photometric monitoring campaigns begin to survey a larger number of star-forming regions, particularly rich young clusters like the ONC, the discovery space is becoming enlarged for larger samples of benchmark PMS objects. For example, the recent PTF surveys of the 25~Ori and North America/Pelican nebulae have already resulted in the discovery of numerous new candidate PMS eclipsing binaries and new FUor/EXor outburst systems \\citep{vaneyken11,covey11}.  In this paper we report the first results of a large-scale, multi-year, high-cadence photometric monitoring survey of the bright Lagoon Nebula (Messier 8) \\ion{H}{2} star-forming region. The massive star cluster illuminating the nebula, NGC 6530, includes a rich population of $\\gtrsim$1100 stellar members spanning the full IMF, from massive OB type stars down to at least the hydrogen burning limit \\citep{damiani04,prisinzano05,damiani06,prisinzano07}. With a nominal age of $\\sim$1 Myr \\citep[e.g.][]{mayne07}, NGC 6530 is thus in many respects an analog of the ONC, but at a distance of $\\sim$1.25 kpc NGC 6530 provides an ONC-like stellar laboratory beyond the immediate solar neighborhood.  An extensive literature has emerged over the past few years, characterizing the PMS population of NGC 6530 from X-ray to infrared wavelengths.  \\citet{damiani04} conducted the first large-scale study of the stellar population of NGC 6530. Using {\\it Chandra} they detected 884 X-ray point sources, finding that 90--95\\% of them constitute cluster members. \\citet{prisinzano05} subsequently performed a complementary deep optical survey of the region, obtaining $BVI_C$ photometry down to $V\\approx 23$ using the Wide Field Imager at the MPG/ESO 2.2m telescope. They matched their catalog to that of \\citet{damiani04} and found 828 common stars, the vast majority of which are cluster members. From their deep color-magnitude diagrams (CMDs) and X-ray membership selection, they determined a cluster distance of 1.25 kpc and a modest cluster extinction of $A_V = 1.1$ mag, and they moreover estimated masses and ages for the stars through comparison with the PMS evolutionary tracks of \\citet{siess00}. \\citet{damiani06} performed a near-infrared (NIR) survey of the cluster to identify stars with NIR excess emission indicative of objects bearing massive protoplanetary disks, and in the process increased the number of PMS cluster members to more than 1100. \\citet{prisinzano07} spectroscopically studied a subsample of 332 cluster members, using H$\\alpha$ emission to classify the stars as classical T~Tauri stars (CTTSs) or weak-lined T~Tauri stars (WTTSs).  Our synoptic survey of the photometric variability properties of NGC 6530 builds on these studies. Our light curves of a $40\\arcmin \\times 40\\arcmin$ region centered on NGC 6530 include the known cluster members as well as a total of $\\sim$50,000 other stars within and surrounding the larger Lagoon Nebula star-forming region. These light curves in concert with the extant literature enable an array of variability studies for the region, including measurement of stellar rotation periods, identification of PMS eclipsing binaries, characterization of accretion induced variations and stellar occultations due to disk obscuration, and discovery of eruptive variables. This first paper presents the results of our systematic search for rotation periods among the members of NGC 6530, the first such survey for stellar rotation periods yet reported for this important cluster.  Understanding the evolution of stellar angular momentum, particularly during the first few Myr of a star's life, remains one of the longest outstanding questions in star formation research \\citep[e.g.][]{vogel81,hartmann86,bouvier86}. As PMS stars contract toward the main sequence, they would be expected to rapidly spin up to near breakup velocity as a consequence of angular momentum conservation. However, numerous surveys of rotation periods among low-mass PMS stars clearly show that these stars typically rotate at a small fraction of breakup, despite significant contraction in stellar radius \\citep[e.g.][]{stauffer87}. A variety of mechanisms for efficiently removing angular momentum from PMS stars during the first $\\sim$10 Myr has been proposed, including magnetic star-disk interaction \\citep[i.e., ``disk-locking\";][]{konigl91,shu94,najita95,ostriker95}, scaled-up solar-type magnetized winds perhaps driven by accretion \\citep[e.g.][]{matt04,matt05a,matt05b,matt08a,matt08b}, and scaled-up solar-type coronal mass ejections \\citep[e.g.][]{aarnio09,aarnio10,aarnio11}. However, the observational support for and theoretical efficacy of these mechanisms remains debated. There is not yet a consensus on the dominant mechanism(s) responsible for governing the angular momentum evolution of low-mass PMS stars.  In contrast, by the time a young cluster of stars reaches the age of the Pleiades ($\\sim$125 Myr), the observational picture is much more clear. By Pleiades' age, a cluster of coeval stars develops two distinct populations in period versus mass, now commonly referred to as the ``I'' and ``C'' sequences \\citep{barnes03,barnes07,barnes10a,barnes10b}. These respectively correspond to curves that trace the upper and lower envelopes of stellar rotation periods. At an age of $\\sim$600 Myr, about the age of the Hyades, nearly all stars inhabiting the ``C'' sequence will have spun down and transitioned onto the ``I'' sequence. This general behavior is attributed principally to changes in the stars' internal structures as a function of mass, and has been found to hold true for clusters at a variety of ages \\citep{patten96, barnes99, allain96, krishnamurthi98, radick87, radick90}. While this ``gyrochronology\" paradigm has yet to be extended back to the PMS stage, it is clear that young low-mass stars must at some stage prior to the main sequence develop specific relationships between rotation period and stellar mass, and that these encode the stellar age. Exploiting the very young age of NGC 6530, we are in a position to better define the initial conditions of PMS angular momentum evolution, and to identify when and how the period--mass--age relationship begins to take shape.  In this paper we report rotation periods for 290 members of NGC 6530. We match these to the literature to produce a catalog with which we investigate correlations between rotation period and other stellar properties, including disk presence, mass, age, spatial distribution within the cluster, and X-ray activity. In \\S\\ref{sec:data} we describe the photometric observations and their reduction, data from the literature, and the construction of our catalog. Our period search and statistical methods are discussed in \\S\\ref{sec:analysis}. In \\S\\ref{sec:results} we examine the distribution of rotation periods, and investigate correlations between rotation period and the other stellar properties in our catalog. We discuss the implications of our results in \\S\\ref{sec:disc}, and in \\S\\ref{sec:summary} we summarize the main conclusions. ",
        "conclusions": "\\label{sec:summary}   We have photometrically monitored $\\sim$50,000 stars in a $40\\arcmin\\times40\\arcmin$ field centered on NGC 6530, the young massive star-forming cluster illuminating the Lagoon Nebula, over 35 nights in the $I_C$-band with a cadence of 1~hr$^{-1}$. These observations are intended to complement recent optical, X-ray, and NIR surveys of the region \\citep{damiani04,prisinzano05,damiani06,prisinzano07}, permitting a comprehensive characterization of the young stellar population in NGC 6530.  From an analysis of periodic variations in our light curves, we measured rotation periods for 290 X-ray selected cluster members of NGC 6530, with masses in the range $0.2 < M/{\\rm M}_\\odot < 2.0$. From the findings of \\citet{damiani04}, we expect $\\sim5\\%$ of our catalog to be contaminated by non-members, or only $\\sim$15 stars in our full photometric catalog. We investigated correlations between rotation period and other stellar properties, including mass, age, spatial distribution within the cluster, the presence of circumstellar disks, and X-ray activity. The major findings of this work are as follows: \\begin{enumerate}  \\item The distribution of rotation periods in NGC 6530 is approximately uniform over the range $0.5 < P < 10$~d; we do not observe obvious bimodality in the period distribution, regardless of whether the distribution is considered in its entirety or limited to narrower ranges of stellar mass. The sharp cutoff in the period distribution at $P \\approx 0.5$~d likely results from the breakup limit for the stars in our sample.  A small number of stars with $P<0.5$~d are present, which should be investigated further as possible pre--main-sequence contact binary systems.  \\item The X-ray luminosities of the stars are roughly flat with rotation period, at approximately the saturation level ($\\log L_X / L_{\\rm bol} \\approx -3.3$). However, the fastest rotators show lower average X-ray luminosities, at a highly statistically significant level, suggestive of so-called ``super saturation.\" At the same time, X-ray luminosity correlates most strongly with the stars' co-rotation radii, suggesting that centrifugal stripping of the coronae may be the fundmental driver of the super saturation phenomenon.  \\item Stars with NIR excesses and H$\\alpha$ emission indicative of warm circumstellar material rotate more slowly on average than stars lacking disk signatures. Disked stars might be presumed to be younger on average, and indeed we find evidence that stars with younger ages as inferred from spatial location within the cluster rotate more slowly on average. However, the statistical significance is low, and indeed we find the opposite association between rotation and age when the ages are inferred from PMS isochrones.  \\item The rotation periods are a function of stellar mass: the lower mass stars rotate more slowly on average than the higher mass stars. This is in the opposite sense of the period-mass relationship observed in the ONC and in all other slightly older clusters.  \\item We show that the slope of the mass-period relationship among slow rotators (defined as the 75\\%-ile rotation periods) in the mass range $0.1 < M/{\\rm M}_\\odot < 0.5$ is a good proxy for the age of a young cluster. Calibrating this empirical mass-period-age relation to the ONC, NGC 2264, NGC 2362, IC348, NGC 2547, and NGC 2516, we find that NGC 6530 is the youngest of all, with a maximum age of 1.65~Myr on an age scale where the ONC is 2~Myr.  \\end{enumerate}  The evidence points strongly to NGC 6530 being in a very early stage of rotational evolution in which the stars are currently evolving toward a state that will presumably resemble the ONC within the next $\\lesssim$1~Myr. Thus NGC 6530 becomes an important new touchstone for theoretical models of angular momentum evolution in young, low-mass stars.  An important question that remains to be resolved is the role of circumstellar disks in the rotational evolution of these stars. The observed correlation between NIR excess and slow rotation has been taken in previous studies as evidence for rotational braking via star-disk interaction. However, it is not clear that theories of star-disk interaction in fact predict this correlation. Additionally, recent detailed modeling of the full spectral energy distributions of young stars with rotation periods in IC348 indicate that, if disks do affect the spin rates of the stars, they must act both to spin down some stars and to spin up others \\citep{leblanc11}. A similarly detailed assessment of the disk torques likely being experienced by the stars in NGC 6530 will be important to determine whether and how disks may yet be acting to shape the mass-period relationship in this very young cluster.  Finally, it remains an important challenge to empirically connect the rotational properties of PMS stars to those of main-sequence stars, and to theoretically connect the dominant mechanisms thought to govern the evolution of angular momentum in the PMS to those on the main sequence. Main sequence angular momentum evolution is principally understood through intrinsic structural changes in the stars that, through stellar winds, lead to distinct period-mass relationships that evolve predictably with time and thus permit reliable age-dating of stars (gyrochronology). In contrast, in the PMS phase the dominant angular momentum evolution mechanisms have generally been thought to be extrinsic to the stars (e.g., disk-locking). Yet it is now clear that low-mass stars already evince clear period-mass relationships in the PMS stage. Evidently, as early as the very young age of NGC 6530, a period-mass relationship that can be projected forward to the main sequence is already taking form, and this relationship already encodes stellar age."
    },
    "1112/1112.0155.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory \\abstract % context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) {The purpose of the present investigation is to probe the dust properties inside a molecular cloud, in particular how particle growth and the presence of ice coatings may change the overall shape of the extinction curve.} % methods heading (mandatory) {Field stars behind a molecular cloud can be used to probe the cloud extinction for both the reddening and the absorption features. By combining multi-colour photometry and IR spectroscopy the spectral class of the star can be determined as can the extinction curve, including the vibrational bands of ices and silicates. } % results heading (mandatory) {Based on observations of field stars           %%a M giant star and an A star behind the dark globule B\\,335, we determine the reddening curve from 0.35 to 24\\,$\\mu$m. The water ice band at 3.1\\,$\\mu$m is weaker ($\\tau$(3.1) = 0.4) than expected from the cloud extinction (A$_{V} \\approx $  10 for the sightline to the most obscured star). On the other hand, the CO ice band at 4.7\\,$\\mu$m is strong ($\\tau$(4.67) = 0.7) and indicates that the mass column density of frozen CO is about the same as that of water ice. We fit the observations to model calculations and find that the thin ice coatings on the silicate and carbon grains (assumed to be spherical) \\textit{lower} the optical  extinction by a few percent. We show that the reddening curves for the two background stars, for which the silicate band has been measured, can be accurately modelled from the UV to $24\\ \\mu$m. These models only include  graphite and silicate grains (plus thin ice mantles for the most obscured star), so there is no need for any additional major grain component to explain the slow decline of the reddening curve beyond the K band. As expected, the dust model for the dense part of the cloud has more  large grains than for the outer regions. We propose that the well established shallow reddening curve beyond the K band has two different explanations: larger graphite grains in dense regions and relatively small grains in the diffuse ISM, giving rise to substantially less extinction beyond the K band than previously thought.} % conclusions heading (optional), leave it empty if necessary { For the sightline towards the most obscured star, we derive the relation A$_{Ks}$ = 0.97$\\cdot$E(J$-$K$_{Ks}$), and assuming that all silicon is bound in silicates, N(2\\,H$_2$+H) $\\approx \\  1.5\\cdot 10^{21} \\cdot A_{V}\\ \\approx \\  9\\cdot 10^{21} \\cdot A_{Ks}$. For the rim of the cloud we get  A$_{Ks}$ = 0.51 $\\cdot$E(J$-$K$_s$), which is close to recent determinations for the diffuse ISM. The corresponding gas column density is N(2\\,H$_2$+H) $\\approx \\  2.3\\cdot 10^{21} \\cdot A_{V}\\ \\approx \\  3\\cdot 10^{22} \\cdot A_{Ks}$.} %% ",
        "introduction": "\\textit{Interstellar reddening} is caused by small dust particles, but the detailed optical and chemical properties of these particles remain largely unknown, in particular for molecular clouds. With the exception of the  220 nm extinction peak (probably caused by small graphite grains), %and the weak DIBs (diffuse interstellar bands) of unknown origin, the UV, optical, and near-IR spectral regions offer little information about the dust particles, except that they must be small compared to the wavelength (even in the UV). In the mid-IR, however, strong absorption bands show up, the most universal of these being the broad band centred at 9.7\\,$\\mu$m due to silicates. Another widespread component in the \\textit{diffuse} ISM (interstellar medium), PAHs (polycyclic aromatic hydrocarbons) are seen in the mid-IR as emission (due to luminescence). In addition, the far-IR thermal emission from the interstellar dust grains gives further constraints on their size distribution, and to some extent, their composition.\\\\ \\\\ Combining all this information, \\cite{2003ARA&A..41..241D} has constructed a dust model        %% Draine 2003 that describes the observed properties for the diffuse ISM. It basically consists of silicate and carbon grains. For molecular clouds, the lower temperature, the higher density and a diluted UV radiation field allow ices to form and/or condense on the grains. In addition, the grains coagulate and the size distribution is changed towards larger grains. These processes alter the optical properties of the grains. The most obvious difference compared to the diffuse ISM is the absorption features from ices. The ISO instruments provided extensive spectroscopic evidence of various ice absorption features, including H$_2$O, CO, CO$_2$, CH$_4$, CH$_3$OH and NH$_3$ \\citep{2004ApJS..151...35G}. In most cases, these absorption features have been observed towards luminous, deeply embedded protostars, and it is hard to separate the circumstellar matter (influenced by dynamical and radiation processes involved in the formation of the star) from the unprocessed cloud component. In addition, the intrinsic SED (spectral energy distribution) of an embedded young star is not in general well known, which means that it is difficult to accurately separate interstellar reddening from intrinsic IR excess. For these reasons, field stars behind a molecular cloud are better probes in investigations of the (more or less) undisturbed cloud. The problem is to find background stars that are heavily obscured to probe the dense regions and still bright enough for IR spectroscopy. In spite of these difficulties, there are some investigations of ice features using background stars: Elias 16 by \\cite{2004ApJS..151...35G} and Spitzer ice-data from three sources behind Taurus and CK2 in Serpens by \\cite{2005ApJ...635L.145K}. More recently, \\cite{2011ApJ...729...92B} has combined ground-based and Spitzer spectroscopy data to analyse the ice and silicate features in quiescent clouds by observing background field stars. \\\\  Apart from the absorption bands, the continuous extinction in the mid-IR for molecular clouds has been the subject of several recent investigations \\citep[e.g.][]{2007ApJ...663.1069F, 2007ApJ...664..357R, 2009ApJ...690..496C, 2011ApJ...729...92B}. However, it is important to note that what is actually determined is the {\\it reddening curve}, and to establish the {\\it extinction curve}, an assumption is made on the extinction ratio for two wavelengths, e.g. A$_H$/A$_K$ = 1.55, which goes back to the frequently cited paper by \\cite{1985ApJ...288..618R}. Their paper combined earlier investigations with mid-IR observations of distant field stars of unknown intrinsic colours. Obviously the data available for their analysis was very limited compared to those in recent investigations, based on large surveys both from ground and space. The reddening curve is relatively straight forward to determine, either by statistical methods (slopes in colour/colour diagrams) or by classifying selected background stars and compare the observed colour to the intrinsic. By observing stars all the way to Spitzer/MIPS 24 $\\mu$m one would expect that the extinction is so close to zero at this wavelength, that it for all practical purposes does not matter if it is assumed to be exactly zero. However, the observations show that the reddening curve levels off in the mid-IR, which is not expected from current dust models, and for this reason the extrapolation to zero wave number remains uncertain. \\\\ In our recent paper \\citep[][hereafter called Paper\\,II]{2010A&A...522A..84O} we confirmed that the reddening curve for molecular clouds in the wavelength range 0.35--2.2 $\\mu$m can be characterized in a functional  form described by \\cite{1989ApJ...345..245C} (henceforth CCM) with only one parameter,  $R_{V}\\,=\\,A_{V}/E_{B-V}$. It was, however, clear that the CCM function did not adequately describe the reddening curve beyond the K band, and in the present paper we extend the investigation into the mid-IR spectral region. For our case study we have compiled our own and archive data for a few stars behind the B\\,335 globule including photometry and spectroscopy from the optical region to 24\\,$\\mu$m. Our main purpose is to characterize the mid-IR extinction. In addition we explore whether both the reddening curve, from UV to 24\\,$\\mu$m, and the absorption features can be explained in a single model. \\\\ \\\\ \\begin{table*} \\begin{minipage}[t]{\\columnwidth} \\caption{Observations and data.} \\label{obs} \\renewcommand{\\footnoterule}{}  % to avoid a line before footnotes \\begin{tabular}{llllll} \\hline %\\textbf{\\textbf{band}}             & U602  & Bb605  & g772  & V606  & r773  & I610  & 2massJ & 2massH & 2massKs&$[3.6\\mu]$&$[5.8\\mu]$\\\\ %$\\lambda\\ \\ \\ [nm]$                & 354.0 & 413.2  & 508.9 & 542.1 & 673.3 & 798.1 & 1258.5 & 1649.7 & 2157.2 &  & \\\\ %$\\Delta\\lambda\\ [nm]$              &  53.7 & 109.4  &  75.3 & 104.8 &  81.1 & 155.2 &  300.  &  300.  &  300.  & 800. & 1500.\\\\ \\hline \\textbf{observation}&\\textbf{band}& \\textbf{observatory} & \\textbf{wavelengths $\\mu m$}& \\textbf{observation date}& \\textbf{reference}\\\\ \\hline imaging          & optical      & NTT-EMMI     & [0.354,\\,0.413,\\,0.509,\\,0.673,\\,0.798] & 2006-06-27--29 \\\\ % & 077.C-0524(A) \\\\ spectroscopy     & optical      & NTT-EMMI     & $0.5\\,<\\lambda<\\,0.8$                   & 2006-06-27--29 \\\\ % & 077.C-0524(A) \\\\ imaging          & NIR          & 2MASS        & [1.258,\\,1.649,\\,2.157]                 &                  & \\cite{2006AJ....131.1163S}\\\\  %% Skrutskie 2006 imaging          & L'           & NOT-SIRCA    & 3.79                                    & 2003-05-17       & unpublished \\\\ spectrometry     & H$_2$O\\,ice  & VLT-ISAAC    & $2.55\\,<\\lambda<\\,4.2$                  & 2004-04-12--13 \\\\    %& 073.C-0436A.1\\\\ spectrometry     & CO\\,ice      & VLT-ISAAC    & $4.45\\,<\\lambda<\\,5.1$                  & 2004-04-12--13 \\\\    %& 073.C-0436A.1\\\\ %% %imaging          & NIR          & Spitzer-IRAC & $3.6,\\ 4.5,\\ 5.8,\\ 8.0$                 & 2004-04-20       & AOR:4926208 IRAC-0000-CB199\\\\ %imaging          & NIR          & Spitzer-IRAC & $3.6,\\ 4.5,\\ 5.8,\\ 8.0$                 & 2004-04-21       & AOR:4926464 IRAC-0000-CB199\\\\ %%%imaging          & NIR          & Spitzer-MIPS & 24.0                                    & 2004-10-16       & AOR:12022016 MIPSC-CB199(B335)\\\\ %%% spectrometry     & silicate     & ISOCAM-CVF   & $5\\,<\\lambda<\\,16.5$                    & 1997-04-25       & Rev-TDT-OSN: 526 003 74\\\\ spectrometry     & silicate     & Spitzer-IRS  & $5\\,<\\lambda<\\,21 $                     & 2005-10-11       & AOR:10705664 VeluCores5\\\\ CO isotopologues & mm           & ESO-SEST     & $^{13}$CO and C$^{18}$O   1--0          & 1993 -- 1998     & Harjunp\\\"{a}\\\"{a} et al (2004)\\\\ %%\\footnote{\\cite{2004A&A...421.1087H}}\\\\               %%Harjunp\\\"{a}\\\"{a} et al, 2004\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} %%ISOCAM/CVF %%Summary Details %%--------------- %% %%Target Name       :B335 CVF %%Inst/AOT          :CAM04 %%Position          :294.254122 7.569690 %%Dimensions        :f.o.v. 192\"x192\" %%Wavelengths       :5.01-16.77\u00b5m %%Rev-TDT-OSN       :526 003 74 %%Observer          :SCABRIT %%Proposal          :CAMFLOWS %%Start Time        :1997-04-25 05:20:07.245 %%End Time          :1997-04-25 06:48:07.245 %%On Target Time [s]:5280 sec %%Quality           :Good ",
        "conclusions": "Based on photometry in the range 0.35 - 24 $\\mu m$ as well as IR spectroscopy we determine the reddening curve towards stars behind the dark globule B\\,335. For the most obscured star we also determine the absorption bands due to water ice, CO ice and silicates. - We  confirm earlier findings that the reddening curve levels off beyond the K band in contrast to current models \\citep[e.g.][]{2001ApJ...548..296W}. From the reddening curve we also draw the following conclusions concerning the grain size distribution, the extinction and the ice bands.\\\\ \\\\ \\texttt{Grain size distribution.}\\\\ \\emph{The grain size distribution features.} Based on a relatively simple assumption for the grain size distribution (a power law and a Gaussian) for silicate and graphite grains we find solutions that closely match the observed reddening curves for the two background stars for which we have silicate band observations. One of the stars probe the central part of the globule and the other the rim. For the most obscured star the model fit has a peak of graphite grains at a radius of 2\\,$\\mu$m, while the graphite peak for the other star is at 0.5\\,$\\mu$m. In neither case there is any need for an additional grain material to explain the reddening curve.\\\\ \\\\ \\emph{The abundance ratio graphite/slicates} differs substantially for the two models (much more graphite in the central region). As a consequence we have two different explanations of the shallow reddening curve beyond the K band: for the denser region it is because of larger graphite or, more general, carbon grains which give rise to a higher mid-IR extinction than predicted by previous models constructed to match the UV to near-IR spectral range \\citep[cf.][]{2001ApJ...548..296W}. For the outer region of the globule the explanation is simply that the {\\it extinction as such is low beyond the K band}. Actually, much of the extinction towards the stars at the rim of the globule may originate in the diffuse medium beyond the cloud, and we suggest that this explanation to the shallow reddening curve beyond the K band is valid not only for the outer regions of the cloud but also for the diffuse ISM.\\\\ \\\\ \\emph{Gas column density.} The column densities ratios $N(2\\cdot H_{2}\\ +\\ H)/A_{V}$ are estimated to be $2.3\\cdot 10^{21}$ for the lower extinction sightline and $1.5\\cdot 10^{21}$ for the higher extinction sightline.\\\\ \\\\ \\texttt{Extinction model.}\\\\ The extinction model for \\emph{the outer part} of  B\\,335 gives the relation  A$_{Ks}$ = 0.51 $\\cdot$E(J$-$K$_s$). This latter relation agrees with recent determinations towards the galactic centre by  \\cite{2006ApJ...638..839N} and \\cite{2011ApJ...737...73F} using different methods. \\cite{2005ApJ...619..931I} get a higher value for the diffuse ISM, A$_{Ks}$ = 0.67$\\cdot$E(J$-$K$_s$), but we show that this difference is not significant and conclude that A$_{Ks}$ $\\approx$ 0.51 $\\cdot$E(J$-$K$_s$) for the diffuse ISM as well as for moderately dense molecular regions. \\\\ \\\\ The  extinction model for \\emph{the denser part} of the B\\,335 gives the relation A$_{Ks}$ = 0.97 $\\cdot$E(J$-$K$_s$). Although this model describes the reddening curve from the UV to 24\\,$\\mu$m, including the silicate, the water-ice, and the CO-ice bands well, we have indications that it represents conditions that may be different from other dense cores. \\\\ \\\\ \\texttt{Ice bands.}\\\\ The \\emph{water ice} band is reproduced by grains coated with pure water ice. The column density of water ice is half of that found in the Taurus clouds (\\cite{2001ApJ...547..872W}'s) for the same extinction.                              %% Whittet et al 2001 The difference is likely due to a larger interstellar radiation field in the B\\,335 region, desorbing ice mantles deeper in the cloud.\\\\ \\\\ The \\emph{CO ice} band differs in position from that of pure CO ice mantles, indicating that CO is mixed in other ices. Based on published laboratory measurements of frozen gas mixtures we show that the CO ice band can be understood if CO is part of such mixtures (where H$_{2}$O is a natural part and N$_{2}$ a possible part). The ice/gas ratio is relatively high (0.8 - 1.) and as this is an average along the $A_{J}\\ =\\ 3.5$ cloud extinction towards the probing star and as we must allow for relatively deep regions free from CO ice on both the front and the rear side of the cloud, the ice/gas CO ratio must be greater than unity in the interior of the cloud, and more than half of the gas is depleted. This confirms earlier findings \\citep{2002A&A...389L...6B} that CO (or CO isotopologues) cannot be used as a mass tracer in dense and cold molecular clouds.\\\\ \\\\ The \\emph{CO$_{2}$ ice} band at 15.2 $\\mu $m indicated in the ISOCAM-CVF spectrum is not confirmed in the Spitzer spectra and we can only give an upper limit to the CO$_{2}$-ice column density.\\\\ \\\\ %%\\begin{table} %%\\caption{Column densities based on Si abundance and all Si in silicates.} %%\\label{coldens} %%\\renewcommand{\\footnoterule}{}  % to avoid a line before footnotes %%\\begin{tabular}{rrrrr} %%\\hline \\hline %%\\textbf{star\\#}  &  \\textbf{A$_V$}   & \\textbf{N(2 H$_2$ + H)/A$_V$} & \\textbf{A$_{K_s}$} & \\textbf{N(2 H$_2$ + H)/A$_{K_s}$}\\\\ %%                 &                   & $\\cdot 10^{-21}\\ [cm^{-2}]$ &                    & $\\cdot 10^{-21}\\ [cm^{-2}]$\\\\ %%\\hline\\\\ %%        2        &  4.8              &  2.30                       & 0.35           & 30.8\\\\ %%        947      &  9.6              &  1.53                       & 1.70           & 8.6\\\\ %%\\hline \\hline %%\\end{tabular} %%\\end{table}"
    },
    "1112/1112.0678_arXiv.txt": {
        "abstract": "Some direct detection experiments have recently collected excess events that could be interpreted as a dark matter (DM) signal, pointing to particles in the $\\sim$10 GeV mass range. We show that scenarios in which DM can self-annihilate with significant couplings to quarks are likely excluded by the cosmic-ray (CR) antiproton data, provided the annihilation is S-wave dominated when DM decouples in the early universe. These limits apply to most of supersymmetric candidates, \\eg~in the minimal supersymmetric standard model (MSSM) and in the next-to-MSSM (NMSSM), and more generally to any thermal DM particle with hadronizing annihilation final states. ",
        "introduction": "\\label{sec:intro} After the finding of an annual modulation consistent with a DM signal in the DAMA/Libra data \\cite{2000PhLB..480...23B,2008EPJC..tmp..167B}, some skepticism had arisen due to the failure of other direct detection experiments to confirm it in the preferred phase space corner. Interestingly enough, this initial caution has been recently tickled by the results of a few other experiments with different setups. CoGeNT \\cite{2011PhRvL.106m1301A} and CRESST \\cite{2011arXiv1109.0702A} have indeed found excess events; CoGeNT has even released hints for an annual modulation \\cite{2011PhRvL.107n1301A}. Yet, XENON-10/100 \\cite{2011PhRvL.107e1301A,2011PhRvL.107m1302A}, CDMS \\cite{2011PhRvL.106m1302A} and SIMPLE \\cite{2011arXiv1106.3014F} do not validate such excesses and therefore challenge the classical DM interpretation of the aforementioned positive results --- \\ie~the elastic collisions of weakly massive interacting particles (WIMPs) off target nuclei \\cite{2011arXiv1110.2721K}. It is still important to recall that these excesses --- and bounds --- rely on data analyses led very close to the current experimental thresholds, where systematic errors can in principle get large and remain to be much better controlled. Nevertheless, these measurements have brought some excitement in the DM community, and many authors have tried to fit the available data with some phenomenological models (\\eg~\\cite{2010PhRvD..81j7302B}). The preferred mass range is around 10 GeV, while difficulties are generally found when trying to get good fits to the signals while avoiding the constraints from other direct detection experiments \\cite{2011arXiv1110.2721K}. Some ideas have been proposed to explain such discrepancies (\\eg~\\cite{2011JCAP...09..022A,2011PhRvD..84d1301F}), but what people usually do is to forget about either the bounds or the signals, depending on whether they want to promote their models or challenge them.   Beside direct detection, indirect detection, that looks for DM annihilation or decay traces originating in high concentration astrophysical sites, provides an interesting way to check the relevance of any phenomenological proposal in this context. Among indirect signatures, CR antiprotons have long been known as an interesting target to look at~\\cite{1984PhRvL..53..624S}. In this proceedings' contribution, we show that they indeed provide quite strong limits on quarkophilic DM candidates, and are even likely excluding those models which have an S-wave dominated annihilation cross section. The following discussion is mostly based upon two articles, Refs. \\cite{2010PhRvD..82h1302L} and \\cite{2012NuPhB.854..738C}, to which we refer the reader for more details and references. ",
        "conclusions": "\\label{sec:concl} We have discussed how CR antiprotons can put limits on light DM candidates. They are generically strong and exclude quarkophilic DM particles in the 5-15 GeV mass range relevant to the direct detection signals, if they annihilate through S-waves with $\\sigv \\geq 10^{-26}\\,{\\rm cm^3/s}$. The only constraint at the origin of this result is the upper bound on the cosmological DM abundance, which translates into a lower limit on the annihilation cross section at the decoupling time, $\\sigv \\gtrsim 2.5\\times 10^{-26} \\,{\\rm cm^3/s}$. Light neutralinos arising in the MSSM with non-unified gaugino masses and in the NMSSM are therefore excluded as potential DM candidates if they fall in the quarkophilic categories discussed above, which is quite often the case. These antiproton limits are complementary to other studies using \\eg~gamma-rays \\cite{2011arXiv1108.2914G}, CMB photons \\cite{2011PhRvD..84b7302G}, and solar neutrinos \\cite{2011NuPhB.850..505K}, which concern more annihilation final states than only quarks."
    },
    "1112/1112.1162_arXiv.txt": {
        "abstract": "{ The light curve of the quasar OJ~287 extends from 1891 up today without major gaps. This is partly due to extensive studies of historical plate archives by Rene Hudec and associates, partly due to several observing campaigns in recent times. Here we summarize the results of the 2005 - 2010 observing campaign in which several hundred scientists and amateur astronomers took part. The main results are the following: (1) The 2005 October optical outburst came at the expected time, thus confirming the general relativistic precession in the binary black hole system. This result disproved the model of a single black hole system with accretion disk oscillations, as well as several toy models of binaries without relativistic precession. In the latter models the main outburst would have been a year later. (2) The nature of the radiation of the 2005 October outburst was expected to be bremsstrahlung from hot gas at the temperature of $3\\times 10^{5}$ $^{\\circ}$K. This was confirmed by combined ground based and ultraviolet observations using the XMM-Newton X-ray telescope. (3) A secondary outburst of the same nature was expected at 2007 September 13. Within the accuracy of observations (about 6 hours), it started at the correct time. Thus the prediction was accurate at the same level as the prediction of the return of Halley's comet in 1986. (4) Further synchrotron outbursts were expected following the two bremsstrahlung outbursts. They came as scheduled between 2007 October and 2009 December. (5) Due to the effect of the secondary on the overall direction of the jet, the parsec scale jet was expected to rotate in the sky by a large angle around 2009. This rotation may have been seen at high frequency radio observations. OJ~287 binary black hole system is currently our best laboratory for testing theories of gravitation. Using OJ~287, the correctness of General Relativity has now been demonstrated up to the third Post-Newtonian order, at higher order than has been possible using the binary pulsars. ",
        "introduction": "In 1982 Aimo Sillanp\\\"a\\\"a put together the historical light curve of OJ~287 based on published measurements. There appeared to be a 12 year outburst cycle (see Figure 1), and moreover, it was obvious that the next cyclic outburst was due very shortly. Indeed, OJ~287 produced the expected event in the following January \\citep{haarala83,sillanpaa85}. Observations showed a sharp decline in the percentage polarization during the outburst maximum, indicating that the outburst was produced essentially by unpolarized light \\citep{smith87}. This is different from ordinary outbursts in OJ~287 which are characterized by an increase of the percentage polarization at the maximum light. Radio outbursts were found to follow the optical outbursts with a time delay of between 2 months and a year, depending on the observing frequency \\citep{valtaoja87}.  \\begin{figure} \\includegraphics[width=4.5cm,angle=270]{Historical.eps} \\caption{The optical light curve of OJ~287 from 1891 to 2010. It includes unpublished data from R.Hudec and M.Basta} \\end{figure}  \\citet{sillanpaa88} suggested that OJ~287 is a binary black hole system where a smaller companion perturbs periodically the accretion disk of a massive primary black hole. The next expected outburst was in 1994; it came as scheduled \\citep{fiorucci 94,sillanpaa96a}. In the binary model there should be two disk crossings per 12 yr orbital period. Thus the 1994 outburst should have an equal pair whose timing was calculated to be at the beginning of 1995 October \\citep{valtonen96,lehto96,sundelius96}. This prediction was also verified \\citep{sillanpaa96b}.   Alternative explanations were also put forward. Quasiperiodic oscillations in an accretion disk were suggested \\citep{igumenshchev99}, and several binary toy models without relativistic precession were also proposed \\citep{katz97,villata98,valtaoja00}. The latter models all predicted the next main outburst of OJ~287 in the autumn of 2006 while the precessing binary model gave a prediction one year earlier, at the beginning of 2005 October \\citep{sundelius97,kidger00}. The second major outburst was expected in late 2007 in all binary models while in the oscillating accretion disk model there was no reason to expect a second major outburst. In the precessing binary model the date was given with high accuracy, with the last prediction prior to the actual event being 2007 September 13 \\citep{valtonen07,valtonen08a}. In the oscillating accretion disk model and in non-precessing binary models the nature of the radiation at these outbursts should have been polarized synchrotron radiation while the precessing binary model predicted unpolarized bremsstrahlung radiation \\citep{lehto96}. In addition, the precessing binary model predicted a series of further outbursts for the interval 2007 - 2010 \\citep{sundelius97}. In this model the companion black hole should also affect the disk of the primary in a predictable way, leading to the wobble of the jet \\citep{valtonen06}.  With these predictions in mind, a multiwavelength campaign of observing OJ 287 during 2005 - 2010 was set up. ",
        "conclusions": "Prior to the 2005 - 2010 multiwavelength campaign there were several ideas about the nature of OJ~287. Fortunately, these models made completely different predictions about the behaviour of OJ~287. The result of a scientific enquiry is seldom as clear-cut as this: the outbursts satisfied the expectations of the precessing binary model in every respect while the alternative ideas failed. Therefore we are confident that OJ~287 can be used as a test laboratory for theories of gravitation. Our first results strongly support general relativity as the correct theory."
    },
    "1112/1112.3351_arXiv.txt": {
        "abstract": "A stellar-mass compact object spiraling into a supermassive black hole, an extreme-mass-ratio inspiral, is one of the targets for future space-based gravitational-wave detectors. Such inspirals offer a unique opportunity to learn about astrophysics and test General Relativity in the strong-field. We here study whether scalar-tensor theories in asymptotically flat spacetimes can be further constrained with these inspirals. In the extreme-mass ratio limit, and assuming analyticity of the coupling functions entering the action, we show that all scalar-tensor theories universally reduce to massive or massless Brans-Dicke theory. We also show that in this limit, black holes do not emit dipolar radiation to all orders in post-Newtonian theory. For massless theories and quasi-circular orbits, we calculate the scalar energy flux in the test-particle, Teukolsky formalism to all orders in post-Newtonian theory and fit it to a high-order post-Newtonian expansion. We then derive the post-Newtonian corrections to the scalar-tensor modified, Fourier transform of the gravitational wave response function and map it to the parameterized post-Einsteinian framework. With the Teukolsky flux at hand, we use the effective-one-body framework adapted to extreme mass-ratio inspirals to calculate the scalar-tensor modifications to the gravitational waveform. We find that such corrections are smaller than those induced in the early inspiral of comparable-mass binaries, leading to projected bounds on the Brans-Dicke coupling parameter that are worse than current Solar System ones. This is because Brans-Dicke theory modifies the weak-field, leading to deviations in the energy flux that are largest at small velocities. For massive theories, superradiance can produce resonances in the scalar energy flux that can lead to quasi-circular floating orbits outside the innermost stable circular orbit and that last until the supermassive black hole loses enough mass and spin-angular momentum. If such floating orbits occur in the frequency band of a LISA-like mission, they would lead to a large dephasing (typically $\\sim 10^6$ rads) that would prevent detection of such modified inspirals using General Relativity templates. A detection that is consistent with General Relativity would then rule out the presence of floating resonances at frequencies lower than the lowest inspiral frequency observed, allowing for the strongest constraints yet on massive scalar tensor theories. ",
        "introduction": "Gravitational waves hold the promise to allow for new and stringent tests of General Relativity (GR) in a previously obscure regime: the strong-field region. This is the regime where gravitational fields are strong and the characteristic velocity of inspiraling bodies is a significant fraction of the speed of light. The inspiral, merger and ringdown of compact objects is the prototypical example of a gravitational wave source that emits in this region. We currently lack experimental verification of GR in the strong-field, as the most stringent tests we possess, thanks to the discovery of the double binary pulsar~\\cite{2005ASPC..328...53B,Kramer:2006nb}, only probe the weak field. For example, the ratio of the total mass to the orbital separation (the gravitational compactness), which controls the magnitude of the orbital velocity, is of ${\\cal{O}}(10^{-5}) \\ll 1$ for the double binary pulsar~\\cite{2005ASPC..328...53B,Kramer:2006nb}, while it is of ${\\cal{O}}(1)$ during binary coalescence.  Much work has been devoted to studying whether gravitational waves from ground-based and space-based detectors could be used to test GR.  In most studies, one chooses a particular theory, or effect that one wishes to constrain, and then derives how gravitational waves would be deformed. Then, if a gravitational wave is detected that is consistent with GR, a constraint on gravitational wave deformations can be placed, which in turn are controlled by the magnitude of the coupling constants of the theory. Such a procedure has been applied to Brans-Dicke theory~\\citep{Will:1994fb,Will:2004xi,Berti:2004bd,Berti:2005qd,Stavridis:2010zz,Arun:2009pq,Keppel:2010qu,Yagi:2009zm}, dynamical Chern-Simons gravity~\\citep{Sopuerta:2009iy,Yunes:2009hc,Pani:2011xj,Yagi:2011xp}, modified quadratic gravity~\\cite{Yagi:2011xp}, phenomenological massive graviton propagation~\\citep{Will:1997bb,Scharre:2001hn,Will:2004xi,Berti:2004bd,Berti:2005qd,Yagi:2009zm,Berti:2011jz}, gravitational Lorentz-violation~\\citep{Mirshekari:2011yq}, gravitational parity-violation~\\citep{Yunes:2010yf,Yagi:2011xp}, violations of Local Position Invariance~\\citep{Yunes:2009bv} and theories with extra-dimensions~\\citep{Yagi:2011yu}.  Most of these studies considered the early inspiral of compact objects within the post-Newtonian (PN) approximation, where fields are expanded in small velocities and weak-fields to high order in perturbation theory. Moreover, they typically considered only the leading-order (Newtonian) correction introduced by the given alternative theory and neglected all higher-order corrections. Such an approximation becomes particularly ill-suited in the late phase of inspiral, when the binary is close to the inner-most stable circular orbit (ISCO). This then suggests that highly relativistic systems might allow us to place different, perhaps more stringent constraints on modified gravity theories.  An example of such a highly relativistic system is an extreme mass-ratio inspiral (EMRI), where a small compact object with mass $m_{\\SCO} \\in (1,10) M_{\\odot}$ spirals into a supermassive black hole with mass $M_{\\MBH} \\in (10^{4},10^{7}) M_{\\odot}$. The EMRI phase of interest to future space-based gravitational wave detection is the late inspiral, when the orbital separation is $\\leq 20 M_{\\MBH}$, the orbital period is $(10^{2},10^{4})$ seconds and low-frequency gravitational waves are emitted in the ($10^{-4},10^{-2}$) Hz frequency band. The merger and ringdown EMRI phases are completely negligible, as they contribute insignificant amounts of signal-to-noise ratio to the signal~\\cite{Yunes:2009ke,2011PhRvD..84f2003C}. The EMRI inspiral alone can last tens to hundreds of years in the detector band, completely outlasting any realistic detection. In a single year of inspiral observation, millions of radians are contained in the signal, which encodes rich information about the spacetime dynamics~\\cite{Ryan:1995wh}. EMRIs are therefore exceptional probes of strong-field GR and one might expect them to be the best probes of strong-field GR modifications.  In this paper, we study whether the detection of equatorial and quasi-circular EMRIs allows for stronger constraints on a large class of scalar-tensor theories than those currently in place due to Solar System observations or those projected with future gravitational wave observations of comparable-mass binary inspirals. We choose to work with equatorial and quasi-circular EMRIs for two reasons. First, EMRI formation scenarios exist where the small compact object is in an equatorial and quasi-circular orbit by the time it enters the detector sensitivity band. One such model is that of~\\cite{Levin:2006uc}, where a stellar-mass compact object is either created in an accretion disk surrounding a supermassive BH or is captured by the disk. The accretion disk is  expected to be in the spin equatorial plane inside a few hundred gravitational radii of the supermassive BH~\\cite{Bardeen:1975zz}. Thus, in this scenario by the time the small compact object reaches sufficiently small separations to produce detectable gravitational waves, it would be in an equatorial and quasi-circular orbit. Second, although such EMRIs are less likely than inclined and eccentric ones (since other EMRI formation channels are more favourable to generic EMRIs), the tests we will discuss below only require a single detection. Nonetheless, the study presented here should lay the foundations for extensions to more generic orbits if needed.  Scalar-tensor theories are prototypical models for non-minimally coupled scalar fields, which are common ingredients in low-energy effective actions of quantum gravity models (see e.g.~\\cite{Fujii_Maeda_book}). Furthermore, scalar-tensor theories can be shown to include $f(R)$ gravity as a particular case (see e.g.~\\cite{Sotiriou:2008rp,DeFelice:2010aj} for some reviews), thus being relevant as potential models for dark energy. Although here we focus on asymptotically flat geometries, our results should also apply to $f(R)$ theories, usually developed in asymptotically de-Sitter backgrounds, under the reasonable assumption that the cosmological background does not affect the local physics (see Ref.~\\cite{Sotiriou:2011dz} for a discussion on this issue).  The most general, stationary, axisymmetric and vacuum solution to generic scalar-tensor theories is the Kerr metric coupled to a constant scalar field. In fact, Sotiriou and Faraoni~\\cite{Sotiriou:2011dz} recently proved, under quite generic conditions, that such Kerr black holes are the endpoint of gravitational collapse in any scalar-tensor theory. This proof, however, fails for a cosmological, non-stationary (linear-in-time) scalar field~\\cite{2011arXiv1111.4009H}, which might lead to violations of the no-hair theorem applied to scalar-tensor theories~\\cite{2011arXiv1111.4009H,Bekenstein:1995un}.  The Kerr metric will be adopted as the background spacetime upon which a small compact object inspirals. Although we mainly work in the Einstein frame, physical observables are obtained by mapping results to the physical Jordan frame. We show that, for a large class of scalar-tensor theories in an asymptotically flat, stationary background and expanding the fields about perturbations proportional to the mass ratio, the equations governing the emission of gravitational and scalar waves only depend on two parameters: a coupling constant $\\alpha$ and the mass of the scalar field $\\mu_s$, both of which are uniquely determined by the specific scalar-tensor model. Thus, one can think of scalar-tensor theories as a continuous (Lie) space (manifold) of theories, where $(\\alpha,\\mu_{s})$ are coordinates.  Brans-Dicke theory is contained in the subset of massless scalar-tensor theories, i.e.~the subset of scalar-tensor theories with $\\mu_{s} = 0$. This theory not only requires a massless scalar field, but also that $\\alpha$ be related to $\\omega_{\\BD}$ by a specific relation (cf.~Eq.~\\eqref{def_alpha} below). Brans-Dicke theory is then parameterized by a single quantity, $\\omega_{\\BD}$, and can thus be thought of as a line in the $(\\alpha,\\mu_{s})$ space of all scalar-tensor theories.  \\section*{Executive Summary of Results} \\label{summary}  We study EMRIs in generic scalar tensor theories. By assuming analyticity of the coupling functions entering the action, we show that any of such scalar-tensor theories in the EMRI limit reduces identically and universally to massless or massive Brans-Dicke theory. To next-order in the mass ratio, this degeneracy in theory space breaks down. Therefore, any bound on Brans-Dicke theory derived from EMRI calculations automatically constrains the linear-in-mass-ratio expansion of all scalar-tensor theories. We also show that in the EMRI limit, black hole binaries in any scalar-tensor theory do not lead to dipolar radiation to all orders in PN theory, provided the background scalar field is constant at spatial infinity.  We then study massless scalar-tensor theories. First, we numerically calculate the scalar energy flux carried to spatial infinity in the point-particle approximation within the Teukolsky formalism and assuming quasi-circular orbits. This calculation of the flux is fully relativistic, albeit valid only to leading order in the mass ratio, but accounting for all orders in PN theory. We fit this flux to a $3$PN and a $3.5$PN expansion and calculate the Fourier transform of the gravitational wave response function in the stationary-phase approximation beyond Newtonian order. We find that the PN corrections to the massless scalar-tensor modifications to the Fourier phase can be mapped to the recently proposed parameterized post-Einsteinian (ppE) framework~\\cite{Yunes:2009ke,2011PhRvD..84f2003C} in a straightforward manner, which we explicitly provide.  We then proceed to determine whether these PN corrections to the modified gravitational wave allow us to constrain the Brans-Dicke coupling parameter better than with comparable-mass binary inspirals. Recall that due to the equivalence discussed earlier, the bounds derived here also apply to all massless scalar-tensor theories to leading-order in the mass ratio. With the Teukolsky flux at hand, we employ the effective-one-body (EOB) framework~\\cite{Buonanno99,Buonanno00}, recently adapted to EMRIs~\\cite{2009GWN.....2....3Y,Yunes:2010zj}, to evolve inspiral with and without the Brans-Dicke correction. We then compare their associated gravitational wave phases, after minimizing the difference over an arbitrary time and phase offset. Contrary to our expectation, we find that EMRIs will not be able to constrain Brans-Dicke theory beyond current Solar System constraints. In fact, the projected bounds we obtain are {\\emph{worse}} than those found by considering gravitational-wave detection of comparable-mass binary inspirals with only the leading-order Brans-Dicke correction~\\citep{Will:1994fb,Will:2004xi,Berti:2005qd,Stavridis:2010zz,Arun:2009pq,Keppel:2010qu,Yagi:2009zm}.  This seemingly surprising result is in fact perfectly reasonable, once one realizes that scalar-tensor theories are not necessarily strong-curvature modifications of GR. At the level of the action, the scalar-field modification acts on the Ricci scalar in the Jordan frame and it does not introduce higher curvature corrections. This is consistent with the fact that scalar-tensor theories include $f(R)$ gravity, which is usually studied as an infrared, rather than an ultraviolet (or strong-field/strong-curvature) correction to GR. In terms of the energy flux, the weak-field nature of massless scalar-tensor theories is evidenced by the scalar flux becoming less and less important than the GR one as the small compact object approaches the horizon. This is why the modified theory introduces a {\\emph{pre}}-Newtonian correction to the gravitational wave phase (a $-1$PN effect relative to GR), instead of a {\\emph{post}}-Newtonian one.  \\begin{figure} \\epsfig{file=mu_omega_Einstein.eps,width=8cm,angle=0,clip=true} \\caption{\\label{fig:floating-regions} Parameter space of massive Brans-Dicke theory for a quasi-circular EMRI with typical neutron star sensitivity $s_\\SCO=0.188$, supermassive black hole spin $a_\\MBH/M_{\\MBH}=0.9$ and mass $M_\\MBH=10^5 M_\\odot$. The red, dashed curve is the boundary of the region in $(\\omega_\\BD,\\mu_{s})$ space that separates floating and non-floating resonances (cf. Eq.~\\eqref{omegaBD_crit_Einstein}). The area below the black curve shows the region that is ruled out by Solar System experiments~\\cite{Perivolaropoulos:2009ak,Alsing:2011er}. The area between the green, dot-dashed curves shows the region where a resonance would occur inside the classic LISA sensitivity band [$(10^{-4},10^{-2})$ Hz]. Due to a large dephasing introduced by floating resonances, any detection of an EMRI in a LISA-like mission would rule out the (open) region delimited by ABC. The dot-dashed blue line shows the values of $(\\omega_\\BD,\\mu_{s})$ below which non-floating resonances lead to dephasings larger than $1$ rad (see Sec.~\\ref{sec:resonant-effect-on-GW}). The ISCO frequency, $\\Omega_\\ISCO$, is shown by a vertical dotted line. The curves $\\omega_\\BD^{c}$ and $\\omega_{\\BD}^{\\rm nf}$ terminate at $\\mu_s\\sim\\Omega_\\orb\\sim\\Omega_\\ISCO\\sim0.22/M_\\MBH$. } \\end{figure}  With the massless case under control, we then proceed to consider massive scalar-tensor theories, where now GR deformations depend both on the mass of the scalar field $\\mu_{s}$ and the scalar-tensor coupling parameter $\\alpha$. Depending on the values of $(\\alpha,\\mu_{s})$, prograde EMRI orbits can experience superradiance and produce resonances in the scalar flux that can momentarily counteract the gravitational wave flux, leading to {\\emph{floating orbits}}~\\cite{Cardoso:2011xi}, i.e.~orbits where the inspiral greatly slows down at a given floating radius. Figure~\\ref{fig:floating-regions} shows the separatrix (red dashed line) between floating and non-floating orbits for an EMRI with total mass $M\\sim10^{5}M_{\\odot}$ in massive Brans-Dicke theory. As indicated in this figure, floating orbits can exist only in the region above and to the right of the red-dashed line. This curve (and all others) terminate when the small compact object approaches the ISCO, approximately when $\\mu_s\\sim\\Omega_\\ISCO$, where $\\Omega_\\ISCO$ is the ISCO frequency (vertical dotted line in Fig.~\\ref{fig:floating-regions}).  We have chosen to show the region of allowed floating orbits in $(\\omega_{\\BD},\\mu_{s})$ coordinates instead of $(\\alpha,\\mu_{s})$, thus restricting attention to massive Brans-Dicke theory via Eq.~\\eqref{def_alpha},  so as to be able to compare directly with Solar System experiments~\\cite{Bertotti:2003rm,Perivolaropoulos:2009ak}. The area under the solid black line in Fig.~\\ref{fig:floating-regions} denotes the region that is currently ruled out by such experiments~\\cite{Bertotti:2003rm,Perivolaropoulos:2009ak}. Observe that as $\\mu_{s} \\to 0$, the Solar System constraint becomes independent of $\\mu_s$ and approaches $\\omega_{\\BD} \\geq 4 \\times 10^{4}$, as expected. For completeness, we also show the region (between green, dot-dashed vertical lines) for which the floating orbits would lead to GWs inside the classic-LISA sensitivity band [$(10^{-4},10^{-2})$ Hz], since the resonant frequency is linearly proportional to $\\mu_{s}$(see Ref.~\\cite{AmaroSeoane:2012km} for a recent review on eLISA-NGO). From this, one can easily observe that there is a large region in $(\\omega_{\\BD},\\mu_{s})$ space that is not ruled out by Solar System experiments, yet where floating orbits should be observed with LISA like instruments, ie.~the region between the greened dot-dashed curves, above the black solid line and below the red dashed line. As we explain below, the modifications to an EMRI GW signal due to hitting a floating orbit are so large that an observation of a signal consistent with GR would rule out the region below the red dashed line, thus improving Solar System constraints by several orders of magnitude.  Floating orbits are not a new invention in the context of EMRIs and GR. Press and Teukolsky~\\cite{1972Natur.238..211P,Thorne:1986iy}, after a suggestion of Misner's, were the first to study floating orbits within GR. The main idea is that for certain orbital configurations, a small object orbiting a supermassive black hole can {\\emph{extract}} angular momentum from the latter. That is, the gravitational waves that impinge on the supermassive black hole horizon are not absorbed, but instead reflected with greater energy and angular momentum that they possessed initially. If such extraction is sufficiently large to balance the energy lost by the system due to gravitational waves traveling out to spatial infinity, the small object might momentarily experience no radiation-reaction force and stall in its orbit.  Press and Teukolsky found that, although such floating orbits are not impossible in GR, for them to exist outside the ISCO of a prograde orbit, one requires an essentially maximally spinning supermassive black hole, with spins exceeding the Thorne limit $|\\vec{S}_{\\rm max}| = 0.998 M_{\\MBH}^{2}$~\\cite{1972Natur.238..211P,Thorne:1986iy,1974ApJ...191..507T}. In massive scalar-tensor theory, more energy can be extracted due to the activation of the scalar energy flux, thus allowing for floating orbits outside the ISCO for supermassive black holes with smaller spins.  Floating orbits, however, are somewhat of a misnomer, as an EMRI that attempts to cross such a floating radius does not completely stall or stagnate as the name implies. Instead, as an EMRI attempts to cross a floating radius, its evolution is greatly slowed down, as it is now driven by the supermassive black hole's energy and spin angular momentum loss. In other words, the radiation-reaction force or the energy flux are weakened by an additional factor of mass-ratio $\\nu$. Thus, while to leading (first) order the energy flux scales with $\\nu^{2}$, during floating it scales with $\\nu^{3}$, i.e.~it becomes a second-order in the mass ratio effect.  The time it takes the EMRI to cross a floating resonance is determined by the supermassive black hole's shedding efficiency [of ${\\cal{O}}(10^{6})$ yrs for a typical EMRI]. This timescale can easily exceed a Hubble time for EMRIs orbiting at separations greater than $326 M_{\\MBH}$. While the EMRI is slowed down, it will produce gravitational waves very different from those in GR, leading to a large dephasing, i.e.~a phase difference relative to the GR phase, of roughly $10^6$ rads over a one year observation for phenomenologically viable $(\\alpha,\\mu_{s})$ parameters. Therefore, if massive, scalar-tensor theories are the correct representation of Nature, such a large dephasing would prevent detection of EMRIs in a GR template-based, matched-filtering gravitational wave search.  On the other hand, the detection of an EMRI gravitational wave that is consistent with GR can then be used to constrain massive scalar-tensor theories. This is because such a detection would imply that a resonance was not present in the range of frequencies sampled by the EMRI. If a resonance were present at a frequency lower than the lowest one detected, the EMRI would have hanged at the resonance frequency for a very long time, evolving very slowly, and never entering the detector's sensitivity band. The requirement that the resonant frequency be lower than the lowest EMRI frequency detected would impose unprecedented constraints on massive scalar tensor theory: the (open) region delimited by ABC in Fig.~\\ref{fig:floating-regions} would be ruled out. Observe that such a constraint is much stronger than current Solar-System ones (solid black line in Fig.~\\ref{fig:floating-regions}). Therefore, even though future space-based gravitational wave detectors will not be able to constrain massless scalar-tensor theories beyond current Solar System levels~\\cite{Bertotti:2003rm,Perivolaropoulos:2009ak}, they will be able to severely bound a large sector of massive scalar tensor theories that is today still phenomenologically viable.  The remainder of this paper presents the details of this calculation. Section~\\ref{sec:review} reviews the basics of scalar-tensor theories applicable to EMRIs. Section~\\ref{sec:GWST} describes the test-particle approximation and the Teukolsky framework used to model EMRIs in scalar-tensor theories. Section~\\ref{Sec:EfluxMassless} calculates the energy flux in the Teukolsky formalism for massless scalar fields, fits the massless flux to a PN expansion and computes the Fourier transform of the response function in the stationary-phase approximation. Section~\\ref{Sec:EfluxMassive} repeats the analysis of Sec.~\\ref{Sec:EfluxMassless} but for massive scalar fields. Section~\\ref{sec:Constraints} describes how to model EMRIs using the EOB-EMRI framework both in GR and in massless scalar-tensor theories, and then calculates the dephasing between a GR waveform and a massless scalar-tensor theory one, allowing us to derive a projected bound. Section~\\ref{sec:Constraints-Massive} derives analytical estimates on how well massive scalar-tensor theories will be constrained with future gravitational wave detectors. Section~\\ref{sec:conclusions} concludes and points to future research.  Throughout the rest of this paper we follow mainly the conventions of Misner, Thorne and Wheeler~\\cite{Misner:1973cw}. We use Greek letters to denote spacetime indices, Latin ones at the middle of the alphabet $i,j,\\ldots$ stand for spatial indices only, and Latin ones at the beginning of the alphabet $a,b,\\ldots$ range over waveform parameters. We also use geometric units with $G=c=\\hbar=1$, except in Section~\\ref{sec:review} where we restore $G$ for clarity. Here $G$ is the gravitational constant as measured by Cavendish-like experiments, i.e. as measured by a distant observer in the Jordan frame. ",
        "conclusions": "\\label{sec:conclusions}  We have investigated whether scalar-tensor theories can be constrained with future EMRI gravitational wave observations. The gravitational and scalar emission from a EMRI system in a generic (albeit analytic) scalar-tensor theory ultimately depends on only two parameters, $(\\alpha,\\mu_s)$. Brans-Dicke theory is recovered in a particular case when $\\mu_s=0$. In such theories, scalar dipole emission can activate and modify the radiation-reaction force via corrections to the energy flux carried away from the binary.  We first studied massless scalar tensor theories ($\\mu_s=0$ case) and numerically calculated the energy flux carried to spatial infinity and into the supermassive black hole horizon by the scalar field perturbation. We obtained this flux in the Teukolsky formalism, by solving certain master equations for a set of given quasi-circular, equatorial geodesics. We then extended the $2.5$PN expansion of this flux to $3.5$PN order through fitting coefficients. We also calculated the Fourier transform of the response function in the stationary-phase approximation beyond Newtonian order. We found that the Brans-Dicke corrected Fourier phase in the PN expansion can easily be mapped to the ppE framework of~\\cite{Yunes:2009ke,2011PhRvD..84f2003C} and we derived the mapping.  Once the massless scalar energy flux had been calculated, we used it in a modified EOB-EMRI framework to estimate the dephasing induced by the massless scalar-tensor terms. We found that EMRIs do not lead to larger scalar-tensor corrections, but instead, the projected bounds on $\\alpha$ (or $\\omega_{\\BD}$) are slightly worse than those derived from comparable-mass, early inspiral studies~\\citep{Will:1994fb,Will:2004xi,Berti:2005qd,Stavridis:2010zz,Arun:2009pq,Keppel:2010qu,Yagi:2009zm}. In EMRIs, the small companion spends all of the observation time very close to the supermassive black hole. Scalar-tensor theories, however, are not necessarily strong-field modifications to GR, but rather weak-field ones. That is, the main scalar-tensor modification to the orbit is through the introduction of pre-Newtonian corrections to the radiation-reaction force due to scalar dipole emission, which is dominant for orbits with large separations.  We then studied massive scalar tensor theories ($\\mu_s\\neq0$ case) and confirmed that certain resonances in the scalar horizon flux arise, leading to floating orbits, and thus, drastic modifications in the gravitational wave phase evolution. During floating, the small compact object continues to inspiral, but at a rate much slower than the standard GR one. The floating inspiral is driven by the shedding of mass and spin angular momentum of the supermassive black hole. The timescale for floating to end depends on the frequency at which floating occurs, but it can easily be greater than $10^{2}$ years. In fact, if the EMRI separation is greater than $\\sim 326 M_{\\MBH}$, this timescale exceeds the Hubble time.  Such large time scales suggest an easy constraint on the parameters of massive scalar-tensor theories. That is, if these resonances existed, then we should have found many systems orbiting at resonant radii, where material that would otherwise be swallowed by the supermassive black hole would accumulate. The prototypical system that comes to mind are accretion disks, where a supermassive black hole accretes material that is driven inwards by viscosity (scalar dipole emission is suppressed in neutron star/pulsar binaries due to similar sensitivities). Focusing on the orbits of accretion disk material in the absence of an EMRI, scalar flux resonances could balance the viscous force that the disk would otherwise experience, leading to a stagnation of orbits and a build up of material at a given radius (see eg.~\\cite{Yunes:2011ws,Kocsis:2011dr}). Such a build up would translate into a sharp feature in the photon energy flux observed, a feature that in fact is not currently needed to fit observations~\\cite{Li:2004aq,Shafee:2005ef,Steiner:2009af,McClintock:2011zq,Kulkarni:2011cy}.  One could argue that this already rules out a certain area of parameter space of massive scalar-tensor theories, but it is not clear whether such tests can be cleanly applied. First, a continuous ring-like, configuration of particles in a circular orbit does not lead to superradiance, as only the $m=0$ modes are excited. More work is necessary to understand what happens to high $(l,m)$ modes that could lead to superradiance when the ring-like configuration is not continuous, with gaps of the size of the orbiting particles. Second, the current observational fits of the accretion disk energy spectrum are somewhat susceptible to uncertainties in accretion disk physics, requiring the inclusion of non-thermal features~\\cite{Steiner:2009af} and deviations from the Novikov-Thorne disk model~\\cite{Kulkarni:2011cy}. Therefore, it is not clear whether accretion disks would feel such floating resonances, and if so, whether accretion disk observations could help us constrain massive scalar-tensor theories in a clean manner.  EMRIs, on the other hand, are much cleaner systems, and thus, their observation could be used to place constraints on massive scalar-tensor theories. For a resonance to affect an EMRI gravitational wave, the former must occur at a radius small enough that an orbiting EMRI can cross it. In the range of scalar masses in which this occurs, the floating orbit introduces an enormous dephasing in the gravitational waveform relative to the GR expectation. Such a large dephasing would prevent the detection of such modified EMRI waves with LISA or any other future detector if one match-filters with GR templates. Therefore, given a future EMRI detection that is consistent with GR, one could rule out a large area of scalar-tensor parameter space. Such projected constraints are several orders of magnitude more stringent than current Solar System ones.  Floating EMRIs emit almost monochromatic gravitational waves, and thus, one might worry that such waves could be detected as a continuous source. These waves are almost monochromatic because the radiation-reaction force becomes second-order in the mass ratio during floating, and since this quantity is much smaller than unity, the inspiral evolution is greatly slowed down. Continuous algorithms might detect such monochromatic waves, but their frequency is much higher than that of white-dwarf binaries, the traditional LISA continuous sources. In fact, there are no continuous, monochromatic sources that would emit in the high-frequency end of the LISA band according to GR. Thus, a continuous source detection at a high gravitational wave frequency could indicate the presence of floating EMRIs.  Future work could concentrate on extending the results of this paper to relax the condition $\\mu_{s} M_{\\MBH} \\ll 1$. Floating orbits can occur near the ISCO of a spinning black hole if $\\mu_s M_\\MBH\\sim0.5$. Our results strongly suggest that, also in this case, floating orbits would lead to a very large dephasing of the gravitational waveform. However, a more detailed quantitative analysis, possibly involving the numerical fluxes we have here presented, would be desirable. One could also investigate whether the non-floating resonances have structure and how this modifies the waveform observable if the resonant width were large enough.  Another possible avenue for future research would be to investigate whether additional constraints can be placed on scalar-tensor theories given other astrophysical observations. For example, if floating orbits exist, then the spin of the supermassive black hole would be decreased after every EMRI coalescence. Statistically speaking, one would then predict a supermassive black hole spin distribution that is peaked close to zero. Therefore, if a sufficient number of supermassive black holes are observed with sufficiently large spins, one would then infer that if floating orbits exist, then they had to be hidden inside the horizon of the supermassive black holes. This can then be used to place a statistical constraint on the coupling parameters of scalar-tensor theories~\\cite{Bence}.  Other future work could focus on determining whether the floating resonances found here are generic to other theories. For example, recently there has been much effort to understand certain quadratic gravity theories~\\cite{Yagi:2011xp}, where the Einstein-Hilbert action is enhanced through the addition of terms composed of the product of a scalar field and all quadratic curvature invariants. Usually, the scalar field in these theories is assumed massless, but if it were endowed with a mass, superradiant-induced, floating resonances could also arise. If so, the same mechanism described here to constrain the mass of scalar-tensor theories could be used to bound the mass of quadratic-gravity scalar fields."
    },
    "1112/1112.3398_arXiv.txt": {
        "abstract": "{HD 181327 is a young Main Sequence F5/F6 V star belonging to the \\bp\\ moving group (age $\\sim$ 12 Myr). It harbors an optically thin belt of circumstellar material at radius $\\sim$90 AU, presumed to result from collisions in a population of unseen planetesimals.} {We aim to study the dust properties in the belt in great details, and to constrain the gas-to-dust ratio.} {We obtained far-IR photometric observations of \\hd\\ with the PACS instrument onboard the Herschel Space Observatory\\thanks{Herschel is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA.}, complemented by new 3.2 mm observations carried with the ATCA\\thanks{The Australian Telescope Compact Array is operated by the Australian Telescope National Facility (ATNF) managed by CSIRO.} array. The geometry of the belt is constrained with newly reduced HST/NICMOS scattered light images that allow the degeneracy between the disk geometry and the dust properties to be broken. We then use the radiative transfer code GRaTer to compute a large grid of models, and we identify the grain models that best reproduce the Spectral Energy Distribution (SED) through a Bayesian analysis. We attempt to detect the oxygen and ionized carbon fine-structure lines with Herschel/PACS spectroscopy, providing observables to our photochemical code ProDiMo.} {The HST observations confirm that the dust is confined in a narrow belt. The continuum is detected with Herschel/PACS completing nicely the SED in the far-infrared. The disk is marginally resolved with both PACS and ATCA. A medium integration of the gas spectral lines only provides upper limits on the [OI] and [CII] line fluxes. We show that the \\hd\\ dust disk consists of micron-sized grains of porous amorphous silicates and carbonaceous material surrounded by an important layer of ice, for a total dust mass of $\\sim0.05M\\dma{\\oplus}$ (in grains up to 1 mm). We discuss evidences that the grains consists of fluffy aggregates. % The upper limits on the gas atomic lines do not provide unambiguous constraints: only if the PAH abundance is high, the gas mass must be lower than $\\sim17M\\dma{\\oplus}$. } {Despite the weak constraints on the gas disk, the age of \\hd\\ and the properties of the dust disk suggest that it has passed the stage of gaseous planets formation. The dust reveals a population of icy planetesimals, similar to the primitive Edgeworth-Kuiper Belt, that may be a source for the future delivery of water and volatiles onto forming terrestrial planets.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.4747_arXiv.txt": {
        "abstract": "We present the results from an ultra-high-resolution 7\\,mm Very Long Baseline Array (VLBA) study of the relativistic jet in the BL Lacertae object \\object{OJ287} from 1995 to 2011 containing 136 total intensity images. Analysis of the image sequence reveals a sharp jet-position-angle swing by $>100^{\\circ}$ during [2004,2006], as viewed in the plane of the sky, that we interpret as the crossing of the jet from one side of the line of sight to the other during a softer and longer term swing of the inner jet. Modulating such long term swing, our images also show for the first time a prominent erratic wobbling behavior of the innermost $\\sim0.4$\\,mas of the jet with fluctuations in position angle of up to $\\sim40^{\\circ}$ over time scales $\\sim2$\\,yr. This is accompanied by highly superluminal motions along non-radial trajectories, which reflect the remarkable non-ballistic nature of the jet plasma on these scales. The erratic nature and short time scales of the observed behavior rules out scenarios such as binary black hole systems, accretion disk precession, and interaction with the ambient medium as possible origins of the phenomenon on the scales probed by our observations, although such processes may cause longer-term modulation of the jet direction. We propose that variable asymmetric injection of the jet flow; perhaps related to turbulence in the accretion disk; coupled with hydrodynamic instabilities, leads to the non-ballistic dynamics that cause the observed non-periodic changes in the direction of the inner jet. ",
        "introduction": "\\label{Intr} The most luminous long-lived sources of radiation in the cosmos---active galactic nuclei (AGN)---are powered by gas falling from an accretion disk onto a super-massive black hole (BH, $\\gtrsim10^7\\,\\rm{M}_{\\sun}$) at their center. AGN classified as blazars are characterized by wild variability of flux of non-thermal radiation from radio to $\\gamma$-ray frequencies. Members of this class include BL~Lacertae objects (BL~Lacs) and, at higher emitted power, flat spectrum radio quasars (FSRQ). The remarkable properties of blazars include superluminal apparent motions that can exceed $40\\,c$ \\citep{Jorstad:2005p264}, as well as substantial changes in flux and linear polarization on time scales as short as minutes. These phenomena are thought to be caused by relativistic jets of highly energized, magnetized plasma that are propelled along the rotational poles of the BH-disk system \\citep[e.g.,][]{Blandford:1977p6285}. AGN jets pointing within $\\lesssim10^{\\circ}$ of our line of sight  beam their radiation and shorten the variability time scales to give blazars their extreme properties.  High-resolution long-term very-long baseline interferometric (VLBI) monitoring of blazars is revealing a growing number of cases where the observed (in projection) jets undergo rapid swings of position angle in their innermost regions while remaining collimated out to kiloparsec scales or farther \\citep[e.g.,][]{2003ApJ...589L...9H,Stirling:2003p154,Tateyama:2004p11950,Bach:2005p462,Lobanov:2005p15363,Mutel:2005p443,Savolainen:2006p13808,Agudo:2007p132}. Time scales ranging between $\\sim2$ and $\\gtrsim20$ years and projected oscillations of the jet position angle with amplitudes from $\\sim25^{\\circ}$ to $\\gtrsim100^{\\circ}$ are typical for the reported cases \\citep[see][for examples]{Agudo:2009p14514}. This phenomenon---which has previously been termed \\emph{jet precession}, although we prefer \\emph{wobbling} given the doubts about its periodicity in some sources \\citep[e.g.,][]{Mutel:2005p443}---has been reported so far only in the innermost (parsec-scale) regions of jets. This already implies that the origin of jet wobbling might be a mechanism intimately tied to the origin of the jet. Bends and helical structures in AGN are also claimed to be produced by changes in direction near the base of the jet \\citep[e.g.][]{Dhawan:1998p15371,2003ApJ...584..135L,Hardee:2005p260,Agudo:2006p331}, which may be related to the wobbling phenomenon as well.  During the past few decades, the BL~Lac object \\object{OJ287} ($z=0.306$) has shown quasi-periodic double-peaked optical flares every $\\sim12$\\,yr, most recently in $[2005,2008]$ \\citep[e.g.,][]{Villforth:2010p11557}. This quasi-periodicity has been used to support binary BH models, as well as other scenarios such as accretion disk or jet instabilities \\citep[see][for a review of these different scenarios]{Villforth:2010p11557}. Consistent with the binary BH hypothesis, \\citet{Tateyama:2004p11950} reported jet wobbling in a 3.5\\,cm long-term Very Long Baseline Array (VLBA) study from 1994 to 2002. They interpreted this as the result of ballistic precession of the jet, also with a periodicity of $\\sim12$\\,yr. Based on the \\cite{Valtonen:2008p5346} binary BH model for \\object{OJ287}, \\citet{Valtonen:2011p15373} have proposed a new model for wobbling of the jet in this source with a periodicity of $\\sim120$\\,yr and a modulation of the jet position angle on a $\\sim12$\\,yr time scale driven by changes in orientation of the inner accretion disk of the proposed primary black hole.  To re-evaluate the origin of jet wobbling in general, and in \\object{OJ287} in particular, we present here the results of a detailed 7\\,mm VLBA study of \\object{OJ287} from 1995 to 2011 that includes 136 VLBI images. Owing to the reduced effect of synchrotron self-absorption in blazar jets at these short wavelengths, compared with centimeter-wave radio bands, and the ultra-high ($\\sim0.15$\\,milliarcsecond) angular resolution of the images, we scrutinize the inner jet regions in much finer detail than is possible at centimeter wavelengths. We adopt hereafter the standard $\\Lambda$CDM cosmology with $H_0$=71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.27$, and $\\Omega_\\Lambda=0.73$. Under this assumption, at the redshift of \\object{OJ287}, 1\\,milliarcsecond (mas) translates to a projected distance of $4.48$\\,pc, and a proper motion of 1\\,mas\\,yr$^{-1}$ corresponds to a superluminal speed of $19.09\\,c$. ",
        "conclusions": "Our 7\\,mm VLBA study has revealed a sharp jet-position-angle swing by $>100^{\\circ}$ during [2004,2006], as viewed in the plane of the sky, that we interpret as the crossing of the jet from one side of the line of sight to the other during a softer and longer term swing of the inner jet. Modulating such long term swing, our images also show that the innermost $\\sim0.4$\\,mas region of the 7\\,mm jet in \\object{OJ287} wobbles following an erratic pattern with amplitudes up to $\\sim40^{\\circ}$ on time scales of $\\sim2$\\,yr. The latter phenomenon, reported here for the first time, rules out scenarios such as the putative orbit of the supermassive black hole responsible for the production of the jet in a binary system, accretion disk precession induced by either a companion compact object or the Lense-Thirring effect, and jet interaction with the external medium as possible origins of the phenomenon. This is based on the long time scales related to these scenarios compared with the $\\sim2$\\,yr time scale and to the apparently erratic nature of the observed wobbling in \\object{OJ287}. The dense time sampling of our observations also allows us to rule out misidentification of jet features across epochs as a possible explanation. This leads us to propose a new possibility to explain the non-periodic jet wobbling reported here through asymmetric fluctuations in injection of plasma flow into the jet coupled with instabilities. This could be related to variations in the accretion process and in the disk-jet connection.  Both our observations and those from other VLBI programs at centimeter wavelengths \\citep{Moor:2011p16214,Valtonen:2011p15373} rule out the $\\sim12$\\,yr periodic precession model proposed by \\citet{Tateyama:2004p11950}, which predicted different evolution of the JPA than observed. In contrast, the 120\\,yr periodicity of the \\citet{Valtonen:2011p15373} precession model seems able to reproduce the JPA evolution observed at centimeter wavelengths, although the 12\\,yr modulation that their model imparts on the predicted JPA changes does not clearly follow their observed JPA variations during the [2005,2011] time range. Such 12\\,yr modulation might be better reproduced by the $3.5$\\,cm JPA observations by \\citet{Moor:2011p16214}, but more data is still needed to clarify the consistency between new Valtonen's JPA evolution-model and the centimeter VLBI jet behavior.  Our interpretation of the apparent mismatch between the JPA evolution reported in this paper (i.e., on scales of a few parsecs) and those observed at centimeter wavelengths (tens of parsecs) leaves open the possibility that different instability modes operate on these two spatial scales. This allows for the JPA on the inner scales to behave in an erratic way, while on the larger scales it can be driven by instabilities coupled to modes with periodic behavior. Periodic models based on binary black hole systems are thus not ruled out by our observations, unless future centimeter VLBI observations of the jet in \\object{OJ287} show more erratic JPA variations similar to those seen at millimeter wavelengths.  The detailed study of the evolution of the JPA of \\object{OJ287} presented here, and the main conclusions drawn above, were only possible with the ultra-high resolution provided by the VLBA at 7\\,mm and long-term monitoring of \\object{OJ287} with dense time sampling. Studies similar to that presented here can be carried out for a number of blazars that have been monitored intensively with millimeter VLBI during the past decade. Performing these studies will be crucial to evaluate how common erratic wobbling is in the jets of AGN, and to draw general conclusions about the possible relation of this phenomenon to the accretion process."
    },
    "1112/1112.6118_arXiv.txt": {
        "abstract": "We present preliminary results of the application of a new sophisticated code which allows high precision integration of orbits of stars belonging to a dense stellar system moving in the vicinity of a massive black hole. This mimics the situation observed in the center of many galaxies, where a nuclear star cluster contains a massive black hole which, in the past, was, likely, an active engine of violent emission of radiation. The main scope of our work is the investigation of the relaxation of the super star cluster on a sufficiently long time, together with the investigation of its feedback with the massive black hole. ",
        "introduction": "It is known, nowadays, that many galaxies harbor a supermassive black hole (SBH) in their innermost regions. In many cases, the massive black hole is surrounded by an ultra compact stellar system, the so called Nuclear Star Cluster (NSC). Only in the case of the Milky Way the NSC is resolved in its stellar components, thanks to the high sensitivity and spatial resolution of modern instruments. The formation of compact nuclei in galaxies has been studied, among others, by \\citet{cap1} which developed  the original \\citet{trem} hypothesis of the M 31 nucleus formation via repeated infalls of globular clusters braked by dynamical friction, as alternative to the in-situ formation hypothesis. Nevertheless, the full $N$-body simulations performed so far (see, e.g., \\citet{cdmio}, \\citet{ant11}) have not been extended enough in time and with a sufficient resolution to reach fairly conclusive statements about it. ",
        "conclusions": ""
    },
    "1112/1112.1495_arXiv.txt": {
        "abstract": "{} {We investigate the performance of an analytic model of the 3D matter distribution, which combines perturbation theory with halo models, for weak-lensing configuration-space statistics.} {We compared our predictions for the weak-lensing convergence two-point and three-point correlation functions with numerical simulations and fitting formulas proposed in previous works. We also considered the second- and third-order moments of the smoothed convergence and of the aperture-mass.} {As in our previous study of Fourier-space weak-lensing statistics, we find that our model agrees better with simulations than previously published fitting formulas. Moreover, we recover the dependence on cosmology of these weak-lensing statistics and we can describe multi-scale moments. This approach allows us to obtain the quantitative relationship between these integrated weak-lensing statistics and the various contributions to the underlying 3D density fluctuations, decomposed over perturbative, two-halo, or one-halo terms.} {} ",
        "introduction": "\\label{Introduction}  The standard paradigm known as the $\\Lambda$CDM cosmology includes dark components: dark matter and dark energy \\citep{Komatsu2011}. Weak lensing of background galaxies by foreground large-scale structures, the so-called ``cosmic shear'',  has been recognized as a potentially powerful tool for probing the distribution of dark matter as well as the nature of dark energy~\\citep{Albrecht2006}. Reports of significant detections of cosmic shear have been made by various groups \\citep{Bacon2000,VanWaerbeke2000,Wittman2000,Hamana2003,Jarvis2006,Semboloni2006,Fu2008,Schrabback2010}.  By analyzing the cosmic shear data, one can directly measure the power spectrum of the matter density fluctuations on cosmological scales, which contain a wealth of cosmological information such as neutrino masses and dark energy equation-of-state parameters~\\citep{Jarvis2006,Semboloni2006,Ichiki2009,Schrabback2010}. Consequently, it is a main goal of cosmology to infer and constrain these quantities from observations. To do this, a number of ambitious surveys are planned, such as the Hyper Suprime-Cam Weak Lensing Survey~\\citep{Miyazaki2006}\\footnote{http://www.naoj.org/Projects/HSC/index.html}, the Dark Energy Survey~(DES)\\footnote{http://www.darkenergysurvey.org/}, the Large Synoptic Survey Telescope~(LSST)\\footnote{http://www.lsst.org/}, Euclid~\\citep{Refregier2010}\\footnote{http://www.euclid-ec.org/}, and the Wide-Field Infrared Survey Telescope~(WFIRST)\\footnote{http://wfirst.gsfc.nasa.gov/}.  Most weak-lensing information is contained in small angular scales and therefore weak-lensing statistics are nonlinear and non-Gaussian~\\citep{Munshi2008,Takada2009,Sato2009,Sato2011a}. If we aim to exploit the full information, we have to treat the nonlinear effects to accurately model the weak lensing statistics. Furthermore, one has to use an appropriate likelihood function with given marginal distributions, otherwise the obtained results would be systematically biased \\citep{Sato2010,Sato2011}.   In a first companion paper \\citep[][hereafter paper I]{Valageas2011f}, we studied the Fourier-space weak-lensing statistics such as the weak-lensing power spectrum and bispectrum, and found that our model proposed by \\citet{Valageas2011d,Valageas2011e}, which combines perturbation theory with halo models, agrees better with ray-tracing simulations than previously published fitting formulas and phenomenological models. In this second paper, we study the real-space weak-lensing statistics, which are more often used for the statistical analysis of actual measurements than Fourier-space statistics, because observations are made in configuration space.  Previous works have already shown that on small scales the halo model provides a good description of the two-, three- and four-point correlations or smoothed moments of the cosmic shear (using some approximations) \\citep{Takada2002,Takada2003,Benabed2006}, whereas a stochastic halo model can recover the probability distribution function of the unsmoothed convergence \\citep{Kainulainen2011,Kainulainen2011a}. Here we include all ``one-halo'', ``two-halo'' and ``three-halo\" terms, as well as one-loop perturbative results, and we compare these with larger-scale simulations. This yields a greater accuracy and allows us to compare these different contributions, from very large to small scales. This should be useful for practical purposes because these different terms have different theoretical accuracies and probe different regimes of gravitational clustering, hence it is important to know their relative impact on weak-lensing probes as a function of angular scale.   This paper is organized as follows. In Sect.~\\ref{Analytic} we briefly recall how configuration-space weak-lensing statistics are computed from polyspectra of the 3D matter density field. We describe our numerical simulations and the data analysis in Sect.~\\ref{Numerical}. Then, we present detailed comparisons between the simulation results, previous models, and our theoretical predictions for two-point functions in Sect.~\\ref{Lensing-two-point} and three-point functions in Sect.~\\ref{Lensing-three-point}. We study the relative importance of the different contributions arising from ``one-halo'', ``two-halo'',  or ``three-halo'' terms in Sect.~\\ref{contributions}. Then, we check the robustness of our model when we vary the cosmological parameters in Sect.~\\ref{Cosmology} and we briefly study multi-scale moments in Sect.~\\ref{multi-scale}. Finally, we conclude in Sect.~\\ref{Conclusion}. ",
        "conclusions": "\\label{Conclusion}     We have investigated the performance of current theoretical modeling of the 3D matter density distribution with respect to weak-lensing statistics, focusing on configuration-space statistics, specifically the convergence two-point and three-point correlation functions and the second- and third-order moments of the smoothed convergence and of the aperture mass. As in paper I, where we studied Fourier-space statistics, we found that a model introduced in previous works \\citep{Valageas2011d,Valageas2011e}, which combines the (resummed) one-loop perturbation theory with a halo model, fares better than some other recipes based on fitting formulae to numerical simulations or more phenomenological approaches. It yields a reasonable agreement with numerical simulations and provides a competitive approach, because it remains difficult and time-consuming to describe a range of scales that spans three orders of magnitude or more by ray-tracing simulations.  One advantage of our approach compared with numerical simulations or fitting formulas is that it allows us to decompose the integrated weak-lensing signal over several contributions that are associated with specific properties of the underlying 3D density field. Thus, we can distinguish perturbative terms, which can be derived from perturbation theory, from nonperturbative terms that are associated for instance with one-halo contributions, which depend on the density profile and mass function of virialized halos. This is useful because i) these different terms suffer from different theoretical uncertainties and ii) it allows one to understand which aspects of the matter distribution are probed by weak-lensing statistics, while angular scales vary.  Like for the Fourier-space statistics studied in paper I and the 3D statistics studied in \\citet{Valageas2011d,Valageas2011e}, we found that including one-loop terms in the perturbative contribution brings a more significant improvement compared with the lowest-order perturbation theory for three-point statistics than for two-point statistics. Then, while large scales are described by these perturbative contributions and small scales by one-halo contributions, the nonperturbative two-halo term that gives an additional contribution to three-point statistics is always subdominant. This is a nice property because this mixed term is more difficult to model and may be less accurate than other contributions (see also paper I).  Consequently, our model provides reliable predictions for weak-lensing statistics, from small to large scales, and for a variety of cosmologies. It could still be improved in various manners. First, the accuracy of the perturbative contribution may be increased by including higher orders beyond one-loop or by using alternative resummation schemes. Second, the underlying halo model could be refined to include substructures \\citep{Sheth2003a,Giocoli2010}, deviations from spherical profiles \\citep{Jing2002,Smith2006}, or the effect of baryons \\citep{Guillet2010}. Next, the model could be generalized to non-Gaussian initial conditions, which yield distinctive signatures in the bispectrum \\citep{Sefusatti2010}.    \\acknowledgement  We would like to thank Takashi Hamana for helpful discussions. M.S. and T.N. are supported by a Grant-in-Aid for the Japan Society for Promotion of Science (JSPS) fellows. This work is supported in part by the French ``Programme National de Cosmologie et Galaxies'' and the French-Japanese ``Programme Hubert Curien/Sakura, projet 25727TL'', the JSPS Core-to-Core Program ``International Research Network for Dark Energy'', a Grant-in-Aid for Scientific Research on Priority Areas No. 467 ``Probing the Dark Energy through an Extremely Wide and Deep Survey with Subaru Telescope'', a Grant-in-Aid for Nagoya University Global COE Program, ``Quest for Fundamental Principles in the Universe: from Particles to the Solar System and the Cosmos'', and World Premier International Research Center Initiative (WPI Initiative), MEXT, Japan. We acknowledge Kobayashi-Maskawa Institute for the Origin of Particles and the Universe, Nagoya University for providing computing resources. Numerical calculations for the present work have been in part carried out under the ``Interdisciplinary Computational Science Program'' in Center for Computational Sciences, University of Tsukuba, and also on Cray XT4 at Center for Computational Astrophysics, CfCA, of National Astronomical Observatory of Japan."
    },
    "1112/1112.6083_arXiv.txt": {
        "abstract": "The composite spectra of quasars are widely used as templates for redshift determination, as well as for measurements of the mean transmission in Ly$\\alpha$-forest studies, and for investigation of general spectral properties of quasars. Possibility of composite spectra utilisation in these fields is related to remarkable similarity of quasar spectra in UV-optical range. But despite of general similarity in spectral shapes, they differ in several parameters, one of which is the spectral index. In the present paper we study the possible effects, related to neglect of this difference. We compiled 16 composite spectra from subsamples of individual SDSS DR7 quasar spectra with different spectral indices $\\alpha_{\\lambda}$ within the wavelength range  1270--1480\\,\\AA, and show that (i) the redshifts measured for a test sample of high signal-to-noise ratio quasar spectra using these composites as templates appear to be systematically higher than those calculated with a traditional template, compiled from spectra with different $\\alpha_{\\lambda}$, with 1.5 times smaller errors in the former case; (ii) the difference in $\\alpha_{\\lambda}$ in individual spectra used for compilation of composites can yield the mean transmission uncertainty up to 20\\%; (iii) a number of emission lines indistinguishable in ordinary composites, but seen in individual high-resolution spectra, can be detected in such composites. It is also shown, that there is no dependence of $\\alpha_{\\lambda}$ on quasar luminosity in SDSS $u$, $g$, $r$ and $i$ bands, and monochromatic luminosity at $1450$\\,\\AA. ",
        "introduction": "\\label{sec:1} \\indent\\indent The spectral properties of quasars in ultraviolet-optical range have been a subject of a number of studies. The smooth part (continuum) of this UV bump, called the ``Big Blue Bump'', is traditionally considered to be thermal emission from an accretion disk, while the emission lines are attributed to this emission reprocessed in surrounding clumped gas in a form of broad and narrow line regions (BLR and NLR) as well as in a dust obscured torus, but many details of this paradigm are still purely understood, as well as the nature of broad absorption lines (BAL), possessed by about 10\\% of quasars. Thus, on the one hand, the detailed study of the quasar spectra plays an important role in understanding of the quasar nature. On the other hand, the accurate quasar spectrum shape, which does not vary significantly from one object to another, is needed as a template for quasar redshift measurements in large redshift surveys. Finally, the precise knowledge of the intrinsic quasar spectrum shape, prior to absorption in intergalactic medium, is required for studying the intergalactic gas, manifested as absorption features in the quasar spectra, first of all as the Hydrogen Ly$\\alpha$-forest.  In all of the fields, mentioned above the composite (mean) spectra are widely used. The possibility of their utilisation is related to the fact that quasar spectra are remarkably similar from one object to another. Due to their high signal-to-noise ratio composite spectra reveal weak features that are rarely detectable in individual quasar spectra. The composite spectra were compiled for a wide set of quasar samples, e.\\,g. the Large Bright Quasar Survey (LBQS) \\citep{francis+91}, the First Bright Quasar Survey (FBQS) \\citep{brotherton+01}, the Sloan Digital Sky Survey (SDSS) \\citep{vandenberk+01,pieri+10}, quasar spectra from the Hubble Space Telescope (HST) \\citep{zheng+97,telfer+02}, and Far Ultraviolet Spectroscopic Explorer (FUSE) \\citep{scott2+04}.  The part of UV-optical quasar spectrum redward of the Ly$\\alpha$ (1215.67\\,\\AA) emission line (free of hydrogen Lyman-series forests of absorption lines) is studied quite well. As it was shown e.\\,g. by \\citet{vandenberk+01} this region up to $\\approx5000$\\AA\\ is described with a smooth continuum, well approximated with a power law $\\sim\\lambda^{\\alpha_{\\lambda}}$ over some limited wavelength range, broad emission and, in some cases, absorption lines. \\citet{vandenberk+01} reported on $\\alpha_{\\lambda}$ to be $-1.54$ within 1350--4230\\,\\AA, that agrees well with the values $-1.54$ and $-1.68$ which were obtained by \\citet{brotherton+01} within 1450--5050\\,\\AA\\ for spectra from LBQS and FBQS respectively, and also with the results of \\citet{carballo+99}, who obtained the values $-1.34\\pm0.15$ and $-2.11\\pm0.16$ within 1300--3000\\,\\AA\\ and 3000--4500\\,\\AA\\  ranges respectively for a sample of radio-loud quasars. Throughout this paper the indicated wavelengths are rest-frame ones, unless other is specified. The most complete to-date list of emission lines redward of the Ly$\\alpha$ emission line found in quasar composite can be found in \\citealt{vandenberk+01}. Some statistical studies were also conducted with samples of individual spectra, e.\\,g. \\citet{tang+12} studied the strong UV and optical emission line properties using a sample of 85 bright quasars.  The part of the spectrum blueward of the Ly$\\alpha$ emission line, used for the Ly$\\alpha$-forest studies, is much worse investigated due to the presence of the Lyman series forests. Direct reconstruction of intrinsic quasar spectral shape in this region is possible only for high resolution spectra, the number of which is not large, that does not allow to conduct statistical studies of the subsamples with different characteristics, e.\\,g. luminosity. Several-parametric modelling of the continuum level within the Ly$\\alpha$ forest region suffers from the degeneration between the mean transmission and parameters of lines and continuum, that is why extrapolation of the continuum from the red part of spectrum seems to be more reliable in some cases (see e.\\,g., \\citealt{press+93,bernardi+03,desjacques_07,greg+10,songaila-04}). But in practice it is only an approximation of the real continuum, because the centre of the ``Big Blue Bump'' is located at $\\approx1000-1300$\\,\\AA. This is confirmed by the values of $\\alpha_{\\lambda}$ obtained from HST and FUSE spectra, $-1.01\\pm0.05$ within 1050--2020\\,\\AA\\ \\citep{zheng+97}, $-0.24\\pm0.12$ within 500--1200\\,\\AA\\ \\citep{telfer+02}, $-1.44^{+0.38}_{-0.28}$ within 630--1155\\,\\AA\\ \\citep{scott2+04}, when comparing with the same values for longer wavelengths mentioned above. Due to lower signal-to-noise ratio in the Lyman series forests region detection of emission lines is more complicated here. For example, \\citet{vandenberk+01} listed only three emission features between Ly$\\alpha$ and Ly$\\beta$ (1025.72\\,\\AA) with large uncertainties in the central wavelength. The other studies of composites and individual spectra have not enlarged significantly their number. In Table~3 we tried to summarise the known emission features, found in optical \\citep{francis+91,brotherton+01,vandenberk+01,tytler+04} and UV \\citep{zheng+97,telfer+02,scott2+04} composites, and in some individual spectra \\citep{brotherton+94,Laor+94,Laor+1995,laor+97,Vestergaard+01,leighly+07,binette+08}, within the range of $\\sim1025-1450$\\,\\AA\\ (between the Ly$\\beta$ and C{\\sc iv} lines), used for studies in the present work.  The spectral properties of quasars described above can be considered only as some generalised ones, because despite of general similarity in spectral shapes, they differ in spectral index (also called spectral slope), intensity of continuum and emission lines, lines' equivalent widths, etc. Some of these parameters are found to have some dependence on each other, e.\\,g., the inverse correlation of equivalent width of some emission lines on monochromatic flux in UV region, known as the Baldwin effect \\citep{baldwin}. Therefore, compilation of composite spectra from subsamples with similar properties (e.\\,g., with similar luminosity) is of interest for studying these properties and there relation to other parameters. E.\\,g., \\citet{richards+11} studied the Baldwin effect constructing composite spectra of quasar subsamples, chosen on the base of C{\\sc iv} (1549\\,\\AA) emission-line properties (equivalent width and blueshift). They also found, that the mean blueshift of C{\\sc iv} line is approximately twice larger for radio-quite quasars, than that of radio-loud ones. \\citet{reichard+03} studied the spectral properties of Broad Absorption Line (BAL) quasars, compiling composites from subsamples with different ionization levels of BALs, and found that there are substantial differences in the emission-line and continuum properties of high- and low-ionization BAL quasars, which, in their opinion, can be related to intrinsic quasar properties such as the continuum spectral index. Also using the composite spectra, \\citet{yip+04,vandenberk+04} studied the spectral properties of quasars depending on luminosity and redshift, and concluded that aside from the Baldwin effect the average spectral properties are similar, e.\\,g. the quasars with different luminosity have similar spectral index.  In the present paper we use the composite spectra compiled for subsamples of spectra with similar spectral index and study the possible effects, related to neglect of the difference in this parameter.  The data used in this work, the sample selection and the composite spectra compilation processes are described in Section~\\ref{sec:2}. The difference in redshift measurements made with the obtained composite spectra as templates compared to that with the SDSS quasar template is studied in Section~\\ref{sec:3}. Section~\\ref{sec:4} gives an estimation of errors introduced into the Ly$\\alpha$ forest mean transmission measurement by neglect the difference in $\\alpha_{\\lambda}$ when using the composite spectra for such studies. In Section~\\ref{sec:5} we present the results of our search for emission lines in our composites. Conclusions are presented in Section~\\ref{sec:6}.  \\section[]{Data}\\label{sec:2} \\subsection{The SDSS DR7 quasar sample}\\label{sec:2.1}  \\indent\\indent  Our sample is taken from the public available release of the sky-residual subtracted spectra for the Sloan Digital Sky Survey (SDSS) Legacy Release 7 \\citep{Abazadjian_2009}. This release (`WH sample' hereafter), which is described in \\citet{wild_10}, contains a total of 106\\,006 spectra, generated using the \\citet{Wild_05} scheme from the spectra of the objects from the \\citet{Schneider_2010} quasar catalogue. This scheme includes a significantly improved technique of the OH sky lines subtraction. The strong OH sky emission lines extend over almost half of the wavelength range ($>6700$\\,\\AA) of the SDSS spectra and is not subtracted optimally enough with the previous SDSS pipelines.  The redshifts in SDSS are measured mainly with the help of two techniques --- emission line measurements and cross-correlation, or one of them. The errors of redshift (due to intrinsic line shifts etc.) could affect the results of our study hence we tried to minimize their influence using the improved redshifts of the objects from the \\citet{Schneider_2010} quasar catalogue, which was generated using the \\citet{hewett_10} scheme, instead of redshifts in the headers of each fits-file. According to \\citet{hewett_10}, their redshifts possess systematic biases, which are for a factor of 20 smaller compared to the SDSS redshift values.  \\subsection{Spectra selection}\\label{sec:2.2}  \\begin{figure} \\centering \\epsfig{figure=zdistr.eps,width=.98\\linewidth} \\caption{Redshift distribution of all objects from the HW sample with $2.3<z<4.6$ and redshift determination confidence level $>0.9$ (\\textit{dash-dot}), all visually selected (\\textit{dotted}) for the present study and those with the rms of the normalization constant $A$ less than 15\\% (\\textit{solid}).} \\label{fig:zdistr} \\end{figure}  \\indent\\indent All the quasar spectra (15\\,154~objects) from the HW sample with redshifts within the range 2.3$-$4.6 and redshift determination confidence level $>0.9$ were firstly selected. Because the SDSS is an automatic survey there is a possibility of pollution of the sample with `wrong' objects whose general photometry or spectroscopy properties required by the automatic selection pipelines are similar to the ones of `real' objects. Thus a visual examination was carried out. Except the non-quasar spectra the spectra of quasars with low signal-to-noise ratio, spectra of BAL quasars and those with the Damped Ly$\\alpha$ (DLA) systems were also excluded during this examination. The presence of the BAL quasars in the sample could introduce additional errors into an estimation of the quasar mean continuum level. The DLA systems are a class of quasar absorbers selected for the presence of H~{\\sc i} column densities $>2\\cdot10^{20}$~cm$^{-2}$ (see e.\\,g.~\\citet{wolfe_05} for review). They are identified as absorption features with the rest equivalent width exceeding 5\\,\\AA. The nature of these systems is still not understood, however due to their high density they are usually related to the galaxy formation and therefore could not be used as representatives of the linear perturbations in the neutral intergalactic medium.  After the visual examination the sample contains 4\\,779~spectra. Redshift distributions of these objects and all the objects initially selected from HW sample are presented in Figure~\\ref{fig:zdistr} (\\textit{dash-dot} and \\textit{dotted} correspondingly).  \\subsection{Subsamples and composites}\\label{sec:2.3}  \\indent\\indent Generating composite spectra of quasars includes three steps: (i) normalization of each spectrum, (ii) setting each spectrum to the rest frame and (iii) calculation of the mean spectrum. Before we smoothed all spectra with a simple moving average over 3 points. Following \\citet{mcdonald+06} we removed the following wavelength regions from our analysis because of calibration problems due to strong sky lines: $5575\\mbox{\\AA}<\\lambda<5583\\mbox{\\AA}$, $5888\\mbox{\\AA}<\\lambda<5895\\mbox{\\AA}$, $6296\\mbox{\\AA}<\\lambda<6308\\mbox{\\AA}$, and $6862\\mbox{\\AA}<\\lambda<6871\\mbox{\\AA}$, where the second region is 1\\AA\\ wider than that presented in \\citet{mcdonald+06}, because we added also Na{\\sc i} (5894.6\\,\\AA) interstellar line to it. Normalization of each spectrum is needed to be applied due to different apparent flux densities. Taking into account the similarity of quasar spectra and following \\citet{press+93} and \\citet{zheng+97} we normalize each spectrum on the (arithmetic) mean flux in all pixels within the rest wavelength range 1450--1470\\,\\AA. This range lies blueward from the C{\\sc iv} emission line and is usually considered to be free of obvious emission and absorption. The last claim appears to be not precise because of weak emission lines, which can be seen in composite spectra, but we neglect this fact, because intensity of these lines is comparable to noise level in individual spectra. For further study to reduce possible uncertainties we used only spectra with the rms of the normalization constant $A$ less than 15\\%; the number of them is 3\\,493 and its redshift distribution is shown in Figure~\\ref{fig:zdistr} (\\textit{solid} line).  Considering the continuum redward of Ly$\\alpha$ emission line to be a power-law $\\sim\\lambda^{\\alpha_{\\lambda}}$, we calculated its index $\\alpha_{\\lambda}$ for each individual spectrum within the range between Ly$\\alpha$ and C{\\sc iv} emission lines. For this purpose we selected the following wavelength ranges from the composite spectrum compiled from all the spectra: 1278$-$1286\\,\\AA\\AA, 1320$-$1326\\,\\AA\\AA, 1345$-$1360\\,\\AA\\AA\\ and 1440$-$1480\\,\\AA\\AA. The distribution of the obtained indices is shown in Figure~\\ref{fig:alpha-distr}.  \\begin{figure} \\centering \\epsfig{figure=alphadistr.eps,width=.98\\linewidth} \\caption{The distribution of spectral indices.} \\label{fig:alpha-distr} \\end{figure}  \\begin{figure*} \\centering \\begin{minipage}{.75\\linewidth} \\epsfig{figure=spec-all.eps,width=.99\\linewidth} \\caption{Composite spectra of 16 subsamples within 1025--1650\\,\\AA. The highest spectrum has the steepest continuum redward of the Ly$\\alpha$ emission line.} \\label{fig:all-spec} \\end{minipage} \\end{figure*}  Using the obtained values of $\\alpha_{\\lambda}$, we selected 16 subsamples each of 200 spectra with $\\alpha_{\\lambda}$ closest to $-0.7-k\\cdot0.1$, $k=0..15$. Then dividing each $j$-th selected spectrum onto its normalization constant $A^{j}$, rebinning them with $\\Delta\\lambda_{rest}=2$\\,\\AA\\ and stacking the spectra into the rest frame, we obtained the mean arithmetic composite spectra. The dispersion $\\sigma^{2}$ of each pixel of the composite spectrum was calculated from the noises $\\sigma_{i}$ of pixels of individual spectra constituting it. These spectra within 1025--1650\\,\\AA\\ are shown in Figure~\\ref{fig:all-spec}. The dashed lines indicate the laboratory wavelengths of the lines identified by \\citet{vandenberk+01} in this range.  \\subsection{Properties of subsamples}\\label{sec:2.4}  \\indent\\indent The first row of Table~\\ref{tab:samples} contains the spectral indices of obtained composite spectra. In this case they were calculated in the way mentioned above, but the parts of the spectra which are the most free from emission lines were selected manually in each spectrum. The part of spectra used for $\\alpha_{\\lambda}$ determination is shown in Figure~\\ref{fig:lena}. The regions selected for continuum fitting slightly vary from spectrum to spectrum and not all four parts used in calculation of indices of the individual spectra were taken into account in the case of composite spectra. This is clearly seen on example of the region with the shortest wavelengths: with the steepening of the spectrum its centre moves from $\\sim1280$\\,\\AA\\ to $\\sim1290$\\,\\AA. It means that determination of the spectral index using the same regions as it was done for individual spectra is only an approximation, but it is sufficient for rough separation of spectra for compilation of composites.  \\begin{figure} \\centering \\epsfig{figure=spec-lena.eps,width=.98\\linewidth} \\caption{The parts of the composite spectra used for continuum approximation (\\textit{short-dashed} line). The highest spectra are the steepest ones. Spectra with indices $-0.91$, $-1.62$ and $-2.14$ are shown with the \\textit{solid} lines along with the fitted power-law continua (\\textit{dashed} lines).} \\label{fig:lena} \\end{figure}  The second row of Table~\\ref{tab:samples} contains the mean redshift of each subsample. It is seen, that the mean redshift slightly increases with $\\alpha_{\\lambda}$, that might be an evidence for redshift evolution of the quasar spectral shape. On the other hand, the absence of such evolution in Fig.~\\ref{fig:z-slo} for the whole sample of 3\\,493 quasars means, that increase of the mean redshift for subsamples is probably a result of some selection effect.  \\begin{table} \\begin{minipage}{80mm} \\centering \\caption{The mean redshift $\\bar{z}$ and the mean monochromatic luminosity $\\langle\\log{l_{1450}}\\rangle$ of subsamples and the spectral indices $\\alpha_{\\lambda}$ of their composites. The quoted error bars of $\\langle\\log{l_{1450}}\\rangle$ and $\\bar{z}$ are the root mean squares for corresponding distributions.}\\label{tab:samples} \\begin{tabular}{c|c|c|c} \\hline n & $\\alpha_{\\lambda}$ & $\\bar{z}$ & $\\langle\\log{l_{1450}}\\rangle$ \\\\ \\hline 1 & $-0.91\\pm0.03$ & $2.90\\pm0.54$ & $42.70\\pm0.21$ \\\\ 2 & $-0.97\\pm0.02$ & $2.89\\pm0.53$ & $42.72\\pm0.21$ \\\\ 3 & $-1.02\\pm0.01$ & $2.87\\pm0.51$ & $42.72\\pm0.21$ \\\\ 4 & $-1.04\\pm0.05$ & $2.85\\pm0.49$ & $42.72\\pm0.24$ \\\\ 5 & $-1.19\\pm0.07$ & $2.82\\pm0.43$ & $42.73\\pm0.24$ \\\\ 6 & $-1.35\\pm0.13$ & $2.84\\pm0.49$ & $42.74\\pm0.21$ \\\\ 7 & $-1.42\\pm0.05$ & $2.82\\pm0.44$ & $42.76\\pm0.23$ \\\\ 8 & $-1.42\\pm0.02$ & $2.76\\pm0.40$ & $42.78\\pm0.21$ \\\\ 9 & $-1.62\\pm0.14$ & $2.77\\pm0.44$ & $42.78\\pm0.21$ \\\\ 10 & $-1.64\\pm0.15$ & $2.80\\pm0.46$ & $42.79\\pm0.27$ \\\\ 11 & $-1.73\\pm0.09$ & $2.75\\pm0.45$ & $42.78\\pm0.25$ \\\\ 12 & $-1.88\\pm0.07$ & $2.75\\pm0.43$ & $42.78\\pm0.22$ \\\\ 13 & $-1.92\\pm0.11$ & $2.74\\pm0.43$ & $42.77\\pm0.24$ \\\\ 14 & $-2.02\\pm0.05$ & $2.72\\pm0.42$ & $42.74\\pm0.23$ \\\\ 15 & $-2.07\\pm0.07$ & $2.74\\pm0.42$ & $42.74\\pm0.23$ \\\\ 16 & $-2.14\\pm0.05$ & $2.74\\pm0.43$ & $42.75\\pm0.24$ \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table}   \\begin{figure*} \\centering \\epsfig{figure=slopes-Mu.eps,width=.98\\linewidth} \\caption{Spectral index ($\\alpha_{\\lambda}$) -- absolute magnitude ($M$) diagrams for \\textit{u,\\,g,\\,r,\\,i}-bands.} \\label{fig:slo-mu} \\end{figure*}  To study any optical-UV luminosity dependence of the spectral index we plotted $\\alpha_{\\lambda}$\\,--\\,$M$ diagrams for absolute magnitudes $M$ in $u$ (3551\\,\\AA), $g$ (4686\\,\\AA), $r$ (6165\\,\\AA), $i$ ($7481$\\,\\AA) bands (Fig.~\\ref{fig:slo-mu}). These magnitudes were calculated within the frame of the flat $\\Lambda$CDM cosmological model with $H_0=70$\\,km/s/Mpc, $\\Omega_M=0.3$, using the apparent psf magnitudes and reddening in corresponding bands from the SDSS DR7, K-correction values from \\citet{richards_06} and correction for Galactic extinction from \\citet{schlegel+98}. As one can see from  Fig.~\\ref{fig:slo-mu} there is no dependence between the spectral index and absolute magnitudes in these bands, however even the $u$-band central wavelength is about 2000\\,\\AA\\ far from the bump peak, thus all these magnitudes are not a good characteristics in case of the UV-bump luminosity. Hence we also plotted the $\\alpha_{\\lambda}$\\,--\\,$\\log{l_{1450}}$ diagram, where $\\log{l_{1450}}$ is the monochromatic luminosity at 1450\\AA, calculated from the mean flux within the wavelength range of $1449-1451$\\,\\AA.  Generally speaking this range is also not the best choice, but this one is the most free of obvious emission and absorption features, thus it should be the most appropriate characteristics the quasar luminosity in continuum within the bump. Any dependence of $\\alpha_{\\lambda}$ on monochromatic luminosity at 1450\\,\\AA\\ can be seen from Fig.~\\ref{fig:nor-slo}, as well as from its mean values listed in the third row of Table~\\ref{tab:samples}.  \\begin{figure} \\centering \\epsfig{figure=z-alpha.eps,width=.98\\linewidth} \\caption{Redshift distribution of spectral indices.}\\label{fig:z-slo} \\end{figure}  \\begin{figure} \\centering \\epsfig{figure=norma-slope-abs.eps,width=.98\\linewidth} \\caption{Spectral index ($\\alpha_{\\lambda}$) -- logarithm of monochromatic luminosity ($\\log{l_{1450}}$) diagram.} \\label{fig:nor-slo} \\end{figure} ",
        "conclusions": "\\label{sec:6}  \\indent\\indent We compiled 16 composite spectra from subsamples of individual SDSS DR7 quasar spectra with different spectral indices $\\alpha_{\\lambda}$ within the wavelength range  1270--1480\\,\\AA\\ and studied the possible effects caused by neglect of this discretization when using composite spectra of quasars in different fields. The main results of the present work are the following:\\\\ \\indent(i) the redshifts measured for a test sample of high signal-to-noise ratio quasar spectra using these composites as templates appear to be systematically higher than those calculated with a traditional template, compiled from spectra with different $\\alpha_{\\lambda}$, with 1.5 times smaller errors in the former case; \\\\ \\indent(ii) the difference in $\\alpha_{\\lambda}$ in individual spectra used for compilation of composites can yield the mean transmission uncertainty up to 20\\%; \\\\ \\indent(iii) a number of emission lines indistinguishable in ordinary composites, but seen in individual high-resolution spectra, can be detected in such composites; \\\\ \\indent(iv) it is confirmed that continuum level within the range between Ly$\\beta$ and Ly$\\alpha$ emission lines, utilized for studies of the Ly$\\alpha$-forest, can be considered as a constant rather than having the same power-law form as that redward from Ly$\\alpha$ emission line, that agrees well with steeper continuum index obtained from the composite UV quasar spectra \\citep{telfer+02,zheng+97}, and was discussed previously \\citep{desjacques_07};\\\\ \\indent(v) it is also shown, that there is no dependence of $\\alpha_{\\lambda}$ on quasar luminosity in SDSS $u$, $g$, $r$ and $i$ bands, and monochromatic luminosity At $1450$\\,\\AA. It confirms results of \\citet{yip+04} and \\citet{vandenberk+04}, who found no relation between luminosity and spectral index, when analysing composite spectra with different luminosity.  Detailed analysis in our previous study of the ranges redward of the Ly$\\alpha$ emission lines, which is free from the Ly$\\alpha$-forest, shown that there is also no evidence for spectral index dependence of equivalent width of emission lines \\citep{torbaniuk+12}. The absence of dependence of the UV-bump shape on luminosity in the bands mentioned above requires further study of this region for understanding the physics behind the difference in the UV-bump shape in spectra of different quasars. It is worth to note, that the absolute magnitudes in given bands used as a characteristic of luminosity include K-correction, which itself is determined within the frame of some model of the spectrum general shape (e.\\,g. with the spectral index of $\\alpha_{\\lambda}=-1.5$).  The proposed approach can be applied for generation of new templates for more precise quasar redshift measurements with the common cross-correlation technique used in redshift surveys, more precise theoretical determination of K-correction and colour-indices, as far as for determination of continuum and mean transmission in Ly$\\alpha$ forest studies."
    },
    "1112/1112.0170_arXiv.txt": {
        "abstract": "This is an account of Allan Sandage's work on (1) The character of the expansion field. For many years he has been the strongest defender of an expanding Universe. He later explained the CMB dipole by a local velocity of  $220\\pm50\\kms$ toward the Virgo cluster and by a bulk motion of the Local supercluster (extending out to $\\sim\\!3500\\kms$) of $450 - 500\\kms$ toward an apex at $l=275$, $b=12$. Allowing for these streaming velocities he found linear expansion to hold down to local scales ($\\sim300\\kms$). (2) The calibration of the Hubble constant. Probing different methods he finally adopted -- from Cepheid-calibrated SNe\\,Ia and from independent RR~Lyr-calibrated TRGBs -- $H_{0}=62.3\\pm1.3\\pm5.0\\kms\\;$Mpc$^{-1}$. ",
        "introduction": "\\label{sec:1} Edwin \\citeauthor{Hubble:29} is generally credited for the discovery of the expansion of the Universe. But as so often in the case of fundamental discoveries, others had contributed. G.~Lemaitre had published a value of the expansion rate (Hubble constant, $H_{0}$) already in \\citeyear{Lemaitre:27}, and H.~P.~Robertson once laconically told Sandage, Hubble found the expansion because I told him. In fact, Robertson had published his value of $H_{0}$ already in \\citeyear{Robertson:28}. Hubbles most astounding achievement is to have convinced the World of the expansion of the Universe with his brilliant monograph \\textit{The Realm of the Nebulae} \\citeyearpar{Hubble:36a}; he had by then much better cards than in 1929 because he had extended with the help of Milton Humason the $\\log$ redshift-apparent magnitude diagram (Hubble diagram) to $19,000\\kms$ \\citep{Hubble:Humason:34}, but his value of $H_{0}$ was still to high by a factor of roughly 8, and correspondingly his expansion age was impossibly short -- a problem which he elegantly managed to bypass. Paradoxically \\citeauthor{Hubble:36b} began to question the reality of the expansion in the same year as his book appeared because he could not make sense of his galaxy counts. His doubts persisted until his death as evidenced in his Darwin Lecture -- posthumously edited by Sandage -- where \\citeauthor{Hubble:53} showed a Hubble diagram including \\citeauthor*{Humason:51} large-redshift clusters out $61,000\\kms$ (Fig.~\\ref{fig:01}) with the remark `no recession factor (applied)', which means that he had corrected the galaxy magnitudes for a single factor of $z$, but not for the $z^{2}$-term required in any expanding model.  \\begin{figure}[t] \\epsscale{0.96} \\epsscale{0.98} \\plotone{fig01.eps} \\caption{Hubbles last Hubble diagram with unrealistically small scatter and the remark no recession factor which documents his doubts about the expansion of the Universe.} \\label{fig:01} \\end{figure}  A definitive description of the expansion had to proceed along two lines. (1) The expansion field had to be mapped in different directions and out to truly cosmic distances -- allowing for deceleration and/or acceleration -- to test whether the expansion is linear, which means that it is observed as the same by any observer in the Universe.\\footnote{% Note: linear expansion does not require a linear Hubble line.}  (2) Only then it is meaningful to search for the cosmic value of $H_{0}$, which in turn would provide the first cosmological test, i.e.\\ the expansion age of the Universe as compared with independent geological and astrophysical age determinations. Sandage has contributed to these two topics more than anybody else, although only about one fourth of his papers are devoted to them.  The remaining possibility that redshifts are \\textit{not} caused by the cosmic expansion has been disproved later by \\citeauthor{Sandage:10} in a series of papers on the difficult Tolman test \\citep[][and references therein]{Sandage:10} which requires that the surface brightness of a galaxy within a \\textit{metric} radius decreases with $z^{-4}$.  \\begin{figure}[t] \\epsscale{0.80} \\plotone{fig02.eps} \\caption{The Hubble diagram of first-ranked cluster galaxies from HMS. The curved Hubble line corresponds to a closed Universe with $q_{0}=2.5$.} \\label{fig:02} \\end{figure} ",
        "conclusions": ""
    },
    "1112/1112.0493_arXiv.txt": {
        "abstract": "We consider the possibility that along the thermal history of the Universe, dark matter (DM) would have been created from Standard Model particles, either through a kinetic mixing portal to an extra $U(1)$ gauge field, or through the Higgs portal. Depending solely on the DM particle mass, on the portal and on the DM hidden sector interaction, we show how the observed DM relic density can be obtained. There are four possible freeze-in/reannihilation/freeze-out regimes, which together result in a simple characteristic relic density phase diagram, with the shape of a ``Mesa\". In the case of the kinetic mixing portal, we show that, unlike other freeze-in scenarios discussed in the literature, the freeze-in regime can be probed by forthcoming DM direct detection experiments. These results are well representative {of} any scenario where a DM hidden sector would be created out of the Standard Model {sector}.         \\vspace{2mm} ",
        "introduction": "There is an all series of gravitational evidences for the existence of dark matter (DM). In particular, from observations of the Cosmic Microwave Background  anisotropies and large-scale structures, it appears that a large fraction of the energy density of the Universe is made of DM, with $\\Omega_{DM} =0.229 \\pm 0.015$ \\cite{Komatsu:2010fb}. To explain such a proportion, the WIMP candidates are particularly interesting because, by producing  thermally the DM from the thermal bath of the Universe, through DM pair creation and annihilation, they lead to a relic density which depends only on the masses and interactions involved at the time of the decoupling of these processes. Moreover it turns out that for interactions of order the electroweak ones, or more generally for couplings of order unity, they lead to the observed relic density for a DM mass of the order of the electroweak scale. This offers the possibility of DM production at colliders, on top of the possibilities of direct and indirect detection.  Another simple mechanism to produce the DM relic density, which was pushed forward recently \\cite{McDonald:2001vt,Hall:2009bx,Cheung:2010gj,Cheung:2010gk}, is the freeze-in mechanism. {Here the relic density, instead of resulting from the thermal freeze-out of the DM annihilation process, stems solely from the DM creation processes, for instance through annihilation into DM particles, $A A\\rightarrow DM DM$, or through a decay process $A\\rightarrow DM B$.} These processes, {that are} assumed to be out of thermal equilibrium, {freeze} when the temperature $T$ drops below the mass of the source particle $A$ or DM particle, {\\em i.e.}~when the production rate gets Boltzmann suppressed. {While being out-of-equilibrium}, this {mechanism} is {actually} also thermal in the sense that the source particle $A$ is assumed to be in thermal equilibrium, so that, here too, the number of particles produced depends only on the masses and couplings involved in the DM creation process. A difference though with the usual freeze-out scenario is that the relic density may also receive a contribution from a \"primordial\" density, which is not washed-out by any process since the DM creation process is out of thermal equilibrium. However there exists a whole regime for which the DM density produced from the Standard Model (SM) sector can be the dominant one.\\footnote{If both sectors are feebly connected, one could argue that the inflaton is likely to belong to only one of the sector, so that it couple dominantly to one sector and reheating occurs dominantly in this sector.} This is the possibility  we will consider in the present work. {Another important} difference with freeze-out is that the {characteristic} coupling required to produce the relic density through freeze-in is tiny, {typically} about 8-10 orders of magnitude below {that} required by freeze-out, implying that the DM basically lies in a hidden sector that is feebly coupled to the SM {one}. {Hence it is {\\em a priori} very difficult to probe experimentally the freeze-in scenario, either at colliders or from direct/indirect detection.} There is nevertheless one way one could consider to test this mechanism in specific scenarios \\cite{Hall:2009bx,Cheung:2010gj,Cheung:2010gk}{: if freeze-in proceeds through $A \\rightarrow DM B$ decay (assuming typically that $A$ and DM particles are both odd under a $Z_2$ symmetry so that DM is stable) and if, in the visible sector, $A$ has substantial coupling to the SM particles, then $A$ particles may be abundantly produced at colliders. Their slow} decay to DM from the feeble interaction could then be probed from the search of long-lived particles at colliders. To be efficiently probed, this {scenario}  requires that $B$ itself is substantially coupled  to SM particles (being typically charged or colored), so that only DM  {particles are} feebly coupled to all other particles.  In the following we show that there is another possibility of testing a DM relic density created through freeze-in, {\\em i.e.}~direct detection. The scenario we consider is {quite} different. {In particular} we do not assume any new particles in the visible sector{, but the ones from the SM. This means that in our framework the A particle is a SM particle, which implies that  the predicted DM relic density depends neither on the mass of any new hypothetical particles in the visible sector, nor on the interactions these particles  may have with the rest of the visible sector.} {The DM} relic density depends only on the DM particle mass, on the connector interaction and on any interaction which may be {relevant} within the hidden sector. One class of models we consider is based on the {possible} existence of {a light mediator particle}. We {will show} that such a light mediator interaction, as tiny as the one required for {DM production through} freeze-in, could be probed by direct detection, because the elastic cross section with nuclei, which exhibits a collinear divergence, is strongly enhanced.  The light mediator DM model we consider in the following as an example is very simple. {It involves, in addition to the SM (the visible sector), a hidden sector composed of  a single particle, which is assumed to be charged under a new $U(1)'$ gauge symmetry.} {This hidden sector is coupled to the SM sector through kinetic mixing, ${\\cal L} \\owns -\\frac{1}{2} \\varepsilon F_Y^{\\mu\\nu} F'_{\\mu\\nu}$ (see Ref.~\\cite{Pospelov:2007mp,Chun:2010ve} for a \"secluded\" DM structure of this type with a massive $\\gamma'$, see also the related model in Ref.~\\cite{Feldman:2006wd}). Thus DM  may interact with the SM particles through a massless mediator, or a light mediator if the $U(1)'$ gauge group is slightly broken (either through the Higgs or St\\\"uckelberg mechanisms).} {On top of the connector interaction, this model involves QED-like interactions between the $\\gamma'$ and the charged DM particle in the hidden sector.} {Another possible motivation for this model is that there are not that many simple mechanisms available to stabilize DM (see e.g.~\\cite{Hambye:2010zb}). Among these, one of the simplest possibilities is to consider the lightest particle charged under an extra, unbroken gauge symmetry. Hence the model we invoke here may also be seen as the simplest and yet generic representative of a larger class of models. In particular we expect that the features we will discuss should hold for many  models of the same kind.} {A further motivation is that, although} the cosmological and astrophysical consequences of such a new long range {gauge interaction}  have been considered extensively in Refs.~\\cite{Ackerman:2008gi,Feng:2009mn,Feng:2008mu}, {this has been done} basically without considering the {possibility of} kinetic mixing connecting both sectors.  Beyond the fact that the freeze-in mechanism {may} be {testable within} this {specific} scenario, the main, {more general} purpose of our work is to determine the various ways a viable DM relic density can be obtained in {the context of particle creation through a portal}. To this end we {have explored} the full parameter space of this model. If both the kinetic mixing  and the hidden sector interactions are feeble, the DM relic density can be calculated by simply counting the number of DM particle produced through the connector. This is precisely the freeze-in scenario. {However}, if instead {either one} or both {of the} interactions are {stronger, the DM relic density is affected by the thermalization of the corresponding processes.} We show {that} this leads to a simple characteristic \"Mesa\" phase diagram for the relic density as a function {of the connector and hidden sector couplings}. {As we explain in details this phase diagram displays four distinct} phases: (i) freeze-in, (ii) simultaneous freezing of the (in-equilibrium) hidden sector interactions and the (out-of-equilibrium) connector, {\\em i.e.}~reannihilation, (iii) freeze-out of the hidden sector interaction, and (iv) freeze-out of {connector} interaction. An interesting feature of this diagram is that{, in each phase, the relic density is essentially characterized by  only one interaction, either the connector  or that in the hidden sector}.  In order to address the generality of this classification, we consider the creation of DM through another simple portal between the SM and a hidden DM sectors, the so-called Higgs portal. We show that DM production from SM particles through the Higgs portal leads to a similar simple characteristic phase diagram which also has characteristic ``Mesa'' shape. This holds even though the mediator, {\\em i.e.}~the Higgs boson, is massive and may even be heavier than the DM particle, unlike for the kinetic mixing where we consider the opposite limit of a mediator (much) lighter than the DM candidate. In this sense the \"Mesa\" diagram is a generic feature of models within which DM is created out of SM particles through a portal (or more generally from particles that are in thermal equilibrium with the SM sector when they create the DM particles).   \\bigskip {The plan of this article is as follows. In Section 2 we first compute the transfer of energy from the visible sector to the hidden sector in the case of the kinetic mixing portal. This is necessary to compute the DM relic density in all regimes, except in the freeze-in one}. We then compute the DM number density produced. {In Section 3 we present and discuss in details the corresponding phase diagram obtained and give analytical approximations for the relic density. In Section 4 we discuss the possibility of direct detection signal through kinetic mixing, and the prospect for testing the freeze-in mechanism. In Section 5 we discuss the compatibility of the relic abundance and direct detection signatures obtained with the various existing cosmological and astrophysical constraints on DM interacting through long range interactions. In Section 6 we present our results for the Higgs portal, and discuss in details the similarities and differences between the cases of massive and massless mediators. Finally, we summarize our results in the last section.} ",
        "conclusions": "We have considered the simple possibility that, along the thermal history of the Universe, DM could have been created out of SM particles through a portal connecting the visible sector (SM) to a dark sector. This framework is particularly predictive as it is induced by the SM particles, that is to say particles whose properties (mass and interactions) and thermal number density are known. As a result, the energy transferred from the visible sector to the hidden sector and the DM relic density produced, depend, besides (obviously) the DM mass, only on the strength of the connector interaction and of the interaction(s) which may thermalize the hidden sector.   Whether we consider the gauge kinetic mixing portal, which involves a massless (or massive) mediator, or the Higgs portal, whose mediator is the Higgs doublet, we have found that there are 4 distinct ways of producing DM, leading to a phase diagram with the characteristic shape of  a ``Mesa\". In two of the phases, the DM relic density depends exclusively on the connector interaction, whereas in the two others it depends essentially on the strength of the hidden sector interaction  (provided the connector interaction is within a certain range).   It is useful to summarize the parametric dependence of the relic density for the various phases, which holds for a connector lighter than the DM particle (both in the case of kinetic mixing and Higgs portals): \\begin{eqnarray} \\hbox{I}:&&\\hspace{-4mm}\\hbox{Freeze-in:}\\quad Y\\sim \\frac{(n^{SM}_{eq})^2 \\langle \\sigma_{connect} v \\rangle}{sH}\\Big|_{T=Max[m_{SM},m_{DM}]} \\quad\\hbox{with}\\quad x_f\\sim \\frac{m_{DM}}{\\hbox{Max}[m_{SM},m_{DM}]}\\,\\,\\,\\, \\\\ \\hbox{II}:&&\\hspace{-4mm}\\hbox{Reannihilation:}\\quad Y\\sim \\frac{x_f}{\\langle \\sigma_{HS} v \\rangle \\tilde{m}^2}\\quad\\hbox{with}\\quad x_f\\sim \\log  \\langle  \\sigma_{eff} v \\rangle \\tilde{m}^2\\\\ \\hbox{III}:&&\\hspace{-4mm}\\hbox{Hidden sector freezeout:}\\quad Y\\sim \\frac{x_f}{ \\langle \\sigma_{HS} v \\rangle \\tilde{m}^2}\\quad\\hbox{with}\\quad x_f\\sim \\log  \\langle  \\sigma_{HS} v \\rangle \\tilde{m}^2 \\\\ \\hbox{IV}:&&\\hspace{-4mm}\\hbox{Connector freezeout:}\\quad Y\\sim \\frac{x_f}{ \\langle \\sigma_{connect} v \\rangle \\tilde{m}^2}\\quad\\hbox{with}\\quad x_f\\sim \\log \\langle  \\sigma_{connect} v \\rangle \\tilde{m}^2\\quad\\quad\\quad \\end{eqnarray} with $\\tilde{m}^2=m_{DM} m_{Pl}$. If none of the interactions thermalize, the freeze-in regime holds. If the hidden sector thermalizes but the connector does not, reannihilation occurs. In this regime both interactions freeze at the same time after going through a period of quasi static equilibrium evolution. We have shown that in this case the relic abundance is inversely proportional to the $\\langle \\sigma_{HS} v \\rangle$ hidden sector cross section, but that the freeze-out temperature $x_f$ is determined by an effective cross section which is the geometric mean of both the hidden sector and connector cross sections $\\langle \\sigma_{eff} v \\rangle=\\sqrt{\\langle \\sigma_{HS} v\\rangle \\langle \\sigma_{connect}v\\rangle}$. If instead, both interactions thermalize one gets an ordinary hidden sector or connector interaction freeze-out regime, depending on which of the two interactions is the fastest. Finally if the hidden sector interaction is so small that it never thermalizes, even if the connector thermalizes, one also lies in the connector freeze-out phase. We have also shown that there are no other possible regimes and that, for a non negligible hidden sector interaction, the four regimes necessarily follow each other in this order.  In the opposite case where the mediator has a mass larger than the DM, as applies for the Higgs portal for $m_{DM}\\lesssim m_h/2$,\\footnote{Or for the kinetic mixing portal with $m_{DM}<m_{\\gamma'}/2$, a possibility we didn't consider here but which would give a same pattern.} one gets the same pattern except that there is an additional decay contribution and that the connector interaction is cut-off at the mediator mass. As a result the decay dominates both the freeze-in and reannihilation regimes, and one has also the possibility that the hidden sector interaction freezes out without thermalization of the connector. Correspondingly one has: \\begin{eqnarray} \\hbox{I}:&&\\hspace{-5mm}\\hbox{Freeze-in:}\\quad Y\\sim \\frac{n^{h}_{eq} \\Gamma_h}{sH}\\Big|_{T=m_h} \\quad \\hbox{with}\\quad x_f\\sim \\frac{m_{DM}}{m_h}\\\\ \\hbox{II}:&&\\hspace{-5mm}\\hbox{Reannihilation:}\\quad Y\\sim \\frac{x_f}{\\langle \\sigma_{HS} v\\rangle \\tilde{m}^2}\\quad\\hbox{with}\\quad x_f\\sim \\log [ \\langle  \\sigma_{HS} v \\rangle \\Gamma_h \\frac{m_{Pl}^2 m_h^{3/2}}{m_{DM}^{5/2}}  ]\\\\ \\hbox{IIIA}:&&\\hspace{-5mm}\\hbox{HS freeze-out (no connector thermal.):}\\quad Y\\sim \\frac{x_f}{ \\langle \\sigma_{HS} v \\rangle \\tilde{m}^2}\\quad\\hbox{with}\\quad x_f\\sim \\xi \\log   \\langle  \\sigma_{HS} v \\rangle \\tilde{m}^2\\,\\,\\,\\,\\,\\\\ \\hbox{IIIB}:&&\\hspace{-5mm}\\hbox{HS freeze-out (connector thermal.):}\\quad Y\\sim \\frac{x_f}{ \\langle \\sigma_{HS} v \\rangle \\tilde{m}^2}\\quad\\hbox{with}\\quad x_f\\sim \\log  \\langle  \\sigma_{HS} v \\rangle \\tilde{m}^2\\\\ \\hbox{IV}:&&\\hspace{-5mm}\\hbox{Connector freeze-out:}\\quad Y\\sim \\frac{x_f}{ \\langle \\sigma_{connect} v \\rangle \\tilde{m}^2}\\quad\\hbox{with}\\quad x_f\\sim \\log  \\langle  \\sigma_{connect} v \\rangle \\tilde{m}^2 \\end{eqnarray} with $\\Gamma_h$ the Higgs boson decay width to DM particles, and $\\xi=T'/T=(g_*^{SM}/g_*^{HS})^{1/4}(\\rho'/\\rho)^{1/4}$. The reannihilation, as well as the hidden sector freeze-out regime (without connector thermalization) give a relic density that depends mainly on the size of one interaction, but also logarithmically on the other one, unlike the other regimes for which the relic density depends only on the strength of one interaction. For a heavy mediator, the reannihilation regime is relevant only for a relatively narrow intermediate DM mass region below the mediator mass, $1~\\hbox{GeV} \\lesssim m_{DM} \\lesssim m_h/3$. Therefore, apart for this region, if the hidden sector interaction thermalizes but not the connector, hidden sector freeze-out always occurs. This is different from the low mediator mass regime where in this case ({\\em i.e.}~if $T'<T$) reannihilation always occurs.  One should also emphasize the fact that one also gets a characteristic Mesa phase diagram in the more complex situation in which there are both a massless and a massive mediator. This applies in particular in the kinetic portal scenario, in the intermediate mass range $1$ GeV$\\lesssim m_{DM} \\lesssim m_Z/2$, see Appendix D.   To discuss the kinetic mixing portal case, we have considered a particularly simple model where the hidden sector consists of  a single particle charged under an unbroken $U(1)'$ gauge group, depending only on $m_{DM}$, on the kinetic mixing coupling and on the extra gauge coupling. In addition to the relic density phase diagram, we have discussed in detail the phenomenology it may imply. Direct detection is particularly interesting for this model because it involves a massless or light mediator, which implies that the cross section is inversely proportional to the recoil energy squared and is therefore strongly enhanced at low recoil energies. {We have shown that the latest Xenon100 data exclude  the regimes of freeze-out (both in the hidden (III) and visible regimes (IV)) for candidates with $m_{DM}> \\sim 5$~GeV (see Fig.9). The same data also probe ({\\em i.e.} exclude) a large fraction of the parameter space corresponding to the reannihilation regime (II), but does not constrain freeze-in (regime I) so that, for the time being, the latter is  allowed for any DM mass. However, although the kinetic mixing parameter required to produce the observed DM abundance through freeze-in is very small, even this regime could be tested by the future Xenon-1T experiment. For one year (four years) of exposure, we have found that candidates in the range $50~\\hbox{GeV}\\lesssim m_{DM} \\lesssim 140$~GeV ($50~\\hbox{GeV}\\lesssim m_{DM} \\lesssim 600$~GeV) could be tested. As for the DAMA and CoGeNT data, they can be accounted for by considering values of the parameters which lay in the reannihilation regime, but which are excluded by the current Xenon-100 data, see Fig.8.} The characteristic $1/E_r^2$ recoil energy spectrum allows to distinguish this model from more standard DM models which, involving a mediator beyond the $\\sim$~MeV scale, predict a cross section on nucleon that is independent of $E_r$.  Other important constraints for the kinetic mixing setup are cosmological, related to the long range interaction caused by the light mediator. The most relevant constraint rests on  the ellipticity of galaxies, which puts an upper bound on $\\alpha'$. There is also another stringent upper bound on $\\alpha'$, which is related to the galactic magnetic fields and which one must satisfy if one wants DM to be present at the Sun location, so as to be probed using direct detection experiments. Both constraints turn out to be fully compatible with the freeze-in regime, but not with the other regimes, except for a DM mass larger  than a few hundreds GeV. We have shown how these constraints get considerably relaxed if instead of considering a massless mediator we make it slightly massive (still lighter than the DM mass scale, so that the relic densities we have obtained are unaffected).     The model we have considered for the Higgs portal consists of a scalar charged under a $U(1)'$ gauge group, and depends only on $m_{DM}$, the $U(1)'$ gauge coupling and the Higgs portal interaction. Here the 4 phases are all allowed by experimental data, even for DM masses above the GeV scale. Direct detection together with collider experiments could potentially cover the connector freeze-out phase and, optimistically, a part of the hidden sector freeze-out regime. The phase diagram we have derived can be consequently considered as the analytic prolongation (towards smaller Higgs portal coupling values) of the Higgs portal freeze-out regime {widely} considered in the literature.  {\\em A priori} such an analytic prolongation could be obtained with the same characteristic mesa shape for any DM candidate which belongs to a SM gauge singlet hidden sector, with (a) a visible sector which consists of particles that are in thermal equilibrium when they produce the DM (as in the case we considered here, where in the visible sector there are only SM particles), and (b) a primordial energy density in the hidden sector which is secondary. In this sense our results are representative of a large class of new DM models."
    },
    "1112/1112.2944_arXiv.txt": {
        "abstract": "We present a comprehensive catalog of 700 confirmed star clusters in the field of M31 compiled from three major existing catalogs. We detect 418 and 257 star clusters in \\textit{Galaxy Evolution Explorer} ($GALEX$) near-ultraviolet (NUV) and far-ultraviolet (FUV) imaging, respectively. Our final catalog includes photometry of star clusters in up to 16 passbands ranging from FUV to NIR as well as ancillary information such as reddening, metallicity, and radial velocities. In particular, this is the most extensive and updated catalog of UV integrated photometry for M31 star clusters. Ages and masses of star clusters are derived by fitting the multi-band photometry with model spectral energy distribution (SED); UV photometry enables more accurate age estimation of young clusters. Our catalog includes 182 young clusters with ages less than 1~Gyr. Our estimated ages and masses of young clusters are in good agreement with previously determined values in the literature. The mean age and mass of young clusters are about 300~Myr and 10$^{4}~M_{\\sun}$, respectively. We found that the compiled [Fe/H] values of young clusters included in our catalog are systematically lower (by more than 1~dex) than those from recent high-quality spectroscopic data and our SED fitting result. We confirm that most of the young clusters kinematics show systematic rotation around the minor axis and association with the thin disk of M31. The young clusters distribution exhibits a distinct peak in the M31 disk around 10 - 12~kpc from the center and follow a spatial distributions similar to other tracers of disk structure such as OB stars, UV star-forming regions, and dust. Some young clusters also show concentration around the ring splitting regions found in the southern part of the M31 disk and most of them have systematically younger ($<$ 100~Myr) ages. Considering the kinematical properties and spatial distribution of young clusters, they might be associated with the well-known 10~kpc star-formation ring structure in the M31 disk. Consequently, we suggest that various properties of young clusters in M31 might be in line with the scenarios that a satellite galaxy had passed through the disk of M31 less than few hundred million years ago. ",
        "introduction": "%  In the galaxy formation scenario based on the cold dark matter models, galaxies build up hierarchically \\citep{whi78,whi91,spr05}. In this scenario, merging and accretion play key roles over cosmic time. From high redshift to the nearby universe, many massive galaxies show evidence of ongoing merging and/or accretion. In this context, large disk galaxies like the Milky Way (MW) and M31 are also thought to have assembled a significant fraction of their mass through interactions with other small galaxies \\citep[e.g.,][]{iba01}. Most of these interactions change the morphology and star formation history of the galaxy.  As a typical spiral galaxy in the nearby Universe \\citep{ham07}, M31 provides a unique and most important opportunity for testing this scenario on external spiral galaxies due to its proximity \\citep[785~kpc;][]{mcc05}. Many recent studies have suggested that M31 is a promising example, exhibiting a hint of a past merger \\citep{blo06,brow06,gor06,iba07,mor08,mcc09}. Most observational and theoretical results concern the halo and outer disk of M31. Especially, previous photometric and spectroscopic observations suggest that merging events have played an important role in the construction and evolution of the halo of M31 (\\citealt{mcc09} and references therein).  Recently, in addition to the well-known 10~kpc ring seen in previous observations (\\citealt{gor06} and references therein), the presence of a second, inner dust ring was discovered in the disk of M31 \\citep{blo06}. While a detailed study of the origin of the ring structure of M31 is needed, the two off-center circular rings suggest that the disk of M31 has been distorted by a very recent passage of its satellite galaxy through the disk \\citep[i.e., a head-on collision with the satellite galaxy about a few tens or hundreds Myr ago,][]{gor06,blo06}. In this case, such a recent event may have enhanced the efficiency of star formation in the disk of M31 \\citep[e.g.,][]{yin09}.  Star cluster systems can be a tracer of galaxy formation and assembly, in the sense that significant star cluster formation is typically produced by major star-forming episodes in a galaxy \\citep{lar00,bro06}. More than 400 globular clusters (GCs) are known in M31 \\citep[e.g.,][]{pea10}, which is about a factor of three more than in the MW. The GC system of M31 has two subpopulations, one is a metal-rich and spatially concentrated subpopulation and the other is metal-poor and spatially extended. The metal-rich GCs show ``bulge-like\" kinematics with rotation \\citep{per02,lee08}. However, unlike in the MW, the metal-poor GCs also show significant rotation \\citep{huc91,per02,lee08}. \\cite{mor04} showed a thin disk population of GCs which constitutes about 27~\\% of the \\cite{per02} sample. Subsequently, it has been shown that at least a subset of these objects are in fact young ($\\le$ 1~Gyr), metal-rich star clusters rather than old metal-poor GCs \\citep{bea04,bur04,fus05,puz05,rey07,cal09,per10}.  In contrast to the MW, a large population of young clusters with ages less than 1 - 2~Gyr is found in M31 \\citep{bur04,bea04,bea05,fus05,puz05,rey07,cal09,pea10}. \\cite{fus05} presented 67 young clusters from the Revised Bologna Catalog \\citep[RBC,][]{gal04} showing blue optical colors [$(B-V)_{0} < 0.45$] and/or high strength of H$\\beta$ spectral index (H$\\beta > 3.5$~\\AA). \\cite{rey07} confirmed these young clusters using \\textit{Galaxy Evolution Explorer} ($GALEX$) ultraviolet (UV) photometry and suggested that the existence of young clusters in the outskirts of the M31 disk is due to the occurrence of significant recent star formation in the thin-disk. More recently, two comprehensive catalogs of young clusters in M31 have been published from the spectroscopic survey of \\cite{cal09} and the Sloan Digital Sky Survey (SDSS) data of \\cite{pea10}. \\cite{cal09} estimated ages and masses of 140 young clusters and \\cite{pea10} defined 156 young clusters with blue colors of $g-r < 0.3$. Most of these clusters are more massive (between 10$^3$ and 10$^5$~$M_\\sun$) than the Galactic open clusters \\citep{cal09}. Furthermore, they have similar characteristics to the blue star clusters in the LMC \\citep{bur04,fus05} and other massive young clusters in Local Group galaxies \\citep{bar09}. However, no such predominant counterparts have yet been discovered in the disk of the MW, except for a handful of massive young clusters identified in the center of the MW \\citep[e.g.,][]{mes09}. The existence of massive young clusters in the outskirts of the M31 disk indicates the occurrence of significant recent star formation in the disk of M31 \\citep{fus05,rey07,cal09}. Assuming that merging and accretion event triggered higher-level star formation in the disk of M31 than in quiescent galactic disks, it is interesting to examine the properties of star clusters related to the M31 disk, elucidating the recent star formation history in M31.  Motivated by the opportunity to study formation and evolution of young clusters in M31, in this paper we construct a comprehensive multi-band catalog of star clusters in M31 compiled from RBC, \\cite{cal09,cal11}, and \\cite{pea10} samples. In particular, we included $GALEX$ UV data, since the UV flux is highly sensitive to young main-sequence stars included in the massive young clusters which radiate more UV flux than in optical passbands \\citep{rey07,kav07}. We select extensive subsamples of young clusters which is complementary with previous catalogs of young clusters in M31 \\citep[e.g.,][]{cal09}. Various properties (age, mass, metallicity, kinematics, and spatial distribution) of young clusters are compared with star-forming (SF) regions and OB type stars in M31, and with the 10 kpc ring structure. This allows us to test whether most young clusters may be the possible outgrowth of a recent accretion of satellite galaxy occurred at the center of the M31 disk.  In Section 2, we describe optical and near-infrared (NIR) data of the M31 star clusters compiled from previous catalogs. Combining additional $GALEX$ UV data and other auxiliary information, we present a final merged catalog of star clusters in M31. In Section 3, by comparing observed spectral energy distributions (SEDs) with simple stellar population models, we estimate ages and masses of star clusters and select young clusters. We present properties of young clusters in Section 4. We also discuss the relationship between young clusters and M31 disk structures in terms of possible recent star formation history in the M31 disk. The conclusions are summarized in Section 5. ",
        "conclusions": "%  We constructed a comprehensive star cluster catalog which contains 700 M31 star clusters compiled from RBC v4 \\citep{gal04}, \\cite{cal09,cal11}, and \\cite{pea10}. We detected 418 and 257 clusters in $GALEX$ NUV and FUV, respectively, above flux limits of 23.7 and 23.6 ABmags and measured their UV magnitudes. Our catalog includes photometry in up to 16 passbands ranging from FUV to NIR as well as ancillary information such as reddening, metallicity, and radial velocities. Our merged catalog is the most extensive and updated one of UV photometry for M31 star clusters, superseding our previous UV catalog \\citep{rey07}. We estimated ages and masses of star clusters by multi-band SED fitting; the UV photometry enables more accurate age estimation of young clusters.  We also extracted a sample of 182 young clusters with ages less than 1~Gyr consisting of 173 clusters with our age estimation and 9 clusters from \\cite{cal09}. Our estimated ages and masses of young clusters are in good agreement with previous literature \\citep[e.g.,][]{bea04,cal09,van09,per10}. We examined the properties of young clusters such as age and mass distribution, metallicity, kinematics, and spatial distribution which provide unique probes of the star formation history of the disk of M31, and the prominent 10~kpc ring structure in particular. The mean age and mass of the young clusters are about 300~Myr and 10$^{4}~M_{\\sun}$, respectively. The mass range of our young clusters in M31 is similar to that of massive clusters found in the LMC and is in between those of Galactic young clusters and old GCs \\citep{bea04,fus05,cal09}. Since most low-mass ($<$ 10$^{5}~M_{\\sun}$) clusters are young objects, we consider they may be disrupted within a few Gyrs \\citep{lam05}.  The [Fe/H] values of young clusters included in our catalog, which are mostly from \\cite{per02}, are systematically lower (by more than 1~dex) than those derived from high-quality spectroscopic data or inferred from our SED fitting. While high precision spectroscopic observations for an extensive sample of young clusters are anticipated, we suggest that most young clusters in M31 might have  moderately enhanced metallicity (i.e., [Fe/H] $>$ $-$1.0). Such a high metallicity of the M31 disk and associated young clusters is consistent with the general result that spiral's stellar disk is metal-rich \\citep{bel00}, when not experiencing substantial gas infall or outflow \\citep[e.g.,][]{dal07}.  By comparing radial velocities of star clusters with a cold disk model, we selected a subsystem of star clusters with thin-disk kinematics, associated with the disk of M31. We confirm that most of the young clusters show systematic rotation around the minor axis and are kinematically associated with the thin-disk of M31. The majority of young clusters is located between 5 and 18~kpc from the center of M31 and shows a distinct peak around 10 - 12~kpc. The distribution of young clusters is closely correlated with that of other tracers of disk structure (OB stars, UV SF regions, and dust).  Considering their kinematical properties and spatial distribution, young clusters are well correlated with the 10~kpc ring structure in M31. This structure, known as the ring of fire or star-formation ring, is off-centered from the galaxy nucleus (\\citealt{blo06} and references therein). \\cite{blo06} discovered an inner dust ring, offset from the center of M31, and suggested that the two rings originated from a recent passage of a satellite. By numerical simulations, \\cite{blo06} inferred that the star formation ring structure results from a head-on collision through the center of the disk of M31 by a companion satellite galaxy. This event can produce density wave rings which triggered massive star formation at the peak of the wave. They postulated a recent (about 210~Myr ago) interaction between M32 and the M31 disk. We confirm that the spatial distribution of young clusters in the M31 disk is well re-presented by a circular 10~kpc star formation ring. Although the age distribution of our sample is somewhat broad, the majority of young clusters is in the age range of 100 - 400~Myr, which appear to be consistent with the prediction of \\cite{blo06}. Therefore, we speculate that a large fraction of young clusters found in the M31 disk might have formed during the recent interaction between a satellite galaxy and the M31 disk, in which the star formation ring also originated. In particular, there is a split of the ring structure in the southern part of the M31 disk which corresponds to a gap in both IR contours and UV SF regions. Young clusters also show concentration outside the ring splitting and, furthermore, most of them have systematically younger ($<$ 100~Myr) ages. Some young clusters in this region might derive from another interaction with a satellite galaxy, related with the Southern Stream emerging from the southwest disk of M31 \\citep{lee08}.   Within the context of merger history of M31, it is not unreasonable that the star formation ring in M31 has been shaped by a recent collision of satellite accompanied by higher star formation rate than that of the MW \\citep{ren05,yin09}. Due to a recent major perturbation of the M31 disk, formation of significant young stellar populations and massive young clusters is expected. Low-level star formation (e.g., in quiescent galactic disks) tends to produce few, if any, massive young clusters. It is worth noting that M31 appears to be representative of the typical population of local spiral galaxies showing evidence of merging in the formation and evolution history \\citep[e.g.,][]{ham07}. On the other hand, the MW is a rather quiescent galaxy without any major interaction over the past few billion years. In this case, the MW disk may have evolved with a secular pattern \\citep[e.g., smooth gas accretion or infall;][]{cro06,ham07} without any violent merging event. This kind of quiescent environment of the Galactic disk can support the nonexistence of the populous massive young clusters found in M31. On the other hand, although head-on collisions between galaxies are rare \\citep[see][]{mad09}, M31 serves as an important local template, to understand more distant collisional ring galaxies (\\citealt{moi09} and references therein)."
    },
    "1112/1112.5049_arXiv.txt": {
        "abstract": "{We report 19 intermediate redshift clusters newly detected in the ROSAT All-Sky survey that are spectroscopically confirmed. They form a part of 911 objects in the \\reflex\\ II cluster catalogue with a limiting flux of 1.8$\\times10^{-12}$ erg/s/cm$^2$ in the 0.1-2.4 keV ROSAT band at redshift $z \\ge 0.2$. In addition we report three clusters from the \\reflex\\ III supplementary catalogue, which contains objects below the \\reflex\\ II flux limit but satisfies the redshift constraint above. These clusters are spectroscopically followed-up by our ESO NTT-EFOSC2 campaigns for the redshift measurement. We describe our observing and data reduction methods. We show how X-ray properties such as spectral hardness ratio and source extent can be used as important diagnostics in selecting galaxy cluster candidates. Physical properties of the clusters are subsequently calculated from the X-ray observations. This sample contains the high mass and intermediate-redshift galaxy clusters for astrophysical and cosmological applications. } ",
        "introduction": "Galaxy clusters are the largest well-defined objects in our Universe characterized by a dynamical equilibrium configuration (e.g. Navarro, Frenk and White 1995). They form an integral part of the cosmic large-scale structure and are ideal probes for the testing of cosmological models (e.g. Allen et al. 2011). However, they are also astrophysical superlatives in many respects causing the largest deflections as gravitational lenses (e.g. Schneider et al. 1992), major mergers of galaxy clusters during the large-scale formation constituting the most energetic events in the Universe apart from the big bang (Sarazin 2002), and imprinting a distinct frequency-dependent feature onto the cosmic microwave background by means of the Sunyaev-Zel'dovich effect (Sunyaev \\& Zel'dovich 1970). Therefore, they are extremely interesting astrophysical laboratories to study various important processes on very large scales.  With technological advances in the past few decades, the detection of an increasing number of galaxy clusters has enabled us to study their properties in detail, and to apply the statistics of the cluster population to a number of cosmological studies.  One of the most efficient methods for detecting galaxy clusters as truly gravitationally bound objects in a close to mass-selected way is the use of X-ray surveys (B\\\"ohringer 2008). In these surveys galaxy clusters are observed as extended X-ray sources owing to the thermal X-ray emission of the very hot intracluster medium (ICM).  The largest sample of X-ray luminous galaxy clusters to date is the one that has been identified in the ROSAT all-sky Survey (RASS), the only all-sky X-ray survey yet conducted with an imaging X-ray telescope (Tr\\\"umper 1993). The southern \\reflex\\ I (B\\\"ohringer et al. 2004) and northern NORAS I (B\\\"ohringer et al. 2000) cluster samples identified in the RASS representing a total of 996 X-ray detected galaxy clusters provide the largest X-ray catalogue compiled, where the \\reflex\\ survey has the most well-understood three-dimensional survey selection function.  With the completion of two recent spectroscopic observing campaigns, we have almost completed the \\reflex\\ II survey by adding 50-60 newly determined redshifts. The \\reflex\\ II survey effectively doubles the number of galaxy clusters from 447 to 911 objects.  The most massive galaxy clusters provide the tightest constraining power for cosmological model testing and are also the most interesting astrophysical laboratories. We therefore communicate here the detailed properties of 22 of the most massive and highest redshift clusters ($z \\ge 0.2$) among the sample of 50-60 X-ray detected clusters with newly determined redshifts from our recent observing campaigns. Our major aim in this paper is to communicate the data of newly discovered X-ray luminous clusters and most importantly to provide the redshift information for all the cluster galaxies with new redshift measurements, so that they can enter public data bases (e.g. NED, SIMBAD) and be used by the community.  In this paper, we provide a brief description of \\reflex\\ II and the present cluster sample in Section 2. Our optical observations and data reduction is outlined in Section 3 and in Section 4 the X-ray properties and their use as the diagnostics of the nature of the sources are described. In the latter section we also provide some comments on individual clusters. Section 5 then provides a summary and some conclusions about the application of these findings.  \\section[]{Description of the cluster sample}  The galaxy clusters of the present sample are the last galaxy cluster candidates of the \\reflex\\ II catalogue for which a final confirmation and redshift measurement was obtained in the last year. The \\reflex\\ II sample as a whole extends the flux limit of \\reflex\\ I from 3$\\times10^{-12}$ erg/s/cm$^2$ to 1.8$\\times10^{-12}$ erg/s/cm$^2$ in the 0.1-2.4 keV ROSAT band. This flux limit was imposed on a fiducial flux calculation assuming cluster parameters of 5 keV for the ICM temperature, a metallicity of 0.3 $Z_{\\odot}$, a redshift of $z$=0, and an interstellar hydrogen column-density according to the 21cm measurements of Dickey and Lockman (1990). This fiducial flux was calculated independent of (prior to) any redshift information and is therefore somewhat analogous to an extinction-corrected magnitude limit without K-correction in optical astronomy. We then calculated the true fluxes and luminosities by taking an estimated ICM temperature from the X-ray luminosity-temperature scaling relation (Pratt et al. 2009) and redshifted spectra into account.  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics[height=6cm]{reflex2_lxz.ps}} \\end{center} \\caption{Luminosity-redshift plot of the \\reflex\\ II cluster catalogue. The limiting flux is 1.8 $\\times10^{-12}$ erg/s/cm$^2$ in the 0.1-2.4 keV ROSAT band. The blue filled circles are the 19 clusters in the \\reflex\\ II catalogue, and the 3 open circles below the flux limit are the clusters in the \\reflex\\ III supplementary catalogue with $z>0.2$.} \\end{figure} \\vspace{0.2cm}  Fluxes were measured by adopting a growth curve analysis in an aperture maximising the signal. The flux values were then corrected to an aperture of $r_{500}$ \\footnote{Defined as the radius where the mean cluster mass density is 500 times the critical density of the Universe at the epoch at which the cluster is observed.} by estimating mass based on the luminosity-mass relation of Pratt et al.(2009) in an iterative way.  The total count rates, from which fluxes and luminosities were determined, were derived with the growth curve analysis (GCA) as described in B\\\"ohringer et al. (2000). The cluster candidates were compiled from a flux-limited sample of all sources in the RASS analysed by the GCA method, combining all information on the X-ray detection parameters, visual inspection of available digital sky survey images, properties in the NASA extragalactic database, \\footnote{The NASA/IPAC Extragalactic Database (NED) is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.} and other available images at optical or X-ray wavelength. In addition, we cross-correlated our data with publicly available SZ catalogues from large surveys such as {\\it Planck}, SPT, and ACT. For a detailed description of the construction of the \\reflex\\ II galaxy cluster catalogue, we refer to B\\\"ohringer et al. (in prep.).  The current cluster sample, comprising 19 objects, was selected from among the clusters of the \\reflex\\ II sample confirmed by means of follow-up in 2010-2011 at redshifts above $z \\ge 0.2$, and the remaining three objects were confirmed from the data taken in 2004. These clusters have luminosities in the range $L_{X}$=1.6--13$\\times 10^{44}$erg/s and estimated masses (based on Pratt et al. 2009) of $M_{500}$=2.3--7.8 $\\times 10^{14}$ $M_{\\odot}$, and cover the redshift range $z$=0.2--0.33. Fig. 1 shows the X-ray luminosity and redshift distribution of the \\reflex\\ II cluster sample and the newly confirmed 19 clusters and 3 supplementary clusters in comparison. The newly added cluster which has an exceptionally high luminosity is RXCJ1914.5-5928. Table 1 lists optical properties, and the essential X-ray properties of the galaxy clusters are found in Table 2. For the derivation of distance-dependent parameters, we use a flat cosmology with $\\Omega_m$=0.3, $h_{70}$=$H_{0}$/70 km/s/Mpc.  \\section[]{Optical spectroscopic data}  \\subsection[]{Observations}  A more accurate determination of physical parameters is possible for spectroscopically confirmed clusters than for those with photometric redshifts, which usually have larger errors. In addition, a well-observed redshift sample of cluster members probes the dynamical state of a cluster that can be studied with the galaxy velocity dispersion within the system.  With this in mind, \\reflex\\ and some NORAS clusters have been spectroscopically followed-up since 1992 with the EFOSC1/2 instruments. The 19 cluster data presented in this paper is based on two spectroscopic observing campaigns, 085.A-0730(A) and 086.A-0055(A), during five nights each in September 2010 and March 2011 with the EFOSC2 instrument at the New Technology Telescope (NTT) in La Silla. The additional three clusters are from observations in 2004 with the EFOSC2 instrument at the ESO 3.6m telescope in La Silla.  Where possible, the multi-object spectroscopy (MOS) mode of EFOSC2 was used, and in some cases if a visual inspection of the pre-imaging data identified a clear BCG, we used a single long-slit (LS) observation on the BCG with at least one other member of the cluster. The different observing modes are marked in Table 1, which also lists the redshifts from the spectroscopic observations.  Each field of our candidate clusters was optically imaged in Gunn r band around the X-ray centre for the target selection and the mask making. The imaging resolution is 0.12 x 0.12 arcsec$^2$, and the field of view is 4.1 x 4.1 arcmin$^2$ for both imaging and spectroscopic observations. When necessary, the field was rotated to help optimise target selection. We used the grism that covers the wavelength range between 4085 $\\AA$ and 7520 $\\AA$ with 1.68 $\\AA$ per pixel at resolution 13.65 $\\AA$ per arcsec. We used slitlets with a fixed width of 1.5 arcsec for the MOS and of 2.0 arcsec for the long-slit observations.  Including at least three bright objects, preferably stars, to orient the field, the slitlets were allocated to the candidate member galaxies. Typically ten slitlets per field were punched onto one mask. Owing to the instrumental set-up, a maximum of five MOS plates can be loaded per night. The typical exposure time for the clusters in this paper ranges from 1500 to 2000 seconds per exposure, whose variation also reflects the observing conditions. The seeing was on average around one arcsec for most of the observations.  \\subsection[]{Data reduction}  We reduced the data using the standard reduction pipeline of IRAF. The reduction scheme is briefly summarised below, which applies equally to MOS and single-slit data.  {\\bf Bias correction} The CCD frames were visually inspected and trimmed at all edges of the images. All bias frames of a similar mean in each run were combined to reduce the rms fluctuations in the master bias, which was then subtracted from the science images.  {\\bf Flatfielding} Unlike the bias that is an additive correction, the flatfielding corrects the multiplicative spatial variations. The correction includes the varying quantum efficiency manifested as the structural variations in the CCD, and vignetting. We combined all dome flats for the combination of our slit, grism, and filter as described in the previous subsection. The spectrum of the quartz lamp was then fitted with a cubic spline function, and normalised. All spectroscopic exposures were then divided by the resulting response frame.  {\\bf Extracting the spectra} After the bias and flatfield corrections, the two science spectra were combined and weighted using IMCOMBINE, which also removes the cosmic-ray events. We used APALL in IRAF to extract the spectra. The one-dimensional collapsed spectra were sky-background subtracted by using the sky region on both sides of an object spectrum as a background estimate. The calibration spectra, which are two He-Ar frames were also extracted from the same spatial position as the science spectra.  {\\bf Wavelength calibration} We identified at least 4-5 known lines in the He-Ar lamp spectrum to fit the wavelength as a function of pixel number with an rms of far smaller than 1 $\\AA$ where the corresponding rms of the long-slit observation is larger owing to the larger slit width. The procedure was performed with the IDENTIFY and REIDENTIFY packages.  {\\bf Astrometric correction} The optical image taken during the spectral observation does not contain sky coordinates. The first guess was taken from the pointing directions used for the telescope, and the image was then astrometrically corrected for based on the accurate location of a number of identified objects in a public optical catalogue.  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{ \\includegraphics[height=6cm]{RXCJ1253.8-2622.ps} } \\resizebox{\\hsize}{!}{ \\includegraphics[height=6cm]{mos.ps} } \\resizebox{\\hsize}{!}{ \\includegraphics[height=6cm,angle=-90]{ap6.v1.ps} } \\end{center} \\caption{The optical image (top), cleaned MOS frame (middle), and an example spectrum (bottom) of RXCJ1253.8-2622. The bottom panel shows the extracted spectrum with the identified absorption lines from the 6th object from the left of the middle panel.} \\end{figure} \\vspace{0.2cm}  \\subsection[]{Redshift determination}  The calibrated spectra undergo several corrections such as skyline removals, a heliocentric correction, and emission line removals before the redshift determination. The redshifts from the emission lines are determined separately after the correlation with the passive galaxy templates, hence we removed them for this step. We then use the RVSAO package, which applies the cross-correlation technique to the input templates of galaxy spectra to measure the object redshift. The \\reflex\\ I templates were used for this analysis, which include 17 galaxy and stellar templates (Guzzo et al. 2009). In essence, the method uses the fast fourier transform to construct the correlation function of the input spectrum together with the template, after filtering to remove the spurious components and binning noise. The significance of the correlation is reflected in the amplitude of the correlation peaks, and the shift in the correlation function corresponds to the velocity of the measured object. Among all correlation functions, we selected the one with the highest R-value as the most closely matching spectrum, and quoted accordingly its redshift measurement. We confirmed a spectroscopic cluster detection if at least three galaxies have their R value greater than 5, and lie within $\\pm$3000 km/s of the mean velocity of the cluster members. We then took the median of those galaxy redshifts as the cluster redshift.  For the long slit observations, a cluster was confirmed if the redshift of the BCG was known and at least another galaxy has a similar redshift within the aforementioned criteria.  A small fraction of our cluster galaxy sample exhibit emission lines, which are identified with EMSAO in the RVSAO package. These are marked by E in Tables A.1--3. The resulting individual galaxy redshifts for the clusters in the \\reflex\\ II sample are listed in the Tables A.1--3 in the appendix.  \\section[]{Properties of the cluster sample}  In Table 2, we list the X-ray properties of our cluster sample. The X-ray flux and luminosity in the 0.1-2.4 keV band are calculated iteratively from the count rate using an estimated temperature and the measured redshift and $N_{H}$ for the count rate to flux conversion and the luminosity K-correction to yield the rest-frame luminosity. We use the luminosity-temperature relation from Pratt et al.(2009), $T_{X}$ = 3.306~$L_{X}^{1/3} h_{70}^{2/3}$, for the temperature estimate, where $T_{X}$ is in units of keV and $L_{X}$ in units of $10^{44}$ erg/s. This flux and luminosity were measured for the largest aperture used by the GCA method. The size of this aperture is also given in Table 2 in units of arcmin. The uncertainties in the flux and luminosity measurements were calculated from the contributions of the Poisson statistics of the source counts and the variance in the background determined in a total area of 4009 arcmin$^2$. The growth curve result was carefully checked for blending of the target source with neighbouring sources and a manual deblending was performed when necessary. The additional flux error introduced by the deblending was estimated and added to the above statistical error. A complex case of deblending is discussed below and shown in Fig. 3.  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{ \\includegraphics[height=4cm]{RXCJ2149m18_sh.ps} } \\resizebox{\\hsize}{!}{ \\includegraphics[height=3.8cm]{RXCJ2149m18_contour.v2.ps} } \\end{center} \\caption{RXCJ2149.9-1859 : (upper panel) photon distribution in the hard band (left) and soft band (right). Most of the emission in the soft band is identified as the Seyfert galaxy occupying a region in the north-east of the cluster centre. (lower panel) Contour plots of the significance in flux in the hard (left) and soft band (right) at 1.5, 2, 3, 4, 5, 6, 7$\\sigma$ per Gaussian beam of 1 arcmin width. Contours are marked with two alternating types of lines for clarity. After deblending, the centre of the cluster moves to where the cross is. Note that the location of the Seyfert galaxy is marked by an asterisk. } \\end{figure} \\vspace{0.2cm}  We also corrected the aperture luminosity to the luminosity inside $r_{500}$ for which we used the $L_{X} - M_{Y}$ relation of Pratt et al.(2009), which translates to $r_{500}$=0.895~Mpc $L_{X}^{0.195} {E(z)}^{-1.056}$, with $L_X$ in units of $10^{44}$ erg/s for the 0.1 to 2.4 keV band. Further details of these calculations will be described in a forthcoming paper on the \\reflex\\ II sample construction by B\\\"ohringer et al.  \\subsection[]{Detection at other wavelengths}  Some of the clusters in our list are sufficiently prominent that seven were previously detected by Abell (1958, 1989) and one was noted by Zwicky, as indicated in Table. 1. One of the clusters, RXCJ1914.5-5928, has been detected by the {\\it Planck} survey (Planck Collaboration 2011a).  \\subsection[]{Additional diagnostic X-ray properties}  Given the low count rates (19-65 counts and two cases with 10 and 16 counts), the low angular resolution of the RASS (90 arcsec half power radius), and the low spectral resolution, the data do not permit a detailed spectral and morphological analysis to be made for these sources. However, it allows us to derive two simple and robust parameters.  One parameter, the spectral hardness ratio, is determined from the count rates in the soft (channel 11-40, $\\sim$ 0.1-0.4 keV) and hard (channel 52-201, $\\sim$ 0.5-2 keV) bands by the formula HR=(H-S)/(H+S). Since for a given temperature, redshift, and assumed metal abundance of 0.3 solar the expected X-ray spectrum of a cluster is known, the hardness ratio for a given interstellar absorption can be calculated. In Fig. 4, we show the distribution of the deviations of the detected from the estimated values of the hardness ratio in units of significance for the total \\reflex\\ II catalogue, the present cluster sample, and we compare these properties to those of a sample of non-cluster sources. \\footnote {These non-cluster sources comprise 1524 of the RASS X-ray sources in the \\reflex\\ sky region with fluxes above the \\reflex\\ II cut that were described as cluster candidates in the \\reflex\\ II sample construction (B\\\"ohringer et al. 2001).} Fewer than one third of the non-cluster sources overlap with the cluster distribution. Therefore, the inspection of the hardness ratios is a powerful diagnostic of the possible AGN (or other non-cluster X-ray sources) contamination of our sample. Indeed all clusters of the present sample deviate by less than 3$\\sigma$ from expectations, except for the cluster RXCJ2149.9-1859, where the contamination is obvious and has been deblended as described below.  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{ \\includegraphics[height=6cm]{hr.ps} } \\end{center} \\caption{Distribution of the measured hardness ratio deviations from the expectation in units of the significance for the \\reflex\\ II clusters is marked with the slanted lines in the histogram. The filled bars show the distribution of the cluster sample in this paper. The blank bars mark 1542 non-cluster sources identified during the construction of the \\reflex\\ II catalogue. It is clear that the hardness deviation of the clusters has an approximate Gaussian distribution with a mean close to zero, while the non-cluster sources have a significant excess towards the negative deviations.} \\end{figure} \\vspace{0.2cm}  The second diagnostic is used to discriminate extended sources from either point or unresolved sources. The most robust test we have applied in the \\reflex\\ project is a Kolmogorov-Smirnov test to check if the source photon distribution was consistent with a point source and the estimated background. The test was conducted in a radial region of 6 arcmin radius around the cluster centre (B\\\"ohringer et al. 2000). We assumed a probability threshold of $<1\\%$ for a reliable designation of a source as extended. We found a significant extent for 11/22 of the clusters in our sample, which is again a very good confirmation given that the clusters have a high redshift and are not always expected to display a significant extent. In comparison, the \\reflex\\ I sample at higher flux levels show a significant extent for 81 \\% of the clusters with unresolved sources mostly at higher redshifts (Fig. 26 in B\\\"ohringer et al. 2001). Fig. 5 shows the distribution of the extent parameter, -$log$(KS probability) for the \\reflex\\ II catalogue, the present sample, and the 1542 non-cluster sources. Only a few percent of the non-cluster sources have an extent parameter larger than 2. They are mostly nearby galaxies, very bright stars with an artificial extent signal, and some sources with spurious extents. Fig. 6 shows the hardness ratio - extent parameter distribution for all \\reflex\\ II clusters and the present sample. Most of the present clusters with an extent parameter smaller than 2, which is used as a conservative threshold, still have parameter values corresponding to a low probability of being a point source.  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{ \\includegraphics[height=6cm]{ks.ps} } \\end{center} \\caption{Fractional histograms of the -$log$ Kolmogorov-Smirnov probability for three categories of objects. The solid line represents the \\reflex\\ II clusters as in Fig. 4. We truncate the values so that the last bin contains the remaining counts for the whole sample, which amounts to 5 \\%. The dotted lines mark 1542 non-cluster sources, and the red dashed lines are for the 22 clusters. The non-cluster sources are strongly concentrated at small extents.} \\end{figure} \\vspace{0.2cm}  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{ \\includegraphics[height=6cm]{hr_ks.ps} } \\end{center} \\caption{Hardness ratio vs. -$log$ probability of the Kolmogorov-Smirnov test for the same clusters in the Figure 2. The larger blue circles indicate the 22 clusters presented in this paper. The lower left sector enclosed by dashed lines is the zone where there is a high chance that a source is not a cluster. We note that this is a conservative limit, and any source that falls into this category must be scrutinised in detail to be confirmed. One cluster indicated by a blue circle that has an exceptionally soft emission among all of the 22 (the large blue circle nearest to the vertical dashed line) is RXCJ2149.9-1859. See section 4.3 for detailed remarks about this source. Two clusters of small extent are clearly compact clusters where there is a central dominant galaxy close to the X-ray centres. } \\end{figure} \\vspace{0.2cm}  \\subsection[]{Remarks on individual clusters}  \\indent \\indent RXCJ0347.4-2129 has a complex emission region with two major components.  RXCJ0521.4-2754 is deblended from a nearby source.  RXCJ1253.8-26 has a complex, diffuse emission region. The structure of this source will be inspected in greater detail using an allocated 7ks Chandra observation.  The centre of RXCJ1914.5-5928 coincides with that of PLCKG337.1-26.0 to within a distance of 4.3 arcmin from the SZ cluster center (Planck Collaboration 2011a). In terms of the flux limit this cluster should have been included in the REFLEX I catalogue. However, the REFLEX I catalogue was strictly based on the detection of X-ray sources, as well as significant galaxy overdensities, in the COSMOS database, which is a relatively shallow with a non-uniform coverage across the sky (B\\\"ohringer et al. 2001). Hence, this cluster was identified as a supplementary cluster on the basis of its X-ray extent. With a more comprehensive identification procedure for the REFLEX II catalogue, it is identified as a part of the new sample.  The cluster, RXCJ2149.9-1859, contains an X-ray point source that is offset by about 2.4 arcmin from the X-ray centre of the cluster. The point source is significantly softer than the cluster emission and originates from a Seyfert galaxy at $z$=0.157. It is located at 21:49:58 in RA and -18:59:24 in Dec and is denoted by an asterisk in the lower panels of Fig. 3. The flux from this source was deblended from the cluster emission. The X-ray surface brightness distribution of this source in the soft (0.1-0.4 keV) and hard band (0.5-2 keV) is shown in Fig. 3. The contribution of the Seyfert galaxy to the cluster X-ray emission is clearly noticeable as a much softer source than the cluster emission north-east of the cluster centre. It was deblended by cutting off the contaminated sector of the cluster emission and filling it with the remaining azimuthal average. After deblending, the new centre is then found to be 1.5 arcmin south of its initial location.  RXCJ2149.0-3228 has been deblended from a source in the south-west of the cluster.  \\begin{table*} \\begin{center} \\caption{Optical properties of all clusters studied} \\label{table1} \\centering \\begin{tabular}{c c c c c c c c} \\hline \\hline \\multicolumn{1}{c}{Cluster} & \\multicolumn{1}{c}{Alt. Name} & \\multicolumn{1}{c}{R.A.} & \\multicolumn{1}{c}{Dec.} & \\multicolumn{1}{c}{$z$} & \\multicolumn{1}{c}{$\\Delta z$} & \\multicolumn{1}{c}{$N_{gals}$} & \\multicolumn{1}{c}{Observing mode}\\\\ (1) & (2) & (3) & (4) & (5) & (6) & (7) & (8)\\\\ \\hline RXCJ0052.4-5746 &      & 13.1093  & -57.7806 & 0.2559 & 2.37 & 3 & MOS \\\\ RXCJ0129.4-6432 &      & 22.3730  & -64.5409 & 0.3264 & 1.98 & 6 & MOS \\\\ RXCJ0139.5-2629 &      & 24.8888  & -26.4886 & 0.2445 & 1.56 & 4 B & LS \\\\ RXCJ0343.5-7152 &      & 55.8775  & -71.8764 & 0.2136 & 1.61 & 3 & MOS \\\\ RXCJ0347.4-2149 & A3168& 56.8724  & -21.8224 & 0.2399 & 1.77 & 5 & MOS \\\\ RXCJ0425.5-0045 & ZwCl0423.2-0053 &  66.3798 & -0.7577  & 0.2240 & 1.75 & 6 & MOS \\\\ RXCJ0521.4-2754 &      & 80.3581  & -27.9065 & 0.3175 & 2.79 & 2 B & LS \\\\ RXCJ1253.8-2622 & A1633& 193.4682 & -26.3828 & 0.2179 & 1.65 & 9 B & MOS \\\\ RXCJ1617.6-0714 &      & 244.4099 & -7.2446  & 0.2523 & 2.36 & 3 B & MOS \\\\ RXCJ1914.5-5928 & PLCKG337.1-26.0 & 288.6333 & -59.4767 & 0.2636 & 1.74 & 4 & MOS \\\\ RXCJ2006.4-1912 &      & 301.6046 &-19.2106  & 0.2769 & 1.83 & 4 & MOS \\\\ RXCJ2041.2-1324 &      & 310.3188 & -13.4149 & 0.3311 & 2.11 & 5 B & MOS \\\\ RXCJ2049.9-3029 & A3715& 312.4885 & -30.4926 & 0.2088 & 1.55 & 3 & MOS \\\\ RXCJ2126.1-3202 & AS0946& 321.5441& -32.0338 & 0.2797 & 1.66 & 2 B & MOS \\\\ RXCJ2149.0-3228 &      & 327.2524 & -32.4744 & 0.2326 & 1.94 & 4 & MOS \\\\ RXCJ2149.9-1859 &      & 327.4500 & -19.0050 & 0.2836 & 3.64 & 2 B & LS\\\\ RXCJ2308.3-0155 &      & 347.0873 & -1.9210  & 0.3047 & 2.07 & 4 & MOS \\\\ RXCJ2317.9-3534 & A3989& 349.4828 & -35.5794 & 0.2041 & 2.04 & 3 & MOS \\\\ RXCJ2351.0-1954 &      & 357.7703 & -19.9132 & 0.2477 & 1.90 & 2 & MOS \\\\ \\hline RXCJ1013.1-1641 & & 153.2794 & -16.6884 & 0.2240 & 2.00 & 3 & LS \\\\ RXCJ1023.2-0636 & A1001& 155.8268 & -6.6046 & 0.2544 & 1.06 & 3 & LS \\\\ RXCJ1355.6-2623 & A1816& 208.9334 & -26.3833 & 0.2035 & 1.43 & 5 & MOS \\\\ \\hline \\end{tabular} \\end{center} Notes: (1) R.A. and Dec. for J2000 in degrees. (2) The centre of RXCJ1914.5-5928 is 4.3 arcmin away from that of PLCKG337.1-26.0. There are three different redshifts associated with this {\\it Planck} cluster (Planck collaboration 2011a, 2011b). We take the redshift of the component with the largest contribution to the Fe line at 0.26 in the Table 2 of Planck Collaboration 2011b. (6) Errors in units of $10^{-4}$ (7) Number of galaxies with concurrent redshifts, B indicates when the BCG redshift is available. LS marks the redshift obtained from a single long-slit observation. The last three entries are based on the observation in March 2004 with the EFOSC2 instrument, which forms part of the \\reflex\\ III supplementary catalogue. \\end{table*} \\begin{table*} \\begin{center} \\caption{X-ray properties of the sample clusters} \\label{table2} \\centering \\begin{tabular}{c c c c c c c c c c c} \\hline \\hline \\multicolumn{1}{c}{Cluster} & \\multicolumn{1}{c}{Count Rate} & \\multicolumn{1}{c}{$F_X$} & \\multicolumn{1}{c}{$L_X$} & \\multicolumn{1}{c}{$L_{X,r_{500}}$} & \\multicolumn{1}{c}{Error} & \\multicolumn{1}{c}{$r_{500}$} & \\multicolumn{1}{c}{Ext} & \\multicolumn{1}{c}{$N_H$} & \\multicolumn{1}{c}{HR} & \\multicolumn{1}{c}{KS prob} \\\\ (1) & (2) & (3) & (4) & (5) & (6) & (7) & (8) & (9) & (10) & (11)\\\\ \\hline RXCJ0052.4-5746 & 0.090 & 1.86&  3.33&  3.26& 28.9 & 0.988&  6.5& 3.40 &0.96 &1.4\\\\ RXCJ0129.4-6432 & 0.090 & 1.80&  5.46&  5.14& 42.1 & 1.043&  8.0& 2.24 &0.50 &3.7\\\\ RXCJ0139.5-2629 & 0.069 & 1.96&  3.19&  3.02& 30.4 & 0.985&  8.5& 1.49 &0.24 &2.0\\\\ RXCJ0343.5-7152 & 0.108 & 2.34&  2.82&  2.66& 20.4 & 0.980&  9.5& 5.31 &0.76 &3.3\\\\ RXCJ0347.4-2149 & 0.153 & 3.11&  4.80&  4.41& 20.3 & 1.070& 13.0& 2.83 &0.59 &5.0\\\\ RXCJ0425.5-0045 & 0.087 & 1.94&  2.60&  2.53& 24.1 & 0.959&  7.0& 6.38 &0.32 &3.6\\\\ RXCJ0521.4-2754 & 0.100 & 1.98&  5.62&  5.45& 29.5 & 1.055&  6.5& 1.77 &0.66 &1.9\\\\ RXCJ1253.8-2622 & 0.100 & 2.25&  2.84&  2.73& 58.8 & 0.979&  8.0& 6.80 &1.00 &1.9\\\\ RXCJ1617.6-0714 & 0.073 & 1.89&  3.28&  3.32& 26.0 & 0.987&  5.5& 12.45&0.63 &2.2\\\\ RXCJ1914.5-5928 & 0.320 & 7.07& 13.16& 12.60& 20.0 & 1.285&  9.0& 5.82 &0.68 &4.2\\\\ RXCJ2006.4-1912 & 0.076 & 1.73&  3.68&  3.58& 52.6 & 0.995&  6.5& 7.22 &1.00 &0.6\\\\ RXCJ2041.2-1324 & 0.093 & 1.94&  6.05&  5.66& 18.3 & 1.061&  8.5& 3.65 &0.86 &1.4\\\\ RXCJ2049.9-3029 & 0.087 & 1.92&  2.23&  2.21& 38.0 & 0.939&  6.5& 6.16 &0.80 &2.1\\\\ RXCJ2126.1-3202 & 0.099 & 2.11&  4.56&  5.49& 25.0 & 1.036&  3.5& 4.54 &1.00 &0.1\\\\ RXCJ2149.0-3228 & 0.075 & 1.49&  2.19&  2.14& 36.0 & 0.922&  6.5& 2.09 &0.28 &0.5\\\\ RXCJ2149.9-1859 & 0.151 & 3.12&  6.94&  6.74& 22.0 & 1.121&  7.0& 3.35 &-0.04&2.1\\\\ RXCJ2308.3-0155 & 0.089 & 1.88&  4.91&  4.72& 22.5 & 1.035&  7.0& 4.29 &1.00 &5.2\\\\ RXCJ2317.9-3534 & 0.101 & 1.98&  2.19&  2.17& 59.4 & 0.938&  6.5& 1.69 &0.18 &0.6\\\\ RXCJ2351.0-1954 & 0.133 & 2.61&  4.33&  4.09& 19.5 & 1.045&  9.0& 1.50 &0.50 &3.2\\\\ \\hline RXCJ1013.1-1641 & 0.071 & 1.54&  2.07&  1.95& 18.3 & 0.917 & 9.0& 5.36 &0.90 &0.1\\\\ RXCJ1023.2-0636 & 0.067 & 1.43&  2.53&  2.84& 19.4 & 0.938 & 4.0& 4.65 &0.94 &1.5\\\\ RXCJ1355.6-2623 & 0.071 & 1.53&  1.67&  1.63& 43.7 & 0.890 & 7.0& 5.02 &0.69 &1.2\\\\ \\hline \\end{tabular} \\end{center} Notes: (2) in s$^{-1}$ (3) in units of $10^{-12}$ erg/s/cm$^2$ in [0.1-2.4] keV band measured in the aperture column (4) $L_{X}$ in $10^{44}$ erg/s in [0.1-2.4] keV band in the aperture (5) $L_{X}$ in $10^{44}$ erg/s in [0.1-2.4] keV band measured in $r_{500}$ (6) error in percentage (7) $r_{500}$ in Mpc (8) source extent in arcmin (9) $N_{H}$ in units of $10^{-20}$ cm$^{-2}$ (10) hardness ratio defined as HR=(H-S)/(H+S) where H is the count rate in the hard band (0.5-2 keV) and S in the soft band (0.1-0.4 keV) (11) -$log$ probability of the Kolmogorov-Smirnov test. \\end{table*}  \\section[]{Summary and conclusions}  We have presented a sample of 22 nearby spectroscopically confirmed clusters selected from the \\reflex\\ II catalogue. These clusters all have redshifts above $z$=0.2 and therefore for the given flux limit of the sample they are very X-ray luminous and have high estimated masses. This implies that these objects are interesting for follow-up studies, in particular for Sunyaev-Zel'dovich effect observations, gravitational lensing measurements, studies of the dark matter and baryon mass fraction, and cosmological studies. For this reason, we wish to make these findings public in advance of the publication of the main catalogue. For the two most prominent clusters of the sample, RXCJ1914.5-5928 and RXCJ2149.9-1859 we have obtained mass estimates, $M_{500}$, of 7.8$\\times 10^{14}M_{\\odot}$ and 5.3$\\times 10^{14}M_{\\odot}$ using the scaling relation of Pratt et al. (2009), $M_{500}$=2.02 $L_x^{0.585}$ $E(z)^{-1.17}$. The most luminous cluster, RXCJ1914.5-5928, was detected among the high significance sample in the {\\it Planck} survey. These two most massive clusters are the most interesting objects, e.g. for lensing studies and with estimated $Y_{X}$ parameters of 8.6$\\times 10^{14}M_{\\odot}$keV and 4.4$\\times 10^{14}M_{\\odot}$keV (using the $Y_{X}$ scaling from Arnaud et al. 2007, $Y_x$ = 0.1934 $M_{500}^{1.825}$ $E(z)^{0.73}$), they should be detectable by on-going Sunyaev-Zel'dovich experiments.  \\begin{acknowledgements} We thank the anonymous referee for helpful comments. \\end{acknowledgements} \\begin{thebibliography}{}  \\bibitem[]{} Abell, G. O., 1958, ApJS, 3, 211A  \\bibitem[]{} Abell, G. O., Corwin, H. G., Olowin, R. P., 1989, ApJS, 70, 1A  \\bibitem[]{} Allen, S., Evrard, A., Mantz, A., 2011, ARA\\&A, 49, 409 % \\bibitem[]{} Arnaud, M., Pointecouteau, E.,  Pratt, G. W., 2007, A\\&A, 473, L37 % \\bibitem[]{} B\\\"ohringer, H., Voges, W., Huchra, J. P. et al., 2000, ApJ, 129, 435 % \\bibitem[]{} B\\\"ohringer, H., Schuecker, P., Guzzo, L. et al., 2004, A\\&A, 425, 367 % \\bibitem[]{} B\\\"ohringer, H., in {\\it The Universe in X-rays}, J.E. Tr\\\"umper, \\& G. Hasinger (eds.), Springer, 2008, p. 395 - 434 % \\bibitem[]{} Colless, M., Peterson, B. A., Jackson, C. et al., 2003, astro-ph/0306581% \\bibitem[]{} Dickey, J., Lockman, F., 1990, ARA\\&A, 28, 215 %  \\bibitem[]{} Guzzo, L., Schuecker, P., B\\\"ohringer, H., et al., 2009, A\\&A, 499, 357  \\bibitem[]{} Jones, D. H., Read, M. A., Saunders, W. et al., 2009, MNRAS, 399, 683% \\bibitem[]{} Kurtz, M., Mink, G., 1998, PASP, 110, 934 % \\bibitem[]{} Navarro, J., Frenk, C., White, S., 1995, MNRAS, 275, 720 %  \\bibitem[]{} Planck Collaboration, 2011a, arXiv:1101.2024  \\bibitem[]{} Planck Collaboration, 2011b, arXiv:1101.2025  \\bibitem[]{} Pratt, G., Croston, J. H., Arnaud, M., B\\\"ohringer, H., 2009, A\\&A, 498, 361 % \\bibitem[]{} Sarazin, C., in {\\it Merging Processes in Galaxy Clusters}, Feretti, L. et al. (eds.), Kluwer Academic Publishers, 2002, p. 1 - 38 % \\bibitem[]{} Schneider, P., J\\\"urgen, E., Falco, E., {\\it Gravitational Lenses}, Springer, 1992 % \\bibitem[]{} Sunyaev, R. \\& Zeldovich, Y.B., 1970, Comments Astrophys. Space Phys., 2, 66  \\end{thebibliography}   \\appendix ",
        "conclusions": ""
    },
    "1112/1112.5425_arXiv.txt": {
        "abstract": "The long wavelength physics in a de Sitter region depends on the initial quantum state. While such long wavelength physics is under control for massive fields near the Hartle-Hawking vacuum state, such initial states make unnatural assumptions about initial data outside the region of causal contact of a local observer. We argue that a reasonable approximation to a maximum entropy state, one that makes minimal assumptions outside an observer's horizon volume, is one where a cutoff is placed on a surface bounded by timelike geodesics, just outside the horizon. For sufficiently early times, such a cutoff induces secular logarithmic divergences with the expansion of the region. For massive fields, these effects sum to finite corrections at sufficiently late times. The difference between the cutoff correlators and Hartle-Hawking correlators provides a measure of the theoretical uncertainty due to lack of knowledge of the initial state in causally disconnected regions. These differences are negligible for primordial inflation, but can become significant during epochs with very long-lived de Sitter regions, such as we may be entering now. ",
        "introduction": "At the classical level, de Sitter spacetime appears to be stable. For conformally-coupled matter rigorous nonlinear stability theorems have been established that show for an open set of initial data the solution evolves to an asymptotically de Sitter solution at late times \\cite{FRIEDRICH:1986fk,Friedrich:1986kx,Friedrich:1991nn}.  Quantum mechanically, the situation is less clear. In a quantum theory correlators do depend on data outside the past light cone of points. It is therefore important to study the effect of choice of initial state on predictions for cosmological correlators, that then in turn determine the spectrum of density perturbations at late times.  There has been much confusion in the literature on this point. Broadly speaking there are two main camps, those who advocate using a dS invariant quantum state, to compute dS invariant correlators; and those who impose more physically motivated initial conditions, breaking the de Sitter symmetries. In the first scenario, the slow roll parameters provide the only means of explicit breaking of the de Sitter isometries. It has been convincingly argued that at least in massive scalar theories, these correlators are infrared finite % \\footnote{For conformally coupled matter, this was studied long ago \\cite{Adler:1972qq,Adler:1973ty}, and for more general matter in \\cite{Drummond:1975yc,Drummond:1977dg,Drummond:1977uy}. These questions have been revisited more recently in \\cite{Marolf:2010nz,Marolf:2010zp,Hollands:2010pr,Marolf:2011sh,Hollands:2011we}.% }. It remains an open question whether the same is true when the quantum fluctuations of gravitons are included.  In the other camp are those who have typically had in mind direct applications to inflation. In that case, an initial state is chosen at some finite time, breaking the de Sitter invariance. Needless to say, the correlators break de Sitter invariance, and in massless scalar theories, exhibit infrared divergences. Infrared cutoffs lead to terms in the correlators that grow as $\\log a(t)$ to some power \\cite{Traschen:1986tn,Tsamis:1992sx,Brandenberger:1993zc,Tsamis:1994ca,Tsamis:1996qm,Tsamis:1996qq,Mukhanov:1996ak,Abramo:1997hu,Abramo:1998hi,Weinberg:2005vy,Weinberg:2006ac,Polyakov:2006bz,vanderMeulen:2007ah,Bartolo:2007ti,Polyakov:2007mm,Adshead:2009cb,Burgess:2009bs,Polyakov:2009nq,Krotov:2010ma} (here $a(t)$ is the scale factor in the Friedmann-Robertson-Walker metric). In this setting there have also been computations involving gravitons at one-loop, which also find the $\\log a(t)$ growth \\cite{Tsamis:1996qk}. This has led to the suggestion that quantum effects may lead to a relaxation of the cosmological constant \\cite{Tsamis:1992sx}. See \\cite{Seery:2010kh} for a review article summarizing many of these approaches. Some also \\cite{Kahya:2009sz,Kundu:2011sg} for some observations closely related to the approach discussed in the present paper.  In this paper we will argue the main results of these two camps are mutually compatible, and that the difference between the two approaches may be viewed as a theoretical uncertainty in the relevant correlator. Since it is likely we do not have causal access to the spacetime region that would allow us to precisely determine the initial state, the best we can do is perform some version of a Bayesian analysis and quantify the theoretical uncertainty in the initial state, and hence the derived correlators.  If we work within the general framework of inflation, the initial state with the minimal number of assumptions that gives rise to our present observable universe, is a state that started as an approximately homogeneous region an inverse $10^{14}$ GeV in size (see \\cite{Kolb:1990vq} for a more detailed discussion of the initial conditions for inflation). To explain the horizon problem, this region must inflate to a size of about 1 m, assuming inflation ends, and reheating produces a thermal gas with temperature around $10^{14}$ GeV (this gives the famous 60 e-foldings of inflation needed for viability). Subsequently the universe evolves according to the Standard Model of cosmology. Thus, at least classically, we can conclude that any state that respects these conditions will give rise to what we see today. According to the Bayesian approach, we should pick the most likely such state (or maximum entropy state). The most obvious such state would involve the homogeneous patch, but would be surrounded by a thermal gas at the Planck temperature. Of course such a state would involve large gravitational back-reaction, and it would be impossible to compute with it using known methods % \\footnote{The landscape of string theory provides a wealth of viable vacua typically separated by Planck scale potential walls. This fits well with the picture of eternal inflation, where a Planck scale curvature foam tends to dominate the volume on any given space like slice. It is worth emphasizing that approximate approaches to performing computations in these situations, such as stochastic inflation \\cite{Starobinsky:1986fx,Starobinsky:1994bd}, where super horizon fluctuations in a fixed de Sitter background are considered in order to compute an initial wavefunction for the universe, or Coleman-de Luccia bubble nucleation \\cite{Coleman:1980aw}, where tunneling to an open universe sets the initial conditions for inflation, both require extensive fine tuning. In \\cite{Starobinsky:1986fx,Starobinsky:1994bd}, because only the dynamics of superhorizon modes is considered, the high curvature spacetime foam is ignored, and replaced by a fixed de Sitter background, which is an inconsistent approximation. Likewise if we wish to apply the results of \\cite{Coleman:1980aw} to eternal inflation in the string landscape we must assume we are in a special state where the semiclassical instanton methods are applicable. The philosophy in the present paper is to sidestep these issues, by making minimal assumptions about the initial state of inflation, in order to better model this generic strongly curved foam outside the initial homogeneous region. % }. In the following we will adopt a compromise that leaves one with computable correlators, but removes most assumptions about the initial region outside the homogeneous region. Therefore we propose to use an initial state with a hard infrared cutoff just outside the horizon at the start of inflation when the homogeneous patch begins to expand.  After this time, the homogeneous patch expands. If the infrared cutoff was kept at a fixed proper length, the patch would soon exceed the size of the cutoff, and the resulting correlators would miss much of the physics of interest to us today. Moreover, to keep the cutoff at the same length, the walls of the box would have to contract faster than the speed of light, which signals an unphysical choice of regulator. Nevertheless, such a cutoff is needed for massless theories in pure dS spacetime to retain the dS invariance of the correlators.  The more physical choice we advocate in the present work is to view the infrared cutoff as a kind of bubble wall. The trajectory of the wall will depend on the details of the bubble, but for any physical choice, at late times, the wall will asymptote to a family of timelike geodesics. Therefore the simplest cutoff that captures these qualitative features is a comoving cutoff. Such a cutoff has been studied before in the literature in \\cite{Lyth:2007jh,Bartolo:2007ti}, where it is referred to as a minimal box. Thus we advocate using a comoving infrared cutoff placed just outside the horizon at the beginning of inflation. This explicitly breaks de Sitter invariance. As we will see, this can lead to important effects at late times in correlators of massless fields.  The dS invariant computations, on the other hand, represent the maximal number of assumptions about the initial state. Choosing the Euclidean vacuum across the entire inflating patch amounts to solving the horizon problem by hand by imposing gaussian fluctuations of fixed size across the entire initial slice. The chief advantage of doing this is that, at least for massive scalar theories (where no additional infrared cutoff is needed), the correlators are dS invariant, and can be much more simply treated to extract predictions for observations today. Moreover, as we shall see, the answers agree with the leading terms of the correlators with a comoving cutoff, provided inflation does not last too long. The difference between the two approaches should be interpreted as a measure of the theoretical uncertainty in the correlator due to our ignorance of the initial state at the start of inflation.  As we shall see, the $\\log a(t)$ terms in the examples we study are subleading versus the dS invariant results provided inflation does not last extraordinarily long. Thus predictions for primordial inflation are largely unchanged.  On the other hand, we may be entering a regime of dS dominance in our present epoch. Assuming the graviton is the only field that experiences these $\\log a(t)$ divergences, we can expect this quantum instability of de Sitter space to become relevant only after a very large number of e-foldings. We note this instability is predicted from a unitary model of quantum gravity based on embedding dS regions in asymptotically AdS spacetimes, dual to conformal field theories \\cite{Lowe:2007ek,Lowe:2010np}. ",
        "conclusions": "\\paragraph{Massive scalar }  We conclude from section \\ref{sub:full-prop} that in the distant future, the loop corrections are insensitive to the choice of infrared cutoff. The result for the comoving cutoff approaches that of the de Sitter invariant Euclidean vacuum results as $\\tau_{2}^{2\\epsilon}$ , i.e. as a negative power of the scale factor.   \\paragraph{Massless scalar, or a light scalar at early times}  Here the result of section \\ref{sub:Early} is applicable. This can be the relevant behavior during the entire course of slow roll inflation. The difference between the result of section \\ref{sub:Early} and section \\ref{sub:full-prop} provides a measure of the theoretical uncertainty in the predictions of the two-point function caused by uncertainties in the initial state at the start of inflation. This difference is proportion to a $\\lambda\\log^{2}\\left(\\tau_{2}/\\tau_{0}\\right)\\log\\left(\\Lambda_{IR}\\tau_{0}\\right)$ relative to the tree level contribution. The slow roll conditions for pure $\\lambda\\phi^{4}$ give the constraint that the initial expectation value for $\\phi$ must be larger than the Planck scale. The condition that the density fluctuations today be sufficiently small $(10^{-5})$ gives the constraint that $\\lambda\\approx10^{-15}$. So we see in this example the effect will be at least $10^{-11}$ smaller than the tree-level piece, even after 60 e-folds of inflation. We conclude therefore, that while these secular logarithms are present during primordial slow-roll inflation, matching the magnitude of our observed density perturbations constrains the parameters of the inflaton model such that the coefficient of the secular log is very small, unless inflation is very long-lived.   \\paragraph{Massless scalar at late times}  The Euclidean vacuum methods break down in this case (see for example \\cite{Hollands:2011we} for a recent discussion of this situation) due to the nonexistence of a de Sitter invariant vacuum for a massless scalar \\cite{Allen:1985ux}. However the comoving IR cutoff approach still provides a well-defined perturbative expansion, up until the point when $\\lambda\\log^{3}\\left(\\tau_{2}/\\tau_{0}\\right)\\sim1$ (see \\eqref{eq:early}). The prediction is that the theoretical uncertainty in correlators becomes large at sufficiently late times. This is the practical sense in which de Sitter spacetime suffers a quantum instability. In the slow roll example, this does not happen until about $10^{5}$ e-folds, so is really only of academic interest as far as primordial inflation goes.   \\paragraph{Graviton at late times}  However it should be noted our result for the massless scalar matches qualitatively with the computation of \\cite{Tsamis:1996qk}. There the full one-loop correction for graviton self-energy was computed with the same kind of comoving infrared cutoff that we have used in the present paper. Compared to the tree-level result the contribution is of order $\\left(\\frac{H}{M_{pl}}\\right)^{4}\\log^{2}\\left(\\tau_{2}/\\tau_{0}\\right)$. Again this is very small over the 60 e-folds needed for primordial inflation. However over the present epoch such a term may become important, if we are indeed evolving toward a period where dark energy dominates. To apply these ideas to our present epoch, we could imagine simply shifting $\\tau_{0}$ to the present time, and placing a comoving infrared cutoff just outside our present horizon. This suggests quantum uncertainties due to the indeterminancy of our initial state would only become large after $\\left(\\frac{M_{pl}}{H_{0}}\\right)^{2}\\approx10^{122}$ e-folds of late time expansion, where now $H_{0}\\approx10^{-33}$ eV is the Hubble parameter today. It is interesting that this provides a mechanism for the possible breakdown of semiclassical physics in a de Sitter region, as predicted by quantum gravity arguments in \\cite{Lowe:2010np}."
    },
    "1112/1112.1426_arXiv.txt": {
        "abstract": "\\noindent In {\\it Hubble Space Telescope} ({\\it HST}) Cycle 17, we imaged the well known globular star cluster 47 Tucanae for 121 orbits using the Wide Field Channel (WFC) of the Advanced Camera for Surveys (ACS) and both the UVIS and IR channels of the newly installed Wide Field Camera 3 (WFC3) instrument (GO-11677, PI -- H.\\ Richer).  This unique data set was obtained to address many scientific questions that demand a very deep, panchromatic, and panoramic view of the cluster's stellar populations.  In total, the program obtained over 0.75 Ms of imaging exposure time with the three {\\it HST} cameras, over a time span of 9 months in 2010.  The primary ACS field was imaged in the two broadband filters $F606W$ and $F814W$ filters, at 13 orientations, for all 121 orbits.  The parallel WFC3 imaging provides a panchromatic (0.4 -- 1.7 micron) and contiguous imaging swath over a 250 degree azimuthal range at impact radii of 6.5 -- 17.9 pc in 47~Tuc.  This imaging totals over 60 arcmin$^2$ in area and utilizes the $F390W$ and $F606W$ broadband filters on WFC3/UVIS and the $F110W$ and $F160W$ broadband filters on WFC3/IR. \\\\  \\noindent In this paper, we describe the observational design of the new survey and one of the methods used to analyze all of the imaging data.  This analysis combines over 700 full-frame images taken with the three {\\it HST} cameras into a handful of ultra-deep, well-sampled combined images in each of the six filters.  We discuss in detail the methods used to calculate accurate transformations that provide optimal alignment of the input images, the methods used to perform sky background offsets in the input stack and the flagging of deviant pixels, and the balance reached between the input pixel drop size onto an output super-sampled pixel grid.  Careful photometric, morphological, and astrometric measurements are performed on the stacks using iterative PSF-fitting techniques, and reveal unprecedented color-magnitude diagrams (CMDs) of the cluster extending to $>$30th magnitude in the optical, 29th magnitude in the UV, and 27th magnitude in the IR.  The data set provides a characterization of the complete stellar populations of 47~Tuc, extending from the faintest hydrogen burning dwarfs through the main-sequence and giant branches, down to very cool white dwarf remnants in the cluster.  The imaging also provides the deepest probe of the stellar populations of the background Small Magellanic Cloud (SMC) galaxy, resolving low mass main-sequence dwarfs with $M \\lesssim$ 0.2~$M_\\odot$. ",
        "introduction": "\\label{introduction}  Observations of the Milky Way's globular star clusters have represented one of the most important aspects of astronomy over the past century \\citep{shapley17,sandage53}.  Clusters are ideal testbeds for theories of stellar structure and evolution as well as dynamical processes, because their constituent stars share incredible similarities, being generally co-eval, co-spatial, and iso-metallic \\citep{harris96,vandenberg85,kalirai10a}.  Photometric measurements of the stars in any given cluster therefore provide a snapshot of how stellar and dynamical evolution operate at fixed conditions (e.g., age, metallicity, and cluster dynamical state), and this picture can be generalized by observing samples of clusters with different conditions.  As one general application, modern day observations of nearby globular star clusters provide stringent tests of stellar-evolution models (e.g., Dotter et~al.\\ 2008), which in turn yield a calibration of population-synthesis techniques.  These techniques are frequently used to interpret light from distant galaxies into fundamental properties such as star formation rates, chemical abundance trends, and mass-to-light ratios (e.g., Bruzual \\& Charlot 2003).  In addition to being excellent tools for stellar structure and evolution theory, globular clusters also play an important role in our understanding of the epoch of baryonic structure formation in the Universe.  Studies of star-forming regions in different environments frequently reveal sites of cluster formation (Whitmore \\& Schweizer~1995; Hillenbrand 1997; Brodie \\& Strader 2006), and therefore much observational and theoretical effort has been devoted to deriving accurate absolute ages of these systems. The most popular technique for this relies on the comparison of theoretical isochrones to the main-sequence turnoff morphology in the color-magnitude diagram (CMD), a feature that is dependent on nuclear reaction rates, uncertainties in chemical composition, the equation of state, and second-order effects such as diffusion, rotation, and turbulence \\citep{renzini88}.  Although the first measurements of the turnoff in a nearby globular cluster were made in the 1950s (Sandage 1953), a huge leap forward was enabled in the 1980s when CCDs came online and again recently by the high-resolution Advanced Camera for Surveys (ACS) on the {\\it Hubble Space Telescope} ({\\it HST}).  For example, \\cite{sarajedini07} executed a systematic and homogeneous survey of 65 Milky Way globular clusters and resolved the main-sequence turnoff of each of them with high photometric precision \\citep{anderson08b}. These data have led to a new understanding of the absolute ages and the age-metallicity relation of the Milky Way's globular-cluster population, and suggest that the most metal-poor clusters formed $\\sim$1~Gyr after the Big Bang \\citep{hansen07,marinfranch09,dotter10}.  Complementing the many ground- and space-based imaging programs that have studied both specific Milky Way globular clusters and the ensemble population, our team has been involved in a decade-long campaign to provide ultra-deep imaging of the complete stellar populations in the nearest clusters over multiple epochs.  Our first study targeted the nearest globular cluster Messier 4 in 2001 (Cycle 9; GO-8679) with 123 orbits of {\\it HST} observing time with the Wide Field Planetary Camera 2,\\footnote{\\cite{bedin09} provide a more recent study of this cluster with ACS.} and the second study targeted the second nearest cluster NGC~6397 in 2005 (Cycle 13; GO-10424) with 126 orbits of ACS imaging.  These two projects have led to a wealth of scientific investigations that include 1) establishing accurate and independent white dwarf cooling ages for the clusters (Hansen et~al.\\ 2004; 2007), 2) constraints on the color-magnitude relation and luminosity and mass function of stars down to the lowest-mass dwarfs capable of core hydrogen burning (King et~al.\\ 1998; Richer et~al.\\ 2004; 2008), 3) new tests on the internal dynamical state of star clusters (Anderson et~al.\\ 2008a; Davis et~al.\\ 2008a; Hurley et~al.\\ 2008), 4) the most accurately constrained external space motions around the Milky Way through unprecedented proper motions (Kalirai et~al.\\ 2004; 2007; Anderson et~al.\\ 2008a), and 5) a wealth of variability and binary-related studies (Davis et~al.\\ 2008b).  Building on the success of these two previous programs, we received 121 orbits of {\\it HST} observing time in Cycle 17 to establish a detailed imaging survey of one of the most luminous clusters in the sky, 47 Tucanae (GO-11677; PI -- H.\\ Richer).  47~Tuc is perhaps the best-studied globular cluster (e.g., Guhathakurta et~al.\\ 1992).  Its rich stellar content has been targeted by large campaigns for studies of extrasolar planets (Gilliland et~al.\\ 2000), x-ray sources (Heinke et~al.\\ 2005), detailed dynamical studies (Meylan \\& Mayor 1986; McLaughlin et~al.\\ 2006), and much more.  The cluster is also the prototype metal-rich population in the Galactic halo and provides stringent tests for theories of stellar evolution (Demarque \\& McClure 1977) and represents a metal-rich anchor for establishing star formation histories and abundances of resolved and unresolved stellar population in other galaxies (Brown et~al.\\ 2003; Bruzual \\& Charlot 2003).  Unlike our previous studies of M4 and NGC~6397, this new project utilizes three {\\it HST} cameras to image the cluster's stellar populations to unprecedented depth over both a wide wavelength range ($\\lambda$ = 0.4 -- 1.7 microns) and over multiple fields spanning $>$60 arcmin$^2$. This data set is aimed at achieving similar science goals to those listed above, yet over this expanded region of parameter space (e.g., the calibration of stellar-evolution models in the IR).  The observations also yield the deepest probe to date of the lower main sequence of the Small Magellanic Cloud (SMC) dwarf galaxy, which is located behind 47 Tuc in projection (see Figure~\\ref{fig:SMCand47Tuc}).  The primary field is imaged at 13 orientations with the ACS/WFC in two visible filters $F606W$ and $F814W$ for all 121 orbits, and the parallel fields are mapped out by the newly installed Wide Field Camera 3 instrument using both the UVIS ($F390W$ and $F606W$) and IR ($F110W$ and $F160W$) cameras.  The wavelength coverage of these filters is shown relative to the spectral energy distribution of three stars that are relevant to the present science program in Figure~\\ref{fig:spectral}; a hot white dwarf, a moderate temperature globular cluster turnoff star, and a cool M dwarf.   \\begin{figure*}[ht] \\begin{center} \\leavevmode \\includegraphics[width=13.0cm, angle=90]{SMC.eps} \\end{center} \\vspace{-1.2cm} \\caption{A wide-field ground-based image of the Small Magellanic Cloud (SMC) in the southern skies reveals two foreground Milky Way globular clusters, NGC~362 just below the SMC and 47~Tuc to the left of the galaxy.  Although the main body of the SMC is separated from 47 Tuc by more than two degrees, a diffuse stellar population persists to greater radii and represents a background source of stars in our study (as demonstrated later).  This image subtends 6.8 $\\times$ 4.5 degrees and was taken with a 300~mm lens in September 2007.  The image was made by combining multiple 10 minute exposures in five visible filters (including H$\\alpha$).  Image is courtesy of St\\'{e}phane Guisard and reproduced here with permission, {\\tt http://www.astrosurf.com/sguisard}. \\label{fig:SMCand47Tuc}} \\end{figure*}     \\begin{figure*}[ht] \\begin{center} \\leavevmode \\includegraphics[width=13cm, angle=270]{11677ACSimage3_2.eps} \\end{center} \\vspace{-0.7cm} \\caption{A visualization of the observational strategy which uses three HST cameras. The picture in the backdrop is a ground based Digital Sky Survey image of 47 Tuc, resampled to a scale of 0.1 arcsec per pixel.  The star-like pointing to the west of the cluster center in the center of the mosaic represents the deep ACS/WFC field, which was imaged in $F606W$ and $F814W$ for all 121 orbits of the program.  The pattern is caused by the requirement to roll the telescope by $\\sim$20~degree in half of the visits.  These multiple roll angles provide very efficient cleaning of image defects and artifacts, and also allow the secondary WFC3 instrument to sweep out an arc around the periphery of the cluster.  These parallel fields sample more than 60 arcmin$^2$ of the cluster's population from $R$ = 6.5 -- 17.9~pc with both the WFC3/UVIS camera in $F390W$ and $F606W$ and the WFC3/IR camera in the $F110W$ and $F160W$ filters.  The two panels on the right show zoomed regions of a small portion of the parallel WFC3 fields.  This mosaic was constructed by drizzling each pixel in the {\\it HST} program onto the ground based Digital Sky Survey image, where WFC3/IR $F160W$ was arbitrarily selected for the parallel fields.\\label{fig:montage}} \\end{figure*}    To support the many on-going scientific projects enabled by this data set, we describe here the detailed observational strategy and data-reduction methods that are used to combine over 700 individual full-frame images into a single well-sampled stack in each filter which is largely free of cosmetic defects.  The methods that are used to measure the photometry, astrometry, and morphology of all stellar sources on these combined images are also described.  The final CMDs exhibit well defined sequences extending down to 29th magnitude in the UV, 30th magnitude in the visible, and 27th magnitude in the IR, and the photometric and astrometric accuracy is assessed using an extensive set of artificial-star tests. The primary features of these CMDs are illustrated in relation to the science goals.  The programs used for image registration, combination, and photometry in this analysis are all freely available and therefore some of the specific techniques used here may be useful for other large {\\it HST} programs. ",
        "conclusions": "\\label{conclusions}  We have presented a detailed analysis of a complex and rich {\\it HST} data set from Cycle~17 GO program 11677.  Over the course of 9 months in 2010, this program spent 121 primary orbits imaging a field with the ACS instrument, and 121 parallel orbits imaging multiple fields with both the UVIS and IR cameras of the newly installed WFC3 instrument.  These observations were obtained to address science questions that require measuring the complete stellar populations of 47~Tuc, from the faintest hydrogen burning dwarfs, through the main-sequence and giant branches, to the coolest and faintest white dwarf remnants.  We describe the observational design of this program and analyze over 700 full frame {\\it HST} images using conventional techniques.  Our analysis used the MultiDrizzle image analysis package to produce co-added images that are well sampled and free of cosmetic defects.  These {\\it single} images in each filter were then subjected to iterative PSF-fitted photometry using the DAOPHOT and ALLSTAR programs to yield the final photometric, astrometric, and morphological catalogs.   The stack-based approach presented in this paper for finding and measuring the faintest stars in the field is fundamentally different from alternative approaches that analyze the raw pixel data in all the images simultaneously.  The other analysis approach is underway and will be presented in an upcoming paper (Anderson et~al.\\ 2012).  The final catalogs and processed images of all fields will be publically released in the near future, upon completion of the science program.  Through our experimentation, we found that three of the most critical steps that affect the quality of the output drizzled image was the accuracy of the input image registration, the calculation of the sky offsets between all input images, and the final balance achieved between the pixfrac and scale parameters.  Currently, the {\\it HST} pipeline only produces drizzled images for input exposures obtained within the same visit (typically those that use the same guide stars), without any empirical corrections for incidental image shifts, rotations, or scale changes.  For pipeline processing of crowded field images, a more careful consideration of the background variations in the sky level is also needed.  Accurately measuring these offsets leads to a smaller standard deviation in the input images and better sigma clipping of deviant pixels.  The pipeline reductions also only produce images at a pixfrac of unity and with no rescaling, leading to a gross loss of resolution.  Based on our experience, the quality of such images can be greatly improved by using the full capabilities of MultiDrizzle in post-processing, and so several of the key steps are explained in detail in Sections~\\ref{MDACS} and \\ref{MDWFC3} (see also ``The MultiDrizzle Handbook'' -- Fruchter \\& Sosey 2009).  Some of this discussion is especially valuable to WFC3 users, as this program likely represents the most comprehensive test of the image-analysis software on data from the new instrument (e.g., the 59 orbit IR Stare field).  Related to photometric methods, first, we found that simple (small) aperture photometry on the final drizzled images yields an excellent CMD that is not grossly different from that achieved with careful PSF-fitting with DAOPHOT.  Our method does improve such first order photometry by building a PSF from isolated, bright stars on the frame, where the input catalogs were screened several times to remove objects with anomalous measurements.  This method has the advantage over techniques where a temporally-varying PSF is derived for each of the individual images in that only a single average PSF is measured on the final stack. A key disadvantage in some implementations of working with drizzled images is that the degree of spatial variation mapped across the field may not be appropriate for {\\it HST}, for example, if not enough stars are present in the field.  In the drizzled images of this program, much of the fine-scale variation of the PSF is averaged out through the combination of many images with individual PSFs at different orientations and detector locations.  The final photometric catalogs contain over 70,000 stars in the ACS primary field and over 50,000 stars in the WFC3 Swath fields. The CMDs from the single drizzled images exhibit tightly defined sequences of the 47~Tuc main-sequence, background SMC main-sequence, and the 47~Tuc white dwarf cooling sequence.  The photometry in the primary ACS field extends to $>$30th magnitude in $F606W$ and reveals both very low-mass hydrogen-burning stars in the cluster and very cool white dwarf remnants.  The parallel WFC3/UVIS data, with exposure times of just 1 orbit per filter per field, exhibit a similar CMD showing the three populations well separated.  The sensitivity of the WFC3/UVIS camera is excellent in these relatively crowded regimes, reaching depths of $F390W$ = 28.5 and $F606W$ = 29.0 in a single orbit of integration.\\footnote{The $F606W$ photometry shown in Figure~\\ref{fig:WFC3CMD} includes only those objects that were also detected in the shallower $F390W$ images.}  The WFC3/IR observations in this program represent the deepest IR probe of a resolved stellar population to date.  Very low mass 47 Tuc stars are measured over 60 arcmin$^2$ with just 3 orbits of integration per field (1 orbit in $F110W$ and 2 orbits in $F160W$). The stellar main-sequence is clearly seen to $F110W$ $\\gtrsim$ 23, and the photometry extends several magnitudes below this.  The morphology of the CMD shows a distinct ``kink'' on the lower main sequence, below which the sequence becomes {\\it bluer} in $F110W - F160W$ with lower luminosity.  The kink is caused by an opacity effect in the atmospheres of cool stars that develop molecular hydrogen at $T_{\\rm eff}$ = 4,300~K.  This feature will provide added leverage to a range of studies where CMDs are modeled to constrain stellar evolution and structure theory, as well as cool atmosphere models.  The WFC3/IR CMD of 47~Tuc also provides encouragement for studies of resolved stellar populations with the James Webb Space Telescope (JWST).  The NIRCam (and NIRISS) instruments on JWST will be sensitive to optical and IR wavelengths extending from 0.6 -- 5.0 microns. Through both a short and long wavelength channel, NIRCam will also be diffraction limited at both 2 and 4 microns.  Relative to HST WFC3/IR, the instrument will have smaller pixels by a factor of 4, a larger field of view by more than a factor of two, and much higher overall total system throughput.  For example, in $\\sim$10,000~seconds of exposure time, the complete stellar populations in dense star clusters down to the hydrogen burning limit will be easily measured out to systems beyond 50~kpc in the Milky Way (e.g., $J$ = 30th magnitude)."
    },
    "1112/1112.6030_arXiv.txt": {
        "abstract": "Quasar distribution is our main source of information about the large-scale distribution of matter at redshifts greater than 1. It can also shed light on the open questions of quasar formation. The present paper analyses the quasar clustering using the two-point correlation function (2pCF) and the largest existing sample of photometrically selected quasars: the SDSS NBCKDE catalogue (from the SDSS DR6). A new technique of random catalogue generation was developed for this purpose, that allows to take into account the original homogeneity of the survey without knowledge of its imaging mask. When averaged over photometrical redshifts $0.8<z_{phot}<2.2$ the 2pCF of photometrically selected quasars is found to be approximated well with the power law $w(\\theta)=\\left(\\theta/\\theta_{0}\\right)^{-\\alpha}$ with $\\theta_{0}=4''.5\\pm1''.4$ and $\\alpha=0.94\\pm0.06$ over the range $1'<\\theta<40'$. It agrees well with previous results by \\citet{myers_2006,myers_2007}, obtained for samples of NBCKDE quasars with similar mean redshift, but averaged over broader $z_{phot}$ range. The parameters of the deprojected 2pCF averaged over the same redshift range and modelled with a power law $\\xi(r)=\\left(r/r_{0}\\right)^{-\\gamma}$, are $r_{0}=7.81^{+1.18}_{-1.16}\\,h^{-1}$\\,Mpc, $\\gamma=1.94\\pm0.06$, which are in perfect agreement with previous results from spectroscopic surveys. We confirm the evidence for an increase of the clustering amplitude with redshift, and find no evidence for luminosity dependence of the quasar clustering. The latter is consistent with the models of the quasar formation, in which bright and faint quasars are assumed to be similar sources, hosted by dark matter halos of similar masses, but observed at different stages of their evolution. Comparison of our results with studies of the X-ray selected AGNs with similar redshift shows that the clustering amplitude of optically selected quasars is similar to that of X-ray selected quasars, but lower than that of samples of all X-ray selected AGNs. As the samples of all X-ray selected AGNs contain AGNs of both types, our result serves as an evidence for different types of AGNs to reside in different environments. ",
        "introduction": "\\label{sec:1}  \\indent\\indent Quasars, the brightest extragalactic objects, play an important role in study of the large-scale structure of the Universe. One of important characteristics of matter spatial inhomogeneity, that may be compared to cosmological theories of structure formation, is the two-point correlation function  (2pCF; e.\\,g. \\citealt{peebles_book}) of quasars, $\\xi(r)$. The main results on quasars' clustering at $z<2.2$ (see, e.\\,g., \\citealt{croom_2005,myers_2006,myers_2007,ross_2009} and references therein) were obtained using the 2-degree Field (2dF) Quasar Survey\\footnote{http://www.2dfquasar.org} (2QZ; \\citealt{Croom_2004}) and the Sloan Digital Sky Survey (SDSS)\\footnote{http://www.sdss.org}, the second stage of which has been completed with the 7th release \\citep{Abazadjian_2009}. The first results on quasar clustering for more complete sample of quasars with $z>2.2$ \\citep{white+2012} appeared recently by virtue of the 9th data release of the SDSS \\citep{dr9} which comprises the Baryon Oscillation Spectroscopic Survey (BOSS).  The study of the matter distribution with the help of extragalactic redshift surveys faces two main problems. Firstly, galaxy and quasar surveys give us information only about the distribution of luminous matter which is biased relative to dark matter \\citep{kaiser_1984}. This bias is expected to be non-linear, scale-dependent and stochastic \\citep{mo+1996,sheth+1999}, but the simplest, linear, scale-independent bias parameter, which relates the 2pCFs of the dark+luminous matter, $\\xi_{m}(r)$, and that of the luminous matter, e.\\,g. quasars,  $\\xi(r)$, as \\begin{equation}\\label{eq:bias} \\xi(r)=b^{2}\\xi_{m}(r), \\end{equation} has proved to be an extremely useful first order approximation \\citep{kaiser_1984,bardeen+1986,cole+1989}. This parameter depends on the nature of extragalactic objects, i.e., on morphological type \\citep{einasto_2007,ross_2007}, star formation rate \\citep{owners_2007}, color-index \\citep{coil_2008}, or the luminosity \\citep{beisbart_2000,zehavi_2005,sorrentino_2006,coil_2008}. The latter dependence is explained within the framework of the hierarchical model of the structure formation, in which more massive, and therefore more luminous galaxies, have to reside in more massive dark matter halos, and thus to be more clustered. It is unclear however whether this relation can be applied to quasar clustering, because of obviously more complicated process of the quasar formation, which in its turn still does not have a concordance model. Several models of quasar formation were presented in the literature (e.\\,g. \\citealt{kauffmann+2000,hopkins+2005,hopkins+2008,croton+2006,thacker+2008}) and based on study of the luminosity dependence of the active galactic nuclei (AGN) clustering, some of them were completely ruled out, while some models are partially supported by observations so far. Hence study of the luminosity dependence of the quasar clustering in different wavelength bands can help to clarify not only the relation between the luminous and dark matter clustering, but also the physics of the quasar formation. It was also shown by several authors (see e.\\,g. \\citealt{croom_2005,weinstein_2004,porciani_2004,myers_2006,myers_2007,daAngela_2008,Mountrichas_2009}) that the linear bias parameter of quasars evolves with redshift.  Secondly, analysis of the distribution of extragalactic objects is related to the observed redshifts of these objects. These redshifts are the only tool for measuring distance between the objects on cosmological scales, but they are `contaminated' by measurement errors and non-Hubble motions. It results in the so called redshift-space distortions, namely \\citet{kaiser_1987} and `Finger of God' effects (see, e.\\,g., \\citealt{daAngela_2008,Mountrichas_2009,ross_2009,ivashchenko_2010} and references therein).  To avoid the problem of the redshift-space distortion the projected, spatial $w(\\sigma)$ and angular $w(\\theta)$, 2pCF are usually used. In the first case the projections $\\sigma$ of the 3-dimensional distances $r$ on the plane perpendicular to the line of sight are used, and the parameters of the 3-dimensional 2pCF $\\xi(r)$ can be reconstructed from the parameters of $w(\\sigma)$. Here the assumption, that the influence of the redshift inexactness causes a negligible effect on $\\sigma$ compared to $r$, is usually made. In the second case only the angular distances on the sky plane are used, and the parameters of $\\xi(r)$ can be reconstructed from $w(\\theta)$ using Limber's equation \\citep{limber_1953}. This process of deprojection of the 2pCF is more complicated than in the first case because it comprises the usage of the redshift distribution function of objects and analytical functions describing the clustering evolution. Nevertheless, taking into account smaller sizes and larger volumes of quasar spectroscopic surveys compared to  surveys of galaxies, the angular 2pCF is still appealing for quasars because it does not include redshift information and thus allows the use of catalogues of photometrically classified quasars, which probe to deeper flux limits, and so typically contain an order of magnitude larger number of objects than spectroscopic quasar catalogues.  Study of the angular clustering of photometrically selected quasars was pioneered by \\citet{myers_2006,myers_2007,myers_2007_2} using the catalogues of NBCKDE photometrically selected quasar candidates from SDSS \\citep{richards_2004,richards_2009} based on the Early and the 4th data releases of SDSS. It was shown that the angular 2pCF of quasars is fitted well with a power law  $w(\\theta)=\\left(\\theta/\\theta_{0}\\right)^{-\\alpha}$, where $\\alpha$ is a slope of the angular 2pCF. Myers et al. noted that the angular 2pCF does not evolve strongly with redshift, suggesting that the bias of quasars relative to underlying dark matter climbs steeply out to redshift 2.2. The reconstructed parameters of the 3D clustering by these authors agree well with the `direct' measurements of 3D clustering of quasars conducted e.\\,g. by \\citet{porciani_2004}, \\citet{croom_2005}, \\citet{daAngela_2008}. It is interesting to compare the  clustering of quasars with those of X-ray point-like sources, a large majority of which are supposed to be AGN. The study of the X-ray AGN clustering has been conducted intensively since the release of the X-ray sky surveys, such as ROSAT All-Sky Survey, XMM-Newton and Chandra, and especially after the cross-correlation of X-ray surveys with optical redshift surveys of AGNs (for some latest results see e.\\,g. \\citealt{Allevato+2011,elyiv_2011}).  In this paper we present our study of the angular and spatial clustering of photometrically selected quasar candidates from the SDSS NBCKDE catalogue \\citep{richards_2009}, based on the 6th release of the SDSS. This catalogue is 3 times larger than that used in \\citealt{myers_2007,myers_2007_2}, which allowed us to place more tight restrictions on the samples, namely to reject objects with photometric redshift determination probability less than 0.5, and to construct the main sample only from objects with redshifts within the `SDSS window', unlike \\citet{myers_2007}, who used the whole redshift range of the catalogue. This or slightly wider `window' is traditionally used in studies of spectroscopic samples because of the better efficiency of the photometric selection algorithm of quasars, and because of the peak of the quasar distribution within this redshift range. Although we follow \\citet{myers_2006} in deprojection of the angular 2pCF and calculation of the bias parameter, our algorithm of the random catalogue generation differs from the traditional technique, used by \\citet{myers_2006}, which is an important issue in verification of the 2pCF results.  The samples, procedure their selection and properties, along with a random catalogue generation method are described in Sec.\\,\\ref{sec:2}. The technique of the angular 2pCF calculation, its results and discussion are presented in Sec.\\,\\ref{sec:3}. Reconstruction of the spatial clustering parameters along with the corresponding technique are described in Sec.\\,\\ref{sec:4}. Sec.\\,\\ref{sec:5} contains the results on the luminosity and redshift dependence of the quasar clustering along with comparison with other studies of optical quasar samples and samples of X-ray selected AGNs. Finally, in Sec.~\\ref{sec:6} we sum up the results.  Throughout the paper we calculate distances within the spatially flat $\\Lambda$CDM model with $\\Omega_{M}=0.28$ and $H_{0}=71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$.  \\vspace*{-3ex} ",
        "conclusions": "\\label{sec:6}  \\indent\\indent Using the SDSS NBCKDE catalogue of photometrically selected quasar candidates \\citep{richards_2009} drawn from the 6th release of the SDSS we studied the angular and spatial clustering of quasars with the help of the 2pCF. For this purpose own technique of the random catalogue generation was investigated and used. The catalogue is 3 times larger than that used in previous studies by \\citet{myers_2006,myers_2007,myers_2007_2}, which allowed us to place more tight restrictions on the samples, namely to reject objects with photometric redshift determination probability less than 0.5, and to construct the main sample only from objects with redshifts within the `SDSS window', $0.8<z<2.2$ \\citep{weinstein_2004}. Hence, even having the main sample with 230\\,829 objects, which is not larger than the samples used by Myers at al., our results should be viewed as a corroboration of these previous works, with the main utilities being our use of an independent technique to generate random catalogues, and also use of a sample with ``higher quality''. The main results of this work are the following: \\begin{enumerate} \\item The angular 2pCF of photometrically selected quasars, averaged over $0.8<z_{phot}<2.2$, is approximated well with the power law $w(\\theta)=\\left(\\theta/\\theta_{0}\\right)^{-\\alpha}$ with $\\theta_{0}=4''.5\\pm1''.4$ and $\\alpha=0.94\\pm0.06$ over the range $1'<\\theta<40'$, that agree well with previous results by \\citet{myers_2006,myers_2007} for similar mean redshift, but averaged over broader $z_{phot}$ range. \\item The parameters of the deprojected 2pCF averaged over $0.8<z_{phot}<2.2$ and modelled with a power law $\\xi(r)=\\left(r/r_{0}\\right)^{-\\gamma}$, are $r_{0}=7.81^{+1.18}_{-1.16}\\,h^{-1}$\\,Mpc, $\\gamma=1.94\\pm0.06$, which are in perfect agreement with previous results from spectroscopic surveys. \\item Splitting the sample into three redshift ranges, we confirmed the evidence for an increase of the clustering amplitude with redshift, found in previous studies of photometrically selected and spectroscopically confirmed quasar samples.  \\item Selecting three subsamples with different $M_{i}$ and $0.8<z_{phot}<2.2$, we have found no evidence for luminosity dependence of the quasar clustering. The marginal decrease in correlation length for the sample of \\textit{faint} quasars can be addressed to the redshift evolution, which is hard to avoid when analysing different luminosity subsamples within a broad redshift range. This result is consistent with the models of the quasar formation, in which bright and faint quasars are assumed to be similar sources, hosted by dark matter halos of similar masses, but observed at different stages of their evolution. \\item Comparison of our results with studies of the X-ray selected AGNs with similar redshift showed that the clustering amplitude of optically selected quasars is similar to that of X-ray selected quasars, but it is lower than that of samples of all X-ray selected AGNs. As far as X-ray selected AGN samples contain both AGN types, this result evidences for different types of AGNs to reside in different environments. \\item We also investigated our sample and found that in addition to the well-known stellar contamination of photometrically selected quasar candidates there is also a small (about 0.1\\%) contamination by artifacts of the automatic selection technique of point sources. These include star formation regions in spiral galaxies and diffraction spikes around bright stars. These artifacts were already extensively discussed in \\citet{myers_2007_2}. The subsamples with high and low Galactic extinction (reddening) showed similar results on 2pCF, suggesting that Galactic reddening does not strongly influence quasar clustering when our technique for random-catalogue-generation is adopted. \\end{enumerate}  \\vspace*{-2ex}"
    },
    "1112/1112.3279_arXiv.txt": {
        "abstract": "The Galactic center stellar disk and the circumnuclear ring provide a unique opportunity  to study in detail the dynamics and physical conditions of distant molecular disks in the nuclei of galaxies. One of the key questions is how these disks form  so close to their host black holes and under what condition they form stars in a tidally stressed environment. We argue that disk formation around a massive black hole is due to partial accretion of extended molecular clouds that temporarily pass through the central region of the Galaxy. The cancellation of angular momentum of the gravitationally focused gas naturally creates a compact gaseous disk. The disk  can potentially become gravitationally  unstable and form stars. We apply these ideas to explain the origin of sub-parsec  megamaser disks found in the nuclei of Seyfert 2 galaxies. We show that  an empirical  scaling relation  between the mass of the black hole and the size of the disk  can be understood in the context of the cloud capture scenario. We conclude that the stellar and gas disks found in our Galactic center act  as a bridge to further our understanding of more distant mega-maser disks in the nuclei of Seyfert 2 galaxies. ",
        "introduction": "Understanding the formation of massive young stars in the immediate vicinity of massive black holes is challenging. Two plausible models for  the origin of the disk of massive stars orbiting Sgr A*, the 4$\\times10^6$ M$_{\\odot}$ black hole at the Galactic center, have been proposed, namely, in-situ star formation and the dynamical migration of star clusters (see Genzel, Eisenhauer and Gillessen 2010; Levin and Beloborodov 2003; Gerhard 2001; Nayakshin et al. 2007). It has become clear that the in-situ star formation model can account for present observations better than the migration scenario in which a young star cluster undergoes dynamical friction and spirals  toward the Galactic center. The question then arises as to what extent disk formation near Sgr A* can provide insight in the formation of sub-parsec gaseous disks orbiting supermassive black holes in AGNs. To address this question, we will first focus on the origin of Galactic center stellar and molecular disks near Sgr A* followed by AGN disks in the context of partial cloud capture by the massive black hole. We then discuss observational signatures such as the scaling relation between the mass of the black hole and the size of the disk as well as the recoil of the black hole in the context of the cloud capture  scenario.  A more detailed account of the results presented here can be found in Wardle and Yusef-Zadeh (2008, Paper I), Yusef-Zadeh et al. (2009, Paper II) and Wardle and Yusef-Zadeh (2011, Paper III). ",
        "conclusions": "In summary, we have investigated the capture  of  gas clouds by massive black holes via the Hoyle-Lyttleton  process. We applied this mechanism to three different evolutionary phases of cloud capture scenario. One is a subset of AGN megamaser disks with high inclinations.  There is no apparent star formation in these disks which lie  within the sphere of influence of their host black holes. Second  application of the cloud capture scenario is to the  2-parsec molecular disk orbiting Sgr A* where the gaseous disk  is young. It has been argued that  this system is showing  early  signatures of star formation. Lastly, sub-parsec stellar disk orbiting Sgr A* where massive star formation has already taken  place a few million years ago.   The lack of star formation in  megamaser disk  systems could be due to strong magnetic fields acting against self-gravity of the disk (Johansen and Levin 2008; Gaburov, Johansen and Levin 2011). Alternatively, the low column density may be responsible. Assuming that  the tidal shear by the black hole is accounted for, then the mass-to-flux ratio is an important parameter for the comparison of  self-gravity and the magnetic field.  Thus, the lack   of  star formation depends on the column density of incoming clouds and the magnetic field of the disks  rotating around  massive black holes.  Future  molecular line and magnetic field observations of thermal gas in these  disks  should be quite useful.  The physical relationship that threads through these different systems is the capture of a cloud by engulfing  the massive black hole. This is an effective way to get rid of the angular momentum of the incoming molecular cloud. The  engulfing scenario can also be applied when a star approaches the tidal disruption radius of a massive black hole. We speculate that these  processes  could operate in galaxy mergers, dense stellar cusps around massive black holes, leading to starburst activity, black hole growth  and powerful  flaring activity in the vicinity of  central massive black holes.   We found a  scaling relation between the radius of the captured disk and the massive black hole which  appears to be supported by the Hoyle-Lyttleton  capture process. The outer radius of the captured cloud is understood in terms of a physical truncation of the disk because the scale determined kinematics of cloud capture depends linearly on the black hole mass. The apparent  inner radius of the captured disk  can be understood if the disk is warped in such a way that the inner disk does not get sufficient exposure to the central source of X-rays, to  produce 22\\,GHz H$_2$O  (Maloney et al. 2002). As for the  disk of stars orbiting Sgr A*, the inner radius  is 0.03pc. In this case, it is possible that the stars that are formed within the inner 0.03pc are influenced by dynamical processes. In fact, there is a cluster of B-type stars with eccentric orbits distributed isotropically within the inner radius but with a completely different distribution than that of the disk of stars. Alternatively, it is possible that  the captured gas within the inner radius is   dissipated to high temperature and does not have sufficient density to overcome the tidal stress.  This idea could be relevant to the formation of the inner radius of the circumnuclear molecular ring where  hot  stars of  the disk are thought  to have ionized the inner 1.7pc of molecular gas and created  a cavity of molecular gas in this region."
    },
    "1112/1112.4490_arXiv.txt": {
        "abstract": "We use the AMUSE-Virgo and AMUSE-Field surveys for nuclear \\hbox{X-ray} emission in early-type galaxies to conduct a controlled comparison of low-level supermassive black hole activity within cluster and field spheroids. While both the Virgo and the Field samples feature highly sub-Eddington \\hbox{X-ray} luminosities ($L_{\\rm X}/L_{\\rm Edd}$ between $\\sim10^{-8}-10^{-4}$), we find that after accounting for the influence of host galaxy stellar mass, the field early-type galaxies tend toward marginally greater ($0.38\\pm0.14$ dex) nuclear \\hbox{X-ray} luminosities, at a given black hole mass, than their cluster counterparts. This trend is qualitatively consistent with the field black holes having access to a greater reservoir of fuel, plausibly in the form of cold gas located near the nucleus. We are able to rule out at high confidence the alternative of enhanced \\hbox{X-ray} activity within clusters. Presuming nuclear \\hbox{X-ray} emission correlates with the total energy and momentum output of these weakly accreting black holes, this indicates that low-level active galactic nucleus feedback is not generally stronger within typical cluster galaxies than in the field. These results confirm that for most cluster early-type galaxies (i.e., excluding brightest cluster galaxies) direct environmental effects, such as gas stripping, are more relevant in quenching star formation. ",
        "introduction": "The gas content and stellar populations of early-type galaxies are observed to depend upon their large-scale surroundings. Field early-type galaxies typically contain more cold (\\ion{H}{1}) gas, sometimes alongside young stellar populations (Oosterloo et al.~2010). Significant star formation in high density environments occurs primarily between $3<z<5$, whereas in low density environments it persists to $1<z<2$ or even lower for low-mass galaxies (Thomas et al.~2005; Treu et al.~2005b; Gobat et al.~2008). Proposed mechanisms of inhibiting star formation within clusters include gas removal (e.g., starvation through ram pressure stripping, tidal stripping, thermal evaporation, or other possibilities; Treu et al.~2003; Moran et al.~2005; and references therein) or morphological quenching (i.e., stabilization of a gas disk through the build-up of a stellar spheroid; Martig et al.~2009). Such processes would operate only at low efficiency within the field.  Another potential mechanism for terminating star formation is feedback from an active galactic nucleus (AGN). Radiation from an efficiently accreting supermassive black hole can launch winds or directly heat the surrounding gas (e.g., Ciotti \\& Ostriker 2007; Proga et al.~2010). The duration of such ``quasar mode'' highly-luminous AGN activity is only $\\simlt10^{8}$~yr (Yu \\& Tremaine 2002), outside of which black hole feeding is highly sub-Eddington. Numerical and observational work indicates that mechanical feedback may be of persistent importance even in weakly accreting systems, in the form of outflows (Blandford \\& Begelman 1999; Pellegrini et al.~2012) or jets (Falcke et al.~2004). The inclusion in simulations of ``radio mode'' feedback for AGNs at the center of rich groups or clusters successfully suppresses star formation and reproduces the red colors of large ellipticals (Croton et al.~2006; Merloni \\& Heinz 2007; Khalatyan et al.~2008). Brightest cluster galaxies (BCGs) often display direct evidence of mechanical feedback in the form of radio jets and inflated bubbles displacing the hot \\hbox{X-ray}-emitting intracluster medium (ICM) (e.g., McNamara \\& Nulsen 2007). At least in luminous early-type galaxies, the rate of accretion onto the supermassive black hole is well correlated with the emerging jet power (e.g., Allen et al.~2006; Balmaverde et al.~2008).  The frequency of AGN activity, within both early and late-type galaxies, may itself depend on local galaxy density, but interpretation is complicated by the sometimes non-overlapping nature of various observational activity indicators. The fractional rate of emission-line galaxies is higher in the field, to a degree exceeding differences in morphological distributions (Dressler et al.~1985). On the other hand, the fractional rate for AGNs selected by \\hbox{X-ray} luminosities is similar between field and cluster samples (Martini et al.~2007; Haggard et al.~2010). Position within the cluster may also play a role (Gavazzi et al.~2011), although Atlee et al.~(2011) find the intracluster radial distribution of \\hbox{X-ray} or IR-selected AGNs to be consistent with that of non-AGNs (while noting luminous \\hbox{X-ray} AGNs may be more centrally concentrated). Within relaxed clusters, Ruderman \\& Ebeling (2005) find an excess of \\hbox{X-ray} point sources both peaked within the central regions and more broadly distributed near the virial radius, which they attribute to black hole activity triggered by interaction with the BCG and by mergers, respectively.  We aim to assess the incidence and magnitude of low-level supermassive black hole activity within representative field versus cluster early-type galaxies, and to determine whether star formation in ordinary cluster early-type galaxies is likely primarily AGN-quenched. We characterize activity using nuclear \\hbox{X-ray} emission, which directly measures high-energy accretion-linked radiative output but more importantly serves as a plausible proxy for mechanical feedback (Allen et al.~2006; Balmaverde et al.~2008). Our samples are the AMUSE\\footnote{AMUSE: AGN Multiwavelength Survey of Early-Type Galaxies}-Virgo and AMUSE-Field surveys. Together, these target 203 optically-selected local early-type galaxies, with both surveys centered around Large {\\it Chandra\\/} Programs of ACIS-S3 snapshot (3--15~ks) observations (Virgo: ID 08900784, 454~ks, PI Treu; Field: ID 11620915, 479~ks, PI Gallo), supplemented with deeper archival {\\it Chandra\\/} coverage. The sample selection, data reduction and analysis, and nuclear \\hbox{X-ray} properties for the AMUSE-Virgo and AMUSE-Field surveys are presented in Gallo et al.~(2008, 2010; hereafter G08, G10) and Miller et al.~(2011; hereafter M11), respectively, from which we obtain values of $L_{\\rm X}$, $M_{\\rm BH}$, and $M_{\\rm star}$ (the nuclear 0.3--10~keV \\hbox{X-ray} luminosity,\\footnote{We assume no intrinsic absorption of the nucleus in these early-type galaxies. $L_{\\rm X}$ would be $\\sim$0.1 (0.5) dex greater for an intrinsic column of 10$^{21}$ (10$^{22}$)~cm$^{-2}$; the former value may be typical of late-type galaxies (Bogd{\\'a}n \\& Gilfanov~2011).} black hole mass, and galaxy stellar mass, which have units of erg~s$^{-1}$, $M_{\\odot}$, and $M_{\\odot}$, respectively). We emphasize that the AMUSE samples are unbiased with respect to nuclear \\hbox{X-ray} properties, and in fact almost all of the objects have $L_{\\rm X}<10^{41}$~erg~s$^{-1}$ and $L_{\\rm X}/L_{\\rm Edd}<10^{-5}$, reaching luminosities well below commonly utilized formal AGN classification limits. Throughout this work, errors are quoted as 1$\\sigma$.  \\begin{figure} \\includegraphics[scale=0.66]{virgorad.ps} \\figcaption{\\small Properties of the M87-associated Virgo galaxies versus distance from M87. The vertical dotted line is $r_{\\rm 200}$, and open symbols are \\hbox{X-ray} upper limits. {\\it Top\\/}: Relative color (dashed/dotted lines are mean for Virgo/Field samples). {\\it Middle \\/}: \\hbox{X-ray} luminosity and residual luminosity. {\\it Bottom \\/}: \\hbox{X-ray} coverage (open) and detections (solid). Crosses show typical uncertainties.} \\end{figure} ",
        "conclusions": "We briefly consider the implications of enhanced $L_{\\rm X}$ in field early-types. As both the Field and Virgo galaxies have highly sub-Eddington luminosities ($10^{-8}<L_{\\rm X}/L_{\\rm Edd}<10^{-4}$), their accretion structures should be physically similar (see, e.g., Soria et al.~2006 for discussion of inefficient flow models). One obvious parameter that could link nuclear activity to the large-scale environment is the fuel supply. Hot gas is subject to offsetting effects in clusters, with ram pressure stripping countered by accretion from the ICM and potential confinement of winds (Brown \\& Bregman 2000), but radio and optical observations establish cold gas and younger stellar populations as more prevalent in field early-type galaxies ($\\S$1). Major tidal interactions, which are less frequent within clusters due to the large galaxy velocity dispersions, could help bring gas to the nucleus, although this is not required as stochastic events may suffice for low-level fueling (as conjectured for Seyferts; Hopkins \\& Hernquist~2006). While accretion in local early-type galaxies does not appear to be limited by gas supply (Soria et al.~2006; Pellegrini 2010), a $\\sim$2$\\times$ larger infall rate (small in absolute terms; e.g., Allen et al.~2006) could plausibly produce the modest observed enhancements in $L_{\\rm X}$ assuming uniform efficiency and outflow fraction.  The Field sample contains galaxies in groups spanning a range in richness, as well as triples, pairs, or isolated galaxies. Groups are intermediate between field and clusters in terms of strength of galaxy-medium interactions, but facilitate strong galaxy-galaxy interactions because the bulk speeds of the galaxies are lower (somewhat similar to cluster outskirts). An apparent smooth decrease in scaled \\hbox{X-ray} luminosities from isolated to group to cluster environments, albeit with large scatter, was interpreted by M11 as tentative evidence of environmental modulation of supermassive black hole fueling. The direct comparison between field and cluster galaxies conducted in this work supports that possibility.  The robust result that these Virgo early-type galaxies are not systematically more \\hbox{X-ray} luminous than their Field counterparts indicates that low-level black hole activity is not generally stronger within typical cluster galaxies. This result is supported by the insensitivity to environment of the fractional rate of \\hbox{X-ray}-identified AGNs (Martini et al.~2007; Haggard et al.~2010). Presuming nuclear \\hbox{X-ray} emission correlates with the total energy and momentum output of these weakly accreting black holes ($\\S$1), this implies low-level AGN feedback is not generally stronger within typical cluster galaxies than in the field. This supports that the older stellar populations in cluster early-type galaxies, and the associated paucity of ongoing star formation, are not directly due to recent black hole activity.  Possibly past AGN outbursts could have already suppressed star formation and expelled gas from cluster galaxies, leaving no obvious current link. Both star formation (Moran et al.~2005) and AGN activity (Ruderman \\& Ebeling 2005) may be triggered at infall near the cluster virial radius, either through environmental or galaxy interactions. The AGN fraction increases with redshift (e.g., Martini et al.~2009; Haggard et al.~2010; Aird et al.~2011), and at $z\\sim3$ may be significantly higher in clusters (Lehmer et al.~2009). However, given that supermassive black hole activity is not generally greater within clusters in the local universe, this scenario requires AGN activity to be terminated or reduced more rapidly within clusters. While various environmental effects may remove gas otherwise available for AGN fueling in clusters, as described above, this depletion also acts to quench star formation and so it is unclear that past AGN activity is required for this task. Our findings are consistent with other arguments that AGN activity is not the proximate cause of star formation quenching, including the observed gradual decline in remnant star formation with decreasing cluster radius which led Moran et al.~(2005) to associate the cessation of star formation with a slow-acting effect such as starvation (see also von der Linden et al.~2010), and the similarity in the distribution of \\hbox{X-ray} luminosities for $9.5<\\log{(M_{\\rm star}/M_{\\odot})}<12$ which led Aird et al.~(2011)\\footnote{Pellegrini (2010) also mention the large scatter in \\hbox{X-ray} luminosities above $M_{\\rm star}\\simeq6\\times10^{9} M_{\\odot}$.} to state no evidence for AGN-quenching.  We note BCGs occupy a privileged position centered within the cluster potential and show direct evidence ($\\S$1) of mechanical AGN feedback. An additional complication is that BCGs in cool core clusters do show recent star formation related to the cooling gas (Hicks et al.~2010), again distinguishing them from typical cluster galaxies. Our cluster sample is exclusive to Virgo, so we do not investigate BCGs in general, but our conclusions about the importance of environmental effects versus AGN feedback likely do not apply to BCGs."
    },
    "1112/1112.1083_arXiv.txt": {
        "abstract": "In this note we point out that, contrary to the standard point of view, slow roll inflation is due to high gravitational friction. We show that the requirement of slow roll coincides with the requirement of a flat scalar field potential in the case of minimally coupled scalar field. In this sense, the search for a successful inflationary theory may be more fruitful by shifting the focus on models with high gravitational friction. We review then a gravitational mechanism, the so called ''Gravitationally Enhanced Friction\" mechanism, such that high gravitational friction is dynamically generated during inflation allowing even steep (i.e. non-flat) scalar potential to inflate. ",
        "introduction": "At large scales, the Universe looks extremely homogeneous isotropic and flat. This can be dynamically obtained by invoking an inflationary period in the early stages of the Universe evolution. The idea is very simple, it basically plays with the fact that any surface is locally flat homogeneous and isotropic:  Suppose the Universe, in the very past, was not homogeneous, not isotropic and spatially curved. Nevertheless, in any small patch, it would look like homogeneous isotropic and flat. Inflation, is literarily a fast (exponential) blowing of this small patch so to obtain, out of it, the today observable ``large'' Universe. Precisely, inflation is realized whenever there is an exponential run away of each point ($i,j$) of the small patch, i.e. \\be\\label{universe} {\\vec R}_{ij}(t)=e^{Ht}\\vec{R}_{ij}(t_0)\\ , \\ee where ${\\vec R}_{ij}(t)$ is the relative physical distance\\footnote{Note that in General Relativity the physical distance differs from the coordinate distance.}, $H$ is a constant, $t$ is the cosmic time and $t_0$ is some initial time. To get an intuition of why inflation is produced by gravitational friction, we will now extensively use a Newtonian-like description of the Universe expansion.  Let us fix the coordinates to be centered in one of the points. If only gravity is the force responsible for the Universe evolution we would have, from Newton's law\\footnote{We suppress from now on the $i,j$ indices and we consider units where the reduced Planck scale is $1$.} \\be \\ddot R=-\\frac{M(R)}{R^2}\\ , \\ee where $R=|\\vec{R}|$ and $M(R)$ is the mass (self gravitating energy) of the Universe. By using (\\ref{universe}) we easily see that \\be H^2=-\\frac{M(R)}{R^3}\\ , \\ee or \\be\\label{4} M(R)=-H^2 R^3\\sim - H^2 V_3\\ , \\ee where $V_3$ is the three-dimensional volume. Eq. (\\ref{4}) implies that an inflating Universe is sourced by a constant ``negative'' gravitating energy per unit volume. Although this looks odd in the context of Newtonian theory, it is not in the context of General Relativity (GR). Let us consider a perfect fluid of energy density $\\rho$ and an isotropic pressure $p$. In GR, any form of energy gravitates, in particular $p$ is a form of energy. Thus, if we discard negative energy density configurations, a negative gravitating energy can be obtained with a negative pressure, as we shall see.  The gravitational potential $U$ is an energy and therefore it is related to time translation of the system. The pressure instead modifies space translations. Therefore, a metric describing a homogeneous and isotropic universe driven by a perfect fluid may be written in ``Newtonian'' coordinates (see \\cite{nicolis}), as \\be\\label{newt} ds^2=-(1-2U)dt^2+(1-2\\Psi)dx\\cdot dx\\ . \\ee From the Einstein equation $G_{\\alpha\\beta}=T_{\\alpha\\beta}=(\\rho+p)u_\\alpha u_\\beta+p g_{\\alpha\\beta}$, where $u_{\\alpha}=\\delta_\\alpha^t$, we have the following (independent) $tt$ and $ii$ components \\begin{eqnarray} \\nabla^2\\Psi=\\frac{\\rho}{2}\\ ,\\cr 6\\ddot\\Psi+2\\nabla^2(U-\\Psi)=3p\\ . \\end{eqnarray} Combining the previous equations and considering a constant energy density, we finally obtain the General Relativistic Poisson equation \\be \\nabla^2 U=\\frac{1}{2}\\left(\\rho+3p\\right)\\ . \\ee For $\\rho>0$ and $\\rho+3p<0$ we can get inflaton if $p$ is constant and negative, as anticipated. In the Newtonian limit, in which the pressure is much smaller than the energy density ($\\rho\\gg p$), we obtain the Poisson equation for Newtonian gravity. ",
        "conclusions": "In this note, we showed that Inflation can only be obtained thanks to a high gravitational friction acting on a scalar field rolling down its own potential. We showed that in the minimally coupled case the high friction requirement is equivalent to the requirement of an almost flat potential. We then reviewed a mechanism that increases gravitationally the friction on the scalar field such to make the scalar field slow rolling even on steep potentials. This is achieved by a non-minimal coupling of the kinetic term of the Inflaton to curvatures in a way that does not introduce any higher derivatives and/or new degrees of freedom rather than a scalar and a spin-2 (the graviton).  \\begin{acknowledgement} I wish to thank Yuki Watanabe for important comments on the first draft of this note. I wish to thank the organizers of the BW2011 Workshop ``Particle Physics from TeV to Plank Scale'' for a stimulating conference. I am supported by the  by Alexander Von Humboldt Foundation.  \\end{acknowledgement}"
    },
    "1112/1112.2878_arXiv.txt": {
        "abstract": "We calculate the maximum mass of neutron stars for three different equations of state (EOS) based on generalized Skyrme functionals that are simultaneously fitted to essentially all the 2003 nuclear mass data (the rms deviation is 0.58 MeV in all three cases) and to one or other of three different equations of state of pure neutron matter, each determined by a different many-body calculation using realistic two- and three-body interactions but leading to significantly different degrees of stiffness at the high densities prevailing in neutron-star interiors. The observation of a neutron star with mass 1.97 $\\pm$ 0.04 $\\mathcal{M}_{\\odot}$ eliminates the softest of our models (BSk19), but does not discriminate between BSk20 and BSk21. However, nuclear-mass measurements that have been made since our models were constructed strongly favor BSk21, our stiffest functional. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.0566_arXiv.txt": {
        "abstract": "{ The interstellar region within the few central parsecs around the super-massive black hole, Sgr~A$^*$ at the very Galactic center is composed by a number of overlapping molecular structures which are subject to one of the most hostile physical environments in the Galaxy. } { Through the study of the morphology and kinematics of the emission from different molecular species as well as the variation of their line ratios we can get key insights on the distribution and interaction between different gas structures and the energetic phenomena taking place in the surroundings of Sgr A$^*$ as well as the physical processes responsible for the heating in this region. } { We present high resolution ($4''\\times3''\\sim0.16\\times0.11$\\,pc) interferometric observations of CN, $^{13}$CN, H$_2$CO, SiO, c-C$_3$H$_2$ and HC$_3$N emission at 1.3~mm towards the central $\\sim4$\\,pc of the Galactic center region. } { {Strong differences are observed in the distribution of the different molecules. The UV resistant species CN, the only species tracing all previously identified circumnuclear disk (CND) structures, is mostly concentrated in optically thick clumps in the rotating filaments around Sgr~A$^*$. H$_2$CO emission traces a shell-like structure that we interpret as the expansion of Sgr~A East against the 50~\\kms~ and 20~\\kms~ giant molecular clouds (GMCs). We derive isotopic ratios $^{12}$C/$^{13}$C$\\sim15-45$ across most of the CND region, but for the northeast arm, where the peak of H$_2$CO is observed and ratios $<10$ are found. The densest molecular material, traced by SiO and HC$_3$N, is located in the southern CND, likely due to shocked gas infalling from the 20~\\kms\\, GMC streamers, and the northeast arm, as a result of the expansion of Sgr\\,A~East % or a connecting point of the 50~\\kms streamer east to the CND. The observed $\\rm c-C_3H_2/HC_3N$ ratio observed in the region is more than an order of magnitude lower than in Galactic PDRs. Toward the central region only CN was detected in absorption. Apart from the known narrow line-of-sight absorptions, a $90$~\\kms\\, wide optically thick spectral feature is observed. We find evidences of an even wider ($>100$~\\kms) absorption feature. Around $70-75\\%$ of the gas mass, concentrated in just the $27\\%$ densest molecular clumps, is associated with rotating structures % and show evidences of association with each of the arcs of ionized gas in the mini-spiral structure.} } { These observations provide a coupled chemical and kinematical picture of the very central region around Sgr~A$^*$. Chemical differentiation has been proven to be a powerful tool to disentangle the many overlapping molecular components in this crowded and heavily obscured region. } ",
        "introduction": "At a distance of $\\sim8.0$\\,kpc \\citep{Reid1993,Groenewegen2008,Ghez2008} the few central parsecs around the super-massive black hole at the dynamic center of the Galaxy \\citep[$4.1\\pm0.6\\times10^6\\,M_\\odot$,][]{Ghez2008} is likely one of the most hostile environments for the interstellar medium (ISM) in the Milky Way. In this region the molecular gas undergoes strong disruptive tidal forces and pervasive X-ray irradiation \\citep[$L_X(20-120 \\rm keV)\\sim2.6\\times10^{35} erg\\,s^{-1}$;][]{B'elanger2006} in the vicinity of the black hole. Likewise, the ISM is subject to supernova remnants (SNR), HII regions, and strong photodissociating UV fields from the central cluster \\citep[$L\\sim1-3\\times10^7L_\\odot$ and ionizing photon flux of $2\\times10^{50}\\rm s^{-1}$;][]{Lacy1980,Davidson1992}.  The super-massive black hole is associated with the compact strong non thermal source Sgr~A$^*$. The mini-spiral of ionized gas consisting of three streamers known as Sgr~A West \\citep{Lo1983,Ekers1983} is likely feeding the central source. Surrounding Sgr~A West, the so called circumnuclear disk (CND) is actually a clumpy ringlike structure rotating counterclockwise around the dynamical center \\citep{Guesten1987,Wright2001,Christopher2005}. The molecular ring extends from its inner sharp edge at a projected distance of 1.6~pc from Sgr~A$^*$, out to a smoother edge at $\\sim2$~pc \\citep{Christopher2005}. However, the CND appears not to be a single symmetric structure forming a coherent ring \\citep{Marshall1995}. It has been suggested that the CND is a warped structure \\citep{Guesten1987} or, due to the incompleteness of the ring, formed by distinct rotating structures \\citep{Jackson1993,Wright2001}. Most of the observed molecular clumps forming the CND, with diameters of $\\sim7''$, are tidally stable and in virial equilibrium \\citep{Christopher2005}. Thus, the rotating molecular structures composing the CND appear to be non transient structures with an estimated lifetime of $\\sim 10^7$~yr \\citep{Christopher2005,Montero-Castano2009}.  \\begin{figure}[!t] \\centering \\includegraphics[width=0.46\\textwidth]{SgrAComposite.pdf} \\caption{Three color composite image on the CND around Sgr~A$^*$ including the emission of SiO (red), CN (green), H$_2$CO (blue). The central glow shows an artistic representation of the pervasive X-ray and UV-fields from the black hole and central star cluster. In the center, the black hole position is illustrated with a size corresponding to the map resolution. \\label{fig.Composite}} \\end{figure}  A number of molecular emission observations have been carried out to study the CND structure, its relation to the Sgr~A West streamers and the possible interaction with the two giant molecular clouds (GMC), namely the 20~\\kms\\, and 50~\\kms\\, GMCs \\citep[see][for references]{Ho1991,Marr1993,Montero-Castano2009}. So far, the only species with high resolution aperture synthesis observations are the HCN $1-0$ \\citep{Guesten1987,Wright2001,Christopher2005}, $3-2$ \\citep{Marshall1995}, and $4-3$ \\citep{Marshall1995,Montero-Castano2009}, the $1-0$ transitions of H$^{13}$CN and HCO$^+$ \\citep{Marr1993,Wright2001,Christopher2005}, the CS $7-6$ \\citep{Montero-Castano2009}, and NH$_3$ (1,1), (2,2), (3,3), (6,6) \\citep{McGary2001,Herrnstein2005}. The emission from the high density tracers like HCN, HCO$^+$ and CN nicely trace all the molecular features of the CND, though significant differences are observed in the line ratios. These differences are attributed to both abundance and density variations across the CND \\citep[e.g. HCO$^+$/HCN,][]{Christopher2005}. Although also tracing the CND, the emission of NH$_3$ significantly differs from the emission of the other tracers. The recent work at low resolution by \\citet{Amo-Baladr'on2011} studies, among other tracers, the emission of SiO and HNCO in the central 12 pc around Sgr~A$^*$. These species show strong abundance variations due to the influence of shocks and UV radiation in this region \\citep{Mart'in2008}. They probe the potential of HNCO/SiO abundance ratio as a measurement of the distance to the central cluster, which provides an insight on the 3-D arrangement of the different molecular components.  In this paper we report the high resolution ($4''\\times3''\\sim0.16\\times0.11$\\,pc) observations of CN, $^{13}$CN, H$_2$CO, SiO, c-C$_3$H$_2$ and HC$_3$N. As illustrated by the three color composite image shown in Figure~\\ref{fig.Composite}, each molecular species traces different molecular components as a result of the varying physical and chemical conditions in the vicinity of Sgr~A$^*$. ",
        "conclusions": "In this paper we present high resolution maps of the CN, $^{13}$CN, H$_2$CO, SiO, c-C$_3$H$_2$ and HC$_3$N emission towards the central $\\sim4$~pc around Sgr~A$^*$. These observations provide a chemical picture of the molecular material in the surroundings of the central super-massive black hole and how the chemically resolved structures relate to the different morphological and kinematical structures. Each of the observed species appears to be tracing different molecular gas structures which are intimately linked to the different heating processes affecting the molecular gas in this region.  Among them, CN is the only species showing emission in all previously described molecular morphological features in the CND, and it shows a similar distribution to that from available HCN and CS emission maps \\citep{Christopher2005,Montero-Castano2009}. Similar to these two molecules, CN is extremely resistant to the high UV irradiation from the central cluster. Moreover, its chemistry is linked to that of HCN. These species are shown to be widespread at the closest distances from Sgr~A$^*$. The emission from H$_2$CO differs drastically from CN and it is mostly located in the outer regions of the CND tracing a $\\sim 37''$ shell-like structure centered $\\sim16''$ away from Sgr~A$^*$ where H$_2$CO is observed to be even brighter than CN (with CN/H$_2$CO line ratios below 0.2). Mostly formed in solid phase, H$_2$CO emission in gas phase can be interpreted as the result of the expansion of Sgr~A East against the 50~\\kms~ and 20~\\kms~ GMCs east and south of the CND, respectively. Moreover, the $^{12}$C/$^{13}$C isotopic ratio, estimated from the $^{12}$CN/$^{13}$CN, supports the differentiation of the northeast arm, where the brightest H$_2$CO emission is observed. While we estimate an opacity corrected carbon isotopic ratio in the range of $15-45$ throughout the CND, consistent with the idea of $^{13}$C enrichment with decreasing galactocentric distance due to nuclear processing \\citep{Wilson1994,Wilson1999}, we find the lowest measured ratios ($<10$) towards the northeast arm. The opacity derived from the hyperfine fit to the CN spectral profiles shows that CN is optically thick throughout most of the CND. On the other hand, the densest ($n>10^6\\rm cm^{-3}$) regions in the CND are located in the southern CND and the northeast arm as traced by HC$_3$N and SiO, although their chemical formation paths are different. The southeast lobe and southern extension, where we find the largest SiO abundance and broader line widths, could be fueled and shocked by the molecular material from two streamers connecting the 20~\\kms\\, GMC, where SiO abundance would be enhanced through grain erosion by low velocity shocks. Similarly, the bright emission in the northeast lobe could be the consequence of the inflowing material from the 50~\\kms~ streamer connecting  the GMC east of the CND. The similar line widths measured in all species allow us to discard high velocity shocks as the origin of SiO. The dense material in the northeast arm could be compressed due to the expansion of Sgr\\,A East against the 50\\,\\kms\\, GMC, where the brightest HC$_3$N emission as well as significant SiO emission are detected. Even though HC$_3$N is efficiently formed through reactions involving CN, the emission of HC$_3$N is deficient close to Sgr~A$^*$ due to the high photodissociation rate of this molecule in UV exposed material. The few regions in the inner CND where HC$_3$N emission is observed, will be tracing the densest cores ($A_v>1-2$) for this molecule to survive. Finally, the carbon chain $\\rm c-C_3H_2$ is observed to follow the emission of HC$_3$N, but with $\\rm c-C_3H_2$/HC$_3$N at least an order of magnitude lower than in other Galactic PDRs. It is therefore unclear why the emission of $\\rm c-C_3H_2$ is so faint in the CND.  In the direction of Sgr~A$^*$ only CN is detected in absorption. It shows two previously identified narrow absorptions at $-157$ and $-23$~\\kms due to galactic arms along the line of sight. A broad $94$~\\kms\\, line at $-32$~\\kms\\, is likely tracing the molecular gas closer to Sgr~A$^*$ absorbing the non-thermal central emission. Moreover, we find hints of a possible wider ($>100$~\\kms) absorption which might be tracing the closest molecular material to the central black hole. However, observations of higher energy transitions of UV-resistant species are required to detect this molecular gas component in emission as shown by the HCN $4-3$ maps by \\citet{Montero-Castano2009}.  We have studied the kinematics of all detected molecular clumps in the CND, where we considered kinematics of each individual velocity component at every selected position. Our study shows most of these clumps are not associated with any rotating structure around Sgr~A$^*$ but with larger scale structures overlapping with the CND. However, even if only $29\\%$ of the molecular selected components appear clearly associated with rotating structures, they represent $70-75\\%$ of the mass of the molecular gas. Only three molecular structures show kinematic evidences of rotation around Sgr~A$^*$. These rotating filaments, namely the southwest lobe, southern extension, and northeast lobe, are rotating in different planes, similar to what is derived from the study of the ionized arcs \\citep{Zhao2009}, and are associated with the ionized mini-spiral structure observed in radio continuum and hydrogen recombination lines.  Moreover, we have shown how the various chemically differentiated structures are directly related to the different kinematical structures. In this comparison we observe how CN is mostly concentrated in the dense molecular optically thick clumps rotating (and likely infalling) filaments around Sgr~A$^*$ while the shell-like structure traced by H$_2$CO is not orbiting around the central black hole, but shows velocities related to the overlapping GMCs.  This paper has shown the potential of detailed high resolution chemical observations to disentangle the complicated molecular structure surrounding Sgr~A$^*$, as well as to establish the leading heating mechanisms affecting the different molecular structures in the region. Similar to the low resolution work by \\citet{Amo-Baladr'on2011}, chemical abundances from deep high resolution imaging of various molecular species as a function of position and distance from  Sgr~A$^*$ could potentially be used to establish the three dimensional distribution of the molecular components in the Galactic center region."
    },
    "1112/1112.5944_arXiv.txt": {
        "abstract": "We present here a comparison of our results from ground-based observations of asteroid (21) Lutetia with imaging data acquired during the flyby of the asteroid by the ESA Rosetta mission. This flyby provided a unique opportunity to evaluate and calibrate our method of determination of size, 3-D shape, and spin of an asteroid from ground-based observations. Knowledge of certain observable physical properties of small bodies (\\eg, size, spin, 3-D shape, and density) have far-reaching implications in furthering our understanding of these objects, such as composition, internal structure, and the effects of non-gravitational forces. We review the different observing techniques used to determine the above physical properties of asteroids and present our 3-D shape-modeling technique KOALA -- Knitted Occultation, Adaptive-optics, and Lightcurve Analysis -- which is based on multi-dataset inversion. We compare the results we obtained with KOALA, prior to the flyby, on asteroid (21) Lutetia with the high-spatial resolution images of the asteroid taken with the OSIRIS camera on-board the ESA Rosetta spacecraft, during its encounter with Lutetia on 2010 July 10. The spin axis determined with KOALA was found to be accurate to within two degrees, while the KOALA diameter determinations were within 2\\% of the Rosetta-derived values. The 3-D shape of the KOALA model is also confirmed by the spectacular visual agreement between both 3-D shape models (KOALA pre- and OSIRIS post-flyby). We found a typical deviation of only 2\\,km at local scales between the profiles from KOALA predictions and OSIRIS images, resulting in a volume uncertainty provided by KOALA \\add{better than 10\\%}. Radiometric techniques for the interpretation of thermal infrared data also benefit greatly from the KOALA shape model: the absolute size and geometric albedo can be derived with high accuracy, and thermal properties, for example the thermal inertia, can be determined unambiguously. The corresponding Lutetia analysis leads to a geometric albedo of 0.19\\,$\\pm$\\,0.01 and a thermal inertia below 40\\,J\\,m$^{-2}$\\,s$^{-0.5}$\\,K$^{-1}$, both in excellent agreement with the Rosetta findings. We consider this to be a validation of the KOALA method. Because space exploration will remain limited to only a few objects, KOALA stands as a powerful technique to study a much larger set of small bodies using Earth-based observations. ",
        "introduction": "} \\indent Perhaps the most striking observable of any asteroid is its shape. In 1993, spacecraft exploration revealed for the first time the stunning non-spherical shape of asteroid (951) Gaspra when NASA's Galileo spacecraft made the first of its two asteroid encounters, on its way to Jupiter. Asteroids had remained point-sources in the sky since the discovery of (1) Ceres in 1801 by Piazzi, almost two centuries before. Only the advent of space exploration and large Earth-based telescopes (\\eg, Arecibo and Goldstone radio telescopes, space-based-optical -- HST, or ground-based near-IR, equipped with adaptive optics -- \\add{Palomar,} Lick, CFHT, Keck, ESO VLT, Gemini) allowed their apparent disks to be spatially resolved, and their irregular shapes to be imaged. The past decade has seen a revolution in our understanding of the physical properties (\\eg, size, 3-D shape, spin axis) of asteroids. This revolution has come about thanks to improved observing facilities and, of equal importance, from improved methods of analysis. \\\\ \\indent Determination of the physical properties for a statistically relevant set of asteroids has many implications for our understanding of these remnants of solar-system formation and, in turn, can be expected to improve our understanding of the history and evolution of the Solar System. For instance, the distribution of spin axes of the larger asteroids (diameter larger than $\\sim$100 km) on the celestial sphere is not expected to be isotropic. Numerical hydrocode simulations have predicted a slight excess in prograde rotators, due to the gas-pebble interaction in the protoplanetary disk \\citep{2010-MNRAS-404-Johansen}. Similarly, the spin state of \\textsl{small} asteroids (diameter not larger than few kilometers) is dominated by the non-gravitational YORP effect \\citep{2011-AA-530-Hanus}. Statistical knowledge of spin coordinates, how they are distributed within and among asteroid families of different ages, will provide strong constraints on the effectiveness of YORP \\citep{2003-Icarus-162-Slivan}.\\\\ \\indent Reconstruction of the 3-D shape (including the size) is required to estimate the volume of an asteroid, which in turn is used to derive its density, possibly the property most fundamental to our understanding of an asteroid \\citep{2002-AsteroidsIII-4.2-Britt}. Observations of the surface of an asteroid, such as \\add{colors}, spectra, or phase effects, can provide clues to the surface composition of the asteroid.  This may or may not be related to the bulk composition of the body \\citep[\\eg,][]{2011-EPSL-305-Elkins-Tanton}. Masses for asteroids can be determined from a spacecraft flyby, from the orbital motion of a natural moon, or even from the perturbations of asteroids on other bodies, such as Mars \\citep{2002-AsteroidsIII-2.2-Hilton}. In most cases, however, the uncertainty in the density is dominated by the uncertainty in the volume, rather than the uncertainty in the mass \\citep{2002-AsteroidsIII-2.2-Merline}. Precise reconstruction of the 3-D shape is therefore of high importance for all asteroids for which a mass has been, or will be, estimated \\citep[\\eg,][]{2002-AsteroidsIII-2.2-Hilton, 2007-AA-472-Mouret, 2011-AJ-141-Baer}. \\\\ \\indent From the comparison of an asteroid's density with the densities of its most-likely constituents, we can constrain the macroporosity \\rem{(the presence of, and extent of, large-scale voids)} \\add{(large-scale voids)} in its interior, probably produced by impacts over its history \\citep{2002-AsteroidsIII-4.2-Britt}. These impacts could have partially disrupted the body, producing large-scale fractures, or even totally disrupted the body, with subsequent reaccumulation of the resulting fragments, leading to a ``rubble-pile'' structure.\\\\ \\indent Evidence of gigantic, but less than totally disruptive, impacts can be seen by high-resolution imaging and also inferred from our shape-modeling of asteroids. The huge impact craters evident in the images of C-type asteroid (253) Mathilde \\citep{1997-Science-278-Veverka} are thought to be about as large as could be sustainable by a body without disruption. There have been suggestions that these craters were created by compaction of low-density target material, rather than explosive ejection typical of hard-rock impacts \\citep{1999-Nature-402-Housen}. Already, our ground-based adaptive-optics imaging has shown what appear to be facets or depression-like features, similar to those seen on Mathilde, in some other large C-type asteroids \\citep[\\eg, (511) Davida in][]{2007-Icarus-191-Conrad}. Alternatively, some of our other images of C-type asteroids, such as (52) Europa, appear to bear no evidence of giant impacts \\citep[see][]{2011-Icarus--Merline}. Evaluation of the prevalence of such large impact events can give us insight into the size and frequency of these impact events over time, and thus into the history of the impacting population.\\\\ \\indent We summarize below the most common of the many observing techniques used to derive size, 3-D shape, and spin-vector coordinates and highlight some of their advantages and drawbacks. Then, in Section~\\ref{sec: koala}, we describe our KOALA multi-data shape-modeling algorithm. In Section~\\ref{sec: comp}, we present a comparison of the results produced by KOALA (from Earth-based observations) with those derived from the ESA Rosetta flyby of asteroid (21) Lutetia. In Section~\\ref{sec: thermal}, we use our KOALA model \\add{in conjunction with mid-infrared data and a thermophysical model} to derive the thermal properties of Lutetia and compare the results with those \\add{derived using thermal observations from the Rosetta spacecraft \\citep{2011-PSS--Gulkis} and from the ground, making use of the shape model from the flyby \\citep{2011-PSS--Rourke}.} We \\add{assess} the accuracy of the KOALA shape-modeling method in Section~\\ref{sec: conclu}.   \\subsection{Optical lightcurve\\label{ssec: lc}} \\indent Historically, spin properties and triaxial-ellipsoid shapes have been studied largely through observations of \\add{rotationally induced variability in disk-integrated brightness} \\rem{disk-integrated photometric variations during rotation} (lightcurves). Indeed, the object's shape, its rotational state (period and spin-vector coordinates), and the scattering properties of its surface can be determined from the analysis of its lightcurves over time (as the viewing/illumination geometry changes).\\\\ \\indent For about a decade, starting with the lightcurve inversion algorithm presented by \\citet{2001-Icarus-153-Kaasalainen-a} and \\citet{2001-Icarus-153-Kaasalainen-b}, lightcurves also have been used extensively to derive 3-D shape models of asteroids (see the examples for asteroids Gaspra and \\v{S}teins in Fig.~\\ref{fig: models}). The 3-D shape models and spin properties of more than 200 asteroids already have been derived \\citep[these are accessible from DAMIT\\,\\footnote{DAMIT:~\\href{http://astro.troja.mff.cuni.cz/projects/asteroids3D/web.php}{http://astro.troja.mff.cuni.cz/projects/asteroids3D/web.php}}, see][]{2010-AA-513-Durech}. These shape models are, however, limited to dimensionless, \\textsl{convex} shapes, with limited spatial resolution. Recently, \\citet{2011-Icarus-214-Durech} have shown that the size of these models can be set by using other types of input data (\\eg, stellar occultation profiles, see Sect.~\\ref{ssec: occ}). The intrinsic \\textsl{convex} nature of the models precludes accurate determination of the volume, and hence density, of the objects, however. \\\\ \\indent Because lightcurve observations require neither large telescope aperture nor specialized instrumentation, they are, and will remain, a major source of information on small bodies. Thousands of lightcurves, for hundreds of asteroids\\,\\footnote{\\eg, the Asteroid Photometry Catalogue (APC) or the Asteroid Lightcurve Database \\citep[LCDB,][]{2009-Icarus-202-Warner} list more than 6000 lightcurves for about 700 asteroids: \\\\ APC: \\href{http://asteroid.astro.helsinki.fi/apc}{http://asteroid.astro.helsinki.fi/apc}\\\\ LCDB: \\href{http://www.minorplanet.info/lightcurvedatabase.html}{http://www.minorplanet.info/lightcurvedatabase.html} }, have been accumulated during the last half century. Amateur astronomers contribute significantly to this ever-growing database of lightcurves\\,\\footnote{\\eg, about 2300 lightcurves for more than 1700 asteroids have been acquired by the CdR group: \\href{http://obswww.unige.ch/~behrend/page_cou.html}{http://obswww.unige.ch/$\\sim$behrend/page\\_cou.html}} \\citep[\\eg,][]{2006-AA-446-Behrend, 2007-AA-465-Durech}.\\\\ \\indent Sparse photometry (\\ie, when the typical separation between measurements is larger than the rotation period, as opposed to historical lightcurves, which are dense in time) can also be used to reconstruct 3-D models \\citep[see][]{2004-AA-422-Kaasalainen}. \\citet{2011-AA-530-Hanus} have used a combination of sparse photometry, together with dense lightcurves, to derive about 100 new shape models, using measurements extracted from large all-sky surveys (such as USNO, Catalina, Siding Spring, and Hipparcos). \\add{Knowledge of the absolute photometry (as opposed to relative photometric measurements, as in a dense lightcurve) is, however, required to make these sparse measurements useful to 3-D shape modeling \\citep[see][for a detailed discussion]{2011-AA-530-Hanus}. } From the upcoming PanSTARRS and Gaia surveys we can expect hundreds of thousands of objects to be modeled using this method \\citep{2005-EMP-97-Durech}.   \\begin{figure} \\includegraphics[width=.5\\textwidth]{fig1.eps} \\caption{% Comparison of four shape models, derived using ground-based observations, with high-spatial-resolution images obtained \\textsl{in situ} by spacecraft (left-to-right): (951) Gaspra lightcurve inversion model from \\citet{2002-AsteroidsIII-2.2-Kaasalainen}, image from NASA Galileo; (2867) \\v{S}teins lightcurve inversion model from \\citet{2008-AA-487-Lamy-b, 2008-AA-487-Lamy-a}, image from ESA Rosetta; (25143) Itokawa radar delay-Doppler model from \\citet{2005-MPS-40-Ostro}, image from JAXA Hayabusa; and (21) Lutetia KOALA model from \\citet{2010-AA-523-Carry} and \\citet{2010-AA-523-Drummond}, image from ESA Rosetta. \\label{fig: models} } \\end{figure}   \\subsection{Stellar occultation\\label{ssec: occ}} \\indent Occultations occur when a solar-system object passes between a star and Earth, causing the asteroid's shadow to cross some portion of Earth as a track.  Because the star is very far relative to the Earth-asteroid separation, the shadow cast by the asteroid is effectively parallel. Thus, the width of the shadow track (perpendicular to the track) gives the maximum width of the asteroid in the cross-track direction. It is usually not possible to get a high density of observers stationed across the track, and thus this dimension may not be so well established. But along the track, the size of the asteroid, as an along-track chord on the asteroid, corresponding to the position of the observer within the shadow, is given by the length of time of the blink-out event.  With many observers, many along-track events can be recorded and the blink-out intervals are converted to chord lengths at the asteroid by using the asteroid's known speed \\citep[see][for a review]{1989-AsteroidsII-Millis}.\\\\ \\indent One advantage of stellar occultations is that very small minor planets can be probed (provided one accounts properly for diffraction effects). Even stellar occultations by small trans-Neptunian objects (TNOs) of a few kilometers diameter can be recorded \\citep[see][]{2009-EMP-105-Roques}. Finally, stellar occultations provide a powerful means for studying thin atmospheres and/or exospheres \\citep[\\eg,][]{2003-Nature-424-Sicardy}.\\\\ \\indent In theory, three occultation events (each recorded by enough observers to secure a unique apparent-ellipse to be fit on the chords) provide enough constraints to determine the triaxial dimensions (ellipsoid) and spin-vector coordinates of an asteroid \\citep{1989-Icarus-78-Drummond,1990-AJ-99-Dunham}. The number of chords that result from an event will often be larger with brighter occulted stars. This is because many more observers can be fielded if the \\rem{requirement for telescope aperture} \\add{required telescope aperture} is modest.  In practice, however, occultations of bright stars by any given asteroid occur rather infrequently. This difficulty in obtaining sufficient chords, plus the \\rem{levels of noise} \\add{noise level} often encountered (especially the systematic errors coming from imperfect knowledge of asteroid and star astrometry, combined with observer timing errors), strongly limit our ability to construct 3-D shape models and derive spin properties from occultations alone. Nevertheless, stellar occultations are an efficient way to provide additional size/shape information, particularly for shape models that otherwise lack a scale, such as those from lightcurves alone \\citep[see~\\ref{ssec: lc} and][]{2011-Icarus-214-Durech}.\\\\ \\indent With the availability of low-cost GPS positioning equipment and CCD cameras (the majority of measurements are made by amateur astronomers), the accuracy of occultation timings has improved greatly over the last decade. From this improved precision, together with the refinement of the orbits of small bodies expected to result from Gaia/PanSTARRS \\citep[allowing an extremely precise prediction of the occultation track on Earth, see][]{2007-AA-474-Tanga}, we can expect that stellar occultations will become ever more useful in the determination of certain physical properties of asteroids, especially for objects having small angular diameters.     \\subsection{Thermal radiometry\\label{ssec: thermal}} \\indent The amount of thermal emission from an asteroid is mainly a function of its physical diameter and surface albedo, and, to a lesser extent, the physical properties of its surface (\\eg, thermal inertia, roughness). Main-belt asteroids are among the brightest sources in the sky in the mid-infrared (5--20\\,microns), so infrared satellites (IRAS, ISO, AKARI, Spitzer, \\add{Herschel,} WISE) have been able to acquire observations of a vast number of these objects \\citep[see][for instance]{2011-ApJ-731-Mainzer}.\\\\ \\indent Estimates for size and albedo are model-dependent, however, and several thermal models co-exist, from the simple Standard Thermal Model (STM) of non-rotating spheres of \\citet{1986-Icarus-68-Lebofsky} to the detailed Thermophysical Model (TPM) of \\citet{1996-AA-310-Lagerros,1997-AA-325-Lagerros}, having a complete description of 3-D shape and surface properties. Unfortunately, the systematic uncertainties involved when applying those models (resulting from their respective assumptions and approximations) are not always properly taken into account in estimating error bars, and results often differ from one determination to another by more than the quoted uncertainties \\citep[see Table 3 in][illustrating the issue]{2009-PSS-57-Delbo}. For instance, it now seems that the database of 2\\,228\\,diameters \\citep{2002-AJ-123-Tedesco-a} estimated using the overly simple STM with IRAS data may be biased by \\add{a few percents \\citep[see the re-analysis of IRAS data by][]{2010-AJ-140-Ryan}.}\\\\ \\indent The radiometric technique to derive sizes and albedos from thermal infrared data benefits greatly from the availability of 3-D shape models. For example, the absolute size of (25143) Itokawa derived by \\citet{2005-AA-443-Muller}, based on a 3-D shape model implementation in a TPM code, agreed \\add{to} within 2\\% of the final in-situ result from the Hayabusa mission \\citep{2006-Science-312-Fujiwara}. In general, mid-infrared observations are highly valuable for scaling dimensionless shape models (similar to the situation with stellar occultations, as mentioned in Sect.~\\ref{ssec: occ}) and mid-infrared data are available for several thousand asteroids. These data even allow determination of the most likely spin-axis solutions in cases where lightcurve inversion techniques lead to more than one possible shape and spin-axis solution \\add{\\citep[see, \\eg,][for a recent example]{2011-AA-525-Muller}}. In cases where the shape-model already comes with size information (or alternatively, if many thermal observations are available for a given target), it is possible to derive the thermal inertia, indicative of the surface characteristics: \\eg, bare rock, ice, boulders, dust regolith \\add{\\citep[see, for instance][among many others]{2005-AA-443-Muller, 2007-Icarus-190-Delbo}}.\\\\ \\indent In some particular cases of extensively observed asteroids, the thermal radiometry can also provide hints on the 3-D shape, through the measure of the apparent projected cross-section of the asteroid on the plane of the sky at each epoch. By comparing the predicted with observed thermal fluxes of Lutetia under many geometries, \\citet{2011-PSS--Rourke} have shown that adding a hill/plateau, whose size remains within the quoted 3-D shape uncertainty, could explain the discrepancies observed for a certain observing geometry.   \\subsection{Radar delay-Doppler echoes} \\indent Radar observations \\rem{involve transmitting a signal toward} \\add{consist in transmitting a radio signal toward} the target and recording the echo. The distribution of the echo power in delay time and Doppler frequency is related to the spin state and 3-D shape of the target \\citep[see the reviews by][]{1989-AsteroidsII-2-Ostro, 2002-AsteroidsIII-2.2-Ostro}. The time and frequency precision currently achievable (Arecibo, Goldstone) allow the study of very small objects, the main limit of radar observations being the distance of the target (echo power scales inversely with distance to the fourth power). This is why most radar studies have concentrated on Near-Earth Objects (NEOs), while dedicated observations of Main-Belt Asteroids (MBAs) have been more limited \\citep[see][]{2002-AsteroidsIII-2.2-Ostro}. \\\\ \\indent The difficulty in constructing 3-D shapes from a series of delay-Doppler ``\\textsl{images}'' is due to the absence of a direct, one-to-one, link between each region of the surface with a pixel in delay-Doppler space. Indeed, all points situated at the same range from the observer, and moving at the same radial velocity (possibly spread over the apparent disk) will contribute to a single delay-Doppler pixel. So delay-Doppler images are many-to-one maps of the shape as highlighted by \\citet{2002-AsteroidsIII-2.2-Ostro}: there is no \\textsl{a priori} way to determine how many regions will contribute to a single pixel, which strongly contrasts with the one-to-one mapping (``WYSIWYG'') achieved in disk-resolved imaging.\\\\ \\indent Radar echoes remain, however, the best way to determine the physical properties of NEOs (\\eg, the very small NEO Itokawa in Fig.\\ref{fig: models}). For instance, the possible migration of the regolith at the surface of a fast-rotating asteroid triggered by YORP spin-up \\citep{2008-Nature-454-Walsh} was suggested by the detailed 3-D shape of the NEO (66391) 1999 KW$_4$ \\citep{2006-Science-314-Ostro}.   \\subsection{Disk-resolved imaging\\label{ssec: hari}} \\indent Since the 1990s, with the advent of the Hubble Space Telescope and large ground-based telescopes equipped with adaptive optics (AO: Lick, CFHT, Keck, ESO VLT, Gemini), we have access to the angular resolution required to resolve the apparent disk of asteroids \\citep[\\eg,][]{1993-Icarus-105-SaintPe2, 1993-Icarus-105-SaintPe1, 1998-Icarus-132-Drummond, 2002-AJ-123-Parker,2006-AdSpR-38-Parker}.\\\\ \\indent From a time-series of disk-resolved images, spin-vector coordinates can be derived (using previous knowledge of the rotation period) by analyzing the changes in the apparent shape of the asteroid profile \\citep[see for instance][]{1997-Icarus-128-Thomas, 2008-Icarus-197-Drummond}, or by following the apparent path taken by an albedo patch on the surface during the rotation \\citep[\\eg,][]{2005-Nature-437-Thomas, 2008-AA-478-Carry}. Triaxial shapes (ellipsoids) can also be derived \\citep[see][for instance]{2009-Icarus-202-Drummond, 2009-Science-326-Schmidt}, and topography (such as the presence of facets or craters) studied from profile deviations to the ellipsoid \\citep[\\eg,][]{1997-Science-277-Thomas, 2007-Icarus-191-Conrad}. With sufficient spatial resolution, imaging of the asteroid disk can allow construction of albedo maps of the surface, allowing the study of composition heterogeneity or localized space weathering effects \\citep[\\eg,][]{1997-Icarus-128-Binzel, 2006-Icarus-182-Li, 2010-Icarus-208-Li, 2008-AA-478-Carry, 2010-Icarus-205-Carry-a}. In the case of asteroids visited by spacecraft, high spatial-resolution and precise photometry can be used to derive precise shape and digital terrain models by using stereophotoclinometric techniques \\citep[see examples in][]{2008-MPS-43-Gaskell}.\\\\ \\indent The size and 3-D shape resulting from disk-resolved images are of great value, being obtained in a direct manner (as compared to an indirect determination of the shape from lightcurve inversion, for instance). The best angular resolution\\,\\footnote{limited by the diffraction, which acts as a low-pass filter with a \\add{cutoff} frequency approximated by $\\Theta$ = $\\lambda$/$\\mathcal{D}$ (radian), $\\lambda$ \\add{being} the wavelength and $\\mathcal{D}$ the diameter of the telescope aperture.\\label{foot: diff}} of current Earth-based telescopes is about 0.04\\arcsec. Due to systematics, however, we have found that our ability to accurately measure sizes degrades rapidly below about 0.10\\arcsec, based on simulations and observations of the moons of Saturn and other objects \\citep{these-carry, 2009-DPS-41-Drummond}.  The sample of asteroids observable (\\ie, having angular sizes that get above about 0.10\\arcsec) is therefore limited to about 200 (over a given 10-year span).\\\\     \\subsection{Interferometry} \\indent Apart from building larger telescopes, one efficient way to improve the angular resolution is to combine light beams from separated telescopes and \\add{to} observe the resulting interference \\add{(fringes)}. In such a mode, each telescope aperture plays the role of a sample aperture within a virtual telescope whose extent is the largest distance between the two telescopes (the spatial resolution $\\Theta$ is still given by the equation in footnote $\\ref{foot: diff}$, except that $\\mathcal{D}$ is now the \\textsl{distance} between the apertures). For instance, with telescopes separated by about 80\\,m, the VLTI provides an order-of-magnitude improvement in angular resolution over a single telescope of the VLT (8\\,m aperture). \\\\ \\indent This improvement in the resolution, however, occurs at the price of a loss in complete spatial information, because the virtual aperture is under-sampled. The very high angular-resolution is limited, at a given instant, to a single line \\rem{in the sky plane,} \\add{on the plane of the sky,} which is \\add{given} by the baseline linking the two apertures. Along that baseline, the signal is directly related to the Fourier transform of the flux distribution on the plane of the sky. To expand the coverage in the spatial-frequency domain, and thus \\add{to} allow the construction of 2-D images of the target, one must increase the number of projected baselines. This is commonly achieved by either adding multiple physical baselines (\\ie, adding more telescopes) and/or by making observations throughout the night, as the Earth's rotation causes a progression of the position angle of the baseline on the sky. However, because asteroids typically complete their rotation in only a few hours, the method using multiple telescopes is more effective at providing higher spatial sampling. Below we describe another promising technique to effectively increase the number of baselines, by using a Fizeau design.\\\\ \\indent For objects with no a priori information, the overall dimensions and spin properties can be determined under the assumption that the shape is well-described by a triaxial ellispoid \\citep{2011-Icarus-211-Li}. For example, the two orthogonal directions of the Fine Guidance Sensor (FGS) on-board HST have been used to study binarity and measure the size of several MBAs \\citep{2001-Icarus-153-Tanga, 2002-AA-391-Hestroffer, 2003-AA-401-Tanga}. \\add{The fringes of interference, however, also contain information on the apparent shape, and can be used in combination with other data to derive 3-D shape models. For instance, \\citet{2006-IP-22-Kaasalainen} used FGS data in combination with optical lightcurves to refine the 3-D shape model of (15) Eunomia from \\citet{2002-Icarus-159-Kaasalainen}.} \\\\ \\indent Interferometry in the mid-infrared has been used recently also, with promising results. In this wavelength regime, the signal is linked to the distribution of temperature (\\ie, \\textsl{emitted} light) on the surface of the target (as opposed to the reflected light seen at visible wavelengths). \\citet{2006-Icarus-181-Delbo} and \\citet{2009-ApJ-694-Delbo} have combined thermal infrared TPM models (see Sect.~\\ref{ssec: thermal}) with interferometry, allowing the size of 3-D shape models to be set \\citep[similar to stellar occultation and thermal-infrared-only, see Sections~\\ref{ssec: occ} and~\\ref{ssec: thermal} and][]{2011-Icarus-215-Matter}.\\\\ \\indent The current limitation of interferometry is set by the sensitivity of available facilities and is driven largely by the integration time, which is usually short. Because the light beams from the two apertures traverse different paths in the atmosphere, their wavefronts encounter different turbulence-dominated perturbations, leading to a shift in their phase. Delay lines are used to ``slow down'' one beam with respect to the other, and \\add{to} match their phase. The technical difficulties in maintaining coherence in this process limits the integration time to few hundredths or thousandths of a second. A new generation of instruments (like PRIMA at VLT) with fringe-tracking systems will overcome this limitation in the near future, \\rem{pushing the technique toward fainter sources} \\add{allowing fainter sources to be targeted}.\\\\ \\indent An interesting compromise between co-axial interferometers and traditional, filled aperture, telescopes can be found in Fizeau-type instruments such as LINC-NIRVANA, being built for the Large Binocular Telescope \\citep[LBT, see][]{2010-SPIE-7733-Hill}, in which improved coverage of the \\add{spatial-frequency plane (commonly called uv-plane, with u and v standing for the two orthogonal unit directions)} can be achieved in less telescope time. This design (Fizeau \\textsl{vs.} pupil-plane interferometer like the VLTI) allows instantaneous filling of the uv-plane up to the frequency set by the 8\\,m apertures (Fizeau interferometers are true imaging devices and produce direct images of the plane of the sky), with an additional filling along one dimension up to the frequencies corresponding to the maximal baseline (22.7\\,m for LBT). Two or three epochs, with different position angles of the baseline on the plane of the sky, will be enough to fill the uv-plane up to the maximal extent of the telescope (achievable for transiting asteroids, especially with high elevation, when the position angle evolves quickly). \\rem{The design of the LBT is therefore close to a 22.7\\,m} \\add{For this purpose, therefore, the LBT is equivalent to a 22.7\\,m} telescope with a mask, corresponding to the configuration of component apertures, placed in the entrance pupil. \\\\ \\indent The opening of interferometric studies to longer wavelengths also has great potential, with the use of millimeter and sub-millimeter arrays, where there are fewer technological limitations than in the optical range. Future facilities such as ALMA, with 50 antennas \\citep[translating into 1\\,225\\,baselines \\textsl{vs}. only 2\\,baselines for MIDI at the VLTI, see][]{2009-ApJ-694-Delbo}, will allow dense spatial coverage, together with an angular resolution of few milli-arcseconds. Simulations have shown that several hundreds of MBAs and TNOs will be observable with ALMA \\citep[see][]{2009-Icarus-200-Busch, 2011-Icarus-213-Moullet}. Interferometry at thermal wavelengths (mid-infrared to millimeter), in combination with other techniques (\\eg, lightcurves, see Sect.~\\ref{ssec: lc}), will thus allow the derivation of 3-D shape models, together with thermal properties, for many asteroids of small apparent diameter. ",
        "conclusions": "} \\indent The flyby of (21) Lutetia by ESA Rosetta provided a spectacular demonstration of the capabilities of the KOALA algorithm. Spin-vector coordinates were found to be accurate within two degrees, 3-D shape modeling to better than 2\\,km~(local topography), and dimensions to 2\\%. Volume estimates provided by KOALA are seen to be accurate to \\add{better than} 10\\%. The thermal properties (albedo and thermal inertia) of Lutetia determined using a thermophysical model in conjunction with the KOALA shape model agree with the Rosetta-flyby-derived values, within the quoted uncertainties. \\add{These levels of accuracy on the spin and 3-D shape/size are typical for large main-belt asteroids, and not specific to Lutetia. Although it was extensively observed from the ground, being a spacecraft target, the number of lightcurves was similar to that of other large main-belt asteroids and the geometry of imaging observations was not particularly favorable (\\ie, mostly close to ``pole-on'').} \\\\ \\indent This ability of KOALA to determine the volume of \\add{main-belt asteroids of size $\\sim$100\\,km} with an accuracy of about 10\\% opens the possibility for study of a larger set of small bodies using Earth-based observations. For instance, it can be expected to help efforts to better understand the densities of asteroids belonging to different taxonomic classes \\citep{2009-Icarus-202-DeMeo}. This will be assisted by mass estimates from about 150--200 asteroids that will be determined from gravitational deflections observed by the upcoming ESA Gaia astrometry mission \\citep{2007-AA-472-Mouret}, and the ever-growing number of \\add{known} binary asteroids. During the first stages of the encounter of Rosetta with its main target, comet 67P/Churyumov-Gerasimenko in 2014, we also plan to use \\add{a} KOALA-style analysis of the first, low-resolution, resolved images to quickly produce a shape/size model. This will support the mission science until a full high-resolution model can be derived.\\\\ \\indent \\add{The} present implementation of KOALA allows the combined use of optical lightcurves (including sparse photometry), profiles from disk-resolved images, and chords from stellar occultations. We continue the development of KOALA toward the use of more data modes (\\eg, interferometry, thermal radiometry) to increase the number of possible targets, and \\add{to} set better constraints on targets observable only with certain techniques. Using different wavelength ranges, in particular in the thermal infrared, opens the \\add{possibility} of deriving additional physical properties, like albedo or thermal inertia. We also plan to incorporate additional cross-checks and constraints on the inversion, such as rigorous attention \\add{to differences in the spatial resolution of input images}, to improve the confidence on the non-convex features and details.\\\\"
    },
    "1112/1112.5343_arXiv.txt": {
        "abstract": "Accretion discs are composed of ionized gas in motion around a central object. Sometimes, the disc is the source of powerful bipolar jets along its rotation axis. Theoretical models invoke the existence of a bipolar magnetic field crossing the disc and require two conditions to produce powerful jets: field lines need to be bent enough at the disc surface and the magnetic field needs to be close to equipartition. The work of Petrucci et al (2008) on the variability of X-ray binaries supposes that transitions between pure accretion phases and accretion-ejection phases are due to some variations of the disc magnetization. This rises the problem of the magnetic field dragging in accretion discs. We revisit the method developed by Lubow et al (1994) by including momentum and mass conservation equations in a time-dependent 1D MHD code. ",
        "introduction": "First works on magnetic field dragging in accretion discs have focused on the \\citet{Blandford_&_Payne} necessary condition for cold jet launching, namely: $tan^{-1}(B_r/B_z)\\geq 30\u00b0 $. Thus, \\citet{1989ASSL..156...99V} , \\citet{Lubow_&_al}, \\citet{Reyes_&_al} developed different methods to address this question in a standard accretion disc (here after SAD). They all conclude that a large effective Magnetic Prandtl number  (here after $Pm=\\nu_v/\\nu_m$ as the ratio of the turbulent viscosity to the magnetic diffusivity) is necessary to reach a sufficient magnetic flux accumulation that satisfy the \\citet{Blandford_&_Payne} criterion. Indeed, in a SAD, bent field lines requires a magnetic Reynolds number $R_m=rU_r/\\nu_m=R_e P_m \\sim r/h $ (with $R_e$ the Reynolds number), whereas in a SAD, $Re \\sim 1$. This is problematic as it tends not to be satisfied according to 3D MRI simulations \\citep{Lesur&Longaretti2009}. More recent works, by \\citet{Lovelace2009}, showed that significant advection of magnetic field could be possible even with $Pm \\sim 1$.\\\\ On the other hand, the existence of jets do require the presence of a large scale vertical field \\citep{McKinney&Blandford}. It has been shown that this field must be close to equipartition in order to drive powerful jets \\citep{Ferreira&Pelletier95}.\\\\ In this work, we revisit the standing problem of field advection in accretion discs, by introducing new physical input from modern works. ",
        "conclusions": "The new advection of magnetic field supposed to occur at the end of the simulation would make a transition with the diffusion of the magnetic field observed in the simulation. Thus, we could assist to a phenomenon of magnetic tide. In this sense, this result is very encouraging for the study of transition from high soft to low hard state in X-ray binaries. However here, the accretion rate remains roughly constant during the simulation. But if we work by imposing the accretion rate at the outer radius, a decreasing of it would enhance the magnetization at the outer radii, which is compatible with a transition between high soft to low hard state."
    },
    "1112/1112.1820_arXiv.txt": {
        "abstract": "Inhomogeneous reionization acts as a source of arcminute-scale anisotropies in the cosmic microwave background (CMB), the most important of which is the kinetic Sunyaev-Zel'dovich (kSZ) effect. Observational efforts with the Atacama Cosmology Telescope (ACT) and the South Pole Telescope (SPT)  are poised to detect this signal for the first time, with projected 1 $\\muKK$-level sensitivity to the dimensionless kSZ power spectrum around a multipole of $l=3000$, $\\Ptot$.  Indeed, recent SPT measurements place a bound of $\\Ptot<2.8~\\muKK$ at 95\\% C.L., which degrades to $\\Ptot<6~\\muKK$ if a significant correlation between the thermal Sunyaev-Zel'dovich (tSZ) effect and the cosmic infrared background (CIB) is allowed. To interpret these and upcoming observations, we compute the kSZ signal from a suite of $\\approx 100$ reionization models using the publicly available code \\cmfast. Our physically motivated reionization models are parameterized by the ionizing efficiency of high-redshift galaxies, the minimum virial temperature of halos capable of hosting stars, and the ionizing photon mean free path -- a parameterization motivated by previous theoretical studies of reionization. We predict the contribution of patchy reionization to the $l=3000$ kSZ power to be $\\PkSZ=$ 1.5--3.5 $\\muKK$.  % Therefore, even when adopting the lowest estimate in the literature for the post-reionization signal of $\\POV \\approx 2~\\muKK$, none of our models are consistent with the aggressive 2$\\sigma$ SPT bound that does not include correlations. This implies that either: (i) the early stages of reionization occurred in a much more homogeneous manner than suggested by the stellar-driven scenarios we explore, such as would be the case if, e.g., very high energy X-rays or exotic particles contributed significantly; and/or (ii) that there is a significant correlation between the CIB and the tSZ.  The later is perhaps not surprising, as massive halos should host both hot gas and star forming galaxies. On the other hand, the conservative SPT bound of $\\Ptot \\lsim 6~\\muKK$ is compatible with all of our models and is on the threshold of constraining physically motivated reionization models.  The largest patchy kSZ signals correspond to an extended reionization process, in which the sources of ionizing photons are abundant and there are many recombinations (absorptions in sinks). We point out that insights into the astrophysics of the early Universe are encoded in both the amplitude and shape of the kSZ power spectrum. ",
        "introduction": "\\label{sec:intro}  Scattering of cosmic microwave background (CMB) photons off of inhomogeneities in the electron density after the recombination epoch imprints ``secondary'' anisotropies in the CMB.  These secondaries are the dominate contribution to the CMB anisotropies at $\\lsim 5'$ scales.   The largest and most studied of these secondaries is the thermal Sunyaev-Zel'dovich (tSZ) effect, which results from Compton scattering off of $\\sim 10^7~$K intracluster gas \\citep{Zeldovich69}, and most of the tSZ arises from $z<1$ (e.g., \\citealt{Komatsu02}).  The second most prominent is the kinetic Sunyaev-Zel'dovich (kSZ) effect and is the subject of this study.  In contrast to the tSZ, the kSZ is not sourced by rare regions in the Universe but arises because most of the volume has a bulk peculiar motion with respect to the CMB frame \\citep{SZ80, OV86, Vishniac87, MF02}.  These motions impart a Doppler shift to the $\\approx 10\\%$ of CMB photons that scatter after recombination.  The kSZ is typically generated at higher redshifts than the tSZ, with a large fraction likely being sourced by inhomogeneities at $z\\gsim 6$, during the epoch of reionization \\citep{GH98, KSD98, VBS01, Santos03, Salvaterra05, McQuinn05, Zahn05}.  Until recently, the prospects for detecting the kSZ were perceived by the CMB community as bleak.  While the unique frequency dependence of the tSZ anisotropies should allow the separation of the kSZ from the tSZ, theoretical projections were that the tSZ signal would be more than an order of magnitude larger than that of kSZ on arcminute scales (e.g., \\citealt{Zahn05}).  As such, instruments targeting the CMB secondaries would face the difficult task of using a limited number of frequency channels ($\\sim 3$) to isolate the much smaller kSZ signal while also using this frequency information to isolate the contamination from dusty galaxies and active galactic nuclei \\citep{HS05}.  However, recent measurements have indicated that the amplitude of the tSZ power spectrum may actually be comparable to that of the kSZ rather than an order of magnitude larger.  This lower value has been attributed to the lower $\\sigma_8$ that is preferred by current cosmological measurements relative to the estimates of a few years ago (the amplitude of the tSZ power spectrum scales as $\\sim \\sigma_8^9$; \\citealt{Komatsu02}), and also to the realization that feedback processes strongly suppress the tSZ in clusters \\citep{Battaglia10, Shaw10, TBO11}.  As a result, separating the two signals is now expected to be an easier task than had been anticipated.  Both the Atacama Cosmology Telescope (ATC)\\footnote{http://www.physics.princeton.edu/act/} and South Pole Telescope (SPT)\\footnote{http://pole.uchicago.edu/} began measuring the CMB at arc-minute scales in 2007 and have recently published upper limits on the kSZ angular power spectrum (defined in Section 2). At present, the greatest challenge for these efforts is separating out several different terms that contribute to the angular power spectrum: the primordial CMB, the tSZ, radio sources, the cosmic infrared background (CIB), in addition to cross-power between the tSZ and radio sources, and between the tSZ and the CIB.   The primordial fluctuations are important at low multiples and the sources (CIB/dusty galaxies and radio sources) are most important at high multipoles. The tSZ-CIB cross-correlation is currently the least understood component.  With ACT's  two frequency bands, \\citet{Dunkley11} placed a limit of ${\\it l}^2/[2\\pi] [ C^{\\rm tSZ}_l + C^{\\rm kSZ}_l] =  6.8\\pm 2.9~\\muKK$ at the angular multipole of $l \\approx 3000$,  where $C^{\\rm tSZ}_l$ and  $C^{\\rm kSZ}_l$ are the tSZ and kSZ angular power spectra, respectively.  Using  SPT's three frequency bands and ignoring the tSZ-dusty galaxy correlation,  \\citet{Reichardt11} recently reported an upper bound ${\\it l}^2/[2\\pi] \\,C^{\\rm kSZ}_l < 2.8~\\muKK$ at 95\\% C.L.. \\citet{Reichardt11} did note that the kSZ constraint is significantly weakened to ${\\it l}^2/[2\\pi] \\,C^{\\rm kSZ}_l< 6~\\muKK$ if this study allows for a significant anti-correlation between the tSZ and the cosmic infrared background. Analysis of the Planck data \\citep{Planck18} suggest that 90\\% of the dust anisotropy power at $l \\sim 3000$ is from $z > 2$ while simulations of the tSZ signal predict that the bulk of the signal is sourced from lower redshifts \\citep{Battaglia11,TBO11}.  While this Planck analysis argues that the tSZ-CIB cross-correlation is small, further work is needed to better characterize the redshift distribution of these signals as current models do not provide good fits to the Planck data \\citep{Shang11}.  The detection of the kSZ signal will likely require a combination of Herschel measurements of the dusty galaxies and multi-frequency millimeter observations by SPT and ACT. Nevertheless, even the recent conservative bound of ${\\it l}^2/[2\\pi] \\,C^{\\rm kSZ}_l< 6~\\muKK$  rules out the most extended reionization models in \\citet{McQuinn05}.   The kSZ is comprised of two components: a component sourced primarily after reionization and referred to as the Ostricker-Vishniac (OV) effect \\citep{OV86, MF02}\\footnote{In some prior studies, ``OV'' has been used to refer to the linear-theory, post-reionization kSZ and the nonlinear theory was referred to as the ``nonlinear kSZ.''  We do not adopt this terminology here.} and a component from during reionization referred to as the patchy kSZ \\citep{GH98, KSD98}.   The OV effect results from scattering off of density inhomogeneities with bulk peculiar velocities, with the fractional contribution per unit redshift peaking at $z\\sim 1$ and monotonically decreasing with redshift above this peak.  The OV is sourced by relatively linear scales, and it is likely that it can be modeled well enough to isolate the kSZ contribution from reionization.   The patchy kSZ signal results from the order unity fluctuations in the ionized fraction that are anticipated by the vast majority of reionization models.  % Reionization also imprints other anisotropies in the CMB, primarily in polarization \\citep{Hu00, Dore07, DS09, Su11}, which are not discussed here. However, other than the large-scale ($l \\lsim 10$) polarization anisotropies, the kSZ is by far the most detectable of the reionization-induced anisotropies.   The patchy kSZ anisotropies were originally calculated for several toy models of patchy reionization \\citep{GH98, KSD98, VBS01, Santos03}.  In the last several years, the kSZ from reionization has been calculated for more physically motivated models of reionization either using numerical simulations or semi-analytic models \\citep{Zahn05, McQuinn05, Iliev07kSZ}.   Numerical simulations of reionization have not achieved large enough box sizes to both resolve this process and also capture the velocity flows which are correlated over hundreds of Mpc and which contribute significantly to the amplitude of the kSZ anisotropies (although, see \\citealt{Iliev07kSZ} for a method to account for flows larger than the box size).  More fundamentally, the heavy computational requirements of numerical simulations do not allow for exploration of the large astrophysical parameter space associated with reionization.  Because of these limitations, recent years have seen the development of so called ``semi-numerical'' simulations of reionization, which incorporate an excursion-set model of reionization \\citep{FHZ04} with large-scale realizations of the density field to produce 3D reionization morphology \\citep{Zahn05, MF07, GW08, Alvarez09, CHR09, MFC11, Zahn11}. In this study, we use the publicly available \\cmfast\\ code\\footnote{http://homepage.sns.it/mesinger/Sim.html} to generate $\\sim$ 100 reionization scenarios in order to explore the parameter space of the patchy kSZ signal \\footnote{When this work was nearing completion, a complementary study by \\citet{Zahn12} appeared in the literature.  This study focused on constraining the duration and timing of reionization with recent SPT results.  We will briefly compare our results to theirs in \\S \\ref{sec:prior}.}.  There are many other potential probes of reionization, including: the CMB Thompson scattering optical depth, $\\tau_e$, (e.g., \\citealt{Komatsu11}); the redshifted 21cm line from neutral hydrogen (see, e.g., \\citealt{FOB06} and references therein); transmission statistics of quasar (QSO) spectra (e.g., \\citealt{WL04_nf, MH04, MH07, Fan06, BH07_quasars, Maselli07, Gallerani08, MMF11, Bolton11}); intergalactic damping wing absorption in gamma-ray burst (GRB) spectra \\citep{Miralda-Escude98, BL04_grbvsqso, Totani06, McQuinn08, MF08damp}; the number density, clustering and evolution of Ly$\\alpha$ emitters (LAEs; \\citealt{MR04,Santos04, HC05,FZH06,MR06, Kashikawa06,McQuinn07LAE, DWH07,MF08LAE,Iliev08, DMW11}). However the CMB-based constraints  % are likely to be the only ones in the next decade that are capable of probing the highest-redshift part of reionization, $z\\gsim10$.  This paper is organized as follows. In \\S \\ref{sec:kSZ}, we summarize our method for estimating the kSZ signal.  Our main results are presented in \\S \\ref{sec:results}, namely the kSZ signal from a large range of reionization scenarios. In \\S \\ref{sec:prior} we compare our results to previous works.  In \\S \\ref{sec:detect}, we put our results in context with recent and forthcoming observations of the kSZ. Finally in \\S \\ref{sec:conc} we present our conclusions.%  Unless stated otherwise, we quote all quantities in comoving units. We adopt the background cosmological parameters \\{$\\Omega_\\Lambda$, $\\Omega_{\\rm M}$, $\\Omega_b$, $n$, $\\sigma_8$, $H_0$\\} = \\{0.73, 0.27, 0.046, 0.96, 0.82, 70 km s$^{-1}$ Mpc$^{-1}$\\}, consistent with the seven--year results of the {\\it WMAP} satellite \\citep{Komatsu11}.   We use the Fourier convention that is the standard in modern cosmological studies in which the $2\\pi$'s appear under the $dk$'s. ",
        "conclusions": "\\label{sec:conc}  Observational efforts such as ACT and SPT are beginning to set interesting constraints on the kSZ power spectrum at $l\\approx3000$.  This is highlighted by the recent SPT constraint of $\\Ptot< 2.8~\\muKK$ (\\citealt{Reichardt11}; $<6~\\muKK$ when allowing for correlations), as well as the ACT measurement of $\\Ptot =  6.8\\pm 2.9~\\muKK$ \\citep{Dunkley11}.  A $\\approx1~\\muKK$ constraint on the signal is projected for the coming year \\citep{Reichardt11}, which should result in the first detection of the kSZ.  Second generation instruments, such as ACTPol and STPPol, will further improve the kSZ measurements in the coming years.  %  This study calculated the patchy kSZ signal for a suite of $\\approx100$ inhomogeneous reionization scenarios, which were generated using the publicly-available code \\cmfast.  % In these models, the contribution of inhomogeneous reionization to the CMB angular power spectrum at $l=3000$ spans only a factor of four, $\\PkSZ \\approx$ 1--4 $\\muKK$, despite the large volume of parameter space covered by the reionization models. % The allowed range narrows further to $\\PkSZ\\approx$ 1.5--3.5 $\\muKK$ when including current constraints on the reionization history from {\\it WMAP} and high-redshift quasars.  We find that for standard models, the bulk of the signal is imprinted in the early to middle stages of reionization, when the (clustered) ionized structures have angular scales comparable to $l\\approx3000$.   The range $\\PkSZ \\approx$ 1.5--3.5 $\\muKK$, combined with estimates for the post reionization kSZ of $\\POV \\approx$ 2--3 $~\\muKK$, results in a total kSZ signal of $\\Ptot = $3.5--6.5 $\\muKK$.  Thus, the recent conservative SPT bound of $\\Ptot <6~\\muKK$, which assumes a free, scale-independent CIB-tSZ correlation, is on the boarder of constraining reionization, even assuming conservatively low values for $\\POV$.  The strongest and easiest to detect patchy kSZ signals correspond to early and extended reionization scenarios.  In such models, the sources of ionizing photons are abundant at early times and/or there are many recombinations (absorptions in sinks) so that reionization was maximally extended.  Furthermore, with the assumption of a slope (but not amplitude) in the CIB-tSZ correlation, \\citet{Zahn12} placed an intermediate upper limit of $\\Ptot \\lsim 4 \\muKK$.  If such a slope is credible, this constraint is sufficient to rule out our reionization models driven by molecularly-cooled halos (with $\\Tvir\\lsim10^4$ K), and place limits of $\\zre\\lsim10$, $\\delz\\lsim2$ (see Fig. \\ref{fig:zre_delz}).   On the other hand, the tightest SPT bound $<2.8~\\muKK$ at 95\\% C.L., which assumes no CIB-tSZ correlation, is incompatible with {\\it all} of our reionization scenarios.  Although later and more rapid reionization models might be in agreement with this constraint, it is difficult to push our models to such extremes since these would conflict with other observations of the high-redshift Lyman alpha forest and galaxy luminosity function (LF). This implies that either: (i) the early to middle stages of reionization occurred in a much more homogeneous manner than suggested by the stellar-driven scenarios we explore, such as would be the case if, e.g., very high energy X-rays or exotic particles contributed significantly; and/or (ii) that there is a significant correlation between the CIB and the tSZ, which invalidates these bounds.  Upcoming combined analyses of microwave and Herschel far infrared data will likely test the later possibility.   We also find that shape of the patchy kSZ power spectrum encodes astrophysical information.  An early and extended reionization in which bubble growth is retarded by recombinations results in a kSZ power spectrum which peaks on smaller scales. The sign of the slope of the patchy kSZ power spectrum at $l\\approx3000$ is useful for discriminating between extended/early and rapid/late reionization processes, which have the same $\\PkSZ$.  In other words, the values of $\\tau_e$ and $\\PkSZ$ do not uniquely determine the redshift and duration of reionization (c.f., \\citealt{Zahn12}).  However, the sign of the kSZ slope at $l\\sim3000$, combined with $\\tau_e$ and $\\PkSZ$, is sufficient to determine the redshift and duration of reionization. Furthermore, we find that positive values of the slope come exclusively from models with low values of $\\mfp \\lsim 10$ Mpc.  Therefore detecting the sign of the slope can constrain the abundance of absorption systems in the early Universe, at redshifts far beyond those of present estimates based on QSO spectra.   We anticipate that these measurements will be complemented by analyses of the Planck EE signal, which is also sensitive to the duration of reionization, and potentially by analysis of the contribution of patchy reionization to the CMB polarization power spectrum.  Reionization probes not based on the CMB will also yield complementary information, using 21cm tomography, QSO spectra and wide-field LAE studies. These tools will soon reveal the details of the dawn of cosmic structure.      \\vskip+0.5in  We especially thank Hy Trac for useful discussions, and Andrea Ferrara for comments on an early draft of this paper. We thank Dave Spiegel for his considerable Matlab assistance without which most of the figures presented here would have been much less visually appealing. AM would like to thank the hospitality of Princeton University where most of this work was completed. MM is supported by NASA through an Einstein Fellowship."
    },
    "1112/1112.1227_arXiv.txt": {
        "abstract": "We examine the gas and stellar metallicities in a sample of HII galaxies from the Sloan Digital Sky Survey, which possibly contains the largest homogeneous sample of HII galaxy spectra to date.  We eliminated all spectra with an insufficient signal-to-noise ratio, without strong emission lines, and without the [OII]~$\\lambda$3727~\\AA\\ line, which is necessary for the determination of the gas metallicity. This excludes galaxies with redshift $\\stackrel{_<}{_\\sim} 0.033$. Our final sample contains $\\sim$700 spectra of HII galaxies.  Through emission line strength calibrations and a detailed stellar population analysis employing evolutionary stellar synthesis methods, which we already used in previous works, we determined the metallicities of both the gas and the stellar content of these galaxies.  We find that in HII galaxies up to stellar masses of $5\\cdot10^9~M_{\\odot}$, enrichment mechanisms do not vary with galactic mass, being the same for low- and high-mass galaxies on average. They do seem to present a greater variety at the high-mass end, though, indicating a more complex assembly history for high-mass galaxies. In around 23 per cent of our HII galaxies we find a metallicity decrease over the last few Gyr. Our results favour galaxy evolution models featuring constantly infalling low-metallicity clouds that retain part of the galactic winds. Above $5\\cdot10^9~M_{\\odot}$ stellar mass, the retention of high metallicity gas by the galaxies' gravitational potential dominates. ",
        "introduction": "HII galaxies are characterised by prominent emission lines. In fact, they are defined as having strong H$\\beta$ and OIII lines, but other ionised hydrogen (H$\\alpha$, H$\\gamma$, etc.) and high excitation metallicity lines (OII, NII, and others) are strong too. These emission lines are driven by photoionisation of interstellar gas by hot massive stars from young stellar populations \\citep{sargent,french}. The star forming activity in HII galaxies is so strong that it can certainly not have been maintained at its present level during a Hubble time \\citep*[see e.g.][]{searle2}. In fact, they are the most extreme case of star forming galaxies, showing the highest excitation emission lines, and thus the highest (relative) star formation rates (SFRs). \\\\ They are also the lowest metallicity galaxies of the interstellar medium, with metallicities around one tenth solar \\citep*[][and others]{vilchez,perezmonteroa,kniazev,kehrigvilchez,izotov,haegele,perezmonteroc}. \\\\ It is now clear that the hypothesis of HII galaxies as galaxies experiencing their first star formation burst can be ruled out, and that their stellar content is predominantly old, older than one Gyr \\citep*[][and others]{raimann,cidfernandes,kong,westera_04,cidasari,asari,hoyos}, Nevertheless, both the high relative SFRs and the low metallicities indicate that HII galaxies are among the least evolved galaxies in existence, which assigns to them a special role as fossil record of galaxy evolution \\citep{lequeux}. They are, therefore, ideal objects for chemical evolution studies.  Several authors have developed theoretical evolutionary models of dwarf irregular galaxies, of which HII galaxies are a sub-category. Some of these models predict the galaxies' present-day stellar - and gas metallicities, among other properties. In these models, several chemical enrichment processes are identified. \\\\ \\citet{maclow} developed single-phased hydrodynamical models for dwarf galaxies. They model the effects of repeated supernova type II (SNII) explosions from starbursts on the interstellar medium, that is, the enrichment by the ejected metals on the one hand, and the gas loss by supernova winds on the other hand, taking into account the gravitational potential of their dark matter haloes. They find that, in galaxies with gas masses below $10^6~M_{\\odot}$, most of the gas is blown away, and in galaxies with masses between $10^7~M_{\\odot}$ and $10^9~M_{\\odot}$, mainly the newly-formed metal-rich gas is blown away, whereas the already present, metal-poor, gas is retained. They do not calculate the final (present-day) metallicities of their model galaxies, but it is obvious that they remain low-metallicity, at least the ones with gas masses below $10^9~M_{\\odot}$. \\\\ \\citet{recchimatta,recchimattb} simulate IZw18, one of the lowest-metallicity galaxies known, using a single star formation burst model and a model with a doubly peaked star formation history. They use a lower SNII heating efficiency than MacLow et al. in their models, but include type Ia supernovae (SNeIa), for which they use a higher heating efficiency. As a consequence, alpha elements are ejected less efficiently than in MacLow et al.'s model, resulting in higher metallicities. The iron-peak elements, which are produced in SNeIa, on the other hand, are still ejected, causing high [$\\alpha$/Fe] ratios. The (gas) metallicity in their single burst model remains very low, [O/H$]\\stackrel{_<}{_\\sim}-1.3$, depending on the galaxy mass. In their (more likely) double burst model, [O/H] can reach values up to -1, depending on the details of the models, that is, the stellar yields used, the stellar initial mass function (IMF), the time between the bursts, the duration of the second burst, the gas density and - metallicity in the star forming region, and the total mass of the stars formed in the second burst. The final metallicities of the double burst model galaxies are between 0.6 and 1 dex higher than after the first burst. \\\\ \\citet{tenoriotagle} include photoionisation and cluster wind in their two-dimensional hydrodynamic calculations. They identify two mechanisms: The storage of clouds into a long-lasting ragged shell inhibiting the expansion of the thermalised wind, and the steady filtering of the shocked wind gas through channels carved within the cloud layer. They conclude that both processes must be at work in HII galaxies. Unfortunately, they make no prediction about the present-day metallicity of their model galaxies. \\\\ In a series of papers, Recchi and collaborators use a new generation of chemo-dynamical models, some of which include infalling clouds. The first two, \\citet{recchimattc} and \\citet{recchihenslera}, are aimed at reproducing the properties of two specific objects, IZw18 and NGC1569, respectively. The third one, \\citet{recchihenslerb} investigates the influence of several factors, such as the cloudiness of the gas distribution, the IMF slope, and stellar yields. They find that models with continuous low level star formation periods in the past followed by a quiescent phase and a recent, stronger burst best reproduce the chemical properties of the studied galaxies. They also find that, for a homogeneous gas distribution, metals get blown away by galactic winds, but cool infalling intergalactic clouds can hamper these galactic winds. The final metallicities in the their various models vary from -2.5 to -0.8.  Although metallicity determinations of the various components of HII galaxies cannot discriminate, how much each of these mechanisms contribute to the chemical evolution, they can quantify the metallicity at various stages of a galaxy's evolution, i.e. at the time of the formation of the old ($>1$ Gyr) stellar populations and at present, and thereby help to rule out some of the evolutionary scenarios, and support others. HII galaxies have the advantage that it is possible to determine independently the metallicities of the young and the old populations from their spectra. As the young populations were formed recently out of the same gas that is responsible for the emission lines, their metallicity can be assumed to be more or less the same as the one of the gas, which can be determined from the emission line strengths. The metallicity of the old populations, on the other hand, can be determined from the continuum and absorption features through population synthesis. \\\\ The main goal of the present work is to measure the chemical enrichment in a homogeneous sample of HII galaxies, and examine possible trends with galactic properties such as total mass. This way, we can evaluate the different models and make statements on the importance of the various enrichment mechanisms.  Until recently, the discovery of HII galaxies was limited to the visual inspection of objective-prism surveys, introducing ill-controlled biases and selection effects \\citep*[for compilations, see][and references therein]{terlevich,kehrig}. \\\\ The Sloan Digital Sky Survey \\citep*[SDSS;][]{york} presents for the first time a comprehensive database of galactic spectra. Data release 7 \\citep{abazajian} contains over 900'000 galaxy spectra that have been selected on clear magnitude limiting criteria. In an unprecedented manner, the SDSS allows to study star forming emission line galaxies based on clear and quantitative criteria, and not mere visual inspection.  In this work, we define clear emission line strength criteria to distinguish galaxies containing strongly excited gas from other emission line galaxies. We only include spectra that contain the [OII]~$\\lambda$3727~\\AA\\ emission line, as this line is indispensable for an adequate determination of the gas metallicity. This criterion excludes galaxies with redshift $\\stackrel{_<}{_\\sim} 0.033$. This way, we selected a homogeneous sample of $\\sim 700$ HII galaxy spectra of high quality from the SDSS. \\\\ In this sample, we determine the metallicities of both the gas and the old stellar content using emission line calibration and stellar population synthesis methods, respectively. It is the first time an independent metallicity determination of the gas and the stars is performed on a clearly defined HII galaxy sample.  The layout of this article is the following: In Sect.~\\ref{spectra}, the catalogue of spectra analysed in this work is presented. Sect.~\\ref{method} gives a detailed description of the method we used to analyse the spectra. Subsection \\ref{OHgas} describes, how we determined the gas metallicity, and Subsection \\ref{FeHstars} is dedicated to the determination of the stellar metallicities. The main results are given in Sect.~\\ref{results} and discussed in Sect.~\\ref{discussion}. A summary and the main conclusions can be found in Sect.~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We performed a  gas - and stellar population analysis of HII galaxies from the Sloan Digital Sky Survey Data Release 7, using their full spectra. We selected a sample which contains 712 HII galaxy spectra, a relatively low number due to our high requirements towards the quality of the spectra and to the criterion that the spectra contain all the spectral lines necessary for our study. We also selected a test sample of 74989 galaxies, using less restrictive requirements, the only criterion being that they show some line emission. Therewith, it is a sample of star forming galaxies. \\\\ We derived independent metallicities for the young populations present in our sample galaxies from the gas emission lines, and for the old populations from full spectra fitting, that is, from the continuum and absorption features from stars.  Our test sample follows the well-known mass-metallicity relation, e.g. the fact that the (gas) metallicity increases with (stellar) galaxy mass, indicating that low mass galaxies are chemically less evolved than high mass galaxies. We do not find any systematical tendency of the difference between the metallicities of the young and the old stellar populations, a quantity we call metallicity enrichment, with galactic mass, indicating that, on average, recycling mechanisms of the interstellar medium should be the same for low - and high mass galaxies.  In our HII galaxy sample, on the other hand, we find no mass dependence of the present day metallicity. For the metallicity enrichment, we find no mass dependence for masses up to $5\\cdot10^9~M_{\\odot}$, whereas above this mass, the enrichment is significantly higher.  We interpret these findings in the light of recent hydrodynamical evolutionary models of dwarf irregular galaxies, of which HII galaxies are a sub-class. We favour models featuring constantly infalling low-metallicity clouds able to retain part of the high metallicity galactic winds, such as the models NGC-4+BC from \\citet{recchihenslera} and NCSM from \\citet{recchihenslerb}, since models without infalling clouds fail to predict the metallicities and the non-dependence on galaxy mass of our sample galaxies. For galaxies with masses above $5\\cdot10^9~M_{\\odot}$, most, if not all, of the metal-rich gas is retained by the galaxy's gravitational potential, as predicted by \\citet{maclow}, which leads to the observed break-down of the constant mass-enrichment relation.  However, the dispersion of both the present day metallicity and the metallicity enrichment does increase with galactic mass, indicating more complex chemical and assembly histories for high mass HII galaxies. \\\\ In 23 per cent of our HII galaxies, the metallicity has decreased in the last few Gyr. We attribute the metallicity decrease mainly to the dilution of the galactic gas by the infalling clouds, which is compatible with the models we favour."
    },
    "1112/1112.3963_arXiv.txt": {
        "abstract": "The known host galaxies of short-hard gamma-ray bursts (GRBs) to date are characterized by low to moderate star-formation rates and a broad range of stellar masses.  In this paper, we positionally associate the recent unambiguously short-hard \\Swift\\ GRB 100206A with a disk galaxy at redshift $z=0.4068$ that is rapidly forming stars at a rate of $\\sim 30$ M$_{\\odot}$ yr$^{-1}$, almost an order of magnitude higher than any previously identified short GRB host.   Using photometry from Gemini, Keck, PAIRITEL, and WISE, we show that the galaxy is very red ($g-K = 4.3$ AB mag), heavily obscured ($A_V \\approx 2$~mag), and has the highest metallicity of any GRB host to date (12 + log[O/H]$_{\\rm KD02}$ = 9.2): it is a classical luminous infrared galaxy (LIRG), with $L_{\\rm IR} \\approx 4 \\times 10^{11} {\\rm L}_{\\odot}$.  While these properties could be interpreted to support an association of this GRB with very recent star formation, modeling of the broadband spectral energy distribution also indicates that a substantial stellar mass of mostly older stars is present.  The current specific star-formation rate is modest (specific SFR $\\approx$ 0.5 Gyr$^{-1}$), the current star-formation rate is not substantially elevated above its long-term average, and the host morphology shows no sign of recent merger activity.   Our observations are therefore equally consistent with an older progenitor, similar to what is inferred for other short-hard GRBs.  Given the precedent established by previous short GRB hosts and the significant fraction of the Universe's stellar mass in LIRG-like systems at $z \\gtrsim 0.3$, an older progenitor represents the most likely origin of this event. ",
        "introduction": "\\label{sec:intro}  Classical gamma-ray bursts (GRBs) segregate into two phenomenological classes on the basis of their prompt-emission properties: long-duration, soft-spectrum GRBs and short-duration, hard-spectrum GRBs \\citep{Kouveliotou+1993}.  The two distributions overlap in both duration and in hardness, but a duration of $T_{90} = 2$~s is commonly taken as the dividing line between ``long'' and ``short'' bursts\\footnote{However, the question of classification has recently become much more complex following the discovery of events with extended emission (\\citealt{Norris+2006}) and deep limits on supernovae from $t>2$~s bursts (\\citealt{Fynbo+2006,Gehrels+2006,GalYam+2006}).  See \\cite{Bloom+2008} for a discourse on the subject of classification.}.  Long bursts constitute the large majority of events detected by all major GRB satellites (about 75\\% of BATSE events and over 95\\% of \\Swift\\ events); because these events have brighter and longer-lived afterglows \\citep{Kann+2011}, they have historically been easier to study.   Observations stretching back almost 15 years associate long-duration bursts exclusively with actively star-forming host galaxies, and specifically with regions bright at ultraviolet (UV) wavelengths \\citep{Fruchter+2006} close to the center \\citep{Bloom+2002} of these galaxies, strongly pointing toward a massive stellar origin.  In a number of cases, evidence for an accompanying supernova has provided direct evidence confirming this conclusion (see \\citealt{WoosleyBloom2006} for a review).  The short GRB sample is much smaller, but its properties are unambiguously different from those of long bursts.  At least two short events (GRBs 050509B and 050724; \\citealt{Bloom+2006,Gehrels+2005,Barthelmy+2005,Berger+2005,Prochaska+2006}) have been associated with massive, evolved galaxies having essentially negligible current star formation, pointing strongly toward a long-lived progenitor \\citep{BloomProchaska2006}.   In addition, the offset distribution for short GRBs relative to their host galaxies extends to much greater distances than for long GRBs --- partially as a simple result of the fact that short-GRB hosts tend to be physically more extended \\citep{Fong+2010}, but short GRBs also generally show minimal correlation with the host blue/UV light or resolved regions of active star formation \\citep{Fox+2005,Fong+2010,Rowlinson+2010b}.   The occurrence of an accompanying bright core-collapse supernova has been definitively ruled out for several short GRBs \\citepeg{Hjorth+2005a,Kocevski+2010}.  The model of short GRBs as mergers of compact binary stars \\citep{Eichler+1989,Paczynski+1991,Narayan+1992} --- either two neutron stars (NS-NS) or a neutron star and a black hole (NS-BH) --- naturally predicts old progenitors, and the observations discussed above appear to strongly support this model.   However, other models permitting older populations do exist \\citepeg{MacFadyen+2005}, and recent observations have painted a picture that appears more complex than it initially seemed after the discovery of the first few short-GRB hosts.  A significant fraction of short GRBs with afterglows actually have no clearly identifiable host \\citep{Berger2010}, suggesting that they have been flung large distances from their original hosts or occur in optically underluminous, distant galaxies in surprisingly large numbers.  Finally, with the exceptions of GRBs 050509B, 050724, and possibly 100117A \\citep{Fong+2011}, every short GRB host to date (at this stage, a fairly formidable sample of $>$20 objects) has had at least modest star formation in relation to its stellar mass \\citepeg{Berger2010}, making it difficult to pin down a minimum progenitor age for the large majority of these events.  The two ``smoking gun'' events of 2005 no longer appear to be particularly representative of the entire class.  The relationship between short and long GRBs is at least vaguely analogous to that between Type Ia and Type II supernovae \\citep[e.g.,][]{Wheeler1981}: Type Ia supernovae are often associated with moderate or old populations (and almost certainly are produced by compact objects, in this case white dwarfs; \\citealt{Nomoto1982,Bloom+2011b,Nugent+2011}), while Type II supernovae are associated exclusively with young populations (and result from massive stellar core collapse). However, recent evidence suggests that Type Ia supernovae can actually be produced fairly rapidly in some cases, given their elevated rate in late-type hosts compared to what might be expected from an exclusively old population \\citep{ScanBildsten2005,Mannucci+2005,Sullivan+2006,Li+2011,Maoz+2011} --- that is, the delay-time distribution is likely to be quite broad, with short-lived stellar progenitors as well as long-lived ones.  It is possible that the short GRB delay-time distribution may be similarly complex.  Confirming this trend --- and in particular, determining whether there is any need for a very young component associated with the youngest (and presumably most massive) stars, which could indicate a physically distinct progenitor \\citepeg{Metzger+2008,Lazzati+2010,Virgili+2011} --- is of clear interest for better understanding the origins of the short-GRB population.  Given the low short-GRB event rate ($< 10$ events yr$^{-1}$) and uncertain selection biases affecting the population ($\\sim 25$\\% of all \\textit{Swift} short GRBs are not detected by the XRT and cannot be localized), teasing such a tendency out of the full population is possible but challenging.  Nevertheless, two recent studies have attempted to accomplish this using different techniques.   \\cite{Virgili+2011} analyzed the short-GRB redshift, fluence, and luminosity distributions compared to predictions from a delay-time convolved star-formation history of the Universe and concluded that a significant fraction of short GRBs may in fact originate from young stars.  \\cite{Leibler+2010} used the properties of the host galaxies of a sample of short GRBs to estimate the age of the predominant stellar population, and similarly concluded that both short-delay and long-delay components are necessitated by the data.  However, even the ``short'' component is consistent with a timescale of order 200 Myr --- much older than the lifetimes of massive stars that produce core-collapse supernovae, long GRBs, and nebular signatures in star-forming galaxies.  The most unambiguous indicator of a very short-delay ($\\lesssim$100 Myr) component would be the discovery of another ``smoking gun'' system, but this time with an extremely \\emph{young} stellar age.  In particular, a short GRB within a starbursting galaxy whose current star-formation rate (SFR) is very large in comparison to its stellar mass (high specific SFR) would be a strong indicator that the event came from a star produced in the starburst, events whose characteristic times rarely exceed a few hundred Myr and are often much less \\citep{DiMatteo+2008,McQuinn+2009,McQuinn+2010}.  In this paper, we discuss the case of GRB 100206A as the first example of a short-duration burst whose host characteristics, at least on the surface, evoke such a system.   In \\S \\ref{sec:obs} we present our observations of the burst and its afterglow, showing it to be an unambiguous short-duration, hard-spectrum GRB with a faint, rapidly-fading afterglow detected only in the X-ray band despite deep, early optical imaging.  We also present observations of two galaxies in or near the X-ray error circle at a variety of wavelengths from optical through the mid-infrared.  In \\S \\ref{sec:anal} we analyze these data in further detail, measuring the metallicity and constraining the underlying stellar population of the brighter host candidate, which we show to be a luminous infrared galaxy (LIRG) at redshift $z=0.4068$ with a current SFR exceeding that of any previous short GRB by almost an order of magnitude.  In \\S \\ref{sec:disc} we argue that {\\it a posteriori} statistical arguments strongly tie this galaxy to the short GRB, and examine the implications for the progenitor of the short-duration burst.  While the current SFR is high, we note that the stellar mass is substantial and dominated by older stars, with no unambiguous evidence for starbursting, merging, or other short-lived features --- consistent with an older progenitor, like that of other short-hard bursts.  We summarize our conclusions in \\S \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  Seen as an isolated event, the age of the progenitor of GRB 100206A is frustratingly inconclusive:  essentially every diagnostic we have applied to the host system is consistent with both a short-lived progenitor ($<$10 Myr) and a long-lived one ($>1$ Gyr).  While the galaxy has a large current SFR ($\\sim 30$ M$_\\odot$ yr$^{-1}$), this is matched by a substantial pre-existing stellar mass ($>10^{11}$ M$_\\odot$).  This, and a largely ordinary disk-like morphology, suggest a steady star-formation history and give no direct indication of the progenitor's age.  Seen in the broader context of other short GRB host-galaxy work, however, our conclusions can be seen as generally supportive of the prevailing paradigm: as discussed in \\S \\ref{sec:implications}, the significant amount of stellar mass present in LIRG systems at $z>0.3$ suggests that it should be no surprise that a few percent of short GRBs should occur within these massive and actively star-forming galaxies, even if their progenitor is exclusively a long-lived object such as a neutron-star binary.   The fraction should be even higher if some of the detection biases are considered: preferential ejection of the progenitor system from low-mass host galaxies (due to supernova ``kicks''), as well as preferential detection of X-ray or optical afterglow from systems in dense environments, would both have much less impact on massive galaxies such as G1 than on more typical lower-mass galaxies (which would be actively selected against), favoring the discovery of this type of host.  Given what has become a quite large short-GRB  sample (46 objects with X-ray or optical localizations and at least 20 likely host galaxies), the occurrence of an event in a LIRG is not only unsurprising for an older progenitor, but could have been predicted.  This, coupled to the lack of \\emph{any} definitive signature tying GRB 100206A to recent star formation in its host galaxy despite careful scrutiny of several possible lines of argument (including the SED-inferred star-formation history and morphology), can be interpreted as evidence that this event is the product of a relatively mature progenitor system in the galaxy.   Given the precedent set by other short GRBs, the most natural explanation for this burst is that it was produced in the same manner as the rest of the population --- for which a compelling statistical case is now building for a  generally old ($>100$ Myr) progenitor, even if the delay-time distribution is indeed broad.  Of course, a younger progenitor is not ruled out.  If this event \\emph{were} the result of a young progenitor (and especially a massive star), it is interesting to note that the high metallicity of this system---at least Solar and perhaps significantly higher---would nevertheless be indicative of an origin distinct from that of long GRBs, which quite notably avoid dusty, massive LIRG-like systems in the local universe.  In that case this event would be more indicative of a broad delay-time distribution for the short-GRB population, rather than an origin from the same population that produces LGRBs due to, for example, off-axis collapsars \\citep{Lazzati+2010}.   While it is possible that a few long-duration GRBs may occur in high-metallicity regions \\citep{Levesque+2010f}, these appear to be relatively rare \\citep{Modjaz+2008,Stanek+2006}, so \\emph{even if} this event were to be associated with the youngest stellar population in this galaxy, it would nevertheless continue to support a distinction between the progenitor systems of short and long GRBs.   \\vskip 1cm"
    },
    "1112/1112.3364_arXiv.txt": {
        "abstract": "Angular momentum is transported outwards through an accretion disc by magnetohydrodynamical (MHD) turbulence thus allowing material to accrete on to the central object. The magneto-rotational instability (MRI) requires a minimum ionisation fraction to drive turbulence in a disc.  The inner parts of the disc around a young stellar object are sufficiently hot to be thermally ionised. Further out, cosmic rays ionise the surface layers and a dead zone forms at the mid-plane where the disc is too cool for the MRI to operate.  The surface density in the turbulent active layer is often assumed to be constant with radius because the cosmic rays penetrate a constant layer.  However, if a critical magnetic Reynolds number, $Re_{\\rm M,crit}$, is used to determine the extent of the dead zone, the surface density in the layer generally increases with radius.  For small critical magnetic Reynolds number of order 1, the constant layer approximation may be a reasonable fit. However, MHD simulations suggest the critical magnetic Reynolds number may be much larger, of order $10^4$. Analytical fits for the surface density in the magnetic active layer show that $\\Sigma_{\\rm m}\\propto Re_{\\rm M,crit}^{-2} R^{\\,9/2}T^{\\,2}$, at temperature $T$ and radius $R$, are a good fit for higher critical magnetic Reynolds number.  For the metallicity variation between our galaxy, the LMC and the SMC, there should be no significant difference in the extent of the dead zone. Observations suggest an increase in the lifetime of the disc with decreasing metallicity that cannot be explained by the dead zone structure (ignoring possible differences in dust abundances). ",
        "introduction": "Low-mass stars are thought to form from the free-fall collapse of a protostellar molecular cloud core to a protostar with a disc on a timescale of a few $10^5\\,\\rm yr$ \\citep{shu87}. Angular momentum transport in accretion discs is driven by turbulence thus allowing material to accrete on to the young star. The magneto-rotational instability (MRI) can drive turbulence if the gas is well coupled to the magnetic field \\citep{balbus91}. However, with a low ionisation fraction the MRI is suppressed \\citep{gammie96,gammie98}.  The inner parts of a disc around a young stellar object are hot enough to be thermally ionised. However, further out, the disc becomes layered with an MRI turbulent (active) layer at each surface and a dead zone at the midplane that is shielded from the ionizing radiation of cosmic rays and X-rays from the star \\citep[e.g.][]{sano00,matsumura03}.  In this work we concentrate on circumstellar discs, but note that dead zone formation is also favourable in circumplanetary discs \\citep{martin11a,martin11b}.    A frequently used assumption in calculations of discs with dead zones is that the surface density of the MRI active surface layer is constant with radius \\citep[e.g.][]{armitage01,zhu09a,terquem08,matsumura09}.  The cosmic rays enter the disc surface and are attenuated exponentially with a stopping depth of surface density around $100\\,\\rm g\\,cm^{-2}$ \\citep{umebayashi81}.  However, this does not allow for recombination effects that may play a significant role in determining the ionisation of the disc.  A more realistic way to determine the dead zone is using a magnetic Reynolds number. MHD simulations show that magnetic turbulence cannot be sustained if the magnetic Reynolds number is lower than some critical value $Re_{\\rm M}<Re_{\\rm M,crit}$ \\citep{fleming00}. However, the critical value is uncertain. Simulations of \\cite{fleming00} that compute the non-linear outcome of instability suggest that $Re_{\\rm M,crit}\\approx 10^4$ without a net magnetic flux through the disc. With a magnetic flux it could be of the order of $100$. \\cite{wardle99} and \\cite{balbus01} considered the linear stability of the discs and find that the effects of Hall electromotive forces are important and the value may be as small as $Re_{\\rm M,crit}=1$. However, conditions for linear instability and for maintaining turbulence are not necessarily the same \\citep{balbus00}. We consider a range of values for the critical magnetic Reynolds number from 1 to $10^4$.  The surface layers of the disc may be ionised by cosmic rays or X-rays.  \\cite{matsumura03} find that cosmic rays determine the extent of the dead zone under typical conditions unless the X-ray energy is very high. We consider only cosmic ray ionisation and neglect X-rays. Electrons in the disc are removed through dissociative recombination with molecular ions and at a slower rate with radiative recombination with heavy metal ions. The ionisation fraction, and hence the extent of the dead zone, is sensitive to the number of metal atoms because they quickly pick up the charges of molecular ions but slowly recombine with the electrons.  The typical life time of a protostellar disc in the solar neighbourhood, where the metallicity is $Z\\approx 0.02$ is of the order of $3-5\\,\\rm Myr$ \\citep[e.g.][]{hernandez09}. For stars near SN 1987A in the Large Magellanic Cloud (LMC) \\cite{demarchi10} find from H$\\alpha$ fluxes that the mean lifetime is of the order of $13\\,\\rm Myr$ and the metallicity there is $Z\\approx 0.008$.  Note however, that these stars are likely to be slightly more massive. Similarly, \\cite{demarchi11} observed the NGC 346 cluster in the Small Magellanic Cloud (SMC), where the metallicity is $Z\\approx 0.002$, and find disc lifetimes of around $20\\,\\rm Myr$. It appears from these results that the life time of the disc increases with decreasing metallicity.  Dead zones affect the accretion rate through the disc and thus lifetime of disc accretion and so we consider how the dead zone varies with metallicity.   In Section~\\ref{sec:layer} we describe the layered disc model where the extent of the dead zone is determined with the critical magnetic Reynolds number. We consider thermal ionisation, cosmic ray ionisation and recombination. In Section~\\ref{disc} we choose a fiducial disc model and find the dead zone and resulting accretion rate through the disc for varying metallicity and critical magnetic Reynolds number. We investigate the extent to which assuming a constant surface density in the active layer can approximate this and how sensitive the dead zone is to the metallicity. In Section~\\ref{an} we find analytical approximations to the surface density in the active layer of the disc. ",
        "conclusions": "A typical assumption used in time-dependent simulations of accretion discs with dead zones is that the surface density in the MRI active layer above the dead zone is constant with radius. However, when the dead zone is identified by the value of the critical magnetic Reynolds number including the effects of recombination, the active layer surface density is generally found to increase with radius. The constant layer is best reproduced with a low critical magnetic Reynolds number.  However, MHD simulations suggest the critical magnetic Reynolds number may be much higher.  For higher critical magnetic Reynolds numbers, $Re_{\\rm M,crit}\\gtrsim 100$ we have found an analytical fit to the active surface density (equations~\\ref{fit} and~\\ref{fit2}) that will be a useful approximation in future time-dependent calculations.  The dead zone structure is very sensitive to the value of the critical magnetic Reynolds number, as seen in Fig.~\\ref{layer}, but the value is still uncertain and needs to be clarified in further work. However, this is complicated by the fact that the $\\alpha$ parameter in the viscosity is also uncertain.   The metallicity variation between our galaxy, the LMC and the SMC is not significant enough to affect the expected size of the dead zone in a disc (ignoring possible differences in dust abundances).  When comparing accretion discs in our galaxy with those in the LMC and SMC for example, the high metallicity limit will be appropriate. In order to explain the observed difference in disc lifetimes with metallicity \\citep[e.g.][]{demarchi10,demarchi11} some other effect must be taken into account."
    },
    "1112/1112.2154_arXiv.txt": {
        "abstract": "A candidate diffuse stellar substructure was previously reported in the halo of the nearby dwarf starburst galaxy NGC~4449 by Karachentsev et al. We map and analyze this feature using a unique combination of deep integrated-light images from the Black Bird 0.5-meter telescope, and high-resolution wide-field images from the 8-meter Subaru telescope, which resolve the nebulosity into a stream of red giant branch stars, and confirm its physical association with NGC~4449. The properties of the stream imply a massive dwarf spheroidal progenitor, which after complete disruption will deposit an amount of stellar mass that is comparable to the existing stellar halo of the main galaxy. The ratio between luminosity or stellar-mass between the two galaxies is $\\sim$~1:50, while the indirectly measured dynamical mass-ratio, when including dark matter, may be $\\sim$~1:10--1:5. This system may thus represent a ``stealth'' merger, where an infalling satellite galaxy is nearly undetectable by conventional means, yet has a substantial dynamical influence on its host galaxy. This singular discovery also suggests that satellite accretion can play a significant role in building up the stellar halos of low-mass galaxies, and possibly in triggering their starbursts. ",
        "introduction": "A fundamental characteristic of the modern cold dark matter ($\\Lambda$CDM) cosmology \\citep{2010gfe..book.....M} is that galaxies assemble hierarchically under the influence of gravity, continually accreting smaller DM halos up until the present day. If those halos contain stars, then they will be visible as satellites around their host galaxies, eventually disrupting through tidal forces into distinct streams and shells before phase-mixing into obscurity \\citep{2010MNRAS.406..744C}. This picture seems to explain the existence of satellites and substructures observed around massive galaxies \\citep{1966ApJS...14....1A,1988ApJ...328...88S,2009Natur.461...66M,2010AJ....140..962M}. However, quantitative confirmation of this aspect of $\\Lambda$CDM has been more elusive, with lingering doubts provoked by small-scale substructure observations \\citep{2011arXiv1104.2929L,2011arXiv1111.2048B,2011arXiv1111.6609F}.   There has been relatively little work on substructure and merging in the halos around low-mass, ``dwarf'' galaxies. Many cases of extended stellar halos around dwarfs have been identified observationally \\citep{2009MNRAS.395.1455S}, but it is not clear if these stars were accreted, or formed in-situ. Star formation in dwarfs is thought to occur in stochastic episodes \\citep{2009ARA&A..47..371T,2011ApJ...739....5W}, which could be triggered by accretion events.  An iconic galaxy in this context is NGC~4449, a dwarf irregular in the field that has been studied intensively as one of the strongest galaxy-wide starbursts in the nearby universe. Its absolute magnitude of $M_V=-18.6$ makes it an LMC-analogue (and not formally a dwarf by some definitions), but with a much higher star formation rate. It is strongly suspected to have recently interacted with another galaxy based on various signatures including peculiar kinematics in its cold gas and HII regions \\citep{1986AJ.....92.1278H,1998ApJ...495L..47H}, but the nature of this interaction is unknown.  An elongated dwarf galaxy or stream candidate near NGC~4449 was first noticed by \\citet{2007AstL...33..512K} from Digitized Sky Survey (POSS-II) plates (object d1228+4358), and is visible in the Sloan Digital Sky Survey (SDSS)\\footnote{http://www.sdss.org/}. Here we present deep, wide-field optical imaging that supplies the definitive detection of this ongoing accretion event involving a smaller galaxy, leading to interesting implications about the evolution of this system and of dwarf galaxies in general.  \\newpage ",
        "conclusions": "We have detected and analyzed a stellar tidal stream in the halo of NGC~4449 which we interpret as the ongoing disruption of a dSph galaxy by a larger dwarf (an LMC/SMC analogue\\footnote{The LMC and SMC may also have a history of interaction, with a stellar mass-ratio of $\\sim$~1:15 \\citep{Bes10}.}). This appears to be the lowest-mass primary galaxy with a verified stellar stream.  We suggest some implications for galaxy evolution. It has been proposed that dSph's orbiting massive galaxies such as the Milky Way were ``pre-processed'' from gas-rich dwarfs by tidal effects within dwarf-galaxy groups \\citep{2009Natur.460..605D}. We may be witnessing such a transformation in-action, with the HI streams surrounding NGC~4449 representing additional tidal debris.  We also suspect it is not just a coincidence that such a novel stream was found first around one of the most intensely star-forming nearby galaxies. The accretion event may well be the starburst trigger. The period of elevated star formation appears to have started $\\sim$~0.5~Gyr ago \\citep{McQ10}, which is suggestively similar to the stream's $\\sim$~1--2~Gyr orbital period\\footnote{Given a projected apocentric radius of $R_{\\rm a}=$~13~kpc, a circular orbit provides a lower-limit for the period of $T=2\\pi R_{\\rm a}/v_{\\rm c}\\simeq$~1.3~Gyr.} (and to any process that is linked to the dynamical time on $\\sim$~30~kpc scales).  Are such accretion events frequent among other dwarf galaxies in recent epochs? We suspect that exact analogues to this stream are not very common, or they would have been noticed already in DSS/SDSS images. However, if the stream had been only a bit fainter, more diffuse, or at a larger radius, it could have been missed, and thus there may be many more dwarf-hosted stellar streams awaiting detection.  In theory, the history of DM halo assembly should be fairly scale-free, and $\\sim$~1:10 mergers are expected to be the most generally dominant contributors to mass growth \\citep{2008ApJ...683..597S}. It is also increasingly recognized that such relatively minor mergers can have important effects on the larger galaxies, such as inciting global disk instabilities \\citep{2011Natur.477..301P}.  If streams as in NGC~4449 are common in dwarfs, they re-ignite classic ideas about galaxy interactions triggering starbursts. Given the high rates of star formation in dwarf galaxies, it is natural to ask if satellites are responsible. Surveys along these lines have produced mixed results \\citep{2001A&A...371..806N,2004MNRAS.349..357B,2008MNRAS.385.1903L}, but until now, low-surface-brightness objects such as dSphs would have been missed.  Regardless of the implications for starbursts, dSph accretion appears to be an increasingly viable avenue for direct assembly of dwarf galaxies' stellar halos---as witnessed by NGC~4449, and by Fornax, which shows traces of swallowing an even smaller dSph \\citep{2005AJ....129.1443C}. Future observational determinations of dwarf stream frequency in combination with theoretical models may provide clues to the general substructure problem."
    },
    "1112/1112.2224_arXiv.txt": {
        "abstract": "We demonstrate feed-forward vibration isolation on a suspended Fabry-Perot interferometer using Wiener filtering and a variant of the common Least Mean Square (LMS) adaptive filter algorithm. We compare the experimental results with theoretical estimates of the cancellation efficiency. Using data from the recent LIGO Science Run, we also estimate the impact of this technique on full scale gravitational wave interferometers. In the future, we expect to use this technique to also remove acoustic, magnetic, and gravitational noise perturbations from the LIGO interferometers. This noise cancellation technique is simple enough to implement in standard laboratory environments and can be used to improve SNR for a variety of high precision experiments. ",
        "introduction": "The next generation of interferometers for gravitational-wave detection, including the Laser Interferometer Gravitational Wave Observatory (LIGO), will have unprecedented sensitivity to astrophysical events~\\cite{RPP:S5}. At low frequencies ($\\sim$10~Hz) it is likely that the displacement noise of the suspended mirrors will be limited at the 10$^{-20}$ $\\mathrm{m} / \\sqrt{\\mathrm{Hz}}$ level by fluctuations in the Newtonian gravitational forces~\\cite{Saulson:GGN, Thorne:GGN, VirgoNN}. The sources of the fluctuations are density perturbations in the environment (e.g. seismic and acoustic) and mechanical vibrations of the nearby experimental apparatus. While care will be taken to mitigate the sources of all of these fluctuations, further reductions of this Newtonian noise may be made by carefully measuring the source terms and subtracting them from the data stream (offline) or in hardware by applying cancellation forces to the mirror.   To demonstrate the efficacy of this technique, we demonstrate below that offline subtraction of seismic noise can be done using static Wiener filtering based on an array of seismic sensors. The online adaptive subtraction is also shown to approach the 'optimal' Wiener limit~\\cite{Haykin:Leaves}. Through nonlinear processes, seismic noise below 10 Hz has been shown to limit the performance of gravitational-wave detectors. This technique will prove to be of substantial value in reducing the non-stationarity of the detectors. ",
        "conclusions": "We have demonstrated the use of Wiener filter based feed-forward seismic noise reduction on a suspended interferometer. We have also implemented a stable, adaptive feed-forward system which has a performance approaching that of the optimal Wiener estimate. These techniques can be simply implemented in any general laboratory requiring vibration isolation using relatively low cost accelerometers and commodity computers and DSP software (e.g. Labview). These 'optimal' feed-forward schemes operate without having to know \\textit{a priori} the transfer function between the disturbance and the primary experiment; they can easily be reconfigured to adapt to new experimental setups.  Similar Wiener filter and LMS based techniques have been utilized in other experiments, both offline and online, for example isolating 2-mirror Fabry-Perot cavities from ground motion~\\cite{FPphasenoise}, reducing acoustic noise in oceanography settings~\\cite{oceanographyAcoustics} and in signal processing to decorrelate degenerate witness channels~\\cite{GramSchmidtDecorrelation}.  In the near future, we will work to use this scheme to reduce the noise in multiple degrees of freedom of the full interferometer. It is clear that this technique can also be applied to remove other sources of environmental noise (e.g. acoustic, magnetic, electronic, etc.). For each noise source in the gravitational wave band, we will inject software gravitational wave signals into the data stream and confirm that they are not distorted.  There will certainly be new challenges associated with each type of noise, but this seems promising as a method which can be employed to reduce the influence of environmental noise in a wide variety of experimental setups."
    },
    "1112/1112.0017_arXiv.txt": {
        "abstract": "We report the confirmation and characterization of a transiting gas giant planet orbiting the M dwarf KOI-254 every \\p\\ days, which was originally discovered by the \\kep\\ mission. We use radial velocity measurements, adaptive optics imaging and near infrared spectroscopy to confirm the planetary nature of the transit events. KOI-254\\,b is the first hot Jupiter discovered around an M-type dwarf star. We also present a new model-independent method of using broadband photometry to estimate the mass and metallicity of an M dwarf without relying on a direct distance measurement. Included in this methodology is a new photometric metallicity calibration based on $J-K$ colors. We use this technique to measure the physical properties of \\koi\\ and its planet. We measure a planet mass of \\msini~$ = \\mp$~\\mjup, radius $R_P = \\rp$~\\rjup\\ and semimajor axis $a = \\arel$~AU, based on our measured stellar mass $M_\\star = \\mstar$~\\msun\\ and radius $R_\\star = \\rstar$~\\rsun. We also find that the host star is metal-rich, which is consistent with the sample of M-type stars known to harbor giant planets. ",
        "introduction": "Of the hundreds of planets discovered outside of the solar system the class known as hot Jupiters, with periods $P<10$~days and $M_P \\gtrsim 0.2$~\\mjup, has provided much of our knowledge about the internal structures and compositions of exoplanets \\citep{charbonneau07,winn08r}. The crucial role hot Jupiters play in present-day exoplanetary science is not because they are common; the occurrence rate of hot Jupiters in the solar neighborhood is only about 1\\% \\citep{marcy05a}. Rather, close-in giant planets are the most easily detectable variety for Doppler and transit surveys owing to the large reflex motions and transit depths they induce. Once found to transit their host stars, their relatively large atmospheric cross-sections make them favorable targets for various follow-up studies to measure planet properties such as atmospheric composition, atmospheric temperature profiles, phase curves, and spin-orbit angles \\citep[e.g.][]{knutson09,crossfield10,mad11,albrecht11}.  While only one in every hundred Sun-like stars harbors a close-in gas giant planet, the occurrence of hot Jupiters is even further depressed around the Galaxy's most numerous denizens, the M dwarfs. This empirical finding has emerged from various Doppler surveys of M dwarfs, which have detected zero close-in giant planets among roughly 300 target stars with masses $M_\\star < 0.6$~\\msun, despite the ready detectability of the large Doppler amplitudes of these short-period giants \\citep{endl03, johnson10c}. This lack of hot Jupiters around M dwarfs, as well as the overall dearth of giant planets around low-mass stars, is likely due to the inefficiency of the planet-formation process within low-mass protoplanetary disks \\citep{laughlin04,ida05b,kennedy08}. This notion is further bolstered by the elevated occurrence of giant planets around stars more massive than the Sun \\citep{johnson07b,johnson10c}.  To date, no planet more massive than $0.1$~\\mjup\\ has been discovered around an M dwarf with a period less than 30 days\\footnote{{\\tt http://exoplanets.org}}. However, orbital companions at either mass extreme have been detected around M dwarfs, and some have been found to transit their stars. At the high-mass end, two transiting brown dwarfs have been found by the \\emph{MEarth} and \\kep\\ transit surveys, respectively \\citep{irwin10,johnson11b}. At the other end of the mass scale, transits of the RV-detected, short-period Neptune Gl\\,436\\,b were detected by ground-based follow-up photometry \\citep{butler04,gillon07}, and the \\emph{MEarth} survey discovered a transiting super earth orbiting the nearby, low-mass star GJ\\,1214 \\citep{charbonneau09}. However, no close-in planet with a mass intermediate to those previous discoveries has been found around an M dwarf. This leaves a gap in the companion-mass continuum that spans more than two orders of magnitude, from $0.1$~\\mjup\\ to 30~\\mjup.  One of the planet candidates discovered by the \\emph{Kepler} mission, KOI-254.01, provides an opportunity to fill in this missing region of parameter space. \\citet{borucki11} reported a transit depth of $\\delta = 3.9093$\\% and a period of $P = 2.455239$ days for the planet candidate. The large transit depth and lack of a visible secondary eclipse, together with the lack of a pixel shift in the stellar position during transit, reduces the \\emph{a priori} probability that the transit is a false-positive to $1.9$\\% \\citep{morton11b}. However, it is nonetheless desirable and prudent to acquire a full suite of follow-up observations to rule out a false positive. We also characterize the stellar characteristics of the low-mass host star in order to accurately and precisely measure the properties of the planet.  This contribution thus builds on the work of \\citet{muirhead11b}, hereafter Paper 1, who presented a spectroscopic analysis of all of the M dwarf KOIs announced by \\citet{borucki11}. We add to that work by including additional data from our multi-band, multi-site follow-up campaign for the specific case of the \\koi\\ planetary system. In \\S~\\ref{sec:observations} we present our observations of \\koi. In \\S~\\ref{sec:analysis} we describe our analysis of the \\emph{Kepler} light curves, KIC photometry, infrared spectra, and radial velocity orbit. We describe our results in \\S~\\ref{sec:lcfit} and conclude with a brief discussion and summary in \\S~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We report the detection of a short-period, Jupiter-mass planet that transits an early M-type dwarf star. The host star has mass $M_\\star = \\mstar$~\\msun\\ and a radius $R_\\star = \\rstar$~\\rsun, which we measure using a new method that draws upon broadband optical and NIR photometry, together with various empirical calibrations of M dwarf properties. Using our estimated stellar parameters, we find that the planet has a period of \\p~days, a mass $M_P = \\mp$~\\mjup, and a radius $R_P = \\rp$~\\rjup. At the planet's semimajor axis of \\arel~AU and given the star's estimated effective temperature of $T_{\\rm eff} = \\teff$~K, the planet has an equilibrium temperature of \\teq~K (see Table~\\ref{tab:starpars} for the full list of parameters and uncertainties).  Based on the rotation of the \\koi, we estimate an age of roughly 0.5~Gyr. Examination of the tabulated theoretical planetary radii computed by \\citet{fortney07} for an age of 0.3~Gyr and stellar insolation equal to that of \\koi\\,b shows that our measured planet radius is consistent with model expectations, but only for a core mass of $\\approx50$~\\mearth. For the 1~Gyr models a 25~\\mearth\\ core is necessary to reproduce the observed radius. This is a remarkably large core mass (heavy element content) and conforms well with the observed correlation between planetary heavy element content and host star metallicity \\citep{burrows07c,torres08,miller11}. This in turn agrees with observations to date that suggest that the formation of a Jupiter-mass planet around a low-mass star requires high metallicity \\citep{johnson09b, sl10,rojas10}.  As a Jupiter-mass planet with a period less than 10~days, \\koi\\,b is the sole example of a hot Jupiter orbiting an M dwarf. More than 300 stars have been monitored in various radial velocity surveys and a planet with $P < 30$~days has not been found around an M-type dwarf, despite the large expected Doppler signal. Indeed, a planet such as \\koi\\,b would be readily detectable in any of the past and present M dwarf Doppler surveys conducted over the past 15 years \\citep[see e.g.][]{endl03,johnson10a,bonfils11}. Given the difficulties associated with photometrically monitoring a large sample of M dwarfs, and the paucity of giant planets around low-mass stars, this fortuitous \\kep\\ detection may serve as our lone example of a hot Jupiter around an M dwarf for some time to come.  \\appendix"
    },
    "1112/1112.4478_arXiv.txt": {
        "abstract": "Redshift space distortion (RSD) is a powerful way of measuring the growth of structure and testing General Relativity, but it is limited by cosmic variance and the degeneracy between galaxy bias $b$ and the growth rate factor $f$. The cross-correlation of lensing shear with the galaxy density field can in principle measure $b$ in a manner free from cosmic variance limits, breaking the $f-b$ degeneracy and allowing inference of the matter power spectrum from the galaxy survey. We analyze the growth constraints from a realistic tomographic weak lensing photo-$z$ survey combined with a spectroscopic galaxy redshift survey {\\em over the same sky area}. For sky coverage $f_{\\rm sky}=0.5$, analysis of the transverse modes measures $b$ to 2--3\\% accuracy per $\\Delta z=0.1$ bin at $z<1$ when $\\sim 10$ galaxies arcmin$^{-2}$ are measured in the lensing survey and all halos with $M>M_{\\rm min}=10^{13}h^{-1}M_\\odot$ have spectra. For the gravitational growth parameter parameter $\\gamma$ ($f=\\Omega_m^{\\gamma})$, combining the lensing information with RSD analysis of non-transverse modes yields accuracy $\\sigma(\\gamma)\\approx0.01$. Adding lensing information to the RSD survey improves $\\sigma(\\gamma)$ by an amount equivalent to a $3\\times$ ($10\\times$) increase in RSD survey area when the spectroscopic survey extends down to halo mass $10^{13.5}$ ($10^{14}$) $h^{-1} M_\\odot$. We also find that the $\\sigma(\\gamma)$ of overlapping surveys is equivalent to that of surveys 1.5--2$\\times$ larger if they are separated on the sky.  This gain is greatest when the spectroscopic mass threshold is $10^{13}$--$10^{14}h^{-1} M_\\odot$, similar to LRG surveys.  The gain of overlapping surveys is reduced for very deep or very shallow spectroscopic surveys, but any practical surveys are more powerful when overlapped than when separated.  The gain of overlapped surveys is larger in the case when the primordial power spectrum normalization is uncertain by $>0.5\\%$. ",
        "introduction": "Measurement of the linear growth of structure of the Universe is essential in that the growth history reflects the nature of dark energy and the underlying gravity model \\citep[e.g.][]{Yamamoto2010}, i.e. whether dark energy is a cosmological constant, or is evolving with time, or if General Relativity (GR) is the correct gravity model that governs the evolution of the Universe. In the linear regime of GR, the growth of perturbations is scale independent. It can be parameterized as the linear growth function $G$, with $P(z)=G^2(z)P_{\\rm CMB}$, where $P(z)$ and $P_{\\rm CMB}$ are the matter density power spectra at redshift $z$ and at the epoch of recombination, respectively. $G(z)$ carries information about the amount of dark energy and dark matter. The growth rate factor $f \\equiv \\frac{\\partial \\ln G}{\\partial \\ln a}$, with $a$ being the scale factor, is another quantify of interest: $f$ can be well approximated as $f=\\Omega_m^{\\gamma}$, with $\\gamma$ in a narrow range near 0.55,  for a wide variety of dark-energy models in General Relativity \\citep{Peebles80,Lahav91,LinderCahn}.  A precise measure of $\\gamma$ therefore enables one to distinguish GR from alternative gravity models. In a braneworld type of modified gravity, for example, $\\gamma$ is different from $\\gamma_{\\rm GR}$ by more than $20\\%$ \\citep{LinderCahn}.  Precise measurements of $G$ and $f$ constrain dark energy and gravity, and Redshift Space Distortion (RSD) has been shown to be a powerful approach to perform this measurement \\citep[e.g.][]{Kaiser87, Cole94, Hamilton00, Peacock01,  Scoccimarro04, Guzzo08, Cabre09, Blake11}. This RSD measurement is, however, only precise in the linear regime.  At late epochs of the Universe, the linear regime (of the velocity field in particular,) is confined to very large scales, $k\\le 0.1h^{-1}{\\rm Mpc}^{-1}$. On these scales, the measurement is usually limited by sample variance, or cosmic variance---we do not have many independent perturbation modes for the measurement because of the finite survey volume observable in a given epoch.  Using multiple tracers of the density field, one can in principle evade sampling variance, and measure the linear growth of structure with unbounded accuracy \\citep{MS09,WSP,Gil-Marin10,BC2011}. The great benefit of multiple tracers is not realized, however, if only the clustering of galaxies is measured because of the following: using RSD, one can only measure $\\beta=f/b$ and the product $fG$. Without any prior knowledge of the galaxy bias, one can not constrain $f$ or $G$ independently. It has been shown by \\citet{BC2011}[BC11] that prior knowledge on galaxy bias significantly improves the constraint on the growth of structure in the case of single survey redshift bin.  In principle, galaxy bias can be measured by cross-correlating weak gravitational lensing convergence with galaxy clustering \\citep{Pen04}. This bias measurement is free of sample variance in the sense that the bias errors from a survey of a fixed number of modes can be reduced without limit as the lensing measurement noise and galaxy shot noise are decreased. This enable us to use large-scale modes that are well in the linear regime for high-accuracy measurement (see BC11). Combining a weak lensing survey (using photometric redshifts for source galaxies) with a spectroscopic redshift survey of lens galaxies can serve this purpose perfectly: the bias of the spectroscopic galaxies is measured by cross-correlation with the lensing signal in transverse modes, while the RSD analysis of the spectroscopic sample is conducted using non-transverse modes over the same volume.  Current and future large surveys are making possible the combination of spectroscopic redshift-space and lensing maps over common volumes. For example, the footprints of the upcoming {\\em Dark Energy Survey (DES)}\\footnote{http://www.darkenergysurvey.org} may overlap with that of an extended {\\em Baryon Oscillation Spectroscopic Survey (BOSS)}\\footnote{http://cosmology.lbl.gov/BOSS} survey near the equator. The future {\\em Euclid} space telescope \\citep{EUCLID} is designed to take spectra and images of galaxies at the same time.  In this work, we will explore the potential improvements in constraint of growth of structure and gravity from overlapping RSD and lensing surveys. This is an extension of BC11, where we consider the case of one single RSD redshift bin with an arbitrarily assigned prior on galaxy bias. In this work, we will consider the more realistic case of tomography using spectroscopic (RSD) and photometric (lensing) surveys covering common sky area, with both types of tracers divided into as many as 20 redshift bins. We will explore how well galaxy bias can be measured using the cross-correlation of galaxy shear and galaxy clustering in this realistic joint tomographic survey. The basic scheme is: \\begin{enumerate} \\item Conduct a galaxy redshift survey and a weak lensing (photo-$z$) survey over the same volume of the Universe. Split both galaxy samples into redshift bins.  \\item Optimally weight galaxies in each bin of the redshift survey to produce a mass density estimator with minimal stochasticity \\citep{Hamaus10,CBS11}.  Measure the 2-point shear-shear (from the lensing survey), and density-density (from the spectroscopic survey) correlations and the shear-density cross correlations between all $z$-bins.  Using these measurements of the covariance in transverse modes, constrain the bias $b$ of the spectroscopic galaxy density estimator and the mass power spectrum $P_m$ in each redshift bin. \\label{step2} \\item In each $z$-bin, split galaxies from the redshift survey into different bias bins. Perform multiple-tracer RSD measurement \\citep{MS09, BC2011} using the redshift-space density field of the binned galaxies. The $b$ and $P_m$ constraints derived from transverse modes in step \\ref{step2} are incorporated to break the $f-b$ degeneracy inherent to RSD, so that separate constraints on $G$ and $f$ can be achieved. Throughout the paper, we will suppress the latin index denoting redshift in equations that involve only a single redshift bin, such as the RSD Fisher matrix. We use Greek indices for galaxy bias bins. \\end{enumerate}  We use the Fisher matrix method to forecast growth constraints resulting from a model survey consisting of spectroscopic and lensing surveys covering a common $f_{\\rm sky}=0.5$ of the sky, reaching the depth of $z=2$.  We split both samples into 20 $z$-bins of width $\\Delta z=0.1$. We employ the halo model for our survey model, assuming that each halo above mass $M_{\\rm min}$ hosts one spectroscopic target galaxy. We set up our forecast methodology for lensing tomography in section \\ref{lensing} and for multi-tracer RSD in section \\ref{RSD}. We summarize our numerical results in section \\ref{result}, and conclude in section \\ref{conclusion}.  Unless noted otherwise, we assume a fiducial flat $\\Lambda$CDM cosmology with the following parameters: $\\Omega_m$=0.272, $\\Omega_{\\Lambda}=0.728$, $\\Omega_b=0.0455$, $\\sigma_8=0.807$, $n_s=0.961$, $H_0=70.2$. \\citep{WMAP7} ",
        "conclusions": "\\label{conclusion} We have shown from Fisher matrix forecast that the constraint on the growth of structure and gravity can be reduced percent-level or even sub-percent level by combining a spectroscopic redshift survey with a photo-$z$ weak lensing survey over the same volume. Whereas BC11 merely assumed that some measure of galaxy bias was available to add to RSD information, we verify here that a realistic tomographic weak lensing survey does in fact yield bias information sufficient to realize a substantial gain in accuracy on the growth parameter $\\gamma$.  Following the suggestions of \\citet{Pen04} we use the shear-galaxy cross-correlation to measure the galaxy bias in the transverse modes---a measurement which is free of sample variance---and apply it to the multi-tracer RSD analysis in a spectroscopic redshift survey \\citep{MS09,BC2011}. The combination of the two surveys make it possible to measure the linear growth function $G$ separately from its derivative $f=d\\ln G / d\\ln a$, whereas RSD alone can only measure the product $fG$.   The performance of multi-tracer RSD measurement depends on the spectroscopic survey depth, the range of galaxy biases in the sample, and the number of linear modes available. The performance of the shear$+$galaxy analysis on the transverse modes depends on: (1) the level of stochasticity between the galaxies and the projected mass, (2) the depth of the lensing survey, or shape noise. When combining two measurements together over the same volume, the results will depend on all those factors that affect each of the survey.  We have demonstrated that for the constraint of the $\\gamma$ parameter, combining two surveys is better than having each of them alone, roughly a factor 1.5 improvement (in survey-area terms) in the regime of likely feasible surveys: source density $N_{\\rm lens}\\approx 10\\,{\\rm arcmin}^{-2}$ in the lensing survey, and galaxy surveys complete for halos in the cluster or small-group range $M_{\\rm min}=10^{13}$--$10^{14}h^{-1} M_\\odot$, similar to LRG surveys. For $M_{\\rm min}>10^{13} h^{-1} M_\\odot$, the lensing$+$RSD survey has constraints many times more powerful than the RSD survey alone. The $\\gamma$ parameterization of growth predicts very little change at $z>1$; a different model for deviations from General Relativity could gain even more from the combination of lensing and RSD surveys.  Having prior constraints on the amplitude of the primordial power spectrum from the CMB is useful in general. Knowing $P_{\\rm CMB}$ to $0.5\\%$, easily within the statistical power of Planck, garners most of the $\\approx2\\times$ gain in accuracy on $\\gamma$ that is possible with perfect {\\it a priori} knowledge of $P_{\\rm CMB}$.  If $P_{\\rm CMB}$ is more poorly known, the gain of having overlapping surveys over the case of separate survey volume is increased.  During preparation of this paper, \\citet{Gaztanaga11} released very similar calculations of the benefit of coincident lensing and spectroscopic surveys.  Their assumed survey configurations and free parameterizations differ substantially from ours, so direct quantitative comparison is not possible. In the particular case of constraints on $\\gamma$, they find overlapping surveys reducing $\\sigma(\\gamma)$ by $\\approx 2.4\\times$ compared to seperated survey volumes, equivalent to a $6\\times$ increase in survey area in the language of our Figure~\\ref{ratio} where we find $\\approx 1.5\\times$ areal gain.  This is qualitatively consistent with our conclusion, but the origin of the substantial quantitative difference is difficult to ascertain given the different assumptions about survey characteristics.  \\citet{Gaztanaga11} also find substantial gains in accuracy of dark energy equation-of-state determination from overlapping surveys.  Our analysis holds this fixed so we would not have detected these gains; we plan to broaden our analysis to the case of unknown distance-redshift relations in the near future.  We notice that bias measurement can in principle also be measured using the same spectroscopic sample from the galaxy bispectrum \\citep{Simpson11}. If the same accuracy of bias can be obtained in this method as using lensing, one can simply use one spectroscopic redshift survey to obtain the same measurement, which might be another attractive survey strategy since no lensing survey is needed.  The lensing survey is, however, a straightforward measure of the galaxy bias, free of assumptions about perturbation theory, second-order bias, and other issues with the bispectrum.  Use of smaller-scale modes are attractive in the sense that one may gain many more modes from the same volume of survey. Growth test statistical accuracy improves rapidly with increasing $k_{\\rm max}$.  Non-linear effects in the density or velocity field and scale-dependent bias may, however, ruin the attempt to achieve percent-level constraint on parameters. Efforts have been made to improve RSD predictions for smaller-scale modes \\citep{Scoccimarro04, Jennings11, Hikage11,Tang11}, e.g. $k_{\\rm max}\\sim0.3$, though it is important that predictions be made for galaxies or halos rather than all mass particles in an $N$-body simulation \\citep{Jennings11}, e.g see \\citet{Reid11} for modeling of halos.  Better understanding of the non-linear biases of different tracers is required before one can confidently select the $k_{\\rm max}$ that admits the most accurate growth constraints."
    },
    "1112/1112.5681_arXiv.txt": {
        "abstract": "We present an analysis of the WISE photometric data for 829 stars in the Sco-Cen OB2 association, using the latest high-mass membership probabilities. We detect infrared excesses associated with 135 BAF-type stars, 99 of which are secure Sco-Cen members. There is a clear increase in excess fraction with membership probability, which can be fitted linearly. We infer that 41$\\pm$5\\% of Sco-Cen OB2 BAF stars to have excesses, while the field star excess fraction is consistent with zero. This is the first time that the probability of non-membership has been used in the calculation of excess fractions for young stars. We do not observe any significant change in excess fraction between the three subgroups. Within our sample, we have observed that B-type association members have a significantly smaller excess fraction than A and F-type association members. ",
        "introduction": "Many young ($\\sim$1-10\\,Myr) stars of all types, ranging from tens of solar masses down to the smallest brown dwarves and in all environments, are surrounded by circumstellar disks \\citep{strom89,lada2000,carpenter06}. A variety of observations have provided us with an overall timeline of disk evolution. The inner portion of the disk ($\\leq 1$\\,AU) dissipates by the age of 10\\,Myr in all but a small fraction of stars \\citep{mamajek04,silverstone06}. From the age of 10\\,Myr onwards, there is an observed decline in the 24$\\mu$ excess relative to the photosphere for stars of B to K type \\citep{gaspar09,rieke05}. It has been postulated that planetesimal stirring though stellar ages of 5 to 20\\,Myr could produce an increase in the strength of the $\\sim$24\\mic excess with age \\citep{kenyon08}, however the current data do not show a statistically significant increase \\citep{carpenter09}.  The Sco-Cen association and its three subgroups-Upper Scorpius (US), Upper Centaurus Lupus (UCL) and Lower Centaurus Crux (LCC)-provide three important constant-age samples within which to study circumstellar disks around young stars. The subgroups have ages of $\\sim$5, $\\sim$16, and $\\sim$17\\,Myr respectively, and are located less than 150\\,pc from the Sun \\citep{zeeuw99}. Previous studies \\citep{carpenter06,carpenter09} have investigated IRAC, IRS and Spitzer photometric data ranging from 4.5 to 70\\,\\mic in the US subgroup. They identified 54 stars with 24\\mic excesses in their sample of 205 targets and found that disks around BAF-type stars appear to be comprised of dusty debris, while disks associated with K and M-type stars are likely optically thick primordial disks which are remnants of the star formation process. The older Sco-Cen subgroups, UCL and LCC have received considerably less attention; only a handful of studies have observed small numbers of UCL and LCC stars \\citep{su06,carpenter09}. The recent study by \\citet{chen11} reports the detection of 41 new disks around F and G-type Sco-Cen stars.  New Bayesian membership probabilities for the high mass members (B to F-type) of the Sco-Cen association are now available \\citep{myfirstpaper}, and preliminary photometric data from the WISE mission have been released \\citep{wise10}. In this paper we present an analysis of the WISE photometry for Sco-Cen members to search for circumstellar disks in three constant-age samples. ",
        "conclusions": "We have analysed the available preliminary WISE photometry for the Sco-Cen stars of the \\citet{myfirstpaper} membership list and detected 135 22\\,\\mic excesses above the expected photosphere emission. We have used Sco-Cen membership probabilities to extrapolate an excess fraction for certain members, {\\color{red} and observe that there is no significant difference in excess fraction between the three association subgroups. This agrees with previous studies \\citep{carpenter09} and may be explained by the revised US age of $\\sim$11\\,Myr \\citep{pecaut11}.} Importantly we find that the excess fraction is significantly lower for the B-type stars in our sample compared to A and F-type association members, which is contrary to the trend seen by \\citet{carpenter09}. It is possible that the lack of a clear trend for the B-type sample could indicate that most of them are in fact young Sco-Cen members despite their kinematics. Another possible explanation relates to multiplicity. B-type stars have a significantly higher multiplicity fraction compared to later type stars \\citep{kouwenhoven05}. The presence of a companion can potentially truncate the inner regions of the circumstellar disk through resonances \\citep{lubow94}, producing a smaller disk fraction. This has been observed in a recent study of binaries in Taurus-Auriga, particularly with close ($<$40\\,AU) companions \\citep{kraus11arxiv}. Among the highest probability members in our sample ($>$90\\%) there are six close ($<$100\\,AU) multiple systems without disk detections and one close multiple system with a detected excess. A closer, more comprehensive, comparison between disk presence and multiplicity information for the Sco-Cen B-type stars may shed light on this issue, but is beyond the scope of this study."
    },
    "1112/1112.5224_arXiv.txt": {
        "abstract": "We use a model of an accretion flow coupled with an emergent flare to interpret the latest 1.3mm VLBI measurements for Sagittarius A*. The visibility data constrained the distances from the flare center to the black hole center as $d_{\\rm EW}\\lesssim20{\\rm R_g}$ and $d_{\\rm NS}\\lesssim80{\\rm R_g}$ in the East-West and North-South directions, respectively. If interpreted by the hot-spot model, the flare was preferred to pass in front of the black hole at a radius much larger than $d_{\\rm EW}$. If interpreted by the episodic jet launched from a nearly edge-on hot accretion flow, the flare was preferred to be ejected with $\\theta_{\\rm j}\\gtrsim40^\\circ$ off the black hole rotating axis.  This method can be generalized to help us understand future sub-millimeter VLBI observations, and study the millimeter/sub-millimeter variabilities in the vicinity of the Galactic Center supermassive black hole. ",
        "introduction": "\\label{intro}  Sagittarius A* (Sgr A*) is the best black hole candidate that can be studied by Very Long Baseline Interferometry (VLBI).   \\citet{Doeleman08} reported 1.3mm VLBI detections of Sgr A* on two baselines, namely ARO/SMT (Arizona) - CARMA (California) and ARO/SMT - JCMT (Hawaii).  These data were fitted by a circular Gaussian component with a full-width of half maximum (FWHM) of $\\sim43\\mu{\\rm as}$, reaching the event-horizon scale of the Galactic Center supermassive black hole.  Several groups soon interpreted these data independently by an accretion flow around a rotating black hole with a high inclination angle \\citep{Broderick09, Huang09a, Huang09b, Dexter09}.  Recently, \\citet{Fish11} reported new measurements on all the three baselines of the ARO/SMT-CARMA-JCMT array over three nights: 2009 April 5-7 (Days 95-97).  The data on Day 95 and Day 96 show high consistency with the data in 2007 April by \\citet{Doeleman08}, fitted by a circular Gaussian component with FWHMs of $\\sim41\\mu{\\rm as}$ and $\\sim44\\mu{\\rm as}$, respectively. But the correlated flux density on all the baselines increased on Day 97.  The Gaussian model with a similar size but higher flux density is needed to fit these data.  However, they pointed out that some model parameters must be changed for Day 97 in order to fit the data on all the three days by a ring model.   This implies that an additional variable component might exist together with the accretion flow associated with a black hole shadow structure.  In this Letter, we introduce a method in Section 2 to constrain the position of the flaring component, as required to interpret the Day 97 data.   In Section 3, we hold general discussions on the resolution of the baselines, challenges to the flare models, and suggestions for future observations.  Summary is provided in Section 4.   Here we adopt black hole mass $M_{\\rm BH}=4.1\\times10^6M_\\odot$ \\citep{Ghez08} and distance of the Galactic Center $d_{\\rm GC}=8{\\rm kpc}$. Thus 1 milli-arcsecond (mas) angular size corresponds to a linear size at the Galactic Center of $200{\\rm R_g}$, where ${\\rm R_g}=GM_{\\rm BH}/c^2$ is the gravitational radius.   \\newpage ",
        "conclusions": "\\label{discuss}  We introduce a simple but useful method to understand the time-variable emission of Sgr A* detected by 1.3mm VLBI \\citep{Fish11}.  We interpret the data on Days 95-96 by an accretion flow in its quiescent state, and the increase of correlated flux density on all the three baselines on Day 97 by the emergence of an extra flare, modeled by a circular Gaussian spot.  The visibility measurements impose a strict constraint on the distance from the flare center to the black hole center in the EW direction ($65^\\circ$ east-of-north) as $d_{\\rm EW}\\lesssim20{\\rm R_g}$. They also place another relatively relaxed constraint on the distance in the NS direction ($25^\\circ$ west-of-north) as $d_{\\rm NS}\\lesssim80{\\rm R_g}$. In the rest of this Section, general discussions on various aspects will follow.   \\subsection{Resolution of Baselines}  The longest baseline length on the tracks of baseline III (ARO/SMT-JCMT) is $\\sim3.55{\\rm G}\\lambda$, corresponding to the highest resolution of $\\theta_{\\rm beam} \\sim70\\mu{\\rm as}$.  The apparent size of the flaring component assumed in this Letter is $\\lesssim66\\mu{\\rm as}$, slightly super-resolved by the current baseline.  Here, we follow the criterion provided in \\citet{Shen97} to judge the degree of resolution.  $\\theta_{\\rm LIM}$, the limit of source size that can be resolved, is defined as \\begin{eqnarray} &\\quad&\\theta_{\\rm LIM}\\quad=\\quad\\left[\\theta_{\\rm LIM}({\\rm Statistical})^4+\\theta_{\\rm LIM}({\\rm Systematic})^4\\right]^{1/4},\\nonumber \\end{eqnarray} where \\begin{eqnarray} &\\quad&\\theta_{\\rm LIM}({\\rm Statistical})\\quad=\\quad\\frac{0.53}{\\sqrt{\\rm SNR}}\\cdot\\theta_{\\rm beam}\\nonumber\\\\ &\\quad&\\theta_{\\rm LIM}({\\rm Systematic})\\quad=\\quad\\frac{0.53}{\\sqrt{|F_\\nu/\\Delta{F_\\nu}}|}\\cdot\\theta_{\\rm beam}. \\end{eqnarray} With a mean signal-to-noise ratio (SNR) of 5.5 on the ARO/SMT-JCMT baseline \\citep{Doeleman08} and $|\\Delta{F_\\nu}/F_\\nu|$ being the sum of the uncertainties of the flux density measurements $\\sim30\\%$, this limit is estimated to be $\\theta_{\\rm LIM}\\sim20\\mu{\\rm as}$.  Thus, we think the flare was resolved in an optimistic manner.   \\newpage  \\subsection{Challenges To The Flare Models} \\label{origin}   The origin of flares in Sgr A* is still controversial.  Various models can be put into two general classes, namely the whole change and the transient structure.  The whole change means enhancement in flux density in the whole emission region, caused by changes in physical quantities of the accretion flow, e.g. increase in heating coefficient or creation of power-law electrons \\citep[e.g.][]{Yuan03, Liu04}.  We can fit the Day 97 data by the Q Model itself with $\\chi^2_{\\rm dof}=1.07$, ${\\rm dof}=N_{\\rm data}=14$, if the total flux density is scaled up to $3$Jy.  This implies those models of whole change are plausible for the Day 97 data, i.e. no extra component is required for data interpretation.  This is consistent with the result of $d_{\\rm EW}\\lesssim20{\\rm R_g}$, indicating that the flare happened very close to, or even inside the emission region of accretion flow.  We are also interested in challenging models of the transient structure, which means an extra component emerged to contribute to structural variability.  Generally speaking, the Day 97 data do not preclude any flare model if the transient structure is located within the preferred flaring region.  A popular model used to interpret short-time variability of Sgr A* in millimeter and near-infrared bands is a hot-spot orbiting in Keplerian angular velocity \\citep[e.g.][]{Broderick05, Dovciak08}.  Such a hot-spot varying in hourly timescale is predicted to be detectable by sub-mm VLBI \\citep{Doeleman09}.   We assume that the flare on Day 97 was caused by a hot-spot orbiting at ${\\rm r_\\parallel}$,  the projective radius from the black hole center.  The normal direction of the orbital plane coincides with the projective rotating axis of the black hole.  According to earlier work, the position angle of Sgr A* is preferred to be in the range $(90^\\circ,180^\\circ)$ or $(-90^\\circ,0^\\circ)$, rather than in the range $(0^\\circ,90^\\circ)$ or $(-180^\\circ,-90^\\circ)$. We then obtain ${\\rm r_\\parallel<40R_g}$, as calculated from the constraints on the flaring region shown in the above Section.  The period of the Keplerian orbit is constrained as ${\\rm T_{Kep}(r=r_\\parallel ) < 8hr}$, which is insensitive to black hole spin.  The observational duration $\\rm T_{obs}$ was $\\sim2{\\rm hr}$ on Day 97,  i.e., ${\\rm T_{obs}>0.25{\\cdot}T_{Kep}(r=r_\\parallel)}$, implying that an apparent light-curve should be detected with a nearly edge-on disk assumed.  However, this is inconsistent with the observed sustaining high flux density.  Therefore, we would either exclude this model, or explain this inconsistency by two arguments.  One is that the measurements of the total flux density had rather large errors. The other is that the radius was shortened by projection effect, i.e. the hot-spot might be passing in front of the black hole at a radius ${\\rm r_\\perp>40R_g}$ with a component aligned with the line-of-sight.  We further generalize this model into a hot-spot moving with sub-Keplerian angular velocity and radial velocity comparable to the light speed, the same as the velocity of the background transonic accretion flow.  In this case, the hot-spot would quickly fall into the event horizon of the black hole and disappear.  We integrate the curve of radial velocity shown in \\citet{Huang09b} to calculate typical timescale of falling at specific starting radii ${\\rm r_\\parallel<40R_g}$.  The lifetime of the hot-spot is constrained as ${\\rm T_{fall}(r=r_\\parallel)<1hr}$,  i.e. ${\\rm T_{obs}>2{\\cdot}T_{fall}}$.  Similar to the Keplerian-rotating model, we would either exclude this model or interpret it as a special case in which the hot-spot started at a larger radius in front of the black hole with ${\\rm r_\\perp>60R_g}$, and plunged into the event horizon.  Furthermore, we consider a different model of episodic jets proposed by \\citet{Yuan09}, in which the flare is contributed by an episodic ejection of plasmoid from a hot accretion flow.  As they calculated,  a plasmoid with initial location of $\\sim10{\\rm R_g}$ to Sgr A*  can accelerate from rest to $\\sim0.8c$ in 35 min.  This predicts that during the 2h observation,  the plasmoid can travel $\\sim250\\rm R_g$ from the black hole center.  To take into account both the constraint on $d_{\\rm NS}$ and the projection effect by a disk with $i\\gtrsim60^\\circ$,  the plasmoid must be ejected in a direction with $\\theta_{\\rm j}\\gtrsim40^\\circ$ off the rotating axis of the black hole.  We wish to make a comment on a model of tidal disruption of asteroid explored by \\citet{Zubovas11}, in which the flare is related to asteroids rather than properties of hot accretion flow.  They predicted a small size for a flaring region $\\lesssim1{\\rm AU}$ or $\\lesssim25{\\rm R_g}$, which is roughly included in our constraints with boundaries of $d_{\\rm EW}$ and $d_{\\rm NS}$.  However, the timescale of an asteroid in a parabolic orbit around Sgr A* is estimated to be $\\lesssim1.5{\\rm hr}$, which is unfavored by the Day 97 flare.    \\newpage  \\subsection{Visibility Prediction For An Array of Five Stations}  \\begin{figure}[ht] \\vspace{-0mm} \\begin{center} \\vspace{-1cm} \\includegraphics[width=12cm]{f3a.eps} \\\\ \\vspace{-1.2cm} \\includegraphics[width=12cm]{f3b.eps} \\\\ \\vspace{-5mm}\\caption{ \\textit{Left Panels}: Non-scattered images of accretion flow coupled with a flare with different central shifts. Boundaries of preferred positions of the flare are shown in dashed white lines.  \\textit{Right Panels}: Visibility amplitudes and closure phases (APEX-ARO/SMT-LMT, in plus signs) corresponding to the left images. See text for more details. \\label{meas975sta}} \\end{center} \\end{figure}   The flaring region of Sgr A* can be constrained better if suitable baselines in the NS direction are available. Theoretically, the separation of $\\Delta S\\approx d_{\\rm NS}\\lesssim80{\\rm R_g}$ between two components, the quiescent accretion flow and the flare, will cause first null/valley point in the visibility profile at $L_0\\gtrsim0.25{\\rm G}\\lambda$ in the NS direction.  We include two additional stations, APEX in Chile and LMT in Mexico, which are promising for contributing to sub-mm VLBI observations in the near future.  The tracks of seven additional baselines, from short to long, ARO/SMT-LMT, CARMA-LMT, APEX-LMT, JCMT-LMT, APEX-ARO/SMT, APEX-CARMA, and APEX-JCMT,  cover from $0.5{\\rm G}\\lambda$ to $7{\\rm G}\\lambda$ in length close to the NS direction.  In an optimistic view, the separation of $\\Delta S\\approx d_{\\rm NS}\\lesssim40{\\rm R_g}$ could be detected if all the baselines work.  In the left panels of Fig.\\ref{meas975sta}, we choose two examples of flares, ${\\rm P_4}$ and ${\\rm P_5}$, both with ${\\rm S_p=0.95Jy}$ and ${\\rm D_p=6.2R_g}$, but with $(x_{\\rm {P4}},y_{\\rm {P4}})/{\\rm R_g}=(-2.04,4.9)$ and $(x_{\\rm {P5}},y_{\\rm {P5}})/{\\rm R_g}=(-16.32, 40.8)$, respectively.  The corresponding visibilities are shown in the right panels.  The image of the Q Model and ${\\rm P_4}$, with $\\Delta S\\approx5{\\rm R_g}$, causes the first valley point at $\\sim4{\\rm G}\\lambda$ on the baseline APEX-LMT.   The image of the Q Model and ${\\rm P_5}$, with $\\Delta S\\approx40{\\rm R_g}$, causes the first valley point at $\\sim0.5{\\rm G}\\lambda$ on the baseline ARO/SMT-LMT, which appears as a drop in amplitude between the baselines ARO/SMT-CARMA and ARO/SMT-LMT.  At the same time, closure phases by a station group including APEX and LMT, also imply the change to symmetry caused by the flare. For example, $|\\phi_{\\rm CL}|$ by the group of APEX - ARO/SMT - LMT are shown in blue plus signs in the right panels of Fig.\\ref{meas975sta}."
    },
    "1112/1112.0309_arXiv.txt": {
        "abstract": "{}{In the last 30 years, the existence of small and cool magnetic loops (height $\\lesssim8$~Mm, $T\\lesssim10^5$~K) has been  proposed and debated to explain the increase of the DEM (differential emission measure) towards the chromosphere.}{We present hydrodynamic simulations of low-lying cool loops to study their conditions of existence and stability, and their contribution to the transition region EUV output.}{We find that stable, quasi-static cool loops (with velocities $<1$~km/s) can be obtained under different and more realistic assumptions on the radiative losses function with respect to previous works. A mixture of the DEMs of these cool loops plus intermediate loops with temperatures between $10^5$ and $10^6$~K can reproduce the observed emission of the lower transition region at the critical turn-up temperature point ($T\\sim2\\times10^5$~K) and below $T=10^5$~K.}{} ",
        "introduction": "In the last 30 years, different theories have been proposed and debated to explain the origin of the solar EUV output at temperatures below 1~MK. The idea that the transition region (hereafter TR) emission originates from the bases of the hot large-scale coronal loops is not confirmed by the measured DEM (Differential Emission Measure) which is orders of magnitude larger than predicted \\citep{gabriel,athay}. The excess observed emission compared to the predictions of the traditional conduction-dominated TR picture has led, in particular, to the suggestion that much of the TR plasma is confined in relatively small and cool magnetic loops (height $\\lesssim8$~Mm, $T\\lesssim10^5$~K), strongly connected to the chromosphere, but thermally insulated from the corona \\citep{dowdy1,dowdy2,feldman1,feldman4}.  The existence of this class of magnetic loops, predicted for decades, remains far from being established. From the observational point of view, they are indeed very difficult to observe. As most of their emission is in the UV, traditional chromospheric diagnostics from ground, such as \\ion{H}{$\\alpha$}, only carry information about their lower boundary. On the other hand, the lack of spectral information with proper temperature sensitivity and spatial resolution from the current space-born EUV instruments capable of observing lines forming between $10^4$ and $10^5$~K has so far prevented any firm conclusions.  Some recent observations with the VAULT instrument \\citep[Very High Angular Ultraviolet Telescope,][]{vault} in the \\ion{H}{i} Ly-$\\alpha$ line at very high spatial resolution, showed loop-like structures with estimated temperatures and densities ($T=10^4-3\\times10^4$~K, $P=0.1-0.3$~dyne~cm$^{-2}$) that could be appropriate for the low-temperature end of cool loops \\citep{patsourakos,vourlidas}. While such interpretation has been debated \\citep{judge}, those observations, together with the persistence of the problem of the excess emission in EUV lines below $10^5$~K, call for further investigations on the matter.  The general properties of static cool loops were first discussed by \\citet{hood} and then studied more in detail from a theoretical point of view by \\citet{an86}, who demonstrated analytically that a mixture of static cool loops with different temperatures can account for the observed lower transition region DEM. This class of loops has specific characteristics compared to the better studied coronal loops. While large-scale hot coronal loops have dimensions of the order of $\\sim10$~Mm, temperatures higher than $10^6$~K, and obey the ``static'' scaling laws described by \\citet[hereafter RTV]{rtv}, cool loops are low-lying (estimated heights of the order of 1.1--5~Mm), nearly isobaric, and with maximum temperature below $10^5$~K. Another major difference with respect to classical coronal loops is in the role played by the three terms appearing in the hydrodynamical equation of plasma energy conservation: the heating rate, the conductive flux, and the radiative losses. Static cool loops are in approximate balance between the heating rate and radiative losses. Thus, inversely to coronal loops, the conductive flux plays a negligible role.  \\cite{cally}, following the study of \\citet{an86}, treated in more detail the stability, structure and evolution of these static cool loops, both analytically and numerically. They reveal, in particular, that the shape of the radiative losses function, $\\Lambda(T)$, poses restrictive conditions on the existence of static cool loops. In the case of a realistic radiative losses function with the presence of the \\ion{H}{i} Ly-$\\alpha$ losses peak around $T=2\\times10^4$~K (e.g., \\citealt{chianti} or \\citealt{colgan}, dotted and long-dashed line in Fig.~\\ref{fig:radlos}, respectively), strictly static cool loops do not exist (see the next sections for details). It is important to note that their analysis is limited to a minimum temperature of log$T=4.29$~K, so that they cannot consider the existence of loops with maximum temperature below this value and cannot deal with Ly-$\\alpha$ losses properly (even in the thin losses case). They obtain static cool loops numerically with maximum temperature around $T=8\\times10^4$~K, assuming a $T^3$ dependence of the radiative losses function on the temperature below $T=10^5$~K (dashed line in Fig.~\\ref{fig:radlos}). They conclude, however, that this temperature is a factor 2--3 too low to explain the observed low-TR emission, even if they do not show any calculated DEM function to support these conclusions. Moreover, they state that ``the steep $T^3$ radiative losses function even if so beneficial for the existence and stability of cool loops rests on shaky foundations''.  After the cited work of \\citet{cally}, our paper is the first attempt to study numerically the conditions of existence of small and cool loops (height $\\lesssim5$~Mm, $T\\lesssim10^5$~K) and, when they exist, their physical structure and their contribution to the TR emission. We show that stable, quasi-static cool loops (with plasma velocities $<1$~km/s) can be obtained through hydrodynamic simulations also assuming different forms of $\\Lambda(T)$, with respect to the work of \\cite{cally}.  \\citet{peter,dem} made the first successful attempt to reproduce the shape of the DEM curve quantitatively and qualitatively, even at temperatures below $\\log T=5.3$~K. They used a forward model in which they synthesize spectra from three-dimensional MHD simulations of the whole Sun atmosphere, from the chromosphere to the corona. We think that their results do not exclude ours. Their simulations may include structures that could be related to the kind of loops we are studying. However, the loops we describe in this paper would be covered by only very few resolution elements in their simulation, and in any case resolving the gradients and the dynamics of the relevant quantities in our loop models would require a much higher resolution. Therefore, we regard our study as complementary to the kind of large-scale simulations by \\citet{dem}.  In the course of our study of loops with the properties just described, we have also found low-lying quasi-static loops with temperatures in the range $10^5-10^6$~K. Following one of the latest loop classifications \\citep{realesolo} we should refer to these loops also as ``cool loops''. In order to avoid confusion we will refer to them as ``intermediate temperature loops''. This class of loops was first studied by \\citet{foukal} and observed by \\citet{brekke}. From observations, they are generally detected in the UV lines and appear to be steady for long times even if more variable and dynamic then hot coronal loops, probably due to the presence of substantial flows \\citep{brekke,digiorgio,realesolo}. Their estimated densities are in the range $10^9-10^{10}$~cm$^{-3}$ \\citep[e.g.,][]{brown}. Even if observed intermediate temperature loops are often related to flows we decided to report however in this paper our findings on their existence (at least analytically) in a quasi-static state. Studying the role of flows in the conditions of their existence would mean introducing a new dimension in the parameter space to be explored, and this is beyond the scope of this paper, concentrated on quasi-static cool loops. Anyway, the quasi-static intermediate temperature loops we find have different physical properties with respect to observed loops with that temperature and this may be correlated with flows.  The paper is structured as follows: in Sec.~\\ref{sec:model}, we describe the numerical model, and the different radiative losses functions adopted in this work are introduced. In Sec.~\\ref{sec:results}, the hydrodynamic simulations and the different loops obtained (cool and intermediate temperature loops) are presented and their properties discussed and analyzed. Section~\\ref{sec:dem}, in particular, is dedicated to the calculated DEMs of these loops and to the comparison with the observed one. Finally, in the conclusions (Sec.~\\ref{sec:concl}), the role of the cool and intermediate temperature loops in the solar atmosphere and the possibility to observe them with current instrumentation will is treated.  \\begin{figure*} \\centering \\includegraphics[clip=true,width=13cm]{fig1_articolo.eps} \\caption{Radiative losses functions considered in this work. Solid line: power-law segments function, equal to $~T^2$ (for log$T<4.95$~K) and $~T^{-1}$ (for log$T>4.95$~K); solid line plus star symbols: a peak mimicking the H Ly-$\\alpha$ losses has been added to the previous function; dashed line: from \\citet{cook}, but with a $T^3$ dependence below $T=10^5$~K; long-dashed line: from \\citet{colgan}; dotted line: from \\citet{chianti}; dot-dashed line: from \\citet{chianti} without the H contribution.} \\label{fig:radlos} \\end{figure*} ",
        "conclusions": "\\label{sec:concl}  We have studied the conditions of existence and stability of cool loops with $T\\lesssim10^5$~K through hydrodynamic simulations, finding that it is possible to obtain quasi-static (velocities lower than $1$~km/s) cool loops, as predicted by \\citet{an86}, stable over hours or more. These loops have been obtained by using different dependences of the radiative losses function on the temperature with respect to the work of \\citet{cally}. We obtained stable quasi-static loops even in conditions judged prohibitive by the same authors on the basis of an analysis of strictly static and stationary loop equations. We examine and discuss the quasi-static solutions we have found, and show that their existence is due indeed to the small departures from static conditions, i.e.\\ to the presence of a small but non-zero conductive flux and velocities, as well as to rather stringent constraints on the pressure gradients. In fact, for low-temperature loops, the requirement of nearly constant pressure implies that these loops can exist only if they are limited to small heights above the chromosphere (a fraction of Mm). We also show that the presence of the peak due to the Ly-$\\alpha$ losses in the radiative losses function does not preclude the existence of cool loops. Moreover, we analyze the contributions of cool loops to the TR DEM, showing that the emission of these kind of loops can account for the observed DEM at $T<10^5$~K, if they were uniformly distributed.  The cool loops found cannot be related with any of the observations present in the literature because of their dimensions and especially of their small heights. We find cool loops with lengths of 5--10~Mm, but with very low heights (in the range $10^{-2}-10^{-1}$~Mm). Loops so low in the solar atmosphere (effectively embedded in the chromosphere) cannot be visible in the Ly-$\\alpha$ line; moreover, if compared with the observations of \\citet{patsourakos}, our loops typically have smaller pressures with respect to their estimates.  Indeed, the shape of the cool loops we have found could give information on the orientation of the magnetic field in the transition region, and in particular on its inclination. Since, as already pointed out, these loops are very shallow, the magnetic flux tube should emerge at transition region temperatures with a quite small angle with respect to the Sun's surface and observations should reveal a predominance of horizontal magnetic field direction at their height levels, assuming that these loops are almost everywhere.  We have also obtained quasi-static loops with maximum temperature in the range $2\\times10^5-10^6$~K, using a realistic radiative losses function. These loops are smaller with respect to coronal loops but have different characteristics compared to the static cool loops proposed by \\citet{an86} and others. These loops in principle could be observed with current telescopes, but in order to resolve them in all their temperature extension, we would need multi-temperature observations, i.e. different UV lines formed at temperatures between $10^5-10^6$~K with resolution of at least 1''. Loops 17 and 18 have the maximum temperature around $2.5\\times10^5$~K that is the upper limit of the region in which \\citet{feldman1} located the ``UFS (unresolved fine structures)''. We find that these intermediate temperature loops follow the scaling laws for coronal loops contrary to results of previous works based on the observational data \\citep[e.g.,][]{brown}. The loops obtained have indeed small pressures that make their parameters obey the RTV scaling laws, but these pressures are 1--2 orders of magnitudes lower than the ones estimated from the observations \\citep{brown}. Observed intermediate temperature loops are usually associated with flows while we are studying possible solutions to the hydrodynamic equations in quasi-static conditions without considering their role. The small values of the pressure we have found for these loops are the only possible in order to find solutions to the hydrodynamic equations in quasi-static conditions. Due to their small pressures, these loops obey to the scaling laws for coronal loops. The presence of substantial flows could influence their pressures.  We are not able to reproduce the DEMs derived from observations with only one set of parameters (a single loop). A similar problem has already been reported in the paper of \\citet{susino} but for coronal loops. We find instead that a combination of cool and intermediate temperature loops, in particular precisely due to their computed pressures, can give a DEM with a shape not too far from observed. This of course does not preclude the possibility that some additional physical processes, not included in our simulations (like that proposed by \\citealt{judge1} or by \\citealt{klimchuk1} and \\citealt{depontieu}, or dynamic structures like cooling, heating loops and spicules), must be taking place in the transition region and contribute to explain the discrepancies with the observed DEM structure. Due to their dimensions, the emission of the quasi-static loops found can be seen as a diffuse component of the transition region.  In our simulations, we have considered only the case of constant heating rate and we have already found quasi-static cool loops in conditions not allowed from strictly static and stationary hydro-dynamical equations. More realistic assumptions could make obtaining stable, quasi-static cool loops even easier. In particular, in this work we have shown how the shape of the radiative losses function below $10^5$~K is very important for the existence of cool loops. Because the radiative losses around log$T\\sim4.2$~K are dominated by the \\ion{H}{i} Ly-$\\alpha$, we plan to explore in a follow-up study the effect on the structure and stability of cool loops of a more realistic treatment of hydrogen radiation losses in the lower TR. To do so, it will be necessary to improve the models by introducing a full treatment of optically thick radiative losses and/or partial ionization, and therefore solving the radiative transfer equations and rate equations of hydrogen (non-local thermodynamic equilibrium). Our first attempt to simulate an optically thick radiative losses function by removing the hydrogen losses from a realistic radiative losses function seems to be promising. Without the hydrogen losses, the function $\\Lambda_{Dea-H}$ becomes lower and we are able to obtain quasi-static cool loops with $T<10^5$~K.  Taking into account finite cross sections for these loops could influence the details of the radiative transfer in the Ly-$\\alpha$ line (by reducing, for instance, the escape probability of Ly-$\\alpha$ photons). Predicting the effect of loop area expansion above the chromosphere is not so easy, although it is difficult to envision substantial expansion factors in such small loops.  Finally, the maximum temperature reachable by cool loops is limited only by the second maximum (or a change of slope) present in the $\\Lambda(T)$ since we have shown that the peak due to the Ly-$\\alpha$ losses is not a limit. The different radiative losses functions used have this maximum (or the slope change) at different temperatures. So, we underline the importance of knowing accurately the $\\Lambda(T)$ also at temperature around $10^5~K$ for the structure of the cool loops and of what we termed ``intermediate loops'', in analogy with the importance of the position and strength of the Ly-$\\alpha$ peak for cool loops."
    },
    "1112/1112.2140_arXiv.txt": {
        "abstract": "{ Since its discovery as an X-ray source with the $Einstein$ Observatory, the hot X-ray emitting interstellar medium of early-type galaxies has been studied intensively, taking advantage of observations of improving quality performed by the subsequent X-ray satellites $ROSAT$, $ASCA$, $Chandra$ and $XMM-Newton$, and comparing the observational results with extensive modeling by means of numerical simulations.  The hot medium originates from the ejecta produced by the normal stellar evolution, and during the galaxy lifetime it can be accumulated or expelled from the galaxy potential well.  The main features of the hot gas evolution are outlined here, focussing on the mass and energy input rates, the relationship between the hot gas flow and the main properties characterizing its host galaxy, the flow behavior on the nuclear and global galactic scales, and the sensitivity of the flow to major galaxy properties as the shape of the mass distribution and the mean rotation velocity of the stars.} ",
        "introduction": "\\label{intro}  X-ray observations performed during the late 1970's first revealed that early-type galaxies (ETGs) emit soft thermal X-rays and then host an interstellar medium (ISM), of a mass up to $\\sim 10^{10}\\,M_{\\odot}$, that had not been discovered previously at other frequencies because of its high temperature ($T\\sim 10^7$ K; see Fabbiano, this volume).  Such a medium was the long sought phase fed by stellar mass losses and predicted by the stellar evolution studies (e.g., \\citep{FG,Mat89}); in fact, both isolated ETGs and those in groups or clusters emit soft thermal X-rays \\cite{Fab89}, which provides simple evidence that the hot ISM is mostly indigenous rather than accreted intergalactic medium.  Besides solving a puzzle, this discovery opened the study of an important component of ETGs: it is a major ingredient of galactic evolution, see for example its role in feeding a central supermassive black hole and maintaining an activity cycle and starformation (e.g., \\citep{FC88,For05,Kor}, and Ciotti \\& Ostriker, this volume), or that in polluting the space surrounding ETGs with metals, via galactic outflows (\\cite{R93}, and Pipino, this volume), or finally in responding to environmental effects, as interaction with neighbors, stripping, sloshing, and conduction (e.g., \\citep{Acr,Sun07,K08}, and Sarazin, this volume).  The most striking X-ray observational feature of ETGs is the wide variation in their luminosity ($\\lx$) values, of $\\gsim 2$ orders of magnitude at any fixed galactic optical luminosity $L_B$, when $L_B>3\\times 10^{10}L_{B,\\odot}$ \\citep[][]{Fab89,OS,Mem, Kim10}. This feature cannot be explained by distance uncertainties, since a variation of the same size is present even in the distance-independent diagram of $\\lx/\\lb$ versus the central stellar velocity dispersion $\\sigma_c$ \\cite[e.g.,][]{Esk}.  A number of reasons have been proposed as responsible for this large variation, of environmental nature (see Sarazin, this volume) or linked to the possibility for the gas content to evolve substantially during the ETG lifetime. In this latter context, many studies investigated with numerical simulations the dynamical evolution of the hot ISM in ETGs \\citep[][]{SW,LM,Da90,Da91,C91,PC98,T09}.  In more recent times, the effect of feedback from a central supermassive black hole (MBH) has revealed as another potential contributor to the variation of the hot ISM luminosity (see Statler, and Ciotti \\& Ostriker, this volume).  Below I briefly review our current knowledge about the feeding and the energetics of the hot gas flows, concentrating on their dynamical state as a function of galactic mass and other major galaxy properties.  Section~\\ref{gasfl} gives an updated summary of the fundamental elements entering the problem, the mass and energy inputs to the flow; Sect.~\\ref{dec} presents the general case for the evolution of the flow, pointing out the different behavior on the global and nuclear galactic scales, and making also use of a representative ETG; the effect of a central MBH, acting as a gravitating point mass, and the expectations for accretion feedback, are presented in Sect.~\\ref{mbh}; finally, Sect.~\\ref{sensit} discusses the observed sensitivity of the flow to major galaxy properties as the flattening of the mass distribution, the mean rotational velocity of the stars, and the shape of the stellar profile. ",
        "conclusions": ""
    },
    "1112/1112.6249_arXiv.txt": {
        "abstract": "The Galactic centre -- as the closest galactic nucleus -- holds both intrinsic interest and possibly represents a useful analogue to star-burst nuclei which we can observe with orders of magnitude finer detail than these external systems. The environmental conditions in the GC -- here taken to mean the inner 200 pc in diameter of the Milky Way -- are extreme with respect to those typically encountered in the Galactic disk. The energy densities of the various GC ISM components are typically $\\sim$two orders of magnitude larger than those found locally and the star-formation rate density $\\sim$three orders of magnitude larger. Unusually within the Galaxy, the Galactic centre exhibits hard-spectrum, diffuse TeV (=$10^{12}$ eV) gamma-ray emission spatially coincident with the region's molecular gas. Recently the nuclei of local star-burst galaxies NGC 253 and M82 have also been detected in gamma-rays of such energies. We have embarked on an extended campaign of modelling the broadband (radio continuum to TeV gamma-ray), non- thermal signals received from the inner 200 pc of the Galaxy. On the basis of this modelling we find that star-formation and associated supernova activity is the ultimate driver of the region's non-thermal activity. This activity drives a large-scale wind of hot plasma and cosmic rays out of the GC. The wind advects the locally-accelerated cosmic rays quickly, before they can lose much energy in situ or penetrate into the densest molecular gas cores where star-formation occurs. The cosmic rays can, however, heat/ionize the lower density/warm $H_2$ phase enveloping the cores. On very large scales ($\\sim$10 kpc) the non-thermal signature of the escaping GC cosmic rays has probably been detected recently as the spectacular `Fermi bubbles' and corresponding `WMAP haze'. ",
        "introduction": "In part, this talk constitutes a general advertisement to the wider astronomical community concerning the utility of $\\gamma$-ray data: analysis of GeV and TeV data -- especially combined with data at other wavelengths -- sometimes allows one to constrain the processes occurring or conditions prevailing in astrophysical regions with much greater precision than when recourse is made only to data in `traditional' wavebands. I will illustrate this general philosophy through an analysis of the inner $\\sim 200$ pc in diameter of our own Galaxy (hereafter the `Galactic centre' or GC). To be sure, this region is important in its own right: as, by definition, the closest galactic nucleus our `up-close' view of the GC is of great interest for what it tells us -- by analogy -- about the activity of galactic nuclei in a cosmological context. But an understanding of the GC -- responsible for $\\sim$ 10\\% of the Galaxy's {\\it massive} star-formation (\\cite{Figer2004}) and, of course, the host of its supermassive black hole -- is a prerequisite for understanding the overall ecology of the Galaxy as I hope to make clear below.  The ISM conditions in the GC are certainly extreme in comparison with those encountered in the Galactic disk: the energy densities/pressures of the various ISM phases are  two orders of magnitude higher than in the Galaxy-at-large, typically 100's of eV per cc. The magnetic field amplitude, for instance, through the region is at least 50 $\\mu$G (\\cite{Crocker2010a}). Equally, the plasma pressures and temperatures, light field, and gas turbulence are all much higher than locally (and all roughly in equipartition with each other). All this is driven by star-formation which peaks sharply over the same inner $\\sim$200 pc region of the Galaxy to a star-formation-rate areal density {\\it three} orders of magnitude than the average value in the disk (e.g., \\cite{Crocker2011b}).  Extending this general picture that the conditions in the GC may more closely resemble those encountered in a starburst environment, my recent work with collaborators (\\cite{Crocker2010b,Crocker2011a,Crocker2011b})) has revealed a number of interesting facts about the GC: i) the star-formation (and resultant) supernova activity in the GC drives a `super'-wind out of the region (more correctly, an outflow rather than a wind as the ejected material does not reach escape velocity); ii) the GC outflow advects co-mingled plasma and non-thermal particle populations to large distance from the plane and -- I claim -- the non-thermal $\\gamma$-ray and microwave signatures of this process have recently been found; and iii)  despite the broad similarity to a star-burst alluded to above, the GC  is basically  in steady state: the current star-formation rate and consequent activity is typical for the time-averaged state of the GC over timescales approaching 10 Gyr. ",
        "conclusions": "\\begin{enumerate} \\item Our modelling shows that the GC environment is an analogue to that found in star-bursts to the extent that we can show that the energy-density in both thermal and non-thermal ISM components is $\\sim$2 orders of magnitude higher than typically encountered in the Galactic disk. \\item Another similarity to star-burst nuclei is the driving of a powerful outflow from the system; in the GC case, at least, this outflow does not reach escape velocity, however. \\item On the other hand, star-formation in the region seems to have been sustained at more-or-less the current rate for many billions of years: most star-formation that has taken place in the system is not really `bursty', though there are, of course, stochastic variations in the star-formation rate when one looks with sufficiently fine grain. \\item We find that cosmic rays do not penetrate into the cores of the giant molecular clouds. This is presumably because their short residence time in the region does not permit this. A contrasting situation where cosmic rays crucially modify ISM chemistry in star-forming regions has recently been claimed for star-bursts (\\cite{Papadopoulos2010}). \\item On the other hand, ionising collisions experienced by the cosmic rays may explaining the anomalously hot and high ionisation state inferred for the (low density)  envelope molecular phase identified in $H_3^+$ absorption studies (\\cite{Goto2008}). \\end{enumerate}  Finally we remark that all the high-energy activity identified above can be explained as a result, ultimately, of the power injected through the process of (massive) star-formation; the super-massive black hole is not {\\it required} to have much influx beyond the inner few pc. A final speculation, in fact, is that a significant consequence of the sustained, star-forming activity of the GC is to {\\it prevent} significant activity of the black hole."
    },
    "1112/1112.1972_arXiv.txt": {
        "abstract": "During its lifetime, the Far Ultraviolet Spectroscopic Explorer (\\textit{FUSE}) was used to observe 99 cataclysmic variables in 211 separate observations. Here, we present a survey of the moderate resolution (R$\\simeq$10,000), far-ultraviolet (905 -- 1188 \\AA), time-averaged \\textit{FUSE} spectra of cataclysmic variables (CVs).  The \\textit{FUSE} spectra are morphologically diverse.  They show contributions from the accretion disk, the disk chromosphere, disk outflows, and the white dwarf, but the relative contribution of each component varies widely as a function of CV subtype, orbital period and evolutionary state, inclination, mass accretion rate, and magnetic field strength of the white dwarf. The data reveal information about the structure, temperature, density and mass flow rates of the disk and disk winds, the temperature of the white dwarf and the effects of ongoing accretion on its structure, and the long-term response of the systems to disk outbursts. The complete atlas of time-averaged \\textit{FUSE} spectra of CVs are available at the Multimission Archive at Space Telescope Science Institute as a High Level Science Product. ",
        "introduction": "\\label{sec_intro}  The Far Ultraviolet Spectroscopic Explorer (\\textit{FUSE}) was a NASA-supported astronomical observatory that operated from 1999 to 2007, acquiring far-ultraviolet (FUV; 905 -- 1188 \\AA) spectra at moderate spectral resolution \\citep[R$\\simeq$10000][]{moos2000, sahnow2000}. \\textit{FUSE} was used to observe a variety of Galactic and extragalactic targets, resulting in numerous scientific contributions in the areas of cosmology, galaxy structure, star formation and evolution, and planet formation.  During its operational lifetime, \\textit{FUSE} was used to observe 99 different cataclysmic variables (CVs) over a total of 211 observations.  CVs are mass-exchanging binary systems in which a late-type donor star transfers mass to a white dwarf (WD). In most CVs, mass transfer occurs via an accretion disk around the WD. The properties of CVs are discussed in detail by \\citet{warner1995}. CVs are divided into a number of sub-classes.  First, they are distinguished between magnetic and non-magnetic systems. In magnetic CVs, the WD field strengths are large enough (B$\\gtrsim$100~kG) to partially or completely disrupt the accretion disk. In non-magnetic CVs, the systems are divided between those systems that show semi-regular outbursts in the accretion disk (dwarf novae; DN) and those that do not (novalikes; NLs).  CVs provide an excellent population of nearby, non-obscured targets in which to study the physics of mass accretion \\citep{warner1995}.  Accretion disks are ubiquitous in astrophysical accretion, from compact binaries and proto-stars to the the central engines of active galaxies. CVs provide a way to probe disk accretion in multiple configurations and under a variety of stimuli.  The far-ultraviolet (FUV) in particular is an excellent window through which to view the accretion disk, disk outflows, and the WD accretor in CVs \\citep[e.g.,][]{ladous1991,cordova1982,drew1985,gansicke2005}.  In high accretion rate systems (DN in outburst and NLs), the accretion disk can reach temperatures of up to $\\simeq$60,000~K \\citep{hartley2005} and the disk SED typically peaks in the FUV.  As a result, the FUV is a sensitive indicator of the inner disk structure and the mass accretion rate onto the WD. The FUV is also rich in spectral lines in a variety of species and ionization states that target the temperature and ionization structure of the disk, disk outflows, and the magnetically-channeled accretion column in magnetic CVs.  In low state systems (quiescent DN and some magnetic CVs), the FUV continuum is often dominated by the WD.  The spectrum of the WD contains clues to the evolutionary and accretion history of the binary.  In the FUV, the wealth of spectral lines allows for determinations of metal abundances on the WD surface, which often show anomalies indicative of CNO processing of the accreted material, an artifact of the binary evolution in the system. Moreover, the physical properties of the interaction region between the accretion disk and the WD in CVs (the boundary layer) remains poorly understood \\citep{ferland1982,godon2005,long2009}. At longer wavelengths it can be difficult to distinguish the spectrum of the WD from that of the disk, but observations at FUV wavelengths give the opportunity to determine the properties of the WD, the quiescent disk, and the boundary layer seperately.  Low-state CVs also show rich emission line spectra \\citep{ladous1990}. The emission lines are usually assumed to originate in a chromosphere above the disk, but the location and properties of the line-emitting region and the vertical structure of the accretion disk are not well determined.  FUV spectroscopy places important constraints on the properties of the line-emitting region.  Most of the analysis of CVs in the FUV to date has proceeded on a case-by-case basis. While intensive studies of individual systems are necessary, it is also important to complement that work with a broad survey of the FUV properties of CVs \\citep[e.g.,][]{verbunt1987,mauche1997, araujo2005, hamilton2007,sion2008}.  By discerning broad trends, we can place the behavior of individual systems in context and determine which characteristics of compact binaries determine their observational properties.  In quiescent systems, a survey of targets can facilitate separation of the WD component from the disk and chromospheric components and help to determine the structure and physical extent of the latter. For high state systems, observations of a wide range of systems allow us to determine the distribution of vertically extended material in accretion disks.  In all systems, a broad survey is necessary to analyze the relative FUV contributions of and trends in the properties of the accretion disk, winds, and the WD as a function of the viewing inclination, binary geometry, mass accretion rate, WD magnetic field strength, and evolutionary history.  To enable a broad survey of the FUV properties of CVs, we present a complete atlas of the \\textit{FUSE} observations of CVs. This is the first large survey of the UV properties of CVs since the IUE atlas of \\citet{ladous1990}, and the first to cover the FUV wavelengths below Ly$\\alpha$. The time-averaged spectra of all the observations are available at the Multimission Archive at STScI (MAST) as a High Level Science Product. (\\textit{FUSE} also observed a number of novae, supersoft X-ray sources, and X-ray binaries.  These are not included here.) In \\S~\\ref{sec_obs}, we summarize the observations and data calibration. In \\S~\\ref{sec_analysis}, we present the time-averaged spectra and examine the properties of CVs in the FUV.  Finally, we present concluding remarks in \\S~\\ref{sec_conclusion}. ",
        "conclusions": "\\label{sec_conclusion}  The \\textit{FUSE} observations of CVs have provided a rich data set that allows observers to probe the physics of accretion and CV evolution among a diverse array of targets. Here, we summarize some of the key characteristics of the \\textit{FUSE} CV spectra.  \\begin{enumerate}  \\item The DN quiescent spectra typically have spectra dominated by the white dwarf accretor, although many systems have strong emission line spectra wherein the WD component is obscured.  WD model fits to the spectra often show the need for additional components contributing to the FUV emission. These components have been interpreted as a hot, rapidly-spinning accretion belt on the surface of the WD, the boundary layer, the hot spot, and/or optically thin disk emission. Several of the systems show abundance anomalies on the WD indicative of previous CNO processing of the accreted material. The emission line dominated DN have FUV spectra that are challenging to model, as a result of our poor understanding of the properties of accretion disks at low accretion rates.  \\item The spectra of DN in outburst show continuum emission from the steady-state accretion disk along with strong, broad spectral features from a disk outflow, varying from absorption to emission as the binary inclination increases. Some systems do not show strong wind lines but instead have narrow, low-velocity absorption lines which may originate in vertically extended material in the outer disk. Fits of steady-state accretion disk models show higher accretion rates for systems above the period gap than below, while excess emission above the predicted model fluxes near the Lyman limit suggest the presence of an additional continuum source in the FUV.  \\item NL CVs are a broad class of objects and the \\textit{FUSE} spectra are similarly broad and diverse in their properties. Several of the low-inclination systems show spectra with broad, blueshifted absorption lines characteristic of strong disk winds.  However, many of the other NLs have relatively narrow absorption line spectra centered on or near their rest velocities which may indicate the presence of vertically-extended disk components but no strong outflow. The higher inclination NLs have emission line spectra from the disk wind as well as possible disk chromospheric components seen in absorption. The continuum emission in NLs is dominated by the accretion disk, although steady-state disk models often deviate from the observations, particularly at the bluest wavelengths.  \\item The \\textit{FUSE} spectra of IPs are characterized by strong emission lines and blue continuum emission. Model fits shows that the continuum emission comes from the bulk of the WD, the hot spots at the base of the magnetically-channeled accretion columns, and possibly the disk. Some of the IPs shows enhanced N and weak C emission lines indicative of CNO processed material in the accretion stream.  \\item The \\textit{FUSE} spectra of polars are relatively faint. They show strong emission lines and either blue continua or a continuum consistent with a cool WD, depending on the accretion rate. Two of the polars show the enhanced N and weak C seen in a few of the IP spectra.  \\end{enumerate}"
    },
    "1112/1112.1005_arXiv.txt": {
        "abstract": "The new spinor-unit field representation of the electromagnetism \\cite{Nash2010} (with quark and lepton sources) is integrated via minimal coupling with standard Einstein gravitation, to formulate a Lagrangian model of the very early universe. A completely new solution to the coupled Einstein-Maxwell equations, with sources, is derived. These equations are generalized somewhat, but not in a way that violates any physical principles.  The  solution of the coupled Euler-Lagrange field equations yields a scale factor $a(t)$ (comoving coordinates) that initially exponentially increases $N$ e-folds  from  $a(0) \\approx 0$ to $a_{1} = a(0) {e}^{N} $ ($N$ = 60 is illustrated), then exponentially decreases, then exponentially increases to $a_{1}$, and so on almost periodically. (Oscillatory cosmological models are not knew, and have been derived from string theory and loop quantum gravity.) It is not known if the scale factor escapes this periodic trap.  This model is noteworthy in several respects: $\\{1\\}$ All fundamental fields other than gravity are realized by spinor fields. $\\{2\\}$ A plausible connection between the \\emph{unit} field $\\mathbf{u}$ and the generalization of the photon wave function with a form of Dark Energy is described, and a simple natural scenario is outlined that allocates a  fraction of the total energy of the Universe  to this form of Dark Energy. $\\{3\\}$ A solution of an analog of the pure Einstein-Maxwell equations is found using an approach that is in marked contrast with the method followed to obtain a solution of the well known Friedmann  model of a radiation-dominated universe. ",
        "introduction": "\\label{sec:introduction}  \\subsection{Description of inflationary cosmology}  Inflationary cosmology  \\cite{Guth:1980zm,Linde:1981mu,Albrecht:1982wi} is a widely accepted and largely successful theory of the physics of the very early universe. Recent observational data from important experiments such as WMAP \\cite{Dunkley:2008ie,Komatsu:2008hk,0067-0049-192-2-18} and the Sloan Digital Sky Survey \\cite{PhysRevD.69.103501,0004-637X-633-2-560,PhysRevD.74.123507}, and others  support \\cite{Dunkley:2008ie,Komatsu:2008hk,Kinney:2008wy} the main predictions of inflation theory. Inflation employs a period of accelerated expansion at the beginning of the Universe to solve the monopole, horizon and flatness problems. The simplest models of inflation are based on the slow-roll of a scalar inflaton field. The simple slow-roll inflaton potential model does not have the  ``exit from inflation\" issue of the basic cosmological constant scenario, since the period of cosmic acceleration is followed by a radiation-dominated era, and then by a transient matter dominated phase \\cite{Starobinsky198099}. After some mean time period, the inflaton  excites % Standard Model degrees of freedom and the hot big bang commences. Quantum fluctuations during the early time evolution of the inflaton field generate primordial curvature fluctuation spectra \\cite{jhep052003013} that evolve into macroscopic perturbations whose time evolution leads to cosmological structure formation. Most models  predict a scale-invariant spectra of energy density perturbations and gravitational waves that act as seeds for structure formation. Moreover, the power spectrum of curvature fluctuations at horizon crossing during inflation may be related to the angular power spectrum of cosmic microwave background (CMB) fluctuations at recombination, and yield predictions of the anisotropies in the cosmic microwave background radiation.  Yet, to date, no inflaton-type field model has gained universal acceptance. In this paper we study a new cosmological model based on the coupled $\\{1\\}$ Einstein, $\\{2\\}$ electromagnetic (with quark and lepton sources) and the $\\{3\\}$ {\\em unit} field $\\mathbf{u}$   equations that exhibits an oscillatory form of inflation/deflation. This model  may also provide some insight into the nature of Dark-Energy. The photon field is represented according to \\cite{Nash2010} using a real bosonic spinor field, and gives rise to the exact conventional interacting local $U(1)-$invariant Maxwell equations after inflation has ended. ``Maxwell\" refers to the electromagnetic sector of the Standard Model. For simplicity, we consider only one generation of leptons and quarks. Also for simplicity, this model employs a flat ($k = 0$) universe Friedmann-Lema\\^{i}tre-Robertson-Walker (FLRW) line element with scale factor $a = a(t)$ (comoving coordinates) for the Einstein sector. However, we do not expect our qualitative results to change for non-zero, constant $k$, or with the inclusion of additional Standard Model degrees of freedom.  The solution of the coupled field equations yields a scale factor $a(t) \\equiv a[F(t)]$, where $F(t)$ is defined in Eq.[\\ref{eq:u}] below, that initially exponentially increases $N$ e-folds  from  $a(0) \\approx 0$ to $a_{1} = a(0) {e}^{N} $ ($N$ = 60 is illustrated; $N$ is strictly a function of initial conditions), then exponentially decreases to $a(0)$, then decreases to $0$, then exponentially increases to $a_{1}$, then exponentially decreases to $0$, and so on almost periodically. The energy density $\\varrho$, discussed below, determines the depth of the potential well of the periodic trap. It is not known if the scale factor escapes this periodic trap. The exponential change is with respect to  $F$. The exact dependence is given by Eq.[\\ref{firstIntegral}] and is  only approximately exponential for large values of $|F|$. Nevertheless, for brevity we use the descriptor ``exponential\". With respect to comoving time $t$ the change is accelerated. The  scale factor  makes the transition from $a(0)$ to $a_{max}$ in a finite comoving time that can be made as small as one chooses by changing initial conditions. Also, $a_{max}$ can be made as large as one wants and $a(0)$ as small by changing initial conditions.  Central to the new cosmological model is the unit field   $\\mathbf{u}$ \\cite{Nash2010}, which is a fundamental spinor field realization of the unit element of the split octonion algebra. \\emph{\\textbf{Throughout this paper we work with the components $u^a$ of $\\mathbf{u}$, rather than with  $\\mathbf{u}$, to avoid the complexities of the non-associativity of the split octonion algebra.}} With respect to a canonical spinor basis of the algebra, its eight components  $\\mathbf{u} \\leftrightarrow u^{a},\\;a=1,...,8,\\;$ comprise a fundamental spinor that transform under a basic spinor representation of $\\overline{SO(4,4)}$. In this paper the unit field is a section $u^{a} = u^{a}(x^{\\alpha})\\;\\alpha=1,...,4$ of a certain spacetime fiber bundle whose technical specification is given in Appendix 1. The unit field may be employed to construct a  local $U(1)-$invariant covariant derivative, enabling a basic (bosonic) spinor field representation of the electromagnetic sector of the Standard Model \\cite{Nash2010}.  In this paper the unit field is parameterized as \\begin{equation} \\label{eq:u} u^a = \\frac{ 1 }{\\sqrt{2}} \\left(0, e^{ -\\frac{F(t)}{2}} ,0,0,0, e^{ \\frac{F(t)}{2}} ,0,0\\right) , \\end{equation} where $F(t) \\equiv 0$ characterizes the Minkowski spacetime of Einstein's special relativity (but not uniquely). This Lagrangian model predicts that \\begin{equation} \\label{firstIntegral} {a(t)}^{3} \\cosh{[F(t)]} = {a(0)}^{3} \\cosh{[F(0)]} = \\textrm{constant}, \\end{equation} i.e., ${a(t)}^{3} \\cosh{[F(t)]}$ is an exact integral of the motion. Since $F(t) \\equiv 0$ characterizes the end of inflation, $\\frac{a(t_{\\textrm{end of inflation}})}{a(0)} =  {\\cosh{[F(0)]}}^{1/3}$, where $\\cosh{[F(t)]} = \\widetilde{\\mathbf{u}} \\mathbf{u} = {[u^{2}(t)]}^{2} + {[u^{6}(t)]}^{2}$, is solely determined by the initial value $F(0)$. (Here, the tilde denotes the transpose.)  \\subsubsection{Radiation}  The Segre classification of the energy-momentum tensor [and, in virtue of the Einstein field equations, of  the Ricci tensor] of a pure radiation field, which is a null Maxwell field, is the algebraic type  A3[(11,2)], with eigenvalue zero. Therefore the ``radiation\" that is modeled in the radiation-dominated  Friedmann equations cannot be a solution of the Einstein-Maxwell equations because the  Friedmann model stress-energy tensor is of type A1[(111),1], which is that of a non-tilted massive perfect fluid with timelike four-velocity. Type A1[(111),1] is also different from the types A1[(111,1)] ($\\Lambda$ cosmology), and A1[(11)(1,1)] (non-null Maxwell field). In the Friedmann universe radiation is an incoherent superposition of waves with random phases and polarizations. If one integrates the  stress-energy tensor  $T_{\\mu \\,\\nu}$ of this  electromagnetic  radiation over solid angle then one finds that the average  $\\big< T_{\\mu \\,\\nu} \\big>$  is of type A1[(111),1]. As an example, define the lightlike Minkowski vector $\\ell_{\\alpha}=(\\sin (\\theta ) \\cos (\\varphi ),\\sin (\\theta ) \\sin (\\varphi ),\\cos (\\theta ),1)$ and use the definition of  $d\\,T_{\\mu \\,\\nu}$ given on page 521 of \\cite{Weinberg1972}, which depends on the product $\\ell^{\\alpha}\\, \\ell_{\\beta}$. Then $ \\frac{1}{4 \\pi} \\int _0^{2 \\pi }\\int _0^{\\pi }\\,\\ell^{\\alpha}\\, \\ell_{\\beta} \\,\\sin (\\theta )\\,d\\theta d\\varphi =\\textrm{diag}\\;\\left(\\frac{1}{3},\\frac{1}{3},\\frac{1}{3},-1\\right) $. $\\big< T_{\\mu \\,\\nu} \\big>$  defines/depends-on  the radiation energy density $\\varrho(t)$ and pressure $p(t)$ functions, and appears in the Einstein equations $G_{\\mu \\,\\nu} = -2\\, \\big< T_{\\mu \\,\\nu} \\big>$ that yield the radiation-dominated  Friedmann equations. The equation of state relating the radiation pressure $p$ and energy density $\\varrho$, after the averaging has been performed, is $p = \\frac{1}{3}\\varrho$.  However, in this paper we solve an analog of the pure Einstein-Maxwell equations. We present a self-consistent model of gravity and radiation, with massive and electrically charged sources, that introduces the  energy density $\\varrho(t)$ and pressure $p(t)$  \\cite{PhysRev.55.364} functions in order to make the total energy momentum tensor homogeneous and isotropic (see Eq.[\\ref{eq:FLUID}]). There is no \\emph{a priori} equation of state relating $p(t)$ and $\\varrho(t)$ because  $p(t)$ is determined dynamically from the field equations (of course, $\\dot{\\varrho}(t)$ is always determined from the field equations).   It should emphasized that we do not ``second quantize\" the quantum fields of massive, electrically charged fermions. A field equation for a massive and electrically charged fermion, information-equivalent to the usual Dirac equation, is defined in terms of a real, eight component spinor field. The charge operator is not diagonal in this representation, as is well known from Schwinger's work with real eight component Fermion spinor fields  \\cite{Schwinger64}. But for simplicity, the real eight component Fermion spinor fields are not ``second quantized.\" This simplifying assumption may be justified by the fact that we study a temperature regime far above  the electron (resp., down quark) rest mass; therefore electrons and positrons (resp., quarks and anti-quarks) never annihilate efficiently.   This paper is organized as follows: in the next section we define our notation and conventions; technical details are mostly relegated to the Appendices. In Section \\ref{sec:Non-equilibriumbirth} the assumption that the universe is created in a non-equilibrium state is discussed, which is argued to imply that in order to make the total energy momentum tensor homogeneous and isotropic an additional term must be incorporated into the Lagrangian. Next the Lagrangian for this problem is defined. The Euler-Lagrange field equations are discussed in Section \\ref{sec:FE}. A representative numerical approximation of the solutions of the field equations for the FLRW scale factor  $a(t)$, for the split octonion unit spinor field $u^a(t)$ and for the electromagnetic split octonion spinor fields are presented. Lastly, we summarize our results. ",
        "conclusions": "\\label{sec:work}  We have studied a new cosmological model that is based on  a Lagrangian and couples Einstein, electromagnetic (with quark and lepton sources) and unit $\\mathbf{u}$ field  \\cite{Nash2010} sectors. We have focussed on a particular solution of coupled Euler-Lagrange field equations that employs a flat ($k = 0$) universe Friedmann-\\textrm{Lema{\\^ \\i}tre}-Robertson-Walker (FLRW) line element with scale factor $a = a(t)$ [comoving coordinates] for the Einstein sector. The solution of the field equations yields a scale factor $a(t)$ that oscillates almost periodically, and admits the first integral ${a(t)}^{3} \\cosh{[F(t)]} = {a(0)}^{3} \\cosh{[F(0)]} = \\textrm{constant} $. Since $F(t) \\equiv 0 \\, \\forall t > t_{\\textrm{end of inflation}}$ characterizes the end of inflation, $\\frac{a(t_{\\textrm{end of inflation}})}{a(0)} =  {\\cosh{[F(0)]}}^{1/3}$, where $\\cosh{[F(t)]} = \\widetilde{\\mathbf{u}} \\mathbf{u} = {[u^{2}(t)]}^{2} + {[u^{6}(t)]}^{2}$, is solely determined by the initial value $F(0)$.   The physics governing the  escape of $F$ and $a$  from the periodic trap has not been discovered. If the value of the energy density $\\varrho(t)$   decreases to the point where \\begin{displaymath} \\lim_{F(t) \\rightarrow 0} \\frac{\\varrho(t)}{{F(t)}^2} = 0 \\end{displaymath} then $\\lim_{F(t) \\rightarrow 0} \\dot{F}(t) = 0 $, and $F \\rightarrow 0$; $a$ is controlled by $F$ (because ${a}^{3} \\cosh{(F)} = \\textrm{constant}$) and will escape the periodic trap and very rapidly transition  to a constant $a(t_{\\textrm{end of inflation}}) \\approx a_{1},\\; (F(t)\\equiv 0 \\forall t > t_{\\textrm{end of inflation}})$. However the governing physics may have more similarity with quantum mechanical tunneling through a barrier than the  decay of the controlling parameter  $\\varrho(t)$ to zero.   \\subsection{Dark Photon} \\label{sec:Photon}  The photon field, carried by a type-2 real bosonic spinor field ${\\psi_{2}}^a$ in this model, is mapped to the electromagnetic four-vector potential $\\mathbb{A}_{\\alpha}$ through a nonlinear (in  $\\mathbf{u}$ )  interaction with the unit field: \\begin{equation} \\label{eqn:A^alph} \\mathbb{A}^{\\alpha} = \\frac{1}{ \\sqrt{ \\, \\widetilde{\\mathbf{u}}\\,\\sigma \\mathbf{u} } } \\, \\widetilde{ \\mathbf{u} } \\, \\sigma  \\;   \\overline{\\tau}^{\\alpha} \\, {\\Psi_{2}} . \\end{equation} $\\mathbb{A}^{\\alpha}$  is a true  $\\overline{SO(3,1; \\mathbb{R})}$ vector field since $\\Psi_2$ is  a type-2 spinor. In general, the type-2 real bosonic spinor photon field ${\\psi_{2}}^a$ is mapped to the electromagnetic spacetime potential $\\mathbb{A}_{B}$ through an interaction with the unit field that includes all eight components: \\begin{eqnarray} \\label{eqn:A^B} \\mathbb{A}^B &=& \\left( \\mathbb{A}_{1}, \\mathbb{A}_{2}, \\mathbb{A}_{3}, \\mathbb{A}_{4}, \\mathbb{C}_{1}, \\mathbb{C}_{2}, \\mathbb{C}_{3}, \\mathbb{C}_{4} \\right)_D \\mathbb{G}^{D B} = \\left( \\overrightarrow{\\mathbb{A}}, \\mathbb{A}_{4}, \\overrightarrow{\\mathbb{C}}, \\mathbb{C}_{4} \\right)_D \\mathbb{G}^{D B} \\nonumber \\\\ &=& \\frac{1}{ \\sqrt{ \\, \\widetilde{\\mathbf{u}}\\,\\sigma \\mathbf{u} } } \\, \\widetilde{ \\mathbf{u} } \\, \\sigma  \\;   \\overline{\\tau}^{B} \\, {\\Psi_{2}} . \\end{eqnarray}  The $(\\mathbb{C}_{1}, \\mathbb{C}_{2}, \\mathbb{C}_{3}, \\mathbb{C}_{4})$ are  ${SO(3,1; \\mathbb{R})}$ scalars; for notational convenience, the  $(\\mathbb{C}_{1}, \\mathbb{C}_{2}, \\mathbb{C}_{3})$ are collected into $\\overrightarrow{\\mathbb{C}}$.  For the representations employed in this paper \\begin{equation} \\mathbb{A}^{B} = \\frac{1}{\\sqrt{2}} \\left\\{ \\begin{array}{c} e^{-\\frac{F}{2}} {\\psi_2}^3+e^{\\frac{F}{2}} {\\psi_2}^7\\\\ e^{-\\frac{F}{2}} {\\psi_2}^4+e^{\\frac{F}{2}} {\\psi_2}^8\\\\ -e^{-\\frac{F}{2}} {\\psi_2}^1-e^{\\frac{F}{2}} {\\psi_2}^5\\\\ e^{-\\frac{F}{2}} {\\psi_2}^1-e^{\\frac{F}{2}} {\\psi_2}^5\\\\ e^{-\\frac{F}{2}} {\\psi_2}^4-e^{\\frac{F}{2}} {\\psi_2}^8\\\\ -e^{-\\frac{F}{2}} {\\psi_2}^3+e^{\\frac{F}{2}} {\\psi_2}^7\\\\ -e^{-\\frac{F}{2}} {\\psi_2}^6+e^{\\frac{F}{2}} {\\psi_2}^2\\\\ e^{-\\frac{F}{2}} {\\psi_2}^6+e^{\\frac{F}{2}} {\\psi_2}^2 \\end{array} \\right. \\end{equation}  In a Minkowski limit  ($F = 0$) the electric and magnetic field frame components $\\{\\mathbb{E}^{\\mu}, \\mathbb{B}^{\\mu}\\}$ are given (after inflation ends) by \\begin{equation} \\left( \\begin{array}{c} \\overrightarrow{\\mathbb{E}}  \\\\ \\mathbb{E}^{4} \\end{array} \\right) = \\left( \\begin{array}{c} - \\frac{1}{c} \\MPL{}{t} \\overrightarrow{\\mathbb{A}} +\\overrightarrow{\\nabla}  \\mathbb{A}_4 +\\overrightarrow{\\nabla} \\times \\overrightarrow{\\mathbb{C}}    \\\\ 0 \\end{array} \\right). \\label{eee} \\end{equation} and \\begin{equation} \\left( \\begin{array}{c} \\overrightarrow{\\mathbb{B}}\\\\ \\mathbb{B}^{4} \\end{array} \\right) = \\left( \\begin{array}{c} \\overrightarrow{\\nabla} \\times \\overrightarrow{\\mathbb{A}} + \\frac{1}{c} \\MPL{}{t} \\overrightarrow{\\mathbb{C}} +\\overrightarrow{\\nabla} \\mathbb{C}_4 \\\\ -\\overrightarrow{\\nabla} \\cdot  \\overrightarrow{\\mathbb{C}} - \\frac{1}{c} \\MPL{}{t}  \\mathbb{C}_4 \\end{array} \\right). \\label{bee} \\end{equation} We note that imposing $ \\overrightarrow{\\mathbb{C}} = \\overrightarrow{\\nabla}  \\mathfrak{C} $ and $ \\mathbb{C}_{4} = - \\frac{1}{c} \\frac{\\partial}{\\partial \\;t}  \\mathfrak{C} $, where $ \\overrightarrow{\\nabla}  \\cdot \\overrightarrow{\\nabla}  \\mathfrak{C} - \\frac{1}{c^2} \\frac{\\partial^2}{\\partial \\;t^2}  \\mathfrak{C}\\,=\\, \\textrm{const}\\times \\textrm{magnetic monopole charge density} \\, = 0 $ removes the $\\left( \\overrightarrow{\\mathbb{C}}, \\mathbb{C}_{4} \\right)$  dependence from  $\\mathbb{E}^{\\mu}$ and $\\mathbb{B}^{\\mu}$. If the $(\\mathbb{C}_{1}, \\mathbb{C}_{2}, \\mathbb{C}_{3}, \\mathbb{C}_{4})$ are decoupled from Standard Model degrees of freedom in this manner, then they comprise one set of \\emph{\\textbf{dark degrees of freedom}}. These dark degrees of freedom receive energy at $t = 0$, just like all of the other degrees of freedom. \\subsection{Fluctuations of the quantized unit field } \\label{sec:quf}   It is widely believed that primordial  curvature perturbations sourced by quantum field  fluctuations during the early time history of the universe evolve into macroscopic perturbations whose time evolution lead to cosmological structure formation. To calculate the quantum fluctuations the  bosonic spinor unit field, which satisfies   nonlinear field equations, must be quantized using an approach that guarantees the renormalizability of the quantum theory and the correct normalization of the unit field. Quantization should be performed after the $\\varrho$ and $p$ functions have been redefined and associated to a field, as in Eq.[\\ref{urhop}]. Since at comoving time t = 0 a non-vanishing average current drift velocity $\\overrightarrow{J}_D$ exists in this model there will be a  non-isotropic term in the total energy-momentum tensor. One  needs to perturb the FLRW metric  for $g_{\\alpha \\beta}$ simply  to obtain the classical solution, then carry out quantization. The investigation is ongoing."
    },
    "1112/1112.4983_arXiv.txt": {
        "abstract": "The technique of degree of randomness is used to model the correlations in sequences containing various subsignals and noise. Kolmogorov stochasticity parameter enables to quantify the randomness in number sequences and hence appears as an efficient tool to distinguish the signals. Numerical experiments for a broad class of composite signals of regular and random properties enable to obtain the qualitative and quantitative criteria for the behavior of the descriptor depending on the input parameters typical to  astrophysical signals. ",
        "introduction": "The correlations functions and the power spectra are common and efficient tools for the study of correlations in the signals. Astrophysical signals typically are superposition of various subsignals, regular and random, by features comparable to each other and of weaker ones, i.e. perturbations or the noise. The procedure of analysing of the needed signal or signals, their separation from the noise is a common problem while dealing with observations and measurements.  The Kolmogorov stochasticity parameter technique enables to quantify the randomness of sequences of number theory or dynamical systems \\cite{Kolm,Arnold,Arnold_UMN,Arnold_MMS,Arnold_FA}.  The technique of the degree of randomness has been applied to the Cosmic Microwave Background (CMB) temperature sky maps and to the X-ray flux data of the clusters of galaxies. The former data were those obtained by the Wilkinson Microwave Anisotropy Probe (WMAP) during 7-year observations \\cite{K,J}, while the X-ray data were obtained by XMM-Newton satellite providing a particularly accurate and complete sky survey (see \\cite{XMM,XMM1}).  In the case of CMB, the Kolmogorov function enabled to separate signals of different origin, e.g. the Galactic and non-Galactic ones, and to detect point sources in the CMB maps \\cite{G2009} (see Fig.1). Concerning the X-ray clusters, it was shown that their X-ray images do possess correlation in the pixelized flux data peculiar  to the gravitational potential of the galaxy clusters \\cite{GD2011}. This technique resembles the methods of the dynamical systems applied to nonlinear problems (e.g. \\cite{GP}).  A crucial step in these studies is the modeling and analysis of generated systems, which enables to reveal the behavior of the stochasticity parameter in the case of a given signal and then to consider the application of this technique for real signals \\cite{EPL2011}. Below we represent the results of numerical experiments for a broad class of signals.  Kolmogorov stochasticity parameter is introduced for a sequence $\\{X_1,X_2,\\dots,X_n\\}$ of  real random variable $X$ sorted in growing order  $X_1\\le X_2\\le\\dots\\le X_n$. Then the theoretical distribution function is \\cite{Kolm,Arnold} \\begin{equation} F(X) = n \\cdot(probability\\, of\\, the\\, event\\, x \\leq X). \\end{equation} the stochasticity parameter is defined as \\begin{equation}\\label{KSP} \\lambda_n=\\sqrt{n}\\ \\sup_x|F_n(x)-F(x)|\\, \\end{equation} where the empirical distribution function is $$F_n(X)= (number\\, of\\, the\\, elements\\, x_i\\, which\\, are\\, less\\, than\\, X),$$ and \\begin{equation} F_n(X)= \\left\\{ \\begin{array}{rl} 0, & X < x_1 \\\\ k / n, & x_k \\leq X < x_{k+1} \\\\ 1, & x_n \\leq X.\\\\ \\end{array} \\right. \\label{eq:empiricdistribution} \\end{equation}  Then for the limit \\begin{equation} \\lim_{n\\to\\infty}P\\{\\lambda_n\\le\\lambda\\}=\\Phi(\\lambda)\\ , \\end{equation} where \\begin{equation} \\Phi(\\lambda)=\\sum_{k=-\\infty}^{+\\infty}\\ (-1)^k\\ e^{-2k^2\\lambda^2},\\, \\Phi(0)=0,\\,   \\lambda>0\\ ,\\label{Phi} \\end{equation} exists at uniform convergence and independent on $F$.  \\begin{figure}[h] \\includegraphics[width=20pc]{phi.eps}\\hspace{2pc}% \\caption{\\label{fig:phi1}The Kolmogorov's function $\\Phi$ for the portion of the 7-year temperature CMB map obtained by WMAP.}% \\end{figure} ",
        "conclusions": ""
    },
    "1112/1112.4517_arXiv.txt": {
        "abstract": "We present a detailed analysis of the radio synchrotron emission induced by WIMP dark matter annihilations and decays in extragalactic halos. We compute intensity, angular correlation, and source counts and discuss the impact on the expected signals of dark matter clustering, as well as of other astrophysical uncertainties as magnetic fields and spatial diffusion. Bounds on dark matter microscopic properties are then derived, and, depending on the specific set of assumptions, they are competitive with constraints from other indirect dark matter searches. At GHz frequencies, dark matter sources can become a significant fraction of the total number of sources with brightness below the microJansky level.  We show that, at this level of fluxes (which are within the reach of the next--generation radio surveys), properties of the faint edge of differential source counts, as well as angular correlation data, can become an important probe for WIMPs. ",
        "introduction": "The nature of non--baryonic Dark Matter (DM) is one of the most puzzling mysteries of modern Cosmology (see e.g. \\cite{Bertone:2004pz,Bergstrom:2000pn} for reviews). A non--gravitational detection of dark matter would be a fundamental step toward the comprehension of its properties. This could be achieved in several ways, the most promising and popular methods are direct searches in underground experiments, DM searches at colliders, and indirect detection methods. The latter techniques are based on the detection of the products of DM annihilations or decays, notably neutrinos, antimatter and photons. Indirect searches are particularly promising if dark matter is in the form of Weakly Interacting Massive Particles (WIMPs). In the early Universe, these particles decouple non--relativistically from the thermal plasma and inherit the correct relic abundance for a thermally averaged annihilation cross section close to $\\langle\\sigma v\\rangle = 3\\times 10^{-26} \\mbox{ cm}^3\\mbox{s}^{-1}.$ Annihilations of WIMPs inside  galaxies are able to inject large amounts of relativistic electrons and positrons, which in turns interact with the galactic magnetic fields producing synchrotron radiation. This is a generic prediction of WIMP models, except for the peculiar case of WIMP candidates annihilating only into neutrinos. For galactic magnetic fields of the order of $\\mu$G and electrons energies below about 10 GeV, the synchrotron emission falls at frequencies around and below the GHz , i.e. in the radio band. Therefore, radio observations are appropriate tools for indirect WIMP searches and they have been extensively discussed in the context of WIMP annihilations inside our galaxy \\cite{Borriello:2008gy,Boehm:2010qt,Bertone:2002je,Boehm:2010kg,Crocker:2010gy,Regis:2008ij,Bergstrom:2008ag,Bertone:2008xr,Cumberbatch:2009ji,Dobler:2007wv,Linden:2010eu,Hooper:2007kb,Fornengo:2011iq}. In particular, low frequencies ($\\lesssim $ 1 GHz) data severely constraint WIMP masses $\\lesssim$ 10 GeV for leptonic annihilation modes \\cite{Fornengo:2011iq}.  Here we focus instead on the radio emission produced by DM annihilations or decays in extragalactic halos. A previous analysis along this direction can be found in Ref.~\\cite{Zhang:2008rs}. Recently, we have shown that for realistic assumptions on the DM clustering and magnetic fields and for a ``thermal'' annihilation cross--section (i.e. $(\\sigma v)=3\\times 10^{-26} \\mbox{ cm}^3\\mbox{s}^{-1}$), the DM signal could account for a significant fraction of the extragalactic--radio emission inferred by the ARCADE 2 Collaboration \\cite{Fornengo:2011cn}. This constitutes a further motivation to seriously consider DM searches with extragalactic radio observations.  In this work, we extend the studies in \\cite{Zhang:2008rs,Fornengo:2011cn} by extensively analyzing the impact of astrophysical and particle physics uncertainties on the DM extragalactic radio emission which is then compared to observational data. We focus on three observables: intensity, differential number counts of sources and angular correlations. By comparing predictions with current data, constraints on the DM annihilation/decay rate and mass are derived. We also compare the DM signals with (data--driven) expectations for astrophysical sources, such as radio loud active galactic nuclei and star forming galaxies, and show that source counts and angular correlation data could be particularly relevant to distinguish the DM contribution from the astrophysical ones.  It is interesting to note that cosmological radio emission induced by WIMP DM decreases very rapidly with redshift (so the main contributions come from structures at $z\\ll1$), and that DM halos can substantially contribute to the counts of sources at brightness below the $\\mu\\mbox{Jy}$ level. Therefore, a specific property of the extragalactic WIMP--induced source population is to peak at low redshift and low brightness. This hypothesis can be tested by telescopes with flux sensitivity at $\\mu\\mbox{Jy}$ level, which will be reached by upcoming radio experiments, in particular by the Square Kilometer Array (SKA).  The paper is organized as follows. In Section \\ref{sec:obsdata} we review the present status of radio observations, focusing on those particularly relevant for DM searches. We discuss data on intensity, angular correlations and counts of sources. The formalism needed to compute the DM and astrophysical signals is presented in Section \\ref{sec:thmod}. Section \\ref{sec:bench} shows our results for both DM and astrophysical models. Our estimates are first derived focusing on few specific DM benchmark cases while a detailed analysis about the impact of astrophysical and particle physics uncertainties is presented in Section \\ref{sec:disc}. The bounds on the microscopic properties of DM are computed in Section \\ref{sec:constraints}, and we derive our conclusions in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we presented an extensive analysis on the cosmological radio emission associated to WIMP dark matter. This emission is given by synchrotron radiation involving electrons and positrons injected by DM annihilations or decays in extragalactic structures. WIMP models with large branching ratios into light leptons lead to large $e^+/e^-$ yields with suppressed $\\pi^0$ production and in turn no bright $\\gamma$--ray emission from $\\pi^0$--decay. For those models the ratio between the DM signal and the background astrophysical emission produced by ``standard'' sources is significantly higher at radio frequencies than in the $\\gamma$--ray band.  For what concerns non--thermal emissions, radio telescopes have probably greater capabilities (i.e., sensitivity and angular resolution) as compared to telescopes at other wavelengths. This allows to study few different and complementary observables. We considered the total intensity, the angular correlation of source distribution and anisotropies, and the source number counts. We found that the emission induced by DM is mostly provided by a large number of small halos at redshift well below one. In general, WIMP models which provide a non--negligible contribution to the total extragalactic intensity take over other astrophysical sources in the number counts at low brightness ($S\\lesssim \\mu$Jy). There are also interesting prospects for the angular correlation at low multipoles and, again, very low brightness. However, it is important to stress that such conclusions are valid under the assumption that the source population given by DM--induced signals can be treated as a separate population with respect to other astrophysical sources, as, e.g., AGN and SFG. On the other hand, any astrophysical source is embedded in a (typically much fainter) DM halo which makes the separation of the DM population awkward.  We discussed the impact on the expected signals of dark matter clustering, magnetic fields, spatial diffusion, and dark matter microscopic properties. We found that bounds in the WIMP parameter space are uncertain by roughly three orders of magnitude. Under optimistic assumptions, they close on light WIMPs with thermal annihilation cross section, and, in general, are competitive with constraints from other indirect dark matter searches.  The extragalactic background induced by DM can account for the ARCADE excess~\\cite{Fornengo:2011cn} and the faint edge of differential source counts and angular correlation data, which will be provided by the next--generation radio surveys, can become an important probe for WIMPs."
    },
    "1112/1112.0428_arXiv.txt": {
        "abstract": "Molecular hydrogen (H$_2$) absorption features observed in the line-of-sight to Q2348$-$011 at redshift z$_{\\rm abs} \\simeq $2.426 are analysed for the purpose of detecting a possible variation of the proton-to-electron mass ratio $\\mu\\equiv m_p/m_e$.  By its structure Q2348$-$011 is the most complex analysed H$_2$ absorption system at high redshift so far, featuring at least seven distinctly visible molecular velocity components. The multiple velocity components associated with each transition of H$_2$ were modeled simultaneously by means of a comprehensive fitting method. The fiducial model resulted in  $\\Delta\\mu/\\mu =(-0.68\\pm2.78) \\times 10^{-5}$, showing no sign that $\\mu$ in this particular absorber is different from its current laboratory value. Although not as tight a constraint as other absorbers have recently provided, this result is consistent with the results from all previously analysed H$_2$-bearing sight-lines. Combining all such measurements yields a constraint of $|\\Delta\\mu/\\mu|$ $\\le 10^{-5}$ for the redshift range $z =$ (2--3). ",
        "introduction": "The spectrum of molecular hydrogen is known to be a testing ground to search for temporal variation of the proton-electron mass ratio, $\\mu\\equiv m_p/m_e$, on cosmological time-scales \\citep{thompson}. The spectral lines of the Lyman and Werner absorption bands of H$_2$ shift into an atmospheric transmission window ($\\lambda >3050$ $\\textrm\\AA$) for absorption systems at redshift $z \\ga 2$, and thus become observable with ground-based optical telescopes. Wavelengths of H$_2$ spectral lines as observed in high redshift absorption systems are compared to the wavelengths measured in the laboratory for which now very accurate calibrations exist (\\citealt{philip2004}; \\citealt{salumbides}; \\citealt{bailly}) as well as for the HD isotopologue (\\citealt{ivanov2008}; \\citealt{ivanov2010}), which is currently also detected at high redshift. Values for the sensitivity coefficients, $K_i$, expressing the shift of each spectral line as a result of a drifting $\\mu$, have been established to sufficient accuracy for H$_2$ (\\citealt{meshkov2006}; \\citealt{ubachs2007}). Hence a condition is accomplished that searches for cosmological $\\mu$ variation via H$_2$ spectra solely depend on the accuracy of the astrophysical data.  From the large number of damped Lyman-$\\alpha$ systems (DLAs) identified only some 20 are known to harbor detectable H$_2$ spectral features (\\citealt{ledoux2003}; \\citealt{srianand2005}; \\citealt{ubachs2007}). As was discussed recently \\citep{ubachs2011} only a handful of those high redshift absorption systems can deliver an H$_2$ spectrum of sufficient quality, \\emph{i.e.} a reasonable signal-to-noise (S/N) ratio on the continuum level to be obtained in reasonable data collection times at large-dish telescopes, and a large number of H$_2$ transitions ($\\ga$ 50). The high quality spectra of the Q0347$-$383 and Q0405$-$443 systems obtained with the ESO Very Large Telescope (VLT), equipped with the high resolution Ultraviolet and Visible Echelle Spectrograph (UVES) initially yielded an indication for a possible drift in $\\mu$, which is expressed via $\\Delta\\mu/\\mu \\equiv (\\mu_{z} - \\mu_{\\rm lab})/ \\mu_{\\rm lab}$ \\citep{reinhold2006}. A more sophisticated re-analysis of the same spectra by the so-called comprehensive fitting method, also invoking an improved thorium-argon (Th-Ar) calibration, reduced the initial finding of a 4-$\\sigma$ effect of a positive $\\Delta\\mu/\\mu$ to a 1.5$-$2.0-$\\sigma$ effect \\citep{king2008}. The same spectra were also re-analysed by \\citet{wendt2008} and by \\citet{thompson2009}, while \\citet{wendt2010} re-observed Q0347$-$383 with VLT-UVES to find a 1.5-$\\sigma$ effect. The highest quality system observed so far is J2123$-$0050. Observations with both the Keck Telescope \\citep{malec2010} and VLT \\citep{weerdenburg2011} were incorporated in a $\\mu$-variation analysis. The two studies yielded constraints on $\\Delta\\mu/\\mu$ which are in very good agreement with each other. The averaged constraint for J2123$-$0050 is $\\Delta\\mu/\\mu = (7.6\\pm3.5)\\times10^{-6}$. The system Q0528$-$250 has been at the heart of $\\mu$-variation analyses since the first attempts by \\citet{varshlovich1993} to derive a constraint on $\\Delta\\mu/\\mu$ from a low-resolution spectrum. \\citet{king2008} derived a tight constraint from a comprehensive fitting analysis of a VLT-spectrum of Q0528$-$250; a recent re-observation of this high-quality system, again at VLT, confirmed the previous conclusion of a tight constraint on $\\Delta\\mu/\\mu$ \\citep{king2011}. Based on the combined results derived from high redshift H$_2$ absorbers it can be concluded that $|\\Delta\\mu/\\mu| < 1 \\times 10^{-5}$ for redshifts $z=2-3$.  The present study presents a detailed analysis of the search for $\\mu$-variation of an absorber system toward the quasar Q2348$-$011. The physical conditions of this system had been investigated before based on data from the VLT (\\citealt{petitjean2006}; \\citealt{ledoux2006}; \\citealt{noterdaeme2007}). Q2348$-$011 was shown to be an exceptional sight-line, with a number of DLAs and sub-DLAs present, and a very complex velocity structure in the major H$_2$ absorbing system at $z \\simeq 2.426$. At least seven distinctly visible absorption features are associated with each H$_2$ transition in this system, while in other quasar sight-lines H$_2$ absorption (if detected) exhibits much simpler profiles consisting typically of one or two features. For the present study dedicated observations were performed at VLT focusing on improvement of S/N and on wavelength calibration of the spectrum. A preliminary $\\mu$-variation analysis of this absorber was presented in Ubachs et al. (2011). The current work improves upon the analysis but the constraint on $\\Delta/\\mu$ is of similar precision as presented previously. It is demonstrated that a complex velocity structure consisting of at least seven distinct H$_2$ velocity features can be disentangled by the method of comprehensive fitting (\\citealt{king2008}; \\citealt{malec2010}; \\citealt{king2011}).  \\section[]{Data}  The new spectrum of the quasar Q2348$-$011 was obtained with the UVES spectrograph on the ESO VLT on four consecutive nights (2008 August 18--21). Q2348$-$011 is not a particularly bright object (R = 18.31) so a long exposure time was required. Fifteen exposures were taken, making up a combined total of 19.25 hours of observation. After each science exposure, without any intervening grating resets, a Th-Ar exposure was recorded for calibration purposes. Seeing was in the range between 0$\\farcs$53 and 2$\\farcs$32. Besides the exposures noted above, the final spectrum incorporates data obtained from the ESO data archive (observations from 2003, October 29--30). They contribute additional 4.50 hours but do not have the individually taken Th-Ar calibration frames.  The raw 2D exposures were bias-corrected, flat-fielded, and the quasar flux extracted using the Common Pipeline Language version of the UVES pipeline\\footnote{ http://www.eso.org/observing/dfo/quality/UVES/pipeline/pipe $\\_$reduc.html.}. The wavelength calibration was established by extracting each Th-Ar frame using the object profile weights from its corresponding quasar exposure. Further details of the wavelength calibration procedure are described in \\citet{murphy2007}. The root-mean-square wavelength calibration residuals were typically $\\approx$70~m~s$^{-1}$ in the Lyman-$\\alpha$ forest portion of the spectrum, where the H$_2$ lines fall. The wavelength scale of each quasar exposure was corrected from air to vacuum, and to the heliocentric reference frame, and the flux was re-dispersed on to a common log-linear scale before being co-added, using {\\sevensize UVES$\\_$POPLER}\\footnote{Developed and maintained by M.~Murphy; see http://astronomy.swin.edu.au/$\\sim$mmurphy/UVES$\\_$popler.}, a code specifically written to re-disperse and combine reduced exposures from UVES.  The spectrum of Q2348$-$011 covers wavelengths 3572--9467 $\\textrm\\AA$, with gaps at 4520--4621 and 7505--7665 $\\textrm\\AA$. The final spectrum has a resolving power of  R $\\sim$ 57,000 and 63,000, and the log-linear dispersion is set to 2.5~km~s$^{-1}$ and 1.5~km~s$^{-1}$ in the blue ($<4000~{\\textrm \\AA}$) and in the red ($>4000~{\\textrm \\AA}$) parts, respectively. CCD pixels were binned by a factor of 2 in both spatial and spectral directions for all exposures in the blue, and no binning was applied for the red exposures. In the blue part where H$_2$ features are detected the average S/N ratio is about 25. ",
        "conclusions": "Molecular hydrogen features present at redshift $z_{\\rm abs} \\simeq 2.426$ in the line-of-sight to the quasar Q2348$-$011 were analysed to detect a possible variation of the proton-to-electron mass ratio on a cosmological time-scale. The constraint derived in the analysis is $\\Delta\\mu/\\mu =(-0.68\\pm2.78) \\times 10^{-5}$, showing no indication that $\\mu$ in this particular absorber is different from $\\mu$ measured in the laboratory. Although being less accurate than other recent H$_2$ constraints on $\\Delta\\mu/\\mu$, the result is consistent with those results, which show that $\\left|\\Delta\\mu/\\mu\\right| < 1\\times10^{-5}$ at $z=2$--3.  In the sample of known high-redshift H$_2$ absorbers, the structure of the absorber in the Q2348$-$011 spectrum is the most complex as it has 7 visually distinguishable velocity components, whilst 12 are justified statistically. The present analysis shows the applicability of the comprehensive fitting method in a case of such a complex structure. However, it also shows that in order to achieve a competitive result on $\\Delta\\mu/\\mu$, the H$_2$ absorber/spectrum selected for analysis must obey several conditions among which are no occurrence of strong additional absorbers and a sufficiently high signal-to-noise ratio."
    },
    "1112/1112.2936_arXiv.txt": {
        "abstract": "The nuclear symmetry energy $E_{sym}(\\rho )$ and its density slope $L(\\rho )$ can be decomposed analytically in terms of the single-nucleon potential in isospin asymmetric nuclear matter. Using three popular nuclear effective interaction models which have been extensively used in nuclear structure and reaction studies, namely, the isospin and momentum dependent MDI interaction model, the Skyrme Hartree-Fock approach and the Gogny Hartree-Fock approach, we analyze the contribution of different terms in the single-nucleon potential to the $E_{sym}(\\rho )$ and $L(\\rho )$. Our results show that the observed different density behaviors of $E_{sym}(\\rho )$ for different interactions are essentially due to the variation of the symmetry potential $% U_{sym,1}(\\rho ,k)$. Furthermore, we find that the contribution of the second-order symmetry potential $U_{sym,2}(\\rho ,k)$ to the $L(\\rho )$ generally cannot be neglected. Moreover, our results demonstrate that the magnitude of the $U_{sym,2}(\\rho ,k)$ is generally comparable with that of $% U_{sym,1}(\\rho ,k)$, indicating the second-order symmetry potential $% U_{sym,2}(\\rho ,k)$ may have significant corrections to the widely used Lane approximation to the single-nucleon potential in extremely neutron(proton)-rich nuclear matter. ",
        "introduction": "During the last decade, the nuclear symmetry energy $E_{sym}(\\rho )$ that essentially characterizes the isospin dependent part of the equation of state (EOS) of asymmetric nuclear matter has attracted much attention from different fields due to its multifaceted influences in nuclear physics and astrophysics \\cite{LiBA98,Dan02,Lat04,Ste05,Bar05,LCK08} as well as some interesting issues regarding possible new physics beyond the standard model \\cite{Hor01b,Sil05,Kra07,Wen09}. For example, the density slope $L$ of the symmetry energy at nuclear matter saturation density $\\rho _{0}$ has been shown to be important in determining several critical quantities such as the size of neutron-skin in heavy nuclei \\cite% {Bro00,Hor01,Fur02,Die03,Che05b,Tod05,Cen09,Sam09,Che10,Roc11}, location of neutron dripline \\cite{Kaz10}, core-crust transition density \\cite% {Hor01,Lat04,Ste05,Oya07,Kub07,XuJ09,New11} and gravitational binding energy \\cite{New09} of neutron stars. The symmetry energy may also have significant influence on gravitational wave emission from compact stars \\cite% {Kra08,Wen09b,Lin11,Gea11,Wen11b}. Furthermore, knowledge on the symmetry energy might be useful to understand the non-Newtonian gravity proposed in grand unified theories and to constrain properties of the neutral weakly coupled light spin-1 gauge $U$-boson originated from supersymmetric extensions of the standard model \\cite{Kri09,Wen09,Zha11,Wen11,Zhe11}.  In recent years, a great deal of experimental and theoretical efforts have been devoted to determining the density dependence of the symmetry energy \\cite{Bar05,LCK08}. Although significant progress has been made, large uncertainties on $E_{sym}(\\rho )$ still exist even around nuclear matter saturation density, e.g., while the value of $E_{sym}(\\rho _{0})$ is determined to be around $30\\pm 4$ MeV, mostly from analyzing nuclear masses, the extracted density slope $L$ scatters in a very large range from about $% 20 $ to $115$ MeV, depending on the observables and methods used in the studies \\cite{Che05,Tsa09,Cen09} (See, e.g., Refs. \\cite% {XuC10,Che10,Tsa11,Che11a,New11b} for a review of recent progress). To reduce the uncertainties on the constraints of $E_{sym}(\\rho _{0})$ and $L$ is thus of critical importance and remains a big challenge in the community.  So far, information on $E_{sym}(\\rho _{0})$ and its density slope $L$ is essentially obtained from theoretical model analysis on the experimental data of heavy ion collisions \\cite{Tsa04,Che05,Fam06,She07,Tsa09}, nuclear mass \\cite{Mye96,Liu10}, excitation energies of isobaric analog states \\cite% {Dan09}, pygmy dipole resonance of neutron-rich nuclei \\cite{Kli07,Car10}, isovector giant dipole resonance of neutron-rich nuclei \\cite{Tri08,Cao08}, and neutron skin thickness \\cite{Cen09,Che10}. In these theoretical models, an energy density functional with a number of parameters is usually assumed apriori, and the model parameters are then obtained from fitting experimental data and the empirical values of some physical quantities. Information on $E_{sym}(\\rho _{0})$ and $L$ is then extracted based on the obtained model parameters. Since all the phenomena/observables are in someway at least indirectly and qualitatively related to the $E_{sym}(\\rho _{0})$ and $L$, it is very useful to directly decompose $E_{sym}(\\rho _{0})$ and $L$ in terms of some relevant parts of the commonly used underlying nuclear effective interaction \\cite{XuC11}. This decomposition of the $% E_{sym}(\\rho _{0})$ and $L$ provides an important and physically more transparent approach to extract information about isospin dependence of strong interaction in nuclear medium from experiments and understand why the predicted symmetry energy from various models is so uncertain \\cite{XuC10b}.  In a recent work \\cite{XuC10}, based on the Hugenholtz--Van Hove (HVH) theorem \\cite{Hug58,Sat99}, it was indeed shown that both $E_{sym}(\\rho _{0}) $ and $L$ are completely determined by the single-nucleon potential in asymmetric nuclear matter which can be extracted from the nucleon global optical model potentials. In that work, the Lane approximation \\cite{Lan62} to the single-nucleon potential in asymmetric nuclear matter, i.e., $% U_{n/p}(\\rho ,\\delta ,k)\\approx U_{0}(\\rho ,k)\\pm U_{sym,1}(\\rho ,k)\\delta $% , has\\ been assumed and also the momentum independent nucleon effective mass has been used, therefore, the higher-order effects such as the contributions from the second-order symmetry potential $U_{sym,2}(\\rho ,k)$ and the momentum dependence of the nucleon effective mass\\ which can also contribute to $L$ have been neglected. So far, to our best knowledge, there is not any empirical or experimental information or even any theoretical predictions on\\ the second-order symmetry potential $U_{sym,2}(\\rho ,k)$. It is thus interesting and important to estimate the $U_{sym,2}(\\rho ,k)$ with some well-established theoretical models.  The main motivation of the present work is to evaluate the $U_{sym,2}(\\rho ,k)$ and estimate its contribution to $L$ based on several popular nuclear effective interaction models which have been extensively used in nuclear structure and reaction studies. Our results indicate that, although the momentum dependence of the nucleon effective mass might not be important, the second-order symmetry potential $U_{sym,2}(\\rho ,k)$\\ might have nonnegligible contribution to the $L$. Furthermore, we find that the magnitude of the $U_{sym,2}(\\rho ,k)$ is generally comparable with that of $% U_{sym,1}(\\rho ,k)$, indicating the second-order symmetry potential $% U_{sym,2}(\\rho ,k)$ may have significant corrections to the widely used Lane approximation $U_{n/p}(\\rho ,\\delta ,k)\\approx U_{0}(\\rho ,k)\\pm U_{sym,1}(\\rho ,k)\\delta $ for the single-nucleon potential in extremely neutron(proton)-rich nuclear matter, e.g., in neutron stars and the neutron-skin region of heavy nuclei. These results imply that it is important to extract experimentally information on $U_{sym,2}(\\rho ,k)$.  The paper is organized as follows. In Section \\ref{theory}, we briefly recall the definition of the symmetry energy and the symmetry potential in asymmetric nuclear matter, and then derive the explicit expressions for the single-nucleon potential decomposition of the symmetry energy and its density slope. The results and discussions are presented in Section \\ref% {results}. A summary is then given in Section \\ref{summary}. For completeness, the theoretical models used in the present paper are briefly described in the Appendix. ",
        "conclusions": "\\label{results}  In the following, we analyze the single-nucleon potential decomposition of $% E_{sym}(\\rho )$ and $L(\\rho )$ as well as the density and momentum dependence of $U_{sym,1}(\\rho ,k)$ and $U_{sym,2}(\\rho ,k)$ using three popular nuclear effective interaction models which have been extensively used in nuclear structure and reaction studies, namely, the isospin and momentum dependent MDI interaction model, the Skyrme Hartree-Fock approach, and the Gogny Hartree-Fock approach. One can find the details about these three models in the Appendix. A very useful feature of these models is that analytical expressions for many interesting physical quantities, such as the single-nucleon potential in asymmetric nuclear matter at zero temperature, can be obtained, and this makes our analysis and calculations physically transparent and very convenient.  For the MDI interaction, we use $3$ parameter sets, i.e., $x=-1$, $x=0$ and $% x=1$ \\cite{Che05}, which\\ give three different density dependence of the symmetry energy, namely, stiff one, moderate one, and soft one, respectively, and have been applied extensively in transport model simulations for heavy ion collisions. For the Skyrme interaction, we mainly use the famous SKM* \\cite{Bart82a} and SLy4 \\cite{Chaba98a} as well as the recently developed MSL0 \\cite{Che10}. In addition, a number of other Skyrme interactions are used for the single-nucleon potential decomposition of $% E_{sym}(\\rho )$ and $L(\\rho )$ at $\\rho _{0}$. For the Gogny interaction, we use the existing D1 \\cite{Dec80}, D1S \\cite{Berger91a}, D1N \\cite{Chap08a}, and D1M \\cite{Gor09a} which have been successfully applied in nuclear structure studies.  \\subsection{Single-nucleon potential decomposition of $E_{sym}(\\protect\\rho % ) $}  \\begin{figure}[tbh] \\centering \\includegraphics[width=8.6cm]{EsymMDI.eps} \\caption{(Color online) Density dependence of $E_{sym}(\\protect\\rho )$, $% E_{1}(\\protect\\rho )=\\frac{1}{3}\\frac{\\hbar ^{2}k_{F}^{2}}{2m_{0}^{\\ast }(% \\protect\\rho ,k_{F})}$ and $E_{2}(\\protect\\rho )=\\frac{1}{2}U_{sym,1}(% \\protect\\rho ,k_{F})$ in the MDI interaction with $x=-1$ (a), $0$ (b), and $% 1 $ (c).} \\label{EsymMDI} \\end{figure}  \\begin{figure}[tbh] \\centering \\includegraphics[width=8.6cm]{EsymSHF.eps} \\caption{(Color online) Same as Fig. \\protect\\ref{EsymMDI} but in the Skyrme Hartree-Fock\\ approach with MSL0 (a), SLy4 (b), and SKM* (c).} \\label{EsymSHF} \\end{figure}  \\begin{figure}[tbh] \\centering \\includegraphics[width=8.5cm]{EsymGogny.eps} \\caption{(Color online) Same as Fig. \\protect\\ref{EsymMDI} but in Gogny Hartree-Fock\\ approach with D1 (a), D1S (b), D1N (c), and D1M (d).} \\label{EsymGogny} \\end{figure}  In Fig. \\ref{EsymMDI}, Fig. \\ref{EsymSHF} and Fig. \\ref{EsymGogny}, we plot the density dependence of $E_{sym}(\\rho )$, $E_{1}(\\rho )$ and $E_{2}(\\rho )$ in the MDI interaction model, the Skyrme-Hartree-Fock approach, and the Gogny-Hartree-Fock approach, respectively. One can see from Fig. \\ref% {EsymMDI} that for different $x$ values, the $E_{1}(\\rho )$ displays the same density dependence while the $E_{2}(\\rho )$ exhibits very different density behaviors, indicating that the different density dependences of $% E_{sym}(\\rho )$ for $x=-1$, $0$ and $1$ are completely due to the different density dependence of $E_{2}(\\rho )$, i.e., the symmetry potential $% U_{sym,1}(\\rho ,k_{F})$. The similar behaviors can also be seen from Fig. % \\ref{EsymSHF} for the Skyrme-Hartree-Fock calculations and Fig. \\ref% {EsymGogny} for the Gogny-Hartree-Fock calculations.  Furthermore, one can see from Fig. \\ref{EsymMDI}, Fig. \\ref{EsymSHF} and Fig. \\ref{EsymGogny} that for all the interactions considered here, the $% E_{1}(\\rho )$ increases with the density and is always positive while $% E_{2}(\\rho )$ can increase or decrease with the density and even become negative at higher densities. These features can be understood since the $% E_{1}(\\rho )=\\frac{1}{3}\\frac{\\hbar ^{2}k_{F}^{2}}{2m_{0}^{\\ast }(\\rho ,k_{F})}$ is determined uniquely by the single-nucleon potential $U_{0}(\\rho ,k)$ in symmetric nuclear matter for which reliable information about its density and momentum dependence has already been obtained from heavy-ion collisions, see, e.g., Ref. \\cite{Dan02}, albeit there is still some room for further improvements, particularly at high momenta/densities, and the nuclear effective interactions are usually constructed to describe reasonably $U_{0}(\\rho ,k)$, especially around $\\rho _{0}$. However, on the contrary, the symmetry potential $U_{sym,1}(\\rho ,k_{F})$, which mainly reflects the isospin dependence of the nuclear effective interaction in nuclear medium, is still not very well determined, especially at high densities and momenta. In fact, it has been identified as the key quantity responsible for the uncertain high density behavior of the symmetry energy as stressed in Ref. \\cite{LCK08} (See also Refs. \\cite{XuC10b,XuC11}). These results show that the observed different density behaviors of $E_{sym}(\\rho ) $ for different interactions are essentially due to the variation of the symmetry potential $U_{sym,1}(\\rho ,k)$.  \\subsection{Single-nucleon potential decomposition of $L(\\protect\\rho )$}  In order to illustrate the single-nucleon potential decomposition of $L(\\rho )$, we show in Fig. \\ref{LMDI}, Fig. \\ref{LSHF} and Fig. \\ref{LGogny} the density dependence of $L(\\rho )$, $L_{1}(\\rho )$, $L_{2}(\\rho )$, $% L_{3}(\\rho )$, $L_{4}(\\rho )$ and $L_{5}(\\rho )$ in the MDI interaction model, the Skyrme-Hartree-Fock approach, and the Gogny-Hartree-Fock approach, respectively. It is seen that the $L_{1}(\\rho )$\\ displays the similar density dependence for all the interactions considered here, just like the $E_{1}(\\rho )$ shown in Fig. \\ref{EsymMDI}, Fig. \\ref{EsymSHF} and Fig. \\ref{EsymGogny} because of $L_{1}(\\rho )=2E_{1}(\\rho )$. The $% L_{2}(\\rho )$ is seen to contribute a small negative value to the $L(\\rho )$% , indicating that the momentum dependence of the nucleon effective mass is generally unimportant. In particular, one can see from Fig. \\ref{LSHF} that $% L_{2}(\\rho )$ vanishes for the Skyrme-Hartree-Fock calculations due to the fact that the nucleon effective mass is momentum independent for the zero-range Skyrme interaction considered here. The $L_{3}(\\rho )$ exhibits different density dependences for different interactions, reflecting the variation of $U_{sym,1}(\\rho ,k_{F})$ with density for different interactions. The $L_{4}(\\rho )$ represents the contribution of the momentum dependence of the symmetry potential to the $L(\\rho )$, and it displays different density dependence for different interactions and can be negative or positive, depending on the interaction used. \\begin{figure}[tbh] \\includegraphics[width=8.6cm]{LMDI.eps} \\caption{(Color online) Density dependence of $L(\\protect\\rho )$, $L_{1}(% \\protect\\rho )=\\frac{2}{3}\\frac{\\hbar ^{2}k_{F}^{2}}{2m_{0}^{\\ast }(\\protect% \\rho ,k_{F})}$, $L_{2}(\\protect\\rho )=-\\frac{1}{6}\\frac{\\hbar ^{2}k_{F}^{3}}{% {m_{0}^{\\ast }}^{2}(\\protect\\rho ,k_{F})}\\frac{\\partial m_{0}^{\\ast }(% \\protect\\rho ,k)}{\\partial k}|_{k_{F}}$, $L_{3}(\\protect\\rho )=\\frac{3}{2}% U_{sym,1}(\\protect\\rho ,k_{F})$, $L_{4}(\\protect\\rho )=\\frac{\\partial U_{sym,1}(\\protect\\rho ,k)}{\\partial k}|_{k_{F}}\\cdot k_{F}$, and $L_{5}(% \\protect\\rho )=3U_{sym,2}(\\protect\\rho ,k_{F})$ in the MDI interaction with $% x=-1$ (a), $0$ (b), and $1$ (c).} \\label{LMDI} \\end{figure}  \\begin{figure}[tbh] \\includegraphics[width=8.6cm]{LSHF.eps} \\caption{(Color online) Same as Fig. \\protect\\ref{LMDI} but in the Skyrme Hartree-Fock\\ approach with MSL0 (a), SLy4 (b), and SKM* (c).} \\label{LSHF} \\end{figure}  \\begin{figure}[tbh] \\includegraphics[width=8.5cm]{LGogny.eps} \\caption{(Color online) Same as Fig. \\protect\\ref{LMDI} but in Gogny Hartree-Fock\\ approach with D1 (a), D1S (b), D1N (c), and D1M (d).} \\label{LGogny} \\end{figure}  It is particularly interesting to analyze the $L_{5}(\\rho )$\\ since it reflects the higher-order $U_{sym,2}(\\rho ,k_{F})$ contribution to the $% L(\\rho )$ and has been neglected in the previous work \\cite{XuC10}. From Fig. \\ref{LMDI}, Fig. \\ref{LSHF} and Fig. \\ref{LGogny}, it is surprising to see that the $L_{5}(\\rho )$ may play an important role to determine the $% L(\\rho )$. In the MDI interaction with $x=-1$, the $L_{5}(\\rho )$ is always positive and increases rapidly with density while the opposite behavior is observed for the MDI interaction with $x=1$. For the MDI interaction with $% x=0$, the $L_{5}(\\rho )$ is positive and moderately increases with density. For the Skyrme-Hartree-Fock calculations, the $L_{5}(\\rho )$ can be positive or negative while it is always negative for the Gogny-Hartree-Fock calculations with D1, D1S, D1N, and D1M. These results indicate that generally the higher-order $U_{sym,2}(\\rho ,k_{F})$ contribution to the $% L(\\rho )$ cannot be neglected and the Lane potential approximation to the single-nucleon potential in asymmetric nuclear matter may cause significant error for the determination of $L(\\rho )$ from the single-nucleon potential decomposition.  \\begin{table*}[tbp] \\caption{The characteristic parameters $\\protect\\rho _{0}$ (fm$^{-3}$), $% E_{0}(\\protect\\rho _{0})$ (MeV), $E_{sym}(\\protect\\rho _{0})$ (MeV), $E_{1}$ (MeV), $E_{2}$ (MeV), $L$ (MeV), $L_{1}$ (MeV), $L_{2}$ (MeV), $L_{3}$ (MeV), $L_{4}$ (MeV), and $L_{5}$ (MeV) at saturation density $\\protect\\rho % _{0}$ for the MDI interaction with $x=-1$, $0$ and $1$, the Gogny-Hartree-Fock predictions with D1, D1S, D1N and D1M as well as the Skyrme-Hartree-Fock predictions with $112$ standard Skyrme interactions. The interactions in different models are in order according to the $L$ value. The corresponding reference is included as the last column.} \\label{TableI}% \\begin{tabular*}{\\textwidth}{@{\\extracolsep{\\fill}}lcccccccccccr} \\toprule Model & $\\rho_0$ & $E_0(\\rho_0)$ & $E_{sym}(\\rho_0)$ & $E_1$ & $E_2$ & L & $L_1$ & $L_2$ & $L_3$ & $L_4$ & $L_5$ & Ref. \\\\ \\hline $\\mathbf{MDI}$ &  &  &  &  &  &  &  &  &  &  &  &  \\\\ MDI x=1 & 0.160 & -16.2 & 30.5 & 18.2 & 12.3 & 14.6 & 36.4 & -6.2 & 36.9 & -18.6 & -33.9 & \\cite{Che05} \\\\ MDI x=0 & 0.160 & -16.2 & 30.5 & 18.2 & 12.3 & 60.2 & 36.4 & -6.2 & 36.9 & -18.6 & 11.6 & \\cite{Che05} \\\\ MDI x=-1 & 0.160 & -16.2 & 30.5 & 18.2 & 12.3 & 105.7 & 36.4 & -6.2 & 36.9 & -18.6 & 57.2 & \\cite{Che05} \\\\ $\\mathbf{Gogny}$ &  &  &  &  &  &  &  &  &  &  &  &  \\\\ D1 & 0.166 & -16.3 & 30.7 & 18.8 & 11.9 & 18.4 & 37.7 & -5.0 & 35.6 & -17.1 & -32.7 & \\cite{Dec80} \\\\ D1S & 0.163 & -16.0 & 31.1 & 17.9 & 13.3 & 22.4 & 35.7 & -7.6 & 39.8 & -7.5 & -37.9 & \\cite{Berger91a} \\\\ D1M & 0.165 & -16.0 & 28.6 & 16.8 & 11.8 & 24.8 & 33.6 & -4.2 & 35.3 & 4.5 & -44.3 & \\cite{Gor09a} \\\\ D1N & 0.161 & -16.0 & 29.6 & 16.5 & 13.1 & 33.6 & 33.0 & -3.7 & 39.2 & 9.2 & -44.2 & \\cite{Chap08a} \\\\ $\\mathbf{Skyrme}$ &  &  &  &  &  &  &  &  &  &  &  &  \\\\ Z-fit & 0.159 & -16.0 & 26.8 & 14.5 & 12.3 & -49.7 & 29.0 & 0 & 36.9 & -13.3 & -102.4 & \\cite{Fried86a} \\\\ Esigma-fit & 0.163 & -16.0 & 26.4 & 14.8 & 11.6 & -36.9 & 29.6 & 0 & 34.9 & -18.1 & -83.2 & \\cite{Fried86a} \\\\ E-fit & 0.159 & -16.1 & 27.7 & 14.1 & 13.6 & -31.3 & 28.2 & 0 & 40.7 & -15.2 & -84.9 & \\cite{Fried86a} \\\\ Zsigma-fit & 0.163 & -15.9 & 26.7 & 15.9 & 10.8 & -29.4 & 31.8 & 0 & 32.4 & -17.1 & -76.5 & \\cite{Fried86a} \\\\ SVII & 0.143 & -15.8 & 27.0 & 11.4 & 15.6 & -10.2 & 22.8 & 0 & 46.7 & -11.7 & -68.0 & \\cite{Giann80a} \\\\ SkSC4o & 0.161 & -15.9 & 27.0 & 12.3 & 14.7 & -9.7 & 24.6 & 0 & 44.0 & 0 & -78.3 & \\cite{Pea00a} \\\\ SVI & 0.143 & -15.8 & 26.9 & 12.0 & 14.9 & -7.3 & 24.1 & 0 & 44.6 & -12.2 & -63.8 & \\cite{Giann80a} \\\\ ZsigmaS-fit & 0.162 & -16.0 & 28.8 & 16.0 & 12.8 & -4.5 & 32.0 & 0 & 38.4 & -17.9 & -57.0 & \\cite{Fried86a} \\\\ v070 & 0.158 & -15.8 & 28.0 & 11.6 & 16.4 & -3.5 & 23.1 & 0 & 49.3 & -34.7 & -41.2 & \\cite{Pea01a} \\\\ v075 & 0.158 & -15.8 & 28.0 & 11.6 & 16.4 & -0.3 & 23.1 & 0 & 49.3 & -27.8 & -45.0 & \\cite{Pea01a} \\\\ SI & 0.155 & -16.0 & 29.2 & 13.2 & 16.0 & 1.2 & 26.4 & 0 & 48.1 & -11.4 & -61.9 & \\cite{Fried86a} \\\\ v080 & 0.157 & -15.8 & 28.0 & 11.6 & 16.4 & 2.2 & 23.1 & 0 & 49.3 & -21.7 & -48.5 & \\cite{Pea01a} \\\\ v090 & 0.157 & -15.8 & 28.0 & 11.6 & 16.4 & 5.0 & 23.1 & 0 & 49.3 & -11.6 & -55.8 & \\cite{Pea01a} \\\\ Skz4 & 0.160 & -16.0 & 32.0 & 17.5 & 14.5 & 5.8 & 35.1 & 0 & 43.4 & 27.9 & -100.6 & \\cite{Mar02a} \\\\ SkSC15 & 0.161 & -15.9 & 28.0 & 12.3 & 15.7 & 6.7 & 24.6 & 0 & 47.0 & 0 & -65.0 & \\cite{Pea00a} \\\\ BSk3 & 0.157 & -15.8 & 27.9 & 10.8 & 17.1 & 6.8 & 21.6 & 0 & 51.3 & -16.6 & -49.6 & \\cite{Gor03a} \\\\ MSk3 & 0.158 & -15.8 & 28.0 & 12.2 & 15.8 & 7.0 & 24.3 & 0 & 47.5 & 0 & -64.8 & \\cite{Ton00a} \\\\ v105 & 0.157 & -15.8 & 28.0 & 11.6 & 16.4 & 7.1 & 23.1 & 0 & 49.3 & 0 & -65.3 & \\cite{Pea01a} \\\\ MSk4 & 0.157 & -15.8 & 28.0 & 11.6 & 16.4 & 7.2 & 23.1 & 0 & 49.3 & 0 & -65.2 & \\cite{Ton00a} \\\\ BSk1 & 0.157 & -15.8 & 27.8 & 11.6 & 16.2 & 7.2 & 23.1 & 0 & 48.7 & 0 & -64.7 & \\cite{Sam02a} \\\\ v110 & 0.157 & -15.8 & 28.0 & 11.6 & 16.4 & 7.5 & 23.1 & 0 & 49.3 & 3.2 & -68.1 & \\cite{Pea01a} \\\\ MSk5 & 0.157 & -15.8 & 28.0 & 11.6 & 16.4 & 7.6 & 23.1 & 0 & 49.3 & 0 & -64.9 & \\cite{Ton00a} \\\\ BSk2p & 0.157 & -15.8 & 28.0 & 11.6 & 16.4 & 7.8 & 23.2 & 0 & 49.3 & -14.9 & -49.7 & \\cite{Gor02a} \\\\ BSk2 & 0.157 & -15.8 & 28.0 & 11.7 & 16.3 & 8.0 & 23.3 & 0 & 49.0 & -14.8 & -49.6 & \\cite{Gor02a} \\\\ MSk8 & 0.158 & -15.8 & 27.9 & 11.0 & 16.9 & 8.3 & 22.1 & 0 & 50.6 & 0 & -64.5 & \\cite{Gor01a} \\\\ v100 & 0.157 & -15.8 & 28.0 & 11.6 & 16.4 & 8.7 & 23.1 & 0 & 49.3 & -3.5 & -60.2 & \\cite{Pea01a} \\\\ MSk7 & 0.158 & -15.8 & 27.9 & 11.6 & 16.4 & 9.4 & 23.1 & 0 & 49.1 & 0 & -62.9 & \\cite{Gor01b} \\\\ MSk6 & 0.157 & -15.8 & 28.0 & 11.6 & 16.4 & 9.6 & 23.1 & 0 & 49.3 & 0 & -62.8 & \\cite{Ton00a} \\\\ SIII & 0.145 & -15.9 & 28.2 & 15.1 & 13.1 & 9.9 & 30.2 & 0 & 39.2 & -14.8 & -44.7 & \\cite{Bei75a} \\\\ MSk9 & 0.158 & -15.8 & 28.0 & 12.2 & 15.8 & 10.4 & 24.3 & 0 & 47.5 & 0.0 & -61.5 & \\cite{Gor01a} \\\\ BSk4 & 0.157 & -15.8 & 28.0 & 13.2 & 14.8 & 12.5 & 26.4 & 0 & 44.4 & -6.5 & -51.7 & \\cite{Gor03a} \\\\ Skz3 & 0.160 & -16.0 & 32.0 & 17.5 & 14.5 & 13.0 & 35.1 & 0 & 43.4 & 10.3 & -75.8 & \\cite{Mar02a} \\\\ BSk8 & 0.159 & -15.8 & 28.0 & 15.3 & 12.7 & 14.9 & 30.6 & 0 & 38.2 & 7.5 & -61.4 & \\cite{Sam04a} \\\\ Dutta & 0.162 & -16.0 & 26.6 & 12.4 & 14.2 & 16.5 & 24.8 & 0 & 42.7 & 0 & -51.0 & \\cite{Dut86a} \\\\ Skz2 & 0.160 & -16.0 & 32.0 & 17.5 & 14.5 & 16.8 & 35.1 & 0 & 43.4 & -10.7 & -50.9 & \\cite{Mar02a} \\\\ BSk6 & 0.157 & -15.8 & 28.0 & 15.2 & 12.8 & 16.8 & 30.4 & 0 & 38.4 & 6.2 & -58.2 & \\cite{Gor03a} \\\\ BSk7 & 0.157 & -15.8 & 28.0 & 15.2 & 12.8 & 18.0 & 30.4 & 0 & 38.4 & 7.5 & -58.3 & \\cite{Gor03a} \\\\ SKP & 0.163 & -16.0 & 30.0 & 12.4 & 17.6 & 19.6 & 24.8 & 0 & 52.7 & -26.1 & -31.9 & \\cite{bro98a} \\\\ BSk5 & 0.157 & -15.8 & 28.7 & 13.2 & 15.5 & 21.4 & 26.4 & 0 & 46.5 & -8.0 & -43.5 & \\cite{Gor03a} \\\\ Skz1 & 0.160 & -16.0 & 32.0 & 17.5 & 14.5 & 27.7 & 35.1 & 0 & 43.4 & -24.7 & -26.1 & \\cite{Mar02a} \\\\ \\hline\\hline \\end{tabular*}% \\end{table*}  \\begin{table*}[tbp] \\renewcommand{\\tablename}{} TABLE I (Continued.)\\newline \\begin{tabular*}{\\textwidth}{@{\\extracolsep{\\fill}}lcccccccccccr} \\toprule Model & $\\rho_0$ & $E_0(\\rho_0)$ & $E_{sym}(\\rho_0)$ & $E_1$ & $E_2$ & L & $L_1$ & $L_2$ & $L_3$ & $L_4$ & $L_5$ & Ref. \\\\ \\hline SIIIs & 0.148 & -16.1 & 32.0 & 15.0 & 17.0 & 28.7 & 29.9 & 0 & 51.0 & -5.0 & -47.3 & \\cite{Giann80a} \\\\ SKT6 & 0.161 & -16.0 & 30.0 & 12.3 & 17.6 & 30.9 & 24.7 & 0 & 52.9 & 0 & -46.7 & \\cite{Ton84a} \\\\ SKT7 & 0.161 & -15.9 & 29.5 & 14.8 & 14.7 & 31.1 & 29.6 & 0 & 44.2 & -14.8 & -27.8 & \\cite{Ton84a} \\\\ SKXm & 0.159 & -16.0 & 31.2 & 12.7 & 18.5 & 32.1 & 25.3 & 0 & 55.6 & -22.0 & -26.9 & \\cite{bro98a} \\\\ RATP & 0.160 & -16.0 & 29.2 & 18.4 & 10.9 & 32.4 & 36.8 & 0 & 32.6 & -20.5 & -16.5 & \\cite{Rayet} \\\\ SkSC14 & 0.161 & -15.9 & 30.0 & 12.3 & 17.7 & 33.1 & 24.6 & 0 & 53.0 & 0 & -44.5 & \\cite{Pea00a} \\\\ SKX & 0.155 & -16.1 & 31.1 & 12.1 & 19.0 & 33.2 & 24.3 & 0 & 56.9 & -23.7 & -24.3 & \\cite{bro98a} \\\\ MSk2 & 0.157 & -15.8 & 30.0 & 11.6 & 18.4 & 33.4 & 23.1 & 0 & 55.3 & 0.0 & -45.1 & \\cite{Ton00a} \\\\ SKXce & 0.155 & -15.9 & 30.1 & 12.0 & 18.2 & 33.5 & 23.9 & 0 & 54.5 & -24.2 & -20.8 & \\cite{bro98a} \\\\ BSk15 & 0.159 & -16.0 & 30.0 & 15.3 & 14.7 & 33.6 & 30.6 & 0.0 & 44.2 & -3.9 & -37.2 & \\cite{Gor08a} \\\\ SKT8 & 0.161 & -15.9 & 29.9 & 14.8 & 15.1 & 33.7 & 29.6 & 0 & 45.4 & 0 & -41.3 & \\cite{Ton84a} \\\\ SKT9 & 0.160 & -15.9 & 29.8 & 14.8 & 15.0 & 33.7 & 29.5 & 0 & 45.0 & 0 & -40.8 & \\cite{Ton84a} \\\\ MSk1 & 0.157 & -15.8 & 30.0 & 12.2 & 17.8 & 33.9 & 24.3 & 0 & 53.5 & 0 & -43.9 & \\cite{Ton00a} \\\\ BSk16 & 0.159 & -16.1 & 30.0 & 15.3 & 14.7 & 34.9 & 30.5 & 0.0 & 44.2 & -2.1 & -37.8 & \\cite{Cha08a} \\\\ Skz0 & 0.160 & -16.0 & 32.0 & 17.5 & 14.5 & 35.1 & 35.1 & 0 & 43.4 & -42.1 & -1.3 & \\cite{Mar02a} \\\\ Skyrme1p & 0.155 & -16.0 & 29.4 & 13.2 & 16.1 & 35.3 & 26.4 & 0 & 48.4 & -11.4 & -28.1 & \\cite{Pet95a} \\\\ BSk17 & 0.159 & -16.1 & 30.0 & 15.3 & 14.7 & 36.3 & 30.5 & 0.0 & 44.2 & -2.1 & -36.4 & \\cite{Gor09b} \\\\ BSk10 & 0.159 & -15.9 & 30.0 & 13.3 & 16.7 & 37.2 & 26.6 & 0 & 50.1 & -10.5 & -29.0 & \\cite{Gor06a} \\\\ SGII & 0.158 & -15.6 & 26.8 & 15.5 & 11.3 & 37.6 & 31.0 & 0 & 33.9 & -16.1 & -11.2 & \\cite{Giai81a} \\\\ BSk12 & 0.159 & -15.9 & 30.0 & 13.3 & 16.7 & 38.0 & 26.5 & 0 & 50.2 & -9.9 & -28.8 & \\cite{Gor06a} \\\\ BSk11 & 0.159 & -15.9 & 30.0 & 13.3 & 16.7 & 38.4 & 26.5 & 0 & 50.2 & -9.8 & -28.6 & \\cite{Gor06a} \\\\ SLy10 & 0.156 & -15.9 & 32.0 & 17.6 & 14.3 & 38.7 & 35.3 & 0 & 43.0 & 15.4 & -55.0 & \\cite{Chaba95a} \\\\ BSk13 & 0.159 & -15.9 & 30.0 & 13.3 & 16.7 & 38.8 & 26.5 & 0 & 50.2 & -9.7 & -28.2 & \\cite{Gor06a} \\\\ BSk9 & 0.159 & -15.9 & 30.0 & 15.3 & 14.7 & 39.9 & 30.6 & 0 & 44.2 & 10.8 & -45.7 & \\cite{Sam04a} \\\\ KDE & 0.164 & -16.0 & 32.0 & 16.5 & 15.4 & 41.4 & 33.1 & 0 & 46.3 & 12.3 & -50.2 & \\cite{Agra05a} \\\\ BSk14 & 0.159 & -15.9 & 30.0 & 15.3 & 14.7 & 43.9 & 30.5 & 0.0 & 44.2 & -2.0 & -28.8 & \\cite{Gor07a} \\\\ SLy230a & 0.160 & -16.0 & 32.0 & 17.6 & 14.4 & 44.3 & 35.2 & 0 & 43.1 & 32.0 & -66.1 & \\cite{Chaba98a} \\\\ KDE0 & 0.161 & -16.1 & 33.0 & 17.2 & 15.8 & 45.2 & 34.4 & 0 & 47.4 & 6.9 & -43.5 & \\cite{Agra05a} \\\\ SLy8 & 0.160 & -16.0 & 31.4 & 17.7 & 13.8 & 45.3 & 35.3 & 0 & 41.3 & 13.7 & -45.1 & \\cite{Chaba95a} \\\\ SLy4 & 0.160 & -16.0 & 31.8 & 17.6 & 14.2 & 45.4 & 35.3 & 0 & 42.5 & 13.8 & -46.2 & \\cite{Chaba98a} \\\\ SLy0 & 0.161 & -16.0 & 31.5 & 17.6 & 13.8 & 45.4 & 35.3 & 0 & 41.5 & 13.6 & -45.0 & \\cite{Chaba95a} \\\\ SLy3 & 0.160 & -16.0 & 32.1 & 17.7 & 14.4 & 45.5 & 35.3 & 0 & 43.2 & 13.8 & -46.8 & \\cite{Chaba95a} \\\\ SKMs & 0.160 & -15.8 & 30.0 & 15.6 & 14.4 & 45.8 & 31.2 & 0 & 43.3 & -19.4 & -9.3 & \\cite{Bart82a} \\\\ SLy230b & 0.160 & -16.0 & 32.0 & 17.6 & 14.4 & 46.0 & 35.3 & 0 & 43.1 & 13.9 & -46.4 & \\cite{Chaba97a} \\\\ SLy7 & 0.158 & -15.9 & 32.0 & 17.7 & 14.3 & 47.2 & 35.4 & 0 & 42.8 & 15.0 & -46.0 & \\cite{Chaba98a} \\\\ SLy6 & 0.159 & -15.9 & 32.0 & 17.7 & 14.2 & 47.4 & 35.4 & 0 & 42.7 & 14.6 & -45.3 & \\cite{Chaba98a} \\\\ SKb & 0.155 & -16.0 & 23.9 & 19.8 & 4.1 & 47.5 & 39.6 & 0 & 12.3 & -21.6 & 17.3 & \\cite{Giai81a} \\\\ SLy5 & 0.160 & -16.0 & 32.0 & 17.6 & 14.4 & 48.3 & 35.3 & 0 & 43.1 & 13.7 & -43.8 & \\cite{Chaba98a} \\\\ SLy2 & 0.160 & -15.9 & 32.3 & 17.6 & 14.7 & 48.8 & 35.2 & 0 & 44.0 & 13.6 & -44.1 & \\cite{Chaba95a} \\\\ SLy1 & 0.160 & -16.0 & 32.5 & 17.6 & 14.9 & 48.8 & 35.2 & 0 & 44.7 & 13.7 & -44.8 & \\cite{Chaba95a} \\\\ SKM & 0.160 & -15.8 & 30.7 & 15.6 & 15.2 & 49.3 & 31.2 & 0 & 45.5 & -18.0 & -9.3 & \\cite{Kri80a} \\\\ SII & 0.148 & -16.0 & 34.2 & 20.1 & 14.0 & 50.0 & 40.3 & 0 & 42.0 & -19.2 & -13.1 & \\cite{BV1972} \\\\ Skzm1 & 0.160 & -16.0 & 32.0 & 17.5 & 14.5 & 54.1 & 35.1 & 0 & 43.4 & -47.9 & 23.6 & \\cite{Mar02a} \\\\ SKT3 & 0.161 & -15.9 & 31.5 & 12.3 & 19.2 & 55.3 & 24.7 & 0 & 57.5 & 0 & -26.8 & \\cite{Ton84a} \\\\ SLy9 & 0.151 & -15.8 & 32.1 & 17.8 & 14.4 & 55.4 & 35.5 & 0 & 43.2 & 17.9 & -41.1 & \\cite{Chaba95a} \\\\ SKT3s & 0.160 & -16.0 & 31.7 & 12.3 & 19.4 & 55.9 & 24.6 & 0 & 58.2 & 0 & -26.9 & \\cite{Ton84a} \\\\ SKT1s & 0.160 & -16.0 & 32.0 & 12.3 & 19.7 & 56.1 & 24.6 & 0 & 59.2 & 0 & -27.7 & \\cite{Ton84a} \\\\ SKT2 & 0.161 & -15.9 & 32.0 & 12.3 & 19.7 & 56.2 & 24.7 & 0 & 59.0 & 0 & -27.5 & \\cite{Ton84a} \\\\ SKT1 & 0.161 & -16.0 & 32.0 & 12.3 & 19.7 & 56.2 & 24.7 & 0 & 59.1 & 0 & -27.5 & \\cite{Ton84a} \\\\ MSkA & 0.153 & -16.0 & 30.3 & 15.1 & 15.3 & 57.2 & 30.1 & 0 & 45.9 & -16.4 & -2.4 & \\cite{Shar95a} \\\\ SkI6 & 0.159 & -15.9 & 29.9 & 19.1 & 10.8 & 59.2 & 38.2 & 0 & 32.4 & 23.1 & -34.4 & \\cite{Naza96a} \\\\ MSL0 & 0.160 & -16.0 & 30.0 & 15.4 & 14.6 & 60.0 & 30.7 & 0 & 43.9 & -13.2 & -1.5 & \\cite{Che10} \\\\ SkI4 & 0.160 & -15.9 & 29.5 & 18.9 & 10.6 & 60.4 & 37.8 & 0 & 31.7 & 21.4 & -30.6 & \\cite{Rein95a} \\\\ LNS & 0.175 & -15.3 & 33.4 & 15.8 & 17.7 & 61.5 & 31.5 & 0 & 53.0 & -12.8 & -10.2 & \\cite{Cao06a} \\\\ \\hline\\hline \\end{tabular*}% \\end{table*}  \\begin{table*}[tbp] \\renewcommand{\\tablename}{} TABLE I (Continued.)\\newline \\begin{tabular*}{\\textwidth}{@{\\extracolsep{\\fill}}lcccccccccccr} \\toprule Model & $\\rho_0$ & $E_0(\\rho_0)$ & $E_{sym}(\\rho_0)$ & $E_1$ & $E_2$ & L & $L_1$ & $L_2$ & $L_3$ & $L_4$ & $L_5$ & Ref. \\\\ \\hline SIV & 0.151 & -16.0 & 31.2 & 25.1 & 6.1 & 63.5 & 50.2 & 0 & 18.4 & -23.5 & 18.5 & \\cite{Bei75a} \\\\ SGI & 0.154 & -15.9 & 28.3 & 19.7 & 8.6 & 63.9 & 39.5 & 0 & 25.8 & -7.0 & 5.6 & \\cite{Giai81a} \\\\ SKOs & 0.160 & -15.8 & 31.9 & 13.7 & 18.2 & 68.9 & 27.4 & 0 & 54.7 & -2.4 & -10.8 & \\cite{Rein99a} \\\\ SkMP & 0.157 & -15.6 & 29.9 & 18.6 & 11.3 & 70.3 & 37.1 & 0 & 34.0 & -13.1 & 12.3 & \\cite{Benn89a} \\\\ Ska & 0.155 & -16.0 & 32.9 & 19.8 & 13.1 & 74.6 & 39.6 & 0 & 39.4 & -21.6 & 17.3 & \\cite{Koh76a} \\\\ SKO & 0.160 & -15.8 & 32.0 & 13.7 & 18.2 & 79.1 & 27.5 & 0 & 54.7 & -4.3 & 1.3 & \\cite{Rein99a} \\\\ SKYT & 0.148 & -15.4 & 33.7 & 19.3 & 14.3 & 80.8 & 38.7 & 0 & 43.0 & -10.8 & 9.9 & \\cite{Ko74a} \\\\ Rsigma-fit & 0.158 & -15.6 & 30.6 & 15.5 & 15.0 & 85.7 & 31.1 & 0.0 & 45.1 & -14.6 & 24.1 & \\cite{Fried86a} \\\\ SK272 & 0.155 & -16.3 & 37.4 & 15.6 & 21.8 & 91.7 & 31.2 & 0 & 65.4 & -17.6 & 12.6 & \\cite{Agra03a} \\\\ Gsigma-fit & 0.158 & -15.6 & 31.4 & 15.5 & 15.9 & 94.0 & 31.0 & 0.0 & 47.6 & -14.6 & 30.0 & \\cite{Fried86a} \\\\ SKT4 & 0.159 & -16.0 & 35.5 & 12.2 & 23.2 & 94.1 & 24.5 & 0 & 69.7 & 0 & 0 & \\cite{Ton84a} \\\\ SK255 & 0.157 & -16.3 & 37.4 & 15.2 & 22.2 & 95.1 & 30.5 & 0 & 66.5 & -20.9 & 19.0 & \\cite{Agra03a} \\\\ SV & 0.155 & -16.0 & 32.8 & 31.4 & 1.4 & 96.1 & 62.8 & 0 & 4.2 & -29.1 & 58.2 & \\cite{Fried86a} \\\\ SKT5 & 0.164 & -16.0 & 37.0 & 12.5 & 24.5 & 98.5 & 25.0 & 0 & 73.6 & 0 & 0 & \\cite{Ton84a} \\\\ SkI3 & 0.158 & -16.0 & 34.8 & 21.1 & 13.8 & 100.5 & 42.1 & 0 & 41.3 & 35.5 & -18.4 & \\cite{Rein95a} \\\\ SkI2 & 0.158 & -15.8 & 33.4 & 17.7 & 15.6 & 104.3 & 35.5 & 0 & 46.9 & 15.7 & 6.2 & \\cite{Rein95a} \\\\ SkI5 & 0.156 & -15.8 & 36.6 & 20.8 & 15.8 & 129.3 & 41.7 & 0 & 47.4 & 35.0 & 5.2 & \\cite{Rein95a} \\\\ SkI1 & 0.160 & -16.0 & 37.5 & 17.7 & 19.8 & 161.1 & 35.5 & 0 & 59.3 & 14.2 & 52.0 & \\cite{Rein95a} \\\\ \\hline\\hline \\end{tabular*}% \\end{table*}  Due to the special interest on the values of $E_{sym}(\\rho )$ and $L(\\rho )$ at $\\rho _{0}$, we list in Table \\ref{TableI} the values of the characteristic parameters $\\rho _{0}$, $E_{0}(\\rho _{0})$, $E_{sym}(\\rho _{0})$, $E_{1}$, $E_{2}$, $L$, $L_{1}$, $L_{2}$, $L_{3}$, $L_{4}$ and $L_{5}$ at $\\rho _{0}$ for the MDI interaction with $x=-1$, $0$ and $1$, the Gogny-Hartree-Fock predictions with D1, D1S, D1N, and D1M as well as the Skyrme-Hartree-Fock predictions with $112$ standard Skyrme interactions (In the table, the interactions in different models are in order according to the $L$ value and the corresponding reference for different interactions is included in the last column). For the MDI interaction model, one can see from Table \\ref{TableI} that the value of $E_{2}(\\rho _{0})$ is comparable with that of $E_{1}(\\rho _{0})$, and all the $E_{0}(\\rho _{0})$, $E_{1}(\\rho _{0})$, and $E_{2}(\\rho _{0})$ are the same for different $x$ values due to the fact that the $U_{0}(\\rho ,k)$ and $U_{sym,1}(\\rho _{0},k)$ are independent of the $x$ parameter by construction \\cite{Che05}. Because of the same reason, the values of $L_{1}$, $L_{2}$, $L_{3}$, $L_{4}$ are all independent of the $x$ parameter too. Therefore, for the MDI interaction, the $x$ dependence of the $L$ parameter is completely determined by the second-order symmetry potential $U_{sym,2}(\\rho ,k_{F})$. Depending on the value of the $x$ parameter, the $U_{sym,2}(\\rho ,k_{F})$ can be positive or negative. In particular, we have $L=L_{1}+L_{2}+L_{3}+L_{4}=48.5$ MeV if we assume $U_{sym,2}(\\rho ,k_{F})=0$. Furthermore, it is seen that the $L_{2}$ contribution is relatively small compared with that of $L_{1}$, $L_{3}$, or $% L_{4}$, indicating that the contribution due to the momentum dependence of the nucleon effective mass is unimportant, and this is consistent with the observation from Fig. \\ref{LMDI}.  For the Gogny interaction, the value of $L$ listed in Table \\ref{TableI} ranges from about $18$ MeV to $34$ MeV. And similarly with the MDI\\ interaction model, the $L_{2}$ is relatively small, ranging from about $-8$ MeV to $-3$ MeV. It is interesting to see that the values of $L_{1}$ and $% L_{3}$ from the Gogny interactions are quite similar with those of the MDI\\ interaction model, and thus the difference of the $L$ parameter from different interactions in these two models is mainly due to the variation of the $L_{4}$ and $L_{5}$. Furthermore, for different Gogny interactions considered here, the value of $L_{4}$ can be positive or negative while the value of $L_{5}$ is always negative, and the $L_{5}$ contribution to the $L$ parameter usually is relatively important.  For the standard Skyrme interactions, we have $L_{2}=0$ MeV. For the $112$ Skyrme interactions considered in Table \\ref{TableI}, it is seen that the value of $L$ ranges from about $-50$ MeV to $160$ MeV, $L_{1}$ from about $% 20 $ MeV to $60$ MeV, $L_{3}$ from about $4$ MeV to $74$ MeV, $L_{4}$ from about $-48$ MeV to $36$ MeV, and $L_{5}$ from about $-102$ MeV to $58$ MeV. Therefore, the different term contributions to the $L$ parameter can change a lot in the standard Skyrme interactions, especially for $L_{3}$, $L_{4}$, and $L_{5}$. \\begin{figure}[tbh] \\includegraphics[width=8.5cm]{L15SHF.eps} \\caption{(Color online) Correlations of $L_{1}(\\protect\\rho _{0})$ (a), $% L_{3}(\\protect\\rho _{0})$ (b), $L_{4}(\\protect\\rho _{0})$ (c), and $L_{5}(% \\protect\\rho _{0})$ (d) with $L(\\protect\\rho _{0})$ for the MDI interaction with $x=-1$, $0$ and $1$, Gogny-Hartree-Fock predictions with D1, D1S, D1N and D1M as well as Skyrme-Hartree-Fock predictions with the $112$ standard Skyrme interactions considered in Table \\protect\\ref{TableI}.} \\label{L15SHF} \\end{figure}  In order to see more clearly and intuitively the different term contributions to the $L$ parameter, we show in Fig. \\ref{L15SHF} the correlations of $L_{1}(\\rho _{0})$, $L_{3}(\\rho _{0})$, $L_{4}(\\rho _{0})$, and $L_{5}(\\rho _{0})$ with $L(\\rho _{0})$ for the MDI interaction with $% x=-1 $, $0$ and $1$, the Gogny-Hartree-Fock predictions with D1, D1S, D1N, and D1M as well as the Skyrme-Hartree-Fock predictions with the $112$ standard Skyrme interactions. One can see from Fig. \\ref{L15SHF} that the results from the MDI interaction model and the Gogny-Hartree-Fock calculations are essentially consistent with the systematics of the Skyrme-Hartree-Fock predictions. Based on these calculated results, we find that a statistical analysis can lead to $L_{1}(\\rho _{0})\\approx 30\\pm 6.5$ MeV, $L_{3}(\\rho _{0})\\approx 46\\pm 9.5$ MeV, $L_{4}(\\rho _{0})\\approx -4\\pm 15$ MeV, and $L_{5}(\\rho _{0})\\approx -35\\pm 30$ MeV. These results indicate that, within the standard Skyrme-Hartree-Fock energy density functional, $% L_{1}(\\rho _{0}) $ and $L_{3}(\\rho _{0})$ are relatively well constrained, and the main uncertainties are due to the $L_{4}(\\rho _{0})$ and $L_{5}(\\rho _{0})$ contributions. Furthermore, it is interesting to see from Fig. \\ref% {L15SHF} that there exhibits an approximately linear correlation between $% L_{5}(\\rho _{0})$ and $L(\\rho _{0})$. If we use the present empirical constraint $L(\\rho _{0})=60\\pm 30$ MeV, then we find the $L_{5}(\\rho _{0})$ can vary from about $-66$ MeV to $24$ MeV, i.e., the value of $% U_{sym,2}(\\rho _{0},k_{F})$ can vary from about $-22$ MeV to $8$ MeV.  \\subsection{The symmetry potential $U_{sym,1}(\\protect\\rho ,k)$}  Shown in Fig. \\ref{Usym1MDI} is the momentum dependence of the $% U_{sym,1}(\\rho ,k)$ at $\\rho =0.5\\rho _{0}$, $\\rho _{0}$ and $2\\rho _{0}$ using the MDI interaction with $x=-1$, $0$, and $1$. For comparison, we also include in Fig. \\ref{Usym1MDI} the corresponding results from several microscopic approaches, including the relativistic impulse approximation (RIA) \\cite{Che05c,LiZH06} using the empirical nucleon-nucleon scattering amplitude determined by Murdock and Horowitz (MH) \\cite{Mur87} with isospin-dependent and isospin-independent Pauli blocking corrections as well as by McNeil, Ray, and Wallace (MRW) \\cite{McN83}, the relativistic Dirac-Brueckner-Hartree-Fock (DBHF) theory \\cite{Dal05}, and the non-relativistic Brueckner-Hartree-Fock (BHF) theory with and without the three-body force (TBF) rearrangement contribution \\cite{Zuo06}. For these microscopic results, one can see that they are all consistent with each other around and below $% \\rho _{0}$ although there still exist larger uncertainties at higher density of $\\rho =2\\rho _{0}$. It is interesting to see that the momentum dependence of the $U_{sym,1}(\\rho ,k) $ from the MDI interaction with $x=0$ are in good agreement with the results from the microscopic approaches. It should be noted that the momentum dependence of the $U_{sym,1}(\\rho ,k)$ at $\\rho _{0}$ is the same for $x=-1$, $0$, and $1$ since $U_{sym,1}(\\rho ,k)$ is independent of the $x$ parameter at $\\rho _{0}$ by construction as mentioned previously. \\begin{figure}[tbh] \\includegraphics[width=8.6cm]{Usym1MDI.eps} \\caption{(Color online) Momentum dependence of the $U_{sym,1}(\\protect\\rho % ,k)$ at $\\protect\\rho =0.5\\protect\\rho _{0}$ (a), $\\protect\\rho _{0}$ (b) and $2\\protect\\rho _{0}$ (c) using the MDI interaction with $x=-1$, $0$, and $1$. The corresponding results from several microscopic approaches are also included for comparison (See the text for details).} \\label{Usym1MDI} \\end{figure} \\begin{figure}[tbh] \\includegraphics[width=8.6cm]{Usym1SHF.eps} \\caption{(Color online) Same as Fig. \\protect\\ref{Usym1MDI} but for the Skyrme Hartree-Fock\\ approach with MSL0, SLy4, and SKM*.} \\label{Usym1SHF} \\end{figure}  Similarly as in Fig. \\ref{Usym1MDI}, we plot in Fig. \\ref{Usym1SHF} and Fig. % \\ref{Usym1Gogny} the the momentum dependence of the $U_{sym,1}(\\rho ,k)$ at $% \\rho =0.5\\rho _{0}$, $\\rho _{0}$ and $2\\rho _{0}$ using the Skyrme Hartree-Fock\\ approach with MSL0, SLy4 and SKM*, and the Gogny Hartree-Fock\\ approach with D1, D1S, D1N and D1M, respectively. On the one hand, it is seen from Fig. \\ref{Usym1SHF} that, for $0<p<700$ MeV/c, the results from the MSL0 and SKM* interactions agree well with those from microscopic approaches while the results from the SLy4 interaction seem to display large deviation from the microscopic results, especially at higher nucleon momenta. On the other hand, one can see from Fig. \\ref{Usym1Gogny} that at $% \\rho =0.5\\rho _{0}$ and $\\rho _{0}$, the results from D1 interaction are consistent with those of the microscopic calculations while the results from D1S, D1N and D1M exhibit large deviation from the the microscopic calculations except at low nucleon momenta ($p\\lesssim 300$ MeV/c). At $\\rho =2\\rho _{0}$, the Gogny Hartree-Fock\\ calculations display different results from the microscopic ones and the strong model dependence appears. \\begin{figure}[tbh] \\includegraphics[width=8.6cm]{Usym1Gogny.eps} \\caption{(Color online) Same as Fig. \\protect\\ref{Usym1MDI} but for Gogny Hartree-Fock\\ approach with D1, D1S, D1N, and D1M.} \\label{Usym1Gogny} \\end{figure}  Overall, one can see from Fig. \\ref{Usym1MDI}, Fig. \\ref{Usym1SHF} and Fig. % \\ref{Usym1Gogny} that the momentum dependence of the $U_{sym,1}$ varies from one interaction to another for the MDI interaction model, the Skyrme Hartree-Fock\\ approach and the Gogny Hartree-Fock\\ approach. For the MDI interaction with $x=-1$, $0$, and $1$, the Gogny interaction with D1, and the Skyrme interaction with MSL0 and SKM*, the value of $U_{sym,1}$ can be negative at higher momentum while for the Gogny interaction with D1S, D1N, D1M and the Skyrme interaction with SLy4, the value of $U_{sym,1}$ is positive at higher momentum.  \\subsection{The second-order symmetry potential $U_{sym,2}(\\protect\\rho ,k)$}  As for $U_{sym,2}$, to our best knowledge, there is not any experimental/empirical information or theoretical predictions so far. Fig.~% \\ref{Usym2MDI} and Fig.~\\ref{Usym2Gogny} show the momentum dependence of $% U_{sym,2}$ at $0.5\\rho _{0}$, $\\rho _{0}$ and $2\\rho _{0}$ in the MDI interaction model and the Gogny Hartree-Fock approach, respectively. Since $% U_{sym,2}$ is independent of the nucleon momentum in the Skyrme Hartree-Fock approach, we only show here its density dependence in Fig.~\\ref{Usym2SHF}. From Fig.~\\ref{Usym2MDI} and Fig.~\\ref{Usym2Gogny}, it is interesting to see that for all interactions in the MDI interaction model and the Gogny Hartree-Fock approach at $\\rho =0.5\\rho _{0}$, $\\rho _{0}$ and $2\\rho _{0}$, $U_{sym,2}$ firstly decreases with nucleon momentum and then essentially saturates when the nucleon momentum is larger than about ${500}$ MeV/c. Especially, the results from the MDI interaction with ${x=1}$ seem to be in quantitative agreement with those from the Gogny Hartree-Fock approach. Furthermore, one can see from Fig.~\\ref{Usym2MDI} and Fig.~\\ref{Usym2Gogny} that the magnitude of $U_{sym,2}$ increases with the density, and this is also true for the Skyrme Hartree-Fock approach as shown in Fig.~\\ref% {Usym2SHF}. Another interesting feature is that for the MDI interaction model and the Skyrme Hartree-Fock approach, $U_{sym,2}$ can be either negative or positive while it is always negative for the Gogny Hartree-Fock approach with the interactions considered here. Therefore, any experimental constraints about $U_{sym,2}$ will be very useful and important for constraining the theoretical models.  \\begin{figure}[tbh] \\includegraphics[width=8.6cm]{Usym2MDI.eps} \\caption{(Color online) Momentum dependence of the $U_{sym,2}(\\protect\\rho % ,k)$ at $\\protect\\rho =0.5\\protect\\rho _{0}$ (a), $\\protect\\rho _{0}$ (b) and $2\\protect\\rho _{0}$ (c) using the MDI interaction with $x=-1$, $0$, and $1$.} \\label{Usym2MDI} \\end{figure} \\begin{figure}[tbh] \\includegraphics[width=8.6cm]{Usym2Gogny.eps} \\caption{(Color online) Same as Fig. \\protect\\ref{Usym2MDI} but for Gogny Hartree-Fock\\ approach with D1, D1S, D1N, and D1M.} \\label{Usym2Gogny} \\end{figure} \\begin{figure}[tbh] \\includegraphics[width=8.3cm]{Usym2SHF.eps} \\caption{(Color online) Density dependence of $U_{sym,2}$ for the Skyrme Hartree-Fock\\ approach with MSL0, SLy4, and SKM*.} \\label{Usym2SHF} \\end{figure}  Comparing Fig.~\\ref{Usym1MDI} with Fig.~\\ref{Usym2MDI} for the MDI interaction model, Fig.~\\ref{Usym1SHF} with Fig.~\\ref{Usym2SHF} for the Skyrme Hartree-Fock approach, and Fig.~\\ref{Usym1Gogny} with Fig.~\\ref% {Usym2Gogny} for the Gogny Hartree-Fock approach, one can see that, at fixed density and momentum, the magnitude of $U_{sym,2}$ is generally comparable with that of $U_{sym,1}$. So, in Eq.~(\\ref{UtauTaylor}), compared with $% U_{sym,1}(\\rho ,k)\\delta $, the contribution from $U_{sym,2}(\\rho ,k)\\delta ^{2}$ could be negligible only if $\\delta $ is small ($\\delta \\ll 1$). Therefore, we conclude from the present model calculations that, when $% \\delta $ is small ($\\delta \\ll 1$), the Lane potential might be a good approximation to the single-nucleon potential $U_{n/p}(\\rho ,\\delta ,p)$. However, the contributions of $U_{sym,2}(\\rho ,k)$ might not be simply neglected when $\\delta $ is close to $1$."
    },
    "1112/1112.0272_arXiv.txt": {
        "abstract": "The SAGE-LMC, SAGE-SMC and HERITAGE surveys have mapped the Magellanic Clouds in the infrared using the Spitzer and Herschel Space Telescopes. Over 8.5 million point sources were detected and catalogued in the LMC alone. Staring mode observations using the InfraRed Spectrograph (IRS) on board Spitzer have been obtained for 1,000 positions in the LMC and $\\sim$250 in the SMC. From the infrared spectroscopy we have identified the nature of the sources for which spectroscopy is available. These IRS staring mode targets represent an important contribution to the SED of these dwarf galaxies. Here we report on our latest results. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.2088_arXiv.txt": {
        "abstract": "Hypervelocity stars are believed to be ejected out from the Galactic center through dynamical interactions of (binary) stars with the central massive black hole(s). In this letter, we report 13 metal-poor F-type hypervelocity star candidates selected from 370,000 stars of the data release 7 of the Sloan Digital Sky Survey. With a detailed analysis of the kinematics of these stars, we find that seven of them were likely ejected from the Galactic center (GC) or the Galactic disk, four neither originated from the GC nor the Galactic disk, and the other two were possibly ejected from either the Galactic disk or other regions. Those candidates which unlikely originated from the GC or the Galactic disk, may be explained by other mechanisms, like the tidal disruption of the Milky Way's dwarf galaxies in the Galactic potential, or the gravitational interactions with a massive black hole at the center of M31 or M32. ",
        "introduction": "The hypervelocity stars (HVSs), recently discovered in the Galactic halo \\citep{Brown05,Hirsch05, Edelmann05}, are moving so fast that they may escape from the Galaxy. A natural explanation of these HVSs is that they are ejected out from the Galactic center (GC) by interactions of stars with the massive black hole (MBH) or the hypothetical binary MBHs as predicted by \\cite{Hills88} and \\cite{Yu03}. The ejection mechanisms of HVSs can be divided into three categories: tidal breakup of binary stars in the vicinity of a single MBH \\citep{Hills88, Yu03, Bromley06}, and the binary stars are probably injected into the vicinity of the MBH from the young stellar disk in the GC \\citep[e.g., ][]{Lu10, Zhang10} or from the Galactic bulge \\citep{Perets09a,Perets09b}; single star encounters with a binary MBH \\citep{Yu03, Sesana06, Sesana07, Merritt06}; or single star encounters with a cluster of stellar mass black holes around the MBH \\citep{Oleary08}.  More than 16 HVSs have been reported in the literature \\citep{Brown09}\\citep[See also: ][]{Tillich09, Kollmeier09, Tillich10}, most of them are 3--4 M$_{\\sun}$ late B-type stars. Assuming a Salpeter initial mass function (IMF), the expected solar mass HVSs are about 10 times more abundant than 3--4 M$_{\\sun}$ HVSs \\citep{Brown09}. \\cite{Kollmeier09} systematically searched for such low mass HVSs in about 290,000 spectra of the Sloan Digital Sky Survey (SDSS), however, they found only 6 metal-poor stars that can be possibly taken as HVS candidates, which might suggest that the IMF of the parent population of these HVSs is top heavy. A top heavy IMF of the HVS parent population is possibly consistent with the disk origination \\citep{Bartko10, Lu10, Kollmeier10}. In order to distinguish the ejection mechanisms of HVSs and put constraints on the origin of the parent population of HVSs, it is quite necessary to search for the low-mass HVSs.  In this letter, we aim to find F/G type low mass HVS candidates from the data release 7 (DR7) of SDSS, which provides a large catalog of stars with precise multi-color photometry and medium resolution ($R\\sim 1800$)spectra \\citep{York00}. We report the finding of 13 old metal-poor F-type HVS candidates, and also discuss their possible origins by detailed analysis of their kinematics. The letter is organized as follows. In Section 2, we introduce the searching process of our  low-mass HVS candidates. Section 3 discusses their possible ejection mechanisms through kinematics. Finally, a brief conclusion is given in Section 4. ",
        "conclusions": "In this letter, we report 13 metal-poor F-type HVS candidates from over 370,000 stars in the SDSS DR7. Through the kinematic analysis, we find that seven of them are likely to have been ejected from the GC or the Galactic disk. Two of them were possibly ejected out from either the Galactic disk or other places. Meanwhile, the other four candidates were unlikely to originate from the GC or the Galactic disk. Those candidates impossibly ejected from the GC or the Galactic disk may be explained by other mechanisms, i.e., being ejected from the tidal break of the Milky Way's dwarf galaxies \\citep{Abadi09, Teyssier09}, or from the center of M31 or M32 by the interactions of stars with the  MBH \\citep{Sherwin08}. In order to understand the origins of these mpHVS candidates well and solve the binary problem, second-epoch observations in future are needed."
    },
    "1112/1112.2041_arXiv.txt": {
        "abstract": "We have imaged the extremely high-velocity outflowing gas in CO (2$-$1) and (3$-$2) associated with the shell-like ultracompact \\htwo region G5.89$-$0.39 at a resolution of $\\sim$3$\\arcsec$ (corresponding to $\\sim$4000 AU) with the Submillimeter Array. The integrated high-velocity ($\\gtrsim$45 km s$^{-1}$) CO emission reveals at least three blueshifted lobes and two redshifted lobes. These lobes belong to two outflows, one oriented N-S, the other NW-SE. The NW-SE outflow is likely identical to the previously detected Br$\\gamma$ outflow. Furthermore, these outflow lobes all clearly show a Hubble-like kinematic structure. For the first time, we estimate the temperature of the outflowing gas as a function of velocity with the large velocity gradient calculations. Our results reveal a clear increasing trend of temperature with gas velocity. The observational features of the extremely high-velocity gas associated with G5.89$-$0.39 qualitatively favor the jet-driven bow shock model. ",
        "introduction": "}  Although molecular outflows have been commonly identified around both low- and high-mass young stellar objects (YSOs) \\citep[e.g.,][]{bac99,arc07}, it is not clear how the bulk of the outflowing gas is accelerated. Past observations show that molecular outflows can often be divided into two components --- the ``classical'' less-collimated component with low-velocity and the highly collimated component with extremely high-velocity (50$-$150 km s$^{-1}$, hereafter EHV) \\citep{bac99,hir06,pal06}. Exploring the physical conditions of the EHV molecular gas in details will be helpful for clarifying its role in star formation processes.  The massive star forming region G5.89$-$0.39 (hereafter G5.89) is associated with energetic CO outflows with velocities up to $\\sim$$\\pm$70 km s$^{-1}$ from its systemic velocity, among the highest gas velocities that have been detected in molecular outflows \\citep[e.g.,][]{cho93}. G5.89 also harbors a shell-like ultracompact (UC) \\htwo region powered by a young O-type star (hereafter Feldt's star) visible in the near-IR \\citep[][]{wc89,fel03}. Based on the deduced gas temperature structure, \\citet{su09} conclude that Feldt's star supplies the majority of energy for heating the dense gas enshrouding the UC \\htwo region.  The outflowing gas associated with G5.89 was first identified by \\citet{har88} in CO (1$-$0) with the IRAM 30-m telescope. Subsequent studies reported outflows with notably different orientations in various tracers. For example, SiO (5$-$4) observations with the Submillimeter Array\\footnote{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics, and is funded by the Smithsonian Institution and the Academia Sinica.} (SMA) resolved a bipolar structure along a position angle (P.A.) of $\\sim$30$^\\circ$ \\citep{sol04}. On the other hand, sub-arcsecond CO (3$-$2) images made with the SMA  \\citep{hun08} detected a well-collimated N-S outflow and tentatively a compact NW-SE outflow associated with the Br$\\gamma$ outflow revealed by \\citet{pug06}. Both CO outflows demonstrate EHV gas, with velocities $\\gtrsim$50 km s$^{-1}$. Furthermore, BIMA observations in CO (1$-$0) and HCO$^+ $ (1$-$0) identified an outflow along the E-W direction \\citep{wat07}. These results imply the existence of multiple YSOs toward G5.89. Indeed, five sub-mm dust sources (denoted as SMA1, SMA2, SMA-N, SMA-E, and SMA-S) were located in the close vicinity of the UC \\htwo region G5.89 \\citep{hun08}.  Here we present SMA observations of G5.89 in CO (2$-$1) and (3$-$2) at resolutions of $\\sim$3$\\arcsec$. We perform large velocity gradient (LVG) calculations to explore the physical conditions of the EHV outflowing gas. We then discuss the nature of the EHV components and their connection to the gas accelerating processes. ",
        "conclusions": ""
    },
    "1112/1112.5457_arXiv.txt": {
        "abstract": "If dark matter interacts with quarks or gluons, the mediator of these interactions is either directly accessible at the LHC or is so heavy that its effects are encoded in contact operators. We find that the self-consistency of a contact operator description at the LHC implies bounds on the mediator scale stronger than those found from missing energy searches.  This translates into spin-independent elastic scattering cross-sections at a level $\\lesssim 10^{-41}$ cm$^2$, with direct implications for the DAMA, CoGeNT, and CRESST-II anomalies.  We then carefully explore the potential of monojet searches in the light mediator limit, focusing on a $Z'$ model with arbitrary couplings to quarks and dark matter. We find that the Tevatron data currently provides the most stringent bounds for dark matter and $Z'$ masses below $100$ GeV, and that these searches can constrain models for the DAMA, CoGeNT, and CRESST-II anomalies only if the mediator can decay to a pair of dark matter particles. ",
        "introduction": "If dark matter (DM) couples to quarks or gluons, its interactions can be discovered or constrained by LHC and Tevatron data in missing transverse energy (MET) searches. This has an interesting connection to direct detection experiments since the same couplings are probed. There is therefore the possibility that a one-to-one correspondence can be established between MET rates at colliders and scattering rates at direct detection experiments. Indeed, the status of DM direct detection, with the three experiments DAMA~\\cite{Bernabei:2008yi}, CoGeNT~\\cite{Aalseth:2011wp}, and CRESST-II~\\cite{Angloher:2011uu} reporting excesses consistent with a 10 GeV DM candidate, cries out for an independent, astrophysics-free verification.  This has recently stimulated significant work aiming at a consistent explanation of these experiments~\\cite{Frandsen:2011ts,Hooper:2011hd,Foot:2011pi,Belli:2011kw,Farina:2011pw,Schwetz:2011xm,Fox:2011px,McCabe:2011sr,An:2011ck,Fornengo:2011sz,Natarajan:2011gz,arXiv:1110.2721,Kelso:2011gd,Arina:2011zh}.  A rather general way of addressing this problem is the construction of an ``effective theory\" of DM that can be probed at colliders~\\cite{Cao:2009uw, Beltran:2010ww,DM1,Fermilab,DM2,DM3,Coolest, Mambrini:2011pw,Lin:2011gj,Kumar:2011dr}. If the mediators of quark-DM or gluon-DM interactions are sufficiently heavy, then their effects can be safely integrated out and described by higher-dimensional contact operators. In direct detection experiments the elastic scattering of WIMPs on nuclei is well-described by a contact interaction so long as the mediator mass is above a few MeV. In contrast, colliders like the LHC probe much higher momentum scales, and the validity of the effective description comes into question when the new physics is lighter than a few TeV.  Here we will see that the very assumption that the DM interaction be consistently described by a contact operator up to LHC energies implies perturbative unitarity bounds on the mediator scale that turn out to be much stronger than those imposed by MET constraints. Importantly, this assumption is also shown to be at odds with both the DAMA and the uncontaminated (surface event-free) CoGeNT regions.     The implication of this result is that if DAMA, CoGeNT, and CRESST-II are indeed detecting DM, then the physics responsible for it must be directly accessible at the LHC and thus not generally amenable to a contact description. This gives hope to the prospect of confirming or ruling out the direct detection anomalies in the near term. To this end we analyze in detail constraints arising from monojet searches at the Tevatron and the LHC in models where the mediator is a new massive vector boson $Z'$ with arbitrary couplings to the standard model quarks and the DM. ",
        "conclusions": "We have seen that the very assumption that the DM-quark interaction remains contact at the LHC provides strong constraints on the scale of new physics. Specifically, perturbative unitarity arguments force the scale suppressing the higher dimensional operator to be above a few TeV. These bounds are notably stronger than those obtained from monojet searches at the LHC.  In light of these results, the DAMA favored region and the (uncontaminated) CoGeNT 90$\\%$ CL signal region \\emph{cannot} be explained by interactions that remain contact at the 7 TeV LHC.  We therefore relaxed the contact interaction assumption and analyzed in detail the constraints arising from monojet searches at the Tevatron and the LHC. We focused on models in which the DM-quark interaction is mediated by a light $Z'$ with generic couplings to the SM quarks and the DM, and presented the bounds in a way that our results are readily applicable to the reader's favorite scenario.  We find that the strongest bounds are currently provided by the Tevatron's CDF experiment for DM and $Z'$ masses below $\\sim100$ GeV, with the LHC taking over for heavier masses.  Translating our monojet bounds to limits on the DM-nucleon elastic scattering cross section we find that these searches can constrain models for the DAMA, CoGeNT, and CRESST-II excesses only if $Z' \\rightarrow \\overline{X} X$ is allowed and if the $Z'$ is a narrow resonance.   \\vspace{1.2cm}  \\begin{center} {{\\bf Note Added}} \\end{center}  After the submission of our arXiv paper, the preprint \\cite{Fox:2012ee} appeared with a discussion claiming that the unitarity constraint of the present paper is in practice unimportant for rather low cutoffs. This conclusion is based on the inspection of the $\\overline{X}X$ invariant mass distribution for a MET search, in the case of a contact operator with $\\Lambda \\sim 600$ GeV. From this the authors find that the number of events that ``violate unitarity\" (i.e., those events with $m_{\\overline{X}X} > \\Lambda/0.4 \\approx 1.5$ TeV) is negligible for light dark matter. However, as we clarified above, the unitarity bound implies that a consistent description requires that new physics appear at invariant masses \\emph{below} $1.5$ TeV. A more telling measure of the validity of the contact operator description would be a comparison between a contact operator and a UV-complete theory leading to the same $\\Lambda$. This comparison was already made in~\\cite{DM3} for monojets, where it was shown that in order for an $s$-channel completion to agree with the contact interaction prediction one should have couplings of order $g\\sim5-10$, in agreement with our Fig.~\\ref{glambda}. With such large couplings, yet higher dimensional operators are expected to become relevant and the monojet prediction is plagued by $O(1)$ uncertainties. We interpret this as a confirmation of our conclusions."
    },
    "1112/1112.6266.txt": {
        "abstract": "CCD photometric observations of the Algol-type eclipsing binary AT Peg have been obtained. The light curves are analyzed with modern techniques and new geometric and photometric elements are derived. A new orbital period analysis of the system, based on the most reliable timings of minima found in the literature, is presented and apparent period modulations are discussed with respect to the Light-Time effect (LITE) and secular changes in the system. The results of these analyses are compared and interpreted in order to obtain a coherent view of the system's behaviour. ",
        "introduction": "The system was observed during 6 nights in summer of 2010 at the Athens University Observatory, using a 20-cm Newtonian telescope equipped with the CCD camera ST-8XMEI and B and R Bessell photometric filters.  % ",
        "conclusions": "{\\bf (a):} The LC solution revealed a third light of $\\sim$7\\%, while the O--C analysis suggested a third body with a minimal mass of 0.57 M$_{\\odot}$. Assuming a MS nature with coplanar orbit and taking into account the Mass-Luminosity relation for dwarfs (L$\\sim$M$^{3.5}$)  the expected light contribution was found $\\sim$1\\%, which is much less than the observed one. However, if the third body orbits the eclipsing binary with an inclination of $\\sim$12$^{\\circ}$ (as a MS star), or it is more evolved, then the observed additional luminosity can be justified. {\\bf (b):} The orbital period secular change found, contrary to what it was expected (the mass should flow from the less massive to the more massive component), resulted in a decreasing period rate. This discrepancy could be explained with a mass loss rate from the system of 6.6$\\times$10$^{-7}$ M$_{\\odot}$/yr, perhaps due to strong stellar winds, or with a magnetic breaking mechanism or even with systemic angular momentum loss, which superimposes the expected mass transfer."
    },
    "1112/1112.3667_arXiv.txt": {
        "abstract": "% {% Condensations in the more than $10^6$\\,K hot corona of the Sun are commonly observed in the extreme ultraviolet (EUV). While their contribution to the total solar EUV radiation is still a matter of debate, these condensations certainly provide a valuable tool for studying the dynamic response of the corona to the heating processes. } {% We investigate different distributions of energy input in time and space to investigate which process is most relevant for understanding these coronal condensations. } {% For a comparison to observations we synthesize EUV emission from a time-dependent, one-dimensional model for coronal loops, where we employ two heating scenarios: simply shutting down the heating and a model where the heating is very concentrated at the loop footpoints, while keeping the total heat input constant. } {% The heating off/on model does not lead to significant EUV count rates that one observes with SDO/AIA. In contrast, the concentration of the heating near the footpoints leads to thermal non-equilibrium near the loop top resulting in the well-known catastrophic cooling. This process gives a good match to observations of coronal condensations. } {% This shows that the corona needs a steady supply of energy to support the coronal plasma, even during coronal condensations. Otherwise the corona would drain very fast, too fast to even form a condensation. } ",
        "introduction": "}   Loops in the corona still present us with numerous open questions. Besides the open fundamental problem of the coronal heating mechanism, it is not clear yet which processes can provide the high densities observed high in the corona, well above what would be expected by hydrostatic equilibrium \\cite[e.g.][]{Peres:1997,Aschwanden+al:2001}. One candidate for this is the thermal non-equilibrium that leads to \\emph{catastrophic cooling}. This process is initiated if the heating is very concentrated at the loop footpoints, and the heating at the loop top can no longer balance the losses from heat conduction and radiation: the loops starts cooling, radiation becomes more efficient, and a runaway process sets in \\citep{Kuin+Martens:1982}. This process is considered to be important in the formation of prominences \\citep{Antiochos+Klimchuk:1991,Karpen+al:2006}. The role of this \\emph{catastrophic cooling} for normal coronal loops has been studied in several numerical experiments that employ one-dimensional models for (semi-circular) loops. \\cite{Mueller+al:2003,Mueller+al:2004} have investigated the cyclic and chaotic behaviour for loops of different lengths and derived observable quantities such as extreme ultraviolet emission line spectra. They relate these events to ``coronal rain'' \\citep{Mueller+al:2005}. \\cite{Antolin+al:2010} proposes that ``coronal rain''  could be a marker of the coronal heating mechanism.   In a recent study \\cite{Klimchuk+al:2010} have been able to show that possibly the catastrophic cooling associated with the thermal non-equilibrium does not play a major role in active region coronal loops. Consequently, \\cite{Klimchuk+al:2010} argue that the heat input for active region loops is probably not concentrated near the loop's footpoints alone. The authors furthermore suggest that the heating is probably not steady.   The conclusions of \\cite{Klimchuk+al:2010} on the relevance of the process of thermal non-equilibrium are still under debate, as has become clear in discussions at the recent Loops- 5 workhop (June 2011).  When the loop cross-section is allowed to vary, and especially when the heating profile is made asymmetric, the nature of catastrophic cooling may be significantly different than concluded by \\cite{Klimchuk+al:2010}.  It may be premature to rule out thermal non-equilibrium for explaining some of the observed properties of coronal loops (Z. Miki\\'c, priv. comm.).   Three-dimensional magneto-hydrodynamic models (3D MHD) of the corona show a horizontally \\emph{averaged} exponential drop in the heating rate with altitude, with a scale height of the order of 5\\,Mm \\citep[e.g.][]{Gudiksen+Nordlund:2002,Gudiksen+Nordlund:2005a,Gudiksen+Nordlund:2005b}, and the heating is found to be intermittent in time and space \\citep{Bingert+Peter:2011}. The assumptions of the \\cite{Klimchuk+al:2010} study concerning the spatial and temporal distribution of the heat input are consistent with the \\emph{average} drop of the heating rate found in the 3D MHD models. This type of 3D MHD model can be considered realistic, since they successfully reproduce numerous properties observed in imaging and spectroscopy, such as the emission measure distribution or the average transition region Doppler shifts \\citep[e.g.][]{Peter+al:2004,Peter+al:2006,Zacharias+al:2011.doppler}. Because these 3D models predict a heating scale height greater than what is needed for catastrophic cooling, this is further evidence that \\emph{extreme} concentration of the heating rate towards the footpoints is not a \\emph{general} feature of coronal heating.  One can consider catastrophic cooling as a process that will happen only in a sub-volume of the corona, where the heating (for some time) is concentrated much more towards the loop footpoints than on average. These catastrophic cooling events are observed and seem to be a common feature. Examples of detailed observational studies of such events are the investigations by \\cite{Schrijver:2001:cooling} in a comparison of filtergrams in Lyman-$\\alpha$ and in the coronal band around 171\\,{\\AA}, by \\cite{deGroof+al:2005} when comparing the 171\\,{\\AA} band with H-$\\alpha$ data, or by \\cite{Antolin+Verwichte:2011} on the role of these condensations for loop oscillations.   Most recently, cooling events have been reported by \\citet{Kamio+al:2011.upflow} using data from the Atmospheric Image Assembly on the Solar Dynamics Observatory \\citep[AIA/SDO,][]{Lemen+al:2011}. In particular, \\citet{Kamio+al:2011.upflow} present cases with peculiar light curves: the 131\\,\\AA\\ and 171\\,\\AA\\ representing plasma at ${\\log}T\\,[\\mbox{K}]{\\approx}5.7$ and 5.9 show narrow single peaks in time. In contrast, the 193 \\,\\AA\\ and 211\\,\\AA\\ channels with a maximum contribution for plasma at ${\\log}T\\,[\\mbox{K}]{\\approx}6.2$ and 6.3 show enhanced emission for much longer times with a more complex light curve, including peaks before and after the peaks in the ``cooler'' channels. A more detailed presentation of this lightcurve and the comparison to our model can be found in \\sect{S:obs}. The lightcurve found by \\citet{Kamio+al:2011.upflow} is counter-intuitive for a situation where the plasma cools, and one would expect a brightening in the channels according to their temperature of maximum contribution. A simplistic model was presented by \\citet{Kamio+al:2011.upflow} to give a first attempt to explain this peculiar behaviour assuming a cooling at constant pressure.  In the present study we investigate the cooling of the plasma in order to understand the peculiar light curves found by \\citet{Kamio+al:2011.upflow}. We perform one-dimensional loop models starting from an equilibrium model that is hot enough that it does not show significant count rates in any of the above channels. We  then follow two scenarios: (1) a \\emph{catastrophic cooling} case where the heating is concentrated at the foot points, and (2) a \\emph{heating off/on} case where the heating is shut off completely for a while and the turned on again. In both cases the plasma will cool, but only the first case will be compatible with the observed cooling events.  After a short outline of the model in \\sect{S:model} we present the hydrodynamic results (\\sect{S:results}) and the observable AIA signatures (\\sect{S:signature}) for both cases, before we discuss a comparison to the observations of \\citet{Kamio+al:2011.upflow} in more detail in \\sect{S:obs} and conclude the paper. ",
        "conclusions": "}   We ran two types of models for cooling in coronal loops and compared them to observations. In the \\emph{heating off/on} model we simply shut down the heating, which leads to cooling and draining of the loop, and then resumed the original heating. In the \\emph{catastrophic cooling} model we kept the total heat input constant but concentrate the heating more towards the footpoints. This leads to a thermal non-equilibrium near the apex, and a condensation forms in a runaway process. While these two processes have been studied in the past, we put them here in the new context of observations at the solar limb in extreme ultraviolet passbands by AIA/SDO.   We found that the \\emph{heating off/on} model will not generate significant count rates to understand observed condensation events above the limb. In contrast, the \\emph{catastrophic cooling} model seems to provide a natural explanation for the observations of condensations, with a good match to the observed light curves. Not only are the observed count rates roughly matched, but also the single peaked nature of the lightcurves in ``cool'' channels (131\\,\\AA\\ and 171\\,\\AA), and the multiple peaks in the ``hot'' channels (193\\,\\AA\\ and 211\\,\\AA) are reproduced.   Good arguments have been presented that the process of thermal non-equilibrium cannot be used as a general process to understand all properties of coronal loops \\citep{Klimchuk+al:2010}. Nonetheless, this process is important in a subvolume of the corona, e.g.\\ where the coronal condensations form, and it is a valuable tool for investigating the dynamic response of the corona to the spatial distribution of the heat input into the corona.     The important conclusion from our study is that, even in condensation events in the corona, a constant supply of energy is needed to keep the coronal pressure. The \\emph{heating off/on} model fails, because the support of the corona is lost when shutting down the heating and the material drains very quickly. To see condensations in the corona, one has to keep the energy input at a sufficient magnitude to support the hot plasma further. In the  \\emph{catastrophic cooling} model, this is achieved by a higher concentration of the heating towards the loop footpoints, which reduces the heat input in the \\emph{high} corona but keeps the \\emph{total} amount of energy supply to the corona at a (more or less) constant level."
    },
    "1112/1112.1381_arXiv.txt": {
        "abstract": "We check how the change in surface conditions between the Sun and red giant branch stars changes the characteristic surface convection length scale to be used in models. We investigate the question in the case of the mixing length theory and of the phenomenology of full spectrum of turbulence. For the observational part, we rely on independent measurements of effective temperatures and interferometric radii of nearby red giants. We find that the local red giant branch cannot be explained taking into account the solar calibrated convective length scale. ",
        "introduction": "\\label{sec1}   In low mass stars the effective temperature and the radius estimate the efficiency of surface convection. The Sun and the red giant branch stars (hereafter RGBs) have very different surface conditions. A red giant surface gravity and energy flux are much smaller than the solar ones. The purpose of the work we present is to check how this changes the surface convection efficiency. First we build calibrated solar models using the mixing length theory (hereafter MLT) or the phenomenology of full spectrum of turbulence by Canuto, Goldman \\& Mazzitelli (1996) (hereafter CGM) for surface convection. Then relying on the same input physics, we build red giant models. The red giant models radii and effective temperatures are compared to an observational sample of 38 objects for which the radii are known directly through interferometry to better than 10 percent. Absolute luminosities and effective temperatures of these objects are also accurately known. In the next section we give the main ingredients of the models affecting the radius and effective temperature and address the solar calibration. In section \\ref{sec3} we describe the RGBs giant sample. Section \\ref{sec4} investigates the changes in characteristic convection length scale from the Sun to RGBs. We conclude in section \\ref{sec5}. ",
        "conclusions": "\\label{sec5}  We modelled the Sun and local RGB stars in order to check if the change in surface conditions implies a change of the characteristic convection length scale $\\rm \\Lambda$ for two local treatments of convection: the MLT and the CGM. In both cases we assume $\\rm \\Lambda=\\alpha H_p$. Therefore we do not consider the original version of the CGM where the characteristic convection length scale is the distance to the boundary with the region stable with respect to convection. At a given absolute luminosity $\\rm \\alpha$ tunes the position of the RGB in effective temperature and in radius. We have accurate data on absolute luminosities, effective temperatures and radii of all RGB stars we consider. The location of the coolest stars or equivalently the largest radii stars of the sample suggest a decrease in surface characteristic length scale with respect to its solar calibrated value. We have shown the decrease to be required for the MLT in the HR diagram and for the CGM in the luminosity radius diagram. However we could have inverted the diagrams with respect to the convection treatments, the result would have been similar. The reader will find many more details in \\cite{Piau11} where we also specifically address the three RGB stars of the sample with asteroseismic mass estimates. The combination of interferometric and asteroseimic data clearly opens up new perspectives in the understanding of stellar fundamental parameters and how they can be used to constrain stellar physics (\\cite{Huber11})."
    },
    "1112/1112.1512_arXiv.txt": {
        "abstract": "{Bow shocks can be formed around planets due to their interaction with the coronal medium of the host stars. The net velocity of the particles impacting on the planet determines the orientation of the shock. At the Earth's orbit, the (mainly radial) solar wind is primarily responsible for the formation of a shock facing towards the Sun. However, for close-in planets that possess high Keplerian velocities and are frequently located at regions where the host star's wind is still accelerating, a shock may develop ahead of the planet. If the compressed material is able to absorb stellar radiation, then the signature of bow shocks may be observed during transits. Bow-shock models have been investigated in a series of papers \\citep{2010ApJ...722L.168V, 2011MNRAS.411L..46V, 2011MNRAS.414.1573V, 2011MNRAS.416L..41L} for known transiting systems. Once the signature of a bow-shock is observed, one can infer the magnetic field intensity of the transiting planet. Here, we investigate the potential to use this model to detect magnetic fields of (hypothetical) planets orbiting inside the habitable zone of M-dwarf stars. For these cases, we show, by means of radiative transfer simulations, that the detection of bow-shocks of planets surrounding M-dwarf stars may be more difficult than for the case of close-in giant planets orbiting solar-type stars.} ",
        "introduction": "The interaction of a planet with the corona of its host star can give rise to the formation of shocks that surround the planet's magnetosphere. Similar to what occurs around the Earth and other planets in the solar system, bow-shocks may develop around exoplanets. This idea had recently been applied to explain the lightcurve asymmetry observed in the near-UV transit of the close-in giant planet WASP-12b.  \\citet{2010ApJ...714L.222F} observed that the near-UV transit lightcurve of WASP-12b shows an early ingress when compared to its optical transit. Such observations indicate the presence of asymmetries in the exosphere of the planet. In particular, \\citet{2010ApJ...722L.168V} suggested that this asymmetry could be explained by the presence of a shock surrounding the planet's magnetosphere.  To test this idea, \\citet{2011MNRAS.416L..41L} performed Monte Carlo radiation transfer simulations of the near-UV transit of WASP-12b. Their results support the hypothesis proposed by \\citet{2010ApJ...722L.168V} as it explains both the observed level of absorption and the time of the (early) ingress observed in the near-UV lightcurve of the planet.  An interesting outcome of this model is that one can probe the presence of planetary magnetic fields. Magnetic fields may provide an import shield against erosion caused by the impacting particles from the stellar wind. In the Earth, the presence of a magnetic field is associated with the development of life. M-dwarf stars are currently the main targets in searches for terrestrial habitable planets. Therefore, it is interesting to analyse whether shocks could be used to trace magnetic field of Earth-like planets orbiting M-dwarf stars.  The present paper is organised as follows. Section~\\ref{sec.shock} summarizes the general characteristic of the bow-shock model that was initially developed for the case of WASP-12b. We also provide an overview of the results obtained through radiative transfer simulations, and discuss how this model could be useful to probe for planetary magnetic fields in other extra-solar planets. In Section~\\ref{sec.mdwarf}, we investigate the potential for detecting bow shocks surrounding the magnetospheres of hypothetical transiting planets orbiting inside the habitable zones of M-dwarf stars.  Section~\\ref{sec.conclusion} presents our conclusions. ",
        "conclusions": "In this paper, we investigated the potential for detecting bow shocks surrounding the magnetospheres of transiting planets orbiting inside the habitable zones of M-dwarf stars. The idea that bow shocks may be present around exoplanets have now been developed in a series of papers \\citep{2010ApJ...722L.168V,2011MNRAS.411L..46V,2011MNRAS.414.1573V,2011MNRAS.416L..41L}. Due to the interaction of the planet with the wind of its host star, bow shocks may form surrounding the magnetosphere of planets, and, if observable, the shock can help constrain planetary magnetic fields.  In order to analyse whether shock could be formed on planets orbiting M-dwarf stars, we applied \\citet{2010ApJ...722L.168V}'s model to two hypothetical low-mass stars with spectral types M1.5 and M4.0. Earth-size planets were assumed to orbit inside the habitable zones of their host-stars and the stars were assumed to have a thermally driven-type wind. We showed that shock can be formed surrounding the planetary magnetosphere. For the case of a planet orbiting at $0.2$~AU from the M1.5 star, we showed that the shock forms at an angle closer to $\\theta = 90^{\\rm o}$, implying that the shock should be occulted from the observer during transit. If this is the case, the presence of the shock may not brake the optical lightcurve symmetry. For the case of a planet orbiting at $0.03$~AU from the M4.0 star, the shock forms at a smaller angle $\\theta = 71^{\\rm o}$.  However, for the shock to be detected by means of an early-ingress, it has to compress the stellar coronal material to a level of optical depth that can cause enough absorption during transit. For the wind characteristics we assumed, the density we computed in the shocked material does not seem to be significantly high to cause an observable early-ingress in the planet near-UV lightcurve. This result was confirmed by means of radiative transfer simulations. This implies that the method applied to close-in giant planets orbiting solar-type stars may not be as useful in detecting magnetic fields of planets orbiting inside the habitable zones of M-dwarf stars.  \\def\\aj{{AJ}}                   % \\def\\araa{{ARA\\&A}}             % \\def\\apj{{ApJ}}                 % \\def\\apjl{{ApJ}}                % \\def\\apjs{{ApJS}}               % \\def\\ao{{Appl.~Opt.}}           % \\def\\apss{{Ap\\&SS}}             % \\def\\aap{{A\\&A}}                % \\def\\aapr{{A\\&A~Rev.}}          % \\def\\aaps{{A\\&AS}}              % \\def\\azh{{AZh}}                 % \\def\\baas{{BAAS}}               % \\def\\jrasc{{JRASC}}             % \\def\\memras{{MmRAS}}            % \\def\\mnras{{MNRAS}}             % \\def\\pra{{Phys.~Rev.~A}}        % \\def\\prb{{Phys.~Rev.~B}}        % \\def\\prc{{Phys.~Rev.~C}}        % \\def\\prd{{Phys.~Rev.~D}}        % \\def\\pre{{Phys.~Rev.~E}}        % \\def\\prl{{Phys.~Rev.~Lett.}}    % \\def\\pasp{{PASP}}               % \\def\\pasj{{PASJ}}               % \\def\\qjras{{QJRAS}}             % \\def\\skytel{{S\\&T}}             % \\def\\solphys{{Sol.~Phys.}}      % \\def\\sovast{{Soviet~Ast.}}      % \\def\\ssr{{Space~Sci.~Rev.}}     % \\def\\zap{{ZAp}}                 % \\def\\nat{{Nature}}              % \\def\\iaucirc{{IAU~Circ.}}       % \\def\\aplett{{Astrophys.~Lett.}} % \\def\\apspr{{Astrophys.~Space~Phys.~Res.}}   % \\def\\bain{{Bull.~Astron.~Inst.~Netherlands}}    % \\def\\fcp{{Fund.~Cosmic~Phys.}}  % \\def\\gca{{Geochim.~Cosmochim.~Acta}}        % \\def\\grl{{Geophys.~Res.~Lett.}} % \\def\\jcp{{J.~Chem.~Phys.}}      % \\def\\jgr{{J.~Geophys.~Res.}}    % \\def\\jqsrt{{J.~Quant.~Spec.~Radiat.~Transf.}}   % \\def\\memsai{{Mem.~Soc.~Astron.~Italiana}}   % \\def\\nphysa{{Nucl.~Phys.~A}}    % \\def\\physrep{{Phys.~Rep.}}      % \\def\\physscr{{Phys.~Scr}}       % \\def\\planss{{Planet.~Space~Sci.}}           % \\def\\procspie{{Proc.~SPIE}}     %  \\let\\astap=\\aap \\let\\apjlett=\\apjl \\let\\apjsupp=\\apjs \\let\\applopt=\\ao"
    },
    "1112/1112.2505_arXiv.txt": {
        "abstract": "We review the high energy properties of Misaligned AGNs associated with $\\gamma$-ray sources detected by {\\it Fermi} in 24 months of survey.  Most of them  are nearby  emission low power radio galaxies (i.e FRIs) which probably have structured jets. On the contrary, high power radio sources (i.e FRIIs)  with GeV emission  are rare. The small number of FRIIs does not seem to be related to their higher redshifts.  Assuming proportionality between the radio core flux and the $\\gamma$-ray flux, several of them  are expected to  be bright enough to be detected above 100 MeV in spite of their distance. We suggest that beaming/jet  structural  differences  are responsible for  the detection rate discrepancy observed between FRIs and FRIIs. ",
        "introduction": "The Large Area Telescope (LAT\\cite{Atwood2009}) on-board the $\\gamma$-ray  ${\\it Fermi}$ satellite detected more than one thousand extragalactic sources\\cite{Ackerman2011} in two years of survey.  The majority are blazars [i.e. BL Lacs and Flat Spectrum Radio Quasars (FSRQs)], confirming that AGNs with the jet oriented in the direction of the observer are the brightest GeV sources. Only a handful of LAT objects have  different counterparts, i.e. Starburst galaxies (SBs), Narrow Line Seyfert 1 galaxies (NLSy1s) and Misaligned AGNs (MAGNs)\\footnote{MAGNs are Radio Galaxies and Steep Spectrum Radio Quasars, i.e. Radio Loud objects with steep radio spectra ($\\alpha>0.5$) and/or showing possibly symmetrical extension in radio maps.}.  The state of the art of extragalactic observations after 24 months of {\\it Fermi} activity is shown in Fig.~\\ref{F1} where the $\\gamma$ spectral slope ($\\Gamma_ {\\gamma}$) of blazars and non-blazar objects is plotted as a function of  the luminosity ($L_{>100~MeV}$).  In spite of their small number, the non-blazar $\\gamma$-ray emitters are extremely appealing, as they  offer a powerful physical tool in approaching  the high energy phenomena.  Starburst Galaxies, for example, are a  laboratory for investigating cosmic ray acceleration\\cite{sb1}. The high energy jet emission discovered in NLSY1s is questioning the paradigm according to which radio-loud AGNs are only hosted in elliptical galaxies\\cite{Fos011}. The detection of $\\gamma$-ray photons in MAGNs is invaluable in revealing  the jet structure complexity\\cite{Abdoa2010}.  \\begin{figure}[h] \\centerline{\\psfig{file=Grandi_1_f1.eps,width=6.0cm}} \\vspace*{6pt} \\caption{Spectral slope versus $\\gamma$-ray luminosity of the sources detected by the LAT in 24 months. The small group of SBs, NLSy1s and MAGNs  preferentially fill the plot region characterized by lower luminosities.\\label{F1} } \\end{figure} ",
        "conclusions": "Although MAGNs are a minority in the GeV sky,  their study is particularly fruitful in providing general insights into the jet structure. We find that more then $10\\%$ of the FRI radio galaxies of the radio sample obtained combining the 3CR, 3CRR, MS4 and 2Jy catalogs have $\\gamma$-ray associations. Their LAT detections seem to be  favored by the presence of different velocity zones in the jet. Only $1\\%$ of FRIIs, belonging to the same 3C-MS4-2Jy sample, are visible above 100 MeV.  The small number of detections does not seem to be due to their larger distances. More likely, it depends on less favorable  jet properties.   \\footnotesize"
    },
    "1112/1112.4609_arXiv.txt": {
        "abstract": "The ATLASGAL 870\\,$\\mu$m continuum survey conducted with the APEX telescope is the first survey covering the whole inner Galactic plane ($60^o>l>-60^o$ \\& $b<\\pm1.5^o$) in submm continuum emission tracing the cold dust of dense and young star-forming regions. Here, we present the overall distribution of sources within our Galactic disk. The submm continuum emission is confined to a narrow range around the galactic plane, but shifted on average by $\\sim$0.07\\,deg below the plane. Source number counts show strong enhancements toward the Galactic center, the spiral arms and toward prominent star-forming regions.  Comparing the distribution of ATLASGAL dust continuum emission to that of young intermediate- to high-mass young stellar objects (YSOs) derived from Spitzer data, we find similarities as well as differences. In particular, the distribution of submm dust continuum emission is significantly more confined to the plane than the YSO distribution (FWHM of 0.7 and 1.1\\,deg, corresponding to mean physical scale heights of approximately 46 and 80\\,pc, respectively). While this difference may partly be caused by the large extinction from the dense submm cores, gradual dispersal of stellar distributions after their birth could also contribute to this effect. Compared to other tracers of Galactic structure, the ATLASGAL data are strongly confined to a narrow latitude strip around the Galactic plane. ",
        "introduction": "\\label{intro}  Since the location of our solar system is within the Galactic disk, studying the Galactic structure of our Milky Way is always a challenging problem. Therefore, we cannot derive such comprehensive and intuitive pictures of our disk as extragalactic studies are able to do for other spiral galaxies (e.g., \\citealt{kennicutt2003,nieten2006,walter2008}).  Nevertheless, based on a diverse set of studies over all wavelengths, in the last few decades we have derived a reasonably comprehensive picture of our Galactic spiral structure (for recent work, see, e.g., \\citealt{benjamin2008,reid2009}). The Galactic plane has been observed in the optical/near-/mid-infrared bands (e.g., \\citealt{dobashi2005,skrutskie2006,churchwell2009,carey2009}) as well as at longer wavelengths, e.g., in CO or cm continuum emission (e.g., \\citealt{dame2001,stil2006}). However, until the arrival of the two (sub)mm Galactic plane surveys ATLASGAL (The APEX Telescope Large Area Survey of the GALaxy at 870\\,$\\mu$m) and BGPS (Bolocam Galactic Plane Survey) \\citep{schuller2009,aguirre2011}, no survey at (sub)mm wavelengths existed that trace the cold dust emission stemming from dense and young star-forming regions at adequate spatial resolution (the COBE (Cosmic Background Explorer) and WMAP (Wilkinson Microwave Anisotropy Probe) data have too coarse resolution to isolate individual star-forming regions). Here, we employ the 870\\,$\\mu$m submm continuum survey ATLASGAL to study the general distribution of the dense dust and gas within our Galactic plane. ",
        "conclusions": "We present a study of the Galactic distribution of submm dust continuum emission from the northern and southern hemisphere. The submm continuum emission, which traces the dense gas emission from star-forming regions, excellently traces the structure within our Galaxy. Spiral arms and prominent structures of star formation are easily distinguished by significantly increased source counts toward these Galactic longitudes. The Galactic latitude distribution is skewed slightly below the Galactic plane, and the mean physical scale height is estimated to 46\\,pc. This scale height corresponds well to other star-formation tracers like CH$_3$OH and H$_2$O maser emission, and it is more confined to the Galactic plane than most of the other populations in our Galaxy. We compare the submm continuum emission with several other tracers of Galactic structure, in particular with the YSO population identified by Spitzer observations. The YSO population has a significantly larger scale height, and we propose that this may be attributed to a combined effect of extinction and dissolving of protostellar clusters after their birth."
    },
    "1112/1112.3029_arXiv.txt": {
        "abstract": "We present some results on the radiative signatures of the one zone hadronic model. For this we have solved five spatially averaged, time-dependent coupled kinetic equations which describe the evolution of relativistic protons, electrons, photons, neutrons and neutrinos in a spherical volume containing a magnetic field. Protons are injected and lose energy by synchrotron, photopair and photopion production. We model photopair and photopion using the results of relevant MC codes, like the SOPHIA code in the case of photopion, which give accurate description for the injection of secondaries which then become source functions in their respective equations. This approach allows us to calculate the expected photon and neutrino spectra simultaneously in addition to examining questions like the efficiency and the temporal behaviour of the hadronic models. ",
        "introduction": "Hadronic models for the high energy emission of Active Galactic Nuclei and other astrophysical non-thermal objects postulate that protons are accelerated by some process in the relativistic outflows connected with these objects. The high energy photon emission comes then either from proton synchrotron radiation \\cite{muckprot} or as a  result of secondary particles produced in photopion collisions -- for example, $\\pi^0$ decay produces very high energy $\\gamma$-rays, while it is possible that these high energy photons will be absorbed by soft photons, initiating thus  intense electromagnetic cascades before escaping from the source \\cite{mann}.  In the present paper we revisit the hadronic model by carefully modeling photopion production which is one of its key processes. Furthermore we do not adopt the usual approach which is to assume an {\\sl {ad hoc}} proton distribution in the source but we go one step back and derive this as a solution of an equation which includes proton injection, losses and escape. This approach allows self consistency and in addition it can adress some important points as is the efficiency of the hadronic models and their expected time variability. ",
        "conclusions": ""
    },
    "1112/1112.4862_arXiv.txt": {
        "abstract": "{  We report the variation of the spectral energy distribution (SED) of blazars as a function of source activity, based on available, simultaneous multi-wavelength (MWL) observations of BL Lac objects. We use a fully automatized $\\chi ^{2}$ minimization procedure, instead of the commonly used eye-ball fit, to model the data sets with a one-zone Synchrotron-Self-Compton (SSC) model. The obtained SSC parameters are then analyzed as a function of source luminosity, and the correlation between parameters is shown. Possibilities of improving the present observational and modeling status of BL Lac objects are also discussed.} ",
        "introduction": "Active Galactic Nuclei (AGNs) are believed to be the most powerful energy sources in the universe. It is usually assumed that at the AGN center there is a super massive black hole (SMBH) surrounded by the accretion disk: a beamed emission perpendicular to the disk is expelled in opposite directions in $\\sim$10\\% of the cases (radio-loud AGNs). The energy required for the beamed emission might be coming from the rotating black hole and the fully ionized accretion disk \\cite{ref1}. AGNs are phenomenologically classified according to their viewing angle to the observer. If the jet is pointing very close to the line of sight to the observer the source is called a blazar. Blazars include Flat Spectrum Radio Quasars (FSRQ) and BL Lac objects.  Blazar emission is dominated by a non-thermal continuum. The spectral energy distribution (SED) shows two peaks. For the subclass of high-peaked BL-Lac objects (HBLs) the lower-energy peak is in the UV/X-ray range and the higher-energy peak is in multi-GeV bands; the so called  low-peaked BL Lac objects (LBLs) have, instead, the low-energy peak in the IR/optical band and the high-energy peak in the multi-MeV band. The low-energy peak is due to synchrotron radiation. According to the leptonic model, the high-energy peak is produced by inverse Compton (IC) upscattering of lower-energy photons off relativistic electrons: in particular the so called synchrotron self Compton (SSC) model suggests that the synchrotron photons are upscattered by the same electron population responsible for the synchrotron emission \\cite{ref2}. ",
        "conclusions": ""
    },
    "1112/1112.3997_arXiv.txt": {
        "abstract": "We have performed deep imaging surveys for Lyman $\\alpha$ emitters (LAEs) at redshift $\\sim 7.3$ in two blank fields, the Subaru Deep Field (SDF) and the Subaru/XMM-Newton Deep survey Field (SXDF), using the Subaru/Suprime-Cam equipped  with new red-sensitive CCDs and a new narrow-band filter, NB1006 ($\\lambda_c=10052$\\AA, FWHM $\\Delta \\lambda=214$\\AA). We identified four objects as LAE candidates that exhibit luminosity excess in NB1006. By carrying out deep follow-up spectroscopy for three of them using Subaru/FOCAS and Keck/DEIMOS, a definitively asymmetric emission line is detected for one of them, SXDF-NB1006-2.  Assuming this line is Ly$\\alpha$, this object is a LAE at $z=7.215$ which has luminosity of $1.2^{+1.5}_{-0.6} \\times 10^{43} $ erg s$^{-1}$ and a weighted skewness S$_\\omega=4.90\\pm 0.86$. Another object, SDF-NB1006-2, shows variable photometry and is thus probably a quasar (QSO) or an active galactic nucleus (AGN). It shows an asymmetric emission line at $10076$\\AA\\, which may be due to either Ly$\\alpha$ at $z=7.288$ or $[$O {\\sc ii}$]$ at $z=1.703$. The third object, SDF-NB1006-1, is likely a galaxy with temporal luminosity enhancement associated with a supernova explosion, as the brightness of this object varies between the observed epochs. Its spectrum does not show any emission lines. The inferred decrease in the number density of LAEs toward higher redshift is $n_{Ly\\alpha}^{z=7.3}/n_{Ly\\alpha}^{z=5.7} = 0.05^{+0.11}_{-0.05}$ from $z=5.7$ to $7.3$ down to $L^{Ly\\alpha}=1.0 \\times 10^{43}$ erg s$^{-1}$. The present result is consistent with the interpretation in previous studies that the neutral hydrogen fraction is rapidly increasing from $z=5.7$ to $7.3$. ",
        "introduction": "Cosmic reionization is one of the most important events in the early history of the universe. It is the drastic phase transition of the intergalactic hydrogen medium (IGM) from neutral to ionized; therefore it is closely related to the formation and evolution of the first stars and galaxies. How and when this cosmic reionization process took place remains an issue of great interest in the studies of the early history of the universe and many observational studies have been made to probe early galaxies during the last decade \\citep{2011arXiv1110.0195I}. The latest measurements of the polarization of the cosmic microwave background (CMB) by the Wilkinson Microwave Anisotropy Probe (WMAP) constrained the optical depth of electron scattering during reionization, suggesting that the probable redshift of reionization, if it took place instantaneously, was $z=11.0 \\pm 1.4$ \\citep{2009ApJS..180..306D}. Measurements of Gunn-Peterson (GP) troughs (Gunn  \\& Peterson 1965) on spectra of high-z quasars sampled by Sloan Digital Sky Survey (SDSS) imply that reionization had been completed by $z\\sim 6$, when the fraction of IGM neutral hydrogen is estimated to be very small, $x^{z\\sim 6.2}_{HI}\\sim 0.01-0.04$ \\citep{2006AJ....132..117F}, where $x_{HI}$ is defined as $x_{HI} = n_{HI}/n_{H}$ with $n_{HI}$ and $n_{H}$ being the number density of neutral hydrogen and all hydrogen in the universe, respectively.  Another promising technique to probe the reionization process is to use Lyman $\\alpha$ emitters (LAEs). Narrow-band (NB)  (bandwidths of $\\sim 100-200$\\AA) imaging can efficiently isolate LAEs, which prominently emit Ly$\\alpha$ from their interstellar gas ionized by massive stars. Comparing Ly$\\alpha$ luminosity functions (LFs) of LAEs before and after cosmic reionization provides an approach to determine the epoch of end of reionization \\citep{1999ApJ...518..138H, 2004ApJ...617L...5M}, because Lyman $\\alpha$ photons are sensitive to neutral IGM abundances. In the reionization era,  Lyman $\\alpha$ photons would be scattered at higher redshifts due to the damping wing from the surrounding neutral IGM. Therefore, the observed abundance of LAEs at this epoch should vary with the neutral fraction of IGM hydrogen. Using LAEs to probe reionization is sensitive to neutral IGM fractions $x_{HI}>10^{-3}$. This is an advantage over the GP test which is not sensitive at $x_{HI}>10^{-3}$. LAE imaging surveys are a good observational method to probe the epoch of the cosmic reionization, although it is difficult to distinguish the effects of reionization from those of galaxy evolution.    Due to their very low number density, wide field deep survey observations are essential to construct a statistical sample of high-redshift LAEs. Wide field coverage also allows us to overcome the uncertainties of cosmic variance, as well as potentially large cosmic variance. The Subaru Prime Focus Camera \\citep[Suprime-Cam]{2002PASJ...54..833M} has ten MIT/LL $2048$ pix $\\times4096$ pix CCDs, located at the prime focus of Subaru Telescope, covering a $34\\arcmin \\times 27\\arcmin$ field-of-view with a pixel scale of $0\\arcsec.202$. The CCDs of Suprime-Cam were replaced to new fully depleted CCDs made by Hamamatsu Photonics in 2008 \\citep{2008SPIE.7021E..52K} (see the next section). Because the practical ranges of redshifts observable for LAEs are limited by OH airglow lines, some dedicated NB filters, corresponding to $z=3.1, 4.8, 5.7$ and $6.6$ were manufactured \\citep[e.g.,][]{2006ApJ...648....7K, 2011ApJ...734..119K, 2008ApJS..176..301O, 2010ApJ...723..869O, 2006PASJ...58..313S, 2005PASJ...57..165T}. Surveys performed using these NB filters revealed that the Ly$\\alpha$ LF exhibits almost no evolution from $z=3.1$ to $5.7$, whereas the Ly$\\alpha$ LF at $z=6.6$ shows an apparent deficit, at least at its bright end, compared to that observed at $z=5.7$, while the rest-UV LFs at $z=6.6$ and $5.7$ are comparable. This result strongly supports the suggestion that the decline in the LAE luminosity function at $z>6$ was caused by the attenuation of the IGM rather than by the galaxy evolution. The derived $L^*$ at $z=6.6$ is $30\\%-60\\%$ lower than that at $z=5.7$, indicating the luminosity attenuation toward higher redshift. On the other hand, \\citet{2010ApJ...725..394H} reported that $L^*$ is unchanged from $z=5.7$ to $6.5$ based on their spectroscopically sample. Then, when and how the end of cosmic reionization was completed is still under debate. Similarly, discovery of an LAE at $z\\sim7.0$, IOK-1 at $z=6.96$, was made using the NB973 filter ($\\lambda_c =9755$\\AA, $\\Delta \\lambda=200$\\AA) imaging \\citep{2006Natur.443..186I, 2008ApJ...677...12O}, revealing that the Ly$\\alpha$ luminosity density gradually decreases from $z=5.7$ even to $7.0$. In fact, the number density of $z=7$ LAEs is only $18-36\\%$ of the density at $z=6.6$. These results indicate a substantial transition in the IGM ionizing state during these early epochs. We might be observing the end of the cosmic reionization process.  To trace the neutral fraction as a function of redshift, it is necessary to survey for LAEs at higher-z. LAEs beyond $z=7$ have been surveyed by NIR instruments covering wide field-of-view and equipped with NB filters including the 4 m Mayall Telescope/NEWFIRM \\citep{2010A&A...515A..97H, 2010ApJ...721.1853T, 2012ApJ...745..122K}. They reported that the Ly$\\alpha$ LF has not evolved beyond $z=6$. These photometric LAE candidates have not yet been spectroscopically confirmed and possible presence of remaining contaminants cannot be ruled out.  Therefore, photometric constraining of LFs at $z>7$ is not yet robust, and spectroscopic confirmation of the Ly$\\alpha$ line is required.  Another effective method is to measure the fraction of Ly$\\alpha$ emitting galaxies among LBGs in order to understand the reionization process. Sizable LBG samples at $z\\sim 7$ found by the HST/WFC3 have been confirmed spectroscopically with the Keck/DEIMOS and the VLT/FORS2 \\citep{2010ApJ...725L.205F, 2012ApJ...744...83O, 2011ApJ...743..132P, 2012ApJ...744..179S, 2011ApJ...730L..35V}. Their results suggest that the Ly$\\alpha$ fraction at $z\\sim 7$ could decline from $z \\le 6$ \\citep{2010MNRAS.408.1628S,2011ApJ...728L...2S} and it might be caused by the increased neutral fraction at $z\\sim 7$. Hence, the current suggested redshift range of the end of the cosmic reionization from $\\sim 6$ to $\\sim 11$ indicates that the details of the cosmic reionization process remain obscure.    In this paper, we present the results of photometric and spectroscopic surveys for $z=7.3$ LAEs and use this sample to constrain the history of cosmic reionization. The remainder of this paper is organized as follows. First we describe the imaging survey for $z=7.3$ LAEs in two blank fields in \\S \\ref{sec_imaging_observation}. We also explain the selection criteria of $z=7.3$ LAE candidates based on narrow-band and broad-band photometry in \\S \\ref{sec_photometric_candidates}. Next, we present deep spectroscopic observations of our photometrically-selected LAE candidates in \\S \\ref{sec_spectroscopy}. We compare the estimated Ly$\\alpha$ luminosity and SFR of $z=7.3$ LAEs with those of $z=6.6$ and $5.7$ LAEs and discuss constraints on cosmic reionization in \\S \\ref{sec_discussion}. In the last section, we summarize our results.  Throughout the present paper, we adopt concordance cosmology with $(\\Omega_m, \\Omega_\\Lambda, h)=(0.3, 0.7, 0.7)$, \\citep{2007ApJS..170..377S} and AB magnitudes with $2\\arcsec$ diameter apertures for the NB1006 surveys unless otherwise specified. ",
        "conclusions": "We have carried out a systematic imaging and spectroscopic survey for LAEs at $z \\sim 7.3$, for the first time, in the SDF and the SXDF, using the Suprime-Cam equipped with new red-sensitive CCDs and a newly developed NB1006 filter. The total survey area and volume are $1719$ arcmin$^2$ and $5.94 \\times 10^5$ Mpc$^3$, respectively. Four $z\\sim 7.3$ LAE candidates were photometrically identified from the NB1006 imaging down to the NB limiting magnitude of $24.6-24.8$, and we obtained spectroscopy for three of them, SDF-NB1006-1, 2 and SXDF-NB1006-2 using Subaru/FOCAS and Keck/DEIMOS. We have detected a definitively asymmetric emission line from SXDF-NB1006-2. If this line is Ly$\\alpha$, it is a LAE at $z=7.215$. SXDF-NB1006-2 has a luminosity of $1.2^{+1.5}_{-0.6} \\times 10^{43} $ erg s$^{-1}$, a star formation rate $11$ M$_\\odot$ yr$^{-1}$ estimated from the Ly$\\alpha$ emission line, and a weighted skewness S$_\\omega = 4.90\\pm 0.86$.  The conclusions of this survey are summarized below.  \\begin{itemize} \\item Compared with the number density and SFR density of LAEs at $z\\le7$, The number density and SFR density of LAEs monotonically decreases from $z=5.7$ to $7.3$ while they do not change from $z=3$ to $6$, and the number density of rest-UV continuum at $z=7.3$ is comparable with the rest-UV cumulative LFs at $z=5.7$ and $6.6$. The derived decreases of the number density are $n_{Ly\\alpha}^{z=7.3}/n_{Ly\\alpha}^{z=5.7} = 0.05^{+0.11}_{-0.05}$ from $z=5.7$ to $7.3$ and $n_{Ly\\alpha}^{z=7.3}/n_{Ly\\alpha}^{z=6.6} = 0.16^{+0.38}_{-0.16}$ from $z=6.6$ to $7.3$ down to $L^{Ly\\alpha} = 1.0 \\times 10^{43}$ erg s$^{-1}$. This trend is consistent with the results of the previous surveys for LAEs at $z\\le7$, and it might be caused by IGM attenuation. \\item The magnitude of SDF-NB1006-2 in the NB1006 images varies and it might be a AGN/QSO at $z=7.3$ or $1.7$ rather than a LAE, in spite of the detected asymmetric emission line (S$_\\omega = 4.59\\pm 1.25$). The third object, SDF-NB1006-1, is likely a galaxy with temporal luminosity enhancement associated with a QSO, an AGN, or a supernova explosion, as the brightness of this object varies between the observed epochs. Its spectrum does not show any emission lines. The magnitude of the candidates in the SXDF, SXDF-NB1006-1 and 2, did not show any significant variation, though the time span to check the variability in SXDF was not long enough to completely trace the typical variability. Hence, it is important to the check whether the candidates are transient objects or not by dividing photometric images into several epochs.  \\end{itemize}  In order to accurately measure the neutral fraction and reveal the relationship of LAEs and LBGs in the evolution of galaxy at $z \\ge 7$ in more detail,  multi-epoch surveys using next generation facilities such as the Subaru/Hyper Suprime-Cam (HSC) that covers $1.5$ deg FoV in diameter,  equipped narrow band filters, are needed to establish larger LAE sample."
    },
    "1112/1112.5443_arXiv.txt": {
        "abstract": "{Previously it has been shown that MOdified Newtonian Dynamics (MOND) can be tested near the saddle points of the Newtonian gravitational potential using the forthcoming LISA Pathfinder (LPF) mission.  These analysis focused on one particular formulation of the MONDian theory, which centered around a particular modified Poisson equation. We show that, in addition to the well known AQUAL formulation, another possibility exists in the form of a driven Poisson equation for the MONDian field.  We consider similar quantitative and qualitative analysis for this formulation and also investigate typical Signal to Noise Ratios (SNR) resulting from these theories for a LPF test. We demonstrate that for a typical 50 km saddle flyby, the SNR is amplified by 25\\% between the two formulations.  We also suggest that the SNR remains enhanced for impact parameters of 1000km or larger.  Finally we show that constraints from a negative result remain as good, but no better between different formulations.  %  } ",
        "introduction": "In the past few decades, attempts to reconcile anomalous observations of galactic dynamics without resorting to dark matter have been one of the driving forces behind modified gravity theories. These theories have centred around using a prescription for a modified force law~\\cite{Milgrom:1983ca}, which suggests the usual Newtonian dynamics at `large' accelerations and modified dynamics at `small' accelerations (large and small here are relative to the Milgrom characteristic acceleration $a_0 \\approx 10^{-10}$ ms$^{-2}$). In doing so, the observed flattening of galaxy rotation curves and the empirical Tully-Fisher relation can both be satisfied.     Recent work~\\cite{Bekenstein:2006fi,Bevis:2009et, Magueijo2011} on an experimental observation of MONDian like theories suggests that anomalously large tidal stresses should be present around gravitational saddle points. The proposed LISA Pathfinder (LPF) mission~\\cite{LISA} should be capable of probing acceleration scales upto $a_0$, this is of special importance because here we would be probing the onset of MONDian behaviour, rather than the acceleration scales dominated by MOND. The purpose of this short paper is extend previous analysis done of the MOND saddle point science case, based upon a scenario where a mission extension is granted. The extension would involve redirecting the spacecraft from Lagrange point L1 to a saddle of the Earth-Moon-Sun system~\\cite{companion} once its nominal mission at L1 is completed. Previous analysis have only considered only non relativistic limits resulting from theories such as TeVeS or ones which are AQUAL-like, here we will examine how these predictions vary using instead a different formulation of MOND.  The structure of this paper is as follows, first we review the non-relativistic limits of the various modified gravity theories in the literature. In section III we solve the Poisson equation for this theory numerically around the Earth-Sun saddle point.  We consider typical signal to noise ratios for detecting MONDian effects with LISA Pathfinder in section V, our main results being summarised in figures \\ref{fig:SNR cont} and \\ref{fig:SNR cont compare}.  Some more general concerns regarding these type II theories are considered in the last section and we will conclude with some thoughts on using all of these results constraining modified gravity theories. We also present some analytical solutions and detail our numerical methods in appendices. ",
        "conclusions": "In this paper we solved the non-relativistic limit of a different MONDian theory formulation (here classed Type II) for the saddle points in the two-body Earth-Sun system.  We probed the additional MONDian effects and forecasted for LPF flyby probable detection possibilities.  We suggest that the chances of detection of a MONDian signal are enhanced compared to Type I theories (eg those resulting from TeVeS), with typical SNRs for a 50km flyby rising from 28 to 35.  Another key result is that the boosted SNRs remain higher even as far out as $b \\approx 1000$km or more, enhancing observational prospects even for catastrophic impact parameters.  We also suggest that in the event no anomalous tidal stresses are observed, the experiment can put constraints on these theories, however they remain no better than in Type I, due to the nature of the regime we would be probing. Further work however remains to be done into ways to differentiate between MONDian formulations.  One significant difference here is that there are no curl type forces present, due to the linear nature of the governing equation.  Whilst a null result would place constraints on these theories, a positive one would do little to guide us towards a particular formulation or even form $\\mu$ or $\\nu$ function.  We are currently working on techniques to constrain or perhaps extract some features of the function from potential data obtained."
    },
    "1112/1112.2922_arXiv.txt": {
        "abstract": "The project Massive Unseen Companions to Hot Faint Underluminous Stars from SDSS (MUCHFUSS) aims at finding hot subdwarf stars with massive compact companions (massive white dwarfs $M>1.0\\,{\\rm M_{\\odot}}$, neutron stars or stellar mass black holes). The existence of such systems is predicted by binary evolution theory and some candidate systems have been found. We classified about $\\simeq1400$ hot subdwarf stars from the Sloan Digital Sky Survey (SDSS) by colour selection and visual inspection of their spectra. Stars with high velocities have been reobserved and individual SDSS spectra have been analysed. In total $201$ radial velocity variable subdwarfs have been dis\\-covered and about $140$ of them have been selected as good candidates for follow-up time resolved spectroscopy to derive their orbital parameters and photometric follow-up to search for features like eclipses in the light curves. Up to now we found seven close binary sdBs with short orbital periods ranging from $\\simeq0.21\\,{\\rm d}$ to $1.5\\,{\\rm d}$ and two eclipsing binaries with companions that are most likely of substellar nature. A new pulsating sdB in a close binary system has been discovered as well. ",
        "introduction": "A large fraction of the sdB stars ($\\simeq50\\%$) are short period binaries \\citep{maxted01,napiwotzki04} with periods ranging from only $0.07\\,{\\rm d}$ to more than $10\\,{\\rm d}$. Close binary sdBs are most likely formed by common envelope (CE) ejection \\citep{han02,han03}. However, it is difficult to determine the nature of the close companions in sdB binaries. Because most of them are single-lined, only lower mass limits have been derived from the binary mass functions, which are in general compatible with late main sequence stars of spectral type M or compact objects like white dwarfs. Only in rare cases (e.g. eclipsing systems) it is possible to distinguish between these two options.  Subdwarf binaries with massive WD companions turned out to be candidates for supernova type Ia (SN Ia) progenitors because these systems lose angular momentum due to the emission of gravitational waves and shrink. Mass transfer or the subsequent merger of the system may cause the WD to reach the Chandrasekhar limit and explode as a SN~Ia. One of the best known candidate systems for the double degenerate merger scenario is the sdB+WD binary KPD\\,1930$+$2752 \\citep{maxted00,geier07}.  \\citet{geier10a,geier10b} analysed high resolution spectra of single-lined sdB binaries. Because the inclinations of these systems are unknown, additional information is required to derive masses. \\citet{geier10a,geier10b} measured the surface gravities and projected rotational velocities. Assuming synchronised orbits the masses and the nature of the unseen companions was constrained. Surprisingly, some companions may be either massive white dwarfs, neutron stars (NS) or stellar mass black holes (BH). However, the assumption of orbital synchronisation in close sdB binaries was shown to be not always justified and the analysis suffers from selection effects \\citep{geier10b}. The existence of sdB+NS/BH systems is predicted by binary evolution theory \\citep{podsi02,pfahl03,yungelson05,nelemans10}. The formation channel includes two phases of unstable mass transfer and one supernova explosion and the fraction of those systems is consistently predicted to be about $1-2\\%$. If the companion were a neutron star, it could be detectable by radio observations as a pulsar. \\citet{coenen11} searched for pulsed radio emission at the positions of four candidate systems from \\citet{geier10b} using the Green Bank radio telescope, but did not detect any signals.  We started a radial velocity (RV) survey (Massive Unseen Companions to Hot Faint Underluminous Stars from SDSS\\footnote{Sloan Digital Sky Survey}, MUCHFUSS) to find sdBs with compact companions like supermassive white dwarfs ($M>1.0\\,M_{\\rm \\odot}$), neutron stars or black holes \\citep{geier11a,geier11b}. The same selection criteria that we applied to find such binaries are also well suited to single out hot subdwarf stars with constant high radial velocities in the Galactic halo like extreme population II and hypervelocity stars \\citep[see Heber et al. these proceedings; ][]{tillich11}. ",
        "conclusions": "The MUCHFUSS project aims at finding hot subdwarf stars with massive compact companions. We identified  1369 hot subdwarfs by colour selection and visual inspection of the SDSS-DR7 spectra. The best candidates for massive compact companions are followed up with time resolved medium resolution spectroscopy. Up to now orbital solutions have been found for eight single-lined binaries. Seven of them have large minimum companion masses compared to the sample of known close binaries, which shows that our target selection works quite well. However, it turns out that our selection strategy also allows us to detect low-mass companions to sdBs in very close orbits. We discovered an eclipsing sdB with a  brown dwarf companion and a very similar candidate system in the course of our photometric follow-up campaign. These early results encourage us to go on, because they demonstrate that MUCHFUSS will find both massive and substellar companions to sdB stars."
    },
    "1112/1112.0320_arXiv.txt": {
        "abstract": "Significant progress has been made in the last few years on understanding how supermassive black holes form and grow. In this paper, we begin by reviewing the spectral signatures of Active Galactic Nuclei (AGN) ranging from radio to hard X-ray wavelengths. We then describe the most commonly used methods to find these sources, including optical/UV, radio, infrared and X-ray emission and optical emission lines. We then describe the main observational properties of the obscured and unobscured AGN population. Finally, we summarize the cosmic history of black hole accretion, i.e., when in the history of the Universe supermassive black holes were getting most of their mass. We finish with a summary of open questions and a description of planned and future observatories that are going to help answer them. ",
        "introduction": "Astrophysical black holes come in a wide range of masses, from $\\gsim$3 $M_\\odot$ for stellar mass black holes \\cite{orosz03} to $\\sim$10$^{10}$ $M_\\odot$ for so-called supermassive black holes \\cite{lynden-bell69,lauer07}. The best evidence for the existence of a supermassive black hole can be found in the center of the Milky Way galaxy, where from dynamical studies, the mass of the Sgr A$^*$ source  was established to be $\\sim$4.4$\\times$10$^{6}$$M_\\odot$ \\cite{genzel97,ghez98}.  Evidence for the existence of supermassive black holes has also been found in other massive nearby galaxies \\cite{kormendy95}, mostly from resolved stellar and gas kinematics. For active galaxies, it has been possible to use the technique known as reverberation mapping \\cite{bahcall72,blandford82,peterson93}. From these observations, a clear correlation has been established between the mass of the central black hole and properties of the host galaxy such as stellar mass in the spheroidal component \\cite{marconi03}, luminosity \\cite{magorrian98}, velocity dispersion \\cite{ferrarese00,gebhardt00} and mass of the dark matter halo \\cite{ferrarese02}. The fact that such correlations exist, even though these components have very different spatial scales, suggests a fundamental relationship between black hole formation and galaxy evolution. Furthermore, it is now well established by simulations \\cite{springel05} that the energy output from the growing central black hole can play a significant role in the star formation history of the host galaxy. In particular, theory suggests that nuclear activity regulates star formation either by removing all the gas \\cite{hopkins06,menci06} or by heating it \\cite{croton06}. It is therefore obvious, that a complete study of galaxy evolution requires a comprehensive understanding of black hole growth.  Most current black hole formation models tell us that the first black hole seeds formed at $z$$\\gsim$15. While the exact mechanism for the formation of the first black holes is not currently known, there are several prevailing theories (see the comprehensive reviews by M. Rees \\cite{rees78} and M. Volonteri \\cite{volonteri10a} for more details). One of the most popular possibilities is that the first black hole seeds are the remnants of the first generation of stars, the so-called population III stars, formed out of primordial ultra-low metallicity gas. These black holes formed at $z$$\\sim$20 and have typical masses $\\sim$100-1,000~$M_\\odot$. This scenario has problems explaining the very high masses, of $\\sim$10$^9$M$_\\odot$, estimated for supermassive black holes in $z$$\\sim$6 optically-selected quasars \\cite{willott10b}. Alternatively, the first black holes could have formed directly as the result of gas-dynamical processes. It is possible for metal-free gas clouds with $T_{\\textnormal vir}$$\\gsim$10$^4$K and suppressed $H_2$ formation to collapse very efficiently \\cite{bromm03}, possibly forming massive black hole seeds with M$\\sim$10$^4$-10$^5$M$_\\odot$ as early as $z$$\\sim$10-15. If instead the UV background is not enough to suppress the formation of $H_2$, the gas will fragment and form ``normal'' stars in a very compact star cluster. In that case, star collisions can lead to the formation of a very massive star, that will then collapse and form a massive black hole seed with mass $\\sim$10$^2$-10$^4$M$_\\odot$ \\cite{devecchi09}.  Given the current masses of 10$^{6-9}$M$_\\odot$, most black hole growth happens in the Active Galactic Nuclei (AGN) phase \\cite{lynden-bell69,soltan82}. With typical bolometric luminosities $\\sim$10$^{45-48}$erg~s$^{-1}$, AGN are amongst the most luminous emitters in the Universe, particularly at high energies and radio wavelengths. These luminosities are a significant fraction of the Eddington luminosity --- the maximum luminosity for spherical accretion beyond which radiation pressure prevents further growth --- for a 10$^{8-9}$~M$_\\odot$ central black hole. A significant fraction of the total black hole growth, $\\sim$60\\% \\cite{treister10}, happens in the most luminous AGN, quasars, which are likely triggered by the major merger of two massive galaxies \\cite{sanders88}. In an AGN phase, which lasts $\\sim$10$^8$ years, the central supermassive black hole can gain up to $\\sim$10$^7$-10$^8$ M$_\\odot$, so even the most massive galaxies will have only a few of these events over their lifetime. Further black hole growth, mostly in low-luminosity (low Eddington rate) AGN, is likely due to stochastic accretion of cold gas, mostly in spiral galaxies \\cite{hopkins06a}.  According to the AGN unification paradigm \\cite{antonucci93,urry95}, a large fraction of these sources, $\\sim$75\\% locally, are heavily obscured by optically and geometrically thick axisymmetric material, which explains many of the observed differences among different types of active galaxies. In addition, luminosity \\cite{lawrence91} and cosmic epoch \\cite{treister06b} play a significant role. One constraint on the fraction of obscured AGN and its evolution comes from the spectral shape of the extragalactic X-ray ``background'' (XRB). Thanks to deep X-ray observations at E$\\lsim$10~keV performed by {\\it Chandra} and {\\it XMM-Newton}, a very large fraction of the X-ray background, $\\sim$80\\%, has been resolved into point sources \\cite{hickox06}, the vast majority of them AGN \\cite{mushotzky00}. Several studies, the first of them $\\sim$20 years ago \\cite{setti89}, have used a combination of obscured and unobscured AGN to explain the spectral shape and normalization of the X-ray background with overall good results. The latest AGN population synthesis models \\cite{treister09b,draper09}  assume an average ratio of obscured to unobscured AGN of $\\sim$3:1 locally, increasing towards lower luminosities and higher redshifts, as well as a fraction of Compton-thick sources (CT; $N_H$$>$10$^{24}$cm$^{-2}$) of $\\sim$5-10\\%, consistent with the value observed at higher energies, E=10-100~keV, of $\\sim$5\\% by INTEGRAL in the local Universe \\cite{treister09b,burlon11}, lower by factors of $\\sim$3 than expectations of previous population synthesis models \\cite{treister05b,gilli07}.  In this paper, we review multiwavelength methods used to trace the growth of SMBHs (\\S2), the known properties of unobscured and obscured AGN (\\S3 and \\S4 respectively), the cosmic history of black hole accretion (\\S5) and prospects for future observations (\\S6). Throughout this paper, we assume a $\\Lambda$CDM cosmology with $h_0$=0.7, $\\Omega_m$=0.27 and $\\Omega_\\Lambda$=0.73, in agreement with the most recent cosmological observations \\citep{hinshaw09}. ",
        "conclusions": ""
    },
    "1112/1112.1326_arXiv.txt": {
        "abstract": "The first high-resolution (R$\\sim$50,000) optical spectrum of the B$-$type star,  LS~III~+52$^{\\circ}$24, identified as the optical counterpart of the hot post$-$AGB candidate IRAS~22023+5249 (I22023) is presented. We report the detailed identifications of the observed absorption and emission features in the full wavelength range (4290$-$9015 \\AA) as well as the atmospheric parameters and photospheric abundances (under the Local Thermodinamic Equilibrium approximation) for the first time. The nebular parameters (T$_{\\rm e}$, N$_{\\rm e}$) are also derived. We estimate T$_{\\rm eff}$=24,000~K, log~g=3.0, $\\xi_{\\rm t}$=7 kms$^{-1}$ and the derived  abundances indicate a slightly metal-deficient evolved star with C/O$<$1. The observed P$-$Cygni profiles of hydrogen and helium clearly indicate on-going post$-$AGB mass loss. The presence of [N~II] and [S~II] lines and the non$-$detection of [O~III] indicate that photoionisation has just started. The observed spectral features, large heliocentric radial velocity, atmospheric parameters, and chemical composition indicate that I22023 is an evolved post$-$AGB star belonging to the old disk population. The derived nebular parameters (T$_{\\rm e}$=7000~K, N$_{\\rm e}$=1.2$\\times$10$^{4}$~cm$^{-3}$) also suggest that I22023 may be evolving into a compact, young low-excitation Planetary Nebula. Our optical spectroscopic analysis together with the recent Spitzer detection of double-dust chemistry (the simultaneous presence of carbonaceous molecules and amorphous silicates) in I22023 and other B-type post-AGB candidates may point to a binary system with a dusty disk as the stellar origin common to the hot post-AGB stars with O-rich central stars. ",
        "introduction": "The discovery of high Galactic latitude cool (F, G, K) and hot (O,B) post Asymptotic Giant Branch (post-AGB) supergiants - e.g., HD 161796 (Parthasarathy \\& Pottasch 1986) and LSII +34$^{\\circ}$26 (Parthasarathy 1993) - indicated that K, G, F, A, O, B post$-$AGB supergiants form an evolutionary sequence in the transition region from the tip of the AGB into the early stages of Planetary Nebulae (PNe). Since then, several cool and hot post$-$AGB candidates have been identified (Pottasch \\& Parthasarathy 1988; Parthasarathy \\& Pottasch 1989; Garc\\'ia$-$Lario et al. 1997a; Parthasarathy et al. 2000a). Gauba \\& Parthasarathy (2003, 2004) analysed the UV spectra and circumstellar dust envelopes of several hot post$-$AGB stars from the above lists, including LS~III~+52$^{\\circ}$24. Yet to date, the high-resolution (R$\\geq$30,000) optical spectra of only a few hot post$-$AGB stars have been studied (Sarkar et al. 2005 and references therein). In order to unveil the evolutionary origins, atmospheric parameters, and chemical compositions of hot post-AGB stars, there is a clear need to carry out detailed spectroscopic studies of more examples of this exotic class of post-AGB stars.  The optically bright B$-$type star, LS~III~+52$^{\\circ}$24 identified with the IR source IRAS~22023+5249 (hereafter, I22023) has far$-$IR colors similar to PNe (see Table 1). It is listed in Wackerling's (1970) Catalog of early$-$type emission$-$line stars. Recent ground$-$based high spatial resolution images in the near-IR have shown H$_{2}$ emission arising close to the central star, possibly in an incipient bipolar morphology (Volk et al. 2004). Indeed, Kelly \\& Hrivnak (2005) show that the excitation mechanism of the H$_{2}$ emission in I22023 is a combination of radiative (fluorescence) and thermal (shock) excitation. Gauba \\& Parthasarathy (2004) reported the presence of weak amorphous (10.8$\\mu$m) and  crystalline (33.6$\\mu$m) silicate features in the I22023's Infrared Space Observatory (ISO) spectrum and classified the object as O-rich. However, more recent and higher sensitiviy Spitzer/IRS spectra show that I22023 display a mixed-chemistry (both C-rich and O-rich dust features) with the presence of the classical aromatic infrared bands (AIBs; e.g., those at 6.2, 7.7, 8.6, and 11.3 $\\mu$m) together with broad 10 $\\mu$m amorphous silicate emission and a strong IR excess (Cerrigone et al. 2009).  In this paper, we explore the first high-resolution (R$\\sim$50,000) optical spectrum of I22023 in order to unveil its evolutionary status and chemical composition and to learn about the stellar origins of this peculiar type of hot post-AGB stars. In Sect. 2 we briefly describe the optical observations of I22023 and the data reduction process. A detailed analysis of the optical spectrum is presented in Sect. 3 while the photospheric and nebular analysis performed are shown in Sect. 4. We finish with a discussion and conclusions in Sect. 5. ",
        "conclusions": "Our analysis of the high-resolution (R$\\sim$50,000) optical spectrum of I22023 together with our detailed line identifications confirm that I22023 is a hot (B-type) O-rich post-AGB star (see below). That I22023 is not a normal population I B star is also suggested by the large heliocentric radial velocity of the star ($-$148.31 $\\pm$ 0.60 kms$^{-1}$)\\footnote{Note that the high radial velocity suggests that I22023 is an old disk low core mass post-AGB star and that there are old disk stars whose chemical composition is close to the solar composition (Furhmann 1998, 2004).}, as measured from the absorption lines present in its high-resolution optical spectrum. Thus, it is more likely that I22023 is a post$-$AGB star belonging to the old disk population. The presence of absorption lines of He~I, C~III, Si~III together with [N~II], [O~I], and [S~II] emission lines indicate a low-excitation nebula surrounding the early B$-$type central star. The observed P$-$Cygni profiles in the lines of hydrogen, He~I, C~II, and Fe~III clearly indicate the presence of a stellar wind with a significant post-AGB mass-loss rate, providing strong evidence for on-going post$-$AGB mass loss.  As a first approximation, using LTE analysis, we estimated T$_{\\rm eff}$ = 24,000~K $\\pm$ 1000~K, log~g = 3.0 $\\pm$ 0.5 and $\\xi_{\\rm t}$ = 7 $\\pm$ 1 kms$^{-1}$. The derived CNO abundances are compared with the average CNO abundances for main$-$sequence B$-$stars from the Ori OB1 association (Kilian 1992) in Table 6. The CNO abundances indicate that I22023 is an evolved star. We estimated C/O $\\sim$ 0.48, implying that the central star is O-rich and that the C/O ratio was not altered during the previous AGB phase. Our Si abundance estimate also suggests that I22023 is only slightly metal-deficient with [Si/H]=-0.12$\\pm$0.10. Our derived abundances can be easily explained if I22023 is the descendant of a low-mass (e.g., below $\\sim$1.5 M$_{\\odot}$) AGB star of roughly solar metallicity. Low-mass stars evolve very slowly and are expected to remain O-rich all the way along the AGB because they experience too few thermal pulses and the third dredge-up is too inefficient\\footnote{Note also that theoretical models predict a higher efficiency of the dredge-up in low metallicity atmospheres with respect to those with solar metallicity (e.g. Lugaro et al. 2003).} to modify the original C/O $<$ 1 ratio (see e.g., Herwig 2005 for a review). However, intermediate-mass (1.5 $<$ M $<$ 3-4 M$_{\\odot}$) AGB stars are converted to carbon and s-process enriched stars (see below) while massive (M $>$ 3-4 M$_{\\odot}$) AGB stars remain also O-rich as a consequence of the Hot Bottom Burning (HBB) activation and experience a completely different s-process nucleosynthesis (see e.g., Garc\\'{\\i}a-Hern\\'andez et al. 2006a, 2007a, and references therein). Note that the low-mass interpretation for I22023 would be consistent with the fact that I22023 is an optically bright O-rich post-AGB star while more massive O-rich post-AGB stars (which experienced HBB in the AGB) usually are completely obscured in the optical range by their thick circumstellar envelopes (e.g., Garc\\'{\\i}a-Hern\\'andez et al. 2007b). In short, attending to the high-resolution spectrum of I22023 only, we may conclude that I22023 is a low-mass O-rich post-AGB star.  Even though the IRAS cool (e.g., F, G) and hot (O, B) post$-$AGB stars show supergiant like spectra, indicating an evolutionary sequence in the transition region from the AGB to the PN stage, they seem to show fundamental differences in their chemical compositions. In the high Galactic latitude hot (O-, B-type) post$-$AGB stars, a severe carbon deficiency (i.e., C/O$<$1) is detected indicating that they left the AGB before the third dredge$-$up (e.g., Conlon et al. 1991; McCausland et al. 1992; Moehler \\& Heber 1998) has enriched the stellar surface with the products of the complex nucleosynthesis (e.g., carbon and heavy s-process elements such as Rb, Zr, Y, Sr, etc.) experienced during the AGB phase (see e.g., Garc\\'{\\i}a-Hern\\'andez et al. 2007a and references therein). A similar carbon deficiency is also detected in the hot post$-$AGB stars in globular clusters  (Moehler et al. 1998; Mooney et al. 2004; Jasniewicz et al. 2004; Thompson et al. 2006). The only exception to this observational evidence is the field hot post$-$AGB star, IRAS~01005+7910, which shows an overabundance of carbon (Klochkova et al. 2002). In contrast, among the IRAS selected cool (F-, G-type) post$-$AGB stars there is a majority of post-AGB stars in which a severe carbon deficiency is not detected. The F-, G-type post$-$AGB stars with the still unidentified 21 micron emission feature show an overabundance of carbon and heavy s$-$process elements (e.g., Van Winckel \\& Reyniers 2000), confirming that they have experienced s-process nucleosynthesis and the third dredge$-$up in the previous AGB phase and that they evolved from intermediate-mass carbon stars.  Interestingly, Cerrigone et al. (2009) found that I22023 is a double-dust chemistry - i.e., it displays the simultaneous presence of both C-rich and O-rich dust features - post-AGB star from their analysis of the recent Spitzer mid-infrared ($\\sim$5-40 $\\mu$m) spectrum of I22023. The mixed-chemistry is deduced from the detection of amorphous silicates emission at $\\sim$10 microns together with the classical aromatic infrared bands (AIBs; e.g., at $\\sim$6.2, 7.7, 8.6, and 11.3 $\\mu$m) usually attributed to carbonaceous compounds. The origin of the double-dust chemistry is still not very well understood and several scenarios, including the presence of a binary central system, a late thermal pulse on the AGB or post-AGB phases, HBB cessation by extreme mass loss, etc., have been proposed to explain the mixed-chemistry phenomenon observed in AGB stars (e.g., Garc\\'{\\i}a-Hern\\'andez et al. 2006b), post-AGB stars (e.g., Waters et al. 1998; Gielen et al. 2011), and PNe (e.g., Perea-Calder\\'on et al. 2009).  The presence of carbonaceous molecules in an O-rich environment such as that in I22023 (the central star is also O-rich!) is surprising and puzzling. Cerrigone et al. (2009) propose that the mixed-chemistry in I22023 (and other hot post-AGB stars) is due to the presence of a circumbinary disk/torus where O-bearing molecules would be preserved from the 3$^{rd}$ dredge-up, while the C-bearing molecules would be formed elsewhere in the outflow. The presence of a binary companion in I22023 and other hot post-AGB stars cannot be ruled out. Indeed, the presence of a close companion (at a distance of $\\sim$0.4\") in the proto-type hot post-AGB object Hen 3-1357 (the ``Stingray Nebula\") is well known (Bobrowsky et al. 1998). In addition, an spectacular incipient bipolar morphology is clearly seen in the HST images of Hen 3-1357. The Spitzer spectrum of Hen 3-1357 (see Perea-Calder\\'on et al. 2009) resembles that of I22023, showing amorphous silicates emission at 10 $\\mu$m together with a strong IR continuum but only a weak carbonaceous emission at 11.3 $\\mu$m is seen; there is a complete lack of the other AIBs at $\\sim$6.2, 7.7, and 8.6 $\\mu$m. In this context, the likely incipient bipolar morphology observed in I22023 (Volk et al. 2004) would support the presence of a binary companion.  On the other hand, the most recent idea to explain the mixed-chemistry phenomenon is from Guzm\\'an-Ram\\'{\\i}rez et al. (2011). These authors propose a chemical model able to form hydrocarbon chains in an UV-irradiated dense torus in order to explain the high detection rate of mixed-chemistry in PNe of the Galactic Bulge. However, the UV radiation field in I22023 (T$_{eff}$=24,000 K) is lower than that in double-dust chemistry PNe (with T$_{eff}$$>$34,000 K) and may be not intense enough to efficiently break the CO molecules. In addition, the Spitzer infrared spectrum of I22023 is very peculiar because the O-rich silicate dust is mostly amorphous and there is no clear evidence for the presence of crystalline silicate features at wavelengths longer than 20 $\\mu$m. This is in strong contrast with the Spitzer spectra of double-dust chemistry PNe (e.g., Perea-Calder\\'on et al. 2009; Guzm\\'an-Ram\\'{\\i}rez et al. 2011) where only crystalline silicates are detected.  An alternative explanation to explain the presence of carbonaceous molecules in I22023 may be non-equilibrium chemistry induced by shocks (Cherchneff 2011). Cherchneff (2011) demonstrates that water can form in C-rich evolved stars, showing that, independently of the stellar C/O ratio, thermal fragmentation of CO occurs in the hot post-shock gas. Our optical spectrum of I22023 shows clear evidences of on-going mass loss - i.e., the presence of a strong and variable stellar wind and shocks - which would support this carbonaceous molecules formation scenario. Indeed, other hot (B-type) post-AGB stars such as IRAS 20462$+$3416 and IRAS 19336-0400 are infrared spectroscopic twins of I22023, showing both an identical (mixed-chemistry) Spitzer spectrum (see Cerrigone et al. 2009) together with clear indications (e.g., P-cygni profiles) of on-going (and variable) mass loss (see Sanchez-Contreras et al. 2008; Arkhipova et al. 2011). In this scenario, the lack of strong infrared features from carbonaceous molecules in other hot and O-rich post-AGB stars such as IRAS 18062$+$2410 (or even the very young PN Hen 3-1357)\\footnote{No P-cygni profiles (i.e., strong stellar winds) are present in IRAS 18062$+$2410 (Arkhipova et al. 2007) and  based on the C IV 1550 \\AA~line in the IUE UV spectrum, the fast wind in Hen 3-1357 was stopped in 1995 (Parthasarathy et al. 1995).} would be related with the inactivity of strong stellar winds with significant mass loss rates; i.e., the absence of strong shocks activating non-equilibrium chemistry.  In summary, we speculate that the simultaneous presence of carbonaceous molecules and amorphous silicates in I22023 and other hot (B-type) post-AGB stars may point to a binary central system with a dusty disk/torus as the stellar origin common to the hot post-AGB stars hosting O-rich central stars. The episodic character of the stellar wind (shocks) and mass loss in these hot O-rich post-AGB stars would favor shock-induced non-equilibrium chemistry as the carbonaceous molecules formation scenario in these O-rich environments. Further monitoring studies (e.g., monitoring of radial velocity, light variations, strengths and profiles of emission and absorption lines) of this star and other hot post-AGB stars are encouraged in order to understand the circumstellar mixed-chemistry, mass loss rate (and evolution) with the ultimate goal of unveiling the stellar origin of this intriguing class of O-rich post-AGB objects."
    },
    "1112/1112.4818_arXiv.txt": {
        "abstract": "\\noindent We study the vortex lines that are a feature of many random or disordered three-dimensional systems. These show universal statistical properties on long length scales, and geometrical phase transitions analogous to percolation transitions but in distinct universality classes. The field theories for these problems have not previously been identified, so that while many numerical studies have been performed, a framework for interpreting the results has been lacking. We provide such a framework with mappings to simple supersymmetric models. Our main focus is on vortices in short-range correlated complex fields, which show a geometrical phase transition that we argue is described by the $CP^{k|k}$ model (essentially the $CP^{n-1}$ model in the replica limit $n\\rightarrow 1$). This can be seen by mapping a lattice version of the problem to a lattice gauge theory. A related field theory with a noncompact gauge field, the `$NCCP^{k|k}$ model', is a supersymmetric extension of the standard dual theory for the XY transition, and we show that XY duality gives another way to understand the appearance of field theories of this type. The supersymmetric descriptions yield results relevant, for example, to vortices in the XY model and in superfluids, to optical vortices, and to certain models of cosmic strings. A distinct but related field theory, the $RP^{2l|2l}$ model (or the $RP^{n-1}$ model in the limit $n\\rightarrow 1$) describes the unoriented vortices which occur for instance in nematic liquid crystals.  Finally, we show that in two dimensions, a lattice gauge theory analogous to that discussed in three dimensions gives a simple way to see the known relation between two-dimensional percolation and the $CP^{k|k}$ sigma model with a $\\theta$ term. ",
        "introduction": "\\noindent Ensembles of random geometric objects lead to some of the most subtle and interesting questions in critical phenomena. A geometric viewpoint is often useful in analyzing phase transitions: for example, one can view an Ising transition as a proliferation of domain walls, or an XY transition as a proliferation of vortices, leading to important dual descriptions \\cite{Kogut review, Jose, banks myerson kogut, peskin, thomas}. But in addition to this, some of the most intruiguing issues concern the critical behaviour of the geometrical objects themselves. In this paper we consider the universal statistical properties of \\emph{vortex lines} in random media in three dimensions, a subject which has received considerable numerical attention \\cite{tricolour percolation, Vachaspati Vilenkin, Scherrer Frieman, dennis, Strobl Hindmarsh}, but which has so far lacked a field theoretic formulation.  Geometrical critical behaviour can occur even in the absence of a conventional thermodynamic phase transition. The best-known example of this is percolation, in which nonlocal geometrical correlation functions -- such as the probability that two sites lie in the same cluster -- yield universal exponents, despite the fact that correlation functions of local observables remain trivial at the critical point \\cite{Essam}. In fact it is precisely because of the short-range correlations of the local degrees of freedom that the percolation universality class is ubiquitous: generic random media are short-range correlated.  Two-dimensional percolation can be thought of as a fluctuating soup of loops \\cite{Cardy SLE review}. On the lattice, these are the cluster boundaries; in the continuum, they are the zero lines of a random real height field $v(x)$. When $v(x)$ is biased towards either positive or negative values, the loops have a finite typical size, but these cases are separated by a critical point at which this lengthscale diverges and the loop ensemble becomes conformally invariant.  Vortex lines in three dimensions constitute an analogous problem. The most generic version of it concerns the statistics of the zero lines of a short-range correlated \\emph{complex} field $w(x)$. Again, these vortices form a fluctuating soup of oriented loops, with the orientation of a loop given by the sense in which the phase of $w(x)$ rotates around the vortex. Although the correlations of $w(x)$ are trivial by construction, there is an interesting geometrical phase transition as a function of bias -- the ratio of the average value of $w(x)$ to the width of its distribution \\cite{tricolour percolation, Strobl Hindmarsh}. As will be explained in detail in the following, there is an \\emph{extended} phase at small or zero bias \\cite{tricolour percolation, dennis, Vachaspati Vilenkin, Scherrer Frieman}, in which infinite vortex lines proliferate (and are found numerically to be Brownian), while at large bias there is a \\emph{localized} phase in which long vortices are exponentially suppressed. In between lies a critical point, where loops have a nontrivial fractal structure (and a fractal dimension close to $5/2$). Similar transitions can also occur in the absence of applied bias, for example in the vicinity of the critical point of the XY model, or for line defects in other order parameters: in particular, one can consider the unoriented ($\\mathbb{Z}_2$) vortices in nematics \\cite{kibble Z2, Strobl Hindmarsh}.    Universal aspects of vortex statistics have been examined numerically in a wide range of contexts, including cosmic strings, beginning with the simulations of Vachaspati and Vilenkin \\cite{Vachaspati Vilenkin, Strobl Hindmarsh, Scherrer Frieman, Scherrer Vilenkin, kibble Z2, Vachaspati}, line defects in random light fields, known as `optical vortices' \\cite{dennis, dennis topology}, vortices in the XY \\cite{Kajantjie, Bittner} and Abelian Higgs models \\cite{Wenzel, Hove, Prokofev Svistunov, Vincent}, in turbulent superfluids \\cite{superfluid turbulence}, and in a Potts-like model introduced to describe polymers \\cite{tricolour percolation}. In addition, similar scaling behaviour has been observed in simulations of lattice models of completely-packed loops \\cite{ortuno somoza chalker, Jaubert}.    Despite this previous work on vortices, numerous questions remain unanswered, because of the lack of a continuum formulation in which the relevant correlation functions can be expressed. In particular, it has not been clear how to classify these problems, or what the criteria are for two instances to be in the same universality class. It has been uncertain, for example, whether critical exponents for the geometrical vortex transition coincide with those for conventional three-dimensional (3D) percolation, or whether universal behaviour for oriented and unoriented vortices coincide (even in the extended phase), or how the geometrical transition for vortices in the XY model relates to the thermodynamic one. Even basic facts such as the appearance of Brownian statistics in the extended phase have not been derived, though previously this has been explained heuristically \\cite{Scherrer Frieman} by analogy with polymers in the melt, long known to be Brownian \\cite{polymers in melt}.  Our effort to address these questions starts from a set of new duality transformations. We begin by showing that the partition function for tricolour percolation, a Potts-like lattice model used for the numerical investigation of vortex statistics \\cite{tricolour percolation, Scherrer Frieman, Strobl Hindmarsh}, can be mapped exactly to that of a supersymmetric lattice gauge theory via a high temperature (graphical) expansion of the latter. By considering the continuum limit of this lattice gauge theory, and, separately, by discussing mappings from vortex problems that are formulated directly in the continuum, we arrive at field theories describing the geometrical properties of vortices.  For oriented vortices, these field theories involve a complex supervector $\\psib=(z^1,\\ldots z^{k+1}, \\chi^1,\\ldots \\chi^k)$, with $k+1$ bosonic and $k$ fermionic components, coupled to a $U(1)$ gauge field $A$. (Supersymmetry \\cite{McKane, Parisi Sourlas SUSY for loops} is a standard tool in problems involving loops or polymers, allowing  representation of geometrical correlation functions.) In order to describe the transition for vortices in short-range correlated complex fields $w(x)$, the gauge field $A$ must be taken to be compact, leading to a model in the universality class of the $CP^{k|k}$ sigma model \\cite{read saleur}. Heuristically, this can be thought of as the replica limit $n\\rightarrow 1$ of the $CP^{n-1}$ sigma model \\cite{di vecchia, Witten CPn-1}, which at $n=2$ is equivalent to the classical Heisenberg (or $O(3)$) model. The ordering transition in the sigma model describes the geometrical transition in the vortex ensemble. $CP^{k|k}$ field theories have been related to lattice models of completely packed loops in both 2D -- by Read and Saleur \\cite{read saleur} and Candu, Jacobsen, Read and Saleur \\cite{candu et al} -- and 3D \\cite{short loop paper, forthcoming long version}. (Strictly speaking we are abusing terminology, since there is a \\emph{separate} $CP^{k|k}$ model for each integer $k\\geq 0$, the model with $k=0$ being trivial. The correspondence with vortices holds for arbitrary $k$, but the range of geometrical correlation functions which can be expressed in the theory grows with $k$.)   The $CP^{k|k}$ description provides the appropriate framework for questions about vortex statistics. Immediate consequences include the explanation for Brownian statistics in the extended phase, which corresponds to the ordered phase in the sigma model language, and qualitative features of the critical point (for example, the fact that the fractal dimension is close to $d_f=5/2$ is easily understood as a symptom of confinement in the gauge theory). Most importantly, it clarifies the relations between the different universality classes, allowing us to classify these problems and explaining why various numerically studied models are in the same universality class and others are not. For instance, it explains the coincidence of exponents measured for the bias induced transition for vortices \\cite{tricolour percolation} and for the completely packed 3D loop models \\cite{ortuno somoza chalker, cardy class C review}, which also permit a mapping to the $CP^{k|k}$ model \\cite{short loop paper}.  In the absence of bias in the distribution for $w(x)$, the gauge field in the supersymmetric theory is non-compact,  yielding the non-compact $CP^{k|k}$ or $NCCP^{k|k}$ model, to borrow the terminology used for non-compact $CP^{n-1}$ models. The appearance of this theory can be simply understood by reference to a well-known duality which maps the XY model to an Abelian Higgs theory with a complex scalar field $z$ coupled to a gauge field \\cite{banks myerson kogut, thomas, peskin, dasgupta halperin, fisher XY duality, kleinert XY duality, zee}. Under this mapping, XY vortices become worldlines of the particles created by $z^*$ (to use a $2+1$ dimensional language). However, this standard dual description does not allow us to probe the long distance behaviour of vortices, because it does not allow us to write down the required geometric correlation functions. The $NCCP^{k|k}$ model is a supersymmetric extension of the Abelian Higgs model with a richer operator content (and possibly a richer phase diagram).  Unoriented vortices (e.g. in a short-range correlated random nematic order parameter) are described by a different field theory, the $RP^{2l|2l}$ model. This is essentially the replica limit of the $RP^{n-1}$ sigma model. A soft spin version of the $CP^{n-1}$ model in the replica limit allows us to clarify the relation between these two geometric problems, as well as the relation to 3D percolation, and also shows that the upper critical dimension for $CP^{k|k}$ is six, meaning that exponents can be calculated in a $6-\\epsilon$ expansion \\cite{Amit, McKane Wallace Zia, Bonfim, Senthil UCD}. The $CP^{1|1}$ model also describes Anderson localization in symmetry class C, which is closely related to classical loop models \\cite{mirlin evers review, glr, beamond cardy chalker, ortuno somoza chalker, cardy class C review}, so we expect six to be an upper critical dimension for this problem too. A similar conclusion has been reached previously by Senthil \\cite{Senthil UCD}, who considered the soft spin formulation of $CP^{1|1}$.   As mentioned above, another important conceptual connection is with 2D percolation. One of the continuum descriptions of this problem is as a $CP^{k|k}$ model with $\\theta$ term, as was shown via a mapping to a supersymmetric spin chain \\cite{read saleur, candu et al}. To connect this up with our 3D results, we give a treatment of 2D percolation as a lattice gauge theory (with an unconventional Boltzmann weight for the gauge field) analogous to our treatment of tricolour percolation in 3D. This leads in a simple way to the continuum description. A similar lattice regularization of $CP^{n-1}$ has been discussed by Affleck \\cite{Affleck CPn-1 and SOS}.  The structure of the paper is as follows. We begin in Sec. \\ref{overview} with a heuristic overview of the key ideas. In Sec. \\ref{short range correlations} we give duality mappings for oriented vortices in short-range-correlated random systems on the lattice (\\ref{tricolour percolation section}) and in the continuum (\\ref{vortex derivation}). For the lattice derivation, we use a specific model, tricolour percolation \\cite{tricolour percolation, Scherrer Frieman}, an elegant realization of the vortex problem which may be mapped to a lattice gauge theory. Sec. \\ref{statistics of vortex lines} gives the basic consequences of the field theory for vortex statistics.  Sec.~\\ref{Other universality classes} considers variations of the vortex problem. First we briefly describe the soft spin version of the $CP^{n-1}$ model in the limit $n\\rightarrow 1$ and the $6-\\epsilon$ expansion (Sec.~\\ref{upper critical dimension}). We then consider a perturbation leading to crossover to conventional percolation (Sec.~\\ref{crossover to percolation}). Here we also discuss the issue of vortex intersections which is important in many realistic vortex problems. Sec.~\\ref{unoriented vortices} considers unoriented vortices, occurring when the fundamental group of the order parameter manifold is $\\mathbb{Z}_2$. Sec.~\\ref{long range correlations} considers the stability of the universal behaviour of vortices in a short-range correlated $w(x)$ to the introduction of long-range correlations. We then consider vortices near the (thermodynamic) critical point of the XY model, their description in terms of the $NCCP^{k|k}$ model (Sec.~\\ref{XY critical point}), and the phase structure of this field theory. This concludes our discussion of vortices. In Sec.~\\ref{percolation} we consider loop models in 2D and their relation to lattice gauge theories with $\\theta$ terms. We close with a short summary in Sec.~\\ref{Summary}. ",
        "conclusions": "\\label{Summary}  \\noindent Line-like topological defects are ubiquitous in three-dimensional systems, and continue to be of theoretical interest; for example, an exciting recent development is the realization that in certain quantum systems vortices can sustain topologically protected zero-energy states \\cite{teo kane}. Here we have addressed the universal fractal geometry of vortices in disordered systems, a topic that has been studied numerically in diverse contexts but has lacked a field-theoretic description. By mapping a lattice model for vortices in a random complex field to a lattice gauge theory, and also via a continuum treatment based on XY duality, we related geometrical correlation functions for vortices to correlation functions in the $CP^{k|k}$ model, and explored simple consequences of this correspondence. For unoriented line defects, such as vortices in nematic fields, we argued that the appropriate field theory is the $RP^{2l|2l}$ model. These models play an important role in the classification of universality classes of geometrical critical phenomena in three dimensions.  We also argued that similar mappings can be fruitfully applied to 2D random curve ensembles, such as contour lines in a random height field, for which a lattice regularization is provided by site percolation and which are described by $CP^{k|k}$ with a $\\theta$ term. Future work will apply similar ideas to other random curve problems in 2D.  Although we have focussed on vortices, we know from the results of \\cite{short loop paper} that $CP^{k|k}$ and $RP^{k|k}$ models also apply to a class of three-dimensional loop models which undergo thermodynamically trivial geometrical phase transitions, and we expect these field theories to be generic descriptions for line defects that are topologically one-dimensional (and cannot branch, or terminate at a dangling end) or for deterministic walks in a random environment such as trajectories in a Lorentz lattice gas. Note that while at first sight the loop models seem remote from vortices, configurations in one of them (that on Cardy's `L' lattice \\cite{cardy class C review}) can be specified in terms of percolation configurations on two interpenetrating lattices, and thus in terms of two height fields (albeit on different lattices) taking values $\\pm 1$. Heuristically, then, there is a relation to zero lines of a complex (i.e. two-component real) function here too.  It is interesting that the distinction between compact and noncompact $CP^{n-1}$ models, much discussed recently in the context of deconfined criticality \\cite{NCCPn-1}, appears even in the replica limit $n\\rightarrow1$. It remains to be seen whether this actually gives a new universality class for vortices, or whether the geometrical transition described by $NCCP^{k|k}$ is inevitably separate from the thermodynamic (inverted XY) transition and reduces to the universality class of the compact $CP^{k|k}$ model (Sec.~\\ref{XY critical point}). On the other hand, we do believe that unoriented vortices (or loops in an appropriate loop model) will show a localized -- extended transition in a distinct universality class from that of oriented vortices -- it remains to perform detailed simulations for this universality class, and to confirm that distinct exponents are indeed obtained."
    },
    "1112/1112.3924_arXiv.txt": {
        "abstract": "When tracking a target particle that is interacting with nearest neighbors in a known way, positional data of the neighbors can be used to improve the state estimate. Effects of the accuracy of such positional data on the target track accuracy are investigated in this paper, in the context of dusty plasmas. In kinematic simulations, notable improvement in the target track accuracy was found when including all nearest neighbors in the state estimation filter and tracking algorithm, whereas the track accuracy was not significantly improved by higher-accuracy measurement techniques. The state estimation algorithm, involving an extended Kalman filter, was shown to either remove or significantly reduce errors due to ``pixel-locking''. {It is concluded that the significant extra complexity and computational expense to achieve these relatively small improvements are likely to be unwarranted for many situations.}{For the purposes of determining the precise particle locations, it is concluded that the simplified state estimation algorithm can be a viable alternative to using more computationally-intensive measurement techniques.} ",
        "introduction": "\\label{sec:intro} Dusty, or complex, plasmas consist of a low-density neutral/ion/electron plasma containing suspended ``dust'' -- negatively-charged macroscopic particles.  Dusty plasmas occur in space as well as in various terrestrial discharges ranging from lightning to industrial applications.\\cite{Merlino2004} In more controlled environments, dusty plasmas offer a unique testbed for exploring very complex and fascinating physical processes on a kinematic level. In all cases, precise determination of properties of the dusty plasma is a top priority.  {In particular, the industrial applications of dusty plasmas could benefit from being able to determine, and subsequently control, physical properties of a dusty plasma.}  Many properties of a dusty plasma can be inferred from the behavior of the dust particles, which can be non-invasively observed over time via laser illumination and a digital camera, as in \\Fig{fig:setup} which is used in experiments on Mach cones\\cite{Samsonov1999} and shock-waves\\cite{Samsonov2004}. The particles acquire a charge due to collisions with ions and electrons in the plasma, causing them to interact with each other and the plasma, with often fascinating results. For example, the dust can form crystalline structures or behave in a liquid-like phase (or exhibit phase transitions\\cite{Samsonov2004}), on timescales amenable to observation at the individual particle level. The precision of one's knowledge of the dust behavior is crucial to the reliability of any conclusions made about physical properties of the system\\cite{hadziavdic2006}.  There are two ways to improve the precision of determining the dynamics of individual particles: more sophisticated measurements, and information processing.  The performance of each are considered in this work.  Determining the tracks of multiple, interacting targets is a non-trivial task. Dust particles are typically observed using a video camera, their locations revealed using computer-based image-processing techniques, and their tracks (connected positions in time) formed using track-association algorithms. Such techniques and algorithms are available in varying levels of sophistication, with commensurate levels of accuracy, precision, resource consumption, etc. That is, more sophisticated techniques typically provide better precision and consume more resources such as time and computer memory.  This article considers tracking performance and, specifically, a trade-off between resources and accuracy/precision.  Motivation for such considerations comes from {the ultimate goal of implementing closed-loop control for such experiments, for which the tracking procedure must be automated.}  Previous work on tracking individual dusty plasma particles has been somewhat unsophisticated.  A common technique is particle tracking velocimetry (PTV), where the distance and direction of particle travel is obtained from consecutive images and used to aid in linking particle locations into tracks.\\cite{Samsonov2000,BoesseAiSR04}  This process can be automated, but more can be done. A higher level in sophistication for particle tracking is offered by the Kalman filter,\\cite{BarShalom,JFRopaedia} which combines a prediction model with measurements to obtain a recursive state estimation algorithm that can be optimal in the minimum-mean-square-error sense for linear dynamics. A simplified Kalman filter has been used on simulated data in the context of a dusty plasma crystal,\\cite{hadziavdic2006} where it was shown to perform very well in many situations. Slightly more sophisticated still is the state-estimation algorithm known as the Extended Kalman filter (EKF)\\cite{BarShalom,JFRopaedia}, which is considered in this work and will be introduced in \\Sec{sec:filter}.  \\begin{figure}[ht] {\\centering {\\includegraphics[width=0.95\\columnwidth]{fig1_experiment_schematic}} \\caption{\\label{fig:setup} Schematic of the dusty plasma experimental setup.\\cite{Durniak2010} The dusty plasma is contained wholly within a chamber as shown in (a), with the dust illuminated by a laser sheet and imaged from above by a digital camera as shown in (b). } } \\end{figure} ",
        "conclusions": "\\label{sec:conclusions} An extended Kalman filter (EKF) algorithm has been designed and implemented for tracking interacting dusty plasma particles.  {The algorithm was tested in virtual experiments (simulations) which allowed for quantification of performance aspects such as the precision of locating particle positions}. Attention was paid to generating images {that closely} resemble real experiments. The algorithm {was tested for a crystal-like phase of the dusty plasma.  This simple scenario provides a benchmark for measuring the performance of more sophisticated algorithms that will be necessary for treating more complex dynamics.} An analytical result showed the validity of an approximation to only nearest-neighbor interactions for this situation. Starting with relatively imprecise measurements of particle positions using the threshold and moment methods, it was shown that the filter algorithm improved the precision, as well as removing or significantly reducing errors due to ``pixel-locking''. The RMS values of position errors exceeded $0.1$ pixels for the measurements, which were improved by the filter to approximately $0.06$ pixels -- approaching the precision obtained in Ref.~\\citenum{Ivanov2007} ($\\approx0.03$ pixels) where a much more sophisticated and time-consuming measurement technique was used (local quadratic kernel method). This suggests that when such ``cheap'' measurements are used in conjunction with an EKF, they may be sufficient for obtaining excellent sub-pixel precision for particle positions -- a task that was previously {achieved through hardware modifications and/or} the aforementioned sophisticated measurement (image-processing) techniques. This particular result bodes well for eventually applying such a particle-tracking filter in real-time, {perhaps as part of a feedback loop.}  It is expected that the algorithm's {will perform quite well}{performance} in experimental situations {}{will be comparable to the results presented here}, but performance will degrade the further a dusty plasma deviates from a crystal-like state. Future work includes applying the algorithm to experimental images, extending the algorithm to cope with highly nonlinear dynamics such as shock-waves and Mach cones, and considering more advanced state estimation and measurement-track association techniques.  The present results will provide a useful performance benchmark for future work."
    },
    "1112/1112.1921_arXiv.txt": {
        "abstract": "{This paper presents new CCD $BVRI$ light curves of the newly discovered RS CVn eclipsing binary V1034 Her in 2009 and 2010, which shapes are different from the previous published results. They show asymmetric outside eclipse and we try to use a spot model to explain the phenomena. Using the Wilson-Devinney program with one-spot or two-spots model, photometric solutions of the system and starspot parameters were derived. Comparing the two results, it shows that the case of two spots is better successful in reproducing the light-curve distortions. For all the spot longitudes, it suggests that the trend towards active longitude belts and each active longitude belts might be switch. Comparing the light curves of 2009 and 2010, it indicates that the light curve changes on a long time scale of one year, especially in phase 0.25. In addition, we also collected the values of the maximum amplitudes of photometric distortion of the short-period RS CVn binary. We found for the first time that there is a trend of increasing activity with decreasing the orbital period. Finally, fitting all available light minimum times including our newly obtained ones with polynomial function confirmed that the orbital period of V1034 underwent up increase. \\\\ ",
        "introduction": "% \\label{sect:intro}  \\indent For V1034 Her (GSC 0983.1044, ROTSE 1 J165241.80+124905.2), it is a newly identified RS CVn-type eclipsing binary, which is also characterized by the light-curve asymmetries (Kaiser et al. 2002; Do$\\check{g}$ru et al. 2009; Ordway \\& Van Hamme 2004). Therefore, it is a very intriguing object to study stellar magnetic activity and test the effect of magnetic braking. However, it is poorly studied by astronomers, so we re-observed it in two different seasons to discuss starspot activity. \\\\ \\indent V1034 Her was discovered by Akerlof et al. (2000) as a 13th mag eclipsing binary with a period of 0.40763 day and an amplitude of 1.022 mag from the ROTSE 1 sky survey. Later, Kaiser et al. (2002) revised the standard magnitude $V=12.90$, and color indices $B-V~=~0.78$. They also derived the linear ephemeris of the system with the period of 0.8153 day and suggested that V1034 Her is a RS CVn type eclipsing detached binary due to the light-curves variations, which were explained using two-spots model on the secondary component. Later, Ordway and Van Hamme (2004) also obtained the $BVRI$ light curves of the system and analyzed them using Wilson-Devinney program with the spot model. Recently, Do$\\check{g}$ru et al. (2009) revised the ephemeris and considered firstly the possibility orbital period variation. The O-C diagram shows an upward parabola, which indicates a secular increase in the orbital period of the system. They also confirmed it the characters of RS CVn type and explained them in terms of two large dark starspots on the primary component.\\\\ \\indent As our on-going program of the study of RS CVn binary with multi-color photometry and high-resolution spectra (Gu et al. 2002, Zhang \\& Gu 2007, Zhang \\& Gu 2008, Zhang et al. 2010a, etc), we present new $B, V, R,$ and $I$ CCD light curves of V1034 Her in two different seasons and analyze them to study the properties of active regions using the version 2003 of Wilson-Devinney code. In addition, we accumulated the maximum amplitudes of photometric distortions of short-period RS CVn binary from the literatures to discuss the relation between the photospheric starspot activity and the orbital period.\\\\ ",
        "conclusions": "In this paper, our new light curves in BVRI bands were analyzed using the 2003 version of the Wilson-Devinney code with a spot model. New absolute physical parameters and starspot parameters were obtained. Secondly, we discussed the relation between photospheric starspot activity and the orbital period for the short-period RS CVn binary. Finally, the variation of the orbital period of V1034 Her were reanalyzed.\\\\ \\subsection{Photometric solution and starspot activity of V1034 Her} \\indent The marked asymmetry of the light curves of V1034 Her might suggest high-level surface activity, which was explained by starspot. From our results, we derived that the case of two-spots model being on the primary star is better successful in representing the distortions of the light curves. For the orbital parameters, the contribution of the primary component of V1034 Her to the total light is 0.71 in $B$, 0.67 in $V$, 0.65 in $R$ and 0.63 in $I$ band. Our new orbital parameters of luminosity ratios, the temperature of the secondary, and the orbital inclination are similar to those derived by Kaiser et al. (2002) and Do$\\check{g}$ru et al. (2009). However, the dimensionless potentials of the primary and secondary components $\\Omega _{1}$ and $\\Omega _{2}$ are a bit smaller than the results derived by Kaiser et al. (2002) and Do$\\check{g}$ru et al. (2009). These may be affected by the starspot. \\\\ \\indent It is well known that active-region longitude is the most reliable spot parameters determined by the traditional Light-curve method. For the result derived by Kaiser et al. (2002), the longitudes of two spots in 2001 are about $118^{\\circ}$ and $335^{\\circ}$, respectively. The spot longitude was transformed to the binary orbital motion system. Later, Do$\\check{g}$ru et al. (2009) derived that the longitudes of the two cool spots were $135^{\\circ}$ and $239^{\\circ}$ in 2006, respectively. For our starspot parameters, the longitude of the spot1 is $70^{\\circ}$ and the spot2 is about $334^{\\circ}$. Therefore, it suggested that active regions tend to appear in two active longitude belts $90^{\\circ}$ and $270^{\\circ}$, and each active longitude might be migrate in the orbital reference frame (see Fig.5). The active longitude phenomenon have also been found using photometry on many other active RS CVn binary systems, such as EI Eri (Berdyugina \\& Tuominen 1998), $\\sigma$ Gem (Henry et al. 1995), HK Lac (Olah Hall \\& Henry 1991), DV Psc (Zhang et al. 2010a) SV Cam (Zeilik et al. 1988) and RT And (Pribulla et al. 2000, Zeilik et al. 1989, Zhang \\& Gu 2007).\\\\ \\begin{figure} \\begin{center} \\includegraphics[angle=0,scale=1.0]{ms961fig5.eps} \\begin{scriptsize} \\caption{The spot longitudes from 2000 to 2009.}\\label{fig:5} \\end{scriptsize} \\end{center} \\end{figure} \\indent Comparasion of the LCs of 2009 and 2010, it indicates the spot1 becomes weaker at phase 0.2 and the spot2 becomes a bit stronger at phase 0.8 in 2010. Therefore, the spot activity changed on long time scale of one year. Indeed, Kaiser et al. (2002) also found the light curves differ on the order of 0.04 magnitudes at phase 0.9 during their observations in 2000 and 2001. These indicates starspot variation on a time scale of one year.\\\\ \\subsection{The orbital period - Activity relation for the short period RS CVn binary} For our observational light curves of V1034 Her, Max.II is around 0.04 mag in $B$, 0.07 mag in $V$, 0.06 mag in $R$ and 0.04 mag in the $I$ band brighter than Max.I, respectively. As you know, the light curve distortions are caused by dark photospheric spots for RS CVn stars. In order to detect the relation of the orbital period and the photospheric activity, we have also collected the values of the maximum amplitude of photometric distortions of other short-period RS CVn binary(Strassmeier et al. 1993; Eker et al. 2008). The values are listed in Table 3, which includes the object name, spectral type, the orbital period, the maximum distortion amplitudes in BVRI bands. For these values, the number of BV bands values are more than the RI bands. To detect the amplitude variation with the orbital period, they are plotted in Figure.6. As can be seen from Fig.6, there is a trend of increasing activity with decreasing the orbital period, and the trend are basically consistent with each other in the B and V bands (see Fig.6). To clarify the change, we also used the linear least-squares method to fit, and the linear formulas were obtained for B and V bands as follows: \\begin{eqnarray} The ~ B~ amplitude(mag)~=0.257(\\pm 0.067) \\nonumber -0.223(\\pm0.106) \\times P.\\\\ The ~ V ~ amplitude(mag)~=0.151(\\pm 0.052) \\nonumber -0.068(\\pm 0.083)\\times P.\\\\ \\end{eqnarray} These fits are also plotted in figure~6 with solid line. They are consistent with the rotation-activity correlation of RS CVn binaries by using chromospheric activity indicators (Montes et al. 1995; Montes et al. 1996), the radio luminosity (Gunn 1998), the extreme ultraviolet emission (Mitrou et al. 1997).\\\\  \\subsection{The period variation of V1034 Her} \\indent The variation of the orbital period of V1034 was reanalyzed on the basis of our new minimum times and those collected from the literatures. The quadratic term indicates that the orbital period of V1034 Her shows a continuous increase at a rate of $0.0107(\\pm 0.0006) second yr^{-1}$. The rate is similar to the result derived by Do$\\check{g}$ru et al. (2009). Since V1034 Her is a detached system, it is impossible that the mass transfer directly from the less massive component to the larger massive primary. So it is probable that the coronal mass flow from the less massive component to the large massive component by stellar wind (G$\\acute{a}$lvez et al., 2007). Because the photometric mass ratio is about 1, so it is difficult to tell the massive component is the primary star or the secondary component. Because the time range of (O-C) information is rather short - about 25 years, it is too early to decide about the character of the period variation. The upward parabola maybe a part of long periodic oscillation cause by a presumed third component or by magnetic activity cycle (Applegate 1992; Lanza et al. 1998). More new observations are needed to confirm it in the next 20-30 year. \\\\ \\begin{table} \\scriptsize\\tabcolsep 0.20truecm \\caption{The maximum distortion amplitudes of short-period RS CVn binary in B, V, R and I bands.} \\begin{tabular}{llllllll} \\hline \\hline \\multicolumn{1}{l}{Object} & \\multicolumn{1}{c}{Spectral type} & \\multicolumn{1}{c}{Period}  & \\multicolumn{4}{c}{The distortion amplitude (mag)}&\\multicolumn{1}{c}{Reference} \\\\ \\cline{4-7} &&(days)&B&V&R&I\\\\  \\hline RT And                     &F8-G0V+K1-3V&0.629    &0.095    & 0.075   &       &  0.05    & Pribulla et al., 2000; Zhang \\& Gu 2007                                                                 \\\\ V1034 Her                  &G5-9V+K1-4V &0.815    &0.040    & 0.070   & 0.06  &  0.04    & Kaiser et al., 2002; Do$\\check{g}$ru et al.,2009; my paper                                                           \\\\ DV Psc                     &K4V+        &0.309    &0.122    & 0.109   & 0.09  &  0.06    & Zhang \\& Zhang 2007; Zhang et al., 2010a                                          \\\\ XY Uma                     &G2-5V+K5V   &0.479    &0.292    & 0.251   & 0.12  &  0.10    & Pribulla et al., 2001; Yuan 2010                                                       \\\\ CG Cyg                     &G0V+K3V     &0.631    &0.170    & 0.140   & 0.05  &  0.05    & Heckert 1996; Af\\c{s}ar et al., 2004; Kozhevnikova et al., 2007a                          \\\\ BH Vir                     &F8V+G5V-M2V &0.817    &0.030    & 0.100   &       &          & Xiang et al., 2000; Kozhevnikova et al., 2007b      \\\\ WY Cnc                     &G5V+K7V    &0.830    &         & 0.110   &       &          & Kozhevnikova et al., 2007b                                                      \\\\ GSC03377-0296              &K3+         &0.422    &0.140    & 0.120   & 0.10  &          & Lloyd et al., 2007                                                              \\\\ DK CVn                     &K7V+M       &0.495    &0.159a   & 0.117a  & 0.10a &  0.08a   & Koff et al., 2005; Terrell et al., 2005                                                           \\\\ SV Cam                     &G0-G5V+K4V  &0.593    &0.150a   & 0.150   &       &          & Pribulla et al., 2003                                                          \\\\ UCAC3 295-68871            &            &0.461    &         &         & 0.10a &          & Solovyov et al., 2011                                                          \\\\ GSC 2038.0293              &G6-9+K1-3   &0.495    &0.090    & 0.070   & 0.06  &          & Bernhard \\& Frank 2006                                                           \\\\ UV Psc                     &G5V-8+K2-3V &0.861    &0.094    & 0.100   &       &          & Akan et al., 1988; Kjurkchieva et al., 2005                                    \\\\ ER Vul                     &G0V+G5V     &0.698    &0.080a   & 0.050   &       &          & Heckert 1991; Qian 2001; Pribulla et al., 2003                                                                   \\\\ BB Scl                     &K3-5V+K4V   &0.477    &         & 0.070   &       &          & Bromage \\& Buchley 1996; Watson et al., 2001                                                            \\\\ UV Leo                     &G0V+G2V     &0.600    &         & 0.050   &       &          & Kjurkchieva \\& Marchev 2007                                                    \\\\ BC Phe                     &G5-8V+K3-5V &0.649    &         & 0.170   &       &          & Cutispoto 1995                                                                 \\\\  \\hline \\end{tabular} The values are denoted by a mark ``a\", which are approximately calculated by the light curves of the literatures. \\end{table}    \\begin{figure*} \\begin{center} \\includegraphics[angle=0,scale=0.405]{ms961fig6a.eps} \\includegraphics[angle=0,scale=0.405]{ms961fig6b.eps} \\caption{The relations between the orbital period and photometric distortion amplitudes in the B and V bands. }\\label{fig:1} \\end{center} \\end{figure*}    \\normalem"
    },
    "1112/1112.4390_arXiv.txt": {
        "abstract": "A new adaptive mesh refinement (AMR) version of the \\zeusddd astrophysical magnetohydrodynamical (MHD) fluid code, \\azeus, is described.  The AMR module in \\azeus has been completely adapted to the staggered mesh that characterises the \\zeus family of codes, on which scalar quantities are zone-centred and vector components are face-centred.  In addition, for applications using static grids, it is necessary to use higher-order interpolations for prolongation to minimise the errors caused by waves crossing from a grid of one resolution to another.  Finally, solutions to test problems in 1-, 2-, and 3-dimensions in both Cartesian and spherical coordinates are presented. ",
        "introduction": "\\label{sec:intro} High-resolution, multidimensional simulations have become indispensable for many complex problems in astrophysics, particularly those involving (magneto-)fluid dynamics.  One of the most important innovations in this area has been the use of dynamic and variable resolution techniques.  Adaptive mesh refinement (AMR), pioneered in the context of the fluid equations by \\citet{bo84} and \\citet[\\citetalias{bc89}]{bc89}, is one such approach.  With AMR, a hierarchy of grids is used to provide high numerical resolution when and where the physics requires it, leaving as much of the volume at lower resolution as possible to minimise computational effort.  This makes AMR an efficient means of studying problems with a very large spatial dynamic range (\\eg, star formation, galaxy evolution), as borne out by the large number of codes which employ it: \\textsl{ORION} \\citep{k99}, \\textsl{FLASH} \\citep{f00}, \\textsl{RIEMANN} \\citep{b01}, \\textsl{RAMSES} \\citep{fht06}, \\textsl{PLUTO} \\citep{m07}, \\textsl{NIRVANA} \\citep{z08}, \\textsl{AstroBEAR} \\citep{cfvmj09}, and \\textsl{ENZO} \\citep{cxnll10} to name several.  Virtually all AMR fluid codes to date are based on a zone-centred grid, with all hydrodynamical variables (density, energy, and momentum components) taken to be located at the centres of their respective zones.  Indeed, AMR was originally designed specifically for zone-centred schemes.  Magnetohydrodynamic (MHD) solvers are designed with either zone-centred or face-centred magnetic field components, depending in part on the mechanism used to preserve the solenoidal condition.  One scheme that has enjoyed somewhat of a renaissance of late is Constrained Transport (CT; \\citealt{eh88}), which places magnetic field components at the centres of the zone-faces to which they are normal.  The staggered mesh introduced in such a scheme has to be specifically accounted for in the AMR modules and in such a way that $\\nabla \\cdot \\vec B$ remains zero everywhere---including within the boundaries---to machine round-off error.  The \\zeus family of codes are one of only a very few astrophysical fluid codes in use that employ a \\emph{fully} staggered grid (\\eg, \\textsl{STAGGER}; \\citealt{kritsuketal2011} and references therein), where the momentum components are also face-centred (Figure \\ref{fig:staggrid}).  While concerns have been expressed over the suitability of its MHD algorithms in certain pedagogical 1-D test problems (\\eg, \\citealt{f02}), the fact remains that in one form or another, \\zeus is among the best tested, documented, and widely-used fluid codes in astrophysics (\\citealt{sn92a,sn92b,smn92,c96,c07,c10,hayesetal06}), and a proper merger with AMR is warranted.  To do this, AMR has to be modified for a fully staggered grid, including the proper treatment of face-centred magnetic field and face-centred momentum.  \\begin{figure}[t!] \\begin{center} \\includegraphics[scale=0.5]{f1} \\end{center} \\caption{\\label{fig:staggrid} Depiction of a single zone on a fully-staggered grid, where scalars ($\\rho$, $e_{\\rm T}$, $e$, $p$) are zone-centred, primitive vectors ($\\vec{v}$, $\\vec{B}$) are face-centred, and derived vectors ($\\vec{E} = - \\vec{v} \\times\\vec{B}$, $\\vec J = \\nabla \\times \\vec B$) are edge-centred.} \\end{figure}  In this paper, we introduce the newest member of the \\zeus family of codes, \\azeus, whose ``maiden simulations\" have already appeared in \\citet{rc11}.  \\azeus is a block-structured AMR version of \\zeusddd \\citep{c96,c10} which preserves the modularity and structure of the underlying \\zeus module.  The AMR scheme of \\citetalias{bc89}, including the changes described in \\citet{bbsw94}, are modified for a fully-staggered grid with additional modifications made to the prolongation procedure to allow for smooth passage of all types of waves between adjacent grids of differing resolution.  \\azeus is currently capable of ideal MHD in 1-, 1.5-, 2-, 2.5-, and 3-D in Cartesian, cylindrical, and spherical polar coordinates using both dynamic and static grids, and with a full suite of physical boundary conditions.  As with all \\zeus-type codes, its operator-split design allows for additional physics (\\eg, gravity, viscosity, radiation, \\etc) to be added without concern over how such additions will affect the MHD algorithm.  How the non-hyperbolic additions are implemented for an adaptive grid is another matter (\\eg, radiation; \\citealt{wa11}).  This paper does not attempt to give a full recount of the basic methodology in either AMR or \\zeus, but focuses instead on the modifications to AMR (not so much to \\zeus) necessary for their merger.  Thus, the reader should be familiar with \\citetalias{bc89} and \\citet{c96,c10}.  In Section \\ref{sec:preamble}, we list the MHD equations solved by \\azeus, and define our conventions and notation.  In Sections \\ref{sec:restrict} and \\ref{sec:prolong}, we enumerate the modifications necessary for restriction and prolongation on a staggered grid, as well as outline the interpolation schemes used to allow the smooth passage of waves across grid boundaries.  Section \\ref{sec:bcs} focuses on boundary conditions, while in Section \\ref{sec:gcreate} we discuss how grids are created and how the proper nesting criterion must be modified for a fully-staggered grid.  Several of the 1-D, 2-D, and 3-D test problems used to validate \\azeus are given in Section \\ref{sec:tests}, followed by a quick summary in Section \\ref{sec:summary}.  Discussion of curvilinear coordinates, use of the vector potential, and a schematic overview of the code are relegated to the appendices. ",
        "conclusions": "\\label{sec:summary} We have described a method for block-based AMR on a fully-staggered mesh, and implemented this method in a new version of \\zeusddd called \\azeus.  In addition to describing the modifications required to AMR to account for the fully staggered grid, we also describe higher order interpolation methods for the prolongation step which we found necessary to allow for static grids.  Static grids are important for problems which, at first order, have a self-similar character and expand over the course of the simulation to ever larger scale lengths.  Such a simulation by \\azeus has already appeared in the literature \\citep{rc11}, which showcases the ability of the code to transmit waves of all types and strengths across grid boundaries, and to do so in cylindrical coordinates.  The higher-order prolongation operator is designed to maintain the conservation of all important physical quantities such as mass, momentum, energy, and magnetic flux.  Numerous test problems were also presented in 1-, 2-, and 3-D, and in both Cartesian and spherical polar coordinates.  These tests demonstrate the ability of the code to produce essentially identical results in ``reversible\" (non-turbulent) problems whether using a single grid or AMR, and give an example of the differences that can occur in ``irreversible\" (turbulent) problems as minute differences caused by the insertion or deletion of a grid amplify.  Finally, the \\azeus website \\href{http://www.ica.smu.ca/azeus}{\\texttt{http://www.ica.smu.ca/azeus}} was introduced on which test problems and simulations will be posted as they become available, and from which the code can be downloaded in the near future."
    },
    "1112/1112.2712_arXiv.txt": {
        "abstract": "The observed stellar mass function (SMF) is very different to the halo mass function predicted by $\\Lambda$CDM, and it is widely accepted that this is due to energy feedback from supernovae and black holes.  However, the strength and form of this feedback is not understood. In this paper, we use the phenomenological model GALFORM to explore how galaxy formation depends on the strength and halo mass dependence of feedback.  We focus on ``expulsion'' models in which the wind mass loading, $\\beta$ is proportional to $1/\\vdisk^n$, with $n=0,1,2$ and contrast these models with the successful Bower et al.\\  model (B8W7), for which $\\beta \\propto 1/\\vdisk^{3.2}$. A crucial development is that our code explicitly accounts for the recapture of expelled gas as the system's halo mass (and thus gravitational potential) increases. While models with high wind speed and mass loading result in a poor match to the observed SMF, a model with slower wind speed matches the flat portion of the SMF between $M_\\star \\sim 10^9-10^{11} \\hMsol$. When combined with AGN feedback, the model provides a good description of the observed SMF above $10^9 \\hMsol$. In order to explore the impact of different feedback schemes further, we  examine how the expulsion models compare with a further range of observational data, contrasting the results with the B8W7 model. In the expulsion models, the brightest galaxies are assembled more recently, and the specific star formation rates of galaxies decrease strongly with decreasing stellar mass. {\\rgb The expulsion models tend to have a cosmic star formation density that is dominated by lower mass galaxies at $z=1-3$, and dominated by high mass galaxies at low redshift. These trends are in conflict with observational data, but the comparison highlights some deficiencies of the B8W7 model also.  The experiments in this paper give us important physical insight to the impact of the feedback process on the formation histories of galaxies, but the strong mass dependence of feedback adopted  in B8W7 still appears to provide the most promising description of the observed universe.} ",
        "introduction": "A central issue in the study of galaxy formation is to understand the connection between the mass of galaxies and the mass of their associated dark matter haloes. This problem is not trivial: whilst the mass function of dark matter haloes predicted by the cold dark matter paradigm has a relatively steep slope ($d\\log N/ d \\log M \\sim -0.9$ over the range of haloes mass relevant to the galaxy formation), the observed mass function of galaxies is characterised by a shallow Schechter function. The abundance of galaxies is nearly independent of stellar mass over the interval $M_\\star \\sim 10^8 -10^{10.5} \\hMsol$, whilst at greater masses the abundance of galaxies declines exponentially. This raises two fundamental questions: i) why is the stellar mass function so flat below $\\sim 10^{10.5}\\hMsol$, when the abundance of haloes is such a strong function of mass, and ii) what physical processes are responsible for the exponential suppression of galaxies with mass greater than $\\sim 10^{10.5}\\hMsol$ \\citep{white_frenk91,benson2003}?  In this paper we investigate the astrophysical processes that map the halo mass function to the galaxy stellar mass function (SMF). There are two common approaches to this problem. Arguably, the most appealing route is to numerically integrate the set of differential equations that describe the rudimentary astrophysical processes (e.g. gravity, hydrodynamics, radiative cooling, star and black hole formation) using a \\emph{local} framework (e.g. a set of particles or a grid). This approach aims to evolve the cosmological density fluctuation power spectrum reflected by the cosmic microwave background using an {\\it ab initio} description of the problem. This system of equations commonly over-predicts the universal abundance of stars --- a problem commonly known as the `over-cooling problem' \\citep[e.g.][]{katz1996, balogh2001, keres2009, schaye2010}.  In order to get closer to a description of the observed universe, an additional set of processes (collectively known as ``feedback'') must be introduced to couple the energy, mass and metals returned by supernovae (SNe) and black holes (AGN) to the surrounding gas. {\\rgb As well as accounting for the stellar mass and gas content of galaxies, feedback enriches the intergalactic medium with metals, potentially resulting in an orthogonal set of observational constraints}. However, the wide range of scales involved in this problem (from AU scales to tens of Mpc) makes it infeasible to model these processes from first principles, forcing recourse to phenomenological, or `sub-grid', treatments (eg., \\citet{springel_hernquist03, springel2005, okamoto2008, schaye_dv08, booth_schaye09}).  Because the properties of model galaxies are remarkably sensitive to the details of sub-grid models, an alternative approach is to establish a set of equations that describes the same astrophysical processes on \\emph{macroscopic} scales, typically averaged over the physical scale of a galaxy. This approach is adopted by phenomenological, or `semi-analytic'\\footnote{The term ``semi-analytic'' commonly used in the literature, but is missleading. With modern computing power it is no longer critical that the resulting equations can be solved analytically. The important distinction is that this type of model provides a macroscopic description of the relevant processes. This approach is key to gaining insight into the problem. Such models are usually referred to as ``phenomenological'' in other areas of physics.} models, such as the GALFORM code considered here (\\citet{cole2000, bower2006}, hereafter Bow06), see also \\citet{kauffmann1993, hatton2003, delucia2006, somerville2008, guo2011}). For example, rather than computing the star formation rate at each point in a galaxy, such models typically compute the total rate of star formation over the entire galaxy, and assume that the rate is a simple function of global galaxy parameters (such as the gas mass, disk size and rotation speed). This leads to an alternative set of differential equations, that provide a macroscopic description of the physics.  A phenomenological model then solves these relatively simple equations within the merging hierarchy of structure formation. Macroscopic treatments are by definition approximate, but their use is ubiquitous in all branches of physics. When applied carefully and within an understood range of validity, the resulting descriptions can lead to significant physical insight.  In this paper, we apply the GALFORM model to seek an understanding of the mass-dependence of the equations of feedback from galaxies. We focus initially on the role supernova-driven winds play in establishing the galaxy stellar mass function on scales below $\\sim10^{10.5}\\hMsol$. Although strong suppression of galaxy formation in low mass haloes is clearly necessary, our understanding of the process remains incomplete, largely because of the difficultly of calculating the effects of feedback from fundamental physical principles. The difficulty primarily arises because the interstellar medium in which SNe explode is inherently multiphase (and magnetized), and does not behave like an ideal gas. Thus, accurately calculating the impact of even a solitary supernova is extremely challenging, as the result depends strongly on the initial density of the medium into which energy is injected. As a sequence of SNe explode, a network of low density channels (or `chimneys') is established that allow the supernova ejecta to escape into the halo of the galaxy \\citep{mckee_ostriker77, efstathiou2000, deavillez2007}.  However, without detailed calculations it impossible to estimate how the mass outflow rate and the specific energy of the outflowing material will depend on the star formation rate and the mass of the host galaxy.  In view of this uncertainty, all currently feasible simulations of galaxy formation parameterise the effect of supernovae and include it as a sub-grid calculation \\citep{springel_hernquist03, dv_schaye08}.  In hydrodynamical simulations, winds are most commonly modeled by adding thermal energy to gas particles (or cells), or by giving gas particles a velocity kick.  Although adding thermal energy seems a promising route, the energy injected is easily radiated away if the temperature of the particles is too low \\citep{katz1996}. Equally the injection of kinetic energy may drive thermodynamic shocks into the surrounding gas, and this thermal energy may also be radiated. These issues arise because the codes treat the complex multiphase interstellar medium (ISM) as a single fluid. Several schemes have been developed to circumvent the problem; one approach is to decouple the relevant particles from the numerical scheme for a period of time, either by preventing cooling for a period \\citep{brooks2007}, or by decoupling the particles from hydrodynamical forces \\citep{springel_hernquist03, okamoto_nemmen2008, oppenheimer2008}. Particles therefore retain the energy injected by feedback for a period of time, easing their escape from the dense interstellar medium of the galaxy. An alternative approach is to stochastically heat or kick the particles so that their energy is sufficiently high that their cooling time is long, or the shocks they drive are sufficiently strong that radiative losses are small \\citep{dv_schaye08, creasey2011}.  Because of the difficulty in interpreting the multi-phase nature of the outflow, the mass loading and velocity of winds are not yet strongly constrained by observation (but see \\citet{martin2005, weiner2009, chan2010, rubin2011} for recent progress). Most hydrodynamical calculations have therefore adopted the simplest possible model, in which the mass loading and velocity of winds are independent of the system in which the feedback event is triggered. This also simplifies implementation of the feedback, since there is no requirement to estimate the environment of the ISM on the fly, for example by determining the mass of the dark matter halo, or the local gravitational potential. An exception is the momentum scaling model of \\citet{oppenheimer2006}, where the wind parameters are set following a radiatively driven wind model \\citep{murray2005}, and a number of similar schemes implemented in the OverWhelmingly Large Simulations \\citep[OWLS][]{schaye2010}.  By contrast, most phenomenological (``semi-analytic'') models of galaxy formation assume that the parameters describing feedback adjust to ensure that the specific energy of outflows are matched to the binding energy of the halo. Conservation of total energy thus ensures that the mass loading of outflows is greater in dwarf galaxies than in larger systems \\citep{dekel_silk86}. Such models are partially motivated by arguments pertaining to the porosity of the interstellar medium: analytic models \\citep[e.g.][]{efstathiou2000} consider the formation of channels in the multiphase ISM and suggest that the porosity of the ISM is self-regulating and determined by the gravitational potential of the disc. In general, phenomenological models further assume that material expelled from the disc is recaptured on a timescale proportional to the dynamical time. This coefficient is allowed to be significantly larger than unity in some models in order to approximate the effect of expulsion of gas from the halo.  Current phenomenological models present a coherent picture for the formation of galaxies in a cosmological context, and provide an excellent explanation of many diverse data-sets. Although the GALFORM model has been developed to explain the observed properties of galaxies (eg., Bow06, \\citet{font2008, lagos2010}), it has also been shown to explain the X-ray scaling relations of groups and clusters \\citep[][hereafter Bow08]{bower2008} and the optical and X-ray emission from AGN \\citep{fanidakis2011}. The models can be used to generate convincing mock catalogues of the observable Universe \\citep{cai2009} and applied to test the procedures used to derive physical parameters from astronomical observations, and to identify the priorities for the next generation of astronomical instruments. These successes have been driven by the inclusion of two key components of the models: i) galaxy winds that scale strongly with halo mass and ii) a ``hot-halo'' mode\\footnote{This mode is often referred to as the `radio' mode, we prefer the term `hot-halo' as it emphasises that this type of feedback is assumed to only be effective when the cooling time of the halo is sufficiently long compared to the dynamical time. See \\citet{fanidakis2011} for further discussion.} of AGN feedback in which the cooling of gas in quasi-hydrostatic haloes is suppressed. By altering the parameterisation of the feedback schemes here, we investigate whether the Bow06 choice is optimal, or whether alternative schemes can similarly reproduce the properties of observed galaxies.  A major advantage of semi-analytic models is that they enable the effects of sub-grid parametrization to be explored quickly and easily \\citep{bower2010}. We exploit this aspect in this study, to explore the effect of various descriptions of outflows of material from galaxies. This requires us to generalise the GALFORM model to include the possibility that gas is expelled from the potential of dark matter haloes. Previous efforts to `calibrate' the semi-analytic method against numerical simulations have included gravity, hydrodynamics and cooling, but not effective feedback \\citep{benson2001, helly2003, delucia2010}. Part of the reason for this is the very different treatments of the winds from galaxies. Our extensions of the code allow us to bring the two approaches into closer alignment, and we include a brief comparison with the Galaxies-Intergalactic Medium Interaction Calculation \\citep[GIMIC][]{crain2009}. This provides a series of relatively high resolution hydrodynamic simulations, featuring relatively high spatial and mass resolution ($\\epsilon = 0.5\\hkpc$ and $m_{\\rm gas} = 1.4\\times 10^6 \\hMsol$ for the highest resolution realisations), and that trace a representative cosmological volume (four spherical volumes with comoving radius $18\\hMpc$ and one with comoving radius $25\\hMpc$). The simulations include radiative cooling \\citep{wiersma_schaye2009}, and sub-grid prescriptions for star formation and the thermodynamics of the ISM \\citep{schaye_dv08}, hydrodynamically-coupled supernova-driven winds \\citep{dv_schaye08} and metal enrichment resulting from stellar evolution \\citep{wiersma2009}. Some important successes include the X-ray scaling relations of $L_\\star$ galaxies \\citep{crain2010}, and the distribution of satellites and stellar halo properties of the Milky Way \\citep{deason2011, font2011}. Our modified GALFORM scheme implements similar physics, and we show that its behviour is very similar to GIMIC. This opens a new avenue, allowing us to use GALFORM to better understand how the parametrisation of feedback impacts upon the formation and evolution of galaxies, and thus guide both the development of sub-grid treatments in hydrodynamical simulations, and the interpretation of observational data.   The structure of this paper is as follows. In \\S2, we discuss the implementation of supernova-wind driven feedback schemes in semi-analytic models, and introduce a new scheme that allows gas and metals to be expelled from low mass haloes and later reaccreted when the binding energy of the halo has increased significantly. In contrast to many previous models, we do not assume that material is always reaccreted after a certain timescale, or that expelled material is always lost from the hierarchy. In \\S3, we present a comparison of different feedback scalings, focusing on the difference between schemes that scale the parameters describing winds with halo mass, and those that adopt fixed parameters.  In \\S4, we explore how feedback from supernovae can be combined with feedback from black holes in order to generate an exponential break in the mass function. In particular, we compare the effect of `hot halo' mode feedback with that of strong quasar-driven winds (similar to those considered by \\citet{springel2005}). We consider a number of additional observational constraints in \\S5. While several of the feedback schemes are able to reproduce the high mass part of the SMF, we show that the specific star formation rate and the downsizing of galaxy formation are important orthogonal constraints. We present a summary of our conclusions in \\S6. Except where otherwise noted, we assume a WMAP7 cosmology, $\\Omega_0=0.272, \\Lambda_0=0.728, \\Omega_b=0.0455, h_0=0.704, n_s=0.967, \\sigma_8=0.81$ (Komatsu et al. 2011). Throughout, we convert observational quantities to the $h$ scaling of the theoretical model so that stellar mass are quoted in $\\hMsol$ etc. ",
        "conclusions": "`Feedback' is a fundamental component of galaxy formation models, allowing us to understand the marked differences between the dark matter halo mass function and the galaxy stellar mass function (SMF). In order to explain the observational data, galaxy winds must strongly suppress the formation of low-mass galaxies. Phenomenological (or ``semi-analytic'') models, such as GALFORM, have been shown to provide a good description of the galaxy mass function together with its evolution, the specific star formation rates of galaxies and the contribution of galaxies of different masses to the cosmic star formation rate density. In the standard GALFORM model, these successes are achieved by adopting a feedback parameterisation that varies strongly with system mass, such that the speed of the outflow tracks the mass of the halo while the mass loading of the wind decreases strongly with halo mass. In contrast most hydrodynamical models, including the GIMIC simulations, adopt a wind that is independent of the system mass.  In this paper, we have modified the semi-analytic code to implement feedback schemes similar to those usually adopted in hydrodynamic models. Our scheme allows the wind speed and mass loading to be adjusted and the consequent effect on the stellar mass function of galaxies, its evolution and other galaxy properties to be determined.  We focus on three particular models. Two models have fixed wind parameters model: the pGIMIC model has high wind speed and modest mass loading similar to the GIMIC hydrodynamical simulations \\citep{crain2009}; the SW model has a slower wind speed (at the same mass loading).  We also consider a model (MS) in which the wind mass loading scales with the inverse of the disk circular velocity. This is similar to the momentum scaling models of \\citet{oppenheimer2008}, although we set the wind speed by keeping the total energy of the wind (rather than its total momentum) fixed across halo masses. Initially, we also considered a model with a $1/\\vdisk^2$ dependence of mass loading (ES), but found this could not reproduce the observed stellar mass function. We contrast these models with a model based on Bow06 and Bow08, but adapted to the WMAP7 cosmology used in this paper (B8W7). For each of these models we consider the effect of including a ``hot-halo'' mode of AGN feedback following the gas expulsion scheme of Bow08.  {\\rgb The main results are as follows: \\begin{itemize} \\item We find that the phenomenological description of feedback that we use in our code reproduces many of the trends seen in hydrodynamic simulations. This give us confidence that the phenomenological approach captures the key physical effects of galactic winds well and allows us to explore the parameter space of winds properties quickly and efficiently. Comparison of the phenomenological model and hydrodynamic simulation also allows us to better understand the key physical processes in galaxy formation. A model which uses a high wind speed and moderate mass loading (and does not include AGN feedback), reproduces the form of the SMF seen in the GIMIC numerical simulations well. Although the strong feedback suppresses the stellar masses of galaxies strongly, the effect is most prominent around $10^{10.5}\\hMsol$, so that the resulting SMF has a dip in the abundance of galaxies at these masses. At lower masses, the SMF rises steeply since star formation is too weakly suppressed, while, at higher masses the wind is not sufficiently energetic to escape the halo.  Using the phenomenological model, we are able to explore how the dip in the mass function depends on the assumed wind speed and mass loading.   \\item The observed mass function is better matched by adopting a relatively slow wind, as seen in the hydrodynamic simulations of Op08. The SW model, with a wind speed of $180\\kms$ and a wind mass loading of 8, results in a good match to the observed galaxy mass function over the range $10^9-10^{11} \\hMsol$. Introducing the AGN scheme developed in Bow08 suppresses the formation of galaxies from cooling in hot gas haloes, resulting in a good match to the high-mass turn-over of the mass function. The combined (AGN plus supernova-driven wind) model provides a good description of galaxy abundance down to $10^9 \\hMsol$ (but rises much more steeply than the observational data at lower masses). Although this is encouraging, we find that the model performs poorly compared to B8W7 in several other aspects. In particular, when we examine the specific star formation rates in this model, we see that the fixed feedback scheme results in a kink in the specific-star formation rates of galaxies at around a stellar mass of $10^{9.5} \\hMsol$, such that low mass galaxies have a factor $\\sim 5$ lower star formation rates than their massive counter-parts.  This is not seen in observational data.  The model also tends to under represent the abundance of massive ($\\Mstar > 10^{11}\\hMsol$) galaxies at $z>1$ compared to the B8W7 model and the observational data.   \\item We also consider a model in which the wind mass-loading scales with the inverse of the circular velocity of the disk.  This model results in a more subtle transition between the regime where material easily escapes the halo and that in which it stalls. With suitable choice of parameters, and the inclusion of the AGN feedback scheme, the model reproduces the observed galaxy mass function above $\\Mstar \\sim 10^9 \\hMsol$. However, although the slope of the SSFR$-\\Mstar$ relation is weaker than that seen in the SW model, the trend is reversed compared to the slightly rising relation seen in observational data. The model tends to under-predict the abundance of massive, $z>1$ galaxies compared to B8W7.  The deficit is, however, relatively small and the MS model could be compatible with the observational data if the random uncertainties in the stellar mass estimates are greater than 0.2~dex.  \\item Feedback from AGN may have two very distinct forms. The ``hot halo'' (or ``radio'') mode feedback suppresses the supply of gas from cooling haloes. This is a key component of the successful models we present. However, it is interesting to investigate if this can be replaced by a ``starburst'' or ``QSO'' mode of feedback. We implement this by enhancing the wind speed and mass loading during starburst events. We find that while this mode may enhance the feedback from quiescent star formation, it does not introduce a characteristic break in the galaxy mass function and cannot be seen as an alternative to the ``hot-halo'' mode feedback. {\\rgb Similar conclusions are reached by \\citet{gabor2011}, when they introduce a comparable feedback scheme in hydrodynamical models.}  \\item We have also investigated the star formation history of the each of the models, and its dependence on the stellar mass. Looking at the integrated star formation rate, the main difference lies in the strength of the decline between $z=2$ and the present day.  This is strongest in the B8W7 model. Nevertheless, given the uncertainties in the observational estimates, both the MS and B8W7 model seem to provide a reasonable description. However, the models show a much greater variety in behaviour when the results are broken down by stellar mass. For this reason, comparison of the total star formation rate may be miss-leading since the observational results assume large extrapolations to account for galaxies that are not directly observed.  It is therefore cleaner to compare the contribution to the total star formation rate density at each epoch as a function of stellar mass. None of the models provide a perfect description of the data. At low redshift, the distribution is broader in the B8W7 model than is the case for the SW or MS models. Observed contribution from low mass galaxies is better described by B8W7, but the model contains too few high star formation rate galaxies at low redshift. At $z\\sim2$, the situation is reversed, with all the models showing a similar contribution from low mass galaxies, and the primary difference being the paucity of high stellar mass galaxies in the MS and SW models.  \\item The differences between the models are emphasised by comparing the mass fractions of baryons in different phases as a function of halo mass.  The conventional model, B8W7, assumes that haloes less massive than $10^{11.5} \\hMsol$ are baryonically closed. Almost all of the baryons in low mass haloes are placed in cool clouds in the halo (because of the short halo cooling time). In contrast the SW and MS, expel much of their baryonic content from the halo. The fraction of baryons retained in the halo has a maximum of 80\\% at $\\Mhalo = 10^{11.7} \\hMsol$ and drops rapidly to a minimum of $\\sim 10\\%$ at $\\Mhalo = 10^{11} \\hMsol$. At the minimum, the majority of baryons are predominantly stars and cold gas. At masses above $10^{12}\\hMsol$, the ``hot-halo'' mode of feedback takes over expelling the halo material, resulting in low baryon fractions as discussed in Bow08, and in good agreement with X-ray observations of groups and clusters. The halo baryon fractions of the MS and SW models offer a good explanation of the low abundance of HI clouds around M31 \\citep{sancisi2008}. Further work is required to compare the model predictions to such data and to the circumgalactic gas haloes inferred from absorption line systems at low redshift \\citep{tumlinson2011} since the ionisation state of the extra-galactic gas must be carefully computed and combined with X-ray limits on the emission from galactic haloes \\citep{crain2010}. In all of these models, the recycling of previously ejected gas plays an important role.   {\\rgb Another potential discriminant between the models is the distribution of metals in the ISM. All the models predict that a major fraction of the metals produced by stars will reside outside galaxies. In the absence of AGN feedback, the B8W7 predicts that the metals will be confined within galaxy haloes, while it will be more wide-spread in the expulsion models (SW and MS). We find, however, that at $z<2$ the major contributor of extra-galactic metals is the AGN powered hot-halo feedback that expels diffuse gas from galaxy groups.  This material is highly enriched compared to the winds from low-mass galaxies: combined with the high fraction of the total baryon budget that has been expelled from groups, this component dominates the diffuse metal content of the low redshift Universe. At $z>3$, AGN feedback makes little contribution to the intergalactic fraction of metals and the observed wide-spread distribution of metals at this epoch favours the expulsion feedback schemes of the MS and SW models. However, the nature of our phenomenological model makes it difficult to predict the dispersal of the expelled material, and it is unclear how far from the parent galaxy groups the metals will be spread \\citep{booth2010}.  Careful consideration of the dynamics of the outflow is needed to make a meaningful comparison with observational data on metal absorption lines.}   \\end{itemize} }  {\\rgb In summary, although we have introduced a feedback scheme that reproduces the results of hydrodynamical simulations well, we find that the original wind scheme of Bow06 produces a better match to observational data on the stellar mass content of haloes, the specific star formation rates of galaxies and the evolution of these quantities. Comparison of the models with the observational data highlights several aspects of the observational data that are not consistent with any of the models considered. Firstly, none of the models reproduce the rapid drop in the normalisation of the observed SMF between $z=0$ and $z=1-2$. Secondly, although the B8W7 model reproduces the observed dependence of the specific star formation rate on stellar mass better than the SW or MS models, none reproduce the observed tendency for lower mass galaxies to have higher SSFR than their massive counter-parts. Both of these trends suggest that the effective wind speed (relative to the halo escape velocity) should increase with redshift. One possibility is that this might arise because of the more clumpy concentration of star forming regions in high redshift galaxies. This is an interesting possibility for further investigation. There is also no intrinsic reason why galaxy winds should be a simple power-law, and the observational suggest we should explore feedback schemes in which the exponent of the mass dependence varies with system. For example, the problem of the SW and MS models over producing galaxies below $10^9\\hMsol$ could be solved by introducing a strong mass dependence to the feedback below this limit. Rather than introducing phenomenological modifications to the feedback scheme, an alternative approach would be to simulate idealised galaxies at high (10 pc), or ultra-high (0.1pc) resolution and to extract suitable parameterisations for the outflow \\citet{hopkins2011,stringer2011b,creasey2011}. }  {\\rgb In future, we will combine the expulsion mode of feedback considered here with developments of the GALFORM code \\citep{lagos2010} to more accurately trace the cold gas content of the universe. We will also make a one-to-one comparison of the formation histories of galaxies from GIMIC, and the forth-coming suite of EAGLE simulations, with the GALFORM code. Our aim is {\\it not only} to develop the GALFORM model as an emulator of numerical simulations (eg., \\citet{bower2010}) but also to use it as a fundamental tool for understanding the key components of successful galaxy formation models. Only by representing the evolution of galaxies as a simple set of coupled differential equations can we claim to have understood the galaxy formation problem, and to have separated the key processes from the details. }"
    },
    "1112/1112.3008_arXiv.txt": {
        "abstract": "{ High contrast imaging is a powerful technique to search for gas \\gp and brown dwarfs orbiting at separation larger than several AU. Around solar-type stars, \\gp are expected to form by core accretion or by gravitational instability, but since core accretion is increasingly difficult as the primary star becomes lighter, gravitational instability would be the a probable formation scenario for yet-to-be-found distant \\gp around a low-mass star. A systematic survey for such planets around M dwarfs would therefore provide a direct test of the efficiency of gravitational instability.} { We search for gas \\gp orbiting around late-type stars and brown dwarfs of the solar neighbourhood.} {We obtained deep high resolution images of 16 targets with the adaptive optic system of VLT-NACO in the $L'$ band, using direct imaging and angular differential imaging.  This is currently the largest and deepest survey for Jupiter-mass planets around M-dwarfs. We developed and used an integrated reduction and analysis pipeline to reduce the images and derive our 2D detection limits for each target. The typical contrast achieved is about 9 magnitudes at 0.5$\\arcsec$ and 11 magnitudes beyond 1$\\arcsec$. For each target we also determine the probability of detecting a planet of a given mass at a given separation in our images.} { We derived accurate detection probabilities for planetary companions, taking into account orbital projection effects, with in average more than 50\\% probability to detect a 3\\mj companion at 10\\,AU and a 1.5\\mj companion at 20\\,AU, bringing strong constraints on the existence of Jupiter-mass planets around this sample of young M-dwarfs. }   \\date{} ",
        "introduction": "The discovery of extrasolar planets around solar analogs, initiated by \\citet{Mayor.1995}, has radically modified our understanding of planetary formation. The new models that strive to describe planetary formation mechanisms need observational constraints on planetary companions on a large range of parameters, such as stellar mass, metallicity, planet mass and orbital parameters. Notably, the most prolific methods for exoplanet detection, radial velocities and transits, only cover separations up to a few AU and barely reaches the large separation range where all the Solar System \\gp reside. This range can be probed by direct high-resolution imaging. Only a few  unambiguously planetary-mass companions to stars \\citep{Chauvin.2004,MaroisHR8799.2008,Kalas.2008,Lagrange.2010} were discovered by direct imaging, but each one of these detections proved extremely valuable to constrain formation models. Until recently, core accretion (CA) was the preferred model to explain both our Solar System planets and the close-in giant planets found by radial velocity (RV) around solar-type  stars. The recently imaged \\gp at large separations from their parent A-type stars \\citep{MaroisHR8799.2008,Kalas.2008} challenge however this simple view as they are very difficult to form {\\it in situ} by CA, though \\citet{Rafikov.2011} shows that under specific condition CA can form Jupiter-mass planets as far as 40-50AU. In the case of HR8799, this could account for the formation {\\it in situ} of the 3 innermost planets, but \\citet{Dodson-Robinson.2009} claim that HR8799bcd formation by CA has to happen at closer separation and be followed by planet scattering at the observed separation, which would not create a stable system. \\citet{Dodson-Robinson.2009} claim that, instead, gravitational instability (GI) in a disc could produce the $bcd$ planets {\\it in situ} while \\citet{Kratter.2010} argue that HR8799bcd can be formed by GI, but only at larger separation and show that under certain conditions they can migrate to the separations where they are currently observed. However the innermost planet of this system, HR8799e \\citep{Marois.2010HR8799e} appears at least as difficult to form by GI as the outermost planet HR8799b would be to form by CA, hinting that planetary formation models still need improvement to be able to account for the first multi-exoplanetary system directly imaged and badly need further observational constraints, especially in different mass range, before drawing any strong conclusion on the respective part they actually play in planetary formation.  Planetary mass objects could also form in the same way as stars, with  \\citet{Bate.2009} showing that a few objects as light as $\\sim$5\\mj can form by core fragmentation, possibly ending in very high mass-ratio binary system, very look-alike to a planetary system. Such planetary-like binary system could also be formed by disc instabilities in a very massive primordial disc \\citep{Stamatellos.2011}, which could produce systems quite similar to 2M1207AB \\citep{Chauvin.2004}, together with numerous brown dwarfs and very low-mass stars.   The dependency of planetary formation modes to the stellar mass is an open field of investigation, with some observational clues that planetary masses, as well as disc masses, scale with stellar masses \\citep[e.g.][]{Forveille.2011,Scholz.2006}. Debris and transitional discs tracing planetary formation have been detected even around M dwarfs and brown dwarfs, with frequencies similar \\citep{Plavchan.2009} or even larger \\citep{Currie.2011} than around solar-type stars.  \\citet[][]{Boley.2009} identifies some indications of bi-modal \\gp formation, but many more detections are needed to confirm this scenario. In that frame, looking for gas giant planets around M-type stars offers a distinct opportunity to test planet formation mechanisms, as models predict that CA is very inefficient \\citep[][KK08]{Kennedy.2008}. KK08 predict a rate of 1$\\%$ for CA formed \\gp around M stars if the rate is of 6$\\%$ for G type stars) or ``all but impossible\" (DR09) in producing \\gp around those stars. The main reason being that there is not enough time to form cores with 5-10M$_{\\oplus}$ before the gas in the protoplanetary disk is dissipated. On the contrary, GI can be quite efficient in forming \\gp around M stars, at separations larger than typically 20\\,AU with \\citet{Boss.2011} claiming that ``even late M dwarf stars might be able to form gas giants on wide orbits\". This hints at the possibility that there are more \\gp at large separations around late-type stars than there are close-in ones.  From an observational point of view, most of the planets found so far around a yet small number of M-type stars  are rather low-mass ($\\leq$ 0.1 \\mj) ones, and in close orbits. The monitoring of M-type stars started recently and these active stars are complicated targets for RV searches; hence, the domain $\\geq$ 3\\,AU is largely unexplored via RV technics. However the trend that massive \\gp are less common around M stars than around more massive stars at short separations seems statistically robust \\citep[e.g.][; Bonfils et al. 2011, submitted]{Forveille.2011}. The picture at the larger separations ($\\geq$ 10\\,AU) is less clear cut. First, RV survey show that planetary frequency around M dwarfs increases as separation increases (Bonfils et al. 2011, submitted). Secondly this trend is supported by evidence from micro-lensing surveys, with \\citet{Gould.2010} showing that the frequency of Jupiter (respectively Saturn) mass planets at $\\sim$10\\,AU is about 10\\% (respectively 20\\%), and correlates well with the extrapolation of RV surveys to large separations.  This separation range started being explored with direct Adaptive Optic (AO) imaging in the recent years, on a small number of late type objects \\citep[17 out of 118 in the recent compilation of][]{Nielsen.2010}, and with limited sensitivities (typically $>$4 \\mj). Others surveys targeted at young late M stars and brown dwarfs were very sensitive to unequal mass binary systems \\citep[$q \\sim 0.2$;][]{Ahmic.2007,Biller.2011} but did not reach mass ratio typical of planetary systems ($q < 0.01$), and therefore did not probe the supposedly more numerous light \\gp ($\\leq$1\\mj) that are supposed to be more frequent at large separation around M dwarfs.  The main limitation of high resolution imaging of M-dwarfs is their intrinsic faintness which makes difficult to achieve good AO correction from the ground. Spatial observations \\citep[e.g.][]{Bouy.2003} and Laser Guide Star observations \\citep[e.g.][]{Biller.2011} do mitigate this issue but make high-resolution deep observations of  M dwarf more challenging than similar observations around solar-type stars. However, since most of  M dwarfs flux is emitted in the infrared, the capabilities of VLT-NAOS adaptive optics system with Infrared Wave-Front Sensor (IRWFS) significantly broaden the sample of M dwarfs which can be used as a good quality natural guide star, thus enabling their observations at high Strehl ratio from the ground. Given the very red Spectral Energy Distribution (SED) of Jupiter mass planets, typically $Ks-L'$=3.3 at 3\\mj and $Ks-L'>$6 at 1\\mj at 30Myr \\citep[according to BT-SETTL 2010 models,][]{Allard.2010}, observing in the $L'$ band offers a greatly enhanced sensitivity to the lowest mass planets. The additional capability of NACO to observe in the $L'$ range therefore makes NACO at VLT a unique instrument to probe for planets around M dwarfs. The high spatial density of M-type stars in the solar neighbourhood makes it easier to build a good sample of young and nearby low-mass stars than it is for solar-type stars. Their youth ensures a more favourable contrast for planetary companions and their intrinsically low flux enable to probe much lower companion masses around M dwarfs than around any other stars for a given contrast. We therefore started a direct imaging survey in $L'$ with NACO to look for \\gp around young nearby M-dwarfs in stellar associations, and probe separations down to a few AU and masses down to $\\sim$0.5 \\mj. We describe our sample and observations in section 2, and detail our data analysis methods in section 3. Our results are presented in section 4. ",
        "conclusions": "We presented the results of the deepest imaging survey of young M dwarfs to date, using $L'$ imaging with NACO at VLT. After developing a dedicated reduction and analysis pipeline, we achieved detections limits in average down to 1.5\\mj beyond 20\\,AU and up to 100-200\\,AU, and 3\\mj at 10\\,AU. On the closest and latest type targets we achieved detection limits well below the mass of Jupiter beyond 10\\,AU, therefore actually starting to probe (on 5 objects) the mass/separation range where planets around M dwarfs are supposed to be be more frequent \\citep[][Bonfils et al, submitted]{Gould.2010}, but found none.  We also probed the high planetary mass range (M$<$13\\mj) at close separations, reaching planetary sensitivity at 5\\,AU or less on 9 out of our 16 targets. In spite of these deep observations we found only one planetary companion, 2M1207B, discovered by \\citet{Chauvin.2004}, in this sample of young M dwarfs. Unfortunately, our sample is currently too small to derive meaningful constraints on the existence of \\gp around late type stars, beyond the simple statement that \\gp more massive than 1\\mj are not common. With the same statistical limitations, our data also support the ``brown dwarfs desert\" hypothesis \\citep{Halbwachs.2000} down to the lowest brown dwarfs masses, since we found no brown dwarfs companions while our survey had in average more than 95\\% de-projected sensitivity to brown dwarfs beyond 15\\,AU and could discover a fraction of them as close as 5\\,AU from the central star. Further deep observations in $L'$ of such targets are necessary to increase our statistical significance and to bring stronger constraints on formation models around low-mass stars."
    },
    "1112/1112.4302_arXiv.txt": {
        "abstract": "We report the AGILE $\\gamma$-ray observations and the results of the multiwavelenght campaigns on seven flaring blazars detected by the mission: During two multiwavelenght campaigns, we observed $\\gamma$-ray activity from two Flat Spectrum Radio Quasars of the Virgo region, e.g. 3C 279 and 3C 273 (the latter being the first extragalactic source simultaneously observed with the $\\gamma$-ray telescope and the hard X ray imager of the mission). Due to the large FOV of the AGILE/GRID instrument, we achieved an almost continuous coverage of the FSRQ 3C 454.3. The source showed flux above 10$^{-6}$ photons/cm$^2$/s (E $>$ 100 MeV) and showed day by day variability during all the AGILE observing periods. In the EGRET era, the source was found in high $\\gamma$-ray activity only once. An other blazar, PKS 1510-089 was frequently found in high $\\gamma$-ray activity. S5 0716+71, an intermediate BL Lac object, exhibited a very high $\\gamma$-ray activity and fast $\\gamma$-ray variability during a period of intense optical activity. We observed high $\\gamma$-ray activity from W Comae, a BL Lac object, and Mrk 421, an high energy peaked BL Lac object. For this source, a multiwavelenght campaign from optical to TeV has been performed. ",
        "introduction": "\\begin{table*} \\caption{List of AGILE flaring Blazars. $^*$This $\\gamma$-ray source is a Blazar candidate, with SDSS J123932.75+044305 as possible optical counterpart.} \\label{tab:sources_list} \\centering \\begin{tabular}{l | c | c | c  } \\hline\\hline Name & Period              &  S/NR  &  ATel  \\\\ & \\emph{start : stop} &        &        \\\\ \\hline S5 0716+714  & 2007-09-04 : 2007-09-23 &  9.6 &  1221      \\\\ & 2007-10-24 : 2007-11-01 &  6.0 &            \\\\ \\hline MRK 0421     & 2008-06-09 : 2008-06-15 &  4.5 & 1574, 1583 \\\\ \\hline W Comae      & 2008-06-09 : 2008-06-15 &  4.0 & 1582       \\\\ \\hline PKS 1510-089 & 2007-08-23 : 2007-09-01 &  5.6 & 1199       \\\\ & 2008-03-01 : 2008-03-20 &  7.0 & 1436       \\\\ \\hline 3C 273       & 2007-12-16 : 2008-01-08 &  4.6 &            \\\\ \\hline 3C 279       & 2007-07-09 : 2007-07-13 & 11.1 &            \\\\ & 2007-12-16 : 2008-01-08 &  4.1 &            \\\\ \\hline 3EG J1236+0457$^*$    & 2007-12-16 : 2008-01-08 &  5.2 &            \\\\ \\hline 3C 454.3     & 2007-07-24 : 2007-07-30 & 13.8 & 1160, 1167 \\\\ & 2007-11-10 : 2007-12-01 & 19.0 & 1278, 1300 \\\\ & 2007-12-01 : 2007-12-16 & 21.3 &            \\\\ & 2008-05-10 : 2008-06-30 & 15.0 & 1545, 1581, 1592 \\\\ & 2008-07-25 : 2008-08-15 & 12.1 & 1634       \\\\ \\hline \\end{tabular} \\end{table*} During the first 1.5 years of observations of AGILE, the main part of the observing plan was reserved to galactic sources, and particularly to the Cygnus field. The satellite was pointed outside the galactic plane only few times, thence the sample is mainly constituted by sources with position not to far from the galactic plane. Nevertheless, due to the AGILE response to ToO requests (from the GASP consortium, or from VERITAS \\cite{veritas}), A few sources have been detected far from the galactic plane. All the optical observatories are located in the northern hemisphere, thence the sources are mainly located in the northern sky, or near the equatorial plane. The list of flaring Blazars detected with AGILE is reported in table \\ref{tab:sources_list}, together with the observing periods. ",
        "conclusions": "AGILE succesfully detects blazars in flare. For some of them, i.e. 3C 454.3, PKS 1510-089, the emission features can be modelled in the framework of leptonic model, including the External Compton component with seeds photons originated by the disk, but reprocessed by the Broad Line Regions. In particular, for 3C 454.3, we observed in December 2007 a so high $\\gamma$-ray flux, that an inverse Compton component with seed photons coming from the hot corona is needed to explain the data. In some other cases (for S5 0716+714, Mrk 421) the correlation of the variability patterns in optical and $\\gamma$-ray, as well as the SED, suggest that the high energy peak of the Spectral Energy Distribution is due to the SSC mechanism."
    },
    "1112/1112.4843_arXiv.txt": {
        "abstract": "We study the collimation of relativistic hydrodynamic jets by the pressure of an ambient medium in the limit where the jet interior has lost causal contact with its surroundings.  For a jet with an ultrarelativistic equation of state and external pressure that decreases as a power of spherical radius, $p \\propto r^{-\\eta}$, the jet interior will lose causal contact when  $\\eta>2$.  However, the outer layers of the jet gradually collimate toward the jet axis as long as  $\\eta < 4$, leading to the formation of a shocked boundary layer. Assuming that pressure-matching across the shock front determines the shape of the shock, we study the resulting structure of the jet in two ways: first by assuming that the pressure remains constant across the shocked boundary layer and looking for solutions to the shock jump equations, and then by constructing self-similar boundary-layer solutions that allow for a pressure gradient across the shocked layer. We demonstrate that the constant-pressure solutions can be characterized by four initial parameters that determine the jet shape and whether the shock closes to the axis. We show that self-similar solutions for the boundary layer can be constructed that exhibit a monotonic decrease in pressure across the boundary layer from the contact discontinuity to the shock front, and that the addition of this pressure gradient in our initial model generally causes the shock front to move outwards, creating a thinner boundary layer and decreasing the tendency of the shock to close. We discuss trends based on the value of the pressure power-law index $\\eta$. ",
        "introduction": "The idea that AGN outflows are highly collimated is supported by observations (e.g. \\citealp{MB84}, \\citealp{Jorstad05}), implying that confinement must occur. The cause of this confinement, which may occur at distances of just a few tens of Schwartzschild radii from the central black hole (\\citealp{JunorBiretta99}), is however not yet well-understood. It is generally accepted that some collimating agent is necessary, but this still allows for a variety of possibilities, both external (e.g. pressure confinement via an ambient static medium \\citealp{Eichler82}, \\citealp{Komissarov97}, inertial confinement via a slow outflow or wind surrounding the jet \\citealp{KomissarovRamPressure94,BL07}) and internal (magnetic confinement via hoop stress due to the toroidal magnetic field component \\citealp{Benford78}, \\citealp{MB84}).  Pressure confinement is of particular interest because accretion disk winds surrounding an AGN provide an ideal external medium for interaction with the jet. Moreover, another collimating agent such as magnetic hoop stress cannot function alone; without an ambient medium to confine the globally-expansive magnetic field, collimation will not occur (\\citealp{MB95,KomBarkVla07,KomVlaKon09,Komissarov11}). Thus pressure confinement may be relevant both on its own and in conjunction with other~processes.  In this context, we note that numerical simulations that study confinement and acceleration of the flow by MHD processes treat the action of the external medium as an applied boundary condition (e.g. \\citealp{KomBarkVla07,KomVlaKon09,Komissarov11}). In this work, we study the detailed physics of jet collimation by the external medium as a first step towards a complete treatment of both the action of the external medium and magnetic effects in collimating and accelerating relativistic jets.  The location and mechanism of jet collimation is interesting because it provides insight into energy dissipation within the jet. In a steady jet, there are two main sites where energy dissipation is likely to occur: at the jet spine, as a result of kink instabilities driven by the toroidal magnetic field (\\citealp{MB98}, \\citealp{Eichler93}), and at the interface between the jet and its environment, as a result of shear instabilities (\\citealp{Micono00,Perucho10,Barkov11}) or collimation shocks, as we discuss here.  In the case of a jet with an opening half-angle of $\\theta$ and a bulk Lorentz factor of $\\Gamma$, the outer edge of the jet remains in causal contact with the jet center only if the opening angle obeys $\\Gamma^{-1} > \\sin \\theta \\approx \\theta$. In this work, we assume a relativistic equation of state and examine the case where causal contact between the jet's spine and edge has been lost. Thus we model the jet as having a transverse structure consisting of two components: an inner region in which the flow is undergoing free expansion as it accelerates, and an outer shocked boundary layer region that results from the loss of causal contact within the jet. The geometrical shape that the jet assumes as it propagates is a direct result of its response to the collimating forces exerted by the ambient medium, and calculating that shape for the case of pressure confinement will be a major focus of this paper.  Our goal is to determine the basic jet structure and geometry under a simple set of collimation and acceleration assumptions, providing a model for the jet's steady-state configuration. This will allow more realistic initial conditions for simulations of instabilities and turbulence, which will in turn provide a model for AGN jets to which we can compare observational signatures.  While we specifically reference AGN jets in this work, our results can easily be extended to other relativistic outflows, such as a gamma-ray burst in the collapsar model, collimated by the stellar envelope it breaks through, as discussed in \\citealp{BL07} (hereafter BL07).  In this paper we use the shock conditions for a hydrodynamic, relativistic jet to derive the basic structure of the jet. This problem has been previously studied in the case of a \"cold\" (inertia-dominated) jet (\\citealp{Komissarov97}, \\citealp{KN09}), but we now focus on the collimation behavior of a \"hot\" (pressure-dominated) jet.  We initially follow a similar approach to that of BL07, but deviate from its methods in our treatment of entropy distribution within the jet. \\citealp{BL09} (BL09) performs a similar analytical inspection in the limit of small angles; we generalize this to all angles.  In \\S 2 we find solutions for the jet shape using the Kompaneets approximation. In  \\S 3 we examine self-similar solutions for the boundary layer when a pressure gradient is allowed to form across the layer, and then revise our solutions from \\S 2 to include this pressure gradient. In \\S 4 we conclude, summarizing the results and discussing future work. ",
        "conclusions": "\\indent We have evaluated the shape of the shocked boundary layer of a hot jet with an ultrarelativistic equation of state, in the case that it is injected into an ambient medium that has a power-law pressure profile of $p \\propto r^{- \\eta}$ with $2 < \\eta < 4$. Using the shock jump conditions and momentum conservation and assuming that the pressure remains constant across the boundary layer, we found that whether the shocked layer closes to the axis --- and where it closes, if it does --- is dependent upon the values of the initial pressure ratio, the power-law index, the initial bulk Lorentz factor, and the product of the initial opening angle and initial Lorentz factor.  We also found that, in the Kompaneets approximation, the contact discontinuity asymptotes to the shape $R_c \\propto z^{\\eta/4}$ in the case where the shock has closed to the axis, but to a conical shape, $R_c \\propto z$, in the case of a non-closing shock or in the region of the jet before the shock has closed.  Due to the expectation of a pressure gradient arising from the centripetal force created by the slight collimation of the jet, we then created a self-similar model of the boundary layer that allows for a pressure gradient across the layer. From this we found solutions for which pressure decreases monotonically across the boundary layer for all values of the power-law index $\\eta$.  An $\\eta$ of 8/3 constitutes a special solution where the pressure goes to zero at a finite distance from the outer wall. For values of $\\eta$ greater or less than 8/3, the pressure approaches zero asymptotically. The curve for pressure steepens  for values of $\\eta$ near 8/3 and approaches a constant value for $\\eta \\rightarrow 2$ or $\\eta \\rightarrow 4$.  The drastic decrease of the pressure inwards from the contact discontinuity for values of $\\eta$ near 8/3 suggests that most of the material is pushed up against the outer wall, creating a hollow-cone structure for the jet that may be observable as edge-brightening. It also indicates that the Kompaneets approximation is not valid for these values of $\\eta$. For values of $\\eta$ near 2 or 4, however, the Kompaneets approximation may be reasonable, since the pressure decrease is very gradual.  To better understand these results, we then revised the Kompaneets solutions to include the pressure gradient of the self-similar boundary-layer model. We held the outer wall fixed and solved for the new position of the shock front by pressure-matching across the shock jump. We found that the addition of the pressure gradient caused the shock front to move outward in order to match the lower pressure, resulting in a thinner boundary layer and preventing the shock from closing to the axis in places where it originally had done so. We also confirmed that the total energy contained within the boundary layer was an increasing function of height $z$, as is required.  The inherent difficulty of examining this problem analytically required us to make several simplifying assumptions over the course of this work. One such assumption, made in \\S 3 to obtain the boundary-layer solutions, is that all the material within the boundary layer is on the same adiabat. This would be true if the majority of the material entered the boundary layer at roughly the same location, however this is not the case in general. But for the boundary-layer solutions where the pressure drops rapidly ($\\eta$ near $8/3$), most of the material must be pushed up against the contact discontinuity early on. This suggests that it must all have entered the boundary layer near the jet base, which demonstrates self-consistency with the same-adiabat assumption for this regime.~Nonetheless, in future work we intend to explore the possibility of treating the jet material on different adiabats.  It is widely believed that magnetic effects contribute to the collimation of relativistic jets (see e.g. \\citealp{KomissarovNumSim99}). While the models discussed here are purely hydrodynamic, they can have important applications both in interpreting the data obtained in three-dimensional general relativistic magnetohydrodynamics simulations that provide a self-consistent description of the jet launching mechanism (e.g. \\citealp{Beckwith08a,Beckwith09,McKinney09}) and in improving the physical content of boundary conditions imposed in simulations that study the role of magnetic fields in large-scale jet collimation (see e.g. \\citealp{KomBarkVla07,KomVlaKon09,Komissarov11}).  In the former case, the Poynting-flux dominated jet is sheathed by an unbound outflow (\\citealp{Hawley06}), which may have important implications for the operation of current-driven instabilities (\\citealp{McKinney09}). However, the numerical resolution used in these simulations is generally insufficient to adequately resolve the sharp density and pressure gradients present in these regions. Combining the models presented here with the conditions present in the ambient medium at the base of the jet (e.g. the disk corona) in these simulations will allow us to assess the extent to which numerical resolution affects simulation outcomes.  In the case of simulations that study the role of magnetic fields in large-scale jet collimation, the work presented here will allow a more complete physical treatment of the effects of the ambient medium. In these simulations (\\citealp{KomBarkVla07,KomVlaKon09,Komissarov11}), this boundary is generally treated as a rigid wall, where fluid quantities (such as gas density and pressure) are simply copied across the boundary, while the normal component of vector quantities (e.g. velocity and magnetic field) is reflected. The models presented here will allow improvement of this treatment by specifying jump conditions for hydrodynamic quantities at this boundary, consistent with the shape of the wall --- thereby including the collimating effect of the ambient medium.  These calculations mark the first step towards a more complete treatment of both the action of the external medium and magnetic effects in collimating and accelerating relativistic jets. The next step in this work is to include the effect of a toroidal magnetic field, allowing for magnetic confinement. Together, this work and its magnetized extension will provide us with an equilibrium model from which we can explore instabilities and radiative mechanisms of the jet near its base, leading to improved constraints on jet dissipation and associated radiative signatures. Through this work we hope to further our understanding of jet collimation and, in particular, boundary-layer behavior."
    },
    "1112/1112.1461_arXiv.txt": {
        "abstract": "We investigate the connection between starburst and active galactic nucleus (AGN) activity by comparing the 3.3 \\micron~polycyclic aromatic hydrocarbon (PAH) emission with AGN properties. Utilizing the slit-less spectroscopic capability of the AKARI space telescope, we observe moderate-luminosity Type I AGN at z$\\sim$0.4 to measure global starburst activity. The 3.3 \\micron~PAH emissions are detected for 7 out of 26 target galaxies. We find no strong correlation between the 3.3 \\micron~PAH emission and AGN luminosity in the limited range of the observed AGN luminosity, suggesting that global star formation may not be closely related to AGN activity. Combining our measurements with the previous 3.3 \\micron~measurements of low redshift Type I AGN in the literature, we investigate the connection between nuclear starburst and AGN activity. In contrast to global star formation, the 3.3 \\micron~PAH luminosity measured from the central part of galaxies correlates with AGN luminosity, implying that starburst and AGN activity are directly connected in the nuclear region. ",
        "introduction": "\\label{intro}  The correlations of black hole mass with galaxy properties imply a close link between black hole growth and galaxy evolution \\citep{FM00, Gt00, gultekin+09, woo10}. Numerous theoretical and observational studies have been dedicated to revealing the nature of the connection (e.g., Kauffmann \\& Haehnelt 2000; Springel et al. 2005; Croton et al. 2006; Woo et al. 2006, 2008; Hopkins et al. 2007; Merloni et al. 2010; Bennert et al. 2011).  Probing the direct connection between starburst and active galactic nuclei (AGN) can shed light on how black hole growth and bulge growth are inter-connected. However, it is challenging to observationally measure both the starburst rate and AGN activity for given galaxies. Optically-identified broad-line AGNs are the best objects for measuring black hole mass and AGN accretion rate. On the other hand the star formation rate of host galaxies is difficult to measure since AGN flux dominates in the UV continuum and hydrogen recombination lines, i.e., \\ha, which are generally used as a star formation indicator in non-AGN galaxies. In the case of Type II AGN, it is relatively easy to measure star formation rate while black hole mass and accretion rate are difficult to determine owing to the dust obscuration of accretion disk and broad-line region.  The polycyclic aromatic hydrocarbon (PAH) emission features emitted by small carbon grains are one of the star formation indicators. Thus, PAH emissions have been used to constrain star formation activity in the host galaxies of Type 1 AGN (e.g., Ogle et al. 2006; Schweitzer et al. 2006; Shi et al. 2007). Based on the ground-based and space observations, various PAH emission features, i.e., 3.3, 6.2, 7.7, 8.6, 11.3, and 12.7 \\micron~lines, have been detected in many local Seyfert 1 galaxies and distant QSOs although not all AGN show PAH emission features \\citep[e.g.,][] {clavel+00, Rod03, imanishi04, ntz07,lz07,lutz+08,shi+09,imanishi10,Oi10,imanishi11}.  Among PAH-detected AGN, there seems to be a correlation between nuclear starburst and AGN activity. Using a sample of low-redshift Palomar-Green (PG) QSOs observed with the Spitzer Infrared Spectrograph (IRS), \\citet{ntz07} showed that 7.7 \\micron~PAH luminosity increases with AGN continuum luminosity at 5100\\AA~ (L$_{\\rm 5100\\AA}$), indicating the starburst-AGN connection. High-redshift QSOs seem to follow the same trend, extending the correlation to much higher luminosity range \\citep{lutz+08}. The starburst-AGN connection is supported by a number of theoretical studies, demonstrating that both nuclear starburst events and AGN activity can be triggered by gas inflow through galaxy interactions (e.g., Kauffmann \\& Haehnelt 2000; Robertson et al. 2006; Ciotti \\& Ostriker 2007; Hopkins et al. 2009).  For lower luminosity AGNs it is less clear how AGN activity is connected to starburst since AGN power is more difficult to measure owing to increasing host galaxy contribution in lower luminosity AGN. Based on the ground-based L-band slit spectra, \\citet{imanishi04} measured the 3.3 \\micron~PAH luminosity of local Seyfert 1 and 2 galaxies, showing a positive correlation between starburst and AGN activity although the correlation is not tight with a large scatter. In their study, radio luminosity and N-band luminosity within 1.5'' of the host galaxies are used as an indicator of AGN power. Using 2-10 kev X-ray luminosity, \\citet{Wt08} showed a similar trend between AGN power and the 3.3 \\micron~PAH luminosity for a larger sample of local Seyfert 1 and 2 galaxies.  Most previous observational studies on PAH features with ground based long-slit spectroscopy or Spitzer IRS observations have probed the inner part of galaxies, indicating a connection between {\\it nuclear} starburst within a few kpc and AGN activity \\citep[e.g.,][]{Wt08}. In contrast, whether global star formation over entire galaxies is related with AGN activity it is not well studied since most previous studies, particularly for local and low-redshift galaxies, have used a relatively small aperture to extract PAH emission features, measuring nuclear PAH emissions. Based on the spatially resolved mid-IR spectroscopy of nearby active galaxies, several studies have reported that PAH emissions are mainly from the disk \\citep[e.g.,][]{laurent+00, lefloch+01}, suggesting that the global PAH luminosity measured over entire galaxies may be much larger than nuclear PAH luminosity.  In this paper, we investigate global and nuclear star formation in the host galaxies of Type 1 AGN, using the 3.3 \\micron~PAH emission measurements based on our new observations with the AKARI space telescope and collected data from the literature. The AKARI telescope provides a unique capability of slit-less spectroscopy to measure global star formation over entire galaxies (e.g., Imanishi et al. 2010). This paper is organized as follows. In \\S \\ref{data}, we describe sample selection, observations and data reduction of the AKARI observations. We present 3.3 \\micron~PAH emission measurements in \\S \\ref{analysis}. Main results on the starburst-AGN connection are presented in \\S \\ref{results}. Discussion and conclusions are followed in \\S~5. Throughout the paper, we assume a Hubble constant of H$_{o}$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda} =0.7$, and $\\Omega_{M} = 0.3$. ",
        "conclusions": "To investigate the starburst-AGN connection, we compared the global and nuclear 3.3 \\micron~PAH emission with AGN luminosity using a sample of Type I AGN. The global 3.3 \\micron~PAH emission luminosities are measured for moderate-luminosity AGN at z$\\sim$0.4, based on our new AKARI slit-less spectroscopy, which covers the entire host galaxies. We find no strong correlation between 3.3 \\micron~PAH and AGN optical luminosities within the limited AGN luminosity range 10$^{44}$ - 10$^{45}$ \\ergs. This result may imply that the global star formation is not strongly related with AGN activity although we cannot rule out a weak correlation with a considerable scatter. We note that since the AGN luminosity range is very limited, whether the global star formation rate correlates with AGN activity is inconclusive.  In contrast, when we compare AGN luminosity of the local Seyfert 1 galaxies and PG QSOs with their nuclear 3.3 \\micron~PAH emission luminosity, we detect a strong correlation, suggesting that that AGN activity is related to the nuclear starburst. These results are consistent with the findings of previous studies on the starburst-AGN connection \\citep{Im03, imanishi04, ntz07, lutz+08, Oi10}.  The slope of the \\LPAH~correlation with AGN luminosity in Figure 5 shows that the \\LPAH~to AGN luminosity ratio decreases toward higher AGN luminosity. At face value, it implies that the nuclear starburst activity is slightly suppressed in the host galaxies of high luminosity AGN (for different predictions on the starburst-to-accretion ratios, see e.g., Kawakatu \\& Wada 2008; Ballantyne 2008). A similar trend was noted by \\cite{lutz+08} in comparing 7.7 \\micron~PAH emission and AGN luminosities. In their study, high luminosity AGN with L$_{\\rm 5100\\AA}$ $>$ 10$^{46}$ \\ergs~ show a much lower starburst-to-AGN ratio \\citep[see Figure 6 in][]{lutz+08}, implying that the black hole growth rate normalized by the starburst rate is a factor of 10 higher in their high redshift QSOs than in the local AGN. In our local sample the decrease of the starburst-to-AGN ratio with increasing AGN luminosity is not as dramatic as in high redshift QSOs, but if confirmed, it implies that for low-luminosity AGN the growth of galaxy centers is faster than black hole growth in the present-day universe. However, it is not clear whether black hole growth is slower or faster than the total host galaxy growth since only nuclear starbursts are compared in our analysis.  We note that there are many AGN with no PAH emission detection. Including the upper limits of these objects significantly weakens the correlation between \\LPAH~and AGN luminosity. Although the true PAH luminosities of these objects may be close to the upper limits within a factor of few, a detailed investigation using a large sample with lower flux limits is necessary to probe the slope and scatter of the correlation."
    },
    "1112/1112.3652_arXiv.txt": {
        "abstract": "\\noindent Current analysis of astronomical data are confronted with the daunting task of modeling the awkward features of astronomical data, among which heteroscedastic (point-dependent) errors, intrinsic scatter, non-ignorable data collection (selection effects), data structure, non-uniform populations (often called Malmquist bias), non-Gaussian data, and upper/lower limits. This chapter shows, by examples, how modeling all these features using Bayesian methods. In short, one just need to formalize, using maths, the logical link between the involved quantities, how the data arize and what we already known on the quantities we want to study. The posterior probability distribution summarizes what we known on the studied quantities after the data, and we should not be afraid about their actual numerical computation, because it is left to (special) Monte Carlo programs such as JAGS. As examples, we show how to predict the mass of a new object disposing of a calibrating sample, how to constraint cosmological parameters from supernovae data and how to check if the fitted data are in tension with the adopted fitting model. Examples are given with their coding. These examples can be easily used as template for completely different analysis, on totally unrelated astronomical objects, requiring to model the same awkward data features. \\newline {\\bf Keywords}: Astrophysics; Cosmology; Bayesian statistics; Regression; Scaling relations; Prediction; Model testing ",
        "introduction": "Astronomical data present a number of quite common awkward features (see \\cite{AandH11review} for a review):  \\begin{itemize}  \\item {\\bf heteroscedastic errors}: error sizes vary from point to point.  \\item {\\bf non-Gaussian data}: the likelihood is asymmetric and thus errors are, e.g $3.4^{+2.5}_{-1.2}$. Upper/lower limits, as 2.3 at 90 \\% probability, are (perhaps extreme) examples of asymmetric likelihood.  \\item {\\bf non uniform populations or data structure}: the number of objects per unit parameter(s) is non-uniform. This is the source of the Malmquist- or Eddington- like bias that affect most astronomical quantities, as parallaxes, star and galaxy counts, mass and luminosity, galaxy cluster velocity dispersions, supernovae luminosity corrections, etc.  \\item {\\bf intrinsic scatter}: data often scatter more than allowed by the errors. The extra-scatter can be due to unidentified sources of errors, often called systematic errors, or indicates an intrinsic spread of the population under study, i.e. the fact that astronomical objects are not identically equal.  \\item {\\bf noisy estimates of the errors}: as every measurement, errors are known with a finite degree of precision. This is even more true when one is measuring complex, and somewhat model dependent, quantities like mass.  \\item {\\bf non-random sampling}: in simple terms, the objects in the sample are not a random sampling of those in the Universe. In some rare occasions in astronomy, sampling is planned to be non-random on purpose, but most of the times non-random sampling is due to selection effects: harder-to-observe objects are very often missed in samples.  \\item {\\bf mixtures}: very often, large samples include the population under interest, but also contaminating objects. Mixtures also arise when one measure the flux of a source in presence of a background (i.e. always).  \\item {\\bf prior}: we often known from past data or from theory that some values of the parameters are more likely than other. In order terms, we have prior knowledge about the parameter being investigated. If we known anything, not even the order of magnitude, about a parameter, it is  difficult even to choose which instrument, or sample, should be used to measure the parameter.  \\item {\\bf non-linear}: laws of Nature can be more complicated than $y=ax+b$.  \\end{itemize}  Bayesian methods allow to deal with these features (and also other ones), even all at the same time, as we illustrate in Sec 3 and 4 with two research examples, it is just matter of stating in mathematical terms our wordy statements about the nature of the measurement and of the objects being measured. The probabilistic (Bayesian) approach returns the whole (posterior) probability distribution of the parameters, very often in form of a Monte Carlo sampling of it.  In this paper we make an attempt to be clear at the cost of being non-rigorous. We defer the reader looking for rigour to general textboox, as \\cite{gelman2004bayesian}, and, to \\cite{AandH10} for our first research example. ",
        "conclusions": "\\label{sec:finalsec}  These two analysis offer a template for modeling the common awkward features of astronomical data, namely heteroscedastic errors, non-Gaussian likelihood (inclusive of upper/lower limits), non-uniform populations or data structure, intrinsic scatter (either due to unidentified source of errors, or due to population spreads), noisy estimates of the errors, mixtures and prior knowledge.  In a Bayesian framework, learning from the data  and the prior it is just matter of formalizing in mathematical terms our wordy statements about the quantities under investigation and how the data arize. The actual numerical computation of the posterior probability distribution of the parameters is left to (special) Monte Carlo programs, of which we don't need to be afraid more than numerical methods (Monte Carlo) used to compute the integral of a function.  The great advantage of the Bayesian modeling is its high flexibility: if the data (or theory) call for a more complex modeling, or call for using distributions different from those initially taken, it is just matter of replacing them in the model, because there are no simplifications forced by the need of reaching the finishing line, as instead is the case in other modeling approaches. Furthermore, if  one is interested in constraining another cosmological model, for example one with a redshift-dependent dark energy equation of state $w = w_0+w_1(1+z)/(1+z_p)$, one should just replace $w$ in the JAGS code with the above equation. This flexible, hierarchical, modeling is native with Bayesian methods.  In both our two examples, the Bayesian approach performs, unsurprisingly, better than than non-Bayesian methods obliged to discard part of the available information in order to reach the finishing line: Bayesian methods just use all the provided information.  As a final note, we remember that the careful researcher, whether using a Bayesian modeling or not, before publishing his own result should check that the numerical computation is adequate for his purpose and that the model is appropriate for the data. Therefore, if you, gentle reader, are using our examples as template, remember to include in your work a sensitivity analysis, by checking that your assumptions (both likelihood and prior) are reasonable. Some prior sensitivity analysis has been performed in both examples to emphasize the importance of showing how much conclusions (the posterior) rely on assumptions (prior). For what concerns the likelihood, i.e. misfitting, the key point consists in updating the model fomulation until it produces data similar to those in hand, as we have shown in great detail for the second example.   More in general, we hope that the template modeling shown in these two examples may be useful for any analysis confronted with modeling the awkward features of astronomical data, among which heteroscedastic (point-dependent) errors, intrinsic scatter, data structure, non-uniform population (often called Malmquist bias) and non-Gaussian data, inclusive of upper/lower limits.   \\medskip {\\it Acknowledgements:} The first part of this chapter is largely based on papers written in collaboration with Merrilee Hurn. It is a pleasure to thank Merrilee for the fruitful collaboration and her wise suggestions along the years, and for comments on an early draft of this chapter. Errors and inconsistencies remain my own."
    },
    "1112/1112.6241_arXiv.txt": {
        "abstract": "We present 1.1 mm observations of the dust continuum emission from the MBM12 high-latitude molecular cloud observed with the Astronomical Thermal Emission Camera (AzTEC) mounted on the James Clerk Maxwell Telescope on Mauna Kea, Hawaii. We surveyed a 6.34 deg$^2$ centered on MBM12, making this the largest area that has ever been surveyed in this region with submillimeter and millimeter telescopes. Eight secure individual sources were detected with a signal-to-noise ratio of over 4.4. These eight AzTEC sources can be considered to be real astronomical objects compared to the other candidates based on calculations of the false detection rate. The distribution of the detected 1.1 mm sources or compact 1.1 mm peaks is spatially anti-correlated with that of the 100 $\\mu$m emission and the $^{12}$CO emission. We detected the 1.1 mm dust continuum emitting sources associated with two classical T Tauri stars, LkH$\\alpha$262 and LkH$\\alpha$264. Observations of spectral energy distributions (SEDs) indicate that LkH$\\alpha$262 is likely to be Class II (pre-main-sequence star), but there are also indications that it could be a late Class I (protostar). A flared disk and a bipolar cavity in the models of Class I sources lead to more complicated SEDs. From the present AzTEC observations of the MBM12 region, it appears that other sources detected with AzTEC are likely to be extragalactic and located behind MBM12. Some of these have radio counterparts and their star formation rates are derived from a fit of the SEDs to the photometric evolution of galaxies in which the effects of a dusty interstellar medium have been included. ",
        "introduction": "\\label{s:intro}  The MBM12 molecular cloud was cataloged as L1453$-$L1454, L1457, and L1458 in Lynds'Dark Cloud Catalog \\citep{lyn62}and is one of the nearby high-latitude molecular clouds \\citep{mbm85}. Until recently, MBM12 was considered to be the nearest star-forming region; however, its distance has been revised to 275 pc by \\citet{luh01} and at 360$\\pm$30 pc from the photometric and spectroscopic results of the Sloan Digital Sky Survey (SDSS) and the Two Micron All Sky Survey (2MASS; \\citet{ander02}). Thus, although it is still nearby, it is farther out than the Taurus dark clouds and Ophiuchus star-forming regions.  A few young stars were reported in MBM12 (\\citet{hearty00}; \\citet{hog03}) and their presence demonstrated that star formation had occurred recently in at least this particular high-latitude molecular cloud \\citep{broeg06}. The high-latitude molecular clouds are known to be translucent, and the low-density environments found in high-latitude clouds create a challenge for star formation by gravitational collapse. The ability of at least a few high-latitude clouds to form cold molecular cores and young stars may be due to a combination of conditions, including variations in the interstellar radiation field, changes in dust-grain size and chemistry, and turbulence in the interstellar medium (ISM).  Regardless of the formation mechanisms, the study of high-latitude star formation environment continues to intrigue researchers, because it is difficult to see how star formation in the low-density environments found in typical high-latitude molecular clouds can proceed via direct gravitational collapse. High-latitude pre-main-sequence (PMS) stars found in isolation could be formed by ejection from their parent cloud or dissipation of the natal cloud. Dissipation timescales for these unbound clouds appear to be a few Myr. It is possible that outflows from low-mass stars that form in low-density environments can drive the dissipation of the parent cloud (\\citet{sd95}; \\citet{fei96}).  In the present paper, we describe the results of our observations of the dust continuum emission from MBM12 at a wavelength of 1.1 mm observed with the AzTEC \\citep{wilson08} on the James Clerk Maxwell Telescope (JCMT). Dust continuum emission at (sub)millimeter wavelength can identify the star forming sites in the molecular clouds and determine the evolutionary stages of young stellar objects. In Section 2, we describe the observations and data reduction procedure. In Section 3, we present information on the sources detected at 1.1 mm and investigate their spectral energy distributions (SEDs). We also describe the assembly of the SEDs and the procedure used to fit the models were assembled and how models were fitted to the data to determine the physical parameters of the detected sources from present observations. In Section 4, we summarize the results of our observations and analysis. ",
        "conclusions": "We presented observations of the dust continuum emission at 1.1 mm from MBM12 performed with AzTEC at the JCMT. We surveyed about 6.34 deg$^2$ field of sky. This is the largest area that has ever been surveyed for the MBM12 region with a (sub)-millimeter telescope. The distribution of 1.1 mm AzTEC sources is spatially anti-correlated with that of the 100 $\\mu$m emission and that of the $^{12}$CO emission. We detected 312 sources with an S/N of over 3 $\\sigma$. The FDR was 319$\\pm$66. Eight sources were detected and their FDR value was smaller than 1. These eight AzTEC sources can be considered to be real astronomical objects as opposed to spurious detections. We detected 1.1 mm sources associated with two classical T Tauri stars, LkH$\\alpha$262 and LkH$\\alpha$264. Observations of their SEDs indicate that LkH$\\alpha$262 is likely to be Class II but may be a late Class I. A flared disk and a bipolar cavity in models of Class I sources lead to more complicated SEDs. Our detections suggest that at least the classical T Tauri stars in MBM12 have disks in a similar mass range with the disks of stars in the Taurus or $\\rho$ Ophiuchus regions. From the present AzTEC observations of the MBM12 region, we find that some of the other detected sources are likely to be located behind MBM12 molecular cloud. These have radio counterparts, and their SFRs were derived from fits of their SEDs to the photometric evolution of galaxies when the effects of a dusty interstellar medium were included. These SFRs ranged from a few hundred to a few thousand $M_\\odot$ yr$^{-1}$. Direct measurements of their red-shifts should be performed with the Large Millimeter Telescope in the near future.  \\vspace*{1cm}  We thank K. Scott, J. Lowenthal, and other AzTEC team members for their contributions to the project. We also thank S. Johnson for helpful suggestions on the manuscript. We thank the anonymous referee for careful reading and helpful comments on the manuscript. This work 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 also has made use of data products from the Two Micron All Sky Survey which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology and the NASA/IPAC Infrared Science Archives which is operated by the Jet Propulsion Laboratory and California Institute of Technology. The James Clerk Maxwell Telescope is operated by the Joint Astronomy Center on behalf of SRC of the United Kingdom, the Netherlands OSR, and the Canada NRC. MK was supported in part by Basic Science Research Program through the National Research Foundation of Korea (NRF) funded by the Ministry of Education, Science and Technology 2009-0066892. This research was supported in part by Mid-career Researcher Program through the National Research Foundation of Korea (NRF) funded by the Ministry of Education, Science and Technology 2011-0028001.  \\newpage"
    },
    "1112/1112.0247_arXiv.txt": {
        "abstract": "{}{}{}{}{}  \\abstract {The nearby Type Ia SN~2011fe has provided an unprecedented  opportunity to derive some of the properties  of the progenitor system, which is one of the key open problems in the supernova (SN) field.} {This study attempts to establish a link between the reasonably well known nature of the progenitor and its surrounding environment. This is done with the aim of enabling the identification of similar systems in the vast majority of the cases, when distance and epoch of discovery do not allow a direct approach. } {To study the circumstellar environment of SN~2011fe we have obtained high-resolution spectroscopy of SN~2011fe on 12 epochs, from 8 to 86 days after the estimated date of explosion, targeting in particular at the time evolution of \\ion{Ca}{ii} and \\ion{Na}{i}.} {Three main absorption systems are identified from \\ion{Ca}{ii} and \\ion{Na}{i}, one associated to the Milky Way, one probably arising within a high-velocity cloud, and one most likely associated to the halo of M101. The Galactic and host galaxy reddening, deduced from the integrated equivalent widths (EW) of the \\ion{Na}{i} lines are $E_{B-V}$=0.011$\\pm$0.002 and $E_{B-V}$=0.014$\\pm$0.002~mag, respectively. The host galaxy absorption is dominated by a component detected at the same velocity measured from the 21-cm \\ion{H}{i} line at the projected SN position ($\\sim$180 \\kms). During the $\\sim$3 months covered by our observations, its EW changed by 15.6$\\pm$6.5 m\\AA. This small variation is shown to be compatible with the geometric effects produced by the rapid SN photosphere expansion coupled to the patchy fractal structure of the ISM. The observed behavior is fully consistent with ISM properties similar to those derived for our own Galaxy, with evidences for structures on scales $\\lesssim$100 AU.} {SN~2011fe appears to be surrounded by a \"clean\" environment. The lack of blue-shifted, time-variant absorption features is fully consistent with the progenitor being a binary system with a main-sequence, or even another degenerate star.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.5238_arXiv.txt": {
        "abstract": "We analyze time series data of the fluctuations of slow solar wind velocity using rank order statistics. We selected a total of $18$ datasets measured by the Helios spacecraft at a distance of $0.32$ AU from the sun in the inner heliosphere. The datasets correspond to the years $1975-1982$ and cover the end of the solar activity cycle $20$ to the middle of the activity cycle $21$. We first apply rank order statistics to time series from known nonlinear systems and then extend the analysis to the solar wind data. We find that the underlying dynamics governing the solar wind velocity remains almost unchanged during an activity cycle. However, during a solar activity cycle the fluctuations in the slow solar wind time series increase just before the maximum of the activity cycle. ",
        "introduction": "\\label{intro} Solar wind is a stream of charged particles, mostly protons and electrons, that are ejected from the upper atmosphere of the sun. \\cite{park58}, who first coined the term `solar wind', showed that the solar wind is in fact the outer coronal atmosphere escaping supersonically into the interstellar space. The solar wind is primarily divided into two components, a fast wind with velocities $\\sim$900 km/s and a slow wind with velocities $\\sim$300 km/s. These two components differ in their origin and characteristics which include flow speed, proton density, temperature, and helium content \\citep{feld05,schw06}. The fast solar wind originates from coronal holes, where the magnetic field lines are open, while slow wind is associated with open magnetic field and is highly structured and highly variable \\citep{krie73,woo97}.  Observations of solar wind velocity distributions made by different spacecrafts have been analyzed in considerable detail. Helios spacecraft probed the heliosphere as close as $0.3$ AU from the sun and these observations have been analyzed using linear and nonlinear time series analysis tools \\citep {mace97,mace98a,mace98b,mace00,reda01,gupt08}. \\cite{mace97} argued that the inner heliosphere is a low dimensional chaotic attractor. \\cite{mace98a} studied a single time series using moving average, singular-value decomposition and surrogate data analysis to characterize the nonlinear behavior of the data. \\cite{mace00} have shown that the Kolmogorov entropy is positive and finite, as it holds for a chaotic system. \\cite{reda01} extended the study by using a nonlinear filter to obtain the noise-reduced data and more reliable estimate of largest Lyapunov exponent and Kolmogorov entropy. \\cite{gupt08} analyzed several Helios time series to understand the local dynamics of fluctuations in the slow solar wind velocity. They suggest the inherent changes in the dynamics over the solar cycle using the variance in largest Lyapunov exponent for each time series.\\\\ \\begin{table}[t] \\caption{The solar wind velocity data measured by the Helios spacecraft in the years 1975 to 1981 with Carrington rotation(CR); Initial time ($T_{i}$); Final time($T_{f}$); Data points($N$). Sunspot numbers are from http://sidc.oma.be/sunspot-data/.} \\vspace{.5cm} \\resizebox{8cm}{!} { \\begin{tabular}{|ccccccc|} \\hline S.No.&CR&Year&$T_{i}$&$T_{f}$&$N$&Sunspot\\\\ &&&(day:h:min) &(day:h:min)& &number\\\\ \\hline 1&1625&1975&069:21:36&071:17:42&3515&18\\\\ 2&1625&1975&077:00:45&078:20:52&2938&22\\\\ 3&1632&1975&260:00:35&261:20:39&2979&14\\\\ 4&1632&1975&267:03:52&268:23:56&3083&0\\\\ 5&1639&1976&085:04:28&087:00:45&2245&38\\\\ 6&1639&1976&092:07:20&094:03:10&1843&23\\\\ 7&1646&1976&282:11:52&284:06:39&1188&7\\\\ 8&1653&1977&099:11:25&101:07:15&2741&0\\\\ 9&1653&1977&106:13:51&108:09:57&3376&31\\\\ 10&1660&1977&296:17:34&298:13:48&1085&30\\\\ 11&1667&1978&114:19:36&116:14:33&2393&106\\\\ 12&1667&1978&121:21:18&123:17:31&1911&89\\\\ 13&1681&1979&137:05:09&139:01:23&1894&148\\\\ 14&1695&1980&145:10:38&147:07:04&2218&229\\\\ 15&1695&1980&152:12:36&154:09:02&2559&152\\\\ 16&1709&1981&159:17:43&161:14:02&3174&58\\\\ 17&1709&1981&167:03:16&168:16:26&2562&119\\\\ 18&1723&1982&175:00:30&176:21:15&1361&112\\\\ \\hline \\end{tabular} } \\label{tab} \\end{table}   In this work, we study how the dynamical system governing the slow solar wind velocity  behaves at different phases of the solar activity cycle. We apply the method of rank order statistics of symbolic sequence \\citep{yang03} to investigate the solar wind time series data. In the next Section, we describe the datasets and their pre-processing. In Section \\ref{symb} we describe the rank order statistics of symbolic sequence technique. In Section \\ref{resu} we apply symbolic analysis technique to time series obtained from known nonlinear dynamical systems like Lorenz, R\\\"ossler, and Chua oscillators and extend the analysis to solar wind data. In the last Section we give conclusions of our study.\\\\ \\begin{figure*} \\centering \\includegraphics[width=10.0cm]{1} \\caption{Pre-processing of an example data set (time series number $16$ in table \\ref{tab}) (a) the original data, (b) the detrended data, (c) the filtered data after nonlinear noise reduction, and (d) the noise removed by nonlinear filter.} \\label{fig1} \\end{figure*} ",
        "conclusions": "\\label{conc} We used rank order statistics to analyze the time series obtained from known nonlinear systems. This method is an effective tool for analyzing short time series even in the presence of high level of external noise and is applied to $18$ solar wind time series obtained by the Helios spacecraft for the years $1975-1982$. Nonlinear filter is used to reduce the noise present in the data. We found that all the time series have least correlation with data sets $13$ and $14$, which are located near the maximum of the solar cycle. Although data set $15$ is also located very close to the data set $14$ (on the solar activity cycle), it shows significant correlation. One possible reason for this may be that data set $15$ is located in the decreasing phase of the solar activity cycle, while data sets $13$ and $14$ are located in the increasing phase of the solar activity cycle. This suggests the presence of larger fluctuations in the solar wind velocity data corresponding to the time just before the maximum of the activity cycle."
    },
    "1112/1112.1077_arXiv.txt": {
        "abstract": "We explore the effects of mergers on the evolution of massive early-type galaxies by modeling the evolution of their stellar populations in a hierarchical context. We investigate how a realistic red sequence population set up by $z\\sim 1$ evolves under different assumptions for the merger and star formation histories, comparing changes in color, luminosity, and mass. The purely passive fading of existing red sequence galaxies, with no further mergers or star formation, results in dramatic changes at the bright end of the luminosity function and color--magnitude relation. Without mergers there is too much evolution in luminosity at a fixed space density compared to observations. The change in color and magnitude at a fixed mass resembles that of a passively evolving population that formed relatively recently, at $z\\sim2$. Mergers amongst the red sequence population (``dry mergers\") occurring after $z=1$ build up mass, counteracting the fading of the existing stellar populations to give smaller changes in both color and luminosity for massive galaxies. By allowing some galaxies to migrate from the blue cloud onto the red sequence after $z=1$ through gas-rich mergers, younger stellar populations are added to the red sequence. This manifestation of the progenitor bias increases the scatter in age and results in even smaller changes in color and luminosity between $z=1$ and $z=0$ at a fixed mass. The resultant evolution appears much slower, resembling the passive evolution of a population that formed at high redshift ($z\\sim3 - 5$), and is in closer agreement with observations. We conclude that measurements of the luminosity and color evolution alone are not sufficient to distinguish between the purely passive evolution of an old population and cosmologically motivated hierarchical growth, although these scenarios have very different implications for the mass growth of early-type galaxies over the last half of cosmic history. ",
        "introduction": "\\label{sec:intro}  Galaxies in the local universe are found in a bimodal distribution in color--mass (or magnitude) space \\citep{Strateva01,Kauffmann03a,Blanton03, Baldry04}. The so-called blue cloud is made up mostly of star-forming, disk galaxies, while the majority of quiescent galaxies with early-type morphologies lie on a tight relation between color and mass known as the red sequence. The tilt in the color--mass relation (CMR) of early-type galaxies can largely be attributed to an increase of metallicity with mass, with more metal rich galaxies having redder colors \\citep{Kodama97, Gallazzi06}. The distinction between red sequence and blue cloud galaxies is particularly clear at $z<1$ \\citep[e.g.,][]{Bell04, Borch06,Faber07,Brown07}, and using optical--NIR colors or mid-IR diagnostics the populations can be separated reasonably well out to $z\\sim3$ \\citep[][]{Labbe05, Williams09,Brammer09,Whitaker11}.  Although most star formation occurs in blue cloud galaxies, the total mass in galaxies on the red sequence has grown by an order of magnitude since $z\\sim2$ \\citep{Arnouts07,Ilbert10, Brammer11} and by a factor of two since $z\\sim1$ \\citep[the epoch we focus on in this paper; e.g.,][]{Bell04, Faber07, Brown07}.  \\citet{Bell07} showed that the increase in mass on the red sequence can largely be explained by an influx of blue cloud galaxies that have recently had their star formation shut down. \\citet{Harker06} and \\citet{Ruhland09} used stellar population synthesis models to show that the scatter in a red sequence that is built up through the quenching of star formation in blue cloud galaxies is consistent with observations of the evolution of the CMR. These works did not make any assumptions on what shuts the star formation down. The merging of gas-rich disk-like galaxies is a likely mechanism for rapid morphological \\citep{Toomre77, Schweizer92,Barnes96} and color transformation, with some kind of feedback process usually invoked to prevent further gas cooling and star formation \\citep[e.g., active galactic nucleus (AGN) ``radio mode\" feedback;][]{Croton06, Bower06, Somerville08}. Most massive early-type galaxies are spheroidal in shape suggesting that they have formed through major mergers \\citep{vanderWel09}. Models in which spheroids form in gas-rich mergers, with a period of related quasar activity, have been fairly successful in matching a number of observations \\citep{Benson03, Croton06, Bower06, Somerville08, Hopkins08b, Hopkins09b}.  We maintain, and will discuss below, that the growth in mass on the red sequence has been well measured only for galaxies near the knee of the luminosity function (LF; $\\sim~L^*$); it is still unclear whether the massive end of the red sequence has seen significant growth since $z=1$. If such growth occurs, it is likely to take place through the merging of gas-poor galaxies that already lie on the red sequence since there are relatively few luminous blue galaxies from which the most massive early-type galaxies could form \\citep{Blanton06}. The role of mergers in the build up of massive early-type galaxies, particularly over the epoch from $z=1$ to $z=0$, is strongly debated. In this paper we look at how a red sequence in place by $z \\sim 1$ evolves with and without further mergers and star formation. We show that observations of color and luminosity evolution that have been taken as support for a limited role for mergers in the past are consistent with the changes we find using hierarchical models.  The outline of the paper is as follows: Section~\\ref{sec:bg} discusses previous work on the evolution of early-type galaxies, the support for a hierarchical picture of early-type evolution, and why early-type galaxies were thought to evolve passively from high redshift. Section~\\ref{sec:models} describes the model set up, building on the simple toy model presented in \\citet{Skelton09a}. In Section~\\ref{sec:results} we present the resultant evolution of the color, luminosity and mass distributions in the models and compare to observations. We consider the implications for massive galaxies in particular and compare the changes in color and magnitude to those of an ancient simple stellar population. In Section~\\ref{sec:discussion} we discuss how the various model choices affect our results and compare to previous work.  The conclusions that can be drawn on early-type evolution are summarized in Section~\\ref{sec:summary}. We adopt a cosmology with $\\Omega_M = 0.3$, $\\Omega_\\Lambda = 0 .7$ and $H_0 = 100 h$~km~s$^{-1}$~Mpc$^{-1}$ with $h = 0.7$ and use Vega magnitudes throughout.  \\subsection{Background} \\label{sec:bg}  Gas-poor or ``dry'' mergers have been observed both in the local universe \\citep{vanDokkum05,McIntosh08,Tal09} and at higher redshifts \\citep{Bell06a, Lotz06, Lin08, Williams11}, and a non-negligible fraction of massive galaxies are found to occur in close pairs that are likely to merge \\citep[][although see \\citealt{Masjedi06}, who measure very low merger rates for luminous red galaxies, LRGs]{Bell06b, Robaina10, Williams11}. Recent measurements place the fraction of $M_* \\gt 5 \\times 10^{10}{{M_{\\sun}}}$ galaxies in close pairs at 1 to 3\\% out to $z=1.2$, implying that galaxies with masses of $\\gt 10^{11}{{M_{\\sun}}}$ have undergone 0.7 mergers on average since this time \\citep[][]{Robaina10}. This merger fraction seems to be sufficient to explain the growth of the red sequence at the massive end since $z=1$ \\citep{Robaina10, Man11}. A particularly large uncertainty in interpreting measurements of the merger fraction as a growth rate is the timescale over which mergers can be identified. The timescale will be different for each measurement technique (e.g., visual or automated morphological selection, close pairs) and is dependent on the properties of the progenitor galaxies and the orbit \\citep{Lotz10a, Lotz10b}. Recent work that calibrates the observability timescale for each method using high-resolution simulations in combination with cosmological models has shown a reassuring consistency between a number of different measurements of the major merger rate and model predictions \\citep{Lotz11}.  Less direct approaches for estimating the extent to which dry mergers build up massive red galaxies, such as measuring their clustering and size evolution, support a degree of merging activity since $z=1$ but are not yet conclusive. Massive early-type galaxies have evolved substantially in size over this period (a factor of two growth for a given mass, \\citealt{vanderWel08}). This can be accounted for by the increase in size of individual red galaxies through dry mergers, combined with an increase in the number density of the red sequence galaxy population over time \\citep{vanderWel09b}.  Minor mergers may be particularly important for the size evolution of early-type galaxies, as well as a viable means of explaining the evolution of the velocity dispersion function since $z\\sim1$ \\citep{Bezanson09,vanderWel09b,Naab09, Hopkins09c, Hopkins10c}.  Clustering analyses using LRGs in the Sloan Digital Sky Survey (SDSS) and 2SLAQ surveys \\citep{Wake08} and the Bo\\\"{o}tes field \\citep{White07} showed little evolution in the clustering amplitude with redshift. Interpreting these results in the light of the halo model suggests that LRGs do not evolve purely passively but that a significant fraction of satellite galaxies (30\\%--50\\%) must either merge with the central LRG or be disrupted into the intracluster light \\citep{Conroy07}. \\citet{Tojeiro10} also examine the clustering evolution of LRGs and find evidence for luminosity growth that is inconsistent with passive evolution. The growth appears to occur preferentially at the faint end of their galaxy selection, supporting an influx of galaxies into the sample through mergers \\citep[see also][]{Tojeiro11b}.  To determine how the massive end of the red sequence evolves, one would ideally want to measure the mass function (MF); however, this requires combining the measured LF with knowledge of the mass-to-light ratio ($M/L$), which also evolves. Stellar population synthesis models are usually used to estimate $M/L$, requiring assumptions on the age and star formation histories of the galaxies. To measure these quantities, the evolution must be known --- this is one of the reasons that passive evolution (the evolution of a simple stellar population formed within a short time at high redshift, with a single metallicity) is often assumed for early-type galaxies.  There is still large uncertainty in both the measurement of the number density evolution and its interpretation. If one corrects the observed LF for the evolution expected for a purely passively fading old population, there is very little remaining evolution at the bright end  since $z\\sim1$ \\citep{Cimatti06, Wake06, Brown07, Cool08, Banerji10}. These results have been used to infer that the number density of massive red galaxies has not changed substantially over the last 8 billion years and that much of the mass is already in place by this time.\\footnote{Perhaps contrary to expectation, one of the largest difficulties in this analysis is determining the bright end of the LF at low redshift to high precision. Bright galaxies in the local universe have extended outer envelopes (possibly caused by dry merging) making it difficult to estimate the total magnitude and accurately subtract the background \\citep[e.g.,][]{Bernardi07a, Lauer07}. Small photometric errors translate into a large error in number density because of the exponential drop of the LF toward the bright end.} If the early-type galaxy population is continuously supplemented by quenched blue cloud galaxies, the assumption of slow luminosity evolution made in the interpretation the evolution of the LF will be incorrect. The build up of stellar mass through mergers may act to compensate faster luminosity and color evolution, resulting in changes that look very much like passive evolution. This is related to the so-called ``progenitor bias'' \\citep{vanDokkumFranx01}; galaxies that have recently moved onto the red sequence through morphological transformation will be missing from early-type galaxy samples at high redshift, so that at early times, any sample of early-type galaxies consists of only the oldest progenitors of today's early-type galaxies. Converting the LF of red galaxies measured by \\citet{Brown07} to a MF using the $M/L$ evolution from Fundamental Plane measurements \\citep{vDvdMarel07} rather than using stellar population models gives rise to stronger evolution \\citep{Robaina10}. Other authors have recently found a similar increase in number density (approximately a factor of four) at the massive end using the galaxy MF \\citep{Matsuoka10}. A recent analysis that uses clustering measurements to select likely progenitors of early-type galaxies at high redshift measures a factor of $\\sim5$ growth in number density, supporting substantial growth through dry mergers \\citep{Padilla11}.  The uniformity of the stellar populations in early-type galaxies suggested by the tight CMR in clusters has long been used as an argument that they formed relatively quickly at high redshift \\citep{Ellis97, BKT98, Andreon11}. However, recent work on the evolution of the red sequence has shown that it is difficult to match both the normalization and change in color with a purely passive, old model \\citep{Kriek08, Ruhland09, Whitaker10}.  The small scatter in the colors of early-types in clusters was thought to impose a limit to the amount of growth that could occur via merging \\citep{BKT98}. In \\citet[SBS09 hereafter]{Skelton09a} we explored how the amount of dry merging expected in a standard hierarchical model affects the scatter and slope of the CMR. We assumed that major gas-rich mergers move galaxies from the blue cloud onto the red sequence, placing them onto a ``creation red sequence'' given by the local relation between color and magnitude. We calculated where red galaxies at low redshift lie in color-magnitude space by combining the fluxes of their progenitors, assuming that merging of galaxies on the red sequence does not involve further star formation. By using the local CMR as a reference, the model essentially assumed that early-type galaxies are evolving passively and at the same rate. With this simple model we showed that dry merging can reduce the scatter and flatten the bright end of the CMR. The resulting relation was found to agree well with the observed red sequence from the SDSS.  In this paper we incorporate more realistic evolution of the stellar populations into the model described above to test how mergers affect the color, luminosity, and mass evolution of the red sequence since $z =1$.  Our aim is not to make a complete, accurate model of the galaxy distribution; the point is rather to show how the progenitor bias expected from hierarchical growth in a standard cold dark matter cosmology affects the evolution of massive galaxies. The models, very much in the spirit of earlier work by \\citet{BKT98, vanDokkumFranx01,Blanton06,Skelton09a}, and others, use stellar population synthesis modeling to follow the luminosity evolution of galaxies, but set the model galaxies into a hierarchical context using merger trees extracted from a semi-analytic model (SAM) of galaxy formation. We assume that the red sequence is built up mainly through major gas-rich mergers. With a reasonably realistic red sequence population in place at $z =1$, we explore the consequences of subsequent gas-poor and gas-rich merging. We compare this to the evolution of the same population undergoing only passive evolution with no mergers after $z=1$. We show that at fixed mass, dry merging and recent additions to the red sequence decrease the change in color and luminosity. The resultant evolution is consistent with that of a population that formed early and evolved passively, {\\it even though there has been significant merging activity}. ",
        "conclusions": "\\label{sec:summary}  We have used galaxy merger histories from a hierarchical model, in conjunction with simple star formation histories and stellar population synthesis models, to explore the color and magnitude evolution of red sequence galaxies in several scenarios.  We use the galaxy merger trees from the S08 SAM combined with stellar evolutionary tracks from the \\citet{BC03} model, making simple assumptions on the relation of mass to metallicity and the redshift at which galaxy formation begins. We assume that major wet mergers are effective at quenching star formation but that star formation may also shut down after some period of time, independently of mergers. With this model setup, we examine the role of recent mergers and the effect of adding younger populations to the red sequence at late times.  A red sequence population with a realistic color--magnitude distribution at $z=1$ can be formed through the passive fading of galaxies that form in a single burst relatively recently ($z_f\\sim2$) or through more extended star formation, shut down by major mergers at different times, in galaxies that began forming earlier ($z_f \\sim 4$). Galaxies that form all their stars at high redshifts have colors that are too red by $z=1$ and produce a CMR that has smaller scatter than observed, but the slow reddening they experience thereafter gives approximately the observed amplitude of color evolution.  A population with approximately the correct color at $z=1$ evolves very rapidly over the last half of cosmic history if no further mergers occur and there is no new mass added to the red sequence (Model I). This results in a color--magnitude relation that is too red at low redshift, with too few bright red galaxies and too much luminosity evolution at fixed space density compared to observations. At a fixed mass, the change in color and magnitude predicted by such a model are approximately the same as the changes predicted for a purely passively evolving population that formed at $z_f\\sim2$.  Dry mergers occurring between $z=1$ and $z=0$ slow down the luminosity and color evolution of the early-type galaxy population. The mass (and light) added through mergers counteracts the fading of the luminosities expected from stellar evolution when no new stars are forming. The result is a slightly bluer CMR at $z=0$ that is closer to the observed relation, and more realistic evolution of the LF and MF (Model II). By allowing star formation to continue in some galaxies after $z=1$, the changes in a population's color and luminosity are reduced even further (Model III). This late morphological transformation is known as progenitor bias \\citep{vanDokkumFranx01}. Both the CMR and the LF evolution of this model are consistent with the observed evolution. The smaller changes in color and magnitude at a fixed mass replicate the behavior of a passively evolving population that formed in a single burst at high redshift ($z_f = 3 - 5$).  At the bright end of the red sequence, the seemingly slow passive evolution of luminosity and color is equally well interpreted as the passive evolution of an ancient population, or as the cosmologically motivated hierarchical growth of an evolving population.  Although they predict very similar evolution, the implications of these scenarios for the growth of stellar mass are very different. Diagnostics other than the LF, such as the number density evolution estimated using dynamical $M/L$ information, the growth rate determined from close-pair statistics and the evolution of the clustering strength of early-type galaxies lend support to the hierarchical picture."
    },
    "1112/1112.3591_arXiv.txt": {
        "abstract": "Within the next few years, GAIA and several instruments aiming at imaging extrasolar planets will see first light. In parallel, low mass planets are being searched around red dwarfs which offer more favourable conditions, both for radial velocity detection and transit studies, than solar-type stars. Authors of the model atmosphere code which has allowed the detection of water vapour in the atmosphere of Hot Jupiters review recent advancement in modelling the stellar to substellar transition. The revised solar oxygen abundances and cloud model allow for the first time to reproduce the photometric and spectroscopic properties of this transition.  Also presented are highlight results of a model atmosphere grid for stars, brown dwarfs and extrasolar planets. ",
        "introduction": "Since spectroscopic observations of very low mass stars (late 80s), brown dwarfs (mid 90s) and extrasolar planets (mid 2000s) are available, one of the most important challenges in modelling their atmospheres and spectroscopic properties lies in high temperature molecular opacities and cloud formation. K dwarfs show the onset of formation metal hydrides (starting around $T_\\mathrm{eff} \\sim 4500$\\,K), TiO and CO (below $T_\\mathrm{eff} \\sim 4000$\\,K), while water vapour forms in early M dwarfs ($T_\\mathrm{eff}\\sim 3900-2000$\\,K), and methane, ammonia and carbon dioxide are detected in late-type brown dwarfs ($T_\\mathrm{eff}\\sim 300-1600$\\,K) and in extrasolar giant planets. The latter are either observed by transit ($T_\\mathrm{eff}\\sim 1000-2000$\\,K depending on the spectral type of the central star and the distance to the star) or by imaging (young planets of $T_\\mathrm{eff}\\sim 300-2000$\\,K depending on their mass and age).  The modelling of the atmospheres of very low mass stars (hereafter VLMs) has evolved (as here illustrated with the development of the \\texttt{PHOENIX} atmosphere code, which has allowed the detection of water vapour in extrasolar planets' atmospheres by Barman et al.\\ 2007, 2008\\nocite{Barman2007,Barman2008}) with the extension of computing capacities from an analytical treatment of the transfer equation using moments of the radiation field \\cite[]{AllardPhDT90}, to a line-by-line opacity sampling in spherical symmetry \\cite[]{ARAA97,NGa,NGb}, and more recently to 3D radiation transfer \\cite[]{SHB2010}. In parallel to detailed radiative transfer in an assumed static environment, hydrodynamical simulations have been developed to reach a realistic representation of the granulation and its induced line shifts for the sun and sun-like stars \\citep[see][]{Freytag2011} by using a non-grey (multi-group binning of opacities) radiative transfer with a pure blackbody source function (scattering is neglected).  \\begin{figure}[!ht] \\plotfiddle{allard_f_fig1.eps}{8.5cm}{0}{54}{50}{-175}{-88} \\caption{ Synthetic spectra compared to the IR SED of VB10 using identical model parameters ($T_\\mathrm{eff}=2800$\\,K, $\\mathrm{log}\\,g=5.0$) and a resolution of 50\\,{\\AA} with different water vapour opacity sources: the Base grid by \\cite{AH95} using \\cite{LudwigH2O}; a test using the 1994 version of the Niels Bohr Institute \\cite[]{NBIH2O}; the NextGen grid by \\cite{AllardH2O94,ARAA97} and \\cite{NGa,NGb}  using the University College London database \\cite[]{UCL94H2O}; and the AMES-Cond/Dusty grid by \\cite{AllardTiOH2O2000,Allard01} using the NASA-Ames Center database \\cite[]{AMESH2O}. All models (except the NextGen/UCL case) underestimate the flux at $K$ (ca.\\ $2.0-2.4\\,\\mu m$) by 0.1 to 0.2 dex.} \\label{f:allard_f1_water} \\end{figure} ",
        "conclusions": "We report progress of the development of a new model atmosphere grid for stars, brown dwarfs and young planets, named BT-Settl. It has been computed using the \\texttt{PHOENIX} code updated for: i) the \\cite{BT2H2O} BT2 water, the \\cite{stdsIAU211} STDS methane, the \\cite{Sharp2007} ammonia and the \\cite{cdsd1000} CDSD-1000 CO$_2$ opacity line lists, ii) the solar abundances revised by \\cite{Asplund09}, and iii) a cloud model accounting for more detailed supersaturation and RHD mixing.  The grid is covering the whole range of stars to young planets 400\\,K $< T_\\mathrm{eff} <$ 70,000\\,K; -0.5 $<\\mathrm{log}\\,g<$ 5.5; and $-4.0 <$ [M/H] $< +0.5$, including values of the $\\alpha$-element enhancement   (supernovae enrichment of the star forming material) between +0.0 and +0.6. Models are available at the \\texttt{PHOENIX} simulator website \\texttt{http://phoenix.ens-lyon.fr/simulator/} and are in preparation for publication to serve among others the GAIA, MUSE and SPHERE/GPI/P1640 instruments to come online in the near futur. Corresponding evolution models are expected for 2012.  We found the previously used NextGen models to systematically overestimate $T_\\mathrm{eff}$ below 3500\\,K by as much as 500\\,K.  The water vapour opacity profile has converged with the most recent line lists reproducing laboratory results, but could not explain this discrepancy. The solution came instead from the revision of the solar abundances which changes the strengh of the water vapour absorption bands, and therefore allows the reproduction of the spectroscopic and photometric properties of M dwarfs as late as M6. Later-type M dwarfs are affected by dust formation and cloud modelling is important to understand their properties.  We find that the Rossow cloud model allows, with revisions to the supersaturation and mixing, the reproduction of the stellar-substellar transition.  A small offset persists however in the M-L transition.  It is possible that all the current cloud models are not efficient enough in producing dust at the onset of cloud formation regime.  Detailed nucleation studies could allow in the futur to resolve this issue. Other uncertainties affect the current cloud modelling such as the assumption of spherical non-porous grains while grains form as fractals in the laboratory. Constraining the models remains therefore very important.  Beyond cloud modelling and molecular opacities, model atmospheres for these objects require reaction rates for the most abundant molecules and/or most important absorbers.  Furthermore, these atmospheres are composed of molecular hydrogen which constitute the main source of collisions.  Also needed are therefore collision rates (by H$_2$) and corresponding damping constants for the broadening molecular lines.  In order to say something about the spectral variability of VLMs, brown dwarfs and planets, 3D global or \"star in a box\" RHD simulations with rotation will be required.  This is our current project supported by the French ``Agence Nationale de la Recherche'' for the period 2010-2015."
    },
    "1112/1112.4652_arXiv.txt": {
        "abstract": "{The MOND paradigm to the missing mass problem requires introducing a functional that is to be identified through observations and experiments. We consider the AQUAL theory as a realization of the MOND.   We show that the accurate value of the Earth GM measured by  the Lunar Laser Ranging measurements   and that  by various artificial Earth satellites,  including the accurate tracking of the LAGEOS satellites,  constrain this functional such that some of the chosen/proposed functional are refuted. }    \\begin{document} ",
        "introduction": "The missing mass problem in galaxies can be resolved either by the Modified  Newtonian Dynamics (MOND), or the Modified Gravity  (MOG) or the dark matter hypothesis.  The first two paradigms require introducing a functional that interpolates the Newtonian regime to the MOND regime. This functional is to be identified by examining the data.  We note that the accurate value of the mass ratio of the Sun/(Earth+Moon) from the Lunar Laser Ranging combined with the Solar GM and the lunar GM from lunar orbiting spacecrafts \\cite{LLR} yields the effective   gravitational mass of the Earth in an Earth-centered reference frame with the precision of one part in $10^8$ : \\begin{equation}\\label{GMLLR} GM^{LLR}_{Earth}(r_{_{LD}}) = 398600.443\\pm 0.004 \\frac{km^3}{s^2} \\end{equation} where $r_{_{LD}}$ represents the Lunar distance: the average distance  between the Earth and Moon. The effective gravitational mass is defined to be the gravitational field multiplied by $r^2$ where $r$ is distance from the center of the earth.  The effective gravitational mass of the Earth has also been measured by various artificial Earth satellites , including the accurate tracking of the LAGEOS satellites orbiting the Earth in nearly circular orbits with semimajor axes about twice the radius of the Earth \\cite{LAGEOS}: \\begin{equation}\\label{GMLAGEOS} GM^{LAGEOS}_{Earth}(2 r_{Earth})=  398600.4419 \\pm 0.0002 \\frac{km^3}{s^2} \\end{equation} where $r_{Earth}$ stands for the radius of the Earth.  In this paper, we study the earth Lunar system constraints on the interpolation function. In section \\ref{MOND}, we first briefly review the MOND paradigm.  We present the choices for the interpolating function. The section \\ref{MONDLaser} proves that a relativistic version of MOND does not affect the null geodesics at the precision of the Lunar Laser Ranging measurements.  The leading effect of MOND is changing the ratio of the `Earth-Moon distance' to the `LAGEOS satellites-Earth distance'.  In section \\ref{constraints}, we translate the consistency of  \\eqref{GMLLR} and \\eqref{GMLAGEOS} to the first Lunar system constraint on the MOND. We show that the Lunar system constraints refute some of the choices of the interpolating function. ",
        "conclusions": "We have translated the GM mass measured by the  Lunar Laser ranging measurements and LAGEOS satellites into a  constraint on the interpolating function of modified gravities. The Lunar constraint is given in \\eqref{LunarCons}. We have applied the Lunar system's constraint on all the families of the interpolating function  considered in the literature and reviewed in \\cite{Famaey:2011kh}. We have proven that the Lunar system constraint is non-trivial for a set of families.  In particular we have proven that the Lunar system constraint refutes the simple interpolating function at the confidence level of   the  $3.69$ sigmas, and completely refutes the Bekenstein choice."
    },
    "1112/1112.4646_arXiv.txt": {
        "abstract": "We investigate the effect of primordial helium on hydrogen reionisation using a hydrodynamical simulation combined with the cosmological radiative transfer code {\\tt CRASH}.  The radiative transfer simulations are performed in a $35.12h^{-1}$ comoving Mpc box using a variety of assumptions for the amplitude and power-law extreme-UV (EUV) spectral index of the ionising emissivity at $z>6$. We use an empirically motivated prescription for ionising sources which, by design, ensures all of the models are consistent with constraints on the Thomson scattering optical depth and the metagalactic hydrogen photo-ionisation rate at $z \\sim 6$.  The inclusion of helium slightly delays reionisation due to the small number of ionising photons which reionise neutral helium instead of hydrogen.  However, helium has a significant impact on the thermal state of the IGM during hydrogen reionisation.  Models with a soft EUV spectral index, $\\alpha=3$, produce IGM temperatures at the mean density at $z\\sim 6$, $T_{0}\\simeq 10500\\rm\\,K$, which are $\\sim 20$ per cent higher compared to models in which helium photo-heating is excluded.  Harder EUV indices produce even larger IGM temperature boosts by the end of hydrogen reionisation.  A comparison of these simulations to recent observational estimates of the IGM temperature at $z\\sim 5$--$6$ suggests that hydrogen reionisation was primarily driven by population-II stellar sources with a soft EUV index, $\\alpha\\la 3$.  We also find that faint, as yet undetected galaxies, characterised by a luminosity function with a steepening faint-end slope ($\\alpha_{\\rm LF} \\leq -2$) and an increasing Lyman continuum escape fraction ($f_{\\rm esc}\\sim 0.5$), are required to reproduce the ionising emissivity used in our simulations at $z>6$.  Finally, we note there is some tension between recent observational constraints which indicate the IGM is $>10$ per cent neutral by volume $z\\sim 7$, and estimates of the ionising emissivity at $z=6$ which indicate only $1$--$3$ ionising photons are emitted per hydrogen atom over a Hubble time at $z=6$.  This tension may be alleviated by either a lower neutral fraction at $z\\sim 7$ or an IGM which still remains a few per cent neutral by volume at $z=6$. ",
        "introduction": "\\label{sec:intro}  The last decade has witnessed the establishment of the two key pieces of observational evidence which presently shape our empirical understanding of the hydrogen reionisation epoch.  The first is the Thomson scattering optical depth inferred from observations of the cosmic microwave background (CMB).  This measurement provides a constraint on the integrated reionisation history, and is consistent with hydrogen reionisation beginning no later than $z=10.6\\pm 1.2$ \\citep{Komatsu11}.  The second is the signature of \\HI \\Lya absorption in the spectra of high redshift quasars; observations of the \\cite{GunnPeterson65} trough indicate that the intergalactic medium (IGM) is largely ionised by redshifts less than $z \\simeq 6$ (\\citealt{Becker01,Fan06}).  These observational data therefore broadly constrain hydrogen reionisation to the redshift range $z\\simeq 6-12$.  Despite this progress, however, a detailed determination of the timing and extent of hydrogen reionisation, as well as the exact nature of the sources responsible for driving this process, remains elusive. Due to the integral nature of the CMB constraint a wide range of extended reionisation histories are compatible with the Thomson scattering optical depth measurement.  The small neutral hydrogen fractions, $x_{\\rm HI}\\sim 10^{-4}$, at which \\Lya absorption saturates also leave room for alternative interpretations of the quasar data (e.g. \\citealt{Songaila04,Becker07}).  Furthermore, \\cite{Mesinger10} has recently pointed out that even at $z\\sim 5$--$6$, the IGM may still harbour large patches of neutral hydrogen; the number of currently known quasar sight-lines is insufficient to fully rule out this possibility with intergalactic \\Lya absorption observations alone.   One route to making further progress is therefore developing detailed simulations (e.g. \\citealt{Ciardi.Ferrara.White_2003,Iliev.Mellema.Shapiro.Pen_2007,Trac.Cen_2007,Finlator09b,Aubert10,Baek_etal_2010}), and semi-numerical/analytical models (e.g. \\citealt{Choudhury.Ferrara_2006,Zahn_etal_2007,Mesinger07, Santos_etal_2010,Shull12,Raskutti12}) which can be compared to these data to make inferences about the reionisation process.  However, most existing numerical simulations do not explore the effect of hard, helium ionising photons on the thermal state of the IGM during hydrogen reionisation (although see e.g. \\citealt{Sokasian2002,Paschos.Norman.Bordner.Harkness_2008,McQuinn2009} for treatments of \\HeII reionisation at $z\\simeq 3$).  This renders the comparison of these models to measurements of the IGM temperature at $z<6$ problematic.  In addition, many numerical models significantly over-predict the number of ionising photons in the IGM relative to observational constraints on the \\HI photo-ionisation rate from the \\Lya forest at $z\\sim 5$--$6$.  These data are consistent with $\\sim 1$--$3$ ionising photons emitted per hydrogen atom over a Hubble time at $z=6$.  As a result, in order for hydrogen reionisation to complete by $z=6$ {\\it and} simultaneously match observational constraints from the CMB and the background photo-ionisation rate at $z\\leq 6$, reionisation must be an extended process where the ionising emissivity increases at $z>6$ (\\citealt{MiraldaEscude2003,Meiksin05,BoltonHaehnelt07,HaardtMadau2011,McQuinn11}). Correctly matching these ``post-reionisation'' constraints therefore has important implications for reionisation and the properties of the ionising sources in the early Universe.  In this work we address these issues using radiative transfer simulations of reionisation performed using the code {\\tt CRASH} \\citep{Ciardi_etal_2001,Maselli.Ferrara.Ciardi_2003,Maselli.Ciardi.Kanekar_2009, Partl_etal_2011}.  {\\tt CRASH} is a 3D Monte Carlo based code which follows the propagation of ionising photons (from both point sources and diffuse radiation) and self-consistently calculates the evolution of the gas temperature and ionisation state of hydrogen and helium in the IGM.  Our approach differs from previous studies in two important ways.  Firstly, we include the effect of helium ionising photons on the progression of hydrogen reionisation.  This is especially important for computing the thermal state of the IGM (e.g. \\citealt{Tittley07,Cantalupo11,PawlikSchaye11}), and it enables us to directly compare our simulations to recent measurements of the IGM temperature at $z=5$--$6$ (\\citealt{Becker11,Bolton12}). Secondly, instead of using a numerical sub-grid model for the sources of ionising photons, the ionising emissivity in our simulations is matched to the CMB and \\Lya forest observational constraints by design.  The goal of this empirical approach is to explore the consequences of satisfying these observational constraints for reionisation models from the outset, instead of tuning free parameters and/or sub-grid prescriptions within the simulations.  The structure of this paper is as follows.  We begin in Section 2 with a discussion of the empirically motivated reionisation models used in our analysis, and continue in Section 3 with a description of our numerical simulations.  In Section 4 we demonstrate that our RT simulations match the observational constraints on the Thomson scattering optical depth and photo-ionisation rate inferred from the \\Lya forest at $z\\sim 6$, before going on to discuss in detail the effect of including helium on the ionisation and thermal state of the IGM in Section 5.  We perform a comparison of our simulations to the observational data in Section 6 and discuss the implications for the the properties of ionising sources at $z>6$.  Finally, we summarise and conclude in Section 7.  An appendix presenting selected numerical convergence tests of our simulations is provided at the end of the paper. Throughout the paper, the following cosmological parameters are used: $\\Omega_{\\Lambda}=0.74$, $\\Omega_m=0.26$, $\\Omega_b=0.024 h^2$, $h=0.72$, $n_s=0.95$ and $\\sigma_8=0.85$, where the symbols have the usual meaning. ",
        "conclusions": "In this work we have investigated the impact of helium on hydrogen reionisation using three dimensional, multi-frequency RT simulations. We performed five simulations using different models for the amplitude and spectral shape of the ionising emissivity during reionisation.  By design, all our models are consistent with measurements of the Thomson scattering optical depth and the metagalactic hydrogen photo-ionisation rate at $z \\sim 6$.  This empirical approach enables us to explore the consequences of satisfying these observational constraints for reionisation.  The main outcomes of this study may be summarised as follows.  \\begin{itemize} \\item The evolution of the volume averaged \\HII fraction, $x_{\\rm HII}$, is very similar for all models with the same hydrogen ionising emissivity independent of the EUV spectral index.  However, the spectral energy distribution has a strong impact on the evolution on the volume averaged \\HeII and \\HeIII fractions, $x_{\\rm HeII}$ and $x_{\\rm HeIII}$. Models with a soft power-law EUV index, $\\alpha=3$, produce a much lower $x_{\\rm HeIII}$ compared to models in which harder photons are present. The inclusion of helium in the RT simulations furthermore slightly delays reionisation due to the small number of ionising photons which reionise neutral helium instead of hydrogen.\\\\  \\item The choice of EUV spectral index has a significant effect on the evolution of the volume averaged IGM temperature during reionisation.  At $z \\simgt 10$, model ${\\mathcal E}$1.2-$\\alpha$1.8-H (without helium) has a volume averaged temperature which is $\\sim 10$ per cent higher than the corresponding model including helium, ${\\mathcal E}$1.2-$\\alpha$1.8, due to the slightly larger volume of the IGM which is photo-ionised by this time.  However, at lower redshift the inclusion of \\HeII photo-ionisation results in a higher volume averaged temperature for ${\\mathcal E}$1.2-$\\alpha$1.8.  In comparison, despite exhibiting behaviour similar to ${\\mathcal E}$1.2-$\\alpha$1.8 and ${\\mathcal E}$1.2-$\\alpha$1-3 for the evolution of the \\HII filling factor, the softer ionising spectrum used in ${\\mathcal E}$1.6-$\\alpha$3 produces volume averaged temperatures which are 20-30 per cent lower than ${\\mathcal E}$1.2-$\\alpha$1.8.  This is partly because the volume filling factor of \\HeIII is smaller in this model, but also because the softer ionising photons produce less photo-heating.\\\\  \\item The temperature (and ionisation fraction) distributions in the simulations exhibit a significant amount of scatter at all redshifts.  This scatter arises from the different reionisation history of each cell in the simulations (i.e. inhomogeneous reionisation) as well as the fact that we do not use monochromatic photons, but rather a spectral energy distribution which can also be hardened by spectral filtering.  This differs significantly from the tight, power-law temperature-density relation expected for an optically thin IGM following reionisation.  We find the temperature-density relation for ionised gas is typically isothermal or mildly inverted during hydrogen reionisation. \\\\  \\item A comparison with recent estimates of the IGM temperature at $z \\sim 5-6$ from \\Lya absorption in the spectra of high redshift quasars suggests that hydrogen reionisation is mainly driven by sources with a soft spectral energy distribution, $\\alpha \\leq 3$. The simulations with harder spectral indices produce temperatures which are larger than the observational constraints.  We conclude that population-II stellar sources are likely to provide most of the ionising photons during reionisation, and the spectral shape of the ionising background must harden at $z<6$ due to the increasing importance of quasars if \\HeII reionisation is to complete by $z\\simeq 3$.  If sources with rather hard spectra, such as mini-quasars or population-III stars were responsible for reionising hydrogen, their contribution must be either small or confined to $z\\geq 9$ to give sufficient time for the IGM temperature to cool and for doubly ionised helium to recombine by $z\\simeq 6$.\\\\  \\item In order to reproduce the ionising emissivity in our simulations at $z>6$, we find that the best fit to the evolution of the galaxy luminosity function presented by \\cite{Bouwens11} at $4<z<8$ requires extrapolation to faint UV magnitudes ($M_{\\rm UV} =-10$), as well as a steepening faint end slope $\\alpha_{\\rm LF} \\leq -2$ and a high Lyman continuum escape fraction $f_{\\rm esc}=0.5$. Faint, low mass galaxies are therefore necessary for providing the required number of photons during reionisation, in agreement with several other complementary studies.\\\\  \\item There is some tension between the empirically motivated ionising emissivity used in our simulations and recent observational constraints on the IGM neutral fraction which indicate that $x_{\\rm HI}>0.1$ at $z\\sim 7.1$.  The ionising emissivity inferred from the \\Lya forest at $z=6$ is equivalent to only $1$--$3$ ionising photons emitted per hydrogen atom over a Hubble time, implying reionisation is extended and that the emissivity must increase at $z>6$ if reionisation is to complete by $z=6$ (\\citealt{MiraldaEscude2003,BoltonHaehnelt07}).  However, an increasing emissivity at $z>6$ is inconsistent with a large neutral fraction at $z\\sim 7$ in our simulations unless the observations are overestimates or the IGM remains a few per cent neutral by volume at $z=6$ (see e.g \\citealt{Mesinger10}).   \\end{itemize}  \\noindent Our results highlight the importance of reproducing post-reionisation constraints such as the IGM temperature and background photo-ionisation rate for constraining reionisation models.  While these simulations were designed mainly to investigate the impact of helium on hydrogen reionisation and the sources of ionising photons at high redshift, the volume used is too small to allow a more detailed discussion on helium reionisation (which is thought to be driven by quasars and to be complete at $z \\sim 2.5-3$) and a more accurate comparison with observational constraints at $z<6$.  We will postpone this further analysis to a future work, together with a more thorough investigation of the impact of unresolved small scale high density peaks.  The latter will be particularly important for regulating the tail-end of the reionisation process and for setting the thermal state of the IGM by absorbing photons close to the \\HI and \\HeII ionisation edges.  Including these effects in numerical models is therefore necessary for refining the comparison of simulations with observations at $z<6$."
    },
    "1112/1112.4470_arXiv.txt": {
        "abstract": "We use a  nearly complete sample of Gamma Ray Bursts (GRBs) detected by the {\\it Swift} satellite to study the correlations between the spectral peak energy \\ep\\ of the prompt emission, the isotropic energetics \\eiso\\ and the isotropic luminosity \\liso. This GRB sample is characterized by a high level of completeness in redshift (90\\%). This allows us to probe in an unbiased way the issue related to the physical origin of these correlations against selection effects. We find that one burst, GRB 061021, is an outlier to the \\epeiso\\ correlation. Despite this case, we find strong \\epeiso\\ and \\epliso\\ correlations for the bursts of the complete sample. Their slopes, normalizations and dispersions are consistent with those found with the whole sample of bursts with measured redshift and \\ep. This means that the biases present in the total sample commonly used to study these correlations do not affect their properties.  Finally, we also find no evolution with redshift of the \\epeiso\\ and \\epliso\\ correlations. ",
        "introduction": "One of the most debated issues in Gamma Ray Bursts (GRBs) concerns the existence of correlations among the spectral parameters of the prompt emission and its energetic and luminosity. Three robust correlations have been identified. Each of them involves the rest frame peak energy \\ep\\ of the $\\nu F_{\\nu}$ prompt spectrum. This quantity strongly correlates with (i) the isotropic energy \\eiso\\ (Amati relation, Amati et al. 2002), (ii) the isotropic peak luminosity \\liso\\ (Yonetoku relation, Yonetoku et al. 2004) and (iii) the collimation--corrected energy $E_\\gamma$ (Ghirlanda relation, Ghirlanda et al. 2004b; see also Liang \\& Zhang 2005). These correlations are valid for long GRBs. There are indications that short bursts (i.e. bursts with observed duration $<$2 s) obey the very same \\epliso\\ relation defined by long events but they are inconsistent with the \\epeiso\\ correlation (Ghirlanda et al. 2009; Amati 2010).  The spectral--energy correlations have relevant implications both for the theoretical understanding of the burst physics and for the application of GRBs as cosmological tools (Ghirlanda et al. 2004a; Ghirlanda et al. 2006). While there is no consensus yet on the interpretation of these correlations (e.g. Eichler \\& Levinson 2004; Levinson \\& Eichler 2005; Rees \\& M\\'esz\\'aros 2005; Thompson, M\\'esz\\'aros \\& Rees 2007), an interesting progress has been recently made by studying the comoving properties of GRBs (Ghirlanda et al. 2012; Panaitescu et al. 2009).  The \\epeiso\\ and \\epliso\\ correlations have been discovered with a very small sample of 12--16 GRBs detected by  {\\it Beppo}SAX and/or BATSE (Amati et al. 2002; Yonetoku et al. 2004). They have been subsequently confirmed by adding bursts detected by other satellites (e.g. {\\it HETE II}, {\\it Konus}/WIND, {\\it INTEGRAL}, {\\it Swift} and {\\it Fermi}). At present the sample of bursts defining the \\epeiso\\ and \\epliso\\ correlations comprises more than 130 events, with only two well known (very peculiar) outliers (but see Ghisellini et al. 2006).  Despite the considerably large number of bursts consistent with these correlations, their physical origin is debated. Some authors claim that they are the result of instrumental selection effects (Nakar \\& Piran 2005; Band \\& Preece 2005; Butler et al. 2007; Shahmoradi \\& Nemiroff 2011). Other studies (Ghirlanda et al. 2008; Nava et al. 2008) quantified the possible instrumental selection biases finding that, even if they do affect the sample, they cannot be responsible for the existence of the spectral--energy correlations. Ghirlanda et al. (2010; 2011a; 2011b) added an important piece of information to this debate. They analysed the time--resolved spectra of both long and short bursts detected by the {\\it Fermi} satellite. They found that the time--resolved \\ep\\ correlates with the luminosity during the temporal evolution of the burst (and for bursts without $z$ they derive a correlation between the time--resolved \\epo\\ and the flux). This is a feature common to short and long GRBs. The time--resolved correlations found within individual GRBs are similar to each other and they are similar to the \\epliso\\ relation defined by the time--integrated spectral properties of different bursts (i.e. the Yonetoku relation). This is a strong argument in favour of the physical origin of the correlations because it is very unlikely that within a single burst the instrumental selection effects (e.g. the detector threshold) or possible effects related to the dependence of the peak energy, energetic and luminosity on the redshift can play any role.   GRBs added to the \\epeiso\\ and \\epliso\\ correlations need to have their redshifts measured. This raised the suspect that these correlations might be biased by the high level of incompleteness in redshift of the samples of GRBs so far used. In this paper we study, for the first time, the  \\epeiso\\ and  \\epliso\\ correlations with a sample of GRBs with a high level of completeness (90\\%). The sample is presented by Salvaterra et al. (2012) and it comprises 58 events detected by  {\\it Swift} with a 1--s peak photon flux $P\\geq2.6$ ph s$^{-1}$cm$^{-2}$ (integrated in the 15--150 keV energy range). Fifty--two out of 58 GRBs have measured redshift $z$, while in 3 cases an upper limit on $z$ can be set from the photometry of the afterglow or of the host galaxy. The completeness level then increases up to $\\sim$95\\% by considering the redshift constraints imposed by the detection of the afterglow or host galaxy in some optical filters.  The complete sample used in this work is particularly appropriate to study the possible evolution of the \\epeiso\\ and \\epliso\\ correlations with redshift. Li (2007), considered a sample of 48 GRBs (from Amati 2006, 2007) and investigated if the \\epeiso\\ correlation evolves with redshift, finding that it does, becoming steeper at higher redshifts. Ghirlanda et al. (2008) repeated the same test with a larger sample of 76 bursts and found no evidence of evolution with redshift. These results, however, were obtained by considering highly incomplete GRB samples.  Our goals are (i) to study how \\ep\\ is correlated to \\eiso\\ and \\liso\\ within the complete sample of bursts and (ii) to study the possible evolution of these correlations with redshift through this unbiased sample. In $\\S$2 we present the spectral properties of the GRBs composing the complete sample. We study the \\epeiso\\ and \\epliso\\ correlations defined by this sample in $\\S3$ and compare their properties with those defined by the (incomplete) larger sample of GRBs with measured $z$ and \\epo. In $\\S4$ we discuss and summarise our conclusions. Throughout the paper we assume a standard cosmology with $h=\\Omega_{\\Lambda}=0.7$ and $\\Omega_{m}=0.3$. ",
        "conclusions": "The definition of a sample of GRBs with a high level of completeness (90\\%) allows to address several open issues regarding population studies of GRBs. In this paper we use this complete sample to study the spectral--energy correlations in an unbiased way. Out of 58 GRBs satisfying the selection criteria 46 have well defined spectral properties and measured redshift. In particular, the rest frame peak energy \\ep\\ can be determined and its relation with \\eiso\\ and \\liso\\ can be investigated. We find that this sample defines two strong correlations: (i) \\epeiso\\ (rank's correlation coefficient $\\rho=0.76$ and chance probability P$_{\\rm chance}=7\\times 10^{-10}$) and (ii) \\epliso\\ ($\\rho=0.65$ and P$_{\\rm chance}=1\\times 10^{-6}$).  The slope of the correlations defined by the complete sample are 0.61 and 0.53 for the Amati and the Yonetoku correlations, respectively. These slopes are consistent within 1$\\sigma$ (2$\\sigma$) with the slopes of the \\epliso\\ (\\epeiso) correlation defined by bursts not satisfying the completeness criteria (complementary sample). From Table \\ref{tab2} we also note that the scatter of the correlations defined with the complete, total and complementary sample are similar among themselves and are about $\\sigma_{\\rm sc}=0.25$ and $\\sigma_{\\rm sc}=0.30$ for the Amati and Yonetoku correlation. We also investigated possible instrumental effects in the correlations found in this complete sample. We defined three sub-samples based on the instrument which derived the burst spectral parameters (see Table 1). In particular, we study the Amati and Yonetoku correlations within the sample of Swift/BAT bursts (10 events), Konus bursts (18 events) and Fermi/GBM bursts (13 events). The results derived from different sub-samples are very similar and consistent within the errors. In fact, we found that the best fit slopes (normalizations) of the Amati correlation are 0.60$\\pm$0.05 (-28.97$\\pm$2.67),  0.59$\\pm$0.08 (-28.41$\\pm$4.19) and 0.60$\\pm$0.11 (-29.19$\\pm$5.77) for the Swift, Konus and Fermi samples respectively. For the Yonetoku correlation we find: 0.44$\\pm$0.11 (-20.76$\\pm$5.95), 0.52$\\pm$0.08 (-24.61$\\pm$4.49) and 0.54$\\pm$0.12 (-25.87$\\pm$6.37).   Our study outlines the presence of one GRB (061021) that lies at 3$\\sigma$ limit (or more) of the tested correlations. In particular, it is an outlier to the Amati correlation.  GRB 061021 has been detected by {\\it Swift}/BAT (Moretti et al. 2006) and by {\\it Konus}/WIND (Golenetskii et al. 2006). It shows a single pulse with a duration of $\\sim10$ s, followed by a weak tail seen up to $\\sim60$ s. On the basis of the lag analysis (Norris \\& Barthelmy 2006) this burst is classified as a long duration event. The time integrated BAT spectrum is well described by a power--law model with photon index --1.31$\\pm 0.06$ (Palmer et al. 2006), suggesting \\epo$>150$ keV. Golenetskii et al. (2006) report the spectral analysis of {\\it Konus}/WIND data separately for the pulse and the tail. The pulse is well described by a power--law with a high--energy exponential cutoff, with spectral parameters $\\alpha=-1.22^{+0.12}_{-0.14}$ and \\epo $=777^{+549}_{-237}$. The soft tail can be modelled as a single power--law, with photon index $-1.93^{+0.32}_{-0.27}$ and a fluence which is nearly 1/3 of the fluence of the main pulse. In order to avoid to underestimate the total energetics, we also include the fluence of the soft tail in the computation of \\eiso. Nevertheless, this burst is on average $\\sim$2 orders of magnitude less energetic as compared to bursts characterized by a similar peak energy. Note that the lack of a spectral analysis integrated over the whole burst duration does not affect the estimate of \\liso, which is well computed by using the peak flux and the spectral shape of the main peak.  Also in this case, the peak luminosity is $\\sim$2 orders of magnitude less than the average \\liso\\ of events with \\ep\\ similar to GRB 061021. Finally, the redshift of this source is $z=0.346$ (Fynbo et al. 2009).  The presence of outliers to the Amati relation is somehow expected from the study of the observational plane \\epo$-Fluence$. Nava et al. (2008; 2011b) studied the distribution of BATSE and GBM bursts in this plane and derived respectively a 6\\% and 3\\% of outliers. We also stress that the definition of the outliers of the \\epeiso\\ correlation refers to the present knowledge of the scatter of this correlation which is assumed to be Gaussian. It might well be that the scatter is still not fully known and/or that it is not Gaussian. The latter possibility seems to be suggested also by the log--normal distribution of the jet opening angles (Ghirlanda et al. 2005) if the source of the scatter of the \\epeiso\\ correlation were, indeed, the jet opening angle.   There are 12 bursts belonging to the complete sample that cannot be used to directly test the correlations, due to the lack of measured redshift and/or \\epo. However, we can still check their consistency with the correlations (Fig. \\ref{fig2}). We find that they are all consistent with both correlations.  By taking advantage of our complete sample, for both correlations we investigate the possible evolution with the redshift of the best fit slopes (Fig. \\ref{fig3}). To this aim we define 4 bins of redshift, chosen in order to have 4 sub--samples characterized by a similar number of objects. We find no relation between the slope of each sub--sample and the redshift. All slopes are consistent with the one defined by the whole complete sample and this is true both for the \\epeiso\\ and the \\epliso\\ correlation.   We also investigate possible differences (in terms of distributions of the most fundamental quantities) between bursts belonging to the complete sample and bursts that do not satisfy the complete sample selection criteria. Our aim is to understand if the different instrumental biases present in the total sample affect in some way the distributions of the prompt emission properties. We find that for all the considered quantities (\\ep, \\eiso\\ and \\liso) the distributions derived from the complete and from the complementary samples are consistent within 2$\\sigma$.  Our results give support to the idea that spectral energy correlations arise because of a robust physical mechanism common to the vast majority of bursts. Outliers exist, but they are a few. Furthermore, the evidence that GRBs do evolve in redshift (Salvaterra \\& Chincarini 2007), while there is no sign of evolution of the slopes of the correlations, also suggests that the spectral peak energy is closely linked to the luminosity/energetics of the burst, another sign that selection effects are not important."
    },
    "1112/1112.3911_arXiv.txt": {
        "abstract": "We overview high-energy neutrinos from galactic sources, transparent to their gamma-ray emission. We focus on young supernova remnants and in particular on RX J1713.7-3946, discussing expectations and upper bounds. We also consider the possibility to detect neutrinos from other strong galactic gamma-ray sources as Vela Junior, the Cygnus Region and the recently discovered Fermi Bubbles. We quantify the impact of the recent hint for a large value of $\\theta_{13}$ on  high-energy neutrino oscillations. ",
        "introduction": "The successes of low energy neutrino astronomy and the discovery of neutrino oscillations added momentum to the search for high energy neutrinos from cosmic sources, initiated long ago with the theoretical proposals of Zheleznykh, Markov and Greisen \\cite{markov}; see Fig.~\\ref{fig0}. The first km$^3$-class detector, IceCUBE and the smaller ANTARES (following MACRO, BAIKAL, AMANDA) have not yet revealed the first signal, yet the excitement remains high. However, nowadays it becomes clear that the search for high energy neutrino sources is difficult. In this paper, we attempt to examine certain generic and specific expectations for such a search, in the hope to contribute to the scientific planning of the future experimental activities. We are interested in the possibility to proceeding further, by building a new telescope of km$^3$-class in the Northern Hemisphere--as Km3NET in the Mediterranean Sea or GVD in Lake Baikal.   The fact that the first explorations of the high-energy neutrino sky, that did not reveal any signal, contradicted  the over-optimistic expectations may lead one to doubt that we can reach reliable predictions. But, even if some degree of diffidence toward theory is healthy, we are convinced that the process of formulating expectations is an important step toward understanding. We would like to clarify our view by proposing some issues: {\\em Do we really need to have high expectations for high energy neutrinos?} Perhaps not, although one should keep in mind that the surprises often happen in astronomy. As recent example in this regard are the discovery of the Fermi Bubbles and the variability of the Crab Nebula. {\\em Are expectations useful?} Of course, yes; good predictions are precious for experiments, but also reasonable expectations eventually contradicted are not useless. {\\em Do we have some relevant and succesful precedent?} Yes, many: High-energy $\\gamma$-ray sources have been predicted in the fifties (Morrison). For solar neutrinos, we had predictions with errorbars since the sixties (Bahcall). For supernova neutrinos, quantitative expectations were available even before SN1987A (Nadyozhin).    \\begin{figure}[b] \\centerline{\\includegraphics[width=7.1cm,height=4.cm,angle=0]{energ2}} \\caption{\\em\\footnotesize  Energy ranges of various types of neutrino telescopes, in $\\log_{10}(E/eV)$; note that they span more than 10 decades in energy. In this paper, we are concerned with the third item, that after IceCUBE we identify with the technology `ice Cherenkov'. \\label{fig0}} \\end{figure}   We adopt a rather common astrophysical attitude: the observation of high-energy neutrinos is very demanding but perhaps possible and would amount to an unambiguous signal of cosmic ray collisions--hopefully, those in their source. We simplify our task by restricting our attention on the specific but important class of sources, those transparent to their gamma rays. We will show how to proceed towards precise expectations and eventually observations of neutrinos, limiting the use of theory inputs and just by the help of  $\\gamma$-ray observations. This aspect is important enough to be emphasized, as we do with Fig.~\\ref{fig1}. In the next section, we quantify the connection between gamma and neutrinos. % ",
        "conclusions": "The promising galactic neutrino sources are intimately tied to $\\gamma$'s above 10 TeV. The main points regarding the predictions of  the high-energy neutrino fluxes from galactic % sources  are simply two: \\begin{itemize} \\item There are only few bright $\\gamma$-ray sources, that correspond to detectable neutrino sources. \\item For these sources,  the expected neutrino signal is modest, even  assuming that the $\\gamma$-ray emission is fully contributed by hadronic processes. \\end{itemize} The studies of $\\gamma$-rays from RX J1713.7-3946 prove that it is possible to proceed toward reliable expectations for high-energy $\\nu$s. The future years will offer us good occasions of progress in this sense: \\begin{enumerate} \\item Multiwavelength observations and theory will help us to understand better the young SNR's and to quantify their hadronic and leptonic emissions. \\item Sub-degree pointing with very high energy $\\gamma$-rays could permit us to further test the paradigm of SNR+MC association for RX J1713.7-3946. \\end{enumerate} These considerations emphasize the relevance of certain lines of progress, and suggest to increase synergies and collaboration between neutrino and gamma ray search.      \\appendix"
    },
    "1112/1112.3120_arXiv.txt": {
        "abstract": "We show that the surface brightness profiles of elliptical galaxies can be parametrized using a linear superposition of two or three components, each of which is described by functions developed in Dhar \\& Williams as the 2D projections of a 3D Einasto density profile. For a sample of 23 ellipticals in and around the Virgo Cluster with total absolute {\\it V}-magnitude $-24$$<$$M_{VT}$$<$$-15$, our multi-component models span a dynamic range up to $10^6$ in surface brightness and up to $10^5$ in radius down to the resolution limit of the HST, have a median rms of 0.032 $\\mgasc$ consistent with the rms of 0.03 from random errors of the data, and are statistically justified at $>$3$\\sigma$. Our models indicate that i) the central component is more concentrated than the outer component; and ii) the central component of massive shallow-cusp ('core') galaxies is much more luminous, extended and concentrated than that of steep-cusp ('cuspy') galaxies, with their near exponential central profiles indicating disk-like systems, whose existence, must be verified spectroscopically.  Galaxy structure can thus be modelled extremely well with a central mass excess for all galaxies. This is not necessarily contrary to the notion of a mass deficit in 'core' galaxies, since mass ejection due to core-scouring by a supermassive black hole (SMBH) binary could have affected the shape of the central components. However, we show that the existence, amount, radial extent and sign of such deficits disagree substantially in the literature, both for a given galaxy and on an average over a sample. We discuss possible implications and suggest that SMBH binaries are unlikely to be the sole mechanism for producing the large 'cores' of massive galaxies.  Using results from the SAURON survey we deduce that under certain conditions of symmetry, inclination angles and degree of triaxiality, the intrinsic (3D) density of light can be well described with a multi-component Einasto model for both steep- as well as shallow- cusp galaxies. This indicates an universality in the functional form describing the 3D density distribution of light in galaxies and dark matter in $\\Lambda$CDM N-body haloes. Finally, Planetary Nebulae and strong lensing observations, and the Einasto index $n$ of $\\Lambda$CDM dark matter haloes indicate that our result -- the outer component of the surface brightness profiles of massive galaxies has 5 $\\lesssim$ $n$ $\\lesssim$ 8 -- could imply i) a common feature of collisionless systems; and ii) that galaxies with such $n$ for their outer component are dark matter dominated. ",
        "introduction": "\\cite{deVauc48} showed that a remarkably simple parametrization of the radial surface brightness profile of galaxies exists for a wide range of ellipticals. He proposed a two parameter fitting function of the form of equation \\eqref{sersic} with $m$$=$4 defining the shape of the distribution and $q$$=$7.67 ensuring that the {\\it effective} or {\\it half-light radius}, $R_{E}$, contains half the total projected light. \\begin{align}\\label{sersic} \\Sigma(R)= \\Sigma_{R_E}\\exp\\left\\{-q\\left[\\left(\\frac{R}{R_E}\\right)^{\\frac{1}{m}} -1\\right]\\right\\} \\end{align} where, $\\Sigma(R)$ is the 2D surface brightness at a plane of sky projected radial distance $R$, $\\Sigma_{R_E}$$=$$\\Sigma(R_E)$, $q$$=$$q(m)$ and $\\Sigma_S(0)$$=$$\\Sigma_{R_E}$$e^q$.  Observing that our Galaxy is made up of multiple subsystems \\cite{Einasto65} proposed a modification, equation \\eqref{einasto}, of the 2D de Vaucouleurs law to model the intrinsic (3D) baryonic mass density $\\rho(r)$ of each subsystem by allowing the shape parameter $n$ to be a free parameter, and $b$$=$$b(n)$. \\begin{align}\\label{einasto} \\rho(r) = \\rho_{s}\\exp\\left\\{-b\\left[\\left(\\frac{r}{r_s}\\right)^{\\frac{1}{n}} -1\\right]\\right\\} \\end{align} where, $\\rho_s$ is the density(3D) at a scale radius $r_s$ and $\\rho(0)$$=$$\\rho_s$$e^b$.  Around the same time \\cite{Sersic68} observed that $m$ in equation \\eqref{sersic}, characterized by $m=4$ in de Vaucouleurs law, was not the same for all galaxies and rendering it a free-parameter provided much better fits to the surface brightness profiles (hereafter, SB). Equation~\\eqref{sersic} is the current standard paradigm for describing the global structural properties of galaxies.  Over the past forty years, the pioneering works of Caon, Capaccioli, Einasto, Ferrarese, Graham, Kormendy, Lauer and their collaborators have shown that no single commonly used three parameter fitting function could model the SB over the entire dynamic radial range. They showed that the SB profiles of all galaxies reveal an inherently multi-component structure such that the outer regions can be described with a Sersic or a Sersic+exponential models while the central regions can be described with power-laws.  Additionally, for some nearby spirals and the giant elliptical M87 (NGC4486), Einasto and collaborators (Rummel, Haud and Tenjes) have shown that if spectroscopic and kinematic constraints are used in addition to the SB data, then the intrinsic 3D mass density including their central regions can be described with a multi-component Einasto model. \\subsection{Motivation for this work} While the power-law+Sersic models are widely accepted as providing an accurate description of the SB profiles, the fit residuals of these models are often larger than measurement errors. Further, these functions do not have the flexibility to model deviations from power-laws within the central regions. Since these fits are often used to draw inferences on galaxy structure, formation and evolution, it is important to address two crucial issues when modelling galaxy structure:  (i)~Models of the SB profiles must be consistent with measurement errors over the entire available radial range; and  (ii)~Model fitting functions must allow one to easily infer the intrinsic 3D luminosity density from the 2D SB profiles.  Our goal in this paper is to address the above issues with a new function derived in \\cite{DW10} (DW10) to model the projected surface mass density profile of Einasto-like systems. We show that a multi-component parametrization with this function (hereafter, the DW-function) provides excellent fits, consistent with measurement errors, to the SB of ellipticals over a large dynamic radial range down to the HST resolution limit.  This suggests that the 3D density distribution of light in galaxies can be described with a multi-component Einasto model.  Such an interpretation is similar to that of Einasto and collaborators, except that a) our work extends to a much larger sample of shallow and steep cusp ellipticals; and b) we deduce the intrinsic 3D Einasto model properties by fitting the 2D SB using a multi-component DW-function.  For the rest of this paper, we draw a distinction between components deduced from fitting and physically distinct kinematic systems or stellar populations. We shall refer to {\\it a component} as a single spherically symmetric closed form fitting function described by at most three parameters -- a scale length, $r_{\\alpha}$, a shape parameter, $\\alpha$, and some normalization, $\\Sigma_{0}$ or $\\Sigma_{\\alpha}$. Our multi-component fits consist of the minimum number of such functions (i.e. DW-functions) required to model the entire dynamic radial range down to the resolution limit of the HST. The minimum number depends on the quality of data, the available degrees of freedom, and the amplitude of systematic patterns in the residuals.  One should hence use caution in interpreting these {\\it fit components} as kinematically {\\it distinct systems or stellar populations} since the physical properties of the components, and even their presence will depend on the choice of parametrization. It is however possible that some of our fitted components do coincide with kinematically identified systems or stellar populations, or combinations thereof, which shall then facilitate interesting interpretations.  Also note that neither fit components nor physically distinct systems have to correspond to structure in the total gravitational potential. Only where $\\rho_{baryons}$$\\gg$$\\rho_{DM}$, or $\\rho_{baryons}$$\\propto$$\\rho_{DM}$ features observed in the SB profiles will trace the total mass density, and hence will provide information about the gravitational potential.  Further, since no galaxy has a truly flat density core, we refrain from using the terms 'core-' and 'cusp-' galaxies and instead refer to them as shallow-cusp and steep-cusp galaxies, respectively. \\subsection{Brief outline and summary of results} In section \\ref{MCstruc} we provide a brief history and basis for believing that galaxies have a multi-component structure. This is followed by an overview, in section \\ref{fittingfuncs}, of the most flexible fitting functions tried to date, and the motivation for this work. The data we use are described in section \\ref{data} and important features of our 2D parametrization and the fitting procedure are discussed in section \\ref{DWfunction}. Section \\ref{fitresults} discusses results of our 2D fits to a sample of 23 ellipticals in and around the Virgo cluster spanning absolute V-magnitudes in the range -24$<$$M_{VT}$$<$-15. Our models produce consistently low residuals, over a large dynamic radial range ($\\sim$$10^5$) and in congruence with the measurement errors.  Section \\ref{comparison} contains a comparison of fits with other parametrizations from the literature. It highlights the main result of this work: {\\it of all the functions tried to date to model the 2D structure of galaxies, the best fitting model is the multi-component DW-function}. This suggests that the intrinsic 3D density structure of galaxies is best represented with a multi-component Einasto model.  In section \\ref{compprops} we discuss the properties of the components deduced from the multi-component fits with the DW-function, and in section \\ref{compsys} we present three cases where the central DW-component of our multi-component fits coincide with spectroscopically identified systems.  We next explore the conditions under which our 2D models can be used to infer that the intrinsic 3D luminosity density distribution can be described with a multi-component Einasto model. In section \\ref{uniqueness} we therefore first discuss the key issues pertaining to non-uniqueness of deprojection, and then in section \\ref{lumdensity} we present our models of the 3D Einasto luminosity density profiles for $14$ galaxies whose deprojections are less likely to suffer from non-uniqueness. Since these $14$ galaxies span a wide luminosity range and belong to both the steep- and shallow-cusp classes, we conclude that the intrinsic structure of all galaxies could be described with multi-component Einasto models as well.  This is followed by section \\ref{massdeficit}, where we discuss the notion of 'mass deficit' in massive ellipticals. Following a brief review of the phenomena of mass ejection by binary SMBHs and results from N-body simulations, we show that {\\it estimates} of the amount, radial extent and sign of the 'observed' deficit in real galaxies, have large disagreements in the literature -- both for a given galaxy as well as on an average for a given sample.  In this paper we have shown that the structure of galaxies can be modelled extremely well with a central mass excess for all galaxies. We suggest that, while mass ejection due to core-scouring by binary SMBHs could have shaped the central and the intermediate DW-components, this phonemena by itself is unlikely to be the sole cause for the observed shallow-cusps; other formation processes are likely to have also contributed to the existence of this feature.  In section \\ref{BaryonsDM} we show that an Einasto shape parameter of 5 $\\lesssim$ $n$ $\\lesssim$ 8 is found in two different systems - (i) the outer baryonic DW-component of the massive ellipticals in our sample, and (ii) in fits with an Einasto profile to the collisionless $\\Lambda$CDM N-body haloes. This range of $n$ may thus be a common feature of collisionless systems. Finally, by comparing results from Planetary nebulae and strong lensing studies we show that it is possible that galaxies with 5 $\\lesssim$ $n$ $\\lesssim$ 8 for the outer component of their SB profiles are likely to be more dark matter dominated.  Section \\ref{conclusion} contains a general discussion and a summary of conclusions drawn from the work presented in this paper. ",
        "conclusions": "\\label{conclusion} Critical to our understanding of galaxies, their formation and evolution, is our ability to accurately quantify the galaxies' structural properties and their variation with mass. The Sersic profile is the single most commonly used function for describing galaxy light distribution. The works of Caon, Capaccioli, Einasto, Ferrarese, Graham, Kormendy, Lauer and their collaborators have revolutionized our understanding of galaxy structure. These authors have shown us that a single Sersic profile does not fit the surface brightness distribution of ellipticals consistent with measurement errors over a radial range larger than $2$-$3$ decades. To extend the radial range of the fit, especially in to the central regions, the core-Sersic ($6$ parameters) and the Nuker ($5$ parameters for the central $10$-$20$ arcsec) profiles have been introduced, with a further addition of a King model ($3$ more parameters) for the nuclear region.  But even with these flexible and multi-parameter models, the fit residuals often exceed the measurement errors in some radial ranges, in spite of the overall $rms$ of residuals being low due to the large regions over which the Sersic profile fits well. Double Sersic models (Hopkins et al. 2009 a,b) provide an improvement, but still the residuals remain larger than measurement errors.  In addition to fit residuals exceeding measurement errors, the Sersic profile has another drawback. The galaxy structure, dynamics and evolution exist in 3D, so it is more meaningful to describe galaxies using 3D functions. The Sersic profile is an intrinsically 2D distribution whose deprojections preassumes an infinite 3D extent. If, for example, a galaxy has a truncation radius in 3D, a Sersic profile will not be able to model that. The parametric forms for deprojecting a Sersic profile (and power-laws, for that matter) are not well behaved near the centre. Further, these analytic functions are either reasonable but not extremely accurate (PS97, LGM99), or are accurate but extremely complicated (BG11).  To overcome the above limitations, we propose that the 3D luminosity density may be described with a multi-component Einasto profile whose parameters can be estimated by modelling the surface brightness using a multi-component DW-function. This is similar to the observations of Einasto and collaborators with a small sample of spiral (and the giant elliptical M87), but extends the idea to a much larger and diverse sample of shallow and steep cusp ellipticals and also allows for a direct parametric description of the 2D surface brightness profiles.  \\medskip  We model the surface brightness (SB) profiles of $23$ ellipticals in and around the Virgo cluster with a multi-component DW-function and summarize our observations as follows:  \\medskip  1.~~~ Multi-component DW-function fits the SB profiles of ellipticals with residuals consistently comparable to measurement errors over large dynamic ranges $\\sim 10^5$ in radii, and $\\sim 10^6$ in SB, with a median sample $rms$ of $0.032~\\mgasc$.  Nine galaxies are well described with a $2$-component (central and outer) DW model, while for fourteen galaxies a third component was required, and confirmed through an F-test. The third component acts as an intermediate component between the central and outer components, except in NGC4621 and NGC4434 where it appears to be embedded within the outer component.  \\medskip  2.~~~ All steep-cusp and shallow-cusp galaxies reveal a central component that is in excess to an inward extrapolation of the outer component. Its shape parameter $n$ is usually less than that of the outer component, which implies, as expected, that the central component is more concentrated than the outer component. Exceptions are NGC4459, NGC4387 and NGC4467 (section \\ref{outcomp}).  \\medskip  3.~~~ The central component of all massive shallow-cusp galaxies have $n \\lesssim 1$, while those of steep-cusp galaxies generally have $n>2$, although there are some cases with $1\\lesssim n<2$ (section \\ref{outcomp}). This indicates that the central components of shallow-cusp galaxies  a)~ are more concentrated than that of the steep-cusp galaxies even though the later are more denser in their central regions; and  b)~ could signal the presence of disk-like systems; however, this must be verified spectroscopically.  \\medskip  4.~~~  The central component of the shallow-cusp galaxies is far more luminous (and massive) and spatially more extended (large $r_{-2}$ or $r_{3E}$) than that of the steep-cusp galaxies. Further, within a factor of two, all galaxies appear to host a similar fraction of total light in their central+intermediate components, with a weak indication that massive galaxies may be hosting a larger fraction. The last point is inconclusive due to the small number of massive ellipticals in our sample.   \\medskip  5.~~~  In most of the shallow-cusp galaxies the outer component makes a comparable contribution to the density in the central regions with respect to that of the central component, while in the steep-cusp galaxies the central component is dominant.  From the modelling point of view (section \\ref{outcomp}), the shallow-cusp feature usually seen in massive galaxies is due to a combination of a larger $r_{-2}$, a low $n \\lesssim 1$ implying a more concentrated component, and the non-negligible contribution to the density in the central regions from the $n\\gtrsim 5$ outer component.  \\medskip  6.~~~ Galaxy formation models indicate that in massive ellipticals, the formation, evolution and subsequent coalesence of binary SMBHs can remove nearly (2-3)$M_{BH}$ of stellar material from the central regions, in a typical merger (refer to iii, section \\ref{Nbodyaftercoal}), leading to mass deficits. Our observation that the central component in shallow-cusp galaxies are already massive, can be used to constrain galaxy formation models by accounting for the amount of mass that must have assembled in these galaxies prior to mass ejection by the binary SMBH's.  It is possible that the current shapes of the central as well as the intermediate DW-components, within the central kiloparsec or so, have been influenced by the evolution of the binary SMBHs; and consequently some deficit could well have occured. However, the amount and sign of such deficits depend on estimating a break radius $R_b$, and assuming that the functional form of the density profile beyond $R_b$ can be extrapolated all the way to the centre. This may not be meaningful and has led to large disagreements in the literature -- both for a given galaxy as well as for averages over a sample. We note that core-scouring by SMBHs is unlikely to be the sole mechanism for producing some of the largest 'cores' (shallow-cusps), and other processes in galaxy formation and evolution are likely to have played their role as well in forming this feature. In such a case extrapolating the density profile inwards from $R_b$, to {\\it estimate} the mass deficit, will yield misleading results.  \\medskip  7.~~~ For $14$ of the $23$ galaxies we could describe the intrinsic 3D luminosity density distribution fairly uniquely with a multi-component Einasto model (section~\\ref{lumdensity}). Since these galaxies span a wide range of luminosity $-24$$<$$M_{VT}$$<$$-15$ and come from both the steep- and shallow-cusp families, it is likely that other galaxies also have an intrinsic structure that can be modelled with a multi-component Einasto profile.  \\medskip  8.~~~ Pure dark matter haloes are known to be well fit with the Einasto function. We have shown here that the same fitting function that describes the intrinsic density distribution of $\\Lambda$CDM N-body haloes can also be used to model the intrinsic baryonic density of ellipticals, but for baryons a multi-component Einasto model is required. There thus appears to be an universality in the functional form of the density profile of baryons (stars) and dark matter.  \\medskip  9.~~~ The Einasto shape parameter $n$ of dark matter haloes in N-body simulations is very similar to the $n$ of the outer component of our massive ellipticals; both are 5 $\\lesssim$ $n$ $\\lesssim$ 8. Further, PNe and strong lensing observations indicate that massive ellipticals are more dark matter dominated than less massive ellipticals. This indicates that our result -- the outer component of the surface brightness profiles of massive galaxies has 5 $\\lesssim$ $n$ $\\lesssim$ 8 -- could imply i) a common feature of collisionless systems; and ii) that galaxies with such $n$ for their outer (major) component are dark matter dominated.  \\medskip  10. ~~~ In section \\ref{massdeficit}, we have shown that the shapes of central and global profiles of shallow- and steep- cusp galaxies differ markedly and could be a result of differing formation pathways. In section \\ref{outcomp} we have shown that the shallow-cusps can be typically described as systems with - i) central DW-components of {\\it low} $n$ ($\\lesssim 1$) and large scale radius $r_{-2}$ or $r_{3E}$, and ii) outer DW-component of {\\it large} $n$ ($\\gtrsim 5$) and central density comparable to that of the central DW-component. And in section \\ref{BaryonsDM} we have shown that large outer $n$ systems are likely to be dark matter dominated.  Hence, if dark matter has played a role in the formation and evolution of massive shallow-cusp galaxies, it will be instructive to explore its role, if any, in shaping the central regions and consequently in forming the shallow cusps.  Finally we note that the galaxies modelled in this paper are all in and around the Virgo Cluster. Galaxy structure, formation and evolution is known to depend on the environment, and even though we suspect that our conclusion about the multi-component Einasto structure should be more widely applicable than Virgo, only more data can confirm that."
    },
    "1112/1112.0123_arXiv.txt": {
        "abstract": "Adopting the SPECTRUM package which is a stellar spectral synthesis program, we have synthesized a comprehensive set of 2,890 Near-InfraRed (NIR) synthetic spectra with a resolution and wavelength sampling similar to the SDSS and the forthcoming LAMOST spectra. During the synthesis, we have applied the `New grids of ATLAS9 Model Atmosphere' to provide a grid of local thermodynamic equilibrium (LTE) model atmospheres for effective temperatures ($T_{eff}$) ranging from 3,500 to 7,500 K, for surface gravities (log $g$) from 0.5 to 5.0 dex, for metallicities ([Fe/H]) from -4.0 to 0.5 dex, for abundances solar ([$\\alpha$/Fe] = 0.0 dex) and non-solar ([$\\alpha$/Fe] = +0.4 dex). So this synthetic stellar library is composed of 1,350 solor scaled abundance (SSA) and 1,530 non-solar scaled abundance (NSSA) spectra, grounding on which we have defined a new set of NIR Ca{\\sc{ii}} triplet indices and an index CaT as the sum of the three. Then, these defined indices have been automatically measured on every spectrum of the synthetic stellar library and calibrated with the indices computed on the observational spectra from the INDO-U.S. stellar library. In order to check the effect of $\\alpha$-element  enhancement on the so-defined Ca{\\sc{ii}} indices, we have compared indices measured on the SSA spectra with those on the NSSA ones at the same terns of stellar parameters ($T_{eff}$, log $g$, [Fe/H]), and luckily, little influences of $\\alpha$-element enhancement has been found. Furthermore, comparisons of our synthetic indices with the observational ones from measurements on the INDO-U.S. stellar library, the SDSS-DR7, and SDSS-DR8 spectroscopic survey have been presented respectively for dwarfs and giants in specific. For dwarfs, our synthetic indices could well reproduce the behaviours of the observational indices versus stellar parameters, which verifies the validity of our index definitions for dwarfs. While for giants, the consistency between our synthetic indices and the observational ones does not appear that good. However, a new synthetic library of NIR Ca{\\sc{ii}} indices has been founded for deeper studies on the NIR waveband of stellar spectra, and this library is particularly appropriate for the SDSS and the forthcoming LAMOST stellar spectra. Later on, we have regressed the strength of the CaT index as a function of stellar parameters for both dwarfs and giants after a series of experimental investigations into relations of the indices with stellar parameters. For dwarfs, log $g$ has little effect on the indices while [Fe/H] and $T_{eff}$ play a role together and the leading factor is probably [Fe/H] which changes the stength of the indices by a positive trend. For giants, log $g$ starts to influence on the strength of the indices by a negative trend for the metal-poor and even impact deeply for the metal-rich, and besides, [Fe/H] and $T_{eff}$ still matter. Ultimately, a supplemental experimentation has been carried out to show that spectral noises do have effects on our set of NIR Ca{\\sc{ii}} indices. However, the infection is weak enough to be ignored if the signal-to-noise ratio (SNR) falls below 20. ",
        "introduction": "Due to more and more spectroscopic surveys, a huger number of spectra are turning available for studies on the physical properties (atmospheric parameters of stars; abundances, composition and ages of composite stellar populations of galaxies) by extracting information hidden in the absorption lines. However, problems of the limited sampling of spectral features caused by the detector pixel and the signal-to-noise ratio (SNR) seriously hamper the detailed fitting of the weak spectral features. So an alternative option is expected to extract physical information from the spectrum with no need of fitting faint absorption lines. Therefore, the concept of ``spectral indices\" has been introduced. This good idea for spectral analysis is to define a set of ``spectral indices\" each of which is an estimate of the equivalent width of a spectral feature measured by taking the ratio of the flux of feature in the central band to that of pseudo-continuum determined by two adjacent continuum bands. So far, various definitions for absorption features have existed as different spectral index systems of which the most successful is the Lick/IDS index system \\citep{b5, b10, b11, b12, b13, b14} which defines a set of 25 spectral indices in the optical waveband. Although it is widely used, Lick/IDS index system lacks for definitions of absorption features in the Near-InfraRed (NIR) waveband since the Lick spectra only cover from 4000 $\\AA$ to 6400 $\\AA$ in wavelength.  In the NIR region ($\\lambda\\lambda$ 7900 $\\sim$ 9100 $\\AA$) of stellar spectra of types A $\\sim$ M, the strongest atomic feature is definitely Ca{\\sc{ii}} triplet ($\\lambda\\lambda$ 8498, 8542, 8662 $\\AA$) which offers several advantages from the observational point of view for being the strongest metal line in the NIR spectrum, exceeding Na{\\sc{i}} ($\\lambda$8190) in strength by at least a factor of 5, and lying in a region relatively free of atmospheric absorptions. The Ca{\\sc{ii}} triplet has been used as a luminosity indicator \\citep{b4}, and its strength is simultaneously a good metallicity indicator in stellar systems where the NIR light is dominated by stars of a restricted range of surface gravities \\citep{b1}. Beside, it has provided a potential tool to study the properties of stellar populations since the old stellar populations emit most of their light in this wavelength region \\citep{b1}.  Progresses have been previously made at studies on the NIR Ca{\\sc{ii}} triplet. By studying a sample of 62 stars of types F0 $\\sim$ M5, \\citet{b4} has derived that the equivalent width (EW) of the NIR Ca{\\sc{ii}} triplet only correlates strongly with log $g$ in stars. This single-valued correlation has later been confirmed by \\citet{b15}. However, by a set of definitions for the Ca{\\sc{ii}} indices on a sample of 106 late-type stars, \\citet{b1} has demonstrated that in the high-metallicity range, the strength of the indices merely depends on log $g$ while in the low-metallicity range, [Fe/H] turns to play a leading role. Afterwards, by refined definitions for the Ca{\\sc{ii}} indices on a larger number of stars with a higher resolution, \\citet{b16} has clarified that $T_{eff}$ has an important impact factor for the low $T_{eff}$ stars and log $g$ plays a more important role for giants than dwarfs. \\citet{b2} has quantitatively estimated the strength of the Ca{\\sc{ii}} indices by introducing a set of ``generic indices\".  Although several works have been made on definitions for the NIR Ca{\\sc{ii}} triplet indices\\citep{b4,b1,b2,b16}, they are all based on samples of less stars with higher resolution and denser wavelength sampling than the Sloan Digital Sky Survey Data Release7 (SDSS DR7; \\citet{b18,b19}) and the forthcoming Large Sky Area Multi-Oject Fiber Spectroscopic Telescope (LAMOST; \\citet{b20}) spectra in the 8300$\\AA$ $\\sim$ 8900$\\AA$ region. Besides, those stars also have inaccurate estimations for stellar parameters. In order to cope with these disadvantages, we are expecting to establish a new index system of the Ca{\\sc{ii}} triplet which would be the most appropriate and convenient for use on the SDSS and the forthcoming LAMOST NIR spectra. We will perform this work on on a larger number of synthetic stellar spectra with approximately the same spectral resolution and wavelength sampling as the SDSS-DR7 and the forthcoming LAMOST spectra.  The aim of this work is to constitute a new synthetic library of a set of the NIR Ca{\\sc{ii}} indices for deeper studies on the NIR waveband of spectra, especially of the SDSS and the forthcoming LAMOST stellar spectra. This library could also be a supplemental part of the NIR indices to the work of \\citet{b3} which has established a synthetic library of several optical indices for SDSS and the forthcoming LAMOST spectra. In Section 2, we generate a large number of synthetic stellar spectra with various stellar parameters, on which a new set of the NIR Ca{\\sc{ii}} indices has been defined and measured. Then, we also investigate the dependence of the Ca{\\sc{ii}} indices on $\\alpha$-element enhancement. In Section 3, we calibrate our synthetic indices with the observational indices measured on the INDO-U.S. stars \\citep{b6}. Then, we compare the calibrated synthetic Ca{\\sc{ii}} indices with the observational INDO-U.S indices. Besides, tests of the calibrated synthetic indices on SDSS-DR7 and SDSS-DR8 \\citep{b21} stars have been carried out. In Section 4, we analyse the relations of the calibrated synthetic Ca{\\sc{ii}} indices with stellar parameters and approximately estimate the stength of the CaT index by stellar parameters. We compare the Ca{\\sc{ii}} triplet indices with the CaHK index in Section 5. Finally, we summarize this paper in Section 6. As a supplement, we have performed a series of experiments in Appendix A to check into the effect of SNR on Ca{\\sc{ii}} indices, and found that such an effect could be ignored until the SNR falls below 20. ",
        "conclusions": ""
    },
    "1112/1112.2126_arXiv.txt": {
        "abstract": "{We present the ``endpoint'' formalism for the calculation of electromagnetic radiation and illustrate its applications in astroparticle physics. We use the formalism to explain why the coherent radiation from the Askaryan effect is not in general Cherenkov radiation, as the emission directly results from the time-variation of the net charge in the particle shower. Secondly, we illustrate how the formalism has been applied in the air shower radio emission code REAS3 to unify the microscopic and macroscopic views of radio emission from extensive air showers. Indeed, the formalism is completely universal and particularly well-suited for implementation in Monte Carlo codes in the time- and frequency-domains. It easily reproduces well-known ``classical mechanisms\" such as synchrotron radiation, Vavilov-Cherenkov radiation and transition radiation in the adequate limits, but has the advantage that it continues to work in realistic, complex situations, where the ``classical mechanisms'' tend to no longer apply and adhering to them can result in misleading interpretations.} ",
        "introduction": "Processes leading to the emission of electromagnetic radiation can be very complex. To get a grasp of these processes, physicists often revert to ``classical named radiation mechanisms'' such as synchrotron radiation, transition radiation or Vavilov-Cherenkov radiation and try to apply these concepts to the problem at hand. However, it turns out that this strategy can often yield misleading results. Usually, the classical mechanisms only apply to idealized problems involving, e.g., infinite particle tracks, and their application to complex realistic situations can cloud the view for the complete picture.  Here, we propose a formalism suitable for the calculation of electromagnetic radiation from any kind of particle acceleration, which lends itself very well to the implementation in Monte Carlo codes, and necessitates no simplifying approximations to the underlying processes. In this ``endpoint formalism'', presented in detail in Ref.\\ \\cite{EndpointTheory}, the trajectory of individual particles is described as a series of points at which particles are accelerated instantaneously, leading to discrete radiation contributions which can then be easily superposed.  The endpoint formalism can be applied in both the frequency and time domains, and correctly reproduces the aforementioned classical named processes. Furthermore, we demonstrate how the application of the endpoint formalism has provided essential insights for radio emission processes in astroparticle physics. ",
        "conclusions": "We have developed an endpoint formalism for the calculation of electromagnetic radiation in arbitrarily complex radiation processes. The formalism is universally applicable and can be implemented easily in calculations both in the time and frequency domains.  Classical named radiation processes such as synchrotron radiation or Vavilov-Cherenkov radiation can be successfully reproduced using the endpoint formalism. However, unlike such descriptions, the endpoint formalism does not make any simplifying assumptions which often break down in realistic situations.  We have successfully used the endpoint formalism to address problems in astroparticle physics. In particular, we were able to calculate the radio emission from extensive air showers using the endpoint formalism, and thereby reconcile discrepancies existing between earlier modelling approaches. Also, the endpoint formalism illustrates that Askaryan radio emission in the lunar regolith does not vanish even if the particle cascade develops close to the surface, as the net growth and decline of the shower will produce radio emission even for a refractive index approaching unity."
    },
    "1112/1112.0315_arXiv.txt": {
        "abstract": "We compare the results of various cosmological gas-dynamical codes used to simulate the formation of a galaxy in the $\\Lambda$CDM structure formation paradigm. The various runs (thirteen in total) differ in their numerical hydrodynamical treatment (SPH, moving-mesh and AMR) but share the {\\it same} initial conditions and adopt in each case their latest published model of gas cooling, star formation and feedback. Despite the common halo assembly history, we find large code-to-code variations in the stellar mass, size, morphology and gas content of the galaxy at $z=0$, due mainly to the different implementations of star formation and feedback. Compared with observation, most codes tend to produce an overly massive galaxy, smaller and less gas-rich than typical spirals, with a massive bulge and a declining rotation curve. A stellar disk is discernible in most simulations, although its prominence varies widely from code to code. There is a well-defined trend between the effects of feedback and the severity of the disagreement with observed spirals. In general, models that are more effective at limiting the baryonic mass of the galaxy come closer to matching observed galaxy scaling laws, but often to the detriment of the disk component.  Although numerical convergence is not particularly good for any of the codes, our conclusions hold at two different numerical resolutions.  Some differences can also be traced to the different numerical techniques; for example, more gas seems able to cool and become available for star formation in grid-based codes than in SPH. However, this effect is small compared to the variations induced by different feedback prescriptions. We conclude that state-of-the-art simulations cannot yet uniquely predict the properties of the baryonic component of a galaxy, even when the assembly history of its host halo is fully specified. Developing feedback algorithms that can effectively regulate the mass of a galaxy without hindering the formation of high-angular momentum stellar disks remains a challenge. ",
        "introduction": "\\label{SecIntro}  Numerical simulations play a central role in studies of cosmic structure formation.  Collisionless N-body simulations have now become the main theoretical tool to predict the non-linear evolution of dark-matter dominated structures once initial conditions are specified. Their high accuracy and huge dynamic range have allowed a detailed comparison of their outcome with observations of the large-scale structure of the universe.  The impressive agreement between these observations and the predictions of the $\\Lambda$ Cold Dark Matter ($\\Lambda$CDM) model has helped to establish it as the current paradigm of structure formation \\citep{Springel2006}.  Simulating the evolution of the visible Universe is much more complex, because it requires understanding the many astrophysical processes which drive the evolution of the baryonic component under the gravitational influence of the dark matter. Numerical hydrodynamics in cosmological simulations has traditionally used either the Lagrangian Smoothed Particle Hydrodynamics technique (SPH; \\citealt{Lucy77, Gingold77, Monaghan92, Katz1996,Springel2010_SPH}) or Eulerian grid-based solvers sometimes aided by Adaptive Mesh Refinement (AMR) techniques (\\citealt{Cen1992,Bryan95,Kravtsov1999,Fryxell2000,Teyssier2002,Quilis2004}).  Both approaches have advantages and disadvantages. It is widely appreciated that SPH is not able to capture shocks with high accuracy and that in certain situations fluid instabilities can be suppressed, at least for standard implementations of SPH \\citep{Agertz07,Creasey2011}.  On the other hand, mesh-based codes are not Galilean invariant and may in some cases generate entropy spuriously through artificial mixing \\citep{Wadsley08}. As a result, even for some simple  non-radiative problems, Lagrangian and Eulerian codes do not converge to the same solution (e.g., \\citealt{Okamoto03,Agertz07, Tasker08, Mitchell2009}). Novel techniques, such as the Lagrangian, moving-mesh {\\sc arepo} code introduced by \\citet{Springel2010}, hold the promise of improving this state of affairs, but their application is still in its infancy.   An even more uncertain ingredient of galaxy formation simulations are the descriptions of poorly understood physical processes such as star formation and feedback. The huge dynamic range between the super-Megaparsec scales needed to follow the hierarchical growth of a galaxy and the sub-parsec scales that govern the transformation of its gas into stars implies that direct simulation of all relevant physical processes is still out of reach of even the most powerful computers and best available algorithms. Star formation and feedback are therefore introduced in cosmological simulations as ``sub-grid'' parameterized prescriptions of limited physical content and lacking numerical rigor.   These difficulties have hampered the progress of simulations of galaxy formation within the $\\Lambda$CDM paradigm, but a few results of general applicability have nevertheless emerged.  For example, the formation of {\\it realistic disk galaxies} in dark matter halos formed hierarchically, as expected in $\\Lambda$CDM, was recognized as a challenge even in early simulations (see, e.g., \\citealt{Navarro1995,NavarroSteinmetz1997}).  Simulated galaxies suffered from ``overcooling'' and from a dearth of angular momentum due to the transfer of angular momentum from the baryonic component to the dark matter during the many merger events that characterize hierarchical assembly (\\citealt{Navarro91,SL99,Navarro00}). Feedback, as a general heating mechanism that can prevent overcooling and regulate the assembly of a galaxy whilst avoiding catastrophic angular momentum losses, emerged as a crucial ingredient of any successful galaxy formation simulation \\citep{W98,NavarroSteinmetz1997,SL99,CS2008,Zavala08}.  Considerable progress has been made since this time: recent simulations have shown that the angular momentum problem can indeed be alleviated when feedback from evolving stars, in particular supernovae (SNe), is included (e.g., \\citealt{Okamoto05,Governato2007,CS2008,Ceverino2009,CS2009,Joung2009,Colin2010,Sales2010,Stinson2010}). Some authors have also investigated alternative feedback mechanisms, such as energy liberated during the formation of supermassive black holes as well as that carried by cosmic rays (e.g., \\citealt{DiMatteo2003, SpringelDiMatteo2005,Jubelgas2008,Booth2009,Fabjan2010, Wadepuhl11}).  Despite progress, difficulties remain. For example, simulations tend to allow far too many baryons to accrete into galaxies to be consistent with the observed stellar mass function of galaxies (see, e.g., \\citealt{Guo2010}). In addition, although stellar disks do form, they are often too concentrated, with steeply-declining rotation curves at odds with observation (e.g., \\citealt{Abadi2003a,Abadi2003b,Stinson2010}).  The difficulties in simulating disk galaxies highlighted above are compounded by the limited guidance afforded by analytic and semi-analytic models of galaxy formation. In such modeling, the properties of disk galaxies are typically envisioned to reflect those of their surrounding halos \\citep{FE80,KWG93,Dalcanton97,Mo98,Bower06,deLucia07}: for example, halos with high net spin and quiet recent merger histories are commonly assumed to be likely sites of disk galaxy formation. However, there are growing indications that these assumptions might be too simplistic. \\citet{CS2009} and \\citet{Stinson2010}, for example, find no clear relation between the presence of disks and the spin parameter of the halo: disks form in halos with low and high spin parameters. Moreover, \\citet{CS2009} report that disk formation is not assured by a quiescent assembly history, since they can be destroyed not only by major or intermediate-mass mergers, but also by secular processes such as a misalignment between the gaseous and stellar disks.  These difficulties have led to little consensus on what determines the morphology of a galaxy; what the main feedback mechanisms are; and what role they play on different mass scales and at different times. Indeed, there is even debate about whether the difficulties in reproducing realistic disks are predominantly a consequence of insufficient numerical resolution (e.g., \\citealt{Governato2004}), inappropriate modeling of the relevant physics (e.g., \\citealt{Mayer2008,Piontek2011}), or a failure of the cosmological model (e.g., \\citealt{Sommer-Larsen2001}).  Progress in this unsettled field requires at least a careful evaluation of the different numerical techniques, resolution, and choice of sub-grid physics adopted by various authors. Most groups do carry out and publish resolution tests and convergence studies of their own numerical setup. However, because of the complexity of the problem, the lack of telling test cases with known solution, and the absence of clear theoretical predictions for individual systems, each group chooses to tune the relevant numerical parameters according to different priorities and/or prejudice, and there has so far been little effort invested in comparing the results of different techniques and codes. Would they give similar results if they followed the formation of a galaxy in the {\\it same} dark matter halo?  The main goal of the Aquila Project is to address this question by comparing the predictions of different codes using {\\it common initial conditions} and a {\\it homogeneous set of analysis tools}. Rather than focusing on whether individual codes perform better or worse than others, we contrast their predictions for the stellar mass, angular momentum content, star formation rates and galaxy size, with observation.   This paper is organized as follows. \\S~\\ref{sec:simulations} describes the initial conditions and the simulation setup. \\S~\\ref{sec:results} compares the galaxy morphology, star formation history, size, angular momentum, and gas content of the $13$ simulated galaxies, together with a brief discussion of the effects of numerical resolution on the results.  Finally, \\S~\\ref{SecConc} summarizes our main findings and lists our main conclusions.     \\begin{figure*} \\includegraphics[width=10cm]{fig1a_LR.eps}  \\vspace{-0.1cm}  \\includegraphics[width=10cm]{fig1b_LR.eps}  \\vspace{-0.1cm}  \\includegraphics[width=10cm]{fig1c_LR.eps} \\caption{Maps of the dark matter distribution in the region surrounding the Aquila halo at $z=0$.  The dark matter is projected in cubic volumes of side length as indicated in each panel.  Pixel brightness corresponds to the dark matter density using a logarithmic scale. } \\vspace{1cm} \\label{DMoverview} \\end{figure*}    \\begin{table*} \\caption{Summary of code characteristics and sub-grid physics} \\label{table_codes} \\begin{tabular}{lcccccc} \\hline\\hline Code & Reference & Type & \\multicolumn{2}{c}{UV background} &Cooling & Feedback \\\\ & & & ($z_{\\rm UV}$) &  (spectrum) & & \\\\ \\hline {\\small G3} ({\\sc gadget3})& [1] & SPH & 6 & [10] & primordial [13]& SN (thermal) \\\\ \\\\  {\\small G3-BH} & [1] & SPH & 6 & [10] & primordial [13]& SN (thermal), BH \\\\ \\\\  {\\small G3-CR} & [1] & SPH & 6 & [10] & primordial [13]& SN (thermal), BH, CR \\\\ \\\\  {\\small G3-CS} & [2] & SPH & 6 & [10] & metal-dependent [14] & SN (thermal) \\\\ \\\\  {\\small G3-TO} & [3] & SPH & 9 & [11] & element-by-element [15] & SN (thermal+kinetic)\\\\ \\\\  {\\small G3-GIMIC} & [4] & SPH & 9 & [11] & element-by-element [15] & SN (kinetic)\\\\ \\\\  {\\small G3-MM} & [5] & SPH & 6 & [10] & primordial [13] & SN (thermal)\\\\ \\\\  {\\small G3-CK} & [6] & SPH & 6 & [10] & metal-dependent [14] & SN (thermal)\\\\ \\\\  {\\small GAS} ({\\sc gasoline})& [7] & SPH & 10 & [12] & metal-dependent [16] & SN (thermal)\\\\ \\\\  {\\small R} ({\\sc ramses}) & [8] & AMR & 12 & [10] & metal-dependent [14] & SN (thermal)\\\\ \\\\  {\\small R-LSFE} & [8] & AMR & 12 & [10] & metal-dependent [14] & SN (thermal)\\\\ \\\\  {\\small R-AGN} & [8] & AMR & 12 & [10] & metal-dependent [14] & SN (thermal), BH\\\\ \\\\  {\\small AREPO} & [9] & Moving Mesh & 6 & [10] & primordial [13] & SN (thermal)\\\\ \\hline \\end{tabular}  {\\sc note:} [1] \\citet{Springel08}; [2] \\citet{CS2005}; \\citet{CS2006}; [3] \\citet{Okamoto2010}; [4] \\citet{Crain2009}; [5] \\citet{M10}; [6] \\citet{kob07}; [7] \\citet{Stinson2006}; [8] \\citet{Teyssier2002}; \\citet{Rasera2006}; \\citet{Dubois2008}; [9] \\citet{Springel2010}; [10] \\citet{HaardtMadau1996}; [11] \\citet{HaardtMadau2001}; [12] Haardt \\& Madau (2005, private communication); [13] \\citet{Katz1996}; [14] \\citet{SutherlandDopita1993}; [15] \\citet{Wiersma2009a}; [16] \\citet{Shen2010}. \\end{table*} ",
        "conclusions": "\\label{SecConc}  The Aquila Project compares the results of simulations of galaxy formation within a Milky Way-sized $\\Lambda$CDM halo. We use $9$ gas-dynamical codes developed and run independently by different groups adopting in each case their preferred implementation of radiative cooling, star formation, and feedback. The codes include SPH-based, grid-based, and moving-mesh techniques; all include supernova feedback in thermal and/or kinetic form and primordial or metal-dependent cooling functions. Two of the codes ({\\sc gadget} and {\\sc ramses}) were run three times with three different sub-grid physics for a total of $13$ different simulations. In addition, each code was run at two different resolution levels in order to investigate numerical convergence. All runs share the same initial conditions and are analyzed with a set of consistent analysis tools. Our main conclusions may be summarized as follows.  {\\tt Galaxy Morphology.} At $z=0$, simulated galaxies exhibit complex morphologies, with spheroids, disks, and bars of varying importance. Morphology shows no obvious dependence on the hydrodynamical method adopted or on numerical resolution, and seems to be mostly related to how and when gas is accreted and transformed into stars. The best indicator of the presence of a disk seems to be the median age of the stars; the later stars form the more prevalent the disk component is.  This suggests that, to be successful at forming disks, codes must be able to preempt the early transformation of gas into stars while at the same time promoting the accretion and retention of late-accreting, high-angular momentum gas.  {\\tt Stellar Mass.}  Despite the common halo-assembly history, we find large code-to-code variations in the final mass of simulated galaxies. The stellar mass ranges from $4\\times 10^{10} \\, M_\\odot$ to $\\sim 3\\times 10^{11} \\, M_\\odot$, depending largely on the adopted strength of the feedback. There is also an indication that the numerical method may play a role: {\\sc arepo} is able to form almost twice as many stars as {\\small G3} although they both adopt the same sub-grid physics. Most simulated galaxies are too massive compared with theoretical expectations based on abundance-matching considerations.  The median stellar age also correlates with galaxy mass, indicating that models that favor late star formation (as needed to form a disk) do so at the expense of allowing too many stars to form overall.  {\\tt Rotation Curves.} All simulated galaxies have rotation speeds in excess of what is expected from the Tully-Fisher relation of late-type spirals. The disagreement worsens as the stellar mass of the simulated galaxy increases, both for galaxies with and without a well-defined stellar disk. At the high mass end, simulated galaxies have unrealistically sharply-peaked, strongly declining rotation curves. Although reasonably flat rotation curves are obtained at low $M_{\\rm stellar}$, the Tully-Fisher zero-point offset persists for these systems.  {\\tt Galaxy Size.} The difficulties matching the Tully-Fisher relation are due to the fact that most simulated galaxies have stellar half-mass radii much smaller than expected from observation given their stellar mass. This is especially true for galaxies where inefficient feedback allows the formation of an overly massive galaxy; these systems are, quite unrealistically, smaller even than dense early-type galaxies. Galaxies where feedback is able to limit the stellar mass to more acceptable levels are also more concentrated than typical spirals, highlighting the difficulty that all codes face to prevent too many low-angular momentum baryons from assembling into the galaxy.  {\\tt Star Formation History.} The excess of low-angular momentum baryons alluded to above may be traced to the inability of feedback schemes to prevent large bursts of star formation at early times. This is clearly seen in the stellar age distribution, which shows that star formation peaks at $z\\sim 4$ and declines steadily afterwards. The same trend holds for essentially all runs, albeit modulated by the particular implementation of feedback adopted by each individual code. Indeed, essentially all models allow more than half of the stars to form in the first $\\sim 3$ Gyrs of evolution. The relative insignificance of star formation in recent times compared to the intensity of the early burst seems to be at the root of many of the difficulties that simulated galaxies have in matching observation.  {\\tt Gas Component.} The properties of the gaseous component of the galaxy at $z=0$ show even wider code-to-code variations than the stars, since it is continuously resupplied by accretion, depleted by star formation, and dislodged by feedback-driven winds. This results in large short-lived fluctuations that lead to poor numerical convergence and large code-to-code scatter.  Most simulated galaxies have gaseous disks with  sizes that compare favorably with observation, although their gas fractions are systematically lower than those of star-forming spirals of comparable mass.  {\\tt Numerical Convergence.} We have investigated numerical convergence by comparing the results at two different numerical resolutions, spanning a range of $\\sim 2$ and $\\sim 8$ in spatial and mass resolution, respectively.  Although numerical convergence is not particularly good for any of the codes, reasonably-good convergence is found for the properties of the stellar component, such as total mass and median age. Less well-converged are the internal properties of the galaxy, such as the half-mass radius, or the fraction of stars in a rotationally-supported disk. For the same reasons cited in the previous item, the properties of the gas are the ones that are most poorly reproduced at the two different resolutions. There is also indication that feedback schemes that are more effective at limiting the stellar mass of the galaxy (such as through the inclusion of AGN-related feedback in addition to supernovae) converge less well than other implementations. Overall, the variations introduced by resolution are small compared to code-to-code variations, which leads us to conclude that none of the trends highlighted above are driven primarily by resolution.  {\\tt SPH vs AMR vs Moving-Mesh.} Our results suggest that numerical hydrodynamics techniques have some influence on the outcome of a simulation. This is most clearly demonstrated by comparing the results of {\\small G3} and {\\sc arepo} which, despite adopting the same sub-grid physics modules, yield galaxies that differ by almost a factor of $2$ in stellar mass. The AMR technique also yields large stellar masses (when similarly inefficient feedback is adopted, as in run {\\small R}), lending support to the view that standard SPH-based codes may underestimate the total amount of gas that can cool and become available for star formation at least in the weak feedback regime. On the other hand, the galaxies formed by {\\small R}, {\\small G3}, and {\\sc arepo} are unrealistically massive and concentrated, so large modifications to the feedback implementation of these codes are needed in order to bring them into agreement with observation. These changes may overwhelm the method-induced differences; for example, {\\small R} makes a prominent disk while {\\small R-AGN} has $5$ times fewer stars and no disk.  It is therefore unclear at this point whether the shortcomings of SPH are actually significant compared with the uncertainties involved in designing a star formation/feedback scheme that can yield realistic galaxies, although it is obviously desirable to avoid numerical inaccuracies as far as possible. Further investigation of this question is needed to clarify this issue.  Aside from these considerations, perhaps the main result of the Aquila Project is that, despite the large spread in properties spanned by the simulated galaxies, none of them has properties fully consistent with theoretical expectations or observational constraints in terms of mass, size, gas content, and morphology. Despite this apparent failure, we believe that the Aquila Project nevertheless yields interesting clues to guide how codes might be modified to yield realistic galaxies. For example, the need (i) to control effectively the overcooling of baryons; (ii) to curtail the early burst in star forming activity; and (iii) to promote the accretion and retention of the late-accreting, high-angular momentum baryons needed to form disks similar to those of normal spirals are of general applicability to all codes.  There seems to be little predictive power at this point in state-of-the-art simulations of galaxy formation; these seem best suited to the identification of the role and importance of various mechanisms rather than to the detailed modeling of individual systems. It may be argued that the strength of this conclusion depends on whether the parent halo of the Aquila runs (Aq-C) is truly destined to harbor a disk galaxy and that there is no hard proof for this. Further, the possibility that Aq-C might be an unrepresentative outlier should also be considered, as suggested by the {\\sc l-galaxies} semi-analytic model (see, e.g., Fig.~\\ref{fig:mstar_mhalo}).  These objections may be addressed when simulations are able to follow a statistically-significant number of galaxies in a cosmologically-representative volume. We might not know what kind of galaxy inhabits an individual halo, but we do know what the {\\it population} of galaxies looks like. Evolving a region large enough to contain at least a few dozen Milky Way-sized galaxies at the resolution achieved here seems like a natural next step, and one that should be achievable in the not too-distant future.  Finally, the complexity of the problem suggests that the best approach to improving galaxy formation simulations may be one where multiple numerical alternatives are developed and explored simultaneously and independently, provided that they are periodically contrasted in controlled experiments such as the one presented here. Given the intricacy of the task and in the absence of a clear front-runner, the diversity of numerical techniques presently available might actually turn out to be a strength rather than a shortcoming."
    },
    "1112/1112.0065_arXiv.txt": {
        "abstract": "Exploration of the time-domain radio sky has huge potential for advancing our knowledge of the dynamic universe. Past surveys have discovered large numbers of pulsars, rotating radio transients and other transient radio phenomena; however, they have typically relied upon off-line processing to cope with the high data and processing rate. This paradigm rules out the possibility of obtaining high-resolution base-band dumps of significant events or of performing immediate follow-up observations, limiting analysis power to what can be gleaned from detection data alone. To overcome this limitation, real-time processing and detection of transient radio events is required. By exploiting the significant computing power of modern graphics processing units (GPUs), we are developing a transient-detection pipeline that runs in real-time on data from the Parkes radio telescope. In this paper we discuss the algorithms used in our pipeline, the details of their implementation on the GPU and the challenges posed by the presence of radio frequency interference. ",
        "introduction": "The High Time Resolution Universe (HTRU) survey currently underway at Parkes Observatory using the 64m dish is an all-sky search for transient and periodic point sources. Its primary goals are to discover new pulsars and fast radio transients and it has so far discovered more than 65 previously unknown pulsars since it began in 2008 \\citep{KeithEtal2011}. The observing backend has 400 MHz of bandwidth over 1024 channels centered at 1381.8 MHz, and data products are produced with a time resolution of 64$\\mu s$. The receiver is a multibeam design and contains 13 separate feed horns pointing at nearby locations on the sky \\citep{StaveleySmith1996}.  The current pulsar and transient search process begins by taking filterbank data (i.e., time series for each frequency channel) from the backend -- stored using 2 bits per sample to give a data rate of around 4 MB/s per beam -- and sending it via a fibre optic link from the telescope down to the Swinburne supercomputer in Melbourne ($\\sim$ 700 km away). There it is written to disk for immediate processing (when possible) and also to tapes for longer term storage and future (re-)processing. The processing itself is then performed on a distributed cluster of multi-core central processing units (CPUs). Once the data are in memory, the current transient search pipeline takes more than 30 minutes to process a 10 minute observation.  While limitations in CPU computing power have necessitated offline processing of the survey data to date, ideally one would like to execute the pipeline in \\textit{real-time} with the observations. If this were achieved, several significant advances become possible. Among them is instant feedback and capacity for follow-up observations of significant events. This would eliminate the long delays incurred by the current off-site processing and undoubtedly increase the survey's discovery power. Another promising possibility is triggered baseband data dumps. These would allow the capture of full-resolution baseband data during events of interest; the baseband data are currently not saved due to the high data rate, but could be kept for a short period in a buffer that is written to disk when deemed worthwhile.  CPU computing performance continues to increase year on year, but is still not up to the task of processing the survey data in real-time (for acceptable monetary, power and floorspace costs). However, the recent appearance of graphics processing units (GPUs) in high-performance computing presents a new opportunity to attempt to break the real-time barrier. These devices can provide an order of magnitude more compute power than CPUs at comparable costs, but pose considerable software challenges due to their unfamiliar architectures [see \\citet{FlukeEtal2011} for an introduction]. It is our aim to implement the transient detection pipeline on a GPU and to exploit the boost in processing power to achieve a real-time processing rate. ",
        "conclusions": "Having successfully implemented each step of the transient detection process on a GPU, we are now working to integrate them into a complete GPU pipeline. Early benchmarks indicate that our target of processing one beam per GPU in real time is well within reach. Once the software pipeline is completely operational, we plan to deploy it on a small cluster of GPU nodes on-site at the Parkes radio telescope.  Our real-time radio transient detection pipeline promises to simplify the data processing procedure by performing it as observations are made. This new data reduction paradigm will enable immediate follow-up observations upon the detection of interesting transient events and provide unprecedented resolution of unique phenomena via baseband dumps."
    },
    "1112/1112.1705_arXiv.txt": {
        "abstract": "We study the relationship between two major baryonic components in galaxy clusters, namely the stars in galaxies, and the ionized gas in the intracluster medium (ICM), using 94 clusters that span the redshift range $0-0.6$. Accurately measured total and ICM masses from \\chandra\\ observations, and stellar masses derived from the Wide-field Infrared Survey Explorer and the Two-Micron All-Sky Survey allow us to trace the evolution of cluster baryon content in a self-consistent fashion. We find that, within $r_{500}$, the evolution of the ICM mass--total mass relation is consistent with the expectation of self-similar model, while there is no evidence for redshift evolution in the stellar mass--total mass relation.  This suggests that the stellar mass and ICM mass in the inner parts of clusters evolve differently. ",
        "introduction": "\\label{sec:intro}    The baryon content in groups and clusters of galaxies is one of the basic quantities that could be used to test models of structure formation. The total baryon mass fraction in massive clusters has been used to estimate cosmological parameters \\citep[e.g.,][]{swhite93a,allen04}. The mass of stars in galaxies is a fossil record of the star formation and galaxy merger history \\citep{conroy07}. The relative proportion in mass of hot (diffuse intracluster medium, ICM) and cold baryons (e.g., stars in galaxies) could be used to gauge the strengths of various feedback mechanisms and star formation efficiency \\citep[e.g.,][]{bode09}.   Although many measurements of the baryon fraction in local clusters exist \\citep[e.g.,][]{lin03b,gonzalez07,andreon10,balogh11}, such studies using large cluster samples at higher redshifts are in short supply. The first attempt has been made with the X-ray selected groups at $z=0.1-1$ detected in the COSMOS field \\citep[][hereafter G09]{giodini09}.  However, the total and ICM masses are not well-measured for individual systems, so G09 could not constrain $M_{\\rm ICM}/M_{\\rm star}$ well, especially at the massive end (see also \\citealt{leauthaud11}).    Here we present a measurement of the baryon mass fraction contained in galactic stars and ICM, as well as the relationship between $M_{\\rm ICM}$ and $M_{\\rm star}$, using 94 clusters at $z\\le 0.6$ -- the largest sample with accurate total and ICM mass measurements to date for this purpose. For the $z>0.1$ systems, we use data from the Wide-field Infrared Survey Explorer (WISE; \\citealt{wright10}) survey to determine $M_{\\rm star}$, while for the local clusters, we rely on data from the Two-Micron All-Sky Survey (2MASS, \\citealt{skrutskie06}). To ensure uniformity of our measurements and to facilitate comparison across the cosmic epochs, all quantities are measured within $r_{500}$, the radius within which the mean overdensity is 500 times the critical density of the Universe $\\rho_c$ at the cluster redshift.    In this paper we have neglected neutral and molecular gas in the baryon budget, as their amount is believed to be small \\citep[e.g.,][]{chung09}. Our analysis does not include the intracluster stars either (as they are well below the surface brightness limit of WISE), although they may contribute a non-negligible fraction to the total cluster stellar mass \\citep[e.g.,][]{gonzalez07}. Throughout this paper we adopt a {\\it WMAP-5} \\citep{komatsu09} $\\Lambda$CDM cosmological model where $\\Omega_M=1-\\Omega_\\Lambda=0.26$, $H_0=71 h_{71}\\,{\\rm km\\,s}^{-1} {\\rm Mpc}^{-1}$. ",
        "conclusions": ""
    },
    "1112/1112.5126_arXiv.txt": {
        "abstract": "We present both phenomenological and more physical photoionization models of the Chandra HETG spectra of the Seyfert-1 AGN NGC 4051. We detect 40 absorption and emission lines, encompassing highly ionized charge states from O, Ne, Mg, Si, S and the Fe L-shell and K-shell. Two independent photoionization packages, XSTAR and Cloudy, were both used to self-consistently model the continuum and line spectra. These fits detected three absorbing regions in this system with densities ranging from 10$^{10}$ to 10$^{11}$ cm$^{-3}$. In particular, our XSTAR models require three components that have ionization parameters of $\\log \\xi =$ 4.5, 3.3, \\& 1.0, and are located within the BLR at 70, 300, and 13,000 R$_g$, respectively, assuming a constant wind density. Larger radii are inferred for density profiles which decline with radius. The Cloudy models give a similar set of parameters with ionization parameters of $\\log \\xi$ = 5.0, 3.6, \\& 2.2 located at 40, 200, and 3,300 R$_g$.  We demonstrate that these regions are out-flowing from the system, and carry a small fraction of material out of the system relative to the implied mass accretion rate. The data suggest that magnetic fields may be an important driving mechanism. ",
        "introduction": "The advent of CCD spectroscopy with ASCA showed that the ``warm absorbers\" \\citep{Halpern84} that are present in the X-ray spectra of some Seyfert galaxies can be described with a combination of He-like O VII and H-like O VIII absorption edges \\citep{Reynolds97}.  The spectral resolution of CCDs is insufficient to clearly reveal individual narrow absorption lines that should be associated with such edges.  The high resolution dispersive spectrometers aboard {\\it Chandra} and {\\it XMM-Newton} have been able to clearly reveal such narrow absorption lines in a number of Seyfert AGN.  Indeed, it is now clear that X-ray warm absorbers span a range of ionizations, generally have blue-shifts on the order of $10^{2-3}~{\\rm km}~{\\rm s}^{-1}$ with respect to their host galaxy, and so must be outflows \\citep{Kaastra00, Kaspi00, Kaspi02}.  With deep observations and careful, self-consistent photoionization studies of high resolution spectra, the nature and constraints on the origin of these outflows can -- in principle -- be determined.  In some warm absorbers, there is evidence that the opacity must derive partly from elements trapped on dust grains \\citep{Lee01}, suggesting that the absorption may partially originate in a parsec--scale ``torus\" \\citep{Antonucci93}.  Similarly distant absorption is a viable description of the absorbers in many spectra \\citep[see, e.g.][]{ Blustin05}.  These outflows are plausibly powered by radiation pressure from the central engine.  In this picture, the outflows may be termed ``disk winds\", but they are unlikely to provide deep insights about the nature of disk accretion for they lack the ability to transfer angular momentum out of the system.  The best spectra of Seyfert warm absorbers may suggest a different origin.  The 900~ksec {\\it Chandra}/HETG spectrum of NGC 3783 is richly detailed, and photoionization analysis suggests that it may originate within several hundred gravitational radii (${\\rm r}_{g} = {\\rm GM}/{\\rm c}^{2}$) of the black hole \\citep{Kaspi02}. Detailed spectral modeling of the X-ray absorption in NGC 4151 finds high column densities and evidence of large bulk motion that may be require large non-radial motions \\citep{Kraemer05}.  This is qualitatively consistent with magnetocentrifugal disk winds of the kind described by \\cite{Blandford82}.  Owing to the fact that these winds can remove angular momentum from accreting gas (thus allowing it to accrete), some warm absorbers (or the most ionized, innermost regions of warm absorbers) may hold clues to the fundamental physics of disk accretion \\citep[see, e.g.][]{Proga03, Fukumura10} for newer theoretical models).  Even stronger evidence of magnetic processes driving disk winds and a possible anti-correlation with radio jets has been found in the observations of some stellar-mass black holes \\citep{Miller06, Miller08, Kallman09, Luketic10,Blum10}.  Understanding the ultimate influence of these warm outflows on host environments is also interesting and timely.  \\cite{Blustin05} considered a sample of Seyfert warm absorbers, and derived typical values for these outflows.  If the winds operate during the entirety of the active Seyfert/AGN phase, then warm absorbers may expel up to $10^{8}~{\\rm M}_{\\odot}$ of material into their host galaxies.  In large elliptical galaxies, this may be unimportant compared to the stellar mass.  However, this may represent a substantial fraction of the gas in the central spheroid of spiral galaxies.  Estimates of the radii at which warm absorbers exist, and their total mass outflow rates, depend on detailed spectroscopy. Factors such as the gas density, covering factor, and filling factor must be estimated using line ratios and a number of physical assumptions.  Indeed, the {\\it geometry} of warm absorber outflows is poorly understood; at present, only lines in NGC 3783 show even modest evidence of P-Cygni profiles \\citep{Kaspi02}, likely indicating that the wind is quasi-spherical.  It is notable, then, that the best spectra provide the best evidence that warm absorbers may be disk winds that could be related to fundamental disk physics.  Clearly, deep observations and sensitive spectra are vital to understanding accretion onto, and mass outflows from, the central engine in Seyfert AGN.  To this end, we have analyzed the summed, deep-exposure {\\it Chandra}/HETG spectrum of the Seyfert-1 galaxy NGC 4051.  This Seyfert is close ($z = 0.00234$), and the mass of the black hole has been constrained using reverberation mapping techniques (${\\rm M}_{BH} = 1.9\\pm 0.8 \\times 10^{6}~{\\rm M}_{\\odot}$, \\cite{Peterson04}; ${\\rm M}_{BH} = 1.7\\pm 0.5 \\times 10^{6}~{\\rm M}_{\\odot}$, \\cite{Denney09}).  The innermost accretion flow in NGC 4051 has been studied using relativistic disk reflection \\citep[e.g.,][or a review, see Miller 2007]{ Ogle04, Ponti06}.  The X-ray warm absorber in this AGN has previously been studied \\citep{Ogle04, Nucita10, Steenbrugge09, Lobban11, Pounds11}. We confirm the multiple absorbing regions with a range of ionization parameters, $1 < \\log \\xi < 5$ described in these papers. However, we do not find the various high velocity components, i.e. $> 2000$ km s$^{-1}$, claimed by \\cite{Steenbrugge09}, \\cite{Lobban11}, and \\cite{Pounds11}. Furthermore, we present evidence for material arising closer to the central engine than previously claimed, i.e. $< 13000 R_g$.  Herein, we present phenomenological and self-consistent, physical, photoionization models for the X-ray warm absorber observed in the time-averaged HETG spectrum of NGC 4051.  We find that the wind likely originates close to the black hole, within a few hundred gravitational radii.  We also find modest evidence of P-Cygni profiles in the H-like charge states, suggesting that the hottest, innermost part of the wind is quasi-spherical. A cosmology of H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, q0 =0 and $\\Lambda_0$ = 0.73 was used throughout this analysis. ",
        "conclusions": "We have used two separate modeling techniques, a phenomenological model and photoionization models, to quantify the absorption and emission features seen in the \\textit{Chandra} HETG spectra of NGC 4051. Amongst the results are three sets of line detections that afford us the ability to diagnose the density of the material. These are the He-like Ne IX and O VII triplets as well as the tentative Fe L-shell doublet, Fe XXII. We also constrain an upper limit to the Fe XXI ratio, another Fe L-shell density diagnostic. In particular, the fact that the Fe XXI and Fe XXII ratios are absorption features means that they are particularly effective probes of the density of the absorbing regions in so much that it is created within the absorbing region itself, whereas the emission features may come from another nearby region.  Following this phenomenological study, we implemented more self-consistent photoionization models via XSTAR and Cloudy. We used electron densities of n$_e=10^{10}$ and n$_e=10^{11}$ cm$^{-3}$ based on the gas density diagnostics, and the data then require three different absorption regions. The two highest ionized regions with  $\\log \\xi > 3.3$, are at the higher density, while the lower density region is less ionized, with $\\log \\xi \\approx 1-2$.  All three components may originate just at or within the BLR. The highly ionized regions likely originate close to the black hole, i.e. less than $300 R_g$ (assuming no radial drop-off in density) and have a large velocity as compared to the lower ionization region that has a lower velocity and resides farther out, i.e. greater than $3300 R_g$. Velocity broadening of the inner-most components would nominally show larger absorption line widths. However, in Section \\ref{res} we demonstrate that a small central source \\citep[$<$10 R$_g$][]{Chartas09} may only impart a small portion of the rotational velocity in the absorption components.  Furthermore, we demonstrate that the mass outflow rate is small compared to the overall mass accretion onto the black hole. It is possible that magnetic fields are partly responsible for driving these winds, at least in the regions closest to the black hole. This is very similar to the galactic black holes, GRO 1655-40 \\citep{Miller08}, H1743-322 \\citep{Miller06b}, 4U 1630-472 \\citep{Kubota07} the dwarf nova, OY Carinae \\citep{Mauche00} and the Seyfert, NGC 4151 \\citep{Kraemer05},  in which magnetic fields are also found to potentially be important in driving outflows.  Finally, we note the data require a relativistic Fe K$\\alpha$ line, which nominally implies a maximal spin 0.998$_{-0.08}$ with an inclination angle of 31.7$^{+0.8}_{-0.7}$$^\\circ$. At a 2 $\\sigma$ level, the data require a relativistic O VIII Lyman $\\alpha$ line. \\cite{Steenbrugge09} and \\cite{Ogle04} also find evidence for a relativistic O VIII Lyman $\\alpha$ line.  \\vspace{0.2in} The authors would like to thank the anonymous referee for their invaluable comments and improvements to this paper.   \\begin{deluxetable*}{c c c c c } \\tablecolumns{5} \\tablewidth{0pc} \\tablecaption{Broad-Band Parameters} \\tablehead{ \\colhead{$\\Gamma$} & \\colhead{Normalization} & \\colhead{Constant} & \\colhead{Diskbb}  & \\colhead{Normalizatoin}  \\\\  & (10$^{-3}$ photons cm$^{-2}$ s$^{-1}$ keV$^{-1}$) &  &(keV) & $\\left( \\left( \\frac{R_{in}/km}{D/kpc} \\right)^2\\cos \\theta \\right)$} \\startdata 1.94 $\\pm$ 0.01 & 4.24 $\\pm$ 0.03 & 1.062 $\\pm$ 0.005 &  0.118 $\\pm$ 0.005 & 6610 $_{-910}^{+1100} $ \\\\ \\enddata \\label{BroadBand} \\tablecomments{This table gives the parameters for the broad-band continuum of both the HEG and MEG spectra. Because of a soft excess component in the MEG spectra, a disk blackbody component was included. The disk blackbody component should not be regarded as physical. The multiplicative constant is applied to the MEG spectrum to correct for the slight discrepancy between the flux in the two spectrographs. Galactic absorption of 1.15 $\\times 10^{20}$ atoms cm$^{-2}$ was also included. $\\chi^2/\\nu$ = 9067/7531} \\end{deluxetable*}    \\begin{deluxetable*}{l l l l l l l l } \\tablecolumns{8} \\tabletypesize{\\scriptsize} \\tablewidth{0pc} \\tablecaption{Emission Line Parameters} \\tablehead{ \\colhead{line center} & \\colhead{$\\sigma $} & \\colhead{Normalization} & \\colhead{EW} & \\colhead{$\\lambda_\\circ$ }& \\colhead{velocity } & \\colhead{Ftest} & \\colhead{ID} \\\\ (\\AA\\/) & (m\\AA\\/) & (10$^{-5}$ ph cm$^{-2}$ s$^{-1}$)  & (m\\AA\\/) & (\\AA\\/) & (km s$^{-1}$) & &} \\startdata  22.135 $\\pm$ 0.003  &    12.8 $_{-    4.6}^{+    4.9}$ &    8.73 $_{-    1.37}^{+    1.40}$ &  132 $\\pm 21$ & 22.101 &   -240 $_{-  50}^{+  40}$ & 5.97$\\times 10^{-08}$  & O VII \\\\ 21.831 $_{-0.027}^{+0.026}$ &    11.6 $_{-    5.6}^{+   10.3}$ &    2.54 $_{-    0.88}^{+    1.30}$ &   37.8 $_{-   13.1}^{+   19.4}$ & 21.807 &   -370 $_{- 370}^{+ 360}$ & 5.37$\\times 10^{-02}$ &    O VII \\\\ 19.013 $_{-0.003}^{+0.006}$ &    12.4 $_{-    7.4}^{+    6.5}$ &    2.52 $_{-    0.92}^{+    0.79}$ &   37.6 $_{-   13.7}^{+   11.9}$ & 18.967 &     18 $_{-  40}^{+  99}$ & 2.10$\\times 10^{-04}$& O VIII \\\\ 18.668 $_{-0.003}^{+0.002}$ &     0.00 $_{}^{+    4.17}$ &    1.07 $\\pm$ 0.31  &   56.8 $\\pm$ 16.6  & 18.627 &    -40 $_{-  49}^{+  31}$ & 5.32$\\times 10^{-03}$ &      O VII \\\\ 16.038 $_{-0.003}^{+0.008}$ &     8.89 $_{-    5.44}^{+    6.34}$ &    1.30 $_{-    0.81}^{+    0.59}$ &   48.0 $_{-   29.9}^{+   21.7}$ & 16.006 &    -100 $_{-  50}^{+ 150}$ & 3.84$\\times 10^{-02}$ &     O VIII \\\\ 13.716 $_{-0.004}^{+0.003}$ &     8.34 $_{-    5.36}^{+    5.37}$ &    0.83 $_{-    0.20}^{+    0.23}$ &   17.9 $_{-    4.3}^{+    5.0}$ & 13.699 &   -330 $_{-  90}^{+  70}$ & 4.76$\\times 10^{-05}$ &   Ne IX \\\\ 13.579 $_{-0.006}^{+0.005}$ &    13.48 $_{-    5.91}^{+    0.01}$ &    0.75 $_{-    0.19}^{+    0.18}$ &   16.3 $_{-    4.2}^{+    3.9}$ & 13.553 &   -117 $_{- 123}^{+ 116}$ & 1.15$\\times 10^{-03}$ &  Ne IX \\\\ 12.155 $_{-0.004}^{+0.005}$ &     0.64 $_{-    0.64}^{+    8.48}$ &    0.41 $_{-    0.13}^{+    0.01}$ &   10.1 $_{-    3.2}^{+    0.2}$ & 12.134 &   -172 $_{-  92}^{+ 113}$ & 4.04$\\times 10^{-03}$ &      Ne X \\\\ 8.467 $_{-0.002}^{+0.001}$ &     0.07 $_{-    0.07}^{+    3.10}$ &    0.21 $\\pm$   0.05  &    5.7 $\\pm$ 1.5  &--   &    -- & 1.80$\\times 10^{-03}$ & (Mg XII ?)\\\\ 4.288 $_{-0.001}^{+0.004}$ &     0.00 $_{}^{+    5.87}$ &    0.23 $\\pm$ 0.07  &    6.6 $\\pm$  1.9  &  4.299 &  -1460 $_{-  40}^{+ 290}$ & 2.49$\\times 10^{-03}$ &     S XV \\\\ 1.941 $_{-0.001}^{+0.002}$ &     1.48 $_{-    1.48}^{+    3.30}$ &    0.54 $_{-    0.12}^{+    0.17}$ &   10.8 $_{-    2.5}^{+    3.5}$ &  1.940 &   -620 $_{- 170}^{+ 230}$ & 4.51$\\times 10^{-04}$ & Fe K$\\alpha$ (narrow) \\\\ 1.927 $_{-0.012}^{+0.011}$ &    62.9 $_{-   11.1}^{+   14.4}$ &    2.31 $_{-    0.42}^{+    0.43}$ &   65.9 $_{-   11.9}^{+   12.2}$ &  1.940 &  -2720 $_{-1830}^{+1720}$ & 1.22$\\times 10^{-06}$ &   Fe K$\\alpha$ (broad) \\\\ 1.781 $\\pm$ 0.004  &    10.00 $_{-    3.38}^{+    0.01}$ &    0.64 $_{-    0.17}^{+    0.16}$ &   17.9 $_{-    4.8}^{+    4.5}$ &  1.750 &   4530 $_{- 650}^{+ 700}$ & 5.23$\\times 10^{-03}$ &     Fe K$\\beta$ \\\\   \\enddata \\label{lineId} \\tablecomments{ The above table lists the properties of the emission lines measured in HETG coadded spectrum of NGC 4051. A range of ionization and velocity shifts are measured. The first column lists the initial wavelength. The second and third list the width as well as the normalization for the corresponding line. Those lines with a $\\sigma$=0 are unresolved. The Equivalent width is given in the fourth column. $\\lambda_0$ is the rest wavelength in NGC 4051 rest frame, which is at a z= 0.00234. The velocity is calculated using these rest wavelengths in column 6. The seventh column lists the significance of detection as 1 minus the confidence level. The final column lists particular ion that is being detected. All the line energies are fit to the data, which explains why the energy is different for the narrow and broad Fe K$\\alpha$ lines.} \\end{deluxetable*}   \\begin{deluxetable*}{l l l l l l l l l} \\tablecolumns{9} \\tablewidth{0pc} \\tabletypesize{\\scriptsize} \\tablecaption{Absorption Line Parameters} \\tablehead{ \\colhead{line center} & \\colhead{$\\sigma $} & \\colhead{Normalization} & \\colhead{EW} & \\colhead{$\\lambda_\\circ$ }& \\colhead{velocity } & \\colhead{Ftest} &\\colhead{N$_z$ }& \\colhead{ID} \\\\ (\\AA\\/) & (m\\AA\\/) & (10$^{-5}$ ph cm$^{-2}$ s$^{-1}$)  & (m\\AA\\/) & (\\AA\\/) & (km s$^{-1}$) & &(10$^{16}$ cm$^{-2}$) } \\startdata 21.617 $_{-0.008}^{+0.006}$ &    13.2 $_{-   13.2}^{+    9.6}$ &   -2.71 $_{-    0.66}^{+    0.52}$ &  -39.9 $_{-    9.7}^{+    7.6}$ & 21.602 &   -497 $_{- 111}^{+  79}$ & 3.04$\\times 10^{-07}$ &      1.2 $\\pm$  0.2  & O VII \\\\ 18.974 $_{-0.004}^{+0.001}$ &    15.7 $_{-    6.0}^{+    4.8}$ &   -3.46 $_{-    0.61}^{+    0.77}$ &  -48.2 $_{-    8.5}^{+   10.7}$ & 18.967 &   -584 $_{-  63}^{+  10}$ & 8.29$\\times 10^{-35}$ &      4.8 $\\pm$  0.8  & O VIII \\\\ 18.650 $_{-0.0012}^{}$ &    33.8 $_{-    5.2}^{+    5.7}$ &   -4.89 $_{-    0.64}^{+    0.61}$ &  -70.8 $_{-    9.3}^{+    8.9}$ & 18.627 &   -331 $_{-  19}^{}$ & 8.37$\\times 10^{-28}$ &     14 $\\pm$  2 & O VII \\\\ 17.790 $_{-0.008}^{+0.003}$ &     1.29 $_{-    1.29}^{+   12.86}$ &   -0.94 $_{-    0.58}^{+    0.17}$ &  -14.8 $_{-    9.1}^{+    2.6}$ & 17.768 &   -332 $_{- 126}^{+  43}$ & 6.11$\\times 10^{-08}$ &      8.5 $_{-   1.3}^{+   1.4}$ & O VII \\\\ 16.020 $_{-0.0074}^{}$ &    14.3 $_{-    7.1}^{+    5.5}$ &   -1.98 $_{-    0.52}^{+    0.85}$ &  -34.5 $_{-    9.1}^{+   14.8}$ & 16.006 &   -438 $_{- 138}^{}$ & 3.94$\\times 10^{-17}$ &     51 $\\pm$  8 & O VIII \\\\ 15.266 $\\pm$ 0.004 &     5.38 $_{-    5.38}^{+    4.79}$ &   -0.46 $_{-    0.15}^{+    0.14}$ &   -9.0 $_{-    3.0}^{+    2.8}$ & 15.261 &   -592 $_{-  76}^{+  77}$ & 4.17$\\times 10^{-03}$ &      0.6 $\\pm$ 0.1  & Fe XVII \\\\ 15.017 $\\pm$ 0.002  &     4.98 $_{-    4.98}^{+    3.63}$ &   -0.76 $_{-    0.16}^{+    0.14}$ &  -15.3 $_{-    3.3}^{+    2.9}$ & 15.014 &   -641 $\\pm$ 45  & 1.27$\\times 10^{-11}$ &      0.3 $\\pm$  0.1  & Fe XVII \\\\ 14.220 $\\pm$ 0.006  &    28.2 $_{-    5.8}^{+    6.7}$ &   -1.55 $_{-    0.32}^{+    0.29}$ &  -32.7 $_{-    6.7}^{+    6.2}$ & 14.203 &   -334 $_{- 137}^{+ 132}$ & 2.06$\\times 10^{-08}$ &      1.7 $\\pm$ 0.3 & Fe XVIII \\\\ 14.042 $_{-0.013}^{+0.010}$ &    23.4 $_{-    6.1}^{+   10.4}$ &   -0.83 $_{-    0.29}^{+    0.25}$ &  -17.9 $_{-    6.3}^{+    5.3}$ & -- & -- & 4.79$\\times 10^{-03}$ &    -- & Unidentified \\\\ 13.841 $\\pm$ 0.002   &     8.41 $_{-    3.57}^{+    4.51}$ &   -0.76 $_{-    0.18}^{+    0.15}$ &  -16.1 $_{-    3.8}^{+    3.1}$ & 13.844 &   -751 $_{-  49}^{+  52}$ & 1.53$\\times 10^{-09}$ &     16.2 $\\pm$ 2.6  & Fe XIX \\\\ 13.658 $_{-0.005}^{+0.003}$ &    14.3 $_{-    5.3}^{+    8.3}$ &   -0.56 $_{-    0.19}^{+    0.17}$ &  -12.1 $_{-    4.2}^{+    3.7}$ & 13.670 &   -971 $_{- 119}^{+  58}$ & 4.47$\\times 10^{-03}$ &      1.4 $\\pm$  0.2  & Fe XIX \\\\ 13.515 $\\pm$ 0.002  &     8.18 $_{-    2.91}^{+    3.19}$ &   -0.73 $_{-    0.14}^{+    0.13}$ &  -16.0 $_{-    3.1}^{+    2.8}$ & 13.518 &   -769 $_{-  51}^{+  53}$ & 8.30$\\times 10^{-11}$ &      1.1 $\\pm$   0.2 & Fe XIX \\\\ 13.450 $\\pm$ 0.002  &    14.6 $_{-    2.5}^{+    2.1}$ &   -1.23 $_{-    0.16}^{+    0.15}$ &  -26.9 $_{-    3.5}^{+    3.3}$ & 13.447 &   -629 $_{-  52}^{+  51}$ & 6.48$\\times 10^{-22}$ &      2.1 $\\pm$  0.3  & Ne IX \\\\ 12.832 $\\pm$ 0.003 &     9.78 $_{-    2.99}^{+    3.34}$ &   -0.66 $_{-    0.12}^{+    0.11}$ &  -15.1 $_{-    2.8}^{+    2.6}$ & 12.847 &  -1046 $_{-  62}^{+  67}$ & 2.33$\\times 10^{-10}$ &      2.2 $\\pm$  0.4  & Fe XX \\\\ 12.281 $_{-0.004}^{+0.003}$ &     3.91 $_{-    3.91}^{+    9.75}$ &   -0.35 $_{-    0.18}^{+    0.09}$ &   -8.3 $_{-    4.3}^{+    2.1}$ & 12.285 &   -779 $_{- 106}^{+  65}$ & 1.07$\\times 10^{-04}$ &      0.6 $\\pm$  0.1  & Fe XXI \\\\ 12.132 $_{-0.002}^{+0.012}$ &     9.12 $_{-    2.72}^{+    6.05}$ &   -0.87 $_{-    0.48}^{+    0.14}$ &  -20.9 $_{-   11.4}^{+    3.4}$ & 12.134 &   -737 $_{-  49}^{+ 303}$ & 7.62$\\times 10^{-25}$ &      5.1 $\\pm$  0.8  & Ne X \\\\ 11.935 $_{-0.0002}^{+0.045}$ &     0.00 $_{}^{+    6.86}$ &   -0.13 $\\pm$ 0.07  &   -3.1 $\\pm$ 1.7  & 11.920 &   -320 $_{-   4}^{+1100}$ & 4.6$\\times 10^{-02}$ &      0.4 $\\pm$ 0.1   & Fe XXII \\\\ 11.771 $\\pm$ 0.002  &     2.90 $_{-    2.90}^{+    4.29}$ &   -0.32 $_{-    0.08}^{+    0.07}$ &   -8.1 $_{-    2.1}^{+    1.7}$ & 11.770 &   -687 $_{-  56}^{+  55}$ & 3.54$\\times 10^{-05}$ &      0.9 $\\pm$  0.1  & Fe XXII \\\\ 11.655 $_{-0.005}^{+0.001}$ &     0.16 $_{-    0.16}^{+    4.44}$ &   -0.19 $\\pm$ 0.06  &   -4.7 $\\pm$  1.6  & -- & -- & 4.40$\\times 10^{-02}$ &   -- & Unidentified\\\\ 11.551 $\\pm$ 0.003  &     6.11 $_{-    6.11}^{+    3.55}$ &   -0.33 $\\pm$   0.08 &   -8.4 $\\pm$  2.1 & 11.547 &   -602 $\\pm$  74 & 1.01$\\times 10^{-04}$ &      4.2 $\\pm$ 0.7 & Ne IX \\\\ 11.415 $_{-0.0001}^{+0.005}$ &     0.00 $_{}^{+    4.50}$ &   -0.18 $\\pm$  0.06 &   -4.6 $\\pm$  1.5 & 11.427 &  -1014 $_{-   1}^{+ 131}$ & 3.48$\\times 10^{-02}$ &      0.2 $\\pm$ 0.1 & Fe XXII \\\\ 11.001 $\\pm$ 0.004 &     8.03 $_{-    4.35}^{+    4.15}$ &   -0.25 $\\pm$  0.08 &   -6.3 $_{-    2.0}^{+    1.9}$ & 11.001 &   -691 $_{- 121}^{+ 120}$ & 1.13$\\times 10^{-02}$ &      9.3 $\\pm$ 1.5 & Ne IX \\\\ 10.237 $_{-0.004}^{+0.003}$ &     5.41 $_{-    5.41}^{+    4.84}$ &   -0.25 $_{-    0.08}^{+    0.07}$ &   -6.4 $_{-    1.9}^{+    1.6}$ & 10.238 &   -738 $_{- 103}^{+  85}$ & 5.60$\\times 10^{-04}$ &     11.5 $\\pm$ 1.8 & Ne X \\\\ 9.838 $_{-0.0003}^{+0.002}$ &     0.00 $_{}^{+    3.03}$ &   -0.18 $\\pm$ 0.05 &   -4.8 $\\pm$ 1.3 &  --&  --& 3.53$\\times 10^{-03 }$&   --  & (Fe XXI?)\\\\ 9.174 $_{-0.004}^{+0.003}$ &     8.19 $_{-    3.33}^{+    4.65}$ &   -0.28 $_{-    0.01}^{+    0.08}$ &   -7.7 $_{-    0.3}^{+    2.3}$ &  9.169 &   -529 $_{- 118}^{+ 110}$ & 6.97$\\times 10^{-04}$ &      1.2 $\\pm$ 0.2 & Mg XI \\\\ 8.421 $\\pm$ 0.001 &     2.44 $_{-    2.44}^{+    1.84}$ &   -0.26 $\\pm$ 0.05 &   -7.2 $\\pm$ 1.4&  8.419 &   -635 $_{-  34}^{+  36}$ & 2.23$\\times 10^{-13}$ &      3.7 $\\pm$ 0.6 & Mg XII \\\\ 6.650 $_{-0.0018}^{+0.0024}$ &     4.32 $_{-    2.60}^{+    2.84}$ &   -0.23 $_{-    0.07}^{+    0.06}$ &   -6.5 $_{-    2.0}^{+    1.6}$ &  6.648 &   -612 $_{-  79}^{+ 109}$ & 3.34$\\times 10^{-05}$ &      1.9 $\\pm$  0.3 & Si XIII \\\\  \\enddata \\label{lineIdabs} \\tablecomments{ This table lists the absorption lines detected in our phenomenological modeling in the HETG spectrum of NGC 4051. All of the features appear to be blue shifted with respect to the host galaxy, indicating a bulk outflow from the galaxy. The first column lists the initial wavelength. The second and third list the width as well as the normalization for the corresponding line. A negative normalization indicates an absorption feature. Those lines with a $\\sigma$=0 are unresolved. The Equivalent width is given in the fourth column. $\\lambda_0$ is the rest wavelength in NGC 4051 rest frame, which is at a z= 0.00234. The velocity is calculated using these rest wavelengths in column 6. The seventh column lists the significance of detection as 1 minus the confidence level. The eighth column lists the column densities derived using the relation $W_{\\lambda}=(\\pi e^2 / m_ec^2) N_j \\lambda^2 f_{ij}$, where $W_\\lambda$ is the equivalent width, $N_j$ is the column density, $\\lambda$ is the rest wavelength, and $f_{ij}$ is the oscillator strength. The final column lists particular ion that is being detected. All the line energies are fit to the data.}  \\end{deluxetable*}   \\begin{deluxetable*}{c c c | c  | c c | c c} \\tablecolumns{8} \\tablewidth{0pt} \\tablecaption{XSTAR Model} \\tablehead{ & & & Absorption & \\multicolumn{2}{|c|}{Emission} && \\\\ \\hline  \\colhead{Density} & \\colhead{Column Density }  & \\colhead{$\\log (\\xi) $} &\\colhead{velocity} & \\colhead{Normalization} & \\colhead{velocity} & \\colhead{$\\Delta \\chi^2$}  &\\colhead{ Ftest}  \\\\ (cm$^{-3}$) &  (10$^{20}$ cm$^{-2}$) & (ergs cm s$^{-1}$) &  (km s$^{-1}$)  & (10$^{-3}$) & (km s$^{-1}$) &&  } \\startdata 10$^{11}$&  8.10$^{+  0.92}_{ -0.99}$&   4.50$^{}_{ -0.06}$&   -680$^{+    20}_{   -60}$&  19.53$^{+  5.22}_{ -5.08}$&    -50$^{+   150}_{  -250}$& 106 & 6.3$\\times 10^{-21}$\\\\ 10$^{11}$& 10.10$^{+  0.63}_{ -0.95}$&   3.28$^{+  0.05}_{ -0.03}$&   -630$^{+    10}_{   -80}$&   1.01$^{+  0.59}_{ -0.44}$&    100$^{+   100}_{  -110}$& 1076 & 1.8$\\times 10^{-193}$\\\\ 10$^{10}$&  1.97$^{+  0.27}_{ -0.25}$&   1.00$^{+  0.06}_{ -0.15}$&   -400$^{+   380}_{  -270}$&   1.08$^{+  0.27}_{ -0.16}$&   -300$^{+    70}_{   -40}$ & 248 & 26.5$\\times 10^{-48}$\\\\  \\hline \\hline Power-law & Normalization & Constant & Diskbb & Normalization  & Phabs \\\\ ($\\Gamma$) & (10$^{-3}$ ph &  & (keV)  & $\\tiny{\\left( \\left( \\frac{R_{in}/km}{D/kpc} \\right)^2\\cos \\theta \\right)}$ &$\\times10^{20}$ \\\\ &cm$^{-2}$ s$^{-1}$ keV$^{-1}$) & &&& (atoms cm$^{-2}$) \\\\ \\hline 1.95 $\\pm$ 0.01&   4.32$^{+  0.04}_{ -0.03}$&  1.059$^{+ 0.005}_{-0.004}$ &  0.133$^{+ 0.002}_{-0.003}$&  3.78$^{+ 0.40}_{-0.37}\\times10^{3}$& 1.15\\\\  \\enddata  \\label{xstar} \\tablecomments{ The above table lists the parameters of the three XSTAR model components. The errors reported here are $1 \\sigma$. $\\chi^2 / \\nu$ 8177/7516=1.09. The actual model is as follows: phabs*(mtable1*mtable2*mtable3(power-law+diskbb)+atable1 + atable2 + atable3)*constant) The multiplicative constant is applied to the MEG spectrum to correct for the slight discrepancy between the flux in the two spectrographs. * Denotes the model reached a maximum in the ionization parameter limits} \\end{deluxetable*}   \\begin{deluxetable*}{c c c |  c | c c | c c} \\tablecolumns{8} \\tablewidth{0pc} \\tablecaption{Cloudy Model} \\tablehead{& & & Absorption & \\multicolumn{2}{|c|}{Emission} && \\\\ \\hline  \\colhead{Density} & \\colhead{Column Density }  & \\colhead{Radius} &\\colhead{velocity} & \\colhead{Normalization} & \\colhead{velocity} & \\colhead{$\\Delta \\chi^2$}  &\\colhead{ Ftest}  \\\\ (cm$^{-3}$) &  (10$^{20}$ cm$^{-2}$) & (10$^{13}$ cm) &  (km s$^{-1}$)  & (10$^{-21}$) & (km s$^{-1}$) &&  } \\startdata 10$^{11}$& 21.8$^{+  4.4}_{ -3.5}$&   1.00$^{+  0.04}_{}$&  -1090$^{+   110}_{  -120}$ &  0.007$\\pm$ 0.002 &   -110$^{+    80}_{  -130}$& 23 & 1.4 $\\times$ 10$^{-3}$ \\\\ 10$^{11}$&  7.68$^{+  0.84}_{ -0.75}$&   5.09$^{+  0.29}_{ -0.26}$ &  -1000$^{+    70}_{   -30}$ &  0.015$\\pm$ 0.006 &    220$^{+   180}_{   -60}$& 642 & 7.6 $\\times$ 10$^{-96}$ \\\\ 10$^{10}$&  6.42$^{+  0.65}_{ -0.81}$&  82.9$^{+  4.7}_{ -4.4}$ &   -700$^{+    30}_{   -20}$ &  2.47$^{+ 0.80}_{ -0.70}$&   -290$^{+    60}_{  -180}$& 484 & 1.9 $\\times$ 10$^{-54}$\\\\  \\hline \\hline Power-law & Normalization & Constant & Diskbb & Normalization  & Phabs \\\\ ($\\Gamma$) & (10$^{-3}$ photons  &  & (keV)  & $\\tiny{\\left( \\left( \\frac{R_{in}/km}{D/kpc} \\right)^2\\cos \\theta \\right)}$ &$\\times10^{20}$ \\\\ &cm$^{-2}$ s$^{-1}$ keV$^{-1}$) & &&& (atoms cm$^{-2}$) \\\\ \\hline 1.97 $\\pm$ 0.01 &   4.42 $^{+  0.04}_{ -0.03}$&  1.06 $\\pm$ 0.01&  0.127 $^{+ 0.002}_{-0.003}$&  4.81 $^{+ 0.77}_{-0.57}\\times10^{3}$ & 1.15  \\\\ \\enddata \\label{cloudy} \\tablecomments{ This table lists the Cloudy model parameters of the three components required by the data. The errors reported here are $1 \\sigma$. $\\chi^2 / \\nu $ =8713/7516=1.16. The actual model is as follows: phabs*(mtable1*mtable2*mtable3(power-law+diskbb)+atable1 + atable2 + atable3)*constant). The multiplicative constant is applied to the MEG spectrum to correct for the slight discrepancy between the flux in the two spectrographs. } \\end{deluxetable*}     \\begin{deluxetable*}{l | l l | l l | l l} \\tablecolumns{7} \\tablewidth{0pc} \\tablecaption{Photoionization Model} \\tablehead{ & XSTAR & Cloudy & XSTAR & Cloudy & XSTAR & Cloudy } \\startdata Density (cm$^{-3}$) & \\multicolumn{2}{c}{$10^{11}$} & \\multicolumn{2}{c}{$10^{11}$} & \\multicolumn{2}{c}{$10^{10}$}\\\\ $\\log \\xi$ (ergs cm$^{-1}$) &  4.50 $^{*}_{ -0.06}$&   5.00 $^{*}_{ -0.03}$&   3.28 $^{+  0.05}_{ -0.03}$&   3.59 $\\pm$ 0.05 &   1.00 $^{+  0.06}_{ -0.15}$&   2.16 $\\pm$ 0.05 \\\\ Radius ($10^{13}$cm) &  1.78 $^{+  0.12}_{*}$&   1.00 $^{+  0.04}_{ *}$&   7.25 $^{+  0.29}_{ -0.38}$&   5.09 $^{+  0.29}_{ -0.26}$& 317 $^{+ 58}_{-20}$&  82.9 $^{+  4.7}_{ -4.4}$\\\\ Radius (R$_g$) &      71 $^{+       4}_{  *}$&       40 $^{+       1}_{       *}$&      290 $^{+      11}_{     -15}$&      203 $^{+      11}_{     -10}$&    12700 $^{+    2300}_{    -800}$&     3310 $^{+     190}_{    -170}$\\\\ $\\Delta$(r) (10$^{10}$cm) &  0.81 $^{+  0.09}_{ -0.10}$&   2.18 $^{+  0.44}_{ -0.35}$&   1.01 $^{+  0.06}_{ -0.09}$&   0.77 $\\pm$ 0.08 &   0.20 $\\pm$ 0.03 &   0.64 $^{+  0.07}_{ -0.08}$\\\\ $C_v$ (10$^{-4}$)&      4.56 $^{*}_{      0.81}$&      21.8 $^{*}_{      4.2}$&       1.39 $^{+      0.17}_{      0.18}$&       1.51 $^{+      0.25}_{      0.22}$&       0.06 $^{+      0.01}_{      0.02}$&       0.77 $\\pm$  0.13 \\\\ $\\dot{M}_{out}$ ($10^{-6}$M$_{\\odot}$yr$^{-1}$ )&      0.30 $^{+      0.09}_{  *}$&       0.73 $^{+      0.28}_{   *}$&       1.42 $^{+      0.35}_{     -0.22}$&       1.20 $^{+      3.39}_{ -     1.82}$&       7.7 $^{+      9.6}_{     -7.4}$&      11.5 $^{+   7.1}_{     -8.6}$\\\\ \\enddata \\label{photo} \\tablecomments{ This table compares the XSTAR and Cloudy model components. The models give similar values, and consistent mass outflow rates. We note that in general the lowest ionization component carries the most mass. We have assumed 1.7 $\\pm$ $0.5 \\times 10^6 M_\\odot$ as the mass of the SMBH. The symbol, *, denotes both upper and lower limits in the models.} \\end{deluxetable*}     \\begin{deluxetable*}{c | c c c c c c c} \\tablecolumns{8} \\tablecaption{Comparing Models} \\tablehead{ &\\scriptsize{Our XSTAR} & \\scriptsize{Collinge}& \\scriptsize{Krongold} &\\scriptsize{Steenbrugge}  &\\scriptsize{Nucita} & \\scriptsize{Lobban} & \\scriptsize{Pounds}  \\\\ & \\scriptsize{Model} & \\scriptsize{et al. (2001)} & \\scriptsize{et al. (2007)}  & \\scriptsize{et al. (2009)}  & \\scriptsize{et al. (2010)}  & \\scriptsize{et al. (2011)}  & \\scriptsize{et al. (2011)} \\\\ Instrument & (HETG) & (HETG) & (RGS )  & (LETGS) ** & (RGS) & (HETG) & (RGS) } \\startdata $\\log \\xi_0$ (ergs cm s$^{-1}$) & -  & -&  - & 0.07 $\\pm$ 0.13 & - & -0.86 $\\pm$ 0.30 & 0.32 $\\pm$ 24 \\\\ v$_0$ (km s$^{-1}$) & &&& -210 $\\pm$ 70 & &-180 $\\pm$ 100  & +120 $\\pm$ 60\\\\ $\\log n_0$ (cm$^{-3}$)&&& &$<$ 6.5  && - & -\\\\ $\\log r_0$ (cm) && &&  $>$ 17.95 &  &$\\lesssim$ 23.4 & -\\\\ $\\log$ N$_{H,0}$ (cm$^{-2}$)& &&&20.08 $\\pm$ 0.18 && 20.48 $\\pm$ 0.13 & 20.0 $\\pm$ 0.1\\\\ \\\\ $\\log \\xi_1$(ergs cm s$^{-1}$) &1.00 $^{+0.06}_{-0.15}$ & - & - & 0.87 $\\pm$ 0.09 & - & 0.60 $\\pm$ 0.30 & 1.43 $\\pm$ 0.22 \\\\ v$_1$ (km s$^{-1}$) &-400 $^{+380}_{-270}$ & -600 $\\pm$ 130 & &-200 $\\pm$ 30 & & -220 $\\pm$ 60  & -530 $\\pm$ 90\\\\ $\\log n_1$ (cm$^{-3}$)  & 10& $<$ 10.6 && $>$ 6.6 & & $<$ 11 & $\\sim$ 6.7 \\\\ $\\log r_1$ (cm) &15.5 $\\pm$ 0.1& - && $<$ 17.48 & &15.8 -- 22.6 &$\\sim$ 18\\\\ $\\log$ N$_{H,1}$ (cm$^{-1}$) & 20.3 $\\pm$ 0.1& $>$ 21.3   && 20.46 $\\pm$ 0.08 & &20.31 $\\pm$ 0.13 &  20.04 $\\pm$ 0.26\\\\ \\\\ $\\log \\xi_{2a}$(ergs cm s$^{-1}$) & - & - & 2.18 $\\pm$ 0.02 * & 2.32 $\\pm$ 0.13 & 2.15 $\\pm$ 0.02 & 1.96 $\\pm$ 0.17 &  2.52 $\\pm$ 0.07 \\\\ v$_{2a}$ (km s$^{-1}$) & &&-492 $\\pm$ 97 &-580 $\\pm$ 50 & 0 &-820 $\\pm$ 30 & -3850 $\\pm$ 60 \\\\ $\\log n_{2a}$ (cm$^{-3}$) & & & $>$ 7.9 &$<$ 6.7&$ \\lesssim$ 7 &- & $\\sim$ 5.7\\\\ $\\log r_{2a}$ (cm) & && $<$ 14.9& $>$ 16.70 & $\\gtrsim$ 16.8 & $\\lesssim$ 20.3 & $\\sim$ 17.8 \\\\ $\\log$ N$_{H,2a}$ (cm$^{-2}$) && & 20.73 $\\pm$ 0.17 & 20.90 $\\pm$ 0.3 & 23.36 $\\pm$ 0.01 &20.70 $\\pm$ 0.16 & 21.26 $\\pm$ 0.10\\\\ \\\\ $\\log \\xi_{2b}$(ergs cm s$^{-1}$) &- & - & - & - &  - &2.16 $\\pm$ 0.21 & 2.77 $\\pm$ 0.15\\\\ v$_{2b}$ (km s$^{-1}$) & && & &&-550 $\\pm$ 60 & -400 $\\pm$ 50 \\\\ $\\log n_{2b}$ (cm$^{-3}$) &  &&&&& - & -\\\\ $\\log r_{2b}$ (cm) & &&& &&$\\lesssim$ 20.4 & -   \\\\ $\\log$ N$_{H,2b}$ (cm$^{-2}$)& & & &&& 21.04 $\\pm$ 0.27 & 21.46 $\\pm$ 0.29\\\\ \\\\ $\\log \\xi_3$(ergs cm s$^{-1}$) &3.28 $^{+0.05}_{0.03}$&- &- &  3.19 $\\pm$ 0.09 & - &2.97 $\\pm$ 0.10 & 2.97 $\\pm$ 0.06\\\\ v$_3$ (km s$^{-1}$) &-620 $^{+10}_{-80}$ & -2340 $\\pm$ 130 & &-4670 $\\pm$ 150 & &-710 $\\pm$ 40 & -5880 $\\pm$ 60\\\\ $\\log n_3$ (cm$^{-3}$) & 11 & $<$ 10.6 & &$<$ 5.4  & &- &  $\\sim$ 5.7  \\\\ $\\log r_3$ (cm)  & 13.86 $\\pm$ 0.02 & - && $>$ 16.85 & & $\\lesssim$ 18.7 & $\\sim$ 17.8 \\\\ $\\log$ N$_{H,3}$ (cm$^{-2}$) &21.00 $^{+0.03}_{-0.04}$ & 19 -- 20.7 && 22.30 $\\pm$ 0.3 & & 21.44 $\\pm$ 0.12 & 22.15 $\\pm$ 0.07\\\\ \\\\ $\\log \\xi_4$(ergs cm s$^{-1}$) & $\\gtrsim$4.5 & - &4.42 $\\pm$ 0.07 * &-  &- & 4.1 $\\pm$ 0.2 & -\\\\ v$_4$ (km s$^{-1}$)  & -680 $^{+20}_{-60}$ && -537 $\\pm$ 130 & & &-5800 $\\pm$ 1200 \\\\ $\\log n_4$ (cm$^{-3}$) & 11 && 6.8 -- 7.3 & & &6.8 -- 8.8   \\\\ $\\log r_4$ (cm) & 13.25 $^{+0.03}$ && 15.1 -- 15.4 & && 15.1 -- 16.1\\\\ $\\log$ N$_{H,4}$ (cm$^{-2}$)& 20.91 $^{+0.05}_{-0.06}$ & &21.42 $\\pm$ 0.12 & & &22.92 $\\pm$ 0.12\\\\ \\\\  \\enddata \\label{compare} \\tablecomments{This table shows the comparison between other absorption X-ray components found in the literature and the work presented in this paper. *\\cite{Krongold07} does not use $\\xi$ but instead uses the ionization parameter of U = $\\frac{Q(H)}{4\\pi r^2 n_e c}$, where Q(H) is the number of ionizing photons, r is the distance from the ionizing source, $n_e$ is the electron number density and c is the speed of light. To convert between the two, we assumed an ionizing luminosity of 10$^{42}$. ** Used Spectrum C for this comparison. } \\end{deluxetable*}     \\clearpage \\begin{figure*}[t] \\begin{center} \\includegraphics[height=7in,angle=270]{f1.ps} \\includegraphics[height=7in,angle=270]{f2.ps} \\includegraphics[height=7in,angle=270]{f3.ps} \\includegraphics[height=7in,angle=270]{f4.ps} \\includegraphics[height=7in,angle=270]{f5.ps} \\includegraphics[height=7in,angle=270]{f6.ps} \\includegraphics[height=7in,angle=270]{f7.ps} \\caption{\\footnotesize{This figure contains the ratio of the data to the fiducial model with $\\chi^2/\\nu$ 9607/7516. This model only includes a powerlaw and disk blackbody component as well as an effective H column density and a multiplicative constant, which corrects for the normalization between the HEG and MEG spectra. Clear absorption and emission features can be seen over this simple model. The wavelength plotted is in the observed reference frame. }} \\label{ratiorange} \\end{center} \\end{figure*}     \\begin{figure}[t] \\begin{center} \\includegraphics[height=3in,angle=270]{f8.ps} \\caption{\\footnotesize{These Ne X and Fe XXI lines are modeled with Gaussian profiles. Here one can see that the absorption is blue shifted with respect to the emission feature seen in the H-Like Lyman $\\alpha$ line of Ne X. The P-Cygni-like profile of the H-like line could be real, but may also be a superposition of distinctive wind components. The line markers are in the rest frame of NGC 4051. The wavelength plotted is in the observed reference frame. }} \\label{pcygNe} \\end{center} \\end{figure}  \\begin{figure}[t] \\begin{center} \\includegraphics[height=3in,angle=270]{f9.ps} \\caption{\\footnotesize{These OVII and OVIII lines are modeled with Gaussian profiles. Here one can see that the absorption is blue shifted with respect to the emission features seen in the H-Like Lyman series of O VIII. The P-Cygni-like profile of the H-like line could be real, but may also be a superposition of distinctive wind components. The line markers are in the rest frame of NGC 4051. The wavelength plotted is in the observed reference frame, and the data was rebinned with a minimum of 30 counts per bin for clarity. }} \\label{pcyg} \\end{center} \\end{figure}    \\begin{figure}[t] \\begin{center} \\includegraphics[height=3in,angle=270]{f10.ps} \\caption{\\footnotesize{This figure shows the OVII lines characterized with the phenomenological model for both absorption and emission. This He-like triplet gives a density estimate of n$_e\\approx$ 10$^{10}$ cm$^{-3}$. The wavelength plotted is in the observed reference frame, while the line identifications are in the rest frame of the host galaxy. }} \\label{OVII} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[height=3in,angle=270]{f11.ps} \\caption{\\footnotesize{ This figure shows the Ne IX line complex that is modeled with Gaussian lines for both the absorption and emission. This He-like triplet gives a density constraint of $n_e<$10$^{12}$ cm$^{-3}$.  Line identifications are in the rest frame of NGC 4051 (z= 0.002336).The wavelength plotted is in the observed reference frame. }} \\label{NeIX} \\end{center} \\end{figure}   \\begin{figure} \\begin{center} \\includegraphics[height=3in,angle=270]{f12.ps} \\caption{\\footnotesize{This figure shows the Fe XXII lines, which are modeled with Gaussian profiles. These lines were used as density diagnostics, constraining the density to be $n_e \\lesssim$ 7 $\\times$ 10$^{12}$ cm$^{-3}$.The wavelength plotted is in the observed reference frame, while the line identifications are in the rest frame of the galaxy. }} \\label{FeXXII} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=3in]{f13.eps} \\caption{\\footnotesize{This figure shows how the ratio of Fe XXI lines 12.38/12.285 change with increasing density. This plot was created using the CHIANTI populations at a temperature of 5 $\\times 10^5$ K.}} \\label{chiantife} \\end{center} \\end{figure}   \\begin{figure} \\begin{center} \\includegraphics[height=3in,angle=270]{f14.ps} \\caption{\\footnotesize{The above plot shows the fluorescence Fe K$\\alpha$ line in NGC 4051, as well as two narrow Fe K $\\alpha$ and $\\beta$ lines. The narrow lines likely originate in distant material, while the broad fluorescence line likely originates in the inner-most parts of the accretion disk. The red line denotes the relativistic fluorescence line characterized with the \\texttt{Kerrdisk} line profile. The blue line does not include this line. The narrow Fe K $\\alpha$ and $\\beta$ features are modeled with Gaussian lines. The line markers are drawn in the rest frame of NGC 4051. The wavelength plotted is in the observed reference frame. }} \\label{fefl} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[height=3in,angle=270]{f15.ps} \\caption{\\footnotesize{This is the O VIII Lyman $\\alpha$ fluorescence line modeled with a  \\texttt{Kerrdisk} component to describe any relativistic O VIII emission. The blue line is model with no relativistic line but includes the XSTAR model. The green line includes a relativistic line with all of its parameters to vary ( spin, inclination angle, inner radius, and normalization). The red line denotes the model in which the O VIII parameters were tied to that of the Fe K$\\alpha$ line. The line is suggested at a 2 $\\sigma$ level if the parameters are free to vary or tied to that of the Fe K$\\alpha$ line. The wavelength plotted is in the observed reference frame, and the data were rebinned with a minimum of 25 counts per bin for clarity. }} \\label{okerr} \\end{center} \\end{figure}   \\begin{figure*}[t] \\begin{center} \\includegraphics[height=7in,angle=270]{f16.ps} \\includegraphics[height=7in,angle=270]{f17.ps} \\includegraphics[height=7in,angle=270]{f18.ps} \\includegraphics[height=7in,angle=270]{f19.ps} \\includegraphics[height=7in,angle=270]{f20.ps} \\includegraphics[height=7in,angle=270]{f21.ps} \\includegraphics[height=7in,angle=270]{f22.ps} \\caption{\\footnotesize{This figure shows the HETG spectra and the best-fit XSTAR model components.  The parameters for this model are given in Table \\ref{xstar}, and the model has $\\chi^2 / \\nu$ = 8177/7516. The absorption features in the residuals as given in Figure \\ref{ratiorange} are well characterized by the XSTAR model, especially in the ranges of 12--18\\AA\\/. The wavelength plotted is in the observed reference frame. }} \\label{xstarrange} \\end{center} \\end{figure*} \\begin{figure*}[t] \\begin{center} \\includegraphics[height=7in,angle=270]{f23.ps} \\includegraphics[height=7in,angle=270]{f24.ps} \\includegraphics[height=7in,angle=270]{f25.ps} \\includegraphics[height=7in,angle=270]{f26.ps} \\includegraphics[height=7in,angle=270]{f27.ps} \\includegraphics[height=7in,angle=270]{f28.ps} \\includegraphics[height=7in,angle=270]{f29.ps} \\caption{\\footnotesize{The above figure plots the HETG spectra and the best-fit Cloudy model. The parameters for this model are given in Table \\ref{cloudy} and the model has $\\chi^2 / \\nu$ = 8713/7516. The wavelength plotted is in the observed reference frame. }} \\label{cloudyrange} \\end{center} \\end{figure*}"
    },
    "1112/1112.1044.txt": {
        "abstract": "The broadband analog-digital heliograph based on the UTR-2 radio telescope is described in detail. This device operates by employing the parallel-series principle when five equi-spaced array pattern beams which scan the given radio source (e.g. solar corona) are simultaneously shaped. As a result, the obtained image presents a frame of 5 $\\times$ 8 pixels with the space resolution 25$\\,'\\times$ 25$\\,'$ at 25 MHz. Each pixel corresponds to the signal from the appropriate pattern beam. The most essential heliograph component is its phase shift module for fast sky scanning by pencil-shape antenna beams. Its design, as well as its switched cable lengths calculation procedure, are presented, too. Each heliogram is formed in the actual heliograph just by using this phase shifter.  Every pixel of a signal received from the corresponding antenna pattern beam is the cross-correlation dynamic spectrum (time-frequency-intensity) measured in real time with the digital spectrum processor. This new generation heliograph gives the solar corona images in the frequency range 8-32 MHz with the frequency resolution 4 kHz, time resolution to 1 ms, and dynamic range about 90 dB. The heliographic observations of radio sources and solar corona made in summer of 2010 are demonstrated as examples. ",
        "introduction": "Comprehensive information on the physical processes that accompany solar activity can be obtained only with engaging of an extensive scope of observation means. At the meter wavelengths, radio spectrographs, heliographs and polarimeters are widely used to meet this purpose \\cite{kruger}. A worse situation is the case for decameter wavelengths. In particular, the heliograph realization is technically a rather labor intensive problem. Obtaining the high-resolution images (being one of the major demands in application of heliographs) at that low frequencies requires the huge size antennas ($>$ 1 km). At the same time, the heliograph can help obtain new useful information on the extremely multifarious types of bursts being rather specific to this range of frequencies.  Any heliograph has essential advantages in the case it ensures the possibility of simultaneous observations at several frequencies. Multi-frequency measurements of the positions of sources of Type II and III bursts allow to immediately determine the velocity of shock waves and this of electron beams which are propagated in the solar corona. Of substantial interest are the measurements of the velocity of electron beams generating Type III and IIIb bursts at different trajectory phase. The positional measurements at several frequencies will probably permit evaluation of the magnitude of visible displacement of burst sources with respect to their true position caused by the refraction and scattering of radio waves in corona.  Thus, the heliograph permits observation of an angular structure of burst sources and its evolution during their lifetime, as well as to identify the bursts with the appropriate activity regions and measure their heights above the photosphere. The position measurements of double Type IIIb-III, III-III bursts and drifting pairs allow to find out whether the both components are excited in the same place of the corona and at which plasma frequency. Of special interest are the investigations of properties and determination of the positions of coronal mass ejections, frequently associated with the Type II bursts. This phenomenon is considered a key one in the problems of solar-terrestrial relationships and space weather.  No less interesting direction of the research efforts are observations of the two-dimensional brightness distributions across the quiet Sun at the decameter waves. At the same time, as is known, just at the heights where the decameter radiation is generated there exists the most interesting corona region where the solar wind originates. The here-presented far from exhaustive list of problems capable of solving with the heliograph makes its construction a rather attractive idea. Especially, in case if there is an opportunity of using the available antenna system suitable for this purpose.  To date, there exist not too many radio telescopes permitting to obtain 2D heliograms of the Sun and other sky radiation sources. In the microwave range, they are represented by the Nobeyama heliograph (Japan) operating at two frequencies, 17 and 35 GHz \\cite{Nishio98}, and the RATAN-600 based heliograph (Russia) operating at 3.75 GHz \\cite{Opeikina95}. Roughly within these frequency ranges, the observations are also made at the Siberian Solar Radio Telescope (Russia) \\cite{Grechnev03} at 5.7 GHz, and at the Owens Valley Radio Observatory interferometer (USA) within 1-18 GHz, whose whole range being split into 86 individual subbands \\cite{Padin91}. The anticipated Chinese heliograph being designed for radio measurements within 0.4-15 GHz \\cite{Yi-Hua06} can be considered belonging under the promising instruments, too. For centimeter and meter wavelength radio astronomy investigation, used are the Nancay radioheliograph (France) \\cite{NRHG93}, operating only at frequencies 169, 327 and 408 MHz, and the Gauribidanur radioheliograph (India) which surveys at some individual frequencies within 40-150 MHz range \\cite{Ramesh98}. Since 1969 till 1984, the Culgoora radioheliograph (Australia) \\cite{Sheridan83} was functioning at frequencies 43, 80, 160 and 327 MHz. Note that very few instruments were used at the decameter wavelengths. In this respect we may mention the Clark Lake Teepee-Tee radio telescope (USA) \\cite{Erickson82} at which the measurements at some individual frequencies within  15-125 MHz were made. Though for now, it is already put out of service. It will be observed also that developing the antenna arrays capable of operating low frequencies (decameter wavelengths) is connected with heavy technological and methodological difficulties. This frequency range is subject to very heavy noise conditions, the ionosphere exerts an essential impact on the behavior of decameter wave propagation, and achieving the narrow beam pattern requires antennas of huge areas. For instance, the Nancay Decameter Array (Nancay, France) \\cite{Lecacheux2000} is effectively used for getting dynamic spectra of solar radio emission, while because of its small sizes practically never used for heliographic observations. On account of these reasons, nowadays only a few decameter range antenna systems exist capable of performing heliographic observations, though they also have limited capacities, and thus need upgrading, or even new antenna systems should be built. Suffice it to mention the promising projects LOFAR \\cite{white,Reich07} and LWA \\cite{Kassim05} which are to be put into service in the near future. Therefore the researches carried out with the UTR-2 based heliograph are of great scientific interest, making up for a deficiency in the knowledge of physical processes on the Sun and/or in the circumsolar space. Note that until now, the UTR-2 radio telescope remains the world-largest and most effective instrument operating within 8-32 MHz, and it is expected to remain that kind performer in long-term future. Accordingly, its intensive design improvement, instrumentation and usage (this including in heliographic mode, too) promises a vast deal of useful astrophysical information.  The high solar activity is frequently accompanied by sporadic radio emission in the meter and decameter wavelength ranges \\cite{kruger,pick2008}. Some of its types (I, II, III and IV) are common for these ranges. At the same time, they have their inherent specificity. For instance, the Type I  bursts, most frequent at meter waves, are practically not observed at the decameter wavelengths. On the other hand, the short-lived stria-bursts, drifting pairs, Type IIIb bursts are met only at frequencies below 60 MHz, though the latter, along with Type III bursts, are the most numerous events of the solar decameter radiation \\cite{Abranin78}.  The solar decameter sporadic radiation has begun to be systematically studied with the work of Ellis and McCulloch \\cite{Ellis67}. Since then, a substantial progress was reached in understanding its nature. Meanwhile, the theoretical formulation of the mechanism and conditions of solar burst generation cannot yet give, in most cases, an adequate representation. Therefore further analysis of new experimental data, including those observed with the heliograph, is of current importance.  The present paper is devoted to the UTR-2 heliograph construction, its general functional scheme and discussion of its most important features. In Section~\\ref{par2} we start with the history of heliograph studies in our institute. The basis of any radioheliograph, the UTR-2 based heliograph making no exception, is the antenna system. In Section~\\ref{par3}, this latter brief description as applied to the heliograph design, as well as the detailed analysis of the modern configuration of a two-dimensional decameter wavelength heliograph based on the UTR-2 T-shaped antenna system (in Sections~\\ref{par4} and \\ref{par5}), are presented. Next, Section~\\ref{par6} is aimed at design description of the most important UTR-2 based heliograph element, namely the phase shifter for fast beam scanning. The phase shifter essential parts are the switched coaxial delay-line cables. In Section~\\ref{par7} the method is described for calculating the lengths of these switched cables. The heliograph control unit helps to generate the pulse train for switching the fast scanning phase shifter and scan markers generation (see Section~\\ref{par9}). Consequently, the UTR-2 array pattern beam is scanning the desired sky area, and the heliograph scan sector format is discussed in Section~\\ref{par8}. Further, in Section~\\ref{par10} we consider the operation features of the heliograph receiver-recorder. It is based on a multichannel digital spectropolarimeter (DSP). In Section~\\ref{par11} we present the preliminary results of test observations made with the considered instrument in summer of 2010. In conclusions, we summarize the heliograph development and prospects for its future application in radio astronomy observations at decameter wavelengths.  \\section[Brief History of Heliograph Studies in the UTR-2] {Brief History of Heliograph Studies in the UTR-2 Observatory}\\label{par2} The first positional observations of solar radiation with the use of a two-dimensional heliograph based on the UTR-2 radio telescope were made during July 31 to August 11, 1976 \\cite{Abranin80}. At that time, the enhanced radiation was associated with emerging of the active McMath No. 14352 region on the eastern limb which, while moving westwards, was intersecting the entire solar disk. For the aforesaid period of observations, numerous Type III and IIIb bursts, Type IIIb-III bursts, as well as short-lived narrow-band stria-bursts being components of the Type IIIb burst chains were recorded. The daily two-hour measurements centered for the local midday were made at the operating frequency 25 MHz, and also in parallel by the ``North-South'' (N-S) antenna array to which output the dynamic spectrograph for the 23.5-25.5 MHz band was mounted. The recorded heliograms permitted constructing the histograms of coordinate distribution of effective centers of Type III and stria-bursts for hour angle ($h$) and declination ($\\delta$) in the selected days for different positions of solar disk active regions. The positional data on the average coordinates of burst sources centers were used to obtain the day-to-day dependence of the radiating region position. The motion velocity of the active region, crossing the solar disk westwards, permitted estimates of the average heights $R_s$ (being radial distances from the Sun's center) of the Type III burst and stria-burst sources during the radio burst. The heights $R_s$ were determined by several methods which yielded close results. In the meter and centimeter wave-lengths, the radial distances of the coronal emitting regions were calculated in papers \\cite{Malinge60,Clavelier68}. The $R_s$ values estimated with the UTR-2 heliographic observations have appeared comparable to each other and corresponded to those first obtained at decameters during the interference measurements at the Clark Lake Radio Observatory (CA, USA) at 30 MHz while investigating a number of solar noise storms \\cite{Gergely75}. The heliograms have permitted the conclusion that the stria-bursts of IIIb chains and Type III bursts appear about the same place on the Sun's image plane.  Further research efforts, assisted by the UTR-2 antenna system operating the mode of a two-dimensional heliograph, were largely concentrated on studying the Type IIId radiobursts with echo components \\cite{Abranin82,Abranin84,Abranin96,Abranin97,Abranin98,Abranin99,Abranin00}. The decameter Type IIId bursts were first recorded at the UTR-2 observations of July 6, 1973. Narrow-band elements of this emission fine structure -- the diffuse stria-bursts -- are a variety of ordinary short-lived stria-bursts, of which the dynamic spectra of Type IIIb bursts are patterned. The principal feature of diffuse stria-bursts consists in that their form largely depends on heliolongitude of the coronal sources region. When the Type IIId source is in the near-limb region, the stria-bursts with steep fronts and sharp peaks are observed monotonically damping in a few seconds. With the active region approach to the central solar meridian, the time splitting occurs, i.e. the echo-like component of bursts with growing delay appears. The echo-burst delay time becomes maximum when the emitting region traverses the central meridian. The appearance of an extra burst echo-component and the delay variation with heliolongitude are qualitatively explained by using a simple model of a point pulsed source placed into a spherically symmetric corona and emitting at the plasma frequency second harmonic \\cite{Riddle74}. Using this kind model in interpreting the origin of Type IIId bursts was suggested in papers \\cite{Abranin82,Abranin84}. However, the theoretically calculated values were poorly coordinated with those UTR-2 measured. The assumption \\cite{Abranin00} was therefore made that the echo-component of stria-bursts is formed not merely due to the reflection from a deeper layer of a spherically symmetric corona, but also as a result of the refraction of radio waves on some large-scale coronal structures. This hypothesis has been supported in \\cite{Afanasiev06} where the author has mathematically simulated the shaping of Type IIId bursts. During the systematic researches of  the  Sun with the UTR-2 array system and positional observations with the two-dimensional heliograph for almost two decades, more than 10 solar Type IIId burst storms were recorded. It will be observed that Type IIId radio bursts are rather rare events, unlike Type III bursts being most mass events of the decameter sporadic solar radiation. With the complex observations using the two-dimensional heliograph, in June 1984 during a week-long Type III storm, about 1000 bursts were recorded at 25 MHz. The data obtained allowed the statistical sample manipulation of decameter events by employing the cluster analysis \\cite{Stepanova95}. For the statistical procedure of sample ordering and clustering, used were such parameters of bursts as time delay between the two sequent bursts, burst du\u00adration at half-power level, maximum intensity and degree of circular polarization.  It should be recalled that way back in the mid-nineties, the radioastronomy observations with the UTR-2 array were made only near to several fixed frequencies: 10.0, 12.5, 14.7, 16.7, 20.0 and 24.8 MHz, while the broadbandness of this antenna system is much greater. On the other hand, in the capacity of a recorder, e.g. for the heliograph, the FTAP-2 facsimile set was used recording the signal on a electrochemical paper with the dynamic range making only 10-16 dB. Such a situation urgently demanded extensive design improvement of the whole UTR-2 hardware facility. When the UTR-2 antenna amplifier system had been upgraded in the late nineties, the measurements became possible in the continuous frequency band of 10-30 MHz and more \\cite{Konovalenko97}. Thereupon, an exigency of installation of appropriate recording equipment arose. With that end in view, for multi-frequency observations a 60-channel spectrum analyzer was put into service which included 60 separate frequency-tunable within 10-30 MHz radio receivers with the 4-10 kHz bandwidth. The next step was applications of digital multichannel (1000 and more channels) spectrum analyzers (digital receivers). These latter already allowed covering the entire UTR-2 array frequency range with the frequency separation of about 4 kHz, as well as recording on digital storage media and perform varied signal computer processing \\cite{Kleewein97,Lecacheux98,Lecacheux2004}.  The very first configuration of the UTR-2 array based heliograph was as follows. By means of a special beam control system of the UTR-2 array it became possible to quickly scan the sky area under investigation. Its image was recorded through signal injection from the outputs of N-S and E-W (``East-West'') antennas into the sum-difference device. The sum and difference signals were amplified by individual receivers with square-law detectors. The outputs of square-law detectors were in opposition, thus providing the difference of their output voltages. This difference voltage then fed the final recorder input. Actually, the radio astronomy research efforts are basically carried out using the digital receiver-recorder, being a digital spectropolarimeter (DSP) capable of signal receiving and processing, and data recording on the PC hard disk. Surprisingly, we still face the fact that though the UTR-2 based heliograph was designed long ago and used in diverse radioastronomy researches, its complete and thorough description had been postponed because of ever-lasting upgrading its hardware environment. For this long period, the heliograph has been essentially reengineered, new possibilities for its application appeared, and things have reached an exigent necessity point to describe the heliograph development at long last review. ",
        "conclusions": "\\label{par12} This paper considers a functional scheme of the decameter wavelength heliograph. The heliograph is based on the UTR-2 antenna system. Any radio source is slowly tracked by the UTR-2 radio telescope standard hardware facility. Within the heliograph field of view, this antenna pattern beam is guided by a specially designed phase shifter. The heliograph is recording the image with the sequential signal injection in each position of the pencil beam of the UTR-2 array out of its 40 positions available in a separate frame of the heliograph. The heliograph image repetition rate can be taken both rather large (suitable for analyzing the spatial characteristics of fast solar bursts), and a small one (convenient for the quiet Sun upper corona research).  The image repetition rate, owing to the motion of a five-beam set of the UTR-2 array within the heliograph scanned area, should be much greater than the Sun tracking rate. The UTR-2 phasing system cannot ensure the beam position change with such a speed since it uses the electromagnetic relays as the switched elements. That is why the UTR-2 phasing system is used only for the Sun position tracking. For fast scanning the scan sector, the UTR-2 phasing system comprises additional phase shifters possessing a fairly high operating speed and long lifetime. As the scanning is performed within small angles, the optimum alternative will be enabling the additional phase shifters between the section outputs and appropriate inputs of the E-W antenna sections phasing system. In this case, at fast scanning, the section pattern position remains the same, and the array beam moves within the section beam pattern width owing to the additional phase shifters. This results in lower received-signal intensity of the source situated near the scan sector edge due to signal strength depression on the section pattern slopes. In recorded-data processing this fact should be taken into account, and the respective corrections made in processing the observation results.  The design features of the decameter wavelength heliograph based on the UTR-2 array are such that we may belong it to the unique radio astronomy instruments. Here, suffice it to mention that among all heliographs used world-wide, this one operates the lowest frequencies for the study of cosmic radio emission. Secondly, it is superbroad-band covering practically the entire range within 10 to 30 MHz. Thirdly, its sensitivity, time and frequency resolutions, dynamic range, signal space selection are record-breaking for the given frequency range, and consequently the problems which can be solved with this instrument have no world counterparts. The heliograph is also unique in its capability of obtaining quiet Sun's corona images for about 1 s period at distances from two to three solar radii, counted from the Sun's center, and practically without skipping in the $(U, V)$-plane. The heliograph's sensitivity allows observation of space-time evolution of low-contrast formations (fine structure of radio bursts, fibrous structure of Type II bursts, etc.). Besides, its dynamic range makes it possible to observe the details simultaneously with the bright radio events (strong Type III bursts, or against the background of Type II and IV bursts). The instrument has undergone a long-continued development, and in retrospect we may think with absolute certainty that the state-of-the-art heliograph performances have justified the expended energies."
    },
    "1112/1112.0535_arXiv.txt": {
        "abstract": "High frequency quasi-periodic oscillations (QPOs) from weakly magnetized neutron stars display rapid frequency variability (second timescales) and high coherence with quality factors up to at least 200 at frequencies about 800-850 Hz. Their parameters have been estimated so far from standard \\minchi2\\ fitting techniques, after combining a large number of Power Density Spectra (PDS), as to have the powers normally distributed (the so-called Gaussian regime). Before combining PDS, different methods to minimize the effects of the frequency drift to the estimates of the QPO parameters have been proposed, but none of them relied on fitting the individual PDS. Accounting for the statistical properties of PDS, we apply a  maximum likelihood method to derive the QPO parameters in the non Gaussian regime. The method presented is general, easy to implement and can be applied to fitting individual PDS, several PDS simultaneously or their average, and is obviously not specific to the analysis of kHz QPO data. It applies to the analysis of any PDS optimized in frequency resolution and for low frequency variability or PDS containing features whose parameters vary on short timescales, as is the case for kHz QPOs. It is equivalent to the standard $\\chi^2$ minimization fitting when the number of PDS fitted is large. The accuracy, reliability and superiority of the method is demonstrated with simulations of synthetic PDS, containing Lorentzian QPOs of known parameters. Accounting for the broadening of the QPO profile, due to the leakage of power inherent to windowed Fourier transforms, the maximum likelihood estimates of the QPO parameters are asymptotically unbiased, and have negligible bias when the QPO is reasonably well detected. By contrast, we show that the standard \\minchi2\\ fitting method gives biased parameters with larger uncertainties.  The maximum likelihood fitting method is applied to a subset of archival Rossi X-ray Timing Explorer (RXTE) data of the neutron star X-ray binary \\4u, for which, we show that the lower kHz QPO parameters can be measured on timescales as short as 8 seconds. As to demonstrate the potential use of the results of the maximum likelihood method, we show that in the observation analyzed the time evolution of the frequency is consistent with a random walk. We then show that the broadening of the QPO due to the frequency drift scales as $\\sqrt{T}$, as expected from a random walk ($T$ is the integration time of the PDS). This enables us to estimate the intrinsic quality factor of the QPO to be $\\sim 260$, whereas previous analysis indicated a maximum value around 200. ",
        "introduction": "The standard method of weighted least squares (i.e. minimum $\\chi^2$, hereafter \\minchi2) model fitting is only equivalent to maximum likelihood estimation (MLE) of the model parameters when the data to be fitted are normally (Gaussian) distributed. As is well known, the distribution of $M$ averaged periodogram ordinates (from a stationary, linear stochastic process) follows a $\\chi^2_{2M}$ distribution with $2M$ degrees of freedom; as $M$ increases, the $\\chi^2_{2M}$ tends towards a normal distribution \\citep{groth75, papadakis93}. In X-ray timing the most commonly used method to reach the so-called Gaussian regime is Bartlett's method \\citep{bartlett48}: segment the original time series into $M$ non-overlapping segments, compute a periodogram for each segment, and average over all segments to produce a spectral estimate to be fitted \\citep{vdk89}. Typically, one uses $M > 50$ to produce approximately Gaussian distributed averages \\citep{papadakis93}. The drawbacks of this method are a loss of frequency resolution by a factor $M$ (reducing the sensitivity for narrow feature detection) and a suppression of the lowest frequencies. There is also a major drawback in the case of the analysis of high frequency quasi-periodic oscillations (QPOs), because their frequency varies rapidly with time, typically on timescales of seconds \\citep{barret05b}. Since a kHz QPO can be a relatively narrow feature (FWHM $\\sim 2-4$ Hz), sub-Hz frequency resolution of the PDS (equivalent to a segment duration of a few seconds) is required to have the QPO profile properly sampled. In order to use \\minchi2\\ fitting one must average a large number of segments, $M$, which then requires a total integration time exceeding hundreds, even thousands of seconds, depending of the strength of the QPO. On such timescale, the drift of the QPO frequency will smear out the QPO profile, leading to biases in the fitted QPO parameters, in particular the width. Methods to minimize the contribution of the frequency drift to the fitted QPO width have been proposed, using for instance the centroid (without fitting) of the excess power, as a best guess for the QPO frequency \\citep{barret05b,barret06}. Although optimized, those methods suffer from the fact that statistical fluctuations of the noise can be stronger than the QPO signal, hence biasing the QPO frequency determination, especially when $M$ is small.  In this paper, we apply the maximum likelihood method to fit spectral models of high frequency QPOs in the non-Gaussian regime. To our knowledge, this is the first time such a method is applied to the analysis of {\\it RXTE} PCA data, although it was discussed early on by \\citet{stella94} for the analysis of continuum PDS recorded with {\\it EXOSAT}. Such a method has been used already in a variety of applications, such as the modeling of solar oscillation spectra \\citep{anderson90,appourchaux11} or fitting broad band PDS from long {\\it XMM-Newton} observations of variable Seyfert 1 galaxies \\citep{v10}. In section \\ref{sec2}, we describe the Maximum Likelihood method. Then in section \\ref{sec3}, the reliability and robustness of the technique is demonstrated using simulations of synthetic PDS. Its superiority over \\minchi2\\ fitting is illustrated in the case of $M=4$, and a comparison between the MLE and the \\minchi2\\ fitting is also presented for typical $M$ values used in QPO analysis. In section \\ref{sec4}, the method is applied to a small subset of the archival {\\it RXTE} data of the neutron star X-ray transient \\4u, and the results of the MLE are used to infer the intrinsic width of the QPO. ",
        "conclusions": "We have demonstrated the application of maximum likelihood fitting to periodogram data showing high frequency QPOs. The method is simple to implement and, importantly, remains valid in the non-Gaussian regime, e.g. when applied to raw periodogram data. This means it can be used in situations where averaging many periodograms is not desirable, e.g. when short time series are to be analysed at the highest frequency resolution. The simulation tests discussed in section $3$ show the method gives accurate and unbiased estimates of the QPO parameters, and can be used to probe QPO properties on timescales as short as $8$ seconds (for strong kHz QPOs). By contrast the standard \\minchi2\\ fitting method gives biased parameters with larger uncertainties when applied to such short sections of data because the spectral estimates are not Gaussian distributed.  We have applied the MLE to an {\\it RXTE} observation of the $\\sim 830$ Hz QPO of the neutron star low-mass X-ray binary \\4u. For one section of the observation we are able to recover the QPO frequency, width and amplitude on timescales down to $8$ seconds. We have shown that the frequency variations are consistent with a random walk in the observation analyzed, and we have inferred the intrinsic width of the QPO to correspond to a quality factor of $\\sim 260$ at 830 Hz: a value significantly higher than previous estimates \\citep{barret05b}. Of course, MLE fitting of time-resolved periodograms is not the only approach available for investigating the variability of the kHz QPOs. Another is to fit time-dependent models directly to the time-frequency data using hierarchical models \\citep{gelman07}. This is beyond the scope of this paper whose aim is to highlight the improvement (in e.g. bias and variance of the fitted QPO parameters, and in some cases time resolution) that can be gained by using Whittle MLE rather than the $\\min(\\chi^2)$ method that is currently the standard procedure.  The method described here was applied for the first time to {\\it RXTE} data. By enabling to measure the QPO parameters on very short timescales with negligible biases, this technique opens the way for new types of analysis to be performed on {\\it RXTE} archive, as we have shown in the case of \\4u\\ for a limited subset of data. In most analysis so far, the QPO properties have been derived from averaging data sets, hence leading to a potential loss of information. The MLE does not require averaging. Hopefully, this will lead to a better understanding of the high frequency variability observed in accreting compact objects, both neutron stars and black holes. This is required to fully exploit the potential of this variability as a probe of strong gravity and dense matter \\citep{klis00}.  \\appendix"
    },
    "1112/1112.2706_arXiv.txt": {
        "abstract": "Although galaxies, groups, and clusters contain $\\sim10$\\% of the baryons, many more reside in the photoionized and shocked-heated intergalactic medium (IGM) and in the circumgalactic medium (CGM).  We update the baryon census in the (\\HI)  \\Lya\\ forest and warm-hot IGM (WHIM) at $10^{5-6}$~K traced by \\OVI\\ $\\lambda1032,1038$ absorption.  Using robust cosmological simulations of heating, cooling, and metal transport, we improve the \\OVI\\ baryon surveys with spatially averaged corrections for metallicity ($Z / Z_{\\odot}$) and \\OVI\\ ionization fraction ($f_{\\rm OVI}$).  Statistically, their product correlates with column density, $(Z / Z_{\\odot}) \\, f_{\\rm OVI}\\ \\approx (0.015)(N_{\\rm OVI} /10^{14}\\,{\\rm cm}^{-2})^{0.70}$, with a \\NOVI-weighted mean of 0.01, which doubles previous estimates of WHIM baryon content.  We also update the \\Lya\\ forest contribution to baryon density out to $z = 0.4$, correcting for the $(1+z)^3$ increase in absorber density,  the $(1+z)^{4.4}$ rise in photoionizing background, and cosmological proper length $d\\ell/dz$. We find substantial baryon fractions in the photoionized \\Lya\\ forest  ($28 \\pm 11$\\%) and WHIM traced by \\OVI\\ and broad-\\Lya\\ absorbers $(25 \\pm 8$\\%).  The collapsed phase (galaxies, groups, clusters, CGM) contains $18\\pm4$\\%, leaving an apparent baryon shortfall of $29\\pm13$\\%.  Our simulations suggest that $\\sim15$\\% reside in hotter WHIM ($T \\geq 10^6$~K).  Additional baryons could be detected in weaker \\Lya\\ and \\OVI\\ absorbers. Further progress requires higher-precision baryon surveys of weak absorbers, down to minimum column densities $N_{\\rm HI} \\geq  10^{12.0}$~\\cd,  $N_{\\rm OVI} \\geq  10^{12.5}$~\\cd, $N_{\\rm OVII} \\geq 10^{14.5}$~\\cd, using high-S/N data from high-resolution UV and X-ray spectrographs. ",
        "introduction": "For low-redshift cosmology and galaxy formation rates, it is important to account for all the baryons synthesized in the Big Bang.   Cosmologists have noted a baryon deficit  in the low-redshift universe (Fukugita, Hogan, \\& Peebles 1998) relative to the predicted density synthesized in the Big Bang.  Although this deficit could arise from an incomplete inventory, it could also challenge our understanding of the thermodynamics of structure formation and the response of the gas to accretion shocks and galactic outflows.  Recent analysis (Komatsu \\etal\\ 2011) of the spectrum of acoustic peaks in the Cosmic Microwave Background (CMB) obtained by the {\\it Wilkinson Microwave Anisotropy Probe} (WMAP) found that baryons comprise a fraction $\\Omega_b = 0.0455 \\pm 0.0028$ of the critical matter-energy density of the universe, $\\rho_{\\rm cr} = (9.205 \\times 10^{-30}~{\\rm g~cm}^{-3}) h_{70}^2$, where $h_{70}$ is the Hubble constant ($H_0$) in units of 70 km~s$^{-1}$~Mpc$^{-1}$.  This 4.6\\% baryon fraction corresponds to a mean comoving density, $\\overline{\\rho}_b = \\Omega_b \\rho_{\\rm cr} = 4.24 \\times 10^{-31}$ g~cm$^{-3}$, and a hydrogen number density $n_H = 1.90 \\times 10^{-7}~{\\rm cm}^{-3}$, assuming a primordial helium mass fraction $Y_p = 0.2477$ (Peimbert \\etal\\  2007).  Baryon inventories in the low redshift universe are uncertain.  They are often complicated by the formation of galaxies and large-scale structure and by the feedback from star formation in the form of ionizing radiation, metals, and outflows. Galaxy surveys have found $\\sim$10\\% of these baryons in collapsed objects such as galaxies, groups, and clusters (Salucci \\& Persic 1999; Bristow \\& Phillipps 1994; Fukugita \\& Peebles 2004).   Over the last 15 years, substantial reservoirs of gas have been found in the intergalactic medium (IGM), in the halos of galaxies, and in the circumgalactic medium (CGM). There are often semantic problems in defining the CGM,  as gas blown out of galaxies or material within the virial radius (Tumlinson \\etal\\ 2011; Prochaska \\etal\\ 2011).  Of the remaining 80-90\\% of cosmological baryons, approximately half can be accounted for in the low-$z$ IGM (Shull 2003;  Bregman 2007;  Danforth \\& Shull 2008;  Danforth 2009) including the warm-hot IGM (or WHIM).   Ultraviolet spectroscopic surveys of \\Lya\\ and \\OVI\\ have identified substantial numbers of absorbers (Danforth \\& Shull 2008;  Tripp \\etal\\ 2008;  Thom \\& Chen 2008), but claimed detections of hotter  in X-ray absorption by \\OVII\\  (Nicastro \\etal\\ 2005) remain controversial (Kaastra \\etal\\ 2006; Yao \\etal\\ 2012).  Unfortunately, X-ray spectra still have not confirmed the potential large reservoir of baryons at $T > 10^6$~K, suggested by cosmological simulations.  An inefficient distribution of collapsed baryons vs.\\ distributed matter is a prediction of nearly all cosmological simulations of large-scale structure formation (Cen \\& Ostriker 1999, 2006;  Dav\\'e \\etal\\ 1999, 2001; Smith \\etal\\ 2011; Tepper-Garcia \\etal\\ 2011).   These N-body hydrodynamical simulations suggest that 10--20\\% of the baryons reside in collapsed objects and dense filaments, with the remaining 80\\%  distributed over a wide range of phases in baryon overdensity ($\\Delta_b = \\rho_b / \\overline{\\rho}_b)$ and temperature ($T$).   Thermodynamic considerations suggest that shock-heated WHIM at $z < 1$ is a natural consequence of gravitational instability in a dark-matter dominated universe.  This hot gas is augmented by galactic-wind shocks and virialization in galaxy halos.  In this paper, we improve the analysis of baryon content in both photoionized and shocked-heated IGM phases and assess the accuracy of observational and theoretical estimates.  Because \\Lya\\ absorption surveys probe the neutral component of diffuse, photoionized filaments, assessing the baryon content requires a correction for the neutral fraction, $f_{\\rm HI}$.   We derive more accurate photoionization corrections based on the $(1+z)^{4.4}$ rise in the metagalactic ionizing background out to $z \\approx 0.7$, as well as the $(1+z)^3$ increase in absorber density and changes in cosmological proper length, $d\\ell/dz$.   We also use cosmological simulations to correct the \\OVI\\ surveys for ionization fraction ($f_{\\rm OVI}$) and metallicity  $(Z / Z_{\\odot})$.  Finally, we use our group's critically evaluated catalog of \\HI\\ and \\OVI\\ absorption lines  (Tilton \\etal\\ 2012) from our  \\HST\\ Legacy Archive project on IGM data taken by the {\\it Hubble Space Telescope} (\\HST) and {\\it Far Ultravolet Spectroscopic Explorer} (\\FUSE).  To constrain the distribution of $Z$ and $f_{\\rm OVI}$, we use cosmological hydrodynamic simulations of IGM heating, cooling, and metal transport to find the column-density weighted average for their product,  $(Z / Z_{\\odot}) \\, f_{\\rm OVI}$. Whereas previous work assumed constant values, $Z/Z_{\\odot} = 0.1$ and $f_{\\rm OVI} =0.2$, our new simulations provide spatially averaged corrections for metallicity, \\OVI\\ ionization fraction, and covariance of  IGM parameters ($T$, $\\rho$, $Z$). Our computed mean value $\\langle (Z / Z_{\\odot}) \\, f_{\\rm OVI} \\rangle = 0.01$ is half the previously assumed value, thereby doubling previous estimates of the baryon census traced by \\OVI.   In Section 2 we describe the simulations and their results for the \\OVI\\ distribution in column density, gas temperature, baryon overdensity, metallicity, and ionization fraction.  In Section 3, we assess the corrections for ionization and metallicity, applied to the \\OVI\\ and \\Lya\\  absorbers, and we derive values of $\\Omega_b$ from recent surveys of IGM and CGM phases.   Variations in these factors are produced by WHIM thermodynamics, gas temperature ($T$), baryon overdensity ($\\Delta_b$), and \\OVI\\ column density ($N_{\\rm OVI}$).   In Section 4 and two Appendices, we summarize the current baryon census with uncertainties on each component. ",
        "conclusions": ""
    },
    "1112/1112.5898_arXiv.txt": {
        "abstract": "Mid-infrared images from the {\\it Spitzer Space Telescope} Galactic Legacy Infrared MidPlane Survey Extraordinaire program reveal that the infrared source IRAS~03063+5735 is a bowshock nebula produced by an early B star, 2MASS~03101044+5747035.  We present new optical spectra of this star, classify it as a B1.5~V, and determine a probable association with a molecular cloud complex at $V_{LSR}$=$-38$ -- $-42$ \\kms\\ in the outer Galaxy near $\\ell=140.59^\\circ$, $b=-0.250$\\degr. On the basis of spectroscopic parallax, we estimate a distance of 4.0$\\pm$1~kpc to both the bowshock nebula and the molecular complex. One plausible scenario is that this a high-velocity runaway star impinging upon a molecular cloud.  We identify the \\HII\\ region and stellar cluster associated with IRAS~03064+5638 at a projected distance of 64 pc as one plausible birth site.  The spectrophotometric distance and linkage to a molecular feature provides another piece of data helping to secure the ill-determined rotation curve in the outer Galaxy.  As a by-product of spectral typing this star, we present empirical spectral diagnostic diagrams suitable for approximate spectral classification of O and B stars using He lines in the little-used yellow-red portion of the optical spectrum. ",
        "introduction": "Short-lived massive stars typically lie within the OB associations in which they were born. However, a significant fraction (17\\%--50\\%) having high (30--200 \\kms) space velocities wander far from such associations and are known as ``runaways'' \\citep{blaauw, stone82, conti77, gies86}.  While many runaways are detected  directly on the basis of large radial velocities or high proper motions, others are inferred by the presence of a tell-tale bowshock produced by their supersonic motion through the surrounding medium.   \\citet{vanb88}, \\citet{vanb95}, and \\citet{noriega97} used \\emph{Infrared Astronomical Satellite} (\\emph{IRAS}) images to locate arc-shaped bowshock features associated with high-velocity O stars.   Bowshocks associated with high-velocity O stars have been detected more easily using the recent generation of mid-infrared surveys with higher angular resolution performed by the \\emph{Midcourse Space Experiment} (\\emph{MSX}) and the \\emph{Spitzer Space Telescope} (\\emph{SST}). Recent studies of bowshocks around massive stars  include \\citet{brown05},\\citet{comeron07}, \\citet{gvar08}, and \\citet{povich}.  In \\citet{kobulnicky10} we used the \\emph{SST} Cygnus-X Legacy Survey data to identify several probable runaway stars and their bowshocks in the vicinity of the Cygnus OB2 association while introducing  a novel technique to exploit bowshock physics in order to estimate mass-loss rates from massive stars.  In this contribution we report the discovery that infrared source IRAS~03063+5735 is a bowshock nebula associated with an early B star in the outer Galaxy near $\\ell=140.59$\\degr, $b=-0.250$ \\degr.  Data from the \\emph{SST} Galactic Legacy Infrared MidPlane Survey Extraordinaire 360 (GLIMPSE 360) mid-infrared survey reveals the bowshock morphology for the first time.   We use optical spectroscopy from the Wyoming Infrared Observatory (WIRO)  2.3~m telescope to determine a spectral type and radial velocity.  These data, in turn, enable an estimate for the distance to the star and an associated molecular cloud complex. ",
        "conclusions": "We have identified an \\emph{IRAS} infrared source IRAS~03063+5735, previously reported as a galaxy, to be a bowshock nebula powered an early B-type (B1.5 V) star. The infrared source was selected from among several hundred sources having \\emph{IRAS} color-color criteria matching those expected from soft UV radiation fields generated by intermediate-mass (i.e., B-type) stars.  Our measurement of the $-58$ \\kms\\ LSR radial velocity of the star, combined with \\tco\\ line maps, allow us to identify a molecular cloud complex at a common velocity and $\\sim$4.0 kpc distance.  The bowshock is plausibly produced by the supersonic motion of a runaway B star as it impinges upon a local maximum in the molecular column density of $\\sim4.5\\times10^{21}$ \\cmmb.  However, a low-elocity star within a local density gradient may also be capable of producing the observed phenomena.  We identify the \\HII\\ region and stellar cluster associated with IRAS~03064+5638 at a projected distance of 64 pc as one plausible birth site for a runaway star.  The association of a star having a spectrophotometric distance with a particular molecular feature adds another data point to help anchor the (poorly defined) rotation curve in the outer Galaxy."
    },
    "1112/1112.1078.txt": {
        "abstract": "Observational work conducted over the last few decades %shows that most if not all galaxies indicates that all massive galaxies have supermassive black holes at their centres. Although the luminosities and brightness fluctuations of quasars in the early Universe suggest that some are powered by black holes with masses greater than 10 billion solar masses \\cite{netzer03,vestergaard08}, the remnants of these objects have not been found in the nearby Universe.  The giant elliptical galaxy Messier 87 hosts the hitherto most massive known black hole, which has a mass of 6.3 billion solar masses \\cite{sargent78,gebhardt11}.  Here we report that NGC 3842, the brightest galaxy in a cluster at a distance from Earth of 98 megaparsecs, has a central black hole with a mass of 9.7 billion solar masses, and that a black hole of comparable or greater mass is present in NGC 4889, the brightest galaxy in the Coma cluster (at a distance of 103 megaparsecs). %Other brightest cluster galaxies can be up to twice as luminous as these two %  galaxies, and may host black holes with even greater masses. These two black holes are significantly more massive than predicted by linearly extrapolating the widely-used correlations between black hole mass and the stellar velocity dispersion or bulge luminosity of the host galaxy \\cite{dressler89,kormendy95,fm00,gebhardt00,gultekin09}. Although these correlations remain useful for predicting black hole masses in less massive elliptical galaxies, our measurements suggest that different evolutionary processes influence the growth of the largest galaxies and their black holes. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.5170_arXiv.txt": {
        "abstract": "We present a general review of the current knowledge of \\ir20551 and its circumnuclear environment. This Ultraluminous Infrared Galaxy is one of the most puzzling sources of its class in the nearby Universe: the near-IR spectrum is typical of a galaxy experiencing a very intense starburst, but a highly obscured active nucleus is identified beyond $\\sim$5~$\\mu$m and possibly dominates the mid-IR energy output of the system. At longer wavelengths star formation is again the main driver of the global spectral shape and features. We interpret all the available IR diagnostics in the framework of simultaneous black hole growth and star formation, and discuss the key properties that make this source an ideal laboratory for the forthcoming {\\it James Webb Space Telescope}. ",
        "introduction": "Two main physical processes characterize the nuclear regions of active galaxies: intense star formation at rates of $\\sim$10$^2$--10$^3$ M$_\\odot$~yr$^{-1}$ (starburst, SB) and accretion on to a supermassive black hole (active galactic nucleus, AGN). The issue of SB and AGN connection in both local and distant galaxies is critical for a proper understanding of galaxy formation and evolution, of star formation history and metal enrichment of the Universe, and of the origin of the extragalactic background at low and high energies. There is indeed increasing evidence of a strong link between the starburst and AGN mechanisms in active systems. The empirical correlation between the mass of black holes (BHs) located at the centre of nearby galaxies (both active and passive/quiescent) and the mass of their spheroids (see Sani et al.~2011 [1] and references therein) suggests that the formation of bulges and the growth of the central BHs are tightly connected. Also the presence of circumnuclear star formation in a substantial fraction of local AGN (Genzel et al.~1998 [2], Cid Fernandes et al.~2004 [3], Schweitzer et al.~2006 [4], Sani et al.~2010 [5]) hints at the relation between the two phenomena. The overall conclusion of these studies is that in 30--50\\% of the cases the accreting supermassive BHs are associated with young (i.e. of age less than a few $\\times$100 Myr) star-forming regions, with clear evidence of an enhanced star formation rate (reaching up to starburst intensities) in most AGN. However, this does not necessarily imply any causal connection between the two physical processes. It could be simply the natural consequence of massive gas fuelling into the nuclear regions, due to either interactions/mergers or secular evolution such as bar-driven inflows. Both star formation and nuclear accretion, in fact, are triggered and subsequently fed by this gas reservoir.\\\\ In the local Universe, the optimal targets to study the AGN/SB interplay are the so-called Ultraluminous Infrared Galaxies (ULIRGs; Sanders \\& Mirabel 1996 [6]). These sources are the result of major mergers, during which the redistribution of the gaseous component drives vigorous starburst events and obscured nuclear accretion. It is now well established that ULIRGs are usually powered by a combination of both processes, giving rise to their huge luminosities ($L_{\\rm{bol}} \\sim L_{\\rm{IR}} > 10^{12} L_\\odot$). However, since the primary radiation field is reprocessed by dust, the identification of the dominant energy supply is often unclear. The simultaneous presence of star formation and AGN signatures in the mid-IR makes this a really favourable band to disentangle the AGN and SB components and explore their environment. In particular, (\\emph{i}) the available spectra of \\textit{bona fide} starburst-dominated and, respectively, unobscured AGN-dominated sources are widely different, and show little dispersion within the separate classes (Risaliti et al. 2006 [7], Brandl et al. 2006 [8]; Netzer et al. 2007 [9], Nardini et al. 2008 [10]). This allowed us to reproduce the AGN/SB contributions with fixed templates, especially over the the 3--8\\mum\\ spectral interval. % (\\emph{ii}) For a given bolometric luminosity, the mid-IR AGN emission is higher than that of a starburst by a factor that rapidly declines with wavelength, ranging from $\\sim$100 at 3--4~$\\mu$m [7] to $\\sim$25 at 5--8~$\\mu$m (Nardini et al.~2009 [11]). Such a large difference is due to the key contribution of the hot dust layers directly exposed to the AGN radiation field. Together with the relatively low dust extinction at these wavelengths, this allows the detection of an AGN even when it is heavily obscured and/or its total luminosity is small compared with the SB counterpart. Based on the above points, we successfully fitted the observed ULIRG spectra with a two-component analytical model, with only two free parameters: the relative AGN/SB contribution and the optical depth of the screen-like obscuration (if any) affecting the compact AGN component.\\\\ To understand whether the link between star formation and nuclear activity is a matter of \\textit{nature} (i.e. feedback processes) or \\textit{nurture} (i.e. host environments), here we investigate the circumnuclear structure of \\ir20551, an ideal laboratory thanks to its unique physical properties (in terms of both relative AGN/SB contribution and AGN obscuration), and to the fairly large multiwavelength dataset available. The paper is organized as follows: in Section 2 we review the present knowledge of the mid-IR properties of \\ir20551. The dust extinction law and gas column density are dealt with in Section 3. A possible general picture and the feasibility of future observations with {\\it James Webb Space Telescope} (\\textit{JWST}) are discussed in Section~4. In Section~5 we summarize our findings and draw the conclusions. Throughout this work we adopt a standard cosmology ($H_0=70$~km/s/Mpc, $\\Omega_m=0.3$, $\\Omega_\\lambda=0.7$). \\begin{figure}[!h] \\centering \\includegraphics[angle=0,width=16cm]{i20551_wfc_nicmos.pdf} \\caption{{\\footnotesize \\textit{HST} images of \\ir20551. Left: composite \\textit{B}+\\textit{I} band image obtained with the WFC F435W and F814W filters (PID 10592, PI A. Evans). The white square identifies the nuclear ($30\\arcsec \\times 30\\arcsec$) region. Right: NICMOS F160W image (PID 11235, PI J. Surace) of the nuclear region. The \\textit{Spitzer}/IRS short-wavelength/low-resolution (SL, $3.6\\arcsec$ width) and short-wavelength/high-resolution (SH, $4.7\\arcsec$ width) slits are shown together with the \\textit{VLT}/ISAAC slit (IS, $1\\arcsec$ width in the \\textit{L}-band). Given a spatial scale of $\\rm 950~pc/\\arcsec$, the SL, SH and IS slits cover regions of 3.4~kpc, 4.5~kpc and 950~pc respectively.}} \\label{fig0} \\end{figure} ",
        "conclusions": "In the present work we have first reviewed the properties of \\ir20551\\, a prototypical local ULIRG observed by our group in the \\textit{L}- and \\textit{M}-band with ISAAC at the \\textit{VLT}. The spectral analysis also includes the 5--8\\mum\\ spectrum obtained by \\textit{Spitzer}/IRS. According to the AGN/SB decomposition method we have developed in several previous papers [7,10,17], \\ir20551\\ turns out to be a composite source, dominated in the mid-IR by hot dust emission associated with deeply embedded BH accretion and characterized by a vigorous circumnuclear starburst which provides the main power supply to the whole system. We have then interpreted the key spectral properties of the source over the $\\sim$3--20\\mum\\ wavelength range (e.g. the reddening of the continuum, the presence of deep absorption features, the lack of high-ionization coronal lines and the detection of \\h2\\ rotational transitions) in the framework of dust and gas spatial distribution and physical conditions. Our main results are the following: (\\textit{i}) the shape of the AGN intrinsic continuum is partly degenerate with the form of the extinction law. This is mainly evident beyond 5\\mum. (\\textit{ii}) Given the gas amount inferred from X-ray observations, the central regions of \\ir20551 seem to have a dust-to-gas ratio much lower than the Galactic interstellar medium. (\\textit{iii}) Aliphatic hydrocarbon and HAC grains are over-abundant with respect to the local molecular clouds. (\\textit{iv}) A large covering of the nuclear engine likely prevents the ionization of the AGN narrow-line region and the excitation of fine-structure lines. Therefore, a screen of cold, dusty gas lies along the line of sight to the AGN, heavily extinguishing its spatially compact primary emission. (\\textit{v}) A large amount ($\\rm M_{H_2}=6.8\\times 10^8 \\rm M_\\odot$) of warm ($T_{\\rm{ex}}=347$~K) molecular hydrogen and PAH grains are associated with the starburst environment on typical scales of a few kpc. The findings have been \\textit{qualitatively} interpreted by means of a simple geometrical configuration as the one sketched in Fig.~\\ref{fig6}. We have finally described the great improvement in terms of sensitivity, spectral coverage and resolution that will be achieved in the near future with the advent of \\textit{JWST}. This will also allow us to separate the different spatial scales and explore in larger detail the connection between the AGN and SB environments and the mutual feedback between the two physical processes."
    },
    "1112/1112.2750_arXiv.txt": {
        "abstract": "Generation and diffusion of the magnetic field on the Sun is a key mechanism responsible for solar activity on all spatial and temporal scales - from the solar cycle down to the evolution of small-scale magnetic elements in the quiet Sun. The solar dynamo operates as a non-linear dynamical process and is thought to be manifest in two types: as a global dynamo responsible for the solar cycle periodicity, and as a small-scale turbulent dynamo responsible for the formation of magnetic carpet in the quiet Sun. Numerous MHD simulations of the solar turbulence did not yet reach a consensus as to the existence of a turbulent dynamo on the Sun. At the same time, high-resolution observations of the quiet Sun from Hinode instruments suggest possibilities for the turbulent dynamo. Analysis of characteristics of turbulence derived from observations would be beneficial in tackling the problem. We analyze magnetic and velocity energy spectra as derived from Hinode/SOT, SOHO/MDI, SDO/HMI and the New Solar Telescope (NST) of Big Bear Solar Observatory (BBSO) to explore the possibilities for the small-scale turbulent dynamo in the quiet Sun. ",
        "introduction": "A dynamo process can be defined as the generation of the magnetic field from an inverse cascade due to motions in a conducting medium. In the convective envelope of the Sun, such motions are the result of turbulent convection in the photosphere and beneath. The turbulent character of these motions is clearly demonstrated by observations. As an example we present high-resolution observations (Figure 1) of quiet sun (QS) granulation obtained recently with the 1.6 meter New Solar telescope (NST) of Big Bear Solar Observatory (BBSO, Goode et al. 2010a,b). The corresponding two-hour movie is available at the BBSO website {\\tt http://bbso.njit.edu/gallery/nst\\_tio\\_20100803.mpg}. Statistical analysis of this data set can be found in Abramenko et al. (2010, 2011) and at {\\tt http://www.bbso.njit.edu/{$\\sim$}avi/2011NSF/}.  \\begin{figure}[!ht] \\centerline{\\epsfxsize=5.0truein \\epsffile{Fig1.eps}} \\caption{Image of solar granulation in a quiet sun area observed on August 3, 2010 at 17:06 UT with the NST of BBSO equipped with a TiO filter. The pixel size is 0.0375 $''$ and the size of the image is 28.3 $\\times$ 26.2 $''$. A movie of the corresponding 2-hour data set further illustrates the turbulent nature of solar plasma flows in the photosphere (see the text for the link).} \\end{figure}  Turbulent dynamo operation implies a gradual gain of magnetic energy inside a volume (Batchelor 1950, Kazantsev 1968). Numerical simulations allow one to probe the turbulent dynamo action (e.g., Boldyrev 2001, Schekochihin et al. 2002, Cho et al. 2003, Boldyrev \\& Cattaneo 2004). Indeed, fron an initial moment for a simulation, one may track the evolution of the magnetic energy accumulated in the volume. In the case an increase of magnetic energy is detected, one can argue for a dynamo action. Unfortunately, this is not a possibility when one deals with observational data. There is no apparent moment of initiation, and the magnetic energy is constantly generated and dissipated keeping the plasma in a state of dynamical equilibrium. We have to find some approach to test this turbulent state and to make inferences on the possibilities of turbulent dynamo action.   Two approaches might be suggested. First, one may test the interplay between the kinetic and magnetic energies, i.e., to explore the spatial power spectra of kinetic and magnetic energy. Second, one might test the structural organization of the magnetic and velocity fields in QS, namely, to explore their multifractal properties. Let us consider a theoretical background for the first approach.  \\begin{figure}[!ht] \\centerline{\\epsfxsize=5.0truein \\epsffile{Fig2.eps}} \\caption{Sketch of a mutual behavior of the kinetic and magnetic power spectra in a turbulent medium. Red and green lines: - a case of high Prandtl number ($Pr=\\lambda_{\\nu}/\\lambda_{\\eta}$) when the magnetic spectrum is more shallow than the velocity spectrum, and at scales larger than $\\lambda_{\\nu}$, the magnetic energy dominates the kinetic energy. In this limit, the situation is favorable for turbulent dynamo action. Blue and green lines: - a case of low Prandtl number, when the  magnetic spectrum shows a dissipation cut-off at scales much larger than $\\lambda_{\\nu}$, magnetic eddies are exhausted at smaller scales, and the dynamo action is questionable. } \\end{figure}   In Figure 2, the green line schematically represents a spatial power spectrum of the turbulent velocities in the QS photosphere. The kinetic energy, being deposited at large scales via the large-scale convective subphotospheric motions, cascades down to smaller scales, producing the Kolmogorov-type spectrum of -5/3. At some scale, $\\lambda_{\\nu}$, the kinetic vorticies start to dissipate, which results as a cut-off of the spectrum.  In order to generate magnetic fields at a scale $\\lambda$, the following condition should be met: the characteristic time for the curling of a magnetic field line by a plasma's vortex (of size $\\lambda$) should be shorter than the characteristic time for the dissipation of a magnetic vortex of the same size. This further requires that the  Prandtl number - the ratio of the kinetic dissipation scale to the magnetic dissipation scale, $\\lambda_{\\nu}/\\lambda_{\\eta}$, - has to exceed unity (Pikelner 1961, Biskamp 1993). At a high Prandtl number the small-scale turbulent dynamo action is both proved theoretically and in simulations (Batchelor 1950, Nakagawa \\& Priest 1973, Boldyrev 2001, Schekochihin et al. 2002, Cho et al. 2003). In this case, the magnetic spectrum (red line in Figure 2) is shallower than the kinetic one (green line in Figure 2), so that magnetic dissipation starts at scales, $\\lambda_{\\eta}$, much smaller than that for the kinetic dissipation scale.   For the opposite case of low Prandtl number (blue line in Figure 2), the magnetic dissipation starts on much larger scales, $\\lambda_{\\eta} > \\lambda_{\\nu}$. Here, the turbulent dynamo action is very problematic, at least, at very small scales ($\\lambda < \\lambda_{\\eta}$) because these small magnetic eddies cannot survive. However, there are studies aimed at probing the turbulent dynamo at low Prandtl number via numerical simulations  (see, e.g., Boldyrev \\& Cattaneo 2004, Iskakov et al. 2007, Pietarila Graham et al. 2010).  Thus, it is useful to see how the real kinetic and magnetic spectra behave in the QS photosphere. ",
        "conclusions": "New generation solar instruments, namely, Hinode/SOT/SP and BBSO/NST demonstrate previously unclear tendencies in the behavior of QS magnetic and velocity fields.  First, as the resolution improves, the shallow magnetic energy spectrum  tends to extend toward higher wave-numbers (smaller scales), the slope remaining the same, whereas the velocity spectrum shows a steady maximum at 1-3 Mm followed by the Kolmogorov-type regime of $E(k) \\sim k^{-5/3}$. This tendency necessarily results in setting in high-Prandtl number turbulence with necessary turbulent dynamo action at small scales below $\\lambda_{TD} \\approx 1 - 0.5$ Mm.  Second, improved resolution of solar data revealed a highly intermittent and multifractal nature of the small-scale magnetic field in QS areas. This property of the magnetic field is extremely favorable for the turbulent dynamo action at all scales where the intermittency/multifractality is present.   The above results were presented and discussed at the 2010 Hinode-4 Meeting in Palermo (Italy). VA is thankfull to Yukio Katsukawa, Jan Stenflo, Saku Tsuneta and Bruce Lites for helpful discussions.  We gratefully acknowledge help of the NST team and support of NSF (ATM-0716512 and ATM-0847126), NASA (NNX08AJ20G, NNX08AQ89G, NNX08BA22G), AFOSR (FA9550-09-1-0655). Hinode is a Japanese mission developed and launched by ISAS/JAXA, collaborating with NAOJ as a domestic partner, NASA and STFC (UK) as international partners. It is operated by these agencies in co-operation with ESA and NSA (Norway)."
    },
    "1112/1112.1565_arXiv.txt": {
        "abstract": "We analyze the sensitivities of a geostationary gravitational wave interferometer mission operating in the sub-Hertz band. Because of its smaller armlength, in the lower part of its accessible frequency band ($10^{-4} - 2 \\times 10^{-2}$ Hz) our proposed Earth-orbiting detector will be less sensitive, by a factor of about seventy, than the Laser Interferometer Space Antenna (LISA) mission.  In the higher part of its band instead ($2 \\times 10^{-2} - 10$ Hz), our proposed interferometer will have the capability of observing super-massive black holes (SMBHs) with masses smaller than $\\sim 10^{6}$ M$_{\\odot}$. With good event rates for these systems, a geostationary interferometer will be able to accurately probe the astrophysical scenarios that account for their formation. ",
        "introduction": "\\label{intro}  The quest for the direct observation of gravitational radiation is one of the most pressing challenges in the physics of this century. Predicted by Einstein shortly after formulating his general theory of relativity, gravitational waves (GW) will allow us to probe regions of space-time otherwise unobservable in the electromagnetic spectrum \\citep{Tho87}.  Although we have excellent indirect evidence of their existence through their effects on the orbital evolution of binary systems containing pulsars, we have not been able yet to directly detect them.  Several ground-based gravitational wave detectors have been operational for awhile, and only recently they have been able to identify the most stringent upper limits to date for the amplitudes of the radiation expected from several classes of sources \\citep{LIGO,VIRGO,GEO,TAMA}. Next generation of Earth-based interferometers and pulsar-timing experiments \\citep{PULSAR}, as well as newly conceived space-born detectors \\citep{PPA98} are expected to achieve this goal.  Ground-based interferometers operate in the frequency band whose lower limit is at about $10$ Hz, mainly because of the large seismic and gravity-gradient noises below this frequency cut-off.  Since the mHz region is potentially very rich in gravitational wave sources, the most natural way to observe them is to build and operate a space-borne detector. The most notable example of a space interferometer, which has been under study for several years jointly by scientists in the United States of America and in Europe, is the Laser Interferometer Space Antenna (LISA) mission. By relying on coherent laser beams exchanged between three remote spacecraft forming a giant (almost) equilateral triangle, LISA aimed to detect and study cosmic gravitational waves in the $10^{-4} - 1$ Hz band.  Although over the years only a few space-based detector mission concepts have been considered as alternatives to the LISA mission (with ASTROD \\citep{ASTROD}, DECIGO \\citep{DECIGO}, and OMEGA \\citep{Hiscock} being the most notable examples), starting in 2011 (with the ending of the NASA/ESA partnership for flying LISA) more mission concepts have appeared in the literature \\citep{PCOS}. Their goals are to meet most (if not all) the LISA scientific objectives (highlighted in the 2010 Astrophysics Decadal Survey {\\it New Worlds, New Horizons (NWNH)} \\citep{NWNH}) at a lower cost.  In this context, motivated by the interest of the Brazilian Space Agency to pursue the development and launch of several geostationary satellites, we decided to analyze the scientific capabilities offered by a geostationary gravitational wave interferometer, henceforth called GEOGRAWI. GEOGRAWI was one of the alternative concepts to the LISA mission submitted in response to NASA's Request for Information \\# NNH11ZDA019L \\citep{RFI}. A detailed analysis of all the submitted projects and a final study report can be found in \\citep{PCOS}. Among them, a proposal by Sean T. McWilliams explores a mission concept very similar to the one discussed in this paper \\citep{OURRFI,Sean2011} called GADFLI. His concept and the one proposed by us, though strikingly similar, were nonetheless independently developed and submitted.  GEOGRAWI, like LISA, has three identical spacecraft interchanging coherent laser beams and forming an (almost) equilateral triangle. However, by being in a geostationary orbit, its accessible frequency band changes to about ($10^{-3} - 10$ Hz). It is worth noting that, although our interferometer reminds that of the OMEGA mission (proposed by R.  Hellings \\citep{Hiscock} about 15 years ago), it is in fact very different from it. The OMEGA spacecraft were six in total, with two at each vertex of a geocentric equilateral triangle of side equal to one million kilometers. Cost of the launching vehicle as well as challenges in performing the phase measurements associated to the large relative velocities between the spacecraft make OMEGA quite different from GEOGRAWI.  The astrophysical sources that GEOGRAWI is expected to observe within its operational frequency band include extra-galactic massive and super-massive black-hole coalescing binaries, the resolved galactic binaries and extra-galactic coalescing binary systems containing white dwarfs and neutron stars, a stochastic background of astrophysical and cosmological origin, and possibly more exotic sources such as cosmic strings. With GEOGRAWI we will be able to test Einstein's theory of relativity by comparing the waveforms detected against those predicted by alternative relativistic theories of gravity, and also by measuring the number of independent polarizations of the detected gravitational wave signals \\citep{TintoAlves2010}.  Since the sensitivities of the geostationary interferometers we considered (see section \\ref{SENSITIVITY} below) are degraded by their shorter armlength in the frequency region ($10^{-4} - 2 \\times 10^{-2}$) Hz, it will be impossible for them to detect the zero-order cyclic spectrum of the white dwarf-white dwarf galactic binary confusion noise. However, as pointed out in \\citep{Edlund}, the periodic nature of the galactic background signal in the data of an interferometer rotating around the Sun will contain higher-order ``cyclic spectra'', which are in principle not affected by the instrumental noise ({\\underbar {if}} this is stationary). This implies that we could still detect the so-called ``confusion noise'', and infer properties of the distribution of the white-dwarf binary systems present in our galaxy \\citep{Edlund}.  In this article we pay particular attention to super-massive black holes, which are the main sources of GWs for interferometers operating at frequencies below $1$ Hz. We show that GEOGRAWI will be able to detect a large number of SMBHs at very high redshifts, and that a significant fraction of them will have masses smaller than $10^{6}$ M$_{\\odot}$ as a consequence of its improved sensitivity at higher frequencies.  This paper is organized as follows. In section \\ref{SENSITIVITY} we provide a description of the mission and a system noise analysis in order to evaluate the sensitivities of the Time-Delay Interferometric (TDI) \\citep{TD2005} response $X$. It is in this section and in \\ref{noise} where we also emphasize that additional instrumental techniques over those baselined for LISA will need to be implemented in order to achieve the derived sensitivities. These are estimated for three different on board subsystem configurations of GEOGRAWI, which we refer to as Geostationary 1 (GEO1), Geostationary 2 (GEO2) , and Geostationary 3 (GEO3). We note that the GEO1 configuration has a reduced sensitivity over that of LISA by a factor of about $70$ in the lower part of its accessible frequency band ($10^{-4} - 2 \\times 10^{-2}$ Hz). On the other hand, in section \\ref{Science} we discuss possible sources of gravitational waves observable by GEOGRAWI in the higher-part of its band ($2 \\times 10^{-2} - 10$ Hz). Particular attention is devoted to SMBHs and we show that the number of observable signals from such systems, with masses smaller than or equal to $10^6$ M$_\\odot$, would be significant. Finally, in Section \\ref{Conclusions} we provide a summary of our work and conclusions.  \\bibliographystyle{model1a-num-names} ",
        "conclusions": "\\label{Conclusions}  The main goal of this article was to analyze the performance and some scientific capabilities of a geostationary gravitational wave interferometer. We have done so by estimating its sensitivities to gravitational radiation when operated under three different onboard subsystem configurations, and analyzed the kind of gravitational wave sources it will be able to detect and study in the sub-Hertz band. We found our proposed Earth-orbiting detector less sensitive than the Laser Interferometer Space Antenna mission by about a factor of seventy in the lower part of its accessible frequency band ($10^{-4} - 2 \\times 10^{-2}$ Hz). On the other hand, in the higher part of the band ($2\\times 10^{-2} - 10$ Hz), it was shown that GEOGRAWI has observational capabilities that are complementary to those of the LISA mission. In relation to this point, we emphasized that our assumed noise performance in the high-frequency region of the accessible band requires additional instrumental capabilities over those baselined for the LISA mission. Although current experimental evidence \\cite{Jordan2006,Dan1,Dan2} indicates that they will be available by the time such a mission will be considered, a more in-depth analysis of these challenges is beyond the scope of this paper.  In the case of binary systems containing SMBHs, our analysis has shown that the GEO1 interferometer could observe a great number of events per year. This is direct consequence of its good sensitivity in the frequency region ($2\\times 10^{-2} - 10$ Hz), where lighter (and larger in number) black-holes radiate in the merger and ring-down phases. Consistently we have also shown that these same systems will be detectable by LISA with a SNR larger than that achievable by GEO1 during the inspiraling phase (see Fig. \\ref{insp}).  Since most of the SNR from the coalescence phase of a binary system containing SMBHs is achieved after about a week of integration time before coalescence, because GEOGRAWI will make most of its observations in the high-part of its accessible frequency band, and because during a time period of one week the amplitude modulation of the GEOGRAWI response will result into a frequency modulation of the observed GW signal that depends on the source location in the sky, one can argue that GEOGRAWI should achieve very good accuracies in the reconstructed GW source parameters of SMBH with masses that are in the range $10^4 - 10^6 \\ M_\\odot$. It is worth noting that McWilliams \\cite{PCOS,McWillians} has performed such an analysis for the GADFLI mission concept (which is similar to GEOGRAWI). He showed that the precision in the estimation of the parameters achievable by GADFLI (which is analogous to our GEO1 configuration) is somewhat better than that for LISA. Although these results imply that GEO1 should achieve comparable accuracies in estimating the parameters characterizing SMBHB, it is our intention to quantitatively study this problem in a forthcoming article.  Since the high-part of the accessible frequency where GEOGRAWI achieves its best sensitivity is at around $1$ Hz, we will investigate also its ability to perform low-latency searches of burst events. Mainly due to its capability of continuously transmitting real-time data to the ground (by being geostationary), it should be able to trigger simultaneous searches for electromagnetic counter-parts."
    },
    "1112/1112.0496_arXiv.txt": {
        "abstract": "Fusion data for $^{13}$C+$^{13}$C, $^{12}$C+$^{13}$C and $^{12}$C+$^{12}$C are analyzed by coupled-channels calculations that are based on the M3Y+repulsion, double-folding potential. The fusion is determined by ingoing-wave-boundary conditions (IWBC) that are imposed at the minimum of the pocket in the entrance channel potential. Quadrupole and octupole transitions to low-lying states in projectile and target are included in the calculations, as well as mutual excitations of these states.  The effect of one-neutron transfer is also considered but the effect is small in the measured energy regime. It is shown that mutual excitations to high-lying states play a very important role in developing a comprehensive and consistent description of the measurements. Thus the shapes of the calculated cross sections for $^{12}$C+$^{13}$C and $^{13}$C+$^{13}$C are in good agreement with the data. The fusion cross sections for $^{12}$C+$^{12}$C determined by the IWBC are generally larger than the measured cross sections but they are consistent with the maxima of some of the observed peak cross sections. They are therefore expected to provide an upper limit for the extrapolation into the low-energy regime of interest to astrophysics. ",
        "introduction": "The fusion of carbon nuclei is an important reaction in the description of type Ia supernovae and other astronomical events in the cosmos like the superburst of an accreting neutron star. There is, unfortunately, a very large uncertainty in the predicted fusion cross sections for $^{12}$C+$^{12}$C that are needed in a stellar environment \\cite{gasq}, partly because the cross sections are difficult to measure down to the low energies that are of interest, and partly because the data contain strong resonance structures that make it difficult to extrapolate the measured cross sections to low energy. In order to get some constraints on the extrapolation it is instructive to analyze the existing fusion data for $^{13}$C+$^{13}$C by Trentalange et al. \\cite{trent} and Charvet et al. \\cite{char},  % and also of the $^{12}$C+$^{13}$C data by Dayras et al. \\cite{dayras} and Notani et al. \\cite{tang1}, because these data do not exhibit strong resonance features.  The $^{13}$C+$^{13}$C fusion data \\cite{trent} are analyzed by coupled-channels calculations that include couplings to the one-phonon quadrupole and octupole excitations in projectile and target, as well as mutual excitations of these states. The effect of one-neutron transfer, which can proceed to the $^{12}$C+$^{14}$C mass partition with positive $Q$ value % is also studied, and so is the effect of the neutron exchange with zero Q value in $^{12}$C+$^{13}$C reactions.  The fusion cross sections are determined by ingoing-wave-boundary conditions (IWBC) that are imposed at the minimum of the pocket in the entrance channel potential.  The calculated cross sections defined this way are fairly smooth as functions of the center-of-mass energy and they are well suited to analyze the fusion data for $^{13}$C+$^{13}$C \\cite{trent,char} and $^{12}$C+$^{13}$C \\cite{dayras,tang1}. The $^{12}$C+$^{12}$C fusion data, on the other hand, contain a lot of structures or resonances, and it is beyond the scope of this investigation to try to reproduce these data in detail. It is, however, of great interest to see how the calculated cross sections, obtained from the IWBC, compare to the data and how they possibly can be used to put constraints on the extrapolation to very low energies.  The coupled equations are solved using either a standard Woods-Saxon \\cite{BW} or the M3Y+repulsion, double-folding potentials \\cite{misi}. The coupled-channels effects on the calculated fusion cross sections are relatively modest compared to the large enhancement of several orders of magnitude that are commonly seen in calculations of heavy-ion fusion reactions. For the carbon systems, the enhancement of fusion compared to the no-coupling limit is typically a factor of 2 at energies far below the Coulomb barrier. However, it is a challenge to explain the data in detail because the coupled-channels calculations are sensitive to channels that have rather high excitation energies. One problem is that these channels are closed at low beam energies but that problem can be solved by imposing the correct, decaying state boundary conditions at large distances between the reacting nuclei \\cite{o16}. Another problem is that the nuclear structure of high-lying states is sometimes very uncertain.  The M3Y+repulsion potential was introduced in Ref. \\cite{misi} to explain the hindrance that has been observed in the fusion of many heavy-ion systems at very low energies \\cite{sys}. The hindrance phenomenon was first observed as a suppression at very low energies of $^{60}$Ni+$^{89}$Y fusion data compared to calculations that were based on a standard Woods-Saxon potential \\cite{niy}. A simple explanation of the phenomenon is the existence of a shallow pocket in the entrance channel potential which forces the fusion cross section to vanish as the center-of-mass energy approaches the minimum energy of the pocket \\cite{misi}. Important issues have been whether the fusion hindrance also occurs in light-ion systems, and how it will affect the extrapolation of measured cross section to the low energies that are of interest to astrophysics \\cite{cljastro}. The evidence for quasi-molecular resonances, observed in the elastic scattering \\cite{bromleyel} and fusion reactions \\cite{bromleyr,michaud} of $^{12}$C+$^{12}$C are usually explained as resonances in a shallow two-body potential between the reacting nuclei. We will show that the analysis of the $^{13}$C+$^{13}$C and $^{12}$C+$^{13}$C fusion data supports the idea of a fusion hindrance and the existence of a shallow pocket in the entrance channel potential.  The fusion cross sections reported in Refs. \\cite{trent,char,dayras} are based on measurements of the characteristic $\\gamma$-rays emitted from some of the evaporation residues that are produced, mostly those associated with the proton, neutron, and $\\alpha$ decay. The total fusion cross sections were obtained with the aid of statistical model calculations which is a major source of the systematic uncertainty. The systematic error is difficult to estimate but it can be quite large. For example, the systematic error of the absolute cross sections quoted in Ref.  \\cite{trent} is $\\pm$15\\%. Since some of the systematic error concerns the overall normalization of the measured cross sections, and not so much the shape (i.~e., the energy dependence of the cross section), it is of interest to adopt an adjustable overall normalization when analysing the data.  The ingredients of the coupled-channels calculations are presented in the next section. The analysis of the $^{13}$C+$^{13}$C fusion data is discussed in section III, where the optimum parameters of the M3Y+repulsion potential, such as the radius parameter of $^{13}$C and the diffuseness associated with the repulsive term, are determined. These parameters are used together with a parameterization of the experimental density of $^{12}$C as input to the coupled-channels calculations of the fusion cross sections for $^{12}$C+$^{13}$C and $^{12}$C+$^{12}$C which are presented in sections IV and V, respectively. The systematics of the low-energy fusion cross sections for the 3 systems of carbon isotopes is discussed in section VI. Finally, the conclusions of this work are presented in section VII. ",
        "conclusions": "We have explored the fusion of different combinations of carbon isotopes within the coupled-channels approach. The fusion cross section is determined by ingoing-wave-boundary conditions that are imposed at the minimum of the pocket in the entrance channel potential. It turned out that the measured cross sections for the fusion of $^{13}$C+$^{13}$C cannot be explained accurately at low energies when a standard Woods-Saxon potential is applied in the calculations.  However, the data can be explained fairly well by applying the shallow potential that is produced by the M3Y+repulsion potential. This is a characteristic feature of the fusion hindrance phenomenon which has been observed in many medium-heavy systems. One reason the shallow potential makes it possible to explain the carbon data by coupled-channels calculations is that the high-lying excitations become closed channels at low center-of-mass energies, not only with respect to reactions but also to fusion.  The coupled-channels calculations are surprisingly sensitive to mutual excitations and this seems to be justified by the comparison to the data. Thus the calculations that include one-phonon excitations alone cannot explain the data so well, neither when the standard Woods-Saxon nor the M3Y+repulsion potential is applied. The data are reproduced much better when mutual excitations are included and the shallow M3Y+repulsion potential is applied. The reason is that the direct coupling to these high-lying states is proportional to the second derivative of the ion-ion potential and this quantity is much larger for the M3Y+repulsion than for the standard Woods-Saxon potential because of the strong repulsion.  The large sensitivity to the excitation of high-lying states is particularly critical when the mutual excitations involve the strong octupole excitation of $^{12}$C. The uncertainty in the strength of the octupole excitation, and also in the strength of transitions to high-lying states that are built on the $2^+$ and $3^-$ states, makes it difficult to predict the fusion cross sections very accurately when $^{12}$C is involved. However, by adopting the nuclear structure properties of the high-lying states predicted by the shell model, one obtains a fairly good description of the $^{12}$C+$^{13}$C fusion data.  The prediction of the cross section for the fusion of $^{12}$C+$^{12}$C, based on the ingoing-wave-boundary conditions and the shell model prediction of the coupling to high-lying states, exceeds most of the measured cross sections and is consistent with the maxima of the observed peak cross sections. Thus it appears that the calculation provides an upper limit of the measured cross sections. The calculation is therefore expected to provide an upper limit for the extrapolation of cross sections into the unexplored territory of very low energies.  On the experimental side it is very important to get a better absolute normalization of the measured fusion cross sections. The current systematic errors are fairly large. This was exploited in the analysis of the $^{13}$C+$^{13}$C fusion data by focusing on the shape of the measured cross section, whereas the absolute normalization was treated as an adjustable parameter. In this way it was possible to calibrate the M3Y+repulsion potential and achieve an excellent fit to the data. It is of great interest to know whether the renormalization that was applied to the data can be justified by more accurate measurements.  {\\bf Acknowledgments}. The authors are grateful to B. A. Brown for providing the shell model calculations. This work was supported by the U.S. Department of Energy, Office of Nuclear Physics, contract no. DE-AC02-06CH11357. One of us (X.T.) is supported by the NSF under Grant No. PHY-0758100 and PHY-0822648, and the National Natural Science Foundation of China under Grant No. 11021504, and the University of Notre Dame."
    },
    "1112/1112.0944_arXiv.txt": {
        "abstract": "{Mass loss from massive stars forms an important aspect of the evolution of massive stars, as well as for the enrichment of the surrounding interstellar medium.} {Our goal is to predict accurate mass-loss rates and terminal wind velocities. These quantities can be compared to empirical values, thereby testing radiation-driven wind models. One specific topical issue is that of the so-called ``weak-wind problem'', where empirically derived mass-loss rates and (modified) wind momenta fall orders of magnitude short of predicted values.} {We employ an established Monte Carlo model and a recently suggested new line acceleration formalism to solve the wind dynamics more consistently.} {We provide a new grid of mass-loss rates and terminal wind velocities of $\\textrm{O}$-type stars, and compare the values to empirical results. Our models fail to provide mass-loss rates for main-sequence stars below a luminosity of log$(L/\\lsun) = 5.2$, where we appear to run into a fundamental limit. At luminosities below this critical value there is insufficient momentum transferred to the wind in the region below the sonic point in order to kick-start the acceleration of the flow. This problem occurs at almost the exact location of the onset of the weak-wind problem. For O dwarfs, the boundary between being able to start a wind, and failing to do so, is at spectral type O6/O6.5. The direct cause of this failure for O6.5 stars is a combination of the lower luminosity and a lack of Fe\\,{\\sc v} lines at the base of the wind. This might indicate that -- in addition to radiation pressure -- another mechanism is required to provide the necessary driving to initiate the wind acceleration.} {For stars more luminous than $10^{5.2}$\\,\\lsun, our new mass-loss rates are in excellent agreement with the mass-loss prescription by Vink et al. 2000 using our terminal wind velocities as input to this recipe. This implies that the main assumption entering the method of the Vink et al. prescriptions -- i.e. that the momentum equation is not explicitly solved for -- does not compromise the reliability of the Vink et al. results for this part of parameter space. Finally, our new models predict terminal velocities that are typically 35 and 45 percent larger than observed values. Such over-predictions are similar to those from (modified) CAK-theory.} ",
        "introduction": "\\label{sec:Introduction}  In this article, we present predictions for mass-loss rates and velocity structures for a grid of $\\textrm{O}$-type stars, using two distinct methods for solving the wind dynamics.  Mass loss forms an integral aspect characterizing massive $\\textrm{O}$-type stars. Because of their short lifetimes, massive stars are important tracers of star formation in galaxies. Furthermore, they enrich the interstellar medium with metals, both during their lives via stellar winds, as well as when they explode at the very end of their evolution. In order to build an evolutionary framework for massive stars, it is essential to map the mass-loss processes during the various evolutionary stages, as the exact rates of mass loss greatly influence the evolutionary tracks \\citep[e.g.][]{1981A&A....99...97M, 1986ARA&A..24..329C}. The effects of mass loss on the evolutionary tracks are at least two-fold: first and foremost the stellar mass is reduced, and secondly, the rotational velocity is strongly affected, as the mass also carries away angular momentum \\citep[e.g.][]{1998A&A...329..551L,2000A&A...361..159M}.  For the continuous stellar winds from massive stars, the outflow is thought to be driven by the transfer of energy and momentum from the radiation field to the atmosphere through the absorption of photons in atomic transitions. The exact amount of momentum and energy transfer has been the subject of both theoretical and observational studies for many decades ~\\citep{1970ApJ...159..879L,1975ApJ...195..157C,1986A&A...164...86P, 1996A&A...305..171P,1997ApJ...477..792D,1999A&A...350..181V,2004A&A...417.1003K,2007A&A...473..603M}. For luminous $\\textrm{O}$-type stars, with log($L/\\lsun) > 5.2$, the theoretical predictions of \\cite{2000A&A...362..295V} seem to be in reasonable agreement with empirical mass-loss rates provided that $\\textrm{O}$-stars are only subject to modest amounts of wind clumping (with clump filling factors of only 5-10). However, for objects with luminosities log($L/\\lsun)$ below approximately 5.2, a severe drop -- by a factor of $\\sim$100 -- in the empirically determined modified wind momentum (basically a multiplication of the mass-loss rate and the terminal velocity) has been revealed. This problem has in literature been referred to as ``the weak-wind problem'' ~\\citep{1996A&A...305..171P,2005A&A...441..735M,2008A&ARv..16..209P,2009A&A...498..837M}.  It deserves proper investigation simply because of the enormity of the effect. It is particularly important to find out whether the problem is caused by the mass-loss diagnostics or the predictions, as both are also applied to more luminous stars, where agreement between diagnostics and predictions has seemingly been achieved. But how certain can we be that this agreement is not a coincidence if we are aware of severe problems at lower luminosity?  Furthermore, we note that the oft-used mass-loss predictions of \\cite{2000A&A...362..295V} are semi-empirical, in the sense that empirical values for the wind velocity structure and terminal velocity are used as input to the modelling. In order to trust our overall knowledge of the mass-loss rates from $\\textrm{O}$-type stars -- at {\\it all} masses and luminosities -- it is pivotal to further scrutinize the \\cite{2000A&A...362..295V} assumptions, most notably that of the adopted wind dynamics.  Recently, \\cite{2008A&A...492..493M} suggested a new parametrization of the line acceleration, expressing it as a function of radius rather than of the velocity gradient, as in \\citeauthor{1975ApJ...195..157C} (\\citeyear{1975ApJ...195..157C}; henceforth CAK) theory. The implementation of this new formalism allows for local dynamical consistency, as one can determine the energy and momentum transfer at each location in the wind through the use of Monte Carlo simulations. Although the formalism was applied with three independent starting conditions that showed convergence to the same wind parameters, it has thus far only been applied to one object, that of an O5 dwarf.  For the adopted line force parameterization \\citeauthor{2008A&A...492..493M} identify an exact solution in case the medium is isothermal. We expand on this result by also accounting for a temperature stratification. To allow for such a study we employ the new line acceleration parameterization but solve for the wind dynamics consistently by applying a numerical  method to solve for the momentum equation.  The purpose of our study is threefold: {\\em    i)} to solve the wind dynamics numerically, and compare the results to those of \\cite{2008A&A...492..493M}, {\\em    ii)} to compute a larger grid of dynamically derived O-star mass-loss rates and wind terminal velocities, and determine the accuracy of the predictions made by \\cite{2000A&A...362..295V}, and {\\em    iii)} to utilize the grid in order to investigate the weak-wind problem.  Our paper is organized as follows. In Sect.~\\ref{sec:Method}, we start off describing the core of our method and the different methods to treat the wind equation. The results are presented in Sect.~\\ref{sec:Results} and discussed in the Sect.~\\ref{sec:discussion}. We end with the conclusions (Sect.~\\ref{sec:Conclusions}). ",
        "conclusions": "\\label{sec:Conclusions}  We have presented new mass-loss rates and terminal wind velocities for a grid of massive $\\textrm{O}$-type stars, improving the treatment of physics in the Monte Carlo method by \\citet{1985ApJ...288..679A} and \\citet{1997ApJ...477..792D} to predict wind properties of early-type stars. Two new types of solutions have been discussed. First, building on the work of \\citet{2008A&A...492..493M}, we present so-called best-$\\beta$ solutions in which one still assumes a $\\beta$-type velocity law (see Eq.~\\ref{eq:betalaw}) in the wind, but in which the terminal velocity and $\\beta$ are no longer adopted but constrained by requiring that they best fit the line force (distribution). Second, we abandon the $\\beta$-type velocity structure and introduce numerical solutions of the wind stratification. Our main conclusions are:  \\begin{enumerate}  \\item For stars more luminous than $10^{5.2}$\\,\\lsun\\, the best-$\\beta$ and hydrodynamical method yield $\\beta$ and $\\vinf$ results in agreement with each other (within 5-20 percent), whilst the mass-loss rates agree within a factor of 2.  \\item Furthermore, both methods are in very good agreement with the mass-loss prescription by \\cite{2000A&A...362..295V} using our terminal velocities in their recipe. This implies that the main assumption entering the method on which the \\citeauthor{2000A&A...362..295V} results are based -- i.e. that the momentum equation is not solved explicitly -- {\\em is not compromising their predicted \\mdot\\ in this luminosity range}. Terminal velocity is an input parameter to the \\citeauthor{2000A&A...362..295V} recipe. If we apply the canonical value $\\vinf = 2.6 \\varv_{\\rm esc}$, the discrepancy between our mass-loss rates and their mass-loss rates is of the order of $0.2 \\textrm{, although occasionally }0.6 \\textrm{ dex}$.  \\item At luminosities $\\la 10^{5.2}$\\,\\lsun\\  our hydrodynamical method fails to produce an outflow because of a lack of line driving at the base of the wind. This critical luminosity coincides with the onset of the `weak-wind problem'.  \\item For O dwarfs the above luminosity criterion translates to a boundary between starting and failing to start a wind at O6/O6.5. The direct cause of the failure to start a wind in O6.5 V stars is the lower luminosity and the lack of Fe\\,{\\sc v} lines at the base of the wind compared to O6 V stars.  \\item The fact that our hydrodynamical method fails to drive a wind at $L \\la 10^{5.2}$ \\lsun\\ may imply that some other mechanism is driving the weak winds or is supplementing the line acceleration at the base of the wind to help drive gas and initiate the wind.  \\item For stars more luminous than $10^{5.2}$\\,\\lsun\\ we predict, using the best-$\\beta$ and hydrodynamical method, terminal velocities that are typically 35 and 45 percent higher than observed. Such over-predictions are similar to what is seen in MCAK-theory \\citep{1995ApJ...455..269L}.  \\item We predict beta values in the range 0.85 to 1.05, with a trend that supergiants have slightly higher $\\beta$ values than dwarfs. This range of $\\beta$ values agrees very well with empirical results by \\citet{2003ApJ...586..996M}.  \\end{enumerate}"
    },
    "1112/1112.1729_arXiv.txt": {
        "abstract": "We test whether the Tinker \\&\\ Chen model of \\mgii\\ absorption due to the gaseous halo around a galaxy can reproduce absorption in quasar pairs (both lensed and physical) and lensed triples and quads from the literature.  These quasars exhibit absorption from a total of 38 \\mgii\\ systems spanning $z=0.043 - 2.066$ with mean redshift $\\langle z \\rangle=1.099$ and weighted mean rest-frame equivalent width of 0.87 \\AA.  Using the Tinker \\&\\ Chen model to generate simulated sight-lines, we marginalize the unknown parameters of the absorbing galaxies: dark matter halo mass, impact parameter, and azimuthal angle on the sky.  We determine the ability of the model to statistically reproduce the observed variation in \\mgii\\ absorption strength between paired sight-lines for different values of the gas covering fraction $f_c$ and the characteristic length scale $\\ell_c$, within which the variation in absorption equivalent widths between sight-lines exponentially decreases.  We find a best-fit $f_c=0.60\\pm0.15\\%$ and $\\ell_c < 8~h^{-1}_{70}~{\\rm kpc}$ (1$\\sigma$ confidence limits), with smaller $f_c$ allowed at larger $\\ell_c$.  At 99.7\\%\\ confidence, we are able to rule out $f_c>0.87$ for all values of $\\ell_c$ and the region where $\\ell_c<1.0~h^{-1}_{70}$ kpc and $f_c<0.3$. ",
        "introduction": "Intervening metal absorbers found in the spectra of background quasars have been known to be associated with galaxies at intermediate redshifts since the work of \\nocite{BB91}{Bergeron} \\& {Boiss{\\'e}} (1991).  As a result, quasar sight-lines have become a pivotal tool in probing the gaseous halos surrounding these galaxies, using (among other transitions) the \\mgii\\ $\\lambda\\lambda2796.352,2803.531$\\,\\AA\\ doublet (e.g., \\nocite{2011ApJ...733..105K}{Kacprzak} {et~al.} 2011).  The physical origin of the gaseous halo and its evolving interplay with its host galaxy, however, are not well understood.  Gas in the halo could be fuel for continued star formation in the galaxy (e.g., \\nocite{MB04}{Maller} \\& {Bullock} 2004), or could be a result of gas blown out by galactic winds (e.g., \\nocite{SSP10}{Steidel} {et~al.} 2010), or both. It is clear that the gas is clumped and not distributed spherically (e.g., \\nocite{CKSiau05}{Churchill}, {Kacprzak} \\&  {Steidel} 2005), but the typical covering fraction is still uncertain.  The spatial extent of the \\mgii\\ absorbing region has been constrained by various studies, although they do not always agree.  The extent could be related to the mass of the galaxy (e.g., \\nocite{TC08}{Tinker} \\& {Chen} 2008, \\nocite{SDP94}{Steidel}, {Dickinson} \\& {Persson} 1994, \\nocite{LS11}{Lovegrove} \\& {Simcoe} 2011) or to the colour and star-formation rate of the galaxy (e.g., \\nocite{ZMNQ07}{Zibetti} {et~al.} 2007, \\nocite{MW09}{M{\\'e}nard} {et~al.} 2009 [but see \\nocite{LC11}{L{\\'o}pez} \\& {Chen} 2011], \\nocite{SSP10}{Steidel} {et~al.} 2010).  The development of new observational and measurement techniques as well as new models of \\mgii\\ absorption is required to piece together the picture of the gaseous halo.   \\nocite{TC08}{Tinker} \\& {Chen} (2008) (hereafter TC08) have constructed a semi-analytical model for gas absorbing in the \\mgii~\\lam2796~\\AA\\ transition within a dark matter (DM) halo using an isothermal density profile.  The rest-frame absorption equivalent width $EW(2796)$ (absorption due to the 2796~\\AA\\ transition of the \\mgii\\ doublet at some projected distance $r$ from the centre of the absorbing gaseous halo, hereafter $EW$) is determined by the amount of cold gas, in the form of discrete clouds, along the line of sight (LOS) through a halo.  They constrained their model using a \\mgii-LRG cross-correlation function and a \\mgii\\ absorbing frequency distribution function.  In \\nocite{CT08}{Chen} \\& {Tinker} (2008), the authors further constrained the TC08 model using a sample of \\mgii\\ absorption features near known galaxies.  The sample consisted of 13 galaxy and absorber pairs and 10 galaxies that do not produce \\mgii\\ absorption lines to within sensitive upper limits. A maximum likelihood analysis showed that the model best describes the observations if the covering fraction is 80-86\\%.  The data had a weighted mean rest-frame equivalent width of 0.34~\\AA.  In this study we present a new test of the TC08 model, using \\mgii\\ absorption features in the lines of sight to lensed and multiple quasar systems.  These systems can provide two or more lines of sight through the same galaxy halo.  By comparing the coincident and anti-coincident absorption systems in multiple lines of sight, constraints on both the covering fraction and coherence scale can be determined (e.g., \\nocite{EIP04}{Ellison} {et~al.} 2004).  Multiple quasar systems may also be useful in measuring the azimuthal profiles of \\mgii\\ absorbers around galaxies, a topic studied statistically by \\nocite{BLK11}{Bordoloi} {et~al.} (2011).  We have compiled a non-exhaustive list from the literature of quasar pairs, triples, and quads (hereafter referred to as asterisms) where \\mgii\\ absorption is observed in one or all lines of sight.  Using a statistical approach, we ascertain how well the TC08 model can reproduce the observed variation in \\mgii\\ absorption between sight-lines for various values of the TC08 model parameters.  We describe our literature sample in \\S\\ \\ref{data}.  Descriptions of the theoretical framework of the TC08 model and the method by which we generate our simulations are given in \\S\\ \\ref{framework}.  Our analysis and results are described in \\S\\ \\ref{analysis}.  We discuss our results in \\S\\ \\ref{discussion}, and summarize our study in \\S\\ \\ref{summary}.  In this paper, we adopt $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $h_{70}=H_0/70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$ \\nocite{wmap3}({Spergel} {et~al.} 2007). ",
        "conclusions": "\\label{discussion}   Our analysis has ruled out significant portions of the parameter space we tested.  All characteristic lengths larger than 20 $h^{-1}_{70}$ kpc are ruled out, regardless of covering fraction.  This result indicates that the underlying form of the TC08 model does not represent the distribution of the \\mgii\\ absorbing gas.  When the characteristic length is so large, it forces both sight-lines in a pair to be scattered by the same value $\\sigma_x$.  This preserves the underlying form of the TC08 model (i.e., the 'unscattered' form), and is thus 'halo-to-halo' scatter; all sight-lines in one gaseous halo are scattered the same amount.  This is different from 'point-to-point' scatter, where all sight-lines in one gaseous halo are scattered by different amounts.  By ruling out large characteristic lengths, we have found that the TC08 model requires the cosmic scatter to be applied point-to-point, and thus the underlying analytical form of the model cannot fully reproduce the observed variation in \\mgii\\ absorption between sight-lines.   Given the above argument, it is not surprising that we have also ruled out covering fractions above 0.87 (for all characteristic lengths).  At covering fractions above 0.87, the probability that any sight-line will be set to zero is low, and thus we are forcing the use of the underlying analytical form of the TC08 model (see Figure \\ref{EWprofile}).  Observations show that not all lines of sight below the estimated $R_g$ show \\mgii\\ absorption, and thus it is not surprising that high covering fractions are ruled out.  There is also a region at low characteristic lengths and low covering fractions that has been ruled out (see the bottom left of Figure \\ref{contour}).  At covering fractions $f_c<0.3$, only the matched sight-line is guaranteed to have absorption; most other sight-lines only have residual absorption governed by the characteristic length, as set out by equations \\ref{WBprime} and \\ref{Wiprime}.  If the separation between sight-lines is much larger than the characteristic length of the simulation, most b, c, or d sight-lines will be set to zero.  In the literature this is clearly not the case, and thus we are able to rule out low covering fractions when the characteristic length is much smaller than the average separation of the data set.  As the characteristic length increases, the exponential factors of equations \\ref{WBprime} and \\ref{Wiprime} are no longer approximately zero.  As a result, even though most b, c, or d sight-lines are randomly chosen to be set to zero (at $f_c<0.3$), the large characteristic length forces many to be non-zero.  This effect is clearly visible in Figure \\ref{contour}, where we can no longer rule out low covering fractions between $2.0 < \\ell_c/ h^{-1}_{70}~ {\\rm kpc} < 20$."
    },
    "1112/1112.3801_arXiv.txt": {
        "abstract": "{Ideal magnetohydrodynamical (MHD) simulations have revealed catastrophic magnetic braking in the protostellar phase, which prevents the formation of a centrifugal disk around a nascent protostar.} {We determine if non-ideal MHD, including the effects of ambipolar diffusion and Ohmic dissipation determined from a detailed chemical network model, will allow for disk formation at the earliest stages of star formation.} {We employ the axisymmetric thin-disk approximation in order to resolve a dynamic range of 9 orders of magnitude in length and 16 orders of magnitude in density, while also calculating partial ionization using up to 19 species in a detailed chemical equilibrium model. Magnetic braking is applied to the rotation using a steady-state approximation, and a barotropic relation is used to capture the thermal evolution.} {We resolve the formation of the first and second cores, with expansion waves at the periphery of each, a magnetic diffusion shock, and prestellar infall profiles at larger radii. Power-law profiles in each region can be understood analytically. After the formation of the second core, the centrifugal support rises rapidly and a low-mass disk of radius $\\approx 10\\,R_{\\odot}$ is formed at the earliest stage of star formation, when the second core has mass $\\sim 10^{-3}\\,M_{\\odot}$. The mass-to-flux ratio is $\\sim 10^4$ times the critical value in the central region.} {A small centrifugal disk can form in the earliest stage of star formation, due to a shut-off of magnetic braking caused by magnetic field dissipation in the first core region. There is enough angular momentum loss to allow the second collapse to occur directly, and a low-mass stellar core to form with a surrounding disk. The disk mass and size will depend upon how the angular momentum transport mechanisms within the disk can keep up with mass infall onto the disk. Accounting only for direct infall, we estimate that the disk will remain $\\lesssim 10~\\mathrm{AU}$, undetectable even by ALMA, for $\\approx 4 \\times 10^{4}~\\mathrm{yr}$, representing the early Class 0 phase. } ",
        "introduction": "\\label{sec:Introduction}  Near the beginning of the star formation process, prestellar cores are observed to be rotating \\citep[e.g.,][]{GoldsmithArquilla1985,GoodmanEtAl1993,CaselliEtAl2002}. At the end of the process, observations show nearly all Class II objects to be surrounded by rotating disks of size $\\gtrsim 100~\\mathrm{AU}$, likely in centrifugal balance undergoing Keplerian rotation \\citep[see, e.g.,][and references therein]{AndrewsWilliams2005,WilliamsCieza2011}. Observational data concerning what happens in between those two snapshots is relatively scarce. Currently, there is no evidence for the presence of centrifugal disks larger than $\\approx 50~\\mathrm{AU}$ around Class 0 or Class I objects \\citep[e.g.,][]{MauryEtAl2010}. However, there are outflows observed around these young objects, commonly linked to disk accretion. Disks must therefore form with a small size and grow over time. The dearth of observational data at smaller scales will be remedied by ALMA. In the meantime, the evolution of angular momentum and the formation of a centrifugal disk have to be studied theoretically and numerically.  In order to form stars of the observed sizes and rotation rates, most of the angular momentum has to be removed from the infalling gas. This is the classical \\textit{angular momentum problem} \\citep[e.g.,][]{Spitzer1978}. The required reduction of angular momentum is often credited to \\textit{magnetic braking}, which acts during the contraction and collapse by linking the core with its envelope and transferring angular momentum \\citep[see][and references therein]{BasuMouschovias1994}. Recent numerical simulations of protostellar collapse under the assumption of ideal magnetohydrodynamics (MHD) have suggested the converse problem, namely that cores may experience \\textit{catastrophic magnetic braking}. This is the result that an extremely pinched magnetic field configuration has enough strength and exerts a long enough lever arm on the inner regions of collapse that magnetic braking can suppress the formation of a centrifugal disk entirely.  Catastrophic magnetic braking was first demonstrated by \\citet{AllenLiShu2003}, who used two-dimensional (2D) axisymmetric calculations. Subsequent ideal MHD simulations by \\citet{MellonLi2008} in 2D and \\citet{HennebelleFromang2008} in three dimensions (3D) showed that catastrophic magnetic braking occurs for initially aligned (magnetic field parallel to rotation axis) rotators in which the magnetic field is strong enough that $\\mu \\lesssim 10$, where $\\mu$ is the initial core mass-to-flux ratio normalized to its critical value for gravitational collapse. Disk formation in ideal MHD then requires at least $\\mu \\gtrsim 10$, which is significantly greater than the typical value $\\sim 2$ for dense cores determined from observations \\citep[e.g.,][]{TrolandCrutcher2008}. Furthermore, \\citet{MellonLi2009} found that the inclusion of non-ideal effects of Ohmic dissipation and ambipolar diffusion in a 2D model was not sufficient to enable disk formation on scales of $\\gtrsim 10$ AU from the center that were resolvable in their simulation. In a newer 2D model that also includes the Hall term of non-ideal MHD, \\citet{LiKrasnopolskyShang2011} again find essentially the same result. All of the above simulations have an inner resolution limit of $\\approx 10$ AU due to the use of a sink cell or an artificially stiff equation of state at high densities. Since these simulations also terminated soon after the central stellar core formed, the main common result of these simulations is that large $\\sim 100$ AU disks (as observed around Class II objects) cannot form during the early Class 0 phase. A 3D ideal MHD model by \\citet{HennebelleCiardi2009} that starts with a non-aligned rotator, also finds catastrophic magnetic braking on these scales but within a smaller range of $\\mu \\lesssim 3$. The effect of numerical reconnection may account for some of the differences amongst these models. All of the above works leave open the possibility that a small disk may actually form within the inner $\\sim 10$ AU during the earliest phases of star formation.  An interesting question is: can a very small disk form in the innermost regions (near the protostar) in the case of ideal MHD? \\citet{GalliEtAl2006} developed a simplified analytic model with a split-monopole magnetic field to show that catastrophic magnetic braking extends to the stellar surface in the flux-frozen limit. \\citet{KrasnopolskyKoenigl2002} and \\citet{BraidingWardle2011} employed self-similar non-ideal MHD models in the thin-disk approximation, using the prescription for steady-state magnetic braking developed by \\citet{BasuMouschovias1994}. Depending on different levels of effective diffusivity, they found either catastrophically-braked or disk-formation solutions. However, the diffusivity was parametrized in a way that preserved self-similarity, and its magnitude and functional form did not make contact with microphysical models. Consequently, a prediction for whether or not disks may actually form was not possible.  Our work fills the gap of numerically studying disk formation in the early Class 0 phase down to scales smaller than a stellar radius. In a previous paper \\citep[][hereafter DB10]{DappBasu2010} we followed the evolution of a collapsing molecular cloud core in axisymmetric thin-disk geometry all the way to protostellar densities ($n>10^{20}~\\mathrm{cm}^{-3}$). We found that Ohmic dissipation alone (in a simplified form) reduces the magnetic field strength sufficiently to render magnetic braking within the first core ineffective. As a result, a centrifugal disk can indeed form around the second core (the protostar).  The only other related MHD models that resolve the stellar core are the 3D non-ideal MHD simulations by \\citet{MachidaEtAl2006,MachidaEtAl2007} and \\citet{MachidaMatsumoto2011}, who also used a simplified treatment of Ohmic dissipation. In most of their published models, the evolution was halted at up to ten days after the second core was formed, due to time-step restrictions, and the formation of a centrifugal disk was not studied. Very recent work \\citep{MachidaMatsumoto2011} follows disk formation in one MHD model, and its relation to our work is discussed in Section \\ref{sec:Discussion}.  In this paper, we increase realism by employing a detailed chemical network including the effect of grains to calculate partial ionization and the resulting Ohmic \\textit{and} ambipolar diffusivities. We use the method presented in \\citet{KunzMouschovias2009}. Both Ohmic dissipation and ambipolar diffusion gain importance with the formation of the first core \\citep[e.g.,][]{LiMcKee1996, ContopoulosEtAl1998,DeschMouschovias2001}. We follow the evolution of the core to stellar sizes, using a multi-fluid approach that includes important inelastic collisions between gas-phase species and grains. We also study the effect of different grain sizes. We model four gas-phase species and up to five different grain sizes with three charge states (singly positively- or negatively charged, and neutral) for each grain size. We demonstrate that Ohmic dissipation and ambipolar diffusion can reduce the efficiency of magnetic braking in the inner regions enough to allow for the formation of a centrifugally-supported disk. There is dramatic magnetic flux loss, and a small disk can form in a very early phase of evolution. We propose a disk formation scenario in which the disk remains $\\lesssim 10~\\mathrm{AU}$ during the Class 0 phase, and subsequently expands significantly due to internal angular momentum transport mechanisms such as magnetorotational instability (MRI) or gravitational torques \\citep[e.g.,][]{BalbusHawley1991,VorobyovBasu2007}.  This paper is structured the following way. In Section \\ref{sec:method_DBK11} we formulate the problem and describe the method of solution. Section \\ref{sec:magn_braking} discusses the implementation of magnetic braking, while Section \\ref{sec:MHD} outlines the derivation of the induction equation. Section \\ref{sec:Chemistry} details the chemical model used to calculate the abundances of each species and Section \\ref{sec:Initial_Conditions} contains the description of our initial conditions. Readers mainly interested in our results may proceed directly to Section \\ref{sec:Results}. The discussion is found in Section \\ref{sec:Discussion}, and we summarize our findings in Section \\ref{sec:Summary_DBK11}. ",
        "conclusions": "\\label{sec:Summary_DBK11}  We present results from a new axisymmetric code using the thin-disk approximation that calculates the collapse of rotating magnetized prestellar cores. While treating vertical ($z$) structure and angular momentum transport in an approximate manner, our code allows us to follow the evolution all the way to stellar sizes and near-stellar densities. We determine the abundances of seven different species, and consider inelastic collisions. Our effective resistivity includes the effect of ambipolar diffusion and Ohmic dissipation for five different (single) grain sizes as well as for an MRN distribution of grain sizes.  By following the complex relationships between the many nonlinear processes at work in the formation of stars (e.g. self-gravity, rotation, chemistry, thermodynamics, non-ideal MHD, grain effects, etc.), we have been able to simultaneously address both the angular momentum problem and the magnetic flux problem of star formation. These two processes are intricately and fundamentally linked.  We demonstrate the formation of a small centrifugally supported disk despite the effects of magnetic braking. The ``magnetic braking catastrophe'' is averted on small scales by substantial magnetic flux loss from the high-density region of the first core. This weakens the magnetic field, preventing it from spinning down the material in that region. The central mass-to-flux ratio increases by a factor of $\\sim 10^{4}$. Shortly after the second collapse, disk formation happens very close to the central object, before the protostar has accreted a lot of mass ($M_{\\bigstar} <  10^{-2}~\\Msun$). This is consistent with the observational evidence of outflows at a very young age and the simultaneous non-detection of disks $\\gtrsim 50~\\mathrm{AU}$ around Class 0 objects \\citep{MauryEtAl2010}. ALMA will allow observers to improve upon these constraints, by probing for disks of size $\\approx 10~\\mathrm{AU}$.  We propose a disk formation scenario in which centrifugal disks form at the earliest stage of star formation and on small scales. We calculate the centrifugal radius for mass shells that are still infalling at the end of our simulation, and estimate the disk to remain $\\lesssim 10~\\mathrm{AU}$ for $\\approx 4 \\times  10^{4}~\\mathrm{yr}$. In this case, a disk would not be observable around a young Class 0 object, even with ALMA. However, indirect indications of the presence of disks --- such as outflows --- can still be expected. The accretion disk slowly grows by continued accretion but more so by internal angular momentum redistribution, and may eventually reach the size $\\approx 100~\\mathrm{AU}$ observed around Class II objects."
    },
    "1112/1112.1035_arXiv.txt": {
        "abstract": "{Observations of transiting extrasolar planets are of key importance to our understanding of planets because their mass, radius, and mass density can be determined. These measurements indicate that planets of similar mass can have very different radii. For low-density planets, it is generally assumed that they are inflated owing to their proximity to the host-star. To determine the causes of this inflation, it is necessary to obtain a statistically significant sample of planets with precisely measured masses and radii.} {The CoRoT space mission allows us to achieve a very high photometric accuracy. By combining CoRoT data with high-precision radial velocity measurements, we derive precise  planetary radii and masses. We report the discovery of CoRoT-19b, a gas-giant planet transiting an old, inactive F9V-type star with a period of four days.} {After excluding alternative physical configurations mimicking a planetary transit signal, we determine the radius and mass of the planet by combining CoRoT photometry with high-resolution spectroscopy obtained with the echelle spectrographs SOPHIE, HARPS, FIES, and SANDIFORD. To improve the precision of its ephemeris and the epoch, we observed additional transits with the TRAPPIST and Euler telescopes. Using HARPS spectra obtained during the transit, we then determine the projected angle between the spin of the star and the orbit of the planet.} {We find that the host star of CoRoT-19b is an inactive F9V-type star close to the end of its main-sequence life. The host star has a mass $M_{*}=1.21\\pm0.05$ \\,$M_\\odot$ and radius $R_{*}=1.65\\pm0.04$\\,$R_\\odot$.  The planet has a mass of $M_P=1.11\\pm0.06$\\,$M_{Jup}$ and radius of $R_P=1.29\\pm0.03$\\,$R_{Jup}$. The resulting bulk density is only $\\rho=0.71\\pm0.06$\\,$g\\,cm^{-3}$, which is much lower than that for Jupiter.} {The exoplanet CoRoT-19b is an example of a giant planet of almost the same mass as Jupiter but a $\\approx $30\\% larger radius. } ",
        "introduction": "The discovery of gas-giant planets orbiting their star at distances $\\leq0.1$\\,AU more than 15 years ago was surprising, and their study is still an active field of research.  A key observation for understanding these objects are those of transiting planets because they allow us to determine their true mass, radius, and full orbital parameters. For some transiting planets, it is even possible  to determine the projected angle between the stellar rotation and orbital spin axes. Thus, transiting planets allow us to obtain very detailed information that is essential to constrain models of  the formation and evolution of extrasolar planets.  The determination of the mass and radius and hence the density of extrasolar planets led to a rather unexpected result that planets of the same mass can have very different radii and thus very different bulk densities. Planets of very low density are possibly inflated owing to their close proximity to the host star. To aid our understanding of why close-in planets are inflated, it is essential to obtain precise values for the masses and radii of many extrasolar planets orbiting stars of different mass, temperature, luminosity, and orbital distances.  Since accurate planetary parameters are crucial for modeling the structure of the planet, it is preferable to have precise measurements of a small sample of transiting exoplanets, rather than low precision measurements of a larger sample. Space-based transit search programs are ideal for this purpose because their high sensitivity and excellent time-coverage enables us to explore a wider range of planet densities, and orbital periods.  Here we present the discovery of \\object{CoRoT-19b}, a planet with the mass of Jupiter but a radius that is substantially larger. By determining the properties of the planet and its host star precisely, we thus add in an important piece to the puzzle of exoplanets. ",
        "conclusions": "\\begin{figure} \\includegraphics[height=.27\\textheight,angle=0]{Allplanets.pdf} \\caption{Mass-radius diagram of the planets and brown dwarfs. CoRoT-19b is marked with a thick red point and by the arrows, all other objects discovered by CoRoT are indicated by the tick blue points, and planets discovered in other surveys are shown as small (black) points. } \\label{Aplanets} \\end{figure}  In Fig.~\\ref{Aplanets} we present the mass-radius diagram for transiting planets and brown dwarfs.  The objects marked with a large blue point are those discovered by CoRoT, the objects with small points are those detected with other instruments. The planet \\object{CoRoT-19b} is indicted by a large red dot and the two arrows. The two CoRoT-planets above \\object{CoRoT-19b} are \\object{CoRoT-1b} (Barge et al. \\cite{barge08}) and \\object{CoRoT-12b} Gillon et al. (\\cite{gillon10}). Their densities are $0.38\\pm0.05$ and $0.31\\pm0.08$ $g\\,cm^{-1}$, respectively, a bit lower than \\object{CoRoT-19b}, which has $0.71\\pm0.06$ $g\\,cm^{-1}$. Thus, \\object{CoRoT-19b} is an example of an inflated extrasolar planet but is less extreme than the others."
    },
    "1112/1112.3340_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}  The distribution of stellar masses that form together, the stellar initial mass function (IMF), is one of the most important astrophysical distribution functions. The determination of the IMF is a very difficult problem because stellar masses cannot be measured directly and because observations usually cannot assess all stars in a population requiring elaborate bias corrections. Indeed, the stellar IMF is not measurable (the {\\sc IMF Unmeasurability Theorem} on p.~\\pageref{kroupa_theorem:IMFtheorem}).  Nevertheless, impressive advances have been achieved such that the shape of the IMF is reasonably well understood from low-mass brown dwarfs (BDs) to very massive stars.  The IMF is of fundamental importance because it is a mathematical expression for describing the mass-spectrum of stars born collectively in ``one event''. Here, {\\it one event} means a {\\it gravitationally-driven collective process of transformation of the interstellar gaseous matter into stars on a spatial scale of about one~pc and within about one~Myr}. Throughout this text such events are referred to as {\\it embedded star clusters}, but they need not lead to gravitationally bound long-lived open or globular clusters. Another astrophysical function of fundamental importance is the star-formation history (SFH) of a stellar system.  The IMF and the SFH are connected through complex self-regulating physical processes on galactic scales, whereby it can be summarised that for late-type galaxies the star-formation rate (SFR) increases with increasing galaxy mass and the deeper gravitational potential. For early-type galaxies the same is true except that the SFH was of short duration ($\\simless\\;$few$\\,$Gyr). Together the IMF and SFH contain the essential information on the transformation of dark gas to shining stars and the spectral energy distribution thereof. They also contain the essential information on the cycle of matter, which fraction of it is locked up in feeble stars and sub-stellar objects, and how much of it is returned enriched with higher chemical elements to the interstellar medium or atmosphere of a galaxy. Knowing the rate with which matter is converted to stars in galaxies is essential for understanding the matter cycle and the matter content in the universe at a fundamental level.  This text is meant to outline essentials in IMF work, to document its form which appears to be invariant for the vast majority of resolved star-formation events and to describe modern evidence for IMF variation and how whole galaxies are to be described as composite or complex populations. The literature on the IMF is vast, and it is unfortunately not possible to cover every research paper on this topic, although some attempt has been made to be as inclusive as possible.  Sec.~\\ref{sec:intro} gives a brief historical review and a short overview of the topic, and pointers to other reviews are provided in Sec.~\\ref{sec:otherreviews}.  \\subsection{Solar neighbourhood}  Given the importance of the IMF a major research effort has been invested to distill its shape and variability. It began by first considering the best-known stellar sample, namely that in the neighbourhood of the Sun.  The seminal contribution by \\cite{S55} whilst staying in Canberra first described the IMF as a power-law, $dN=\\xi(m)\\,dm=k\\,m^{-\\alpha}$, where $dN$ is the number of stars in the mass interval $m,m+dm$ and $k$ is the normalisation constant.  By modelling the spatial distribution of the then observed stars with assumptions on the star-formation rate, Galactic-disk structure and stellar evolution time-scales, Salpeter arrived at the power-law index (or ``slope'') $\\alpha=2.35$ for $0.4 \\simless m/M_\\odot \\simless 10$, which today is known as the ``Salpeter IMF''.\\footnote{As noted by \\cite{Zinnecker11}, Salpeter used an age of~6~Gyr for the MW disk; had he used the now adopted age of~12~Gyr he would have arrived at a ``Salpeter index'' $\\alpha\\approx 2.05$ instead of~2.35.  }  This IMF form implies a diverging mass density for $m\\rightarrow 0$, which was interesting since dark matter was speculated, until the early~1990's, to possibly be made-up of faint stars or sub-stellar objects.  Studies of the stellar velocities in the solar-neighbourhood also implied a large amount of missing, or dark, mass in the disk of the Milky Way (MW) \\citep{B84}. Careful compilation in Heidelberg of the Gliese {\\it Catalogue of Nearby Stars} beginning in the 1960's \\citep{JW97}\\footnote{The latest version of the catalogue can be found at http://www.ari.uni-heidelberg.de/datenbanken/aricns/, while http://www.nstars.nau.edu/ contains the Nearby Stars (NStars) database.}, and the application at the beginning of the 1980's of an innovative photographic pencil-beam survey-technique reaching deep into the Galactic field in Edinburgh by \\citet{RG82} significantly improved knowledge of the space density of low-mass stars (LMSs, $m/M_\\odot \\simless 0.5$).  Major studies extending Salpeter's work to lower and larger masses followed, showing that the mass function (MF) of Galactic-field stars turns over below one solar mass thus avoiding the divergence. Since stars with masses $m\\simless 0.8\\,M_\\odot$ do not evolve significantly over the age of the Galactic disk, the MF equals the IMF for these. While the work of \\cite{MS79} relied on using the then-known nearby stellar sample to define the IMF for $m<1\\,M_\\odot$, \\cite{Sc86} relied mostly on a more recent deep pencil-beam star-count survey using photographic plates. \\cite{Sc86} stands out as the most thorough and comprehensive analysis of the IMF in existence, laying down notation and ideas in use today.  The form of the IMF for low-mass stars was revised in the early 1990's in Cambridge, especially through the quantification of significant non-linearities in the stellar mass--luminosity relation and evaluation of the bias due to unresolved binary systems \\citep{KTG90,KTG91,KTG93}. This work lead to a detailed understanding of the shape of the stellar luminosity function (LF) in terms of stellar physics. It also resolved the difference between the results obtained by \\cite{MS79} and \\cite{Sc86} through rigorous modelling of all biases affecting local trigonometric-based and distant photometric-parallax-based surveys, such as come from an intrinsic metallicity scatter, evolution along the main sequence and contraction to the main sequence.  In doing so this work also included an updated local stellar sample {\\it and} the then best-available deep pencil-beam survey. As such it stands unique today as being the only rigorous analysis of the late-type-star MF using simultaneously both the {\\it nearby trigonometric parallax} and the {\\it far, pencil-beam} star-count data to constrain the one underlying MF of stars.  This study was further extended to an analysis of other ground-based pencil-beam surveys being in excellent agreement with measurements with the HST of the LF through the entire thickness of the MW (Fig.~\\ref{fig:MWlf} below).  These results ($\\alpha\\approx 1.3, 0.1-0.5\\,M_\\odot$), were confirmed by \\citet{RGH02} using updated local star-counts that included Hipparcos parallax data.  Indeed, the continued long-term observational effort on nearby stars by Neill Reid, John Gizis and collaborators forms one of the very major pillars of modern IMF work; continued discussion of controversial interpretations have much improved and sharpened our general understanding of the issues.  A re-analysis of the nearby mass function of stars in terms of a log-normal form in the mass range $0.07-1\\,M_\\odot$ was provided by Gilles Chabrier, finding agreement to the deep HST star-count data once unresolved multiple stars and a metal-deficient colour-magnitude relation for thick-disk M~dwarfs are accounted for \\citep{Ch03}. The log-normal form with a power-law extension to high masses is indistinguishable from the older canonical two-part power-law form (see Fig.~\\ref{fig:canIMF} below). The immense analysis by \\cite{Bochanski10} of the LF and MF of~$15\\times 10^6$ field low-mass dwarfs ($0.1\\simless m/M_\\odot \\simless 0.8$) derived from Sloan Digital Sky Survey Data Release~6 photometry using the photometric parallax method again finds good consistency with the previous work on the stellar LF and MF.  The Galactic-field stellar IMF for $0.08\\simless m/M_\\odot\\simless 1$ can thus be regarded as being reasonably well-constrained. It converges to a finite mass density such that low-mass stars cannot be a dynamically-significant contributor to the MW disk. The dynamical evidence for significant amounts of dark matter in the disk was finally eroded by \\cite{KG91} and \\cite{FF94}.  \\subsection{Star clusters}  In contrast to the Galactic-field sample, where stars of many ages and metallicities are mixed, clusters offer the advantage that the stars have the same age and metallicity and distance. And so a very large effort has been invested to try to extract the IMF from open and embedded clusters, as well as from associations.  On the theoretical side, the ever-improving modelling of stellar and BD atmospheres being pushed forward with excellent results notably by the Lyon group (Isabelle Baraffe and Gilles Chabrier), has allowed consistently better constraints on the faint-star MF by a wide variety of observational surveys \\footnote{\\label{foot:pms} Here it should be emphasised and acknowledged that the intensive and highly fruitful discourse between Guenther Wuchterl and the Lyon group has led to the important understanding that the classical evolution tracks computed in Lyon and by others are unreliable for ages less than a few~Myr \\citep{Tout99,WT03}. This comes about because the emerging star's structure retains a memory of its accretion history. In particular \\cite{WK01} present a SPH computation of the gravitational collapse and early evolution of a solar-type star documenting the significant difference to a pre-main sequence track if the star is instead classically assumed to form in isolation.}. Furthermore, the development of high-precision $N$-body codes that rely on complex mathematical and algorithmic regularisation of the equations of motions through the work of Sverre Aarseth and Seppo Mikkola \\citep{Aarseth99} and others has led to important progress on understanding the variation of the dynamical properties\\footnote{\\label{foot:dynprop}The {\\it dynamical properties} of a stellar system are its {\\it mass} and, if it is a multiple star, its orbital parameters ({\\it semi-major axis, mass ratio, eccentricity}, both inner and outer if it is a higher-order multiple system). See also Sec.~\\ref{sec:bins}.} of stellar populations in individual clusters and the Galactic field.  In general, the MF found for clusters is consistent with the Galactic-field IMF for $m<1\\,M_\\odot$, but it is still unclear as to why open clusters have a significant deficit of white dwarfs \\citep{Fellhauer_etal03}. Important issues are the rapid and violent early dynamical evolution of clusters due to the expulsion of residual gas and the associated loss of a large fraction of the cluster population, and the density-dependent disruption of primordial binary systems \\citep{KAH}.  Young clusters have thus undergone a highly complex dynamical evolution which continues into old age \\citep{BaumMakin03} and are therefore subject to biases that can only be studied effectively with full-scale $N$-body methods, thus imposing a complexity of analysis that surpasses that for the Galactic-field sample. The pronounced deficit of BDs and low-mass stars in the 600 Myr old Hyades are excellent observational proof how dynamical evolution affects the MF \\citep{BKM08,Roeseretal11}. While estimating masses of stars younger than a few~Myr is subject to major uncertainty$^{\\rm \\ref{foot:pms}}$, differential reddening is also a complicated problem to handle when studying the IMF (Sec.~\\ref{sec:biases}).  \\vspace{2mm} \\centerline{ \\fbox{\\parbox{\\columnwidth}{ {\\sc Important Note}: In order to constrain the shape and variation of the IMF from star-count data in clusters high-precision $N$-body modelling must be mastered in addition to stellar-evolution theory and knowledge of the properties of multiple stars.}}}  \\vspace{2mm}  \\subsection{Intermediate-mass and massive stars}  Intermediate-mass and particularly massive stars are rare and so larger spatial volumes need to be surveyed to assess their distribution by mass.  For intermediate-mass and massive stars, the \\citet{Sc86} IMF is based on a combination of Galactic-field star-counts and OB association data and has a slope $\\alpha\\approx 2.7$ for $m \\simgreat 2\\,M_\\odot$ with much uncertainty for $m\\simgreat 10\\,M_\\odot$.  The previous determination by \\cite{MS79} also implied a relatively steep field-IMF with $\\alpha \\approx 2.5, 1 \\le m/M_\\odot \\le 10$ and $\\alpha \\approx 3.3, m>10\\,M_\\odot$. \\cite{ES06} point out that structure in the MF is generated at a mass if the SFR of the population under study varies on a time-scale comparable to the stellar evolution time-scale of its stars near that mass. Artificial structure in the stellar IMF may be deduced in this case, and this effect is particularly relevant for the field-IMF. Perhaps the peculiar structure detected by \\cite{GBH05} in the field-MF of stars in the Large Magellanic Cloud may be due to this effect.  An interesting finding in this context reported by \\citet{M98, Massey03} is that the OB stellar population found in the field of the LMC has a very steep MF, $\\alpha \\approx 4.5$, which can be interpreted to be due to the preferred formation of small groups or even isolated O and B-type stars (e.g. \\citealt{Selier11}).  Another or an additional effect influencing the deduced shape of the IMF via a changing SFR \\citep{ES06} is the dynamical ejection of OB stars from dynamically unstable cores of young clusters. This may lead to such a steep IMF because, as dynamical work suggests, preferentially the less-massive members of a core of massive stars are ejected \\citep{ClarkePringle92, P-AK06a}.  This process would require further study using fully-consistent and high-precision $N$-body modelling of young clusters to see if the observed distribution of field-OB stars can be accounted for with this process alone, or if indeed an exotic star-formation mode needs to be invoked to explain some of the observations.  The IMF for intermediate-mass and massive stars deduced from star counts in the field of a galaxy may therefore yield false information on the true shape of the stellar IMF, and stellar samples with well defined formation histories such as OB associations and star clusters are more useful to constrain the IMF. Photometric surveys of such regions, while essential for reaching faint stars, do not allow a reliable assessment of the mass distribution of massive stars since their spectral energy distribution is largely output at short wavelengths beyond the optical. Spectroscopic classification is therefore an essential tool for this purpose.  Phil Massey's work at Tucson, based on extensive spectroscopic classification of stars in OB associations and in star clusters, demonstrated that for massive stars $\\alpha =2.35\\pm 0.1, m\\simgreat10\\,M_\\odot$ \\citep{M98}, for a large variety of physical environments as found in the MW and the Magellanic Clouds, namely OB associations and dense clusters and for populations with metallicity ranging from near-solar-abundance to about $1/10$th metal abundance (Fig.~\\ref{fig:kroupa_figmassey} below).  The significant differences in the metallicity between the outer and the inner edge of the Galactic disk seem not to influence star-formation, as \\citet{YKT08} find no apparent difference in the stellar MFs of clusters in the extreme outer Galaxy compared to the rest of the disk down to $0.1\\,M_\\odot$.  \\subsection{The invariant IMF and its conflict with theory} \\label{sec:introd_IMF} The thus empirically constrained stellar IMF can be described well by a two-part power-law form with $\\alpha_1=1.3$ for $m\\simless 0.5\\,M_\\odot$ and with the {\\it Salpeter/Massey index}, $\\alpha_2=2.35$ for $m\\simgreat 0.5\\,M_\\odot$, being remarkably invariant.\\footnote{Note that \\cite{Sc98} emphasises that the IMF remains poorly constrained owing to the small number of massive stars in any one sample. This is a true albeit conservative stand-point, and the present authors prefer to accept Massey's result as a working IMF-universality hypothesis.  \\label{pk_fn1}.} This IMF is referred to as the {\\it canonical IMF} and is given in Sec.~\\ref{sec:canIMF} as the simpler two-part power-law form by Eq.~\\ref{eq:imf} and as the log-normal plus power-law form by Eq.~\\ref{eq:chabimf}.  The empirical result leads to the statement of the following hypothesis:  \\vspace{2mm} \\centerline{ \\fbox{\\parbox{\\columnwidth}{{\\sc Invariant IMF Hypothesis}: There exists a universal parent distribution function which describes the distribution of stellar masses in individual star-forming events.  }}} \\label{hyp:InvIMFHyp} \\vspace{2mm}  \\noindent The {\\it Invariant IMF Hypothesis} needs to be tested with star-count data in galactic fields and in individual clusters and OB associations for possible significant deviations. But it is mandatory to take into account the biases listed in Sec.~\\ref{sec:biases} when doing so.  The {\\it Invariant IMF Hypothesis} is not consistent with theory (see also Sec.~\\ref{sec:theor} and~\\ref{sec:IMFvar}): There are two broad theoretical ansatzes for the origin of stellar masses: \\vspace{0mm} \\begin{description} \\item {\\sc The Jeans-Mass Ansatz}: According to the Jeans-mass argument (e.g. \\citealt{Jeans02}, \\citealt{L98}, \\citealt{Bate2005a}, \\citealt{BCB06}, Eq.~\\ref{eq:MJeans} below) star-formation at lower metallicity ought to produce stars of, on average, heavier mass and thus an effectively top-heavy IMF (i.e. the ratio between the number of massive stars and low mass stars ought to increase).  At lower metallicity the cooling is less efficient causing larger Jeans masses as a requirement for gravitational collapse to a proto-star and thus larger stellar masses.  That warmer gas produces an IMF shifted to larger masses has been demonstrated with state-of-the art SPH simulations (e.g. \\citealt{KSJ07}). \\item {\\sc The Self-Regulatory Ansatz}: Another approach is formulated by \\cite{AF96} who argue that the {\\sc Jeans-Mass Ansatz} is invalid since there is no preferred Jeans mass in a turbulent molecular cloud. Instead they invoke the central-limit theorem\\footnote{\\label{foot:cenlimth} Citing from \\cite{BJ04}: ``According to the central limit theorem of statistics, if the mass of a protostellar condensation $M_c = f_1 \\times f_2 \\times \\dots \\times f_N$, then the distribution of $M_c$ tends to a lognormal regardless of the distributions of the individual physical parameters $f_i (i = 1, \\dots N)$, if $N$ is large. Depending on the specific distributions of the $f_i$, a convergence to a lognormal may even occur for moderate $N$.''  The central limit theorem was invoked for the first time by \\cite{Zinnecker84} to study the form of the IMF from hierarchical fragmentation of collapsing cloud cores.  } together with self-regulated assembly and they suggest that the final stellar masses are given by the balance between feedback energy from the forming star (accretion luminosity, outflows) and the rate of accretion from the proto-stellar envelope and circum-stellar disk. As the proto-star builds-up its luminosity increases until the accretion is shut off. When shut-off occurs depends on the accretion rate.  Indeed, \\cite{BJ04} explain the observed power-law extension of the IMF at large stellar masses as being due to a distribution of different accretion rates.  The self-regulating character of star formation has been studied profusely by the group around Christopher McKee (e.g. \\citealt{Tan06}) and has been shown to lead to decreasing star-formation efficiencies with increasing metallicity \\citep{Dibetal10,Dib11, Dibetal11}.  In low-metallicity gas the coupling of the photons to the gas is less efficient causing a less effective opposition against the accreting material. And, at lower metallicity the cooling is reduced causing a higher temperature of the gas and thus a higher speed of sound with a larger accretion rate. The final stellar IMF is expected to be populated by more massive stars in metal-poor environments. \\end{description} Both approaches can be refined by studying a distribution of physical conditions in a given star-forming cloud, but both lead to the same conclusion, namely that low-metallicity and high-temperature ought to produce top-heavy stellar IMFs. This leads to the following robust theoretical IMF result:  \\vspace{2mm} \\centerline{ \\fbox{\\parbox{\\columnwidth}{{\\sc The Variable IMF Prediciton}: Both the Jeans-mass and the self-regulation arguments invoke very different physical principles and yet they lead to the same result: The IMF ought to become top-heavy under low-metallicity and high-temperature star-forming conditions. \\label{box:varIMFpred}}}} \\vspace{2mm}  \\noindent Star-formation in the very early universe must have therefore produced top-heavy IMFs \\citep{L98,Bromm,Clark11}. But the samples of simple-stellar populations spanning all cosmological epochs (globular clusters to current star-formation in embedded clusters) available in the Local Group of galaxies have until recently not shown convincing evidence supporting {\\sc The Variable-IMF Prediction}. This issue is addressed in more detail in Sec.~\\ref{sec:currentstate}.   \\subsection{A philosophical note} \\label{sec:philosophy}  Much of the current discussion on star-formation, from the smallest to the largest (galaxy-wide) scales, can be categorised into two broad conceptual approaches which are related to the {\\sc Jeans-Mass Ansatz} {\\it vs} the {\\sc Self-Regulatory Ansatz} of Sec.~\\ref{sec:introd_IMF}:  \\begin{description} \\vspace{-3mm} \\item {\\sc Approach~A} is related to the notion that star formation is inherently stochastic such that the IMF is a probabilistic distribution function only. This is a natural notion under the argument that the processes governing star-formation are so many and complex that the outcome is essentially stochastic in nature. Followers of this line of reasoning argue for example that massive stars can form in isolation and that the mass of the most massive star cluster forming in a galaxy depends on the time-scale over which an ensemble of star clusters is considered (the {\\it size-of-sample effect}, even at very low SFRs a galaxy would produce a very massive star cluster if one waits long enough, i.e. if the sample of clusters is large enough). Approach~A can be formulated concisely as {\\sc nature plays dice when stars form}.  \\item {\\sc Approach~B} is related to the notion that nature is inherently self-regulated and deterministic. This is a natural notion given that physical processes must always depend on the boundary conditions which are a result of the physical processes at hand. An example of such would be gravitationally-driven growth processes with feedback in media with limited resources.  The emerging phenomena such as the distribution of stellar masses, of star-cluster masses and/or of how phase-space is populated to make binary stellar systems are, at least in principle, computable. They are computable in the sense that statistical mathematics provides the required tools such that the distribution functions used to describe the outcomes are subject to constrains. For example, a young stellar population of mass $M_{\\rm ecl}$ is excellently described by $\\xi(m)$ with the condition $m\\le m_{\\rm max}(M_{\\rm ecl})$. However, purely random sampling from $\\xi(m)$ even under this constraint will not reproduce a realistic population if nature follows {\\sc Optimal Sampling} (p.~\\pageref{box:optdistr}). This is because {\\sc Optimal Sampling} will never allow a cluster to be made up of $M_{\\rm ecl}/m_{\\rm max}$ stars of mass $m_{\\rm max}$, while constrained random sampling would.  Approach~B can be formulated concisely as {\\sc nature does not play dice when stars form}. \\end{description}  \\noindent Depending on which of the two notions is applied, the resulting astrophysical description of galaxies leads to very diverging results. Either a galaxy can be described as an object in which stars form purely stochastically such that the galaxy-wide IMF is equal to the stellar IMF (Approach~A). In this case a thousand small groups of 20~pre-main sequence stars will have the same stellar IMF as one very young star-cluster containing 20000 stars.  Or an embedded cluster or a galaxy are understood to be highly self-regulated systems such that the galaxy-wide IMF differs from the stellar IMF (Approach~B, Sec.~\\ref{sec:comppop}). According to this notion, a thousand small groups of 20~pre-main sequence stars would not contain a single star with $m>5\\,M_\\odot$, while a very young star cluster of~20000 stars would contain many such stars.  The different {\\sc Approaches} have very different implications for understanding the matter cycle in the universe.  \\subsection{Hypothesis testing} \\label{sec:hyptest}  The studies aimed at constraining the stellar IMF observationally typically have the goal of testing the {\\sc Invariant IMF Hypothesis} (p.~\\pageref{hyp:InvIMFHyp}) either in individual star-forming events such as in a star cluster, or on galaxy-wide scales. Here, it is important to be reminded of the following:  \\vspace{2mm} \\centerline{ \\fbox{\\parbox{\\columnwidth}{{\\sc Elementary Logics of Hypothesis Testing}: Negation of a hypothesis~I does not imply that the alternative hypothesis~II is correct \\label{box:hyptest}}}} \\vspace{2mm}  \\noindent By showing that hypothesis~I is consistent with some data does not imply that an alternative hypothesis~II is therewith ruled out.  A case in point is the discussion about dark matter and dark energy: Ruling out the standard cosmological (LCDM) model does not imply that any particular alternative is correct \\citep{Kroupa10, Kroupa12}. Conversely, ruling out a particular alternative does not imply that the LCDM model is correct \\citep{Wojtaketal11}.  Concerning the IMF, if a purely stochastic model ({\\sc Approach~A}, Sec.~\\ref{sec:philosophy}) is consistent with some observational data, then this does not imply that the alternative ({\\sc Optimal Sampling}, which is related to {\\sc Approach~B}) is falsified.  A case in point is provided by the following example relevant for the tests of the IGIMF theory on p.~\\pageref{box:igimfpred}): The masses of an ensemble of observed dwarf galaxies are calculated from spectral energy distribution modelling using the universal canonical IMF. These masses are then applied in testing a possible variation of the galaxy-wide IMF in terms of the UV and H$\\alpha$ flux ratios. If the universal IMF calculations allow for fluctuating SFRs whereas the variable IMF calculations do not, then the (wrong) conclusion of such an approach would plausibly be that nature appears to play dice, because observational data naturally contain measurement uncertainties which act as randomisation agents.  The consistent approach would instead be to compute the galaxy masses assuming a variable galaxy-wide IMF to test whether the hypothesis that the IMF varies systematically with galaxy mass can be discarded. The logically consistent procedure would be to calculate all fluxes within both scenarios independently of each other and assuming in both that the SFR can fluctuate, and to test these calculations against the observed fluxes. The result of this consistent procedure are opposite to those of the above inconsistent procedure in that the data are in better agreement with the systematically variable IMF, i.e. that nature does not play dice.  A final point to consider is is when a hypothesis ought to be finally discarded. Two examples illustrate this: Consider the Taurus-Auriga and Orion star forming clouds. Here the number of stars with $m\\simgreat 1\\,M_\\odot$ is significantly below the expectation from the purely stochastic model (see~box {\\sc IGIMF Predictions/Tests} on p.~\\pageref{box:igimfpred}).  This unambiguously falsifies the stochastic model.  But the data are in excellent agreement with the expectation from the IGIMF theory.  Is it then meaningful to nevertheless keep adopting the stochastic model on cluster and galaxy problems?  Another example, being perhaps more relevant to Sec.~\\ref{sec:philosophy}, is the issue with the current standard cosmological model. It is ruled out by \\cite{PN10, Kroupa12}. Should it nevertheless be adopted in further cosmological and related research?    \\subsection{About this text}  As is evident from the above introduction, the IMF may be well defined and is quite universal in each star-forming event out of which comes a spatially and temporally well correlated set of stars. But many such events will produce a summed IMF which may be different because the individual IMFs need to be added whereby the distribution of the star-formation events in mass, space and time become an issue. It thus emerges that it is necessary to distinguish between simple stellar populations and composite populations. Some definitions are useful:  \\vspace{4mm} \\centerline{ \\fbox{\\parbox{\\columnwidth}{ {\\sc Definitions}: \\\\[0.5mm] $\\bullet$ A {\\it simple population} obtains from a spatially ($\\simless\\,$few pc) and temporarily ($\\simless\\,$Myr) correlated star formation event (CSFE, also referred to as a {\\it collective star formation event}, being essentially an embedded star cluster). The mass of a CSFE may range from a few~$M_\\odot$ upwards.\\\\[0.5mm] $\\bullet$ A {\\it composite} or {\\it complex population} consists of more than one simple population.\\\\[0.5mm] $\\bullet$ The {\\it stellar IMF} refers to the IMF of stars in a simple population. \\\\[0.5mm] $\\bullet$ A {\\it composite} or {\\it integrated IMF} is the IMF of a composite or complex population, i.e. a population composed of many CSFEs, most of which may be gravitationally unbound. The galaxy-wide version is the IGIMF (Sec.~\\ref{sec:comppop}).  \\\\[0.5mm] $\\bullet$ The PDMF is the present-day MF of a stellar population not corrected for stellar evolution nor for losses through stellar deaths. Note that a {\\it canonical PDMF} is a PDMF derived from a canonical IMF.  \\\\[0.5mm] $\\bullet$ A {\\it stellar system} can be a multiple star or a single star. It has {\\it dynamical properties} (Footnote~\\ref{foot:dynprop} on p.~\\pageref{foot:dynprop}).  \\\\[0.5mm] $\\bullet$ The {\\it system} luminosity or mass function is the LF or MF obtained by counting all stellar systems.  The {\\it stellar} LF or {\\it stellar} MF is the true distribution of all stars in the sample, thereby counting all individual companions in multiple systems. This is also referred to as the {\\it individual}-star LF/MF. \\\\[0.5mm] $\\bullet$ Note 1$\\,$km/s$=1.0227\\,$pc/Myr, 1$\\,$g$\\,$cm$^{-3}=1.478\\times 10^{22}\\,M_\\odot\\,$pc$^{-3}$ and 1$\\,$g$\\,$cm$^{-2}=4788.4\\,M_\\odot\\,$pc$^{-2}$ for a solar-metallicity gas.  \\\\[0.5mm] $\\bullet$ The following additional abbreviations are used: SFR$=$star formation rate in units of $M_\\odot\\,$yr$^{-1}$; SFH$=$star formation history$=$SFR as a function of time; SFRD$=$star formation rate density in units of $M_\\odot\\,$yr$^{-1}\\,$pc$^{-3}$. \\label{box:definitions} }}} \\vspace{4mm}  \\noindent This treatise provides an overview of the general methods used to derive the IMF with special attention on the pitfalls that are typically encountered. The binary properties of stars and of brown dwarfs are discussed as well because they are essential to understand the true shape of the stellar IMF. While the stellar IMF appears to have emerged as being universal in star-formation events as currently found in the Local Group of galaxies, the recent realisation that star-clusters limit the mass spectrum of their stars is one form of IMF variation and has interesting implications for the formation of stars in a cluster and leads to the insight that composite populations must show IMFs that differ from the stellar IMF in each cluster.  With this finale, this treatise reaches the cosmological arena.  \\subsection{Other IMF reviews} \\label{sec:otherreviews}  The seminal contribution by \\cite{Sc86} on the IMF remains a necessary source for consultation on the fundamentals of the IMF problem.  The landmark review by \\cite{Massey03} on massive stars in the Local Group of galaxies is an essential read, as is the review by \\cite{ZH07} on massive-star formation.  Other reviews of the IMF are by \\cite{Sc98}, \\cite{K02}, \\cite{Chrev03}, \\cite{BLZ07}, \\cite{El09} and \\cite{Bastian2010}. The proceedings of the ``38th Herstmonceux Conference on the Stellar Initial Mass Function'' \\citep{GilHow98} and the proceedings of the ``IMF\\@50'' meeting in celebration of Ed Salpeter's 80th birthday \\citep{IMF50} contain a wealth of important contributions to the field. A recent major but also somewhat exclusive conference on the IMF was held from June~20th to 25th~2010 in Sedona, Arizona, for researchers to discuss the recently accumulating evidence for IMF variations: ``UP2010: Have Observations Revealed a Variable Upper End of the Initial Mass Function?''. The published contributions are a unique source of information on this problem \\citep{UP2010}. A comprehensive review of extreme star formation is available by \\cite{turner09}. ",
        "conclusions": "\\label{sec:concs}  Spectacular advances have been achieved over the past two decades in the field of IMF research and this affects a vast area of astrophysics.  The discrepant observed nearby and distant luminosity functions of stars in the local Galactic disk have been unified with one IMF, and the strong peak in the luminosity function near $M_V=12$ is well understood as a result of the changing internal constitution of stars as a function of their mass.  The unification of the luminosity functions also ultimately lead to a unification of the observed discrepant binary populations in different environments through dynamical processing in star-forming regions.  The observationally determined stellar IMF is described by Eq.~\\ref{eq:imf} and/\\-or~Eq.~\\ref{eq:chabimf}). Its form has traditionally been understood to be invariant in contradiction to the {\\sc Variable IMF Prediction} (p.~\\pageref{box:varIMFpred}).  The IMF of stellar systems and of the Galactic field is provided by Eqs.~\\ref{eq:sysimf} and~\\ref{eq:KTGimf}, respectively.  By studying the IMF star-formation theory is being tested.  According to the {\\sc IMF Origin Conjecture} (p.~\\pageref{conj:IMForigin}) the vast majority of stellar masses ($0.1 \\simless m/M_\\odot$ $\\simless$ ${\\rm few}\\,$) do not appear to be affected by competitive accretion or proto-stellar interactions (Sec.~\\ref{sec:cloudcore}). According to the computational star-formation research the most massive star in a star-formation event correlates physically with the mass of the event.  This $m_{\\rm max}-M_{\\rm ecl}$ relation is a natural outcome of the competitive accretion-growth and the fragmentation-induced starvation of stellar masses, and the computations reproduce the general shape of the IMF. Different ideas (competitive accretion, coagulation and simply the distribution of gravitationally unstable regions in turbulent clouds) all lead to virtually the same type of theoretical power-law IMF. The IMF appears to be mostly form invariant within space-time correlated star-formation events (CSFEs, i.e. in individual embedded star clusters) for star-formation rate densities $SFRD \\simless 0.1\\,M_\\odot\\,{\\rm pc}^{-3}\\,{\\rm yr}^{-1}$ within a spatial scale of about a~pc. But, the IMF varies (trivially) among such individual CSFEs of stellar mass $M_{\\rm ecl}$ through the $m_{\\rm max}-M_{\\rm ecl}$ relation. This relation follows from the notion of {\\sc Optimal Sampling}.  Computational star-formation has achieved a remarkable degree of realism, although the binary population has not yet emerged properly and the computations cannot yet reach the birth of populous star clusters.  That {\\sc Optimal Sampling} appears to be describing a freshly born stellar population would invalidate the concept of a stellar population being a purely random representation of the IMF, leading to {\\sc Open Question II} (p.~\\pageref{quest:openII}). The remarkable similarity of observationally determined IMF power-law indices ({\\sc Open Question III}, p.~\\pageref{quest:openIII}) may well be an indication that nature follows {\\sc Optimal Sampling}.  Two other {\\sc Open Questions} related to star formation and the IMF are furthermore stated (p.~\\pageref{quest:openI}, and \\pageref{quest:openIV}).  Evidence for a variation of the shape of the IMF has emerged, being consistent with the long-previously predicted IMF variation.  The evidence for top-heavy IMFs for $SFRD \\simgreat 0.1\\,M_\\odot\\,{\\rm pc}^{-3}\\,{\\rm yr}^{-1}$ (Eq.~\\ref{eq:toph_imf_density}) comes from either unresolved clusters or from populations that are very difficult to observe but appears to be increasingly established. The long-sought after evidence for a systematically varying stellar IMF with metallicity (Fig.~\\ref{fig:varIMFmetal}) may have emerged.  A principal-component-type analysis of the PDMFs of GCs has now for the first time yielded a formulation of a systematically varying IMF with star-forming cloud density and metallicity (Eq.~\\ref{eq:toph_master}).  Among other intriguing recent results are that BDs appear to be a distinct population from that of low-mass stars; their pairing properties have a different energy scale. This is well reproduced by computational star formation.  BDs and stars thus follow different mass distributions which do not join.  A continuous log-normal function across the VLMS/BD mass scale does not therefore correctly describe the IMF.  Furthermore, the IMF does appear to have a physical maximum stellar mass that has now been found empirically. Stars with $m\\simgreat150\\,M_\\odot$ do not appear to form, unless they implode invisibly shortly after being formed. Populous star clusters with super-canonical ($m>150\\,M_\\odot$) stars are termed to be super saturated (p.~\\pageref{box:saturated}). There is no statistically meaningful observational evidence for the formation of massive stars in isolation, in agreement with star-formation computations. The invariance of the observationally derived IMF precludes the exotic IMF associated with isolated massive star formation.  By realising that CSFEs are the true fundamental building blocks of a galaxy, all such events with their IMFs need to be added up to arrive at the integrated galactic initial mass function. This IGIMF varies in dependence of the SFR of the galaxy. According to the IGIMF theory, galaxies with low SFRs have a smaller ratio between the number of massive stars and low-mass stars than galaxies with high SFRs. This is the ``IGIMF effect''.  This formulation allows computation of the IGIMF as a function of time for galaxies with different SFRs.  One implication of this is that equally-old galaxies can have very different chemical compositions ranging from unevolved to evolved, and that the cosmological supernova~type~II rate per star would be significantly different and dependent on galaxy type than if an invariant stellar IMF is assumed. Galaxy-formation and evolution computations with this latter assumption are not likely to be correct. Indeed, it transpires that only with the IGIMF theory is a fundamental time-scale of about~3~Gyr uncovered on which all late-type galaxies consume their current gas supply. Very simple star-formation laws for galaxies emerge (Eq.~\\ref{eq:sflaws}).  Also, only with the IGIMF theory are the stellar-mass buildup times of dwarf galaxies consistent with the Hubble time. The top-heavy IMF of the Galactic Bulge deduced from chemical evolution research follows immediately from the IGIMF theory.  Thus, with the IGIMF theory it has now become possible to calculate how the observationally well-constrained star-formation on pc~scales propagates through to galactic scales of cosmological relevance. A unification of scales has therewith been achieved which was quite unthinkable only a few years ago.  Many details still need to be worked out though. For example, direct verification of the IGIMF effect is needed. This can be achieved by directly counting the number of massive stars in dwarf galaxies with low SFRs, or, by counting the number of massive stars in a giant molecular cloud and comparing this number with the number of late-type stars formed there.  Also, the existence of and exact form of the $m_{\\rm max}-M_{\\rm ecl}$ relation is important not only for the calculation of the IGIMF but also for understanding to which degree star-formation is self-regulated on a pc~scale. Such work will establish which of the two {\\sc Philosophical Approaches} of Sec.~\\ref{sec:philosophy} describe the astrophysics of star-formation and of galaxies.  The fact that the Orion South Cloud has formed a significant deficit of massive stars with half the cloud producing thousands of stars without a single massive star \\citep{Hsu12} and the fact that the $m_{\\rm max}-M_{\\rm ecl}$ relation is excellently mapped with a small scatter at the lowest masses \\citep{KM11} already constitute major evidence that the IGIMF theory and its fundamental assumptions appear to be given by nature.  Concerning the wider picture, the stellar and BD IMFs fit-in quite continuously into the universal mass function of condensed structures as pointed out by \\cite{BinggeliHascher07}.  \\vspace{5mm}  {\\bf Acknowledgements:} {\\small PK would like to thank Christopher Tout and Gerry Gilmore for very stimulating and important contributions without which much of this material would not have become available. PK is especially indebted to Sverre Aarseth whose friendly tutoring (against ``payments'' in the form of many bottles of {\\it lieblichen} German white wine) eased the numerical dynamics work in 1993/94.  Douglas Heggie be thanked for fruitful discussions with PK on {\\sc Optimal Sampling} in Heidelberg in Sept.~2011.  We thank Sambaran Banerjee for very helpful comments on the manuscript. PK would also like to thank M.R.S.~Hawkins who had introduced him to this field in~1987 whilst PK visited the Siding-Spring Observatory as a summer vacation scholar at the ANU. Mike gave PK a delightful lecture on the low-mass LF one night when visiting his observing run to learn about modern, state-of-the-art Schmidt-telescope surveying with photographic plates {\\it before} PK embarked on post-graduate work.  This research was much later supported through DFG grants KR1635/2, KR1635/3 and a Heisenberg fellowship, KR1635/4, KR1635/12, KR1635/13 and currently KR1635/25. MM acknowledges the Bonn/Cologne International Max-Planck Research School for support.}  \\clearpage \\lhead{} \\chead{}"
    },
    "1112/1112.3389_arXiv.txt": {
        "abstract": "The rotation of erupting filaments in the solar corona is addressed through a parametric simulation study of unstable, rotating flux ropes in bipolar force-free initial equilibrium. The Lorentz force due to the external shear field component and the relaxation of tension in the twisted field are the major contributors to the rotation in this model, while reconnection with the ambient field is of minor importance, due to the field's simple structure. In the low-beta corona, the rotation is not guided by the changing orientation of the vertical field component's polarity inversion line with height. The model yields strong initial rotations which saturate in the corona and differ qualitatively from the profile of rotation \\vs\\ height obtained in a recent simulation of an eruption without preexisting flux rope. Both major mechanisms writhe the flux rope axis, converting part of the initial twist helicity, and produce rotation profiles which, to a large part, are very similar in a range of shear-twist combinations. A difference lies in the tendency of twist-driven rotation to saturate at lower heights than shear-driven rotation. For parameters characteristic of the source regions of erupting filaments and coronal mass ejections, the shear field is found to be the dominant origin of rotations in the corona and to be required if the rotation reaches angles of order 90 degrees and higher; it dominates even if the twist exceeds the threshold of the helical kink instability. The contributions by shear and twist to the total rotation can be disentangled in the analysis of observations if the rotation and rise profiles are simultaneously compared with model calculations. The resulting twist estimate allows one to judge whether the helical kink instability occurred. This is demonstrated for the erupting prominence in the ``Cartwheel CME'' on 9~April 2008, which has shown a rotation of $\\approx115\\degrees$ up to a height of $1.5~R_\\odot$ above the photosphere. Out of a range of initial equilibria which include strongly kink-unstable (twist $\\Phi=5\\pi$), weakly kink-unstable ($\\Phi=3.5\\pi$), and kink-stable ($\\Phi=2.5\\pi$) configurations, only the evolution of the weakly kink-unstable flux rope matches the observations in their entirety. ",
        "introduction": "\\label{s:intro}   The geoeffectiveness of solar coronal mass ejections (CMEs) depends primarily on two parameters, the velocity and the magnetic orientation of the CME at the impact on the Earth's magnetosphere. The higher the CME velocity and the closer its front side magnetic field to a southward orientation, the more intense the interaction will typically be. Therefore, understanding the physics that determines these CME parameters at 1~AU is one of the key issues in space weather research. This involves the formation and main acceleration of the CME in the solar corona, as well as its propagation through the interplanetary space. The particulars of the trigger process also play a role in some events. It appears that typically the corona is the place where the basic decisions are made: will the CME be fast or slow, and will it keep the orientation given by the source, \\ie, will its magnetic axis remain oriented nearly parallel to the photospheric polarity inversion line (PIL), or will it rotate substantially?  In the present paper we employ the technique of MHD simulation to carry out a first systematic, but in view of the problem's complexity necessarily incomplete investigation of a number of processes that cause and influence changes of CME orientation in the corona. Such changes can be described as a rotation of the CME volume, more specifically of the magnetic axis of the flux rope in the CME, about the direction of ascent. This \\emph{rotation} should be distinguished from the possible rotation of the flux rope about its own axis, referred to as the \\emph{roll effect} \\cite{martin:2003,Panasenco&al2011}, which we do not address here.  Understanding the rotation of erupting flux ropes in the corona is also relevant for the question which processes trigger the eruptions, as a substantial rotation may indicate the occurrence of the helical kink instability (KI); see, \\eg, \\inlinecite{rust:1996}, \\inlinecite{Romano&al2003}, and \\inlinecite{rust:2005}. This instability is one of the candidate mechanisms for the initiation of CMEs \\cite{Sakurai1976,fan:2003,kliem:2004}. It commences when the twist of the rope exceeds a critical value, $\\Phi=2{\\pi}N>\\Phi_\\mathrm{cr}$, where $N$ is the winding number of the field lines about the rope's magnetic axis. The dynamical evolution of the instability has shown very good quantitative agreement with a number of well observed events, which range from confined filament eruptions to the fastest CME on record \\cite{torok:2005,williams:2005}. However, \\inlinecite{isenberg:2007} have pointed out an alternative mechanism for the rotation of line-tied flux ropes, which relies on the presence of an external toroidal field component, $B_\\mathrm{et}$, due to sources external to the current in the flux rope and pointing along the rope, \\ie, an external shear field component. The mechanism can easily be understood in the simplified picture of a current loop in vacuum field. When the loop legs move out of their equilibrium position to a more vertical orientation, the cross product of the loop current with the shear field component yields a sideways Lorentz force on the legs, which is antisymmetric with respect to the vertical line that passes through the apex of the loop. This torque forces the rising top part of the loop to rotate. The effect is also found in a full fluid description \\cite{lynch:2009}. For a given chirality of the erupting field, it yields the same direction of rotation as the helical kink. Hence, a comparative study of these two mechanisms is required before firm conclusions about the occurrence of the KI can be drawn from observations of flux rope rotations, which is a further main objective of this paper.  Since the rotations caused by the KI and by the external shear field point in the same direction, they are difficult to disentangle. In fact, from a more general perspective, they are of similar nature. Both cause a writhing of the flux rope which, by conservation of magnetic helicity, reduces the twist of the rope field lines about the writhing axis. Consequently, one can expect that observed flux rope rotations are often consistent with a range of $\\Phi$--$B_\\mathrm{et}$ parameter combinations which give the writhing of the flux rope by the helical kink and by the shear field different individual but similar combined strengths.  Other causes of flux rope rotation include magnetic reconnection with the ambient field \\cite{Jacobs&al2009,Shiota&al2010,Cohen&al2010,Thompson2011,Vourlidas&al2011} and the propagation through the overlying field. The latter comprises any asymmetric deflection of the rising flux from radial ascent, \\eg, by adjacent coronal holes (see, \\eg, \\opencite{Panasenco&al2011}), and the interaction with the heliospheric current sheet \\cite{Yurchyshyn2008,Yurchyshyn&al2009}.  One may conjecture that the generally changing orientation of the PIL with height in the corona acts similarly to the heliospheric current sheet at larger heights, \\ie, that the upper part of the rising flux continuously adjusts its orientation to align with the PIL. If this were the dominant effect, the rotation of erupting flux could be predicted rather straightforwardly from extrapolation of the photospheric field, since the overlying field is often close to the potential field. However, this conjecture is not valid in the lower corona where $\\beta\\ll1$, and where the main part of the total rotation often occurs. We demonstrate this in the Appendix.  The amount of rotation depends on the individual strengths of the four potentially contributing processes. Each of them is controlled by more than a single parameter. This is immediately obvious for the torque by the shear field, which must depend on the height \\emph{profile} $B_\\mathrm{et}(z)$, and for the reconnection, which is sensitive to the \\emph{structure} of the ambient field, \\ie, whether the field is bipolar, quadrupolar, or multipolar and whether the orientation of the line between the resulting new footpoints of the erupting flux differs strongly from the original orientation. The rotation by the KI does not only depend on the initial flux rope twist, $\\Phi-\\Phi_\\mathrm{cr}$, but also on the strength and height profile of the overlying field \\cite{Torok&al2010}. If the overlying field decreases only slowly with height, then the upward expansion develops slowly and, accordingly, its contribution to the relaxation of the field line tension is initially weak. The relaxation is then primarily accomplished by a strong rotation at small heights. In the opposite case of very strong upward expansion, the rotation is distributed across a large height range, which also increases the likelihood of further changes by the onset of reconnection (see \\opencite{Lugaz&al2011} for an example). The effect of the heliospheric current sheet can be expected to depend on the angle with the top section of the flux rope's axis, on the horizontal elongation of the CME (whether its horizontal cross section is very elliptical or more nearly circular), and on the magnetic pressure of the CME relative to the pressure of the interplanetary plasma.  Moreover, the total rotation experienced by an erupting flux rope likely depends also on the dynamics of its evolution. For example, a torque strongly localized at low heights, operating on a still small loop, may hurl the flux around more efficiently than a torque which is distributed across a large height range. As another example, in a complex (multipolar) coronal environment the sequence and strength of reconnection with the ambient field may strongly depend upon the height profiles of the rope's angular and rise velocities caused by other processes, \\eg, by an ideal MHD instability. The relative velocity between reconnecting flux systems controls how strongly the reconnection with the ambient field is driven. Hence, quantitative studies of flux rope rotation face a very high degree of complexity.  Here we focus on two mechanisms that can cause strong rotations in the corona, the helical kink instability and the torque exerted by an external shear field component. By comparing a parametric study of both mechanisms in a force-free, line-tied flux rope equilibrium with the data of a well observed, strong\\-ly rotating erupting prominence, we demonstrate that their contributions can be disentangled to some degree. We also demonstrate the very strong influence of the ambient potential field's height profile on the amount of rotation by the KI, and briefly address the influence of reconnection between the CME flux rope and the ambient field on the rotation.  \\begin{figure} \\begin{center} \\includegraphics[width=0.9\\textwidth]{euvi_cor1_20080409samples.eps} \\end{center} \\caption[] {Images and plots of the prominence eruption at 10:16~UT, as seen by the EUVI-Ahead telescope in the 304~{\\AA} channel, and at 10:55 and 11:25~UT, as seen in white light by COR1-Ahead. The prominence apex has reached heights of 0.56, 1.6, and $2.3~R_\\odot$ above the photosphere at these times. The right panels display the reconstructed three-dimensional position of the marked prominence threads, using a reprojection to a viewpoint at the position of radial CME propagation, Stonyhurst longitude $98\\degrees$ west (relative to Earth) and latitude $24\\degrees$ south, where the counterclockwise rotation is apparent. The axes are in units of solar radii.} \\label{f:obs} \\end{figure}  This investigation was stimulated by the analysis of the strong rotation in a prominence eruption and CME on 9~April 2008, occasionally referred to as the ``Cartwheel CME'', in \\citeauthor{Thompson&al2011} (\\citeyear{Thompson&al2011}, in the following Paper~I). Their stereoscopic reconstruction revealed the height-rotation profile of the erupting filament/prominence in the core of a CME for the first time \\cite{Thompson&al2009}. This profile provides a strong constraint for the numerical modeling. In combination with the further observations of the event, it allows us to infer the major causes of the rotation and the range of source parameters compatible with the data. The analysis of Paper~I has given the following results relevant for the present study. The prominence erupted from the remnants of NOAA active region (AR) 10989 close to the west limb and appeared as a flux rope -- a single, weakly to moderately twisted loop -- throughout the height range covered by the STEREO EUVI and COR1 telescopes \\cite{howard:2008}, \\ie, up to $4~R_\\odot$ from Sun center. It rotated counterclockwise by $\\approx\\!115\\degrees$ up to a heliocentric height of $2.5~R_\\odot$, where the rotation leveled off. Two thirds of this rotation were acquired within $0.5~R_\\odot$ from the photosphere. The data indicate a subsequent gentle backward rotation by $\\approx\\!15\\degrees$ in the height range up to $3.3~R_\\odot$. In addition, the analysis of STEREO COR2 data in \\inlinecite{Patsourakos&Vourlidas2011} demonstrated that a flux rope structure is consistent also with the three-dimensional shape of the CME at a heliocentric distance of $13~R_\\odot$, where it had changed its orientation by a total of $150\\degrees\\pm7\\degrees$ from the original one, most likely by further counterclockwise rotation. At this stage the erupting flux was very closely aligned with the heliospheric current sheet above the active region. The prominence was initially accelerated mainly horizontally along the filament channel. This gradually turned into a radial propagation at a position $\\approx$\\,98W24S as seen from Earth, $15\\degrees\\mbox{--}20\\degrees$ away from the original location. The prominence experienced most of its upward acceleration in the heliocentric height range up to $\\sim\\!2.5~R_\\odot$ and reached a velocity of $\\sim\\!400$~km\\,s$^{-1}$ in the COR2 field of view. At the same time, the leading edge of the CME accelerated to over 700~km\\,s$^{-1}$ \\cite{Landi&al2010}. Representative images of the prominence from STEREO \\emph{Ahead}, which had the best perspective at the structure, and the corresponding three-dimensional reconstructions of the location of several prominence threads are compiled in Figure~\\ref{f:obs} (from Paper~I). The rotation (height-rotation) profile and the rise (time-height) profile are included below in the observation-simulation comparisons (Figures~\\ref{f:rotation} and \\ref{f:rise}, respectively).  As already noted above, we focus our attention here on the coronal evolution of this event, leaving the interaction with the heliospheric current sheet for future investigation. Moreover, we exclude the possible slight backward rotation by $\\approx15\\degrees$ in the COR1 height range from our modeling, since we are interested in the generally important effects which cause significant rotations in the corona. This part of the rotation is not fully certain, and, if real, it was likely caused by the particular structure of the large-scale coronal field above the active region, which nearly reversed its horizontal direction at heights $\\gtrsim\\!0.3~R_\\odot$ above the photosphere (Paper~I). Thus, we will consider a saturation of the modeled rotation at angles near 115\\degrees\\ and heights $h\\approx(1.5\\mbox{--}2.3)~R_\\odot$ above the photosphere to be in agreement with the observation data. Furthermore, we will disregard the initial nearly horizontal motion of the prominence along the PIL.  The combined effects of flux dispersal and foreshortening in the course of the source region's rotation to the solar limb made it impossible to obtain a well-defined estimate of the distance between the main flux concentrations in the bipolar region at the time of the eruption, which is a parameter of strong influence on the height profile of the ambient potential field. Only a relatively wide range of $\\sim\\!(40\\mbox{--}150)$~Mm could be estimated by extrapolating the region's evolution in the course of its disk passage through the final three days before the event. It will be seen that this range still sets a useful constraint on the modeling.  In the following we model the radial propagation of the prominence in the Cartwheel CME in the coronal range of heights as the MHD evolution of an unstable force-free and line-tied flux rope (Section~\\ref{s:model}). A parametric study of the resulting rotation and rise, focusing on the rotation caused by the helical kink instability and by the external shear field, is compared with the observation data, to constrain the parameters in the source of the event and to study whether the relative importance of these mechanisms can be disentangled and individually estimated (Section~\\ref{s:comparison}). The discussion in Section~\\ref{s:discussion} addresses the simplifying assumptions made in the modeling and differences to earlier relevant work. Section~\\ref{s:conclusions} gives our conclusions. The Appendix relates the rotation of erupting flux ropes in low-beta plasma to the changing orientation of the PIL with height. ",
        "conclusions": "\\label{s:conclusions}   Our parametric study of force-free flux ropes which erupt from simple bipolar source regions with no overlying current sheet and rotate about the direction of ascent yields the following conclusions.  1) Both the force by an external shear field component $B_\\mathrm{et}$ \\cite{isenberg:2007} and the relaxation of twist $\\Phi$ (\\eg, \\opencite{Torok&al2010}), are potentially very significant contributors to the rotation.  2) For parameters typical of CME source regions, in particular if the sources of the external stabilizing field (usually the main flux concentrations next to the PIL) have a smaller distance than the footpoints of the erupting flux, the shear field yields the dominant contribution to the rotation for a wide range of shear field strengths. The relaxation of twist remains the weaker contributor under these conditions, even if it is sufficiently high to trigger the helical kink instability. However, since twist always exists in force-free flux ropes, it always causes at least some rotation. Strong rotations ($\\gtrsim90\\degrees$) can be produced by the twist alone, but only for considerably larger distances between the sources of the external stabilizing field than typically observed.  3) The rotation in low-beta plasma is not guided by the changing orientation of the PIL with height. For the geometrical conditions typical of CME source regions, it is opposite in direction (see the Appendix).  4) For a given chirality of the configuration, the external shear field and the twist cause flux rope rotation in the same direction, which is clockwise for right-handed field and counterclockwise for left-handed field if seen from above.  5) The two processes are related to each other when considered in terms of magnetic helicity. Both convert initial twist helicity of the flux rope into writhe helicity. The same total rotation, and rotation profiles which are very similar in a substantial part of the total height range of rotation, result in a range of $B_\\mathrm{et}$--$\\Phi$ combinations.  6) The rotation due to twist relaxation tends to act mainly low in the corona, in a height range up to only a few times the distance between the footpoints of the erupting flux. The rotation by the shear field tends to be distributed across a larger height range.  7) The mere fact that erupting flux rotates does not by itself imply that the helical kink instability occurred.  8) The relative contributions to the total rotation by the shear field and by the twist can be disentangled by comparing both the observed rotation and rise profiles with the corresponding curves from a model, since these profiles possess a different dependence upon the $B_\\mathrm{et}$--$\\Phi$ parameter combination. The resulting estimate for the twist allows one to judge the occurrence of the helical kink instability.  9) Magnetic reconnection contributes only weakly (much less than the shear field) to the total rotation in the simple bipolar source regions considered.  \\medskip From the comparison with the simulation of rotating flux in \\inlinecite{lynch:2009} we conclude:  10) The rotation profile differs strongly between configurations with and without a flux rope at the onset of the eruption.  \\medskip The comparison with the stereoscopic observations and three-dimensional reconstruction of the erupting prominence in the 9~April 2008 ``Cartwheel CME'' additionally shows the following.  11) The rotation profile obtained in Paper~I from the stereoscopic reconstruction of STEREO data is equally well reproduced by our model up to heights $\\approx\\!1.5~R_\\odot$ above the photosphere for a range of $\\Phi$--$B_\\mathrm{et}$ combinations which include a strongly kink-unstable case ($\\Phi=5\\pi$, $B_\\mathrm{et}/B_\\mathrm{ep}=0.42$), a weakly kink-unstable case ($\\Phi=3.5\\pi$, $B_\\mathrm{et}/B_\\mathrm{ep}=0.67$), and a kink-stable case ($\\Phi=2.5\\pi$, $B_\\mathrm{et}/B_\\mathrm{ep}=1.06$). However, the strongly kink-unstable configuration is ruled out by the simultaneous consideration of the rise profile, and several features of the kink-stable model argue strongly against this configuration. These are the implied high value of the shear field, the rotation profile at greater heights, and the unrealistic start height of the unstable rise of $\\approx\\!0.2~R_\\odot$. Hence, the occurrence of a weak helical kink instability in the Cartwheel event is very likely.  \\medskip Our results add to the complexity of the phenomenon of flux rope rotation in eruptions which is already known from investigations that focused on the influence of reconnection (\\eg, \\opencite{Jacobs&al2009}; \\opencite{Shiota&al2010}; \\opencite{Cohen&al2010}; \\opencite{Thompson2011}; \\opencite{Lugaz&al2011}). An overall very complicated dependence on several parameters and on the structure of the ambient field is revealed. Thus, the quantitative prediction of the rotation is a difficult task. The parametric study performed here indicates for simple bipolar source regions that the strength of the external shear field is the primary parameter determining the total rotation. The twist and the height profile of the external poloidal field are of relatively minor importance as long as they stay in the typical ranges indicated by the observations. We did not yet study a possible influence of the height profile of the external shear field. The external shear field of filament channels may be estimated to sufficient precision from a simple linear force-free field extrapolation. It will be worth testing whether numerical modeling starting from such fields, embedded in current-free outer field, yields rotations in agreement with observations of eruptions from bipolar source regions.  Several investigations indicate that erupting flux ropes align with the heliospheric current sheet in the course of their interplanetary propagation (\\eg, \\opencite{Bothmer&Schwenn1998}; \\opencite{Yurchyshyn2008}; Paper~I). This suggests that the coronal rotation merely decides whether a parallel or an antiparallel alignment will result at 1~AU. However, since complex physics is involved and since rotations on the order of 90\\degrees\\ may not be rare, the quantitative study of the effects that determine the rotation in the corona remains of high scientific and practical interest."
    },
    "1112/1112.3162_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{intro}  \\begin{figure*}[h] \\begin{center} \\includegraphics[width=16.5cm]{testcube.pdf} \\caption{Properties of the test noise cube. \\it Left: \\rm Right ascension-velocity slice going through the centre of the cube. Artefacts are visible as faint vertical stripes. \\it Right: \\rm Histogram of voxel values for the entire cube (black line) and a Gaussian distribution with $\\sigma=1.6$ mJy beam$^{-1}$ (red line).} \\label{fig0} \\end{center} \\end{figure*}  \\begin{figure*}[h] \\begin{center} \\includegraphics[width=16.5cm]{skyN_4.0_scat.pdf}  \\includegraphics[width=16.5cm]{skyN_4.0_cont_all.pdf}  \\includegraphics[width=16.5cm]{skyN_4.0_cont.pdf} \\caption{Distribution of detections in the test \\hi\\ cube on all projections of the adopted parameter space (see text). \\it Top\\rm: Positive (blue) and negative (red) detections. \\it Middle\\rm: True (black circles) and false (grey crosses) positive detections. \\it Bottom\\rm: Same as middle panels, but showing detections with $R>0.99$ only. We also show constant surface-density contours of positive (blue) and negative (red) detections estimated from the distributions in the top panels as described in the text.} \\label{fig1} \\end{center} \\end{figure*}   \\begin{figure*}[h] \\begin{center} \\includegraphics[width=16.5cm]{rfi_4.0_cont.pdf} \\caption{Constant surface-density contours of positive (blue) and negative (red) detections for the datacube with RFI. Grey crosses indicate sources with $R>0.99$.} \\label{fig2} \\end{center} \\end{figure*} In the coming years, a number of interferometric neutral-hydrogen (\\hi) surveys will begin \\citep[e.g.][]{2009pra..confE..14S,2009pra..confE..10V}. They will observe \\hi\\ within large cosmic volumes and detect tens of thousands of sources, many of which will be resolved both on the sky and in velocity.  These surveys will rely on automated source finders to detect objects present in the data. In particular, the detection of the faintest objects will require detection criteria close to the noise level. However, as fainter and fainter ``true'' objects are detected an increasing number of ``false'' detections will be included in the source catalogue (i.e., detections that are, in fact, noise peaks). Quantifying this effect is crucial to enable a proper scientific exploitation of the final \\hi\\ catalogues.  A quantity often used for this purpose is the reliability $R$ of a source catalogue. This is defined as:  \\begin{equation} R=\\frac{T}{T+F}, \\label{eq1} \\end{equation}  \\noindent where $T$ and $F$ are the number of true and false detections, respectively. Normally, the price to pay for detecting faint objects is a decrease in $R$.  In some cases, a single value of $R$ may be used to characterise an entire source catalogue. However, it is more informative to study the $n$-dimensional function $R(p_1,p_2,...,p_n)$ where the $p_i$'s are a set of source parameters. For example, $R$ may be given as a function of objects' total flux and line width.  There are many ways of measuring $R$. \\cite{2004MNRAS.350.1210Z} estimate the reliability of the HIPASS catalogue \\citep{2004MNRAS.350.1195M} as a function of source total flux, peak flux and line-width by re-observing a sub-sample of the detected objects. They label confirmed detections as true and non-confirmed detection as false, and adopt a formalism equivalent to Eq. \\ref{eq1} to estimate $R$. Unfortunately, this empirical procedure may not always be practical and it requires sources to be re-observed with at least the same data quality of the original observations.  Another technique is to create a dataset where model sources are injected on top of modelled (or observed) noise and run a source finder as one would do with the real data \\citep[e.g.][]{2007ApJS..169..401K,2007AJ....133.2087S}. Detections corresponding to an input source are labelled as true and detections not corresponding to an input source are labelled as false. This approach gives a correct estimate of $R$ only if the model noise is a good approximation of the real noise and if model sources are representative of the objects actually contained in the data.  Here we discuss yet another method to measure $R$ based on the detection of ``negative'' sources, i.e., sources with negative total flux. This technique has been used in various forms by several authors working in different fields \\citep[e.g.][]{2004ApJ...600L..99D,2004ApJ...612L..93Y,2009MNRAS.400..743K}. In this paper we develop it further with the aim of making it useful for future \\hi\\ surveys.  The main idea is to assume that true sources are ``positive'' (i.e., they have positive total flux) and that the noise is symmetric (we discuss the applicability of these assumptions in Section \\ref{fut}). It follows that the number of false positive detections equals the number of negative detections. The reliability can then be defined as:  \\begin{equation} R=\\frac{P-N}{P}, \\label{eq2} \\end{equation}  \\noindent where $P$ and $N$ are the number of positive and negative detections, respectively. It is trivial to verify that Eq. \\ref{eq2} is equivalent to Eq. \\ref{eq1} under the aforementioned assumptions.  The advantage of this method is that $R$ is measured directly from the data with no additional observational or modelling effort. In what follows we demonstrate this technique by applying it to a test \\hi\\ cube. We describe the cube and the source finder used for this purpose in Sections \\ref{dataset} and \\ref{sourcefinder}, respectively. In Section \\ref{res} we illustrate the results. In Section \\ref{fut} we discuss possible caveats and improvements of this technique. We draw conclusions in Section \\ref{concl}. ",
        "conclusions": "\\label{concl}  We discuss a method to determine the reliability of sources detected in \\hi\\ cubes. We assume that true sources are positive and that the noise is symmetric. It follows that the number of negative detections equals the number of positive false detections. Negative detections can therefore be used to estimate the reliability $R$ of positive detections as a function of their position in a chosen source parameter space.  We demonstrate this method by running a smooth-and-clip source finder on a test \\hi\\ cube containing real interferometer noise and real \\hi\\ sources. We show that sources with $R>0.99$ are true. The volume of parameter space where this simple method gives $R>0.99$ is essentially the same which we would have selected knowing which source is true and which false in this test cube.  We discuss the applicability of this method to \\hi\\ cubes with artefacts. We show that at least in the analysed case of a cube with RFI the method performs well. The reason is that artefacts in interferometric images tend to be both positive and negative, so that they do not necessarily invalidate the noise symmetry assumption."
    },
    "1112/1112.3448_arXiv.txt": {
        "abstract": "Despite the sphericity of the collapsing stellar core, the birth conditions of neutron stars can be highly non spherical due to a hydrodynamical instability of the shocked accretion flow. Here we report the first laboratory experiment of a shallow water analogue, based on the physics of hydraulic jumps. Both the experiment and its shallow water modeling demonstrate a robust linear instability and nonlinear properties of symmetry breaking, in a system which is one million times smaller and {about} hundred times slower than its astrophysical analogue. \\begin{description} \\item[PACS numbers]97.60.Bw, 95.30.Lz, 47.20.-k, 01.50.Pa \\end{description} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.4268_arXiv.txt": {
        "abstract": "In this paper, we present a simulation method within the two-component spherical collapse model to investigate dark energy perturbations associated with the formation of dark matter halos. The realistic mass accretion history of a dark matter halo taking into account its fast and slow growth is considered by imposing suitable initial conditions and isotropized virializations for the spherical collapse process. The dark energy component is treated as a perfect fluid described by two important parameters, the equation of state parameter $w$ and the sound speed $c_s$. Quintessence models with $w>-1$ are analyzed. We adopt the Newtonian gauge to describe the spacetime which is perturbed mainly by the formation of a dark matter halo. It is found that the dark energy density perturbation $\\delta_{DE}$ depends on $w$ and $c_s$, and its behavior follows closely the gravitational potential $\\Phi$ of the dark matter halo with $\\delta_{DE}\\approx -(1+w)\\Phi/c_s^2$. For $w>-1$, the dark energy perturbation presents a clustering behavior with $\\delta_{DE}>0$ during the entire formation of the dark matter halo, from linear to nonlinear and virialized stages. The value of $\\delta_{DE}$ increases with the increase of the halo mass. For a cluster of mass $M\\sim 10^{15} M_{\\odot}$, $\\delta_{DE}\\sim 10^{-5}$ within the virialized region for $c_s^2 \\in [0.5, 1]$, and it can reach $\\delta_{DE}=O(1)$ with $c_s^2=0.00001$. For a scalar-field dark energy model, we find that with suitably modeled $w$ and $c_s$, its perturbation behavior associated with the nonlinear formation of dark matter halos can well be analyzed using the fluid approach, demonstrating the validity of the fluid description for dark energy even considering its perturbation in the stage of nonlinear dark matter structure formation. ",
        "introduction": "Since the discovery of the accelerating expansion of the universe, tremendous progress has been achieved in dark energy studies\\cite{Amanullah10}\\cite{Percival10}\\cite{Reid10}\\cite{Larson11}\\cite{Copeland06}\\cite{Frieman08}. Future generations of cosmological observations are expected to be able to provide us tight constraints on properties of dark energy \\cite{Zhan06}\\cite{Barnard08}\\cite{Zhan09}\\cite{Zhao10}\\cite{Mortonson10}. The full realization of their constraining power, however, depends on our thorough understandings on how different dark energy models affect the expansion of the universe and the formation and evolution of cosmic structures differently \\cite{Ma99}\\cite{Klpyin03}.  For dynamical dark energy models, apart from the background field that drives the accelerating expansion of the universe, dark energy perturbations exist intrinsically. Their behaviors on large scales and the effects on the linear power spectrum of the matter density perturbations and the anisotropy of the cosmic microwave background radiation (CMB) have been extensively investigated and taken into account in deriving cosmological constraints from CMB observations \\cite{GZhao05}\\cite{Li08b}\\cite{Komatsu11}\\cite{Fang08}. On small scales of the nonlinear structure formation, they should also be considered. Therefore a fully consistent simulation to study the impacts of dark energy on the structure formation should simulate the evolution of the dark energy component explicitly together with the matter components (dark matter and baryonic matter). The additional dark energy component can however complicate the numerical simulation significantly. Therefore in many studies, the dark energy is not included as an independent component in simulations. Instead, approximations or simplifications suitable to specific models are often adopted so that the effects of dark energy perturbations on the structure formation can be partially taken into account \\cite{Maccio04}\\cite{Baldi09}\\cite{Baldi10a}\\cite{Baldi10b}\\cite{Baldi11}\\cite{Alimi10}. In a very recent series of papers regarding coupled dark matter-dark energy models, $f(R)$ gravity models and scalar-tensor theories, the corresponding scalar field component has been explicitly included in simulations \\cite{BLi09}\\cite{BLi11a}\\cite{BLi11b}\\cite{BLi11c}\\cite{GZhao11}. Such simulations not only can study the structure formation self-consistently, but also can provide the detailed dynamical evolution of the relevant scalar field component, which may contain additional information in differentiating different models.  To study the effects of dark energy and particularly the behavior of dark energy perturbations along with the formation of nonlinear structures, the analyses under the spherical collapse model are extensively performed. Although it is a simplified model comparing with the full cosmological simulations, it allows us to isolate different effects and to explore a broad parameter space efficiently. Studies along this line have shown that the simple extrapolation from the amplitude of the dark energy perturbation obtained in the linear stage of structure formation can grossly overestimate the level of the dark energy perturbation in the nonlinear epoch of the structure formation \\cite{Dutta07}\\cite{Mota08}\\cite{Wang09}. For minimally-coupled quintessence dark energy models with $c_s^2\\sim 1$, the dark energy perturbation induced by the formation of cluster-scale dark matter halos remains at the level of $< 10^{-5}$, in contrast with the amplitude of $\\sim 10^{-2}$ obtained by extrapolating the linear analyses to late stages of dark matter halo formation where nonlinear processes actually happen\\cite{Dutta07}\\cite{Mota08}\\cite{Wang09}\\cite{Maor05}\\cite{Nunes06}\\cite{Abramo07}\\cite{Basse11}\\cite{Sefusatti11}\\cite{Devi11}\\cite{He09}\\cite{Creminell10}\\cite{He10}. Such a large difference indicates that the virialization of dark matter halos plays an important role in determining the behavior of dark energy perturbations in the halo region. Previous analyses take into account the virialization process by adding in a simple prescription into the analytical formulation to prevent the halo material from collapsing all the way to the center (e.g., \\cite{Mota08}\\cite{Wang09}). Although such a modeling can catch the basics of the virialization process, it cannot describe it very realistically. Particularly, it cannot handle properly the shell crossing that is important for the virialization process.  To overcome the shortcomings of the artificial treatment of the virialization process, in this paper, we carry out numerical simulations for a set of two-component spherical systems. Following Lu et al. \\cite{Lu06}, by specifying proper initial conditions, the realistic mass assembling history for a dark matter halo is considered in the spherical 1-D simulations. The tangential velocity is also taken into account,  which can affect the final density profile of a dark matter halo significantly. With these implementations, the formation of dark matter halos can be suitably simulated, and the corresponding dark energy perturbations can be investigated. In our simulations, the dark energy component is regarded as an ideal fluid specified by two important parameters, the equation of state parameter $w$ and the sound speed parameter $c_s$. We investigate the dependence of dark energy perturbations on various quantities, such as $w$, $c_s$ and the mass of the dark matter halo. The validity of the fluid approach taking into account the dark energy perturbations is also analyzed.  The rest of the paper is organized as follows. In Sec.~\\ref{sec:formula}, we present the formulations related to the two-component spherical system. Sec.~\\ref{sec:simulation} describes the methodology of simulations. The results and discussions are shown in Sec.~\\ref{sec:result} and Sec.~\\ref{sec:discussion}, respectively. ",
        "conclusions": "}  With two-component 1-D numerical simulations taking into account the fast and slow growth of dark matter halos, we analyze the behavior of dark energy perturbations induced by the formation of spherical dark matter halos. By comparing the results calculated directly from the scalar field evolution and that from the fluid description for a quintessence dark energy model with a double-exponential potential, we show that the fluid approach can be used to analyze the dark energy behavior very well even in the nonlinear stage of structure formation. In the fluid treatment for dark energy, the equation of state $w$ and the sound speed $c_s$ are the two important parameters that affect the dark energy perturbations significantly. Our studies find that in general, the dark energy perturbation arising from the formation of a dark matter halo has a much more extended profile than that of the dark matter halo except for the case $c_s\\approx 0$ where the dark energy component follows the dark matter closely. A relation of $\\delta_{DE}\\approx [-(1+w)/c_s^2] \\Phi$ with an accuracy about $1-2\\%$ within the virial radius of a halo is revealed, which provides us a potential means to estimate the power spectrum of dark energy perturbations from that of the dark matter potential field. From $w=-0.9$ to $w=-0.7$, $\\delta_{DE}$ increases about three times from $\\sim 2.5\\times 10^{-5}$ to $7.5\\times 10^{-5}$ for $c_s^2=0.5$. On the other hand, varying $c_s^2$ from $0.5$ to $0.05$, the dark energy perturbation increases by an order of magnitude from $\\sim 10^{-5}$ to $\\sim 10^{-4}$. Our analyses show that our simulation treatment and the linear evolution of dark energy perturbations are valid for $c_s^2$ down to $c_s^2\\sim 0.00001$. For even smaller $c_s^2$, more accurate analyses taking into account nonlinear dark energy perturbations and their feedback effects on the formation of dark matter halos are needed (e.g.,\\cite{Creminell10}). The dependence of the dark energy perturbation on the mass of the halo is also analyzed. It is shown that $\\delta_{DE}$ increases by $\\sim 20$ times from $M\\approx 10^{13}\\hbox{ M}_{\\odot}$ to $M\\approx 10^{15}\\hbox{ M}_{\\odot}$.  The numerical analyses presented in this paper are done for dark energy models with $w >-1$. Our simulations show that for such models, the dark energy perturbation induced by the formation of dark matter halos has a clustering behavior with $\\delta_{DE}>0$ during the entire evolutionary process. This can also be seen clearly from the approximate relation $\\delta_{DE}\\approx [-(1+w)/c_s^2] \\Phi$ (note $\\Phi<0$ in dark matter halo regions). On the other hand, for $w< -1$, dark energy voids are expected from this relation although we do not study these cases explicitly in our simulations. These results are in good agreement with studies shown in, e.g., \\cite{Abramo07}\\cite{Basse11}. However, there are other analyses that point to the existence of dark energy void with $\\delta_{DE}<0$ during the quasi-linear stages of dark matter halo formation even for models with $w> -1$ (e.g., \\cite{Dutta07}\\cite{Mota08}\\cite{Wang09}). It is noticed that all the latter studies are performed using the Lemaitre-Tolman-Bondi (LTB) metric that is in the synchronous gauge. On the other hand, our simulation analyses presented here adopt the Newtonian gauge. It is well known that cosmological energy density perturbations are gauge dependent (e.g., \\cite{Bardeen80}\\cite{Ma95}\\cite{Christopherson10}). Concerning the dark energy density perturbation in the LTB synchronous gauge $\\delta_{DE}(Syn)$ and in the Newtonian gauge $\\delta_{DE}(New)$, we have the relation \\cite{Christopherson10}\\cite{Ma95} \\begin{equation} \\delta_{DE}(Syn) = \\delta_{DE}(New)-\\alpha\\frac{\\dot {\\bar {\\rho}}_{DE}}{\\bar {\\rho}_{DE}}, \\label{eq:syn_new_1} \\end{equation} where $\\alpha(t,\\vec x)$ is related to the coordinate transformation for the time component associated with the two gauges, $t(Syn)=t(New)+\\alpha$, and $\\partial \\alpha /\\partial t=\\Phi$. We have ${\\dot {\\bar {\\rho}}_{DE}}/{\\bar {\\rho}_{DE}}=-3(1+w)H$. Further with the approximate relation between $\\delta_{DE}(New)$ and $\\Phi$, we then obtain \\begin{equation} \\delta_{DE}(Syn) \\approx \\delta_{DE}(New)-3(1+w) H c_s^2 \\int \\frac{\\delta_{DE}(New)}{1+w} dt. \\label{eq:syn_new_2} \\end{equation} Thus if the second term is larger than the first term, a dark energy void with $\\delta_{DE}(Syn)<0$ can occur in the LTB synchronous gauge even for $w> -1$ with $\\delta_{DE}(New)>0$. In other words, our results about the dark energy clustering for $w> -1$ presented in this paper do not conflict with those showing the existence of dark energy voids for the same models. Instead, the differences merely reflect the differences of the specific gauge used in different analyses."
    },
    "1112/1112.1398_arXiv.txt": {
        "abstract": "We report the simultaneous extreme ultraviolet observation of magnetic reconnection inflow and outflow in a flare on 2010 August 18 observed by the Atmospheric Imaging Assembly on board the {\\it Solar Dynamics Observatory}. We found that during the rise phase of the flare, some plasma blobs appeared in the sheet structure above the hot loops.  The plasma blobs were ejected bidirectionally along the sheet structure (outflow), at the same time as the threads visible in extreme ultraviolet images moved toward the sheet structure (inflow).  The upward and downward ejection velocities are 220-460~km~s$^{-1}$ and 250-280~km~s$^{-1}$, respectively.  The inflow speed changed from 90~km~s$^{-1}$ to 12~km~s$^{-1}$ in 5 minutes. By using these velocities, we estimated the nondimensional reconnection rate, which we found to vary during this period from 0.20 to 0.055. We also found that  the plasma blobs in the sheet structure collided or merged with each other before they were ejected from the sheet structure. We hypothesize that the sheet structure is the current sheet and that these plasma blobs are plasmoids or magnetic islands, which could be important for understanding the dynamics of the reconnection region. ",
        "introduction": "Solar flares are the most energetic phenomenon in the solar atmosphere. Huge amounts of magnetic energy ($10^{29}$\\,-\\,$10^{32}$\\,erg), that has been stored in the corona, is released over a short period of time ($10^2$\\,-\\,$10^4$~sec). It is widely believed that magnetic reconnection, a physical process in which a magnetic field in a highly conducting plasma changes its connectivity due to finite resistivity, plays a central role in the rapid dissipation of magnetic energy in solar flares, as well as other explosive phenomena in astrophysical, space, and laboratory plasmas \\citep{yam10}.  Many studies have reported the observational features of magnetic reconnection in solar flares. These studies include observations of cusp-shaped loops \\citep{tsu92}, loop-top hard X-ray sources \\citep{mas94}, plasmoid ejection \\citep{shi95,tsu97,ohy98}, reconnection inflows \\citep{yok01,nar06,li09}, reconnection outflows \\citep{wan07}, supra-arcade downflows \\citep{mc99,inn03,asa04,sav11} and current-sheet-like structures \\citep{sui03,lin05,liu10}. Yokoyama et al.(2001) found a pattern of bright threads that merged toward an X-point (where the magnetic field was hypothesized to be reconnecting) in extreme ultraviolet (EUV) images, and identified the motion as reconnection inflow. Narukage \\& Shibata (2006) further worked on the statistical study of inflow by EUV observations, and found that the observed inflows are closely associated with plasmoid ejections or coronal mass ejections.   The classic magnetic reconnection model for solar flares, the CSHKP model \\citep{car64,stu66,hir74,kp76}, successfully explains some features of the long-lived (more than 1 hour) flare. This model has been improved to include a generalized configuration, leading to a unified model proposed by Shibata et al. (1995). Many of the observations presented above relating to the global structure of flares have been successfully explained by this magnetic reconnection model. For reviews of our current knowledge of magnetic reconnection in solar flares see, for example, Shibata (1999).   There are still fundamental difficulties in the theory of magnetic reconnection in flares, especially on the physics that determine the reconnection rate. The reconnection rate is defined as the reconnected magnetic flux per unit time, where the dimensionless form is the Alfv\\'en Mach number $M_{A}=V_{\\rm inflow}/V_{A}$, where $V_{\\rm inflow}$ is the inflow speed into the reconnection point and $V_{A}$ is the Alfv\\'en speed, respectively. Observationally the value of the reconnection rate is estimated to be 0.001 - 0.1 \\citep{iso02,iso05,nar06}. In the classic reconnection theory by Sweet (1958) and Parker (1957), the reconnection rate $M_{A}$ is given by $R_m^{-1/2}$, which is as low as $\\sim 10^{-7}$ in the solar corona because the magnetic Reynolds number (Lundquist number) $R_m = LV_A/\\eta \\sim 10^{14}$, where $L (\\sim10^{9}$~cm) is the typical size of the flare and $\\eta$ ($\\sim10^{13}T^{-3/2}\\sim 10^4~{\\rm cm^2~s^{-1}}$ for $T\\sim 10^6$~K) is the magnetic diffusivity for Spitzer resistivity. This value is too small to explain the timescales and the energy release rates of solar flares. However, Petschek (1964) predicted that the reconnection rate $M_A$ can be of the order of $0.01$\\,-\\,$0.1$, if the case when slow shocks emanate from a small diffusion region. This value is plausible to explain the observed reconnection rate in solar flares. MHD simulations have shown that localized resistivity can facilitate fast reconnection (e.g., Ugai \\& Tsuda 1977; Yokoyama \\& Shibata 1994), where fast reconnection is defined as the reconnection whose reconnection rate is close to unity and almost independent of the magnetic Reynolds number, such as Petschek's model. Observational support for Petschek's model is given by Tsuneta (1996).  It is believed that localized resistivity may be realized through kinetic effects in the coronal plasma, such as microscopic instabilities and wave-particle interaction. Indeed, magnetospheric observation and numerical simulations have shown that fast reconnection sets in when the current sheet thickness becomes comparable to the ion inertia length or Larmor length \\citep{tre01}. However, such microscopic scales are tiny ($\\sim100$~cm) in the solar corona compared with the characteristic scale of flares  ($\\sim10^9$~cm). How to connect such a huge gap  in the micro and the macro scales is one of the fundamental problems in the reconnection physics in solar flares.   To transcend this scale gap, Shibata \\& Tanuma (2001) suggested the fractal reconnection model. In this model, plasmoids (regions of magnetically confined plasma) play an important role not only in achieving fast magnetic reconnection but also in linking the small and large scale physics. A current sheet in an antiparallel magnetic field is unstable for the tearing instability, and can create plasmoids \\citep{FKR, ste83, lou07}. Once they are formed, the plasmoids tend to coalesce with each other to form bigger plasmoids \\citep{taj87}. When a plasmoid is ejected from the current sheet, rapid inflows are locally induced to fill the evacuated region. This leads to the intermittent enhancement of the reconnection rate. In addition, such plasmoids may be created on various scales. In this way plasmoids can fill the scale gap between the micro and macro scales of reconnection \\citep{taj97}.   Many observations report evidence that plasmoids are created in flare current sheets \\citep{ohy98,kli00,kar07}. Asai et al. (2004) found several dark downflows (dark voids)  in EUV images that were ejected intermittently during in the decay phase and also in the impulsive and main phases of a flare. Moreover, a nonthermal emission was associated with the motion of each dark void. Nishizuka et al. (2010) examined multiple plasmoids ejections from a flare using soft X-ray images. They found that the ejected plasmoids were strongly accelerated during multiple peaks in the hard X-ray emission in a flare. Both studies imply that the reconnection rate was enhanced by plasmoid ejections or downflows.   From the above examples, it can be seen that the observed reconnection rate is important to give a constraint on magnetic reconnection models for solar flares. One way to estimate the reconnection rate observationally is to simultaneously detect both reconnection inflow speed and outflow speed. Because two-dimensional magnetic reconnection theories predict that the reconnection outflow speed $V_{{\\rm outflow}} \\sim V_{A}$ (e.g. Priest \\& Forbes 2000), the reconnection rate $M_{A}$ is of the order of $V_{{\\rm inflow}} / V_{\\rm outflow}$.   We simultaneously detected both reconnection inflows and outflows for the flare that appeared on 2010 August 18 in EUV wavelengths and use these values to estimate the reconnection rate of the flare. We also found successive ejections of plasma blobs and their interaction in the current sheet. This study indicates a relationship between the energy release process of the flare and the dynamics of the plasma blobs in the current sheet. ",
        "conclusions": "We have examined in detail the morphology and dynamics of the the magnetic reconnection region in the limb flare on 2010 August 18. The overall characteristics, namely large scale eruption, inflow and outflow and hot loops below, are consistent with the classical CSHKP model. We could examine the fine scale dynamics of the reconnection region owing to the high spatial and temporal resolutions of AIA. Figure~5 shows the schematic picture of this flaring region. We observed simultaneous reconnection inflows and outflows and measured their velocities. A lot of plasma blobs appeared in the sheet and collided with each other or were ejected from it. We consider that the sheet structure was the reconnecting current sheet and that the plasma blobs were the magnetic islands or plasmoids created by the tearing instability. Figure~3(d) is similar to the distance-time diagram obtained from the simulation result of Shen et al. (2011) (see Figure~6 in their paper) where the plasmoids are ejected bidirectionally, which supports our assertions. Some plasmoids were made up of the multi-thermal plasma because they are recognized in the six wavelength images. These plasmoids could contain hot ($>$7~MK) plasma as they were visible in the AIA channels sensitive to hot plasma (94 and 131~{\\AA}), but these channels are also sensitive to cooler plasma \\citep{fos11}, which makes diagnosis of the temperature difficult. There is a possibility that they were heated by the coalescence of plasmoids \\citep{kli00}.   We derive the reconnection rate from the observed values by using the relation $M_{A} \\sim V_{{\\rm inflow}} / V_{\\rm outflow}$. The upward outflow velocity was found to change from $\\sim$460~km~s$^{-1}$ to $\\sim$220~km~s$^{-1}$, assuming that the plasmoid velocity ejected upward ($\\sim$460~km~s$^{-1}$) is almost the same as the outflow velocity. It should be noted that we did not find another upward ejection after the second ejection shown in Figure~3. Therefore, we use 460~km~s$^{-1}$ as the outflow velocity in the period of time between 5:10:30 UT and 5:11:00 UT and 220~km~s$^{-1}$ after 5:11:00 UT (see Figure~3). It should be noted that the apparent motion in the EUV images is due to the inflow but may not be the true plasma motion \\citep{che04}. Spectroscopic observation of reconnection region is necessary to quantitatively address the flow velocity \\citep{lin05, har06, wan07}. We assume that inflow speed was close to the apparent motion speed.  Because the length of the visible edge of the sheet did not change drastically, we neglect any effects for this motion on the inflow speed \\citep{yok01}. If projection effects are neglected, the inflow speed can be estimated as $V_{{\\rm inflow}} \\sim V_{{\\rm pattern}}$. The pattern velocity changed from 90~km~s$^{-1}$ to 12~km~s$^{-1}$, and therefore, the reconnection rate $M_{\\rm A}$ changed from 0.20 to 0.055 within the 5 minute period starting from when the current sheet became bright in EUV.  This value is comparable with the predicted value from Petschek's (1964) reconnection model and other observed values \\citep{nar06}.   We found that the disappearance of the plasmoids in the later phase coincides with the reduction of the reconnection rate from 0.20 to 0.055 (see Figure~3 and 4). Theoretically, plasmoid ejections induce strong inflow, which can promote the thinning of the current sheet and enhance the reconnection rate \\citep{shi01}. Therefore the reduction of the reconnection rate may be due to the suppression of the plasmoid activities. Although we lack detailed information of the time and space variations of the inflow velocity near the current sheet, which is necessary to confirm the above assertion, this result probably suggests that the reconnection rate of this flare was enhanced by the plasmoid dynamics."
    },
    "1112/1112.1021_arXiv.txt": {
        "abstract": "We present in this paper the results of searching the $Kepler$ Q2 public dataset for the secondary eclipses of 76 hot Jupiter planet candidates from the list of 1,235 candidates published by \\citet{Borucki2011}. This search has been performed by modeling both the $Kepler$ pre-search data conditioned light curves and new light curves produced via our own photometric pipeline. We derive new stellar and planetary parameters for each system, while calculating robust errors for both. We find 16 systems with 1-2$\\sigma$, 14 systems with 2-3$\\sigma$, and 6 systems with $>$3$\\sigma$ confidence level secondary eclipse detections in at least one light curve produced via the Kepler pre-search data conditioned light curve or our own pipeline, however, results can vary depending on the light curve modeled and whether eccentricity is allowed to vary or not. We estimate false alarm probabilities of 31\\%, 10\\%, and 6\\% for the 1-2$\\sigma$, 2-3$\\sigma$, and $>$3$\\sigma$ confidence intervals respectively. Comparing each secondary eclipse result to theoretical expectations, we find that the majority of detected planet candidates emit more light than expected due to thermal blackbody emission in the optical $Kepler$ bandpass, and present a trend of increasing excess emission with decreasing maximum effective planetary temperature. These results agree with previously published optical secondary eclipse data for other hot Jupiters. We explore modeling biases, significant planetary albedos, non-LTE or other thermal emission, significant internal energy generation, and mis-identification of brown dwarfs, low-mass stars, or stellar blends as possible causes of both the excess emission and its correlation with expected planetary temperature. Although we find no single cause is able to explain all of the planet candidates, significant planetary albedos, with a general trend of increasing planetary albedos with decreasing atmospheric temperatures, is able to explain most of the systems. Identifying systems that we deem likely to be low-mass stars or stellar blends, we estimate an 11\\% false positive rate in the current $Kepler$ planet candidate sample of hot Jupiters. We also establish robust upper limits on the eclipse depth for the remaining systems, and find that the emission of a significant fraction of these systems is consistent with the planets having very low albedos, i.e., at least 30\\% of all systems have $A_{g}$~$<$~0.3 at 1$\\sigma$ confidence levels. This result augments the current number of constrained exoplanetary albedos and extends the sample of low albedo determinations to planets with temperatures as low as 1200~K. Finally, we note that continued observations with the $Kepler$ spacecraft, and improved techniques for the removal of systematic noise in the $Kepler$ data, are needed to better characterize these systems. ",
        "introduction": "\\label{introsec}  Measuring the secondary eclipses of transiting exoplanets at optical wavelengths is a powerful tool for probing their atmospheres, in particular their albedos, brightness temperatures, and energy redistribution factors. The $Kepler$ mission has recently uncovered over a thousand new transiting planet candidates \\citep{Borucki2011}, which provide an unprecedented and uniform sample of high photometric precision light curves among which secondary eclipse signals can be detected.  In the past decade, many surprising discoveries regarding the atmospheric properties of hot Jupiters have been made. For example, many hot Jupiters appear to have temperature inversions, with numerous proposed explanations, but no definitive evidence for exactly which physical processes are involved \\citep{Hubeny2003,Fortney2006,Burrows2007,Fortney2008,Spiegel2009,Zahnle2009,Knutson2010,Madhusudhan2010}. Other results have found that the atmospheric composition of different planets vary significantly, or that they present a wide range of heat circulation efficiencies between their day and night sides \\citep[see][and references therein]{Baraffe2010}.  Most of the observations yielding to those discoveries have been done in the mid-infrared (3.6 - 24 $\\mu$m) with $Spitzer$. Observations at shorter wavelengths are more scarce, especially in the visible, but most of them point towards the predominance of very low geometric albedo, ($A_g$ $<$ 0.3 at the 3$\\sigma$ level upper limits), atmospheres in hot Jupiters \\citep{Charbonneau1999,Leigh2003a,Leigh2003b,Rodler2008,Rodler2010,Rowe2008,Alonso2009a,Alonso2009b,Alonso2010,Snellen2009,Christiansen2010,Welsh2010,KippingBakos2011a,KippingSpiegel2011,Desert2011a,Desert2011b,Langford2011}, in contrast to the up to $A_g=0.5$ albedos observed in the colder gas giants in our Solar System \\citep{Karkoschka1994}. Those results are in fair agreement with early theoretical models \\citep[e.g.,][]{Marley1999,Sudarsky2000,Seager2000}, which predict significant absorption of the incident stellar radiation in the visible by sodium and potassium, followed by re-emission in the infrared. Other molecules, such as TiO, VO, and HS, have also been suggested as possible strong optical absorbers \\citep[e.g.,][]{Hubeny2003,Fortney2008,Zahnle2009}.  However, three recent studies suggest higher geometric albedos for two planets. \\citet{Berdyugina2011} have published a value of $A_g$ = 0.28$\\pm$0.16 for HD 189733b via polarized reflected light\\footnote{We note that \\citep{Wiktorowicz2009} reported a non-detection of polarized light from this planet, and placed an upper limit to the polarimetric modulation of the exoplanet at $\\Delta$P $<$ 7.9$\\times$10$^{-5}$.}, while \\citet{KippingBakos2011a} and \\citet{Demory2011} respectively suggest albedos of $A_g$ = 0.38$\\pm$0.12 and $A_g$ = 0.32$\\pm$0.03 for Kepler-7b based on measurements of the emission of the planet with $Kepler$ during secondary eclipse. Some plausible explanations for such high albedos include Rayleigh scattering and the presence of clouds or hazes in the atmospheres of those planets \\citep{Demory2011}. Also, in the case of the hottest planets, some amount of thermal emission could be contributing to the measured emission levels in the reddest edge of the observed visible wavelength windows (e.g., at $\\lambda \\sim$ 0.8 $\\mu$m).  Although new theoretical work is necessary to determine the cause of apparently high albedos in some hot Jupiters, the key answer to whether these results are typical or not relies on more observations, since the current discussions are based on a statistically insufficient sample of only three planets. The purpose of this work is to significantly increase that sample by searching for the emission of hot Jupiters among the publicly available $Kepler$ light curves of planet candidates reported by \\citet{Borucki2011}. Given the photometric precision of the $Kepler$ data and the wavelength coverage of the $Kepler$ passband (0.4 - 0.9 $\\mu$m), these datasets provide unprecedented quality data to detect the secondary eclipses of those planets in the visible and statistically determine the albedos of hot Jupiters. Furthermore, as \\citet{Borucki2011} do no explicitly state how they modeled their light curves or obtained their parameters, a re-modeling of the data will perform an independent test on the methods they employed.  In addition to providing estimations of the planetary albedo, measuring the timing and duration of the secondary eclipse, when coupled with the primary eclipse, can directly measure the orbital eccentricity of a system \\citep[e.g.,][]{Knutson2007}. Also, if there is a significant flux contrast between the day and night side of the planet, one may be able to measure the varying amount of emitted light by the planet in the light curve, and directly measure the day-to-night contrast ratio \\citep{Harrington2006,Knutson2007}. Even a robust upper limit on the eclipse depth can narrow the range of possible planetary albedos and yield useful information on the statistics of exoplanetary albedos.  In \\S\\ref{datasec} we present our target selection criteria and describe how we reprocess the $Kepler$ light curves from the pixel-level data. In \\S\\ref{modelsec} we describe how we model the data using the JKTEBOP code, and obtain robust errors on all parameters while accounting for potential systematic noise. We present our derived physical parameters of both the planet candidates and their host stars in \\S\\ref{paramsec}, and in \\S\\ref{trendssec} we examine possible trends in our results. In \\S\\ref{indivsec} we discuss individual candidates of interest, and finally in \\S\\ref{concsec} we summarize our findings and examine possible future directions for the study of this sample. ",
        "conclusions": "\\label{concsec}  We have analyzed the $Kepler$ Q2 light curves of 76 hot Jupiter transiting planet candidates using both the $Kepler$ PDC data and the results from our own photometric pipeline for producing light curves from the pixel-level data. Of the 76 initial candidates only 55 have light curves with high enough photometric stability to search for secondary eclipses. For the remaining targets this search is hindered by either intrinsic variability of the host star or residual systematics in the light curve analyses. We have found that significant systematics in the $Kepler$ light curves due to small photometric apertures and large centroid motions hinder analyses if not properly removed or accounted for, and that a re-reduction of the photometry is best done at the pixel-stamp level. We also stress the importance of taking into account how $Kepler$ light curves are produced from the pixel level data when considering any detection of low-frequency signals, as they can vary significantly depending on the technique employed.  We have also re-determined the stellar and planetary parameters of each system while deriving robust errors that take into account residual systematic noise in the light curves. We detect what appear to be the secondary eclipse signals of $\\sim$20-30 of the targets in our list at $> 1 \\sigma$ confidence levels, and also derive robust upper limits for the secondary eclipse emission of all the remaining objects in our sample. All of our sample present excess emission compared to what is expected via blackbody thermal emission alone, as well as a trend of increasing excess emission with decreasing expected maximum effective planetary temperature, in agreement with previously reported secondary eclipse detections of hot Jupiters in the optical, which can be attributed to the 'appearance' of increasing albedos with decreasing planetary temperatures. By performing statistical analyses of those results we arrive to the following main conclusions.  \\begin{enumerate}  \\item Assuming no contribution from reflected light, i.e., A = 0.0, the majority of the detected secondary eclipses reveal thermal emission levels higher than the maximum emission levels expected for planets in local thermodynamical equilibrium.  \\item While the extra emission from many of the planets can be accounted for by varying the amount of reflected light, the emission from  $\\sim$50\\% of the detected objects ($>$1$\\sigma$) can not be accounted for even when assuming perfectly reflective planets (i.e., A = 1.0). These planets must either have much higher thermal emission in the optical compared to a blackbody, as predicted by theoretical models, have very large non-LTE optical emission features, have underestimated host star masses, radii, or effective temperatures, or are in fact false-positives and not planets but rather brown dwarfs, very low-mass stars, or stellar blends. Follow-up observations of these systems are necessary to confirm this conclusion. The most outstanding potentially false positive systems are KOI 64.01, 144.01, 684.01, 843.01, 1541.01, and 1543.01. Given that this is 6 of the 55 systems that we modeled, or 11\\% of the sample, it appears to agree with the expected $\\sim$10\\% false positive rate of the initial 1,235 candidates estimated by \\citet{Morton2011}.  \\item Although we do not identify a sole cause of the observed trend of increasing excess planetary emission with decreasing expected maximum effective planetary temperature, the hypothesis of increasing planetary albedo with decreasing planetary temperature is able to explain many of the systems. This would be physically plausible as the upper atmospheres of Jupiter-like planets transition from very low albedos at high temperatures, as observed for hot Jupiters, to higher albedos at lower temperatures, as observed for cool Jupiters in our own solar system. We note that further observations via Kepler and multi-wavelength ground-based facilities of both the planetary candidates and their host stars are still needed to fully explain this trend.  \\item From the emission upper limits placed on planet candidates for which we do not detect secondary eclipse signals, we conclude that a significant number, at least 30\\% at the 1$\\sigma$ level, of those targets must have very low-albedos, ($A_{g}$ $<$ 0.3), which is a result that is consistent with the majority of previous observations and early theoretical hot Jupiter model predictions. All previous observational upper limits had been placed on hot Jupiters with expected atmospheric effective temperatures higher than $\\sim$1650 K. Our results extend that temperature limit to planets with expected effective temperatures higher than 1200 K and can help further establish what chemicals place a critical role in the atmospheric properties of hot Jupiters.  \\item From the inspection of individual targets we conclude that the majority of our secondary eclipse depths for candidates with previously published eclipse detections are consistent with the results from those other studies. We note that several of those other studies have access to $Kepler$ data from quarters after Q2, while our analysis is limited to just the public Q2 data, and therefore our results have larger detection errorbars in some cases.  \\item Our results are based on only one quarter of $Kepler$ data, but 12-24 quarters should become available in the future. Therefore, we expect that future studies of these targets will be able to improve our secondary eclipse detections and upper limits by factors of 3-5, or much greater if noise systematics in the light curves can be further reduced.  \\end{enumerate}"
    },
    "1112/1112.2214_arXiv.txt": {
        "abstract": "We present the metallicity distribution function (MDF) for 24,270 G and 16,847 K dwarfs at distances from 0.2 to 2.3 kpc from the Galactic plane, based on spectroscopy from the Sloan Extension for Galactic Understanding and Exploration (SEGUE) survey. This stellar sample is significantly larger in both number and volume than previous spectroscopic analyses, which were limited to the solar vicinity, making it ideal for comparison with local volume-limited samples and Galactic models. For the first time, we have corrected the MDF for the various observational biases introduced by the SEGUE target selection strategy. The SEGUE sample is particularly notable for K dwarfs, which are too faint to examine spectroscopically far from the solar neighborhood. The MDF of both spectral types becomes more metal-poor with increasing $|Z|$, which reflects the transition from a sample with small [$\\alpha$/Fe] values at small heights to one with enhanced [$\\alpha$/Fe] above 1 kpc. Comparison of our SEGUE distributions to those of two different Milky Way models reveals that both are more metal-rich than our observed distributions at all heights above the plane. Our unbiased observations of G and K dwarfs provide valuable constraints over the $|Z|$-height range of the Milky Way disk for chemical and dynamical Galaxy evolution models, previously only calibrated to the solar neighborhood, with particular utility for thin- and thick-disk formation models. ",
        "introduction": "\\label{sec:intro}  Measuring the metallicity distribution of stars in the Milky Way disk is imperative for understanding its chemical and dynamical evolution. Cool stars, such as G and K dwarfs, have lifetimes comparable to the age of the Galaxy, providing a complete fossil record of chemical development. By measuring the metallicity distribution of these cool stars, we provide constraints on the disk's star-formation history, and how it varies with respect to time and location. The metal-poor end of the metallicity distribution function (MDF) reveals information about the earliest era of star formation in the Galaxy, such as the mass function of the earliest stars, the extent of chemical pre-enrichment, and the rates and yields of core collapse supernovae. The metal-rich end of the MDF reflects recent Galaxy conditions, such as the present-day mass function and the frequency and yields of both type Ia and II supernovae. The MDF provides information about the regulation of star formation and the merger and accretion history of the Galaxy.  There is a wide range of chemical and dynamical evolution simulations that model the structure of the Galactic disk. These vary with respect to properties such as the star-formation history, initial mass function, prominence of inflows and outflows, and the likelihood of mergers. All use observed samples to test their predictions of the metallicity distribution. Until recently, we lacked an adequate observational sample to quantitatively test these models beyond the solar neighborhood. In this work, we use the Sloan Extension for Galactic Understanding and Exploration (SEGUE, \\citealt{yanny09}) survey to determine an unbiased MDF of cool dwarfs over a large volume of the disk, from 0.2 to 2.3 kpc from the plane of the Galaxy.  Initial examination of the chemistry of local G dwarfs revealed the ``G-dwarf problem'' \\citep{vandenbergh62, pp75, wg95, rpm96, favata97}. Early models of star formation and chemical evolution, such as the simple closed box model of \\citet{schmidt}, predicted many more low-metallicity G dwarfs than were actually observed. It was initially suspected that the G-dwarf problem arose from observational biases, namely that metal-rich stars are brighter and thus likely to be over-represented in a magnitude-limited sample. However, the deficiency of low-metallicity stars persisted in later observational samples that corrected for these biases. Work such as the Geneva-Copenhagen survey (GCS) of the solar neighborhood \\citep{jorgensen00, nordstrom04, holmberg07, holmberg09, casagrande11} indicated that for a large volume-complete and kinematically-unbiased sample, the simple closed box model over-predicted the number of cool metal-poor dwarfs even more than originally reported.  We know that many of the assumptions in the simple closed box model are unphysical, in particular instantaneous recycling and the absence of gas flows. More recent models have updated and experimented with the model parameters to better match the observed metallicity distribution, for example, adding inflows of low-metallicity material \\citep{larson72, chiappini97}, varying the formation time scales of different Galaxy components \\citep{chiappini01}, or commencing star formation from pre-enriched material \\citep{truran71}. Models have also experimented with the initial mass function (IMF) \\citep{chiappini00, romano05}; a proposed metallicity-dependent IMF will produce more high- than low-mass stars at early times, resulting in fewer cool metal-poor stars.  This exploration of the G-dwarf problem prompted observations of the MDF of cooler stars. K and M dwarfs have even longer lifetimes than G-dwarf stars. There is a chance that more metal-rich G dwarfs have evolved off of the main sequence, leading to a metallicity bias in the MDF, which will not occur for cooler spectral types. In addition, comparing the MDFs of different spectral types is particularly useful for constraining the variation in the IMF. For example, if more low-mass stars are created as the metallicity of the environment increases, there should be relatively fewer metal-poor K dwarfs than G dwarfs. Work such as \\citet{mould82}, \\citet{favata97}, \\citet{flynnmorell97}, \\citet{rpmaciel98}, \\citet{kotoneva02}, and \\citet{woolf12}, found that, just as with G-dwarf stars, the simple closed box model predicted more metal-poor stars than were observed for both K and M dwarfs, implying that the IMF and star-formation history for these different spectral types are similar to one another.  Measurements of the G- and K-dwarf MDF have been mostly confined to the solar neighborhood, limiting our understanding of disk properties and evolution processes to the local volume. However, the chemical and dynamical structure of the disk does not appear to be uniform with respect to Galactic height, but rather composed of two components, a thin and thick disk, of unclear interdependence. \\citet{gilmore83} first detected the thick disk in the Milky Way when they determined that the stellar number density as a function of height above the plane was best fit by two components\\footnote[1]{The likely existence of a second stellar component close to the Galactic plane was previously noted by \\citet{yoshii82}, who referred to it as a halo component, even though its inferred density was 10 times that of the local halo.  The \\citet{yoshii82} normalization relative to the local thin disk (0.01--0.02) and scale height ($\\sim$2 kpc) were commensurate with the values later determined by \\citet{gilmore83} (0.02, and 1.5 kpc, respectively). See also \\citet{yoshii87}.}, one with a scale height of approximately 300 pc, and the second with a scale height of 1350 pc.  Analyses by \\citet{gilmore85, wg95}, and \\citet{chiba00} established that the two populations were distinct in [Fe/H]. The so-called thick disk was metal-poor, with a peak metallicity around [Fe/H]$=-$0.6, in contrast to the thin disk, with a peak metallicity typically around [Fe/H]$=-$0.2. Further observation of solar neighborhood samples revealed that the two populations were distinct in kinematics \\citep{soubiran03}, age \\citep{fuhrmann98}, and $\\alpha$-abundance \\citep{fuhrmann98, prochaska00, bensby03, bensby05, reddy06}. However, \\citet{norris_ryan} and recent work by \\citet{bovy11, bovy11b} and \\citet{liu12}, focusing on [$\\alpha$/Fe] vs. [Fe/H] and kinematics, questioned whether or not the two components were actually separable from one another, proposing that the Milky Way disk has a ``thicker disk component,'' rather than two distinct structures. The debate over the basic structure of the Milky Way disk emphasizes the need for an unbiased spectroscopic sample that extends over a large volume of the disk, such that we can investigate the metallicity structure of both of the proposed components with a uniform large data set. Furthermore, it is critical to use our own Galaxy to disentangle the mechanisms behind disk development because the thick disk appears to be a regular feature in numerous galaxies \\citep{burstein79, dalcanton02, yoachim08a, yoachim08b}, including those at redshifts as high as z$\\sim$3 \\citep{elmegreen06}. In this work, we investigate and constrain the chemical structure of the Milky Way as a whole, including both thin- and thick-disk components.  \\subsection{Previous Analyses of the Chemical Structure of the Disk}  Past observations of the MDF have been limited in accuracy, sample size, and, most importantly, volume. Due to the low luminosity of cool stars, analyses such as \\citet{pp75}, \\citet{gilmore85}, \\citet{wg95}, \\citet{rpm96}, \\citet{flynnmorell97}, \\citet{rpmaciel98}, \\citet{jorgensen00}, and \\citet{kotoneva02} relied on photometric calibrations to determine metallicities. Photometric metallicity determinations are susceptible to errors from reddening corrections, have reduced sensitivity at low metallicity, and depend strongly on the adopted calibration to spectroscopic estimates, which vary from work to work. Using spectroscopic measurements increases the accuracy and precision of metallicity determinations, in addition to providing kinematic information such as radial velocities, albeit with the significant added cost of increased observing time. Previous spectroscopic analyses of these low luminosity targets were limited to hundreds of stars along individual lines of sight \\citep{wg95, favata97, fuhrmann98, fuhrmann04, allendeprieto04, luck05, luck06, luck07, fuhrmann08, arnadottir09, fuhrmann11, katz11}.  Two recent surveys have improved analyses of the MDF of the Milky Way disk by increasing the sample size and sky coverage. The GCS is a magnitude-limited survey, with $\\sim$14,000 F and G dwarfs with metallicities estimated from Str{\\\"o}mgren photometry \\citep{jorgensen00, nordstrom04, holmberg07, holmberg09, casagrande11}. The RAdial Velocity Experiment (RAVE) further expands the sample of cool stars, with around 17,000 F and G dwarfs and spectroscopic metallicities \\citep{siebert11}. Although they have uniform data sets over a large region of the sky, neither of these have dwarf stars far beyond the plane of the Galaxy. The GCS sample is limited to 200 pc from the Sun \\citep{casagrande11}. The dwarfs in the RAVE sample probe to a maximum distance of approximately 1 kpc; the typical distance for their cool dwarf sample ranges from 50--250 pc \\citep{zwitter10, steinmetz12}. Previous analyses of individual line of sights were similarly limited, extending no further than 50 pc from the Sun. Both the GCS and RAVE samples exhibit MDFs that peak at around solar metallicity, despite their different methods for estimating [Fe/H] \\citep{casagrande11, coskunoglu11}, indicating that they are dominated by thin-disk stars, as expected for a solar neighborhood sample. Without a large sample of cool stars far beyond the solar neighborhood, we cannot accurately constrain the chemical and dynamical development of the Milky Way disk as a whole.  In contrast to GCS and RAVE, the early work on G dwarfs of \\citet{gilmore85} reached Galactocentric height above the plane, $|Z|$, of around 1.6 kpc, beyond the local volume. This analysis, however, consisted of only ten lines of sight and estimated metallicity from ultraviolet excess, based on broadband photometry. More recent work by \\citet{katz11} examines subgiants and giants as high as 5 kpc above the plane of the Galaxy. Although they use spectroscopic metallicities, this sample probes only $\\sim$400 stars along two lines of sight, limiting their statistical power to measure the extremes of the MDF or its variation over their sample volume.  SDSS and SEGUE provide a uniform sample over a large portion of the Milky Way; previous analyses have sought to use this data to constrain the chemical structure beyond the solar neighborhood. \\citet{ivezic08} and \\citet{bond10} examined over 2 million F and G dwarfs in SDSS, utilizing broadband \\emph{ugr} photometry to estimate both the metallicity and distance for each star. This large sample probes far beyond the local volume, ranging from $|Z|$ of 0.5 to 7 kpc, and is dominated by thick-disk and halo stars. They find that the disk metallicity decreases from around [Fe/H] of $-$0.6 close to the plane of the Galaxy to plateau at around $-$0.8. They also detect a clear break between the disk and halo stars, which have a mean metallicity of around $-$1.4 \\citep{bond10}. Although they have a large and unbiased stellar sample, uncertainties in the photometric metallicity calibration affect the absolute scale of the MDF and manifest as errors in the derived photometric parallax relationships. They also cannot exclude binaries from the sample; an undetected companion has a significant effect on stellar photometry \\citep{schlesinger10}. Finally, although they analyze a range of stellar types, they focus their analysis on F stars, specifying a color range of $(g-r)$ from 0.2 to 0.4. F dwarfs have shorter lifetimes than G dwarfs, and are thus an incomplete sample of the disk chemical evolution.  Avoiding reddening and calibration uncertainties, past studies utilized SDSS spectroscopy to measure the MDF. \\citet{allendeprieto06} examined a sample of 22,770 F and G dwarfs from SDSS Data Release 3, ranging from 1 to 8 kpc in $|Z|$. They focus their analysis on more metal-poor stars, examining the metallicity differences between the thick disk and halo. \\citet{allendeprieto06} find that the MDF of G dwarfs between 1 and 3 kpc of the plane, associated with the thick disk, have a peak around [Fe/H]$=-$0.7, significantly more metal-rich than those above this height, which are halo stars and exhibit a peak [Fe/H] of $-$1.6. With the release of SEGUE in Data Release 6, the sample size and coverage in Galactocentric height and radius have improved significantly since this work, allowing us to more accurately determine the MDF of G dwarfs with respect to spatial position. Furthermore, \\citet{allendeprieto06} do not account for observational biases that originate from the SDSS target selection algorithm, which did not sample the stellar range of metallicity and age equally.  The recent work of \\citet{lee11b} utilize the SEGUE spectroscopic G-dwarf sample to examine the properties of the $\\alpha$-separated thin- and thick-disk components. Similar to \\citet{ivezic08}, and \\citet{bond10}, they find that the thick-disk component is more metal-poor than the thin-disk component, and dominates far from the plane of the Galaxy. Although they use up-to-date SSPP parameters, they do not examine the entire G-dwarf sample available in SEGUE, selecting only stars specifically targeted by SEGUE as G dwarfs, rather than all stars that fulfill the criteria. \\citet{lee11b} also do not account for the target selection biases in SEGUE, and are thus biased towards metal-poor stars, although these observational biases are less significant in their sample than those of \\citet{allendeprieto06}.  Most importantly, rather than examining the MDF of the disk as a whole, they examine the relative metallicity structure of two $\\alpha$-separated components. We examine the MDF of the SEGUE sample without making any chemical or kinematic separation of the populations.  \\subsection{This Work}  With a large number of stars and extensive sky and volume coverage, the complete SEGUE survey provides an ideal sample to examine the chemical abundance distribution in the Galaxy \\citep{yanny09}. The SEGUE data consist of SDSS $ugriz$ photometry and spectroscopy for 240,000 stars over a range of 14$<g<$20.3 in $\\sim$3500 square degrees on the sky. There are 50,210 and 26,834 SEGUE stars that fulfill the G- and K-dwarf photometric criteria, reaching Galactic distances from the plane, $|Z|$, of around 3.5 kpc and ranging from Galactic radius, R, of 5 to 13 kpc. Not only is this the largest spectroscopic sample available, it covers a much more extensive volume of the Milky Way disk than all previous analyses. Furthermore, it is the largest and most extensive spectroscopic sample of K-dwarf stars, whose faint magnitudes make it difficult to probe far from the solar neighborhood. Previous samples of K-dwarfs were limited in volume, spatial-coverage, sample size, and atmospheric parameter accuracy.  We utilize this sample to determine an unbiased MDF of both G- and K-dwarf stars throughout a large volume of the disk and determine how they change with respect to height above the plane. By using both spectral types, we examine a larger volume of the disk, constraining the disk chemical structure both close and far from the plane of the Galaxy, whereas analyses of hotter stars in SEGUE are limited to thick-disk dominated space. In addition, we compare the MDF of the two spectral types to analyze whether or not the star-formation and chemical-evolution history vary with respect to stellar mass.  SEGUE selects stars as spectroscopic targets using a series of photometric and proper motion criteria to isolate stars of different spectral types. Because the SDSS and SEGUE survey addressed many different scientific questions, it had a large variety of target types. Each of these isolate a different portion of parameter space, resulting in samples which do not sample the stellar range of metallicity and age equally. The net effect for the entire sample is an observational bias in favor of metal-poor stars. This is the first work to systematically determine and account for the various observational biases in the SEGUE G- and K-dwarf sample. By comparing SDSS photometry and SEGUE spectroscopy for each line of sight, we assign weights to each SEGUE star to account for the various target-selection biases. Computing the MDF with this weighted sample allows us to examine how the underlying population of stars throughout a large volume of the disk varies with respect to spatial position.  In this work, we first discuss how we extract the G and K dwarf stars from the SEGUE database (\\S\\,\\ref{sec:sample}) and estimate the distance to each star using isochrones (\\S\\,\\ref{sec:distances}). We then constrain the different observational biases that may arise in the SEGUE sample (\\S\\,\\ref{sec:contamination}) and detail our technique to correct for those biases that arise from the SEGUE target selection algorithm (\\S\\,\\ref{sec:accounting_biases}). Utilizing these target selection weights, we determine complete and unbiased metallicity distribution functions for both G and K dwarfs that accurately reflect the underlying chemical structure (\\S\\,\\ref{sec:mdf}). In \\S\\,\\ref{sec:discussion}, we compare the observed metallicity structure to previous analyses and two Galaxy models, the stellar population synthesis model TRILEGAL 1.4 \\citep{girardi05} and the radial migration model of \\citet{schonrich09a, schonrich09b}. These comparisons reveal how critical an accurate picture of the disk chemical structure is for understanding how the disk formed and evolved and how valuable our observations are for understanding how the Milky Way developed. ",
        "conclusions": "Our SEGUE sample consists of 16,847 K dwarfs with $S/N\\geq$10, surface gravities larger than $\\log g=$4.1, appropriate $r_0$ magnitudes, and E($g-r$) less than 0.2 mag. This sample has $-$3.4$\\leq$[Fe/H]$\\leq$0.6, 4.1$\\leq \\log g \\leq$5.0, and 4150$\\leq$T$_{eff}\\leq$6315 K. It contains stars with distances from 0.3 to 5.8 kpc and covers Galactic radii from 4.8 to 13.4 kpc and heights from the plane up to $\\pm$4.0 kpc.  We have 24,270 G dwarfs with the appropriate $S/N$, surface gravity, extinction, and $r_0$. This sample ranges from [Fe/H] of $-$3.1 to 0.6, surface gravities up to 5.0, and effective temperatures from 4210 to 6300 K. It extends to larger distances than the K-dwarf sample, from $\\sim$0.4 to 10.2 kpc, covering from 3.7 to 16.5 kpc in Galactic radius and -7.7 to 9.7 kpc in height above the Galactic plane.  We have performed the first complete analysis of the different possible biases in this SEGUE sample.  The color and magnitude selection criteria of the two categories result in a sample that is quite clean, with few stars of inappropriate color being shifted into the sample by photometric errors or undetected binarity. We also remove lines of sight from our sample with E$(g-r)$ greater than 0.2 mag and establish that uncertainties in the \\citet{sfd98} extinction measurements have little effect on our distance estimates.  We also constrain how stars of different evolutionary phase affect our sample. Combining an analysis of $\\log g$ with MgH and Galaxy models, we determine that the SEGUE G and K dwarfs will have a negligible amount of contamination from evolved stars. We will lose old, metal-rich stars from the G-dwarf sample due to evolution; our $(g-r)$ color cut biases the G-dwarf sample against stars with [Fe/H]$\\geq$0.  We estimated the effect of different uncertainties in the G- and K-dwarf sample on the MDF (see Fig.\\,\\ref{fig:mdf_per_error}). The uncertainties are dominated by distance errors, both random and systematic. None of the various uncertainties, from photometric errors to undetected binarity, induce a bias in the metallicity of the sample. In our MDF, we include every different source of uncertainty in the structure, combining them in quadrature (\\S\\,\\ref{sec:monte_carlo}).  Although issues ranging from surface gravity to extinction have little effect on our derived MDF, the target selection algorithm of SEGUE does bias the sample in metallicity space, favoring metal-poor stars. The same number of stars are observed in fields at low Galactic latitude, which have a high stellar density and are dominated by high-metallicity stars, as in those at high Galactic latitude, where the stellar density is much lower and the stars are primarily metal poor. Thanks to the quantitative target-selection algorithm of SEGUE, we have both quantified and corrected for these different biases using a series of weights. The uncertainties from these values are integrated into our MDF uncertainties via a bootstrap analysis (\\S\\,\\ref{sec:monte_carlo}).   \\label{sec:discussion}  \\subsection{The MDF for Different Spectral Types} \\label{sec:st_variation}  Past work by \\citet{favata97} found a distribution of nearby K-dwarf stars that was significantly more deficient in metal-poor stars than their G-dwarf sample. They conjectured that this result could originate from either a metallicity bias in their original catalog or metal-enhanced cool star formation. However, later work by \\citet{rpmaciel98} found their K-dwarf sample in good agreement with their G dwarfs. Similarly, \\citet{kotoneva02} found a much broader distribution of K dwarfs with respect to [Fe/H] than \\citet{favata97}. Likely, the narrow distribution of K dwarfs in \\citet{favata97} was due to small sample size and a tendency to target brighter stars for spectroscopy.  We compare the MDF of the weighted SEGUE G and K dwarfs over the overlapping distance range in Figure\\,\\ref{fig:overall_comp}.  Our analysis confirms the work of \\citet{rpmaciel98}, which found the distribution of metallicity in G and K dwarfs in the solar neighborhood to be in good agreement. For the overlap region for the G and K dwarfs, i.e., between 1.59 and 1.84 kpc, both spectral types peak at around [Fe/H] of $-$0.5, consistent with the peak metallicity for the $\\alpha$-enhanced sample identified by \\citet{lee11b}. The peak agreement of the two spectral types implies that the chemical-enrichment and star-formation history for the G and K dwarfs are quite similar.  The MDFs with respect to $|Z|$ of the two spectral types differ at the extremes of the metallicity distribution (Figure\\,\\ref{fig:mdf_z_kdgd_dcut}). For heights below 1 kpc, the K dwarfs have a larger population of metal-rich stars. As discussed in \\S\\,\\ref{sec:evolutionary}, our color criteria for G dwarfs will remove the most metal-rich stars with ages greater than 5 Gyr. Above 1 kpc, there are few stars at or above solar metallicity, so the discrepancy between the two spectral types is no longer noticeable.  When we examine the metal-poor end for the G- and K-dwarf samples at different $|Z|$, we find a discrepancy (Figure\\,\\ref{fig:mdf_z_kdgd_cumu}). For $|Z|$ distances below 1.5 kpc, the fractional contribution of stars with [Fe/H]$<-$1.0 agree within 1$\\sigma$ (Figure\\,\\ref{fig:mdf_z_kdgd_cumu}). The percentage of stars below [Fe/H] of $-$1.0 is consistent for the two spectral types up to $|Z|=$1.5 kpc. However, at $|Z|$ greater than 1.5 kpc, the percentage of metal-poor ([Fe/H]$<-$1.0) K dwarfs is significantly larger than the metal-poor G dwarfs. For $|Z|>$1.5 kpc, the fraction of G dwarfs increases from $\\sim$6\\% at smaller $|Z|$ to 9$_{-5}^{+7}$\\%. The K-dwarf percentage increases significantly more, from around 15\\% for 0.2$\\leq |Z|<$1.5 kpc to 32$_{-12}^{+15}$\\% above $|Z|$ of 1.5 kpc (see Tables\\,\\ref{tab:gd_zbin_st}, \\,\\ref{tab:kd_zbin_st}, \\,\\ref{tab:gd_zbin_vc}, \\,\\ref{tab:kd_zbin_vc}). For smaller $|Z|$, the median metallicity for both G and K dwarfs agree to within 0.04 dex. For the highest $|Z|$ bin, the median metallicity for the G dwarfs is $-$0.55, whereas the K-dwarf median [Fe/H] is $-$0.74, much more metal-poor.  The origins of this discrepancy are currently unclear. The optimized version of the SSPP ensures that the [Fe/H] estimates for the G and K dwarfs are consistent with one another; there is no systematic offset in SSPP metallicity between the two spectral types of this size (Appendix\\,\\ref{sec:feh_det}). One possible cause of the MDF differences at large $|Z|$ is our assumed binarity fraction. If there is a larger probability of an undetected companion for K dwarfs than G dwarfs, we will be underestimating the distance to more K stars. This will cause more metal-poor stars to fall into the K-dwarf sample.  Examination of the G- and K-dwarf MDF at different heights above the Galactic plane reveals that the K dwarfs are more metal-rich at low heights and more metal-poor at large $|Z|$. This makes the total K-dwarf MDF appear broader than the G-dwarf MDF over the overlapping distance range. As the fraction below [Fe/H] of $-$1.0 from both spectral types is consistent within 2$\\sigma$, we leave a more in-depth study of the metal-poor end of the G- and K-dwarf MDF for a future paper.  \\subsection{Extending the G-Dwarf Problem}  The original G-dwarf problem stemmed from the observational results that a simple closed box model of chemical evolution over-predicted the number of metal-poor stars. In the top row of Figure\\,\\ref{fig:simple_model_zbin}, we compare our total MDF of G and K dwarfs to a simple closed box model with instantaneous recycling, based on \\citet{pagel97}. As expected, we find that our empirical distributions of both G and K dwarfs are deficient at the low-metallicity end when compared to these models, confirming that the G-dwarf problem persists in the large Galactic volume of our survey and to lower temperatures. For the total sample, the fraction of K stars with [Fe/H] less than $-$0.8 is $\\sim$17$_{-3}^{+5}$\\%, whereas the model predicts 28\\%.  Similarly, the model predicts 27\\% of G dwarfs should have [Fe/H]$< -$0.8, but we observe $\\sim$10$_{-2}^{+3}$\\%.  The G-dwarf problem was originally reported for stars within the solar neighborhood. Using the SEGUE sample, we find that the discrepancy with the simple closed box model is evident throughout the disk (Figure\\,\\ref{fig:simple_model_zbin}). The discrepancies are on the order of 10--20\\% between the observed distribution and the simple closed box model, enhanced at larger distances from the plane, and consistent between the spectral types.  \\subsection{Comparison of the G and K Dwarf MDF with Previous Work} \\label{sec:comp_totmdf}  Our total MDF for G and K dwarfs covers a substantial portion of both the thin and thick \\emph{in-situ} Milky Way disk. As a consequence, the peak metallicity for both spectral types is lower than previous analyses, such as RAVE \\citep{siebert11, coskunoglu11, bilir12} and GCS \\citep{holmberg07, casagrande11}, which are dominated by thin-disk, more metal-rich stars. Similarly, the local MDFs for smaller samples, such as \\citet{rpmaciel98} and \\citet{kotoneva02}, are more metal-rich than ours. These analyses have no stars in common with the SEGUE sample and probe a lower $|Z|$ range. Although the lowest $|Z|$ bin of G and K dwarfs ranges from 0.2 to 0.5 kpc, the median distance from the Galactic plane for these SEGUE stars is around 0.4 dex. Thus, our low $|Z|$ MDF is not directly comparable with local samples, and we defer this analysis to a later paper.  \\citet{allendeprieto06} measure the MDF of G-dwarf stars, identified by effective temperature, between $|Z|$ of 1 and 3 kpc. This volume of space is dominated by the thick disk. \\citet{allendeprieto06} used the SDSS calibration star sample to examine the MDF between 1 and 3 kpc, a volume of space dominated by the thick disk. The size and homogeneity of the data enabled measuremnts of the change in the observed sample MDF with respect to distance from the plane. However, this sample pre-dated the SEGUE survey, and the calibration star selection introduced significant metallicity and evolutionary bias. The target selection of their sample causes the MDF of \\citet{allendeprieto06} to drop off steeply above [Fe/H] of $-$0.4; it may also cause the MDF to peak at around [Fe/H] of $-$0.7, more metal-poor than our MDF above $|Z|$ of 1 kpc.  Whereas the GCS, RAVE, and \\citet{allendeprieto06} sample probe individual components of the disk, the photometric analyses of \\citet{ivezic08} and \\citet{bond10} examine a larger volume of space, from $|Z|$ of 0.8 to 7 kpc. At heights below 2 kpc, their MDF is slightly more metal-poor than ours, by around 0.1 to 0.2 dex, perhaps due to inaccuracies in their photometric metallicity calibration. The contrasting MDFs of these different samples emphasizes that the chemical structure of the disk varies significantly with respect to spatial location. Only with a uniform, extensive, spectroscopic sample such as SEGUE can we obtain accurate constraints on the chemistry of the disk as a whole.  \\subsubsection{\\citet{katz11}}  Past work by \\citet{katz11} sought to constrain the vertical metallicity gradient of the thick disk alone. Their sample consists of approximately 400 stars along two lines of sight; they calculate atmospheric parameters using the ETOILE spectral matching program. Utilizing a series of Besan\\c{c}on Galaxy models, they distinguish the thick-disk stars from the thin-disk population.  The MDF of the \\citet{katz11} sample is slightly more metal-rich than our G- and K-dwarf sample at all $|Z|$ (Figure\\,\\ref{fig:katz_gdwarf}). However, the two samples show general similarities to one another in structure and peak metallicity. Although the \\citet{katz11} sample probes beyond the solar neighborhood, it is limited to only two lines of sight. Thus, it is promising for their analysis that we see general agreement in the metallicity distributions, but their limited sample size and volume make it difficult to measure the variation in the MDF at different locations in the disk.  \\subsubsection{\\citet{lee11b}}  \\citet{lee11b} use the [$\\alpha$/Fe] ratios of the SEGUE G-dwarf stars to chemically separate the thin- and thick-disk components. They use a different sample than ours, extracting only stars specifically targeted as G dwarfs with $S/N >$30, $\\log g \\geq$4.2, and distances within 3 kpc. We compare our G-dwarf metallicity distribution over the spectral distance range to the $\\alpha$-poor and $\\alpha$-enhanced metallicity distributions in Figure\\,\\ref{fig:lee_gdwarf}. In this figure, the two $\\alpha$-separated components are normalized by the total number of stars in the \\citet{lee11b} sample. Farther away from the plane of the Galaxy, the $\\alpha$-rich component makes up a larger part of the sample, indicating that this $\\alpha$-rich population is an important contribution to our MDF far from the plane.  At low heights above the plane, our observed MDFs contains a mix of the $\\alpha$-poor and $\\alpha$-rich samples from \\citet{lee11b}. At $|Z|$ of 1.0 kpc, our MDF skews more metal-poor, matching the \\citet{lee11b} $\\alpha$-rich sample better. This suggests that, at these heights, the G- and K-dwarf sample is dominated by the $\\alpha$-enhanced disk stars. Based on this comparison with the results of \\citet{lee11b}, our observed vertical variation in the MDF reflects the transition from a predominantly $\\alpha$-poor population at low $|Z|$ to one dominated by $\\alpha$-rich stars at high $|Z|$, as well as a change in mean metallicity with distance from the plane.  \\subsection{Comparison with Galaxy Models} \\label{sec:galmodel_mdfcomp}  Models of chemical and dynamical evolution have evolved substantially since the simple closed box model of \\citet{schmidt}. Despite these improvements, models of Galaxy development remain hamstrung by the lack of extensive observational samples and are calibrated on samples from the local volume. Until SEGUE, there were few samples available to test and refine their predictions for the Milky Way properties far from the plane of the Galaxy. With our G- and K-dwarf MDFs, we can test existing Galaxy models and provide critical constraints on the metallicity structure with respect to location, to improve the model predictions and examine the star-formation history of the disk far beyond the local volume. In particular, the range of $|Z|$ heights makes the sample ideal for examining the transition between the thin and thick disk.  We compare our MDFs with predictions from two Galaxy models. Those of Sch{\\\"o}nrich \\& Binney (SB, 2009a,b) combine both chemical and dynamical evolution, whereas the TRILEGAL 1.4 model is based on empirical relationships. The models differ importantly in their modeling of the thick disk.  While the TRILEGAL model treats the thick disk as a separate component of the Galaxy and allows different input parameters for its spatial distribution and local density, the SB models implement the radial migration of disk stars, forming a thick disk in the solar neighborhood via outward migration of older, inner disk stars. Thus, the SB model's metallicity distribution at high $|Z|$ will be constrained by the chemistry of old, inner disk stars, while TRILEGAL's thick disk metallicity distribution is fixed.  We are also currently working on a comparison with the Besan\\c{c}on Galaxy model \\citep{robin03} (Schlesinger et al., in preparation). We also have provided numerical tables of our MDFs (Tables \\,\\ref{tab:total_gkmdf}, \\,\\ref{tab:gd_zbin_st}, \\,\\ref{tab:kd_zbin_st}, \\,\\ref{tab:gd_zbin_vc}, \\,\\ref{tab:kd_zbin_vc}) for comparison with future Galaxy models (e.g., \\citealt{calura12}).  \\subsubsection{The Sch{\\\"o}nrich \\& Binney Galaxy Models} \\label{sec:schoenrich_models}  The models of \\citet{schonrich09a, schonrich09b} rely on assumed star-formation laws, nucleosynthetic yields, and a Salpeter initial mass function \\citep{salpeter55}.  These models base their star-formation rate on the Kennicutt law \\citep{kennicutt98}; \\citet{schonrich09a, schonrich09b} model the expected nucleosynthetic yields in conjunction with inflows, resulting in a predicted age-metallicity relationship, rather than building off an empirical version.  In addition, these models include dynamical behavior, specifically, radial mixing, which alters the modeled chemical structure. \\citet{schonrich09b} compare the modeled vertical density structure to that of \\citet{juric08} and \\citet{ivezic08}, finding that the model accurately predicts the density structure of various Galaxy components. Finally, these models are built using the BaSTI isochrones \\citep{pietr04}, which they artificially widen to simulate the photometric errors of SDSS. Thus, the expected uncertainties derived from applying our techniques on these models should accurately reflect the photometric errors of the SEGUE sample. We do not factor in the other uncertainties in the SEGUE sample, such as [Fe/H] and distance, into the SB models.  The SB model has been calibrated in [Fe/H] with the Geneva-Copenhagen survey \\citep{nordstrom04, haywood08} and with updated parameters from \\citet{casagrande11}. \\citet{schonrich09a} have shown that their model accurately reflects the properties of stars in the solar neighborhood. They compare their model predictions to other observed samples, such as \\citet{fuhrmann98}, \\citet{bensby03}, \\citet{reddy06}, and \\citet{haywood08}, to constrain chemical parameters. Although many of these samples contain thick-disk stars, they are all limited to the solar neighborhood, which is thin-disk dominated. The GCS sample is typically within 200 pc \\citep{casagrande11}. Similarly, the farthest distance probed by the \\citet{reddy06} sample, which was selected based on Hipparcos data, is 150 pc. Thus, although the SB Galaxy model predicts the structure of the disk as a whole, it is only calibrated on the local volume.   \\subsubsection{The TRILEGAL Galaxy Models} \\label{sec:trilegal_models}  Whereas the SB models simulate the disk's chemical and dynamical evolution, the TRILEGAL 1.4 models are more empirical in scope, utilizing a derived luminosity function, built from an assumed star-formation rate, age-metallicity relation, and initial mass function. These parameters vary slightly for different components of the Galaxy.  The TRILEGAL 1.4 program allows us to simulate different lines of sight and assume a range of Galaxy parameters. We manufacture a TRILEGAL model for each SEGUE line of sight, using a limiting magnitude of 23.0 in $r$, no dust extinction, and no binary contamination over a circular 7 square degree area of the sky.  We specify an exponential thin disk with a two-step star-formation rate. The thin disk is assigned a scale length of 2800 pc and a scale height which increases with stellar age, as described in \\citet{girardi05}. The assumed age-metallicity relationship for this component is based on the empirical determination from \\citet{rp00}, and includes $\\alpha$-enhancement at lower [M/H]. The thick disk is also modeled as an exponential, with a scale height of 800 pc and a scale length of 2800 pc. In contrast to the thin disk, a constant star-formation rate and a Z of 0.008 ([Fe/H]$=-$0.7 for an $\\alpha$-rich population) and $\\sigma$[M/H] of 0.1 dex (which is typically smaller than metallicity measurement errors) is assumed for the thick disk. In addition to the thin and thick disk, we specify an oblate spheroid halo and no central bulge. These properties are combined with the Padova evolutionary tracks to produce a derived luminosity function. Note that the TRILEGAL models predicts the total metallicity, [M/H], rather than [Fe/H].  This model shows good agreement between the observed and predicted star counts when compared to various stellar catalogs, such as 2MASS and Hipparcos \\citep{girardi05}. There has been little chemical calibration of the TRILEGAL model.  \\subsubsection{Accuracy of Galaxy Models} \\label{sec:model_accuracy}  We compare the G-dwarf metallicity distribution to the SB and TRILEGAL Galaxy models in Figures\\,\\ref{fig:scho_model_comp} and\\,\\ref{fig:tri_model_comp}, respectively. For both models, we combine all of the SEGUE-modeled lines of sight and extract all simulated stars that fulfill the color, magnitude, and distance criteria for our G-dwarf spectral type.  Both of these models are calibrated to past analyses, which are typically limited to the local volume. Thus, we expect agreement between the model and SEGUE sample close to the plane of the Galaxy. However, Figures\\,\\ref{fig:scho_model_comp} and\\,\\ref{fig:tri_model_comp} indicate that, although there is substantial overlap in the range of [Fe/H], both models are more metal-rich and sharply peaked than our observed sample. The observed sample is broadened significantly by correlated errors in [Fe/H] and distance, in addition to general uncertainites in [Fe/H] (see \\S\\,\\ref{sec:monte_carlo} and Figure\\,\\ref{fig:correlated_mdf_err}). The models do not simulate these uncertainties and thus exhibit a narrower MDF.  At higher distances from the plane, the metallicity discrepancy between the SB model and the SEGUE sample increases. Whereas the SEGUE distributions become more metal-poor with increasing $|Z|$, the models remain at an [Fe/H] of $-$0.1 for SB.  In fact, the SB distributions become slightly more metal rich at greater distances from the plane. We find similar discrepancies for the K-dwarf sample. We compare the SB modeled lines of sight with the $\\alpha$-separated stellar populations from \\citet{lee11b} in Figure\\,\\ref{fig:scho_model_comp}. The model agrees well with the $\\alpha$-poor distribution alone but are unable to accurately represent the $\\alpha$-rich component.  Rather than [Fe/H], the TRILEGAL model predicts [M/H], which includes contributions from [$\\alpha$/Fe]. The modeled [M/H] will thus be systematically higher than our [Fe/H] at large $|Z|$, where [$\\alpha$/Fe] is expected to be large. We model the relationship between [$\\alpha$/Fe] and [Fe/H] using a simple linear relationship based on the work of \\citet{lee11b}, and remove the expected [$\\alpha$/Fe] contribution from the [M/H] value of each simulated TRILEGAL star. We assume that stars of solar metallicity have [$\\alpha$/Fe] of 0. As [Fe/H] decreases, [$\\alpha$/Fe] increases to 0.20 at [Fe/H]$=-$0.5, and 0.40 at [Fe/H]$=-$1.0. Without this correction, the distribution is similar to that of SB, where the model is systematically more metal rich than the observations at all heights. When we account for [$\\alpha$/Fe], the match between our observed G-dwarf MDF and the model improves, however, TRILEGAL is still offset in [Fe/H], by around 0.1--0.2 dex, at all heights (Figure\\,\\ref{fig:tri_model_comp}). We find that the model does not adequately simulate the structure of the $\\alpha$-rich population at large heights. In particular, from 1.5 to 2.3 kpc in $|Z|$, the TRILEGAL distribution is very similar to that of the \\citet{lee11b} $\\alpha$-poor component. Note that we can improve the match between the simulation and observations by adjusting the contribution from [$\\alpha$/Fe].  The two models predict stars at large $|Z|$ to be more metal rich than we observe. The current inability of these Galaxy models, in particular the SB model, to predict the chemical structure of the disk points to the danger of assuming that the local volume reflects the properties over the span of the disk. Past analyses of the metallicity distribution, used to calibrate these Galaxy models, were limited to the solar neighborhood; they were dominated by thin-disk stars. While these samples contain thick-disk stars passing through the local volume, they may not provide a complete picture of the component's structure. There appears to be a significant disagreement between the picture of the disk structure from local kinematic samples and our \\emph{in-situ} observations. Thus, our unbiased metallicity distributions of the SEGUE G and K dwarfs, and how they change with respect to R and $|Z|$, provide critical constraints for refining and improving Galaxy models, such that we gain a better understanding of the different chemical and dynamical processes occuring in the disk and no longer rely on local samples alone."
    },
    "1112/1112.0027_arXiv.txt": {
        "abstract": "Microlensing has proven an effective probe of the structure of the innermost regions of quasars, and an important test of accretion disk models. We present light curves of the lensed quasar \\hefourlong\\ in the $R$ band and in the ultraviolet, and consider them together with X-ray light curves in two energy bands that are presented in a companion paper. Using a Bayesian Monte Carlo method, we constrain the size of the accretion disk in the rest-frame near- and far-UV, and constrain for the first time the size of the X-ray emission regions in two X-ray energy bands. The $R$-band scale size of the accretion disk is about $10^{15.53}$\\,cm ($\\sim$$46r_g$), slightly smaller than previous estimates, but larger than is predicted by the standard thin disk model. In the UV, the source size is weakly constrained, with a strong prior dependence. The UV to $R$-band size ratio is consistent with the thin disk model prediction, with large error bars. In soft and hard X-rays, the source size is smaller than $\\sim$$10^{15}$\\,cm ($\\sim$$13r_g$) at 90\\% confidence. We do not find evidence of structure in the X-ray emission region, as the most likely value for the ratio of the hard X-ray size to the soft X-ray size is unity. The simulations do not support the idea that quasar flux variability is due to coherent changes in accretion disk area. Finally, we find that the most likely value for the mean mass of stars in the lens galaxy is $\\sim$$0.3M_\\odot$, consistent with other studies. ",
        "introduction": "\\label{sec:intro}  The structure of the innermost regions of quasars and active galactic nuclei, where the X-ray and continuum optical and ultraviolet emission originates, is not well understood. The thin accretion disk model \\citep{Shakura:1973p337, Novikov:1973p343} is widely used to explain the optical and \\ac{UV} continuum, but has had difficulty reproducing its observed spectral shape \\citep[e.g.,][]{Blaes:2001p560}. Some progress has been made using quasar variability \\citep[e.g.,][]{Kelly:2009p895, MacLeod:2010p1014, Dexter:2011pL24} and ``disk reverberation'' \\citep[e.g.,][]{Collier:2001p1527}, but the issue remains unclear. The structure of the X-ray emitting regions is even more mysterious, since the maximum temperature at the inner disk edge is too low to produce X-rays. Instead, a jet or a hot corona above the disk have been suggested as the X-ray source, giving rise to direct and disk-reflected spectral components \\citep[see, e.g., the review by][]{Reynolds:2003p389}. Decisive tests of these models have proven elusive.  Gravitational microlensing of strongly lensed quasars has emerged as an effective technique for measuring the spatial structure of quasar accretion disks. The stars in the foreground lens galaxy magnify (or demagnify) the multiple quasar images, causing flux variability that is not correlated as would be expected from intrinsic quasar variability. The ratio of the angular size of the emission region to the Einstein radius of a star in the lens galaxy determines the strength of the variability, with large source sizes smoothing out the light curves. Since the accretion disk temperature falls with radius, the source is larger at longer wavelengths, and the microlensing effect is weaker. For a standard thin disk, where the effective temperature scales as the $\\beta = 3/4$ power of the radius, the characteristic size of the disk varies with wavelength as $\\lambda^{1/\\beta}$. Several studies have used microlensing to put limits on accretion disk sizes, generally finding that they are larger at fixed wavelength than predicted by the thin disk model, and much larger than sizes calculated using the optical flux \\citep{Pooley:2007p19, Anguita:2008p327, Morgan:2010p1129}. In particular, \\citet{Morgan:2010p1129} find that disk sizes scale with black hole mass in a manner consistent with the $M_\\mathrm{BH}^{2/3}$ scaling expected for a roughly constant Eddington fraction. There are also studies of the dependence of the accretion disk size on wavelength, which generally find that the size increases with wavelength \\citep{Poindexter:2008p34, Bate:2008p1955, Eigenbrod:2008p933, Floyd:2009p233, Blackburne:2011p34, Mosquera:2011p145}. One means of solving the size discrepancy is to have a flatter temperature profile, more like the $T \\propto R^{1/2}$ profile of an irradiated disk \\citep[see][]{Morgan:2010p1129}, but so far the uncertainties in the slope estimates from these studies are too large to distinguish between these models. It does appear that scattering on large scales and contamination from broad line emission have too small an effect to resolve the problem \\citep{Dai:2010p278, Morgan:2010p1129, Mosquera:2011p145}.  In general, the observed effects of microlensing are stronger in X-rays than at optical wavelengths, indicating that the X-rays are emitted from a very compact region \\citep[e.g.,][]{Pooley:2007p19,Morgan:2008p755, Chartas:2009p174, Dai:2010p278}. Recently \\citet{Chen:2011pL34} demonstrated that the hard X-rays from the quasar \\qtwotwolong\\ are even more strongly microlensed than the soft X-rays, indicating that there may be temperature structure within the electron gas that presumably generates the X-ray emission.  At a redshift $z_S = 1.689$, \\hefourlong\\ \\citep[][hereafter \\hefour]{Wisotzki:2002p17, Morgan:2005p2531} is lensed by a foreground early-type galaxy at a redshift of $z_L = 0.46$ into four images arranged in a nearly perfect cross configuration. In this paper, we present light curves of the four images of the quasar in the $R$ band and in the \\ac{UV}, and additionally use light curves in two X-ray energy bands from a companion paper, Chen~et~al.~(2012, in preparation).  We use these light curves to constrain the size of the source at all four wavelengths. The whole is greater than the sum of the parts: the simultaneous analysis of all four wavelengths yields better constraints than any separate ``monochromatic'' analysis. In addition, combining the \\ac{UV} emission, which is thought to originate in the inner regions of the accretion disk, with the optical emission from larger radii gives us leverage to probe the change in size of the disk with wavelength. Likewise, because we have both hard and soft X-ray light curves, we can for the first time put a quantitative constraint on the ratio of sizes in the two energy bands. Finally, we put weak constraints on the mean mass of the stars in the lensing galaxy. We outline the data in Section~\\ref{sec:data}, describe our analysis method in Section~\\ref{sec:method}, and discuss the results in Section~\\ref{sec:results}. ",
        "conclusions": ""
    },
    "1112/1112.2678_arXiv.txt": {
        "abstract": "We have studied the effects of thermal conduction on the structure of viscous and resistive advection-dominated accretion flows (ADAFs). The importance of thermal conduction on hot accretion flow is confirmed by observations of hot gas that surrounds Sgr A$^*$ and a few other nearby galactic nuclei. In this research, thermal conduction is studied by a saturated form of it, as is appropriated for weakly-collisional systems. It is assumed the viscosity and the magnetic diffusivity are due to turbulence and dissipation in the flow. The viscosity also is due to angular momentum transport. Here, the magnetic diffusivity and the kinematic viscosity are not constant and vary by position and $\\alpha$-prescription is used for them. The govern equations on system have been solved by the steady self-similar method. The solutions show the radial velocity is highly subsonic and the rotational velocity behaves sub-Keplerian. The rotational velocity for a specific value of the thermal conduction coefficient becomes zero. This amount of conductivity strongly  depends on magnetic pressure fraction, magnetic Prandtl number, and viscosity parameter. Comparison of energy transport by thermal conduction with the other energy mechanisms implies that thermal conduction can be a significant energy mechanism in resistive and magnetized ADAFs. This property is confirmed by non-ideal magnetohydrodynamics (MHD) simulations. ",
        "introduction": "The observational features in active galactic nuclei (AGN) and X-ray binaries can be successfully explained by the standard geometrically thin, optically thick accretion disc model (Shakura \\& Sunyaev 1973).  The motion of the matter in the standard thin model of the accretion disc is approximately Keplerian, and the energy released in the accreting gas is radiated away locally. In the past two decades, another type of accretion flow has been considered that the energy released due to heating processes in the flow may be trapped within accreting gas. As, only the small fraction of the energy released in the accretion flow is radiated away due to inefficient cooling, and most of the energy is stored in the accretion flow and advected to the central object. This type of accretion flow is called as advection-dominated accretion flow (ADAF). The models of ADAF have been studied by a number of researchers (e.g. Ichimaru 1977; Rees et al. 1982; Narayan \\& Yi 1994; Abramowicz et al. 1995; Blandford \\& Begelman 1999; Ogilvie 1999).  The observations of black holes confirm existence of hot accretion flow that contrasted with classical cold and thin accretion disc model (Shakura \\& Sunyaev 1973). Hot accretion flows can be seen in supermassive black holes of galactic nuclei and during quiescent of accretion onto stellar-mass black holes in X-ray transients (e.g., Lasota et al. 1996; Esin et al. 1997, 2001; Narayan et al. 1998a; Menou et al. 1999;  Di Matteo et al. 2000; see Narayan et al. 1998b; Melia \\& Falcke 2001; Narayan 2002; Narayan \\& Quataert 2005 for reviews). \\textit{Chandra} observations provide constraints on the density and temperature of gas at or near the the Bondi capture radius in Sgr A$^*$ and several nearby galactic nuclei (Loewenstein et al. 2001; Baganoff et al. 2003; Di Matteo et al. 2003; Ho et al. 2003). Tanaka \\& Menou (2006) exploited these constraints to calculate mean free path of the observed gas. They suggested  accretion in such systems are under weakly collisional condition. Moreover, they suggested thermal conduction as a possible mechanism by which the sufficient extra heating is provided in hot accretion flows. In following, Johnson \\& Quataert (2007) studied the effects of electron thermal conduction on the properties of a spherical hot accretion flows. Their model is applicable for Sgr A$^*$ in the Galactic centre. Because, electron heat conduction is important for low accretion rate systems. They also found a supervirial temperature in the presence of thermal conduction. Similar to Tanaka \\& Menou (2006), they assumed a steady state model, but they solved their equations numerically. Abbassi et al. (2008) presented a set of self-similar solutions for ADAFs with a toroidal magnetic field in which the saturated thermal conduction has a important role in the energy transport in the radial direction.  Since the observational evidences and magnetohydrodynamics (MHD) simulations have expressed the toroidal magnetic field and the magnetic diffusivity are important in accretion flows (see Faghei 2011a and references therein),  Faghei (2011a) examined the self-similar solutions of viscous and resistive ADAFs in the presence of a toroidal magnetic field. But, he did not consider the effects of thermal conduction in his model. While, the recent studies of resistive accretion flows have represented that thermal conduction can play an importance role in such systems (e.g. Sharma et al. 2008; Ghanbari et al. 2009). Sharma et al. (2008) studied the effects of thermal conduction on magnetized spherical accretion flows using global axisymmetric MHD simulations. In their model, when the magnetic energy density becomes comparable to the gravitational potential energy density, the plasma due to resistivity is heated to roughly the virial temperature, the mean inflow becomes highly subsonic, and most of the energy released  by accretion is transported to large radii by thermal conduction and the accretion rate became much smaller than Bondi accretion rate.  Moreover, they found that for a larger values of conductive parameter, energy transport through thermal conduction becomes the dominant energy transport mechanism at small radii. Ghanbari et al. (2009) studied a two-dimensional advective accretion disc bathed in the poloidal magnetic field of a central accretor in the presence of thermal conduction. They did not consider toroidal component of magnetic field and for simplicity assumed the resistivity as a constant. They studied induction equation of magnetic field in a steady state that is not according to anti-dynamo theorem (e.g. Cowling 1981) and is useful only in particular systems where the magnetic dissipation time is much longer than the age of the system. On the other hand, this assumption implies that the flow is in balance between escape and creation of the magnetic field.  In this paper, we adopt the presented solutions by Tanaka \\& Menou (2006) and Faghei (2011a). Thus, we will investigate the influences of thermal conduction on a viscous and resistive ADAF in the presence of a toroidal magnetic field. Moreover, it is assumed that magnetic diffusivity in the present model is not constant, and escaping and creating of magnetic field are unbalanced. From some aspects will be shown that the present model is in according with the observations and the resistive MHD simulations. The paper is organized as follows. In section 2, the basic equations of constructing a model for quasi-spherical magnetized advection dominated accretion flow with thermal conduction will be defined. In section 3, self- similar method for solving equations which govern the behaviour of the accreting gas was utilized. The summary of the model will appear in section 4. ",
        "conclusions": ""
    },
    "1112/1112.0482_arXiv.txt": {
        "abstract": "{The outer Galaxy, where the environmental conditions are different from the solar neighbourhood is a laboratory in which it is possible to investigate the dependence of star formation process on the environmental parameters.} {We investigate the X-ray properties of NGC\\,1893, a young cluster { $ (\\sim 1-2$ Myr)} in the outer part of the Galaxy ( galactic radius $\\geq$ 11 kpc) where we expect differences in the disk evolution and in the mass distribution of the stars, to explore the X-ray emission of its members and compare it with that of young stars in star forming regions  near to the Sun.} {We analyze 5 deep \\textit{Chandra} ACIS-I observations with a total exposure time of 450 ks.  Source events of the 1021 X-ray sources have been extracted with the IDL-based routine ACIS-Extract. Using spectral fitting and  quantile analysis of X-ray spectra, we derive X-ray luminosities and compare the respective properties of Class II and Class III members. We also evaluate the variability of sources using the Kolmogorov-Smirnov test and we identify flares in the lightcurves. } {The X-ray luminosity of NGC\\,1893 X-ray members is in the range $10^{29.5}-10^{31.5}$ erg s$^{-1}$. Diskless stars are brighter in X-rays than disk-bearing stars, given the same bolometric luminosity. We found that 34\\% of the 1021 lightcurves appear variable and that they show $0.16$ flare per source, on average. Comparing our results with those relative to the Orion Nebula Cluster, we find that, accounting for observational biases, the X-ray properties of NGC\\,1893 and the Orion ones are very similar. } {The X-ray properties in NGC\\,1893 are not affected by the environment and the  stellar population in the outer Galaxy may have the same coronal properties of nearby star forming regions. The X-ray luminosity properties and the X-ray luminosity function appear to be universal and can therefore be used for distance estimations  and for determining stellar properties as already suggested by Feigelson and collaborators.} ",
        "introduction": "\\label{intro_x} The formation of stars, the evolution of circumstellar disks and eventually the formation of planets are key research topics in modern astrophysics. So far, several star forming regions of different ages and in different environment conditions have been studied in order to explore the parameter space as well as possible. Although several questions have been answered, some aspects of the star formation mechanism and initial stellar evolution are still obscure.  A key question is whether and how star formation depends on the environmental conditions. Since the Jeans mass depends on gas temperature, chemical abundances, and density \\citep[cf.][]{Elmegreen_02, Elmegreen_04}, some dependencies on the physical conditions of star-forming clouds are expected. Moreover, the observed variation in the initial mass function (IMF) at low stellar masses across different environments \\citep{Reyle_01, Gould_98} suggests that gravity is not the only important parameter in the cloud fragmentation and subsequent evolution. As a consequence, other processes, e.g. turbulence and/or magnetic field strength in the cloud have been considered \\citep{Shu_87,Padoan_99,Maclow_04}.  In this context, a natural way to explore the differences due to the environment is to look at the outer Galaxy. Molecular clouds and very young stellar associations are present also at large distances from the Galactic Center \\citep{Snell_02}, implying that star formation is occurring in these regions. Conditions in the outer Milky Way are thought to be less favorable to star formation, when compared with those in the solar neighborhood: the average surface and volume densities of atomic and molecular hydrogen are much smaller than in the solar neighborhood or in the inner Galaxy \\citep{Wouterloot_90}, the interstellar radiation field is weaker \\citep{Mathis_83}, prominent spiral arms are lacking and there are fewer supernovae to act as external triggers of star formation. Metal content is, on average, smaller \\citep{Wilson_92}, decreasing radiative cooling and therefore increasing cloud temperatures and consequently pressure support. Moreover, the pressure of the inter-cloud medium in the outer Galaxy is smaller \\citep{Elmegreen_89}.  Given these relevant differences, the study of star forming regions in the outer Galaxy is likely to be helpful to ascertain the role of environmental conditions in the star formation process. First of all, the environmental conditions could affect the evolution of protoplanetary disks: \\citet{Yasui_09} measured, in the extreme outer Galaxy, a disk frequency significantly lower than in the inner part of the Milky Way. The eventual dependence of disk lifetimes on metallicity could affect the formation of planets and may give also insight into the ``planet-metallicity correlation'', since stars known to harbor giant planets appear to preferentially be metal rich \\citep[][and references therein]{Santos_03,Fischer_05}. Moreover, a recent theoretical works \\citep{Ercolano_09} shows that the disk lifetime in regions of different metallicity can give insight into the general mechanism of disk dispersal, allowing to discriminate between the two currently dominant models of disk dispersal: photoevaporation and planet formation.%  In principle, the metal abundance differences between the inner and the outer Galaxy could also affect the coronal properties. Indeed, the thickness of the outer convection zone in low mass stars is a function of metallicity \\citep{Pizzolato_01}. This, in turn, likely changes the dynamo process itself, i.e. the  X-ray emission engine. Moreover, differences in the composition of X-ray emitting plasma may produce other differences: X-ray radiative losses, for example, are dominated by radiation of highly ionized atoms and therefore we may expect changes in the emission measure, as well as, in individual lines. Finally, if there was a difference in the disk fraction, we would expect an indirect effect on the global X-ray emission:  Class II stars have lower X-ray luminosities than Class III  stars \\citep{Stelzer_01, Flaccomio_03b, Stassun_04_a, Preibisch_05_b,Flaccomio_06, Telleschi_07} thus a different fraction of disked stars may affect the global X-ray luminosity function. The assumption of a universal X-ray luminosity function has been proposed by \\citet{Feigelson_05b} and used to estimate the distance \\citep{Kuhn_10} and total populations of several young clusters \\citep[e.g.][]{Getman_06, Broos_07}. In this context, the study of a cluster lying in a very different environment could be helpful in demonstrating the universality of X-ray properties of young star forming regions and justify the use of X-ray luminosity functions to determine the properties of young clusters.  In order to investigate the X-ray properties of a cluster in the outer Galaxy and compare them to similar cluster in the inner regions of the Galaxy, we have identified, as a suitable target, the young ($\\sim$ 1-2 Myr) cluster NGC\\,1893, whose galactocentric distance is \u2265 11 kpc \\citep{Prisinzano_11}. NGC\\,1893 is in a spiral arm different than that of the Sun and it is located near the edge of the Galaxy, where the environmental condition are quite different from those of the solar neighborhood. Situated in the Aur OB2 association toward the Galactic anti-center, NGC\\,1893 is associated with the HII region IC 410 and the two pennant nebulae, Sim19 and Sim130 \\citep{Gaze_52}, known as ``Tadpole''. It contains a group of early-type stars with some molecular clouds but only moderate extinction. Several studies \\citep{Vallenari_99, Marco_02, Sharma_07, Caramazza_08, Prisinzano_11} demonstrated that star formation is still ongoing and therefore we can study the X-ray properties of the PMS stars, with particular focus on the differences between the Class II and the Class III stars X-ray emission, given that the disk evolution may depend on the environment \\citep{Yasui_09}. The morphology, the age distribution of the cluster and the star formation history have been studied in detail by \\citet{Sanz-Forcada_11}. They find that the  cluster has not relaxed yet to a spherical distribution, and it still has different episodes of stellar formation. The analysis of the age and disc frequency of the objects in NGC\\,1893, in relation with the massive stars and the nebulae, revealed an ongoing stellar formation close to the dark molecular cloud and the two smaller nebulae Sim 129 and Sim 130 \\citep{Sanz-Forcada_11}. While both massive and low mass stars seem to form in the vicinity of the denser molecular cloud, Sim 130 and specially Sim 129 harbour the formation of low mass stars \\citep{Sanz-Forcada_11}. A parameter that is still very uncertain for this cluster is the metallicity. As we discussed in  \\citet{Prisinzano_11}, the literature studies about the metallicity of NGC\\,1893 have ambiguous results. In a paper focused on the Galactic metallicity, \\citet{Rolleston_00} found for 8 NGC\\,1893 members a slight indication of under solar abundances, \\citet{Daflon_04} show a marginal indication of subsolar metallicity for 2 members of the cluster, but these studies do not show a strong evidence of subsolarmetallicity, therefore, as assumed also in \\citet{Prisinzano_11}, we adopt for NGC\\,1893 a solar metallicity.  This study is part of a large project aimed at investigating the star formation processes in the outer Galaxy and is based on a large observational campaign on the young cluster NGC\\,1893. The first results, concerning the joint Chandra-Spitzer large program ``The Initial Mass Function in the Outer Galaxy: the star forming region NGC\\,1893'' (P.I. G. Micela), have been  presented in \\citet{Caramazza_08}. In that earlier work we  have  analyzed the {\\it Spitzer} IRAC maps of NGC\\,1893 joined to the {\\it Chandra} ACIS-I long exposure observation, giving the most complete census of the cluster to date. The IRAC observations of NGC\\,1893 allowed us to identify the members of the cluster with infrared excesses (Class 0/I and Class II stars), while the ACIS observation was used to select the Class III objects that have already lost their disks and have no prominent infrared excess. Although that first catalog of members was incomplete, the presence of 359 members indicates that NGC\\,1893 is quite rich, with intense star forming activity despite the ``unfavorable'' environmental conditions in the outer Milky Way. In \\citet{Prisinzano_11}, we continued the study of  the properties of NGC\\,1893 by using deep optical and JHK data and compiling a catalog extending from X-rays to NIR data. In this second work, we have assessed the membership status of each star, finding 415 diskless candidate members plus 1061 young stellar objects with a circumstellar disk or Class II candidate members, 125 of which are also H$_\\alpha$ emitters. Moreover, optical and NIR photometric properties have been used to evaluate the cluster parameters. Using the diskless candidate members, the cluster distance has been found to be 3.6$\\pm$0.2 kpc and the mean interstellar reddening E(B-V)=0.6$\\pm$0.1 with evidence of differential reddening across the whole surveyed region. These previous studies show that NGC\\,1893 contains a conspicuous population of pre-main sequence stars with a disk fraction of about 70\\%, similar to that found in clusters of similar age in the solar neighborhood. This demonstrates that, despite the expected unfavorable conditions for star formation, very rich young clusters can form also in the outer regions of our Galaxy.  In the present work, as part of the same project,  we present the in-depth analysis of  450 ks long {\\it Chandra} ACIS-I observation, already introduced in \\citet{Caramazza_08}, focusing on the X-ray properties of the 1021 detected sources in light of the previous results of \\citet{Caramazza_08}  and \\citet{Prisinzano_11}. To put our work in context, we will compare the present study with the 850 ks long {\\it Chandra} ACIS-I observation of the Orion Nebula Cluster (ONC), termed Chandra Orion Ultradeep Project (COUP) \\citep{Getman_05}. The COUP with 1616 detected sources provided one of the most  comprehensive dataset so far acquired on the X-ray emission of PMSs. Several studies were conducted on the 1616 COUP detected sources, in order to characterize the X-ray emission and to understand the coronal processes in the PMS stars. For example, \\citet{Preibisch_05_b} have analyzed the dependence of the X-ray luminosity on the stellar properties, finding for the Orion stars that  $L_X$ is a function of the mass, \\citet{Wolk_05} have studied the variability of the solar mass stars, in order to have indication of the emission of the young Sun, \\citep{Favata_05} have investigated the geometry of flare loops, resulting in the first observational indications of magnetic structure connecting the star and the disk. These are just few of the several results of the COUP published in more than 20 papers. Since COUP gives the most comprehensive view of magnetic activity in young stars ever achieved in the nearest rich cluster of very young stars, it is used here as a touchstone for comparison with NGC\\,1893.  In the present paper, we analyse the properties of the X-ray sources in NGC1893, focusing,  in particular, on the X-ray luminosity of low mass stars. We  compare the X-ray properties of the NGC\\,1893 members with infrared excesses (hereafter Class II stars), as defined in \\citet{Caramazza_08} and in \\citet{Prisinzano_11} with the members without infrared excesses (hereafter Class III stars), that are classified as members of the cluster by means of their X-ray emission \\citep[see ][for details]{Caramazza_08, Prisinzano_11}. The outline of the present paper is the following: in \\S \\ref{extraction} we describe the data reduction,  the photon extraction procedure, and present our X-ray catalog. In \\S \\ref{spe_prop} we describe the spectral analysis that we pursued by means of spectral fitting and quantile analysis of X-ray spectra. In \\S \\ref{spectral_properties} we compare the X-ray properties of Class II and Class III of NGC\\,1893 and in \\S \\ref{variability} we present the variability properties of the members. In \\S \\ref{ngc1893_x_results} we discuss our results and in \\S \\ref{summary} we summarize our work. ",
        "conclusions": ""
    },
    "1112/1112.1092_arXiv.txt": {
        "abstract": "We present optical to far-infrared photometry of 31 reddened QSOs that show evidence for radiatively driven outflows originating from AGN in their rest-frame UV spectra. We use these data to study the relationships between the AGN-driven outflows, and the AGN and starburst infrared luminosities. We find that FeLoBAL QSOs are invariably IR-luminous, with IR luminosities exceeding $10^{12}$L$_{\\odot}$ in all cases. The AGN supplies 76\\% of the total IR emission, on average, but with a range from 20\\% to 100$\\%$. We find no evidence that the absolute luminosity of obscured star formation is affected by the AGN-driven outflows. Conversely, we find an anticorrelation between the strength of AGN-driven outflows, as measured from the range of outflow velocities over which absorption exceeds a minimal threshold, and the contribution from star formation to the total IR luminosity, with a much higher chance of seeing a starburst contribution in excess of 25\\% in systems with weak outflows than in systems with strong outflows. Moreover, we find no convincing evidence that this effect is driven by the IR luminosity of the AGN. We conclude that radiatively driven outflows from AGN can have a dramatic, negative impact on luminous star formation in their host galaxies. We find that such outflows act to curtail star formation such that star formation contributes less than $\\sim 25\\%$ of the total IR luminosity. We also propose that the degree to which termination of star formation takes place is not deducible from the IR luminosity of the AGN. ",
        "introduction": "The last decade has seen substantial progress in understanding how galaxies assemble their stars and central supermassive black holes (SMBHs). Recent results almost all point to the same conclusion; that the bulk of galaxy assembly occurred at $z\\gtrsim0.7$, and that a significant fraction of it took place in obscured `bursts' of intense star formation and SMBH accretion. Indirect evidence for this comes from, for example, the stellar ages of low-redshift galaxies \\citep{heavens04}, the existence of evolved elliptical galaxies at high redshifts \\citep{dun96,ell97,rak07}, and studies of stellar mass assembly over wide redshift ranges \\citep{glaz04,fon06,per08,march09,ilb10}. Direct evidence mostly comes from extragalactic infrared and millimetre imaging surveys, which find that the number density of IR-luminous galaxies increases dramatically with increasing redshift (e.g. \\citealt{leflo05,chapman05,babbedge06,coppin06,shupe08,aus09,eales10}), and that much of the growth period of SMBHs was shrouded in dust \\citep{mart05,alex08}.  The importance of obscured starburst and AGN activity in assembling galaxies suggests that the two phenomena may affect each other. A link between them is implied by, for example, the tight correlation between the mass of the SMBH and both stellar bulge mass \\citep{mag98,ferrarese00,tre02,marc03,shields03,har04} and dark matter halo mass \\citep{ferrarese02}, and the coeval presence of both starburst and AGN activity in many IR-luminous systems at all redshifts (see e.g. \\citealt{blain02,lag05,lonsdale06,hernan09}).  \\begin{deluxetable*}{clcccccccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{The Sample. \\label{tablesample}} \\tablewidth{0pt} \\tablehead{ \\colhead{ID} & \\colhead{Name} & \\colhead{z} & \\colhead{SDSS} & \\multicolumn{3}{c}{2MASS (or UKIDSS) } & \\multicolumn{4}{c}{WISE} \\\\ &                          &      & $m_z$ & $m_J$ &    $m_H$ &    $m_K$ & 3.4 & 4.6 & 12 & 22 } \\startdata 1  & SDSS J011117.34+142653.6 & 1.15 & 17.37 & \t\t        16.05 &                  15.53 &                     14.78 & 13.40 & 12.15 & 9.22  & 6.80   \\\\ 2  & SDSS J024254.66-072205.6 & 1.22 & 19.03 & \t\t        --    &                  --    &                     --    & 14.36 & 13.04 & 10.42 & 8.97 (2.2$\\sigma$)    \\\\ 3  & SDSS J030000.57+004828.0 & 0.89 & 16.13 &\t\t        15.09 &                  14.59 &                     14.11 & 12.68 & 11.40 & 8.31  & 5.96    \\\\ 4  & SDSS J033810.84+005617.7 & 1.63 & 18.38 & 17.77\\tablenotemark{a} & 17.32\\tablenotemark{a} &    16.93\\tablenotemark{a} & 16.06 & 14.42 & 11.07 & 8.34    \\\\ 5  & SDSS J081312.60+432640.0 & 1.09 & 18.82 &                  --    &   \t         --    &                     --    & 15.15 & 13.81 & 10.65 & 8.51    \\\\ 6  & SDSS J083522.76+424258.3 & 0.81 & 17.24 & \t\t        15.95 &   \t         15.70 &                     15.05 & 13.41 & 12.17 & 9.43  & 6.88    \\\\ 7  & SDSS J084044.41+363327.8 & 1.23 & 16.11 & \t\t        15.02 &   \t         14.39 &                     13.89 & 12.74 & 11.46 & 8.74  & 6.11    \\\\ 8  & SDSS J091103.49+444630.4 & 1.30 & 19.19 & \t\t        --    &   \t         --    &                     --    & 14.43 & 12.65 & 8.97  & 6.57    \\\\ 9  & SDSS J091854.48+583339.7 & 1.32 & 18.89 &\t                --    &   \t         --    &                     --    & 15.44 & 14.35 & 12.38 & 8.76 (1.3$\\sigma$)  \\\\ 10 & SDSS J100605.66+051349.0 & 0.97 & 18.28 & 16.62\\tablenotemark{a} & 16.39\\tablenotemark{a} &    15.26\\tablenotemark{a} & 13.37 & 12.06 & 9.14  & 6.72    \\\\ 11 & SDSS J101927.36+022521.4 & 1.36 & 18.00 & \t\t        16.37 &                  15.22 &    15.04\\tablenotemark{a} & 14.03 & 12.78 & 9.49  & 6.74    \\\\ 12 & SDSS J102036.10+602338.9 & 0.99 & 17.91 & \t\t        16.56 &\t              $<$15.60 &                     15.40 & 13.99 & 12.57 & 9.54  & 7.14    \\\\ 13 & SDSS J102358.97+015255.8 & 1.08 & 18.91 &\t\t        --    & 17.40\\tablenotemark{a} &    16.81\\tablenotemark{a} & --    & --    & --    & --      \\\\ 14 & SDSS J105748.63+610910.8 & 1.28 & 19.55 & \t\t        --    &                  --    &                     --    & 15.02 & 13.94 & 11.06 & 8.93    \\\\ 15 & SDSS J112526.12+002901.3 & 0.86 & 17.68 & \t \t        16.51 &               $<$16.32 &                     15.36 & --    & --    & --    & --      \\\\ 16 & SDSS J112828.31+011337.9 & 0.89 & 18.02 & 16.55\\tablenotemark{a} & 16.55\\tablenotemark{a} &    15.91\\tablenotemark{a} & 14.69 & 13.50 & 10.60 & 8.22    \\\\ 17 & SDSS J112901.71+050617.0 & 1.29 & 18.93 & \t\t        --    & 16.15\\tablenotemark{a} &    15.54\\tablenotemark{a} & 14.67 & 13.74 & 11.55 & 8.58 (1.4$\\sigma$)    \\\\ 18 & SDSS J114556.26+110018.4 & 0.93 & 18.53 & \t\t        --    & 16.85\\tablenotemark{a} &    16.32\\tablenotemark{a} & --    & --    & --    & --      \\\\ 19 & SDSS J115436.60+030006.3 & 1.39 & 17.46 &\t\t        16.05 &                  15.42 &                     15.23 & 13.87 & 12.38 &  9.27 & 7.29    \\\\ 20 & SDSS J115852.86-004302.0 & 0.98 & 18.81 &\t\t\t17.06 &                  16.61 &                     15.55 & --    & --    & --    & --      \\\\ 21 & SDSS J120049.55+632211.8 & 0.89 & 18.47 &                  --    &                  --    &                     --    & 15.12 & 14.25 & 11.24 & 8.37    \\\\ 22 & SDSS J120627.62+002335.3 & 1.11 & 18.55 & 17.59\\tablenotemark{a} & 16.76\\tablenotemark{a} &    15.98\\tablenotemark{a} & 14.50 & 13.23 & 10.65 & 8.51    \\\\ 23 & SDSS J121441.42-000137.8 & 1.05 & 18.52 & 17.58\\tablenotemark{a} & 16.86\\tablenotemark{a} &    15.98\\tablenotemark{a} & 14.15 & 12.76 & 10.04 & 8.11    \\\\ 24 & SDSS J123549.95+013252.6 & 1.29 & 18.82 &                  --    & 16.41\\tablenotemark{a} &    16.19\\tablenotemark{a} & 15.00 & 13.89 & 10.93 & 8.61    \\\\ 25 & SDSS J132401.53+032020.6 & 0.93 & 18.48 & 16.68\\tablenotemark{a} & 16.58\\tablenotemark{a} &    15.83\\tablenotemark{a} & 14.41 & 13.34 & 10.84 & 7.28 (1.8$\\sigma$)     \\\\ 26 & SDSS J142703.62+270940.3 & 1.17 & 17.96 &                  --    &                  --    &                     --    & 14.32 & 12.93 & 10.04 & 7.73    \\\\ 27 & SDSS J155633.77+351757.3 & 1.50 & 17.60 &                  15.91 &                  14.91 &                     14.79 & 13.19 & 11.72 &  8.75 & 6.67   \\\\ 28 & SDSS J173753.97+553604.8 & 1.10 & 19.82 &                  --    &                  --    &                     --    & 15.73 & 14.36 & 11.43 & 8.77    \\\\ 29 & SDSS J210712.77+005439.4 & 0.92 & 19.69 &                  --    &                  --    &                     --    & 13.24 & 11.60 &  8.55 & 6.34     \\\\ 30 & SDSS J221511.93-004549.9 & 1.48 & 16.59 &                  15.65 &                  14.89 &                     14.69 & --    & --    & --    & --     \\\\ 31 & SDSS J233646.20-010732.6 & 1.29 & 18.68 &                  17.24 &                  --    &                     16.23 & --    & --    & --    & -- \\enddata \\tablecomments{Positions, redshifts and SDSS magnitudes are taken from the SDSS DR6. For the SDSS magnitudes we used the PSF magnitudes in the AB system. For 2MASS, we used the `default' magnitudes in the Vega system, taken from the public 2MASS All-Sky Point Source Catalogue. WISE magnitudes from the operational database as of June 2011. An `--' indicates that the source is not in the catalogue in that band. } \\tablenotetext{a}{UKIDSS LAS magnitude} \\end{deluxetable*}  Recently, interest into the relationship between starburst and AGN activity has been stimulated by what first appeared to be irreconcilable differences between observational and theoretical results on the assembly history of galaxies. Foremost among these were the difficulties that models faced in explaining the observed galaxy luminosity function (LF) at low and high redshifts simultaneously; if the models were tuned to match the local galaxy LF then they underpredicted the number of massive galaxies observed at high redshifts, whereas the models that matched the number of high redshift galaxies gave poor fits to the local galaxy LF \\citep{kauff99,cole00,somer01,benson03}. Other problems included: (1) the prediction that a higher fraction of rich galaxy clusters should harbour cooling flows than is observed \\citep{peterson03,xu02,peterson06}, (2) the difficulties that early semi-analytic models faced in reproducing the large number of IR-luminous galaxies observed at high redshift (e.g. \\citealt{baugh05}), and (3) the expectation that the mass return rate from stars would cause central SMBHs to be larger than is observed \\citep{ciotti91,ciotti07}.   \\begin{deluxetable*}{lccccccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{MIPS fluxes, absorption strength measures (see Equation \\ref{eqn:bal} \\& Table \\ref{absmeasures}), and IR luminosities. \\label{tableasandlirs}} \\tablewidth{0pt} \\tablehead{ ID&\\multicolumn{3}{c}{MIPS Fluxes (mJy)}           &\\multicolumn{4}{c}{Absorption Strengths (km s$^{-1}$)} &\\multicolumn{3}{c}{Infrared Luminosities (log (L$_{\\odot}$))} & $\\chi^{2}_{red}$ \\\\ & 24$\\mu$m     & 70$\\mu$m      & 160$\\mu$m      & AS$_{0}$ & AS$_{0}^{Gib}$ &AS$_{2}$ & AS$_{4}$ & L$_{Tot}$\\tablenotemark{a} & L$_{AGN}$\\tablenotemark{b} & L$_{SB}$\\tablenotemark{c} & } \\startdata 1  & $16.9\\pm0.8$ & $26.4\\pm7.9$  & $15.4\\pm17.1$  &        0 & --    &        0 &      348 & 13.10$^{+0.03}_{-0.04}$ & 13.06$^{+0.03}_{-0.02}$ & 11.97$^{+0.42}_{-0.40}$ &  0.7 \\\\ 2  & $6.4 \\pm0.3$ & $10.8\\pm4.5$  & $16.2\\pm14.6$  &        0 & 0     &      241 &     2030 & 12.70$^{+0.06}_{-0.07}$ & 12.62$^{+0.03}_{-0.09}$ & 12.00$^{+0.26}_{-0.91}$ &  2.6 \\\\ 3  & $29.4\\pm2.9$ & $56.3\\pm10.6$ & $-8.8\\pm16.0$  &     7001 & 15310 &     7098 &    12722 & 13.12$^{+0.03}_{-0.04}$ & 13.11$^{+0.02}_{-0.12}$ & $<11.95$                &  1.8 \\\\ 4  & $2.4 \\pm0.3$ & $3.6\\pm6.5$   & $18.2\\pm18.0$  &        0 & 0     &        0 &      749 & 12.76$^{+0.22}_{-0.15}$ & 12.62$^{+0.07}_{-0.11}$ & $<$12.23$^{+0.49}$      &  1.1 \\\\ 5  & $5.4\\pm0.5$  & $10.5\\pm3.8$  & $12.7\\pm14.7$  &        0 & --    &        0 &      143 & 12.51$^{+0.08}_{-0.05}$ & 12.41$^{+0.05}_{-0.10}$ & 11.86$^{+0.37}_{-0.54}$ &  0.7 \\\\ 6  & $12.3\\pm0.6$ & $33.1\\pm7.2$  & $16.2\\pm14.9$  &        0 & --    &        0 &     1335 & 12.69$^{+0.03}_{-0.03}$ & 12.60$^{+0.06}_{-0.13}$ & 11.90$^{+0.39}_{-0.21}$ &  1.9 \\\\ 7  & $26.8\\pm1.3$ & $43.7\\pm7.8$  & $21.4\\pm9.7$   &     5958 & 4863  &     8707 &    12164 & 13.33$^{+0.06}_{-0.02}$ & 13.30$^{+0.09}_{-0.01}$ & 12.04$^{+0.35}_{-0.34}$ &  1.4 \\\\ 8  & $16.3\\pm0.6$ & $51.1\\pm7.8$  & $22.0\\pm7.9$   &     3623 & 0     &     4095 &     8156 & 13.25$^{+0.08}_{-0.16}$ & 13.23$^{+0.04}_{-0.20}$ & 12.29$^{+0.09}_{-1.01}$ &  1.7 \\\\ 9  & $1.0\\pm0.3$  & $1.7\\pm6.6$   & $24.4\\pm14.3$  &      770 & --    &     1437 &     3479 & 12.29$^{+0.49}_{-0.22}$ & 12.09$^{+0.05}_{-0.26}$ & $<12.72$                &  1.0 \\\\ 10 & $15.3\\pm0.8$ & $49.2\\pm16.4$ & $73.6\\pm21.0$  &        0 & 0     &      267 &     2540 & 12.97$^{+0.01}_{-0.10}$ & 12.85$^{+0.02}_{-0.11}$ & 12.37$^{+0.17}_{-0.80}$ &  2.3 \\\\ 11 & $16.0\\pm0.8$ & $61.4\\pm9.9$  & $28.7\\pm11.5$  &     3360 & 1803  &     4255 &     8361 & 13.27$^{+0.06}_{-0.13}$ & 13.20$^{+0.09}_{-0.16}$ & 12.40$^{+0.19}_{-0.22}$ &  1.9 \\\\ 12 & $8.8\\pm0.9$  & $23.6\\pm8.1$  & $10.9\\pm12.4$  &     5206 & 5253  &     7372 &    10775 & 12.75$^{+0.05}_{-0.06}$ & 12.70$^{+0.07}_{-0.07}$ & 11.69$^{+0.32}_{-1.12}$ &  0.6 \\\\ 13 & $0.9\\pm0.1$  & $3.7\\pm7.6$   & $14.1\\pm11.7$  &        0 & --    &        0 &     2202 & 12.11$^{+0.20}_{-0.24}$ & 11.85$^{+0.03}_{-0.20}$ & 11.75$^{+0.50}_{-1.28}$ &  0.2 \\\\ 14 & $3.4\\pm0.3$  & $10.6\\pm6.6$  & $36.7\\pm15.4$  &        0 & --    &        0 &      509 & 12.65$^{+0.12}_{-0.21}$ & 12.39$^{+0.07}_{-0.10}$ & 12.30$^{+0.26}_{-1.25}$ &  1.1 \\\\ 15 & $6.8\\pm0.4$  & $15.5\\pm5.7$  & $10.2\\pm15.1$  &     4630 & 3716  &     5357 &     9705 & 12.52$^{+0.04}_{-0.05}$ & 12.48$^{+0.02}_{-0.07}$ & 11.41$^{+0.66}_{-1.21}$ &  0.3 \\\\ 16 & $3.6\\pm0.4$  & $4.4\\pm6.5$   & $8.9\\pm12.7$   &     4348 & 3769  &     5185 &     9020 & 12.28$^{+0.06}_{-0.05}$ & 12.16$^{+0.08}_{-0.04}$ & $<$11.66$^{+0.27}$      &  1.5 \\\\ 17 & $4.0\\pm0.3$  & $14.0\\pm5.4$  & $21.8\\pm11.5$  &     1572 & 673   &     1770 &     5249 & 12.60$^{+0.14}_{-0.11}$ & 12.40$^{+0.01}_{-0.04}$ & 12.15$^{+0.27}_{-0.25}$ &  3.0 \\\\ 18 & $2.2\\pm0.3$  & $1.5\\pm6.4$   & $-13.4\\pm12.2$ &     1498 & --    &     1500 &     4674 & 12.08$^{+0.11}_{-0.03}$ & 12.07$^{+0.10}_{-0.05}$ & $<11.57$                &  0.4 \\\\ 19 & $9.3\\pm0.5$  & $13.5\\pm5.6$  & $15.0\\pm16.5$  &    10369 & 10422 &    12585 &    17989 & 13.12$^{+0.06}_{-0.10}$ & 13.01$^{+0.12}_{-0.03}$ & $<$12.17$^{+12.52}$     &  1.3 \\\\ 20 & $3.7\\pm0.3$  & $-5.5\\pm7.5$  & $15.2\\pm14.2$  &        0 & --    &      377 &      972 & 12.38$^{+0.15}_{-0.06}$ & 12.31$^{+0.05}_{-0.09}$ & $<$11.31$^{+0.91}$      &  0.7 \\\\ 21 & $4.6\\pm0.4$  & $19.0\\pm7.3$  & $9.9\\pm14.0$   &        0 & 37    &      637 &     4627 & 12.12$^{+0.10}_{-0.08}$ & 12.02$^{+0.04}_{-0.13}$ & $<$11.46$^{+0.55}$      &  2.9 \\\\ 22 & $3.7\\pm0.3$  & $4.0\\pm6.5$   & $10.5\\pm14.4$  &     6322 & 5476  &     7889 &    12930 & 12.48$^{+0.10}_{-0.12}$ & 12.34$^{+0.13}_{-0.02}$ & $<$11.78$^{+0.42}$      &  0.6 \\\\ 23 & $4.9\\pm0.4$  & $36.8\\pm9.4$  & $78.9\\pm23.1$  &     1400 & 1531  &     1400 &     3914 & 12.83$^{+0.05}_{-0.17}$ & 12.47$^{+0.06}_{-0.18}$ & 12.58$^{+0.17}_{-0.43}$ &  1.4 \\\\ 24 & $4.7\\pm0.3$  & $30.0\\pm7.5$  & $22.7\\pm10.8$  &     1427 & 734   &     2365 &     5068 & 12.76$^{+0.13}_{-0.09}$ & 12.56$^{+0.06}_{-0.10}$ & 12.31$^{+0.28}_{-0.21}$ &  1.5 \\\\ 25 & $7.2\\pm0.3$  & $0.6\\pm7.9$   & $34.4\\pm13.2$  &      309 & 0     &      901 &     4472 & 12.45$^{+0.08}_{-0.07}$ & 12.38$^{+0.09}_{-0.21}$ & $<$11.66$^{+0.40}$      &  1.8 \\\\ 26 & $5.6\\pm0.3$  & $36.8\\pm6.3$  & $69.0\\pm14.4$  &      732 & --    &     2141 &     4917 & 12.96$^{+0.03}_{-0.04}$ & 12.59$^{+0.10}_{-0.02}$ & 12.69$^{+0.11}_{-0.10}$ &  0.7 \\\\ 27 & $16.5\\pm0.5$ & $31.5\\pm5.1$  & $35.4\\pm18.6$  &    11891 & 13053 &    13932 &    19000 & 13.43$^{+0.02}_{-0.10}$ & 13.39$^{+0.03}_{-0.11}$ & 12.39$^{+0.35}_{-0.35}$ &  0.8 \\\\ 28 & $3.5\\pm0.5$  & $1.8\\pm7.6$   & $27.3\\pm11.4$  &        0 & --    &        0 &      573 & 12.32$^{+0.13}_{-0.15}$ & 12.14$^{+0.06}_{-0.13}$ & $<$11.83$^{+0.37}$      &  1.1 \\\\ 29 & $20.1\\pm0.4$ & $44.1\\pm8.2$  & $60.3\\pm20.8$  &     2422 & 0     &     3742 &     7250 & 13.06$^{+0.02}_{-0.03}$ & 13.02$^{+0.02}_{-0.03}$ & 12.09$^{+0.15}_{-0.86}$ &  0.4 \\\\ 30 & $12.2\\pm0.6$ & $33.6\\pm5.8$  & $42.1\\pm21.0$  &        0 & 551   &        0 &     1926 & 13.35$^{+0.03}_{-0.04}$ & 13.18$^{+0.01}_{-0.02}$ & 12.84$^{+0.08}_{-0.12}$ &  1.4 \\\\ 31 & $0.9\\pm0.2$  & $2.2\\pm6.7$   & $20.2\\pm16.3$  &        0 & --    &        0 &     1318 & 12.74$^{+0.11}_{-0.39}$ & 12.00$^{+0.17}_{-0.39}$ & 12.67$^{+0.11}_{-0.83}$ &  0.8 \\enddata \\tablecomments{Errors on the absorption strengths from uncertainties in the fits are 100-200 km s$^{-1}$, but see \\S\\ref{discusssboutflow} \\& the appendix. The AS$_{0}^{Gib}$ values are an alternative measure of AS$_{0}$ (Table \\ref{absmeasures}), taken from \\citet{gibson09} (a `--' indicates that the object is not in the \\citet{gibson09} sample). Errors are the 90\\% confidence intervals, see \\S\\ref{sectfitting}.} \\tablenotetext{a}{Total infrared luminosity, integrated over 1-1000$\\mu$m in the rest-frame.} \\tablenotetext{b}{Infrared luminosity of the AGN component, integrated over 1-1000$\\mu$m in the rest-frame.} \\tablenotetext{c}{Infrared luminosity of the starburst component, integrated over 1-1000$\\mu$m in the rest-frame.} \\end{deluxetable*}   Several solutions have been proposed to resolve these issues. Among the most promising is an idea termed `AGN feedback'. AGN feedback is the exertion of influence of an SMBH on kpc to Mpc scales to curtail star formation in the host galaxy, and/or further accretion onto the SMBH itself\\footnote{The idea that SMBHs may exert `negative' feedback on inflowing gas to regulate their own growth and that of the host galaxy predated its use to reconcile galaxy evolution models with observations, see e.g. \\citealt{haiman96,silk98}}. This can occur in four ways. Radiation from regions immediately local to the SMBH can (1) heat gas in the ISM so it cannot collapse to form stars, and/or (2) radiatively drive gas out of the galaxy (or, equivalently, stop it falling in from the IGM), thus emptying the galaxy of fuel for further star formation. Additionally, matter expelled from regions immediately local to the SMBH can (3) heat gas in the ISM, and/or (4) drive gas out of the galaxy via kinetic pressure.  Most discussions of AGN feedback condense the four feedback mechanisms described above into two simplified paradigms; `quasar' mode feedback and `radio' mode feedback. Quasar mode feedback assumes that radiation from an accretion disk terminates star formation in the host galaxy, usually by coupling some fraction of the QSO luminosity to kinetic energy injected into the ISM. Quasar mode feedback is thought to be short-term, acting for only the $\\sim10^{8}$ years over which the quasar is accreting rapidly. Radio mode feedback on the other hand is assumed to occur via a relativistic jet that transfers momentum to the ISM. Radio mode feedback can occur over longer timescales ($\\sim10^{9}$ years) than quasar mode as it requires only a low accretion rate to produce jets sufficiently powerful to affect the ISM. Nevertheless, radio mode and quasar mode feedback are not mutually exclusive.  \\begin{figure*} \\includegraphics[width=78mm,angle=0]{SDSS_J011117_dual_plot_jy.ps} \\includegraphics[width=78mm,angle=0]{SDSS_J024254_dual_plot_jy.ps}\\\\ \\includegraphics[width=78mm,angle=0]{SDSS_J030000_dual_plot_jy.ps} \\includegraphics[width=78mm,angle=0]{SDSS_J033810_dual_plot_jy.ps}\\\\ \\includegraphics[width=78mm,angle=0]{SDSS_J081312_dual_plot_jy.ps} \\includegraphics[width=78mm,angle=0]{SDSS_J083522_dual_plot_jy.ps} \\caption{Observed-frame fits to the optical through MIPS SEDs for objects 1-6 in Table \\ref{tablesample}. The solid line is the combined best-fit model, while the dashed and dotted lines are the starburst and AGN components, respectively. A description of the models is in \\S\\ref{subsmodels}. The number in the top right hand corner of each plot is the ID number in Table \\ref{tablesample}.}\\label{figexamplefits1} \\end{figure*}  \\begin{figure*} \\includegraphics[width=78mm,angle=0]{SDSS_J084044_dual_plot_jy.ps} \\includegraphics[width=78mm,angle=0]{SDSS_J091103_dual_plot_jy.ps}\\\\ \\includegraphics[width=78mm,angle=0]{SDSS_J091854_dual_plot_jy.ps} \\includegraphics[width=78mm,angle=0]{SDSS_J100605_dual_plot_jy.ps}\\\\ \\includegraphics[width=78mm,angle=0]{SDSS_J101927_dual_plot_jy.ps} \\includegraphics[width=78mm,angle=0]{SDSS_J102036_dual_plot_jy.ps} \\caption{Observed-frame fits to the optical through MIPS SEDs for objects 7-12 in Table \\ref{tablesample}. Details are the same as for figure \\ref{figexamplefits1}.}\\label{figexamplefits2} \\end{figure*}  \\begin{figure*} \\includegraphics[width=78mm,angle=0]{SDSS_J102358_dual_plot_jy.ps} \\includegraphics[width=78mm,angle=0]{SDSS_J105748_dual_plot_jy.ps}\\\\ \\includegraphics[width=78mm,angle=0]{SDSS_J112526_dual_plot_jy.ps} \\includegraphics[width=78mm,angle=0]{SDSS_J112828_dual_plot_jy.ps}\\\\ \\includegraphics[width=78mm,angle=0]{SDSS_J112901_dual_plot_jy.ps} \\includegraphics[width=78mm,angle=0]{SDSS_J114556_dual_plot_jy.ps} \\caption{Observed-frame fits to the optical through MIPS SEDs for objects 13-18 in Table \\ref{tablesample}. Details are the same as for figure \\ref{figexamplefits1}.}\\label{figexamplefits3} \\end{figure*}  \\begin{figure*} \\includegraphics[width=78mm,angle=0]{SDSS_J115436_dual_plot_jy.ps} \\includegraphics[width=78mm,angle=0]{SDSS_J115852_dual_plot_jy.ps}\\\\ \\includegraphics[width=78mm,angle=0]{SDSS_J120049_dual_plot_jy.ps} \\includegraphics[width=78mm,angle=0]{SDSS_J120627_dual_plot_jy.ps}\\\\ \\includegraphics[width=78mm,angle=0]{SDSS_J121441_dual_plot_jy.ps} \\includegraphics[width=78mm,angle=0]{SDSS_J123549_dual_plot_jy.ps} \\caption{Observed-frame fits to the optical through MIPS SEDs for objects 19-24 in Table \\ref{tablesample}. Details are the same as for figure \\ref{figexamplefits1}.}\\label{figexamplefits4} \\end{figure*}  \\begin{figure*} \\includegraphics[width=78mm,angle=0]{SDSS_J132401_dual_plot_jy.ps} \\includegraphics[width=78mm,angle=0]{SDSS_J142703_dual_plot_jy.ps}\\\\ \\includegraphics[width=78mm,angle=0]{SDSS_J155633_dual_plot_jy.ps} \\includegraphics[width=78mm,angle=0]{SDSS_J173753_dual_plot_jy.ps}\\\\ \\includegraphics[width=78mm,angle=0]{SDSS_J210712_dual_plot_jy.ps} \\includegraphics[width=78mm,angle=0]{SDSS_J221511_dual_plot_jy.ps} \\caption{Observed-frame fits to the optical through MIPS SEDs for objects 25-30 in Table \\ref{tablesample}. Details are the same as for figure \\ref{figexamplefits1}.}\\label{figexamplefits5} \\end{figure*}  \\begin{figure} \\includegraphics[width=78mm,angle=0]{SDSS_J233646_dual_plot_jy.ps} \\caption{Observed-frame fit to the optical through MIPS SEDs for object 31 in Table \\ref{tablesample}. Details are the same as for figure \\ref{figexamplefits1}.}\\label{figexamplefits6} \\end{figure}  For individual galaxies, AGN feedback is predicted to have a significant impact, both from analytic \\citep{fabian99,wyithe03,saz05,fabian06,pope09,power11,kav11} and numerical studies \\citep{omma04,ciotti07,ant08,tort09,ost10,wagner11,hambrick11}. For galaxy mergers, quasar mode feedback is predicted to have a profound influence \\citep{dimatteo05,sprdimher05,debuhr11,snyd11}. Models for the cosmological evolution of galaxies and clusters have incorporated one or both of these feedback paradigms, often with noticeable improvements in reproducing observations. These include semi-analytic models using quasar mode \\citep{granato04,menci08}, radio mode \\citep{bower06,bower08,gonzalez09}, or both \\citep{somer08}, and numerical models with quasar mode \\citep{booth09,sales10,mccar11,teyss11,chat11}, or both forms of feedback \\citep{croton06,deluc06,sija07,kitz07,monaco07,puch08,short09,gasp11}.  Observationally, evidence for AGN feedback is mounting. Studies of early type galaxies suggest that some form of feedback may have terminated their star formation (\\citealt{schaw07,rose09}, but see also \\citealt{shin11}). Far-IR observations have also found powerful, plausibly AGN driven outflows in OH$^{-}$ and CO 1-0 \\citep{ferug10,chung11,sturm11} in local ULIRGs and QSOs that could exhaust the fuel supply for star formation within $\\sim$10$^{7}$ years, while X-ray observations have found mildly relativistic outflows in radio-quiet and radio-loud AGN, with an origin close to the SMBH \\citep{chartas03,braito07,tombesi10}. Further observations have provided indirect evidence for `quasar mode' feedback in some galaxies \\citep{dunn10,alex10,muller11,rupke11}, and for radio-mode feedback in galaxies and clusters \\citep{best06,macnam07,mit09,nesv10,werner10,ehlert11,shab11,ma11,wern11}. Controversies do, however, remain. For example, \\citet{jahnke11} suggest that there is no need for AGN feedback to explain the Bulge-SMBH mass relation, \\citet{ammons11} suggest that only a small fraction of intermediate luminosity AGN at $z\\sim1$ can be undergoing galaxy-wide AGN feedback, and  \\citet{lutz10} find that there is no simple inverse relation between star formation rate and AGN luminosity in X-ray selected AGN. It is also worth mentioning that the effectiveness of AGN feedback may depend on the mode of star formation, since spatially compact, `bulge' star formation is likely easier to turn off than spatially extended, `disk' star formation.  Our group has been examining the role of AGN feedback by looking at systems in which such feedback may be ongoing. To do so, we have been studying the `FeLoBAL' class of QSOs \\citep{hazard87}, which comprise part of the Broad Absorption Line (BAL) QSO population\\footnote{BAL QSOs show broad, deep absorption troughs in their rest-frame UV spectra. They come in three subtypes. High Ionization BAL QSOs (HiBALs) show absorption in CIV $\\lambda$1549\\AA, NV $\\lambda$1240\\AA, SiIV $\\lambda$1394\\AA\\ \\& Ly$\\alpha$. Low Ionization BAL QSOs (LoBALs) additionally show absorption in MgII $\\lambda$2799\\AA\\ and other lower ionization species. Finally, the FeLoBAL QSOs, in addition to showing all the absorption lines seen in LoBALs, also show weak absorption from any {\\it excited} term of FeII (e.g. FeII* $\\lambda$2400,2600\\AA), \\& also FeII $\\lambda$2750\\AA, FeIII $\\lambda$1895,1914,1926\\AA. See e.g. \\citealt{lyn67,weymann91,green96,becker97,sch99,hall02,lacy02,trump06,gallagher07,case08,gibson09,leigh11}.}. We selected this population for three reasons. First, the UV absorption troughs are unambiguous signatures of radiatively driven outflows powered by an AGN. They have velocities of up to 0.2$c$, widths of at least a few thousand km s$^{-1}$), and are usually very deep or black, implying high column densities moving at high velocities, which plausibly implies high mass-loss rates \\citep{arav94,proga00,chartas03,crenshaw03}. Therefore, they cannot be driven by even the most extreme starbursts. Furthermore, in at least some FeLoBAL QSOs the outflows may extend up to several Kpc into the host galaxy \\citep{arav08,moe09,bautista10,dunn10}. Second, they are invariably reddened objects with high IR luminosities \\citep{farrah07}, and sometimes harbour intense starbursts \\citep{farrah10}, suggesting they may be a transition phase in the lifetime of an AGN (FeLoBAL features are also occasionally seen in ULIRGs, see e.g. \\citealt{farrah05}). Third, recent results have shown that FeLoBAL QSOs are much more common at $z\\gtrsim0.5$ than was originally thought, by up to a factor of ten \\citep{dai08,f2ms,allen11}.  In this paper, we compare the strengths of the AGN-driven outflows to the luminosities of obscured star formation in a sample of 31 FeLoBAL QSOs, selected purely on the basis of their rest-frame UV spectral properties. We combine data from the Spitzer space telescope \\citep{werner04} with data from the Sloan Digital Sky Survey (SDSS, \\citealt{york00}), the Two Micron All-Sky Survey (2MASS, \\citealt{skrutskie06}), the UKIRT Infrared Deep Sky Survey (UKIDSS, \\citealt{law07}), and from the Wide-Field Infrared Survey Explorer (WISE, \\citealt{wright10,jarr11}) in order to measure the luminosities of both obscured AGN activity and star formation. We measure the strengths of the AGN-driven outflows from the UV absorption troughs in the SDSS spectra. We assume a spatially flat cosmology, with H$_{0}$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega = 1$, and $\\Omega_{\\Lambda} = 0.7$. We use the term \"IR luminosity\" to refer to the luminosity integrated over 1-1000$\\mu$m in the rest-frame. We quote luminosities in units of bolometric Solar luminosities, where L$_{\\odot} = 3.826\\times10^{26}$ Watts. ",
        "conclusions": "We have presented a sample of 31 FeLoBAL QSOs from the SDSS at $0.8<z<1.8$. These QSOs have broad, deep absorption troughs in their rest-frame UV spectra that are unambiguous signatures of radiatively driven outflows powered by an AGN. Previous work has suggested that FeLoBAL QSOs are IR-luminous, and sometimes harbor star formation rates of up to a few thousand solar masses per year. Furthermore, there is evidence that the AGN-driven outflows can extend up to several kiloparsecs into the host galaxies. FeLoBAL QSOs are thus ideal laboratories for studying the effects of radiatively-driven outflows from an AGN on obscured star formation.  We selected our sample purely on the basis of their rest-frame UV spectral properties, and assembled for them optical through far-IR photometry from the SDSS, 2MASS, UKIDSS, WISE and Spitzer. We then fit these data with radiative transfer models for the IR emission from AGN and starbursts to extract best-fit AGN and starburst IR luminosities. We then compared these luminosities to the strength of their outflows as inferred from the absorption properties of the MgII$\\lambda$2799\\AA\\ line in their SDSS spectra. Our conclusions are:  1 - FeLoBAL QSOs are luminous in the IR. All of our sample have total IR luminosities in excess of $10^{12}$L$_{\\odot}$. Nearly one-third of the sample exceed $10^{13}$L$_{\\odot}$. The dominant power source behind the IR emission is, in most cases, AGN activity. A pure AGN is either the most likely power source, or consistent within the 90\\% confidence interval, for eleven of the sample. A starburst component is required for the remaining objects, but in only twelve of these objects is the starburst more luminous than $10^{12}$L$_{\\odot}$, and in only three objects is the starburst more luminous than the AGN. The mean AGN contribution to the total IR luminosity is $\\sim76\\%$. The spread in AGN contribution to the total IR luminosity is however wide, spanning 0.2 to 1.0. Overall, the IR properties of FeLoBAL QSOs appear consistent with those of the general red/dusty QSO population. We do not however find convincing evidence that FeLoBAL QSOs are more likely to be a ULIRG to QSO transition phase than the general red QSO population.  2 - We find no convincing relationship between the strength of the outflows and the total IR luminosity, or between the strength of the outflows and either the starburst or AGN luminosities. Conversely, we find a clear relationship between the strength of the outflows and the contribution from star formation to the total IR luminosity. If we divide our sample in two at an outflow strength boundary of AS$_{2}=3500$km s$^{-1}$ and construct probability distribution functions for the starburst contribution for the two subsamples, then the low absorption strength subsample shows a higher chance of a higher starburst contribution than the sample as a whole, and a much higher chance of a higher starburst contribution than the high absorption strength subsample (Figure \\ref{figmergedpdf_agnfrac}). We quantify this by extracting the probabilities of obtaining starburst contributions to the total IR luminosity in excess of 25\\% (Table \\ref{probabilities2}, rows 1-3). We find, at $>5\\sigma$ significance, a higher chance of obtaining a starburst contribution in excess of $25\\%$ in the low absorption strength sample compared to the high absorption strength sample.  3 - This anticorrelation between outflow strength and the contribution from star formation to the total IR luminosity is straightforwardly interpreted as the outflow from the AGN curtailing star formation in the host galaxy. There are however several other ways that it could arise. The most obvious alternative is that stronger outflows reflect an increase in the IR emission from the AGN, but have no effect on the starburst; if this is the case then the starburst contribution would decline with stronger outflows, but without there being a direct relationship behind the decline. To test this possibility we divided our sample into two at an AGN IR luminosity boundary of L$_{AGN}<10^{12.5}$L$_{\\odot}$ and constructed probability distribution functions for the starburst contribution for these two subsamples (Figure \\ref{figmergedpdf_agnfracbylagn}). If it is the case that outflow strength is a proxy for AGN luminosity, then we should see a bigger difference between these two PDFs than that seen between the PDFs in Figure \\ref{figmergedpdf_agnfrac}. Instead, the PDFs show a smaller difference. Furthermore, we find that (a) we are (marginally) more successful in finding systems with a large starburst contribution to the total IR emission by selecting on weak outflows than we are by selecting on low AGN luminosity, and (b) the probabilities of a starburst contribution in excess of 25\\% are statistically indistinguishable between the high and low AGN luminosity subsamples, but are different at $>5\\sigma$ between the high and low absorption strength subsamples (Table \\ref{probabilities2}). We considered several further alternative possibilities (\\S\\ref{discinterp}) but did not find any of them convincing.  4 - We therefore conclude that strong, radiatively driven outflows in FeLoBAL QSOs can have a dramatic, negative effect on obscured star formation in their host galaxies. This is the most direct evidence yet obtained for 'quasar mode' feedback in the QSO population at high redshifts. We find that outflows with an absorption strength in MgII$\\lambda$2799\\AA\\ of greater than $AS_{2}=3500$km s$^{-1}$ (Table \\ref{absmeasures}) act to curtail luminous star formation in their host galaxies. We also find that this effect is at least largely relative; the starburst luminosity is reduced to less than about 25\\% of the total luminosity. Finally, we propose that the magnitude of this effect is not deducible from the IR luminosity of the AGN."
    },
    "1112/1112.4448_arXiv.txt": {
        "abstract": "{High energy electrons accelerated during solar flare are abundant in the solar corona and in the interplanetary space. Commonly, the number and the energy of non-thermal electrons at the Sun is estimated using hard X-ray (HXR) spectral observations (e.g. RHESSI) and a single-particle collisional approximation. } {To investigate the role of the spectrally evolving Langmuir turbulence on the population of energetic electrons in the solar corona.} {We numerically simulate the relaxation of a power-law non-thermal electron population in a collisional inhomogeneous plasma including wave-particle, and wave-wave interactions.} {The numerical simulations show that the long-time evolution of electron population above 20 keV deviates substantially from the collisional approximation when wave-particle interactions in non-uniform plasma are taken into account. The evolution of Langmuir wave spectrum towards smaller wavenumbers, due to large-scale density fluctuations and wave-wave interactions, leads to an effective acceleration of electrons. Furthermore, the time-integrated spectrum of non-thermal electrons, which is normally observed with HXR above 20 keV, is noticeably increased due to acceleration of non-thermal electrons by Langmuir waves.} {The results show that the observed HXR spectrum, when interpreted in terms of collisional relaxation, can lead to an overestimated number and energy of energetic electrons accelerated in the corona.} ",
        "introduction": "Solar flares provide a number of theoretical challenges for various aspects of electron dynamics, including their acceleration and transport. One of the important aspect of solar flares is the high acceleration efficiency of electrons, seen via their Hard X-ray emission (HXR) in the dense atmosphere footpoints \\citep[e.g.][]{1990ApJS...73..343B,1999Ap&SS.264..129S,2002SSRv..101....1A,2003AdSpR..32.1001L,2006SSRv..124..233L,2010A&A...519A.114B,2011A&A...526A...3B,2011A&A...533L...2B,2011SoPh..270..493C,2011ApJ...739...96K}. It was realized early on that the solar corona can store large amounts of magnetic energy \\citep{1958IAUS....6..123S,2002A&ARv..10..313P}, that can be released to increase the kinetic energy of the surrounding plasma.  While the exact location and properties of the electron acceleration site are still under debate, spatially resolved HXR observations, notably with RHESSI, point to the coronal origin of electron acceleration and subsequent transport of energetic electrons towards the denser lower corona \\citep[e.g.][as the recent reviews]{2011SSRv..159..107H,2011SSRv..159..301K}. HXR emission from the dense regions is then used as the model dependent diagnostic of electron acceleration in the corona. The low dynamic range of current hard X-ray instruments is insufficient to observe the evolution of the energetic electron distribution {\\it enroute} between the coronal and the footpoint parts. The exception to this rule is probably the radio observations of emissions from energetic electrons \\citep[e.g.][]{2011ApJ...731L..19F}.  The propagation of energetic electrons between the coronal and the footpoint sources is often treated as a simple collisional transport. In this approximation the energy of an electron $E=m{\\mr v}^2/2$ decreases with time due to binary Coulomb collisions, so that for non-thermal particles ${\\mr v}(t) \\gg{\\mr v}_{Te}$, one can write \\begin{equation}\\label{dedx_TT} \\frac{d{\\mr v}}{dt}=-\\frac{\\Gamma} {{\\mr v}^{2}}, \\end{equation} where $\\Gamma=4\\pi e^4 n\\ln \\Lambda/m^2$, $n$ is the number density of the surrounding plasma, $e$ is the electron charge, $m$ is electron mass, ${\\mr v}_{Te}$ is the electron thermal velocity, and $\\ln \\Lambda \\simeq20$ is the Coulomb logarithm. This treatment obviously ignores all collective interactions in plasma, but gives a simple relation between the accelerated and the resulting mean electron spectra responsible for the observed HXR spectra \\citep[e.g.][]{2003ApJ...595L.115B}\\footnote{Normally, only spatial dependency $x={\\mr v}t$ instead of temporal $t$ is considered.}.  Realizing the importance of various streaming instabilities and the subsequent wave-particle interactions, a number of  models have included the generation of Langmuir waves. Indeed, Langmuir wave turbulence can be effectively generated when the bump on tail distribution is formed \\citep{1973plas.book.....K,1995lnlp.book.....T}. Firstly, the initially stable $\\partial f/\\partial {\\mr v} < 0$ electron distribution function $f({\\mr v})$ can be turned unstable i.e. $\\partial f/\\partial {\\mr v} >0$, when the fast electrons overtake the slower ones. In the case of collisional relaxation of the electron population, since slower particles lose their energy faster, $\\tau _{col}(E)\\simeq m{\\mr v}^3/(2\\Gamma)$, a ``gap'' distribution  can appear from an initially power-law distribution on collisional time-scales. This ``gap'' distribution of electrons creates a positive slope $\\partial f/\\partial{\\mr v} >0$, \\citep{1972AZh....49..334S,1984ApJ...279..882E,1987ApJ...321..721H}, which could lead to growth of Langmuir waves. The beam-plasma instability acts to decrease the gradient of the electron distribution in velocity space, while Coulomb collisions tend to restore it.  The net result is a plateau-like distribution,  with the plateau height slowly decreasing as energy is lost to the background plasma via collisions. It has been shown \\citep[e.g.][]{1987ApJ...321..721H,1987A&A...175..255M,2011A&A...529A.109H} that even for the time independent injection of electrons into plasma, the electrons should effectively generate Langmuir waves, however the time or space integrated electron spectrum (to compare with the HXR observations) is only weakly affected by this process. The overall evolution of the beam-plasma system is complicated by various non-linear processes, which can affect the spectrum of Langmuir turbulence \\citep[see][as a review]{1995lnlp.book.....T}. Due to these processes, the energy of Langmuir waves can be transferred from one range of phase velocities to another, hence suppressing the growth of Langmuir waves. The removal of Langmuir waves out of resonance with the electron beam can potentially lead to nonlinear stabilization of the beam-plasma instability \\citep{1975PhFl...18.1769P,1979ApJ...233..717V,1985A&A...142..219R,1991PhFlB...3.1968M}. While these earlier works were predominantly analytical estimates, the recent weakly turbulent treatments \\citep[e.g.][]{2001PhPl....8.3982Z,2002PhRvE..65f6408K,2011ApJ...727...16Z} are self-consistent and capture a complex interplay between wave-particle and wave-wave processes in collisionless plasma. For beam relaxation in a plasma, these processes can lead to the appearance of electrons with velocities above the maximum injected. At short scales, beam-plasma interaction Vlasov or PIC simulations \\citep[e.g.][]{2009A&A...506.1437K,2010SoPh..267..393T,2011PhPl...18e2107D,2011PhPl...18e2903T} are particularly useful to capture phase-space effects of beam-plasma interaction at short time scales. However, the interpretation of solar flare HXR emission requires knowledge of electron evolution at much longer time-scales, inaccessible by Vlasov or PIC approaches.  In this paper, we investigate the evolution of non-thermal electrons and Langmuir waves using weak-turbulence theory in a collisional, inhomogeneous solar coronal plasma. We show that the evolution of the Langmuir wave spectrum caused by plasma inhomogeneities and wave-wave processes leads to effective acceleration of high energy electrons. The resulting HXR spectrum is found to be strongly affected by the evolution of Langmuir turbulence, so that if the HXR spectrum is interpreted in terms of a collisional model, the accelerated electron spectrum will be overestimated. ",
        "conclusions": "The above results show the significant effect that Langmuir waves have on the corresponding electron distribution function. A constant density gradient, or random long-wavelength density fluctuations, produce enhancement of the electron distribution in the range 30-60 ${\\mr v}_{Te}$. On the other hand, the wave-wave interactions with short wavelength fluctuations $q \\sim 2k$ show a more pronounced effect at velocities ${\\mr v}\\simeq 60{\\mr v}_{Te}$. Unlike the case of a uniform plasma, the time averaged electron spectrum is noticeably different from collisional relaxation. As the time averaged spectrum can be directly compared to Hard X-ray observations, the role of Langmuir waves could be observable. The resulting HXR spectrum above $\\sim 20$ keV is higher than expected from collisional relaxation, and if interpreted as thick-target relaxation, the number of energetic electrons injected is overestimated.  \\begin{figure} \\centering \\includegraphics[width=7cm]{fig6.eps} \\caption{The effective energy loss rate of an electron $\\langle{\\partial E}/{\\partial t}\\rangle_{eff}$ using Equation (\\ref{eq:dE_dt}), at 0.5s. Collisional losses  $\\langle{\\partial E}/{\\partial t}\\rangle_{eff} =-\\Gamma\\sqrt{m^3/2}E^{-1/2}$ from Equation \\ref{dedx_TT} (black dashed line), collisional losses as in Figure \\ref{fig:collisions} (black line); Constant density gradient as in Figure \\ref{fig:Grad} (blue line); Random density fluctuations as in Figure \\ref{fig:diffusion} (red line); Non-linear interactions in inhomogeneous plasma as in Figure \\ref{fig:random_nonlin} (green line).} \\label{fig:dE_dt} \\end{figure} Given the evolution of the electron distribution, $f({\\mr v},t)$, we can calculate the effective energy change rate. Evidently the actual energy loss/gain rate is given by the kinetic equations (\\ref{eqk1}-\\ref{alphaL}), but {\\it assuming} a single particle evolution often used in the literature \\citep[e.g.][]{2009A&A...508..993B} the effective energy loss rate $\\langle dE/dt\\rangle_{eff}$ can be defined as the one satisfying a continuity equation: \\begin{equation}\\label{eq:continuity} \\frac{\\partial f(E,t)}{\\partial t}+\\frac{\\partial }{\\partial E}\\left(\\langle\\frac{d E}{d t}\\rangle_{eff}f(E,t)\\right)=0 \\end{equation} where $f(E,t)=m{\\mr v}(E)f({\\mr v}(E),t)$ is the energy distribution function. The effective energy loss/gain rate satisfying this equation is \\begin{equation}\\label{eq:dE_dt} \\langle\\frac{d E}{d t}\\rangle_{eff}=\\frac{1}{f(E,t)}\\int_{E}^{\\infty}\\frac{\\partial f(E',t)}{\\partial t}dE'. \\end{equation} If the evolution of particles is governed by collisional only, the effective loss rate is simply $\\langle{\\partial E}/{\\partial t}\\rangle_{eff} =-\\Gamma\\sqrt{m^3/(2E)}$ (Fig. \\ref{fig:dE_dt}). The evolving Langmuir wave turbulence notably changes  $\\langle{\\partial E}/{\\partial t}\\rangle_{eff}$, decreasing the effective energy loss at energies of a few keV, increasing the loss rate in the intermediate range near 10 keV, and leading to acceleration and/or reduced energy loss rate at tens of keV and above. This effective energy change rate increases the flux of electrons above 20 keV by a factor of a few to a few tens (shown in Figures \\ref{fig:Grad}-\\ref{fig:random_nonlin}).  We emphasize that the acceleration of energetic electrons is due to spectral evolution of Langmuir waves. The Langmuir wave generation in a uniform plasma only weakly changes the time integrated HXR spectrum and, hence, is not likely to be detected in current HXR measurements. The important exception is the generation of waves by spatially localized electron beams \\citep[e.g.][]{2011A&A...529A.109H}, when electrons are accelerated in bunches rather as a continuous stream as expected in some acceleration models. The acceleration of energetic electrons in non-uniform plasma is due to Langmuir waves generated by the non-thermal electrons themselves and does not increase the total energy of the beam-plasma system. On the contrary, the energetic electrons lose their energy not only via collisions but also via generation of Langmuir waves and their subsequent absorbtion via collisions. Therefore, although the low energy electrons lose their energy faster, the higher energy electrons are gaining energy and hence can generate HXR emission for a longer time. This ``self-acceleration'' of some electron population reduces the number of accelerated electrons necessary to explain the observed HXR fluxes. This process also produces a similar effect to a re-acceleration from externally applied electric fields as discussed by \\citet{2009A&A...508..993B}, but evidently does not require additional energy input.  A few important comments are due on the application of these results to the interpretation of solar flare X-ray spectrum. The overall effect of electron re-acceleration due to evolution of Langmuir turbulence is beam-density dependent. The higher the beam density, the stronger the generation of Langmuir waves, and therefore the effects considered should be more evident in HXR data. Therefore, the stronger the flare, the larger the overestimation of non-thermal electrons, when interpreted by a collisional thick-target model. Further, these simulations need to be extended to account for the spatial transport of energetic electrons in solar flare conditions, a relativistic treatment of the electrons included, and the role of whistlers \\citep[e.g.][]{2002SoPh..211..135S} and kinetic Alfven waves \\citep[e.g.][]{2010A&A...519A.114B} considered. The relaxation produces an enhanced level of Langmuir waves, which could lead to plasma microwave emission \\citep{1984ApJ...279..882E,1987ApJ...321..721H} and needs to be investigated further. However, as the spectrum of Langmuir turbulence evolves with time, the wave-wave and density gradient effects need to be included.  In summary, the evolving Langmuir turbulence spectrum has a noticeable effect on both the instantaneous and time-averaged distribution functions leading to the overestimation of electron flux (and hence energy) in non-thermal electrons $>20$~keV,  if interpreted in terms of a standard collisional thick-target model. Therefore, if Langmuir waves are effectively generated in flares and their spectrum evolves, the number of accelerated electrons could be substantially smaller than inferred from HXR thick-target interpretation."
    },
    "1112/1112.5354.txt": {
        "abstract": "Modern adaptive optics (AO) systems for large telescopes require tomographic techniques to reconstruct the phase aberrations induced by the turbulent atmosphere along a line of sight to a target which is angularly separated from the guide sources that are used to sample the atmosphere. Multi-object adaptive optics (MOAO) is one such technique. Here, we present a method which uses an artificial neural network (ANN) to reconstruct the target phase given off-axis references sources. We compare our ANN method with a standard least squares type matrix multiplication method and to the learn and apply method developed for the CANARY MOAO instrument. The ANN is trained with a large range of possible turbulent layer positions and therefore does not require any input of the optical turbulence profile. It is therefore less susceptible to changing conditions than some existing methods. We also exploit the non-linear response of the ANN to make it more robust to noisy centroid measurements than other linear techniques. ",
        "introduction": "Adaptive optics (AO) systems require guide sources to sample the turbulent atmosphere above the telescope. If a guide star is located very close to the target or we can use the target itself, then this star can be used to directly measure the phase aberrations along the line of sight to the target. However, if there is no guide star bright enough or we would like to observe multiple or extended objects in the field then we require multiple guide stars to sample the turbulent atmosphere. Another reason for using multiple guide sources is when artificial guide stars are employed. In this case each guide star only illuminates a cone within the turbulent volume above the telescope. If the light cones of these guide stars overlap with the cylinder illuminated by the target we can use tomographic techniques to reconstruct the phase aberrations along the line of sight to the target. The majority of modern AO systems (with the exception of extreme AO for extrasolar planet imaging) make use of tomographic reconstruction techniques. Three major varieties of tomographic AO currently under investigation are laser tomography AO (LTAO) \\cite{Louarn04}, multi-conjugate AO (MCAO) \\cite{Beckers89} and multi-object AO (MOAO) \\cite{Hammer02,Assemat07}.  In the case of MOAO a number of target directions are observed simultaneously and corrected independently by one deformable mirror (DM) per channel. The guide stars (natural and laser) are distributed around the field and are monitored with open loop wavefront sensors (WFS) (i.e. without a DM). The information from each guide star is then combined in such a way as to estimate the phase aberrations for each target. CANARY \\cite{Morris10,Gendron11} is the first on-sky test of tomographic MOAO and is thus a perfect test bench for both the opto-mechanical technology that needs to be developed and for the algorithms that are required for the control of the instrument.\\comment{ For this reason we plan to test the reconstructor with CANARY and the parameters for our simulations have been chosen to mirror those of CANARY.}   Optical turbulence profilers show that the atmosphere can be considered to be made up of a number of independent very thin turbulent layers. The altitude and strength of these layers can change with time and so the vertical profile of the optical turbulence will develop and evolve with time \\cite{Avila06}. Median profiles are used in simulations for performance analysis but it should be remembered that a median profile is not representative of any real profile. The tomographic reconstructor must be able to handle these changes of turbulent profiles.    Here we present a new method which uses an Artificial Neural Network (ANN) to combine the information from the WFSs and output the integrated reconstructed phase aberrations from the target to the telescope. ANNs are trained by exposing them to a large number of inputs together with the desired output. In theory this training data should cover the full range of possible scenarios. However, this is obviously not possible and so given enough training the ANN will provide a best guess to the solution. When the ANN is confronted with a superposition of a number of the independent training sets it can then predict an output by combining a number of the synaptic pathways. In this way we do not need to train the ANN with every possible turbulent profile.    We propose to train an ANN off-line with simulated data. The reconstructor is named CARMEN (Complex Atmospheric Reconstructor based on Machine lEarNing). The idea is to train the reconstructor to be able to handle any turbulent profile that it might be exposed to. We do this by carefully selecting the optimum training routines. This is a train and apply technique; once trained with the correct parameters (for example, the number and geometry of the guide stars) it will work for any optical turbulence profile. Therefore, we train CARMEN with a large number of independent turbulence profiles. We train it with the off-axis WFS slopes and the desired on-axis target Zernike coefficients. When the network is implemented and shown the off-axis WFS data it will estimate what the on-axis Zernike coefficients will be. We have chosen to output Zernike coefficients at this stage to limit the number of outputs required (i.e 27 values, assuming we predict up to $6^{\\mathrm{th}}$ order Zernikes, instead of a value of the order of the number of actuators in the DM). It would be possible to predict a higher number of degrees of freedom and we would expect the performance to increase accordingly. No {\\it a priori} knowledge of the atmosphere is required and no input from the user or re-training is required if the atmospheric turbulence profile changes during observing. This is an alternative approach to most other tomographic reconstructors.      ANNs have been applied to the field of AO in the past. Angel {\\it et al.} (1990) \\cite{Angel90}, Sandler {\\it et al.} (1991) \\cite{Sandler91} and Lloyd-Hart {\\it et al.} (1992) \\cite{Lloyd-Hart92} present successful results using neural networks for wavefront sensing in the focal plane. Montera {\\it et al.} (1996) \\cite{Montera96} experimented with an ANN to reduce WFS centroiding error and to estimate the Fried parameter $r_{0}$ and the WFS slope measurement error. They found that the ANN performed as well as but not better than a standard linear approach for estimating the WFS slopes and for the estimation of the Fried parameter, r0, however the ANN was very good at estimating the WFS slope measurement error. ANNs have also been investigated for spatial and temporal predictions of the slope measurements. Lloyd-Hart \\& McGuire (1995) \\cite{Lloyd-Hart95} use an ANN to make a temporal prediction of the WFS slopes. The AO latency is then reduced allowing for a better correction. Weddell \\& Webb (2006, 2007) \\cite{Weddell06, Weddell07} developed this idea and used off-axis WFS measurements to temporally predict the on-axis slopes. However, this was limited to low-order Zernike modes (tip/tilt) only. More recently neural networks have  been used to model open loop DMs for MOAO \\cite{Guzman10}. An accurate DM model is required for open-loop AO as the DM is not seen by the WFS.  The difference between our proposal and the work of Lloyd-Hart \\& McGuire (1995) \\cite{Lloyd-Hart95} and Weddell \\& Webb (2006) \\cite{Weddell06} is that we will train the network in simulation rather than on-sky. This allows us to select and control what the network learns and means that we can predict to a higher order. One advantage of the on-sky training is that it will inherently be trained to the concurrent turbulence profile. However, if this profile were to change then, like other reconstructors that need to be re-calcualted, it would need to be re-trained.    In section 2 we discuss two of the existing tomographic reconstructors. These are the least squares matrix vector multiplication (LS) and learn and apply (L+A) techniques. We will use these two reconstructors as a benchmark to test CARMEN. In section 3 we describe the neural network, present the optimum training scenario used and briefly state other approaches which were investigated. In section 4 we present the tomographic results of CARMEN and the two other reconstructors from simulation of three test atmospheric profiles. We also show the effects of photon and read noise in the wavefront sensor. In section 5 we discuss some of the issues related to implementing CARMEN in a real AO system. ",
        "conclusions": "We have presented and tested in simulation a novel and versatile tomographic reconstruction technique using an artificial neural network. We train the network with a number of simulated datasets designed to sample the full range of possible input signals. The data set comprises of a single turbulent layer positioned at a number of different altitudes in order to show the network as many different overlap fractions as possible. After testing several different training scenarios and network architectures we found that the simplest, with a single hidden layer, is the best. The reconstructor has been compared in simulation to a standard LS technique and to L+A.  We compare with LS and L+A only as a benchmark to show that the performance of CARMEN is on a par with other accepted reconstructor techniques. It is possible to optimise these reconstructors even more to obtain a better correction but we also believe that we can optimise CARMEN more by allowing the training process more time. Therefore, we do not want to draw conclusions about the magnitude of the correction at any one time.  We have shown that the strength of CARMEN is two fold. By using an ANN we are able to train and apply a reconstructor which can adapt to wide range of atmospheric conditions. We tested CARMEN with the test case atmospheres used for the CANARY project and also with three more extreme atmospheres consisting of only two layers but with a 50\\% split in the fractional turbulence strength and the high layer set to three different values separated by 5~km. We find that no change to CARMEN was needed even when the atmosphere changes drastically. We have also tested CARMEN with an atmosphere consisting of a large number of layers of varying fractional strength. Again we see that the ANN is able to successfully correct a large fraction of the turbulence induced phase aberrations. We varied the total integrated turbulence strength and the outer scale within a reasonable range of values (0.05~m $< r_{0} <$ 0.25~m, $D/2 < L_{0} < 25D$) and see that CARMEN follows the same trends as the two other reconstructors tested. If the atmosphere were to change dramatically during an observation the other reconstructors can be re-conditioned to deal with it. However, this does take time and is not required for the neural network approach.  The second strength of CARMEN is its ability to process shot noise corrupted centroid measurements. We have shown through Monte Carlo simulation that CARMEN is able to reconstruct the on-axis Zernike coefficients from noisy off-axis guide sources better than LS and L+A reconstructors. For example, using the CANARY median test case LS and L+A result in residual WFE of 543 and 547~nm respectively, CARMEN achieves a residual WFE of 368~nm. From analysis of the variance of the Zernike residuals we see that the majority of this improvement comes from the low order modes.  We have shown that, in simulation, ANNs can be used for tomographic reconstruction and can compete with other methods. The next step will be to test CARMEN in a more realistic situation. We intend to test CARMEN in the lab and later on-sky with CANARY."
    },
    "1112/1112.4815_arXiv.txt": {
        "abstract": "With adaptive optics imaging at Keck observatory, we have discovered a substellar companion to the F6 Pleiades star HD~23514, one of the dustiest main-sequence stars known to date ($L_{IR}/L_{*}\\sim2\\%$). This is one of the first brown dwarfs discovered as a companion to a star in the Pleiades. The 0.06~M$_\\odot$ late-M secondary has a projected separation of $\\sim$360~AU. The scarcity of substellar companions to stellar primaries in the Pleiades combined with the extremely dusty environment make this a unique system to study. ",
        "introduction": "HD~23514 (HII~1132) is an F6 star in the Pleiades. It is one of only a handful of youthful stars, ages 35 to $\\sim$100 Myr, that are surrounded by warm dust particles located in the terrestrial planet zone and which absorb a percent or so of the bolometric luminosity of their central star \\citep{Melis:2010}. We have observed HD~23514 with an adaptive optics system at the Keck observatory using the NIRC2 camera to check the binarity of this unusually dusty star.  To the best of our knowledge, to date no substellar companion to a stellar primary has been found in the Pleiades. However, while this paper was being reviewed, \\citet{Geissler:2011} announced the discovery of a substellar companion to the Pleiad HII~1348. Studies of substellar objects in the Pleiades focus on free floating brown dwarfs or companions to low mass stars \\citep{Stauffer:2007, Lodieu:2007, Bouy:2006, Martin:2003, Pinfield:2000}. With an adaptive optics survey of 144 solar type stars in the Pleiades, \\citet{Bouvier:1997} was sensitive  to objects with masses down to $\\sim$0.03~M$_\\odot$ at separations of $\\sim$1\\arcsec. However, their lowest mass companion ($\\sim$0.15~M$_\\odot$) is still above the substellar boundary (they did not observe HD~23514). Our work has revealed one such substellar companion in the Pleiades. ",
        "conclusions": "Using adaptive optics imaging at Keck, we have discovered a substellar companion to the dusty Pleiades star HD~23514. Based on comparison with evolutionary models, we estimate the companion to have a mass of about $0.06\\pm0.01$\\msun\\, and temperature $2600\\pm100$~K. To our knowledge, HII~1348 and HD~23514 are the only Pleiades stars known to have brown dwarf companions \\citep{Geissler:2011}.  HD~23514 is one of the dustiest main-sequence stars known to date with $L_{IR}/L_{*}\\sim2\\%$. The warm dust present in the system is located within a few AU of the primary. In contrast, the companion has a projected separation of $\\sim$360~AU. At that distance, the companion's gravity is expected to have little influence over the evolution of the warm dust seen around the primary. However, if a tertiary object were present in the system very close to the star, the Kozai mechanism can affect its orbit and possibly enhance collisions that generate dust."
    },
    "1112/1112.4738_arXiv.txt": {
        "abstract": "In two preceding papers \\cite{guidAsy,guidQSS} we have shown that, when reaction networks are well-removed from equilibrium, explicit asymptotic and quasi-steady-state approximations can  give algebraically-stabilized integration schemes that rival standard implicit methods in accuracy and speed for extremely stiff systems. However, we also showed that these explicit methods remain accurate but are no longer competitive in speed as the network approaches equilibrium. In this paper we analyze this failure and show that it is associated with the presence of fast equilibration timescales that neither asymptotic nor quasi-steady-state approximations are able to remove efficiently from the numerical integration. Based on this understanding, we develop a partial equilibrium method to deal effectively with the approach to equilibrium and show that explicit asymptotic methods, combined with the new partial equilibrium methods, give an integration scheme that plausibly can deal with the stiffest networks, even in the approach to equilibrium, with accuracy and speed competitive with that of implicit methods. Thus we demonstrate that such explicit methods may offer alternatives to implicit integration of even extremely stiff systems, and that these methods may permit integration of much larger networks than have been possible before in a number of fields. ",
        "introduction": "Introduction}   Problems from many fields of science and technology require the solution of large coupled reaction networks describing the flow of population between various sources and sinks.  Some important examples include  reaction networks in combustion chemistry \\cite{oran05}, geochemical cycling of elements \\cite{magick}, and thermonuclear reaction networks in astrophysics \\cite{hix05,timmes}.  The systems of differential equations that are commonly used to model these reaction networks usually exhibit stiffness, which we may think of loosely as arising from multiple timescales in the problem that differ by many orders of magnitude \\cite{oran05,gear71,lamb91,press92}.  The most straightforward way to solve such a system might appear to be an explicit numerical integration scheme (for an explicit method, advancing a timestep requires only information known from previous timesteps).  However, the textbooks routinely state \\cite{oran05,lamb91,press92} that such systems cannot be integrated efficiently using explicit forward finite-difference methods because of stability issues:  for an explicit algorithm, the maximum stable timestep in a stiff system is typically set by the fastest timescales, even if those timescales are not of central interest. The traditional solution of the stiffness problem is to replace explicit integration with implicit integration (integration methods that require the values for derivatives at timesteps not yet evaluated).  The astrophysical CNO cycle for conversion of hydrogen to helium (Fig.~\\ref{fig:cnoCycle}) is an instructive example of the stiffness issue. In the CNO cycle the fastest rates typically are $\\beta$-decays with half-lives of order 100 seconds, but tracking main-sequence hydrogen burning may require integration of the hydrogen-burning network for as long as billions of years ($\\sim 10^{16}$ seconds). With explicit forward differencing the largest stable integration timestep is set by the fastest rates and will be of order $10^2$ seconds, so $\\sim 10^{14}$ explicit integration steps could be required, even for the idealized case of constant temperature and density.  Conversely, this same numerical integration requires at most a few hundred steps using implicit methods. For this reason, it is generally thought that explicit methods are not viable for extremely stiff networks. To quote {\\em Numerical Recipes} \\cite{press92}, ``For stiff problems we {\\em must} use an implicit method if we want to avoid having tiny stepsizes.'' \\putfig {cnoCycle} {0pt} {\\figdn} {0.90} {fig1.eps} {The CNO cycle. On the left side the main branch of the cycle is illustrated with solid arrows and a side branch is illustrated with dashed arrows. On the right side, the main branch of the CNO cycle is illustrated in more detail ($p$ stands for \\isotope1H and $\\alpha$ for \\isotope 4{He}).}  Although implicit methods typically are stable for stiff systems, they require substantial additional computational overhead relative to explicit methods.  For large coupled sets of equations, this normally takes the form of iterative solutions that require the inversion of large matrices at each step. The required matrix inversions dictate that, except where simplifications based on matrix structure can be exploited (for example, sparse-matrix methods), implicit algorithms may scale as poorly as quadratically or cubically with the size of the network. Thus, implicit methods can be expensive to implement for large networks, particularly those that are coupled in real time to a broader problem such as hydrodynamical evolution.  Despite the generally negative view of explicit methods for stiff systems sketched in the preceding paragraphs, they would be attractive options for complex networks if they could take larger timesteps because of their overall simplicity and highly-favorable linear scaling with network size relative to implicit methods. To increase the integration step size for explicit methods obviously requires overcoming formidable numerical stability problems that are associated with the stiffness. In principle, this might be accomplished by using approximate analytical solutions to reduce the stiffness of the equation set to be solved, thereby improving the overall stability of the network. In the first two papers of this series \\cite{guidAsy,guidQSS}, we have shown in a variety of examples that this can be accomplished for systems that are not near equilibrium by using asymptotic and quasi-steady-state approximations to stabilize the numerical integration. However, as we shall now discuss, the nature of stiffness is different for networks that are near equilibrium. This implies that the modifications required to stabilize standard explicit methods when stiffness is encountered far from equilibrium must be altered when that same network approaches equilibrium. ",
        "conclusions": "Summary and Conclusions}  Numerical integration for extremely stiff reaction networks has typically been viewed as requiring implicit schemes, because normally explicit timesteps of efficient size would be unstable.  An alternative to implicit integration modifies the equations  using approximate solutions to reduce the stiffness, and then integrates the resulting equations explicitly. For example, asymptotic and quasi-steady-state methods found a measure of success integrating moderately stiff chemical-kinetics systems explicitly \\cite{oran05,mott00,mott99}, but have generally failed for extremely stiff systems, giving inaccurate results with non-competitive integration timesteps. In this paper, and the two preceding it \\cite{guidAsy,guidQSS}, we have presented evidence suggesting that algebraically-stabilized explicit integration can give correct results and timesteps competitive  with that of implicit methods for even the stiffest of networks.  Central to this new view of the efficacy of explicit integration for stiff systems is the understanding that in reaction networks there are at least three different sources of stiffness:  \\begin{enumerate}  \\item {\\em Negative populations}, which can occur for an initially small positive population if an explicit timestep is too large; this anomalous negative population leads to exponentially growing terms that destabilize the network.  \\item {\\em Macroscopic equilibration}, where the right sides of the differential equations  $dY = F = \\Fplus{} - \\Fminus{}$ approach a constant derived from the difference of two large numbers (the total flux in $\\Fplus{}$ and total flux out $\\Fminus{}$), and numerical errors in taking this difference destabilize the network if the timestep is too large.  \\item {\\em Microscopic equilibration}, where on the right sides of Eq.~(\\ref{equilDecomposition}) the net flux in specific forward-reverse reaction pairs $(f^+_i - f^-_i)$ tends to zero as the system approaches equilibrium; this leads to large errors if the timestep is too large because the net flux is derived from the difference of two large numbers and the timescale equilibrating the populations is short compared with the desired numerical timestep.  \\end{enumerate} The algebra required to stabilize the explicit solution depends on which of these forms of stiffness is operative. In Refs.~\\cite{guidAsy,guidQSS} we have shown that asymptotic and quasi-steady-state approximations give correct results, with timestepping highly competitive with implicit methods even for the stiffest of thermonuclear networks, as long as the system exhibits only the first two types of stiffness listed above. However, in those papers we have also shown that asymptotic and quasi-steady-state methods give correct results, but with timesteps that are much larger than for purely explicit methods but still far from competitive with implicit methods, if there is any significant microscopic equilibration in the system.  In this paper we have addressed in depth the role of microscopic equilibration in stiffness and presented a partial equilibrium scheme that can be used to augment asymptotic and quasi-steady-state methods when equilibrium becomes important in a network. We have presented strong evidence that this partial equilibrium scheme can remove the timestepping deficiencies encountered by asymptotic and quasi-steady-state methods in the approach to equilibrium, making these methods competitive with implicit methods even near equilibrium. The methods that we have presented build on earlier work \\cite{oran05,mott00,mott99}, but we find that our versions of asymptotic and quasi-steady-state methods are far more accurate, and our versions of partial equilibrium are faster by multiple orders of magnitude, than found in previous work when applied to extremely stiff astrophysical thermonuclear networks.  The present paper, in conjunction with the previous ones on asymptotic methods \\cite{guidAsy} and quasi-steady-state methods \\cite{guidQSS}, suggests that algebraically-stabilized explicit integration methods are capable of stable timesteps competitive with those of implicit methods even for extremely-stiff reaction networks. Explicit methods can compute timesteps faster than implicit methods in large networks, implying that algebraically-stabilized explicit algorithms may compete favorably with implicit integration, even in complex, extremely-stiff applications. Because explicit methods scale linearly with network size, these findings could permit the coupling of physically more realistic reaction networks to fluid dynamics simulations in a variety of fields.   \\begin{ack} We thank Austin Harris for help with some of the calculations, Elisha Feger, Tony Mezzacappa, Bronson Messer, Suzanne Parete-Koon, and Kenny Roche for useful discussions, and Suzanne Parete-Koon for a careful reading of the manuscript. Research was sponsored by the  Office of Nuclear Physics, U.S. Department of Energy. \\end{ack}   \\clearpage \\newpage  \\setcounter{section}{0} \\appendix"
    },
    "1112/1112.5813_arXiv.txt": {
        "abstract": "{Following the great success of the current Imaging Atmospheric Cherenkov Telescopes, the next generation gamma-ray telescope arrays, Cherenkov Telescope Array (CTA), is being prepared. CTA will have an order of magnitude higher sensitivity and wider energy range than current instruments. Large samples of very high energy (VHE; $>30$ GeV) sources will be obtained with CTA. Here we discuss potential AGN population studies for the CTA era concerning the overall source counts and the statistics of high redshift sources. Based on our latest blazar gamma-ray luminosity function (GLF) model, we find that CTA will detect $\\sim50$ and $\\sim160$ blazars with 1 year and 10 years of blank field sky survey, respectively. CTA is also expected to find a blazar at $z\\sim1.4$ ($\\sim20$ blazars above $z=1$) based on the blazar GLF. Furthermore, we also examine the detectability of high redshift {\\it Fermi} blazars. By extrapolating the {\\it Fermi} blazars' spectra with a power-law, we find that CTA has a capability to detect a {\\it Fermi} blazar at $z=2.49$.}  \\FullConference{AGN Physics in the CTA Era - AGN2011,\\\\ May 16-17, 2011\\\\ Toulouse, France}   \\begin{document} ",
        "introduction": "Following the great success of the current Imaging Atmospheric Cherenkov Telescopes (IACTs), the next generation gamma-ray telescope arrays, Cherenkov Telescope Array (CTA), is being prepared \\cite{CTA}. It is designed to have an order of magnitude higher sensitivity and wider energy range compared to arrays.  Current IACTs have already found $\\sim$100 very high energy (VHE) sources, including more than 40 blazars up to redshift $z=0.536$\\footnote{{http://www.mpp.mpg.de/\\%7Erwagner/sources/}\\\\ \\ \\ \\ \\ \\ \\ \\ \\ \\  {http://tevcat.uchicago.edu/}}. 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 and {\\it Fermi} are blazars \\cite{har99,abd10_catalog}. In the CTA era, it is naturally expected that the number of blazars detected in the VHE range increases dramatically and that a much deeper VHE universe will be explored. Therefore, it would be possible to carry out a population study of VHE blazars with CTA, which would provide a crucial key to understanding AGN populations, high-energy phenomena around supermassive black holes in AGNs, the cosmological evolution of AGNs, and the distribution of the Extragalactic background light (EBL).  The purpose of this paper is to consider the prospect of future blazar surveys by CTA. For this purpose, the blazar gamma-ray luminosity function (GLF) and spectral energy distribution (SED) are needed. The blazar GLF has been studied in detail in many papers (see e.g. \\cite{it09,itm10}). A new blazar GLFs \\cite{it09,itm10} has recently developed on the basis of the latest determination of the X-ray luminosity function of AGNs \\cite{ued03,has05} featuring a so-called luminosity density evolution. Another new feature of these GLFs is taking into account the blazar SED sequence. The blazar sequence is a feature seen in the mean SED of blazars that the synchrotron and inverse Compton (IC) peak photon energies decrease as the bolometric luminosity increases \\cite{fos98,kub98}.  The key parameters in GLF have been carefully determined to match the observed flux and redshift distribution of EGRET blazars by a likelihood analysis. Predictions for the {\\it Fermi} mission are in good agreement with the recent {\\it Fermi} observed data \\cite{itm10,ino11}.  We utilize the model by \\cite{itm10} to predict the expected number and distributions of physical quantities of VHE blazars in a future CTA sky   survey. We also discuss the expected high redshift blazar studies by CTA based on the first year {\\it Fermi} source catalog. Throughout this paper, we adopt the standard cosmological parameters of $(h, \\Omega_M , \\Omega_\\Lambda) = (0.7, 0.3, 0.7)$. ",
        "conclusions": "In this paper, we discussed the expected number of blazars and highest redshift accessible with future CTA observation based on a recent blazar GLF model.  For a blank field sky survey, a wide and shallow search will enable us to carry out an efficient survey. CTA will detect $\\sim50$ and $\\sim160$ blazars with 1 year and 10 years blank field sky survey, respectively. With CTA, we should be able to find a blazar at $z\\sim1.4$ ($\\sim20$ blazars above $z=1$). We also estimated the detectable {\\it Fermi} blazars. By simply extrapolating the GeV spectrum of the {\\it Fermi} blazars in the first catalog, CTA is expected to detect a blazar at $z=2.49$. This will enable us to study the cosmological evolution of VHE blazars.   Although we discussed only blazars in this paper, other classes of AGNs  are also expected as potential sources for CTA, such as core emission from radio galaxies \\cite{ino11b}, kpc jet emission \\cite{har11}, low luminous AGNs \\cite{tak11}. By observing various AGN populations and high redshift sources, CTA will help to unify AGN populations and study the cosmological evolution of the VHE universe."
    },
    "1112/1112.0431_arXiv.txt": {
        "abstract": "{ The ages of individual Red Giant Branch stars can range from 1~Gyr old to the age of the Universe, and it is believed that the abundances of most chemical elements in their photospheres remain unchanged with time (those that are not affected by the first dredge-up). This means that they trace the interstellar medium in the galaxy at the time the star formed, and hence the chemical enrichment history of the galaxy. } { Colour-Magnitude Diagram analysis has shown the Carina dwarf spheroidal to have had an unusually episodic star formation history and this is expected to be reflected in the abundances of different chemical elements. } { We use the VLT-FLAMES multi-fibre spectrograph in high-resolution mode (R$\\approx$20000) to measure the abundances of several chemical elements, including Fe, Mg, Ca and Ba, in a sample of 35 individual Red Giant Branch stars in the Carina dwarf spheroidal galaxy.  We also combine these abundances with photometry to derive age estimates for these stars. This allows us to determine which of two distinct star formation episodes the stars in our sample belong to, and thus to define the relationship between star formation and chemical enrichment during these two episodes. } { As is expected from the star formation history, Carina contains two distinct populations of Red Giant Branch stars: one old ($\\gtrsim$~10~Gyr), which we have found to be metal-poor ([Fe/H]$ < -1.5$), and $\\alpha$-rich ([Mg/Fe]$ > 0$); the other intermediate age ($\\approx 2-6$~Gyr), which we have found to have a metallicity range ($-1.8 < $[Fe/H]$ < -1.2$) with a large spread in [$\\alpha$/Fe] abundance, going from extremely low values ([Mg/Fe]$ < -$0.3) to the same mean values as the older population ($<$[Mg/Fe]$>\\sim 0.3$). } { We show that the chemical enrichment history of the Carina dwarf spheroidal was different for each star formation episode.  The earliest was short ($\\sim 2-3$~Gyr) and resulted in the rapid chemical enrichment of the whole galaxy to [Fe/H]$\\sim -1.5$ with both SNe II and SNe Ia contributions.  The subsequent episode occured after a gap of $\\sim 3-4$~Gyr, forming $\\sim$~70\\% of the stars in the Carina dSph, but it appears to have resulted in relatively little evolution in either [Fe/H] or [$\\alpha$/Fe]. } ",
        "introduction": "{ The main advantage of nearby galaxies in the Local Group is that it is possible to study their resolved stellar populations. This means that stars of particular ages can be picked out of a Colour-Magnitude Diagram (CMD) for a more detailed study of specific epochs in the evolution of a galaxy. An important and bright population of stars that trace the chemical evolution over almost the entire star formation history (SFH) of a galaxy are the Red Giant Branch (RGB) stars \\citep[see][and references therein]{Tolstoy2009}.  In the context of a $\\Lambda$-CDM universe dwarf galaxies play a key role in galaxy formation \\citep[e.g.][]{Kauf1993}. The most simple scenario assumes that the small and unevolved dwarf spheroidal (dSph) galaxies that we see around the Milky Way should be related to the type of objects that originally merged to form the Milky Way. This suggests that the stellar populations at the ages when mergers occurred should be similar in small dSph galaxies and the Milky Way.  The comparison of stellar abundances in dSph galaxies around the Milky Way and the stars in the Milky Way itself have shown that the merging of these nearby dwarf galaxies to create the Milky Way could only have happened at very early times, within the first Gyr of star formation occuring in either system \\citep[e.g.][]{Una1996,She2001,Tolstoy2003,Venn2004,Tolstoy2011}. After this time the differences in $\\alpha$-element enrichment (the position of the ``knee'') become too different between large and small systems.  The Carina dSph galaxy has long been known to have a most unusual and episodic SFH. The presence of RR~Lyrae variable stars first indicated an ancient stellar population \\citep{Saha1986}, and anomalous Cepheid variable stars \\citep{Mat1998} showed that a young ($<$~1~Gyr) population is also present. Anomalous Cepheids are believed to be metal-poor classical Cepheids \\citep{Bono1997}, but they could also be the result of mass transfer and possibly coalescence in a low-mass binary systems \\citep{Zinn1976}. The Carina dSph is still the only galaxy where such distinct main sequence turnoffs can clearly be seen in deep CMDs \\citep[first shown by][]{SH1996}. This is the unequivocal signature of periods of active star formation separated by similarly long periods of no star formation at all.  There have been numerous analyses of the CMDs of the Carina dSph, including the distinct Main Sequence turn-offs, in a variety of filters, for different fields of view to determine the full SFH \\citep[e.g.,][]{Migh1997,HK1998,Dolphin2002,Mon2003,Riz2003,Bono2010}.  There is still quite some discrepancy between the SFHs obtained; specifically precisely when the different episodes of star formation occurred.  These differences usually relate to the field of view of the observations modelled and thus how well different features in the CMD are populated \\citep[see][]{Cig2010} and also the different methods used to determine the SFH.  However, all studies agree that there were at least two major episodes of star formation in the Carina dSph which were separated by a long quiescent era. They also agree that the dominant stellar population ($\\sim 70\\%$ of the stars) comes from an intermediate age episode of star formation.  Despite its complex SFH the Carina dSph has an extremely narrow RGB. This is a result of the age-metallicity degeneracy, where the metal-rich, young stars have almost the same colour as the metal-poor older stars. The thinness of the RGB is also a consequence of the dominant intermediate age population which formed over a relatively short time (a few Gyr) with only a small spread in [Fe/H]. The thin RGB has made reliable photometric determinations of the metallicity and its evolution challenging \\citep[e.g.,][]{Riz2003, Bono2010}.  The first spectroscopic studies of the metallicities of individual stars in the Carina dSph came from the empirical calibration of the strong Ca~II triplet (CaT) lines \\citep{Arm1991,DaCos1994,SH1999}.  These studies typically reported a mean metallicity between [Fe/H]=$-1.5$ and $-2.0$ dex, with a small dispersion.  The largest sample was $\\sim$~50 RGB stars. The advent of multi-fibre instruments like FLAMES at the VLT led to a dramatic increase in the sizes of these CaT samples \\citep[e.g., 437 stars from][]{Koch2006}, leading to a much better defined and statistically robust metallicity distribution function \\citep{Helmi2006,Stark2010}. The spread in iron \\citep[e.g., a FWHM of 0.92 dex derived from CaT by][]{Koch2006} is much larger than the one suggested by photometric evidence.  The first detailed, high resolution, abundance study of a sample of five individual RGB stars in the Carina dSph used VLT/UVES \\citep{She2003}.  These results were analyzed in terms of galactic evolution with the help of both photometric ages and detailed spectroscopic abundances \\citep{Tolstoy2003}.  Recently a study of iron and $\\alpha$-element abundances from UVES spectra of a further 10 RGB stars has been made \\citep{Koch2008} and a study of a large range of elements in 9 stars with UVES and MIKE on Magellan spectra \\citep{Venn2011}. As the total number of RGB stars studied in the Carina dSph has increased so has the apparent star to star scatter in the abundance ratios of most elements, although the dispersion in elemental ratios was already apparent from the start \\citep{She2003,Koch2008}.  There have been several attempts to self-consistently model the SFH and the chemical evolution of the Carina dSph \\citep[e.g.,][]{Lan2003,Lan2004,Lan2006,Revaz2009}. This is challenging mostly because the Carina dSph is so unusual compared to the other dSph in the Local Group. It is hard to find a unique explanation for the clear gaps seen in the SFH \\citep[e.g.][]{Pas2011}.  It is hoped that detailed chemical abundances of a large sample of individual RGB stars, with age estimates, in the Carina dSph will better define the effects of star formation on the chemical evolution. These large samples will also better quantify the scatter seen in the abundances of a variety of chemical elements. }  \\par This paper is organized as follows: in Sect.~\\ref{sample}, we describe the sample selection, observations and data reduction. In Sect.~\\ref{atmparam_section}, we explain how we derived the atmospheric parameters (T$_{\\rm eff}$, $log~g$, v$_{t}$ and [Fe/H]) and the individual abundances. The abundances are discussed in detail in Sect.~\\ref{results} and the ages in Sect.~\\ref{ages_sect}. They are interpreted in terms of the SFH of the Carina dSph in Sect.~\\ref{disc}. ",
        "conclusions": "We determined detailed high resolution abundances for a variety of elements, including Mg, Ca and Ba, for most of our 35 RGB stars in the Carina dSph. This is a significant increase in the number of RGB stars analysed from previous studies. We confirm that the abundances of individual RGB stars, at a given [Fe/H] in the Carina dSph are typically more scattered than in the Milky Way or other dSph, as has already been noted by \\citet{Koch2008}. This scatter exists at all [Fe/H] and is clearly representative of the global properties of the Carina dSph.  We also determined ages for our sample of RGB stars and for $\\sim$20 additional RGB stars in the Carina dSph with detailed abundances from the literature. This allows us for the first time to distinguish significant numbers of stars from different episodes of star formation in the Carina dSph. The ages we determined are broadly consistent with the two episodes of star formation expected from CMD analysis. Our ages are not accurate enough to do more than place our stars in one star formation episode or another.  As in other dSphs, Carina contains old metal-poor, alpha-rich stars, and among these stars there appears to be a general trend of decreasing [Mg/Fe] with increasing [Fe/H], which is presumably due to the onset of SNe Ia enrichment (although a stochastic star formation could also be invoked). This population of stars is associated to the oldest episode of star formation in the Carina dSph. We have also identified stars associated to a second episode of star formation. Their enrichment history is not so easy to interpret. The most metal poor stars in this later, longer lasting and more intense episode of star formation are somewhat more metal poor than the most metal rich stars in the oldest sample. This suggests that there has not been a simple chemical evolution path following from one episode to the next in a fixed reservoir of gas.  An important step would be to carry out a combined model of the detailed SFH and the abundances, as has been done for Sculptor (de Boer et al. 2011, in prep), and of course to better populate the entire [Fe/H] range of Carina with accurate and detailed stellar abundances.  Despite a large increase in the number of RGB stars with accurate abundances in the Carina dSph, we still have not been able to fully disentangle the chemical evolution history."
    },
    "1112/1112.2744_arXiv.txt": {
        "abstract": "{We report on the results of the first \\xmm\\ systematic ``excess variance\"  study of all the radio quiet, X-ray un-obscured AGN. The entire sample consist of 161 sources observed by \\xmm\\ for more than 10 ks in pointed observations, which is the largest sample used so far to study AGN X-ray variability on time scales less than a day. } {Recently it has been suggested that the same engine might be at work in the core of every Black Hole (BH) accreting object. In this hypothesis, the same variability should be observed in all AGN, once rescaled by the M$_{\\rm BH}$ (M$_{\\rm BH}$) and accretion rate ($\\dot{m}$).} {We systematically compute the excess variance for all AGN, on different time-scales (10, 20, 40 and 80 ks) and in different energy bands (0.3-0.7, 0.7-2 and 2-10 keV).} {We observe a highly significant and tight ($\\sim0.7$ dex) correlation between $\\sigma^2_{\\rm rms}$ and M$_{\\rm BH}$. The subsample of reverberation mapped AGN shows an even smaller scatter (only a factor of 2-3) comparable to the one induced by the M$_{\\rm BH}$ uncertainties. This implies that X-ray variability can be used as an accurate tool to measure M$_{\\rm BH}$ and this method is more accurate than the ones based on single epoch optical spectra. This allows us to measure M$_{\\rm BH}$ for 65 AGN and estimate lower limits for the remaining 96 AGN. On the other hand, the $\\sigma^2_{\\rm rms}$ vs. accretion rate dependence is weaker than expected based on the PSD break frequency scaling. This strongly suggests that both the PSD high frequency break and the normalisation depend on accretion rate in such a way that they almost completely counterbalance each other (PSD$_{\\rm amp} \\propto\\dot{m}^{-0.8}$). A highly significant correlation between $\\sigma^2_{\\rm rms}$ and 2-10 keV spectral index is observed. The highly significant correlations between $\\sigma^2_{\\rm rms}$ and both the L$_{\\rm Bol}$ and the FWHM$_{\\rm H\\beta}$ are consistent with being just by-products of the $\\sigma^2_{\\rm rms}$ vs. M$_{\\rm BH}$ relation. The soft and medium $\\sigma^2_{\\rm rms}$ is very well correlated with the hard $\\sigma^2_{\\rm rms}$, with no deviations from a linear one to one correlation. This suggests that the additional soft components (i.e. soft excess, warm absorber) add a minor contribution to the total variability. Once the variability is rescaled for M$_{\\rm BH}$ and $\\dot{m}$, no significant difference between narrow-line and broad-line Seyfert 1 is observed.} {The results are in agreement with a picture where, to first approximation, all local AGN have the same variability properties once rescaled for M$_{\\rm BH}$ and $\\dot{m}$.} ",
        "introduction": "Rapid variability is one of the major observational properties of accretion onto Black Holes (BH). Since the early observations, it was realised that the fast variability observed in AGN implied that the primary source had to be both very compact and able to emit with a large luminosity, thus requiring an extremely efficient engine, providing the first compelling argument for the presence of BH in the heart of AGN (e.g. Lynden-Bell 1969; Rees 1984). Early studies of the AGN X--ray variability properties suggested that more luminous sources are ``less\" variable (Barr \\& Mushotzky 1986). Later on, the {\\it EXOSAT} long looks allowed the first study of the observed variations with the use of Fourier analysis methods, like the power spectral density function (PSD), for the brightest AGN.  Lawrence \\& Papadakis (1993) and Green, McHardy \\& Lehto (1993) found that the majority of the {\\it EXOSAT} AGN light curves exhibited ``red noise\" variations, i.e.  their PSD were well fitted by power-laws with slopes steeper than 1. They also found that the PSD amplitude at a fixed frequency  was anti-correlated with the source luminosity,  confirming previous results.  The availability of well sampled, very long (up to $\\sim 10-15$ years in some cases) {\\it RXTE} light curves, and their use together with high signal to noise, continuous {\\it XMM-Newton} light curves has revolutionised our view of AGN X--ray variability in the last 15 years. For a dozen objects, the combined use of these light curves allowed the accurate determination of their PSD over a broad range of time scales. The past suggestions that the steep PSD at high frequencies had to flatten below a certain time scale (see e.g. Papadakis \\& McHardy, 1995) were soon confirmed: all the best quality PSD are best fitted by a steep power law of slope $\\sim -2$ which ``breaks\" to a flatter ($\\sim -1$) slope at frequencies below the so-called  ``break frequency\" $\\nu_{\\rm b}$ (e.g. Edelson \\& Nandra, 1999; Uttley et al. 2002; Markowitz et al. 2004; Kelly et al. 2010). In some AGN the $-1$ slope part could be measured for more than 3 decades in frequency, with no lower frequency breaks being detected, thus indicating a behaviour similar to Galactic X--ray black hole binary candidates (BHB) in their so-called ``soft state\" (e.g. McHardy et al, 2004). In at least one case, namely Ark 564, two frequency breaks have been detected (Papadakis et al 2002). In fact, the PSD of this object may be better fitted by ``Lorentzian\" functions (as opposed to \"broken\" power-laws; McHardy et al. 2007), similar to the functions that are widely used in the parametrisation of the BHB PSD.  The detailed PSD analysis of the high quality {\\it RXTE} and {\\it XMM} light curves revealed a remarkable similarity between the AGN and BHB power-spectra, although the time scales sampled in the two classes of objects were vastly different. This result implied a deep link between all BH accreting objects, reinforcing initial suggestions that both systems host the same engine, and that the same emission and variability mechanisms operate in them. The difference in time-scales could be explained by the difference in the mass of the compact object in the center of these objects. The results from these first studies made it clear that the observed AGN X--ray variations are most probably determined by BH mass (M$_{\\rm BH}$) rather than luminosity (Hayashida et al. 1998; Czerny et al. 2001; Uttley et al. 2002), and even suggested that  X-ray variability  might be used as a tool to estimate M$_{\\rm BH}$ in these objects.  McHardy et al. (2006) were the first to demonstrate clearly that AGN are scaled-up versions of BHB. Using PSD results for almost a dozen objects, spanning a range of $\\sim 8$ orders of magnitude in M$_{\\rm BH}$ and $\\sim 3$ orders of magnitude in accretion rate (hereinafter we will use the term accretion rate as a sinonimus of Eddington ratio),  they demonstrated that the PSD ``break time scales\" increase proportionally with M$_{\\rm BH}$, and decrease with increasing accretion rate.  Koerding et al. (2007) showed that the three physical parameters, namely M$_{\\rm BH}$, accretion rate and break frequency, are intimately related, both in AGN and BHB in the soft state, determining a ``variability plane\" of accreting BHs. Interestingly, if a constant offset is introduced, this plane can also be extended to the ``hard state\" objects as well. Casella et al. (2008) even used this ``variability plane\" to estimate the M$_{\\rm BH}$ of Ultra Luminous X-ray sources (ULXs).  Although the PSD analysis is a powerful tool to characterise the variability properties of AGN and BHB, it nevertheless requires long, uninterrupted observations from especially tailored monitoring campaigns to fully exploit its potential. Such observations are available at the moment for no more than two dozen AGN. However, the last few years, shorter, high signal-to-noise X--ray observations for tens of objects have populated the archive of X--ray observatories such as {\\it Chandra}, and, perhaps even more importantly {\\it XMM-Newton} (due to its high sensitivity, and broad energy band pass). A convenient analysis tool for such short data sets is the so-called ``normalized excess variance\", $\\sigma^2_{\\rm rms}$ (Nandra et al. 1997). Although it does not offer the same wealth of information like the PSD analysis, it can certainly be used to confirm the PSD results using larger data samples. In fact, the availability of larger samples can also allow the discovery of new correlations between the X--ray variability amplitude and other AGN physical parameters.  The first AGN excess variance surveys were performed using  \\asca\\ light curves. The results confirmed the variability vs. luminosity anti-correlation (Nandra et al. 1997; George et al. 2000) but also indicated, for the first time, that the variability amplitude (i.e. the excess variance) correlates with the X--ray spectral index, and anti-correlates with the FWHM of the H$\\beta$ line (Turner et al. 1999). Similar studies also indicated, for the first time, that the so called ``Narrow Line Seyfert 1\" galaxies were systematically ``more variable\" than the classical AGN of equal luminosity (Leighly 1999). Later on, the variability vs. luminosity anti-correlation was also confirmed on long time scales (Markowitz \\& Edelson 2001), and it was soon suggested that the excess variance vs. luminosity relation might be just a by-product of a more ``fundamental\" relation, that of the ``variability vs. M$_{\\rm BH}$ relation\" (Lu \\& Yu, 2001; Bian \\& Zhao, 2003; Papadakis, 2004). These suggestions were put forward even before the McHardy et al (2006) and Koerding et al (2007) M$_{\\rm BH}$ (and accretion rate) scaling relations were published. These relations, together with the standard accretion theory of $\\alpha$-disks (Shakura \\& Sunyaev 1973), which predicts that all the disc characteristic time scales should depend linearly on M$_{\\rm BH}$ (e.g. Treves, Maraschi \\& Abramowicz 1988), reinforce the hypothesis of the ``variability vs. M$_{\\rm BH}$\" relation being the fundamental relation which determines all the other observed relations, although alternative suggestions are still being considered (Liu \\& Zhang 2008).  The largest collection so far of excess variance measurements in local AGN has been presented by O'Neill et al. (2005). These authors used all the available 40~ks long \\asca\\ light curves to calculate the excess variance of 46 AGN. They found a strong anti-correlation of variability with M$_{\\rm BH}$ and with luminosity. However, the latter correlation disappeared once the $\\sigma^2_{\\rm rms}$ vs. M$_{\\rm BH}$ correlation was taken into account. O'Neill et al. (2005) found a weaker correlation ($\\sim96$ \\%) between $\\sigma^2_{\\rm rms}$ and 2-10 keV spectral index, $\\Gamma$, than observed previously (Turner et al. 1999). Miniutti et al. (2009) confirmed the variability vs. M$_{\\rm BH}$ trend extending the variability estimation to smaller M$_{\\rm BH}$ AGN. Finally, Zhou et al. (2010) tried to calibrate accurately, and estimate the intrinsic scatter of the $\\sigma^2_{\\rm rms}$ vs. M$_{\\rm BH}$ relation, using high quality {\\it XMM-Newton} light curves of AGN with M$_{\\rm BH}$ measured through the ``reverberation mapping\" technique. They found that the intrinsic scatter of this relation is even smaller than the one implied by the M$_{\\rm BH}$ uncertainties. Moreover, no dependence of the variability on either accretion rate or spectral index was observed.  In this paper we present the results from the excess variance measurements of a sample of 161 AGN. This is about three times larger than the \\asca\\ sample of O'Neill et al. (2005) and, at present, the largest sample in which the short time variability (less than a day) has been systematically investigated. In addition, we estimated the excess variance in various energy bands, and on different time scales. The first aim of our work is to investigate whether the results obtained from the detailed PSD analysis of good quality light curves of a few AGN are applicable to the ``majority\" of the X--ray studied AGN as well. Note that the McHardy et al (2006) and Koerding et al (2007) scaling relations were based on the accurate PSD results for a relatively small number of objects (for example there were just 10 AGN in the McHardy et al. 2006 sample). The excess variance, being the integral of the PSD over the frequency window sampled in the light curve that is used to compute it, can be a powerful tool that can be used to investigate the applicability of these relations to a much larger sample of AGN. This can be done with the excess variance vs. M$_{\\rm BH}$ and excess variance vs. accretion rate relations, which we study in detail in this work, putting particular emphasis on the comparison between the PSD model predictions and the observed relations.  The second major aim of our work is to investigate how accurately we can ``weigh\" the central M$_{\\rm BH}$ in AGN with $\\sigma^2_{\\rm rms}$, and we provide accurate recipes to measure M$_{\\rm BH}$ from X-ray variability using light curves of various lengths. Moreover, we present the results from the study of the variability amplitude with the X--ray spectral index and FWHM H$\\beta$. Our results indicate that the same variability mechanism operates in all AGN, and that significant differences in the variability amplitude of NLS1 and ``typical\" broad line Seyferts (BLS1) can be fully understood once the observed variability is properly ``normalized\" to the M$_{\\rm BH}$ and accretion rate of these objects. ",
        "conclusions": "We studied the variability properties, computing the excess variance, of all the radio quiet, X-ray unobscured, AGN observed by \\xmm\\ for more than 10 ks in pointed observations. The \\xmm\\ high sensitivity and long orbit allows us to measure the variability amplitude of 161 AGN and probe the variability on shorter timescales down to 10 ks. We sample timescales going from 80 to 10 ks, with the latter being significantly smaller than the time scales considered in all previous works (like e.g. O'Neill et al. 2005; Zhou et al. 2010). CAIXAvar is the largest sample in which short time-scale variability (on time-scales smaller than about a day) has been studied and is composed primarily by local AGN. Thus the relations and variability properties investigated in this work are representative of a population of AGN characteristic of the local universe.  \\begin{itemize}  \\item{} We systematically explore the energy dependence of the excess variance, comparing the variability of the hard (2-10 keV) with the soft (0.3-0.7 keV) and medium (0.7-2 keV) bands. Surprisingly, although different spectral components, such as the soft excess and warm absorbers ``dominate\" the emission at softer energies, the soft band variability amplitude  is very well correlated with the hard band variability. In particular, no deviations from the linear one-to-one correlations are observed, suggesting that these components add a minor contribution to the total variability.  \\item{} We observe a highly significant and tight (0.4 dex) correlation between $\\sigma^2_{\\rm rms}$ and M$_{\\rm BH}$ for the reverberation sample. A larger scatter (0.7 dex) is observed for CAIXAvar. This is due to the larger uncertainties on the M$_{\\rm BH}$ of the non-reverberation estimates. This result implies that our best-fit Rev sample results could be used to measure M$_{\\rm BH}$ from $\\sigma^2_{\\rm rms}$ measurements. Provided that these estimates are as accurate as the variability estimates for the objects in our Rev sample, the M$_{\\rm BH}$ estimates should be more accurate than the ones based on single epoch spectra. Particular care though should be given to the results from the use of this method if the M$_{\\rm BH}$ turns out to be smaller than $\\sim 5\\times 10^{6}-5\\times 10^{7}$ M$_{\\odot}$ (depending on the length of the light curve in use), as in this regime the variability - M$_{\\rm BH}$ relation may have a larger intrinsic scatter.  \\item{} The $\\sigma^2_{\\rm rms}$ vs. M$_{\\rm BH}$ correlation has a small scatter (smaller than expected, see i.e. McHardy et al. 2006; Koerding et al. 2007) because the PSD normalisation decreases with accretion rate as PSD$_{\\rm amp}\\propto \\dot {m}^{-0.8}$. The combination of this relation with the PSD break frequency scaling relations of McHardy et al. (2006) and Koerding et al. (2007) results in the PSD high frequency integral being weakly dependent on accretion rate, being thus a good tool to estimate M$_{\\rm BH}$ (Nikolajuk et al. 2004; 2006; 2009).  We note that the universal constancy of the high frequency tail (rather than the break frequency) of the PSD in {\\it all} BH accreting sources, seems to suggest a deep link with some common physical process experienced by all accreting BH.  \\item{} We measure M$_{\\rm BH}$ or provide stringent upper limits for 65 AGN, the majority of which (44) have no either reverberation mapping or stellar velocity dispersion M$_{\\rm BH}$ estimates and for 6 of which we did not find any M$_{\\rm BH}$ estimate in literature. We also estimate lower limits to the M$_{\\rm BH}$ for the remaining 96 AGN.  \\item{} We observe a significant correlation between  $\\sigma^2_{\\rm rms}$ and $\\dot{m}$, which dissapears when we ``normalize\" $\\sigma^2_{\\rm rms}$ to the M$_{\\rm BH}$ of the objects, in agreement with the previous results of O'Neill et al. (2005) and Zhou et al. (2010). We argue that, the lack of significant intrinsic scatter in the $\\sigma^2_{\\rm rms}$ vs. M$_{\\rm BH}$ relation, and the lack of a significant correlation between variability (normalised to M$_{\\rm BH}$) and $\\dot{m}$ are in agreement with the hypothesis that: all AGN have the same PSD shape, with the break frequencies scaling with M$_{\\rm BH}$ and $\\dot m$ as in McHardy et al. (2006) and Koerding et al. (2007), but only {\\it if} PSD$_{\\rm amp}\\propto \\dot {m}^{-0.8}$, in a fashion similar to what is observed in BH binaries (Axelsson et al. 2005; Gierlinski et al. 2008).  \\item{} We observe that NLS1 are more variable than broad line AGN. This is due to their smaller M$_{\\rm BH}$ and higher accretion rate. Once rescaled, no difference from broad line AGN is required to explain their different variability properties.  \\item{} We observe significant correlations between $\\sigma^2_{\\rm rms}$ and both the bolometric luminosity and the FWHM of the H$\\beta$. Both these correlations disappears when the $\\sigma^2_{\\rm rms}$ vs. M$_{\\rm BH}$ is considered, suggesting that the formers are just by-products of the variability vs. M$_{\\rm BH}$ dependence.  \\item{} We observe, for the first time at very high significance ($>99.99$ \\%), a correlation between $\\sigma^2_{\\rm rms}$ and 2-10 keV spectral index. This correlation provides indirect support to previous suggestions that the spectral index in AGN depends on accretion rate. In particular, the CAIXAvar $\\sigma^2_{\\rm rms}-\\Gamma$ relation can be fully explained if $\\Gamma \\propto \\dot {m}^{0.1}$, as suggested by Sobolewska \\& Papadakis (2009).  \\end{itemize}  In summary, all our results are in agreement with a picture where, to first approximation, all AGN have {\\it the same} variability properties once rescaled properly for  M$_{\\rm BH}$ and accretion rate. This confirms and extends the work of McHardy et al. (2006) and Koerding et al. (2007) who studied the PSD of less than two dozen sources. In our case, this picture is based on the study of a significantly larger number of objects. Our work demonstrates the significance of studies which are based on the estimation of the excess variance, which although as a statistic carries far less information than a full PSD analysis, its use allows the consideration of large AGN samples, hence establishing the validity of the current paradigm for the X-ray variability of AGN, and even allowing the unexpected detection of important ``complications\", such as the PSD amplitude dependence on accretion rate."
    },
    "1112/1112.2572_arXiv.txt": {
        "abstract": "*{We review the present-day methods of mode identification applied to main sequence pulsators focusing on those that make use of multicolour photometry and radial velocity data. The effects which may affect diagnostic properties of these observables are discussed. We also raise the problem of identification of high $\\ell$ modes which can dominate oscillation spectra obtained from space-based projects.}  \\vspace{-8mm} \\abstract{We review the present-day methods of mode identification applied to main sequence pulsators focusing on those that make use of multicolour photometry and radial velocity data. The effects which may affect diagnostic properties of these observables are discussed. We also raise the problem of identification of high $\\ell$ modes which can dominate oscillation spectra obtained from space-based projects.}  \\vspace{-8mm} ",
        "introduction": "\\label{sec:1}  The accuracy of seismic model increases with the number of well identified pulsation frequencies. Depending on the character of modes we can probe different parts of a star and derive constraints on various parameters of model and theory. Since asteroseismology has entered the space era, the number of detected peaks in the frequency spectra has grown immensely. However, to explore these data, the unequivocal identification of mode geometry is required. In the case of main sequence pulsators this is not an easy task because their oscillation spectra do not exhibit equidistant or regular patterns. An alternative is the usage of the photometric and spectroscopic variations. In this paper, we summarize potentials and limits in using these observables for mode identification.  In Section\\,\\ref{sec:2} we recall the photometric diagrams for $\\beta$ Cephei, SPB (Slowly Pulsating B-type) and $\\delta$ Scuti star models. We mention the most important effects which may affect diagnostic properties of such diagrams, i.e., convection, rotation and model atmospheres. In Section\\,\\ref{sec:3} we describe a method which, beside mode identification, can yield constraints on models and theory. A prospect for identification of high $\\ell$ modes is discussed in Section\\,\\ref{sec:4}. The last Section contains conclusions.  \\vspace{-8mm} ",
        "conclusions": "\\label{sec:5} \\vspace{-3mm}  In this short review, we mentioned the basic properties and uncertainties of mode identification from multicolour photometry and radial velocity data. We focused on main sequence pulsators because of their irregular oscillation spectra.  Then, we checked a prospect for identification of the high degree modes. Such modes are highly expected in the rich frequency spectra obtained from, e.g., the CoRoT, Kepler and BRITE data. Identification of the high $\\ell$ modes only from multicolour photometry is rather hopeless. However, it turned out that it is possible to get some constraints on intermediate $\\ell$ modes from the diagrams using amplitudes and phases of the photometric and radial velocity changes. Unfortunately, there is no prospect for extracting modes with $\\ell>30$ from these data. Here only the analysis of the line-profile variations may help.   \\begin{acknowledgement} The authors acknowledge partial financial support from the Polish MNiSW grant No. N N203 379 636. \\end{acknowledgement}  \\vspace{-12mm}"
    },
    "1112/1112.1693_arXiv.txt": {
        "abstract": "We develop an efficient, non-parametric Bayesian method for reconstructing the time evolution of the dark energy equation of state $w(z)$ from observational data.  Of particular importance is the choice of prior, which must be chosen carefully to minimise variance and bias in the reconstruction.  Using a principal component analysis, we show how a correlated prior can be used to create a smooth reconstruction and also avoid bias in the mean behaviour of $w(z)$.  We test our method using Wiener reconstructions based on Fisher matrix projections, and also against more realistic MCMC analyses of simulated data sets for Planck and a future space-based dark energy mission. While the accuracy of our reconstruction depends on the smoothness of the assumed $w(z)$, the relative error for typical dark energy models is $\\lesssim10\\%$ out to redshift $z=1.5$. ",
        "introduction": "The nature of dark energy (DE), the source driving the acceleration of the universe in the framework of general relativity, has remained a mystery since the acceleration was first discovered \\cite{RiessPerl}. With the accumulating observational data, including supernovae (SN), cosmic microwave background radiation (CMB) and large scale structure (LSS) data, we hope to understand whether DE is a cosmological constant or whether it is dynamical.   Dynamical dark energy is characterised by having an equation-of-state $w \\neq -1$, and which in general could be a function of redshift $w(z)$ or, equivalently, the scale factor $w(a)$.  One of the key goals in dark energy observations is to constrain the equation of state and its evolution.   The time evolution of $w$ can, in principle, be reconstructed from data using either parametric or non-parametric methods (e.g. \\cite{Sahni:2006pa}). Non-parametric methods have the advantage that they do not assume an $ad~hoc$ functional form of $w(z)$, which could lead one to miss evidence for a different kind of evolution.   Unfortunately, non-parametric methods generically contain many degrees of freedom which can lead to parameter degeneracies, making them difficult to explore.  Many efforts have been made in the literature to develop non-parametric methods, which can vary greatly depending on the nature of the data and the choices of the functions to be reconstructed.  One particular focus has been using SN and other standard candle data to reconstruct various ways of paramterizing the background cosmology; the techniques either fit directly to SN magnitudes or their luminosity distances $D_L(z)$, or to more indirect quantities such as dark energy density, $\\rho_{DE}(z)$, the expansion history, $H(z)$, and the equation of state $w(z)$, or much more indirect quantities such as the potential of the DE scalar field(s) \\cite{Huterer:1998qv,Saini:1999ba,Chiba:2000im,Huterer:2000mj,Weller:2001gf,Alam:2003sc,Daly:2003iy,Shafieloo:2005nd,Dick:2006ev,Shafieloo:2007cs,Clarkson:2010bm}.  Unfortunately, many of these techniques are specific to luminosity distance measurements, or more generally to measures of the background expansion rate, making it difficult to include other types of data.  The background measurements can be limited by our knowledge of other parameters, in particular the value of the dark energy density today, $\\Omega_{DE}$ \\cite{Chiba:2000im,Alam:2003sc}.  Thus, it is useful to include measurements which could be sensitive to the growth of structure; such measurements, including weak lensing, galaxy clustering and the integrated Sachs-Wolfe effect (ISW), can help to break the degeneracies and can potentially be the basis for consistency tests of the DE paradigm \\cite{Mortonson:2008qy, Mortonson:2010fy}.  When comparing background and structure evolution data, it is useful to focus on more indirect parameterisations like the equation of state, $w(z)$.  Non-parametric approaches typically would expand $w(z)$ in terms of bins or basis functions.  A powerful tool is the principal component approach, which uses forecasts of the future experiments in the form of Fisher matrix projections to find linear combinations of these basis functions which are independent and well determined \\cite{ Huterer:2002hy, Crittenden:2005wj, Albrecht:2007qy, Tang:2008hm, Albrecht:2009ct, Kitching:2009yr, Mortonson:2010px,Clarkson:2010bm}.  One reconstruction approach is to make a truncation of the best constrained principal components, but the number of modes kept is an open question.  This often is decided without reference to the data, potentially missing evidence because it is not expected.   This simple truncation can lead to significant biases and unrealistically small errors in the reconstruction where the data are poor or are absent entirely \\cite{ Huterer:2002hy}.    It is clear that some bias in reconstructions is inevitable, particularly  when one attempts to reconstruct $w(z)$ in regimes where the data are poor or are absent entirely.  However, it is worth investigating whether a better reconstruction method can be found.  In a fully Bayesian approach, one can explicitly specify a prior on the behaviour of $w(z)$, which then determines which modes are kept in the reconstruction, based on where the evidence from the data outweighs the prior.  But the crucial question is then to understand the best way to specify priors in order to minimise the bias, while at the same time reducing the impact of data noise in the reconstruction. Here we extend previous work, where we assumed that $w(z)$ could be treated as a Gaussian random field with a proposed correlation function \\cite{Crittenden:2005wj}; this correlation function effectively enforces a smoothness on the reconstruction.   Our approach is close in spirit to recently proposed Gaussian Process (GP) method \\cite{Holsclaw:2010nb,Holsclaw:2010sk}, which applied a similar prior to reconstructions of $w(z)$ using SN data, though there are important differences as described below.   In this paper we discuss how to choose a prior and implement it in a computationally efficient method to reconstruct $w(z)$ non-parametrically from the observed data. As we will show, it is very accurate (with relative error $\\lesssim$ 10\\% for a range of dark energy models), computationally cheap and fast, straightforward to implement, and can be used to analyse any kind of cosmological data. After discussing general issues regarding the choice of prior correlation functions, we test our reconstruction accuracy using the simulated mock data for Planck \\cite{Planck} and a future space-based mission.  We then conclude with a discussion of how to extend this method to other parameterization and how to calculate more realistic priors. ",
        "conclusions": "Here we have outlined a Bayesian method for reconstructing $w(a)$, which essentially involves multiplying the likelihood of a non-parametric description of $w(a)$ by a suitably defined prior.  This prior is defined in such a way as to make the reconstructions smooth and independent of the binning choice, eliminating flat directions in parameter space which are a generic problem for non-parametric approaches.  With the right choice of prior, we can also eliminate the dependence on the fiducial model;  reconstructions where the data are poor or non-existent are instead determined by the assumed smoothness of the function, tending towards the mean of the measured $w(a).$   However, application of a prior always implies some bias; in this case, the bias is largest for models which are quickly changing in redshift or scale factor.  The advantages of this approach are that it makes one's priors explicit and it is easy to implement; once one decides on the prior parameters, e.g. the correlation length and the strength of the prior, it is straightforward and fast to evaluate the prior likelihood for whatever choice of binning is used.  This can be combined with the data in an MCMC analysis, searching for the highest posterior solution for the $w$ bin amplitudes; indeed, by breaking the parameter degeneracies, the prior will speed up MCMC analyses.  With the accumulating high-quality data of SN, CMB and LSS, our method is an ideal tool to help reveal the nature of dark energy in the near future.  Here we have focused on the methodology only, but in future work we will apply this prior to present data."
    },
    "1112/1112.4922.txt": {
        "abstract": "Magnetic fields are everywhere in nature and they play an important role in every astronomical environment which involves the formation of plasma and currents.  It is natural therefore to suppose that magnetic fields could be present in the turbulent high temperature environment of the big bang.   Such a primordial magnetic field (PMF)  would be expected to manifest itself in the cosmic microwave background (CMB) temperature and  polarization anisotropies, and also in the formation of large- scale structure.  In this review we summarize the theoretical framework which we have developed to calculate the PMF power spectrum to  high precision. Using this formulation, we summarize calculations of  the effects of a PMF which take accurate  quantitative account of the time evolution of the cut off scale. We review the constructed numerical program, which is without approximation, and an improvement over the approach used in a number of  previous works for studying the effect of the PMF on the cosmological perturbations. We demonstrate how  the PMF is an important cosmological physical process on small scales. We also summarize the current  constraints on the PMF  amplitude $B_\\lambda$ and the power spectral index $n_B$ which have been deduced from the available CMB observational data by using our  computational framework. ",
        "introduction": "%_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ Many astrophysical and cosmological phenomena in the universe are affected by magnetic fields over a broad range of scales. Indeed,  magnetic fields with a strength of $B \\sim 1.0~ \\mu G$ have been detected \\cite{Wolfe:1992ab,Clarke:2000bz,Xu:2005rb} even on scales as large as that of galactic clusters. Such magnetic fields are frozen-in to the ionized baryons.  The magnetic energy density scales as $B^{2} \\propto a^{-4}$ while the baryon density scale as $\\rho \\propto a^{-3}$.  Hence, we can relate the field strength to the cosmic density $B^3 \\propto \\rho^2$. Considering  that clusters of galaxies observed today would have isotropically-collapsed  relative to the background density, a cluster magnetic field  of $B \\sim 1.0~\\mu G$ would correspond to a primordial magnetic field (PMF) of order $\\sim 1.0~n G$  at the epoch of  photon last scattering, $z \\sim 1100$.  The origin of such a cosmological primordial magnetic field has been  an area of active research by many authors. One expects that such a field would have a random distribution of orientations and field strength.  If  the PMF has a nearly scale invariant spectrum, an origin from vector potentials generated during the inflation epoch  is one of the best candidates \\cite{Turner:1987bw,Ratra:1991bn,Bamba:2004cu}. Cosmological phase transitions could also produce a PMF with a bluer spectrum \\cite{Vachaspati:1991nm,Kibble:1995aa,Ahonen:1997wh,Joyce:1997uy}. Several authors have also discussed magnetic field generation on  smaller scales during or after the epoch of photon last scattering  ($z ^<_\\sim  1100$) \\cite{Takahashi:2005nd,Hanayama:2005hd,ichiki:2006sc}. Since each model for the generation of the PMF  involves  different length scales, the spectral index of the PMF power spectrum, $n_\\mathrm{B}$ correlates with the generation models. Therefore, constraining $n_\\mathrm{B}$ can lead to constraints on models for the generation of the  PMF.  A primordial magnetic field (PMF) affects electrons and protons through the Lorentz force. Subsequently,  ionized baryons affect photons through  Thomson scattering. In this way the PMF affects photons indirectly in the early universe. A number of studies in the literature \\cite{ Subramanian:1998fn, Mack:2001gc, Subramanian:2002nh, Lewis:2004ef, Yamazaki:2004vq, Kahniashvili:2005xe, Challinor:2005ye, Dolgov:2005ti, Gopal:2005sg, Yamazaki:2005yd, Kahniashvili:2006hy, Yamazaki:2006mi, Yamazaki:2006bq, Yamazaki:2006ah, Giovannini:2006kc, Yamazaki:2007oc, Paoletti:2008ck, Finelli:2008xh, Yamazaki:2008bb, 2008nuco.confE.239Y, Sethi:2008eq, Kojima:2008rf, 2008PhRvD..78f3012K, Giovannini:2008aa, Yamazaki:2009na, 2010PhRvD..81j3519Y, 2010AIPC.1269...57Y} have analyzed effects of a  fiducial PMF of order $\\sim 1~nG$ on the matter power spectrum as well as the  temperature fluctuations and polarization anisotropies of the cosmic microwave background (CMB). Many others have been, also been attempting to constrain the strength of the PMF by various means, e.g. the non-gauusianity of the temperature fluctuations of CMB\\cite{Brown:2005kr,Seshadri:2009sy,Caprini:2009vk}, Faraday rotation effects\\cite{Kosowsky:1996yc,Kolatt:1997xu,Kosowsky:2004zh,Campanelli:2004pm}, CMB anomalies\\cite{Chen:2004nf,Bernui:2008ve,2008PhRvD..78f3012K}, and effects on  large scale structure (LSS) \\cite{Sethi:2003vp,Sethi:2004pe}. These studies have  have indicated  that the  PMF effects are  mainly manifest  on the smallest scales in the linear regime, i.e. before the formation of nonlinear structure. This is of interest because, there are possible discrepancies between observation and theory in the linear regime on small angular scales which could be attributable to  a PMF. However, to accurately estimate the influence  of not only a PMF but also non-linear effects, it is crucial to have the proper tools to constrain the  PMF  from cosmological observations.  The purpose of this review is to summarize the development of such tools.  We have developed methods \\cite{Yamazaki:2006mi, Yamazaki:2006bq, Yamazaki:2006ah, Giovannini:2006kc, Yamazaki:2007oc, Yamazaki:2008bb} to analyze  the  effects of a PMF on the matter and CMB power spectra.  The   parameters characterizing  the PMF have then been constrained from a fit \\cite{Yamazaki:2009na} to the observational data. By simultaneously fitting  the matter and CMB contributions, we have shown that  a more comprehensive and accurate model for  the PMF effects in cosmology can be obtained.   This allows for  improved constraints on the  parameters characterizing the PMF.  In section II of this review we will introduce the basic  cosmological physics of the PMF  and in sections III and IV we will discuss its effects  on the matter power spectrum and CMB.  In section V we will discuss the results of a search for constraints on parameters of the PMF from observations of the matter power spectra and CMB. In   section IV  we will also explore the  possibilities for  new physics to be obtained from a better understanding of  the PMF, e.g. contributions of the PMF to the  gravitational wave mode (BB-mode). %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ ",
        "conclusions": "%_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ Effects of a PMF on the early universe have  been studied by a number  of researchers including ourselves.  A summary of some of the relevant topics and their implications for different  PMF field strengths is given in Table 2.  For example, in the first row, if a PMF  originated during a post inflation epoch, its strength is probably $|\\mathrm{B}_\\lambda| \\ll 1 nG$, while if it was generated during inflation its field strength  at the surface of last photon scattering was likely $|\\mathrm{B}_\\lambda| \\sim 1 nG$. Furthermore the BB mode of the CMB from the tensor component is strongly affected by the gravitational wave background. Therefore, the BB mode also shows characteristics of the gravitational wave background from the PMF and an inflation origin.  If the PMF has a field strength order of $1~nG$, then from figure \\ref{fig3} and \\ref{fig4}, we expect that the gravity wave background  (without the PMF) dominates the BB mode for smaller $\\ell < 200$, while, the PMF dominates this mode for $\\ell > 200$. Hence, if the gravity wave background without the PMF is generated by inflation, the origin of the gravity wave background on large scales will be inflation, while  the background for smaller scales is from the PMF.  Based upon the current status of the field, we can summarize the following: 1) A PMF mainly affects the CMB temperature fluctuations and polarization anisotropies for small angular scales; 2)  Fits to observations of the CMB polarization anisotropies for small angular scales are generally better when effects of the PMF are included; and 3)  A PMF can either increase or decrease the matter density fluctuations on scales  less than that  of galaxy clusters. \\begin{table}[b] \\begin{tabular}{|m{10em}||m{17em}|m{17em}|} \\hline \\multicolumn{1}{|c||}{Topics} & \\multicolumn{1}{c}{$|\\mathrm{B}_\\lambda| \\ll 1 nG$}& \\multicolumn{1}{|c|}{$|\\mathrm{B}_\\lambda| \\sim 1 nG$}\\\\ \\hline \\hline %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ the origin of PMF & post-inflation & Inflation \\\\ \\hline %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ the gravitational wave background & Inflation & Large scale\u0081Finflation Small scale\u0081FPMF \\\\ \\hline %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ the formation of LSS & necessity of some generation model for the magnetic field or  vorticity  & necessity for proper initial conditions considering the magnetic field and vorticity \\\\ \\hline %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ the magnetic field in clusters of galaxies & necessity of some global magnetic field generation mode + amplification model  & PMF origin \\\\ \\hline %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ the re-ionization and the first objects & --- & Formation of the  first stars and the re-ionization are stimulated by a PMF.\\\\ \\hline \\end{tabular} \\caption{Effects of the PMF on each physical processes} \\end{table} %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ \\subsection{Comparison with other determinations} %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ As noted above, in our previous numerical estimate \\cite{Yamazaki:2007oc},  the curvature perturbations of the scalar mode  were too small to stabilize the numerical calculations for large scales and early times. This is caused by an unwanted cancellation between contributions from the  primordial magnetic field and the primordial radiation fields which seemed to imply excess power for the lowest mulitpoles of the CMB. In order to solve this problem, we have adopted the  numerical approach  and  initial conditions for the matter contributions from    Ref.~\\cite{Kojima:2009ms}.  This formulation  has been developed independently and gives consistent results for large scales and early  times. We have now confirmed that our results are consistent not only with results of Ref.~\\cite{Kojima:2009ms}, but also those of  other groups which employ semi-numerical methods [e.g.~Ref.~\\cite{Paoletti:2008ck}]. %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ \\subsection{Influence of the PMF on Large Scale Structure} %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ In Ref.~\\cite{Yamazaki:2009na} we deduced  the effects of a PMF on the  matter energy density fields by considering a stochastic PMF that depends upon scale.   We then  quantitatively discussed the effect of a PMF on the seeds of LSS in the early universe. We have also considered more general effects of the PMF than those considered in  previous work. For example, we included not only the magnetic field tension but also the increases in pressure and energy density perturbations from the field. Furthermore, by considering the correlation between the PMF and the matter density fluctuations, and utilizing  mathematically exact stochastic PMF power spectrum sets,  we have obtained reasonable and accurate models for the evolution of baryons, CDM, and photons, and therefore, the large scale structure. We have shown \\cite{Yamazaki:2009na} that the PMF can play a very  different role on the evolution of density perturbations depending upon  how the PMF and matter fluctuations are  correlated. We have also considered the fact that after  decoupling, the CDM is  influenced indirectly by the PMF through gravitational interaction.  We reported in Ref.~ \\cite{Yamazaki:2007oc}  that a PMF on small angular scales scales ($\\lambda \\approx$1Mpc) and a field strength of  $B_\\lambda \\sim 1$ nG \\cite{Mack:2001gc,Lewis:2004ef,Yamazaki:2004vq,Yamazaki:2006bq,Yamazaki:2006ah,Durrer:1999bk} provides a new interpretation for the excess of CMB anisotropies for high multipoles. We showed that if a PMF with such strength was present, it is very likely that it has affected the formation of large scale structure. In Ref.~\\cite{Yoshida:2003sy} it was suggested that in order to avoid false coupling between  the baryons and the CDM on small scales, one should use  independent transfer functions for the baryons and CDM. A PMF would be another source of this difference in the transfer function for baryons and the CDM. Since the density perturbations in the early universe have evolved to the present LSS, the evolution of the LSS with a PMF becomes  different from that without a PMF. We have shown that the baryon and CDM energy density perturbations evolve  very differently in the presence of a PMF;  and with the PMF taken into consideration, the evolution of large scale structure  becomes more complicated. %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ \\subsection{Future Observations} %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ %_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/_/ Magnetic fields affect the  generation and evolution of a wide variety  of cosmological and stellar objects in the universe, e.g. galaxy clusters, galaxies,  primordial stars, etc.,  (cf.~Table 2). Therefore, it is important to constrain the parameters characterizing the  PMF  in order to study their effects on physical processes in the universe. In this regard, forthcoming   missions, e.g. QUIET, PolarBear, and Planck are expected  to provide much more precise data on the CMB temperature and polarization anisotropies.   At the same time data on the formation of large scale structure will continue to be  measured more precisely both by future space-based observations and from  continuing ground-based observations such as the SDSS and 2dF projects.  As summarized in this review, at the present time, we now understand the   main effects of a  PMF  on both the CMB and the matter power spectrum.  The important task remaining in the future   will be to apply our new and correct numerical methods for  calculating all PMF fluctuations of the scalar, vector and tensor modes in the TT and BB power spectra (see figures \\ref{fig9}, \\ref{fig10} and \\ref{fig11}) to fit the new observations as they emerge.   Only through a critical comparison with such future data can  we hope to ultimately clarify the  magnitude and origin  of the PMF in the early universe.  Although, some degeneracy between the standard cosmological  and PMF parameters is expected in the present data, we believe that by simultaneously  investigating both the  CMB anisotropies and matter power spectra one can  resolve these degeneracies. The ultimate establishment of the existence and properties of a PMF will permit a better understanding of its  generation and evolution and also provide new insight into the formation of LSS as well as a possible new probe of the physics of the early universe."
    },
    "1112/1112.1696.txt": {
        "abstract": "The stellar haloes of galaxies can currently be studied either through observations of resolved halo stars or through surface photometry. Curiously, the two methods appear to give conflicting results, as a number of surface photometry measurements have revealed integrated colours that are too red to be reconciled with the halo properties inferred from the study of resolved stars. Several explanations for this anomaly have been proposed -- including dust photoluminescence, extinction of extragalactic background light and a bottom-heavy stellar initial mass function. A decisive test is, however, still lacking. Here, we explain how observations of the halo of a nearby galaxy, involving a combination of both surface photometry and bright star counts, can be used to distinguish between the proposed explanations. We derive the observational requirements for this endeavour and find that star counts in filters $VI$ and surface photometry in filters $VIJ$ appears to be the optimal strategy. Since the required halo star counts are already available for many nearby galaxies, the most challenging part of this test is likely to be the optical surface photometry, which requires several nights of exposure time on a 4--8 m telescope, and the near-IR surface photometry, which is most readily carried out using the upcoming {\\it James Webb Space Telescope}. ",
        "introduction": "The stellar haloes of galaxies contain the fossil record of galaxy assembly, and observations of resolved halo stars in the Milky Way, Andromeda and other nearby galaxies \\citep*[e.g.][]{Mouhcine et al.,Ibata et al.,Helmi et al.,Rejkuba et al.,Durrell et al., Tanaka et al.} have provided a wealth of information on such structures. However, stellar haloes can also be studied through surface photometry, i.e. observations of the integrated light from large numbers of unresolved stars within each system.  Both techniques have their strong sides as well as shortcomings. Star counts can trace structures much further out into the halo, but this method is currently limited to the brightest halo stars (with the Milky Way halo as a notable exception) and to systems within $\\sim 10$ Mpc. Surface photometry can be applied for more distant galaxies, is sensitive to less luminous halo stars and diffuse light from the interstellar medium, but works only for the relatively bright inner regions of halos.  At the current time, the results from these two techniques are difficult to completely reconcile. A number of attempts to study the haloes of edge-on disk galaxies through optical and near-infrared (NIR) surface photometry have revealed halo colours that are too red to be consistent with the halo populations captured through the study of resolved stars \\citep[e.g.][]{Sackett et al.,Lequeux et al. a,Lequeux et al. b, Rudy et al.,James & Casali,Lequeux et al. c,Zibetti et al.,Zibetti & Ferguson,Bergvall et al. b}. The most recent observations indicate that this red excess turns up at extremely faint surface brightness levels \\citep[$\\mu_i\\approx 27$--29 mag arcsec$^{-2}$;][]{Zibetti et al.,Bergvall et al. b} and at wavelengths around the $i$-band ($\\approx 7600$ \\AA). Older measurements \\citep{Rudy et al.,James & Casali} suggest that it may also continue into the near-infrared (12000--22000 \\AA). On the other hand, surface photometry carried out at wavelengths shortward of the $i$-band seem to turn up nothing anomalous \\citep[][]{Zibetti et al.,Jablonka et al.,Bergvall et al. b}.  It has been suggested that these `red halos' could be due to mundane observational effects like instrumental scattering \\citep{de Jong} or flux residuals left over from an incomplete sky subtraction \\citep{Zheng et al.,Jablonka et al.}. In some of the claimed detections, this may well be the case. However, \\citet{Bergvall et al. b} correct for instrumental scattering (point spread function) effects and are still left with a pronounced red excess in their study of halos around stacked, edge-on low surface brightness disks. They also show that while the halos of these disks have extremely red colours, the disks themselves do not display any anomalous colours at the corresponding surface brightness levels. This strongly argues against a sky flux residual as the main culprit, since this would result in strange colours in both components.  The red halo phenomenon is not necessarily confined to disk galaxies only. Near-infrared excesses have been reported in the faint outskirts of blue compact galaxies \\citep{Bergvall & \u00d6stlin,Bergvall et al. a}, although \\citet{Zackrisson et al. b} and \\citet{Micheva et al.} have questioned the need for extreme stellar population properties to explain these. There are also older reports of extremely red optical colours in the outskirts of cluster elliptical galaxies \\citep{Maccagni et al.,Molinari et al.}. However, no similar anomalies were seen in the recent study by \\citet{Tal & van Dokkum} of the outskirts of 42000 stacked early-type galaxies (mostly located in galaxy groups).  If one assumes that the red excess in the halos of edge-on disks is not due to instrumental scattering or sky flux residuals, this still leaves room for a number of possible explanations  -- including a bottom-heavy stellar initial mass function, diffuse light from the interstellar medium and extinction of extragalactic background light \\citep{Zackrisson et al. a,Zackrisson et al. c,Bergvall et al. b}. These solutions can all give rise to similar halo colours. Given the extremely challenging task of improving current surface photometry measurements with existing telescopes, it has so far remained unclear if the degeneracy between these really can be broken. In this paper, we describe a method to discriminate between these scenarios using a combination of star counts and surface photometry.  In Section.~\\ref{explanations}, we review a number of explanations for the red excess that have previously been demonstrated to be broadly consistent with the available data. The principle of our proposed test is described in  Section~\\ref{test}, and the observational requirements derived in Section~\\ref{requirements}. Section~\\ref{discussion} discusses a number of potential target galaxies that meet the distance requirements of the test. We also provide a worked example for the required exposure times in the case of the edge-on disk galaxy NGC\\,4565, which is deemed to be one of the most suitable candidates. Our findings are summarized in Section~\\ref{summary}. ",
        "conclusions": "\\label{discussion} The need to obtain matching star counts and surface photometry data requires a nearby target (Sect.~\\ref{starcounts}), with a large angular halo diameter as a result ($\\approx 5$--10$\\arcmin$). Since the the HST currently represents the best option for the star counts (with a field of view of $3.3\\arcmin \\times 3.3\\arcmin$ for the ACS instrument), the proposed test will realistically be limited to a small region of the halo. While this means that halo substructure (in the form of satellite galaxies and tidal streams) may well interfere with the proposed measurements, this should not affect the outcome of the proposed test, since such structures are expected to affect the star counts and surface photometry in similar ways. In fact, the modern view of a stellar halo is that of large collection of substructures that slowly dissolve into the diffuse component, which is created by both accretion events and in-situ star formation \\citep[e.g.][]{Bullock & Johnston,Zolotov et al.}. At this time, it is unclear whether the red excess originates in the diffuse halo component, in halo substructure or in both.  Most of the massive ($V_{\\rm rot}$$>$100 km/s), edge-on (inclination $>$ 85 degrees) disk galaxies that have reliable distance estimates that meet the distance criteria for $HST$ star counts ($D\\leq 16$ Mpc) have already been targeted by the GHOSTS\\footnote{Galaxy Halos, Outer disks, Substructure, Truncations, Star Clusters} survey \\citep[][]{de Jong et al.,Radburn-Smith et al.}, which means that the star counts data are already available. To implement the proposed test, one therefore simply has to obtain matching surface photometry data. Below, we discuss a number of potential targets for which star counts are already available from GHOSTS or elsewhere --  NGC 5907, NGC 7814, NGC 3109 and NGC 4565. We argue that NGC 4565 is likely to be the best target for this endeavour, and present a worked example of the required exposure times on telescopes of various sizes. As will be demonstrated, a 4--8 m class telescope represents the most practical option for the $V$ and $I$-band surface photometry.  The large diameters of these halos at the relevant isophotes ($\\approx 5$--10$\\arcmin$) imply that the surface photometry measurements either have to be carried out with wide-field imagers (e.g. Subaru/Suprime-Cam\\footnote{http://www.naoj.org}, VISTA\\footnote{http://www.vista.ac.uk/}, CFHT/MegaCam\\footnote{http://www.cfht.hawaii.edu/Instruments/Imaging/MegaPrime/}) - or in sky-chopping mode. Both approaches are admittedly very challenging at the extremely faint isophotes where the red excess is likely to turn up and a combination of both may provide the best route in the long run.  While wide-field imagers save observing time, one tends to have better overview and control of the bias level and flatfielding accuracy in single chip imaging with single amplifier readout. Additionally, the single chip/amplifier setup avoids or at least significantly reduces the effects of filter-dependent large-scale variations of the pixel response, crosstalk and amplifier gain differences. On the other hand,  due to the necessary proximity of the target galaxy, using a single chip would imply that the level of the background sky cannot be estimated from the target observations. Therefore chopping mode observations are necessary in order to obtain separate imaging of the sky. Chopping mode is not a traditional observing strategy in the optical. However, it is routinely used to reach surface brightness levels $\\sim 10000$ times below the sky in the near-IR regime \\citep[e.g.][]{Cairos et al.,Noeske et al.,Vaduvescu et al.}, where the temporal sky variations are rapid and can easily result in residual low frequency variations of the background in the sky-subtracted frame. The optical sky is in contrast much more stable and there is no obvious reason why chopping mode in the optical would not be at least as successful as it is in the comparatively more troublesome NIR regime. For this reason, our exposure time estimates are based on the assumption that the surface photometry observations are carried out in sky-chopping mode.  \\subsection{NGC\\,5907} The edge-on disk galaxy NGC 5907, at a TRGB distance of 16.8$\\pm$0.8 Mpc \\citep{Radburn-Smith et al.}, by far represents the most famous case of a claimed `red halo' \\citep[e.g.][]{Sackett et al.,Lequeux et al. b,Rudy et al.,Lequeux et al. c,James & Casali,Zepf et al.}. As discussed in Sect.~\\ref{surfphot}, the observed $V-I$ profile reaches an anomalous value already at $\\mu_I\\geq 26$ mag arcsec$^{-2}$ \\citep{Lequeux et al. b,Lequeux et al. c}. Interest in this object declined with the discovery of what appeared to be the remnants of a disrupted dwarf galaxy on one side of the halo \\citep{Shang et al.} and suggestions that this ring-like feature, in combination with contamination by foreground stars, could have resulted in the spurious halo detection \\citep{Zheng et al.}. To some extent, this is a matter of semantics, since tidal features are nowadays regarded as ubiquitous features of stellar halos \\citep[e.g.][]{Bullock & Johnston}. The details of how this tidal feature supposedly explains the extreme halo colours reported by others remain unclear, especially since the ring itself appears to have normal colours \\citep{Zheng et al.}.  While the colour anomaly of NGC\\,5907 is well reported, its distance makes it a very difficult target for our test. To obtain accurate point source photometry 2 mags below TRGB would require at least 15 HST orbits per pointing. Several HST pointings would probably be required, as \\citet{Martinez-Delgado et al. a} have reported a complex system of streams around this galaxy, and halo measurements on and off the streams would be required to determine the origin of the colour anomaly. The 6 orbit HST observation of the GHOSTS survey was barely deep enough to provide a TRGB distance to this galaxy and contains too few stars for a reliable luminosity profile. It is furthermore located poorly on the major axis near the disk truncation at the end of the disk and gives no clean halo measurement.  \\subsection{NGC\\,7814} The edge-on galaxy NGC 7814, at a distance of 14.4$\\pm$0.7 Mpc \\citep{Radburn-Smith et al.} has also been argued to exhibit an anomalous $V-I$ colour in its outskirts \\citep{Lequeux et al. a}. This object is extraordinary in the sense that the excess seemingly turns up at relatively bright isophotes ($\\mu_I\\approx 23$--24 mag arcsec$^{-2}$). No follow-up observations have apparently been undertaken. However, due to the very prominent bulge in this galaxy, and the fact that the colours were not measured in the polar direction but in the outer disk region, it is not clear which component of the galaxy is responsible for this.  The prominent bulge and the relatively large distance make NGC\\,7814 a difficult target for our test. The HST star counts suffer from crowding at about $\\mu_I\\approx$ 24 mag arcsec$^{-2}$ at this distance, exactly where the colour anomaly becomes significant if the \\citet{Lequeux et al. a} measurements are correct. Extensive HST imaging data does exist for this galaxy with 4 ACS and WFC3 fields along the minor axis out to about 14 arcmin radius ($\\sim$60kpc) and several more fields along the major axis. While absolute matching of the star counts and integrated light profile in the central region of the galaxy may not be possible, the relative profiles could definitely be compared over at least 3 magnitudes in surface brightness where \\citet{Lequeux et al. a} claim a strong colour gradient. However, a cursory look at the pipeline-reduced Sloan Digital Sky Survey\\footnote{www.sdss.org} mosaics (data release 8) does not give any obvious support for an anomalous or even rising $g-i$ colour at the distance where \\citet{Lequeux et al. a} note a $V-I$ excess.  \\subsection{NGC\\,3109} NGC 3109, at a distance of $1.32\\pm0.06$ Mpc \\citep*{Hidalgo et al.}, represents one of the few high-inclination disk galaxies located sufficiently close to make groundbased star counts a realistic alternative for our proposed test. Star counts reaching $\\approx 2$ mag below the TRGB in the $V$ and $I$ bands are already available in regions located in the direction perpendicular to the disk \\citep{Hidalgo et al.}, allowing the surface densities of resolved stars to be traced $\\approx 10\\arcmin$ ($\\approx 4$ kpc) into the halo (out to $\\mu_V\\approx 30$ mag arcsec$^{-2}$). There is, however, no reported colour anomaly in the halo of this galaxy, and its relatively low Galactic latitude ($b=23.1\\deg$) makes surface photometry measurements prone to the effects of Galactic cirrus and scattered light from foreground stars.  \\subsection{NGC\\,4565} NGC\\,4565 is a near-perfect edge-on galaxy that is ideally suited for our research. It is at a distance of $D=11.9\\pm 0.3$ Mpc as determined from the TRGB from {\\it HST} imaging \\citep{Radburn-Smith et al.}, which is near enough that the RGB population can be comfortably be measured with HST, while it has an angular diameter small enough ($D_{25}\\sim$16 arcmin)  that a large fraction of its halo can be observed within a single image of a groundbased camera, allowing good flatfielding checks. NGC\\,4565 has a high galactic latitude ($b=86.4\\deg$) and therefore Galactic cirrus and scattered light from foreground stars will have minimal effect on the integrated light observations. There are no obvious tidal streams in existing deep images, but an H{\\sc i} tidal bridge exists to a small dwarf galaxy to the North. However, no anomalous halo colours have been reported for this galaxy as of yet.  Existing {\\it HST} images obtained with the {\\it Advanced Camera for Surveys} ({\\it ACS}) and {\\it Wide Field Camera 3} ({\\it WFC3} cameras as part of the GHOSTS survey with long exposure times ($>$7000 s)  in F606W (equivalent to $V$-band) and F814W ($I$).  So far 5 fields have been obtained, two of which along the minor axis, and one more observation with the {\\it ACS} and {\\it Wide-Field Camera 3} ({\\it WFC3}) cameras working in parallel has been scheduled. The observations are deep enough to reach $\\approx 2.0$ mag below the TRGB, which is sufficient for the proposed test. Starting a few kpc above the disk the color-magnitude diagrams are completely dominated by RGB and AGB stars with no young stars indicated. Even at 30 kpc above the midplane of the galaxy, the estimated contamination of unresolved background galaxies is still less than 2\\%.  To match these {\\it HST} star counts with $V$ and $I$-band surface photometry data, new ground-based observations would be required. Based on the \\citet{Wu et al.} $i_{6660}$-band surface brightness profile (similar to the $R$ band), we estimate that the surface brightness level of the NGC\\,4565 halo in the most distant {\\it GHOSTS} field (at $\\approx 6$\\arcmin{ } from the disk) to be $\\mu_R\\approx28$ mag arcsec${}^{-2}$. This can be converted into corresponding $\\mu_V$ or $\\mu_I$ values, by assuming a shape for the observed halo spectrum in this wavelength range. The \\citet{Marigo et al.} model predicts that an old and metal-poor stellar population with a standard IMF would display $V-R\\approx 0.5$ \\citep[see also the discussion in][]{Jablonka et al.} and $R-I\\approx 0.5$, but if this were the case, there would be no red excess and the test would not be meaningful. \\citet{Zackrisson et al. a} and \\citet{Bergvall et al. b} present models for old stellar populations with bottom-heavy IMFs that are able to reproduce the anomalous $gri$ halo colours of stacked edge-on disks. These models predict $V-R\\approx 0.8$, $R-I\\approx 0.9$. EBL extinction models capable of fitting the $gri$ data can produce normal $g-r$ colours yet $r-i$ as red as those predicted by a bottom-heavy IMF \\citep{Zackrisson et al. c}, suggesting $V-R\\approx 0.5$, $R-I\\approx 0.9$. However, since the $R$-band surface brightness is fixed, the observations become more challenging if $V-R$ is red (implying a fainter $V$-band profile) rather than blue (brighter $V$-band profile). In order not to produce estimates that are overly optimistic, we therefore adopt the more extreme estimates on both colours: $V-R\\approx 0.8$ and $R-I\\approx 0.9$. The $\\mu_R\\approx28$ mag arcsec${}^{-2}$ isophote may then correspond to $\\mu_V\\approx28.8$ mag arcsec${}^{-2}$ and $\\mu_I\\approx27.1$ mag arcsec${}^{-2}$. This is somewhat brighter than the ideal requirements derived in Sect.~\\ref{surfphot}, but sufficient to uncover a red halo if it is similar to the stacked halo detected by \\citet{Zibetti et al.}, $\\mu_I\\approx 26.2$ mag arcsec$^{-2}$.  By integrating over 10\\arcsec\\ strips across this {\\it GHOSTS} field ($3.3\\arcmin \\times 3.3\\arcmin$), we estimate that one can reach the required levels ($\\mu_V\\approx28.8$ mag arcsec${}^{-2}$,  $\\mu_I\\approx27.1$ mag arcsec${}^{-2}$) with an error of $\\sigma=0.1$ mag in approximately 2 hours for the $V$ band and 4 hours for the $I$-band with the 8.2m {\\it VLT} and the {\\it FORS2} instrument (see Appendix A for details). For comparison, the corresponding times on smaller telescopes (with diameters $D=4$, 2, 1 and 0.5m, respectively) would be $T=9$, 37, 148 and 592 hours for the $V$ band and $T=17$, 67, 269 and 1076 hours for the $I$ band. This programme would seem best suited for telescopes with diameters in the 4--8 m range, since smaller telescopes require excessive amounts of observing time. There is, however, otherwise nothing that precludes the use of smaller telescopes for measurements of this type. As demonstrated by \\citet{Martinez-Delgado et al. a,Martinez-Delgado et al. b}, tidal features in halos have in fact been detected at similar isophotes using long observing campaigns at very small telescopes ($D\\leq 0.5$m).  Should the test tend to favour an excess of $I$ band flux over a deficit of $V$ band flux as the likely explanation for the anomalous halo colour, follow-up $J$-band observations may be required to distinguish between ERE and a bottom-heavy IMF. As argued in Sect.~\\ref{surfphot}, a normal stellar halo should display $V-J\\approx 1.5-2.0$. This could also be the $V-J$ colour expected in the case of ERE, since neither the $V$ or $J$ filters are expected to be strongly affected by this mechanism. To rule out a normal $V-J$ colour in the NGC\\,4565 field, one may therefore need to reach as deep as the the $\\mu_J=\\mu_V-2.0$ isophote, which in this case corresponds to $\\mu_J\\approx 26.8$ mag arcsec$^{-2}$. The {\\it JWST} appears to be the best option for this. The estimates provided in Sect.~\\ref{surfphot} and Appendix A indicate that one should be able to reach this level in less than two hours. Despite having a relatively small field of view ($2.2\\arcmin \\times 2.2\\arcmin$), the improved sensitivity of the {\\it JWST/NIRCam} instrument compared to existing {\\it HST} cameras will also allow resolved halo stars to be mapped across the halo, in just a fraction of the {\\it HST }observing time spent on existing {\\it GHOSTS} fields."
    },
    "1112/1112.3143_arXiv.txt": {
        "abstract": "Statistical weak lensing by large-scale structure -- cosmic shear -- is a promising cosmological tool, which has motivated the design of several large upcoming surveys.  Here, we present a measurement of cosmic shear using coadded Sloan Digital Sky Survey (SDSS) imaging in 168 square degrees of the equatorial region, with $r<23.5$ and $i<22.5$, a source number density of 2.2 per arcmin$^2$ and median redshift of $z_\\mathrm{med}=0.52$.  These coadds were generated using a new method described in the companion Paper I \\citep{2011arXiv1111.6958H} that was intended to minimise systematic errors in the lensing measurement due to coherent PSF anisotropies that are otherwise prevalent in the SDSS imaging data.  We present measurements of cosmic shear out to angular separations of 2 degrees, along with systematics tests that (combined with those from Paper I on the catalogue generation) demonstrate that our results are dominated by statistical rather than systematic errors. Assuming a cosmological model corresponding to WMAP7 \\citep{2011ApJS..192...18K} and allowing only the amplitude of matter fluctuations $\\sigma_8$ to vary, we find a best-fit value of $\\sigma_8=0.636^{+0.109}_{-0.154}$ (1$\\sigma$); without systematic errors this would be $\\sigma_8=0.636^{+0.099}_{-0.137}$ (1$\\sigma$). Assuming a flat $\\Lambda$CDM model, the combined constraints with WMAP7 are $\\sigma_8=0.784^{+0.028}_{-0.026}$(1$\\sigma$)$^{+0.055}_{-0.054}$(2$\\sigma$) and $\\Omega_mh^2=0.1303^{+0.0047}_{-0.0048}$(1$\\sigma$)$^{+0.009}_{-0.009}$(2$\\sigma$); the 2$\\sigma$ error ranges are respectively 14\\ and 17\\ per cent smaller than WMAP7 alone.  Aside from the intrinsic value of such cosmological constraints from the growth of structure, we identify some important lessons for upcoming surveys that may face similar issues when combining multi-epoch data to measure cosmic shear. ",
        "introduction": "As a result of gravitational lensing, large scale inhomogeneities in the matter density field produce small but systematic fluctuations in the sizes, shapes, and fluxes of distant objects that are coherent across large scales. This effect was first suggested as a tool for constraining the form of the metric in 1966 by \\cite{1966ApJ...143..379K}. In a more modern context, the two-point statistics of lensing fluctuations allow the only truly direct measurement of the matter power spectrum and the growth of structure at late times, when dark energy has caused an accelerated expansion of the universe \\citep{1998AJ....116.1009R,1999ApJ...517..565P} and affected the growth of structure. Many studies have pointed out that high signal-to-noise cosmic shear measurements would be extraordinarily sensitive probes of cosmological parameters \\citep[e.g.,][]{2002PhRvD..65f3001H, 2004PhRvD..70l3515B}, which led to its being flagged as one of the most promising probes of dark energy by the Dark Energy Task Force \\citep{2006astro.ph..9591A}. Direct measurements of the growth of structure also offer the opportunity to test alternative models of gravity \\citep[e.g.,][]{2011arXiv1109.4535L}.  Cosmic shear measurements were attempted as early as 1967 \\citep{1967ApJ...147..864K}, but until the turn of the millennium \\citep{2000MNRAS.318..625B,2000astro.ph..3338K,2000A&A...358...30V,2000Natur.405..143W}, no astronomical survey had the statistical power to detect it. The difficulty of the measurement is a consequence of the near-homogeneity and isotropy of the universe. An order unity distortion to galaxy images requires an integrated line-of-sight matter over-density of: \\begin{equation} \\Sigma_{\\rm{crit}} =  \\frac{c^2}{4 \\pi G} \\frac{d_S}{d_L \\: d_{LS}} \\end{equation} where $d_S$, $d_L$, and $d_{LS}$ are the angular diameter distances from the observer to the background source, from the observer to the lens, and from the lens to the background source, respectively. A fluctuation in the surface density $\\Delta \\Sigma$ leads to a shear distortion $\\gamma \\sim \\Delta \\Sigma / \\Sigma_{\\rm{crit}}$.  Averaged over large ($\\sim$ 100 Mpc) scales, typical line-of-sight matter fluctuations are only $10^{-3}\\Sigma_{\\rm{crit}}$.  The primary source of noise in the shear measurement, the random intrinsic dispersion in galaxy shapes, is orders of magnitude larger; typically the shape noise results in a dispersion in the shear of $\\sigma_{\\gamma}=0.2$. Worse, even in modern ground-based astronomical imaging surveys, the coherent distortions -- or point-spread function (PSF) -- induced by effects of the atmosphere, telescope optics, and detectors are typically several times larger than the cosmological signal (e.g., \\citealt{2011arXiv1110.4913H} and Paper I in this series). Estimating the distances to the background sources is both crucial \\citep{2006ApJ...636...21M} and difficult \\citep{2008ApJ...682...39M,2010MNRAS.401.1399B}; errors there will modulate the amplitude of the signal through $\\Sigma_{\\rm{crit}}$, biasing inference of the growth of structure.  These obstacles define the observational problem. While the existence of cosmic shear has been established by the first studies to detect the effect, the full potential of cosmological lensing remains to be exploited. Few data sets capable of achieving the signal strength for a cosmologically competitive measurement presently exist -- the Canada-France-Hawaii Telescope Legacy Survey (CFHTLS; \\citealt{2006ApJ...647..116H,2006A&A...452...51S,2007MNRAS.381..702B,2008A&A...479....9F}), the Cosmological Evolution Survey (COSMOS; \\citealt{2007ApJS..172..239M,2010A&A...516A..63S}), and the subset of the SDSS imaging studied here.  However, several large surveys are planned for the immediate and longer-term future that will substantially expand the amount of available data for cosmological weak lensing studies.  In the next few years, these include Hyper Suprime-Cam (HSC, \\citealt{2006SPIE.6269E...9M}), Dark Energy Survey (DES\\footnote{\\texttt{https://www.darkenergysurvey.org/}}, \\citealt{2005astro.ph.10346T}), the KIlo-Degree Survey (KIDS\\footnote{\\texttt{http://www.astro-wise.org/projects/KIDS/}}), and the Panoramic Survey Telescope and Rapid Response System (Pan-STARRS\\footnote{\\texttt{http://pan-starrs.ifa.hawaii.edu/public/}}, \\citealt{2010SPIE.7733E..12K}).  Further in the future, there are even more ambitious programs such as the Large Synoptic Survey Telescope (LSST\\footnote{\\texttt{http://www.lsst.org/lsst}}, \\citealt{2009arXiv0912.0201L}), Euclid\\footnote{\\texttt{http://sci.esa.int/science-e/www/area/index.cfm?fareaid\\ =102}}, and the Wide-Field Infrared Survey Telescope (WFIRST\\footnote{\\texttt{http://wfirst.gsfc.nasa.gov/}}).   For this work, we have combined several methods discussed in the literature as viable techniques for measuring cosmic shear while removing common systematic errors. In Paper I \\citep{2011arXiv1111.6958H}, we began with the PSF model generated by the Sloan Digital Sky Survey (SDSS) pipeline over $\\sim 250$ deg$^2$ that had been imaged many times, and employed a rounding kernel method similar to that proposed in \\cite{2002AJ....123..583B}.  The result, after appropriate masking of problematic regions, was 168 square degrees of deep coadded imaging with a well controlled, homogeneous PSF and sufficient galaxy surface density to measure a cosmic shear signal. The usable area in $r$ band was only 140 square degrees because of a PSF model error problem on the camcol 2 charge-coupled device (CCD), which is suspected to be an amplifier non-linearity problem.  In this work, we use the catalogue from Paper I to produce a cosmic shear measurement that is dominated by statistical errors. Section \\ref{sec:model} enumerates the primary sources of systematic error when measuring cosmic shear using our catalogue (the properties of which are summarized briefly in Sec.~\\ref{sec:catalogues}), and describes our approaches to constraining each of them. In Section \\ref{sec:analysis_tools}, we outline our % correlation function estimator and several transformations of it that are used for systematics tests. Our methods for estimating covariance matrices for our observable quantities (both due to statistical and systematic errors) are described in Sec.~\\ref{sec:allcov}.  % Finally, section \\ref{sec:constraints} presents the constraining power of this measurement alone for a fiducial cosmology, and in combination with the 7-year {\\slshape Wilkinson Microwave Anisotropy Probe} \\citep[WMAP7,][]{2011ApJS..192...18K} parameter constraints to produce a posterior probability distribution over $\\Omega_m h^2$, $\\Omega_b h^2$, $\\sigma_8$, $n_s$, and $w$. We show that in addition to its value as an independent measurement of the late-time matter power spectrum, this measurement provides some additional constraining power over WMAP7 within the context of $\\Lambda$CDM.  We conclude with some lessons for the future in Sec.~\\ref{sec:conclusions}.  While this work was underway, we learned of a parallel effort by Lin et al. (2011). These two efforts use different methods of coaddition, different shape measurement codes, different sets of cuts for the selection of input images and galaxies, and analyze their final results in different ways; what they have in common is their use of SDSS data (not necessarily the same sets of input imaging) and their use of the SDSS {\\sc Photo} pipeline for the initial reduction of the single epoch data and the final reduction of the coadded data (however, they use different versions of {\\sc Photo}).  Using these different methods, both groups have extracted the cosmic shear signal and its cosmological interpretations. We have coordinated submission with them but have not consulted their results prior to this, so these two analysis efforts are independent, representing versions of two independent pipelines. ",
        "conclusions": "\\label{sec:conclusions}  Using coadded imaging constructed from SDSS Stripe 82 data, we constructed a weak lensing catalogue of 1\\,328\\,885 galaxies covering 168 square degrees (Paper I), and showed that the additive shear systematics arising from the PSF are negligible compared to the cosmic shear signal. In this paper, we carried out a cosmic shear measurement that resulted in a 20\\ per cent constraint on $\\sigma_8$ (with all other cosmological parameters fixed).  This adds constraining power beyond that from WMAP7, and serves as an important independent data point on the amplitude of the matter power spectrum at late times. In particular, the primary CMB anisotropies presently provide only a modest constraint on $\\Omega_mh^2$, and (due to the effect of matter density on the growth of structure) there is then an elongated allowed region in the $(\\Omega_mh^2,\\sigma_8)$ plane; see Fig.~\\ref{fig:LCDMconstraints}. The WMAP7-allowed region is ideally oriented for lensing to play a role: the lensing signal at the high-$\\Omega_mh^2$, high-$\\sigma_8$ end of the ellipse leads to a much higher lensing signal than low $\\Omega_mh^2$, low $\\sigma_8$. The low amplitude of cosmic shear observed in this paper eliminates the high-$\\Omega_mh^2$, high-$\\sigma_8$ solutions, and leads to a WMAP7+SDSS lensing solution of $\\sigma_8=0.784^{+0.028}_{-0.026}$(1$\\sigma$)$^{+0.055}_{-0.054}$(2$\\sigma$) and $\\Omega_mh^2=0.1303^{+0.0047}_{-0.0048}$(1$\\sigma$)$^{+0.0091}_{-0.0092}$(2$\\sigma$); the 2$\\sigma$ error ranges are respectively 14\\ and 17\\ per cent smaller than for WMAP7 alone.  We have also carefully evaluated other sources of uncertainty such as the source redshift distribution, intrinsic alignments, and shear calibration, to ensure that our measurement is dominated by statistical errors rather than systematic errors.  This achievement is important when considering that (i) the SDSS data were never designed with this application in mind, and indeed includes several features (e.g. the minimal amount of cross-scan dithering) that cause significant difficulty, and (ii) with the multitude of upcoming multi-exposure lensing surveys in the next few years, it is important to cultivate new data analysis techniques (such as the one used here) that are capable of producing homogeneous data with tight control over PSF anisotropies. As a quantitative measure of the extent of PSF correction possible with SDSS data, we take the RMS residual spurious shear at a particular scale estimated from the star-galaxy correlations, \\begin{equation} \\gamma_{\\rm rms,eq}(\\theta) = \\frac{\\sqrt{{\\cal R}_{\\rm psf}\\,\\xi_{+,\\rm sg}(\\theta)}}{{\\cal R}}. \\end{equation} From Fig.~\\ref{fig:sgxf_w}, we see that this is $\\sim 2\\times 10^{-3}$ at the smallest scales (1--6 arcmin), is $<10^{-3}$ at scales $\\theta>0.1\\,$degree, and drops to $3.7\\times 10^{-4}$ in the final bin (1.2--2.0 degrees).\\footnote{We used ${\\cal R}_{\\rm psf}=0.9$ and ${\\cal R}=1.776$, as described in the text.} There is almost no difference between the $++$ and $\\times\\times$ signals, suggesting that the spurious additive ellipticity signal contains similar amounts of $E$- and $B$-modes\\footnote{Recall that $\\xi_{++}(\\theta)-\\xi_{\\times\\times}(\\theta)$ and $P_E(\\ell)-P_B(\\ell)$ are $J_4$ Hankel transforms of each other.}; something similar was seen in the SDSS single-epoch data via run-by-run comparisons of ellipticity measurements on the same galaxies \\citep[Fig. 8]{2006MNRAS.370.1008M}. This is good news for the use of the $B$-mode as a diagnostic of PSF systematics, although an understanding of the generality of this pattern remains elusive.  A major lesson learned from this project is the importance of {\\em masking bias}, in which the intrinsic orientation of a galaxy affects whether it falls within the survey mask. This is likely the main reason why we had to implement the $\\langle e_1\\rangle$ projection. While we have clearly not exhausted the range of options for removing this bias at the catalogue level, future surveys should be designed to produce more uniform data quality via an appropriate dithering strategy and suppress the masking bias at the earliest stages of the analysis.  Our major limitation in the end was the source number density, which was driven by the fact that our PSF-matching procedure was limited by the worst seeing in the images that we use, and therefore we had to eliminate the images with seeing worse than the median.  This means that the coadds were not as deep as they could have been, and the final effective seeing was 1.31 arcsec (full-width half maximum).  In principle this will be an obstacle to applying this technique in the future, but in fact, that statement depends on context.  For example, for a survey such as HSC or LSST where we expect typically $\\sim 0.7$ arcsec seeing, and with plans to preferentially use the best-seeing nights for $r$ and $i$-band imaging that will be used for shape measurement, it is conceivable that nearly all images intended for lensing will have seeing in the 0.6--0.8 arcsec range.  In that context, a PSF-matched coadd that has the rounding kernel applied may actually not result in much loss of information about the shapes of most useful galaxies, and will have the advantage of the removal of PSF anisotropies.  Moreover, even for surveys for which the loss of information that results from this method may not be suitable for the final cosmological analysis, this method may still serve as a useful diagnostic of the additive PSF systematics."
    },
    "1112/1112.4139_arXiv.txt": {
        "abstract": "{The Gaia satellite will observe about one billion stars and other point-like sources. The astrometric core solution will determine the astrometric parameters (position, parallax, and proper motion) for a subset of these sources, using a global solution approach which must also include a large number of parameters for the satellite attitude and optical instrument. The accurate and efficient implementation of this solution is an extremely demanding task, but crucial for the outcome of the mission.} {We aim to provide a comprehensive overview of the mathematical and physical models applicable to this solution, as well as its numerical and algorithmic framework. } {The astrometric core solution is a simultaneous least-squares estimation of about half a billion parameters, including the astrometric parameters for some 100 million well-behaved so-called primary sources. The global nature of the solution requires an iterative approach, which can be broken down into a small number of distinct processing blocks (source, attitude, calibration and global updating) and auxiliary processes (including the frame rotator and selection of primary sources). We describe each of these processes in some detail, formulate the underlying models, from which the observation equations are derived, and outline the adopted numerical solution methods with due consideration of robustness and the structure of the resulting system of equations. Appendices provide brief introductions to some important mathematical tools (quaternions and B-splines for the attitude representation, and a modified Cholesky algorithm for positive semidefinite problems) and discuss some complications expected in the real mission data.} {A complete software system called AGIS (Astrometric Global Iterative Solution) is being built according to the methods described in the paper. Based on simulated data for 2 million primary sources we present some initial results, demonstrating the basic mathematical and numerical validity of the approach and, by a reasonable extrapolation, its practical feasibility in terms of data management and computations for the real mission.} {} ",
        "introduction": "\\label{sec:intro}  The space astrometry mission Gaia, planned to be launched by the European Space Agency (ESA) in 2013, will determine accurate astrometric data for about one billion objects in the magnitude range from 6 to 20. Accuracies of 8--25~micro-arcsec ($\\mu$as) are typically expected for the trigonometric parallaxes, positions at mean epoch, and annual proper motions of simple (i.e., apparently single) stars down to 15th magnitude. The astrometric data are complemented by photometric and spectroscopic information collected with dedicated instruments on board the Gaia satellite. The mission will result in an astronomical database of unprecedented scope, accuracy and completeness becoming available to the scientific community around 2021.  The original interferometric concept for a successor mission to Hipparcos, called GAIA (Global Astrometric Interferometer for Astrophysics), was described by \\citet{gaia1996} but has since evolved considerably by the incorporation of novel ideas \\citep{iaus248:EH} and as a result of industrial studies conducted under ESA contracts \\citep{gaia2001}. The mission, now in the final integration phase with EADS Astrium as prime contractor, is no longer an interferometer but has retained the name Gaia, which is thus no acronym. For some brief but up-to-date overviews of the mission, see \\citet{iaus248:LL} and \\citet{iaus261:LL}. The scientific case is most comprehensively described in the proceedings of the conference The Three-Dimensional Universe with Gaia \\citep{2005ESASP.576.....T}.  In parallel with the industrial development of the satellite, the Gaia Data Processing and Analysis Consortium \\citep[DPAC;][]{iaus248:FM} is charged with the task of developing and running a complete data processing system for analysing the satellite data and constructing the resulting database (`Gaia Catalogue'). This task is extremely difficult due to the large quantities of data involved, the complex relationships between different kinds of data (astrometric, photometric, spectroscopic) as well as between data collected at different epochs, the need for complex yet efficient software systems, and the interaction and sustained support of many individuals and groups over an extended period of time.  A fundamental part of the data processing task is the astrometric core solution, currently under development in DPAC's Coordination Unit 3 (CU3), `Core Processing'. Mathematically, the astrometric core solution is a simultaneous determination of a very large number of unknowns representing three kinds of information: (i) the astrometric parameters for a subset of the observed stars, representing the astrometric reference frame; (ii) the instrument attitude, representing the accurate celestial pointing of the instrument axes in that reference frame as a function of time; and (iii) the geometric instrument calibration, representing the mapping from pixels on the CCD detectors to angular directions relative to the instrument axes. Although the astrometric core solution is only made for a subset of the stars, the resulting celestial reference frame, attitude and instrument calibration are fundamental inputs for the processing of all observations. Optionally, a fourth kind of unknowns, the global parameters, may be introduced to describe for example a hypothetical deviation from General Relativity.  We use the term `source' to denote any astronomical object that Gaia detects and observes as a separate entity. The vast majority of the Gaia sources are ordinary stars, many of them close binaries or the components of wide systems, but some are non-stellar (for example asteroids and quasars). Nearly everywhere in this paper, one can substitute `star' for `source' without distortion; however, for consistency with established practice in the Gaia community we use `source' throughout.  While the total number of distinct sources that will be observed by Gaia is estimated to slightly more than one billion, only a subset of them shall be used in the astrometric core solution. This subset, known as the `primary sources', is selected to be astrometrically well-behaved (see Sect.~\\ref{sec:primsel}) and consists of (effectively) single stars and extragalactic sources (quasars and AGNs) that are sufficiently stable and point-like. We assume here that the number of primary sources is about $10^8$, i.e., roughly one tenth of the total number of objects, although in the end it is possible that an even larger number will be used.  In comparison with many other parts of the Gaia data processing, the astrometric core solution is in principle simple, mainly because it only uses a subset of the observations (namely, those of the primary sources), which can be accurately modelled in a relatively straightforward way. In practice, the problem is however formidable: the total number of unknowns is of the order of $5\\times 10^8$, the solution uses some $10^{11}$ individual observations extracted from some 70~Terabyte (TB) of raw satellite data, and the entangled nature of the data excludes a direct solution. A feasible approach has nevertheless been found, including a precise mathematical formulation, practical solution method, and efficient software implementation. It is the aim of this paper to provide a comprehensive overview of this approach.  \\begin{figure*} \\begin{center} \\includegraphics[scale=1.0]{fig01.pdf} \\caption{ \\label{fig:blockdiagram} Schematic representation of the main elements of the astrometric data analysis. In the shaded area is the mathematical model that allows to calculate the position of the stellar image in pixel coordinates, and hence the expected CCD data, for any given set of model parameters. To the right are the processes that fit this model to the observed CCD data by adjusting the parameters in the rectangular boxes along the middle. This paper is primarily concerned with the geometrical part of the analysis contained in the dashed box. However, a brief outline of the CCD data modelling and processing (bottom part of the diagram) is given in Sect.~\\ref{sec:signalmodel} and Appendix~\\ref{sec:CTI}.} \\end{center} \\end{figure*}   Concerning notations we have followed the usual convention to denote all non-scalar entities (vectors, tensors, matrices, quaternions) by boldfaced characters. Lower-case bold italics ($\\vec{a}$) are used for vectors and one-dimensional column matrices; upper-case bold italics ($\\vec{A}$) usually denote two-dimensional matrices. Following \\citet{murray1983} the prime ($\\,'$) signifies both matrix transpose ($\\vec{a}'$, $\\vec{A}'$) and scalar multiplication of vectors; thus $\\lVert a \\rVert=(\\vec{a}'\\!\\vec{a})^{1/2}$ defines the magnitude of the vector $\\vec{a}$ as well as the Euclidean norm of the column matrix $\\vec{a}$. Angular brackets denote normalization to unit length, as in $\\langle\\vec{x}\\rangle=\\vec{x}\\lVert\\vec{x}\\rVert^{-1}$. In this notation, no special distinction is thus made between vectors as physical entities (also known as geometric, spatial or Euclidean vectors) on one hand, and their numerical representations in some coordinate system as column matrices (also known as list vectors) on the other hand. Moreover, list vectors can of course have any dimension: $\\vec{a}\\in\\mathbb{R}^n$. In the coordinate system whose axes are aligned with unit vectors $\\vec{x}$, $\\vec{y}$, and $\\vec{z}$, the components of the arbitrary vector $\\vec{a}$ are given by $a_x=\\vec{x}'\\vec{a}$, $a_y=\\vec{y}'\\vec{a}$, and $a_z=\\vec{z}'\\vec{a}$; if the coordinate system is represented by the vector triad $\\tens{S}=\\left[\\vec{x}~\\vec{y}~\\vec{z}\\right]$, these components are given by the column matrix $\\tens{S}'\\vec{a}$ \\citep[cf.\\ Appendix~A in][]{murray1983}. This notation, although perhaps unfamiliar to many readers, provides a convenient and unambiguous framework for representing and transforming spatial vectors in different coordinate systems. For a vector-valued function $\\vec{f}(\\vec{x})$, $\\partial\\vec{f}/\\partial\\vec{x}'$ denotes the $\\dim(\\vec{f})\\times\\dim(\\vec{x})$ matrix whose $(i,j)$th element is $\\partial f_i/\\partial x_j$. Quaternions follow their own algebra (see Appendix~\\ref{sec:quaternions} for a brief introduction) and must not be confused with vectors/matrices; quaternions are therefore denoted by bold Roman characters ($\\quat{a}$). When taking a derivative with respect to the quaternion $\\quat{a}$, $\\partial x/\\partial\\quat{a}$ denotes the $4\\times 1$ matrix of derivatives with respect to the quaternion components; $\\partial x/\\partial\\quat{a}'$ is the transposed matrix. Tables of acronyms and variables are provided in Appendix~\\ref{sec:vars}. ",
        "conclusions": "\\label{sec:concl}  A fundamental part of the scientific data processing for the Gaia mission is the astrometric core solution, which will be run during and after the mission (ca.\\ 2013--2020) with successively larger datasets and eventually encompassing at least some 100~million primary sources. This solution is central to the performance of the mission as a whole, since it not only provides the astrometric results for these primary sources, but also the reference frame (in the form of the instrument attitude as a function of time) and the geometric calibration of the instrument, for use by a large number of other processes in the overall scientific reduction of the Gaia data.  In order to accomplish the astrometric core solution, a software system known as AGIS (Astrometric Global Iterative Solution) is being built within Coordination Unit 3 (CU3) of the Gaia Data Processing and Analysis Consortium (DPAC). As detailed in this paper, the necessary mathematical models and numerical algorithms are well understood, and have been developed with sufficient rigour to allow the potential accuracy of Gaia to be fully exploited. Most critical parts of this system have been implemented, and numerous test runs have demonstrated the theoretical validity of the global iterative approach, as well as its practical feasibility in terms of data management and computations. While a number of additional features will have to be included in the software before it can be considered ready for the flight data, and many more complications will undoubtedly be discovered during the actual analysis of these data, all the fundamental parts of AGIS are already in place.  In 2011, with roughly two years left until the launch of the satellite and time to mature the concepts and software presented here into a robust operational system, we have no reason to doubt that AGIS will be able to compute an accurate astrometric solution, consistent with the ambitious goals of the Gaia mission."
    },
    "1112/1112.0006_arXiv.txt": {
        "abstract": "We present the rest-frame $J$- and $H$- band luminosity function (LF) of field galaxies, based on a deep multi-wavelength composite sample from the MUSYC, FIRES and FIREWORKS survey public catalogues, covering a total area of 450 arcmin$^2$. The availability of flux measurements in the Spitzer IRAC  3.6, 4.5, 5.8, and 8 $\\mu$m channels allows us to compute absolute magnitudes in the rest-frame $J$ and $H$ bands up to $z=3.5$ minimizing the dependence on the stellar evolution models. We compute the LF in the four redshift bins $1.5<z<2.0$, $2.0<z<2.5$, $2.5<z<3.0$ and $3.0<z<3.5$. Combining our results with those already available at lower redshifts, we find that (1) the faint end slope is consistent with being constant up to $z=3.5$, with $\\alpha=-1.05\\pm0.03$ for the rest-frame $J$ band and $\\alpha=-1.15\\pm0.02$ for the rest-frame $H$ band; (2) the normalization $\\phi^*$ decreases by a factor of 4-6 between $z=0$ and $z\\simeq 1.75$ and by a factor of 2-3 between $z\\simeq 1.75$ and $z=3.25$; (3) the characteristic magnitude $M^*$ shows a brightening from $z=0$ to $z\\simeq 2$ followed by a slow dimming to $z=3.25$. We finally compute the luminosity density (LD) in both rest-frame $J$ and $H$ bands. The analysis of our results together with those available in the literature shows that the LD is approximately constant up to $z\\approx 1$, and it then decreases by a factor of 6 up to $z=3.5$. ",
        "introduction": "In the current concordance model, galaxies are the result of continuous mergers of dark matter halos driving baryonic matter assembly. In the last decades, simulations on halo occupation models have been able to quite accurately plot the formation of dark matter clusters. However, big uncertainties still remain at the time of translating dark matter haloes to what can actually be detected with our telescopes.  To this respect, the luminosity function (LF) of galaxies, i.e. the number density of galaxies per unit flux, is an extremely powerful tool to study the galaxy population and its evolution with cosmic time. Specifically, the analysis of the LF at different rest-frame wavelengths can give us information on different aspects of our present view of the Universe. The UV-optical LF allows for the study of the content and the evolution of the star formation rates with cosmic time. On the other hand the near infra-red (NIR) LF, being less sensitive to the absorption by dust and dominated by the light of older stars, is a better estimator of the overall stellar mass assembly of galaxies and of its rate of growth with time, revealing itself as a good test-bench for halo models.  The local NIR LF is still not yet well determined. Although a number of measurements have been derived so far, there seems to be uncertainties especially for the faint end slope $\\alpha$. Estimates of the slope $\\alpha$ range from $\\approx -0.8$ (\\citealt{bell2003}; \\citealt{eke2005}), to $\\alpha \\approx -1.2$ (\\citealt{jones2006}), with a median value around -1 (\\citealt{mobasher93}; \\citealt{glazebrook1995}; \\citealt{cowie96}; \\citealt{gardner97} and  \\citealt{szokoly98}; \\citealt{kochanek2001};  \\citealt{cole2001} and  \\citealt{hill2010}).  At even larger redshift, the LF determinations (most of which are done in the rest-frame $K_S$ band) still suffer from significant uncertainties (\\citealt{saracco2006}). The faint-end slope seems to be always compatible with $\\alpha=-1$ (\\citealt{drory2003}; \\citealt{pozzetti2003}; \\citealt{dahlen2005}; \\citealt{saracco2006}; \\citealt{cirasuolo2010}) although these measurements suffer from the large uncertainties given by the limits in the depth of the photometric catalogues available so far. There seems to be a general consensus however that the NIR LF does not significantly evolve to $z \\approx 1$ with respect to the local LF ( \\citealt{cowie96}; \\citealt{pozzetti2003}; \\citealt{drory2003}; \\citealt{feulner2003}; \\citealt{dahlen2005}).  A  brightening of the characteristic magnitude is instead found around $z \\approx 1.2-1.5$ together with a decrease of the normalization, decrease that is seen up to $z=3$ (\\citealt{saracco2006}; \\citealt{cirasuolo2010}).  In this paper we present the rest-frame $J$- and $H$-bands LFs and luminosity density (LD) of field galaxies, obtained from three deep photometric redshift surveys, namely MUSYC, FIRES and FIREWORKS, complemented by deep \\emph{Spitzer} 3.6, 4.5, 5.8, and 8 $\\mu$m data. As discussed in e. g. \\citet{berta2007}, the combination of the Planck spectral peak from low-mass stars, the minimum in the $H^-$ opacity in stellar atmospheres and the molecular absorptions in the spectra of cold stars produce a maximum for the emission in the rest-frame NIR portion of galaxy spectra located at $1.6\\mu$m (the so called \\emph{$1.6\\mu$m bump}). Furthermore, the AGN light can contribute significantly to the rest-frame $K$ band. Specifically, the contribution from the dust torus of the AGN can be in the rest-frame $K$ band a factor of 10 larger than in the rest-frame $J$ band, and a factor of ~4 larger than in the rest-frame $H$ band (e.g., \\citealt{polletta2008}). The adoption of the rest-frame $J$ and $H$ bands makes thus the measurement of the LFs and LDs less sensitive to potential dust-obscured AGN contamination compared to measurements of the LFs in the rest-frame $K_s$, yet allowing us to sample a wavelength range dominated by stellar emission and very little affected by obscuration by dust. The combination of depth and wavelength coverage in the mid-IR out to 8$\\mu$m allows us to directly probe the rest-frame $J$ and $H$ bands out to $z\\simeq 3.5$, relying more on observational data rather than on stellar population models, which are still significantly uncertain in the rest-frame NIR, due to different implementations of the TP-AGB phase \\citep{maraston2005,conroy2009}. The total surveyed area sums to 450 arcmin$^2$ with complete U-to-8$\\mu$m coverage, reducing thus the effects of cosmic variance, which we estimate to give on average a 15-20\\% contribution.  This paper is organized as follows. In section 2 we present the data sets used for this work together with a description on how we recover photometric redshifts and we select galaxies from the full sample. Section 3 presents the three methods adopted to estimate the LF and its associated uncertainties. In section 4 we present our results, while our conclusions are summarized in section 5.  Throughout this work, the adopted cosmology is $\\Omega_{\\Lambda}=0.7$, $\\Omega_m=0.3$ and $H_0=70$ Km/s/Mpc. All magnitudes are expressed in the AB system. ",
        "conclusions": "In the present work, we used a composite sample constructed from deep multi-wavelength publicly available photometric catalogues from the MUSYC, FIRES and FIREWORKS survey. The availability of \\emph{Spitzer}-IRAC data in the 3.6, 4.5, 5.8, and 8$\\mu$m channels allows us to robustly estimate the LFs and LDs in the rest-frame $J$  and $H$ bands with a minimum dependence on the SED templates up to $z=3.5$. Ours is the first measurement of the rest-frame $H$-band LF at $z>0$ to-date. We determined the LF with three independent methods, namely the $1/V_{\\rm max}$, the SWML ,and the STYML methods, finding that they agree well with each other. Uncertainties introduced by the cosmic variance were estimated using two distinct methods, one by \\citet{moster2010} and the other by \\citet{driver2010}. We find that, for our data, the two approaches broadly agree.  Our rest-frame $J$-band LF is consistent with previous determination by \\citet{saracco2006}, although the recovered Schechter parameters $M^*$ and $\\alpha$ are consistent only at the 2$\\sigma$ level. This might be due to the limited range in rest-frame magnitudes probed by the sample in \\citet{saracco2006} and by large errors due to cosmic variance given their small surveyed area. Our $J$-band LF is consistent also with the LF measured by \\citet{pozzetti2003} at $z\\sim1.5$.  We analyzed the evolution with redshift of the Schechter function parameters, making full use of the data available from the literature. We found that the faint end slope $\\alpha$ of the LF is nearly constant over the whole redshift range, with $\\alpha_J=-1.05\\pm0.03$ and $\\alpha_H=-1.15\\pm0.02$ in the $J$ and $H$ bands, respectively. The characteristic density $\\phi^*$ decreases by a factor of $\\approx 6$ from $z\\sim0$ to $z=1.75$, and by a factor of $\\approx 3$ from $z\\sim 1.75$ to $z=3.25$. We introduced a parameterization based on an exponential form for the evolution of $\\phi^*$ as a function of $z$. The fit of this function to the available data shows good agreement, especially for the rest-frame $J$ band, where more data from the literature are available in the redshift range $z \\in[0, 1]$, complementing our measurements at $z>1.5$.  The characteristic magnitude $M^*$ is found to brighten from $z=0$ to $z\\sim2$ by $\\sim0.8$~mag, whereas $M^*$ gets fainter with increasing redshift at $z \\gtrsim 2$. We adopted a \\cite{schechter1976}-like expression for its description, resulting in a good representation of the observed evolution.  We computed the LD in the rest-frame $J$ and $H$ bands, using the Schechter parameters previously determined. The LD is nearly constant up to $z\\approx 1$ and decreases as a power-law by a factor of $\\approx 6$ from $z\\approx1$ to $z=3.25$.  We compared the evolution with redshift of the LFs and LDs in the rest-frame $J$ and $K$ bands. The Schechter parameters $\\phi^*$ and $M^*$ in the rest-frame $K$ band show different evolutions with redshift with respect to the $J$ band, although these differences are only marginally significant at the 2-$\\sigma$ level. Specifically, the decrease of the characteristic density with redshift appears faster in the $K$ than in the $J$ band at $z \\gtrsim 1.5$. The characteristic magnitude in the rest-frame $K$ band brightens with redshift over the whole redshift interval $0<z<4$, whereas $M^*_{\\rm J}$ gets brighter from $z=0$ to $z\\sim2$, and then slowly gets fainter out to $z=3.5$ (although a constant value of $M^*_{\\rm J}$ at $z>1.5$ is consistent within the uncertainties). Most of these differences cancel out when computing the LDs, with similar evolutions with redshift of the rest-frame $J$- and $K$-band LDs out to $z\\sim2.5$. Evidence for a faster decrease with increasing redshift of the rest-frame $K$-band LD at $z\\gtrsim2.5$ seems present. However, these differences are significant only to a 95\\% level (2 $\\sigma$). Large errors at $z>2$ in both the $J$ and $K$ bands prevent to firmly assess differences between the evolution with redshift of the $J$- and $K$-band LFs and LDs.  In order to further constrain the rest-frame $J$- and $H$-band LFs, a larger area is needed to better probe the bright end and reduce the impact of field-to-field variations and low number statistics. Better photometric redshift estimates are also needed to improve the LF measurements at the bright end, by, e.g., reducing the impact of catastrophic outliers. The recently publicly released NEWFIRM Medium-Band Survey \\citep{whitaker2011} will provide the dataset to significantly improve on all of these aspects. Ongoing very deep ground-based surveys, such as the Ultra-VISTA survey\\footnote{http://www.eso.org/sci/observing/policies/PublicSurveys/scienePublicSurveys.html}, and space-based surveys with WFC3 on the \\emph{Hubble Space Telescope}, such as the CANDELS \\citep{grogin2011,koekemoer2011} and 3D-HST (Brammer et al. 2011, in prep.) surveys, will allow for much improved constraints of the faint end of the rest-frame NIR LF and of the contribution of low-luminosity galaxies to the total NIR LD."
    },
    "1112/1112.0577.txt": {
        "abstract": "By logging encounters between planetesimals and planets we compute the distribution of encounters in a numerically integrated two planet system that is migrating due to interactions with an exterior planetesimal belt. Capture of an irregular satellite in orbit about a planet through an exchange reaction with a binary planetesimal is only likely when the binary planetesimal undergoes a slow and close encounter with the planet. In our simulations we find that close and slow encounters between planetesimals and a planet primarily occur with the outermost and not innermost planet. Taking care to consider where a planet orbit crossing binary planetesimal would first be tidally disrupted, we estimate the probability of both tidal disruption and irregular satellite capture. We estimate that the probability that the secondary of a binary planetesimal is captured and becomes an irregular satellite about a Neptune mass outer planet is about 1/100 for binaries with masses and separations similar to transneptunian planetesimal binaries. If young exoplanetary debris disks host a binary planetesimal population then outwards migrating outer planets should host captured irregular satellite populations. We discuss interpretation of emission associated with the exoplanet Fomalhaut b in terms of collisional evolution of a captured irregular satellite population that is replenished due to planetary migration.  %The probability of satellite capture %is only weakly dependent on the outer planet mass and the ratio of the two planet masses. ",
        "introduction": "The majority of irregular satellites about the gas giant planets in our Solar system are in eccentric and high inclination orbits  suggesting that likely originated in heliocentric orbit and were later captured into orbit within a planet's Hill sphere (see the review by \\citealt{jewitt07}).  Irregular satellite populations could have at one time been a source of heavy dust production \\citep{bottke10}.  If irregular satellite populations exist in distant planetary systems then exoplanets could be discovered by  detecting scattered stellar light from their associated dust clouds \\citep{kennedy11}. %The planet Fomalhaut b has been detected in two visible bands only  \\citep{kalas08}. The color of the visible light from the detected object, Fomalhaut b,  in orbit about the nearby star Fomalhaut, is consistent with that of reflected light from the star \\citep{kalas08} and so does not originate from the planet itself but could arise from a planetary ring system (as discussed by \\citealt{kalas08}),  dust cloud associated with irregular satellites (as proposed by \\citealt{kennedy11}) or circumplanetary disk akin to the birth site of the Galilean satellites (e.g., as explored by \\citealt{canup02}).  We focus on extrasolar planets that are in proximity to  a debris disk as this disk could host planetesimals that could be captured by the planets in the system,  becoming irregular satellites. Neither the Fomalhaut system or the HR8799 system are likely to have strongly scattered or thick planetesimal disks, the Fomalhaut system \\citep{kalas05,kalas08} because the disk is thin (with disk aspect ratio only $h/r \\sim 0.013$ \\citep{quillen07}, and the HR8799 system \\citep{marois10} because the planets are closely spaced and the time until the planet's cross each others orbits  is probably short \\citealt{god09,fabrycky10,moore11}) implying that the planets have not yet experienced close encounters. The Fomalhaut and HR8799 systems can be considered to be analogs to our solar system prior to the late heavy bombardment epoch of our solar system that involved close approaches between the giant planets (the `Nice' model; \\citealt{nice}) and subsequent scattering of the planetesimal disk.  In our Solar system, the irregular satellites were likely captured during planet/planet close approaches during the late heavy bombardment era  \\citep{nes07}. %The irregular satellite population since then has evolved due to collisions \\citep{bottke10}. However, the recent discovery of a large transneptunian binary population \\citep{noll08} has motivated study of the exchange reaction mechanism for irregular satellite capture \\citep{agnor06,vok08,philpott10,nogueira11}. Irregular satellites in orbit around the giant planets prior to the epoch of late-heavy bombardment were likely lost during close approaches between planets \\citep{nes07}. Little is known about possible irregular satellite populations that could have existed prior to the late heavy bombardment era.  Both exoplanets Fomalhaut b and HR8799 b are located just interior to dusty disks that are generated from collisions of planetesimals and so are called debris disks \\citep{marsh05,su09}. Planetesimals originating from these disks can be scattered by the outer planet into orbits that cross the orbits of interior planets. The interaction between planetesimals and planets allows the outer planet to migrate outwards into the planetesimal disk via the exchange of angular momentum between planets and planetesimals \\citep{fip84,ida00}.  As the outermost planet migrates outwards the population of planet crossing planetesimals can be continually replenished.  \\citet{kennedy11} showed that a dust producing irregular satellite population could be long lived, however they did not explore mechanisms for irregular satellite capture about exoplanets. If the planetesimal disks in these systems contain  binary planetesimals, then the exchange reaction mechanism for irregular satellite capture \\citep{agnor06} is a viable mechanism for the capture of irregular satellites about planets from the population of planet-orbit crossing planetesimals. Currently, it is estimated that binaries account for 30\\% of transneptunian or Kuiper belt objects (KBOs) with inclinations $<5^\\circ$,  and $\\sim5$\\% of the rest of the KBOs \\citep{noll08}. If binary planetesimals are common in extrasolar planetesimal disks then the migration of an outer planet into a planetesimal disk would cause binary planetesimals to cross the orbits of giant planets in the system, facilitating the capture of irregular satellites.  A planetesimal binary is disrupted inside the Hill sphere of a planet when the distance to the planet is sufficiently small that the planet's tidal force overcomes the gravitational attraction of the binary. The disruption can leave an object in a bound orbit around the planet  if the velocity of the binary components is sufficiently large and oriented so as to decrease the energy of the lower mass planetesimal when the binary is disrupted \\citep{agnor06,vok08}. The probability of irregular satellite capture via binary exchanges depends both on the statistics of close approaches and their energies \\citep{vok08,philpott10,nogueira11}. In this paper we consider the statistics of close approaches between planetesimals and planets in systems containing two migrating planets.   Our goal is to better understand mechanisms for irregular satellite capture that might operate in young exoplanetary systems and prior to the epoch of late heavy bombardment in our Solar system.  %Think about Peretz work!  %\\citet{agnor06} suggested that Triton was captured by Neptune in an exchange reaction between a binary planetesimal pair and Neptune.  However \\citet{vok08} showed using initial conditions guided by the ``Nice'' model following the epoch of planet/planet encounters that the probability of irregular satellite capturing binary exchange reactions was low and that the orbital properties of the predicted irregular satellite populations differed from those observed. \\citet{philpott10} reconsiders the binary exchange reaction mechanism primarily for Jupiter as there may not have been any close approaches between Jupiter  and the other planets \\citep{nice}. ",
        "conclusions": "Using numerical integrations of planets interacting with a planetesimal belt we have explored the statistics of close encounters of planetesimals with planets during planetary migration. We have focused exclusively on the problem of a two planet system in proximity to a cold (in terms of velocity dispersion) planetesimal belt  exterior to both planets that induces planetary migration. If the planetesimal belt contains binary planetesimals then close encounters with planets can disrupt these binaries.   To predict the location of tidal disruption (at what radius and which planet), we have taken into account the order of close encounters with planets. We find that the velocity distribution (for velocities with respect to the planet) of tidally disrupting encounters differs for the two planets.  Lower velocity encounters are only likely with the outer planet. As a consequence we predict that the probability that irregular satellites are captured due to binary exchange reactions is higher, by about two orders of magnitude, with an outer planet than an inner one.   We estimate that planetesimal binaries similar to those in the Kuiper belt would have a probability of about 1/100 of disrupting and leaving a bound object about an outer migrating planet.  Our numerical estimate is consistent with that recently estimated by \\citet{nogueira11} for the capture of Triton into orbit about Neptune. In migrating systems we infer that only the outer planet could gain a significant population of captured irregular satellites and other scenarios would be needed to account for irregular satellite populations around inner giant planets.  We note that the simulations considered here were restricted by numerical considerations to a narrow range of planet and disk masses, only considered cold and exterior planetesimal disks, only considered two giant planet systems and situations lacking planetary encounters.   Future studies can expand the types of systems numerically integrated to predict the probability of satellite capture for more diverse systems.  We estimate that the tidal disruption radius in units of the planet's Hill radius is approximately set by the binary separation in units of its own Hill radius.  Thus if optical emission from exoplanets is associated with a dust cloud generated by irregular satellites, as explored by \\citet{kennedy11}, then there is a relation between the size of the dust cloud, the planet mass and the properties of the primordial binary planetesimal population. We estimate that capture is allowed only if the incoming binary velocity, with respect to the planet is only a few times the velocity of the binary orbit prior to disruption. We have used this estimate and our numerically measured encounter distributions to estimate the probability that the secondary of a binary planetesimal could be captured into orbit about a planet.   To improve upon our derived capture probabilities it is likely that the four body problem (binary planetesimal encounter with a planet in orbit about a star) must be integrated.  \\citet{kennedy11} suggested that irregular satellites can be captured if Fomalhaut b crosses into the planetesimal belt.  Here we explore a related scenario. We consider  Fomalhaut b not crossing into the planetesimal belt, but  migrating outwards and capturing planetesimals due to binary exchange reactions with binaries that have been perturbed into planet orbit crossing trajectories.  A crude estimate of the mass in irregular satellites captured per unit time suggests that the mass captured per unit time could be larger than the estimated dust production rate if a collisional cascade operates among the irregular satellite population. Consequently there could be sufficient mass captured as irregular satellites to fuel such a collisional cascade.    In a migration scenario, we find that the outer planet is likely to gain two orders of magnitude more irregular satellites than inner planets.  It is tempting to consider this scenario as accounting for the detection of Fomalhaut b but not other, possibly more massive planets that could exist interior to  it in the system \\citep{kenworthy09}. %\\citet{bottke10} find that the irregular satellite populations in our solar system are ``the most collisionally %evolved populations in the solar system.'' %It is interesting to consider that irregular satellites %around exoplanets  might  be detectable through their associated dust production. Further study is needed to determine whether a migration scenario for Fomalhaut b that included irregular satellite capture would be consistent with the $\\sim 0.1$ eccentricity of the dust ring and planet and the location and shape of the ring edge.   %todo %Gravitational focusing jujitsu xxxx %Scott Gaudi: estimate rate of capture of KB binaries to Neptune now based on Centaur numbers and lifetimes. This was essentially already done by Nogeuira et al.  \\vskip 0.3 truein Acknowledgements.   This work was in part supported by NSF through award AST-0907841. We thank NVIDIA for gift of graphics cards.  We thank Scott Gaudi for interesting comments."
    },
    "1112/1112.5765_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "We have discussed the increased  stability of strange quark matter given  improved understanding of strange quark mass and clustering feature of cold quark matter. Our result implies that nuclear matter could be converted to strange quark matter by compression, e.g. in relativistic heavy ion  collisions, or more likely, in the center of neutron stars.  Many laboratory searches for small strangelets have produced no positive signature to date~\\cite{Sandweiss:2004bu}. We hope that  persistence of impact formations over geologic timescales which turn rocky planets into integrating detectors will  open the search for large strangelets. \\\\[3mm]  \\noindent This work was supported by  the grant from the U.S. Department of Energy, DE-FG02-04ER41318."
    },
    "1112/1112.1599_arXiv.txt": {
        "abstract": "We present Sunyaev-Zel'dovich (SZ) effect observations of a sample of 25 massive relaxed galaxy clusters observed with the Sunyaev-Zel'dovich Array (SZA), an 8-element interferometer that is part of the Combined Array for Research in Millimeter-wave Astronomy (CARMA). We perform an analysis of new SZA data and archival \\chandra\\ observations of this sample to investigate the integrated pressure---a proxy for cluster mass---determined from X-ray and SZ observations, two independent probes of the intra-cluster medium. This analysis makes use of a model for the intra-cluster medium introduced by Bulbul (2010) which can be applied simultaneously to SZ and X-ray data.  With this model, we estimate the pressure profile for each cluster using a joint analysis of the SZ and X-ray data, and using the SZ data alone. We find that the integrated pressures measured from X-ray and SZ data are consistent. This conclusion is in agreement with recent results obtained using \\textit{WMAP} and \\textit{Planck} data, confirming that SZ and X-ray observations of massive clusters detect the same amount of thermal pressure from the intra-cluster medium. To test for possible biases introduced by our choice of model, we also fit the SZ data using the universal pressure profile proposed by Arnaud (2010), and find consistency between the two models out to $r_{500}$ in the pressure profiles and integrated pressures. ",
        "introduction": "The Sunyaev-Zel'dovich (SZ) effect \\citep{sunyaev1972} is a spectral distortion of the cosmic microwave background (CMB) caused by the scattering of CMB photons off the hot electrons of the intra-cluster medium (ICM). Over the past two decades, SZ observations with both single-dish and interferometric instruments have become routine \\citep[e.g.,][]{birkinshaw1991,carlstrom1996,holzapfel1997,carlstrom2002}, and SZ surveys are now producing catalogs of newly-discovered clusters out to high redshift \\citep[][]{vanderlinde2010,marriage2011,williamson2011,ade2011}. SZ measurements are complementary to the X-ray measurements which have long been used to study clusters, but only in recent years have sufficiently large samples of objects been observed in the SZ to permit a rigorous comparison between these two techniques \\citep[e.g.,][]{reese2002,bonamente2006,laroque2006}.  The SZ effect causes a perturbation $\\Delta T$ of the CMB temperature $T_{CMB}$ given by \\begin{equation} \\frac{\\Delta T}{T_{CMB}} = f(x) \\int \\! \\sigma_T n_e \\frac{k T_e}{m_e c^2} d\\ell = f(x) y , \\label{eq:compy} \\end{equation} where $f(x)$ is the frequency dependence of the SZ effect \\citep[e.g.,][]{laroque2006}; $\\sigma_T$ is the Thomson cross-section; $n_e$, $T_e$ and $m_e$ are the number density, temperature, and mass of the electrons, respectively; $k$ is the Boltzmann constant; $c$ the speed of light; and the integral is along the line of sight $\\ell$. At a given frequency, the amplitude of the effect depends linearly upon the Compton $y$ parameter, which is defined implicitly in Equation~\\ref{eq:compy}. Note that the $y$ parameter is proportional to the ICM pressure integrated along the line of sight.  At frequencies below 218~GHz, the SZ effect causes  a decrement in the CMB temperature in the direction of the cluster. The integral of $y$ over the solid angle $\\Omega$ subtended by the cluster, known as the (cylindrical) integrated Compton $y$ parameter $Y_{\\rm cyl}=\\int \\! y \\, d \\Omega$, is expected to be a good proxy for cluster total mass since it traces the thermal energy content of the cluster gas.  Alternatively, the Compton $y$ parameter can be integrated spherically, \\begin{equation} Y_{\\rm sph}(r_{500}) = \\frac{1}{D_A^2} \\left( \\frac{k \\sigma_T}{m_e c^2} \\right)\\int  n_e T_e dV \\label{eq:Y} \\end{equation} where the volume $V$ is a sphere centered on the cluster and $D_A$ is the angular diameter distance.  X-ray data can also be used to constrain the density and temperature---and thus the pressure---of the ICM.  Over the past decade, several groups have investigated the consistency between X-ray and SZ pressure measurements. Early measurements of the SZ signal from \\wmap\\ by, e.g., \\cite{lieu2006}, \\cite{bielby2007}, detected an SZ signal at a lower level than expected. \\cite{atrio-barandela2008} showed that the isothermal beta model leads to an electron pressure profile that exceeds the measured values at large radii by a factor of few, and that the baryon profile is consistent with a model based on the \\cite{navarro1997} matter profile. \\cite{diego2010} also showed that contamination by compact radio sources may have led to underestimates of the SZ effect flux decrements in the \\wmap\\ data. More recent comparisons of \\chandra\\ X-ray data to stacked data from \\wmap\\ and \\planck\\ \\citep{melin2011,Planck2011-sz} have found consistency between SZ and X-ray measurements for large samples of clusters. \\cite{komatsu2011} also analyzed a sample of massive nearby clusters individually resolved by \\wmap, again finding good agreement with X-ray predictions.  In this paper, we present Sunyaev-Zel'dovich Array (\\sza) observations of the \\cite{allen2008} sample of massive relaxed galaxy clusters, together with archival \\chandra\\ X-ray observations that are available for all clusters in this sample. The sensitivity and resolution of our data permit us to measure the pressure profile and the integrated pressure out to $r_{500}$---the radius within which the average cluster density is 500 times the critical density---for each cluster individually, without the need to resort to scaling relations between the X-ray luminosity and mass \\citep[as was done by][for example]{melin2011,Planck2011-sz}. We use the \\citet{bulbul2010} model of the cluster pressure, density, and temperature. Since this model has a consistent parameterization for all thermodynamic quantities, it is especially well-suited for joint X-ray and SZ analysis.  As a cross-check against model-dependent biases, we also fit the SZ data using the model of \\citet{arnaud2010} based on the numerical simulations of \\cite{nagai2007b}.  We find consistency to within our measurement uncertainties both between the X-ray and SZ measurements, and between the \\citet{bulbul2010} and \\citet{arnaud2010} models.  The paper is structured as follows: Section~\\ref{sec:sample} describes our observations and our sample, Section~\\ref{sec:analysis} presents our joint analysis technique, Section~\\ref{sec:Y} describes our method of measuring the integrated $Y_{\\rm sph}(r_{500})$ parameter (defined in Equation~\\ref{eq:Y}), Section~\\ref{sec:discussion} presents and discusses our results, and our conclusions are presented in Section~\\ref{sec:conclusions}. ",
        "conclusions": ""
    },
    "1112/1112.4719_arXiv.txt": {
        "abstract": "The structure of compact star core filled by gluon matter plasma is investigated within the fluid-like QCD framework. The energy momentum tensor, density and pressure relevant for gluonic plasma having the nature of fluid bulk of gluon sea are derived within the model. It is shown that the model provides a new equation of state for perfect fluid with only a single parameter of fluid distribution  $\\phi(x)$. The results are applied to construct the equation of state describing the gluonic plasma dominated compact star core. The equations of pressure and density distribution are solved analytically for small compact star core radius. The phase transition of plasma near the core surface is also discussed. ",
        "introduction": "\\label{sec:intro}  Recent experiments in the last decades on relativistic nuclear collisions shed light on the phenomena of hot plasma formed by dense quarks and gluons. Those experiments suggest that the quark gluon matter behaves more like a deconfined quark-gluon plasma (QGP) liquid \\cite{star,phenix,shuryak}.  This fact immediately encourages some models based on either the (relativistic) hydrodynamic \\cite{bouras,romatschke}, or pure Quantum Chromodynamics (QCD) approaches \\cite{Alford}. Within the hydrodynamic framework, the plasma is dominated by either quark \\cite{romatschke} or gluon matter \\cite{bouras}. In particular, in the quark matter dominated plasma dissipative ideal hydrodynamics has been used to fit some experimental data at high energy heavy ion program at the Relativistic Heavy Ion Collider (RHIC) \\cite{rhic}. The successful fit requires the models to take into account very small value of the ratio shear viscosity over entropy \\cite{teaney,huovinen,kolb,kolb2,hirano,baier}. However the puzzle must still be confirmed by the next coming experiments at the Large Hadron Collider (LHC) \\cite{lhc}. While the gluon matter plasma is motivated by the discoveries of jet-quenching in heavy-ion-collision at RHIC indicating the shock waves in form of March cones \\cite{adams,adare}.  On the other hand, in pure QCD, QGP is described as a quark soup before hadronization which is a phase of QCD, and exists at extremely high temperature and/or density. It is argued that this phase consists of almost free quarks and gluons. Therefore, the phase transition from the deconfined QGP to the hadronic matters or vice versa gets particular interest in this approach. It unfortunately turns out to the many body problems with large color charge which cannot be calculated analytically using perturbation. As a result, the main theoretical tools to explore QGP within QCD is lattice gauge theory. The lattice calculation predicts that the phase transition occurs at approximately 175 MeV \\cite{gottlieb,petreczky}. Many works on both analytical and numerical calculations have been done at a sophisticated level providing some interesting predictions like color glass superconductivity (CGS), color-flavor locked (CFL) phase and so on.  Anyway, describing the QGP as either quark or gluon matter dominated plasma within hydrodynamic or pure QCD should be understood as conjectures in the world of QGP. The important point is, since the QGP contains many quark-anti-quarks and gluons, it is considerable to treat it using the well-established QCD. At the same time from hydrodynamic point of view, the experiments strongly suggest that QGP behaves like a fluid. In particular, in the scenario of viscous gluonic plasma this is required to form and propagate shock waves \\cite{adler2,adams2,back}. Moreover it dissolves into an almost perfect dense fluid \\cite{zajc}. Therefore it is plausible to describe it as a strongly interacting fluid system. In this sense, there are approaches based on unifying or hybridizing the charge field with flow field \\cite{heinz,holm,choquet,blaizot,bistrovic,manuel}. Recently, some works have constructed the models in a lagrangian with certain non-Abelian gauge symmetry to the matter inside the fluid \\cite{marmanis,sulaiman}.  In this paper, we follow the scenario of gluon matter plasma as suggested by hydrodynamic approach, while it is still governed by QCD. In this framework, the strongly interacting system of QCD is considered as a macroscopic fluid system rather than a result of many particles interactions. It is shown that one can derive the total energy momentum tensor of gluon plasma within the model. Further, the equation of state (EoS) relevant for plasma dominated compact star core interiors is constructed. In the model, the density ($\\rho$) and pressure ($p$) are determined and related each other through the fluid field distribution $\\phi(x)$. Hence, the EoS can be solved through the ordinary differential equation (ODE) of $\\rho$ to show the phase transition within the model.  The paper is organized as follows. First we briefly introduce the underlying model of gauge invariant fluid lagrangian and discuss the relevant physical scale and region within the model. Then, the energy momentum tensor, density and pressure in the model are derived and investigated. Subsequently it is followed by constructing the relevant EoS for compact star core interiors. Before summarizing the paper, the analytical solution for the pressure and density distributions of small compact star core radius is given and discussed. ",
        "conclusions": "The QGP, in particular gluon matter, dominated compact start core interior has been investigated using the fluid QCD lagrangian. From the lagrangian one can derive the energy-momentum tensor and subsequently the density and pressure distributions from the first principle. In the model, both density and pressure are related each other and form an extended EoS of perfect fluid through the field distribution $\\phi(r)$. The relation leads to the linear EoS known in the conventional cosmological model, but allowing the EoS parameter $w$ to be a function of temperature $T$ and core radius $r$. The analytical solutions for the total pressure and density distribution equations are also given for small normalized compact star core radius, $\\rp \\ll 1$.  The model describes the dense gluon plasma inside the compact star core as a bulk dynamics of gluonic matters before the cooling phase near the compact star core surface as depicted in Fig. \\ref{fig:d-r}. It should be noted that the figure provides a general contour for pressure or density distribution inside the core. This suggests that near its surface the hadronization is getting involved to form surfaces of small compact star cores,  On the other hand, the phase transition for certain size of compact star core in term of temperature changing is shown in Fig. \\ref{fig:d-t}. It can be observed that the phase transition occurs at the scale of hadronization around the core surface, while it keeps fluidity as the temperature is getting higher in the core interiors."
    },
    "1112/1112.1416_arXiv.txt": {
        "abstract": "The scale radius of dark matter halos is a critical parameter for specifying the density distribution of dark matter, and is therefore a fundamental parameter for modeling galaxies.  We develop here a novel, observationally motivated probe to quantitatively infer its value.  We demonstrate that disturbances in the extended atomic hydrogen gas disks of galaxies can be used to infer the scale radius of dark matter halos.  Our primary metric is the phase of the $m=1$ mode of the disturbance in the outskirts of the gas disk, which we take to be produced by a tidal interaction.  We apply the method to the Whirlpool Galaxy, which has an optically visible satellite.  We explore potential degeneracies due to orbital inclination and initial conditions and find our results to be relatively insensitive to these considerations.   Our method of tracing the dark potential well through observed disturbances in outer gas disks is complementary to gravitational lensing. ",
        "introduction": "The cold dark matter (CDM) paradigm of structure formation is successful at recovering the basic skeletal structure of the universe -- the large-scale distribution of galaxies (Springel et al. 2006).  However, the agreement between theory and observation is less secure when this model is applied to galactic (and sub-galactic) scales.   While the existence of dark matter halos in galaxies was observationally inferred in the 1970's (Rubin et al. 1977; Bosma 1978), there are few observational probes that can be used to infer the details of the distribution of dark matter.  Simulations have not yet been able to resolve the mitigating effects of gas cooling (Gnedin et al. 2004) versus outflows (Governato et al. 2010), which can act to increase (decrease) the distribution of dark matter in galaxies.  In the coming decades, the CDM paradigm will be increasingly vetted against detailed observations of the innards of galaxies that will delve even more deeply into galaxy evolution.  Navarro, Frenk \\& White (1996; 1997; NFW) found that dark matter halos formed through dissipationless hierarchical clustering have an universal density profile.  Due to the absence of scale in this heirarchical structure formation scenario, dark matter haloes are scaled versions of each other, with the total mass being the scaling parameter.  The NFW density profile is given by: \\begin{equation} \\rho(r) = \\frac{\\delta_{c}\\rho_{c}}{(r/R_{s})\\left(1+r/R_{s}\\right)^{2}} \\end{equation} where $\\rho_{c}$ is the critical density for closure of the universe, and $\\delta_{c}$ is a characteristic overdensity for the halo.  The scale radius, $R_{s}=R_{200}/c$ is a characteristic radius, and $c$ is the concentration parameter that relates this characteristic radius and $R_{200}$, which is the radius at which the mean enclosed dark matter density is 200 times the critical density.   The characteristic overdensity of the halo, $\\delta_{c}$, is simply a function of the concentration parameter $c$ (NFW).  The NFW profile differs from the Hernquist (1990) profile in its asympotic behavior, scaling as $r^{-3}$ rather than $r^{-4}$ for $r/r_{s} \\gg 1$.  The scale radius of the dark matter halo along with the concentration parameter specify the density distribution for a given halo mass in the NFW formalism.  Currently, these parameters are prescribed from large cosmological N-body simulations (Bullock et al. 2001; Maccio et al. 2008).  It is not presently clear how we can determine these values for individual galaxies.  What we seek to do here is to develop an independent, observationally motivated probe of the scale radius of dark matter halos, and apply it to the well-known Whirlpool Galaxy (M51) to infer its scale radius.  The  density switches its behavior at the scale radius, from $r^{-1}$ for $r \\ll R_{s}$ to $r^{-3}$ for $r \\gg R_{s}$, and passes through the transitional $r^{-2}$ region for $r \\sim R_{s}$.  Thus, the scale radius of dark matter halos is a critical parameter for specifying the density distribution of dark matter in galaxies, and therefore a fundamental parameter for modeling galaxies.  The extended atomic hydrogen (HI) disks of galaxies provide an unique probe of galaxy evolution.  They are ideal tracers of tidal interactions with satellites and the galactic gravitational potential well.  We recently developed a novel method whereby one can infer the mass, and relative position (in radius and azimuth) of satellites from analysis of observed disturbances in outer gas disks, without requiring knowledge of their optical light (Chakrabarti \\& Blitz 2009, henceforth CB09; Chakrabarti \\& Blitz 2011, henceforth CB11; Chakrabarti, Bigiel, Chang \\& Blitz 2011, henceforth CBCB; Chang \\& Chakrabarti 2011, henceforth CC11).   We applied this method to M51 and inferred that its satellite has a mass  one-third that of the primary galaxy, with a pericentric approach distance of 15$~\\rm kpc$.  We found these estimates to be corroborated by observations (Smith et al. 1990) and recent simulation studies (Salo \\& Laurikainen 2000; Dobbs et al. 2010).  Moreover, at the time when our simulations achieve the best-fit to the HI data, the azimuth of the satellite in the simulations agrees closely with its observed location.  CBCB note that the derivation of these numbers is uncertain at the factor of two level due to variations in the initial conditions of the simulated M51 galaxy, orbital inclination and orbital velocity of the satellite.  We call this method \"Tidal Analysis\".  We frame the problem here in the following way -- for galaxies like M51, for which we have demonstrated that our Tidal Analysis method works, can we determine the scale radius of the dark matter halo of M51? We take the satellite mass and pericentric distance from our earlier studies as an input.  It is worth noting that such a procedure has viable and immediate applicability.  The THINGS survey produced HI maps of local spiral galaxies (Walter et al. 2008), many of which display disturbances in the outskirts (Bigiel et al. 2010).   Attempting to constrain the dark matter halo in this way is similar in spirit to earlier work on analysis of stellar tidal tails by Mihos, Dubinski \\& Hernquist (1998).  If the satellite mass and pericentric distance of the satellites that produce these disturbances can be characterized, as for M51 (CBCB), then as we show below, the phase of the $m=1$ mode directly yields the scale radius of the dark matter halo.  This method of probing the dark matter mass distribution is independent of the stellar light and is complementary to strong gravitational lensing, which primarily probes regions interior to the Einstein radius, $\\sim 10~\\rm kpc$ in spirals (Wright \\& Brainerd 2000; Treu \\& Koopmans 2002).  The weak lensing signal has also been exploited by stacking the surface density contrast to produce mean density profiles for clusters (Sheldon et al. 2008) and total galaxy masses (Mandelbaum et al. 2006).  This paper is organized as follows:  in \\S 2, we demonstrate that the phase of the $m=1$ mode can be used to quantitatively infer the scale radius of dark matter halos and apply the method to M51.  We discuss caveats and future work in \\S 3, and conclude in \\S 4. ",
        "conclusions": "$\\bullet$  We find that the shape of disturbances in the fragile, extended HI disks of galaxies reflects the underlying density profile of the dark matter halo in spiral galaxies.  We employ the phase of the $m=1$ mode as our primary metric of comparison to observed HI data.  Building on prior results from CBCB where the satellite mass and pericentric approach distance was determined for M51, we find that the scale radius of the dark matter halo can be determined from the phase of the $m=1$ mode.  $\\bullet$ The three regimes of interest for the phase of the $m=1$ mode that delineate the variation of the scale radius and how it is reflected in the phase are:  $d/dr(\\phi_{1})<0$ for $r/R_{s} < 1$, a transitional region at $r /R_{s}\\sim 1$, and $d/dr(\\phi_{1})>0$ for $r/R_{s} >1$ til the edge of the gas disk.   Long baseline HI observations that can reveal this transitional region between the $r^{-1}$ shallow profile (that would produce more loosely wrapped spirals) and the steeper $r^{-3}$ profile (that would produce more tightly wrapped spirals) will allow us to quantitatively infer the scale radius of dark matter halos of $\\it{specific~spiral~galaxies}$.  The slight uncertainty in the phase that arises due to variations in initial conditions and orbits does not preclude us from inferring the value of the scale radius to within a factor of $\\sim 1.5$.  We have discussed preliminary work that suggests that other effects (such as intrinsic processes arising from non-spherical halos) will not significantly affect our results here.  $\\bullet$ Application of this method to the Whirpool Galaxy yields a scale radius of $17~\\rm kpc$, which is consistent with the expected range of scale radii of $\\sim 10^{12} M_{\\odot}$ halos in dissipationless cosmological simulations.  $\\bullet$  We have also demonstrated that for $r/R_{s} > 1$, our results depend essentially on the scale radius and not on the concentration parameter.  These results suggest that we can utilize observed disturbances in the extended HI disks of galaxies to characterize the density profile of the dark matter halo in spiral galaxies.  In the future, we will apply this method to a large sample of spirals to obtain a statistical determination of the scale radii of spiral galaxies in the Local Volume.  \\bigskip \\bigskip"
    },
    "1112/1112.5145_arXiv.txt": {
        "abstract": "We report detection of thermal emission from the exoplanet WASP-19b at 3.6, 4.5, 5.8 and 8.0 \\micron. We used the InfraRed Array Camera on the Spitzer Space Telescope to observe two occultations of WASP-19b by its host star. We combine our new detections with previous measurements of WASP-19b's emission at 1.6 and 2.09 \\micron\\ to construct a spectral energy distribution of the planet's dayside atmosphere. By comparing this with model-atmosphere spectra, we find that the dayside atmosphere of WASP-19b lacks a strong temperature inversion. As WASP-19 is an active star (\\rhk\\ = $-4.50 \\pm 0.03$), this finding supports the hypothesis of \\citet{2010ApJ...720.1569K} that inversions are suppressed in hot Jupiters orbiting active stars. The available data are unable to differentiate between a carbon-rich and an oxygen-rich atmosphere. ",
        "introduction": "By observing the occultation of an exoplanet by its host star, we can measure the emergent flux from the planet's dayside atmosphere. Such measurements are challenging due to the low planet-to-star flux ratio (for the best cases, typically a few tenths of one per cent in the near infrared) and sources of noise both instrumental and stellar in origin \\citep[e.g.: ][]{2008ApJ...673..526K, 2011MNRAS.416.2096S}. To date, most such measurements \\citep[including the first,][]{2005ApJ...626..523C,2005Natur.434..740D} have been made with the {\\it Spitzer Space Telescope}, though ground-based facilities are now an essential complement \\citep[e.g.:][]{2009A&A...493L..35D,2009A&A...493L..31S}.  With a single occultation we can measure the corresponding brightness temperature and determine the eccentricity of the planet's orbit \\citep[e.g.][]{2005ApJ...626..523C}, which is necessary for the accurate determination of the stellar and planetary radii in transiting systems \\citep{2012MNRAS.422.1988A} and is important for studies of the formation and tidal inflation of short-period, giant planets \\citep[e.g.:][]{2011ApJ...742...72N,2010ApJ...713..751I}.  With photometric measurements at various wavelengths, we can construct a spectral energy distribution of the planet's dayside atmosphere. From this we can infer properties such as the planetary albedo, the dayside energy budget and the efficiency of dayside-to-nightside energy redistribution \\citep[e.g.][]{2005ApJ...632.1132B}. As the atmospheric depth probed depends on the molecular opacity in the observation band, a spectrum can describe an atmosphere's vertical temperature structure. The observation that some planet atmospheres exhibit strong temperature inversions, or stratospheres, and others do not \\citep[see][for a summary]{2010ApJ...720.1569K} led to a suggestion that inversions are present in atmospheres hot enough to maintain high-opacity absorbers in the gas phase in the upper atmosphere \\citep[e.g.][]{2008ApJ...678.1419F}. However, this was challenged by recent contrary results \\citep{2008ApJ...684.1427M,2010ApJ...711..374F} and theory \\citep[e.g.][]{2009ApJ...699.1487S}. As an alternative, \\citet{2010ApJ...720.1569K} suggested that those planets orbiting chromospherically active stars lack inversions because the associated high UV flux destroys the high-opacity, high-altitude compounds that would otherwise induce inversions. This hypothesis is based on a small sample and further measurements across a wider parameter space are vital to test it. Specifically, there is a paucity of measurements for planets orbiting active stars. We suggest composition could be a key factor, with low-metallicity planets lacking high concentrations of the high-opacity absorbers, such as TiO and sulphur, thought to be responsible for inversions \\citep[e.g.][]{2008ApJ...678.1419F,2009ApJ...701L..20Z}.  As the spectral coverage increases, so too does the information about an atmosphere that we can discern. \\citet{2011Natur.469...64M} used seven measurements of the emission of WASP-12b \\citep{2011AJ....141...30C,2011ApJ...727..125C} to show that the dayside atmosphere is the first known to be carbon-dominated. Thus we are entering an era in which we can make statistically-sound inferences about the composition of exoplanet atmospheres. Further, \\citet{2011Natur.469...64M} demonstrated that the planet lacks a prominent thermal inversion and has very inefficient day-night energy circulation.  In this paper, we present {\\it Spitzer} measurements of WASP-19b's dayside thermal emission at 3.6, 4.5, 5.8 and 8.0 \\micron. Discovered by the Wide Angle Search for Planets \\citep{2006PASP..118.1407P,2010ApJ...708..224H}, the shortest-period hot Jupiter, WASP-19b, is a 1.17-\\mjup\\ planet in a near-circular, 19-hr orbit around a G8V star. \\citet{2011ApJ...730L..31H} showed the planet's orbital axis to be aligned with the spin axis of its host star. The planet's thermal emission was previously measured at 1.6 and 2.09 \\micron\\ \\citep{2010A&A...513L...3A, 2010MNRAS.404L.114G}, thus we bring the measurement tally to six bands. We use all six thermal emission measurements to characterise the planet's atmosphere and refine the system parameters by combining these data with pre-existing photometry of the transit and radial-velocity data. ",
        "conclusions": "\\subsection{Atmosphere model}  We interpret our observations of the hot Jupiter WASP-19b using the exoplanetary atmospheric modeling and retrieval method developed in \\citet{2009ApJ...707...24M, 2010ApJ...725..261M, 2011ApJ...729...41M}. We model a plane-parallel atmosphere of WASP-19b observed in thermal emission at secondary eclipse. The dayside spectrum of the planet is generated using line-by-line radiative transfer, with constraints of hydrostatic equilibrium and global energy balance, and includes the dominant sources of infrared opacity expected in gaseous atmospheres at high temperature. Our sources of opacity include molecular absorption due to H$_2$O, CO, CH$_4$, CO$_2$, NH$_3$, TiO, and VO \\citep{2008ApJS..174..504F, 2005Rothman, 2010Icar..205..674K} and H$_2$--H$_2$ collision-induced absorption \\citep{1997A&A...324..185B, 2002A&A...390..779B}. The concentrations of the species and the pressure--temperature ($P$--$T$) profile constitute the free parameters in the model \\citep{2009ApJ...707...24M}. We explore the parameter space of the model using a Markov-chain Monte Carlo scheme \\citep[see][]{2010ApJ...725..261M, 2011ApJ...729...41M}, and constrain regions of parameter space consistent with the measured planet-to-star flux density ratios at different levels of fit. Our goal is to constrain the existence of a possible temperature inversion, the dayside-to-nightside redistribution efficiency, the concentrations of the different molecular species, and the C/O ratio \\citep[e.g.][]{2011Natur.469...64M} in the dayside atmosphere of WASP-19b. In what follows, we discuss model solutions that explain the data within the 1-$\\sigma$ observational uncertainties, as shown in Fig.~\\ref{fig:atmos}.  \\subsection{Temperature inversion}  The data indicate the lack of a strong temperature  inversion in the dayside atmospheres of WASP-19b. The observations and two models fitting the data are shown in Fig.~\\ref{fig:atmos}. The lack of a temperature  inversion in WASP-19b is evident from the data even without detailed modeling. Firstly, the brightness temperature in the 4.5 $\\micron$ channel, $T_{\\rm b,4.5}$, is lower than that in the 3.6 $\\micron$ channel, $T_{\\rm b,3.6}$ (Table~\\ref{tab:sys-params}). The presence of a temperature  inversion is often indicated by a $T_{\\rm b,4.5}$ value considerably higher than the $T_{\\rm b,3.6}$ value, due to strong CO emission and some H$_2$O emission in the 4.5 $\\micron$ channel \\citep{2008ApJ...678.1436B, 2008ApJ...678.1419F, 2010ApJ...725..261M}. Secondly, the brightness temperatures at 1.6 and 2.09 \\micron\\ are larger than those in all four IRAC channels, which are at longer wavelengths. Since the bands at 1.6 and 2.09 \\micron\\ are windows in molecular opacity, they are expected to probe temperatures in deeper layers of the atmosphere compared to any of the IRAC channels. Therefore, the high temperatures in the 1.6 and 2.09 \\micron\\ bands compared to all the IRAC channels imply temperature decreasing outwards in the atmosphere, and hence the absence of a temperature inversion or, at most, the presence of one too weak to be detectable with the available data. Two model $P$--$T$ profiles without temperature inversions and the corresponding model spectra are shown in Fig.~\\ref{fig:atmos}. All the IRAC data can be explained by molecular absorption in the atmosphere due to the temperature decreasing outwards.  The lack of a temperature  inversion in WASP-19b offers a new constraint on existing classification schemes of irradiated giant exoplanets. WASP-19b falls in the category of highly irradiated hot Jupiters which have been predicted to host temperature  inversions due to TiO and VO, assuming solar abundances \\citep{2008ApJ...678.1419F} - the so called `TiO/VO hypothesis'. However, our finding of a lack of a strong temperature inversion in WASP-19b implies that either TiO and VO are depleted or an entirely different process is at play. If the composition is oxygen-rich, the lack of a temperature inversion in WASP-19b can be explained if TiO and VO are depleted in the upper atmosphere due to gravitational settling, which can be significant if the vertical mixing is weak \\citep{2009ApJ...699.1487S}. On the other hand, if the composition is carbon-rich, TiO and VO would be naturally scarce \\citep{2011ApJ...743..191M}.  \\citet{2010ApJ...720.1569K} posit that the high UV flux impinging on those planets orbiting chromospherically active stars could destroy the high-opacity, high-altitude compounds that would otherwise lead to temperature inversions. With \\rhk\\ = $-4.50 \\pm 0.03$, WASP-19 has a similar activity level to that of the handful of other stars around which hot Jupiters without temperature inversions (\\rhk\\ = $-4.9$ to $-4.5$) are known to orbit. Those planets thought to have inversions orbit quieter stars, with \\rhk\\ = $-5.4$ to $-4.9$. Thus, our finding that WASP-19b has no inversion supports the hypothesis of \\citet{2010ApJ...720.1569K} and usefully adds to the handful of such systems known. {\\it Spitzer} routinely measures the thermal emission of planets at 3.6 and 4.5 \\micron. For planets with temperature inversions, CO and water switch from absorption to emission, resulting in a higher flux in the 4.5 \\micron\\ band, in which these molecules have features. \\citet{2010ApJ...720.1569K} proposed a model-independent, empirical metric for classifying hot Jupiters, which we denote with $\\zeta$. This is the gradient of the measurements at 3.6 and 4.5 \\micron, i.e. $(\\iractwo - \\iractwo) / 0.9 \\micron$, minus the gradient of the black body that is the best-fit to the two measurements. A positive $\\zeta$-value would suggest an inverted atmosphere and a strongly negative $\\zeta$-value would indicate a non-inverted atmosphere; \\citet{2010ApJ...720.1569K} suggest a delineation around $\\zeta = -0.05$ per cent \\micron$^{-1}$. For WASP-19b we calculated $\\zeta = -0.031 \\pm 0.043$ per cent \\micron$^{-1}$. This value is at the border between inversion and no-inversion and therefore is consistent with our finding that WASP-19b does not have a strong inversion.  We note that in their activity-inversion plot (their figure~5), \\citet{2010ApJ...720.1569K} omit XO-3 but include TrES-4, HAT-P-7 and WASP-18, all of which have $\\teff \\gtrsim 6200$. XO-3b has a temperature inversion and XO-3 has a \\rhk\\ index indicative of activity, which seems to contradict the activity-inversion hypothesis. The other three planets have inversions and orbit quiet stars. \\citet{2010ApJ...720.1569K} concluded that in fact XO-3 is likely to be chromospherically quiet, based on a visual inspection of their spectrum and having noted that the \\rhk\\ calibration is unreliable for stars with $\\teff \\gtrsim 6200$. Perhaps then TrES-4, HAT-P-7 and WASP-18 are also suspect since they also have $\\teff \\gtrsim 6200$.  \\subsection{Atmospheric composition} We find that the observations can be explained by models with oxygen-rich as well as carbon-rich compositions. The absorption in the near-IR (1.6 and 2.09 \\micron) bands is minimal due to the lack of major molecular features. The constraints on the composition come primarily from the IRAC data, which together encompass features of CO, H$_2$O, CH$_4$, and CO$_2$. The near-IR data, however, are critical to constraining the temperature of the lower atmosphere and thus are key in anchoring the model-atmosphere spectra to the measured SED. Two models with different C/O ratios, C/O = 0.5 (oxygen-rich) and C/O = 1 (carbon-rich), are shown in Fig.~\\ref{fig:atmos}. As demonstrated in \\citet{2011ApJ...743..191M}, CO is a dominant carbon-bearing molecule in both C-rich and O-rich regimes. Consequently, the 4.5 $\\micron$ absorption in both models in Fig.~\\ref{fig:atmos} is caused primarily by CO absorption; in the O-rich model CO$_2$ contributes additional absorption in this channel. The absorption in the 3.6, 5.8, and 8.0 micron IRAC channels in the O-rich model is caused primarily by H$_2$O absorption, whereas absorption in the C-rich model is caused by a combination of H$_2$O and CH$_4$ absorption; H$_2$O is depleted by a factor of 100 and CH$_4$ is enhanced by a factor of 1000 with respect to the O-rich model, both of which are chemically feasible \\citep{2011ApJ...743..191M}. The principle difficulty in differentiating between the two models with the current data are the large uncertainties in the 5.8 and 8.0 $\\micron$ IRAC data. For example, a high 5.8 $\\micron$ point would indicate low water absorption, and hence high C/O, as demonstrated in \\citet{2011Natur.469...64M}. New observations in the near infrared can differentiate between spectra from the carbon-rich and oxygen-rich compositions \\citep[e.g.][]{2011Natur.469...64M}. The water abundance can be measured via transmission spectoscopy of the 1.4 \\micron\\ water band using the G141 grism of HST/WFC3; these observations were recently peformed for WASP-19b by \\citet{2009hst..prop12181D}. We could measure, or at least place useful constraints on, the TiO abundance with ground-based occultation observations in the $z$ and $J$ bands.  \\begin{figure} \\includegraphics[width=0.5\\textwidth]{spectra_wasp19b.eps} \\caption{Spectral energy distribution of WASP-19b relative to that of its host star. The blue dots are our fitted planet-to-star flux density ratios. The observation-band transmission curves are shown along the abscissa. The planet-to-star flux density ratios indicate the absence of a strong temperature inversion. Two model-atmosphere spectra are shown, each of which fits the data and lacks a temperature inversion. One model has a C/O ratio of 0.5 and is shown as a green line. The other has a C/O ratio of 1.0 is depicted by a red line. The corresponding band-integrated model fluxes are shown as dots with the same colour scheme. The models' pressure--temperature profiles down to 100 bar are shown in the inset. Assuming zero albedo, the maximum dayside-to-nightside redistribution efficiency is 20 per cent for the C-rich model and 34 per cent for the O-rich model. The three dotted lines show planetary black bodies with temperatures of 1800, 2250 and 2900 K (chosen to give a sense of scale to the spectra) divided by a stellar spectrum. \\label{fig:atmos}} \\end{figure}    \\subsection{Orbital eccentricity}  For a circular orbit, mid-occultation occurs half an orbital period after mid-transit. We find the occultation to occur only $4^{+35}_{-30}$ s later than this and constrain both \\ecos\\ and \\esin\\ to a small region around zero. Hence, the orbit is very nearly circular, though the available data do permit a small, non-zero eccentricity providing that the major axis of the orbit is near-aligned with our line of sight, such that the occulation time is not affected (i.e. $|\\omega| \\approx 90$; Figure~\\ref{fig:eom}). We place a 3-$\\sigma$ upper limit on eccentricity of $e < 0.027$.  \\begin{figure} \\includegraphics[width=0.5\\textwidth]{ecc-omega.eps} \\caption{ The MCMC posterior distributions of $e$ and $\\omega$. The white, blue and black contours are, respectively, the 1-, 2- and 3-$\\sigma$ confidence limits. The shading of each bin is proportional to the logarithm of the number of MCMC steps within. The available data exclude large values of $e$ for any orbital orientation, and only very small values of $e$ are permitted unless $|\\omega| \\approx 90$. \\label{fig:eom}} \\end{figure}  Hot Jupiters are considered to have been moved inwards to close orbits through planet--planet scattering or by the Kozai mechanism and tidal circularisation \\citep[e.g.:][]{2011Natur.473..187N,2011ApJ...735..109W,2011A&A...533A...7B}. However, for the very shortest-period systems, such as WASP-19b, it is unlikely that they could have been moved directly to their current orbit, since that would have required careful fine-tuning to avoid destruction by collision with the star.  Thus most likely WASP-19b was first moved to an orbit near $\\sim$\\,2 Roche radii and has since spiralled inwards through tidal orbital decay (see \\citealp{2011ApJ...732...74G} and the discussion of WASP-19b specifically in \\citealp{2011ApJ...730L..31H}).  From tidal theory, the circularisation of a hot Jupiter's orbit is thought to proceed much faster than the infall, and this is consistent with the the observation that hot Jupiters tend to be in circular orbits. Thus the suggestion that WASP-19b has undergone significant tidal decay, from $\\sim$\\,2 Roche radii to the current 1.2 Roche radii, leads to the expectation that the current eccentricity will be essentially zero, in line with our results."
    },
    "1112/1112.5684_arXiv.txt": {
        "abstract": "We explore heavy-element nucleosynthesis in the explosion of massive stars which are triggered by a quark--hadron phase transition during the early post-bounce phase of core-collapse supernovae. The present study is based on general-relativistic radiation hydrodynamics simulations with three-flavor Boltzmann neutrino transport in spherical symmetry, which utilize a quark-hadron hybrid equation of state based on the MIT bag model for strange quark matter. The quark--hadron phase transition inside the stellar core forms a shock wave propagating toward the surface of the proto-neutron star. This shock wave results in an explosion and ejects neutron-rich matter from the outer accreted layers of the proto-neutron star. Later, during the cooling phase, the proto-neutron star develops a proton-rich neutrino-driven wind. We present a detailed analysis of the nucleosynthesis outcome in both neutron-rich and proton-rich ejecta and compare our integrated nucleosynthesis with observations of the solar system and metal-poor stars. For our standard scenario we find that a ``weak\" $r$-process occurs and elements up to the second peak ($A \\sim 130$) are successfully synthesized. Furthermore, uncertainties in the explosion dynamics could barely allow to obtain the strong $r$-process which produces heavier isotopes including the third peak ($A \\sim 195$) and actinide elements. ",
        "introduction": "Nucleosynthesis from core-collapse supernovae is certainly responsible for the production of intermediate mass elements including the so-called alpha elements and a certain fraction of iron isotopes and their neighbors (the Fe-group nuclei). The lighter alpha elements from oxygen to silicon have contributions from hydrostatic burning which takes place during stellar evolutions, while the heavier ones and the Fe-group isotopes originate from explosive burning \\citep{WoosleyWeaver:1995,Thielemann:etal:1996,Woosley:etal:2002, Nomoto:etal:2006,WoosleyHeger:2007,HegerWoosley:2010,Thielemann:etal:2011}. A major open question is related to the source of heavy elements beyond iron.  There have been strong expectations that the innermost layer of ejecta, close to the forming neutron star, remains neutron-rich and is a possible site for the $r$-process. For many years the late neutrino-driven wind (NDW), following the actual explosion, seemed an adequate $r$-process site in core-collapse supernovae that formed a neutron star \\citep{QianWoosley:1996,Takahashi:etal:1994,Woosley:etal:1994}. Such investigations were followed up by many parametrized calculations that explored the sensitivity of the most relevant $r$-process parameter, the neutron-to-seed ratio, via variations of entropy $S$, electron fraction (or number of proton per nucleon) $Y_e$ and expansion timescale $\\tau$ \\citep{Hoffman:etal:1997,MeyerBrown:1997,Freiburghaus:etal:1999,Farouqi:etal:2010}. However, steady state wind models \\citep{Thompson:etal:2001,Wanajo:2006} showed that it is very hard to attain the required entropies. These results have been recently confirmed by fully hydrodynamical simulations \\citep{Arcones:etal:2007} that in addition showed that the presence of a reverse shock does not have any major impact on the neutron-to-seed ratio.  Furthermore, recent investigations noticed that the early NDW turns matter proton-rich, producing specific Fe-group isotopes and in the subsequent $\\nu p$-process nuclei with a mass number up to $A\\sim 80$ -- $90$ \\citep{Liebendoerfer:etal:2003,Pruet:etal:2005,Froehlich:etal:2006a, Froehlich:etal:2006b,Pruet:etal:2006,Wanajo:2006}. While it was initially hoped that there exists a chance that the late NDW still turns neutron-rich (after its initial, early proton-rich phase) existing core-collapse calculations seem to indicate that the wind becomes even more proton-rich in the long-term evolution \\citep{Fischer:etal:2010,Huedepohl:etal:2010,Fischer:2012}, although very recent investigations \\citep{MartinezPinedo:2012,Roberts:2012} related to the spectral evolution of $\\nu_e$ and $\\bar\\nu_e$ in medium might change this somewhat.  It was realized that the best chance to eject neutron-rich matter, is to utilize material stemming from the initial collapse and compression, where electron captures turned it neutron-rich early in the explosion, before neutrino interactions have the chance to convert it into proton-rich matter in the NDW. Electron-capture supernovae, which explode without a long phase of accretion onto the proto-neutron star, apparently provide such conditions \\citep{Wanajo:etal:2011}. However, the $Y_{e}$'s obtained under such conditions do not support a strong $r$-process, which successfully reproduces the platinum peak of $r$-elements around $A=195$.  An alternative explosive scenario as a possible site for the $r$-process has been suggested by \\citet{Jaikumar:2007}, where after a supernova explosion a quark--hadron phase transition inside the neutron star can eject very neutron-rich material from the neutron star surface. This scenario has been referred to as a \\emph{quark-nova}. In the present paper, we follow a different approach and investigate core-collapse supernova explosions that are triggered by a quark-hadron phase transition during the early post-bounce phase \\citep[for details, see][]{Sagert:etal:2009,Fischer:etal:2011} and their nucleosynthesis features. Under such conditions, zones which are initially neutron-rich, can be (promptly) ejected without experiencing the strong effects of the neutrino flux which comes from the central proto-neutron star.  In the following sections, we discuss the supernova explosion model and the nuclear physics inputs utilized for the nucleosynthesis calculations in Section 2, the conditions experienced and the resulting ejecta in Section 3, and give a detailed analysis of the ejecta composition in Section 4, followed by conclusions, also in comparison to alternative $r$-process sites, in Section 5. ",
        "conclusions": "Supernova nucleosynthesis is well understood for the outer ejected mass zones, which can be well approximated by a shock wave with appropriate energy passing through the layers of the progenitor \\citep{WoosleyWeaver:1995,Thielemann:etal:1996,Woosley:etal:2002, Nomoto:etal:2006,WoosleyHeger:2007,HegerWoosley:2010,Thielemann:etal:2011}. What remains uncertain is the composition of the innermost ejecta, directly linked to the explosion mechanism, i.e., the collapse and explosion phase. In the present paper, we analyzed these mass zones of core-collapse supernova explosions triggered by a quark--hadron phase transition during the early post-bounce phase \\citep{Sagert:etal:2009,Fischer:etal:2011}. A number of aspects are important for understanding these results. The very innermost ejecta are strongly affected by the NDW. Recent investigations noticed that this NDW turns matter proton-rich, producing specific Fe-group isotopes and in the subsequent $\\nu p$-process nuclei with masses up to $A=80-90$ \\citep{Liebendoerfer:etal:2003,Pruet:etal:2005, Froehlich:etal:2006a,Froehlich:etal:2006b,Pruet:etal:2006,Wanajo:2006}. Even in the long-term evolution proton-rich conditions prevail \\citep{Fischer:etal:2010,Huedepohl:etal:2010}. Thus, there seems to exist no chance to produce $r$-process matter in these innermost regions, despite many interesting parameter studies for NDW ejecta in terms of entropy $S$, electron fraction $Y_e$, and expansion timescale $\\tau$ \\citep{Hoffman:etal:1997,MeyerBrown:1997,Freiburghaus:etal:1999,Farouqi:etal:2010} or hydrodynamic studies, partially with parameter variations \\citep{Arcones:etal:2007,Kuroda:etal:2008,PanovJanka:2009,Roberts:etal:2010,ArconesMontes:2011}. Very recently, \\citet{MartinezPinedo:2012} and \\citet{Roberts:2012}, explored charged-current weak processes (i.e., electron-neutrino and anti-neutrino captures on neutrons and protons), consistent with the EOSs in neutrino-driven supernova explosion. They find modifications of the reaction $Q$-values due to medium effects, which increase spectral differences between $\\nu_e$ and $\\bar\\nu_e$. This results in slightly neutron-rich ejecta. Nevertheless, since the entropy per baryon is still low a strong $r$-process is unlikely to occur.  In order to obtain $r$-process conditions, a better chance to eject neutron-rich matter is provided, when neutron-rich matter stems from the initial collapse and compression, where electron captures made it neutron-rich, early in an explosion, before neutrino interactions have the chance to turn it proton-rich in the NDW. Core-collapse supernovae, exploding via the quark--hadron phase transition (the focus of the present study), or electron-capture supernovae \\citep{Wanajo:etal:2011}, which explode without a long phase of accretion onto the proto-neutron star, both lead to a rather prompt ejection of prior compressed and neutronized matter. However, the $Y_e$ obtained under such conditions does not support a full $r$-process. This kind of outcome also characterizes the conditions we find in the prompt and quasi-prompt ejecta of the present study, which did not experience a strong neutrino flux. However, the $Y_e$-values attained are not smaller than $0.33$. Whether uncertainties in the input or explosion physics can change this down to values close to $0.23$, necessary for obtaining the third $r$-process peak, remains to be shown in an upcoming article where we will include charged-current reaction rates that are consistent with the EOS following \\citet{MartinezPinedo:2012} and \\citet{Roberts:2012}.  A related, but different, phenomenon has been discussed in \\citet[, and references therein]{Jaikumar:2007}. In their scenario, the quark--hadron phase transition does not happen shortly after core collapse and is not the cause of the supernova explosion. Instead a regular neutrino-driven supernova explosion occurs first and leads to a deleptonized neutron star with low-$Y_e$ and a steep density gradient at the surface. Their scenario of a ``quark-novae\", occurring via a quark--hadron phase transition in an existing neutron star, is reported to produce $r$-process elements via the prompt ejection of extremely neutron-rich material from the neutron star surface. In this sense, quark-novae differ from the scenario explored in the current study. They result in a more compact object already in $\\beta$-equilibrium, such that the expansion timescale of the ejected material is much faster, compared to the scenario discussed in the present paper. However, in \\citet{Jaikumar:2007} the authors apply a very simplified description for the quark--hadron phase transition and ignore neutrino transport in general, as well as weak processes in the hadronic part of their neutron star. The latter may be essential for producing neutrinos which may change $Y_e$ during shock passage and hence lead to different nucleosynthesis results than what has been discussed in \\citet{Jaikumar:2007}. It remains to be shown if the obtained favorable conditions for the $r$-process will remain, applying more sophisticated input physics.  Given our results described above, we conclude that supernova explosions triggered via the quark--hadron phase transition during the early post-bounce phase can contribute to a weak $r$-process, consistent with observations of low-metallicity stars \\citep{Honda:etal:2006} and with LEPP abundances \\citep{Travaglio:etal:2004}. Such conditions apparently do not occur in regular core-collapse supernovae. However, strong $r$-process conditions, which also produce the third $r$-process peak and the actinides, have in simulations only materialized in neutron star mergers \\citep{Freiburghaus:etal:1999b,Goriely:etal:2011}, fast rotating core-collapse supernovae with strong magnetic fields and jet ejecta \\citep{Cameron:2003,Nishimura:etal:2006,Fujimoto:etal:2008} or accretion disks around black holes \\citep{Surman:etal:2008,WanajoJanka:2012}. For the explosion model discussed here, investigation of uncertainties in the hydrodynamics and the nuclear physics input showed that a strong $r$-process to produce up to the third peak elements is rarely possible to obtain.  Neutron star mergers have been shown to be powerful sources of $r$-process matter, in fact ejecting a factor of $100$ -- $1000$ more $r$-process material than required on average from core-collapse supernovae, if those would have to explain solar $r$-process abundances. This would actually support the large scatter of Eu/Fe in comparison to, e.g., O, Mg, Si, S, and Ca/Fe, where the latter are clearly produced in supernovae. The only problem is that it might be hard to explain the early appearance of $r$-process matter for metallicities at and below $\\rm{[Fe/H]}=-3$. Neutron star mergers will appear after the first supernovae have already produced Fe and studies by \\citet{Argast:etal:2004} showed that one expects $r$-process products only for metallicities $\\rm{[Fe/H]} = -3 \\sim -2$. Some recent studies, which include the fact that our Galaxy is possibly the result from smaller merging subsystems (with different star formation rates) have been expected to show a way out of this dilemma. If this cannot solved, we need another strong $r$-process source already at low metallicities, and possibly jets from rotating core collapses with strong magnetic fields could be the solution \\citep[see, e.g,][]{Nishimura:etal:2006,Fujimoto:etal:2007,Winteler:2012}."
    },
    "1112/1112.0068_arXiv.txt": {
        "abstract": "Cosmic ray (CR) transport and acceleration is essential for many astrophysical problems, e.g., CMB foreground, ionization of molecular clouds and all high energy phenomena. Recent advances in MHD turbulence call for revisions in the paradigm of cosmic ray transport. We use the models of magnetohydrodynamic turbulence that were tested in numerical simulation, in which turbulence is injected at large scale and cascades to to small scales. We shall address the issue of the transport of CRs, both parallel and perpendicular to the magnetic field and show that the issue of cosmic ray subdiffusion is only important for restricted cases when the ambient turbulence is far from that suggested by numerical simulations. Moreover, on scales less than injection scale of turbulence, CRs's transport becomes super-diffusive. We also shall discuss the nonlinear growth of kinetic gyroresonance instability of cosmic rays induced by large scale compressible turbulence. This gyroresonance of cosmic rays on turbulence is demonstrated an important scattering mechanism in addition to direct interaction with the compressible turbulence. The feedback of the instability on large scale turbulence cannot be neglected, and should be included in future simulations. ",
        "introduction": "The propagation and acceleration of cosmic rays (CRs) is governed by their interactions with magnetic fields. Astrophysical magnetic fields are turbulent and, therefore, the resonant and non-resonant (e.g. transient time damping, or TTD) interaction of cosmic rays with MHD turbulence is the accepted principal mechanism to scatter and isotropize cosmic rays \\citep[see][]{Schlickeiser02}. In addition, efficient scattering is essential for the acceleration of cosmic rays. For instance, scattering of cosmic rays back into the shock is a vital component of the first order Fermi acceleration \\citep[see][]{Longairbook}. At the same time, stochastic acceleration by turbulence is entirely based on scattering. The dynamics of cosmic rays in MHD turbulence holds the key to all high energy astrophysics and related problems.  We live in an exiting era when we are starting to test fundamental processes taking place at the beginning of the Universe, at the event horizon of black holes, when the nature of dark matter and dark energy is being probed etc. In the mean time, with the launching of the new facilities like Fermi, we have much more observational data available than ever before.  Using computers many researchers make sophisticated complex models to confront the observations in unprecedented details. This makes it urgent that we understand the key physical processes underlying astrophysical phenomena, can parameterize them and, if necessary, use as a subgrid input in our computer models.   At present, the propagation of the CRs is an advanced theory, which makes use both of analytical studies and numerical simulations. However, these advances have been done within the turbulence paradigm which is being changed by the current research in the field. Instead of the empirical 2D+slab model of turbulence, numerical simulations suggest anisotropic Alfv\\'enic modes (an analog of 2D, but not an exact one, as the anisotropy changes with the scale involved) + fast modes or/and slab modes \\citep{CL02_PRL}.  At the same time, one should not disregard the possibilities of generation of additional perturbations on small scales by CR themselves. For instance,  the slab Alfv\\'enic perturbation can be created, e.g., via streaming instability \\citep[see][]{Wentzel74, Cesarsky80}. These perturbations are present for a range of CRs energies (e.g., $\\lesssim 100$GeV in interstellar medium for the streaming instability) owing to non-linear damping arising from ambient turbulence (\\citealp[][henceforth YL02, YL04]{YL02, YL04}, \\citealp{FG04}). Instabilities induced by anisotropic distribution of CRs were also suggested as a possibility to scatter CRs \\citep[]{Lerche, Melrose74}. Gyroresonance instability induced by large scale turbulence compressions, can be an important feedback processes that creates slab modes to efficiently scatter CRs \\citep{LB06, YL11}.  Propagation of CRs perpendicular to mean magnetic field is another important problem for which one needs to take into account both large and small scale interactions in tested models of turbulence. Indeed, if one takes only the diffusion along the magnetic field line and field line random walk \\citep[FLRW][]{Jokipii1966, Jokipii_Parker1969, Forman1974}, compound (or subdiffusion) would arise. Whether the subdiffusion is realistic in fact depends on the models of turbulence chosen \\citep{YL08}. In this paper we again seek the answer to  this questions within domain of numerically tested models of MHD turbulence.  In what follows, we discuss the cosmic ray transport in large scale turbulence in \\S2. In \\S3, we shall study the perpendicular transport of cosmic rays on both large and small scales. We shall also discuss the issue of super-diffusion and the applicability of sub-diffusion. In \\S4, we address the issue of gyroresonance instability of CRs and its feedback on large scale compressible turbulence. Summary is provided in \\S5. ",
        "conclusions": "\\begin{itemize} \\item Compressible fast modes are most important for CR scattering. CR transport therefore varies from place to place. \\item Large scale mirror is essential for pitch angle scattering ( including 90 degree). \\item Subdiffusion does not happen in 3D turbulence . \\item Our results are tested using input from turbulence simulations. \\item Small scale slab waves are generated in compressible turbulence by gyroresonance instability, dominating the scattering of low energy CRs ($<100$GeV). Feedback of CRs on turbulence should be included in future simulations. \\end{itemize}"
    },
    "1112/1112.5151_arXiv.txt": {
        "abstract": "We present Bolocam observations of two galaxy cluster candidates reported as unconfirmed in the Planck early Sunyaev-Zel'dovich (eSZ) sample, PLCKESZ G115.71+17.52 and PLCKESZ G189.84-37.24. We observed each of these candidates with Bolocam at 140~GHz from the Caltech Submm Observatory in October 2011. The resulting images have white noise levels of $\\simeq 30$~\\uKamin~in their central regions. We find a significant SZ decrement towards PLCKESZ G115.71. This decrement has a false detection probability of $5.3 \\times 10^{-5}$, and we therefore confirm PLCKESZ G115.71 as a cluster. The maximum SZ decrement towards PLCKESZ G189.84 corresponds to a false detection probability of 0.027, and it therefore remains as an unconfirmed cluster candidate. In order to make our SZ-derived results more robust, we have also analyzed data from the Wide-field Infrared Survey Explorer (WISE) at the location of each cluster candidate. We find an overdensity of WISE sources consistent with other clusters in the eSZ at the location of PLCKESZ G115.71, providing further evidence that it is a cluster. We do not find a significant overdensity of WISE sources at the location of PLCKESZ G189.84. ",
        "introduction": "The thermal Sunyaev-Zel'dovich (SZ) effect describes the scattering of CMB photons with the hot electrons in the intra-cluster medium (ICM) \\citep{sunyaev72}. The surface brightness of the SZ effect is independent of redshift, and the integrated SZ signal is approximately equal to the total thermal energy of the ICM divided by the square of the angular diameter distance, making the SZ signal an excellent cluster mass proxy with little redshift evolution~\\citep{barbosa96, holder00}. Consequently, surveys using the SZ effect to detect clusters have an almost mass-limited selection function and can yield powerful constraints on cosmological parameters \\citep{haiman01,holder01}. Three such SZ surveys are currently ongoing, conducted by the South Pole Telescope (SPT), the Atacama Cosmology Telescope (ACT), and the Planck satellite, and each has published initial cluster catalogues \\citep{vanderlinde10, williamson11, marriage11, planck_esz}  The Planck survey has released a list of 189 cluster candidates called the early SZ (eSZ) sample \\citep{planck_esz}. Among this list, 169 are previously known clusters, and 12 were confirmed by the Planck team prior to the publication of the eSZ. Subsequent SZ observations by SPT and the Arcminute Microkelvin Imager (AMI) have confirmed 5 and 1 additional eSZ candidates, respectively \\citep{story11, hurley-walker11}. Prior to our original submission of this manuscript, there had been no published confirmations of the remaining 2 cluster candidates in the eSZ sample, PLCKESZ G115.71+17.52 and PLCKESZ G189.84-37.24.\\footnote{ AMI observed, but did not detect, PLCKESZ G115.71 \\citep{hurley-walker11}. Furthermore, \\citet{muchovej12} published an SZA detection of PLCKESZ G115.71, and an SZA non-detection of PLCKESZ G189.84, after our original submission of this manuscript.} We note that \\citet{planck_esz} lists possible contamination from dust emission near PLCKESZ G115.71 and a high level of contamination from Galactic emission towards G189.84. Additionally, \\citet{planck_esz} specifically notes that G189.84 is a low reliability source, and is the only unconfirmed eSZ candidate without a ROSAT All-Sky Survey (RASS) association. ",
        "conclusions": "In October 2011 we performed dedicated 140~GHz Bolocam observations of two cluster candidates from the Planck eSZ sample, PLCKESZ G115.71 and PLCKESZ G189.84. These observations total $\\simeq 10$~hours each, and produce images with a noise RMSs of $\\simeq 30$~\\uKamin. From these images we have confirmed PLCKESZ G115.71 as a cluster at high significance, with a false detection probability of $5.3 \\times 10^{-5}$. We are not able to confirm PLCKESZ G189.84 as a cluster, although our image does show a weak decrement with a false detection probability of 0.027. Both our detection of PLCKESZ G115.71, and our non-detection of PLCKESZ G189.84, indicate SZ signals significantly lower than expected based on Bolocam observations of other eSZ cluster candidates which are known to be true clusters. However, we note the presence of positive flux immediately south of the SZ decrement for PLCKESZ G115.71, which may be partially filling in the SZ decrement from the cluster. We have also examined WISE data at the positions of these cluster candidates, and find a significant overdensity of sources towards PLCKESZ G115.71, providing further evidence that it is a cluster."
    },
    "1112/1112.5221_arXiv.txt": {
        "abstract": "We present observations of a large-scale disc of neutral hydrogen (\\HI) in the nearby Fanaroff $\\&$ Riley type-I radio galaxy NGC~3801 with the Westerbork Synthesis Radio Telescope. The \\HI\\ disc (34 kpc in diameter and with $M_{\\rm HI} = 1.3 \\times 10^{9} M_{\\odot}$) is aligned with the radio jet axis. This makes NGC~3801 an ideal system for investigating the evolution of a small radio source through its host galaxy's cold ISM. The large-scale \\HI\\ disc is perpendicular to a known inner CO disc and dust-lane. We argue that the formation history of the large-scale \\HI\\ disc is in agreement with earlier speculation that NGC~3801 was involved in a past gas-rich galaxy-galaxy merger (although other formation histories are discussed). The fact that NGC~3801 is located in an environment of several \\HI-rich companions, and shows indications of ongoing interaction with the nearby companion NGC~3802, strengthens this possibility. The large amounts of ambient cold ISM, combined with X-ray results by \\citet*{cro07} on the presence of over-pressured radio jets and evidence for an obscuring torus, are properties that are generally not, or no longer, associated with more evolved \\FRI\\ radio sources. We do show, however, that the \\HI\\ properties of NGC~3801 are comparable to those of a significant fraction of nearby low-power compact radio sources, suggesting that studies of NGC~3801 may reveal important insight into a more general phase in the evolution of at least a significant fraction of nearby radio galaxies. ",
        "introduction": "\\label{sec:intro} Over the past decade it has become clear that radio sources emanating from super-massive black-holes in the centres of galaxies are - rather than merely a by-product of galaxy evolution - key players in the formation and evolution of galaxies throughout the Universe. It is now fairly well established that the properties of Active Galactic Nuclei (AGN) are intimately linked with the properties of their host galaxies. While empirical evidence has established the co-evolution of massive black holes and their host galaxies at various redshifts (e.g. \\citealt*{fer00}, \\citealt{ale05}, see also \\citealt*{sil98}), numerical simulations demonstrate that AGN-induced outflows are important feedback processes for clearing the circum-nuclear regions and halting the growth of the super-massive black-holes, regulating the correlations found between the black-hole mass and the host galaxy's bulge properties and preventing the formation of too many massive galaxies in the early Universe (e.g. \\citealt*{dim05}, \\citealt*{spr05}, \\citealt{hop05}, \\citealt{wag11}).  An important phase of AGN feedback is when synchrotron jets emanating from the central black-hole region start propagating into the host galaxy's inter-stellar medium (ISM). At high redshifts, radio sources are observed to be aligned with optical continuum from warm gas and stars (\\citealt*{cha87}, \\citealt{mcc87}) and the gas in these regions is characterised by highly disturbed kinematics (\\citealt*{mcc96}, \\citealt*{vil99}, \\citealt{hum06}). This indicates that radio sources vigorously interact with the host galaxy's ISM, entrain and expel gas and induce or quench star formation \\citep[e.g.][]{bic00}. At intermediate and low redshifts, powerful radio jets are know to induce fast ($>1000$\\,\\kms) outflows of ionised and neutral gas \\citep[e.g.][]{tad91,cla98,oos00,mor03,mor05a,mor05,mor07,emo05,hol06,hol08,hol11} and can shock-heat the molecular gas \\citep{gui11}, exerting significant feedback on the host galaxy's ISM.  While propagating radio jets thus have a profound impact on their surrounding ISM, the reverse is also true \\citep[see simulations by e.g.][]{sax02b,sax02a,jey05,bic06,sut07}. It is well established that external factors, in particular the inter-stellar and inter-galactic medium (ISM and IGM), shape the morphological and physical properties of radio sources. For example, there is strong evidence that radio sources in general show a degree of intrinsic asymmetry (apart from orientation effects), which is most prominent in compact sources \\citep[][]{sai95,sai01,sai03a}. Compact radio sources also show a larger degree of polarisation asymmetry than extended sources, suggesting a link between the propagating radio jets and surrounding ISM \\citep*{sai03b}. Also, the impact of supersonic jets onto the ISM/IGM can create cocoons of (shocked) gas and synchrotron emission with `hot-spot' morphologies, while entrainment of gas by slower jets can lead to bright radio sources with edge-darkened lobes \\citep[see e.g. reviews by][]{bri84,fer98,bic06}. These external factors may thus contribute to the classical dichotomy in radio sources between Fanaroff $\\&$ Riley type-I  \\citep[\\FRI;][]{fan74} objects, with low-power, sub-relativistic jets that have an edge-darkened morphology, and type-II (\\FRII) sources, with powerful, relativistic jets ending in a bright hot-spot \\citep[see e.g.][]{gop00}.   A particularly interesting object in this respect is the nearby ($z = 0.0115$) radio galaxy NGC~3801. NGC~3801 is classified as an \\FRI\\ radio galaxy, based on its total radio power \\citep[$P_{\\rm 1.4\\,GHz} = 2.9 \\times 10^{23}$\\,W\\,Hz$^{-1}$;][]{con88}, radio source morphology \\citep{jen82,cro07} and weak emission-line spectrum \\citep{hec86}. The radio source of NGC~3801 is small \\citep[$\\sim 11$ kpc;][]{jen82} and contained well within the optical boundaries of a peculiar early-type host galaxy, with a patchy dust feature along the major stellar axis and a perpendicular prominent inner dust-lane \\citep{hec86,ver99}, as well as a distorted and box-shaped faint outer envelope \\citep{hec86}. The inner dust-lane coincides with a kpc-scale CO disc, oriented perpendicular to the radio axis (\\citealt{das05}, see also \\citealt{oca10}).  X-ray observations by \\citet{cro07} showed evidence for shocked X-ray gas in the central region of NGC~3801, created by supersonic, over-pressured jets that are not in equilibrium with the hot ISM. These X-ray studies also revealed a high absorbing column towards the nucleus of NGC~3801, likely due to an obscuring torus \\citep{cro07}. In addition, single-dish observations showed that NGC~3801 is rich in neutral hydrogen (\\HI) gas \\citep{hec83,dup96}. Based on its cold gas content and peculiar optical host galaxy morphology, it has been suggested that NGC~3801 could have been involved in a gas-rich galaxy-galaxy merger \\citep{das05,cro07}. As speculated by \\citet{cro07}, these are all properties that are generally not associated with \\FRI\\ sources, but are believed to be common among more powerful, high-excitation radio-loud AGN (see also \\citealt{hec86}, \\citealt{bau92}, \\citealt*{har06}, \\citealt{cro08}, as well as Sect. \\ref{sec:classical}).  In this paper we present radio synthesis observations of the \\HI\\ gas in NGC~3801. Our observations reveal that the \\HI\\ gas is distributed in a large-scale disc of cold gas. The large-scale \\HI\\ disc is aligned with the radio axis, indicating that the radio jets propagate directly into the cold gas disc. Based on our \\HI\\ results, we discuss the formation history of NGC~3801 and investigate further how unique NGC~3801 is compared to other nearby \\FRI\\ radio galaxies.  Throughout this paper we assume H$_{0} = 71$\\,\\kmsMp, which puts NGC~3801 at a distance of 48.6 Mpc and 1 arcsec = 0.24 kpc. ",
        "conclusions": "\\label{sec:conclusions}  We have shown that the \\HI\\ gas in NGC~3801 -- detected in earlier single-dish observations by \\citet{hec83} and \\citet{dup96} -- is distributed in a regularly rotating large-scale \\HI\\ disc. We also detect \\HI\\ gas in the close companion NGC~3802 (which is in apparent ongoing interaction with NGC~3801) as well as several other galaxies and gas clouds in the immediate environment of NGC~3801. The small radio source in NGC~3801 is propagating directly into the large-scale \\HI\\ disc, making NGC~3801 an important system for investigating in detail the evolution of a small radio source through its host galaxy's cold ISM. Because of the similarity between the \\HI\\ properties of NGC~3801 and those of other nearby low-power compact radio sources, studying NGC~3801 may reveal important insight into the early evolution of at least a significant fraction of low-power radio sources."
    },
    "1112/1112.0318_arXiv.txt": {
        "abstract": "{Since its launch on October 2002, the {\\it INTEGRAL} satellite has revolutionized our knowledge of the hard X--ray sky thanks to its unprecedented imaging capabilities and source detection positional accuracy above 20 keV. Nevertheless, many of the newly-detected sources in the {\\it INTEGRAL} sky surveys are of unknown nature. The combined use of available information at longer wavelengths (mainly soft X--rays and radio) and of optical spectroscopy on the putative counterparts of these new hard X--ray objects allows pinpointing their exact nature. Continuing our long-standing program running since 2004, here we report the classification, through optical spectroscopy, of 19 more unidentified or poorly studied high-energy sources detected with the IBIS instrument onboard {\\it INTEGRAL}.}  \\FullConference{The Extreme and Variable High Energy Sky - extremesky2011,\\\\ September 19-23, 2011\\\\ Chia Laguna (Cagliari), Italy}   \\begin{document} ",
        "introduction": "One of the main objectives of the {\\it INTEGRAL} \\cite{w03} is a regular survey of the whole sky in the hard X--ray band. This makes use of the unique imaging capability of the IBIS instrument \\cite{u03} which allows the detection of sources at the mCrab level with a typical localization accuracy of a few arcmin. The latest (4$^{\\rm th}$ one) IBIS catalogue published up to now by \\cite{b10} shows that, among more than 700 sources detected above 20 keV, a substantial fraction ($\\sim$30\\%) had no obvious counterpart at other wavelengths and therefore could not yet be associated with any known class of high-energy emitting objects.  With the task of identifying them, our group started in 2004 an optical spectroscopy program which used data from ten telescopes worldwide up to now and which proved extremely successful, leading to the identification of about 150 {\\it INTEGRAL} objects\\footnote{The list of {\\it INTEGRAL} sources identified through optical and near-infrared spectroscopy can be found at {\\tt http://www.iasfbo.inaf.it/extras/IGR/main.html}}, that is about three quarters of the unidentified sources sample of the first three IBIS catalogues.  It is found that more than half of the sources identified by us are Active Galactic Nuclei (AGNs), followed by X--ray binaries ($\\sim$30\\%) and Cataclysmic Variables (CVs; $\\sim$13\\%); see \\cite{m10} and references therein.  Here we present the continuation of this follow-up and identification work on {\\it INTEGRAL} objects with the use of data collected in more than one year, between March 2010 and May 2011. ",
        "conclusions": "Here we presented the precise identification of a set of 19 sources of unknown or uncertain nature and which were detected by IBIS onboard {\\it INTEGRAL}. We found that nearly three quarters of them have an extragalactic nature, with redshifts in the range 0.024--3.12, while the remaining ones are Galactic sources. It is noteworthy that the majority of the latter ones are (possibly magnetic) CVs or Symbiotic stars.  These preliminary results further confirm the {\\it INTEGRAL} capabilities of detecting AGNs and hard X-ray emitting CVs."
    },
    "1112/1112.1919_arXiv.txt": {
        "abstract": "We develop statistical methods for identifying star streams in the halo of the Milky Way galaxy that exploit observed spatial and radial velocity distributions. Within a great circle, departures of the observed spatial distribution from random provide a measure of the likelihood of a potential star stream. Comparisons between the radial velocity distribution within a great circle and the radial velocity distribution of the entire sample also measure the statistical significance of potential streams. The radial velocities enable construction of a more powerful joint statistical test for identifying star streams in the Milky Way halo. Applying our method to halo stars in the Hypervelocity Star (HVS) survey, we detect the Sagittarius stream at high significance. Great circle counts and comparisons with theoretical models suggest that the Sagittarius stream comprises 10\\% to 17\\% of the halo stars in the HVS sample. The population of blue stragglers and blue horizontal branch stars varies along the stream and is a potential probe of the distribution of stellar populations in the Sagittarius dwarf galaxy prior to disruption. ",
        "introduction": "\\label{sec: Intro}  The timescale for halo stars to exchange their energy and angular momentum is long compared to the age of the Milky Way galaxy. Thus, the stellar halo of the Milky Way provides a unique laboratory for studying hierarchical galaxy formation. Theoretical simulations show that the remnant debris of hierarchical galaxy formation should be visible as star streams in the Milky Way halo \\citep{Johnston96, Harding01, Abadi03b, Bullock05, Font06}. Indeed, observers have discovered the tidal debris of the Sagittarius dwarf galaxy \\citep{Ibata94} wrapping around the Milky Way \\citep{Majewski03} and other stellar overdensities throughout the stellar halo \\citep[e.g.,][]{Belokurov06}. The vast majority of these detections come from star counts without radial velocity measurements.  We describe a powerful technique for identifying tidal star streams that exploits the additional information provided by a complete radial velocity survey. With large stellar radial velocity surveys such as RAVE \\citep{Steinmetz06, Zwitter08} and SEGUE \\citep{Yanny09} now available, approaches including radial velocities enable better constraints on structure in the halo.  To identify statistically significant tidal structures, we examine the surface number distribution and the radial velocity distribution of stars in great circles. Both the surface density and the radial velocity distributions probe the origin of structure in the halo. A radial velocity test for structure applies independently of the underlying spatial distribution. Our technique opens up the possibility of finding tidal star streams based on structure in velocity space. Combining our spatial and radial velocity statistical tests provides a more powerful joint test for identifying star streams.  Previous studies of positions and velocities have focused on known tidal streams, such as Sagittarius \\citep[e.g.,][]{Yanny09, Ruhland11}, or have analyzed disconnected patches scattered across the sky \\citep{Starkenburg09, Schlaufman09}.  Known tidal debris, however, is contiguous over areas of hundreds or thousands of square degrees \\citep{Johnston99,Ibata01}.  As a first demonstration of our technique, we search for structure in the HVS radial velocity survey, a complete, non-kinematically selected sample of late B-type stars covering a contiguous area of more than 8400 deg$^{2}$ on the sky \\citep{brown05, brown06b, brown09a, brown10}. The HVS survey targets blue stragglers and blue horizontal branch (BHB) stars within the halo and is distinct in its uniform velocity data for a well-defined sample of potential halo stars.  We detect significant structure in the sample of halo stars within the HVS survey.  Most of this structure is attributed to the Sagittarius stream. We compare our detections with the \\citet{Law10a} {\\it N}-body model for the Sagittarius dwarf, derive estimates for the fraction of HVS stars within the Sagittarius stream, and assess the relative fractions of blue stragglers and BHB stars along the stream.  In \\S \\ref{sec:IDStreams}, we describe the statistical technique. In \\S \\ref{sec:DataSection}, we apply this technique to the HVS survey. In \\S \\ref{sec:SgrBlueStars}, we compare the survey stars with the $N$-body model of \\citet{Law10a} in order to explore the nature and surface density of stars within the Sagittarius stream. We conclude in \\S \\ref{sec:Conclusion}.   \\begin{figure} \\begin{center} \\includegraphics[width=3.25in]{f1.eps} \\caption{ Solid angle, $\\Omega $, subtended by a great circle segment with arc length $\\protect\\theta $, width $\\Delta $, and pole at galactic coordinates $(l,b)$. \\label{fig: graphic} } \\end{center} \\end{figure} ",
        "conclusions": "\\label{sec:Conclusion}  We construct a statistical technique to detect star streams in the Galactic halo based on both the surface number density and the radial velocity distribution within great circles. The addition of radial velocities to the analysis increases our ability to identify star streams and to show that some overdensities in number counts are superpositions of several distinct structures in velocity space. The technique is straightforward yet powerful, can be generalized to diverse sets of data, and quantifies the statistical significance of tidal star streams. Detection is possible even when a structure contributes only a few percent of the stars in a sample.  We apply the method to the HVS survey and detect the Sagittarius stream at high significance. Comparison with theoretical models confirms that blue stragglers and BHB stars are members of the Sagittarius stream. Great circle counts and comparisons with theoretical models suggest that the Sagittarius stream comprises 10\\% to 17\\% of the halo stars in the HVS sample.  The ratio of blue stragglers to blue horizontal branch stars varies along the length of the stream, with roughly equal numbers of these two stellar types at RA = 0\\degree--50\\degree\\ and a 2:1 mix of blue stragglers to blue horizontal branch stars at RA = 130\\degree--180\\degree. Our conclusions support previous indications of variations in stellar population related to the original structure of the dwarf galaxy.  This technique is easily applied to other large radial velocity surveys including RAVE \\citep{Steinmetz06,Zwitter08} and SEGUE \\citep{Yanny09}. Identifying tidal star streams in these radial velocity surveys will ultimately improve constraints on the Milky Way's formation history, dark matter halo mass distribution, and the distributions of stellar populations within the progenitors."
    },
    "1112/1112.1128_arXiv.txt": {
        "abstract": "We  study the interaction between dark sectors by considering the momentum transfer caused by the dark matter scattering elastically within the dark energy fluid. Describing the dark scattering analogy to the Thomson scattering which couples baryons and photons, we examine the impact of the dark scattering in CMB observations. Performing global fitting with the latest observational data, we find that for a dark energy equation of state $w<-1$, the CMB gives tight constraints on dark matter-dark energy elastic scattering. Assuming a dark matter particle of proton mass, we derive an elastic scattering cross section of $\\sigma_D < 3.295 \\times 10^{-10} \\sigma_T$ where $\\sigma_T$ is the cross section of Thomson scattering. For $w>-1$, however, the constraints are poor. For $w=-1$, $\\sigma_D$ can formally take any value. ",
        "introduction": "There has been convincing evidence indicating that our universe is expanding at an increasing rate and it is mainly composed of dark matter (DM) and  dark energy (DE) at the present moment. The simplest explanation of DE is the cosmological constant, that fits observations very well. However, the observed value of the cosmological constant  falls far below the value predicted by any sensible quantum field theory. This is the so called cosmological constant problem. Furthermore using the cosmological constant to explain the DE unavoidably leads to the coincidence problem, namely, why are the vacuum and matter energy densities of precisely the same order today\\cite{Weinberg00}\\cite{Carroll01}?  Considering that DE and DM contribute significant fractions of the contents of the universe, it is natural, in the framework of field theory, to consider the interaction between them. The covariant description of the interaction between DE and DM is given by \\cite{Kodama} $\\nabla_{\\mu}T^{\\mu\\nu}_{(\\lambda)} = Q^{\\nu}_{(\\lambda)}$, where $Q^{\\nu}_{(\\lambda)}$ is a four vector governing the energy momentum transfer between DE and DM.  The subindex $\\lambda$ refers to DM or DE component respectively. For the whole system, the energy and momentum are conserved, and the transfer vector satisfies $\\sum_{\\lambda}Q^{\\nu}_{(\\lambda)} = 0$. The four vector $Q^{\\nu}_{(\\lambda)}$ describing the interaction between DE and DM can be phenomenologically decomposed into two parts with respect to a given observer $\\lambda'$  with four velocity $U^{\\mu}_{(\\lambda')}$ \\cite{He.prd83}, namely, $Q^{\\mu}_{(\\lambda)} = Q_{(\\lambda)}|_{\\lambda'}U^{\\mu}_{(\\lambda')} + F^{\\mu}_{(\\lambda)}|_{\\lambda'}$, where $Q_{(\\lambda)}|_{\\lambda'} = -U_{(\\lambda')\\nu}Q^{\\nu}_{(\\lambda)}$ is the energy transfer rate of the $\\lambda$ component observed by  the observer situated on the $\\lambda'$ component. $F^{\\mu}_{(\\lambda)}|_{\\lambda'} = h^{\\mu}_{(\\lambda')\\nu}Q^{\\nu}_{(\\lambda)}$ is the corresponding momentum transfer between the two components observed by the observer in the $\\lambda'$ frame.  The possibility that DE and DM interact with energy exchange between them has been studied extensively \\cite{He.prd83}-\\cite{Baldi.MNRAS11}. It has been shown that the energy transfer between DE and DM can provide a mechanism to alleviate the coincidence problem \\cite{He.prd83}-\\cite{He.prd80}. Complementary observational signatures of the energy exchange between DE and DM have been obtained from the cosmic expansion history by using the WMAP, SNIa, BAO and SDSS data etc\\cite{He.prd83}\\cite{He.prd80}-\\cite{Xu:PLB11} and the growth of cosmic structure \\cite{He.jcap07}-\\cite{Baldi.MNRAS11}.  Besides the energy transfer between DE and DM, their interaction may also impart a transfer of momentum. Considering the extremely low DE density, and the nonrelativistic velocities of DM motions, elastic scattering appears naturally between the DM and DE fluid.  We do not need to restrict to a particular DE model since the macroscopic physics is independent of the microphysics in the scattering process. The DM scattering elastically within the DE fluid results in the momentum transfer between them. The implications of the elastic scattering between DE and DM was explored in \\cite{Simpson}. It was found that the growth of structure was suppressed by a drag term arising from elastic scattering between dark sectors. It is of great interest to extend the study in \\cite{Simpson} to investigate the cosmological signal of the elastic scattering between the DM and DE fluid. In this work we will explore the implications of the momentum transfer between dark sectors in the cosmic background radiation and  constrain the scattering cross section between DE and DM by using the WMAP observation and the joint analysis of CMB, SNIa and BAO data. ",
        "conclusions": "In this paper we generalized the study of the interaction between DE and DM. Instead of considering the energy exchange between DE and DM in the literatures, in this work we investigated the momentum transfer by exploring the elastic scattering between dark sectors. Adopting the phenomenological description of the DM elastically scattering within the DE fluid \\cite{Simpson}, which is analogous to the Thomson scattering for baryons and photons, we examined the impact of the dark scattering on CMB observations. We found that in addition to the growth of structure \\cite{Simpson},  the elastic scattering between the DM and DE fluid can also leave cosmological signals in CMB. The imprint of the dark scattering in CMB highly depends on the equation of state of DE. When the DE equation of state $w<-1$, we can constrain the dark scattering well from CMB observations. However when $w>-1$, the cosmological signal of the scattering is weak. The impact of the dark scattering is suppressed when $w\\sim -1$.  From our result, we learnt that the CMB probes are not sufficient to constrain the dark scattering. It would be interesting to carry out the complementary constraints by combining the CMB test with the large scale structure tests as employed in the study of the energy transfer between dark sectors \\cite{He.jcap07}-\\cite{Baldi.MNRAS11}. Progress in this direction will be reported in the future."
    },
    "1112/1112.3893_arXiv.txt": {
        "abstract": "{ Cataclysmic Variables (CVs) are a distinct class of X-ray binaries transfering mass from a donar star to a compact star accretor, a white dwarf. They constitute a laboratory for accretion physics, mechanisms and disk theory together with dynamics of outflows and interaction with surrounding medium. Our understanding of the X-ray emission from CVs has shown substantial improvement in the last decade with the instruments onboard present X-ray telescopes. We have a better understanding of the boundary layers and the accretion shocks using high sensitivity instruments and/or grating X-ray spectroscopy attaining high spectral resolution yielding temperature and velocity diagnostics. In addition, high time resolution utitilizes accretion mode diagnostics with characterization of the variability of the gas flow in discs. Here, I briefly review the X-ray properties of CVs with a highlight on Intermediate Polars, Dwarf Novae and Classical/Recurrent Novae. ",
        "introduction": "A Cataclysmic Variable (CV) is a close interacting binary system in which a white dwarf (WD) accretes material from its late-type low mass main sequence companion. CVs can be studied in two main classes. The accretion occurs through an accretion disc in cases where magnetic field of the WD is weak or nonexistent ( $B$ $<$ 0.01 MG), such systems are referred as nonmagnetic CVs characterized by their eruptive behavior (see Warner 1995, Warner 2003 for a review). The other class is the magnetic CVs (MCVs) divided into two sub-classes according to the degree of synchronization of the binary. Polars have 230 MG$>$$B$$>$20 MG where this strong field causes the accretion flow to directly channel onto the magnetic pole/s of the WD inhibiting the creation of an accretion disk (see Cropper 1990; Warner 2003) causing the WD rotation to synchronize with the binary orbit. The second class of MCVs are the Intermediate Polars which have less field strength of 1-20 MG compared with Polars and are thus asynchronous systems (Patterson 1994; Warner 2003). The space density of CVs is (0.5-10)$\\times$10$^{-6}$ pc$^{-3}$ calculated using the X-ray luminosity function (a truncated power law) depending on the limiting X-ray flux (see Pretorius \\& Knigge 2011). The cumulative luminosity density of CVs in the 2-10 keV band is (0.8-1.4)$\\times$10$^{27}$ erg s$^{-1}$ M$_{\\odot}^{-1}$ (Sazonov et al. 2006, this value is similar in 17-60 keV range). ",
        "conclusions": "\\noindent {\\bf Simone Scaringi: }\\ Does the high frequency break move to higher frequencies when looking at higher energies ?\\\\ \\noindent {\\bf Solen Balman: }\\ No not expected. The $RXTE$, $XMM$-$Newton$ EPIC and OM-UV light curves show similar breaks within errors.\\\\ \\noindent {\\bf Vitaly Neustroev: }\\ Where do the orbital modulations in Intermediate Polars come from ?\\\\ \\noindent {\\bf Solen Balman: }\\ They are due to cold/warm absorption from the bulge at the accretion stream impact zone on the outer edge of the disk."
    },
    "1112/1112.4157_arXiv.txt": {
        "abstract": "We combine the output of hydrodynamical simulations of Population III star cluster formation with stellar evolution models, and calculate the evolution of protostars experiencing variable mass accretion rates due to interactions within a massive disk. We find that the primordial protostars are extended `fluffy' objects for the bulk of their pre-main-sequence lifetimes. Accretion luminosity feedback from such objects is high, but as shown in previous work, has a minimal effect on the star cluster. The extended radii of the protostars, combined with the observation of close encounters in the simulations, suggests that mergers will occur in such systems. Furthermore, mass transfer between close protostellar binaries with extended radii could lead to massive tight binaries, which are a possible progenitor of gamma ray bursts. ",
        "introduction": "It is becoming increasingly clear that Population III star formation is not as different from that in the present day as was once thought. Several authors have now shown that instead of a single massive star being formed at the centre of a primordial halo, a small cluster is formed (e.g.\\ \\citealt{Machida08}; \\citeauthor{Clark11a}~2011a,b; \\citealt{Turk09,Stacy10,Prieto11}).  There are frequent dynamical interactions within these clusters that can strongly alter the accretion rates of the protostars \\citep{Greif11}. As such, Population III star formation begins to resemble simulations of local clustered star formation \\citep[e.g][]{Bonnell11,Offner10,Girichidis11}. For a review of present day clustered star formation see \\citet{MacLow04}.  A parallel strand in the study of Population III star formation has been the investigation of the stellar evolution of the protostar as it grows in mass.  As the protostar grows, the consequences of its protostellar evolution become increasingly significant. For example, during the protostellar stage of the evolution, radiation from the protostar alters the thermal and chemical structure of the circumstellar disk. As a result, the nature of disk fragmentation via the gravitational instability is affected by the protostellar radiation (\\citeauthor{Clark11b}~2011b; \\citealt{Smith11b}). As the star grows in mass and approaches the main sequence, it also starts to emit UV radiation. This creates an HII region which can prevent further mass accretion onto the star and limit the final stellar mass (\\citealt{McKee08}; \\citeauthor{Hosokawa11b}~2011b; \\citealt{Stacy11c}). The stellar evolution also determines the  protostellar radius, which is an important quantity for understanding the dynamic evolution of a young cluster, as this determines the likelihood of protostellar collisions and mergers (for an analysis of present day protostellar clusters see \\citealt{Baumgardt11}).  Much of the pioneering work on the stellar evolution of Population III protostars was done by \\citet{Stahler86a} and \\citet{Omukai01,Omukai03} who numerically solved the interior structure equations assuming various constant accretion rates and showed that the evolution quantitatively changes with different rates. However, as was pointed out in \\citet[][hereafter S11]{Smith11b}, the protostars in a Population III minihalo do not accrete at a constant rate, but at a highly variable one, as the environment from which they accrete is not uniform. Time-variable mass accretion is also ubiquitous in present-day star formation \\citep[e.g.][]{Schmeja04,Vorobyov06,Machida11} and its effects on protostellar evolution have begun to be explored (\\citealt{Wuchterl01,Baraffe09,Baraffe10}; \\citeauthor{Hosokawa11}~2011a; \\citealt{Hartmann11}). In this paper, we seek to combine the strands of hydrodynamic simulations of fragmentation with stellar modelling and study how time-variable mass accretion affects the evolution of primordial protostars. ",
        "conclusions": "We have combined the accretion rates measured from Population III protostars in hydrodynamic simulations with a stellar evolution code to determine the structure of the protostars. The measured accretion rates had a large degree of variability and were particularly high during the adiabatic stage of the protostars' evolution. This increased the stellar radius of the protostars, causing them to become puffy extended objects. In the early stages of the protostar's life this intermittently increased the accretion rate of the protostar, but as shown in \\citet{Smith11b} this has only a small effect on the dynamic evolution of the halo. In some cases the separation of the protostars during their dynamic evolution was sufficiently close that mass transfer and mergers may occur between protostars. Mass transfer between massive protostellar systems could lead to approximately equal-mass massive binaries that are a possible precursor of GRBs. The extended size and relatively large tidal radii of Population III protostars implies that interactions between them are likely to dissipate more energy than if they were pure point masses."
    },
    "1112/1112.6364_arXiv.txt": {
        "abstract": "The most recent results of searches at the LHC for the Higgs boson $h$ have turned up possible hints of such a particle with mass $m_h^{}$ about 125\\,GeV consistent with standard model (SM) expectations. This has many potential implications for the SM and beyond. We consider some of them in the contexts of a simple Higgs-portal dark matter (DM) model, the SM plus a real gauge-singlet scalar field $D$ as the DM candidate, and a couple of its variations. In the simplest model with one Higgs doublet and three or four generations of fermions, for $D$ mass \\,$m_D^{}<m_h^{}/2$\\, the invisible decay \\,$h\\to DD$\\, tends to have a substantial branching ratio. If future LHC data confirm the preliminary Higgs indications, $m_D^{}$ will have to exceed~$m_h^{}/2$.  To keep the DM lighter than~$m_h^{}/2$, one will need to extend the model and also satisfy constraints from DM direct searches. The latter can be accommodated if the model provides sizable isospin violation in the DM-nucleon interactions. We explore this in a two-Higgs-doublet model combined with the scalar field~$D$. This  model can offer a 125-GeV SM-like Higgs and a~light DM candidate having isospin-violating interactions with nucleons at roughly the required level, albeit with some degree of fine-tuning. ",
        "introduction": "The latest searches for the standard model (SM) Higgs boson performed by the ATLAS and CMS Collaborations at the LHC have come up with tantalizing hints of the particle~\\cite{lhc}. They observed modest excesses of events compatible with a SM Higgs $h$ having mass $m_h^{}$ in the range from 124 to 126~GeV, but the statistical significance of the excesses is not enough for making any conclusive statement on the Higgs existence or nonexistence~\\cite{lhc}. Interestingly, these numbers are consistent with \\,$m_h^{}=125_{-10}^{+8}$~GeV\\, from the SM complete fit to electroweak precision data plus constraints from direct Higgs searches at LEP and Tevatron~\\cite{gfitter1}. Needless to say, if upcoming measurements confirm these preliminary findings at the LHC to be glimpses of the SM Higgs, or a SM-like Higgs, the implications will be far-reaching for efforts to extend the SM, as physics beyond the SM is still necessary to account for, among other things, the observed evidence for dark matter. Particularly, all new-physics models would have to include such a~spinless boson as one of their ingredients.  Simultaneously with the quest for the Higgs, a number of underground experiments have for years been searching directly for weakly interacting massive particle~(WIMP) dark matter~(DM) by looking for the recoil energy of nuclei caused by WIMPs colliding with nucleons. Intriguingly, some of these searches have turned up excess events which may have been WIMP signals. Specifically, the DAMA, CoGeNT, and CRESST-II Collaborations~\\cite{Bernabei:2010mq,Aalseth:2010vx,cresst} have acquired data that seem to be pointing to light WIMPs of mass in the region roughly from~5~to~30~GeV and spin-independent WIMP-nucleon scattering cross-sections of order \\,$10^{-42}$ to $10^{-40}$~cm$^2$,\\, although the respective ranges preferred by these experiments do not fully agree with each other. In contrast, other direct searches, especially by the CDMS, XENON, and SIMPLE Collaborations~\\cite{Akerib:2010pv,Angle:2011th,Aprile:2011hi,Felizardo:2011uw}, still have not produced any WIMP evidence. Although presently for WIMP masses under 15\\,GeV the latter null results are still controversial~\\cite{Hooper:2010uy} and future WIMP searches with improved sensitivity may eventually settle this issue definitively, there may be alternative explanations worth exploring which can reconcile these disagreeing findings on~DM.  Since the various DM searches employed different target materials for WIMP detection, one of the possible resolutions to the light-DM controversy that have been proposed is to allow large isospin violation in the WIMP-nucleon interactions~\\cite{Bernabei:2001ve,Kurylov:2003ra,Feng:2011vu}. It turns out that the tension between the conflicting direct-search results can be partially eased if the effective WIMP couplings $f_p^{}$ and $f_n^{}$ to the proton and neutron, respectively, satisfy the ratio~\\,$f_n^{}/f_p^{}\\simeq-0.7$\\,~\\cite{Kurylov:2003ra,Feng:2011vu,Kopp:2011yr}.  In view of these developments in the hunts for the Higgs and for~DM, it is of interest to consider some of their implications within the context of a relatively simple framework. For if the two sectors are intimately connected, detecting the signs of one of them could shine light on still hidden elements of the other.  The most economical model possessing both a Higgs boson and a WIMP candidate is the~SM+D, which is the SM expanded with the addition of a real gauge-singlet scalar field $D$ dubbed darkon acting as the~WIMP~\\cite{Silveira:1985rk,Burgess:2000yq,He:2008qm,He:2010nt,He:2011de}. This model predicts that the Higgs decay for \\,$m_h^{}\\sim120$-130\\,GeV\\, is highly dominated by the invisible mode \\,$h\\to DD$\\, if the darkon mass~\\,$m_D^{}<m_h^{}/2$,\\, except when $m_D^{}$ is close to, but less than,~$m_h^{}/2$\\,~\\cite{He:2008qm,He:2010nt,He:2011de}. Thus, if a~Higgs with \\,$m_h^{}\\sim125$\\,GeV\\,~and characteristics within the SM preference manifests itself unambiguously in LHC data, the parameter space of the SM+D with a light darkon will be strongly diminished. On the other hand, for \\,$m_D^{}\\ge m_h^{}/2$\\, the model is consistent with the existence of such a~Higgs, as its decay pattern is unaffected  at leading order by the presence of the darkon, although limited portions of this $m_D^{}$ zone are already excluded by direct-search data.  In order to have a Higgs compatible with the LHC indications as well as a light-WIMP candidate, one must therefore expand the~SM+D. This motivates us to study in this paper a~slightly extended model we call THDM+D, which is a two-Higgs-doublet model combined with a darkon.\\footnote{\\baselineskip=11pt% Various aspects of the THDM+D were previously addressed in~Refs.\\,\\cite{He:2008qm,Bird:2006jd,He:2007tt,Cai:2011kb}.} It can offer such a~Higgs and an ample amount of viable parameter space for the darkon. The model can also supply isospin-violating WIMP-nucleon interactions at about the desired level, although this will require fine-tuning to some extent. In addition, the potential presence of tree-level flavor-changing couplings of the neutral Higgs bosons in the THDM+D implies that the Higgs-mediated top-quark decays \\,$t\\to(u,c)DD$,\\, if kinematically allowed, would contribute to the decays \\,$t\\to(u,c)$ plus missing energy. If these decays have rates that are sufficiently amplified to be measurable at the~LHC, they can be another avenue to probe the darkon~\\cite{He:2007tt}.  In the next section, we consider in more detail some implications of the possible discovery of a~Higgs having SM-like properties for the SM+D with either three or four sequential generations of fermions (hereafter referred to as SM3+D or SM4+D, respectively). In~Sec.\\,\\ref{thdm+d} we describe the THDM+D with a SM-like Higgs and study its prediction for WIMP-nucleon cross-sections in the limit that isospin violation is negligible in the WIMP-nucleon effective couplings. In~Sec.\\,\\ref{ivdm} we treat the THDM+D case with isospin-violating WIMP-nucleon interactions. We give our conclusions in~Sec.\\,\\ref{concl}. ",
        "conclusions": "}  The preliminary indications of a Higgs boson with SM-like properties in the latest LHC data, if confirmed by future measurements, will have important implications for WIMP dark matter models. We have explored a number of such implications for some of the simplest darkon models. For the simplest one, SM3+D, most of the light-darkon mass range will be ruled out if an SM3-like Higgs with mass near 125\\,GeV is found. Such a discovery would also be at odds with the SM4 Higgs expectations, thus disfavoring the SM4+D. In contrast, the type-III two-Higgs-doublet model enlarged with the addition of a darkon has an abundance of allowed parameter space in its DM sector. It can accommodate an~SM3-like Higgs and simultaneously offers a WIMP candidate in harmony with the light-WIMP hypothesis inspired by the clues from a number of DM searches, although it is in conflict with the null results of other searches. The model can also provide substantial isospin violation in the WIMP-nucleon interactions which can alleviate some of this tension in the data on light DM. However, this could be achieved only with some amount of fine-tuning in several of the relevant parameters. Nevertheless, upcoming searches for the Higgs at the LHC and future DM direct searches can test further the THDM+D which we have considered. Additional signals of the darkon in this scenario may be available from flavor-changing top-quark decays into lighter up-type quarks with missing energy which are potentially observable at the~LHC."
    },
    "1112/1112.3294_arXiv.txt": {
        "abstract": "We report on the characterization of candidate light sensors for use in the next-generation Imaging Atmospheric Cherenkov Telescope project called Cherenkov Telescope Array, a major astro-particle physics project of about 100 telescopes that is currently in the prototyping phase. Our goal is to develop with the manufacturers the best possible light sensors (highest photon detection efficiency, lowest crosstalk and afterpulsing). The cameras of those telescopes will be based on classical super-bi-alkali Photomultiplier tubes but also Silicon Photomultipliers are candidate light sensors. A full characterisation of selected sensors was done. We are working in close contact with several manufacturers, giving them feedback and suggesting improvements. ",
        "introduction": "\\label{aba:sec1} The atmosphere is opaque for cosmic rays. But when particles with very high energy (VHE) of at least some tens of GeV hit the atmosphere they produce extended showers of secondary particles. These emit blueish Cherenkov light flashes due to their high speed, exceeding that of light in the atmosphere. The faint flashes can be collected by a telescope and projected onto it's fine pixelised camera. This method is called Imaging Atmospheric Cherenkov Telescope (IACT) technique and it is currently the most successful one with over 120 sources of VHE gamma-rays discovered \\cite{ TevCat}. The camera consists of hundreds of ultra fast and highly sensitive light sensors, fast enough to follow the development of the shower\\cite{ Ostankov01}. All the current telescopes use Photomultiplier Tubes (PMT) --- with the exception of the FACT telescope\\footnote{Their approach is to use Silicon Photomultipliers (SiPM) from Hamamatsu (MPPC)\\cite{ Anderhub11}}. ",
        "conclusions": "The PMT development for CTA is in progress and already now has almost achieved the specified targets. During the next two years further optimization of PMTs from both Hamamatsu and Electron Tubes Enterprises is foreseen.  SiPM development is ongoing. We consider their use interesting in many applications when the PDE of SiPM will become 1.5-2 times higher than that of PMTs at comparable costs. New MEPhI SiPMs are getting there with peak PDE $\\simeq$ 50\\% at $P_{crosstalk}$ $\\simeq5\\%$. Soon these are going to become a commercial product. In a couple of years this type of sensor could become a serious alternative to the PMT in CTA."
    },
    "1112/1112.2151_arXiv.txt": {
        "abstract": "The discovery of radio pulsars in compact orbits around \\sgr\\ would allow an unprecedented and detailed investigation of the spacetime of the supermassive black hole. This paper shows that pulsar timing, including that of a single pulsar, has the potential to provide novel tests of general relativity, in particular its cosmic censorship conjecture and no-hair theorem for rotating black holes. These experiments can be performed by timing observations with 100\\,$\\mu$s precision, achievable with the Square Kilometre Array for a normal pulsar at frequency above 15\\,GHz. Based on the standard pulsar timing technique, we develop a method that allows the determination of the mass, spin, and quadrupole moment of \\sgr, and provides a consistent covariance analysis of the measurement errors. Furthermore, we test this method in detailed mock data simulations. It seems likely that only for orbital periods below $\\sim 0.3$\\,yr is there the possibility of having negligible external perturbations. For such orbits we expect a $\\sim 10^{-3}$ test of the frame dragging and a $\\sim 10^{-2}$ test of the no-hair theorem within 5 years,  if \\sgr\\ is spinning rapidly. Our method is also capable of identifying perturbations caused by distributed mass around \\sgr, thus providing high confidence in these gravity tests. Our analysis is not affected by uncertainties in our knowledge of the distance to the Galactic center, $R_0$. A combination of pulsar timing with the astrometric results of stellar orbits would greatly improve the measurement precision of $R_0$. ",
        "introduction": "\\label{sec:intro}  One of the most intriguing results of general relativity (GR) is the uniqueness theorem for the stationary black hole solutions of the Einstein-Maxwell equations \\citep[see][and references therein]{heu98}. This uniqueness theorem states that (under certain conditions) all stationary electrovac\\footnote{Electrovac spacetimes are the solutions of the Einstein-Maxwell equations.} black hole spacetimes with a non-degenerate horizon are described by the Kerr-Newman metric. It implies that in GR all stationary black holes are parameterized by only three parameters: mass ($M$), spin ($S$) and electric charge (``black holes have no hair''). All uncharged black hole solutions are described by the Kerr metric and, therefore, uniquely determined by $M$ and $S$. Astrophysical black holes are believed to be the result of a gravitational collapse. During this collapse all the properties of the progenitor, apart from mass and spin, are radiated away by gravitational radiation while the gravitational field asymptotically approaches its stationary configuration \\citep{pri72a,pri72b}. The outer spacetime of an astrophysical black hole should therefore be described by the Kerr metric.\\footnote{Strictly speaking, this is only true for a certain approximation since, to some extent, astrophysical black holes will be influenced by nearby masses (accretion, orbiting objects). We will address this issue for \\sgr\\ in this paper at the end of the discussion Section.} Since the Kerr metric has a maximum spin at which it still exhibits an event horizon, Penrose's cosmic censorship conjecture (CCC) within GR \\citep{pen79} requires the dimensionless spin parameter $\\chi$ to satisfy \\begin{equation} \\chi \\; \\equiv \\; \\frac{c}{G}\\,\\frac{S}{M^2} \\; \\le 1 \\; \\end{equation} A measured value for $\\chi$ that exceeds 1 would pose a serious problem for our understanding of spacetime, since this would indicate that either GR is wrong or that a region may be visible to the outside universe, where our present understanding of gravity and spacetime breaks down.  As a result of the no-hair theorem, all higher multipole moments ($l \\ge 2$) of the gravitational field of an astrophysical black hole can be expressed as a function of $M$ and $S$ \\citep{han74}. In particular, the dimensionless quadrupole moment $q$ satisfies the relation \\begin{equation} q \\; \\equiv \\; \\frac{c^4}{G^2}\\,\\frac{Q}{M^3} \\; = \\; -\\chi^2 \\;. \\end{equation} A measurement of the quadrupole moment, in combination with a mass and a spin measurement, would therefore provide a test of the no-hair theorem for Kerr black holes.  Some of the clearest evidence for the existence of black holes comes from the monitoring of $\\sim 30$ stellar orbits in the center of our Galaxy \\citep{sog+02,gsw+08,get+09}, where the shortest orbital period, $P_{\\rm b}$, is 16 years. Known by its radio nomenclature of \\sgr, current estimates put the mass of this black hole to around $4 \\times 10^6\\,\\msun$. A black hole of that size at a distance of 8\\,kpc is an ideal laboratory for black-hole physics, strong-field gravity and, in particular, a test of the no-hair theorem for Kerr black holes \\citep[e.g.][]{psa08,pj10}. It has been shown by \\cite{wil08} that the discovery of stars in highly eccentric ($e\\sim 0.9$) orbits very close to \\sgr\\ ($P_{\\rm b} \\la 0.1\\,{\\rm yr}$) could be used to test the general relativistic no-hair theorem. This experiment requires an astrometric precision at the level of 10\\,$\\mu$as, which seems achievable with the upcoming infrared-astrometry experiments, such as GRAVITY \\citep{epb+09}. At a distance of 8\\,kpc, an angle of 10\\,$\\mu$as corresponds to a length scale of $\\sim 10^7$\\,km.  On the other hand, if close-in pulsars could be found and tracked in their orbits, even for those with poor timing precision the time-of-arrival (TOA) for a (integrated) pulsar signal can be measured with an uncertainty of a few milliseconds, corresponding to a light-travel distance of $\\sim 10^3$\\,km. Moreover, a phase-connected solution with an appropriate timing model allows a determination of the pulsar orbit, which provides even more precision. Hence, as shown by \\cite{wk99} and \\cite{kbc+04}, a pulsar in orbit around the supermassive black hole in the Galactic center (GC) would provide an ideal probe to measure the mass, the spin, and the quadrupole moment of \\sgr, and consequently test the no-hair theorem for Kerr black holes. In their discussion on ``bumpy black holes'', \\cite{vh10} suggest that a pulsar in orbit around a black hole could be used for mapping its multipole moment structure. In a recent publication \\cite{asm10} have discussed the importance of frame dragging and the quadrupole moment for stars and pulsars in orbit around \\sgr, based on numerical integration of geodesics in a Kerr spacetime. For pulsars, however, the results in this paper are only indicative.\\footnote{The \\cite{asm10} results for a pulsar are not derived from a consistent covariance analysis based on a timing model that incorporates all the relevant effects simultaneously. Moreover, pulsar timing is treated as a radial velocity measurement experiment, which is incorrect. In fact, pulsar timing makes use of phase-connected solutions for the rotational phase of a pulsar leading to a precision in the parameter estimations that can be several orders of magnitude better \\citep{lk05}. Also, the timing precision assumed by \\cite{asm10} seems too optimistic for a GC pulsar, as will become clear from the analysis presented in this paper.} \\citet{wcp+09,wjc+09} have shown that if a pulsar is found in a very close orbit around \\sgr\\ (e.g.~$P_{\\rm b} \\la 1\\,{\\rm day}$), observers at the Earth can receive additional pulses travelling along a path that is strongly bent by the gravitational potential of the black hole. The exploitation of this information would provide unique constraints on the strongly curved spacetime geometry near \\sgr.   In this paper we will demonstrate how the mass, the spin (magnitude and orientation) and the quadrupole moment of \\sgr\\ can be determined from timing a pulsar in a sufficiently tight orbit around the supermassive black hole. Our analysis is based on simulated TOAs and a timing model that allows for a phase-connected solution, consistently accounting for the relativistic effects in the motion of the pulsar and the propagation of the radio signals. Based on this, we can determine the expected precision for the individual parameters while accounting for all the correlations between the parameters. In all discussions and simulations we focus on the leading order in the individual effects of interest. We are well aware that in an actual timing model for orbits close to \\sgr\\ we need to account for higher order effects in the orbital motion and signal propagation. For many of the effects discussed below, higher order terms have already been calculated \\citep[e.g.][]{ds88, wex95, kop97b, kg05}. However, to estimate the expected precision and covariances in the parameter determination, it is sufficient to use a timing model that combines just the leading terms of all contributions relevant here. The second order terms contribute at a $\\sim\\beta_{\\rm O}^2$ level, where $\\beta_{\\rm O}$ is the orbital velocity parameter introduced by \\cite{dt92}. For a test particle in orbit around a mass $M$ \\begin{eqnarray} \\beta_{\\rm O} &=& \\left(\\frac{2\\pi GM}{c^3P_{\\rm b}}\\right)^{1/3} \\nonumber \\\\ &\\simeq& 0.0158 \\left(\\frac{M}{4\\times 10^6\\,M_\\odot}\\right)^{1/3} \\left(\\frac{P_{\\rm b}}{1\\,{\\rm yr}}\\right)^{-1/3}\\;. \\end{eqnarray} As an example, for an orbital period of 0.1\\,yr we find $\\beta_{\\rm O}^2 \\approx 10^{-3}$. At this point it is worth mentioning, that for the most relativistic binary pulsar, the double pulsar, one finds $\\beta_{\\rm O}^2 \\simeq 4.3\\times 10^{-6}$ \\citep{kw09}, which nicely illustrates how much more relativistic a pulsar in a $P_{\\rm b} \\la 1\\,{\\rm yr}$ orbit around \\sgr\\ would be.  Previous studies suggest that about 1000 pulsars can be expected to be orbiting \\sgr\\ with periods less than 100 years \\citep{cl97,gso+03,pl04}\\footnote{The result from the GC survey by \\cite{mkf+10} indicated that the actual number of such pulsars may be less.}, and some of them may be associated with remnants of the observed S-star population in the neighborhood. Deep pulsar searches towards the GC region were already conducted with a few radio telescopes (Effelsberg, Green Bank, Parkes) at frequencies up to 15\\,GHz \\citep[e.g.][]{kkl+00,mkf+10}. Five pulsars were found no more than 200\\,pc away from \\sgr\\ \\citep[][Kramer et al.~private communication]{jkl+06,dcl09}, which is consistent with the estimated large pulsar population within that region. However, they are not close enough to the supermassive black hole to probe its gravitational field. In this paper we will focus on timing observations of such pulsars, and show how far they could take us in probing the gravitational field of \\sgr, provided the system is found to be sufficiently free of external perturbations. The main purpose therefore is the development of the methodology, and the estimation of its potentials in testing the Kerr nature of \\sgr\\ based on mock data simulations. For further elaboration on either the existence of pulsars in close orbits around \\sgr\\ or the search for them we refer to the rich literature \\citep{lk05} and future work in progress. In Section~\\ref{sec:timing} we discuss the expected timing precision and the orbital periods required for our measurements. Section~\\ref{sec:mass} presents the various relativistic effects that can be used to determine the mass of \\sgr\\ and, based on simulations, the expected precision in the mass measurement. In Section~\\ref{sec:spin} we show how the spin can be extracted from the timing measurement, how this information can be used to test the CCC, and how the presence of a distributed mass in the vicinity of \\sgr\\ would affect this measurement. Section~\\ref{sec:quadrupole} provides the details on the quadrupole measurement and the no-hair theorem test. In Section~\\ref{sec:discussion} we summarize and discuss our findings. ",
        "conclusions": "\\label{sec:discussion}  In this paper we have developed a method to determine the mass, the spin, and the quadrupole moment of \\sgr\\ using a pulsar in a compact orbit around this super-massive black hole. Our investigation is based on a consistent timing model, that includes all the relativistic and precessional effects that can be used to extract these parameters of \\sgr. Based on simulated timing data for a pulsar in orbit around \\sgr, we have shown in a consistent covariance analysis, that even with a moderate timing precision ($\\sim 100\\,\\mu{\\rm s}$), one can expect to be able to determine the mass, the spin, and the quadrupole moment of \\sgr\\ with high precision, provided the orbital period is well below one year. As a result of our simulations, for a compact orbit (orbital period of a few months) one can expect to measure the spin with a precision of $10^{-3}$, or even better. We have shown how the method would allow the identification of an object whose frame-dragging exceeds that of an extreme Kerr black hole, and therefore would provide a test of the CCC. Furthermore, for such orbits the determination of the quadrupole moment of \\sgr\\ seems feasible at a few percent precision level or even better, depending on the size and orientation of the pulsar orbit and the spin of \\sgr. In combination with the precise spin measurement from the Lense-Thirring effect, this yields a high precision test of the no-hair theorem of stationary black holes.  Moreover, we have shown that, in general, our analysis will be able to unveil the presence of external perturbations caused by the presence of distributed mass, therefore providing high confidence in a spin and quadrupole determination based on pulsar timing. Nevertheless, further studies are required to see whether the spin and quadrupole moment can still be extracted if the timing data is ``contaminated'' by external perturbations. If perturbations arise from a smooth concentration of dark matter particles in the vicinity of \\sgr, we may be able to learn something about the properties of dark matter that clusters around \\sgr, assuming GR is correct.  Finally, we need to emphasize that the tests presented are not affected by an uncertainty in the distance to the GC. On the contrary, a mass determination via pulsar timing would give a greatly improved value for $R_0$ if combined with the astrometric measurements in the near infrared.  Once a pulsar is detected in a compact orbit around \\sgr, continuous timing will allow more and more measurements and tests as the timing baseline grows with time. In the following we summarize the most important steps in this experiment: \\begin{itemize} \\item After timing one orbit, all Keplerian parameters will be well known, and also the pericenter advance will be measured with good precision. This will already provide a good estimate of the mass of \\sgr. \\item Timing a few more orbits would then allow the determination of additional post-Keplerian parameters, like the Shapiro parameters ($r_{\\rm Sh}$, $\\sin i$) and the amplitude of the Einstein delay ($\\gamma_{\\rm E}$). These parameters allow a robust determination of \\sgr\\ mass and the inclination of the pulsar orbit with respect to the line-of-sight. \\item At this stage, a measurement of a change in the projected semi-major axis ($\\dot{x}$) caused by the Lense-Thirring will allow an early test of the CCC, and mapping of the region in the $\\chi_\\theta-\\chi_\\lambda$ plane where the solutions for Kerr black holes are. \\item Around the same time the mass measurement should reach a precision that allows the extraction of the Lense-Thirring contribution to the precession of the pericenter ($\\dot{\\omega}_{\\rm LT}$), giving a line-like region in the $\\chi_\\theta-\\chi_\\lambda$ plane. \\item After a few years of timing the second time derivatives of $\\omega$ and $x$ should be known with high precision, allowing a precise determination of the \\sgr\\ spin (magnitude and direction). At this stage we also have a test for the ``cleanness'' of the system, and whether the spin is below the Kerr bound ($\\chi = 1$). \\item At the same time the obtained parameters for the pulsar orbit and the \\sgr\\ spin can be used to model the periodic features in the timing residuals, which are caused by the quadrupole moment of \\sgr. This allows to a determination of the quadrupole moment and a test of the no-hair theorem. \\end{itemize}  A potential problem for the timing of a pulsar in a compact orbit around \\sgr\\ is posed by the relativistic spin precession, as pointed out by \\cite{mamw10}. This change in the pulsar orientation with respect to a distant observer not only causes a variation of the pulse profile, which makes precise timing more difficult, but also can turn the pulsar emission away from our line-of-sight \\citep{wrt89,kra98}. To leading order the spin-precession is given by the de Sitter precession rate, which for $M_{\\rm BH} \\gg M_{\\rm PSR}$ reads \\citep{bo75b} \\begin{equation} \\Omega_{\\rm dS} \\simeq \\frac{3\\pi}{P_{\\rm b}}\\,\\beta_e^2 \\simeq (0.13^{\\circ}{\\rm /yr})\\,\\frac{1}{1 - e^2} \\left(\\frac{P_{\\rm b}}{1\\,{\\rm yr}}\\right)^{-5/3} \\;. \\end{equation} Consequently, for orbital periods below one year relativistic spin precession is expected to play an important role in the timing observations.  We note in passing, that the Pugh-Schiff precession rate caused by frame dragging \\citep[$\\Omega_{\\rm FD} \\sim  2\\pi \\beta_e^3\\chi/P_{\\rm b}$,][]{pug59,sch60} is only relevant in the case of very compact, highly eccentric orbits ($\\Omega_{\\rm FD} \\sim 1^{\\circ}$/yr for $P_{\\rm b} = 0.1$\\,yr, $e=0.8$ and $\\chi=1$), and could provide an independent test of the rotation of \\sgr.  The no-hair theorem test can also be affected by the accretion disc around \\sgr. To get an idea of the strength of this effect, one can estimate the influence by calculating the fraction of the quadrupolar potential of the disc to that of the black hole. This ratio is given by \\begin{equation} \\mathscr{R}\\sim M_{\\rm disc} r^2_{\\rm disc}/(M_{\\rm BH}r^2_{\\rm g})\\;, \\end{equation} where $M_{\\rm disc}$ and $r_{\\rm disc}$ are the mass and outer radius of the disc, and $r_{\\rm g} \\equiv GM_{\\rm BH}/c^2$ is the black hole gravitational radius. Following the advection-dominated accretion flow (ADAF) model of \\cite{ych+09} and adopting, as an upper limit, the disc scale of $\\approx 1\\,{\\rm arcsec}$ determined from X-ray observation \\citep{bbb+01}, we obtain $\\mathscr{R} \\approx 0.4\\%$, which indicates that the quadrupole moment measurement of \\sgr\\ would not be biased by the contribution of the disc above the $1\\%$ precision level. Furthermore, in a very resent publication, based on current X-ray and millimeter observations, \\cite{psa11} concludes that for compact orbits, like the ones discussed in this paper, hydrodynamic drag forces from plasma in the vicinity of \\sgr\\ are expected to be negligible."
    },
    "1112/1112.0012_arXiv.txt": {
        "abstract": "*{We report on the discovery of the first BL~Herculis star displaying period-doubling behaviour as predicted by the hydrodynamic models of Buchler \\& Moskalik \\cite{BM92}. The star, with $P_0=2.4$\\thinspace d, is located in the Galactic bulge and was discovered with OGLE-III photometry. We present new nonlinear convective models, which, together with recent evolutionary tracks, put constraints on the stellar parameters. In particular, we estimate the mass and metallicity of the object.}  \\abstract{We report on the discovery of the first BL~Herculis star displaying period-doubling behaviour as predicted by the hydrodynamic models of Buchler \\& Moskalik \\cite{BM92}. The star, with $P_0=2.4$\\thinspace d, is located in the Galactic bulge and was discovered with OGLE-III photometry. We present new nonlinear convective models, which, together with recent evolutionary tracks, put constraints on the stellar parameters. In particular, we estimate the mass and metallicity of the object.} ",
        "introduction": "BL~Herculis stars are a subgroup of type II (or Population II) Cepheids, pulsating with periods between $1$ and $4$ days (see \\cite{Wal} for a recent review). They show a singly-periodic large-amplitude light variation. Such behaviour was qualitatively reproduced with early nonlinear radiative models, e.g. Buchler \\& Moskalik \\cite{BM92, MB93}. In several of their models with periods between $2$ and $2.6$ days, Buchler \\& Moskalik \\cite{BM92} found period-doubling behaviour -- oscillations with periodic alternations of deep and shallow minima -- a phenomenon not observed in any BL~Her star at that time. Here we report on the discovery of the first BL~Her star clearly showing period-doubling behaviour.  We summarize the observations and new pulsation models for this star. A more detailed analysis, as well as discussion of another BL~Her star in which period-doubling behaviour is strongly suspected can be found in Smolec et al. \\cite{SmolecEtal11}. ",
        "conclusions": ""
    },
    "1112/1112.3136.txt": {
        "abstract": "Many observers in the past gave detailed descriptions of the telescopic aspect of Venus during its extremely rare transits across the Solar disk. In particular, at the ingress and egress, the portion of the planet's disk outside the Solar photosphere has been repeatedly perceived as outlined by a thin, bright arc (''aureole''). Those historical visual observations allowed inferring the existence of Venus' atmosphere, the bright arc being correctly ascribed to the refraction of light by the outer layers of a dense atmosphere. On June 8$^{th}$, 2004, fast photometry based on electronic imaging devices allowed the first quantitative analysis of the phenomenon. Several observers used a variety of acquisition systems to image the event -- ranging from amateur-sized to professional telescopes and cameras -- thus collecting for the first time a large amount of quantitative information on this atmospheric phenomenon. In this paper, after reviewing some elements brought by the historical records, we give a detailed report of the ground based observations of the 2004 transit. Besides confirming the historical descriptions, we perform the first photometric analysis of the aureole using various acquisition systems. The spatially resolved data provide measurements of the aureole flux as a function of the planetocentric latitude along the limb. A new differential refraction model of solar disk through the upper atmosphere allows us to relate the variable photometry to the latitudinal dependency of scale-height with temperature in the South polar region, as well as the latitudinal variation of the cloud-top layer altitude. We compare our measurements to recent analysis of the Venus Express VIRTIS-M, VMC and SPICAV/SOIR thermal field and aerosol distribution. Our results can be used a starting point for new, more optimized experiments during the 2012 transit event. ",
        "introduction": "Since the mid 18$^{th}$ Century, observers have reported unusual features of the telescopic image of Venus near the inferior conjunction, promptly attributed to its atmosphere. Some of them are at the reach of modest instruments (although at small angular distance from the Sun), such as the cusp extension first described by \\citet{Schroeter791}, which tends to transform the thin crescent of Venus into a ring of light \\citep[e.g.][]{Russell99, Dollfus65}.  One of the most relevant features pertaining to ground-based studies of the Venus atmospheric structure has been observed only during transits, close to the phases of ingress (between 1st and 2nd contact) ot egress (between 3rd and 4th contact) as a bright arc outlining -- in part or entirely -- the portion of Venus' disk projected outside the solar photosphere. Traditionally the first account of this phenomenon was attributed to Mikhail V. Lomonosov (1711--1765) who reported his observations of the transit at St. Petersburg Observatory on May 26, 1761~\\citep{Marov05}. Actually the poor performance of his small refractor hints that most probably other observers (such as Chappe d'Auteroche, Bergman, and Wargentin) were the first genuine witness on the same date \\citep{Link69, Pasachoff12}. However, Lomonosov correctly attributed the putative phenomenon to the presence of an atmosphere around the planet, refracting the sunlight in the observer's direction.  In the following, adopting a denomination widely used in the historical accounts, we will often call this arc ``aureole''. Since both the aureole and the cusp extension occur close to the planet terminator, they are also collectively known as ''twilight phenomena''. For detailed historical reviews the interested reader can refer to \\citet{Link69} and \\citet{Edson63}. Of course, in this context we neglect the initial motivation for the observation of transits: the determination of the Astronomical Unit by the solar parallax (proposed by E. Halley in 1716), whose interest is purely historical today.  Venus transits are rare, as they occur in pairs 8 years apart, each pair separated by 121.5 or 105.5 years, alternating between descending node (June pairs: 1761/1769, 2004/2012), and ascending node (December pairs: 1631/1639, 1874/1882, 2117/2125). As a consequence, data concerning the aureole are correspondingly sparse and, up to the last event, they have been obtained by simple visual inspection, mainly by using refracting telescopes of modest aperture (typically up to 15-20 cm). This limitation was mainly due to the constraint of organizing complex expeditions including the delicate transportation of all the instruments.  The 2004 event represents a giant leap in the observation of Venus transits, as the modern imaging technologies available allow for the first time a quantitative analysis of the atmospheric phenomena associated to the transit of Venus.  We performed measurements of the aureole on the original images obtained through several different instruments, and compared them both with a simple refraction model and with observations obtained in the past. This work summarizes the aspect of Venus close to the Sun's limb during the June 8, 2004 transit as observed by ground--based instruments. A recent paper \\citep{Pasachoff11} deals with imaging using NASA's then operating Transition Region and Coronal Explorer solar observatory (TRACE).  Since the ground-based observations were not specifically organized beforehand, data at our disposal are rather heterogenous. In order to bridge the gap between past visual observations through small telescopes and today's technologies, we decided to consider accounts obtained both with professional instruments and low-cost amateur telescopes, either by CCD imaging or by direct image inspection by experienced observers. In fact, as shown in recent studies of distant solar system objects based on stellar occultation campaigns (e.g. Widemann et al., 2009) while CCD imaging offers today the most valuable quantitative measurements, small telescopes and visual observations allow a significant increase of constrains on the phenomena  - and in the case of Venus, the most direct comparison to past reports. The results obtained from the analysis of the most significant image sets, representing a certain variety of instrument size and quality, are illustrated in this paper.  The paper is organized as follows. First, we describe the conditions of the 2004 event, and the reconstruction of limb geometry (section~\\ref{S:geom}). We then provide an extensive review of the measured spatial and temporal variations for the brightness of the aureole in wide or narrow-band photometry (section~\\ref{S:obs} and \\ref{S:color}). We then address the basic physical principles of the atmospheric differential refraction model producing the aureole (section~\\ref{S:model}) and we use it for the interpretation of the observations (\\ref{S:application}). In Section \\ref{S:VeX}, the modeling of the aureole is compared to recent analysis of the Venus Express observations (VIRTIS-M, VeRa and SPICAV/SOIR) regarding the thermal field and cloud-deck altitude and haze distribution pertaining to this study, as well as ground-based mid-infrared spectroscopy of non-LTE CO$_2$ emission. The comparison is discussed. % In conclusion, we set grounds for experiments during the coming 2012 event. ",
        "conclusions": "Our results represent the first successful model of the aureole of Venus, observed during solar transits. We are able to reproduce the main features of the lower mesosphere observational constraints : longitudinally averaged thermal structure near 117 km, slanted opacity of aerosols and their meridional variation, with consistent physical parameters.  The results obtained by this first set of measurements are encouraging and suggest that more accurate planning will produce more precise data. Improvements are also possible concerning the refraction model, which could be completed by a more realistic, gradual transition from transparency to absorption at the cloud deck level based on most recent results on the latitudinal distribution of upper haze aerosols. Also, the contribution of scattered light in addition to refracted light could be included. In fact, the absorption curves in \\citet{Wilquet09} Fig. 3 are wavelength dependent, i.e. the opaque layer altitude in our model can be function of the color. As a result the aureole should not be perfectly ''grey'', suggesting that future observations capable of accurate multicolor photometry should investigate this possibility.Even a grey absorption would cause the solar flux in the aureole to look dimmer than without aerosols, which would be wrongly interpreted as a drop in scale height, i.e. in local temperature at tangent point altitude, in the inversion code.  In the two hemispheres, the cold collar temperature structure and the polar regions are very similar~\\citep{Tellmann09}, so the latitudinal variations obtained for the 2004 transit around the South polar region might tentatively apply to the North polar data we expect to acquire in June 2012. Recent interpretations also attest a general N-S symmetry in the latitudinal variation of aerosols extinction \\citep{Wilquet11b}. Given the difference in brightness between the arc and the solar photosphere, accurate measurements of this phenomena are challenging. The degree of turbulence, the magnification, the amount of light scattering in the optics are all elements that contribute in determining the effective visibility of the aureole in a given instrument, and the accuracy of the measurements. The transit of June 2012 will be our last opportunity for observing the aureole using this spatially resolved technique, until the next pair of transits of Venus, which will be in the ascending node, on 11 December 2117 and 8 December 2125."
    },
    "1112/1112.4659.txt": {
        "abstract": "text{% \\noindent ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.3668.txt": {
        "abstract": "We discuss the peculiar nature of the nucleus of M94 (NGC 4736) in the context of new measurements of the broad H$\\alpha$ emission from $HST$-STIS observations.  We show that this component is unambiguously associated with the high-resolution X-ray, radio, and variable UV sources detected at the optical nucleus of this galaxy.  These multi-wavelength observations suggest that NGC 4736 is one of the least luminous broad-line (type 1) LINERs, with $L_{\\rm bol} = 2.5 \\times 10^{40}$~erg s$^-1$.  This LINER galaxy has also possibly the least luminous broad line region known ($L_{H \\alpha} = 2.2 \\times 10^{37}$~erg s$^{-1}$).  We compare black hole mass estimates of this system to the recently measured $\\sim 7\\times 10^{6}~M_\\sun$ dynamical black hole mass measurement.  The fundamental plane and $M - \\sigma^*$ relationship roughly agree with the measured black hole mass, while other accretion based estimates (the $M-FWHM(H\\alpha)$ relation, empirical correlation of BH mass with high-ionization mid IR emission lines, and the X-ray excess variance) provide much lower estimates ($\\sim 10^5 M_\\sun$).  An energy budget test shows that the AGN in this system may be deficient in ionizing radiation relative to the observed emission-line activity.  This deficiency may result from source variability or the superposition of multiple sources including supernovae. ",
        "introduction": "%galaxies and active BHs Most high mass galaxies are known to host massive black holes (BH), some passively lurking in their centers while others are actively accreting surrounding material (e.g., G\\\"{u}ltekin et al. 2009a and references therein).  The mechanism that causes a BH's activity to turn on and off is still largely unknown. %There is no clear discernment of what causes BH's activity to turn on and off. Understanding the structure of the active galactic nuclei (AGN) at their lowest luminosities is crucial to determining the physical and possibly evolutionary links between the most luminous galaxy centers and the passive ones.  However, at low luminosities it is difficult to disentangle the various emission mechanisms that could be concurrently present in galaxy centers.  As a consequence, the dominant power source of a large majority of actively line-emitting galaxies remains ambiguous (e.g., Ho 2008 for a review).  %bpt diagrams, kinda failing in recognizing LINERs as AGN Diagnostic diagrams \\citep{bpt, vei87, kau03, kew06} are relatively successful in separating out bona-fide accretion sources (Seyferts) from nuclei whose emission-line activity is mainly powered by young, hot stars (H {\\sc ii} galaxies), based on emission line ratios.  At least 50\\% of the strong line emitters fall easily onto the H {\\sc ii} locus, however, only less than 10\\% are of the Seyfert type \\citep{ho97V,kau03,con06}. %!!! Citation needed here or after the next sentence, probably Ho+ 97 (Palomar paper V) = YES. A large fraction of the objects ``in between'' these two categories, that exhibit relatively low levels of ionization (i.e., low values of [\\ion{O}{3}]/H$\\beta$), maintain reasonably strong forbidden line activity (i.e., high values of [\\ion{N}{2}]/H$\\alpha$), and are classified as low ionization nuclear emission regions (LINERs).  The reminder are usually called Transition objects (Ts).  Whether Ts and Ls are powered, at least partly, by accreting BHs, and thus could be called AGN, is a matter of continuous debate \\citep{ho08}.  %recognizing AGN in LINERs There are some typical emission characteristics that are considered to be particularly good indications that accretion onto a massive BH is an important, if not the dominant source of ionization in some Ls, and possibly Ts. The detection of broad H$\\alpha$ emission, regardless of its strength or luminosity, is generally considered to be $the$ clue to AGN emission.  Some LINERs (and maybe Ts as well) exhibit these features, however, the majority of these systems show only narrow emission, which could be generated by shocks, post-starbursts, or other processes unrelated with accretion.  Observations outside of the optical wavelengths often reveal AGN signatures in ambiguous and even starburst galaxy nuclei.  X-rays are particularly good tracers of accretion, however, they are not efficient in distinguishing AGN at $L_X \\la 10^{42}$ erg s$^{-1}$, where contamination by X-ray binaries can be significant. %!!!Changed 10^42 to 10^40...  AC changed it back to 42. X-rays are also unlikely to detect heavily absorbed AGN (i.e., Compton-thick; $N_{\\rm H} > 1.5 \\times 10^{24}$ cm$^{-2}$).  On the other hand, mid-IR high ionization emission lines like [\\ion{Ne}{5}] $\\lambda$14.32 $\\mu$m, 24.32$\\mu$m (97.1 eV) appear to be a trustworthy indicator of AGN activity \\citep{weedman05, armus06, abel08} due to the extreme conditions (i.e., very hard ionizing radiation) required to produce them; because these features have considerably lower optical depth, their detection can also reveal Compton thick AGN \\citep[e.g., NGC 1068]{sturm02}.  This technique has now been applied to reveal new and large numbers of optically unidentified AGN \\citep{satyapal08, goulding09}, providing thus sensitive improvements on previous AGN censuses.  These studies, along with X-ray and radio studies of nearby galaxies \\citep[e.g.][]{nagar05,zhang09,desroches09} suggest that a majority of LINERs and a large fraction of transition galaxies might in fact host accreting black holes. %!!!ACS added final sentence.  Need to add references.  LMK if any %!!!confusion. The presence of an accreting black hole does not however guarantee that the accretion power is the dominant source of ionization of those galaxy nuclei. A more recent assessment of the energy budget of LINERs by \\citet{eracleous10} argues that in 85\\% of LINERs the AGN ionizing photons are not sufficient for producing the observed nebular emission, and thus other power sources are likely to dominate.   %advertise the sequence Some new potentially powerful insights into the excitation mechanism of the low luminosity AGN (LLAGN), and in particular the ambiguous sources, come from recent studies of large statistical samples of nearby galaxy nuclei, which reveal a potential {\\it H {\\sc ii} $\\rightarrow$ S $\\rightarrow$ T $\\rightarrow$ LINER $\\rightarrow$ Passive Galaxies} evolutionary sequence in the process of BH growth within galaxies \\citep{con08, con09}. This sequence traces trends in (1) increasing host halo mass, (2) increasing environmental density, (3) increasing central BH mass and host stellar mass, (4) decreasing BH accretion rate, (5) aging of the stellar population associated with their nuclei, and (6) decreasing in the amount of dust obscuration, which might translate into a decrease in the amount of material available for star-forming or accretion.  In this picture, Seyferts and Ts are transition phases between the initial onset of accretion, usually swamped by the star-forming gas and associated dust, which is seen optically as an H {\\sc ii} system, and the final phase of accretion observed as LINERs of already massive BHs.  While this idea is supported by various other independent observational studies of low luminosity AGN and starburst galaxies \\citep{sch07, sch10}, along with state of the art hydrodynamical models for the life-cycles of the most luminous AGN (e.g., Di Matteo et al. 2005, Hopkins et al. 2006), it is very probable that not $all$ sources fit into this scenario. It is also important that this evolutionary sequence idea is tested on samples that span a narrow distribution in Hubble types;  because most of the H {\\sc ii}s and Ts are relatively late-type disk galaxies that, likely, never experienced a (recent) major merger, it is possible that the trigger of such a sequence is different from the merger that initiates a similar life-cycle at high luminosities.  % [although recent work by \\citet{Dierickx11}]. %Edited following sentences ACS Nevertheless, investigating the challenges certain objects bring to this idea is useful for identifying and quantifying the caveats associated with this sequence.  These objects may also be at interesting stages in their galaxy evolution.   %more on when/where/for what objects the sequence might fail?   %introducing M94: NGC 4736 (or M94, UGC 7996) is a captivating example of an ambiguous galaxy nucleus, which poses challenges to the general understanding of AGN phenomena, including the above mentioned sequence.  This object is one of the closest ($d = 4.3$ Mpc; Tully \\& Fisher 1988) nearly face-on spiral, with a SAab Hubble Type.  Its proximity enables study of details that would be unobservable in more distant systems.  Its nucleus has a low-luminosity LINER spectrum, but has been also included in catalogs of Transition objects \\citep{fil85}, or Seyfert 2s (e.g., Spinelli et al. 2006). %The ground-based observations and corresponding modeling %of the underlying optical absorption-line spectrum are suggestive of a %$10^9$ year old starburst \\citep{tan96}. %ACS. Commented this out, since I'm not sure its really true on the scale we're looking at. The AGN nature of this object has been constantly debated, an aging starburst being a compelling alternative \\citep{maoz95, mao05, era02, koerding05}.  The galaxy presents a ring of H {\\sc ii} regions at a radius of $\\sim 50$\\arcsec, red arcs at $\\sim 15$\\arcsec, a high surface brightness nuclear region, and high far-infrared bulge emission \\citep{kinney93, smith94}. %ACS: What are red arcs? Its intricate structure of off-nuclear compact source detections in X-ray \\citep{era02}, radio \\citep{koerding05} and UV \\citep{mao05}, that do not necessarily match with each other, certainly increase the ambiguity associated with the nature of the main nuclear ionization mechanism.  A common implied scenario in all of these studies is that this system is probably in the final stages of a merger.  We reexamine here NGC 4736 in the context of additional evidence for its AGN nature, which is the detection of a broad H$\\alpha$ component in its nuclear spectrum, as observed by the {\\it Hubble Space Telescope} ($HST$) with the Space Telescope Imaging Spectrograph (STIS).  We also gather multi-wavelength data and show that source of the broad H-alpha emission line is coincident with a compact X-ray and radio source.   These observations suggest that NGC~4736 hosts a broad line region of significantly low luminosity, which makes this object one of the least luminous LINERs with strong evidence for BH accretion.  Interestingly, BH mass indicators calibrated on rapidly accreting Seyfert galaxies give highly discrepant mass estimates in this more quiescent system.  %, suggesting changes to the broad line emission mechanism.%estimates of the BH mass from different indicators differ by 2 orders of magnitude. The $\\sim 10^7 M_{\\sun}$ value given by the $M-\\sigma^*$ is two orders of magnitude higher than the values obtained via estimators based on the observed emission (X-ray variability, scaling relations, mid-IR emission) which, although physically independent of each other, give a consistent result of $\\sim 10^{5} M_{\\sun}$. We are thus facing the following conundrum: either a) the standard AGN BH mass indicators do not necessarily apply to sources emitting in this low luminosity regime as, probably, the emission mechanism is fundamentally different from that associated with higher luminosity AGN, or b) the emission signatures do not trace accretion onto the central BH.  We discuss possible resolutions of this discrepancy in \\S5, where we propose some rather exotic scenarios.    %We are thus facing the following conundrum: either (i) the BH is of intermediate mass ($\\sim10^5 M_{\\sun}$) and the $M-\\sigma^*$ method fails in this particular case, or (ii) the BH is of high mass ($\\sim10^7 M_{\\sun}$) and the emission mechanism is fundamentally different from that associated with higher luminosity AGN.  We discuss possible resolutions of this discrepancy in \\S5.  %% These observations expose the least luminous AGN LINER discovered so %% far, with a bolometric luminosity $L_{\\rm bol} \\approx 2 \\times %% 10^{40} erg s^{-1}$.  While an AGN interpretation appears the most %% plausible, alternative scenarios are also possible.  The most %% interesting thing about this source is probably that various BH mass %% estimates differ by $> 2$ orders of magnitude.  The $\\sim 10^7 %% M_{\\sun}$ value given by the $M-\\sigma^*$ is much higher than the %% values obtained via estimators based on the observed emission (X-ray %% variability, scaling relations, mid-IR emission) which, although %% physically independent of each other, give a consistent result of %% $\\sim 10^{4-5} M_{\\sun}$.  We are thus facing the following conundrum: %% either (i) the BH is of intermediate mass ($\\sim10^5 M_{\\sun}$) and %% the $M-\\sigma^*$ method fails in this particular case, or (ii) the BH %% is of high mass ($\\sim10^7 M_{\\sun}$) and the emission mechanism is %% fundamentally different from that associated with higher luminosity %% AGN. ",
        "conclusions": "%summary We have presented here an exhaustive multiwavelength analysis of the nuclear emission properties of NGC 4736, prompted by new measurements of a broad H$\\alpha$ emission component detected in its high resolution {\\it HST}-STIS optical spectrum.  This broad H$\\alpha$ component, with a luminosity of $9\\times10^{37}$ ergs s$^{-1}$, is one of the lowest luminosity broad line known. This broad H$\\alpha$ is coincident with a compact bright X-ray and radio source.  Our measurements of this object's spectral energy distribution reveal a bolometric luminosity of $L_{\\rm bol} \\approx 2.5 \\times 10^{40}$ erg s$^{-1}$, that categorizes NGC 4736 as one of the least luminous LINERs with strong evidence for BH accretion. Our comparison of five independent BH mass estimates reveals a discrepancy of two orders of magnitude between the value $\\sim 10^7 M_{\\sun}$ predicted by the $M - \\sigma^*$ relation and the value $\\sim 10^{5} M_{\\sun}$ toward which methods based on AGN emission activity in optical, mid-IR and X-rays, seem to converge;  the fifth method is provided by the fundamental plane relation, which however, due to its large associated uncertainties, does not offer any additional constraint to this comparison.  We conclude that this system's BH mass cannot be reliably estimated via standard AGN BH mass indicators because the nuclear emission in this system is not entirely tracing the accretion onto the central BH.  Our assessment of the energy budgets of the ionizing and emission-line power suggests a possible deficit in the AGN ionization and production of a broad H$\\alpha$ emission feature which can be made up by a peculiar kind of Type IIn SN that matches well the nuclear emission of NGC 4736 over the whole electromagnetic spectrum, and supports this galaxy nucleus' general aging starburst-like appearance."
    },
    "1112/1112.0288_arXiv.txt": {
        "abstract": "To measure the magnetic field strength in the solar corona, we examined 10 fast ($\\geq 1000$ km s$^{-1}$) limb CMEs which show clear shock structures in SOHO/LASCO images. By applying piston-shock relationship to the observed CME's standoff distance and electron density compression ratio, we estimated the Mach number, Alfven speed, and magnetic field strength in the height range 3 to 15 solar radii ($R_s$). Main results from this study are: (1) the standoff distance observed in solar corona is consistent with those from a magnetohydrodynamic (MHD) model and near-Earth observations; (2) the Mach number as a shock strength is in the range 1.49 to 3.43 from the standoff distance ratio, but when we use the density compression ratio, the Mach number is in the range 1.47 to 1.90, implying that the measured density compression ratio is likely to be underestimated due to observational limits; (3) the Alfven speed ranges from 259 to 982 km s$^{-1}$ and the magnetic field strength is in the range 6 to $105 mG$ when the standoff distance is used; (4) if we multiply the density compression ratio by a factor of 2, the Alfven speeds and the magnetic field strengths are consistent in both methods; (5) the magnetic field strengths derived from the shock parameters are similar to those of empirical models and previous estimates. ",
        "introduction": "The solar corona is the plasma atmosphere of the Sun, extending millions of kilometers into space. One of the major issues of the solar corona is to measure the magnetic field, which determines the coronal structure and dynamics from the upper chromosphere out into the heliosphere. Evidence for magnetic field in the corona has been found in several kinds of observations, but only a few of them give magnetic field information since the coronal plasma is optically thin (Wiegelmann 2008). These observations are possible only in a limited spatial extent and the magnetic field has to be derived with some uncertainties in interpretation.  One of the main techniques to estimate the coronal magnetic field is the optical observations of vector magnetic fields in the photosphere and their extrapolation into the corona. Lin et al. (2000) presented the magnetic field strength in the inner corona based on Stokes V circular polarization profiles. Solanki et al. (2003) reported three-dimensional magnetic field topology in an interaction region near the base of the solar corona from the measurement of Stokes vector. Their results of magnetic field strength, $B$ are from tens to several hundreds $G$ in a very limited region ($r < 0.5 R_s$). The extrapolations of the photospheric magnetic field into the solar corona depend on the assumptions such as low $\\beta$ plasma, which may not be valid in the outer corona (Gary 2001). Radio data may also be used as a diagnostic of coronal magnetic structure. Lee et al. (1999) used radio observations of an active region to examine the coronal magnetic field obtained via a nonlinear force-free field extrapolation of a photospheric vector magnetogram. Ramesh et al. (2010) presented that the estimated values of $B$ at two different distances at 1.5  and $1.7 R_s$ from the observations of circularly polarized thermal radio emission from solar coronal streamers are 6 $\\pm 2 G$ and 5 $\\pm 1 G$, respectively. It should be noted that the above techniques can be applied to measure the magnetic fields only in the inner corona ($r < 2 R_s$).  Several studies have been performed to measure magnetic fields in the corona from the band splittings in type II radio bursts (Smerd et al. 1974). Vrsnak et al. (2004) measured the interplanetary magnetic field strength using the band splitting of coronal and interplanetary type II bursts. But, they considered only four bins in the distance range of 25 to 225 $R_s$ due to difficulties in the observation, such as a low signal-to-noise ratio and small amount of data. Cho et al. (2007) used band splitting of coronal type II radio bursts and obtained coronal magnetic field strength of 1.3 to 0.4 $G$ in the height range of 1.5 to 2 $R_s$. Faraday rotation techniques have been occasionally used in estimating the magnetic field strengths at several solar radii (Patzold et al. 1987; Spangler 2005; Ingleby et al. 2007).  It is known that the observations in solar corona approximately follow the empirical formulas $B(r)=0.5 (r/R_{s} -1)^{-1.5} G$ for active regions (Dulk $\\&$ McLean 1978), and $B(r)=2.2 (r/R_{s})^{-2} G$ for quiet regions (Mann et al. 1999). Since these formulas were made from the observations of lower corona below $\\sim 10 R_{s}$, we may wonder whether  these radial dependencies of $B$ are still effective in the upper coronal region up to several tens of solar radii. The observations by Large Angle and Spectrometric COronagraph (LASCO) on aboard the SOlar and Heliospheric Observatory satellite (SOHO) (Brueckner et al. 1995) enable us to study the upper corona in the range $1.5 R_s < r < 30 R_s$.  These observations have shown that the speeds of coronal mass ejections (CMEs) vary from a few hundred to more than 2500 km s$^{-1}$ (Yashiro et al. 2004; Gopalswamy 2010). CMEs with speeds in excess of the ambient Alfven speed, which is characterized by the magnetic field and plasma density, may drive fast mode MHD shocks. As a shock signature associated with CMEs, streamer deflections have been suggested and observed (Gosling et al. 1974; Michels et al. 1984; Sheeley et al. 2000). Recently, several studies have shown that the CME-driven shocks could be directly observed in white-light coronagraph images, under suitable conditions from the analysis of whith light (Vourlidas et al. 2003; Gopalswamy et al. 2009a; Ontiveros \\& Vourlidas 2009; Bemporad \\& Mancuso 2010; Gopalswamy 2010) and EUV (Gopalswamy et al. 2011; Ma et al. 2011; Kozarev et al. 2011) images as well as from magnetohydrodynamic (MHD) simulations (Stepanova \\& Kosovichev 2000;  Manchester et al. 2004). If these structures are truly the shocks, there should be noticeable rapid rises in pressure, temperature and density of the flow and applicable to the piston-shock relationship (Eselevich \\& Eselevich 2010).  In piston-driven shocks such as the Earth's bow shock, there are several parameters that can govern the shock shapes: the size and shape of the obstacle, the electron density compression ratio, the upstream Mach number, and the standoff distance (the distance between the obstacle and its shock nose). Russell and Mulligan (2002) applied the relation between shock standoff distance to a CME near-Earth to explain the curvature of the driving interplanetary CME (ICME). Gopalswamy \\& Yashiro (2011) measured the standoff distance of a CME-driven shock in the corona using SOHO/LASCO and the Solar TErrestrial RElations Observatory (STEREO) images and derived the coronal magnetic field in the range of 6 to 23 $R_s$.  In this paper, we consider a large number of shock-driving CMEs identified in the SOHO/LASCO images and use the standoff distance technique (Gopalswamy \\& Yashiro 2011) for measuring the coronal magnetic field. Following the case study of Gopalswamy \\& Yashiro (2011), we estimate the magnetic field radial distributions in the upper corona using the standoff distance technique. We also use the density compression ratio across the shock to determine the magnetic field and compare the results with those from the standoff distance technique. In addition, we compare the magnetic field distributions with those in previous studies.  We also examine the physical properties of the upstream medium from different techniques for consistency. The paper is organized as follows. The data and methodology are described in Section 2. In section 3, we present the estimations of the coronal Alfven speeds and magnetic field strengths using the standoff distance and the density compression ratio methods. A brief summary and conclusions are given in Section 4. ",
        "conclusions": "To measure the magnetic field strength in the solar corona, we examined 10 fast ($\\geq 1000$ km s$^{-1}$) limb CMEs which show clear shock structures in SOHO/LASCO images. By applying the piston-shock relationship to the observed CME's standoff distance, we obtained the coronal Alfven speed to be in the range from 259 to 982 km s$^{-1}$ and the magnetic field strength in the range from 6 to $105 mG$ in the heliocentric distance range of 3 and $15 R_s$. The magnetic field strength are consistent with the empirical models (Dulk \\& McLean 1978) and other studies (Patzold et al. 1987; Spangler 2005; Cho et al. 2007; Ingleby et al. 2007; Bemporad \\& Mancuso 2010; Gopalswamy \\& Yashiro 2011). These results confirm that  the standoff distance ratio provides us with a useful tool to derive the magnetic fields in the wide range of solar corona ($\\sim 30 R_s$).  The Alfven speeds and magnetic field strengths derived from the density compression ratio are about two times higher than the above results. We speculate that the density compression ratio obtained from white-light observations might be underestimated since the observed density is based on the electrons integrated over the line-of-sight, while the measurements of standoff distance ratio have no such weakness. To inspect the line-of-sight effect on the density compression ratio, we adopted the method proposed by Eselevich \\& Eselevich (2011), which uses the differential brightness dP from the LASCO tB image and the shock size $l$ along the line-of-sight. Assuming that the CME and Shock's configuration is symmetric, we measured the length of tangential line at CME\u2019s nose, which is considered as equivalent to the length along the line-of-sight, $l$. As the result, we found that if we choose Newkirk density model (Newkirk 1961) for upstrem density, the density compression ratio, which is calculated from dP and $l$, has very good correlation with our previous result from Van de Hulst inversion method with a correlation coefficient of 0.96 as shown in Figure 11.  We also note that when we multiply the density compression ratio by a factor of 2,  the Alfven speeds and the magnetic field strengths are consistent with those from the standoff distance technique. This supports the idea that the diffuse structures surrounding the CME front, as shown in Figure 1, can be interpreted as shock structures - shock sheath to be precise. There are two main observational results that support the existence of shocks in the low corona. Type II radio bursts in the metric (Cliver et al. 1999) and longer wavelengths (Gopalswamy et al., 2005) are good indicators near the Sun. The white light observations of diffuse features surrounding the CME flux ropes confirm this, as inferred from streamer deflections (Gosling et al. 1974) and other white-light signatures (Vourlidas et al. 2003; Sheeley et al. 2000; Gopalswamy et al. 2008a; Gopalswamy et al. 2009a) and MHD simulations (Manchester 2009).  Finally, we would like to stress on the fact that this study is a new attempt (together with Gopalswamy \\& Yashiro 2011) to measure the magnetic field strengths in the upper corona up to $30 R_s$ by applying the piston-shock relationship to CME coronagraph observations. This method can be applied to CMEs showing clear shock structures surrounding the CME front so that it can provide us a useful method to derive the magnetic fields in the  solar corona.  It is a unique method to derive magnetic fields in the upper solar corona (10 $\\sim 20 R_s$)."
    },
    "1112/1112.2591_arXiv.txt": {
        "abstract": "We explore the application of {\\it XMM-Newton} Optical Monitor (XMM-OM) ultraviolet (UV) data to study galaxy evolution. Our sample is constructed as the intersection of all Abell clusters with $z < 0.05$ and having archival XMM-OM data in either the $UVM2$ or $UVW1$ filters, plus optical and UV photometry from the Sloan Digital Sky Survey and {\\it GALEX}, respectively. The eleven resulting clusters include 726 galaxies with measured redshifts, 520 of which have redshifts placing them within their parent Abell clusters. We develop procedures for manipulating the XMM-OM images and measuring galaxy photometry from them, and confirm our results via comparison with published catalogs. Color magnitude diagrams (CMDs) constructed using the XMM-OM data along with SDSS optical data show promise for evolutionary studies, with good separation between red and blue sequences and real variation in the width of the red sequence that is likely indicative of differences in star formation history. This is particularly true for $UVW1$ data, as the relative abundance of data collected using this filter and its depth make it an attractive choice. Available tools that use stellar synthesis libraries to fit the UV and optical photometric data may also be used, thereby better describing star formation history within the past Gyr and providing estimates of total stellar mass that include contributions from young stars. Finally, color-color diagrams that include XMM-OM UV data appear useful to the photometric identification of both extragalactic and stellar sources. ",
        "introduction": "The power of ultraviolet (UV) data has been demonstrated by numerous publications driven by data collected with the {\\it Galaxy Evolution Explorer} satellite \\citep[GALEX,][]{martin2005}. UV emission comes primarily from young stars, with the youngest and most massive stars dominating the far-UV while stars with lifetimes of about a Gyr dominate the near-UV. The addition of GALEX UV photometry to optical photometry from the Sloan Digital Sky Survey \\citep[SDSS;][]{york2000} thus allows for improved classification of millions of stars and galaxies \\citep{bianchi2007}. In fact, GALEX UV photometry is invaluable for identification of hot stars and white dwarfs within the Milky Way \\citep{bianchi2011}. It is also vital in estimating the dust attenuation in galaxies and thus properly evaluating star formation rates \\citep{treyer2007,johnson2007,salim2007}.  One of the simple and yet powerful tools in applying UV data to the understanding of galaxies and their evolution is the color-magnitude diagram (CMD). Galaxies are known to separate well into red and blue populations \\citep[e.g.,][]{baldry2004}, forming a tight red sequence and a broader but still well-described blue cloud. Colors formed using near-UV data paired with optical data produce a larger separation between red and blue galaxies, thereby producing cleaner assignments of galaxies \\citep{wyder2007}. More significantly, CMDs constructed using near-UV and optical data reveal a population of objects intermediate to the red sequence and blue cloud and dubbed the ``green valley'' \\citep{wyder2007}. These may represent an important transitory population consisting of star-forming galaxies and active galactic nuclei (AGN) fading into quiescence and an eventual home as ``dead and red'' galaxies on the red sequence \\citep{schiminovich2007,martin2007}. Alternatively, they may be rejuvenated red sequence galaxies to which a ``frosting'' of younger stars has been added \\citep{trager2000}. Along these same lines, the spread in the UV-optical color at a given absolute magnitude for early-type galaxies can be explained by a significant fraction ($\\sim30\\%$) having experienced star formation within the past Gyr \\citep{kaviraj2007}.  The {\\it XMM-Newton} satellite \\citep{jansen2001} is also equipped with an Optical Monitor telescope \\citep[][XMM-OM]{mason2001}. This instrument provides concurrent observation of XMM's X-ray targets in a selection from among six filters, three of which approximate the Johnson $U~B~V$ filters and three of which reside in the UV ($UVW2$, $UVM2$, and $UVW1$). Although of smaller aperture and with a smaller field of view than GALEX, XMM-OM provides higher resolution images with a point spread function that is roughly a factor of three finer than that of GALEX. The $UVM2$ filter has an effective wavelength and width comparable to GALEX $NUV$, and thus can directly emulate the science performed with that filter. The XMM-OM data might also prove useful in evaluating the age and abundance of old stellar populations and hence elliptical galaxies. The region of the spectrum of a stellar population between about 2000$\\mbox{\\AA}$ and 3200$\\mbox{\\AA}$, and hence the $UVW1$ filter, should be dominated by main sequence stars near the main sequence turn-off. Photometry in this wavelength range can help break the age-metallicity degeneracy that plagues optical filter measurements by placing a strong lower limit on the metallicity of that population \\citep{dorman2003}.  In this paper, we investigate the usage of XMM-OM data to study galaxies. We ultimately strive to apply the XMM-OM photometry to understand the star formation histories of cluster galaxies, and evaluate such information in relation to evolutionary processes believed to be at work in cluster environments. XMM is a popular choice for X-ray observations of galaxy clusters on account of its sensitivity to the extended emission of the hot intracluster medium from thermal bremsstrahlung, implying that there is a large and nearly untapped database of associated UV observations of galaxy clusters made by XMM-OM. Prior publications capitalizing on XMM-OM data associated with observations of galaxy clusters have focused on solely the brightest cluster galaxies and used the UV data to evaluate star formation in such objects and cluster cooling flows \\citep{mittaz2001,hicks2005,donahue2010}. Our work utilizes the full field of view of XMM-OM in order to compile a database with hundreds of galaxies having UV photometry to pair with SDSS data. Unless otherwise noted, all magnitudes are in the AB system and we assume a standard $\\Lambda$CDM cosmology with $\\Omega_m = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0 = 70$ km~sec$^{-1}$~Mpc$^{-1}$.  We describe the sample construction and data sources in Section \\ref{sec-data}. Our procedures for processing the images and measuring galaxy photometry are described in Section \\ref{sec-proc}, followed by the construction of multiwavelength catalogs in Section \\ref{sec-catalog} and their usage as checks of calibration in Section \\ref{sec-calcheck}. The remainder of the paper is devoted to applications of the XMM-OM UV data, including construction and assessment of CMDs (Section \\ref{sec-cmd}), evaluation of star formation history (Section \\ref{sec-sfh}), inclusion of XMM-OM UV data in photometric redshift and galaxy stellar mass estimation (Section \\ref{sec-pzandm}), and selection of unusual objects via color-color diagrams including XMM-OM UV data (Section \\ref{sec-nonclus}). A summary is provided in Section \\ref{sec-conclude}. ",
        "conclusions": "\\label{sec-conclude}  The archive of XMM observations that have targeted galaxy clusters is large, and for many of these observations the XMM-OM was also collecting data in UV filters. By taking the intersection of such cluster observations with the coverage of the SDSS and GALEX surveys and imposing a redshift cutoff of $z=0.5$, we have isolated a sample of 11 nearby Abell clusters with XMM-OM data in either of the $UVM2$ or $UVW1$ filters. We matched the astrometry of these images to that of the SDSS survey, made corrections for degradation of the XMM-OM system as a function of time, and stacked observations in fields that were observed in multiple campaigns. The accuracy of our image manipulation and photometry procedures were confirmed by their consistency with the OMCat and GALEX catalogs. In general, the $UVM2$ data are of comparable depth to the $NUV$ data of the GALEX AIS and the $UVW1$ data are comparable to the $u$ data of the SDSS.  After folding in available redshift data, we arrived at a sample of 726 galaxies for use in investigating the potential value of XMM-OM UV photometry for galaxy evolution studies. Of these, 520 belong to the 11 Abell clusters of the sample. The benefits of CMDs created using GALEX $NUV$ data are likely paralleled by XMM-OM $UVM2$ data, although our sample does not provide large numbers of such objects to test this claim. However, CMDs constructed using the $UVW1$ filter and SDSS $r$ show good promise for galaxy evolution studies. They show a strong red sequence with a wider real dispersion than CMDs created using SDSS $u$ and $r$ data, likely reflecting variations in star formation history among red sequence galaxies. They also provide greater separation between red and blue sequences than those obtained using only the SDSS data. Our sample is dominated by red sequence objects and we are unable to determine whether $UVW1$ data might prove useful in identifying green valley objects, although the aforementioned characteristics of $UVW1$-incorporating CMDs provides a basis for some optimism on this subject. The \\citet{blanton2007} \\kcorrect{} package may also be applied to investigate star formation histories and produce estimates of galaxy stellar masses that include the contributions of younger stellar populations. Finally, color-color diagrams using $UVW1$ data show promise for the identification of objects ranging from cool to hot stars, and extragalactic sources such as quasars."
    },
    "1112/1112.1840_arXiv.txt": {
        "abstract": "A search for an excess of muon--neutrinos from dark matter annihilations in the Sun has been performed with the AMANDA--II neutrino telescope using data collected in 812 days of livetime between 2001 and 2006 and 149 days of livetime collected with the AMANDA--II and the 40--string configuration of IceCube during 2008 and early 2009. No excess over the expected atmospheric neutrino background has been observed. We combine these results with the previously published IceCube limits obtained with data taken during 2007 to obtain a total livetime of 1065 days. We provide an upper limit at 90\\% confidence level on the annihilation rate of captured neutralinos in the Sun, as well as the corresponding muon flux limit at the Earth, both as functions of the neutralino mass in the range 50~GeV--5000~GeV. We also derive a limit on the neutralino--proton spin--dependent and spin--independent cross section. The limits presented here improve the previous results obtained by the collaboration between a factor of two and five, as well as extending the neutralino masses probed down to 50 GeV. The spin--dependent cross section limits are the most stringent so far for neutralino masses above 200~GeV, and well below direct search results in the mass range from 50~GeV to 5~TeV. ",
        "introduction": "Introduction}  There is an impressive corpus of astrophysical and cosmological observations that indicate that a yet unknown form of non--luminous matter constitutes about 23\\% of the content of the universe. This matter can not be baryonic though, as the abundance of baryons in any form is strongly constrained by the inferred primordially synthesized abundances of deuterium and helium, as well as the observed small-scale anisotropy in the cosmic microwave background (see, e.g., the review in Ref.~\\cite{PDG:10}). However, the observational limits on non--baryonic dark matter in the form of relic stable particles from the Big Bang are not so strong and, indeed, a wealth of models exist that propose candidates with interaction cross sections, masses and relic densities which are compatible with observations. Generically, any particle physics model that provides a stable weakly--interacting massive particle (WIMP) is of interest from the point of view of the dark matter problem. One class of models that are extensively studied for their interest in particle physics are supersymmetric extensions of the Standard Model. Several flavours of supersymmetry are currently under scrutiny at the LHC, and both ATLAS and CMS have already probed parts of the parameter space of specific benchmark models~\\cite{ATLAS,CMS}. In this paper we focus on the Minimal Supersymmetric extension to the Standard Model (MSSM), which provides a WIMP candidate in the lightest neutralino, $\\widetilde{\\chi}^{0}_{1}$, a linear combination of the supersymmetric partners of the electroweak neutral gauge-- and Higgs bosons. Assuming $R$--parity conservation, the neutralino is stable and a good dark matter candidate. Accelerator searches and relic density constraints from WMAP data allow a lower limit to be set  on the mass of the MSSM neutralino~\\cite{Genieveve:04a}. Typical lower limits for $m_{\\widetilde{\\chi}^{0}_{1}}$ from such studies lie around  20 GeV, depending on the values chosen for tan$\\beta$, the ratio of the  vacuum expectation values of the two neutral Higgses. Theoretical arguments based on the requirement of unitarity set an upper limit on  $m_{\\widetilde{\\chi}^{0}_{1}}$ of 340~TeV~\\cite{Griest:90a}. Within these limits, the allowed parameter space of minimal supersymmetry can be exploited to build realistic models which provide relic neutralino densities of cosmological interest to address the dark matter problem. \\par  Relic neutralinos in the galactic halo can lose energy through scatterings while traversing celestial bodies, like the Sun, and become gravitationally bound into orbits inside. The build-up of the accreted particles is limited by annihilations, which ultimately create high--energy neutrinos~\\cite{Silk:85a}. In this paper we present two independent searches for a neutrino flux from the annihilations of neutralinos captured in the center of the Sun, performed with the AMANDA--II and IceCube neutrino telescopes. The overwhelming majority of triggers in AMANDA--II and IceCube, $\\mathcal{O}$(10$^{10}$/year), are due to atmospheric muons reaching the depth of the array. The analyses are therefore based on the search for up--going muon tracks from the direction of the Sun when it is below the horizon, and have been optimized in order to maximize the sensitivity to the predicted signal from MSSM neutralinos. The results can, however, be interpreted  generically in terms of any other dark matter candidate which would produce a similar neutrino spectrum at the detector. The first analysis (Analysis A) uses the data taken with AMANDA--II during 812 days of livetime in stand--alone operation between March 2001 and October 2006. In 2007 the detector was switched to a new data acquisition system (DAQ) which included full waveform recording, and this allowed us to integrate the detector as a subsystem of IceCube.  Analysis B uses 149 days of livetime collected between April 2008 and April 2009 (when AMANDA--II was decommissioned) using the new AMANDA--II DAQ and data from both detectors, AMANDA--II and IceCube. \\par  We have previously published a search for dark matter accumulated in the Sun using data taken in 2007 with the 22--string configuration of IceCube~\\cite{Abbasi:09b}. In Section~\\ref{sec:combination} we combine these three independent analyses (Analysis A, Analysis B and the results from Ref.~\\cite{Abbasi:09b}) which cover the period 2001--2008, presenting competitive limits on the muon flux from neutralino annihilations in the Sun and limits on the spin--dependent and spin--independent neutralino--proton cross section.  \\begin{figure}[t] \\includegraphics[width=\\linewidth]{surfacemap03.pdf} \\caption{Surface location of the IceCube 40--string configuration (filled squares) and AMANDA--II (filled circles)} \\label{fig:geometry} \\end{figure} ",
        "conclusions": ""
    },
    "1112/1112.1071_arXiv.txt": {
        "abstract": "{ \\hyperion is a new three-dimensional dust continuum Monte-Carlo radiative transfer code that is designed to be as generic as possible, allowing radiative transfer to be computed through a variety of three-dimensional grids. The main part of the code is problem-independent, and only requires an arbitrary three-dimensional density structure, dust properties, the position and properties of the illuminating sources, and parameters controlling the running and output of the code. \\hyperion is parallelized, and is shown to scale well to thousands of processes. Two common benchmark models for protoplanetary disks were computed, and the results are found to be in excellent agreement with those from other codes. Finally, to demonstrate the capabilities of the code, dust temperatures, SEDs, and synthetic multi-wavelength images were computed for a dynamical simulation of a low-mass star formation region. \\hyperion is being actively developed to include new features, and is publicly available (\\url{http://www.hyperion-rt.org}). }{} ",
        "introduction": "\\label{sec:introduction}  The investigation of astrophysical sources via multi-wavelength observations requires understanding the transfer of radiation through dust and gas in order to reliably derive geometrical, physical, and chemical properties. While a small subset of problems can be solved analytically, most require a numerical solution.  A variety of techniques have been developed to this effect, and until recently, most relied on a direct numerical solution to the differential equation of radiative transfer. Examples include 2$^{\\rm nd}$ order finite differences \\citep{steinacker:03:405}, 5$^{\\rm th}$ order Runge-Kutta integration \\citep{steinacker:06:920}, and variable Eddington tensors \\citep{dullemond:00:1187, dullemond:02:853}.  While these techniques can often provide results accurate to arbitrary precision, and have been very successful for one-dimensional problems, they become increasingly complex when applied to two- and three-dimensional problems. Within the last decade, the \\textit{Monte-Carlo} technique applied to radiative transfer has become an increasingly popular alternative that is well suited to arbitrary three-dimensional density distributions. Rather than solving the equation of radiative transfer directly, this method relies on random sampling of probability distribution functions to propagate \\textit{photon packets} through a grid of constant density and temperature cells.  In its simplest implementation, Monte-Carlo radiative transfer considers single-frequency photon packets that are emitted by sources, and are propagated through a grid of constant density cells. Each photon packet travels a certain optical depth, before being either scattered, or absorbed and immediately re-emitted (to conserve energy). These photon packets can then continue to travel and interact until they escape the grid. The optical depth, type of interaction, and frequency of re-emitted photon packets, as well as scattering/re-emission directions are all sampled from probability distribution functions. Photon packets escaping from the grid can be used to compute synthetic observations, such as spectral energy distributions (SEDs) and images. By repeating this process for large numbers of photon packets, the signal-to-noise of the synthetic observations can be increased.  In cases where local thermodynamic equilibrium (LTE) can be assumed, the energy absorbed in each cell can be used to compute the equilibrium temperature. Since the emissivity depends on the temperature, the radiative transfer has to be run iteratively to obtain a self-consistent solution for the temperature.  Various methods have been developed to improve the performance of Monte-Carlo radiative transfer codes, which would otherwise be inefficient in many cases. For example, rather than simply using discrete absorptions in cells to compute the equilibrium temperature, \\citet{Lucy:99:282} suggested summing the optical depth to absorption of photon packet paths through each grid cell, requiring far fewer photons to achieve a comparable signal-to-noise in the derived temperatures. \\citet{bjorkman:01:615} developed an algorithm whereby the temperature of a cell is updated immediately each time a photon is absorbed, and the frequency of the re-emitted photon is sampled from a difference emissivity function that introduces a correction to account for the fact that previous photons were sampled from emissivities for different temperatures. This algorithm is sometimes referred to simply as `immediate re-emission', but this is a misnomer since other algorithms also immediately re-emit a photon following an absorption to conserve energy \\citep[e.g.][]{Lucy:99:282}; instead, what this algorithm introduces is the \\textit{immediate temperature correction} aspect.  For synthetic observations, the \\textit{peeling-off} technique \\citep{Yusef-Zadeh:84:186} greatly improves the signal-to-noise of SEDs and images: every time a photon packet is emitted, scattered, or re-emitted after an absorption, the probability that the emitted or scattered photon would reach the observer is used to build up the observations. Thus each photon contributes multiple times to the SEDs or images. Another example is that of \\textit{forced first scattering} \\citep[e.g.][]{mattila:70:53, wood:99:799}, which can be used in optically thin cases to force photon packets to interact with the dust (this requires weighting the photons accordingly to avoid any bias due to the forcing).  In standard Monte-Carlo, the number of interactions a photon packet will undergo before traveling a certain distance increases approximately with the square of the density, so that in very optically thick cells, photon packets may become effectively trapped. This problem can be avoided by locally making use of the diffusion approximation to solve the radiative transfer in these regions. For example, \\citet{min:09:155} developed a modified random walk procedure based on the diffusion approximation that can dramatically reduce the number of steps required for a photon to escape from a grid cell of arbitrarily large optical depth.  A large collection of dust continuum Monte-Carlo radiative transfer codes has been developed  \\citep[e.g.][]{wolf:99:839, harries:00:722, gordon:01:269, misselt:01:277, wood:01:299, wood:02:1183, wood:02:887, wolf:03:99, stamatellos:03:941, whitney:03:1079, whitney:03:1049, whitney:04:1177, dullemond:04:159, jonsson:06:2, pinte:06:797, min:09:155}, each including some or all of the above optimizations, as well as other optimizations not mentioned here. While some of the early codes assumed spherical or axis-symmetric geometries for simplicity, many have since been adapted to compute fully arbitrary three-dimensional distributions. In addition to dust continuum radiative transfer, some codes can also compute non-LTE line transfer \\citep{carciofi:06:1081, carciofi:08:1374}, or photoionization \\citep[e.g.][]{ercolano:03:1136, ercolano:05:1038, ercolano:08:534}.  This paper presents \\hyperion, a new dust-continuum Monte-Carlo radiative transfer code that is designed to be applicable to a wide range of problems. \\hyperion implements many of the recent optimizations to the Monte-Carlo technique discussed above, was written from the start to be a parallel code that can scale to thousands of processes, and is written in a modular and extensible way so as to be easily improved in future. It can treat the emission from an arbitrary number of sources, can include multiple dust types, and can compute the anisotropic scattering of polarized radiation using fully numerical scattering phase functions. It uses the \\citet{Lucy:99:282} iterative method to determine the radiative equilibrium temperature, but does not use the \\citet{bjorkman:01:615} temperature correction technique, as the latter is much more challenging to parallelize efficiently. Thanks to the modular nature of the code, the radiative transfer can be computed on a number of different three-dimensional grid types, and additional grid types can be added in future. \\hyperion can compute SEDs and multi-wavelength images and polarization maps. The code is released under an open source license, and is hosted on a service that allows members of the community to easily contribute to the code and documentation.  Section \\ref{sec:overview} gives an overview of the implementation of the code. Section \\ref{sec:parallel} discusses the efficiency of the parallelized code. Section \\ref{sec:benchmarks} presents the results for two benchmark models of protoplanetary disks. Finally, Section \\ref{sec:simulation} demonstrates the capabilities of the code by computing temperatures, SEDs, and synthetic images for a simulation of a star-formation region. The availability of the code and plans for the future are discussed in Section \\ref{sec:future}. ",
        "conclusions": "\\hyperion is a new radiative transfer code that has been designed to be applicable to a wide range of problems. Radiative transfer can be carried out through a variety of three-dimensional grids, including cartesian, cylindrical and spherical polar, and adaptive cartesian grids. The algorithms used by the code, including the photon packet propagation, raytracing, convergence detection, and diffusion approximations are described in detail (\\S\\ref{sec:overview}).  The code is parallelized and scales well to hundreds or thousands of processes. As expected, the maximum speedup depends on the fraction of the execution time spent in the serial part of the code (such as file input/output), which in turn depends on the problem being computed (\\S\\ref{sec:parallel}).  The code was tested against the \\citet{Pascucci:04:793} and \\citet{Pinte:09:967} benchmarks for circumstellar disks and was found to be in excellent agreement with the respective published benchmark results (\\S\\ref{sec:benchmarks}).  As a case study, \\hyperion was used to predict near-infrared to sub-millimeter observations of a star formation region using a dynamical simulation of a region of low-mass star formation \\citep{offner:09:131}. One interesting finding was that despite the absence of outflows in the simulation, three-dimensional accretion flows caused beaming of radiation that produced fairly symmetric structures usually associated with bipolar outflow cavities.  \\hyperion is published under an open-source license at \\url{http://www.hyperion-rt.org}, using a version control hosting website that makes it easy for users to contribute to the code and documentation."
    },
    "1112/1112.2008_arXiv.txt": {
        "abstract": "{  We present a theoretical analysis of the optical light curves (LCs) for short-period high-mass transiting extrasolar planet systems. Our method considers the primary transit, the secondary eclipse, and the overall phase shape of the LC between the occultations. Phase variations arise from (i) reflected and thermally emitted light by the planet, (ii) the ellipsoidal shape of the star due to the gravitational pull of the planet, and (iii) the Doppler shift of the stellar light as the star orbits the center of mass of the system. Our full model of the out-of-eclipse variations contains information about the planetary mass, orbital eccentricity, the orientation of periastron and the planet's albedo. For a range of hypothetical systems we demonstrate that the ellipsoidal variations (ii.) can be large enough to be distinguished from the remaining components and that this effect can be used to constrain the planet's mass. To detect the ellipsoidal variations, the LC requires a minimum precision of $10^{-4}$, which coincides with the precision of the Kepler mission. As a test of our approach, we consider the Kepler LC of the transiting object HAT-P-7. We are able to estimate the mass of the companion, and confirm its planetary nature solely from the LC data. Future space missions, such as PLATO and the James Webb Space Telescope with even higher photometric precision, will be able to reduce the errors in all parameters. Detailed modeling of any out-of-eclipse variations seen in new systems will be a useful diagnostic tool prior to the requisite ground based radial velocity follow-up.  } ",
        "introduction": "Two observational methods have dominated the study of extrasolar planets so far: radial velocity (RV) measurements and transit light curve (LC) analyses. Both have advantages and disadvantages. While RV determinations provide estimates of the planetary mass $(M_\\mathrm{p})$, the eccentricity $(e)$ and the semi-major axis $(a)$, they do not constrain the inclination $(i)$ of the orbital plane with respect to the observer, thus only lower limits to $M_\\mathrm{p}$ can be determined. The transit method, on the other hand, provides information on $i$, the ratio of the planetary to the stellar radius $(R_\\mathrm{p}/R_\\mathrm{s})$, and the duration of the transit $(D)$. So far, only a combination of both strategies yields a full set of orbital and physical parameters for extrasolar planets.\\par  Currently, two space-missions are targeted at the detection of transiting extrasolar planets: CoRoT launched in 2006 \\citep{1997ASPC..119..259D} and Kepler launched in 2009 \\citep{1997ASPC..119..153B}. Their instruments are monitoring hundreds of thousands of stars, discovering thousands of planet candidates, whose RV follow-up will take many years (and may never be completed for the faintest candidates). Without RV follow-up, the most fundamental parameter of an extrasolar planet, its mass, remains undetermined from CoRoT and Kepler observations alone. The mass is the crucial parameter classifying an object as a planet, brown dwarf or a star. \\par  High-accuracy photometry has already been used for a number of systems to show that the planetary thermal emission, as well as the reflection of the stellar light from the planet, are detectable. In particular, Welsh et al. (2010) report the discovery of ellipsoidal variations in the Kepler LC of HAT-P-7. This is an effect more commonly known from close stellar systems, where phase-dependent light variation arises from the gravitationally distorted stars. In HAT-P-7, the planet is close enough and massive enough to induce the same effect.  Other authors \\citep{2003ApJ...588L.117L, 2011MNRAS.415.3921F} have already presented treatments of reflected light, thermal emission, ellipsoidal variations and Doppler beaming. Our approach to the problem includes two main differences: (i) we investigate the sensitivity of our technique on the adopted phase function (Section 3.2), and (ii) we attempt to solve separately for the reflected light and thermal emission from the planet. For hot Jupiters around hot stars, the temperature of the planet increases dramatically and reflected light alone fails to explain the LC.  In this paper, we investigate whether the presence of ellipsoidal variations in a planetary system can be used to place meaningful constraints on the mass of the planetary companion. In Sect. \\ref{sec:theory} we describe the basic phenomena included in our simplified model of a planetary system's LC. In Sect. \\ref{results} we compare the magnitudes of the relative contributions and consider the likely range of planetary systems for which ellipsoidal variations ought to be detectable with both CoRoT and Kepler. We reanalyze the HAT-P-7 LC and demonstrate that the ellipsoidal variations can be used to estimate the mass of the planet to within 10\\% of the radial-velocity measured mass. We summarize our conclusions in Sect. \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  The mathematical tools presented in this article can be used for a complete parametrization of transiting exoplanet systems on the basis of high-accuracy LCs only. In our model, for a specific range of systems, RV measurements are not necessary to constrain the mass of the planet ($M_\\mathrm{p}$), orbital eccentricity ($e$), the orientation of periastron ($\\omega$), and the geometric albedo of the planet ($\\alpha_\\mathrm{g}$). Our model also incorporates the characterization of the ratio of planetary and stellar radius ($R_\\mathrm{p}/R_\\star$), orbital period ($P$), and the orbital inclination ($i$). In order for this method to be fully applicable, the planet must be massive enough and orbit its host star close enough as to distort the stellar structure significantly.\\par With the current Kepler and potential future missions such as PLATO, we are able to measure the mass of the hot Jupiters, where the orbital periods must be $\\leq~1.5$, $\\leq~4$, and $\\leq~10$ days in order to measure planetary masses of $0.5\\,M_\\mathrm{J}$, $1\\,M_\\mathrm{J}$, and $5\\,M_\\mathrm{J}$ respectively. Furthermore, for host stars with spectral types earlier than FV, thermal emission component is detectable in the LC (assuming black bodies). This will affect the accuracy and even the detection of ellipsoidal variation. \\par As we show, the Kepler mission provides an accuracy suitable enough for our procedure to be applied for the characterization of extrasolar planets. We confirm the planetary natures of the hot Jupiters HAT-P-7 from Kepler data alone. Our method yields masses of $M_\\mathrm{p} = 1.27^{+0.10}_{-0.45}$ $M_\\mathrm{J}~\\times~M_\\star/M_\\odot$ for HAT-P-7b. Our results ($M_\\mathrm{P}$, $\\alpha_\\mathrm{g}$, $\\alpha_\\mathrm{bol}$, $e$, $T_{d}$, $T_{n}$) are in good agreement with \\citet{2010ApJ...713L.145W} and we also confirm the ellipsoidal variations of HAT-P-7b system. Using ellipsoidal variations we calculate the planetary mass from the photometric LC itself, without any RV measurments. Our technique will benefit from future space missions such as PLATO and the James Webb Space Telescope \\citep{2009PASP..121..952D} with $\\mathrm{RMS}~\\sim~3\\cdot 10^{-5}$ and $\\lesssim 10^{-6}$, respectively."
    },
    "1112/1112.2006.txt": {
        "abstract": "Evaporative evolution of stellar clusters is shown to be relaxation limited when the number of stars satisfies $N>>N_c$, where $N_c\\simeq 1600$. {For a Maxwell velocity distribution that extends beyond the escape velocity,} this process is {\\em bright} in that the Kelvin-Helmholtz time scale, $f_H^{-1}t_{relax}$, is shorter than the Ambartsumian-Spitzer time scale, $f_N^{-1}t_{relax}$, where $f_H>f_N$ denote the fractional changes in total energy and number of stars per relaxation time, $t_{relax}$. The resulting evaporative lifetime $t_{ev}\\simeq 20.5 t_{relax}$ for isolated clusters is consistent with Fokker-Planck and N-body simulations, where $t_{relax}$ is expressed in terms of the half-mass radius. We calculate the grey body factor by averaging over the anisotropic perturbation of the potential barrier across the tidal sphere, and derive the tidal sensitivity ${d\\ln t_{ev}}/{dy}\\simeq -1.9$ to $-0.7$ as a function of the ratio $y$ of the virial-to-tidal radius. Relaxation limited evaporation applies to the majority of globular clusters of the Milky Way with $N=10^4-10^6$ that are in a pre-collapse phase. It drives streams of stars into the tidal field with a mean kinetic energy of 0.71 relative to temperature of the cluster. Their $S$ shape morphology leads in sub-orbital and a trails in super-orbital streams separated by $3.4\\sigma/\\Omega$ in the radial direction of the orbit, where $\\Omega$ denotes the orbital angular velocity and $\\sigma$ the stellar velocity dispersion in the cluster. These correlations may be tested by advanced wide field photometry and spectroscopy. ",
        "introduction": "Globular clusters represent some of the oldest stellar systems as indicated by their stellar ages of at least 12 Gyr \\citep{pee68,van96,han02} consistent with the ages of white dwarfs in the thin disk of the Milky Way \\citep{noh92,sal00,del08,ise08,ren10}, the oldest halo field stars such as HE 2523-0901 of 13.2 Gyr$\\pm2.7$ Gyr (13.4$\\pm$1.8 Gyr based on U/r  chronometry alone, \\cite{fre07}) and the age of the Universe of 13.75 $\\pm 0.11$ Gyr \\citep{jar10}. The population of clusters of the Milky Way as observed today is but a shadow of a primordial population \\citep{che90,mcl08,pri08}, that may have formed in the era of reionization \\citep{van96}. Their ages hereby exceed their relaxation times by a factor of a few to an order of magnitude. They are essentially virialized close to thermal equilibrium \\citep{mey97} and some 80\\% satisfy the King model (KM, \\cite{kin66,che90}), which bears out well also in numerical simulations \\citep{cha10}. The exception is a minor fraction of clusters in the direction of the galactic halo with a centrally peaked core, commonly designated as post core-collapse (PCC) \\citep{che89,tra93,joh08}.  Prior to core collapse, the evolution of stellar clusters is well-described by two-body interactions in the small angle scattering approximation \\citep{spu96,aar98}, that naturally gives rise to evaporation of the cluster \\citep{amb38,spi40}. This process is amenable to detailed modeling by Fokker-Planck and N-body simulations \\citep{hen61,spu96}, there results of which can be interpreted statistically (reviewed in \\cite{mey97}) in terms of evaporation, tidal effects, stellar mass loss, mass segregation and heating and crossings of the galactic disk \\citep{gne97,wat00,jos01,fre02,heg01,cha10,heg11}.  There is an increasing interest in these various processes to develop efficient semi-analytic approaches to the birth and evolution of large numbers of stellar clusters as tracers of cosmological evolution of galaxies \\citep{kru11}. Evaporation stands out as a common process in many-body systems that includes for instance atmospheric science (e.g. \\cite{cha63,hay65,cha67,shi86,joh08,joh10}), where it extracts the lightest elements from an otherwise nearly thermal state (e.g. \\cite{cha96,cha02}).  Evaporation is a radiation process that defines a finite Kelvin-Helmholtz time scale due to energy extraction by the escaping stars. Key to their evolution, therefore, is an accurate determination of the energy loss carried off by stars escaping the cluster a determination of the time scale of evaporation, in particular whether it follows the relaxation time scale or the crossing time scale \\citep{heg01} or possibly a combination thereof \\citep{bau01}. The energy extraction by escaping stars is potentially useful in classifying post core-collapse clusters \\citep{kup08b}.  Evaporation of stellar clusters has received considerable attention (reviewed by \\cite{joh93}) with an emphasis on mass loss typically more so than energy loss (but see \\cite{wie67}). Yet, the {  energy} distribution function of all stars, those remaining in the globular cluster and those expelled by evaporation, continuously spreads towards increasingly negative and positive values associated with contraction and heating. The kinetic energies of evaporated stars in the halo hereby carry a thermal imprint of the history of the clusters. Since the {  energy} distribution function is unbounded in the point particle limit, there is the potential for extended lifetimes already at small binary mass fractions. However, observational evidence for a dynamically significant population of compact binaries has not been demonstrated, nor in the correlation between the observed population of luminous X-ray binaries and the binary star-star interaction rate \\citep{poo03} that points to relatively small binary populations that do not necessarily affect the lifetime of the clusters.  In this paper, we explore the parameter range for relaxation-limited evaporation{, the effect of energy loss,} and some observational consequences for the present state of globular clusters in the Milky Way.  We discuss the thermodynamics of evaporating globular clusters with tidal perturbation in the large-N limit of equal mass stars as a function of a tidal parameter \\begin{eqnarray} y=\\frac{R}{r_t}, \\label{EQN_t03} \\end{eqnarray} representing the ratio between the virial radius $R$ (close to the half-mass radius, to be discussed further in \\S2 below) and the tidal radius $r_t$ of the cluster. For old clusters such as of the Milky Way, $r_t$ is much larger than the radius of the cluster. The tidal sphere has the effect of lowering the potential barrier for stars to escape along orbits into the the halo of the galaxy and, generally, introduces strong anisotropy in the directions of the escaping stars defined by the Lagrange points $L_1$ and $L_2$ \\citep{hay67}. Existing approaches to incorporate tidal effects commonly use the minimum energy for escape at $L_1$ \\citep{spi87,lee87,che90,tak97,ros97}, effectively applying a spherically symmetric extension thereof to the tidal sphere. However, since last scatterings arise typically from within the half-mass radius \\citep{bau02}, the escape rate is effectively determined by the probability of escape derived from spherical averaging over the anisotropic tidal field for randomly oriented ejections. We thus derive the evaporative lifetime up to the moment of core-collapse to be \\begin{eqnarray} t_{ev} \\simeq 20.53  \\left(1-1.90y+2.78y^2-2.40y^3+0.84y^4+O(y^5)\\right)t_{relax}, \\label{EQN_t04} \\end{eqnarray} where the relaxation time $t_{relax}$ expressed in terms of the half-mass radius $R_h$ (the coefficient is $13.85$ when $t_{relax}$ is expressed in $R$). While our approach is {a largely explorative extension to \\cite{amb38} focused on the sensitivity of the evaporative lifetime to energy loss, some} confidence can be gained by comparing (\\ref{EQN_t04}) with known results reviewed in \\cite{spi87,tak97,mey97,heg03}.  Early in their evolution over the first few Gyr, globular clusters may have experienced considerable mass loss in stellar winds \\citep{aar98,bau03} due to stellar evolution (e.g. \\cite{app86}), that created a crucial factor in the likelihood of survival of a cluster after its formation \\citep{che90}. A common approximation used in Fokker-Planck and N-body simulations is an exponential evolution on a time scale $t_m$ \\citep{aar85,che90,mey97}. The same approach is readily incorporated in a thermodynamic model as a regular perturbation provided that $t_m>>t_{relax}$. Evaporation has also been found to narrow the mass spectrum due to a relatively enhanced rate of escape of the lightest stars,  by which older clusters tend evolves similarly to a system of equal mass stars \\citep{che90}. The aforementioned analogy with atmospheric science is useful here, noting that evaporation leads to an early escape of light molecules such as hydrogen and helium from the terrestrial atmosphere leaving only the relatively heavy molecules of nitrogen and oxygen to survive. Early evolution of star clusters therefore deviates from the evaporation driven evolution of single mass clusters. However, for the old clusters of the Milky Way at the present epoch, these considerations are less critical, and will not be considered here further.  As an application, we focus on high velocity stars in and around globular clusters that may provide ``real-time\" snapshot of last scattering events \\citep{hen61}. These last scattering events from within the half-mass radius (in the limit of isolated clusters, \\cite{bau02}) and most frequently derive from the core, where the density is highest and the local relaxation time is the shortest \\citep{aar74,joh93}. This includes tidal tails produced by ejections that, indirectly, serves as a diagnostic for the cluster temperature as, e.g., in NGC 6254 \\citep{leo00} and 47 Tuc \\citep{lan10a,lan10b}, where the velocity dispersion rises to about 60\\% at the tidal radius relative to the velocity dispersion inside the cluster. Observational data on tidal tails hereby provides a signature that may be used to calibrate numerical models when applied to individual clusters, e.g., M15 \\citep{gri95,dru98}, Palomar 5 \\citep{ode01,mey01,deh04}, NGC 5466 \\citep{bel06} and Palomar 14 \\citep{sol11} and their tails by $N-$body simulations (e.g. \\cite{heg01,lee06}). This may include further processes such as additional heating by past disk crossings for clusters near the center of the galaxy, but possibly less so for the outer clusters \\citep{gne97,leo00}. Some of these tidal tails display an $S-$shaped morphology \\citep{heg01,bel06,kro08} along the Lagrange points $L_1$ and $L_2$ on short scales, e.g., in NGC 5904 \\citep{leo00,jor10},  and along the orbit of the cluster on long scales, e.g., NGC 5644 \\citep{fel07}, where the leading (trailing) tail is inside (outside) the cluster orbit around the galactic center. The temperature of the clusters varies broadly, from about 1 km/s in Pal 12 \\citep{leo00} to about 15 km/s in $\\omega$-Centauri \\citep{van06}, with evidently different prospects for the manifestation of these tidal tails.  The planned Large Synoptic Survey Telescope \\citep{lsst09}, the Baryon Oscillation Spectroscopic Survey (BigBOSS, a multi-object spectrograph, \\cite{sch09}) the Japan Astrometry Satellite Mission for INfrared Exploration (JASMINE \\cite{gou10}), and GAIA \\citep{lin95} seem to be ideally suited for studying these tidal tails by large field of view photometry and spectroscopy. This development offers novel opportunities for testing the physics of evaporation by measuring correlations between the kinematics of stars in tails to the thermal state of the clusters.  In \\S2, we define notation and the basic quantities describing the thermodynamic state of virialized star clusters, their crossing and relaxation time. In \\S3, we introduce the general set-up for the evolution equations. in \\S4, the parameter range for relaxation-limited evaporation is identified. In \\S5, we calculate the evaporation coefficients for the number of stars and energy emitted per unit of relaxation time. In \\S6, we extend our model by regular perturbation analysis to include tidal interactions. In \\S7, we discuss the remaining relaxation-limited evaporative lifetime of the globular clusters in the Milky Way.  We give model predictions for observational test on tidal tails in \\S8 and summarize our results in \\S9. ",
        "conclusions": "Stellar clusters are open systems that radiate stars and energy by evaporation{, mostly by diffusion in momentum space subject to a singular perturbation by rare but finite jumps due to hard 2- or possibly 3-body interactions}. We identify a parameter range for relaxation-limited evaporation applicable to clusters in their pre-collapse phase (generally of KM morphology) with $N>>N_c$ and small to moderate tidal ratios $y<1$, and develop a model for the evaporative lifetime for single-mass clusters. {For a Maxwell velocity distribution that extends beyond the escape velocity, evaporation is} {\\em bright} as characterized by an inequality between the Kelvin-Helmholtz (\\ref{EQN_tKH}) and Ambartsumian-Spitzer time scales (\\ref{EQN_fev}), satisfying \\begin{eqnarray} t_{KH}<t_{AS} \\label{EQN_KHAS} \\end{eqnarray} due to $f_H>f_N$, wherein evaporation is largely driven by energy-loss rather than mass loss. The evolution isolated clusters, as a reference to more realistic modeling, is hereby described by a finite-time singularity solution \\begin{eqnarray} \\{N, R, T \\}\\propto  \\left(1-\\frac{\\tau}{\\tau}_{ev}\\right)^{\\{0.0905,~0.6351,~-0.5405\\}}. \\label{EQN_RT} \\end{eqnarray} It shows that the temperature and the radius evolutions (\\ref{EQN_RT}) are relatively less shallow than (\\ref{EQN_SN}), but still evolve only moderately away from the time of complete evaporation according to (\\ref{EQN_RT}). Here, the indices $0.0905$ and 0.6351 in (\\ref{EQN_RT}) are close to the values 0.085 and 0.695, respectively, observed in N-body simulations \\citep{heg03}. We do not attribute the relatively small $t_{ev}$ (relative to $t_{AS}$) to a radial stratification of the relaxation time with $f_H=0$ as proposed in \\citep{joh93}. It also {at odds with} an older claim that $f_H<f_N$ (with commensurably $t_{KH}>t_{AS})$ by \\cite{hoe58,kin58,kin58b}. Similarly, stars escaping by ejections from last scattering events from within the cluster should be contrasted with low velocity evaporation arising from diffusive escape through a low temperature outer boundary, commonly used in the Fokker-Planck approximations as if the cluster were opaque throughout up to its tidal boundary \\citep{spi87}.  {The agreement of (\\ref{EQN_t04}) with numerical simulations supports the idea that energy loss in accord with (\\ref{EQN_KHAS}) is important. For a Maxwell velocity distribution, (\\ref{EQN_t04}) is shorter than that based on particle loss alone by a factor of about three due energy loss (\\ref{EQN_fEb}). It would be of interest to identify the various scattering processes that may give rise to a relatively large $f_H$, see e.g. \\cite{pet70,ret79,joh93,mey01a,ash04}, and determine the implied shortening of the evaporative lifetime following \\S5.1. These studies fall outside the scope of the present work, however. Instead, we focus on some observational tests in \\S8.}  The presence of a tidal field modifies the net rate of evaporation by two effects. It generally enhances evaporation by lowering the threshold on the kinetic energy for stars to spill over the tidal radius, but the escape probably is attenuated by the generally anisotropic tidal field as defined by the Lagrange points $L_1$ and $L_2$ described by the grey body factor (\\ref{EQN_grey}). For moderate strengths of the tidal field, the resulting evaporation rate can be expanded in a regular perturbation in the tidal ratio (\\ref{EQN_t03}).  In studying the globular clusters of the Milky Way, we note that they (1) amply satisfy the criterion for relaxation-limited evaporation (\\ref{EQN_Nc}), (2) their tidal rations satisfy $y<<1$, (3) most of them are of KM morphology indicative of a pre-core collapse state. For this reason, the perturbative expansion (\\ref{EQN_t04}) for their remaining relaxation-limited evaporative lifetimes is believed to be applicable, and it is found to be consistent with their current ages.  The energy of the escaping stars is directly correlated to the temperature of the cluster (\\ref{EQN_j}), and the outflows thus produced develop a bi-polar $S-$shaped morphology in the plane of the Lagrange points $L_1$ and $L_2$ and the orbit of the cluster. For those tidal tails where the $S-$shaped tail can be resolved observationally, the de-projected orbital displacement of the tails is representative for the outflow velocity by the correlations (\\ref{EQN_V2}-\\ref{EQN_j}).  If periodically destroyed by disk crossings, these tails hereby re-appear on a time scale of 10 Myr with a length scale of a few times the tidal radius for clusters that are not too small in mass and velocity dispersions. Detailed observational studies are proposed to test (\\ref{EQN_V2}) on a wide field of view around globular clusters using upcoming high resolution photometric and spectroscopic stellar survey instruments. The results may also be used for comparison with alternative mechanisms for inducing tidal tails, such by tidal heating or shocks in response to disk crossings \\citep{heg01} and velocity anisotropy \\citep{tak00}. We remark that destruction by tidal fields of an initial distribution might have been more severe around spiral galaxies than around the relatively larger elliptical galaxies, which tend to have relatively large numbers of globular clusters in proportion to their luminosity \\citep{har81,bur10}.  The upcoming large field of view photometric and spectroscopic instruments LSST, BigBOSS, GAIA and JASMINE seem to be ideally suited for detailed measurements on tidal tails in energy (\\ref{EQN_V2}), stellar count (\\ref{EQN_n}) and morphology (\\ref{EQN_j}) to advance an observational test of the present thermodynamic model.  {\\bf  Acknowledgment.} The author gratefully acknowledges detailed constructive comments from the anonymous referee, J. Scalo and H. van Beijeren."
    },
    "1112/1112.3462_arXiv.txt": {
        "abstract": "{ Recent observations of the late ($t=10^8$--$10^9$ s) emission of the supernovae (SNe) associated to GRBs (GRB-SN) show a distinctive emission in the X-ray regime consistent with temperatures $\\sim 10^7$--$10^8$ K. Similar features have been also observed in two Type Ic SNe SN 2002ap and SN 1994I that are not associated to GRBs. } { We advance the possibility that the late X-ray emission observed in GRB-SN and in isolated SN is associated to a hot neutron star just formed in the SN event, here defined as a neo-neutron star. } { We discuss the thermal evolution of neo-neutron stars in the age regime that spans from $\\sim 1$ minute (just after the proto-neutron star phase) all the way up to ages $<$ 10--100 yr. We examine critically the key factor governing the neo-neutron star cooling with special emphasis on the neutrino emission. We introduce a phenomenological heating source, as well as new boundary conditions, in order to mimic the high temperature of the atmosphere for young neutron stars. In this way we match the neo-neutron star luminosity to the observed late X-ray emission of the GRB-SN events: URCA-1 in GRB980425-SN1998bw, URCA-2 in GRB030329-SN2003dh, and URCA-3 in GRB031203-SN2003lw. } { We identify the major role played by the neutrino emissivity in the thermal evolution of neo-neutron stars. By calibrating our additional heating source at early times to $\\sim 10^{12}$--$10^{15}$ erg/g/s, we find a striking agreement of the luminosity obtained from the cooling of a neo-neutron stars with the prolonged ($t=10^{8}$--$10^{9}$ s) X-ray emission observed in GRB associated with SN. It is therefore appropriate a revision of the boundary conditions usually used in the thermal cooling theory of neutron stars, to match the proper conditions of the atmosphere at young ages. The traditional thermal processes taking place in the crust might be enhanced by the extreme high-temperature conditions of a neo-neutron star. Additional heating processes that are still not studied within this context, such as $e^+e^-$ pair creation by overcritical fields, nuclear fusion, and fission energy release, might also take place under such conditions and deserve further analysis. } { Observation of GRB-SN has shown the possibility of witnessing the thermal evolution of neo-neutron stars. A new campaign of dedicated observations is recommended both of GRB-SN and of isolated Type Ic SN. } ",
        "introduction": "The investigation of the thermal evolution of neutron stars is a powerful tool for probing the inner composition of these objects. The cooling of neutron stars has been investigated by many authors, where many different microscopic models were assumed \\citep[see][]{Schaab1996,Page2004,Page2006,Page2009,Blaschke2000, Grigorian2005,Blaschke2006,Negreiros2010}. Most of the research on the thermal evolution of compact stars focus on objects older than 10 to 100 years, which is comprehensible if one considers that the thermal data currently available to us is for pulsars with estimated ages the same as or greater than 330 years \\citep{Page2004,Page2009}. We discuss the thermal evolution of young neutron stars, in the little explored time window that spans ages greater than 1 minute (just after the proto-neutron star regime \\citep{Prakash2001}) to ages $\\leq 10$--$100$ years, when the neutron star becomes isothermal \\citep[see][for details]{Gnedin2001}.  We discuss the possibility that the late X-ray emission (URCA hereafter \\footnote{The names URCA-1 and URCA-2 mentioned here were given to these sources when presented for the first time at the MG10 meeting held in Rio de Janeiro in the town of Urca. The location of the MG10 meeting was very close to the ``Cassino da Urca'' where George Gamow and Mario Schoenberg conceived the process of neutrino emission for the cooling process of neutron stars, which also took the name from the town of Urca, the Urca process \\citep[see e.g detailed history in][]{2005tmgm.meet..369R,1970mwla.book.....G}}) following a few GRBs associated with supernovae (SNe); e.g. URCA-1 in GRB980425-SN1998bw \\citep{2004AdSpR..34.2715R,2005tmgm.meet.2451F,GraziaMG11}, URCA-2 in GRB030329-SN2003dh \\citep{2004AIPC..727..312B,2005tmgm.meet.2459B}, and URCA-3 in GRB031203-SN2003lw \\citep{2005ApJ...634L..29B,2007ESASP.622..561R,2008ralc.conf..399R} (see Fig.~\\ref{fig:xray_lc} for details), might actually be created from young ($t \\sim$ 1 minute--$(10$--$100)$ years), hot ($T \\sim 10^7$--$10^8$ K) neutron stars that are remnants of the SN \\citep{2007ESASP.622..561R} and which we here call neo-neutron stars. Relevant also are the observations of the isolated Type Ic supernova SN 1994I \\citep{2002ApJ...573L..27I} and SN 2002ap \\citep{2004A&A...413..107S}, which present late emissions similar to the ones observed in URCA-1, URCA-2, and URCA-3.  Here we propose a revision of the boundary conditions usually employed in the thermal cooling theory of neutron stars, in order to match the proper conditions of the atmosphere at young ages. We also discuss the importance of the thermal processes taking place in the crust, which also have strong effects on the initial stages of thermal evolution. We stress that we are not calling the validity of the current treatment of the atmosphere of compact stars into question but, instead, we point out the need of extending them to appropriately describe the conditions of neo-neutron stars. ",
        "conclusions": "The major role played by the neutrino emissions from the crust of a neo-neutron star in its initial stages of the object is illustrated in Fig.~\\ref{fig:cool_crust_1}. In addition, by calibrating our additional heating source at early times to $H_0 \\sim 10^{12}$--$10^{15}$ erg/g/s, we find striking agreement of the luminosity obtained from the cooling of a neo-neutron stars with the prolonged ($t=10^{8}$--$10^{9}$ s) X-ray emission observed in GRB associated with an SN (see Fig.~\\ref{fig:cool_crust_1} for details). This could indicate that something might be missing in our current understanding of the crust of neutron stars. It might be that, as is the case for the atmosphere, we need to develop our current models for the crust further, so as to describe properly the properties of neo-neutron stars. The traditional thermal processes taking place in the crust might be enhanced by the extremely high temperature conditions of neo-neutron star and, additional heating processes not yet studied within this context could also take place under such conditions and deserve further analysis.  Particularly interesting in this respect are the processes of $e^+e^-$ pair creation expected to occur in the interphase between the core and the crust during the neutron star formation leading to the appearance of critical fields \\citep[see][for details]{ 2007IJMPD..16....1R,2008pint.conf..207R,2010JKPS...57..560R,2010AIPC.1205..143R,2010AIPC.1205..127P,2010PhR...487....1R,2011PhRvC..83d5805R,2011PhLB..701..667R,2011IJMPD..20.1789R,2011IJMPD..20.1995R,2011arXiv1107.2777R,2011arXiv1104.4062R}  It is also worth mentioning that the additional heating source needed at early times, $H_0 \\sim 10^{12}$--$10^{15}$ erg/g/s (or $H_0 \\sim 10^{-6}$--$10^{-3}$ MeV/Nucleon/s), is in striking agreement with the heat released from nuclear fusion reactions, radiative neutron captures, and photodisintegrations in the early stages of neutron star mergers found by \\cite{2011ApJ...738L..32G,2011A&A...531A..78G}. Both fission and $\\beta$-decays have also been included there: i.e neutron-induced fission, spontaneous fission, $\\beta$-delayed fission, photofission, and $\\beta$-delayed neutron emission.  All this suggests the exciting possibility that we are, for the first time, observing a nascent hot neutron star. This possibility alone warrants further study of this subject, so we might obtain a more concrete picture of the thermal evolution of neo-neutron stars. A proposal has been recently submitted by E.~Pian et al. to the Chandra satellite to observe whether a similar prolonged X-ray emission also exists in GRB100316D that is associated with SN2010bh \\citep{pian2011}. We also encourage dedicated observations of isolated SNe in view of the similarities between URCA-1--URCA-3 and the Type Ic supernova SN 1994I \\citep{2002ApJ...573L..27I} and SN 2002ap \\citep{2004A&A...413..107S}."
    },
    "1112/1112.1698_arXiv.txt": {
        "abstract": "We explore the circumgalactic medium (CGM) of two simulated star-forming galaxies with luminosities $L \\approx 0.1$ and $1 L^\\star$ generated using the smooth particle hydrodynamic code {\\sc GASOLINE}. These simulations are part of the Making Galaxies In a Cosmological Context (MaGICC) program in which the stellar feedback is tuned to match the stellar mass--halo mass relationship.  For comparison, each galaxy was also simulated using a ``lower feedback'' (LF) model which has strength comparable to other implementations in the literature.  The ``MaGICC feedback'' (MF) model has a higher incidence of massive stars and an $\\approx 2\\times$ higher energy input per supernova.  Aside from the low mass halo using LF, each galaxy exhibits a metal-enriched CGM that extends to approximately the virial radius.  A significant fraction of this gas has been heated in supernova explosions in the disk and subsequently ejected into the CGM where it is predicted to give rise to substantial \\iovi\\ absorption.  % The simulations do not yet address the question of what happens to the \\iovi\\ when the galaxies stop forming stars.  Our models also predict a reservoir of cool \\ihi\\ clouds that show strong \\lya\\ absorption to several hundred kpc.  Comparing these models to recent surveys with the {\\it Hubble Space Telescope}, we find that only the MF models have sufficient \\iovi\\ and \\ihi\\ gas in the CGM to reproduce the observed distributions.  In separate analyses, these same MF models also show better agreement with other galaxy observables (e.g.\\ rotation curves, surface brightness profiles, and HI gas distribution).  We infer that the CGM is the dominant reservoir of baryons for galaxy halos.% ",
        "introduction": "In the paradigm of cold dark matter (CDM) cosmology, gravitational collapse leads to the formation of virialized and bound  dark matter halos.  It is predicted that baryons fall into these halos with the dark matter, comprising a mass fraction up to the cosmological ratio.  Recent analyses indicates that stars comprise less than 25\\% of the baryons that should have collapsed into halos \\citep{Conroy2009,Mandelbaum2009,Moster2010,Guo2010,More2011}. This presents a ``missing baryon'' problem reflected in the unknown whereabouts of the baryons, which did not cool to form stars. Substantial amounts of gas are found in a diffuse and highly ionized medium which is referred to as halo gas or the circumgalactic medium (CGM).  The CGM is known to manifest in at least two phases: (i) a warm/hot, collisionally ionized gas with temperature near the virial temperature ($T \\gtrsim 10^6$K); and (ii) a cooler phase (perhaps predominantly photoionized) with temperature $T \\lta 10^4$K. The latter phase has been detected in  21cm emission in our Galaxy in the population known as the high velocity clouds \\citep[e.g.][]{Wakker1997} lying at distances of tens of kpc \\citep{Thom2006}. The hotter phase may be observed via bremsstrahlung emission with X-ray telescopes, in the intracluster medium of massive galaxy clusters \\citep[e.g.][]{Allen2001}, but so far extended diffuse emission has only been detected around one disc galaxy, the massive NGC1961 \\citep{Anderson2011}.  Smaller disc galaxies have X-ray bright outflows \\citep{Strickland2004}, but halo emission remains undetected.  The diffuse nature of the CGM makes positive detections of emission rare.  The majority of empirical constraints, therefore, come from absorption-line analysis of galactic halos, along coincident sightlines to distant quasars and galaxies \\citep[e.g.][]{Bowen1996,Lanzetta1995,Chen2008,Rubin2010a,Steidel2010}. These observations reveal a cool and frequently metal-enriched phase traced by \\ihi\\ Lyman series absorption \\citep[e.g.][]{Chen2003} and low-ionization metal-line transitions \\citep[e.g.][]{Chen2008,Barton2009,Rubin2011}. Importantly, galaxies of essentially all luminosity and spectral type exhibit significant \\ihi\\ absorption out to impact parameters $R \\approx 300$\\,kpc \\citep{Wakker2009,Prochaska2011}, implying a massive ($M \\sim 10^{10} \\msun$) extended CGM.  A particularly useful tracer of the hot phase in quasar absorption line analysis is the \\iovi\\ doublet, which occurs in highly ionized and enriched regions of the Universe.  If collisionally ionized, the gas has temperature $T \\sim 10^5-10^6$\\,K. \\iovi\\ is observed along the majority of sightlines through our Galactic halo and is believed to trace coronal material on scales of 10--100kpc \\citep{Sembach2003}.  This gas is also associated with the CGM of local $L>0.1L^\\star$ galaxies  \\citep{Stocke2006,Wakker2009}.  \\cite{Prochaska2011} have further demonstrated that the extended CGM ($R \\sim 300$\\,kpc) of $L \\gtrsim 0.1 L^\\star$ galaxies has a high covering fraction to \\iovi\\ and can account for all of the \\iovi\\ detected in the present-day universe. Most recently, \\cite{Tumlinson2011} report a nearly 100\\%\\ incidence of strong \\iovi\\ absorption for the halos of $L \\sim L^\\star$ star-forming galaxies, with an ionized metal mass that likely exceeds that of the galaxies' interstellar media, demonstrating that the CGM is a major reservoir of highly ionized metals at $z \\sim 0$.  The observations reveal a multi-phase, highly ionized and metal-enriched CGM around present-day galaxies \\citep[also see][]{Thom2008,Savage2011,Savage2011b,Wakker2012}. The high degree of enrichment demands that a large mass of gas (and metals) is transported from galaxies and/or their progenitors  to the CGM.  Thus, these observations provide robust constraints on the processes of gas accretion and feedback which may be compared directly against models of galaxy formation.  Previous theoretical work on the CGM for individual galaxies has largely focused on analytic or semi-analytic treatments of idealized gas and dark matter/temperature profiles \\citep{Mo1996,Telfer2002,Maller2004,Binney2009}.  Although these works provide crucial insight into the nature of the CGM, they lack proper cosmological context (e.g.\\ mergers) and have generally not included the role of feedback from star-formation in the galaxy and its satellites.  The latter is now considered critical to match a wide-range of galaxy observables including the luminosity function, the stellar mass--halo mass relationship, and the baryonic Tully--Fisher relationship \\citep{Guo2011}. Galactic-scale winds likely resulting from stellar feedback are a nearly generic feature of star forming galaxies \\citep{Shapley2003,Weiner2009,Rubin2010} and may significantly influence properties of the CGM.  A number of groups have now used cosmological simulations to examine the nature and enrichment of the IGM at $z \\sim 0$ and offer comparisons to quasar absorption line observations \\citep[e.g.][]{Cen2006,Dav'e2010,Oppenheimer2011,shs+11,Cen2012}. These studies have not explicitly examined the physical nature of the CGM for individual galaxies and likely have insufficient resolution to perform such analysis.  A few authors have examined the distribution of gas within individual galactic halos, but have not yet treated the distribution of metals or the impact of feedback \\citep{kh09,stewart11}. At $z \\sim 3$, \\cite{Kawata2007} have analyzed numerical simulations of outflows in L$^\\star$ progenitors, and found that \\iovi\\  better reflects the strength of galactic winds than \\ihi.  Most recently, \\cite{fpk+11} and \\cite{Shen2011} have examined the metal enrichment of the CGM with the latter demonstrating that metals can be ejected to large distances at early times from L$^\\star$ progenitor galaxies while their gravitational potential remains low.   We have recently begun the MaGICC project to use sufficient stellar feedback to simulate galaxies that match the stellar mass--halo mass relationship. Early results of the project have shown that the ejection of low angular momentum gas via outflows \\citep{Brook2011a} redistributes angular momentum via large scale galactic fountains \\citep{Brook2011b} and thus play a crucial role in forming disc galaxies, particularly those without classical bulges.  In this way, we form galaxies that match scaling relations between rotation velocity, size, luminosity, colour, stellar mass,  halo mass, HI mass, baryonic mass and metallicity (Brook \\etal 2012, submitted).  \\citet{Macci`o2012} also showed that these galaxies have cored dark matter density profiles.  Here, we test the baryon cycle of these simulations with the constraints provided by observations of the CGM, and in particular the column densities of H\\,I and O\\,VI as observed at $z=0$. ",
        "conclusions": "We have studied CGM gas in high resolution galaxy formation simulations.  The stellar feedback is constrained using the stellar mass--halo mass relationship.  This constraint leads to a stellar feedback implementation that decreases star formation and ejects hot, metal-enriched gas into the CGM.  Inside the galaxies, the stellar feedback results in slowly rising rotation curves, flattened dark matter density profiles and HI mass-Luminosity, baryonic Tully--Fisher relations, and mass-metallicity relations that resemble observed galaxies \\citep{Brook2012}.  We compared these galaxy models to one with more standard feedback implementation \\citep{Stinson2010, Scannapieco2011}. The lower feedback simulations match none of these commonly observed internal properties of galaxies.  An examination of the CGM in our model reveals that:  \\begin{itemize} \\item The MaGICC feedback implementation produces extended, metal-enriched gaseous coronae that extend to impact parameters of $\\sim$300 kpc, even around low mass (L$_*\\sim0.1$L$_\\odot$) halos.  These results match observed absorption line features \\citep{Wakker2009,Prochaska2011, Tumlinson2011}. \\item The hot coronae extend beyond $r_{vir}$, though most of their mass remains contained within $r_{vir}$ of the simulated galaxies. The coronae extend furthest outside $r_{vir}$ in low mass systems. \\item  A reasonably strong feedback implementation, compared with those in the literature (our ``lower feedback'' case), is insufficient to match \\ihi\\ absorption features at $z=0$ and severely underestimates \\iovi\\,, performing particularly poorly in low mass ($\\sim0.1$ L$^\\star$) simulations. \\item The amount of hot gas in the LF and MF runs did not differ significantly, but rather: \\begin{itemize} \\item (i) The MaGICC feedback ejected ten times more oxygen out of the disc than the lower feedback in the high mass case and a factor of $10^4$ more in the low mass case. \\item (ii) The MaGICC feedback halo contains ten times the cold gas as the lower feedback case in both high and low mass simulations.  The cold gas results from cooling out of the hot halo gas and cold gas entrained in outflows. \\end{itemize} \\item The MaGICC feedback simulations match the observed \\ihi\\ and \\iovi\\ absorption columns at a range of masses. \\item The total mass of the coronae is several times higher than the stellar mass in the simulated MaGICC feedback galaxies.  Thus, the simulations of individual galaxies predict that the ``missing baryons'' are found in the CGM of galaxies. \\end{itemize}  In our simulations, we only considered galaxies defined as star forming by \\citet{Prochaska2011} and \\citet{Tumlinson2011}.  One of the key results of those works is the absence of OVI detection in quiescent galaxies.  Our simulations are unable to address the question of what happens to the OVI when galaxies stop forming stars.   Our examination focused on the radial distribution of \\ihi\\ and \\iovi\\ gas in the CGM of a simulated $L \\approx L^\\star$ galaxy and compared these results to observed surface density profiles.  There are, of course, additional observables related to the entire population of \\iovi\\ absorbers identified in quasar absorption line studies \\citep[e.g.][]{Danforth2008,Thom2008a,Tripp2008};  any truly successful model must also reproduce these resuts. One key constraint is the observed distribution of line-widths (a.k.a. Doppler parameters or $b$-values) for the \\iovi\\ gas \\citep{Thom2008a,Tripp2008} which gauge, in a statistical fashion, the contributions of thermal broadening, turbulence, and gas dynamics to the motions of the gas.  On average, the $b$-values for \\iovi\\ absorbers are substantial ($b > 20$\\kms) and appear to correlate with the strength of \\iovi\\ absorption.  This has posed a considerable challenge to models where \\iovi\\ is predominantly photoionized \\citep[e.g.][]{od09}, as such gas has a small thermal width.  Another valuable kinematic constraint is the common observation of co-aligned \\ihi\\ and \\iovi\\ absorption relative to a common redshift \\citep[approximately half of observed \\iovi\\ systems;][]{Thom2008}. Indeed, this appears to be at odds with the apparent separation of strong \\ihi\\ absorption (at smaller radii) with the extended \\iovi\\ absorption of our modeled CGM.  On the other hand, we do observe non-negligible \\ihi\\ absorption to large radii (Figure~\\ref{fig:6galhiradprof}) which may trace the \\iovi\\ gas.  A proper comparison of our model with these observables will require the generation and detailed analysis of line-profiles, as well as the examination of the CGM for galaxies spanning a wide range of masses.  This will be the focus of our next paper, which will consider whether the extended CGM of $z \\sim 0$ galaxies can reproduce all of the observed statistics for \\iovi\\ as argued by \\cite{Prochaska2011}.  We further note that the \\ihi\\ surface densities at $\\rho < 50$\\,kpc are sufficiently large that the CGM will yield significant absorption from lower ionization states of heavy metals (e.g.\\ MgII, SiII, SiIII).  A proper estimate of the column densities for these ions, however, will require a full treatment of radiative transfer (i.e.\\ to account for self-shielding by optically thick \\ihi\\ gas).  The fact that the resultant extended, metal enriched CGM around our simulated galaxies matches the observations over a wide range of mass provides strong support for a vigorous baryon cycle in which outflows and subsequent cooling of halo gas play a key role in forming disc galaxies."
    },
    "1112/1112.3654_arXiv.txt": {
        "abstract": "As the number of observed Gamma-Ray Bursts (GRBs) continues to grow, follow-up resources need to be used more efficiently in order to maximize science output from limited telescope time. As such, it is becoming increasingly important to  rapidly identify bursts of interest as soon as possible after the event, before the afterglows fade beyond detectability.  Studying the most distant (highest redshift) events, for instance, remains a primary goal for many in the field.\u00a0 Here we present our Random forest Automated Triage Estimator for GRB redshifts (\\rategrbz{}) for rapid identification of high-redshift candidates using early-time metrics from the three telescopes onboard \\textit{Swift}. While the basic \\rate{} methodology is generalizable to a number of resource allocation problems, here we demonstrate its utility for telescope-constrained follow-up efforts with the primary goal to identify and study high-$z$ GRBs.  For each new GRB, \\rategrbz{}  provides a recommendation---based on the available telescope time---of whether the event warrants additional follow-up resources.  We train  \\rategrbz{}   using a set consisting of 135 \\Swift{} bursts with known redshifts, only 18 of which are $z > 4$.  Cross-validated performance metrics on this training data suggest that $\\sim$56\\% of high-$z$ bursts can be captured from following up the top 20\\% of the ranked candidates, and $\\sim$84\\% of high-$z$ bursts are identified after following up the top $\\sim$40\\% of candidates.  We further use the method to rank $ 200+$ \\Swift{}  bursts with unknown redshifts according to their likelihood of being high-$z$. ",
        "introduction": "\\label{sec:intro}  As the most luminous electromagnetic explosions, gamma-ray bursts (GRBs) offer a unique probe into the distant universe---but only if their rapidly fading afterglows are observed before dimming beyond detectability \\citep[e.g.,][]{wijers98,escude98,lamb00,kawai08,McQuinn08}.  Since the launch of the \\Swift{} satellite in November 2004 \\citep{gehrels04}, more than 170 long duration \\Swift{} gamma-ray bursts have had measured redshifts, but only a handful fall into the  highest redshift range that allow for the probing of the earliest ages of the universe, up to less than a billion years after the Big Bang (Fig. \\ref{fig:reddist}).  With a limited budget of large-aperture telescope time accessible for deep follow-up, it is becoming increasingly important to rapidly identify these GRBs of interest in order to capture the most interesting events without spending available resources on more mundane events.  Along with quasars \\citep[e.g.,][]{Mortlock11} and NIR-dropout lyman-break galaxies \\citep[e.g.,][]{Bouwens10,Bouwens11}, GRBs have been established as among the most distant objects detectable in the universe, with a spectroscopically confirmed event at $z=8.2$ \\citep[GRB 090423;][]{tanvir09a,salvaterra09a} and a photometric candidate at $z \\sim 9.4$ \\citep[GRB 090429B;][]{Cucchiara11}. Such observations can provide valuable constraints on star formation in the early universe, illuminate the locations and properties of some of the earliest galaxies and stars, and probe the epoch of reionization. \\citep[e.g.,][and references therein]{Tanvir07}. Further, the relatively simple spectra of GRB afterglows compared to other cosmic lighthouses makes it easier to both identify their redshifts and extract useful spectral features such as neutral hydrogen absorption signatures for the study of cosmic reionization. \\citep[e.g.,][]{escude98,Barkana04,Totani06,McQuinn08}.  However, such benefits can only be realized if spectra are obtained with large-aperture telescopes before the afterglow fades beyond the level required to obtain a useful signal, typically within a day after the GRB.  As such, there has been a long-standing effort to extract a measure of a GRB's redshift from its early time, high-energy signal, with a primary goal of the rapid identification of high-$z$ candidates. This might appear in principle to be a straightforward exercise; for instance, distant GRBs should on average appear fainter and longer-duration than nearby events due to distance and cosmological time dilation, respectively.  In practice, however, the large intrinsic diversity of GRBs, as well as thresholding effects, confounds the straightforward use of early-time observations in divulging redshift and other important properties.  While much effort has gone into tightening the correlations between high-energy properties in order to homogenize the sample for use as a luminosity (and hence distance/redshift) predictor \\citep[e.g.,][]{amati02,ghirlanda04,firmani06,schaefer07}, there has been significant debate as to whether some of these relations are actually due to thresholding effects specific to the detectors rather than intrinsic physical properties of the GRBs \\citep[e.g.,][]{friedman05,butler07b,butler09}.  Regardless, whether or not these inferred relationships are actually physical or simply detector effects would not affect their utility as a {\\it detector-specific} parameter prediction tool.  By restricting ourselves to \\Swift{} events only, we avoid the uncertainty of whether certain correlations remain when using different detectors.  With this in mind, we set out to search for indications of high-redshift GRBs in the rich, mostly homogeneous dataset provided by 6+ years of GRB observations by the three telescopes onboard \\Swift{}  (BAT; \\citealt{barthelmy05}, XRT; \\citealt{burrows05}, UVOT; \\citealt{roming05}).  Past studies exploring high-$z$ indicators have used hard cuts on certain features such as UVOT afterglow detection, burst duration, and inferred hydrogen column density \\citep[e.g.,][]{grupe07,vandenberk08,ukwatta09}, regression on such features \\citep{koen09,koen10}, and combinations of potential GRB luminosity indicators \\citep{xiao09,xiao11}. In this work, we take a different approach by utilizing supervised machine learning algorithms, specifically Random Forest classification, to make follow-up recommendations for each event automatically and in real time. Particular attention is paid to careful treatment of performance evaluation by using cross-validation (\\S\\ref{sec:results}), a robust methodology to guard against over-fitting and the circular practice of testing hypotheses using the same data that suggested (and constrained) them.   The primary driving force of this study is simple: \\emph{given limited follow-up time available on telescopes, we want to maximize the time spent on high-$z$ GRBs}\\footnote{For the purposes of this study, ``high-redshift'' corresponds to all $z>4.0$: a compromise between only keeping the most interesting events and having enough data to train on.  However, we have explored performance of different redshift cuts; see \\S\\ref{sec:finalclassifier}.}. To this end, we provide a deliverable metric, explained in \\S\\ref{sec:deliverables}, to assist in the decision making process on whether to follow up a new GRB.  Real-time distribution of this metric is available for each new \\Swift{}  trigger via website\\footnote{\\url{http://rate.grbz.info/}} and RSS feed\\footnote{\\url{http://rate.grbz.info/rss.xml}}.  The structure of this paper is as follows: in \\S\\ref{sec:obs} we outline the collation of the data, and describe the particular GRB features utilized in redshift classification.  In \\S\\ref{sec:methods}, the Random Forest algorithm is detailed, along with some specific challenges posed by this particular data set.  Performance metrics of the classifiers are presented in \\S\\ref{sec:results}, and in \\S\\ref{sec:application} we discuss the results of testing the classifiers on additional GRBs, both with and without known redshifts.  Finally, our conclusions are given in \\S\\ref{sec:conclusions}.   \\begin{figure}[t!] \\centerline{\\plotone{Fig1.eps}} \\  \\caption[Redshift distribution of long-duration \\Swift{}  GRBs.] {Redshift distribution of the 135 long-duration \\Swift{} GRBs in our sample (Table \\ref{tab:trainingredshifts}).  For the purposes of this study, ``high'' redshift is defined as those bursts with redshifts larger than $z=4$, which corresponds to approximately 1-$\\sigma$ above the mean of the distribution.  In our sample, 18 bursts fall into this category (black, and in inset). In determining age since the Big Bang, we assume a cosmology with $h=0.71$, $\\Omega_m=0.3$, and $\\Omega_\\Lambda=0.7$. Solid lines show the cumulative number of GRBs as a function of redshift for high-$z$ bursts (grey) and all bursts (black). } \\label{fig:reddist} \\end{figure} ",
        "conclusions": "\\label{sec:conclusions}  In this paper, we presented the \\rategrbz{}  method for allocating follow-up telescope resources to high-redshift GRB candidates using Random Forest classification on early-time \\Swift{}  metrics. The \\rate{}  method is generalizable to any prioritization problem that can be parameterized as ``observe'' or ``don't observe'', and accommodates statistical challenges such as small datasets, imbalanced classes, and missing feature values. The issue of resource allocation is becoming increasingly important in the era of data-driven transient surveys such as PTF, Pan-STARRS, and LSST which provide extremely high discovery rates without a significant increase in follow-up resources. With enough training instances of any object of interest for a given transient survey, the \\rate{} method can be applied to prioritize follow-up of future high-priority candidates.  In the \\rategrbz{}  application, our robust, cross-validated performance metrics indicate that by observing just 20\\% of bursts, one can capture $56\\% \\pm 6\\%$ of $z > 4$ events with a sample purity of $37\\% \\pm 4\\%$.  Further,  following up on half of all events will yield nearly all ($96\\% \\pm 4\\%$) of the high-$z$ events.  The method provides a simple decision point for each new event: if the prioritization value $\\widehat{\\mathcal{Q}}$ is smaller than the percent of events a user wishes to allocate resources to, then follow-up is recommended. These rapid predictions, combined with the more traditional photometric dropout technique from simultaneous multi-filter NIR observatories (such as PAIRITEL, GROND, and the upcoming RATIR), offer a robust tool in more efficiently informing GRB follow-up decisions. To facilitate the dissemination of high-redshift GRB predictions to the community, we have set up a website (\\url{http://rate.grbz.info}) with $\\widehat{\\mathcal{Q}}$ values for past bursts, and an RSS feed (\\url{http://rate.grbz.info/rss.xml}) to provide real-time results from our classifier on new events."
    },
    "1112/1112.4903_arXiv.txt": {
        "abstract": "We derive the equations of motion, the periastron shift, and the gravitational radiation damping for quasicircular compact binaries in a massive variant of the Brans-Dicke theory of gravity. We also study the Shapiro time delay and the Nordtvedt effect in this theory. By comparing with recent observational data, we put bounds on the two parameters of the theory: the Brans-Dicke coupling parameter $\\omega_{\\rm BD}$ and the scalar mass $m_s$. We find that the most stringent bounds come from Cassini measurements of the Shapiro time delay in the Solar System, that yield a lower bound $\\omega_{\\rm BD}>40000$ for scalar masses $m_s<2.5\\times10^{-20}\\mathrm{eV}$ (or Compton wavelengths $\\lambda_s=h/(m_sc)>5\\times 10^{10}$~km), to 95\\% confidence. In comparison, observations of the Nordtvedt effect using Lunar Laser Ranging (LLR) experiments yield $\\omega_{\\rm BD}>1000$ for $m_s<2.5\\times10^{-20}\\mathrm{eV}$. Observations of the orbital period derivative of the quasicircular white dwarf-neutron star binary PSR J1012+5307 yield $\\omega_{\\rm BD}>1250$ for $m_s<10^{-20}\\mathrm{eV}$ ($\\lambda_s>1.2\\times 10^{11}$~km). A first estimate suggests that bounds comparable to the Shapiro time delay may come from observations of radiation damping in the eccentric white dwarf-neutron star binary PSR J1141-6545, but a quantitative prediction requires the extension of our work to eccentric orbits. ",
        "introduction": "\\label{sec:mbd_theory}  \\subsection{The generic scalar-tensor theory with a single scalar field}  A general class of scalar-tensor theories containing a single scalar field in addition to the tensor field was studied by Bergmann and Wagoner \\cite{Bergmann:1968ve,Wagoner:1970vr}. We can characterize the Bergmann-Wagoner theory via the following postulates:  \\noindent 1) The principle of general covariance is imposed, leading to tensorial equations.  \\noindent 2) The field equations are derived from the action \\begin{align} S=\\int\\big(\\mathscr{L}_\\mathrm{G}+\\mathscr{L}_\\mathrm{M}\\big)\\;d^4x, \\end{align} where $\\mathscr{L}_\\mathrm{G}$ and $\\mathscr{L}_\\mathrm{M}$ are the Lagrangian densities for the gravitational and matter fields, respectively.  \\noindent 3) We postulate that the long-range forces of nature are mediated by the three lowest spin bosons, and assume that the electromagnetic field is the only vector field.  This leaves a scalar degree of freedom $\\phi$ and a tensor degree of freedom (the metric $g_{\\mu\\nu}$) to describe the dynamics of the gravitational field.  \\noindent 4) The field equations are of at most of second differential order, and the tensor and scalar fields are nonminimally coupled; this leads us to the general form \\begin{align} \\mathscr{L}_\\mathrm{G}=(-g)^{\\f{1}{2}} \\big[h(\\phi)R+l(\\phi)g^{\\mu\\nu}\\phi_{,\\mu}\\phi_{,\\nu}+\\lambda(\\phi)\\big] \\end{align} for the gravitational Lagrangian density, where $h(\\phi)$, $l(\\phi)$ and $\\lambda(\\phi)$ are arbitrary functions of the scalar field $\\phi$.  \\noindent 5) We postulate a principle of \\emph{mutual coupling}, in which the matter Lagrangian density depends on the gravitational fields according to \\begin{align} \\mathscr{L}_\\mathrm{M}=\\mathscr{L}_\\mathrm{M}(\\psi^2(\\phi)g^{\\mu\\nu},\\Psi)\\,, \\end{align} where $\\psi(\\phi)$ is a fourth arbitrary function of $\\phi$, and $\\Psi$ represents the collective matter fields. This guarantees consistency with the strong equivalence principle \\cite{Fujii:2003pa}.  Now let us make the conformal transformation $g_{\\mu\\nu}\\rightarrow\\psi^2(\\phi)g_{\\mu\\nu}$, and in doing so move into a conformal frame in which the matter fields do not couple directly (but only indirectly, via the metric) to the scalar field; this is commonly referred to as the Jordan frame \\cite{Fujii:2003pa,Clifton:2011jh}.  Furthermore, without loss of generality we can redefine the scalar field such that $h(\\phi)\\rightarrow\\phi$. These two redefinitions recast the action into the form \\begin{eqnarray} \\label{action_general} S&=&\\f{1}{16\\pi}\\int \\left[ \\phi R -\\f{\\omega(\\phi)}{\\phi} g^{\\mu\\nu}\\phi_{,\\mu}\\phi_{,\\nu}+M(\\phi) \\right] (-g)^\\f{1}{2}d^4x \\nonumber \\\\ &+&\\int\\mathscr{L}_\\mathrm{M}(g^{\\mu\\nu},\\Psi)d^4x, \\end{eqnarray} which has the additional advantage that the resulting weak-field equations for $g_{\\mu\\nu}$ and $\\phi$ decouple from one another. The generic theory now contains two undetermined functions: the \\emph{cosmological function} $M(\\phi)$ and the \\emph{coupling function} $\\omega(\\phi)$ (in the language of \\cite{Will:1993ns}). The effect of the coupling function on compact binary dynamics has been studied extensively, and it can lead to interesting consequences if ``spontaneous scalarization'' occurs \\cite{Damour:1992we,Damour:1993hw,Damour:1995kt,Damour:1996ke,Damour:1998jk}. Here we focus on the cosmological function, which has three major effects in the generic theory. Firstly, in the resulting field equations for $g_{\\mu\\nu}$ it plays the role of a cosmological constant. Secondly, it endows the scalar with mass: this manifests itself most clearly in the fact that solutions for $\\phi$ for isolated systems contain Yukawa-like terms $e^{-m_sr}/r$, where $m_s$ is the mass of the scalar field, which in turn gives the field a characteristic range $\\ell \\sim1/m_s$ \\cite{Will:1993ns}. Finally, the cosmological function may introduce nonlinearities in the dynamics of the scalar field.  \\subsection{The matter action and the field equations} Let us now turn to the matter action. Throughout this paper we will make the assumption that all bodies in our system can be treated as point masses. Einstein, Infeld and Hoffmann (EIH) \\cite{Einstein:1938yz} developed a method for obtaining the equations of motion for a system of gravitating point-like masses. In their approach, one begins by obtaining the local gravitational field of a single body (in a comoving frame), under the assumption that the body is small and nearly spherical. One then proceeds to match the interbody gravitational fields onto the obtained local field of the single body under inspection; imposing self consistency yields the EIH equations of motion. The same equations of motion can be obtained with significantly less effort, albeit at the sacrifice of some rigor, by taking the stress-energy tensor to be a distribution of delta functions and neglecting any infinite self-energy terms as they arise \\cite{Will:1993ns}. In scalar-tensor theory, however, we must deal with the additional complication that the inertial mass and internal structure of a gravitating body will depend on the local value of the scalar field (i.e. the local value of the effective gravitational ``constant''). Variations in internal structure may act back on the motion of the body, leading to violations of the (weak) equivalence principle. Eardley \\cite{1975ApJ...196L..59E} showed that these effects could be accounted for by simply supposing that the masses of the bodies are in general functions of the scalar field, such that the matter action for a system of point-like masses can be written as \\begin{align} \\label{matter_action} S_\\mathrm{M}=-\\sum_a\\int m_a(\\phi)d\\tau_a, \\end{align} where the particles (labeled by $a$) have inertial masses $m_a(\\phi)$, and $\\tau_a$ is the proper time of particle $a$ measured along its worldline $x^\\lambda_a$. The distributional stress-energy tensor $T^{\\mu\\nu}$ and its trace $T$ hence take the form \\begin{align} T^{\\mu\\nu}(x^\\lambda)=(-g)^{-\\f{1}{2}}\\sum_am_a(\\phi) \\frac{u^\\mu u^\\nu}{u^0}\\delta^4(x^\\lambda-x^\\lambda_a)\\,, \\\\ T=g_{\\mu\\nu}T^{\\mu\\nu}=-(-g)^{-\\f{1}{2}} \\sum_a \\frac{m_a(\\phi)}{u^0} \\delta^4(x^\\lambda-x_a^\\lambda)\\,. \\end{align} Far from the system, the scalar will take on its cosmologically imposed value, denoted by $\\phi_0$. The relationship between the effective gravitational constant, $G$, and the scalar field $\\phi$ is therefore (in our chosen system of units) $G=\\phi_0/\\phi$. In the post-Newtonian limit, we expand $\\phi$ about its asymptotic value and define the small perturbation $\\varphi$ such that $\\phi=\\phi_0+\\varphi$. In this case, we can write the variation of the inertial masses $m_a$ with $\\phi$ as \\begin{align} \\label{mass_expansion} m_a(\\phi)= & m_a(\\mathrm{ln}\\;G)= m_a(\\phi_0) \\left[1+s_a\\Big(\\f{\\varphi}{\\phi_0}\\Big)\\right. \\nonumber \\\\ & -\\left. \\f{1}{2}\\big(s'_a-s_a^2+s_a)\\Big(\\f{\\varphi}{\\phi_0}\\Big)^2+O\\Big(\\Big(\\f{\\varphi}{\\phi_0}\\Big)^3\\Big)\\right]\\,, \\end{align} where we have defined the ``first and second sensitivities'' $s_a$ and $s'_a$ to be\\footnote{White-dwarfs typically have sensitivities $s_a\\sim10^{-4}$, neutron stars have sensitivities $s_a\\sim0.2$, and black holes have $s_a=1/2$: see \\cite{Zaglauer:1992bp} for detailed calculations.} \\begin{align} \\label{sensitivity} s_a=-\\f{\\partial(\\mathrm{ln}\\;m_a)}{\\partial(\\mathrm{ln}\\;G)}\\Big|_{\\phi_0}, \\quad s'_a=-\\f{\\partial^2(\\mathrm{ln}\\;m_a)}{\\partial(\\mathrm{ln}\\;G)^2}\\Big|_{\\phi_0}\\,. \\end{align}  The full action is now given by \\begin{align} \\label{action} S= & \\f{1}{16\\pi}\\int \\big[ \\phi R -\\f{\\omega(\\phi)}{\\phi} g^{\\mu\\nu}\\phi_{,\\mu}\\phi_{,\\nu}+M(\\phi) \\big](-g)^\\f{1}{2}d^4x \\nonumber \\\\ & -\\sum_a\\int m_a(\\phi)d\\tau_a. \\end{align} By varying the action \\eqref{action} with respect to the tensor and scalar fields, respectively, we obtain the full field equations of the generic theory described above: \\begin{widetext} \\begin{eqnarray} \\label{generic_tensor_fe} R_{\\mu\\nu} - \\frac{1}{2}g_{\\mu\\nu}R -\\f{1}{2}\\phi^{-1}M(\\phi)g_{\\mu\\nu}&=& 8\\pi\\phi^{-1}T_{\\mu\\nu} +\\omega(\\phi)\\phi^{-2}(\\phi_{,\\mu}\\phi_{,\\nu}-\\frac{1}{2}g_{\\mu\\nu}\\phi_{,\\alpha}\\phi^{,\\alpha})+\\phi^{-1}(\\phi_{,\\mu\\nu}-g_{\\mu\\nu}\\Box_g\\phi)\\,,\\\\ \\label{generic_scalar_fe} \\Box_g\\phi+  \\frac{\\phi \\frac{dM(\\phi)}{d\\phi}-2M(\\phi)}{\\big(3+2\\omega(\\phi)\\big)} \\nonumber &=&\\f{1}{3+2\\omega(\\phi)}\\Big(8\\pi T^*-\\f{d\\omega(\\phi)}{d\\phi}\\phi_{,\\alpha}\\phi^{,\\alpha}\\Big)\\,, \\end{eqnarray} \\end{widetext} where we have defined $T^*=T-2\\phi\\f{\\partial T}{\\partial\\phi}$ and $\\Box_g$ is the curved space d'Alembertian, defined by \\begin{align} \\Box_g=(-g)^{-\\frac{1}{2}}\\partial_\\nu((-g)^\\frac{1}{2}g^{\\mu\\nu}\\partial_\\mu)\\,. \\end{align} A detailed derivation of this result can be found in \\cite{Fujii:2003pa}.  \\subsection{Massive Brans-Dicke theory: The field equations and their weak-field limit}  As we recalled earlier, the effects of a generic coupling function on the dynamics of compact binaries have been studied fairly extensively by Damour and Esposito-Far\\'ese \\cite{Damour:1992we,Damour:1993hw,Damour:1995kt,Damour:1996ke,Damour:1998jk}. Here we are primarily interested in the effects of a nonzero mass of the scalar field. In the limit where $m_s\\to 0$, our final result for the dipolar and quadrupolar flux can be shown to match Eq.~(6.40) in \\cite{Damour:1992we} (the monopole contribution vanishes for circular orbits).  It would be interesting to study a theory with generic functional forms for both $\\omega(\\phi)$ and $M(\\phi)$, but for simplicity here we will consider a constant coupling function: $\\omega(\\phi)=\\omega_{\\rm BD}=\\mathrm{constant}$, as in the usual Brans-Dicke theory \\cite{Brans:1961sx}. The scalar field equation then reduces to \\begin{align} \\Box_g\\phi+\\f{\\phi\\f{dM(\\phi)}{d\\phi}-2M(\\phi)}{3+2\\omega_{\\rm BD}}=\\f{8\\pi T^*}{3+2\\omega_{\\rm BD}}\\,. \\end{align} In order to get a handle on the effects of the cosmological function $M(\\phi)$, let us expand the metric about a Minkowski background $\\eta_{\\mu\\nu}$ and the scalar field around its (cosmologically determined) constant background value $\\phi_0$. Following closely the method of \\cite{Will:1993ns}, we define small perturbations $\\varphi$, $h_{\\mu\\nu}$ and $\\theta_{\\mu\\nu}$ such that \\begin{align} \\label{perturb} & \\phi=\\phi_0+\\varphi, \\nonumber \\\\ & g_{\\mu\\nu}=\\eta_{\\mu\\nu}+h_{\\mu\\nu},\\;g^{\\mu\\nu}=\\eta^{\\mu\\nu}-h^{\\mu\\nu}, \\nonumber \\\\ & \\theta^{\\mu\\nu}=h^{\\mu\\nu}-\\frac{1}{2}h\\eta^{\\mu\\nu}-\\Big(\\frac{\\varphi}{\\phi_0}\\Big)\\eta^{\\mu\\nu}, \\end{align} Let us also expand $M(\\phi)$ in a Taylor series about $\\phi_0$: \\begin{align} \\label{Mphi} M(\\phi)=M(\\phi_0)+M'(\\phi_0)\\varphi+\\f{1}{2}M''(\\phi_0)\\varphi^2+\\dots \\end{align} We require that the expanded field equations are consistent at all orders in $(v/c)^n$. Substituting the weak field perturbations \\eqref{perturb} into the field equations \\eqref{generic_tensor_fe} and \\eqref{generic_scalar_fe} and examining the leading-order terms under the assumption of asymptotic flatness, we find that $M(\\phi_0)=M'(\\phi_0)=0$. We are therefore left with the quadratic term, that endows the scalar field with mass. To see this, let us substitute $M(\\phi)=\\f{1}{2}M''(\\phi_0)\\varphi^2$ into the scalar field equation, yielding \\begin{align} \\Box_g\\phi-m_s^2(\\phi-\\phi_0)=\\f{8\\pi T^*}{3+2\\omega_{\\rm BD}}\\,, \\end{align} where we have defined the mass of the scalar field \\begin{align} m_s^2\\equiv -\\f{\\phi_0}{3+2\\omega_{\\rm BD}} M''(\\phi_0)\\,. \\end{align} We will see shortly that $m_s$ is precisely the parameter appearing in Yukawa-type corrections $\\sim e^{-m_sr}$ to the Newtonian gravitational potential, as well as the ordinary mass parameter in the Klein-Gordon equation. Since the scalar field is expected to be small, we will neglect cubic and higher-order terms in $M(\\phi)$, that would introduce additional nonlinearities into the scalar field equation.  In summary, with our choice of coupling and cosmological functions, the field equations of the massive Brans-Dicke theory read \\begin{widetext} \\begin{eqnarray} \\label{mbd_tensor_fe} &&R_{\\mu\\nu} - \\frac{1}{2}g_{\\mu\\nu}R = - \\f{3+2\\omega_{\\rm BD}}{4\\phi_0\\phi}m_s^2(\\phi-\\phi_0)^2g_{\\mu\\nu} + \\frac{8\\pi}{\\phi}T_{\\mu\\nu} +\\frac{\\omega_{\\rm BD}}{\\phi^2} \\left(\\phi_{,\\mu}\\phi_{,\\nu}-\\frac{1}{2}g_{\\mu\\nu}\\phi_{,\\alpha}\\phi^{,\\alpha}\\right)+ \\frac{1}{\\phi}(\\phi_{,\\mu\\nu}-g_{\\mu\\nu}\\Box_g\\phi)\\,, \\\\ &&\\Box_g\\phi-m_s^2(\\phi-\\phi_0)=\\f{8\\pi T^*}{(3+2\\omega_{\\rm BD})}\\,. \\label{mbd_scalar_fe} \\end{eqnarray} \\end{widetext} \\subsection{The weak-field limit} Let us use the weak-field perturbations \\eqref{perturb} to obtain the field equations in the weak-field limit. Expanding the left hand side of \\eqref{mbd_tensor_fe} and imposing the harmonic gauge condition $\\theta^{\\mu\\nu}_{\\;\\;,\\nu}=0$ we find \\begin{align} R_{\\mu\\nu}-&\\f{1}{2}g_{\\mu\\nu}R=-\\f{1}{2}\\Box_\\eta\\theta_{\\mu\\nu}+\\f{\\varphi_{,\\mu\\nu}}{\\phi_0}-\\eta_{\\mu\\nu}\\Box_\\eta\\Big(\\f{\\varphi}{\\phi_0}\\Big)\\,, \\end{align} where $\\Box_\\eta$ is the flat-space d'Alembertian, and we neglected quadratic and higher-order terms. The tensor field equation can hence be written as \\begin{align} \\label{weak_tensor_fe} \\Box_\\eta\\theta^{\\mu\\nu}=-16\\pi\\tau^{\\mu\\nu}\\,, \\end{align} where $\\tau_{\\mu\\nu}=\\phi_0^{-1}T_{\\mu\\nu}+t_{\\mu\\nu}$. We have collected the quadratic and higher-order terms in the perturbations $\\varphi$ and $\\theta_{\\mu\\nu}$ into the gravitational stress-energy pseudotensor $t_{\\mu\\nu}$. By virtue of the gauge condition on $\\theta^{\\mu\\nu}$, we have the useful result that \\begin{align} \\tau^{\\mu\\nu}_{\\;\\;,\\nu}=0\\,. \\end{align} Following a similar procedure for the scalar field equation, we expand $\\Box_g\\phi$ in the weak-field perturbations: \\begin{align} \\Box_g\\phi= & (1+\\f{1}{2}\\theta+\\f{\\varphi}{\\phi_0})\\Box_\\eta\\f{\\varphi}{\\phi_0}-\\theta^{\\mu\\nu}\\f{\\varphi_{,\\mu\\nu}}{\\phi_0}-\\f{\\varphi_{,\\alpha}\\varphi^{,\\alpha}}{\\phi_0^2} \\nonumber \\\\ & +O(\\theta^3,\\;\\theta^2\\varphi,\\varphi^2\\theta\\dots). \\end{align} Substituting this back into the scalar field equation we find, as anticipated, the standard Klein-Gordon equation \\begin{align} \\label{weak_scalar_fe} (\\Box_\\eta-m_s^2)\\varphi=-16\\pi S\\,, \\end{align} where we have defined the source $S$ as \\begin{align} \\label{scalar_source} S\\equiv & -(6+4\\omega_{\\rm BD})^{-1}\\Big(T-2\\phi\\frac{\\partial T}{\\partial\\phi}\\Big)\\Big(1-\\frac{1}{2}\\theta-\\frac{\\varphi}{\\phi_0}\\Big) \\nonumber \\\\ & -\\frac{1}{16\\pi}\\big(\\theta^{\\mu\\nu}\\varphi_{,\\mu\\nu}+\\phi_0^{-1}\\phi_{,\\alpha}\\phi^{,\\alpha}-m_s^2\\phi_0^{-1}\\varphi^2-\\frac{1}{2}m_s^2\\theta\\varphi\\big) \\nonumber \\\\ & +O(\\theta^3,\\theta^2\\varphi,\\theta\\varphi^2,\\varphi^3)\\,. \\end{align} ",
        "conclusions": "In this paper we set constraints on massive Brans-Dicke (or Bergmann-Wagoner) theories with an action of the form (\\ref{action_general}) with $\\omega(\\phi)=\\omega_{\\rm BD}$, assuming that only a mass term is present in the expansion of $M(\\phi)$ around some cosmologically imposed value $\\phi_0$.  In particular we computed the orbital period derivative for quasicircular binaries. From an observational standpoint it will be important to generalize our work to eccentric (and possibly spinning) binaries, that could yield more stringent constraints on scalar-tensor theories. It will also be interesting to explore possible bounds on massive scalar-tensor theories that could result from Earth- and space-based gravitational-wave observations of compact binaries, along the line of Refs.~\\cite{Will:1994fb,Berti:2004bd,Berti:2005qd}.  A second obvious generalization of our work will consist in relaxing our assumptions on the form of $\\omega(\\phi)$ and $M(\\phi)$. More generic assumptions on these functions are necessary for a deeper understanding of binary dynamics in the context of modified gravity models that try to explain cosmological observations. It will be interesting to verify whether time-varying boundary conditions on the scalar field may lead to interesting binary dynamics \\cite{Horbatsch:2011ye}.  Last but not least, full numerical relativity simulations of compact binaries in scalar-tensor theories are under investigation by several groups (see e.g. \\cite{Salgado:2008xh,Healy:2011ef}). Numerical progress in evolving binary dynamics in alternative theories is important, as it could reveal strong-field effects that may be inaccessible to post-Newtonian or perturbative calculations.  \\noindent"
    },
    "1112/1112.1717_arXiv.txt": {
        "abstract": "Recent X-ray observations have enabled the study of reverberation delays in AGN for the first time. All the detections so far are in sources with a strong soft excess, and the measured delay is between the hard (1--3 keV) direct continuum and the soft excess (0.5--1 keV), interpreted as the reflection continuum smeared by relativistic effects. There is however an inherent ambiguity in identifying and studying the details of the lines in the soft excess. Here we report the first detection of reverberation in the iron K band in any AGN. Using XMM-Newton observations of NGC~4151, we find delays of order 2000 s on time-scales of $10^5$ s between the $5-6$ keV band and $2-3$ and $7-8$ keV bands, with a broad lag profile resembling a relativistically-broadened iron line. The peak of the lag spectra shifts to lower energies at higher frequencies, consistent with the red wing of the line being emitted at smaller radii, as expected from reflection off the inner accretion disk. This is a first detection of a broad iron line using timing studies. ",
        "introduction": "X-ray observations of active galactic nuclei (AGN) have shown that most of the radiation is emitted from very close to the central super-massive black hole. Broadening of X-ray spectral features caused by relativistic effects has been the main tool in studying these systems (\\citealt{2007ARA&A..45..441M} for a review). Recent work has moved beyond spectral modelling to detecting time delays between the different spectral components. Short time delays have first been seen in the Narrow Line Seyfert 1 galaxy 1H0707-495 (\\citealt{2009Natur.459..540F,2010MNRAS.401.2419Z}). Thanks to its strong soft excess and high iron abundance, small delays of a few tens of seconds are detected between the 1--4 and 0.5--1 keV bands where the soft band lags the hard band (hence soft lag). The natural explanation for the lag is reprocessing by the reflecting accretion disc in the vicinity of the black hole. Similar delays have also been seen in several other objects (\\citealt{2011MNRAS.416L..94E,2011MNRAS.417L..98D,2011MNRAS.418.2642Z}). A common feature among all these sources is a strong soft excess in the spectrum below 1 keV. This is interpreted, in line with the lags, as a relativistically-blurred reflection continuum, that responds to variations in a primary illuminating hard power-law.  The interpretation of the soft lag (soft excess lagging the hard continuum) as a light crossing effect close to the black hole was questioned by \\cite{2010MNRAS.408.1928M}, who suggested that such small delays could be an artefact of a much larger reverberating system of clouds ($\\sim$1000 instead of $\\sim$ a few gravitational radii \\rg). Although this interpretation fails when a complete set of observables (variability power, flux and lag spectra) are considered (\\citealt{2011MNRAS.412...59Z}), it stands on the ambiguity in interpreting the soft excess itself (e.g. \\citealt{2006MNRAS.365.1067C,2011arXiv1107.5429D}). The smoothness of its spectrum makes identifying the relativistically-broadened emission lines therein difficult unless the object has some extreme properties (e.g. high iron abundance and strong Fe-L in 1H0707-495). Therefore, a key challenge for any model is its ability to explain the reverberation in the iron K region (6--7 keV) where such an ambiguity is not an issue. The vital task is to find and observe an object that has reverberation signatures in this band. The importance of reverberation studies in the K band have been discussed previously, along with theoretical calculations of the expected behaviour of the broad iron line with time (\\citealt{1990Natur.344..747S,1992MNRAS.259..433M,1999ApJ...514..164R}). However, it remained difficult to detect these reverberation delays in real data.  The count rate at the bands of interest is a key parameter in reverberation studies. All previous lag detections have been in the soft band ($< 1$ keV) in objects with a strong soft excess. Those objects however, show steep spectra where the flux drops significantly in the iron-K band. Therefore, a different class of objects need to be explored if Fe-K reverberation is to be studied. In this work we show that NGC~4151, which is $\\sim300$ times brighter than 1H0707-495 at 6.4 keV, shows inter-band delays that indicate that a significant fraction of the flux in this band originates in the vicinity of the black hole.   NGC 4151 is a well-studied Seyfert galaxy ($z=0.0033$) that harbours one of the brightest AGN in the X-ray band. It has been observed and studied by all major X-ray satellites. The X-ray to $\\gamma$-ray spectrum is dominated by a Comptonisation component that breaks at $\\sim 100$ keV. There is a strong narrow iron line at 6.4 keV and a reflection hump peaking at 30 keV (\\citealt{1996MNRAS.283..193Z,2002MNRAS.334..811S}). These components are absorbed below 4 keV by a high column ($N_H\\sim10^{22} {\\rm cm}^{-2}$) of gas in the line of sight (\\citealt{1994ApJ...436L..27W}). The spectrum below 2 keV is very complex. High resolution Chandra observations revealed that it is spatially-resolved and it is associated with a highly ionised plasma coincident with the optical narrow-line region (\\citealt{2000ApJ...545L..81O,2011ApJ...729...75W}). This component does not vary much and will not be discussed here.  \\begin{figure} \\centering \\includegraphics[width=210pt,clip ]{figures/unfolded_pn.eps} \\caption[Unfolded spectra of NGC 4151]{Unfolded spectra (in $EF(E)$ form) of NGC 4151 from the XMM-Newton PN detector. The labels include: Obs. number (Obs. ID : total exposure time). } \\label{fig:unfold_spec} \\end{figure}  NGC 4151 has been observed with XMM-Newton several times in short observations constrained by its visibility. The unfolded spectra are shown in Fig. \\ref{fig:unfold_spec} for the eight longest observations ($\\sim19-62$ ks) from a total of 13 XMM-Newton. There is clearly significant spectral variability. The difference between the spectra can be as simple as a uniform apparent increase in flux, or it could be due to changes in the absorber on time scales of years (e.g. difference between observations 1 and 7 in Fig. \\ref{fig:unfold_spec}).  In this work, light curves from eight archival XMM observations are used (Obs. IDs and length of observations are shown in Fig. \\ref{fig:unfold_spec}). ODF data were reduced using \\textsc{SAS v. 11.0.0}. Light curves are produced in equal energy steps above 2 keV with 1024 s time bins. Instrumental corrections were applied using the task \\texttt{epiclccorr}. The lag was then estimated at the observed times following \\cite{1992ApJ...398..169R}. We determine the uncertainty in the lags via Monte Carlo simulations.  We shift each point in the light curves by drawing a random number from a Gaussian distribution with mean equal to the count rate and standard deviation equal to the uncertainty in the count rate.  This is performed 100 times, and the 1-sigma uncertainty in the lag is taken from the resulting distribution. The errors are very similar to those estimated in the standard way (\\citealt{1999ApJ...510..874N}). Lag-energy plots are produced by calculating the lag at the frequency of interest between the energy band and the light curve in the whole 2--10 keV band excluding the current band. This maximises the signal to noise and essentially measures the delays relative to the average (e.g. \\citealt{2011MNRAS.412...59Z}). ",
        "conclusions": "All previous lag detections in AGN have been in the soft band ($<1$ keV). In this work we reported the first detection of inter-band delays in the Fe-K band in an AGN. The lag is seen prominently in low flux observations and it is much smaller at high fluxes. The lag is in a sense that the $5-6$ keV band lags both $2-3$ and $8-9$ keV bands by $\\sim2000$ seconds (this value is frequency-dependent). The interpretation of these delays is consistent, to first order, with a picture in which an illuminating power-law component, dominating at $<\\sim3$ and $>7$ keV is reflected by an accretion disk at distances of $\\sim 10$ \\rg (The light crossing time at 1 \\rg is $5\\times10^{-5} M$ seconds, where $M$ is the black hole mass in solar units). The measured delays are due to light-travel time between the illuminating corona and a reflecting accretion disk. The main difference between low and high flux observations is a difference in the reflection fraction. Because the energy bands used do not contain a single component (direct power-law only or reflection only), the reduced reflection fraction causes the delays to be smaller in the high flux observations because they are diluted by the direct component where there is no clear lag between the energy bands.  The shape of the lag spectrum resembles a broad iron line peaking around 5 keV. Although the flux spectrum has a strong narrow line component, several arguments indicate that it is not the origin of the delays. First, the delays are small in value ($\\sim 2000$ s), which is much smaller than the distance scales of the region producing the narrow line (the molecular torus or the broad line region). Earlier work, constrained the region producing the narrow line to be on scales light years (\\cite{2000ApJ...545L..81O}). The lag spectrum peaks well below 6.4 keV, or in other words, the most lagged band is not 6.4 keV corresponding to the peak of the line as would be expected if the narrow component is producing the lag. Furthermore, the shift of the peak with frequency (Fig. \\ref{fig:gauss_cont}) rules out models where the delays are due to path-difference close to the line-of-sight (e.g. \\citealt{2010MNRAS.408.1928M}). In those models, no frequency-dependent behaviour is expected, as the shape of the distant reflection spectrum is not expected to change on these time-scales.  The light-crossing time-scale for NGC~4151 with a mass of $4\\times10^7$ \\Mo at say 6 \\rg is $\\sim1200$ s, which is of the same magnitude as those measured here. The thermal time-scale at this radius for a standard thin disk (\\citealt{1973A&A....24..337S}), which could be the variability time-scale if ionisation/density changes are involved, is $\\sim2\\times10^5$ seconds assuming a viscosity parameter $\\alpha$ of 0.1. Again, this time-scale is also consistent with the observations. Spectral fitting with a relativistic model gives an inner disk radius of $1.4\\pm0.1$ \\rg with an emissivity index of $4\\pm0.3$, a profile flatter than usually observed. If the lag spectra are fitted with a relativistic line model using the inclination and emissivity from the flux spectra, we find that the low and high frequency variability originate at $\\sim 10$ and $\\sim4$ \\rg respectively. This also presumably indicate that the height of the illuminating source is not very large (a few \\rg ).  The lag spectrum of NGC~4151 consists of broad line-like feature that peaks between 4--6 keV. We have found that this peak shifts closer to 6.4 keV (i.e. higher energies) when lower frequencies are probed. This is consistent with the redder parts of the line emitted at smaller radii, and so they vary on smaller time-scales, and their light crossing time is smaller. Building an exact model however, requires a full treatment of the problem, including general relativistic effects and the interplay between the direct and reflected components in the observed data. To understand the lags, we plan to explore these details in the near future. This important result highlights what can be achieved with current telescopes in the brightest AGN NGC~4151. Larger collecting areas in future X-ray missions (e.g. LOFT and ATHENA), will make it possible to perform detailed studies in frequency, lag and energy parameter spaces in this and many other objects."
    },
    "1112/1112.1703_arXiv.txt": {
        "abstract": "We carry out several isolated galaxy evolution simulations in a fixed dark matter halo gravitational potential using the new version of our N-body/Smoothed Particle Hydrodynamics (SPH) code {\\tt GCD+}. The new code allows us to more accurately model and follow the evolution of the gas and stellar components of the system including powerful supernovae feedback and its effects on the inter-stellar medium  (ISM). Here we present the results of six simulations of an M33-sized late-type disc galaxy each with varying values for our model parameters which include the star formation efficiency $(C_{\\rm *})$, the energy released per supernovae explosion ($E_{\\rm SN}$) and the energy released per unit time from stellar winds ($E_{\\rm SW}$). We carry out both a pixel-by-pixel and radial ring analysis method for each of our galaxies comparing our results to the observed Schmidt-Kennicutt Law and vertical gas velocity dispersion versus radius relation amongst others. We find that our models with higher feedback more closely resemble the observations and that feedback plays a pivotal role in obtaining {\\it both} the observed Schmidt-Kennicutt and gas velocity dispersion relations. ",
        "introduction": "\\label{intro-sec}  The process of galaxy formation and evolution remains one of the most outstanding problems in modern astrophysics. From a cosmological point of view, great progress has been made including the development of a theory; $\\Lambda$CDM cosmology which accurately explains the present day large scale distribution of galaxies from the primordial small scale density fluctuations in the early universe. The actual processes which govern individual galaxy formation and evolution are far less well understood however. One of these important processes is that of star formation itself. \\citet{sch59} suggested that the star formation rate (SFR) density and the interstellar gas density in our Galaxy were related by a simple power law relation of the form $\\rho _{\\rm SFR}\\propto (\\rho _{\\rm gas})^{n}$, where $n$ is the power law index. \\citet{sch59} originally suggested a value for $n$ of 2, so that the rate of star formation varies with the square of the gas density. Studying the form of the observed star formation law, \\citet{rk89,ken98} analysed the SFR {\\it surface} density and the gas {\\it surface} density for a much larger sample of galaxies and over a wider range of gas densities, determining the  power law index $n$ to be equal to 1.4 \\citep[see also][]{ken98AR}. We now refer to the empirical relation between SFR and gas surface density as the Schmidt-Kennicutt (SK) relation. Recently, the SK relation has been studied in much more depth \\citep[e.g.][]{blw08,lwb08} including observing the relation for individual components of the gas (HI and molecular) and as a function of environment. Several observational studies found evidence for a two-slope SK relation \\citep[e.g.][]{blw08}. \\citet{blw08} showed that a break exists at gas surface densities of 9 M$_{\\odot}\\textup{yr}^{-1}$, below which the slope of the SFR steepens. This break corresponds to the transition between the HI and molecular phase of the gas. Theoretical models also predict such breaks in the SK relation. For example \\citet{kmt09} made a theoretical model consisting of three prescriptions: molecular self-shielding, internal feedback and turbulent regulation in molecular clouds which address the physics responsible for determining the SFR on local scales and successfully managed to reproduce the observed trends in the SK relation.  The SK relation is of great importance as it is used for the calibration of numerical codes of galaxy formation and evolution, with the SFR often explicitly required to obey a SK type law \\citep[e.g.][]{sh03,bto07,sd08}. Some more recent studies however, do not force the SFR to follow a SK-like relation, and so its adherence to this law is a prediction of the model \\citep[e.g.][]{rk08,sdk08,mmb11}.  Recently \\citet{tam09} showed that the gas velocity dispersion is also an important factor to consider for any galaxy evolution model. Using HI line width maps obtained  for 11 disc galaxies of The HI Nearby Galaxy Survey \\citep[THINGS,][]{wbb08} they determined that all of their sample galaxies showed a systematic radial decline in their HI velocity dispersion, with a characteristic value of $\\sim$ 10 km s$^{-1}$ by $r_{25}$ $-$ roughly the radial extent of the star forming disc. The radial decline in velocity dispersion was well correlated with the radial decline of SFR and therefore they suggested that high gas velocity dispersion is induced by supernovae feedback.  The implications of such studies are that the gas velocity dispersion should now be used as an additional constraint for any new numerical simulation codes of galaxy evolution. Recently, \\citet{dbp11} performed isolated disc galaxy evolution simulations to investigate how the properties of the ISM, molecular clouds and SFR depend on the level of stellar feedback, comparing their results against both the observed SK relation and gas velocity dispersion. They found that supernovae were vital in reproducing the observed velocity dispersions and scale heights of the individual components of the ISM. Recently, \\citet{pgc11} analysed both the SK and HI velocity dispersion relations for a fully consistent cosmologically simulated late-type dwarf spiral galaxy.  In this paper we initiate a project to create an accurate and self-consistent numerical model for late-type spiral galaxies. To this end we have been developing our numerical simulation code {\\tt GCD+} \\citep{kg03a}. The new version of the code includes new features and improvements to a number of aspects of the modelling of galaxy formation and evolution, including new star formation and feedback recipes. The overall effect of these changes is that the code can now capture and accurately follow the effects of powerful supernovae explosions on the evolution of the galaxy. The details of the new model are described in the next section. In this study, we set up an isolated disc galaxy consisting of gas and stellar discs with no bulge component, in a static dark matter halo gravitational potential. To calibrate our new model, we here focus on M33{\\it-like} bulgeless galaxies. M33 is a suitably chosen late-type bulgeless spiral galaxy. Its close proximity in the Local Group means it has been a favourite choice for detailed observational campaigns \\citep[e.g.][]{hcs04,vcg10}. Its smaller size compared to the Milky Way also means that we can simulate it to much higher resolution using the same number of particles. M33 also lacks a supermassive black hole at its centre \\citep{mfj01}, a factor which helps simplify the modelling and analysis. Therefore, M33{\\it-like} galaxies are ideal target galaxies for testing and calibrating the numerical modelling of the physical processes in galaxy evolution. Note that our aim is $\\it{not}$ to reproduce completely the observed properties of M33. It is likely that M33 had a recent interaction with M31 \\citep{mii09,ppm09} and so it may not be a typical galaxy. We focus more on the {\\it general} observed properties for M33{\\it-sized} disc galaxies, such as the SK relation and the gas velocity dispersion.  The paper is organised as follows. In Section 2, we describe the galaxy model together with  some new features of our numerical simulation code. In Section 3, we present our results and analyses. We summarise our findings and present our conclusions and planned future work in Section 4.   \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Simulation parameters} \\renewcommand{\\footnoterule}{} \\begin{tabular}{@{}llllllllll@{}} \\hline $C_{\\rm NFW}$ & $M_{\\rm vir}$ & $r_{\\rm vir}$ & $M_{\\rm Star}$ & $M_{\\rm Gas}$ % & $R_{\\rm d,star}$ & $R_{\\rm d,gas}$ & $z_{\\rm d,star}$ \\\\ & (M$_{\\odot})$ & (kpc) & (M$_{\\odot})$ & (M$_{\\odot})$ & (kpc) & (kpc) & & \\\\ \\hline 10 & $5.0\\times10^{11}$ & 163 & $4.0\\times10^{9}$ & $2.0\\times10^{9}$ & 1.4 & 4.0 & 0.3  \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} ",
        "conclusions": "\\label{conc}  We aim to build the most accurate and self-consistent numerical model for late-type spiral galaxies to date. To this end we have been developing our numerical simulation code {\\tt GCD+}. We include new star formation and feedback recipes which allow our simulations to follow the effects of powerful SNe explosions on galaxy evolution.  In this study we set up an isolated M33-sized disc galaxy consisting of gas and stellar discs with no bulge component, in a static dark matter halo gravitational potential. We carried out six main simulation runs in which we varied three of our model parameters; the star formation efficiency $(C_{\\rm *})$, the thermal energy released per SNe explosion ($E_{\\rm SN}$) and the thermal energy released per unit time from stellar winds ($E_{\\rm SW}$). We studied the effect of varying these parameters on the Schmidt-Kennicutt Law and vertical gas velocity dispersion versus radius relation comparing to observations from \\citet{blw08} and \\citet{tam09} respectively. We find that our models with higher energy feedback (Runs 3$-$6)  more closely resemble the observations. Powerful supernova explosions cause the formation of large holes or bubbles in the gas distribution in the galaxy and stir the ISM turbulence and are thus essential for producing the observed gas velocity dispersions. The gas velocity dispersion is also affected (albeit less strongly) by the star formation efficiency. Our model with high energy feedback from SNe and stellar winds and intermediate star formation efficiency (Run 6) most closely resembled the observations. This run represented our `best compromised model' since we had fine-tuned the star formation efficiency and feedback parameters. In our pixel-by-pixel analysis of the star formation law, we found that our model with high star formation efficiency and low energy feedback (Run 2) produced a KS relation with an unrealistically steep slope. Models with low energy feedback had a distinct lack of pixels with both low $\\Sigma_{\\rm H}$ and $\\Sigma_{\\rm SFR}$, resulting in a failure to accurately reproduce the observed star formation law. By increasing the energy feedback parameters, $E_{\\rm SN}$ and $E_{\\rm SW}$ (as in Run 4) we managed to rectify this problem and found that the pixels which had low values for both $\\Sigma_{\\rm H}$ and $\\Sigma_{\\rm SFR}$ were predominantly located inside bubble regions, which were more numerous and larger in size for the higher feedback runs. Our results therefore suggest that SNe can play a pivotal role in shaping the star formation law. We used both a pixel-by-pixel and radial ring method to analyse the star formation law in our galaxies and found that overall, both methods gave broadly consistent results.  In this paper we employed static potentials for the dark matter halo and stellar disc. The main benefits of not using a live halo include the prevention of numerical effects such as artificial disc heating due to the scattering of baryonic particles by more massive dark matter particles which may be an issue in lower resolution cosmological simulations. Our numerical simulations thus do not suffer from these global disc instabilities. However, we underestimate the perturbation from accreting satellites \\citep[e.g.][]{pbt11} and globular clusters \\citep[e.g.][]{vcf09}.  In this paper we demonstrate that reproducing both the SK relation and the observed gas velocity dispersions puts stronger constraints on the sub-grid modelling for galaxy evolution simulations. Encouraged by the success of this work, we are currently in the process of running higher resolution simulations. This should result in us being able to probe deeper into the origins of the $\\Sigma_{\\rm GAS}$$-$$\\Sigma_{\\rm SFR}$ relation. It should also mean that our models will be less sensitive to the employed value of the star formation efficiency, given we could resolve molecular cloud formation processes \\citep{sdk08}. After calibrating our code at this higher resolution, we then aim to run new extremely high resolution cosmological simulations. The fully self-consistent nature of the new higher resolution cosmological simulation will be a powerful tool  to unravel the mysteries behind  the formation and evolution of late type spiral galaxies like M33."
    },
    "1112/1112.4583_arXiv.txt": {
        "abstract": "In December 2009, the bright blazar, 3C~454.3 exhibited a strong outburst in the optical, X-ray and gamma-ray regions. We performed photometric and polarimetric monitoring of this outburst in the optical and near-infrared bands with TRISPEC and HOWPol attached to the Kanata telescope. We also observed this outburst in the infrared band with {\\it AKARI}, and the radio band with the 32-m radio telescope of Yamaguchi University. The object was in an active state from JD~2455055 to 2455159. It was 1.3~mag brighter than its quiescent state before JD~2455055 in the optical band. After the end of the active state in JD~2455159, a prominent outburst was observed in all wavelengths. The outburst continued for two months. Our optical and near-infrared polarimetric observations revealed that the position angle of the polarization (PA) apparently rotated clockwise by 240~degrees within 11~d in the active state (JD~2455063---2455074), and after this rotation, PA remained almost constant during our monitoring. In the outburst state, PA smoothly rotated counterclockwise by 350~degrees within 35~d (JD~2455157---2455192). Thus, we detected two distinct rotation events of polarization vector in opposite directions. We discuss these two events compared with the past rotation events observed in 2005, 2007 and 2008. ",
        "introduction": "Blazar is a class of active galactic nuclei, whose relativistic jets are considered to be directed along the line of sight (e.g. \\cite{Blandford79}). Radiation of blazars has three main properties. First, the blazars emit electromagnetic radiation in a broad range from the radio to gamma-ray bands. Their emission consists of two major components (e.g. \\cite{Kubo98}). The low energy component is synchrotron radiation observed from the radio to the optical, sometimes extending to the X-ray bands. The high energy component is inverse-Compton scattering from the X-ray to the gamma-ray bands. Second, blazars exhibit rapid and violent variability in all wavelength bands \\citep{Antonucci93}. The variability has various timescales from less than a day (e.g. \\cite{Aharonian07}; \\cite{Albert07}; \\cite{Sasada08}; \\cite{Raiteri08a}) to longer than years (e.g. \\cite{Sillanpaa96}). Third, blazars possess relativistic jets (e.g. \\cite{Lister09}), which are responsible for high polarization observed at optical/near-infrared (NIR), and radio wavelengths. Since the polarization can be a probe of the magnetic field in the jet, polarimetric observations are important to study the structure of the jet.  The temporal variation of the polarization vector is complex in general. \\citet{Jones85} reported that the polarization behavior was erratic in blazars. On the other hand, several papers reported that the polarization vector exhibited systematic variations, for example, positive correlations between the flux and the degree of polarization (PD) (e.g. \\cite{Smith86}). Recently, \\citet{Marscher08} have reported that the polarization vector in BL~Lac smoothly rotated when the object was bright. From this result, they proposed that this rotation indicated an emission zone passing through a helical magnetic field in the jet. \\citet{Abdo10a} reported a rotation of polarization in 3C~279, and proposed that the rotation event is attributed to a bent jet. However, there are only a few rotation events which have been reported to date.  3C~454.3 is one of the most famous blazars. The object is classified as Flat Spectrum Radio Quasars (FSRQs), and its redshift is $z=0.859$ \\citep{Jackson91}. Although the object had been quiet in the optical band until 2001, the object has kept showing the active behavior since then \\citep{Villata06}. In 2005, the object showed an exceptional outburst. In this outburst, the object brightened from the radio to the gamma-ray bands (\\cite{Fuhrmann06}; \\cite{Pian06}; \\cite{Giommi06} ; \\cite{Villata07}). After this outburst, similar outbursts were detected in 2007 and 2008. In 2005 and 2007 outbursts, rotation events of the optical polarization vector were reported by \\citet{Jorstad10} and \\citet{Sasada10}. In Dec. 2009, a prominent outburst was reported for this object in the gamma-ray band by {\\it Fermi}/LAT and AGILE (\\cite{Striani09a}; \\cite{Striani09b}; \\cite{Escande09}; \\cite{Striani10}; \\cite{Pacciani10}; \\cite{Ackermann10}), in the X-ray band by {\\it INTEGRAL}/IBIS \\citep{Vercellone09}, by {\\it Swift}/XRT \\citep{Sakamoto09}, by {\\it Swift}/BAT \\citep{Krimm09}, in the optical bands (\\cite{Villata09a}; \\cite{Bonning09}; \\cite{Sasada09}). The object was the brightest source in the GeV-gamma-ray sky for over a week \\citep{Ackermann10}. The object showed flux variability over timescales less than three hours and very mild spectral variability with an indication of gradual hardening preceding major flares. The minimum Doppler factor was 13, estimated by using these results. \\citet{Bonnoli11} also estimated the Doppler factors $\\sim$~25 during the outburst by constructing a multi-wavelength spectral model.  We performed monitor observations of 3C~454.3 from May 2009 to February 2010 in a multi-color photometric and polarimetric mode using the Kanata telescope. We also observed in the radio and infrared (IR) bands. In this paper, we report on the behavior of the 2009 outburst in these bands, and the detections of two rotation events in the polarization vector. Directions of these two rotations were different, suggesting a complex magnetic field in the jet. This paper is arranged as follows: In section~2, we present the observation method and analysis in the radio, infrared, optical and X-ray bands. In section~3, first we report the light curves and the spectral indexes of the X-ray and optical bands. Then, we report the temporal behavior of polarization in the optical and NIR bands. After that, we report the spectral energy distribution from the radio to optical band. In section~4, we discuss two rotation events which we detected, by comparing them with the past rotation events observed in 2005, 2007 and 2008. The conclusion is drawn in section~5. ",
        "conclusions": "We performed the photometric and polarimetric monitoring of 3C~454.3 during 2009---2010. The object showed the active and outburst states in the optical and near-infrared light curves and two rotation events in the temporal variations of polarization. In the past and our polarimetric monitoring, five rotation events were reported in 3C~454.3. The rotations were observed in both clockwise and counterclockwise directions in the $QU$ plane. The complex model of the structure of the magnetic field in the jet is needed to explain these rotation events. \\\\ \\\\ This work was partly supported by a Grand-in-Aid from the Ministry of Education, Culture, Sports, Science, and Technology of Japan (22540252,). This research has made use of data from the University of Michigan Radio Astronomy Observatory which has been supported by the University of Michigan and by a series of grants from the National Science Foundation, most recently AST-0607523. The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica. M.~Sasada and M.~Yamanaka have been supported by the JSPS Research Fellowship for Young Scientists."
    },
    "1112/1112.4851_arXiv.txt": {
        "abstract": "The Palomar Transient Factory (PTF) provides multiple epoch imaging for a large fraction of the celestial sphere. Here we describe the photometric calibration of the PTF data products that relates the PTF magnitudes to other magnitude systems. The calibration process utilizes Sloan Digital Sky Survey (SDSS) $r\\sim16$\\,mag point source objects as photometric standards. During photometric conditions, this allows us to solve for the extinction coefficients and color terms, and to estimate the camera illumination correction. This also enables the calibration of fields that are outside the SDSS footprint. We test the precision and repeatability of the PTF photometric calibration. Given that PTF is observing in a single filter each night, we define a PTF calibrated magnitude system for $R$-band and $g$-band. We show that, in this system, $\\approx 59\\%$ ($47\\%$) of the photometrically calibrated PTF $R$-band ($g$-band) data achieve a photometric precision of 0.02--0.04\\,mag, and have color terms and extinction coefficients that are close to their average values. Given the objects' color, the PTF magnitude system can be converted to other systems. Moreover, a night-by-night comparison of the calibrated magnitudes of individual stars observed on multiple nights shows that they are consistent to a level of $\\approx0.02$\\,mag. Most of the data that were taken under non-photometric conditions can be calibrated relative to other epochs of the same sky footprint obtained during photometric conditions. We provide a concise guide describing the use of the PTF photometric calibration data products, as well as the transformations between the PTF magnitude system and the SDSS and Johnson-Cousins systems. ",
        "introduction": "\\label{Introduction}  The Palomar Transient Factory\\footnote{http://www.astro.caltech.edu/ptf/} (PTF; Law et al. 2009; Rau et al. 2009) is a synoptic survey designed to explore the transient sky. The project utilizes the $48''$ Samuel Oschin Schmidt Telescope on Mount Palomar. The telescope has a digital camera equipped with 11 active CCDs, each 2K$\\times$4K pixels (Rahmer et al. 2008). Each PTF image covers 7.26\\,deg$^{2}$ with a pixel scale of $1.01''$\\,pix$^{-1}$. From the beginning of the PTF survey in March 2009 until January 2011, most of the images were taken with the Mould $R$ filter. Starting January 2011, we performed the PTF main survey in $g$-band during dark time and in $R$-band during bright time. In addition, a few nights around times of full Moon are used for taking images with narrow-band  H$\\alpha$ filters. A PTF system overview and review of the first year's performance are given in Law et al. (2010).   Accurate photometric calibration is a non-trivial task since one must know both the atmospheric transmission (e.g., Padmanabhan et al. 2008; Burke et al. 2010) and the optical/detector system response as a function of wavelength, time, and sky position. Moreover, observations of flux standards are required, and the true spectral energy distribution of these standards needs to be known. In practice simplifying assumptions are made in order to achieve solutions which are good to the few percent level. For example, in many cases, it is customary to use the calibrated magnitudes of the standard stars instead of their spectral energy distribution, and to assume that the atmospheric transmission is a smooth function of airmass. We note that relative-photometry approaches (e.g., Gilliland \\& Brown 1988; Gilliland et al. 1991; Honeycutt 1992; Everett \\& Howell 2001; Ofek et al. 2011a) can achieve better accuracy from the ground, but are limited by scintillation noise (e.g., Young 1967; Gilliland \\& Brown 1988), Poisson noise and flat-fielding errors.   There are several approaches to photometric calibration. For example, the Sloan Digital Sky Survey (SDSS; York et al. 2000) images are calibrated using an auxiliary 20-inch telescope, which determines the photometric condition on a nightly basis (Hogg et al. 2001). This telescope measures the extinction and photometric zero point using a network of standard stars (Smith et al. 2002), which, in turn, are tied to the standard star BD$+17$:$4708$. Furthermore, the photometric uniformity in the SDSS is achieved by the algorithm described in Padmanabhan et al. (2008). The SDSS photometry is uniform to better than 2\\% in all bands (Adelman-McCarthy et al. 2008).   Ofek (2008; see also Pickles \\& Depagne 2010) has suggested using Tycho-2 (H{\\o}g et al. 2000) stars to photometrically calibrate astronomical images. The SDSS $griz$ magnitudes of such stars are based on synthetic magnitudes of stellar spectral templates fitted with the Hipparchos $B_{T}$, $V_{T}$ magnitudes and the 2MASS (Skrutskie et al. 2006) $JHK$ magnitudes. However, this strategy requires images containing unsaturated Tycho calibration stars which lie in the $\\approx9$--$12$ magnitude range. The saturation limit for the PTF $g$- and $R$-band survey is typically\\footnote{The PTF camera electronics were modified on 2009 Oct 22 to increase the dynamic range. Before this date, the saturation limit was around 15th magnitude.} around 14 magnitude.  Herein, we describe the method we have developed for the photometric calibration of PTF data taken in the $g$ and $R$ bands. The photometric calibration of the PTF H$\\alpha$ survey will be described elsewhere. We note that the relative photometric calibration of PTF currently\\footnote{For relative photometry, we use a scheme similar to that proposed by Honeycutt (1992) with some modifications outlined in Ofek et al. (2011a).} achieves precision as good as 3\\,mmag, in given fields, at magnitude 15 (e.g., van Eyken et al. 2011; Ag\\\"{u}eros et al. 2011; Law et al. 2011; Levitan et al. 2011; Polishook et al. 2011). The PTF relative photometry pipeline will be described in Levitan et al. (in prep.).  Our method of photometric calibration for PTF data is similar to the ``classical'' method of observing a small number of calibration stars (e.g., Landolt 1992) through various airmasses and assuming photometric conditions -- i.e., the atmosphere transmission properties are constant in time and are a continuous function of airmass. The main difference is that we are using SDSS stars as standard stars and we typically observe $\\sim10^{5}$ SDSS stars with high S/N in each CCD per night. Another important difference is that PTF observations are done in a single filter each night. Therefore, in order to relate the PTF calibrated magnitudes to one of the common absolute systems one need to apply color terms. The paper is organized as follows. In \\S\\ref{Process}, we describe the PTF photometric-calibration method. In \\S\\ref{Illum}, we discuss the illumination correction, while \\S\\ref{fog} describes a related problem that plagues early PTF data. In \\S\\ref{Data} we describe the data products and what calibration information is stored with the image-product data. The performance of the PTF magnitude calibration is given in \\S\\ref{Perf}, and the derived photometric parameter statistics are discussed in \\S\\ref{Pars}. We provide the color transformation between the PTF magnitude system and other systems in \\S\\ref{Color}. Finally, we conclude the paper in \\S\\ref{Sum}. Unless specified otherwise, the statistics given here are based on all PTF data obtained from March 2009 to July 2011. ",
        "conclusions": "\\label{Sum}  A wide variety of scientific applications require photometric calibration of astronomical images. In this paper, we provide a description of the current PTF photometric calibration process and products. We show that roughly half of the PTF data were taken under conditions that allow calibration of the images to an accuracy of 0.02--0.04\\,mag relative to SDSS. Most importantly, the rest of the data, taken under non-photometric conditions, can be calibrated using data taken in photometric conditions. At the bright end ($\\approx15$\\,mag), the PTF photometric calibration is stable to a level of $\\approx0.02$\\,mag. Future improvements to the PTF photometric calibration process that are important for the PTF type-Ia and IIp supernovae cosmology projects will be described elsewhere.  The PTF photometric calibration allows us to measure the calibrated magnitude of transients outside the SDSS footprint. We have used the PTF dataset to compile a photometric catalog. The first version of this catalog is described in Ofek et al. (in prep.). We note that our tests show that PTF data can deliver relative photometry, in single images, with precision as good as 3\\,mmag (e.g., van Eyken et al. 2011; Ag\\\"{u}eros et al. 2011). This capability is being built into the PTF relative photometry pipeline and will be described in Levitan et al. (in prep.)."
    },
    "1112/1112.0019_arXiv.txt": {
        "abstract": "Spiral density wave theory attempts to describe the spiral pattern in spiral galaxies in terms of a long-lived wave structure with a constant pattern speed in order to avoid the winding dilemma. The pattern is consequently a rigidly rotating, long-lived feature. We run N-body simulations of a giant disc galaxy consisting of a pure stellar disc and a static dark matter halo, and find that the spiral arms are transient features whose pattern speeds decrease with radius, in such a way that the pattern speed is almost equal to the rotation curve of the galaxy. We trace particle motion around the spiral arms. We show that particles from behind and in front of the spiral arm are drawn towards and join the arm. Particles move along the arm in the radial direction and we find a clear trend that they migrate toward the outer (inner) radii on the trailing (leading) side of the arm. Our simulations demonstrate that because the spiral arm feature is co-rotating, the particles continue to be accelerated (decelerated) by the spiral arm for long periods, which leads to strong and efficient migration, at all radii in the disc. ",
        "introduction": "Spiral density wave theory, first postulated by \\citet{L60} and carried forward by \\citet{LS64}, has been the most widely accepted theory of the spiral arm for almost 50 years. This theory considers the spiral arm to be the crest of a stellar density wave that rotates around the centre of the Galaxy at an angular pattern speed, $\\Omega _p$, that does not vary with galacto-centric radius. Spiral density wave theory naturally explains the long-lived nature of the spiral arm. However, short-lived spiral arms were proposed by \\cite{GLB65}, \\cite{JT66} and \\cite{T81}. These studies argue that spiral structure can grow from local instabilities and shearing rotation predicted from a flat rotation curve of disc stars that help to develop the spiral features \\citep{A84}. Although these spiral arms are short-lived features that continuously disappear, they are regenerated by local growth of instability. In these transient spiral theories, the spiral structure itself can be self-excited in small over-dense regions as a consequence of shot noise \\citep{T90}. If the over-density increases above a few percent, the growth becomes very non-linear and results in rapid growth of the spiral arm density (e.g. \\citealt{T81}; \\citealt{DT94}; \\citealt{B03}; \\citealt{Se10}). Non-linear growth of transient rigid waves is also demonstrated in the context of a shearing sheet \\citep{F05}. While shearing is important for non-linear growth of spiral arms, it is also reported to be the main factor in the appearance of spurs and spiral arm bifurcation \\citep{DB08}.  More recently, the development of high-resolution numerical N-body/SPH simulations has played an important role in the study of spiral structure formation and evolution. Many numerical simulations including self-gravity and with no addition or forcing of spiral potential show short-lived transient spiral arms, which come and go throughout the evolution such that the spiral pattern of the disc is never removed (e.g. \\citealt{SC84};  \\citealt{BAM09}; \\citealt{QDBM10}; \\citealt{Fu11}; \\citealt{WBS11}). \\citet{SC84} note that the spiral pattern in N-body simulations generally fades over time because the spiral arm structure heats the disc kinematically and causes Toomre's Q parameter to rise, hence the disc becomes too stable against the development of non-axisymmetric structure. They suggest that continuous addition of a kinematically cold population of stars is crucial to maintain the spiral arm pattern. \\citet{CF85} demonstrate this with gas clouds in circular orbit that are allowed to become future sites of star formation. Recently \\citet{Fu11} suggest that the rapid disappearance of the pattern probably results from a low number of particles in these simulations, which leads to higher Poisson noise and more rapid growth of strong spirals, which then heat the disk and erase the spiral pattern. They demonstrated that the spiral pattern can be sustained without dissipation or continuous addition of cold components if a sufficiently high number of particles ($>10^6$) is used to prevent a large Poisson noise and therefore fast disc heating. Still, there are no numerical simulations that reproduce long-lived spiral arms like those predicted by spiral density wave theory \\citep{Se11}, and this is the case for both spiral and barred-spiral galaxies (e.g. \\citealt{BAM09}; \\citealt{QDBM10}; \\citealt{WBS11}).  In spiral density wave theory, the angular pattern speed of the spiral arm is constant as a function of radius. The flat rotation curve of spiral galaxies implies the decreasing angular velocity of stars. Therefore, there is one radius where the stars rotate at the same velocity as the arm. This is called the co-rotation radius, inside which stars rotate faster than the spiral arm and outside which they rotate more slowly. The co-rotation radius marks a point of resonance, where stars very close to this radius have been shown to undergo large changes in their angular momentum and migrate radially to another part of the disc. This was first proposed by \\citet{LBK72}, who showed that radial migration around co-rotation can occur without increasing the random motions of stars. \\citet{SB02} extended this theory, and showed that transient arms provide a good mechanism for mixing, where each transient arm possesses a unique co-rotation radius around which nearby stars can migrate. The main conclusions of this paper were that only around the radius of co-rotation is there large angular momentum transport of individual stars without heating. There have been many studies on the effect of co-rotation resonances on structure, radial heating and mixing in spiral galaxies (e.g. \\citealt{CS85}; \\citealt{AL97}; \\citealt{SB02}; \\citealt{MQ06}; \\citealt{R08}a; \\citealt{RDS08}b; \\citealt{SBl09}; \\citealt{MFC11}) and in barred-spiral galaxies where a bar resonance is also included (e.g. \\citealt{MF10}; \\citealt{MA11}).    Very recently, there have been observational studies with evidence of a pattern speed that varies with radius, which is different from the constant pattern speed of spiral density wave theory. The evidence comes from a range of techniques and sources that include observations of no offset between different star formation tracers in nearby galaxies \\citep{FR11}, and radially decreasing spiral arm pattern speeds deduced by the Tremaine-Weinberg method (e.g. \\citealt{MRM06}; \\citealt{MRM08}; \\citealt{MRM09}; \\citealt{SW11}), which inevitably leads to short-lived arms as a consequence. Numerical simulations of an interacting galaxy in \\citet{DTP10} also shows a pattern speed that decreases as a function of radius. During the preparation of this work, \\citet{WBS11} performed simulations of an isolated galaxy, and found a pattern speed that traces very well the rotation curve, which indicates that there is co-rotation of the spiral pattern with the stars at all radii.  In this study, we analyse the pattern speed of the spiral arm feature and the dynamics of particles around the arm using our N-body simulations of a pure stellar disc. We determine the nature of the spiral arm pattern speed, and monitor how stellar motion is affected by the arm. We also find that the spiral arms in the simulations are transient and that the pattern speed is always similar to the circular velocity. We offer a qualitative insight into how phenomena such as radial migration might occur with the transient co-rotating spiral arms, and how the co-rotating spiral arm may evolve.  In Section 2 we explain how we set up the model and the initial parameters that we choose. In Section 3 we present the results of our analysis, compare them with previous studies and discuss their implications. In Section 4 we summarise the significance of the results and remark upon the value of the simulations and future work. ",
        "conclusions": "We have performed N-body simulations of a pure stellar disc, and then performed a dynamical analysis of the spiral arms and particles around the spiral arms and traced their evolution. We come to the following conclusions:  \\begin{enumerate} \\item{} We find in our simulations that spiral arms are transient recurring features: we observe the continuous disappearance of spiral arms and the reappearance of new ones. This transient nature has always been found in numerical simulations, and is not consistent with spiral density wave theory. \\item{} We performed two analyses on the pattern speed of the spiral arms: the stellar density trace method; and spectrogram analysis. The latter shows a possible configuration of two or three wave modes  that could mean that wave modes interfere constructively and destructively with each other.  Regardless of whether this is the case, the stellar density trace method gave us clear evidence of an apparent spiral arm pattern speed that co-rotates with the stars at a large range of radii, which is reflected in points iii), iv) and v) below. This is consistent with what is found in \\citet{WBS11}, and contrary to classic spiral density wave theory, which allows a constant pattern speed as a function of radius, and hence there is \\emph{only one} co-rotation radius. Although we cannot reject that multiple modes are creating this apparent co-rotation, there is still the question of why the apparent pattern speed is always similar to co-rotation. \\item{} Particles are shown to join the spiral arm from both sides at all radii. This is further evidence for the co-rotating nature of the spiral arm feature, because the arm must move at a similar speed to the particles in order for them to join the arm from both sides at all radii. \\item{} Particles migrate radially along the spiral arm at all radii. Stars behind the arm are accelerated by the arm and slide along the arm to larger radii. Stars in front of the arm are dragged back by the potential, and slide down the arm to smaller radii. \\item{} Migrating particles do not actually cross the spiral arm. The co-rotating nature means that the particles stay on their respective sides of the spiral arm, so they are accelerated (decelerated) until the spiral arm is disrupted. This means that radial migration is more efficient than that discussed by \\citet{SB02}, with little contribution to disc heating. \\end{enumerate}  Conclusions (i) and (ii) are already known from previous studies, and we include them to show the consistency of this study. Conclusions (iii), (iv) and (v) are new results which to our knowledge have not been reported before.  A possible limitation of this study is that it focuses only on the stellar component. However, our preliminary work with simulations including gas, star formation and supernova feedback also show a consistent picture of transient co-rotating spiral arms and strong radial migration, which will be presented in a forthcoming paper.  In this study, we have not addressed the mechanism of formation and destruction of the spiral arm structure. In this simulation, spiral arms appear to grow from small shot (Poisson) noise in the (initial) density distribution and increase in prominence from there, possibly through mechanisms similar to swing amplification.  Although co-rotating spiral arms are expected to be short-lived due to winding, we observe that the spiral arms are disrupted before they are wound up very tightly. We also observe kinks i.e. change in pitch angle, in the spiral arms. The mechanism that controls the transient shape of the spiral arms and drives the disruption is unclear, and requires further work.  In this study, we focused on the evolution between the eras of formation and destruction, and found that the spiral arms in our simulations are not consistent with the long-lived, rigidly rotating spiral arms of spiral density wave theory. The transient co-rotating spiral arm found in N-body simulations should be tested by observations of both our Galaxy \\citep{Se11} and external galaxies (e.g. \\citealt{MRM09}; \\citealt{FR11}; \\citealt{SW11}). Although this paper focuses on the non-barred spiral galaxies, it is worthwhile to explore barred-spiral galaxies where the spiral arms are affected by the bar (e.g. \\citealt{SS87}), and thus could be longer lived (e.g. \\citealt{DT94} \\citealt{BT08}; \\citealt{BAM09}, \\citealt{QDBM10}). We will study further the spiral arms in our simulations and endeavour to find observational consequences of the co-rotating spiral arm."
    },
    "1112/1112.0533_arXiv.txt": {
        "abstract": "The galactic neighborhood, extending from the Milky Way to redshifts of about 0.1, is our unique local laboratory for detailed study of galaxies and their interplay with the environment. Such study provides a foundation of knowledge for interpreting observations of more distant galaxies and their environment. The Astro 2010 Science Frontier Galactic Neighborhood Panel identified four key scientific questions: 1) What are the flows of matter and energy in the circumgalactic medium? 2) What controls the mass-energy-chemical cycles within galaxies? 3) What is the fossil record of galaxy assembly from first stars to present? 4) What are the connections between dark and luminous matter? These questions, essential to the understanding of galaxies as interconnected complexes, can be addressed most effectively and/or uniquely in the galactic neighborhood. The panel also highlighted the discovery potential of time-domain astronomy and astrometry with powerful new techniques and facilities to greatly advance our understanding of the precise connections among stars, galaxies, and newly discovered transient events. The relevant needs for laboratory astrophysics will be emphasized, especially in the context of supporting NASA missions. ",
        "introduction": "The Galactic Neighborhood (GAN) Panel was charged to identify key scientific questions (as well as a potential major discovery area), which could be effectively addressed in the upcoming decade, regarding galaxies and their surroundings out to redshifts $z \\approx 0.1$. This local volume of the universe contains a diverse array of objects (e.g., galaxies of vastly different masses, morphologies, and star formation rates). But the basic constituents of the GAN objects may be divided into three classes: 1) stars (including their remnants); 2) gaseous systems [the interstellar medium (ISM), circumgalactic medium (CGM), and intergalactic medium (IGM)]; 3) dark components [massive black holes (MBHs), typically seen at centers of galaxies, and dark matter]. The study of the interconnection among these constituents is a key part of the GAN science.  Why is the GAN science important in astronomy and astrophysics? The GAN is where astronomical phenomena can be examined in great detail. Because of their proximity, objects can be observed with unparalleled sensitivities, on small physical scales, across the electromagnetic spectrum, and within a relatively well-determined galactic or intergalactic environment. For example, stellar populations can be resolved into individual stars only in the GAN. Such observations are often necessary in order to firmly identify underlying astrophysical processes and, occasionally, even new physics (e.g., the discovery of dark matter around galaxies). The understanding of astronomical phenomena and astrophysical processes in the GAN thus represents the cornerstone for properly interpreting observations of distant universe. The GAN further provides a test bed to check the validity of various assumptions, approximations, or recipes that are often needed in modeling/simulating the structure formation and evolution of the universe. Locally calibrated empirical relations (such as the peak to width relation of Type Ia SN light-curves and star formation laws) are also very useful tools for studying distant  galaxies. Moreover, local measurements (e.g., star formation rates and total stellar masses of galaxies) provide important anchors in determining how the universe has evolved.  The panel strived to identify key science questions that the GAN can potentially offer the most powerful and unique constraints to. In doing so, a broad range of questions were synthesized and ranked without regard for cost, current agency plan, or specific proposed instrumentation, although the panel was mindful of various limitations, both technical and financial. The four embraced questions, as well as the two identified potential discovery areas, exploit the use of the GAN as a venue for studying interconnected astrophysical systems, for constraining complex physical processes, and for probing small scales.  High signal-to-noise observations, which can often be obtained for GAN objects, need to be interpreted with correspondingly accurate physical data. They are sometimes best obtained by experiment, and sometimes by theoretical calculation, jointly referred to as ``laboratory astrophysics''. Its importance is highlighted in the summary of the GAN \\cite{panel}: ``The prospects for advances in the coming decade are especially exciting in these four areas, particularly if supported by a comprehensive program of theory and numerical calculation, together with laboratory astrophysical measurements or calculations.'' In the following, I briefly describe my interpretation of the laboratory astrophysics needs as well as the questions/discovery areas, particularly with  consideration of the observing capabilities likely available from existing and upcoming NASA missions. ",
        "conclusions": "While progress in addressing the four science questions and  in the areas of discovery potential can be made with existing and upcoming facilities, reaching the full science goals will require powerful new observing capabilities: Imaging/spectroscopy abilities in UV and X-ray will be essential to the understanding of the interconnected, multiphase nature of galaxies and their surroundings, while enhanced capability at longer wavelengths from ratio to optical will be particularly important to probing the processes that transform accreted gas into stars, to measuring the fossil record, and to finding the connections between dark and luminous matter. In many of these efforts, the laboratory astrophysics can play a significant or even essential role!"
    },
    "1112/1112.5120_arXiv.txt": {
        "abstract": "A commonly cited drawback of SPH is the introduction of spurious shear viscosity by the artificial viscosity term in situations involving rotation. Existing approaches for quantifying its effect include approximate analytic formulae and disc-averaged behaviour in specific ring-spreading simulations, based on the kinematic effects produced by the artificial viscosity. These methods have disadvantages, in that they typically are applicable to a very small range of physical scenarios, have a large number of simplifying assumptions, and often are tied to specific SPH formulations which do not include corrective (e.g., Balsara) or time-dependent artificial viscosity terms. In this study we have developed a simple, generally applicable and practical technique for evaluating the local effect of artificial viscosity directly from the creation of specific entropy for each SPH particle. This local approach is simple and quick to implement, and it allows a detailed characterization of viscous effects as a function of position.  Several advantages of this method are discussed, including its ease in evaluation, its greater accuracy and its broad applicability. In order to compare this new method with existing ones, simple disc flow examples are used.  Even in these basic cases, the very roughly approximate nature of the previous methods is shown.  Our local method provides a detailed description of the effects of the artificial viscosity throughout the disc, even for extended examples which implement Balsara corrections.  As a further use of this approach, explicit dependencies of the effective viscosity in terms of SPH and flow parameters are estimated from the example cases. In an appendix, a method for the initial placement of SPH particles is discussed which is very effective in reducing numerical fluctuations. ",
        "introduction": "\\label{sec:intro}  Continued advances in computing power and availability have led to an increasingly important role for numerical simulations in understanding hydrodynamical phenomena, and astrophysics provides a particular range of interesting applications.  High-resolution, multidimensional studies frequently include an ambitious amount of physics in order to reproduce the dynamics of a system with sufficient accuracy. However, numerical artefacts and limitations are necessarily present and must be understood as fully as possible to appreciate the actual physics in the systems being represented.  Non-physical features enter \\textit{ab initio} into all hydrodynamic simulations simply because of the discretisation of a `continuous' fluid (which is itself discrete, of course, but on very small scales). The two types of method developed for this task have very different structures: Lagrangian methods follow moving fluid elements, while Eulerian methods compute flow properties on a grid of points; additionally, there are `hybrid' approaches, the so-called ALE (arbitrary Lagrangian-Eulerian) schemes, e.g. the BETHE-hydro code \\citep{2008ApJS..179..209M}.  In this paper, we focus on the Lagrangian smoothed particle hydrodynamics (SPH) method, which has been widely used for modelling astrophysical phenomena \\citep{1977MNRAS.181..375G,1977AJ.....82.1013L}.  SPH has both advantages and disadvantages with respect to other numerical techniques.  One main advantage is that it automatically adapts to following dynamic flows and arbitrary geometries, without the need for mesh refinement or other readjustment techniques required with Eulerian-based codes.  Furthermore, the energy equation is solved in the local comoving frame of each fluid element, giving a direct implementation of the first law of thermodynamics.  In modelling fluids by means of SPH particles, most SPH codes have included an artificial viscosity term, both for making the ensemble of discrete particles behave more like a continuum (in continuous regions of the flow) and for handling shock discontinuities which may arise during the course of a simulation. Methods for shock-handling with Riemann solvers have also been implemented in SPH \\citep{2002JCoPh.179..238I,2003MNRAS.340...73C}; these approaches have both advantages and disadvantages. Here, we focus on the artificial viscosity methods as used in standard SPH formulations. As well as having beneficial effects, artificial viscosity can also introduce unwanted numerical artefacts, among which spurious shearing torques in rotating flows are particularly troublesome (see, for example, \\citet{1994ApJ...431..754F}). Techniques such as the Balsara correction (described in the next section) have been developed for counteracting this. Typically, SPH prescriptions for artificial viscosity are a composite of both bulk and shear viscosity components, but act as a source of dissipation primarily related to shear when modelling rotating fluid flows. It is important to have some quantitative measure of how large the effects of this are, whether one is dealing with problems having no intrinsic shear viscosity or whether there is also a real, physical shear viscosity; in either case, any numerical viscosity should be suitably negligible (unless one is trying to use an artificial viscosity to actually model a physical viscosity, see e.g., \\citet{2010MNRAS.405.1212L}). We focus here on two standard artificial viscosity prescriptions, with and without the Balsara correction.  The main objective of this paper is to advocate monitoring local entropy generation as a means for quantifying the residual effects of artificial viscosity in regular flows not involving shocks. Our method is completely general, but we test it out by applying it to frequently-used idealized test problems involving isothermal discs and compare its performance on these with that of other alternative methods. Two main types of approach have been used previously, both of which measure the kinematic effect of the artificial viscosity and relate it to an effective viscosity coefficient, $\\nu$. Firstly, there are the empirical `ring-spreading' tests \\citep{1952ZNatA...7...87L,1974MNRAS.168..603L,1994ApJ...431..754F,1996MNRAS.279..402M,1999JCoAM.109..231S}. These are based on the analytic relations for the kinematic contribution of the artificial viscosity to disc-averaged behaviour, with assumptions of an isothermal equation of state and a small viscosity (so as not to generate large radial flow) that is constant across the disc. In practice, these tests can only be performed in two dimensions, as the thermal pressure of the fluid is zeroed to isolate the viscosity term in the equations of motion. These methods also require post-simulation fitting of averaged results. Secondly, analytic approximations have been developed relating $\\nu$ to SPH and disc parameters \\citep{1996MNRAS.279..402M,2010MNRAS.405.1212L}. These relations are generally derived for specific formulations of SPH artificial viscosity with the quadratic term set to zero.  Various additional assumptions include having constant smoothing lengths (in the continuum limit) and particular SPH particle kernel shapes, as well as considerations that the viscosity is active for both receding and approaching particles.  Both of these kinds of existing approach involve restrictive assumptions and neither may be applied to artificial viscosity prescriptions that include often-used `corrective' terms, such as the Balsara method or time-dependent approaches (discussed below), which aim to reduce `excess' shearing.  The technique presented here involves evaluating the effect of artificial viscosity in rotating flows directly from local entropy production (i.e. using the energy equation rather than the equation of motion). It does not assume any particular equation of state or rotation profile and also does not assume constancy of $\\nu$ or of smoothing lengths, or any particular artificial viscosity prescription.  Moreover, it requires no special simulation setup and is directly calculable from standard SPH particle quantities. Impetus for this approach arises from the nature of the SPH method itself and its direct application of the first law of thermodynamics, noting that a clear understanding of entropy production is of great importance in numerical simulations. The technique is simple and quick to implement, and it allows a detailed characterization of the effects of the artificial viscosity as a function of position.  Several advantages of this method over existing ones are discussed, including its ease in evaluation, its direct interpretation, its greater accuracy and its broad applicability, e.g., to arbitary rotation profiles and equations of state.  For all of the simulations described here, we have utilised the public-release version of the Gadget-2 code \\citep{2005MNRAS.364.1105S}.  While several artificial viscosity formulations of varying sophistication have been developed over the years, in this paper we focus mainly on the {\\it method} of estimating the effects of this viscosity and therefore utilize only the relatively simple form that is standard in Gadget-2 (with and without the Balsara treatment, described below). However, the method itself is independent of viscosity prescription, and our aim is to apply it to a comparison of several other formulations in a forthcoming paper. Here, we begin by briefly reviewing in \\S\\ref{sec:newtonhydro} the form of SPH equations used and considerations concerning the artificial viscosity; in \\S\\ref{sec:rotflowartvisc}, the ring-spreading test is described; in \\S\\ref{sec:motiv}, both the motivation for and the general formulation of our local criterion for measuring the effective viscosity in rotating flows are described; in \\S\\ref{sec:general_spec}, example applications of our method in thin discs are given, and various simulations performed, for which the results of the local measures of viscosity are compared with those of the ring-spreading tests and existing analytic formulations; these results are then discussed in \\S\\ref{sec:discuss}.  Additionally, in Appendix \\ref{appendixa}, we discuss the SPH particle setup used for our initial conditions, which is presented as a simple improvement to aid in reducing numerical effects and increasing the efficiency of convergence. ",
        "conclusions": "\\label{sec:discuss}  An entropy-based local method has been described to measure the effective kinematic viscosity in the flow due to SPH artificial viscosity, which can be simply related to the Shakura-Sunyaev $\\alpha$ for appropriate thin discs.  This approach has several advantages over that of estimating a global $\\nu$ from an isothermal ring-spreading test.  First, it calculates the local viscosity of an annulus from the (viscous) energy or entropy expression, which contains only the term dependent on $\\Pi_{ab}$, and therefore only viscous effects are measured.  This procedure then yields $\\nu$ and $\\alpha$ as functions of radius, so that variations in the effective values across a disc can be determined; global values may be obtained easily via averaging, if desired.  In addition, results are generated immediately, without the necessity of finding fits for analytic curves as in ring-spreading, and also thereby eliminating arbitrariness from the estimates. Moreover, the viscosity values are calculated from quantities which are already calculated in the SPH algorithm.  Finally, whereas traditional ring-spreading tests are restricted to use with an isothermal equation of state (or, in practice, even more restricted to cold, 2D systems) because it is only then that one has the analytic solution for comparison, this local method can be used with any equation of state.  The effective viscosity may be obtained for rotating flows in general as a function of viscous dissipation, such as from Eq.~\\ref{general_visc}or \\ref{nu_szmil} (both derived using the very general Eq.~\\ref{visc_en_entr}). The exact form would simply differ from that of Eq.~\\ref{newnucrit}, in which a Keplerian profile ($\\Omega = \\Omega_{{\\rm K}}$) has been assumed.  Further work will also investigate an analogous approach applicable to estimating the effective local shear viscosity for systems in linear motion.  We have compared two existing methods (standard ring-spreading and analytic approximations) with the new, local entropy method in studying the behaviour of the SPH artificial viscosity in rotating flows. In cases where the particular assumptions of the former methods were applicable, average results of all approaches were shown generally to be in fair agreement, with far greater detail coming from the local entropy method. Further examples changing the disc structure and the SPH artificial viscosity values were given.  Examples of explicit expressions which approximate the dependence of the numerical $\\alpha$ on various flow and SPH parameters were given for both the Balsara and non-Balsara corrected forms of artificial viscosity. Procedures such as these may provide useful guides when setting up simulations.  However, further work must be done to test the generalizability of these specific forms, i.e. Eqs.~\\ref{alphalphrelbal}-\\ref{alphalphrel}, across wider parameter space and to other SPH codes. The means for estimating $\\nu$ and $\\alpha$ in each case, however, applies generally with the local method.  We note that the relations between SPH variables and the effective viscosity in disc flows described here can serve as useful examples of the type of analysis possible with the local entropy measures.  The ability for hydrodynamic schemes to represent flows with particular physical viscosity characterizations, such as the widely-used Shakura-Sunyaev discs, may be investigated more directly and in more detail with the local method.  Moreover, this analysis may be useful in the further development and evaluation of physically-motivated viscosity terms in SPH, to which existing methods cannot typically be applied. For example, it was shown here that the $\\alpha$ in the case of Balsara corrected artificial viscosity remained finite and linearly dependent on smoothing length. Studies continue to develop schemes for minimizing unwanted effects of artificial viscosity in rotating flows, such as the recent work by \\citet{2010MNRAS.408..669C}, and we suggest that the simple approach outlined here for measuring the local viscous effects directly would be extremely useful in quantitatively assessing these developments. (Indeed, we will return to discussion of this in a following paper.)  Finally, several studies have described and investigated numerical implementations of physical viscosity prescriptions derived directly from the analytic Navier-Stokes equations (e.g., \\citet{1994ApJ...431..754F,2006A&A...453.1027L}), with separate bulk and shear viscosity terms (unlike the standard SPH artificial viscosity, which does not differentiate the two); this is what one must do in order to incorporate real viscosity behaviour into SPH calculations. However, some numerical considerations remain for such physical viscosity prescriptions, in particular regarding spatial derivatives and dependence on particle disorder (see, e.g., \\citet{2009NewAR..53...78R}, as well as Appendix \\ref{appendixa} of the present paper). While one could think of using the SPH artificial viscosity as a model for a real, physical one (see, e.g., the discussion in Sec. 3.2.3. of \\citet{2010MNRAS.405.1212L}), this provides only a rough approximation, and it is best to think of the present work in terms of establishing acceptable limits for the setup parameters so that the artificial viscosity does not interfere with the effects of the real one."
    },
    "1112/1112.2700_arXiv.txt": {
        "abstract": "We present a novel method to determine eccentricity constraints of extrasolar planets in systems with multiple transiting planets through photometry alone. Our method is highly model independent, making no assumptions about the stellar parameters and requiring no radial velocity, transit timing or occultation events. Our technique exploits the fact the light curve derived stellar density must be the same for all planets transiting a common star. Assuming a circular orbit, the derived stellar density departs from the true value by a predictable factor, $\\Psi$, which contains information on the eccentricity of the system. By comparing multiple stellar densities, any differences must be due to eccentricity and thus meaningful constraints can be placed in the absence of any other information. The technique, dubbed ``Multibody Asterodensity Profiling'' (MAP), is a new observable which can be used alone or in combination with other observables, such as radial velocities and transit timing variations. An eccentricity prior can also be included as desired. MAP is most sensitive to the minimum pair-combined eccentricity e.g. $(e_1 + e_2)_{\\mathrm{min}}$. Individual eccentricity constraints are less stringent but an empirical eccentricity posterior is always derivable and is freely available from transit photometry alone.  We present a description of our technique using both analytic and numerical implementations, followed by two example analyses on synthetic photometry as a proof of principle. We point out that MAP has the potential to constrain the eccentricity, and thus habitability, of Earth-like planets in the absence of radial velocity data, which is likely for terrestrial-mass objects. ",
        "introduction": "\\label{sec:intro}  In February of 2011, 1235 \\emph{Kepler} transiting candidate planets were announced by \\citet{borucki:2011}, amongst which the majority are expected to be genuine \\citep{morton:2011}. At the latest counting, the score has since risen to 1781 \\citep{rowe:2011} and is expected to continue rising. Due to the unprecedented yield of new transiting planet candidates, follow-up with radial velocity (RV) measurements is generally not feasible due to both the typical faintness of the targets and the intensive nature of the required telescope time for so many targets. Historically, radial velocity has emerged as the tool of choice to confirm transiting candidates and so the \\emph{Kepler} team have devoted considerable effort to find ways to confirm candidates without the need for RV. This has led to some pioneering techniques such as blend analysis \\citep{torres:2011,fressin:2011} and confirmation through transit timing variations (TTV) \\citep{holman:2011,lissauer:2011}.  However, even though transiting candidates have been shown to be confirmable without RV, its absence means that the orbital eccentricity, $e$, of the planets cannot be determined (unless very strong TTVs are detected). One remaining avenue to constrain $e$ is to detect an occultation event. Occultations occur exactly half an orbital period after the transit event for a circular orbit\\footnote{There also exists a small light travel time across the system} but become offset for eccentric orbits, thus offering a potential diagnostic of eccentricity (details on the precise obtainable constraints are provided in \\citealt{thesis:2011}). Such occultations are due to a combination of reflected light from the planet and thermal emission (which is very small in \\emph{Kepler}'s visible bandpass e.g. \\citealt{darkest:2011}). Despite \\emph{Kepler's} ground-breaking photometric precision, most transiting planets discovered so far have not exhibited such events (one notable exception is the high albedo planet Kepler-7b, see \\citealt{kippingbakos:2011a}). Therefore, in the majority of cases we are left without any way of characterizing the orbital eccentricity.  Our principal motivation for addressing this problem stems from the fact that based upon the estimated frequency of Earth-like planets \\citep{howard:2010,catanzarite:2011,wittenmyer:2011}, it seems probable that \\emph{Kepler} will detect numerous habitable-zone Earth-radius planets. After the initial detection, the natural question we will be ``can this world sustain life?''. A high orbital eccentricity causes a planet to spend the majority of its orbit outside the habitable-zone and thus leads to potentially marginal or transient habitability \\citep{williams:2002,dressing:2010}. For example, eccentricities in excess of $0.3$ for a habitable-zone planet around a Sun-like star cause equilibrium temperatures to vary by more than 100\\,K. The ability to measure, or at least constrain, eccentricity therefore would greatly benefit the assessment of an exoplanet's habitability.  Recently, \\citet{moorhead:2011} (M11 from here on in) have studied the problem in an effort to glean some information about the eccentricity of a transiting planet. Their approach is that if one knows the stellar density a-priori, say from stellar spectroscopy combined with evolution models, then one can compute the maximum allowed transit duration for a planet on a circular orbit using (note that this approach was originally outlined in \\citealt{ford:2008}):  \\begin{equation} \\tilde{T}_{\\mathrm{max}}^{\\mathrm{circ}} \\simeq \\frac{R_* P}{\\pi a} \\end{equation}  where $\\tilde{T}$ is the duration between the planet's centre crossing the stellar limb to exiting under the same condition, $R_*$ is the stellar radius, $P$ is the planet's orbital period and $a$ is the planet's semi-major axis. The above equation assumes an equatorial transit and hence is the maximum duration possible. M11 discuss how if the observed transit duration exceeds this quantity (i.e. $\\tilde{T}/\\tilde{T}_{\\mathrm{max}}^{\\mathrm{circ}}>1$), this indicates that the orbit must be eccentric. The two weaknesses in this approach are that: 1) The spectroscopically determined value of $R_*$ has both a large statistical uncertainty and a large and unknown systematic uncertainty (for example, \\citet{brown:2011} state that the Kepler Input Catalogue effective temperatures and radii estimations are reliable for Sun-like stars, but are ``untrustworthy'' for stars with $T_{\\mathrm{eff}}<3750$\\,K). 2) A planet can be eccentric yet still cause $\\tilde{T}/\\tilde{T}_{\\mathrm{max}}^{\\mathrm{circ}}<1$ i.e. only planets transiting near to apoapse, the slowest part of the orbit, will be identified as eccentric, which is for $\\omega\\sim270^{\\circ}$. The first weakness means the technique is model-dependent and that an individual system may not be reliable due to possible systematic errors in $R_*$. The second weakness limits the scope of application of the technique to planets with $b\\sim0$ and $\\omega\\sim270^{\\circ}$, which it should be noted is the least probable value of $\\omega$ from geometric priors \\citep{kane:2008}.  M11 show that the $R_*$ uncertainty weakness can be overcome by adopting a statistical perspective. Even though an individual system may not be reliable, the bulk of systems should be and so any overall distributions which emerge should be reliable. The valuable technique of M11 allows us to actually say something about the eccentricities of the \\emph{Kepler} candidates as a whole.  But what about individual systems? Or those which don't both fortuitously transit near $\\omega\\sim270^{\\circ}$ and have $b\\sim0$? If the star has been studied with asteroseismology, then this prior can solve the riddle (as pointed out in M11). This is scenario is discussed in more detail later in \\S~\\ref{sub:SAP}. However, to date, relatively few targets have had asteroseismology studies conducted. We are therefore left with the quandary that it is usually impossible to say anything empirical about the orbital eccentricity of a transiting planet through photometry alone\\footnote{Unless an occultation is observed}. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Recurring Patterns in MAP} \\label{sub:patterns}  In the 2D posteriors shown in the example systems KOI-SOG, KOI-SOC, KOI-S01 and KOI-S02, it is possible to visually detect some recurring patterns. In this subsection, we will discuss the reason for these patterns.  The most obvious pattern is that the plots tend to exhibit higher densities for $e_x \\simeq e_y$. This results in a diagonal high density region extending from the bottom-left to the top-right, getting narrower as it goes up. Why does this happen?  Consider two planets on circular orbits. The light curve derived stellar densities will be identical meaning MAP will find that a system with two circular orbits to be a highly compatible solution. However, it is not the only compatible solution. Consider a MAP realization with $e_1 = 0$ and $e_2 = 0.1$. In this case, a rather broad range of $\\omega_2$ values can reproduce a stellar density which is still equivalent (or within error) to $\\rho_{*,1}^{\\mathrm{circ}}$. As $e_2$ increases to higher eccentricities (i.e. as we move along one of the axes), the range of $\\omega_2$ values which can still reproduce such a result diminishes. Consequently, the integrated probability density over all $\\omega_2$ values is small and thus the region becomes a low-probability sector.  Now consider MAP realizations where both $e_1$ and $e_2$ depart from zero. In this case, a situation where $e_2 = 0.9$ and $e_1 = 0.1$ is not grossly different from the case just described where $e_2$ is very large and $e_1$ is zero. Thus these regions tend to be low probability. However, if $e_1 \\simeq e_2$ and both values are large then the range of $\\omega_1$ and $\\omega_2$ values which can reproduce $\\rho_{*,1}^{\\mathrm{circ}} \\simeq \\rho_{*,2}^{\\mathrm{circ}}$ is broader. Thus the integrated probability density becomes larger too. This can be seen be simple inspection of Equation~\\ref{eqn:rhoratio}. For this reason, a common pattern seen in MAP is the tail extending along the $e_x = e_y$ axis.  \\subsection{MAP in Combination with Other Observables} \\label{sub:MultiMAP}  Multibody Asterodensity Profiling (MAP) can provide constraints on the orbital eccentricity of transiting planets in multiple planet systems without the need for radial velocities, occultations or transit timing variations. However, there may some cases where these observables are available. In such a case, MAP can be combined with these other observables to further refine the constraints on the orbital eccentricities. Our n-MAP technique utilizes an MCMC routine to explore the parameter space of possible orbital eccentricities. This MCMC works by assigning a $\\chi^2$ merit-function to each point and then proceeding via the Metropolis-Hastings rule. Using n-MAP alone, the merit function is:  \\begin{align} \\chi^2 &= \\chi_{\\mathrm{MAP}}^2 \\end{align}  where $\\chi_{\\mathrm{MAP}}^2$ is given in Equation~\\ref{eqn:chi2}. If additional observables are available, such as TTVs or RVs then one may simply append the associated merit functions:  \\begin{align} \\chi^2 &= \\chi_{\\mathrm{MAP}}^2 + \\chi_{\\mathrm{RV}}^2 + \\chi_{\\mathrm{TTV}}^2 + ... \\end{align}  Once appended, the n-MAP routine is simply executed as before. In our next paper, we will provide an analysis of this technique on a real transiting planet system featuring both radial velocities and transit timing variations.  \\subsection{Differences to the \\citet{moorhead:2011} Technique}  A direct comparison to the method of M11 is not fair because the techniques operate under different conditions and assumptions. Despite this, we will here outline a few differences between the two methods. One advantage of the M11 method is that it works for all transiting planets whereas MAP requires $>1$ transiting planet in a given system. One advantage of MAP over the M11 technique is that MAP is highly model independent\\footnote{There is a very weak model dependency in MAP via the adopted limb darkening law, but this can also be fitted for with sufficient signal to noise}, assuming only the planets orbit the same star whereas M11 require a value for the stellar radius determined via the more model sensitive route of stellar evolution.  In principle, MAP exploits more information in the light curve than that of the M11 technique. M11 compare the observed transit duration to the maximum theoretical value for a circular orbit; in other words M11 make use of one metric for the duration. Typically, this metric is the transit duration defined as the time for the planet to move from its centre crossing the stellar limb to exitting under the same condition, $\\tilde{T}$, since this is independent of the derived planetary radius. In contrast, MAP uses both $\\tilde{T}$ and the ingress duration, $T_{12}$, to derive the light curve derived stellar density assuming a circular orbit, $\\rho_{*}^{\\mathrm{circ}}$. Thus, it can be appreciated that MAP uses the same information as M11 plus some extra information. In the limit of ignoring this ingress information, the fundamental transit information used by both techniques would be identical.  As a result of M11 negating the ingress duration information, the impact parameter is unresolved. For this reason, M11 adopt the conservative assumption that the limiting case is for $b=0$. \\citet{ford:2008} alternatively discuss how a prior in $b$ could be adopted by assuming an isotropic distribution in orbital inclination. In contrast, MAP takes the distribution in $b$ from the data itself, essentially characterized by the ingress duration. Due to the $b=0$ conservative assumption made by M11, only transit durations longer than this limiting case can ever be detected. Unlike MAP, this limits the method to detecting eccentric planet transiting near apocentre, the slowest part of the orbit, and also the least likely geometric configuration to detect a transiting planet in \\citep{kane:2008} (this is also pointed out in \\citealt{tingley:2005}). This bias, discussed in M11, requires de-biasing any eccentricity statistics deduced and of course reduces the overall sample size since only a subset of eccentric planets are detected. For this reason, we anticipate MAP would provide a more powerful diagnostic of the statistics of eccentric planets, but of course is only applicable in multiple systems.  \\subsection{Single-body Asterodensity Profiling (SAP)} \\label{sub:SAP}  MAP makes no assumption about the properties of the parent star. For stars with poor characterization, this is an advantage since the stellar properties are frequently subject to unknown systematic uncertainties. However, in some cases the stellar properties are well-characterized and this is information which MAP ignores. An example of this is a star which has been studied with asteroseismology, leading to a highly precise determination of $\\rho_*$.  Single-body Asterodensity Profiling (SAP) is the logical extension of MAP which can include this information. In this work, we argue in favour of not using SAP due to the frequently unreliable stellar parameters and the strong model dependency of those derived results, which is in stark contrast to the MAP technique. However, for sake of completion we will discuss here a possible implementation of SAP.  In each MCMC trial of the n-MAP algorithm, we create a \\{$e_k,\\omega_k$\\} trial vector. This may be used to construct a $\\{\\Psi_k\\}$ vector. For a single planet, this would be simply be a 1-dimensional vector: $\\{\\Psi_1\\}$. This value may be used to infer the true stellar density based upon the derived value of $\\rho_{*,1}^{\\mathrm{circ}}$ via:  \\begin{align} \\rho_{*,1}^{\\mathrm{SAP}} &= \\rho_{*,1}^{\\mathrm{circ}}/\\Psi_1 \\end{align}  For each MCMC trial, we can use the same value of $\\rho_{*,1}^{\\mathrm{circ}}$, namely the median of the light curve derived posterior. We now simply define $\\chi^2_{\\mathrm{SAP}}$ as the number of standard deviations between this trial value of $\\rho_{*,1}^{\\mathrm{MAP}}$ and the empirically determined value of $\\rho_*$, say $\\rho_{*}^{\\mathrm{seismology}}$:  \\begin{align} \\chi_{\\mathrm{SAP},k}^2 &= \\frac{[\\rho_{*,k}^{\\mathrm{SAP}} - \\rho_{*}^{\\mathrm{seismology}}]^2}{[\\sigma(\\rho_{*,k}^{\\mathrm{SAP}})]^2 + [\\sigma(\\rho_{*}^{\\mathrm{seismology}})]^2} \\end{align}  where we have replaced the 1 subscript (for planet ``1'') with $k$ to make it a general equation. Clearly for $n$ planets we have $n$ contributions to the total $\\chi^2$ function from SAP:  \\begin{align} \\chi_{\\mathrm{SAP}}^2 &= \\sum_{k=1}^n \\chi_{\\mathrm{SAP},k}^2 \\end{align}  This may be combined with the MAP, RV or TTV merit functions as desired to produce finer constraints on the orbital eccentricity.  \\subsection{Overview} \\label{sub:overview}  In this work, we have presented a new method to photometrically constrain the orbital eccentricities of transiting planets. The method is only applicable to multiple transiting planet systems and relies on the key assumption that all of the transiting planets orbit the same star. The new method works by comparing the light curve derived stellar density between each planet and thus is dubbed ``Multibody Asterodensity Profiling'' (MAP). MAP requires no prior information on the star's properties and thus is highly robust against systematic uncertainties.  MAP constitutes a new observable for which the likelihood of a given orbital configuration can be computed in a $\\chi^2$-sense. Thus, MAP be be combined with other pieces of information about the eccentricity of the system e.g. transit timing variations, radial velocities, occultations. In this work, we have adopted uniform priors for the orbital eccentricities but more realistic priors based upon dynamics or planet formation can also be invoked (see \\S\\ref{sub:directpriors}).  In its simplest form, MAP can be applied by employing some simple analytic expressions (a-MAP) to deduce the combined eccentricities of two planets (e.g. Equation~\\ref{eqn:doubleA}). In this sense, MAP exhibits impressive sensitivity to a system with even a moderate to low eccentricity. For example, for a synthetic system with one planet with $e=0.15$ (KOI-S01, see \\S\\ref{sub:KOI-S01}), Q0-Q2 Kepler data can infer a significant eccentricity at the 13-$\\sigma$ level.  To determine the individual eccentricities require the use of numerical methods, or n-MAP. n-MAP is shown to recover the same constraints for the pair-combined eccentricities as a-MAP does, but additionally provides individual constraints as well as a fuller picture of the inter-relationships between the various eccentricity terms. However, we find MAP is more sensitive to the minimum pair-combined eccentricities than individual terms. Nevertheless, an empricial determination of the posterior for each term is always derivable. The highly model independent nature of MAP means that these posteriors will narrow ad-infinitum as more data and signal-to-noise is accumulated. Further, the application of MAP to dozens of systems raises the potential of discerning truly unbiased statistics on the eccentricity distribution of planets in multiple systems.  Ultimately, MAP has the potential to characterize the eccentricity of the first truly habitable Earth-like planet, which could be found by \\emph{Kepler}. In such a case, the Earth-mass planet will likely be too challenging to detect with radial velocities but MAP will be an ever-present tool provided the photometric time series is of reasonable quality and the system has more than one transiting planet (true of $\\sim$50\\% of all transiting planet candidates found by \\emph{Kepler}). In conclusion, MAP offers a method to diagnose both interesting dynamical systems and potentially interesting astrobiological targets too.  The n-MAP algorithm is available as a Fortran 90 script upon request."
    },
    "1112/1112.6229_arXiv.txt": {
        "abstract": "{It has been proposed that ultracompact minihalos (UCMHs) might be formed in earlier epoch. If dark matter is in the form of Weakly Interacting Massive Particles (WIMPs), UCMHs can be treated as the $\\gamma$-ray sources because of the dark matter annihilation within them. In this paper, we investigate the contributions of UCMHs formed during three phase transitions (electroweak symmetry breaking, QCD confinement and $e^{+}e^{-}$ annihilation) to the extragalactic $\\gamma$-ray background. Moreover, we use the Fermi-LAT observation data of the extragalactic $\\gamma$-ray background to get the constraints on the current abundance of UCMHs produced during phase transitions. We also compare these results with those obtained from Cosmic Microwave Background (CMB) observations finding that the constraints from the Fermi-LAT are more stringent than CMB.} ",
        "introduction": "It has been known that the present structures of cosmology originate from the density perturbations ($\\delta \\sim 10^{-5}$) in earlier epoch. Recently, the authors of \\cite{0908.0735} have proposed that a new class of dark matter structure, ultracompact minihalos (UCMHs) would be formed if the density perturbations are between $3 \\times 10^{-4}$ and 0.3 which are not enough to produce the primordial black holes (PBHs). Although within the conventional cases the density perturbations are not larger enough to form these objects, they could be enhanced through the inflation potential or during the phase transitions in the early Universe \\cite{PRL_1997}. On the other hand, if dark matter is in the form of Weakly Interacting Massive Particles (WIMPs), UCMHs would become the sources of $\\gamma$-ray due to the dark matter annihilation within them. In paper \\cite{scott_PRL_2009}, the authors have investigated the integrated $\\gamma$-ray flux of nearby UCMHs formed during three transitions: electroweak symmetry breaking (EW), QCD confinement and $e^{+}e^{-}$ annihilation. They find that for the same mass fraction, the flux from UCMHs formed during $e^{+}e^{-}$ within the 100pc would have been observed by EGRET or Fermi-LAT. The authors of \\cite{josan_PRD_2010} have obtained the constraints on the abundance of UCMHs using the point source sensitivity of $Fermi$ above 100 MeV from the neighborhood: $f_{UCMHs} \\sim 10^{-7}$ for $M_{UCMHs} \\sim 10^{3} M_{\\odot}$. The constraints on the UCMHs formed during $e^{+}e^{-}$ annihilation using the WMAP-7 years data and the forcast for the Planck-3 are obtained by authors in \\cite{yyp,epjp}.  \\par The extragalactic $\\gamma$-ray background has been observed by $Fermi$ satellite \\cite{fermi_extra} and its origination has not known exactly now. The main sources would be the usual, unresolved astrophysical objects, such as the Active Galactic Nuclei (AGN), normal galaxies and the clusters of galaxies \\cite{dermer_2007,vasiliki_02}. Other possible sources may be the dark matter annihilation \\cite{mack_94,elsosser_04,ullio_02}. In our paper, we consider the contributions of dark matter annihilation from UCMHs, and get the constraints on the current abundance of them using the $Fermi$ observations.  In addition, the annihilation of dark matter has an effect on the cosmological evolution \\cite{lezhang_PRD_2006} e.g. recombination and reionization. So the property of dark matter can be constrained by the CMB observations. UCMHs would have the similar effect because of their higher density, and the CMB data can give constraints on their nature. Here, we also use the CMB data to investigate the current abundance of UCMHs.     \\par  This paper is organized as following: we calculate the extragalactic $\\gamma$-ray background from UCMHs in section II. In section III, we get the constraints on the current abundance of UCMHs using the $Fermi$ and CMB data, and we conclude in section IV. ",
        "conclusions": "We have investigated the extragalactic $\\gamma$-ray background from a new class of dark matter structure named UCMHs and gotten the constraints on the current abundance of them using the $Fermi$ observations. On the other hand, because the dark matter annihilation within the UCMHs has effect on the evolution of cosmological ionization fraction, we also use the CMB data to get the constraints on the fraction of UCMHs. In our paper, we only show the constraints on the UCMHs formed during $e^{+}e^{-}$ due to the same contributions of UCMHs produced during EW and QCD. \\par From the results, we find that the constraints are weaker for CMB data than $Fermi$ data. About the dark matter mass $M_{\\chi} \\sim 10 GeV$, the current fractions of UCMHs are about $\\sim 10^{-5}$ and $\\sim 10^{-4}$ respectively corresponding to $Fermi$ and CMB data. For the $Fermi$ observations, we have calculated two channels: $\\tau^{+} \\tau^{-}$ and $b \\bar b$. We find that for lower mass of dark matter $M_{\\chi} \\lesssim 100 GeV$, the constraints are comparable, while for larger one the constraints from $b \\bar b$ channel are more stringent. If we consider the contributions from the halos and subhalos, the results are not changed nearly especially for larger dark matter mass $M_{\\chi} \\gtrsim 100 GeV$. \\par  For this new dark matter structures, it is hope that other present or the future observations can have the power to give new constraints on them. For the recent observations of positron, such as PAMELA \\cite{pamela}, Fermi \\cite{fermi} and ATIC \\cite{atic}, may be the signal from the dark matter annihilation (for a review one can see \\cite{0909.1182}). The characteristic of UCMHs would explain these observations. On the other hand, the extra energy injection into the medium can also have effect on the 21cm signals. So we expect that this future observations can also give the new and stronger constraints on the UCMHs.  Moreover, the constraints on the abundance of UCMHs can be translated into constraints on the primordial curvature perturbation~\\cite{1110.2484}. We will consider them in the near future."
    },
    "1112/1112.3126_arXiv.txt": {
        "abstract": "We report the results of photometric observations of the blazars Mrk 421 and 3C 454.3 designed to search for intraday variability (IDV) and short-term variability (STV). Optical photometric observations were spread over eighteen nights for Mrk 421 and seven nights for 3C 454.3 during our observing run in 2009-2010 at the 1.04 m telescope at ARIES, India. Genuine IDV is found for the source 3C 454.3 but not for Mrk 421.  Genuine STV is found for both sources. Mrk 421 was revealed by the MAXI X-ray detector on the International Space Station to be in an exceptionally high flux state in 2010 January - February.  We performed a correlation between the X-ray and optical bands to search for time delays and found a weak correlation with higher frequencies leading the lower frequencies by about ten days. The blazar 3C 454.3 was found to be in high flux state in November-December 2009. We performed correlations in optical observations made at three telescopes, along with X-ray data from the MAXI satellite and public release $\\gamma$-ray data from the Fermi space telescope. We found strong correlations between the $\\gamma$-ray and optical bands at a time lag of about four days but the X-ray flux is not correlated with either. We briefly discuss the possible reasons for the time delays between these bands within the framework of  existing models for  X-ray and $\\gamma$-ray emission mechanisms. ",
        "introduction": "Blazars are the radio$-$loud Active Galactic Nuclei (AGN) that are classified as BL Lacertae objects (BL Lacs) if they have largely featureless optical spectra or as flat spectrum radio quasars (FSRQs) if they have prominent emission lines. All blazars are characterized by broadband non-thermal emission extending over the complete electromagnetic spectrum, strong polarization from radio to optical wavelengths, and displays of violent variability on timescales that can extend from a fraction of an hour to many years. In the usual orientation-based unified model of radio$-$loud AGN,  blazar jets make an angle of $\\leq$ 10$^\\circ$ from the line of sight and the emission from these jets is thus Doppler boosted and dominates what we observe (Urry \\& Padovani 1995).  Blazar variability timescales have often been arbitrarily divided into three classes: timescales from a few tens of minutes to less than a day are called intraday variability (IDV) (Wagner \\& Witzel 1995) or micro-variability or intranight variability, those from several days to a few months are short timescale variability (STV), while long timescale variability (LTV) covers changes from several months to many years (Gupta et al.\\ 2004).  The spectral energy distribution (SED) of blazars have two peaks (e.g., Giommi, Ansari \\& Micol 1995; Fossati et al.\\ 1998).  The locations of those peaks can be used to classify blazars into LBLs (Low Energy Peaked Blazars)  with the first hump in the near infrared (NIR) or optical band and the second hump usually peaking at GeV $\\gamma$-ray energies, while HBLs (High Energy Peaked Blazars) are those with first peak in the UV or X-ray band and the second  peak located at  up to TeV energies (e.g., Padovani \\& Giommi 1995). The high polarization of the radio to optical emission suggests that the lower energy peak is produced via the synchrotron process but the high energy emission mechanism in blazars is not yet fully understood, though it is probably due to the inverse Compton (IC) mechanism.  \\subsection{\\bf Mrk 421}  With its redshift $z=0.031$, Mrk 421 (B2 1101$+$384) ($\\alpha_{2000.0}$=11h 04m 27.2s and $\\delta_{2000.0}$=$+38^{\\circ} 12^{'} 32^{\\prime \\prime}$) is among the closest blazars, at a distance of 134 Mpc (H$_{0} =$ 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m} =$ 0.27, $\\Omega_{\\lambda} =$ 0.73).  It is classified as an HBL because the energy of its synchrotron peak is higher than 0.1 keV. It is the brightest TeV $\\gamma-$ ray emitting blazar in northern hemisphere.  Mrk 421 was first noted to be an object with a blue excess which turned out to be an elliptical galaxy with a bright point like nucleus (Ulrich et al.\\ 1975). The object showed optical polarization and the spectrum of the nucleus was seen to be featureless and so it was classified as a BL Lac.   Mrk 421 was detected in the GeV band by the EGRET instrument on the Compton Gamma$-$ray Observatory (CGRO) (Lin et al. 1992; Michelson et al. 1992). It was the first known extragalactic TeV $\\gamma-$ray emitter (Punch et al.\\ 1992), and has been repeatedly confirmed as a TeV source by  ground-based $\\gamma$-ray telescopes (Aleksic et al. 2011; Acciari et al. 2011 and references therein).  Mrk 421 has been extensively observed at all wavelengths and some noticeable studies in low-energy bands include an exhaustive compilation of radio data at 22 and 37 GHz over about 25 years (Ter{\\\"a}sranta et al.\\ 2004, 2005). The source is  characterized by strong variability in the optical region (e.g., Miller 1975; Liu et al.\\ 1997) including LTV  of $\\sim$ 4.6 mag  (Stein et al.\\ 1976) and extreme rapid optical variability  exemplified by a 1.4 magnitude brightness change in only 2.5 hours (Xie et al.\\ 1988).  NIR data extending over about three decades were given by Fan \\& Lin (1999) and Gupta et al.\\ (2004) reported detection of IDV and STV in the blazar in the IR.  Weakly correlated X-ray and NIR variability was reported in the source by Gupta et al.\\ (2008).  It is worth noting that Mrk 421 is almost always very active at high energies. % In spring/summer 2006, Mrk 421 reached its highest X-ray flux recorded until that time with a peak flux $\\sim$85 mCrab in the 2.0--10.0 keV band, with the corresponding first peak of the SED often occurring at $>$10 keV (Tramacere et al.\\ 2009).   A hard X-ray flare from Mrk 421 was detected by Super-AGILE on 2008 June 10 (Costa et al.\\ 2008) which was followed by detection in $\\gamma-$rays (Pittori et al. 2008) by the AGILE/GRID (Gamma-ray Imaging Detector).   Ushio et al.\\ (2009) presented the observations of X-ray variability of Mrk 421 with Suzaku.  Recently, strong X-ray flares were detected from the source in 2010 January and February with the Monitor of All-sky X$-$ray Image (MAXI) instrument on the International Space Station (ISS).  The February 2010 flare reached 164 mCrab and is the strongest among those reported from the object (Isobe et al.\\ 2010).  The time-average energy spectrum of Mrk 421 during the flaring stage has been measured at high energies with HESS using large-zenith-angle observations (Aharonian et al.\\ 2005) and with MAGIC (Albert et al.\\ 2007). X$-$ray and TeV flares were observed around May 16, 1994 (Takahashi et al.\\ 1994; Kerrick et al.\\ 1995) and around April 25, 1995 (Takahashi et al.\\ 1995). The observed correlated variability between X-rays and TeV $\\gamma-$rays (Maraschi et al.\\ 1999; Fossati et al.\\ 2008) can be explained in the synchrotron-self Compton (SSC) framework (Ghisellini et al.\\ 1998), whereas the external Compton (EC) scenario is unlikely to  apply in HBLs due to the low density of ambient photons. There was a detection of a rapid variability timescale of TeV $\\gamma-$ray emission from Mrk 421 ($\\sim$10 min; Gaidos et al.\\ 1996) that may require a very large Doppler factor $\\delta \\geq$ 50 (Ghisellini \\& Tavecchio 2008).   Mrk 421 has been a target of several simultaneous multi-wavelength monitoring campaigns (Takahashi et al.\\ 2000; Rebillot et al.\\ 2006; Fossati et al.\\ 2008; Lichti et al.\\ 2008). Despite all the studies of this source, we found that there have been a very few investigations of IDV and STV of Mrk 421 in optical bands.  So we decided to pursue the present study which addresses both the IDV and STV of Mrk 421. Our optical observations are synchronized with the X$-$ray observations from MAXI.  The present optical observations give simultaneous information in an additional spectral window and make it possible to search for the correlations and time delays between optical and X$-$ray bands. This work could give us useful input for multi$-$wavelength modelling of blazars and lead to a better  understanding of the cause of their variability.  \\subsection{\\bf 3C 454.3}  3C454.3 (PKS 2251+158) ($\\alpha_{2000.0}$=22h 53m 57.75s $\\delta_{2000.0}$=$+16^{\\circ} 08^{'} 53.56^{\\prime \\prime}$) is a well known flat spectrum radio quasar (FSRQ) at redshift z=0.859. It has displayed pronounced variability at all wavelengths and has been extensively observed over the years in most energy bands, from radio (e.g., Bennett 1962) through microwave (Bennett et al.\\ 2003), optical (e.g., Sandage 1966; Raiteri et al. 1998), X$-$ray (e.g., Worrall et al.\\ 1987; Tavecchio et al.\\ 2002), low energy $\\gamma-$ray (Blom et al.\\ 1995; Zhang et al.\\ 2005) and high energy $\\gamma-$ray (Hartmann et al.\\ 1993, 1999).  3C 454.3 entered a bright phase starting at 2000, and has shown remarkable activity in the past decade (Fuhrmann et al.\\ 2006; Villata et al.\\ 2006, 2007, Raiteri et al.\\ 2007, 2008). In spring 2005 it underwent major outbursts, reaching an R-band magnitude of 12.0 that was the largest apparent optical luminosity ever recorded from this blazar (Villata et al.\\ 2006). A Whole Earth Blazar Telescope monitoring effort continued after the optical outburst and followed the subsequent radio activity (Villata et al.\\ 2007) and then the faint state in the 2006-2007 observing season (Raiteri et al.\\ 2007). In this last period, the relatively low contribution of the synchrotron emission from the jet meant that the ``little blue bump'', believed to arise from Fe-line emission from the broad line region, as well as the ``big blue bump'', due to the nearly thermal emission from the accretion disc were recognizable. An increase in activity occurred at X-ray and radio wavelengths as well, with the 230 GHz radio variations having a delay of $\\sim$ 2 months with respect to the optical variability (Raiteri et al.\\ 2008). After the quiescent state, the quasar underwent a new stage of high optical activity (Raiteri et al. 2008) that continued to the end of 2008. UV excesses observed at some times are likely to be a signature of the thermal radiation from the accretion disk feeding the central black hole.  Swift telescope observations of 3C 454.3 (Giommi et al.\\ 2006) and hard X-ray observations (Pian et al.\\ 2006) covering this active phase in 2005 detected large fluxes between 10 and 100mCrab. Starting with the 2007 July AGILE detection above 100 MeV (Vercellone et al.\\  2008), 3C454.3 has been very active in $\\gamma$-rays. It was subsequently monitored by AGILE in 2007-2009 and showed repeated flares that usually coincide with periods of intense optical and enhanced X-ray activity (Chen et al.\\ 2007, Donnarumma et al.\\ 2009). During this time span, very bright $\\gamma-$ray emission was detected (Tosti et al.\\ 2008), with an excellent correlation between the  $\\gamma$ ray and NIR/optical variations (Bonning et al.\\ 2009).  The source is listed as 1FGL J2253.9$+$1608 in the First FERMI-LAT active galactic nucleus (AGN) catalog (Abdo et al.\\ 2010). 3C454.3  is the first source for which daily resolved broadband spectral energy distribution (SEDs) with GeV data have been obtained (Abdo et al.\\ 2009). Strong Ly$\\alpha$ radiation has been seen from 3C454.3 (Bonnoli et al. 2010), indicating the presence of an external photon source for the Compton scattering aside from torus emission (Sikora et al.\\ 2009). Modeling of SEDs has been performed by Finke \\& Dermer (2010), and Pacciani et al.\\ (2010).  This FSRQ showed strong activity at optical frequencies in 2008-2009. (Villata et al.\\ 2008; Sasada et al.\\ 2009; Gupta et al.\\ 2009). Foschini et al.\\ (2010) claim $\\gamma-$ray variability from 3C454.3 on timescales as short as a few hours from the LAT data. The continuous monitoring by the Fermi-LAT showed that the source activity faded considerably in early 2009 and then rose back up from June onward. It underwent an exceptional outburst in 2009 November$-$ 2010 January when it became brightest $\\gamma-$ray source in the sky for over a weak. (Striani et al.\\ 2009, 2010; Escande \\& Tanaka 2009).  Here we present our optical observations from ARIES, India during the outburst flare in November-December 2009, along with data from the optical monitoring program during this period with the Small and Moderate Aperture Research Telescope System (SMARTS) in Chile as well as optical data from the 1.5 m telescope of KANATA observatory in Japan. In order to see if there is related variability of the blazar 3C 454.3 in optical, X-ray and gamma bands, we correlate the above optical data with the X-ray data (in 2-10 KeV) obtained from MAXI as well as with 1-300 Gev fluxes made public by the Fermi Science Support Center. These results should be useful for understanding the $\\gamma$-ray and X-ray emission mechanisms and could also offer a test of the existing models (leptonic and hadronic) for the $\\gamma$-ray emission.  The paper is structured as follows.  In Section 2, we give brief descriptions of observations and data reduction methods. In Section 3, we discuss the techniques we used to search for variability properties and we provide the results in Section 4.  A discussion and our conclusions are given in Section 5. ",
        "conclusions": ""
    },
    "1112/1112.1123_arXiv.txt": {
        "abstract": "Some modern models of neutron star evolution predict that initially large magnetic fields rapidly decay down to some saturation value $\\sim {\\rm few}\\times 10^{13}$~G and weaker magnetic fields do not decay significantly \\citep{Pons09}. It is difficult to check the predictions of this model for initially highly magnetized objects on the time scale of a few million years. We propose to use Be/X-ray binaries for this purpose. We apply several methods to estimate magnetic fields of neutron stars in these accreting systems using the data obtained by the RXTE satellite \\citep{Galache05}. Only using the most modern approach for estimating the magnetic field strengths of long  period NSs as proposed by \\citet{Shakura11} we are able to obtain a field distribution compatible with predictions of the theoretical model of field decay of \\citet{Pons09}. ",
        "introduction": "Neutron stars (NSs) are the final stage of the evolution of massive stars. There are different types of young isolated NSs: radio pulsars (see a review in \\citealt{ATNF}), compact central X-ray sources in supernova remnants \\citep{Pavlov04, Luca08}, magnetars: anomalous x-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs) \\citep{Mereghetti08, Esposito11}, close-by cooling radioquiet isolated NSs -- the Magnificent seven (M7) \\citep{Haberl07, Kaplan08}, rotating radio transients (RRATs) \\citep{Keane10}. Another large class of observed systems with NSs is formed by X-ray binaries.  They are subdivided into high mass X-ray binaries (HMXB) and low mass X-ray binaries (LMXB) (see a review on X-ray binaries and their evolution  in \\citet{Postnov06, Bhattacharyya10}).  In order to deeper comprehend the evolution and to follow links between different types of NSs we need a better understanding of the basic processes related to NS thermal and magneto-rotational evolution. For the initially highly magnetized NSs ($B\\gtrsim \\rm{few} \\times 10^{13}$~G), modern models predict that  the magnetic field decays significantly on the time scale of several hundred thousand years \\citep{Aguilera08a}. For magnetars there is a strong observational evidence for magnetic field decay (see, for instance, \\citet{Pons07} and references therein). These sources are very young NSs with ages of the order of tens of thousand years or less and with magnetic fields $10^{14}-10^{15}$~G.  Unfortunately, it is difficult to use magnetars for testing models of magnetic field decay on a time scale $\\gtrsim 10^6$~yrs because they become extremely elusive and avoid detection at ages more than about a hundred thousand years.  For ``standard'' values of the initial magnetic fields ($10^{11}$ -- $10^{13}$~G) typical for radio pulsars the field decay in these models is insignificant: the magnetic field diminishes just by a factor of 2. Therefore, the population of radio pulsars is ill-suited to test these model predictions.  \\citet{Popov10} carried out a comprehensive population synthesis of young NSs (radio pulsars, magnetars, near-by cooling NSs) and showed that a model with magnetic field decay can explain the basic properties of all the three populations considered using the same initial magnetic field distribution.  All types of these objects have ages less than $\\sim$ million years.  It is important to find a way to confront model predictions and observations on a longer time scale. NSs in some HMXBs can be suitable for testing the evolution of magnetic fields on\t the time scales of several million years. Compact objects in these systems have average ages $\\sim$~a few million years (their massive companions usually have a total lifetime  less than 10-30 million years).  In the case of wind-fed accretion \\citep{Negueruela10}, the accretion rate is not large enough to affect significantly the evolution of the field (as opposed to systems with intense disc accretion which produce millisecond pulsars).  Thus, by estimating magnetic fields of NSs in HMXBs with accretion from a wind, it is possible to study the magnetic field evolution on the time scale of a few million years.  The most numerous homogeneous class of HMXBs are Be/X-ray binaries \\citep{Liu06, Reig11}. Now there are about 70 objects of this type known in the Galaxy, 78 in the Small Magellanic Cloud (SMC), and 16 in the Large Magellanic Cloud \\citep{Raguzova2005}.  We consider 40 Be/X-ray binaries in the SMC to estimate magnetic fields of NSs by means of different methods.  The article is structured as follows. In Sec. 2 we describe the model of the  magnetic field decay that we compare our estimates with. In Sec. 3 we present different methods to obtain estimates of  magnetic fields of  accreting NSs. In Sec.  4 we present the observational data used in the paper. Sec.  5 contains our results, and in Sec.  6 we discuss them. ",
        "conclusions": "In Fig. 4 we show magnetic field distributions. Here the theoretical distribution for the initial polar magnetic field $B_0$ is generated randomly for $10^5$ NSs from a log-normal distributon with the mean value $<\\log{B_0/[G]}>=13.25$ and standard deviation $\\sigma_{\\log{B_0}}=0.6$ \\citep{Popov10}. Note, that formally this initial field distribution is not absolutely self-consistent with our model of field decay, because we use parameters slightly different from those used in numerical calculations by \\citep{Popov10}. However, this imperfection cannot significantly influence our conclusions, because the main assumption is related to the asymptotic field value.  The evolved fields  are calculated  according to eq. (\\ref{E:Bevol}) using the values: $\\tau _{Hall}=10^4(B_{0}/10^{15}\\rm{G})^{-1}$~yrs and $\\tau _{Ohm}=10^6$~yrs. The  NS ages are chosen uniformly in the interval $[0, t_{max}]$, where $t_{max}=10^7 \\rm$~yrs is the maximum life time of Be/X-ray binaries. Magnetic field estimates obtained in this paper are shown in four different panels in Fig. \\ref{fig:mfe}. In the  panel A with a solid line we show the  magnetic field distribution based on the assumption of the onset of pulsations when $R_A=R_{co}$. The magnetic field distribution according to the model by Bisnovatyi-Kogan is plotted in the panel B. Distributions for the models by Illarionov, Kompaneets and the hybrid model are shown in the panels C and D, correspondingly. Theoretical models for the initial and the evolved magnetic fields are shown in all four panels by dotted and dot-dashed lines, correspondingly.  \\begin{figure} \\includegraphics[width=200pt]{hist1} \\includegraphics[width=200pt]{hist5} \\includegraphics[width=200pt]{hist2} \\includegraphics[width=200pt]{hist3} \\caption{Magnetic field distributions for different models (solid lines). Panel A -- onset of pulsations when $R_A=R_{co}$. Panel B -- the model by Bisnovatyi-Kogan. Panel C -- the model by Illarionov and Kompaneets. Panel D -- the hybrid model. Theoretical models for the initial and evolved magnetic fields are shown in all the four panels by dashed and dash-dotted lines, respectively. All distributions are normalized to the total number of observed objects. Error bars are given according to the Poissonian values. } \\label{fig:mfe} \\end{figure}   To compare the theoretical predictions with results based on  the hybrid model we use the Kolmogorov-Smirnov test. According to this test, the hypothesis of the same distribution of the theoretical and estimated values cannot be rejected at 80\\% significance level. However,  there is a small number of objects for the hybrid model distribution in the bin $10^{13}<B<10^{13.5}$~G (Fig. \\ref{fig:mfe}D). Discrepancies can be related to uncertainty in evolution of the magnetic field for $t>10^6$~yrs. Probably, the magnetic field decay continues but more slowly than before.  Likewise, different highly magnetized objects can evolve to different values of the asymptotic magnetic field (see Sec. 2).   In this study we assumed that NSs in binaries have the same initial magnetic distribution as isolated compact objects. This is not so obvious in the case of magnetars, as in the stardard scenario their magnetic fields are generated via the dynamo mechanism \\citep{dt1992}. The dynamo mechanism is effective in the case of very rapid rotation of a protoNS. Isolated massive stars hardly can have rapidly rotating cores. Rapid rotation can be reached in close binary systems due to accretion or/and tidal synchronization. On the other hand, all known magnetars are isolated objects, and authors studied evolutionary channels in binary systems which result is formation of isolated compact objects with rapid initial rotation \\citep{bp2009, pp2006}. In these models objects with extremely large field tend to form after a binary is destroyed to explain that all known magnetars are isolated.  To test the possibility of absence of extreme magnetars in binary systems for the case of the hybrid model (which is the most successful in explaining the properties of the population under study), we plot an additional histogram (Fig. \\ref{fig:mfe2}) for a modified initial magnetic field distribution: no NSs with $B_0>10^{14}$~G are formed. As we see, estimates made on the basis of observational data are even in better agreement with the prediction for the evolved field in the model by \\citet{Pons09} with the cutoff at $B_0=10^{14}$~G, as the discrepancy in the bin $10^{13}<B<10^{13.5}$~G is decreased. However, we note that our calculations are based on an approximation to the actual detailed model by \\citet{Pons09}, and for the specific choice of parameters, which can vary ($\\tau_{Hall},\\tau_{Ohm}, B_{min}$). We do not overestimate the accuracy of our results, and so do not make further speculations.  \\begin{figure} \\includegraphics[width=350pt]{hist4} \\caption{Magnetic field distributions. The distribution based on observational data after application of the hybrid model is shown with a solid line. Poissonian error bars are indicated. Results for the theoretical model for evolved magnetics fields is shown with a dash-dotted line. Finally, a modification of the theoretical model for which a cut-off at $10^{14}$~G for initial fields was used is shown with a dashed histogram. All distributions are normalized to the total number of observed objects.}  \\label{fig:mfe2} \\end{figure}  An additional consistency check for magnetic field estimates can be made using the Corbet diagram \\citep{Corbet84}. We apply this test to the hybrid model. This diagram relates spin periods and orbital periods of NSs. We simulate this distribution. Orbital periods of NSs are chosen uniformly in the interval [10, 1000] days. Using an orbital period and evolved magnetic field we can estimate the spin period of a NS. We calculate the evolved magnetic field as discussed above. These values of spin and orbital periods are plotted in Fig. 6 as dots. We plot the observational data \\citep{Reig11} as diamonds in this figure. One can see good agreement between theoretical and observational data. According to the non-parametric Kolmogorov-Smirnov test, the hypothesis that both subsets arise from the same distribution cannot be rejected at 77\\% significance level.   The methods that we use to estimate  magnetic fields have limitations, especially the most simple ones. For instance, the hypothesis of the equilibrium period may not be valid for some systems.  In our estimates we use $\\dot M$ obtained for the moment when pulsations are detected for the first time. But for NSs in systems with high eccentricity which spend much time at the propeller stage, it can take some time to reach the equilibrium. So, their observed periods are larger than the equilibrium value. This results in underestimating of the magnetic field.  Likewise, we should consider that both disc and wind accretion estimates (see, Fig. \\ref{fig:M1}) cannot be correct at the same time because just one regime is realized in any given source. On average, estimates for disc accretion are more realistic for shorter spin periods. Still, we think that estimates based on old models are not realistic in general, as they produce too large field values for long period sources.  In the model by \\citet{Shakura11} there is a critical value of $\\dot M$. For accretion rates above this $\\dot M_{cr}$ an envelope around a compact object can cool rapidly. This results in an enhanced accretion rate, and no equilibrium is possible in this regime. This critical value depends on several parameters and can be different in different systems, but roughly $\\dot M_{cr}\\lesssim 10^{17}$~g/s for magnetic fields $\\gtrsim 10^{13}$~G. Marginally, this condition is fulfiled for our systems. We can slightly overestimate the accretion rate, however, if the model by \\citet{Shakura11} is applicable, then the magnetic field estimate depends on the accretion rate only mildly.  It is instructive to compare the results of indirect magnetic field estimates with direct magnetic field measurements through cyclotron lines, as it was made for the long-period object GX 301-2. For this object, magnetic field strengh is estimated by \\citet{LaBarbera2005} as  $\\sim 5\\times 10^{12}$~G. This value is consistent  with the magnetic field estimate made by \\citet{Shakura11} using their model. On the other hand, the classical models used by \\citet{Doroshenko10} provide the magnetic field estimate nearly two orders of magnitude larger than it is given by the cyclotron line measurements.  However, the cyclotron line can be formed high above the NS surface that allows the measured magnetic field strength to be considerably weaker than that on the surface \\citep{Doroshenko10}.  \\begin{figure} \\includegraphics[width=400pt]{corbet} \\caption{Corbet diagram for the observed data (open diamonds) and for the hybrid model (dots).} \\label{fig:mfe3} \\end{figure}    Finally, we comment on a possible evolution of the angle between spin and magnetic axis. For radio pulsars it is possible to describe the angular momentum evolution using the magneto-dipole formula \\citep{Pacini67}: \\begin{equation} \\label{E:angle} \\frac{dI \\omega}{dt}=-\\frac{1}{6}\\frac{B^2\\omega ^3 R^6}{c^3}\\rm{sin}^2\\beta, \\end{equation} Here $\\beta$ is the angle between spin and magnetic axis. From this formula one can estimate the effective magnetic field $B_{eff}=B \\rm{sin}\\beta$. So, the effect of the variation of the effective magnetic field can be related not only to the magnetic field decay, but also to evolution of the angle between magnetic and spin axis. Therefore, it is difficult to distinguish between these effects. However, there are different methods to estimate magnetic field of magnetars and to demonstrate that this is really the field decay, not the variation of the angle $\\beta$. Note, that our methods are only weakly sensitive to the angle between spin and magnetic axis, and we estimate the absolute value of magnetic field.  Still, using Be/X-ray systems we can try to study evolution of the angle between these axis. According to two most popular models (magneto-dipole and current losses), the angle evolves very fast on the time scale of spin period evolution. The angle $\\beta$ evolves to the value $\\beta=0^{\\circ}$ in the magneto-dipole model and to the value $\\beta=90^{\\circ}$ in the current losses model \\citep{Beskin93}.  But this effect is not observed for PSRs. It is possible to use Be/X-ray binaries to study evolution of the angle on the time scale up to $10^7$~yrs jointly with the magnetic field decay. We plan to use this approach in future studies."
    },
    "1112/1112.1409_arXiv.txt": {
        "abstract": "We present the results of a \\emph{Chandra} observation of the galaxy cluster Abell 689 (z=0.279).  Abell 689 is one of the most luminous clusters detected in the ROSAT All Sky Survey (RASS), but was flagged as possibly including significant point source contamination.  The small PSF of the \\emph{Chandra} telescope allows us to confirm this and separate the point source from the extended cluster X-ray emission. For the cluster we  determine a bolometric luminosity of L$_{\\rm bol}$ = (3.3$\\pm$0.3)$\\times$10$^{\\rm 44}$ erg s$^{\\rm -1}$ and a temperature of kT=5.1$^{+2.2}_{-1.3}$ keV when including a physically motivated background model.  We compare our measured luminosity for A689 to that quoted in the Rosat All Sky Survey (RASS) and find L$_{\\rm 0.1 - 2.4, keV}$ = 2.8$\\times$10$^{44}$ erg s$^{-1}$, a value $\\sim$10 times lower than the ROSAT measurement.  Our analysis of the point source shows evidence for  significant pileup, with a pile-up fraction of $\\simeq$60\\%. SDSS spectra and HST images lead us to the conclusion that the point source within Abell 689 is a BL Lac object. Using radio and optical observations  from the VLA and HST archives, we determine $\\alpha_{\\rm ro}$=0.50, $\\alpha_{\\rm ox}$=0.77 and $\\alpha_{\\rm rx}$=0.58 for the BL Lac, which would classify it as being of `High-energy peak BL Lac' (HBL) type.  Spectra extracted of A689 show a hard X-ray excess at energies above 6 keV that we interpret as inverse Compton emission from aged electrons that may have been transported into the cluster from the BL Lac. ",
        "introduction": "\\label{intro} Studies of clusters of galaxies, including measurements of their number density and growth from the highest density perturbations in the early Universe, offer insight into the underlying cosmology \\citep{2003ApJ...590...15V,2004MNRAS.353..457A}.  However, in order to use clusters as a cosmological probe three essential tools are required \\citep{2010A&A...514A..80D}: (a) an efficient method to find clusters over a wide redshift range, (b) an observational method of determining the cluster mass, and (c) a method to compute the selection function or the survey volume in which clusters are found. These requirements are met by large surveys with well understood selection criteria. Arguably the most effective method of building large, well defined cluster samples has been via X-ray selection. The high X-ray luminosities of clusters make it relatively easy to detect and study clusters out to high redshifts.  Many cluster samples have been constructed based upon large X-ray surveys such as the Einstein Medium Sensitivity Survey \\citep[EMSS;][]{1990ApJS...72..567G} and the ROSAT All Sky Survey \\citep[RASS;][]{1992eocm.rept....9V}.  However due to the relatively poor angular resolution of these X-ray observatories, observations of clusters were susceptible to point source contamination.  Indeed, within the ROSAT Brightest Cluster Sample \\citep[BCS;][]{1998MNRAS.301..881E} and its low-flux extension \\citep[eBCS;][]{2000MNRAS.318..333E}, 9 out of 201 clusters and 8 out of 99 clusters respectively were flagged as probably having a significant fraction of the quoted flux from embedded point sources. Being able to resolve these point sources is of crucial importance for the reliable estimation of cluster properties, and indeed the nature of the point source contamination is of independent interest.  The study of galaxy clusters has been transformed with the launch of powerful X-ray telescopes such as \\emph{Chandra} and XMM Newton, which have allowed the study of the X-ray emitting intracluster medium (ICM) with unprecedented detail and accuracy.  With the launch of this new generation of X-ray telescopes, we are able to uncover interesting features in the morphologies of individual clusters.  In particular, \\emph{Chandra's} high angular resolution provides the means to examine individual cluster features with great detail.  Abell 689 \\citep[hereafter A689;][]{1958ApJS....3..211A}, at z=0.279 \\citep{1995MNRAS.274.1071C}, was detected in the RASS in an accumulated exposure time of 317s.  It is included in the BCS, with a measured X-ray luminosity of 3.0$\\times$10$^{45}$ erg s$^{-1}$ in the 0.1--2.4 keV band.  This luminosity is the third highest in the BCS, and thus A689 meets the selection criteria for various highly X-ray luminous cluster samples \\citep[e.g.][]{2002ApJS..139..313D}.  However this cluster was noted as having possible point source contamination, and for this reason has often been rejected from many flux limited samples.  In this study we present results of \\emph{Chandra} observations designed to separate any point sources and determine uncontaminated cluster properties.  The outline of this paper is as follows.  In $\\S$~\\ref{reduction} we discuss the observation and data analysis.  Results of the X-ray cluster analysis is presented in $\\S$~\\ref{xray analysis}. In $\\S$~\\ref{point source} we present our analysis of the X-ray point source through X-ray, optical and radio observations.  We interpret our results in $\\S$~\\ref{discussion} and the conclusions are presented in $\\S$~\\ref{conclusions}.  Throughout this paper we adopt a cosmology with $\\Omega_M$= 0.3, $\\Omega_{\\Lambda}$ = 0.7 and H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, so that 1$^{\\prime\\prime}$ corresponds to 4.22 kpc at the redshift of A689.  We define spectral index, $\\alpha$, in the sense S $\\propto$ $\\nu^{-\\alpha}$. ",
        "conclusions": "\\label{conclusions} We have used a 14 ks \\emph{Chandra} observation of the galaxy cluster A689 in order to determine the nature of the cluster's point source contamination and to analyze the cluster properties excluding the central point source.  Our main conclusions are as follows. \\newcounter{qcounter} \\begin{list}{ \\arabic{qcounter}.~}{\\usecounter{qcounter}} \\item Using background subtraction of both local and blank-sky backgrounds, we obtain temperatures which are high relative to the luminosity-temperature relation. \\item We construct a physically motivated model for the background and include this model in a fit to the cluster spectrum.  If the particle background in the cluster is allowed to exceed that in the local and blank-sky backgrounds we obtain a temperature of 5.1$^{+2.2}_{-1.3}$ keV.  However, there is no reason for there to be a higher particle rate in the specific region of the CCD in which the cluster lies.  A hard excess needed to bring the temperature to a reasonable value must have a different origin. \\item We confirm the presence of a point source within A689 as suspected in the BCS.  When excluding the point source and using our derive background model we find a luminosity of L$_{\\rm 0.1 - 2.4, keV}$ = 2.8$\\times$10$^{44}$ erg s$^{-1}$, a value $\\sim$10 times lower than quoted in the BCS. \\item From the X-ray analysis of the point source we find a ``flat-topped'' point source with a pileup fraction of $\\simeq$ 60\\%. \\item Optical observations of the cluster from SDSS and HST lead us to conclude that the point source is a BL Lac type AGN. \\item We classify the BL Lac as an `High-energy peak BL Lac' with $\\alpha$$_{rx}$=0.58$\\pm$0.04. \\item We interpret the hard X-ray excess needed to bring the cluster temperature to a reasonable value as inverse-Compton emission from aged electrons that may have been transported into the cluster from the BL Lac. \\end{list}  We have shown here not only the importance of resolving and excluding point sources in cluster observations, but also the effect these point sources can have when determining the ICM properties of galaxy clusters. The detailed analysis we have performed here may not however be suitable for all clusters as it is unclear whether this analysis can be performed at higher redshifts.  Separating the point source and cluster emissions becomes increasingly difficult at high redshifts, however Chandra has proved capable at resolving point source emission in clusters to z$\\sim$2 \\citep[e.g.][]{2009A&A...507..147A}. Resolving point source and cluster emission out to high redshift becomes more important with redshift due to the evolution of the number density of point sources within clusters \\citep{2009ApJ...694.1309G}.  The area used to model the background components associated with the high energy cluster regions will decrease spatially with redshift, and separating the between thermal and inverse Compton emissions at higher redshift will become increasingly difficult \\citep{2003MNRAS.341..729F}, due to the increase in the energy density of the CMB with redshift."
    },
    "1112/1112.3917_arXiv.txt": {
        "abstract": "{ We report on the detection of very-high energy (VHE, $E>100$~GeV) $\\gamma$-ray emission from \\mbox{\\object{NGC 1275}}, the central radio galaxy of the Perseus cluster of galaxies. The source has been detected by the MAGIC telescopes with a statistical significance of $6.6~\\sigma$ above $100$~GeV in $46$~hr of stereo observations carried out between August 2010 and February 2011. The measured differential energy spectrum between $70$~GeV and $500$~GeV can be described by a power law with a steep spectral index of $\\Gamma=-4.1\\pm0.7_{stat}\\pm0.3_{syst}$, and the average flux above $100$~GeV is $F_{\\gamma}=(1.3\\pm0.2_{stat}\\pm0.3_{syst})\\times10^{-11}~\\mathrm{cm^{-2}~s^{-1}}$. These results, combined with the power-law spectrum measured in the first two years of observations by the \\emph{Fermi}--LAT above $100$~MeV, with a spectral index of $\\Gamma \\simeq -2.1$, strongly suggest the presence of a break or cut-off around tens of GeV in the \\mbox{\\object{NGC 1275}} spectrum. The light curve of the source above $100$~GeV does not show hints of variability on a month time scale. Finally, we report on the nondetection in the present data of the radio galaxy \\mbox{\\object{IC 310}}, previously discovered by the \\emph{Fermi}--LAT and MAGIC. The derived flux upper limit $F^{U.L.}_{\\gamma} (>300~\\mathrm{GeV})=1.2\\times10^{-12}~\\mathrm{cm^{-2}~s^{-1}}$ is a factor $\\sim 3$ lower than the mean flux measured by MAGIC between October~2009 and February~2010, thus confirming the year time-scale variability of the source at VHE. } ",
        "introduction": "\\label{sec:section1} \\mbox{\\object{NGC 1275}} ($z=0.0179$), the central dominant galaxy of the Perseus cluster, harbors one of the closest active galactic nuclei (AGN), already included in the original Seyfert list~\\citep{se43}. The AGN is a very bright radio source showing an extended jet with Fanaroff-Riley~I morphology (e.g.~\\citealp{ve94,bu10}). The optical emission of the nucleus is variable and strongly polarized from 3$\\%$ to 6$\\%$~\\citep{ma79,ma83}, implying that the relativistic jet contributes significantly to the optical continuum ~\\citep{an80}. \\begingroup \\let\\thefootnote\\relax\\footnotetext{* Corresponding authors: S. Lombardi (saverio.lombar\\-di@pd.infn.it), P. Colin (colin@mppmu.mpg.de), D. Hildebrand (dorothee.hildebrand@phys.ethz.ch), and F. Zandanel (fabio@iaa.es)} \\endgroup The source has also been classified as a BL~Lac object~\\citep{ver78}. However, the jet increases its inclination from $10^\\circ$ to $20^\\circ$ on milliarcsecond scales up to $40^\\circ$ to $60^\\circ$ at arcsecond scales~\\citep{kr92}. Due to its brightness and proximity this source is ideally suited to study the physics of relativistic outflows and the ``feedback'' effects of the jet on the cluster environment (e.g.~\\citealp{fa08,gal09}).  In fact, \\mbox{\\object{NGC 1275}} is one of the closest $\\gamma$-ray emitting AGN. It was first unambiguously detected in the high-energy (HE, $100$~MeV~$<E<$~$100$~GeV) $\\gamma$-ray range by the \\emph{Fermi} Large Area Telescope (LAT)~\\citep{Fermi_1}, during the first four months of \\emph{all-sky-survey} observations, with an average flux above $100$~MeV of $F_{\\gamma}=(2.10\\pm0.23)\\times10^{-7}~\\mathrm{cm^{-2}~s^{-1}}$. The differential energy spectrum  between $100$~MeV and $25$~GeV was described well by a power law with a spectral index of $\\Gamma$~=~$-2.17\\pm0.05$. While no variability was observed during these four months of observations, subsequent results based on the first year of \\emph{Fermi}--LAT observations~\\citep{Fermi_2} show evidence of flux variability on time scales of months. Furthermore, the average $\\gamma$-ray spectrum show a significant deviation from a simple power law, indicating an exponential cut-off at the break photon energy of $E_{c}$~=~$(42.2\\pm19.6)$~GeV.  More recently, the results obtained from the first two years of \\emph{Fermi}--LAT observations~\\citep{Fermi_3} have given clear evidence for variability on time scales of days above $800$~MeV, revealing that several major flaring events occurred during the two-year observation period. A harder-when-brighter correlation between flux and spectral index was also found. Brighter and therefore harder $>$~GeV states are then promising for triggering observations at very high energy (VHE, $E>100$~GeV). Finally, present upper limits at VHE provided by MAGIC-I~\\citep{al10a} and VERITAS~\\citep{acc09} combined with the \\emph{Fermi}--LAT results mentioned above suggested that \\mbox{\\object{NGC 1275}} may have a break or cut-off in the spectrum around tens of GeV.  The Perseus galaxy cluster contains another $\\gamma$-ray source, the radio galaxy \\mbox{\\object{IC 310}}, located at $\\sim 0.6^{\\circ}$ from \\mbox{\\object{NGC 1275}}. It was discovered in 2010 by the \\emph{Fermi}--LAT in the $30$~GeV~--~$300$~GeV energy range~\\citep{nsv10} and for energies $>260$~GeV by MAGIC~\\citep{al10b}. The combined MAGIC and \\emph{Fermi}--LAT spectrum is consistent with a flat spectral energy distribution (SED) stretching without a break over more than three orders of magnitude in energy ($2$~GeV~--~$7$~TeV). The spectrum at VHE measured by MAGIC has a spectral index of $\\Gamma$~=~$-2.00\\pm0.14$, and the mean flux above $300$~GeV, from October 2009 to February 2010, was $F_{\\gamma}=(3.1\\pm0.5)\\times10^{-12}~\\mathrm{cm^{-2}~s^{-1}}$. Strong hints of a week to a year time-scale variability were seen in the MAGIC data.  In this letter we present the results of the observations of \\mbox{\\object{NGC 1275}} at VHE performed with the MAGIC telescopes between August 2010 and February 2011, which resulted in the first detection of the source above $100$~GeV. The same observational campaign also provided results on the variability at VHE of \\mbox{\\object{IC 310}}. This letter is accompanied by a separate paper dedicated to the study of the Perseus cluster environment, focusing on possible VHE $\\gamma$-ray emission induced by cosmic~rays~\\citep{al11b}. The multiwavelength emission of \\mbox{\\object{NGC 1275}} from radio to VHE will be addressed in future work. ",
        "conclusions": "\\label{sec:section4} The MAGIC telescopes have detected VHE $\\gamma$-ray emission from \\mbox{\\object{NGC 1275}}, the central radio galaxy in the Perseus cluster, at a statistical significance of $6.6~\\sigma$ from observations performed between August 2010 and February 2011. The corresponding average flux above $100$~GeV is $F_{\\gamma}=(1.3\\pm0.2_{stat}\\pm0.3_{syst})\\times10^{-11}~\\mathrm{cm^{-2}~s^{-1}}$. % This is the fourth nearby radio galaxy detected at VHE, after \\mbox{\\object{M 87}}~\\citep{ah03,ah06,al08_M87,acc08}, \\mbox{\\object{Cen A}}~\\citep{ah09}, and \\mbox{\\object{IC 310}}~\\citep{al10b}. The MAGIC observation yields a spectrum that can be fitted between $70$~GeV and $500$~GeV by a simple power law with a spectral index of $\\Gamma=-4.1\\pm0.7_{stat}\\pm0.3_{syst}$. This result, combined with previous \\emph{Fermi}-\u2013LAT results~\\citep{Fermi_1,Fermi_2,Fermi_3}, showing a power-law spectrum with a spectral index of $\\Gamma \\simeq -2.1$ above $100$~MeV, strongly suggests the presence of a break or cut-off around tens of GeV in the \\mbox{\\object{NGC 1275}} spectrum. No evidence of variability on month time scale has been found above $100$~GeV. Finally, the variability on a year time scale of the source \\mbox{\\object{IC 310}}~\\citep{nsv10,al10b} has been confirmed. The upper limit above $300$~GeV presented here is in fact about a factor 3 lower than the flux measured by MAGIC between October 2009 and February 2010."
    },
    "1112/1112.4640_arXiv.txt": {
        "abstract": "Low-luminosity active galactic nuclei (LLAGNs) represent the bulk of the AGN population in the present-day universe and they trace the low-level accreting supermassive black holes. In order to probe the accretion and jet physical properties in LLAGNs as a class, we model the broadband radio to X-rays spectral energy distributions (SEDs) of 21 LLAGNs in low-ionization nuclear emission-line regions (LINERs) with a coupled accretion-jet model. The accretion flow is modeled as an inner ADAF outside of which there is a truncated standard thin disk. We find that the radio emission is severely underpredicted by ADAF models and is explained by the relativistic jet. The origin of the X-ray radiation in most sources can be explained by three distinct scenarios: the X-rays can be dominated by emission from the ADAF, or the jet, or the X-rays can arise from a jet-ADAF combination in which both components contribute to the emission with similar importance. For 3 objects both the jet and ADAF fit equally well the X-ray spectrum and can be the dominant source of X-rays whereas for 11 LLAGNs a jet-dominated model accounts better than the ADAF-dominated model for the data. The individual and average SED models that we computed can be useful for different studies of the nuclear emission of LLAGNs. From the model fits, we estimate important parameters of the central engine powering LLAGNs in LINERs, such as the mass accretion rate and the mass-loss rate in the jet and the jet power - relevant for studies of the kinetic feedback from jets. ",
        "introduction": "One of the central paradigms in extragalactic astronomy is that today's galaxies host supermassive black holes (SMBHs) at their centers and have a symbiotic evolution with them \\citep{magorrian98, ferrarese05}. The central black holes seeded in proto-galaxies at high redshift grow in mass during cosmic history through a sequence of mergers of massive black holes and accretion episodes (e.g., \\citealt{volonteri05,fani11}). The massive black holes grow in such a way that they co-evolve with the galactic bulges \\citep{kor95, ferrarese00, geb00, ferrarese05} but do not seem to do so with the galaxy disks \\citep{kor11a} or dark matter haloes \\citep{kor11b}.  The growth of SMBHs is dominated by the quasar phase \\citep{marconi04,mclure04,sijacki07,merloni08} during which they accrete from their large gas reservoir via geometrically thin accretion disks which are radiatively efficient \\citep{shakura73,novikov73} as suggested by different lines of evidence \\citep{shields78,koratkar99,shang05}. The quasar population has a strong cosmological evolution with its density decreasing by a factor of $\\sim 100$ from $z \\sim 2$ -- where the density reaches its peak -- to the present time (e.g., \\citealt{osmer04,silverman05}).  In the present-day universe, SMBHs are underfed compared to the ones at high $z$ and are ``sleeping''. Most of SMBH activity at low $z$ is dominated by the weak end of the AGN luminosity function in the form of low-luminosity AGNs (LLAGNs; \\citealt{ho95,ho97,nagar05,ho08}). Though the LLAGN phase dominates the time evolution of SMBHs, it contributes little to their mass growth \\citep{hopkins06b,sijacki07,merloni08,xu10}.  The bulk of the LLAGN population ($\\approx 2/3$; \\citealt{ho08,ho09}) reside in low-ionization nuclear emission-line regions (LINERs; \\citealt{heckman80,ho97apjs}). LLAGNs are extremely sub-Eddington systems which are many orders of magnitude less luminous than quasars, with average bolometric luminosities $\\left \\langle L_{\\rm bol} \\right \\rangle \\sim 10^{40}-10^{41} \\ {\\rm erg \\ s}^{-1}$ and an average Eddington ratio of $L_{\\rm bol}/L_{\\rm Edd} \\sim 10^{-5}$ \\citep{ho09} where $L_{\\rm Edd}$ is the Eddington luminosity. The observational properties of LLAGNs are quite different from those of more luminous AGNs. Regarding the broadband spectral energy distributions (SEDs), LLAGNs seem not to have the thermal continuum prominence in the ultraviolet (UV) -- the ``big blue bump'' -- which is one of the signatures of the presence of an optically thick, geometrically thin accretion disk \\citep{ho99, nemmen06, wu07, ho08, erac10}. Regarding the emission-lines, LLAGNs typically have weak and narrow Fe K$\\alpha$ emission \\citep{terashima02} and a handful of LINERs display broad double-peaked H$\\alpha$ lines (e.g., \\citealt{sb03}); these properties of the emission-line spectrum are consistent with the absence of a thin accretion disk, or a thin accretion disk whose inner radius is truncated at $\\gtrsim 100 GM/c^2$ \\citep{chen89a, chen89b}. Last but not least, with the typical fuel supply of hot diffuse gas (via Bondi accretion) and cold dense gas (via stellar mass loss) available in nearby galaxies, LLAGNs would be expected to produce much higher luminosities than observed on the assumption of standard thin disks with a $10\\%$ radiative efficiency \\citep{ho09}. Taken together, this set of observational properties favors the scenario in which the accretion flow in LLAGNs is advection-dominated or radiatively inefficient.  Advection-dominated accretion flows (ADAFs\\footnote{In this paper, we consider ADAFs and radiatively inefficient accretion flows (RIAFs) to be the same kind of accretion flow solution. For a clarification regarding the terminology, see \\citealt{yu11}.}; for reviews see \\citealt{nar98,yuan07,nar08}) are very hot, geometrically thick, optically thin flows which are typified by low radiative efficiencies ($L \\ll 0.1 \\dot{M} c^2$) and occur at low accretion rates ($\\dot{M} \\lesssim 0.01 \\dot{M}_{\\rm Edd}$). SMBHs are thought to spend most of their lives in the ADAF state \\citep{hopkins06b,xu10}, the best studied individual case being Sgr A* (e.g., \\citealt{yuan07}).  In many LLAGNs, another component in the accretion flow besides the ADAF is required in order to account for different observations, including a ``red bump'' in the SEDs \\citep{lasota96,quat99,nemmen06,yu11} and as mentioned before the double-peaked Balmer emission lines (e.g., \\citealt{sb03,erac09}): the emission from a thin accretion disk whose inner radius is truncated at the outer radius of the ADAF. In many LLAGNs the accretion flow may begin as a standard thin disk but somehow at a certain transition radius the accretion flow abruptly switches from a cold to a hot ADAF mode. The details of how this transition might happen are still not well understood \\citep{manmoto00ssd,yuan04,nar08}, but it seems to be analogous to the transition between the different spectral states in black hole binary systems \\citep{remi06,done07}.  \\citet{maoz07} challenged the scenario of the central engines of LLAGNs consisting of ADAFs and truncated thin disks. Maoz argues that LLAGNs in LINERs have UV/X-ray luminosity ratios similar on average to those of brighter Seyfert 1s and based on that observation he posits that thin disks extending all the way down to the radius of the innermost stable circular orbits (ISCO) persist even for LLAGNs, despite their smaller accretion rates. \\citet{yu11} showed that the SEDs compiled by \\citet{maoz07} are naturally fitted by ADAF models with the addition of a truncated thin disk and also discussed on theoretical grounds why the ADAF model has a superior explanation capability than the pure thin disk model. The ADAF model is the only model that can naturally account for the set of observational properties of LLAGNs within a self-contained theoretical framework. Hence, it is the physical scenario adopted in this work.  ADAFs are relevant to the understanding of AGN feedback since they are quite efficient at producing powerful outflows and jets, as suggested by theoretical studies including analytical theory \\citep{ny94,meier01,nemmen07} and numerical simulations \\citep{mckinney04,hawley06,sasha10}). This is in line with the different observational studies which demonstrate that LLAGNs are generally radio-loud \\citep{ho02, ho08} and accompanied by powerful jets (e.g., \\citealt{heinz07, merloni08}). In fact, the so-called ``radio mode'' of AGN feedback invoked in semi-analytic and hydrodynamic simulations of galaxy formation (e.g., \\citealt{bower06,croton06,sijacki07,fani11}) would presumably correspond to the ADAF accretion state actively producing jets as explicitly incorporated in some works \\citep{sijacki07, oka08}.  It is clear that an understanding of the physical nature of LLAGNs in LINERs will shed light on the nature of black hole accretion, outflows and consequently black hole growth and AGN feedback in present-day galaxies. One of the best ways of exploring the astrophysics of black hole accretion and outflows is by using multiwavelength observations of black hole systems and comparing the spectra predicted by specific models with the data. The goal of this work is therefore to probe the physics of accretion and ejection in the LLAGN population, by modeling their nuclear, broadband, radio to X-rays SEDs which provide constraints on physical models for the emission of the accretion flow and the jet. Furthermore, with a large enough sample of systems, we can derive from the fits to the data the parameters that characterize the central engines and build a census of the ``astrophysical diet'' of low-state AGNs.  In this work, we will present how the currently favored model for the central engines of weakly accreting black holes is able to account for the observed SEDs of LLAGNs. We will also present the typical accretion/ejection -- or feeding/feedback -- properties of LLAGNs in LINERs as inferred from our detailed SED fits. With these parameters, we will then be able to draw a connection between the SMBHs in nearby active galaxies and the one in the Milky Way, Sgr A*. We will also present the inferred production site of the continuum emission and other properties inferred from the fits.  The structure of this paper is as follows. In Section \\ref{sec:sample} we describe the sample of 21 LLAGNs in LINERs that we used, the data and classifications. Section \\ref{sec:models} describes the physical model that we adopted in order to interpret and fit the SEDs and derive the central engine parameters. In Section \\ref{sec:seds} we describe the model fits to the SEDs, explaining in more detail the SED models for 5 objects which we consider the ``reference'' or fiducial LLAGN examples. Section \\ref{ap:seds} presents the model fits to the other 16 SEDs in our sample. In Section \\ref{sec:pars} we describe the distributions of central engine parameters resulting from our detailed model fits. We present the average SED resulting from our fits -- including the predicted emission in wavebands which can be observed with future facilities such as ALMA and the James Webb Space Telescope -- and the inferred nature of the broadband emission in each wavelength in \\textsection \\ref{sec:avgsed}. We discuss in \\ref{sec:disc} which model is favored to explain our LLAGN sample, in particular which component in the flow dominates the X-ray emission. We conclude by presenting a summary of our results in \\ref{sec:end} and comparing the SMBH ``diet'' of massive black holes in LINER AGNs and Sgr A*. ",
        "conclusions": "\\label{sec:disc}  In our AD models, the X-rays are produced predominantly by inverse Compton scattering by the ADAF of seed synchrotron photons produced in the accretion flow itself; in contrast, in our JD models, the X-rays are dominated by the optically thin synchrotron emission in the base of the jet. In this regard, the \\emph{Chandra} X-ray spectrum is one of the most important constraints for the origin of high-energy photons in LLAGNs.  There are only two objects for which the AD or the JD can be ruled out, within the observational uncertainties and the explored model parameter space: NGC 1553 (\\textsection \\ref{sec:1553}) and NGC 4552 (\\textsection \\ref{sec:4552}). For NGC 1553, the JD model does not reproduce the hardness of the X-ray spectrum whereas for NGC 4552 the AD model is not able to account for the corresponding X-ray softness. For the remaining objects, by inspecting the SED fits we can see that both the AD and the JD models are able to reproduce with varying degrees of success the whole SED from radio to X-rays, in particular the X-ray portion.  In order to quantify which model systematically fits better the broadband SEDs we use as a goodness-of-fit parameter a modified $\\chi^2$. We define this modified $\\chi^2$ simply as \\begin{equation} \\chi^2 \\equiv \\sum_i \\left( y_{\\rm obs} - y_{\\rm model} \\right)^2 \\end{equation} where $y=\\log \\nu L_\\nu$ and the sum is carried over the observed data points. We sample the X-ray continuum using 3 data points from 0.5 keV to 10 keV and we do not take into account upper limit data points. The values of $\\chi^2$ for each model are displayed in Table \\ref{tab:models}.  Considering the 20 objects for which we were able to obtain both JD and AD fits, we find that $\\chi^2 (JD) < \\chi^2 (AD)$ for 9 objects and $\\chi^2 (JD) > \\chi^2 (AD)$ for the remaining 8. In other words, for 9 objects (10 including NGC 4552 for which we did not find an appropriate AD fit) the JD results in better SED fits as opposed to 8 for which the AD model is a better fit. The results from this goodness-of-fit test show that 52\\% of the SEDs in our sample are better fitted by the JD model whereas 38\\% are better reproduced by the AD scenario. For the 3 objects with $\\chi^2 (JD)/\\chi^2 (AD) \\approx 1$ (NGC 1097, NGC 3998 and NGC 4374) it is hard to tell -- based on the broadband SED alone -- which model provides a better fit.  We should point out a number of caveats involved when considering which scenario better accounts for the nature of the X-ray emission. From an observational point of view, even though we are carrying the largest systematic analysis of the SEDs of LLAGNs in LINERs done up to date, the size of our sample is limited. We need to extend our modeling to a larger number of LLAGNs in order to increase the statistics and draw further conclusions. Secondly, for many sources we only have a few data points available to fit, so clearly for these objects we need to obtain more measurements to better constrain the models. In addition, in our modeling we find that one of the main observables that help to pinpoint the dominant source of X-rays is the hardness of the X-ray spectrum. ADAF models tend to predict harder X-ray spectra while jet models predict softer ones. In most of the LLAGNs in our sample there is a significant observational uncertainty in the X-ray hardness derived from the \\emph{Chandra} observations which make it difficult to determine which whether the AD or JD is favored based on the X-ray spectrum.  From a theoretical perspective, there are considerable theoretical uncertainties involved in the the ADAF-jet models. For instance, the values of the model parameters that affect the microphysics and dynamics of the ADAF and jet are not sufficiently constrained. In particular, $\\delta$ and $p$ have a large impact on the hardness of the X-ray spectrum predicted by ADAF/jet models but there is a considerable uncertainty in the range of possible values of these parameters despite recent progress (e.g., \\citealt{sharma07,medvedev06}). Considering the overall goodness-of-fit of the AD and JD models, it is important to point out that our jet model is basically a phenomenological one. Uncertainties affecting for example the jet model could translate in larger values of $\\chi^2$. For instance, when we find $\\chi^2 (JD) < \\chi^2 (AD)$, that could be possibly because the JD model is less well constrained than the AD one, i.e, we have more freedom in the jet fitting compared to the ADAF model.  Taking into account these caveats, it is still premature to favor one model over the other based only on the SED. The nature of the X-ray emission in LLAGNs has been debated in the last few years by several authors  favoring in some cases the AD or JD models based on the analysis of individual sources (e.g., \\citealt{falcke00,yuan02sgr,yuan02_4258,yuan03,wu07,markoff08,miller10}) or in a statistical sense based on the fundamental plane of black hole activity \\citep{merloni03,falcke04,yc05,yuan09,plotkin11}. We suggest different routes to clarify this issue in the future.  The AD and JD models should predict different characteristic radio and X-ray variability timescales (e.g., \\citealt{ptak98}), hence we should be able to test which component is dominant in X-rays by carrying out a simultaneous monitoring of the variability of radio and X-rays. By comparing the variability pattern predicted by jet/ADAF models we could pinpoint the X-ray dominant component. This promising strategy is quite similar to what the observational campaigns carried out for M81 \\citep{markoff08,miller10} and the radio-quiet Seyfert 1 NGC 4051 \\citep{king11} and could be applied to many other LLAGNs. We note that in the case of M81, \\cite{miller10} favor a scenario in which the X-rays originate in the transition region between the thin disk and the ADAF; NGC 4051 is a different case since it is a relatively bright AGN accreting at the $\\sim 0.05 L_{\\rm Edd}$ level \\citep{peterson04} and its jet-disk coupling is quite distinct compared to LLAGNs.  The \\emph{Fermi} Gamma-ray Space Telescope can be quite helpful in this regard. LLAGNs are potential sources of $\\gamma$-rays \\citep{takami11} and in fact one of the source in this sample (NGC 4486) has already been detected by the \\emph{Fermi} LAT \\citep{m87fermi}. By using the model parameters that we derived in our radio-to-X-rays fits we should be able to predict the $\\gamma$-ray spectrum \\citep{mahadevan97,takami11} and compare with future \\emph{Fermi} detections. In this way we should be able to compare AD/JD predictions and through $\\gamma$-ray observations place further constraints on the production site responsible for the high-energy emission and the jet-disk connection in LLAGNs.  Observations in the mm and sub-mm are also very helpful because they constrain the ADAF synchrotron emission and hence also the X-ray emission. This follows because the synchrotron photons in the ADAF are inverse-Compton-scattered to X-rays.  Finally, further refinements and a better understanding of the fundamental plane of black hole activity \\citep{merloni03,falcke04,gultekin09,yuan09,plotkin11} should make the fundamental plane a better tool for constraining the radiative processes shaping the radio and X-ray emission in sub-Eddington black hole source and by consequence LLAGNs."
    },
    "1112/1112.2857_arXiv.txt": {
        "abstract": "{Recent studies carried out with SOHO and Hinode high-resolution spectrometers have shown that the plasma in the off-disk solar corona is close to isothermal. If confirmed, these findings may have significant consequences for theoretical models of coronal heating. However, these studies have been carried out with diagnostic techniques whose ability to reconstruct the plasma distribution with temperature has not been thoroughly tested.} {In this paper, we carry out tests on the Monte Carlo Markov Chain (MCMC) technique with the aim of determining: 1) its ability to retrieve isothermal plasmas from a set of spectral line intensities, with and without random noise; 2) to what extent can it discriminate between an isothermal solution and a narrow multithermal distribution; and 3) how well it can detect multiple isothermal components along the line of sight. We also test the effects of 4) atomic data uncertainties on the results, and 5) the number of ions whose lines are available for the DEM reconstruction.} {We first use the CHIANTI database to calculate synthetic spectra from different thermal distributions: single isothermal plasmas, multithermal plasmas made of multiple isothermal components, and multithermal plasmas with a Gaussian DEM distribution with variable width. We then apply the MCMC technique on each of these synthetic spectra, so that the ability of the MCMC technique at reconstructing the original thermal distribution can be evaluated. Next, we add a random noise to the synthetic spectra, and repeat the exercise, in order to determine the effects of random errors on the results. We also we repeat the exercise using a different set of atomic data from those used to calculate synthetic line intensities, to understand the robustness of the results against atomic physics uncertainties. The size of the temperature bin of the MCMC reconstruction is varied in all cases, in order to determine the optimal width.} {We find that the MCMC technique is unable to retrieve isothermal plasmas to better than $\\Delta \\log T \\simeq 0.05$. Also, the DEM curves obtained using lines calculated with an isothermal plasma and with a Gaussian distribution with FWHM of $\\log T \\simeq 0.05$ are very similar. Two near-isothermal components can be resolved if their temperature separation is $\\Delta \\log T=0.2$ or larger. Thus, DEM diagnostics has an intrinsic resolving power of $\\log T=0.05$. Atomic data uncertainties may significantly affect both temperature and peak DEM values, but do not alter our conclusions. The availability of small sets of lines also does not worsen the performance of the MCMC technique, provided these lines are formed in a wide temperature range.} {Our analysis shows the present limitations in our ability to identify the presence of strictly isothermal plasmas in stellar and solar coronal spectra.} ",
        "introduction": "Detailed measurements of the temperature distribution of the plasma of solar and stellar coronae are  of critical importance to understand how they are heated to multimillion degree temperatures, so it is no surprise that the thermal structure of the solar corona has been studied ever since the publication of the seminal paper of Edlen (1942), who discovered its high temperature for the first time. These studies have known a revival in recent years, thanks to the availability of a large body of observations obtained with the new instrumentation carried on board the SOHO, TRACE, STEREO, Hinode and SDO satellites.  Plasma loops are a fundamental component of the solar corona, since they are ubiquitous both in quiet and active solar plasmas. Steady-state models of loops predict that the temperature of these structures varies along the loop (Rosner \\etal 1978), and that loops with different length can have very different temperatures; nanoflare-based models predict finely-structured loops filled with multithermal plasma all along their length (Patsourakos \\& Klimchuk 2005). A consequence of these models is that quiet coronal plasmas observed {\\em outside the solar limb} should be multithermal because they are the sum of a multitude of loops with different lengths and temperatures along the line of sight. The question of clearly distinguishing between multithermal and isothermal plasma along the line of sight is thus very important in the framework of coronal heating models (e.g. Klimchuk 2006, Reale 2010). However, recent measurements have provided increasing evidence of a nearly isothermal nature of both coronal hole and quiet Sun plasmas when they are observed above the solar limb, contrary to model predictions. Several authors (e.g. Feldman \\etal 1998 and 1999, Warren 1999, Doschek \\etal 2001, Landi 2008) reported single-temperature determination which are remarkably similar even if taken with different instruments and at different times during the solar cycle. A constant temperature was also found to characterize an entire streamer by Landi \\etal (2006). Even a complex active region, when observed at the solar west limb, seemed to be composed by three near isothermal plasmas (Landi \\& Feldman 2008). Further studies carried out with Hinode/EIS observations, on the contrary, provided evidence that the quiescent solar corona is not {\\em strictly} isothermal, but it is characterized by a rather narrow temperature structure with a tail that extends to higher temperatures (Warren \\& Brooks 2009, Brooks \\etal 2009, Hahn \\etal 2011). The aim of the present paper is to assess how reliable such isothermal plasma claims are.  These measurements have been carried out either with the Emission Measure (EM) loci technique or with traditional Differential Emission Measured (DEM) techniques (see Feldman \\& Landi 2008 and references therein). The former allows to determine the temperature and the total emission measure (EM) of an isothermal plasma, and it has been the technique used in most of the studies that resulted in the isothermal corona results. This technique, while useful in the case of strictly isothermal plasmas, fails at providing an assessment of the uncertainty of its results. Traditional DEM techniques rely on several different methods to reconstruct the thermal structure of multithermal plasmas determining the DEM curve. However, most of them rely on assumptions on the smoothness of the distribution, that make it very difficult to study isothermal or near-isothermal plasmas.  Landi \\& Klimchuk (2010) carried out tests aimed at evaluating the ability of the EM loci technique to discriminate between isothermal and multithermal plasmas, when the spectral line intensities used for the reconstruction included uncertainties. They also developed a quantitative method to determine how multithermal could a plasma distribution be in order to be compatible with a set of observed line intensities that was also compatible with an isothermal condition.  In the present work we extend the Landi \\& Klimchuk (2010) study to the application of DEM diagnostic techniques to spectral line intensities. In particular, we aim at: 1) determining the ability of such techniques to retrieve isothermal plasma from a set of spectral line intensities, with and without random noise; 2) understanding to what extent are DEM diagnostic techniques able to discriminate between a truly isothermal plasma and a multithermal plasma with a narrow distribution; 3) quantifying how well can such technique detect and resolve multiple isothermal components along the lines of sight; 4) investigate the effects of atomic data uncertainties on the results; and 5) study how DEM reconstructions are affected by the number of ions whose lines are available.  Similar studies with different goals were recently carried out in preparation for the SDO and Hinode missions by Golub \\etal (2004) and Weber \\etal (2004). Both studies were focused on determining how well, and with what uncertainties, could a DEM curve known {\\em a-priori} be reconstructed using data from the limited number of broad-band filter intensities provided by Hinode/XRT and narrow-band filters in SDO/AIA. They also discussed how efficiently could DEM curves from the large number of pixels expected from such images be determined and visualized. Unlike the present work, such studies were aimed at imaging instruments and did not use spectral lines intensities; also, no guidelines were provided with regard to the interpretation of almost isothermal results. Further, the results of such tests strongly depend on the specific temperature response function of each instrument, so that they can not be easily extended to other instruments.  Line intensities can potentially provide better constraints on DEM curves due 1) to the better temperature resolution provided by their emissivities; and 2) to the finer sampling of the temperature range of the solar corona provided by the large number of ions included in the spectral range of the available spectrometers (CDS and SUMER on SOHO, EIS on Hinode). In the present work, we focus on spectral lines; given the larger amount of available constraints, we can hopefully obtain further insight on the intrinsic abilities of DEM diagnostic techniques at studying isothermal or near-isothermal plasmas. Our results can be applied to analyses of spectra observed by any of the available spectrometers observing the solar upper atmosphere, such as Hinode/EIS, SOHO/CDS, SOHO/SUMER, SOHO/UVCS, EUNIS, SERTS, HRTS and, to a certain extent, even the lower-resolution SDO/EVE. In fact, different spectrometers working at different wavelength ranges may still observe lines from the same element, so that they will be sampling the same temperature interval regardless of their passband. Thus, our results will be of more general use than Golub \\etal (2004) and Weber \\etal (2004).  Many different DEM diagnostic techniques have been developed in the past (see the reviews of Phillips \\etal 2008 and Harrison \\& Thompson 1992). Harrison \\& Thompson (1992) carried out a comparative analysis aimed at determining which, among six different DEM diagnostic techniques, was most successful at reproducing pre-defined DEM curves using a set of emission lines. More recently, Mark Weber and Paul Boerner led a similar study in preparation to the SDO mission ({\\em http://www.lmsal.com/~boerner/demtest/}), where, however, the results are not clearly discussed and no clear conclusions are drawn; also, such a study does not appear to have been published in the literature. Most importantly, Boerner \\& Weber did not address the ability at discriminating between isothermal and non-isothermal plasmas, which is the focus of the present work.  After inspecting the material in the Boerner \\& Weber website, we found that Monte Carlo Markov Chain (MCMC) method developed by Kashyap \\& Drake (1998) provided the most robust and accurate results. Also, this method dispenses from many of the assumptions common to other DEM diagnostic techniques, and provides an assessment of the uncertainties of the final DEM curve. Thus, we chose to focus on this method in the present study. Kashyap \\& Drake (1998) and Kashyap \\etal (2004) briefly described tests made to ensure and assess the robustness and reliability of the MCMC technique, but they did not focus on the ability of the MCMC technique to discriminate between an isothermal and multithermal plasma. Here we also provide a set of empirical guidelines for the choice of the input parameters and for the interpretation of MCMC reconstructions.  Section~\\ref{method} introduces the main plasma structure diagnostic techniques, including MCMC, and discusses their assumptions; it also explains the methodology we follow in our tests. Section~\\ref{results} describes the results of our tests for several different ad-hoc distributions, and Section~\\ref{conclusions} reviews the results of this study. ",
        "conclusions": "\\label{conclusions}  In this work we have tested the MCMC diagnostic technique by Kashyap \\& Drake (1998) in order to: determine its capability to reconstruct a given plasma thermal distribution (isothermal or multicomponent) from a set of individual line intensities; to study how the presence of random errors affects the MCMC reconstructions; to investigate the optimal size of the temperature bin width $W$ as a compromise between retaining as much temperature structure as possible and minimizing the effect of noise in the solution; and to test the effect of different atomic data sets on the final results. We did this by applying the MCMC technique to sets of lines whose intensities were calculated using a known thermal distribution first, and then randomized the intensities within 20\\% to simulate experimental uncertainties.  Our study shows that the optimal bin width $W$ is $W=0.05$, as smaller values cause the MCMC technique to provide spurious results and a high level of noise in the final reconstruction, while larger values oversmooth the results. Even when the emitting plasma is strictly isothermal, the MCMC technique is unable to distinguish between a truly isothermal solution and a Gaussian DEM with FWHM=0.05; also, the MCMC technique is able to separate multiple near-isothermal EM components only if their separation in temperature is $\\Delta \\log T=0.20$. Atomic data uncertainties can affect the results by providing less accurate peak EM values and shifting the EM peak temperature, but the $\\Delta \\log T=0.20$ resolving power of the MCMC technique is unaffected. The number of available ions does not affect the quality of the reconstruction, provided these ions are formed over a temperature range larger than the range where the plasma EM is significant. A smaller ion formation temperature range decreases the temperature resolution achieved by the MCMC technique.  \\bigskip  The work of Enrico Landi is supported by NASA grants NNX10AM17G and NNX11AC20G. Fabio Reale acknowledges support from Italian Ministero dell'Universit\\'a e Ricerca and from Agenzia Spaziale Italiana (ASI), contract I/015/07/0. Paola Testa was supported by NASA contract NNM07AB07C to the Smithsonian Astrophysical Observatory. We thank the referee for valuable comments that helped us improve the manuscript."
    },
    "1112/1112.0063_arXiv.txt": {
        "abstract": "We present empirical relations between the global dust reddening and other physical galaxy properties including the \\hal\\ luminosity, \\hal\\ surface brightness, metallicity and axial ratio for star-forming disc galaxies.  The study is based on a large sample of $\\sim$ 22 000 well-defined star-forming galaxies selected from the Sloan Digital Sky Survey (SDSS).  The reddening parameterized by color excess \\ebv\\ is derived from the Balmer decrement.  Besides the dependency of reddening on \\hal\\ luminosity / surface brightness and gas phase metallicity, it is also correlated with the galaxy inclination, in the sense that edge-on galaxies are more attenuated than face-on galaxies at a give intrinsic luminosity. In light of these correlations, we present the empirical formulae of \\ebv\\ as a function of these galaxy properties, with a scatter of only 0.07 mag.  The empirical relation can be reproduced if most dust attenuation to the \\hii\\ region is due to diffuse interstellar dust distributing in a disc thicker than that of \\hii\\ regions. The empirical formulae can be incorporated into semi-analytical models of galaxy formation and evolution to estimate the dust reddening and enable comparison with observations more practically. ",
        "introduction": "Dust is a crucial component of galaxies in modifying the observed properties of galaxies by absorbing and scattering starlight, and also re-emits the absorbed energy in the mid- and far-infrared bands.  The extinction\\footnote{The nomenclature conventionally adopted is that, ``extinction'' describes the absorption and scattering of light from a point source behind the dust screen; the ``attenuation'' describes the reduction and reddening of the light in a galaxy where stars and dust are mixed in geometrical distribution.  In this work, we do not discriminate between these two terms when referring to the dust attenuation within the galaxy, except for the Galactic ``extinction curve'' and the ``attenuation curve'' for the starburst galaxies.  The extinction and reddening can be converted to each other with an assumed extinction curve.} cross section of dust generally decreases with increasing wavelength, i.e., the extinction is more severe at shorter wavelength, particularly at ultra-violet (UV) band. As a result, the spectral energy distributions (SED) of galaxies appear redder, which is so-called ``reddening''.  The effects of dust should be properly accounted for when interpreting the observations of the galaxies, for example the luminosity or star formation rate (SFR) of a galaxy.  For a galaxy with spectroscopic data covering suitable wavelengths, it is usually possible to estimate the amount of dust extinction from the spectrum.  For example, the average dust reddening of \\hii\\ regions can be derived from Balmer decrement, and be used to correct for the dust attenuation to the emission-line luminosity.  In other cases, one has to rely on empirical relations to get an estimate of dust extinction in a statistical way.  Additionally, it is necessary to incorporate dust into the galaxy formation and evolution model.  Currently there are two ways to account for dust effects: full radiative transfer calculations assuming dust properties and geometric distribution, or adopting simple recipes obtained empirically.  Since the dust formation in the galactic environment is a complicated chemical process, it has not yet been implemented in the current galaxy evolution models.  The dust properties and distribution are input in the models and should be tested with observations.  Alternatively, many authors have adopted the empirical relations to compute the perpendicular optical depth of a galactic disc, and then assign a random inclination angle for each galaxy to get the final dust corrections (e.g. Guiderdoni \\& Rocca-Volmerange 1987; Kauffmann et al. 1999; Somerville \\& Primack 1999; De Lucia, Kauffmann \\& White 2004; De Lucia \\& Blaizot 2007; Kang et al. 2005; Kitzbichler \\& White 2007).  Therefore, a well-defined empirical recipe of dust reddening will help the comparison of model predictions with observations.  Previous studies have suggested that dust reddening in star-forming galaxies is correlated with the SFR, which can be well estimated with the intrinsic \\hal\\ luminosity (e.g. Kennicutt 1998; Calzetti et al. 2010). Galaxies with large SFR (or high luminosity) show strong global extinction in the emission lines (e.g., Wang \\& Heckman 1996; Charlot \\& Fall 2000;  Calzetti 2001; Stasi\\'nska \\& Sodr\\'e 2001; Afonso et al. 2003; Kewley et al. 2004; Zoran, Barkana, \\& Thompson 2006; Calzetti et al. 2007; Garn et al. 2010), and also large infrared to ultraviolet flux ratios (e.g., Iglesias-Paramo et al. 2006 and reference therein).  In addition, it has been recognized that dust reddening might also be a function of gas-phase metallicity, i.e. the reddening and extinction increase with metallicity.  For instance, the average reddening to \\hii\\ regions in an individual galaxy is correlated with the metallicity in the disc (e.g., Quillen \\& Yukita 2001; Boisser et al. 2004).  For the case of whole galaxies, the relationship between reddening and metallicity also exists (e.g. Heckman et al. 1998; Buat et al. 2002; Asari et al. 2007).  It is also supported by the fact that the low-metallicity galaxies such as blue compact galaxies are usually less reddened (e.g. Kong 2004).  Recently, Garn \\& Best (2010) also found a significant correlation between dust extinction and metallicity, however, they claimed that the dependency of reddening on stellar mass is more fundamental. They built a sample of about 90 000 star-forming galaxies using Sloan Digital Sky Survey (SDSS) data, and compare the relationship between the dust extinction of \\hal\\ luminosity and SFR, metallicity as well as stellar mass, respectively.  They concluded that the dust extinction $\\ahal$ can be best predicted from the stellar mass, with a scatter of $\\sim$0.3 mag.  Besides, the dependency of dust attenuation on inclination has been also investigated.  While the correlation between dust attenuation of optical stellar continuum and disc inclination has been well established for disc galaxies (e.g. Driver et al. 2007; Shao et al. 2007; Unterborn \\& Ryden 2008), that of emission lines is still uncertain (e.g. Yip et al. 2010).  These works focused on how the dust reddening depends on one property of galaxies, luminosity, metallicity or stellar mass, which are generally correlated with each other (e.g. Garnett \\& Shields 1987; Skillman et al. 1989; Zaritsky et al. 1994; Zahid et al. 2011; Tremonti er al. 2004).  Disentangling the dependencies of dust reddening on these galaxy properties can not only help us to derive a more accurate empirical formulae for dust reddening, but also provide more insight to the dust formation within galaxy environment.  In this work, we study a large sample of star-forming disc galaxies in the local universe with a median redshift of 0.07, selected from the spectroscopy database of SDSS. We obtain empirical formulae for dust reddening as a function of intrinsic \\hal\\ luminosity / surface brightness, gas-phase metallicity and the axial ratio of the disc.  In Section~\\ref{sec:sample} we describe our sample selection, and in Section~\\ref{sec:methods} describe the spectral analysis and methods used to estimate the parameters such as the dust reddening, intrinsic \\hal\\ luminosity, surface brightness, metallicity and axial ratio etc.  We present the resultant expression for the empirical formulae of dust reddening, and compare our results with previous work in Section~\\ref{sec:relation}.  In Section~\\ref{sec:toymodel} we reproduce the observed trend with a toy model of parallel-slab disc, which gives some insights into the geometry of dust distribution and dust-to-gas ratio. Conclusions are given in Section~\\ref{sec:conclusion}.  Throughout the paper, we will assume a cosmology of $H_0=72$ \\kms Mpc\\per, $\\Omega_M=0.3$ and $\\Omega_\\Lambda=0.7$. ",
        "conclusions": "\\label{sec:conclusion}  We present the empirical formulae of intrinsic reddening as a function of intrinsic \\hal\\ luminosity or surface brightness, metallicity and disc inclination (Equation~\\ref{eq:powfit}, \\ref{eq:powfit_br} and Table~\\ref{tab:lha} and \\ref{tab:br}) for a large sample of $\\sim$22 000 well-defined star-forming disc galaxies selected from the SDSS.  With the empirical formulae, the reddening parameterized by \\ebv\\ could be predicted within 1$\\sigma$ uncertainty of 0.07 mag.  The \\debv, defined as the observed reddening estimated from the Balmer decrement minus predicted value, does not correlate with the three parameters used in the formulae. We also find that \\debv\\ is independent of the stellar mass, 4000 \\AA\\ break strength or electron density within the galaxy.  The observed trend between the dust reddening and \\hal\\ surface brightness could be reproduced by a plane-parallel slab toy model, in which the dust is scaled with gas, and all the distributions of the dust, gas and star-formation regions are smooth and follow exponential laws in the disc.  We find from the comparisons of the models to the observation, the relative vertical scale of dust distribution to \\hii\\ regions distribution may vary with \\hal\\ surface brightness or metallicity. The higher intensity of star formation could drive the dust to distribute in a thicker layer, or the dust disc is thicker than the disc of the \\hii\\ regions when the metallicity is higher.  However, the real galaxies may be not fully described by this simple toy model.  For example, when the spiral arms are present, or the SF regions are clumpy, the assumptions of smooth distributions will be improper.  In addition, our result implies a different scaling law of dust-to-gas ratio as a function of metallicity $Z$ from the linearly relation (Equation~\\ref{eq:rhod1}).  Our empirical relations, which suggest dust reddening is partially dependent on metallicity ($\\propto Z^{0.6\\sim0.7}$), can be an observational constraints on the output of radiative-transfer model calculations, and hence provides constraints on the assumed dust-to-gas ratio in the model.  Since metallicity-dependence has been introduced, we believe that our empirical formulae should be applicable to the early evolutionary stage of a galaxy, when the gas metallicity might be low.  In the starburst galaxies at high redshift, dust surrounding star-forming region has been proven to be prone to sub-milimeter and far-infrared observations.  But most of those galaxies are relatively dim in optical, implying severe extinction and reddening to these galaxies.  The observed galaxies at high redshifts are affected by reddening, in the way that the galaxies with heavy attenuation and reddening may escape from the detection. The importance of the reddening correction depends on the wavelength band that is used.  For example, UV continuum luminosity is usually used to infer SFR in high-redshift galaxies.  It is very sensitive to the attenuation/reddening.  Calzetti et al. (2000) find that the extinction to emission lines and to stellar continuum in optical band are different, and the continuum reddening is only about 40 percent of that for emission lines.  However, the extinction to UV continuum is likely similar to the \\hii\\ region because UV continuum is also emitted by young hot stars, which are likely surrounded by \\hii\\ regions.  As a result, our results may be used to estimate the reddening for these high-redshift galaxies with caution, that the metallicity should be well-calibrated and be within the range covered by our sample.  Following Equation~\\ref{eq:powfit} and the discussions in Section~\\ref{sec:relation}, SFR in luminous star formation or high metallicity galaxies might have been under-estimated (e.g. Heckman et al. 1998; Hopkins et al. 2001; Panuzzo et al. 2007), or those galaxies are even completely lost in optical surveys.  The quantitative relationship of reddening with the metallicity, luminosity or surface brightness, and the inclination for disc galaxies can also be incorporated into the current semi-analytic models of galaxy formation and evolution (e.g. Croton et al. 2006; De Lucia et al. 2004; Kang et al. 2005). In those models, a simple chemical enrichment scheme has been included, and the metal abundance has been actually predicted.  With the known metallicity and the star formation rate (intrinsic luminosity) provided by the models, the reddening can be estimated accurately with our formulae, and then used to shape the spectral energy distribution and predict the emerging luminosity more realistically."
    },
    "1112/1112.5193_arXiv.txt": {
        "abstract": "The Characterisation data model \\citep{2011arXiv1111.2281L} is a standard of the International Virtual Observatory Alliance (IVOA) that describes observational datasets in the multi-dimensional parameter space. Defining three properties: coverage, resolution, and sampling along different physical axes (e.g. spatial, spectral, flux) with variable level of details for the description, this model has been used in several IVOA contexts: Simple Spectral Access Protocol, Spectrum Data Model, ObsTAP (Table Access Protocol for the Core Components of the Observation Data Model \\citep{2011arXiv1111.1758L}. Here we propose a  new version  which addresses more completely the most detailed level of description (level 4) dealing with variation maps of coverage, resolution, and sampling. It also introduces new specific axes in order to cover various photometric measurements, velocity and polarimetry. Special care is given for composed data sets. These improvements and add-ons follow the evolution of needs expressed to the uptake of VO tools in various observation domains for data discovery but also for data analysis requirements. This is also introduced to tackle use-cases designed to analyse scientifically datasets in the V0 context together with calibration/provenance information. ",
        "introduction": "The Observation data model project appeared at the first Data Model forum held at the May 2003 IVOA meeting in Cambridge,UK. Rapidly some main classes appeared to be necessary to organise the metadata: Dataset or Observation, Identification, Physical Characterisation, Provenance (either instrumental or software) and Curation. A description of the early stages of this development can be found in the Observation DM IVOA Note \\citep{OBSNOTE}. In parallel, an effort dedicated to spectra was lead by the DM Working group. The Spectrum data model represents all necessary metadata for one specific type of observational data: simple spectra. For the overall Observation Data Model, the physical characterisation has been identified to be on first priority already in 2004. It was completed as an IVOA recommendation after 4 years of discussion which included computer scientists, astronomers and data providers under the lead of J. McDowell. In theory full characterisation should be done by  describing the  transformation of incoming radiation into the actual measurements.     In practice the effects of this observation process can be split into various levels and pieces of information. The IVOA Characterisation data model organises metadata as a 3D matrix spanning independently the various physical axes: spatial, spectral, time, flux or whatever observable quantity (1st dimension), and describing for each of these axes 3 kind of properties, namely coverage (or sensitivity), resolution and sampling. The third dimension is the level of description of these items, from coarse average or typical values of parameters down to variation maps. This scheme allows to support selection of data sets for data discovery as well as data analysis. ",
        "conclusions": ""
    },
    "1112/1112.2911_arXiv.txt": {
        "abstract": "We investigate the cosmic thermal evolution with a vacuum energy which decays into photon at the low-redshift. We assume that the vacuum energy is  a function of the scale factor that increases toward the early universe. We put on the constraints  using  recent observations of both type Ia supernovae (SNIa)  by Union-2 compilation and  the cosmic microwave background (CMB) temperature at the range of the redshift $0.01 < z <  3$. From SNIa,  we find that the effects of a decaying vacuum energy on the  cosmic expansion rate should be  very small but could be possible for $z < 1.5$. On the other hand, we obtain the severe constraints for parameters from the CMB temperature observations.  Although the temperature can be still lower than the case of the standard cosmological model, it should only affect the thermal evolution at the early epoch. ",
        "introduction": "One of the biggest cosmological mysteries is the accelerating cosmic expansion which was discovered by the observations of distant type Ia supernovae (SNIa) starting from more than 10 years ago. For the origin of the accelerating expansion, the following possibilities have been proposed: the modified gravity such as $f(R)$ gravity\\cite{fr-grav},  brane-world cosmology~\\cite{Umezu2006}, inhomogeneous cosmology~\\cite{Inhomogeneous}, and existence of unknown  energy called as dark energy which is equivalent to a cosmic fluid with a negative pressure.  It is strongly suggested that the dark energy amounts to 70\\% of the total energy density of the universe from the astronomical observations such as SNIa \\cite{rf:Perlmutter,Union-2}, the anisotropy of cosmic microwave background (CMB) ~\\cite{QUaD,WMAP7}, and the baryon acoustic oscillation~\\cite{BAO}. Although there are various theoretical models of dark energy  (details are shown in a review \\cite{DEreview}), its physical nature is still unknown. {% Considering  above-referenced observations, we can not exclude the constant $\\Lambda$ term} which is  the most simple model of dark energy.  The interacting dark energy with CMB photon has been discussed, where% a decaying vacuum into photon is related to a primordial light-element abundances~\\cite{Freese1986}  and the CMB intensity~\\cite{Overduin1993}. From the point of the thermodynamical evolution in the universe, a decaying vacuum energy modifies the temperature-redshift relation~\\cite{Lima1995,Jetzer2011}. This modification affects the cosmic thermal evolution after a hydrogen recombination~\\cite{rf:Kimura}, such as formation of  molecules~\\cite{rf:Hashimoto,rf:Puy} and the first star~\\cite{rf:Hashimoto}. In addition to this model, the angular power spectrum of CMB could be also modified~\\cite{Nakamura2008}. In the previous analysis, thermal history at the higher redshift has been studied. However, as seen in \\cite{rf:Kimura,rf:Hashimoto}, the time variation of temperature at low redshift differs from the model without a decaying-vacuum  In the present study, we update the observational consistency of a vacuum energy (hereafter we denote as $\\Lambda$) coupled with CMB photon. Hereafter we call this model D$\\Lambda$CDM (model). To examine the consistency with observations at low redshift,  we focus on the cosmic evolution comparing both the type Ia supernovae and the redshift dependence of CMB temperatures. This is because the time dependence of the CMB temperature for a D$\\Lambda$CDM differs from CDM model with a constant $\\Lambda$ term of which hereafter we call S$\\Lambda$CDM (model). In \\S 2 formulation of  D$\\Lambda$CDM is reviewed. In \\S 3 the $m-z$ relation is investigated for D$\\Lambda$CDM  in a flat universe. Parameters inherent in this model are constrained from the CMB temperature observations in \\S 4. Concluding remarks are given in \\S 5. ",
        "conclusions": "We have investigated the consistency of the decaying-$\\Lambda$ model with both the $m-z$ relation of SNIa and CMB temperature. We obtain the upper bound of our model parameters from both observations. First we have constrained our model from SNIa: we can obtain the upper-bound of parameters,  which region is wider compared with previous constraints~\\cite{rf:Puy,Nakamura2008}. From this result, a decaying-$\\Lambda$ has minor effects on the cosmic expansion at low-$z$ and its parameters cannot be well constrained .  On the other hand, we can obtain more severe parameter constraints from CMB temperature. We acquire the limit of $\\Omega^{}_{\\Lambda2}$ as  {% $\\Omega_{\\Lambda2} < 1.7\\times10^{-3}$ at 2$\\sigma$ C.L. } From this result, the temperature can become lower by a few percents. That should affect the hydrogen recombination. We note that we cannot get the bound of $m$. Because change in $\\Omega^{}_{\\Lambda2}$ cancels the effects of the variation of $m$.  Our parameter is mainly obtained by the observations at $z>2$. It has been suggested that the observations of the temperature at higher-redshift is important to constrain the cosmological models with energy flow into photon~\\cite{Nakamura2008}. {% The upper-limit of $\\Omega^{}_{\\Lambda2}$ obtained from the present constraints is still larger than the previous constraints of $\\Omega^{}_{\\Lambda2} < 1.7\\times10^{-4}$ at $95.4\\%$ C.L.~\\cite{Nakamura2008}. The analysis using the CMB temperature fluctuation of the recent observation seems to be unable to give the tight constraints.} We plan to perform more detailed analysis, for instance using CMB anisotropy of the newest  WMAP data~\\cite{WMAP7} ."
    },
    "1112/1112.0313_arXiv.txt": {
        "abstract": "We study the growth of the red sequence through the number density and structural evolution of a sample of young and old quiescent galaxies at $0<z<2$.  The galaxies are selected from the NEWFIRM Medium-Band Survey (NMBS) in the Cosmic Evolution Survey (COSMOS) field.  We find a large population of massive young recently quenched (``post-starburst'') galaxies at $z>1$ that are almost non-existent at $z<1$; their number density is $5\\times10^{-5}$ Mpc$^{-3}$ at $z$$=$2, whereas it is a factor of 10 less at $z=0.5$. The observed number densities of young and old quiescent galaxies at $z>1$ are consistent with a simple model in which all old quiescent galaxies were once identified as post-starburst galaxies.  We find that the overall population of quiescent galaxies have smaller sizes and slightly more elongated shapes at higher redshift, in agreement with other recent studies. Interestingly, the most recently quenched galaxies at $1<z<2$ are not larger, and possibly even smaller, than older galaxies at those redshifts.  This result is inconsistent with the idea that the evolution of the average size of quiescent galaxies is largely driven by continuous transformations of larger, star-forming galaxies: in that case, the youngest quiescent galaxies would also be the largest.  Instead, mergers or other mechanisms appear to be required to explain the size growth of quiescent galaxies from $z=2$ to the present. ",
        "introduction": "\\label{sec:intro}  Galaxies with quiescent stellar populations form a well-defined relation in color-magnitude (color-mass) space known as the ``red sequence'', observed both in the local universe \\citep[e.g.,][]{Kauffmann03,Baldry04}, as well as out to high redshift \\citep[e.g.,][]{Franzetti07, Cassata08,Kriek08b,Williams09,Brammer09,Whitaker10,Whitaker11}. This red sequence is thought to grow over cosmic time by the continuous quenching and migration of star-forming galaxies, thereby changing the properties of these galaxies beyond simple passive evolution.  There is overwhelming evidence that the structures of these red sequence galaxies have undergone a rapid evolution since $z\\sim2$ \\citep[][etc.]{Daddi05,Cimatti08,vanDokkum08a}, where galaxies of the same stellar mass have significantly smaller sizes in the past. Although it has been shown that red sequence galaxies will experience some growth through minor mergers and accretion \\citep[e.g.,][]{Bezanson09, Naab09, Newman11}, the simplest explanation of this size growth is the continuous addition of recently quenched galaxies \\citep[e.g.,][]{vanderWel09a}.  These galaxies that quench at later times are expected to have larger sizes and rounder shapes because the universe was less dense and therefore gas-rich, dissipative processes were less efficient \\citep[e.g.,][]{Robertson06,Khochfar06}. The continuous transformation of these larger, star-forming galaxies at later times could be responsible for part of the size growth of the red sequence \\citep[e.g.,][]{vanDokkum08a, vanderWel09a, Hopkins09, Szomoru11}.  We can directly address the influence of transformations on the structural evolution of red sequence galaxies by looking to those galaxies that were most recently added. In this scenario, these recently quenched galaxies will have larger sizes than older galaxies on average at any given redshift. These young quiescent galaxies are easily identified when they pass through the ``post-starburst'' evolutionary phase because their spectral energy distributions (SEDs) are dominated by A-type stars, resulting in luminosity-weighted ages of $\\sim1$ Gyr \\citep[e.g.,][]{LeBorgne06}. The structures of the population of young quiescent galaxies relative to older quiescent galaxies provide crucial clues for the growth of the red sequence.  \\begin{figure*}[t!] \\leavevmode \\centering \\plottwo{f1a.eps}{f1b.eps} \\caption{(Left) Quiescent galaxies are selected based on their $U$--$V$ and $V$--$J$ colors, further divided into sub-samples of young (blue) and old (red) galaxies, where the greyscale represents the density of points. The separation of young and old quiescent galaxies is motivated by the age sequence of quiescent galaxies. (Right) The $U$--$V$ and $V$--$J$ colors of the entire NMBS parent sample, color-coded by the average best-fit age from \\citet{BC03} stellar population synthesis models.  The size of the symbol reflects the number of galaxies in the bin.  The \\citet{BC03} model track for a single burst stellar population with ages ranging from 0.5--10 Gyr is also shown. } \\label{fig:selection} \\end{figure*}  \\begin{figure}[t!] \\leavevmode \\centering \\includegraphics[width=0.9\\linewidth]{f2_small.eps} \\caption{Young and old quiescent galaxies are selected based on their $U$--$V$ and $V$--$J$ colors (see Figure~\\ref{fig:selection}). The composite rest-frame SED and running median of young (dark gray/blue) quiescent galaxies is distinct from that of old (light gray/red) quiescent galaxies (top panel).  The best-fit \\citet{BC03} models and the $U$, $V$ and $J$ transmission curves are shown in the bottom panel.} \\label{fig:compositesed} \\end{figure}  Different evolutionary scenarios for the growth of the red sequence have been notoriously difficult to test as recently quenched galaxies are hard to discriminate from older galaxies at high redshifts.  Owing to the excellent redshifts and accurate rest-frame colors of the NEWFIRM Medium-Band Survey \\citep[NMBS;][]{Whitaker10,Whitaker11}, we are now able to isolate samples of recently quenched galaxies. In this paper, we present the number density and structural evolution samples of young and old quiescent galaxies out to $z=2.2$.  We further discuss mechanisms to produce the size growth of quiescent galaxies.  We assume a $\\Lambda$CDM cosmology with $\\Omega_{M}$=0.3, $\\Omega_{\\Lambda}$=0.7, and $H_{0}$=70 km s$^{-1}$ Mpc$^{-1}$ throughout the paper.  All magnitudes are given in the AB system. ",
        "conclusions": "\\label{sec:summary}  We study the growth of the red sequence through the number density and structural evolution of a sample of young and old quiescent galaxies selected from the NEWFIRM Medium-Band Survey.  Due to the higher spectral resolution photometry of the NMBS, we can now --- for the first time --- study the properties of large samples of massive young quiescent galaxies at $z>1$.  We do not see many massive young galaxies that have recently quenched their star formation locally or even at low to intermediate redshifts.  In fact, we only begin to see a dramatic rise in the population of recently quenched post-starburst galaxies at $z>1$, finding an order of magnitude increase in their number densities between $z=0.5$ and $z=2$. The observed number densities of young and old quiescent galaxies are consistent with a simple model where all quiescent galaxies at $z>1$ have passed through the post-starburst evolutionary phase.  As the simple model used in this paper does not take into account growth due to mergers, we repeat the analysis selecting galaxies (both quiescent and star-forming) at a constant number density of $2\\times10^{-4}$ Mpc$^{-3}$, following \\citet{vanDokkum10}.  When taking into account mass growth on the red sequence by both mergers and the transformation of star-forming galaxies, we find an even more pronounced rise in the number density of young quiescent galaxies from $<10^{-6}$ Mpc$^{-3}$ at $z<1$ to $4\\times10^{-5}$ Mpc$^{-3}$ at $1.5<z<2$.  Somewhat surprisingly, the younger quiescent galaxies are not larger, and perhaps even somewhat smaller, than the older galaxies at a fixed redshift. The compact nature of recently quenched galaxies has implications for both the size growth of the red sequence and the mechanism by which galaxies quench their star formation.  If there is a faint extended halo around post-starburst galaxies at large radii that we are not resolving, the measured size will underestimate the true radial distribution of the stellar light \\citep[e.g.,][]{Hopkins09}. Although we measure more centrally concentrated surface brightness profiles for younger quiescent galaxies, there does not seem to exist a substantial fraction of stellar light in the outskirts that we are missing.  The surface brightness profiles from the stacked images of both young and old quiescent galaxies are similar out to 5 kpc.  With a careful analysis of quiescent galaxies, \\citet{Szomoru10, Szomoru11} similarly rule out the existence of a faint extended envelope around the observed galaxies down to the surface brightness limit of $H\\sim28$ mag arcsec$^{-2}$.  The size growth of quiescent galaxies does not seem to be a consequence of how we measure galaxy sizes.  The apparent lack of a size difference may not be intrinsic if the star formation history varies over the galaxy. In the simplest scenario where no major structural evolution occurs, massive blue galaxies may have a younger population in the central regions. This population may not necessarily constitute a substantial fraction of the galaxy's stellar mass, but it is very luminous.  This would lead to a post-starburst galaxy that is both bright and blue in the center but older in the outskirts.  Aging of the stellar populations will lessen the differences in brightness between the center and the rest of the galaxy and lead to an increasing size. This scenario therefore predicts gradients in the stellar populations of quiescent galaxies, with an older stellar population in a faint extended halo \\citep[e.g.,][]{Pracy11}. However, both \\citet{vanDokkumBrammer10} and \\citet{Guo11} have found the opposite, the inner regions of quiescent galaxies at $z\\sim2$ have redder rest-frame colors than the outer regions. We find no strong evidence for color gradients in this data set, as the ratio of the mean size of the old and young quiescent galaxies is approximately independent of the passband.  The small size of recently quenched post-starburst galaxies directly contradicts the theory that the average size of quiescent galaxies increases over time due to the continuous quenching and migration of larger, star-forming galaxies \\citep[e.g.,][]{vanderWel09b}. Rather, as both young and old galaxies appear to have similar sizes, the size growth of quiescent galaxies from $z=2$ to the present is likely the result of minor mergers and accretion of smaller galaxies.  The compact nature of these massive post-starburst galaxies is puzzling in the context of a coherent evolutionary picture.  Similarly massive star-forming galaxies at the same epoch are likely experiencing their last burst of star formation and will shortly quench and migrate to the red sequence. Without invoking quenching mechanisms, the structures of star-forming and quiescent galaxies must be connected in some way.  However, a simple size-mass relation is inadequate to describe the size variations of the entire galaxy population at $z=2$.  At all epochs, star-forming galaxies are the largest among galaxies of a given stellar mass \\citep[e.g.,][]{Williams10, Wuyts11b}. Furthermore, about half of massive star-forming galaxies at $z\\sim2$ are observed to have clumpy structures \\citep[e.g.,][]{ForsterSchreiber11,Swinbank10}. For these star-forming galaxies to be the progenitors of small quiescent galaxies, a compact core must already be present if the star formation quenches due only to the exhaustion of the gas reservoir.  Albeit for a small sample, \\citet{Kriek09b} find that at least half of star-forming galaxies at $z\\sim2$ require major structural changes beyond simple passive evolution to transform into these compact quenched systems. The process that makes these galaxies small may also quench the star formation.  Alternatively, the samples of star-forming galaxies might be incomplete.  Understanding the structural evolution of galaxies may hold the clues to why they transition from the active to passive phase.  Gas exhaustion may not be the sole driver of galaxies migrating to the red sequence, as it is during this epoch that the cosmic star formation rate density peaks \\citep[e.g.,][]{Hopkins06}, and galaxy halos are predicted to have large accretion rates \\citep[e.g.,][]{Dekel09}.  Moreover, the structures of galaxies are observed to be related to the stellar populations, implying a causal connection between the quenching and the morphological transition.  Among proposed quenching mechanisms that result in (post-)starbursts are gas-rich major mergers \\citep[e.g.,][]{Mihos94} after which point the available gas is exhausted from supernovae winds or Active Galactic Nuclei (AGN) feedback \\citep[e.g.,][]{Croton06,PHopkins06}, shock heating in massive halos \\citep{Birnboim07,Yan09}, or cold streams resulting in clump migration and coalescence \\citep[e.g.,][]{Dekel09b}. The transformation mechanism responsible may drive cold gas to the central disk of the halo to trigger a starburst, thereby changing the structure of the galaxy and subsequently quenching star formation.  It is clear that more information about the stellar populations of massive star-forming and quiescent galaxies is necessary before we can begin to understand the connection between the active and passive phases of galaxies.  Future detailed studies of the environment and morphologies of both post-starburst and star-forming galaxies at high redshift may be able to disentangle the relative importance of these mechanisms."
    },
    "1112/1112.1514_arXiv.txt": {
        "abstract": "{Determining the mass transfer in a close binary system is of prime importance for understanding its evolution. SS Leporis, a symbiotic star showing the Algol paradox and presenting clear evidence of ongoing mass transfer, in which the donor has been thought to fill its Roche lobe, is a target particularly suited to this kind of study.} {Since previous spectroscopic and interferometric observations have not been able to fully constrain the system morphology and characteristics, we go one step further to determine its orbital parameters, for which we need new interferometric observations directly probing the inner parts of the system with a much higher number of spatial frequencies.} {We use data obtained at eight different epochs with the VLTI instruments AMBER and PIONIER in the H- and K-bands. We performed aperture synthesis imaging to obtain the first model-independent view of this system. We then modelled it as a binary (whose giant is spatially resolved) that is surrounded by a circumbinary disc.} {Combining these interferometric measurements with previous radial velocities, we fully constrain the orbit of the system. { We then determine the mass of each star and significantly revise the mass ratio. The M giant also appears to be almost twice smaller than previously thought. Additionally, the low spectral resolution of the data allows the flux of both stars and of the dusty disc to be determined along the H and K bands, and thereby extracting their temperatures.}} {We find that the M giant actually does not  \\protect{\\it stricto sensus} fill its Roche lobe. The mass transfer is more likely to occur through the accretion of an important part of the giant wind. We finally rise the possibility for {  an enhanced mass loss from the giant, and we show that an accretion disc should have formed around the A star}.} ",
        "introduction": "\\label{part:intro}  Symbiotic stars are interacting binaries composed of a hot star accreting material from a more evolved red giant companion. They are excellent laboratories for studying a wide spectrum of poorly understood physical processes, like the late stage of stellar evolution, the mass loss of red giants, and the mass transfer and accretion in binary systems \\citep{mikolajewska_2007}. Their study has important implications for a wide range of objects, like Type Ia supernovae, barium stars, the shaping of planetary nebulae, and compact binaries like cataclysmic variables \\citep{podsiakolski_2007}.  SS Leporis (17 Lep; HD 41511; HR 2148) is a prime example of such a long-period interacting system, even though it does not belong to the most common symbiotic systems, because the hot star is not the usual compact white dwarf. As such, SS Lep is a symbiotic system in the first phase of mass transfer, while most symbiotic stars are in their second episode of mass transfer, following the first one that produced the white dwarf.  SS Lep has been known for many decades to present symbiotic features, and its optical spectrum shows at least three components \\citep{struve_1939, molaro_1983, welty_1995}. The spectral lines of an A star are largely obliterated by shell features that dominate at shorter wavelengths, while an M star spectrum becomes increasingly obvious at longer wavelengths. \\citet{welty_1995} estimated an M4 III spectral type for the cool companion, while even earlier types have been estimated by previous authors. The shell is absorbing light primarily from the A star, indicating some mass loss from the hotter star. The system, however, presents the so-called Algol paradox, as the most evolved star is also the least massive, which implies that the cool star must have lost a large quantity of matter and that the hot companion has accreted part or most of it. Moreover, the regular ``outbursts'' \\citep{struve_1930, welty_1995} and the UV activity \\citep{polidan_1993} of the A star shell are clear testimony t ongoing mass-transfer episodes. From interferometric observations, \\citet{verhoelst_2007} inferred that the mass transfer occurs because the M giant fills its Roche lobe.  The binary system is additionally surrounded by a large circumbinary dust disc and/or envelope \\citep{jura_2001}. Interferometric observations confirmed this fact by revealing its presence in the inner part of the system \\citep{verhoelst_2011}, further noticing that the structure must be in a disc-like geometry to be compatible with the low extinction towards the central star. \\citet{jura_2001} suggest that the circumbinary disc contains large grains that are formed by coagulation and, based on the large and rather unique 12 $\\mu$m excess of SS Lep despite its rather low luminosity,  that the disc may be losing mass by a wind at a rate of $8\\times10^{-9}$ M$_\\odot$yr$^{-1}$.  The orbital characteristics and circumbinary disc of SS Lep very closely resemble those of the post-AGB binaries with stable, Keplerian circumbinary dust discs \\citep{vanwinckel_2003}, and, as such, SS Lep may be considered as a system linking binary M giants and post-AGB systems -- the M-giant should indeed very soon evolve into a post-AGB star. In those post-AGB binaries, the spectra are rich in crystalline features while the spectrum of SS Lep appears entirely amorphous, which, if a link is indeed in order, would imply further dust processing in the disc. Whether the disc can survive long enough if it is indeed losing mass through a wind is, however, still an open question.  \\citet{cowley_1967} argued that the system consists of a B9 V primary and an M1 III secondary in a 260-d eccentric orbit ($e=0.132$). She developed a scenario in which the secondary fills its Roche lobe near periastron and mass transfer proceeds for a short time thereafter. From spectra covering 3.5 orbits, \\citet{welty_1995} proved this scenario unlikely, as their revised orbit provided a similar orbital period but a significantly reduced eccentricity  $e=0.024$. They estimated a mass ratio of $1/q = 3.50 \\pm 0.57$, where the error was very likely severely underestimated given the poor fit of the single Mg II line they used to measure the radial velocity of the A star.  Recently, \\citet{leeuwen_2007} has reevaluated the parallax of SS Lep from Hipparcos data, obtaining $\\pi = 3.59\\pm0.31$\\,mas, that is, a distance of $279\\pm24$\\,pc, so smaller than the previously and generally used value of $330^{+90}_{-60}$\\,pc. The most recent parameters of SS Lep as collected from the literature until the present work, is presented in Table~\\ref{tab1}.  We report here interferometric observations that have allowed model-independent image synthesis (Section \\ref{part:image}) and a more precise modelling of the system (Section \\ref{part:model}). Section \\ref{part:binary} focuses on the binary by determining mainly the orbital parameters and the M star diameter. We finally discuss the mass transfer process in Section \\ref{part:masstransfer}.  \\begin{table}[t!]\\label{tab1} \\caption{\\label{tab:param} Previously estimated parameters of SS Lep} \\centering \\begin{tabular}{lc|lcc} \\hline \\hline & System &          & A star & M star   \\\\ \\hline $d$ [pc]        & $279\\pm24$ $^a$ &  {\\it SpT}                       & A1V $^{b}$    & M6III $^{c}$ \\\\ $P$ [d]          & $260.3\\pm1.8$ $^b$     & $T_{\\rm eff}$ [K]  & 9000 $^{c,d}$   & 3500 $^{c,d}$ \\\\ $e$          & $0.024\\pm0.005$ $^b$ &$\\theta$ [mas]   & $0.53\\pm0.02$  $^c$  & $3.11\\pm0.32$  $^c$\\\\ $i$           & $30^\\circ \\pm 10^\\circ$ $^c$   &$R$ [${\\rm R}_\\odot$]        & $\\sim$15 $^{c}$      & $110\\pm30$ $^{c}$ \\\\ $f(M)$& \t$0.261\\pm0.005$ $^b$ &$F$ [\\%] & $11\\pm7$ $^c$     & $60\\pm50$ $^c$ \\\\ $1/q$            & $4\\pm1$ $^{b,c}$  & $M$ [${\\rm M}_\\odot$] & 2$\\sim$3 $^{a,b}$ & 0.4$\\sim$1 $^{a,b}$ \\\\ \\hline  \\hline \\end{tabular} \\tablefoot{$d$ is the distance, $P$ the orbital period, $e$ the eccentricity, $i$ the inclination, $f(M)$ the mass function, $q=M_M/M_A$ the mass ratio. For the stars, {\\it SpT} is the spectral type, $T_{\\rm eff}$ the temperature, $\\theta$ the apparent diameter, $R$ the linear radius, $F$ the flux contribution at $2.2\\,\\mu m$, and $M$ the mass. References: (a) \\citealt{leeuwen_2007}; (b) \\citealt{welty_1995}; (c) \\citealt{verhoelst_2007}; (d) \\citealt{blondel_1993}.} \\end{table} ",
        "conclusions": "We have presented here the results of VLTI observations, and we focused on the binary. After having computed the characteristics of the orbit, we demonstrated that the mass ratio is lower than previously thought and that the M giant does not fill its Roche lobe. However the system is possibly in the configuration of a wind Roche lobe overflow, where a substantial part of the giant's stellar wind can be accreted by the A star. We also have good reasons to think that the A star is actually surrounded by an accretion disc, although this needs to be investigated further.  We still lack a low-resolution spectro-photometry of SS~Lep between $1.6$ and $2.5$\\,$\\mu$m to compute the absolute luminosity of each component and extract more specific information. The current data also present good hints of unmodelled material escaping the system. It would be of uttermost importance to obtain additional interferometric observations of SS Lep -- at large but also small baselines -- to study the binary interactions in more details (M star distortion, mass transfer, circumbinary disc) and to determine if the A star is surrounded by an accretion disc.    \\acknowledgement{The authors would like to thank O.\\ Absil and G.\\ Dubus for their help, as well as the referee N.\\ Elias, whose careful review of the text helped them improve the papers clarity and quality. PIONIER is funded by the Universit\\'e Joseph Fourier (UJF, Grenoble) with the programme TUNES-SMING, the Institut de Plan\\'etologie et d'Astrophysique de Grenoble (IPAG, ex-LAOG), and the Institut National des Science de l'Univers (INSU) with the programmes ``Programme National de Physique Stellaire'' and ``Programme National de Plan\\'etologie''. PIONIER is equipped with a detector provided by W. Traub (JPL, Caltech). The authors want to warmly thank the VLTI team. This work is based on observations made with the ESO telescopes. It made use of the Smithsonian/NASA's Astrophysics Data System (ADS) and of the Centre de Donnees astronomiques de Strasbourg (CDS). All calculations and figures were performed with the freeware \\texttt{Yorick}\\footnote{\\texttt{http://yorick.sourceforge.net}}. The image reconstruction was performed with the software \\texttt{MIRA}\\footnote{\\texttt{http://www-obs.univ-lyon1.fr/labo/perso/ eric.thiebaut/mira.html}}. }"
    },
    "1112/1112.3985_arXiv.txt": {
        "abstract": "We present the first results from AMUSE-Field, a \\cxo~survey designed to characterize the occurrence and intensity of low-level accretion onto supermassive black holes (SMBHs) at the center of local early-type field galaxies. This is accomplished by means of a Large~Program targeting a distance-limited ($<$30~Mpc) sample of 103 early types spanning a wide range in stellar masses. We acquired new ACIS-S observations for 61 objects down to a limiting (0.3--10~keV) luminosity of $2.5\\times 10^{38}$ \\es, and we include an additional 42 objects with archival (typically deeper) coverage. A nuclear \\hbox{X-ray} source is detected in 52 out of the 103 galaxies. After accounting for potential contamination from low-mass \\hbox{X-ray} binaries, we estimate that the fraction of accreting SMBHs within the sample is $45\\pm7$\\%, which sets a firm lower limit on the occupation fraction within the field. The measured nuclear X-ray luminosities are invariably highly sub-Eddington, with $L_{\\rm X}/L_{\\rm Edd}$ ratios between $\\sim$$10^{-4}$--$10^{-8}$. As also found in a companion survey targeting Virgo early types, the active fraction increases with increasing host galaxy stellar mass, reflective of ``Eddington incompleteness'' within the lower-mass objects. For the Field sample, the average nuclear X-ray luminosity scales with the host stellar mass as $M_{\\rm star}^{0.71\\pm0.10}$, with an intrinsic scatter of $0.73\\pm0.09$ dex. Qualitatively similar results hold for morphologically homogeneous (type E) or uniform sensitivity (new observations only) subsets. A majority of the AMUSE-Field galaxies (78\\%) inhabits groups, enabling us to investigate the influence of group richness upon nuclear activity. We see no evidence for a positive correlation between nuclear \\hbox{X-ray} luminosity, normalized to host properties, and galaxy density. Rather, while the scatter is substantial, it appears that the Eddington-scaled \\hbox{X-ray} luminosity of group members may be slightly lower than for isolated galaxies, and that this trend continues to cluster early-types. ",
        "introduction": "Convincing evidence that supermassive black holes (SMBHs) are able to form in the early universe is provided by observations of high-redshift quasars (e.g., Volonteri \\& Rees 2006; Vestergaard \\& Osmer 2009; Willott et al.~2010; Treister et al.~2011). The seeds for these early SMBHs may have been produced from supermassive (e.g., Begelman 2010) or Population III stars, or from direct gas collapse (e.g., Volonteri \\& Natarajan 2009), perhaps in massive protogalaxy mergers (e.g., Mayer et al.~2010; Volonteri 2010). While high-redshift quasars generally display optical/UV and X-ray properties similar to those of their local analogues (e.g., Shemmer et al.~2006), {\\it Spitzer\\/} observations indicate many are young sources (some lack hot dust) that are growing rapidly (Jiang et al.~2010). In contrast, similarly massive SMBHs at more moderate redshifts are growing more slowly (e.g., Netzer et al.~2007). Indeed, the growth of SMBHs appears to be ``anti-hierarchical'' in the sense that active accretion is concentrated in higher/lower-mass SMBHs at earlier/later cosmological times (e.g., Heckman et al.~2004; Merloni \\& Heinz 2007; Shankar et al.~2009; Gallo et al.~2010; Goulding et al.~2010; Kelly et al.~2010; Lamastra et al.~2010; Schulze \\& Wisotzki 2010; Schawinski et al.~2010).  The peak of the quasar space density around $z\\sim2$ (e.g., Brown et al.~2006; Hopkins et al.~2007; Kelly et al.~2010)\\footnote{The quasar peak redshift is luminosity-dependent, displaying a similar ``downsizing'' effect (e.g., Hasinger et al.~2005; Croom et al.~2009) to that found for SMBH growth.} and the $\\sim$10$^{8}$~yr quasar lifetime (e.g., Yu \\& Tremaine 2002) suggest (e.g., Soltan 1982) that ``inactive'' galactic nuclei typically also host SMBHs accreting at low levels and/or radiating inefficiently in a post-quasar stage which may result after much of the available fuel has been consumed or expelled (e.g., Hopkins et al.~2005). The distinctions between active galactic nuclei (AGNs), low-luminosity AGNs (Ho 1999, 2008), and formally inactive galaxies are somewhat arbitrary, but as a general guideline (which we will adhere to in this work) AGNs display bolometric, Eddington-scaled luminosities in excess of a few percent, low-luminosity AGNs are in the range $10^{-4}<L/L_{\\rm Edd}<10^{-2}$, and inactive nuclei are highly sub-Eddington, with \\eddra \\simlt $10^{-4}$.  The Milky Way is one example of a formally inactive galaxy. It is known to host a quiescent central SMBH with mass of $3.6\\times10^{6} M_{\\odot}$ (e.g., Sch{\\\"o}del et al.~2009) from which low-level, persistent accretion-powered \\hbox{X-ray} emission has been detected (\\xeddra$\\simeq 10^{-11}$; Baganoff et al.~2001). Even low levels of SMBH accretion-powered activity in nearby galaxies may be detected efficiently in sensitive, high spatial resolution \\hbox{X-ray} observations, although contamination from bright \\hbox{X-ray} binaries must be properly accounted for when dealing with nuclear X-ray luminosities comparable to the Eddington limit for a few solar masses (e.g., Zhang et al.~2009; Gallo et al.~2010). This problem can be substantially alleviated by restricting the search for highly sub-Eddington nuclei to early-type galaxies, as they conveniently avoid the brighter high-mass \\hbox{X-ray} binaries produced concurrently with star formation (e.g., King et al.~2001; Ghosh et al.~2009).  From \\hbox{X-ray} and other observations as well as numerical simulations, a coherent scenario for the history and observed characteristics of quiescent or weakly accreting SMBHs in early-type galaxies is emerging. Although hot halo gas or stellar winds provide readily available fuel and gas infall may proceed at near the Bondi rate, outflows can drive off a substantial fraction of the accreting mass and the radiative efficiency is inferred from observed spectral energy distributions to be highly sub-Eddington (e.g., Di Matteo et al.~2000; Pellegrini 2005; Soria et al.~2006b).  The growth of SMBHs is believed to be closely tied to the evolution of their hosts (as evidenced by, for example, the correlations between SMBH mass and bulge luminosity or central stellar velocity dispersion; see, e.g., overview by Ferrarese \\& Ford 2005). Radiative and mechanical AGN feedback act to regulate SMBH accretion and quench star formation (e.g., Ciotti et al.~2009, 2010; Pellegrini et al.~2012), and mechanical feedback continues to play an important role in low-activity ``radio-mode'' systems (e.g., Croton et al.~2006; Merloni \\& Heinz~2007; Marulli et al.~2008). The assembly of early-type galaxies apparently proceeds hierarchically, despite the older ages and shorter star-formation timescales of massive ellipticals (e.g., De Lucia et al.~2006; Eliche-Moral et al.~2010), with cold-gas-rich/poor (``wet''/``dry'') mergers tending to produce discy/boxy types (e.g., Kang et al.~2007; Hopkins \\& Quataert~2011).  Naturally the properties of higher-mass SMBHs (in more massive galaxies) may be more easily constrained; the dominance of compact stellar nuclei at lower luminosities (e.g., Ferrarese et al.~2006a) presented initial observational challenges to establishing whether lower-mass early-type galaxies necessarily even contained a SMBH. The presence of ``missing light'' or cores within luminous ellipticals may be explained by binary black holes disrupting the interior region in the process of merging (e.g., Kormendy et al.~2009 and references therein), and so preserved nuclear star clusters could potentially indicate a lack of a SMBH binary phase (e.g., Kang et al.~2007). However, it is now clear that SMBHs may coexist with nuclear star clusters (e.g., Seth et al.~2008; Graham \\& Spitler~2009; Kormendy et al.~2009).  To quantify the rate of low-level SMBH activity over a well defined sample of nearby early-type galaxies, our group carried out the AGN MUltiwavelength Survey of Early-Type Galaxies in the Virgo cluster. AMUSE-Virgo (ID 8900784, \\cxo~Cycle~8, PI: Treu, 454~ks) targeted the 100 spheroidal galaxies which compose the \\hst~ACS Virgo Cluster Survey (VCS; C{\\^o}t{\\'e} et al.~2004) with \\cxo~ACIS-S and \\spi~MIPS (Gallo et al.~2008, 2010, hereafter G08, G10; Leipski et al.~2012). The VCS sample was selected based solely on optical properties and spans a wide range in stellar mass, from $10^{8}-10^{12}~M_{\\odot}$. At the average Virgo distance of 16.5~Mpc (Mei et al.~2007), the AMUSE-Virgo {\\it Chandra\\/} snapshot observations (about 5 ks each) reached a $2\\sigma$ limiting 0.5--7~keV luminosity of $1.3\\times10^{38}$~erg~s$^{-1}$ (i.e., close to the Eddington limit for 1 \\msun). A nuclear \\hbox{X-ray} source was detected by G10 in 32/100 objects, which, after taking into consideration the minor but non-negligible effects of low-mass \\hbox{X-ray} binary (LMXB) contamination, sets a lower limit of 24--34\\% (at the 95\\% confidence level) to the occupation fraction of SMBHs in the nuclei of cluster early-type galaxies. Within the AMUSE-Virgo survey the fraction of \\hbox{X-ray}-active SMBHs is seen to increase with host stellar mass,\\footnote{That SMBH activity in general is detected more frequently with larger host stellar mass or luminosity is known since the first Palomar sample (e.g., Ho et al.~1997).} but in a manner consistent with arising from {``Eddington incompleteness\"}, i.e., the intrinsic inability of any luminosity-limited survey to reach the same Eddington-scaled luminosity level across a wide range of inferred black hole masses. Despite this effect, which obviously penalizes lower mass black holes, the average Eddington-scaled \\hbox{X-ray} luminosity of the AMUSE-Virgo galaxies is found to scale inversely with black hole mass as $\\langle L_{\\rm X}/L_{\\rm Edd}\\rangle \\propto M_{\\rm BH}^{-0.62}$, so that nearby lower-mass SMBHs are relatively more \\hbox{X-ray}-active {(a trend G10 refer to as ``downsizing in black hole accretion'')}. In all cases the \\hbox{X-ray} luminosity is highly sub-Eddington, with $10^{-8}{\\simlt}L_{\\rm X}/L_{\\rm Edd}{\\simlt}10^{-5}$.  While this and other studies (e.g., Pellegrini 2005; Zhang et al.~2009; Boroson et al.~2011) have provided useful and detailed census of low-level SMBH activity within the local universe, including as a function of the host galaxy properties (Pellegrini 2010), {\\it the impact of environment upon the occurrence and intensity of SMBH activity remains unclear}. The properties of early-type galaxies in more sparsely populated regions are distinct from their counterparts in clusters: relative to cluster sources, field early-type galaxies face reduced ram pressure stripping (e.g., Acreman et al.~2003; Gavazzi et al.~2010; Shin et al.~2012) and on average contain more cold gas and tend to have younger stellar populations (e.g., Thomas et al.~2005; Oosterloo et al.~2010). Within clusters, tidal interactions play an important role (e.g., Gnedin 2003), for example as harassment from high-speed encounters (e.g., Treu et al.~2003); such effects are only infrequently relevant to isolated field galaxies, although within groups the lower galaxy velocity dispersions can help facilitate mergers. Early work with {\\it ROSAT\\/} links the environment with \\hbox{X-ray} emission, but to a debated degree; for example, Brown \\& Bregman (2000) found that cluster early-type galaxies tend to be more \\hbox{X-ray} luminous than their field counterparts, and hypothesized that this results from environmental suppression of outflows or the availability for accretion of hot intergalactic gas within clusters, whereas O'Sullivan et al.~(2001) found similar $L_{\\rm X}(L_{\\rm B})$ relations for cluster, group, and field galaxies. While such results are intriguing, their utility is necessarily restricted by the limited angular resolution of {\\it ROSAT\\/}, which makes it difficult to disentangle diffuse thermal \\hbox{X-ray} or off-nuclear \\hbox{X-ray} binary emission from that linked to the SMBH.  The goal of this work is to investigate low-level SMBH activity within a well defined sample of nearby field spheroidal galaxies (E and E-S0 galaxies), and to characterize the dependence (if any) of such activity upon environmental effects, such as living in isolation versus belonging to a group of galaxies, or belonging to a poor versus a rich group. Toward this end, we were awarded a Large {\\it Chandra\\/} program building upon and extending AMUSE-Virgo to non-cluster environments: AMUSE-Field (ID 11620915, \\cxo~Cycle 11, 479~ks, PI Gallo) carried out new observations of 61 field early-type galaxies located within 30~Mpc and spanning a wide range of stellar and black hole masses. The new observations are supplemented with an additional 42 objects with archival coverage, for a total of 103 objects.  This paper is structured as follows: $\\S$2 describes the sample selection and {\\it Chandra\\/} data reduction; $\\S$3 provides the \\hbox{X-ray} results; $\\S$4 investigates the dependence of nuclear \\hbox{X-ray} luminosity on stellar mass; $\\S$5 discusses environmental influences upon SMBH accretion; $\\S$6 summarizes and concludes. $L_{\\rm X}$, $M_{\\rm star}$, and $M_{\\rm BH}$ are the nuclear 0.3--10~keV \\hbox{X-ray} luminosity, galaxy stellar mass, and black hole mass, respectively, and are given in units of erg~s$^{-1}$, $M_{\\odot}$, and $M_{\\odot}$ throughout. Errors are $1\\sigma$ unless otherwise noted. ",
        "conclusions": "This work presents first results from the AMUSE-Field survey, a Large {\\it Chandra\\/} program designed to characterize low-level SMBH activity within nearby field spheroids. The Field sample contains 103 early-type non-cluster galaxies within 30~Mpc, selected entirely based on their optical properties and spanning a wide range in absolute magnitude and stellar mass. For 61 objects we obtained new ACIS-S observations to a uniform limiting luminosity of $2.5\\times10^{38}$~erg~s$^{-1}$, and for the remaining 42 objects we make use of high-quality archival coverage. The primary results of this paper are as follows:  1. The detection fraction of nuclear \\hbox{X-ray} sources is 52/103. A handful of these may be contaminated by LMXBs; we estimate the fraction of accreting SMBHs is $45\\pm7$\\%. This provides a firm lower limit on the SMBH occupation fraction within field spheroids.  2. The ratio of \\hbox{X-ray} to Eddington luminosities within the Field sample is in all cases low, from $\\sim10^{-4}-10^{-8}$. This range is similar to that found in the AMUSE-Virgo study of cluster early-types (G10). Independent of environment, SMBH accretion within local early-types is generally highly sub-Eddington.  3. We quantify the correlation between nuclear \\hbox{X-ray} luminosity and host stellar mass, finding a best-fit relation of $(\\log{L_{\\rm X}}-38.4)=(-0.04\\pm0.12)+(0.71\\pm0.10)\\times(\\log{M_{\\rm star}}-9.8)$. The calculated slope is incompatible with zero dependence and marginally indicative ($2.9\\sigma$) of greater $L_{\\rm X}$ per unit $M_{\\rm star}$ in lower-mass galaxies.  4. After scaling by black hole mass or accounting for the influence of stellar mass, typical values of $\\log{L_{\\rm X}/L_{\\rm Edd}}$ or ${\\Delta}L_{\\rm X,M*}$ of group members may be slightly lower than for isolated galaxies, and this newly-identified trend appears to continue to cluster early-types. If confirmed, a progressive decrease in relative \\hbox{X-ray} luminosity with increasing richness mandates an environmental modulation of SMBH fueling.  Several complementary projects that build on this dataset are underway. A direct comparison of the nuclear \\hbox{X-ray} luminosity as a function of black hole mass within field and cluster galaxies, controlling for stellar mass, is discussed in Miller et al.~(2012). An additional later work will present the catalog of off-nuclear \\hbox{X-ray} sources within the field sample, and compare the distribution of ULXs to that found for Virgo. We are also currently obtaining high-resolution {\\it HST\\/}/ACS imaging of Field galaxies with \\hbox{X-ray} detections, to check for and characterize any nuclear star clusters."
    },
    "1112/1112.4870_arXiv.txt": {
        "abstract": "{We report thorium abundances for 77 metal-poor stars in the metallicity range of $-3.5<\\mathrm{[Fe/H]}<-1.0$, based on ``snapshot'' spectra obtained with VLT-UT2/UVES during the HERES Survey. We were able to determine the thorium abundances with better than $1\\sigma$ confidence for 17 stars, while for 60 stars we derived upper limits. For five stars common with previous studies, our results were in good agreement with the literature results. The thorium abundances span a wide range of about 4.0 dex, and scatter exists in the distribution of log(Th/Eu) ratios for lower metallicity stars, supporting previous studies suggesting the r-process is not universal. We derived ages from the log(Th/Eu) ratios for 12 stars, resulting in large scattered ages, and two stars with significant enhancement of Th relative to Eu are found, indicating the ``actinide boost'' does not seem to be a rare phenomenon and thus highlighting the risk in using log(Th/Eu) to derive stellar ages. ",
        "introduction": "The observed chemical abundances of metal-poor stars provide rich information on star formation and nucleosynthesis in the early stages of evolution of our galaxy and other galaxies. In particular, the study of heavy elements beyond the iron group in metal-poor stars have greatly improved the understanding of the neutron-capture nucleosynthesis processes in the early Universe, e.g., placing constraints on the astrophysical site(s) of r-process. Neutron-capture elements in r-process enhanced metal-poor stars have been reported to match the scaled solar r-process abundance pattern at least in the range of $\\mathrm{Z}=52\\sim76$ \\citep[see e.g.,][]{cowan:02,sneden:03,honda:04}. In contrast, lighter elements ($\\mathrm{Z}<50$) abundances deviate significantly from the Solar-pattern, indicating the existence of different initial production ratios for the r-process elements \\citep{sneden:00,sneden:03,honda:04}. In old, metal-poor stars, thorium and uranium with relatively long half-life are the only heavier r-process elements we could observe today, and offer a way to better understand the heaviest products in r-process nucleosynthesis \\citep{roederer:09}.  The detection of the radioactive-decay elements in r-process enhanced metal-poor stars allowed a new approach to derive the age of the oldest stars, by means of comparing the observed ratios of radioactive decay elements over stable elements with the corresponding initial values at the time when the star was born\\footnote{Accurately speaking, it was the time when the radioactive decay elements were produced, not necessarily to be the same as when the star was born.}, hence providing a lower limit on the age of the universe. The actinide thorium is one such radioactive species. It is produced in the r-process, and has been widely used, particularly log(Th/Eu), for cosmochronometry to estimate stellar ages \\citep[e.g.,][]{sneden:96,hill:02,honda:04,frebel:07,hayek:09}. Some authors have argued that log(Th/Eu) is not a reliable chronometer to derive stellar ages, because of the relatively large uncertainties \\citep[e.g.,][]{goriely:99,goriely:01,honda:04,frebel:07}. Due to its longer half-life of $14.05$ Gyr as compared with, say, $4.5$\\,Gyr of uranium, some authors have suggested that the log(U/Th) pair would be a better chronometer, because of their close atomic number and thus relatively small nuclear physics uncertainties \\citep{cayrel:01,wanajo:02,beers:05}. Additionally, because Pb and Bi mostly originate from $\\alpha$-decay of the Th and U isotopes, the measurement of Pb or Bi abundances can offer a consistency check on the calculated initial abundances for this pair \\citep{cowan:99}. However, it is very difficult to detect U lines in stellar spectra, due to the blending of the weak U lines (e.g., 3859 \\AA~) and the relatively low amount of uranium in an old star with age of $\\sim12-15$\\,Gyr \\citep[see][]{plez:04,roederer:09}. Alternatively, Hf behaves similarly to third-peak elements and the log(Th/Hf) ratio has been suggested by \\citet{kratz:07} as a promising tool for chronometer studies.  Many efforts have been made to determine thorium abundances in metal-poor stars using high resolution spectra. So far, more than 30 metal-poor stars ($\\mathrm{[Fe/H]}<-1$) have been reported with thorium abundances in previous studies, and the error is typically of order of $0.15$ dex. It appears that among the r-process enhanced stars with measured Th or other actinide elements (e.g.,U), except for an actinide normal group, which means no obvious enhancement of actinide element abundances with respect to the scaled solar r-process pattern, an actinide boost group (e.g., CS31082-001, HE1219-0312, CS30306-132) also came to be known, although the former group seems to be more common among r-process enhanced metal-poor stars (e.g., CS22892-052, CS29497-004). This imples that for elements in the range of $\\mathrm{Z}\\ge90$, significantly different chemical yields might be produced due to the various conditions of the star formation regions \\citep{hill:02,roederer:09,mashonkina:10}.  To obtain accurate age estimates and to explore r-process nucleosynthesis in metal-poor stars in detail, high-precision Th abundances are needed. Further, large samples of stars with Th abundances are also important to study the distribution of thorium abundances in metal-poor stars. The HERES survey offers a good opportunity to perform such a study. More than 22 elemental abundances (not including Th) for the sample stars have been reported in previously published papers in this series\\citep[e.g.,][]{barklem:05,zhang:10}, and thorium abundances have been studied furtherly with better quality spectra in a few stars \\citep[][]{hayek:09,mashonkina:10}.  Here we report thorium abundances for 77 stars from the sample of \\citet{barklem:05} (hereafter Paper II), using a modified version of the analysis method described in that paper. A brief description of the sample is given in sect. \\ref{sect:obs}. The abundance analysis is described in sect. \\ref{sect:anay}. In sect. \\ref{sect:thorium}, thorium abundances are presented and compared with previous works, and the implications for stellar age estimates are discussed. Section \\ref{sect:conclusion} presents the conclusions. ",
        "conclusions": "\\label{sect:conclusion} Using the ``snapshot'' spectra from HERES survey, we derived the thorium abundances of 77 metal-poor stars, 17 of which have detected Th abundance, while for the rest, only upper limits are available. Thorium abundances cover a wide range of about 4.0 dex, and a scatter exists in the distribution of log(Th/Eu) ratios, supporting previous studies suggesting that the r-process is not universal. For the five common stars, our results are in good agreement with previous studies, which gives us the confidence to present a reliable Th distribution and to discuss r-process pattern for such a large sample of metal-poor stars.  With avaliabe abundances, we explored the correlation between Th and other r-process elements. We confirmed the relatively large scatter in log(Th/Ba) and log(Th/Sr) distributions, and found that a better consistence exists within r-process stars, which confirms that these stars formed from the gas experienced similar nucleosynthesis.  Using log(Th/Eu) as the chronometer, we derived the ages for r-process metal-poor stars, for which both Th and Eu are available. Two stars might be ''actinide boosted'', considering an error to 11.7 Gyr in our age estimation. These stars might have experienced very different chemical enrichments during their formation and evolution. At present the only possible explanation would be that, the r-process elements in these stars were implanted long after the formation of the star, still holding the principle of a universal r-process pattern. It suggests that ``actinide boost'' might not be a rare phenomenon, thus questioning the reliability in using log(Th/$\\mathrm{X}_{\\mathrm{stable}}$) for cosmochronometry to derive stellar ages. For this kind of stars, U abundances are needed to confirm the ''actinide boost'' feature and log(Th/U) can be used to give more accurate age estimations.   However, it should be cautioned again that, $1\\sigma$ detection significance was adopted in this work, due to the weakness of Th II line and the limited quality of our present data, which may lead to significant uncertainty in the thorium abundance. Thus, the detection of Th in stars without enhancement of other n-capture elements and the ''actinide boost'' phenomenon in some stars are not confirmed conclusion. Better quality spectra are needed for higher precision Th and other r-process heavy elements abundance determinations. Thus here we would not like to give any conclusion, but rather to draw attention on those very interesting objects from this work, i.e., HE~2219$-$0713, HE~0305$-$6024, HE~0105$-$6141 and HE~1332$-$0309. Actually we ourselves have already submitted a proposal for the observation time using ESO VLT/UVES for obtaining better quality data of these objects, if some of them can be confirmed, it will be a very important discovery.  \\begin{table*} \\centering \\caption{Stars with detected Th abundances classified into different subclasses of metal-poor stars.} \\label{tab:thdetect} \\begin{tabular}{lcrrrrrr} \\hline\\hline Name & $\\mathrm{T}_\\mathrm{eff}$ &  log$g$        & [Fe/H] & log$\\epsilon$(Th) & [Th/Fe] & log(Th/Eu)& \"Age\" \\\\ & [K]                       & [cm s$^{-2}$]  &        &                   &         &           & [Gyr] \\\\ \\hline &&&&&&&\\\\ \\multicolumn{8}{c}{\\underline{r-I stars}} \\\\ &&&&&&&\\\\ HD~20          & 5445 &  2.39 & $-1.58$ & $-1.06$  & 0.43 & $-0.79$ & 23.8 \\\\ HE~0105$-$6141 & 5218 &  2.83 & $-2.55$ & $-0.87$  & 1.59 & 0.49    & $-38$ \\\\ HE~0240$-$0807 & 4729 &  1.54 & $-2.68$ & $-1.90$  & 0.69 & $-0.47$ & $6.5$ \\\\ HE~0442$-$1234 & 4604 &  1.34 & $-2.41$ & $-1.91$  & 0.41 & $-0.53$ & 9.3 \\\\ HE~0524$-$2055 & 4739 &  1.57 & $-2.58$ & $-1.91$  & 0.58 & $-0.32$ & $-0.5$   \\\\ HE~2252$-$4157 & 5090 &  2.87 & $-1.93$ & $-1.32$  & 0.52 & $-0.43$ &  4.7  \\\\ &&&&&&&\\\\ \\multicolumn{8}{c}{\\underline{r-II stars}} \\\\ &&&&&&&\\\\ CS~22892$-$052 & 4884 &  1.81 & $-2.95$ & $-1.76$  & 1.10    & $-0.86$ &  24.7 \\\\ CS~29497$-$004 & 5013 &  2.23 & $-2.81$ & $-1.17$  & 1.55    & $-0.49$ &  7.5 \\\\ CS~31082$-$001 & 4922 &  1.90 & $-2.78$ & $-1.00$  & 1.69    & $-0.39$ &  2.8 \\\\ HE~2224$+$0143 & 5198 &  2.66 & $-2.58$ & $-1.47$  & 1.02    & $-0.45$ &  5.6 \\\\ HE~2327$-$5642 & 5048 &  2.22 & $-2.95$ & $-1.45$  & 1.41    & $-0.25$ &  $-3.7$ \\\\ &&&&&&&\\\\ \\multicolumn{8}{c}{\\underline{r-normal stars}} \\\\ &&&&&&&\\\\ HE~0051$-$2304 & 4537 & 1.22 & $-2.41$ & $-1.94$ & 0.38 & $-0.19$ & $-6.5$ \\\\ &&&&&&&\\\\ \\multicolumn{8}{c}{\\underline{unclassified stars}} \\\\ &&&&&&&\\\\ HE~0005$-$0002 & 4726 &  1.58 & $-3.09$ & $-2.03$  & 0.97    & -  & - \\\\ HE~0353$-$6024 & 5320 &  3.15 & $-3.17$ & $-1.06$  & 2.02    & -  & - \\\\ HE~1332$-$0309 & 5125 & 2.40 & $-2.46$ & $-0.97$ & 1.40    & 0.5      & $-38.7$ \\\\ HE~2219$-$0713 & 4789 &  1.68 & $-2.91$ & $-1.97$  & 0.85    & -  & - \\\\ HE~2311$+$0129 & 5188 &  2.65 & $-2.78$ & $-1.32$  & 1.37    & -  & - \\\\ \\hline \\end{tabular} \\end{table*}"
    },
    "1112/1112.3107_arXiv.txt": {
        "abstract": "The origin of complex organic molecules such as amino acids and their precursors found in meteorites and comets is unknown.  Previous studies have accounted for the complex organic inventory of the Solar System by aqueous chemistry on warm meteoritic parent bodies, or by accretion of organics formed in the interstellar medium.  This paper proposes a third possibility: that complex organics were created \\textit{in situ} by ultraviolet light from nearby O/B stars irradiating ices already in the Sun's protoplanetary disk.  If the Sun was born in a dense cluster near UV-bright stars, the flux hitting the disk from external stars could be  many orders of magnitude higher than that from the Sun alone. Such photolysis of ices in the laboratory can rapidly produce amino acid precursors and other complex organic molecules. I present a simple model coupling grain growth and UV exposure in a young circumstellar disk.  It is shown that the production may be sufficient to create the Solar System's entire complex organic inventory within $10^6\\ \\rm{yr}$. Subsequent aqueous alteration on meteoritic parent bodies is not ruled out. ",
        "introduction": "The early Solar System was replete with complex organic molecules, which can be seen today in preserved ancient bodies such as meteorites and comets.  Upwards of 100 different amino acids have been detected on chondritic meteorites such as Murchison, Allende and Orguiel.  Many of these have no known natural terrestrial occurrence \\citep{egb01}, and are believed to be of extra-terrestrial origin.  Comets and the interstellar medium (ISM) are also rich inventories of complex organic molecules, including amino acid precursors \\citep{sny06}.  Amino acids and other complex organics have have implications for the origins of life on Earth, so understanding their formation and history is an area of great interest.  If these molecules are formed and distributed easily in a variety of disks and conditions, then pre-biotic compounds may be common throughout distant planetary systems.  Of the complex organics in the solar system, amino acids are particularly interesting to study.  These compounds are necessary for life, and may have been biotic precursors on Earth.  Carbonaceous chondrites are typically 5\\% or more carbon by mass, most of which is in aromatic polymers and thousands of other pre-biotic organic compounds \\citep{sgg10,ha81}.  Laboratory studies have found Murchison to have some 10--30 ppm by mass in identified amino acids \\citep{ss90b}.  The simplest amino acid (glycine, NH$_2$CH$_2$COOH) has recently been detected in comets, but so far has eluded detection in the ISM \\citep{egd09,sny06,kch03}.  Potential amino precursors such as formic acid (HCOOH) have been detected in comets but not the interstellar medium \\citep{hps03,blw00}.  However, amino acids in ice form in the interstellar medium (ISM) would be more stable but harder to detect than their corresponding gas forms in comets \\citep{ebd01,ctb90}, so it is likely that they exist in the ISM but have not yet been detected.   The origin of these molecules is unknown.  Two broad explanations exist for the amino acids in our solar system.  First, they may be produced \\textit{endogenically}, by chemical synthesis within the solar system itself. A variety of energy sources for this exist, including infall heating, radiogenic heating, lightning, and shocks.  For the amino acids present on meteoritic parent bodies, the best-studied endogenic process is Strecker synthesis.  This is a method by which amino acids are formed in aqueous environments, such as the warm sub-surface aquifers that could have been present on asteroids as they were heated by $^{26}$Al and other radionuclides during their first few Myr \\citep{egb01}.  Amino acids are produced in the laboratory by this process, and sufficient $^{26}$Al existed to heat parent bodies to liquid water temperatures \\citep{mgg02}.  However, a serious problem exists with this model.  Isotopic measurements of the chondritic amino acids consistently show deuterium \\textit{enhancements} of $\\delta D$ = 600-2000\\promille, while measurements of the water in these same bodies show deuterium \\textit{depletions} of roughly 100\\promille.  \\textit{In situ} synthesis of aminos does not preferentially change the D/H ratio, so this argues that the water and organics come from distinct sources \\citep{ler97}.  Second, the molecules may have been inherited \\textit{exogenically} from the interstellar medium which formed the solar nebula.  The ISM is known to be rich in complex chemistry, with over 150 different gas-phase species detected to date in molecular clouds \\citep{hv09}.   Amino acids have not been detected in the ISM, but sugars, alcohols, polycyclic aromatic hydrocarbons (PAHs), and other complex molecules have been.  The formation of these compounds in the ISM is thought to be due to a combination of processes, including gas phase, gas-grain, and UV-grain reactions. Once molecules are formed, stable compounds can be incorporated into young disks as new YSOs condense out of the ISM. This pathway is supported by measurements showing the organic composition of comets and the ISM to be quite similar \\citep{blw00,isc00}.  Laboratory results have shown that simple ices can be turned into a zoo of amino acids and other complex organics simply by the presence of UV flux followed by a warming stage \\citep{nab08,ncy07,bds02,mms02}.  Though these studies are performed at fluxes higher than the ISM, their total photon dosages are comparable to the current problem.  Amino acids similar to those seen in meteorites are created.  Moreover, chemical reactions (including photolysis) which occur at cold temperatures $<70\\ \\rm{K}$ can preferentially increase D/H because deuterium's higher mass allows it to bond more readily than H at low temperatures.  D/H enrichments are seen throughout the ISM, and the D/H enrichments in amino acids suggests that they too have a low-temperature origin compared with the water in meteorite parent bodies \\citep{sbd01}. Continued UV exposure may sometimes destroy the same large molecules the UV created earlier, but for a proper range of timescales and fluxes, cold-temperature photolysis of ices may be a more plausible pathway to forming the meteoritic amino acids than warm aqueous synthesis.   The ISM formation model is not without problems.  First, the high optical depth within dense molecular clouds blocks nearly all external UV light.  The only source of UV within the clouds is that caused by cosmic rays interacting with gas in the clouds, resulting in a UV flux of $10^{-4}$--$10^{-5}$~G$_0$, where G$_0$ is the interstellar flux at the Sun today \\citep{pt83}.  Second, some organic molecules, such as the CHON particles in comets like Halley, may be easily destroyed by shocks during infall, although some debate exists on this point \\citep{vvd09,mws93,zg90}.  Finally, some of the common meteoritic amino acids, such as $\\alpha$-aminoisobutyric acid (AIB) and isovaline, have not yet been produced by irradiation in the laboratory \\citep{hmd08}, and it is not known whether this is due to these species only being formed during Strecker synthesis \\citep[as hypothesized by][]{egb01}, or simply the ice experiments not using the proper initial mixtures.  Nevertheless, overall the ISM formation model paints a plausible picture and may well play a role in the solar system's organic history.  This paper proposes a third pathway for the formation of these complex organic molecules, which has not been examined previously.  In our model, the solar nebula forms within a large molecular cloud similar to those in Orion.  The Sun and its disk form completely, and condense and begin to grow.  Within the next 1~Myr, nearby O and B stars turn on, bathing the young disk in UV light.  These UV photons photo-evaporate gas from the disk \\citep{tb05,jhb98}, but also irradiate small ice grains exposed in the disk's outer skin layer.  The simple ices grains, such as H$_2$O, CO$_2$ and NH$_3$, are exposed to the UV and begin to photolyze into more complex species.  Each individual grain is exposed only briefly to UV light; they spend most of their time in the disk's dark, turbulent interior.  Grains continue to grow, gas which has not formed planets is lost due to photo-evaporation and viscous loss, and within 5~Myr a gas-free debris disk is left with its ices enriched in complex organics.  These organics can then be incorporated into comets, asteroids, and the planets. Organics produced in this way could complement, and perhaps greatly exceed, those produced by the other two pathways. This method is similar to the exogenic ISM production in that it relies on UV photochemistry, but at far high flux ($10^6$~G$_0$ vs.  $10^{-5}$~G$_0$), at warmer temperatures, for a much shorter time.  The model allows for subsequent aqueous alteration as seen in the meteoritic record.  Work by \\citet{rob02} proposed that the organic material was created by X-ray irradiation of the disk by the Sun during its T Tauri phase.  They did not present a disk model, but used high-precision D/H measurements of organic and non-organic material to show that D/H fractionation varied with heliocentric distance, as would irradiation. \\citet{rpr06} compared the meteoritic and interstellar D/H values, along with their C-H bond dissociation energies, and concluded that the solar system's D/H enrichment was created \\textit{in situ}, rather than inherited from the ISM. They proposed that the young Sun might provide the necessary UV source; their model did not study the timescales or fluxes involved.  The source of the Earth's organic inventory (as opposed to the Solar nebula's) is a parallel question which has received some attention.  It is believed that although some organics were probably synthesized on the young Earth \\textit{in situ} by lightning or spark discharges \\citep{mu59}, shock heating in the terrestrial atmosphere \\citep{cs92}, warm ocean vents \\citep{cor90}, or any number of other terrestrial processes, far more were probably delivered from external sources such as comets, asteroids, and interplanetary dust particles \\citep{cs92}.  This paper does not further address the origin of life or the Earth's inventory, except to acknowledge that increasing global abundance of organics in the solar nebula probably results in increased delivery to the Earth as well.  This paper uses a simple model to describe the production of complex organics in UV-illuminated disks in a variety of cases.  The results are necessarily general, and do not explain the abundances or species in one particular sample or disk, but provides an initial assessment of the problem.  External UV photolysis has been ignored in almost all previous models, yet it may be one of the most important sources of energy in both the young Solar System and other proto-planetary disks.  The problem is set up in \\S\\ref{sect:background}, and the model described in \\S\\ref{sect:model}.  \\S\\ref{sect:results} contains results, which are discussed in \\S\\ref{sect:discussion}.  Conclusions are in \\S\\ref{sect:conclusions}, and Appendix~A contains a derivation of the simple grain growth model used here. ",
        "conclusions": "\\label{sect:conclusions}  \\b There was sufficient flux from nearby O/B stars in the Sun's birth environment to convert the the majority of the ice mass in the young Solar System and other proto-planetary disks into complex organic molecules within a few $10^5\\ \\rm{yr}$.  Even when considering the disks' high optical depth and rapid grain growth, both of which sequester surface area, enough ice remained exposed for photolysis to continue rapidly.  \\b UV photolysis from external sources is likely to be the source of the meteoritic amino acids, which have D/H ratios difficult to reconcile with the warm, aqueous environments required from Strecker synthesis of these species.  \\b Organic production is predicted to be substantially less for disks that have evolved for more than $10^5$--$10^6$~years before O/B stars have turned on, because the grain growth in this time locks up surface area against photolysis.  \\b Externally driven photochemistry is most important in the outer disk, where grain growth is slow.  The inner disk is warmer, more massive, and sees more rapid growth, and thus processes such as Fischer-Tropsch reactions may dominate the production of organics here.  \\b External UV flux in 1~Myr provides roughly $100\\times$ as much flux as from the central star, and $10^4\\times$ as much energy as gravitational collapse.  This energy is easily usable for chemical reactions because it is directly incident on ice grain surfaces.  \\b Photo-evaporation can occur at the same time as photolysis of ices, but these two processes operate largely independently.  Photo-evaporation removes only the smallest grains, and photolysis can continue long after the gas disk has been removed.  Dense star clusters have historically been thought of as a hazard to planet formation, because the rapid timescales of disk destruction limit the conditions under which planets can form.  Our earlier results showed that under some conditions photo-evaporation may in fact \\textit{speed} the formation of planetesimals \\citep{tb05}.  Our new results here add to this irony, showing that planetary systems that can form in such `hostile' environments may also be be among the richest in organic molecules and and the pre-cursors of life."
    },
    "1112/1112.3088_arXiv.txt": {
        "abstract": "Following the eruption of a filament from a flaring active region, sunward-flowing voids are often seen above developing post-eruption arcades.  First discovered using the soft X-ray telescope aboard \\textit{Yohkoh}, these supra-arcade downflows (SADs) are now an expected observation of extreme ultra-violet (EUV) and soft X-ray coronal imagers and spectrographs (e.g, \\textit{TRACE}, \\textit{SOHO}/SUMER, \\textit{Hinode}/XRT, \\textit{SDO}/AIA).  Observations made prior to the operation of AIA suggested that these plasma voids (which are seen in contrast to bright, high-temperature plasma associated with current sheets) are the cross-sections of evacuated flux tubes retracting from reconnection sites high in the corona.  The high temperature imaging afforded by AIA's 131, 94, and 193~$\\mbox{\\AA}$ channels coupled with the fast temporal cadence allows for unprecedented scrutiny of the voids.  For a flare occurring on 2011 October 22, we provide evidence suggesting that SADs, instead of being the cross-sections of relatively large, evacuated flux tubes, are actually wakes (i.e., trailing regions of low density) created by the retraction of much thinner tubes.  This re-interpretation is a significant shift in the fundamental understanding of SADs, as the features once thought to be identifiable as the shrinking loops themselves now appear to be ``side effects\" of the passage of the loops through the supra-arcade plasma.  In light of the fact that previous measurements have attributed to the shrinking loops characteristics that may instead belong to their wakes, we discuss the implications of this new interpretation on previous parameter estimations, and on reconnection theory. ",
        "introduction": "Introduction}  Magnetic reconnection is a ubiquitously occurring process in the universe and is widely accepted to be important for energy release in solar eruptions.  However, it is expected to occur in regions of low emission measure, and therefore, observations have tended to be indirect (e.g., \\citeauthor{mckenzie_2002} \\citeyear{mckenzie_2002}).  Retracting loops and sunward-flowing plasma voids above post-eruption flare arcades are often observed throughout long duration events using instruments sensitive to emissions from high temperature plasmas or white-light scattering from density structures in the corona (\\textit{TRACE}, \\textit{SOHO}/SUMER, \\textit{SOHO}/LASCO, \\textit{Hinode}/XRT, \\textit{SDO}/AIA) (see \\citeauthor{savage-mckenzie_2011}~\\citeyear{savage-mckenzie_2011}, Figure~1 therein, for example images of these features).  Both supra-arcade downflowing loops (referred to as SADLs in \\citeauthor{savageEA_2010}~\\citeyear{savageEA_2010}, for example) and supra-arcade downflows (SADs: plasma voids) have been interpreted as the outflows created during the re-organization of the magnetic field during the reconnection process.  To reconcile the difference in observational appearance between the plasma voids (SADs) and shrinking loops, a geometrical explanation was offered based on the line of sight to SADs or SADLs above curved polarity inversion lines (\\citeauthor{savage-mckenzie_2011}~\\citeyear{savage-mckenzie_2011}; \\citeauthor{warren-obrien-sheeley_2011} \\citeyear{warren-obrien-sheeley_2011}).  SADs were suggested to be the cross-sections of the retracting post-reconnection flux tubes, viewed along a line of sight tangential to their axes, while the shrinking loops were viewed from a line of sight orthogonal to their axes.  Many voids and shrinking loops have been reported in observations taken prior to the launch of \\textit{SDO}/AIA (\\citeauthor{mckenzie-hudson_1999}~\\citeyear{mckenzie-hudson_1999}; \\citeauthor{mckenzie_2000}~\\citeyear{mckenzie_2000}; \\citeauthor{innes-mckenzie-wang_2003a}~\\citeyear{innes-mckenzie-wang_2003a}; \\citeauthor{asaiEA_2004}~\\citeyear{asaiEA_2004}; \\citeauthor{sheeley-warren-wang_2004}~\\citeyear{sheeley-warren-wang_2004}; \\citeauthor{khan-bain-fletcher_2007}~\\citeyear{khan-bain-fletcher_2007}; \\citeauthor{reeves-seaton-forbes_2008}~\\citeyear{reeves-seaton-forbes_2008}; \\citeauthor{mckenzie-savage_2009}~\\citeyear{mckenzie-savage_2009}; \\citeauthor{savageEA_2010}~\\citeyear{savageEA_2010}; \\citeauthor{savage-mckenzie_2011}~\\citeyear{savage-mckenzie_2011}).  In this Letter, we present new evidence from recent AIA observations of the 2011 October 22 flare (SOL2011-10-22T10:00:00L045C065) that support an alternative interpretation of void-type SADs as \\textit{wakes} behind thin retracting loops, rather than \\textit{cross-sections} of much larger, evacuated loops.  For the purpose of this reference, we will refer to the density depletions that trail behind the shrinking loops as ``wakes\" (i.e., disturbances in the density of the current sheet, apparently resulting from the passage of the shrinking loops).  The density within the thin loops is unknown due to their small size; however, the fact that they are emitting implies that they are not evacuated.  The key observation is depicted in Figure~\\ref{sad_raster} and in the movies accompanying this Letter.    \\begin{figure}[!ht] \\begin{center}  \\framebox{ \\includegraphics[width=.7\\textwidth]{sad_raster.pdf} }  \\caption{Sequence of AIA 131~$\\mbox{\\AA}$ images (reverse-scaled) focusing on the evolution of one SAD.  The top panels show the descent of the leading edge of a large void.  Panel~(a) indicates the position of two initial SADs that appear to merge in subsequent frames.  Around 12:06~UT, two thin retracting loops appear at the leading edge of the void as indicated in the bottom panels, which have been differenced and scaled to enhance motion.  Panel (h) provides the clearest example image of the two leading loops. We refer the reader to the online movies accompanying this Letter for clearer evidence of this process.  The white box in the overlay panel contains two large SADs shown to be regions of density depletion in Figure~\\ref{emmap.fig}.  (Animations of this figure are available in the online Astrophysical Journal.)} \\label{sad_raster} \\end{center} \\end{figure}   To be clear, SADs were also associated with the idea of a wake in the previous interpretation:  The shape of the plasma voids has been likened to that of a tadpole -- typically consisting of an oblong, tear-drop shaped trough at the head of the downflow followed by a thin, often waving ``tail\" (\\citeauthor{savage-mckenzie_2011}~\\citeyear{savage-mckenzie_2011}, Figure~1 therein).  In the traditional interpretation, the head of the flow is the cross-section of a large evacuated flux tube, while the trailing tail has been explained as the wake of this large loop.  In the new interpretation, the tail is still considered part of the wake, but now the head of the flow is also \\textit{completely} a wake and \\textit{not} a loop cross-section.  An unresolved problem with the previous line-of-sight interpretation is the fact that the shrinking loop areas have been consistently measured to be much smaller than the void areas.  Loop areas measured with \\textit{Yohkoh}/SXT, \\textit{Hinode}/XRT, and TRACE range from $\\sim$2 --~18~Mm$^{2}$ (median:  $\\sim$2~Mm$^{2}$) versus the range for void areas of $\\sim$2 -- 89~Mm$^{2}$ (median: $\\sim$25~Mm$^{2}$) \\citep{savage-mckenzie_2011}.  The new interpretation we present here would resolve the disparity in sizes between shrinking loops and downflowing plasma voids; however, these new findings call into question previously-reported parameter estimations such as fluxes and shrinkage energies.  \\clearpage  \\section*{\\label{sads_20111022:obs}Observations}   On 2011 October 22 at about 9:18 UT, \\textit{SDO}/AIA \\citep{lemenEA_2011} observed the slow eruption of a filament from AR 11314 (GOES M1.3) between $\\sim$9:30 and 10:00~UT.  The eruption and arcade development were located between 10 and 20$^{\\circ}$ in front of the west limb.   The passbands of interest for the observations reported herein are 131, 94, and 193~$\\mbox{\\AA}$ because the hot plasma in the supra-arcade region emits sufficiently at those wavelengths and the flows are clearly observable.   The plasma temperatures to which the narrow band AIA 193, 94, and 131~$\\mbox{\\AA}$ bandpasses admit significant response under flaring conditions are approximately 6~\\&~20~MK, 7-10~MK, and 11-14~MK, respectively \\citep{odwyerEA_2010}.  The angular resolution of the telescope is $\\sim$0.6~arcseconds per pixel (corresponding to  $\\sim$435~km).  AIA continuously observes the full Sun in all wavelengths at a 12~second cadence.   \\begin{figure}[!ht] \\begin{center}  \\includegraphics[width=.95\\textwidth]{goes_secchi_flows_9.pdf}  \\caption{\\textit{Top}:  GOES light curve with downflow initial detection times overlaid at top (red: shrinking loops; blue: shrinking loops leading voids).  The shaded region indicates the presence of a fan of hot plasma in the supra-arcade region as seen primarily in the AIA~131, 94, and 335~$\\mbox{\\AA}$ channels.  The large peaks in the light curve after 14:00~UT are due to solar activity from other areas of the solar disk. \\textit{Bottom}:  Sequence of \\textit{STEREO-A}/SECCHI images taken with the 195 and 304~$\\mbox{\\AA}$ filters.  The image times per wavelength with respect to the GOES light curve are indicated by the arrows.  The white dashed curve represents the solar limb as seen from Earth's perspective.  A color version of this figure is available online.} \\label{goes_secchi_flows} \\end{center} \\end{figure}   The top panel of Figure~\\ref{goes_secchi_flows} shows a GOES light curve with downflow initial detection times overlaid at top.  The downflows are classified as either shrinking loops only or loops leading large voids (the latter will be explained below).   The shaded region in Figure~\\ref{goes_secchi_flows} indicates the presence of a fan of hot plasma in the supra-arcade region as seen in the AIA~131, 94, and 335~$\\mbox{\\AA}$ channels (and to a lesser degree in the 193~$\\mbox{\\AA}$ bandpass).  The shrinking loops are primarily observed prior to the presence of this plasma while the voids are only observed flowing through the fully-developed fan.  As the plasma is filling the region after $\\sim$11:30~UT, wakes (for want of a better term) become marginally observable behind the shrinking loops.  If the wakes are present prior to this time period, they are undetectable.    The orientation of the footpoint ribbons and post-eruption arcade, as seen with the \\textit{STEREO-A}/SECCHI 195 and 304~$\\mbox{\\AA}$ filters, is depicted in the bottom panel of Figure~\\ref{goes_secchi_flows}.  The times for each wavelength panel with respect to the GOES light curve are indicated by the arrows.  This sequence of images shows that there is no significant change in the orientation of the arcade's axis throughout the event.   In addition, the arcade's axis appears to be simple and straight, not allowing for a single line of sight (as viewed from Earth) to be simultaneously tangent \\textit{and} perpendicular to the polarity inversion line.  \\clearpage ",
        "conclusions": "Discussion}   The increased temporal cadence, continual full disk field of view, and temperature coverage of the \\textit{SDO}/AIA observatory has provided new information upon which to base the present re-interpretation of SADs.  Namely, SADs should no longer be considered cross-sections of newly-reconnected, large evacuated flux tubes; instead, these descending voids are apparently density depletions left in the wake of thin flux tubes retracting from a reconnection site in the supra-arcade region (Figure~\\ref{diagram_reinterpret_sads_sadls}).  There is the possibility of a high-beta plasma in the sheet, which is suggested by the observed vortical eddy-like motions, though concrete supporting evidence for high beta is still lacking.  Alternatively, the density depletions which we observe could be considered to result from the wave-interaction effects modeled by \\cite{costaEA_2009}, \\cite{schulzEA_2010}, and \\cite{maglioneEA_2011}; or from plasma flows initiated/accelerated by the passage of shrinking loops.  Both of these effects are found in models of low-beta plasma.  The latter, field-aligned flows resulting in rarefaction behind the shrinking loops, is the subject of modeling currently underway.   \\begin{figure*}[!ht] \\begin{center}  \\includegraphics[width=.95\\textwidth]{diagram_reinterpret_sads_sadls.pdf}  \\caption{(a)  Based on new evidence, the interpretation of SADs as the cross-sections of large, evacuated retracting flux tubes is no longer supported.  Rather, they appear to be the wakes of much thinner shrinking loops.  (b) Only shrinking loops are observed during the early phase of the eruptive event while SADs  become apparent after there is a significant increase in hot plasma (c), presumably surrounding the current sheet, in the supra-arcade region.  Note that the viewing orientation does not need to change to observe the two features; however, the increase in supra-arcade plasma is necessary to observe SADs (i.e., the wakes).  A color version of this figure is available online.} \\label{diagram_reinterpret_sads_sadls} \\end{center} \\end{figure*}   In light of this development, previous parameter estimations of SAD characteristics require amending and re-evaluation.  For example, the sizes, fluxes, and shrinkage energy of post-reconnection flux tubes based on SADs measurements (\\citeauthor{savage-mckenzie_2011} \\citeyear{savage-mckenzie_2011}; \\citeauthor{mckenzie-savage_2011} \\citeyear{mckenzie-savage_2011}) are all over-estimates because the leading loop is much smaller in diameter than the void.  The relationship of a void area to the size of the associated flux tube, or to its flux, is not known at present.  However, we note that the smaller SADs tend to travel straight down into the arcade while the larger ones approach the arcade from an angle.   The presence of apparent wakes, and other complex flows revealed by the AIA data, are unexpected in the magnetically structured plasma of the supra-arcade region, and are not predicted in any current model of post-CME current sheet formation of which we are aware.  They raise questions about current sheet characteristics such as the plasma beta that surely must be investigated, but which are beyond the scope of the present Letter.  As aforementioned, magneto-hydrodynamic simulations of wake creation are underway.  While these observations indicate a fundamental difference from the original interpretation of SADs, the plasma voids are still inherently associated with reconnection outflows and thus remain probes of the reconnection process:  1)  The presence of downflows in the flare decay phase illustrates that reconnection continues for many hours after the initial eruption.  Their presence in the impulsive phase, coinciding with non-thermal emissions (hard X-ray and microwave bursts), demonstrates that the SADs are associated with energy release (\\citeauthor{khan-bain-fletcher_2007} \\citeyear{khan-bain-fletcher_2007}; \\citeauthor{asaiEA_2004} \\citeyear{asaiEA_2004}).  2)  The discreteness of SADs indicates that the reconnection is highly localized (i.e., ``patchy\").  Similarly, temporal variation in the appearance of SADs indicates their production is bursty.  The reconnection in each location turns on and off on short timescales, independently of other locations \\citep{linton-longcope_2006}.  3)  The speeds of SADs and SADLs are approximately 50$-$500~km~s$^{-1}$, and are constant or slightly decelerating.  For comparison, typically assumed Alfv\\'{e}n speeds are on the order of 1000~km~s$^{-1}$.  Reconnection models typically predict outflow speeds of (0.3$-$1.0)~$\\times~v_{A}$ \\citep{linton-longcope_2006}.  Deceleration may be expected, due to buildup of downstream magnetic pressure, though drag mechanisms may also be considered \\citep{savage-mckenzie_2011}.   The revised relationship between SADs and shrinking loops is depicted in Figure~\\ref{diagram_reinterpret_sads_sadls}.  The only distinguishing characteristic necessary between the two observational circumstances is the amount of hot plasma in the supra-arcade region surrounding the current sheet.  A change in orientation may indeed occur; however, SADs will \\textit{only} be observed in the presence of this plasma.  The sizes of the shrinking loops ($\\sim$2$-$3 pixels [0.9$-$1.3~Mm]) are consistent with measurements for shrinking loops (i.e., SADLs) given in \\cite{savage-mckenzie_2011}; however, many of the SADs are too small to detect a leading flux tube with \\textit{SDO}/AIA, which is currently the highest resolution solar observatory capable of coronal measurements.  Therefore, the relevant size scales for post-eruption current sheet reconnection (i.e., the physical size of reconnection patches) appear to be less than $\\sim$435~km (\\textless~1~pixel).  This long-duration event was also observed by RHESSI and \\textit{SOHO}/LASCO; analysis of those data, which pertains to aspects of the flare other than the fundamental nature of the SADs, will be described separately in a forthcoming paper."
    },
    "1112/1112.4661_arXiv.txt": {
        "abstract": "Cosmological measurements are affected by the energy density of both active and sterile massive neutrinos. We extend here a recent analysis of current cosmological data to non minimal cosmologies. Several possible scenarios are examined: a constant $w \\neq -1$ dark energy equation of state, a non flat universe, a time varying dark energy component and coupled dark matter dark energy universes or modified gravity scenarios. When considering cosmological data only, (3+2) massive neutrino models with $\\sim 0.5$~eV sterile species are allowed at $95\\%$ CL. This scenario has been shown to reconcile reactor, LSND and MiniBooNE positive signals with null results from other searches. Big Bang Nucleosynthesis bounds could compromise the viability of (3+2) models if the two sterile species are fully thermalized states at decoupling. ",
        "introduction": "Solar, atmospheric, reactor, and accelerator neutrinos have provided compelling evidence for the existence of neutrino oscillations, implying non-zero neutrino masses (see Ref.~\\cite{GonzalezGarcia:2007ib} and references therein). The present data require the number of massive neutrinos to be equal or larger than two, since there are at least two mass squared differences ($\\Delta m^2_{\\rm atmos}$ and $\\Delta m^2_{\\rm solar}$) driving the atmospheric and solar neutrino oscillations respectively. Unfortunately, oscillation experiments only provide bounds on the neutrino mass squared differences, i.e. they are not sensitive to the overall neutrino mass scale.   Cosmology provides one of the means to tackle the absolute scale of neutrino masses. Neutrinos can leave key signatures in several cosmological data sets. The amount of primordial relativistic neutrinos changes the epoch of the matter- radiation equality, leaving an imprint on both Cosmic Microwave Background (CMB) anisotropies (through the so-called Integrated Sachs-Wolfe effect) and on structure formation, while non relativistic neutrinos in the recent Universe suppress the growth of matter density fluctuations and galaxy clustering, see Ref.~\\cite{Lesgourgues:2006nd}. Cosmology can therefore weigh neutrinos, providing an upper bound on the sum of the three active neutrino masses, $\\sum m_\\nu \\sim 0.58$~eV at $95\\%$ CL~\\cite{Komatsu:2010fb}. The former bound is found when CMB measurements from the Wilkinson Microwave Anisotropy Probe (WMAP) are combined with measurements of the distribution of galaxies (SDSS-II-BAO) and of the Hubble constant $H_0$ (HST)~\\footnote{For other recent analyses, see also Refs.~\\cite{Reid:2009nq,Hamann:2010pw}.} in the assumption of a flat universe with a cosmological constant, i.e. a $\\Lambda$CDM cosmology.   There is no fundamental symmetry in nature forcing  a definite number of right-handed (sterile) neutrino species, as those are allowed in the Standard Model fermion content. Indeed, cosmological probes have been extensively used to set bounds on the relativistic energy density of the universe in terms of the effective number of neutrinos $N_\\nu^{\\rm eff}$ (see, for instance, Refs.~\\cite{Mangano:2006ur,Hamann:2007pi,Reid:2009nq,Hamann:2010pw,Mangano:2010ei,Komatsu:2010fb,darkr}. If the effective number of neutrinos $N_\\nu^{\\rm eff}$ is larger than the Standard Model prediction of $N_\\nu^{\\rm eff}=3.046$ at the Big Bang Nucleosynthesis (BBN) era, the relativistic degrees of freedom, and, consequently, the Hubble expansion rate will also be larger causing weak interactions to become uneffective earlier. This will lead to a larger neutron-to-proton ratio and will change the standard BBN predictions for light element abundances. Combining Deuterium and $^4$He data, the authors of Ref.~\\cite{Mangano:2006ur} found $N_\\nu^{\\rm eff}=3.1^{+1.4}_{-1.2}$ at $95\\%$ CL.  Models with one additional $\\sim 1$~eV massive sterile neutrino, i.e. the so called (3+1) models, were introduced to explain LSND short baseline (SBL) antineutrino data~\\cite{Aguilar:2001ty} by means of neutrino oscillations. A much better fit to SBL appearance data and, to a lesser extent, to disappearance data, is provided by models with two sterile neutrinos (3+2)~\\cite{Sorel:2003hf,Karagiorgi:2006jf} which can also explain both the MiniBooNE neutrino~\\cite{AguilarArevalo:2007it} and antineutrino data~\\cite{AguilarArevalo:2009xn} if CP violation is allowed~\\cite{Karagiorgi:2009nb}. More recently, a combined analysis including the new reactor antineutrino fluxes~\\cite{Mueller:2011nm,Mention:2011rk} has shown that (3+2) models provide a very good fit to short baseline data~\\cite{Kopp:2011qd}. While these models with extra sterile species show some tension with BBN bounds on $N_\\nu^{\\rm eff}$, the extra sterile neutrinos do not necessarily have to feature thermal abundances at decoupling, see Refs.~\\cite{Melchiorri:2008gq,Acero:2008rh}, where the usual full thermalization scenario for the sterile neutrino species was not assumed. Up to date cosmological constraints on massive sterile and active neutrino species have been presented in Refs.~\\cite{rt,us} in the context of a $\\Lambda$CDM universe. It is well known that bounds on active neutrino species are relaxed if the dark energy equation of state is different from $-1$~\\cite{Hannestad:2005gj,Hamann:2011ge,Calabrese:2011hg} and/or interactions between the dark matter and dark energy sectors are switched on~\\cite{LaVacca:2008mh,Gavela:2009cy}. In the same line, the authors of Ref.~\\cite{Kristiansen:2011mp} have found that in models with non zero curvature and two extra sterile neutrinos the cosmological constant scenario is ruled out at $95\\%$ CL.  Here we extend the minimal cosmological scenario considered in our previous study~\\cite{us} and compute the bounds on the masses of the active and the sterile neutrino states as well as on the number of sterile states in the presence of a constant equation of state $w\\neq1$, a time varying dark energy fluid, a non vanishing curvature component and interactions among the dark sectors. The paper is organized as follows. Section \\ref{sec:i} describes the details of the analysis carried out here, including the cosmological parameters and datasets. The four different cosmological scenarios explored here are analyzed and the most important degeneracies among the neutrino parameters are carefully explored. Section \\ref{sec:seciii} summarizes our main results and conclusions. ",
        "conclusions": "\\label{sec:seciii} Neutrino oscillation experiments indicate that neutrinos have non zero masses and open the possibility for a number of extra sterile neutrino species. LSND and MiniBooNE antineutrino data require these extra sterile species to be massive. Much effort has been devoted in the literature to constrain the so called (3+1) (three active plus one sterile) and (3+2) (three active plus two sterile). Recently, global fit analyses incorporating new reactor antineutrino fluxes have shown that (3+2) models with $0.5-1$~eV sterile species provide a very good fit to short baseline data. Cosmology can set bounds on both the active and sterile neutrino masses as well as on the number of sterile neutrino species. We have explored here the current constraints on these parameters in natural extensions of the minimal $\\Lambda$CDM cosmology. Namely, we have explored the neutrino constraints in scenarios without a cosmological constant as the dark energy fluid, with a non vanishing curvature, or with coupled dark matter-dark energy fluids. Figure \\ref{fig:final} summarizes our results in the $m_\\nu-N_{{\\nu}_s}$ and $N_{{\\nu}_s}-m_{{\\nu}_s}$ planes. Notice that models with two massive  $0.5-1$~eV sterile neutrinos plus three sub-eV active states are perfectly allowed at $95\\%$ CL by current Cosmic Microwave Background, galaxy clustering and Supernovae Ia data. Interestingly, these models are precisely the ones which, with the new reactor fluxes prediction, improve considerably the global fit to short baseline data. While we have not checked directly the results of Ref.~\\cite{Kristiansen:2011mp} in which models with non zero curvature and two extra sterile neutrinos seem to exclude $w=-1$ at $95\\%$ CL, it seems plausible that non standard $\\Lambda$CDM cosmologies with sterile neutrino species provide also a very good fit to cosmological data. Big Bang Nucleosynthesis constraints could compromise the viability of these models if the two sterile neutrino states are fully thermal. If the BBN constraints are obtained using the helium fraction measurements $Y_p$ from \\cite{aver} exclusively, two extra sterile neutrino states are perfectly allowed ($N_{{\\nu}_s} < 2.3$ at 95\\% CL). Even after combining these data with Deuterium measurements, two massive neutrino states are only marginally excluded  ($N_{{\\nu}_s} < 1.7$ at 95\\% CL). The tightest bound on the number of sterile neutrino species arises when helium measurements from Ref.~\\cite{it} are combined with Deuterium data. Further developments in BBN determinations of light element abundances may have a large impact in further constraining the number of sterile neutrino species $N_{{\\nu}_s}$.  \\begin{figure} \\begin{tabular}{c} \\includegraphics[width=7.5cm]{nnus_mnu_all.eps}\\\\ \\includegraphics[width=7.5cm]{nnus_mnus_all.eps}\\\\ \\end{tabular} \\caption{The upper panel shows the 68\\%  and 95\\% CL bounds from ``run 1'' in the $m_\\nu-N_{{\\nu}_s}$ plane. The lower plane shows the analogous in the $N_{{\\nu}_s}-m_{{\\nu}_s}$ plane. The masses of the sterile and active neutrinos are both in eV. Red, blue, magenta and green contours denote $w$CDM, $w(a)$CDM, $\\Omega_k$CDM and $\\xi$CDM cosmologies, respectively.} \\label{fig:final} \\end{figure}"
    },
    "1112/1112.0274_arXiv.txt": {
        "abstract": "The formation of giant planets requires accumulation of $\\sim$10 Earth mass in solids; but how do protoplanets acquire their mass?  There are many, often competing processes that regulate the accretion rate of protoplanets. To assess their effects we present a new, publicly-available toy model.  The rationale behind the toy model is that it encompasses as many physically-relevant processes as possible, but at the same time does not compromise its simplicity, speed, and physical insight.  The toy model follows a modular structure, where key features -- \\eg\\ planetesimal fragmentation, radial orbital decay, nebula turbulence -- can be \\textit{switched} on or off.  Our model assumes three discrete components (fragments, planetesimals, and embryos) and is zero dimensional in space.  We have tested the outcomes of the toy model against literature results and generally find satisfactory agreement.  We include, for the first time, model features that capture the three-way interactions among small particles, gas, and protoplanets.  Collisions among planetesimals will result in fragmentation, transferring a substantial amount of the solid mass to small particles, which couple strongly to the gas.  Our results indicate that the efficiency of the accretion process then becomes very sensitive to the gas properties -- especially to the turbulent state and the magnitude of the disk headwind (the decrease of the orbital velocity of the gas with respect to Keplerian) -- as well as to the characteristic fragment size. ",
        "introduction": "\\labelH{sec:Intro} Giant, gas-rich planets are not formed in one day, but their formation should be  complete within the $\\sim$10$^6$ yr over which the gas-rich phase of protoplanetary disks last \\citep{FedeleEtal2010}.  As there are about 100 $e$-foldings required to grow a sufficiently massive core ($\\sim$$10\\ M_\\oplus$; \\citealt{PollackEtal1996}) out of ISM dust grains ($m\\approx10^{-17}$ g), this nevertheless represents a fast process. Not only is the process fast: the planet formation mechanism also seems to be very efficient.  Over the years surveys like \\textsc{Harps} and \\textsc{Kepler} have discovered a wide number and variety of exoplanets \\citep{MayorEtal2010,HowardEtal2011}.  Planet formation may be ubiquitous.  The view that the initial stage of giant planet formation proceeds similarly to that of terrestrial planet formation is commonly held.  This is the core accretion paradigm \\citep{MizunoEtal1978,PollackEtal1996}: a rocky core is formed first, whereafter it binds the nebular gas, when it becomes sufficiently massive.  Another paradigm for the formation of massive planets is the disk-instability model, where the gaseous disk is sufficiently massive to become gravitationally unstable \\citep[\\eg][]{MayerEtal2007,Boley2009,Boss2011}.  Here, we will consider the core accretion paradigm and focus exclusively on the problem of forming a $\\sim$$10\\ M_\\oplus$ core quickly enough.  An important intermediate step is the formation of planetesimals -- gravitationally interacting bodies usually thought to be of km-size or larger \\citep{Safronov1969}.  These are the building blocks for planets.  Indeed, a typical assumption in protoplanet growth studies is that the majority of solids (\\ie\\ the rock and ice) of the protoplanetary disk resides in planetesimals.  But how planetesimals form in the first place remains ill-understood. Models favoring direct sticking by dust particles \\citep{Weidenschilling1997}, have recently been confronted with a variety of obstacles: the fragmentation barrier at m-sizes \\citep[\\eg][]{BirnstielEtal2010i}; the bouncing barrier at mm-sizes \\citep{GuettlerEtal2010,ZsomEtal2010,ZsomEtal2011}; or a charging barrier for even smaller (and fluffier) particles \\citep{Okuzumi2009}.  Perhaps these results indicate that planetesimal formation requires special conditions: \\eg\\ sticky material properties (like ice; \\citealt{WadaEtal2009}); particle pileups in drift-free regions \\citep{YoudinChiang2004,KretkeLin2007,BrauerEtal2008i}; concentration of chondrule-size particles and their subsequent gravitational collapse \\citep{CuzziEtal2010,Chambers2010}; or the shearing instability that operates on a dense layer of dm-m size boulders \\citep{YoudinGoodman2005,JohansenEtal2007,JohansenEtal2009}.  The transition to planetesimal size bodies, instigates the runaway and oligarchic growth phases \\citep{IdaMakino1993,KokuboIda1996,KokuboIda1998,OrmelEtal2010}.  Runaway growth is fast and a protoplanet seed forms \\citep{WetherillStewart1989}.  However, at some point the runaway body faces a backlash from its own viscous stirring, thereby allowing embryos\\footnote{We will employ the phrasing `oligarchs', `core', `embryo', and `protoplanet' mostly as synonyms.  More correctly, an embryo or protoplanet consist of a rocky (or icy) core with a gaseous atmosphere.} in neighboring zones to catch up in terms of mass.  What follows is a two component system of embryos and planetesimals, with the latter containing most of the solid mass.  The core formation challenge then amounts to transferring the mass reservoir of planetesimals onto these embryos.  There have been recent reviews of the runaway/oligarchic growth stage \\citep[\\eg][]{GoldreichEtal2004,LevisonEtal2010,OrmelEtal2010i} and we will not repeat these here, but the underlying problem is that during oligarchic growth embryos dynamically stir the planetesimals faster than they can accrete them.  For this reason, the next $e$-folding in embryo mass always takes longer than the previous. Allowing for embryo atmospheres \\citep{InabaIkoma2003}, which enhances the radius at which planetesimals will be captured, alleviates the problem somewhat, but generally modelers require massive disks -- more massive than the minimum mass solar nebula (MMSN; \\citealt{Weidenschilling1977i,HayashiEtal1985}) benchmark -- to form massive protoplanets within the time that the gas disk disperses.  \\begin{figure*}[t] \\plotone{fig1.eps} \\caption{Sketch of processes affecting the dynamical state (\\textit{red} arrows) and the surface density (\\textit{black} arrows) of the components. Each of the three components is characterized by a surface density $\\Sigma_k$, size, $s_k$ (or $R_E$ for the embryo's), and excitation state $e_k$ (the embryos' eccentricity is by definition 0).  \\textit{Dashed} arrows denote external processes like gas drag and drift motions.  The random motion (eccentricity) of the planetesimal and fragment population is increased by viscous stirring and turbulent stirring, and damped by gas drag.  The surface density in embryos increases due to accretion.   The surface density in fragments and planetesimals changes to accretion, mutual collisions (erosion and fragmentation), radial drift, and inflow of solids from external regions.  Collisions within the planetesimal and fragment population changes the characteristic size ($s_P$ and $s_F$) of these bodies (\\textit{blue} arrows).} \\labelH{fig:sketch} \\end{figure*} However, planetesimal-planetesimal collisions cause a copious amount of debris, especially for massive disks.  The question is whether these fragments -- or small particles in general -- are more conducive to growth than their km-size progenitors.  This is not an easy question to answer since smaller particles interact more strongly, and in more diverse ways, with the gas.  On the positive side, gas damping is much stronger, reducing their eccentricities.  Similarly, the dissipative nature of gas drag enhances capture rates.  However, orbital decay timescales can be extremely short \\citep{Weidenschilling1977} and diffusion by turbulence could severely dilute their midplane densities.  One of the foci of this work is to capture these key physical processes into a common framework.  The two component approximation with which oligarchic growth is characterized renders it ideal to be studied by simple numerical or even analytical models \\citep{GoldreichEtal2004,Chambers2006,Chambers2008,BruniniBenvenuto2008,GuileraEtal2011}. With a `toy model' (TM) we mean anything where the mass distribution is approximated by two or three (in case fragments are included) components.  This in contrast to models that include the full mass spectrum \\citep[\\eg][]{Weidenschilling1997,KobayashiEtal2010,BromleyKenyon2011i}, which may be referred to as `sophisticated'.  The goal of a TM may be to provide physical insight or to allow for a quick (but still sufficiently accurate) computation.  Population synthesis models, for example, \\citet{IdaLin2004}, \\citet{MordasiniEtal2009}, and their sequels employ a TM for the embryo growth stage, if not for all stages.  But do TMs capture all relevant physics?  There is often a tendency for the (above) TM to include more sophisticated features, \\eg\\ a radial dimension to follow the migration of protoplanets or multiple mass bins to capture the collisional cascade.  In this work we will attempt to adhere more strictly to what we consider characterizes a TM. Specifically this demands the model to be: \\begin{enumerate} \\item Quick.  CPU speeds should be seconds, rather than hours, let alone days. \\item General. That is, it should include all relevant physical process that could potentially affect protoplanet growth; \\item Transparent. It must be clear which effect causes what. \\end{enumerate} Combining the second and third properties is challenging, since a more complete model, characterized by more features, is often also a more complex one.  For that reason, we have opted for a modular approach, where features can be \\textit{switched} on or off.  As a result of the first and third property, our TM contains many simplification, perhaps \\textit{over}simplifications.  However, in our view, a TM is complementary, not in competition, to more sophisticated models of planet formation.  That is, the TM provides a first glance of what result may be expected.  It can explore the parameter space most efficiently and assesses whether or not novel features do matter.  It may not get the details correct; its results may even be off by several factors of unities.  But here it is where it can work in conjunction with sophisticated models.  In this paper, the emphasis lies on the presentation and validation of the toy model to follow protoplanet growth.  In \\se{toy-model} we present the toy model.  In \\se{tests} we validate it against several literature studies. In \\se{applic} the novel features of the toy model are explored.  We summarize and conclude in \\se{summ}. ",
        "conclusions": "\\label{sec:summ} We have presented a new model to follow the protoplanet growth stage in planet formation.  The model is a toy model: it assumes that the mass distribution can be approximated by three components (embryos, planetesimals, and fragments) and does not involve a radial dimension. To obtain the equations that govern the evolution, we have employed simple physically-based recipes.  Yet the toy model is intended to be general: a wide array of physical relevant processes can be included at the discretion of the user.  In particular, we have focused on a more accurate treatment of fragments, for which coupling to gas is important.  Due to our modular setup (the switches) the toy model is intended to expose the key mechanisms that drive protoplanet growth.  One should be aware of the limitations, as we have outlined in \\se{overview}.  The (at most) three component approximations entails that the toy model cannot model the runaway growth phase (the phase preceding oligarchy) and does a bad job in modeling the fragmentation cascade.  The oligarchic assumptions will also break down when the mass in embryos ($\\Sigma_E$) starts to exceed that of planetesimals and fragments -- a situation that may be reached sooner rather than later in cases where drift is important.  Furthermore, as of now, type I migration, or planetesimal driven migration \\citep[][Ormel \\& Ida in prep]{KirshEtal2009,CapobiancoEtal2011} are not included either. Our model is in it strictest sense only locally valid.  We have shown in \\se{tests} that our results are in good agreement with those of \\citet{KobayashiEtal2010,KobayashiEtal2011}.  These calculations indicate that planetesimal fragmentation (a natural outcome of oligarchic growth) is quite harmful to protoplanet growth, due to their strong orbital decay.  This conclusion contrasts with many earlier works, which suggest that small planetesimals, or fragments, are very conducive for growth \\citep{Rafikov2004,GoldreichEtal2004,Chambers2006,KenyonBromley2009,GuileraEtal2011}.  However, \\citet{Chambers2008} already warned that fragmentation acts like a double-edged sword -- increasing collision rates but at the same time depleting the disk of solids due to orbital decay -- whereas \\citet{KobayashiEtal2011} finds that fragmentation is detrimental to embryo growth.  For this reason, they prefer bigger planetesimals since these are stronger and less collisional -- both effects reduce the fragmentation rate.  But the disadvantage of big planetesimals is that they are damped poorly and growth is relatively slow.  But it may still be premature to provide a final answer to the question whether small planetesimals, or small fragments, promote or hinder (fast) growth.  Very small particles have the additional handicap that turbulence stirs them up as we saw in \\fg{jupiter}b and \\fg{ext-accret}.  Few works concerning protoplanet growth have accounted for turbulent stirring (\\citealt{BromleyKenyon2011i} are a recent exception), which is somewhat peculiar since $\\alpha_\\mathrm{ss}$ is a well known -- and critical -- parameter for the pre-planetesimal growth state \\citep[\\eg][]{BirnstielEtal2010i,ZsomEtal2011}.  Indeed, from an observation perspective mm/cm particles dominate (the mass) of protoplanetary disks and are the likely precursors of the first generation of big planetesimals c.q.\\ protoplanet seeds \\citep{JohansenEtal2007,JohansenEtal2009,CuzziEtal2010}.  Therefore, we better investigate the accretion behavior of small particles (\\citealt{JohansenLacerda2010,OrmelKlahr2010}, \\fg{ext-accret}). Here we reemphasize that our treatment of atmosphere enhancements (\\texttt{AE}) and nebular drag effects (\\texttt{ND}) is somewhat `ad-hoc' (see our comments towards the end of \\se{atmos-enh}).  In a follow-up study we intend to return to the question regarding the efficiency of accretion of (big) planetesimal \\vs\\ (small) fragment.  The result of this work offers some tentative clues, though.  First, the efficiency scales with the magnitude of the headwind $v_\\mathrm{hw}$ -- the deviation of the orbital velocity of the gas from Keplerian.  It seems very hard to grow big protoplanets ($\\sim$10 $M_\\oplus$) out of small particles when $v_\\mathrm{hw}$ is not reduced or with some form of replenishment of solids.  Interestingly, a similar conclusion may hold for planetesimal formation itself \\citep{CuzziEtal2010,BaiStone2010ii}.  A second important point is that the accretion efficiency is strongly size dependent (\\fg{ext-accret}). (Sub)km-size planetesimals accrete well, but are prone to collisional fragmentation.  The resulting m-size fragments fare much worse though in terms of accretion efficiency,  despite the fact that they have a large specific accretion rate.  Strong radial drift for these particles is the limiting factor. The accretion efficiency of smaller particles strongly depends on the turbulent state of the disk.  Overall, these sensitivities indicate that accretion is very much dependent on the condition regarding the gaseous state (pressure gradient, turbulence) and on the material strength of the particles, which to a large extent determines the typical fragment size.  In the end, any growth scenario must be tested against observational constraints, \\eg\\ the size distribution of the asteroid \\citep{MorbidelliEtal2009,Weidenschilling2011}, dust production rates in debris disks \\citep{ShannonWu2011,KenyonBromley2008,KenyonBromley2010}, or the exoplanet census \\citep{AlibertEtal2011,IdaLin2010}.  For detailed comparison, however, our toy model with its discrete, three component approximation is a rather crude tool.  But it is a great tool to conduct a preliminary exploration of the vast parameter space. This provides the user with a first, crude comparison with the observational data and a means to identify the relevant physical processes that are at work, thereby guiding further sophisticated modeling efforts.  We believe that such an interplay between a toy- and sophisticated model is a very powerful asset."
    },
    "1112/1112.6301.txt": {
        "abstract": " ",
        "introduction": "%\u0081yintroduction\u0082\u00cc\u008d\u009c\u0091g\u0082\u00dd\u0081z\\\\ %\u0081E\u008cn\u008aO\u0098f\u0090\u00af\\\\ %\u0081E\u008cn\u008aO\u0098f\u0090\u00af\u0082\u00cc\u0092\u00bc\u0090\u00da\u008a\u00cf\u0091\u00aa\\\\ %\u0081E\u0083R\u0083\u008d\u0083i\u0083O\u0083\u0089\u0083t\\\\ %\u0081E\u0083o\u0083C\u0083i\u0083\u008a\u0093\u00b5\u0083}\u0083X\u0083N\u0083R\u0083\u008d\u0083i\u0083O\u0083\u0089\u0083t\\\\ %\u0081E\u0090\u00e6\u008ds\u008c\u00a4\u008b\u0086\u0081i\u0090\u00a2\u008aE\u0082\u00cc\u0081\u0095\u0089\u00e4\u0081X\u0082\u00cc\u0081j\\\\ %\u0081E\u0096{\u008c\u00a4\u008b\u0086\u0081i\u0096\u00da\u0093I\u0081E\u0088\u00d3\u008b`\u0081E\u0088\u00ca\u0092u\u0082\u00c3\u0082\u00af\u0082\u00f0\u0096\u00be\u008am\u0082\u00c9\u0082\u00b7\u0082\u00e9\u0081j\\\\ ",
        "conclusions": ""
    },
    "1112/1112.0568_arXiv.txt": {
        "abstract": "As a neutron star (NS) is tidally disrupted by a black hole (BH) companion at the end of a BH--NS binary inspiral, its magnetic fields will be stretched and amplified.  If sufficiently strong, these magnetic fields may impact the gravitational waveforms, merger evolution and mass of the remnant disk.  Formation of highly-collimated magnetic field lines in the disk+spinning BH remnant may launch relativistic jets, providing the engine for a short-hard GRB.  We analyze this scenario through fully general relativistic, magnetohydrodynamic (GRMHD) BHNS simulations from inspiral through merger and disk formation. Different initial magnetic field configurations and strengths are chosen for the NS interior for both nonspinning and moderately spinning ($a_{\\rm BH}/M_{\\rm BH}$=0.75) BHs aligned with the orbital angular momentum. Only strong interior ($B_{\\max} \\sim10^{17}$G) initial magnetic fields in the NS significantly influence merger dynamics, enhancing the remnant disk mass by 100\\% and 40\\% in the nonspinning and spinning BH cases, respectively.  However, detecting the imprint of even a strong magnetic field may be challenging for Advanced LIGO.  Though there is no evidence of mass outflows or magnetic field collimation during the preliminary simulations we have performed, higher resolution, coupled with longer disk evolutions and different initial magnetic field configurations, may be required to definitively assess the possibility of BHNS binaries as short-hard GRB progenitors. ",
        "introduction": "With the first direct detection of gravitational waves (GWs) expected in the next few years, numerical relativity simulations will be crucial for distinguishing different GW sources from one another.  Mergers of black hole-neutron star (BHNS) binaries are among the most promising sources of gravitational waves detectable by ground-based laser interferometers like LIGO~\\cite{LIGO1,LIGO2}, VIRGO~\\cite{VIRGO1,VIRGO2}, GEO~\\cite{GEO}, LCGT~\\cite{LCGT}, and AIGO~\\cite{aigo}, as well as by the proposed space-based LISA-like interferometers~\\cite{LISA} and DECIGO~\\cite{DECIGO} and third-generation ground-based detectors such as the Einstein telescope~\\cite{ET1,ET2}.  Analysis of gravitational waveforms from BHNS mergers may spark new insights into the behavior of matter at nuclear densities.  Theoretical models indicate that a neutron star-neutron star (NSNS)~\\cite{HMNS1,HMNS2,ShiUIUC,STU1,STU2,ST,SST,2008PhRvD..78h4033B} or BHNS~\\cite{FBSTR,FBST,SU1,SU2,SST,ST07,eflstb08,elsb09} binary merger may result in a hot, massive disk around a black hole, whose temperatures and densities could be sufficient to trigger a short-hard gamma-ray burst (SGRB).  Indeed, SGRBs have been associated with galaxies with extremely low star formation rates (see~\\cite{SGRB_local} and references therein for a review), indicating that the source is likely to involve an evolved population, rather than main sequence stars. The number of detectable BHNS mergers in the observable universe is still an open question, due to uncertainties in population synthesis calculations. The estimated event rate of BHNS mergers observable by an Advanced LIGO detector typically falls in the range ${\\cal R}\\sim$ 0.2--300~${\\rm yr}^{-1}$~\\cite{KBKOW,aetal10}.  Motivated by the significance of BHNS binaries both as detectable GW sources and SGRB candidates, many simulations of BHNS systems have been performed in past years in a Newtonian or post-Newtonian framework (see, e.g., \\cite{Lee00,RSW,Rosswog,Koba,RKLRasio}) and in conformally-flat relativistic gravitation~\\cite{FBSTR,FBST}. Recently, several groups have performed dynamical simulations of BHNS binary inspirals and mergers in full GR~\\cite{LRA,SU1,SU2,ST07,yst08,skyt09,kst10,eflstb08,elsb09,dfkpst08,dfkot10,fdkt11,cabllmn10,sep11,eps11,kost11,st11,fdksst11}.  Over the past few years, we have studied BHNS mergers beginning with the construction of quasiequilibrium circular orbit initial data~\\cite{BSS,TBFS05,TBFS06,TBFS07a,TBFS07b} and following up with full GR dynamical simulations~\\cite{eflstb08,elsb09}. Our GR simulations, and those by other groups, suggest that for initially nonspinning BHs, the remnant disk mass is substantial for $q \\equiv M_{\\rm BH}/M_{\\rm NS} \\lesssim 3$ and tends to increase with decreasing $q$. For a fixed $q$, the disk mass also increases for smaller NS compaction and for more rapidly spinning BHs aligned with the orbital angular momentum of the binary.  For sufficiently high spins, small mass ratios, and/or lower NS compactions, a substantial disk can form following the merger, favoring BHNS mergers as plausible central engines for SGRBs.  However, these simulations have yet to account for magnetic field effects -- a crucial component in many SGRB models involving a disk around a spinning BH (see, e.g.~\\cite{vk03a,vk03b,SGRB_review}).  For NS surface field strengths $B \\lesssim 10^{16}$G, magnetic fields are unlikely to affect the dynamics of the BHNS inspiral and merger~\\footnote{However, it has been suggested that during the BHNS inspiral the interaction of BH and NS magnetic fields could drive a Poynting flux, which might give rise to observable electromagnetic signals~\\cite{ml11}.}. This was shown to be the case for NSNS binary inspirals in~\\cite{IT00,vkmg03,mvg05,2009MNRAS.399L.164G}. Despite this conclusion, magnetic fields may significantly influence the post-merger dynamics, as the fields are likely to be amplified during and after merger. Magnetic fields could stir turbulence in the remnant disk, resulting in angular momentum transport and accretion onto the BH. They could also lead to matter outflow and jets along the BH remnant spin axis \\cite{2011ApJ...732L...6R}, another ingredient required by most SGRB models~\\cite{SGRB_review,vk03a,vk03b}.  In this paper, we present a new set of fully relativistic BHNS simulations that probes how magnetic fields influence the dynamics and outcome of the merger using our new adaptive mesh refinement (AMR) GRMHD code~\\cite{els10,APPENDIXPAPER}. Fixing the BH:NS mass ratio at $q=3$, we consider the cases where the BH possesses no spin (A cases) and moderate spin $a_{\\rm BH}/M_{\\rm BH}=0.75$ aligned with the orbital angular momentum (B cases). Since the {\\em internal} magnetic field strength and configuration in a NS is not known, we vary the strength and geometry of the internal fields to study their effects. We find that for low and moderate field strengths $\\lesssim 10^{16}\\rm G$, magnetic fields do not significantly alter the inspiral and merger dynamics, which is consistent with the result reported in~\\cite{cabllmn10}. Here, by low and moderate magnetic fields we mean those with field strengths small or moderate when compared to the virial value of about $10^{18}$G; well below this value the field is dynamically unimportant. However, when the central field strength approaches $\\sim 10^{17}$G, corresponding to a magnetic to gas pressure ratio of $\\sim 0.5$\\%, the merger dynamics and remnant disk mass are affected significantly. Yet even with such strong internal magnetic fields, the emitted GWs are not appreciably different from the unmagnetized case, at least for the preliminary set of models considered here. In the late inspiral and merger phases, tidal deformation and disruption of the NS play key roles in distinguishing GWs from BHNSs and BHBHs.  During merger, most of the magnetized NS matter is captured by the BH.  Only when the NS interior is seeded with strong magnetic fields ($B_{\\rm max} \\sim 10^{17}$G, near the center of the NS) is a significant impact on the dynamics observed, resulting in a disk that has up to twice the rest mass as the corresponding unmagnetized case. In all cases the disk accretion rate onto the BH decreases with time immediately after merger, before settling down to a quasistationary state.  Most of the magnetic field lines are tightly wound within the remnant disk, and no evidence of magnetic field collimation around the final spinning BH is observed by the time we terminate our simulations. The remnant disk is hot ($T\\sim 1$MeV) and massive ($M_{\\rm disk} \\sim 0.02M_\\odot$ and $\\sim 0.1M_\\odot$ for cases A and B, respectively).  The magnetic fields threading the remnant disk may be amplified and tangled on a longer timescale than we simulate, stirring up MHD turbulence. Based on extrapolation of the accretion rates near the end of our simulations, the lifetime of the disk is roughly 0.3($M_0/1.4M_\\odot$)s. While there is no evidence of outflows during these preliminary simulations, longer disk evolutions, higher resolution and different B-field geometries may be required to definitively assess the possibility of BHNS binaries as short-hard GRB progenitors.  The following sections are organized as follows. Sections~\\ref{sec:basic_eqns} and~\\ref{sec:numerical} review the basic equations, including our initial data, gauge conditions, matter evolution equations, and diagnostics, as well as their implementation in our GRMHD code. Section~\\ref{sec:results} presents the results of our magnetized BHNS merger simulations. Finally, we summarize our findings and comment on future directions in Sec.~\\ref{sec:summaryandfuturework}. ",
        "conclusions": "\\label{sec:summaryandfuturework}  We present preliminary results from magnetized simulations of BHNS binaries with BH:NS mass ratio 3:1. We treat both initially nonspinning (cases A) and moderately-spinning (cases B) BHs. For those magnetic field configurations we consider, only initial NS magnetic fields with maximum (internal) strength of $\\approx 10^{17}$G -- corresponding to average magnetic to gas pressure ratio of 0.5\\% -- significantly impact the inspiral and merger dynamics. During merger, most of the magnetized NS matter is captured by the BH.  After disruption, the dynamics are followed for about 30--50 $(M_0/1.4M_\\odot)$ms.  Only in the cases in which magnetic fields are strongest are magnetic effects dynamically significant, increasing the final disk mass by up to a factor of two.  The strong magnetic fields help push out the outer layers of the NS during tidal disruption, resulting in a gravitational wave mismatch of 0.004 for the Advanced LIGO broadband configuration. It may be challenging for the upcoming generation of gravitational wave detectors to observe effects from such magnetic fields. Further studies with different field geometries, black hole spins and higher resolution may be required to confirm this finding.  Some SGRB models require a massive, hot, magnetized disk around a BH with collimated magnetic fields to launch jets that generate $\\gamma$-rays (see, e.g.~\\cite{vk03a,vk03b}). In our BHNS simulations, the remnant disk is hot ($T\\sim 1$MeV) and massive ($M_{\\rm disk} \\sim 0.02M_\\odot$ and $\\sim 0.1M_\\odot$ for cases A and B, respectively), and possesses magnetic fields that are tightly wound.  However, evidence for magnetic collimation around the BH or magnetic field-induced turbulence in the disk is not observed. Future analyses will focus on mode growth studies and B-field decomposition in poloidal and toroidal components with respect to the centre of mass of the NS, similar to previous studies of single rotating or magnetized stars \\cite{Zink,Lasky,Ciolfi}. The lack of collimation may be due to the short disk evolution time before our simulation is terminated. The absence of turbulence in the disk may be due to insufficient resolution in the disk, thereby suppressing instabilities like MRI. Therefore, future work will focus on evolving the disk at higher resolution, coupled with longer disk evolutions and different initial magnetic field configurations to more thoroughly assess the possibility of BHNS binaries as short-hard GRB progenitors."
    },
    "1112/1112.0980.txt": {
        "abstract": "Baryon Acoustic Oscillations (BAOs) are a feature imprinted in the galaxy distribution by acoustic waves traveling in the plasma of the early universe. Their detection at the expected scale in large-scale structures strongly supports current cosmological models with a nearly linear evolution from redshift $z \\approx 1000$ and the existence of dark energy. Besides, BAOs provide a standard ruler for studying cosmic expansion. In this paper we focus on methods for BAO detection using the correlation function measurement $\\hat{\\xi}$. For each method, we want to understand the tested hypothesis (the hypothesis $\\mathcal{H}_0$ to be rejected) and the underlying assumptions. We first present wavelet methods which are mildly model-dependent and mostly sensitive to the BAO feature. Then we turn to fully model-dependent methods. We present the most often used method based on the $\\chi^2$ statistic, but we find it has limitations. In general the assumptions of the $\\chi^2$ method are not verified, and it only gives a rough estimate of the significance. The estimate can become very wrong when considering more realistic hypotheses, where the covariance matrix of $\\hat{\\xi}$ depends on cosmological parameters. Instead we propose to use the $\\Delta l$ method based on two modifications: we modify the procedure for computing the significance and make it rigorous, and we modify the statistic to obtain better results in the case of varying covariance matrix. We verify with simulations that correct significances are different from the ones obtained using the classical $\\chi^2$ procedure. We also test a simple example of varying covariance matrix. In this case we find that our modified statistic outperforms the classical $\\chi^2$ statistic when both significances are correctly computed. Finally we find that taking into account variations of the covariance matrix can change both BAO detection levels and cosmological parameter constraints. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.2609_arXiv.txt": {
        "abstract": "{} {We test the robustness of published time delays for 11 lensed quasars by using two techniques to measure time shifts in their light curves.} {We chose to use two fundamentally different techniques to determine time delays in gravitationally lensed quasars: a method based on fitting a numerical model and another one derived from the minimum dispersion method introduced by Pelt and collaborators. To analyse our sample in a homogeneous way and avoid bias caused by the choice of the method used, we apply both methods to 11 different lensed systems for which delays have been published: JVAS B0218+357, SBS 0909+523, RX J0911+0551, FBQS J0951+2635, HE 1104-1805, PG 1115+080, JVAS B1422+231, SBS 1520+530, CLASS B1600+434, CLASS B1608+656, and HE 2149-2745} {Time delays for three double lenses, JVAS B0218+357, HE 1104-1805, and CLASS B1600+434, as well as the quadruply lensed quasar CLASS B1608+656 are confirmed within the error bars. We correct the delay for SBS 1520+530. For PG 1115+080 and RX J0911+0551, the existence of a second solution on top of the published delay is revealed. The time delays in four systems, SBS 0909+523, FBQS J0951+2635, JVAS B1422+231, and HE 2149-2745 prove to be less reliable than previously claimed.} {If we wish to derive an estimate of $H_{0}$ based on time delays in gravitationally lensed quasars, we need to obtain more robust light curves for most of these systems in order to achieve a higher accuracy and robustness on the time delays.} ",
        "introduction": "The determination of the time delay between different images of gravitationally lensed quasars is an important step in different kinds of studies: in deriving $H_{0}$, the expansion rate of the Universe \\citep[e.g.][]{2008AA...488..481V}, for microlensing studies \\citep[e.g.][]{2006AA...455L...1P}, and for detailed studies of the structure of a lensed quasar \\citep[e.g.][]{2002MNRAS.334..905G,2008ApJ...676...80M}.  However, previous time delay determinations have been far from homogeneous, not only because they are based on different methods, but also because of their varying levels of reliability.  Their accuracy depends, among other factors, on the amplitude and shape of the quasar's intrinsic variations, the perturbations of the light curves by microlensing effects, the photometric error bars, the typical time sampling of the monitoring, the total time span of the observing campaign, and the intervals between observing seasons.  Moreover, the published error bars are generally internal errors only and the way in which they are determined varies from study to study. Unfortunately, once a time delay has been published, the value may be used for years without verification, % even when the authors of the original article caution the reader that the result is not very well-constrained \\citep[e.g.][]{2005AA...431..103J}.  Hence, we are of the opinion that it would be useful to re-evaluate published time delays in a number of systems using the same methods as for all of the delays, as well as for the estimate of the error bars.  An idea of the robustness of the time delay values can thus be obtained, not only internally with our methods, but also by comparing our results with published ones. We do not claim that our methods are superior to other ones. However, a critical reanalysis of published results using two fundamentally different approaches allows us to sort the results in terms of the reliability and independence of the method and to determine which lensed systems may be useful for determining $H_{0}$. Our main purpose is thus to examine whether the published light curves allow the determination of reliable time delays. If the answer is positive, we attempt to estimate realistic error bars, and correct some of the published values for small systematic errors.  For the determination of $H_{0}$ by means of gravitational lensing to be competitive with more classical methods \\citep[e.g.][]{2009ApJ...699..539R}, we need to reach at least a comparable accuracy of $\\sim5\\%$ in $H_{0}$. As time delays from different lensed systems should be combined to obtain $H_{0}$, we can assume that the error in the individual time delays contributes to the statistical error in $H_{0}$. Since the time delay uncertainties are only one of several sources of error in $H_{0}$ determinations (to be added e.g. to uncertainties in the dark matter distribution in the lens), they should in any case not exceed $5\\%$.   Section 2 and 3 present the two methods used for time delay determination.  These methods are applied to the 11 lensed systems, for which the main results are described in Section 4. The lenses of our sample are those for which accurate astrometry has been determined by means of the deconvolution of near-infrared Hubble Space Telescope images (Sluse et al, submitted to A\\&A). A summary of these results is presented in Section 5, together with our conclusions. ",
        "conclusions": "We have presented an improved numerical method, the numerical model fit method, to calculate time delays in doubly or quadruply lensed quasar systems, and applied it to 11 systems for which time delays had been published previously. This allowed the validity of these time delay values to be evaluated in a coherent way. The use of a minimum dispersion method allowed us to check the independence of the results from the method. The results are summarized in Table \\ref{tab:Conclusions}.  We caution that some published time delay values should be interpreted with care: even if we have been able to confirm some values (time delays for JVAS B0218+357, HE 1104-1805, CLASS B1600+434, and for the three delays in the quadruply lensed quasar CLASS B1608+656) and give an improved value for one system, SBS 1520+530, many of the published time delays considered in our analysis have proven to be be unreliable for various reasons: the analysis is either too dependent on some data points, leads to multiple solutions, is sensitive to the addition of random errors, or is incoherent between the two different methods used.  Given the accuracy that is needed for time delays to be useful to further studies, we note that it will be necessary to perform long-term monitoring programs on dedicated telescopes to obtain high-quality light curves of lensed quasars, not only for new systems but also for the majority of the lenses in this sample, for which the time delay has been considered to be known. The COSMOGRAIL collaboration, which has been observing over 20 lensed systems for several years now, will soon be improving the time delay values for some of these systems for which the accuracy is unsatisfactory.     \\begin{table*}[htbp] \\begin{center} \\begin{tabular}{lcp{3mm}ll} System & Our Result / Comments & & Published Time Delay (days) & Reference \\\\ \\hline JVAS B0218+357 & $\\Delta t_{AB} = 9.9_{-0.9}^{+4.0}$ & & $\\Delta t_{AB} = 10.1_{-1.6}^{+1.5}$ & \\citet{2000ApJ...545..578C} \\\\ & or & & $\\Delta t_{AB} = 12\\pm3$ & \\citet{1996IAUS..173...37C} \\\\ & $\\Delta t_{AB} = 11.8\\pm2.3$ & & $\\Delta t_{AB} = 10.5\\pm0.4$ & \\citet{1999MNRAS.304..349B} \\\\ \\hline SBS 0909+523 & unreliable & & $\\Delta t_{BA} = 49\\pm6$ & \\citet{2008NewA...13..182G} \\\\ & & & $\\Delta t_{BA} = 45_{-1}^{+11}$ & {\\citet{2006AA...452...25U}} \\\\ \\hline RX J0911+0551 & 2 solutions: & & $\\Delta t_{BA} = 150\\pm6$ & \\citet{2001PhDT.........4B}\\\\ & $\\Delta t_{BA} \\sim 146$ or $\\sim 157$ & & $\\Delta t_{BA} = 146\\pm4$ & \\citet{2002ApJ...572L..11H}\\\\ \\hline FBQS J0951+2635 & unreliable & & $\\Delta t_{AB} = 16\\pm2$ & \\citet{2005AA...431..103J} \\\\ \\hline HE 1104-1805 & & & $\\Delta t_{BA} = 152_{-3.0}^{+2.8}$ & \\citet{2007ApJ...660..146P} \\\\ & & \\multirow{3}{1cm}{$\\left\\{\\rule{0mm}{6mm}\\right.$} & $\\Delta t_{BA} = 161\\pm7$ & \\citet{2003ApJ...594..101O} \\\\ & impossible to distinguish & & $\\Delta t_{BA} = 157\\pm10$ & \\citet{2003AcA....53..229W} \\\\ & but identical within error bars & &  $\\Delta t_{BA} = 162.2_{-5.9}^{+6.3}$ & \\citet{2008ApJ...676...80M} \\\\ \\hline PG 1115+080 & dependent on method & & $\\Delta t_{CA} \\sim 9.4$ & \\citet{1997ApJ...475L..85S} \\\\ &  & & $\\Delta t_{CB} = 23.7\\pm3.4$ & \\citet{1997ApJ...475L..85S} \\\\ & & & $\\Delta t_{CB} = 25.0_{-3.8}^{+3.3}$ & \\citet{1997ApJ...489...21B}\\\\ \\hline JVAS B1422+231 & contradictory results & & $\\Delta t_{BA} = 1.5\\pm1.4$ & \\citet{2001MNRAS.326.1403P} \\\\ & between methods: & & $\\Delta t_{AC} = 7.6\\pm2.5$ & \\\\ & BAC or CAB? & & $\\Delta t_{BC} = 8.2\\pm2.0$ & \\\\ \\hline SBS 1520+530 & $\\Delta t_{AB} = 125.8\\pm2.1$ & & $\\Delta t_{AB} = 130\\pm3$ & \\citet{2002AA...391..481B} \\\\ & & & $\\Delta t_{AB} = 130.5\\pm2.9$ & \\citet{2005AA...440...53G}\\\\ \\hline CLASS B1600+434 & $\\Delta t_{AB} = 47.8\\pm1.2$ & & $\\Delta t_{AB} = 51\\pm4$ & \\citet{2000ApJ...544..117B} \\\\ \\hline CLASS B1608+656 & $\\Delta t_{BA} = 31.6\\pm1.5$ & & $\\Delta t_{BA} = 31.5_{-1}^{+2}$ & \\citet{2002ApJ...581..823F} \\\\ & $\\Delta t_{BC} = 35.7\\pm1.4$ & & $\\Delta t_{BC} = 36.0\\pm1.5$ & \\\\ & $\\Delta t_{BD} = 77.5\\pm2.2$ & & $\\Delta t_{BD} = 77.0_{-1}^{+2}$ & \\\\ \\hline HE 2149-2745 & unreliable & & $\\Delta t_{AB} = 103\\pm12$ & \\citet{2002AA...383...71B} \\\\ \\hline \\end{tabular} \\end{center} \\caption{Summary of Time Delays for 11 Lensed Systems.} \\label{tab:Conclusions} \\end{table*}"
    },
    "1112/1112.2579_arXiv.txt": {
        "abstract": "We are at the dawn of a new golden age for radio astronomy, with a new generation of facilities under construction and the global community focussed on the Square Kilometre Array as its goal for the next decade. These new facilities offer orders of magnitude improvements in survey speed compared to existing radio telescopes and arrays. Futhermore, the study of transient and variable radio sources, and what they can tell us about the extremes of astrophysics as well as the state of the diffuse intervening media, have been embraced as key science projects for these new facilities. In this paper we review the studies of the populations of radio transients made to date, largely based upon archival surveys. Many of these radio transients and variables have been found in the image plane, and their astrophysical origin remains unclear.  We take this population and combine it with sensitivity estimates for the next generation arrays to demonstrate that in the coming decade we may find ourselves detecting $10^5$ image plane radio transients per year, providing a vast and rich field of research and an almost limitless set of targets for multiwavelength follow up. ",
        "introduction": "The Universe is a violent and dynamic environment, in which the explosions of massive stars can outshine an entire galaxy, supermassive black holes swallow stars whole, merging neutron stars cause ripples in the fabric of spacetime and bursts of ultra-high energy radiation which can be detected at vast distances, and particles are accelerated to energies far surpassing anything possible in laboratories on the Earth. The extremes of physics -- density, temperature, pressure, velocity, gravitational and magnetic fields -- experienced in these environments provide a unique glimpse at the laws of physics operating in extraordinary regimes.  Such `extreme astrophysics', is a high priority for global research in the 21st century\\footnote{See e.g. the ASTRONET report {\\em A Science Vision for European Astronomy} available at {\\bf www.astronet-eu.org}}. In the second half of the 20th century, astronomy moved beyonds its optical origins into the radio, infrared, ultraviolet, X-ray and gamma-ray bands. The most extreme environments in the Universe betray themselves by their copious high-energy emission, in the X-ray and gamma-ray bands, and also by their radio emission. Strong magnetic fields and shock waves as matter collides in the ambient medium, result in radio bursts from electrons (and positrons) spiralling around, and sometimes being channelled along, these magnetic field lines. Wide-field orbiting X- and gamma-ray observatories, with all-sky monitors, have revealed our Galaxy to be a hotbed of explosive and relativistic effects, which go largely unnoticed by traditional optical telescopes both because they are missed in the narrow fields of view, and are obscured by interstellar dust.  By far the best way to observe and understand the violent Universe from the ground is via the radio emission which is ubiquitously associated with these violent events. The emission arises as the shocked particles glow with synchrotron radiation, streaming along and around magnetic field lines, or are induced to move and emit coherently, producing extremely bright bursts. However, in the past radio astronomy, like optical astronomy, has suffered from the narrow fields of view which are a consequence of large dishes, and cannot survey the sky sufficiently rapidly to detect rare and rapid events, which may be the most significant astrophysically.  All of this is about to change, with a new generation of radio telescopes, taking advantage of new technology and the vast advances made in data transport and computer processing, connecting large numbers of small detectors, able to deliver enormous fields of view. Finally the age of the radio all-sky monitors is here. What is more, these radio all-sky monitors will have several advantages over the traditional X-ray and gamma-ray monitors: they will be able to localise events immediately with arcsec precision, and will be able to go significantly deeper / further.  As we shall demonstrate in this review, our current understanding of the population of radio transients implies that we should detect thousands, or even millions, of such events with `next generation' radio facilities such as The Low-Frequency Array (LOFAR), The South African Karoo Array Telescope (MeerKAT) and The Australian SKA Pathfinder (ASKAP), and ultimately the Square Kilometre Array (SKA). As well as drawing our attention to the extremes of astrophysics, these sources are potentially observable to vast distances (possibly as far as the Epoch of Reionisation at redshift $z > 6$) and will turn out to be amongst our most valuable probes of the intergalactic and intercluster medium on large scales. They may furthermore, in some cases, turn out to be the electromagnetic signatures of events which produce detectable gravitational waves, which would be a breakthrough connection. ",
        "conclusions": "The benchmark study of Bower et al. (2007) is five years old and remains the standard for our understanding of the transient sky at GHz facilities. Several subsequent radio transient searches have delivered upper limits or detections which are consistent with this, but neither significantly refine the statistics nor detect sources in anything close to real time, crucial for multiwavelength follow up and our understanding of the astrophysics of these sources. The new generation of radio facilities and approved key science projects should change all of this, with near real time reporting of thousands of transient events per year. This in turn will provide an extremely rich resource for multiwavelength follow up, and the possible opening of a new window on the high-energy Universe. Prepare for the flood.   \\bgroup\\small \\noindent{\\em Note in manuscript:} Since this manuscript was written, Frail et al. (2011) have reported that the Bower transient rates, on which the estimates in this paper were largely based, may have been overestimated by up to an order of magnitude. In this case the rates predicted here may be similarly overestimated, but still correspond to between thousands and millions per year, depending on the facility. \\egroup"
    },
    "1112/1112.2786_arXiv.txt": {
        "abstract": "We present results of a search for bright Lyman break galaxies at $1.5\\le z\\le 2.5$ in the GOODS-S field using a NUV-dropout technique in combination with color-selection. We derived a sample of 73 LBG candidates. We compare our selection efficiencies to BM/BX- and BzK methods (techniques solely based on ground-based data sets), and find the NUV data to provide greater efficiency for selecting star-forming galaxies. We estimate LBG candidate ages, masses, star formation rates, and extinction from fitting PEGASE synthesis evolution models. We find about 20\\% of our LBG candidates are comparable to infrared luminous LBGs or sub-millimeter galaxies which are thought to be precursors of massive elliptical galaxies today.  Overall, we can show that although BM/BX and BzK methods do identify star-forming galaxies at $z\\sim 2$, the sample they provide biases against those star-forming galaxies which are more massive and contain sizeable red stellar populations.  A true Lyman break criterion at $z\\sim 2$ is therefore more directly comparable to the populations found at $z\\sim 3$, which does contain a red fraction. ",
        "introduction": "Since their discovery in the early 1990s \\citep[][]{1993AJ....105.2017S}, Lyman break galaxies (LBGs) have been used very successfully to study galaxy evolution processes in the high redshift Universe. Using the dropout technique, multi-wavelength surveys have built up large samples of LBGs at redshifts between $z\\sim 3 - 8$, where the Lyman break is accessible with solely ground-based data sets \\citep[e.g.][]{1993AJ....105.2017S,1995AJ....110.2519S,1999ApJ...519....1S,2003ApJ...592..728S,2003ApJ...593..630L,2004ApJ...616L..79B}. The studies have shown that at $z\\gsim 4$, a significant fraction of the ionizing UV flux is emitted by low luminosity LBGs \\citep[e.g.][]{2003ApJ...593..630L, 2006ApJ...648..299S, 2007ApJ...670..928B,2010Natur.467..940L}.  Analysis of the stellar populations showed that the population of LBGs at $z\\sim5$ consists of significantly less massive and younger galaxies than comparable samples at $z\\sim3$ \\citep[e.g.][]{2007MNRAS.377.1024V,2009ApJ...693..507Y}. \\citeauthor[][]{2007MNRAS.377.1024V} suggested that these LBGs are experiencing their first significant buildup of stellar mass \\citep[see also][]{2009ApJ...697.1493S}. At $z\\sim3$, the LBG population consists of more massive galaxies that have formed stars over longer periods \\citep[e.g.][]{2001ApJ...562...95S,2008ApJS..175...48R}. A considerable fraction of the $z\\sim 3$ LBGs ($\\sim 15\\%$) consists of massive $M_* > 10^{11} M_\\odot$ infrared luminous galaxies \\citep[][]{2006ApJ...648...81R,2010MNRAS.401.1521M}; these ILLBGs are considered one possible class of progenitors of today's massive elliptical galaxies.  Less is known about the LBGs at $z\\lsim 2.5$. At this redshift, the Lyman break occurs in the UV, which can only be observed from space. First studies using deep images observed by the GALEX satellite obtained samples of LBGs at $z\\sim 1$ using FUV dropouts \\citep{2006A&A...450...69B,2007MNRAS.380..986B,2009ApJ...702..506H}. The majority of LBGs ($\\sim$ 75\\%) at this redshift are disk-dominated, while the rest consist of interacting/merging galaxies \\citep[][]{2006A&A...450...69B}.  At $z\\sim2$, a number of groups have worked around UV restrictions by using selection techniques which solely rely on ground based data sets as a proxy method for identifying star-forming galaxies. The BM/BX-method \\citep[][]{2004ApJ...607..226A,2004ApJ...604..534S} is based on {\\it U$_n$GR} color selection.  The color cuts were determined from a sample of spectroscopically-verified $1.9<z<2.7$ star-forming galaxies, which were fitted with dust-reddened SEDs. The resulting best-fits were redshifted to $z=1.5, 2.0, 2.5$ and used to calculate characteristic colors; for more details see \\cite[][]{2004ApJ...607..226A}.  A second ground-based color method is the BzK-selection \\citep[][]{2004ApJ...617..746D}. The choice of filters covers the flat UV part of the spectra and flanks the Balmer break at rest-frame 4000~\\AA. This method is sensitive to selecting and distinguishing between star-forming and passively evolving galaxies at $z\\gsim 1.4$. Using the K20 data from the GOODS fields, \\citeauthor[][]{2004ApJ...617..746D} identified 25 star-forming and 7 old galaxies and reported a high selection efficiency with only 13\\% contamination by interlopers.  Follow-up studies showed that the optically (BX) selected star-forming galaxies at $z\\sim 2$ have typical stellar masses on the order of $3\\times10^{10}M_{\\odot}$ and  typical star formation rates of $\\sim50\\,M_{\\odot}$~yr$^{-1}$ \\citep[e.g.][]{2004ApJ...603L..13R,2005ApJ...626..698S}. Space densities are  $\\sim 3\\times10^{-3}$~Mpc$^{-3}$ with a correlation length of $r_0 = 4.2\\pm0.4 h^{-1}$~Mpc \\citep[][]{2005ApJ...633..748R,2005ApJ...619..697A}, indicating that typical star-forming galaxies at $z\\sim2$ are hosted by $\\sim 2\\times10^{12}M_{\\odot}$ dark matter halos.  Studies of the optical-NIR selected star-forming galaxies by \\citet[][]{2005ApJ...631L..13D} show that a typical $BzK$ galaxy at $z\\sim 2$ is likely to be a ULIRG with $L_{IR} \\sim (1-2)\\times 10^{12} L_{\\odot}$ and star formation rates SFR $\\sim 200 - 300 M_{\\odot} $~yr$^{-1}$. The space density of these galaxies is $\\sim(1-2)\\times 10^{-4}$~Mpc$^{-3}$, which is 2-3 times higher than at $z\\sim0.1$ and 1 \\citep[e.g.][]{2003AJ....126.1607S,2005ApJ...632..169L} and makes them relatively common objects at $z\\sim 2$. Their ULIRG nature is also supported by morphological studies using HST imaging \\citep[][]{2004ApJ...600L.127D}, indicating that many of the $BzK$-selected galaxies show hints of merger events driving strong star formation events.  Dynamical studies of star-forming galaxies at $z\\sim 2$ \\citep[][]{2009ApJ...706.1364F}, together with the results from \\citet[][]{2006A&A...450...69B}, indicate that between redshifts of $z\\sim 2$ and $z\\sim 1$ is when disk galaxies with significant masses start to contribute to the co-moving density of galaxies. It is also the phase when the star formation intensity reaches its peak \\citep[][and references therein]{2004ApJ...615..209H}.  From these solely ground-based selection techniques and the development of new instruments like the IFU SINFONI at the VLT and OSIRIS at Keck, we have learned much about star formation at the critical epoch $z\\sim 2$. However, since the selection processes are different to the techniques used for true LBGs at $z \\gsim 3$, comparisons between the samples are not straightforward. A consistent selection method using the Lyman break, independent of intrinsic galaxy properties, should be used to select consistent samples of star-forming galaxies at $z\\sim 2$. Observations require the space-based facilities like GALEX or the UV-channels of the newly installed WFC3 on board HST. First studies include \\citet[][]{2010ApJ...720.1708H} and \\citeauthor{2009ApJ...697.1410L}\\,  The latter analyzed a very deep ($\\sim$140ksec) GALEX NUV image of the Subaru Deep Field \\citep[SDF][]{2004PASJ...56.1011K} and other filters to select true LBGs at redshifts $1.6<z<2.7$, and compiled a sample of $\\sim$8000 LBG candidates down to $V = 25.3$ with $NUV \\approx 27$, including 24 spectroscopically confirmed at $1.5<z<2.7$. For $z\\sim2$ LBGs, they reported a factor $\\sim 1.8$ higher summed luminosity density compared to $z\\sim 2$ BX and $z\\sim 3$ LBGs, and an increase in the luminosity density by a factor of 3-6 between $z\\sim 5$ and $z\\sim 2$.  \\citeauthor[][]{2010ApJ...720.1708H} examined HST+ WFC3 images (F225W, F275W, and F336W) of a 50 arcmin$^2$ field of the GOODS-S field, as part of an extended multi-color survey. The group used $U_{225}-, U_{275}-, {\\rm and}\\, U_{336}-$dropouts down to $AB=26.5$ combined with color criteria to select 473 LBG candidates at $z\\sim 1.7$, 2.1 and 2.7 respectively. The fitted Schechter function parameters agree well with predictions from other observations with respect to redshift.  These two studies demonstrate the feasibility of selecting samples of true LBGs at $z\\sim 2$ consistent with higher redshift $z\\gsim 3$ LBG samples. Although they are a first step in better understanding this crucial population of star forming galaxies at an epoch formally known as the ``redshift desert'', little currently is known about their properties. The studies either lack data coverage especially in the NIR and mid-IR \\citep[][]{2009ApJ...697.1410L} or field of view available \\citep[][]{2010ApJ...720.1708H}.  In this paper, we present first results from a dropout+color selection of a sample of true $z\\sim 2$ LBGs based on deep GALEX observations of the GOODS-S. Our specific purpose is to target a rarer, brighter sample of LBGs than found in the small, deep survey of \\citet[][]{2010ApJ...720.1708H}. In addition, it complements the wider, shallower study of \\citet[][]{2009ApJ...697.1410L} because we use a wider range of filters and target a different field, to minimize the effects of cosmic variance. We will show that we can obtain a sample of bright $z\\sim 2$ LBGs similar to the higher redshift population observed at $z \\gsim 3$. Here we report first results of estimates of ages, stellar masses, SFRs, and extinction for our purely UV-selected $z\\sim 2$ LBG sample. Using the deep GALEX UV data available for the CDF-S, we were able to utilize the complete size of the GOODS-S field which offers a wide range of data from the UV to mid-IR and also tripled the  \\citeauthor[][]{2010ApJ...720.1708H} sample size.  The paper is organized as follows: \\S2 summarizes the data used, \\S3 describes our search and defines the selection criteria, and \\S4, \\S5, and \\S6 present photometric redshifts, SED modeling and the results. Our summary and conclusions are in \\S7. Throughout the paper we use a cosmology with $H_0 = 71$ km s$^{-1}$~Mpc$^{-3}$, $\\Omega_M = 0.3, {\\rm and}\\, \\Omega_{\\lambda} = 0.7$. All distances quoted are comoving. ",
        "conclusions": "We present results of a search for Lyman break galaxies at $1.5\\leq z \\leq 2.5$, an epoch once called the redshift desert. Over the last few years, studies employing the the BM/BX- \\citep[][]{2004ApJ...604..534S,2004ApJ...607..226A} and BzK-methods \\citep[][]{2004ApJ...617..746D,2004ApJ...600L.127D} used observed-frame optical color-selection criteria to find and study star-forming galaxies in this redshift range -- which excluded the Lyman break. Such methods identify galaxies which, though star-forming, are not completely consistent with LBG samples at higher redshift ($z\\gsim 3$). The difference is that the dropout method allows for a wider range in color redward of the Lyman limit, which in practice means including redder, more evolved and potentially more massive galaxies.  Our NUV-dropout plus optical color selection technique targeted star-forming galaxies at $1.5\\leq z \\leq 2.5$, and is similar to the one used by \\citet[][]{2009ApJ...697.1410L}. Using the deep publicly available data set of the GOODS-S (CDF-S) field covering the UV to mid-IR, we identified a sample of 73 LBG candidates. Our selection method resulted in a success rate of $\\sim 80$\\% ($\\sim$95\\% for $1.0\\leq z \\leq 2.5$).  The NUV dropout technique allows for the detection of more evolved and/or dust-reddened galaxies not detected with the BM/BX-method, while also decreasing the number of interlopers compared to the BzK-method. Although the BM/BX and BzK methods can select much fainter objects than we can here, due to the limitation of GALEX in sensitivity, they exclude the more evolved and/or dust-reddened galaxies completely. The fact that 40\\% of the dropout-selected LBGs are not detected by the BM/BX and/or BzK methods provides for additional UV/optical-selected star formation at $z\\sim 2$, especially if hidden by large amounts of dust.  We find that our $z\\sim 2$ galaxy sample is consistent with LBGs at $z\\gsim 3$, and therefore allows for consistent follow-up studies at redshifts spanning $z\\sim 1.0 - 7$.  We calculated useful photometric redshifts using \\HyperZ\\ and included a small set of spectroscopic redshifts for our LBG candidate sample (Figure~\\ref{redshift}, red filled circles). The normalized rms scatter for our LBG candidate sample is $\\delta z_{rms} = 0.198$, with a smaller mean offset of $\\langle dz_{offset} \\rangle = 0.095$.  Using the combination of spectroscopic (where available) and photometric redshifts, we fitted synthetic models from a library of PEGASE spectra to our LBG sample, and determined ages ranging from 4~Myr to 3~Gyr, masses between $9\\times10^8M_\\odot$ to about $2\\times 10^{10}M_{\\odot}$, and SFRs up to $\\approx 270~M_\\odot $~yr$^{-1}$.  Our LBG candidates have moderate stellar masses and SFRs consistent with $z\\sim 3$ LBGs \\citep[e.g.][]{2001ApJ...562...95S}. Ages and dust extinction values of our sample are also more consistent with LBGs at $z\\sim 3$, in comparison with the BX-selected galaxies at $z\\sim 2$ of \\citep[e.g.][]{2005ApJ...626..698S}, although our sample shows a lack of the very young star-forming galaxies found in the BM/BX and $z\\sim 3$ samples (Figure~\\ref{age}). Our sample represents the moderate mass range of the UV luminosity distribution of LBGs at this redshift. This can  lead to selection biases against objects with much higher masses. We therefore should exercise care when comparing our results to other surveys.  We also identified a small fraction of our LBG sample ($\\sim$20\\%) as potential ILLBGs and/or sub-millimeter galaxies using Spitzer 8.0 and 24~$\\mu$m data. Although there is some overlap with the other $z\\sim 2$ LBGs, the $R-[3.6]$ vs. $[3.6]$ plot in Figure~\\ref{spitzer} shows that the ILLBGs (blue and orange diamonds) represent the UV-bright end of the LBG distribution for a given value of $R-[3.6]$.  Due to their red colors at optical wavelengths (compared to the BM/BX selected), most of these galaxies are missed by the BM/BX and BzK methods ($\\sim10$\\% detected).  The ILLBG SED fits result, on average, in higher total stellar masses, SFRs, and extinctions than our overall sample.   Compared with higher and lower redshift LBGs from the literature, we find that our $z\\sim 2$ sample has an average extinction which is higher than LBGs at $\\gsim 3$ \\citep[e.g][]{2001ApJ...562...95S,2007MNRAS.377.1024V,2009ApJ...693..507Y}. A similar result holds in comparison with ILLBGs at $z\\gsim 3$ \\citep[e.g.][]{2006ApJ...648...81R}, with the exception of massive, IR luminous, ULIRG-like LBGs which are selected in the infrared \\citep[][]{2007ApJ...669..749S}. We find that the masses and extinction values are in better agreement with results for $z\\sim 1$ LBGs \\citep[][]{2011A&A...526A..10N}.  The median total IR luminosity of our subsample of $L_{TIR} = 4\\times 10^{11} L_{\\odot}$ indicates that these ILLBGs are comprised mainly of LIRGs. Only $\\sim15$\\% (2 out of 15 ILLBGs) are comparable to ULIRGs with $L_{TIR} \\geq 10^{12} L_{\\odot}$. From the total IR luminosities, we estimate a median SFR half that compared to the one derived from the SED model fits.  The role of ILLBGs in the overall scheme of star formation at $z>1.2$ and their contribution to the cosmic IR background merits further investigation, as ILLBGs could contribute significantly.  A simple morphological classification shows that the majority of our LBG candidates appear disk-like, and many show indications for  mergers and interactions. Similar results were found for the ILLBGs, with 80\\% having disk-like morphologies. This is consistent with results published for $z\\sim 1$ LBGs \\citep[][]{2006A&A...450...69B} and $z\\sim 1$ LIRGs \\citep[e.g.][]{2005ApJ...632L..65M,2005ApJ...625...23B,2007ApJ...662..322H}.  Overall, our study shows that the selection of star-forming galaxies at $1.5\\le z\\le 2.5$ using a true Lyman break dropout technique allows for the selection of a robust sample of LBGs, which can used for comparable follow-up studies with LBGs at higher redshifts.  A Lyman break selection like we used here allows redder, more massive LBGs to be included in comparison to other selection techniques such as the BM/BX and BzK methods."
    },
    "1112/1112.4476.txt": {
        "abstract": "New observational facilities are becoming increasingly capable of observing reflected light from transiting and directly imaged extrasolar planets. In this study, we provide an analytic framework to interpret observed phase curves, geometric albedos, and polarization of giant planet atmospheres. We compute the observables for non-conservative Rayleigh scattering in homogeneous semi-infinite atmospheres using both scalar and vector formalisms. In addition, we compute phase curves and albedos for Lambertian, isotropic, and anisotropic scattering phase functions. We provide analytic expressions for geometric albedos and spherical albedos as a function of the scattering albedo for Rayleigh scattering in semi-infinite atmospheres. Given an observed geometric albedo our prescriptions can be used to estimate the underlying scattering albedo of the atmosphere, which in turn is indicative of the scattering and absorptive properties of the atmosphere. We also study the dependence of polarization in Rayleigh scattering atmospheres on the orbital parameters of the planet-star system, particularly on the orbital inclination. We show how the orbital inclination of non-transiting exoplanets can be constrained from their observed polarization parameters. We consolidate the formalism, solution techniques, and results from analytic models available in the literature, often scattered in various sources, and present a systematic procedure to compute albedos, phase curves, and polarization of reflected light. ",
        "introduction": "Observations of reflected light provide constraints on the scattering and absorption processes in planetary atmospheres. In the solar system, geometric albedo spectra and phase curves have been used to infer chemical compositions and cloud properties in several planetary and satellite atmospheres (Hansen \\& Hovenier 1974; Karkoschka 1994; Satoh et al. 2000; Sromovsky et al. 2001; Irwin et al. 2002; Karkoschka \\& Tomasko 2011). Observational techniques are becoming increasingly capable of detecting reflected light from exoplanetary atmospheres, though currently available data are often limited to broadband visible photometry (Snellen et al. 2009, Borucki et al. 2010, Demory et al. 2011). Our goal in the present work is to provide an analytic framework to aid in the interpretation of measurements of geometric albedos and phase curves of irradiated exoplanets.  Several theoretical studies in the past decade have advocated the use of reflected light observations for atmospheric characterization of extrasolar giant planets. Numerical models have been reported for directly imaged planets on wide orbits, as well as for closer-in giant planets with very high incident stellar irradiation (Seager \\& Sasselov 1998; Marley et al. 1999; Seager et al. 2000;  Goukenleuque et al. 1999; Sudarsky et al. 2000,2005; Stam et al. 2004; Sengupta \\& Maiti 2006; Burrows et al. 2008; Buenzli \\& Schmid 2009; Cahoy et al. 2011; de Kok et al. 2011; Kostogryz et al. 2011; Marley \\& Sengupta 2011). Atmospheres of extrasolar planets known to date, predominantly of extrasolar giant planets, experience a much wider range of equilibrium temperatures ($T_{\\rm eq}  \\sim$ 500 - 3000 K) than that of solar-system giant planets ($T_{\\rm eq} \\lesssim$ 125 K). Sudarsky et al. (2000) suggested a classification scheme of albedo spectra for giant planets, based on which highly irradiated planets ($T_{\\rm eq} \\sim$ 900 - 1500 K) are expected to have lower optical albedos due to strong Na and K absorption. At even higher temperatures, the presence of silicate condensates might increase the albedos. Albedo-based classification schemes for exoplanetary atmospheres were also discussed by  Burrows et al. (2008), Cahoy et al. (2011), and Cowan \\& Agol (2010) as a function of different system parameters.  Recently, observational constraints on the geometric albedos, visible phase curves, and potential polarization parameters, for several exoplanetary atmospheres have been reported. Some of the earliest constraints on geometric albedos ($A_g$) of extrasolar giant planets were reported for the non-transiting planets $\\tau$ Boo, with upper limits of 0.3 (Charbonneau et al. 1999), 0.22 (Collier-Cameron et al. 2000) and 0.39 (Leigh et al. 2003a), $\\upsilon$ And with an upper limit of 0.42 (Collier-Cameron et al. 2002), and HD75289 with an upper limit of 0.12 (Leigh et al. 2003b). However, in recent years, with the discoveries of transiting exoplanets and with the advent of high-precision transit photometry, geometric albedo constraints have been obtained for a substantial number of exoplanets. Rowe et al. (2008) observed the transiting hot Jupiter HD 209458b with the {\\it MOST} satellite and reported a 3-sigma upper-limit of $A_g < 0.17$. Other hot Jupiters with $A_g$ estimates include TrES-3 (Winn et al. 2008), CoRoT-1b (Snellen et al. 2009), CoRoT-2b (Snellen et al. 2010), HAT-P-7b (Christiansen et al. 2010), Kepler-5b and 6b (Kipping \\& Bakos 2010; Desert et al. 2011), HD~189733b (Berdyugina et al. 2011), and Kipping \\& Spiegel (2011), all of which have $A_g$ upper-limits less than 0.3. On the other hand, Kipping \\& Bakos (2010) and Demory et al. (2011) reported independent detections of a high geometric albedo for Kepler-7b of $A_g = 0.38 \\pm 0.12$ and  $A_g = 0.32 \\pm 0.03$. Recently, several groups have also reported attempts to measure polarization parameters of exoplanetary atmospheres in the visible (Berdyugina et al. 2008, 2011a,b; Wiktorowicz 2009).  In this work, we derive an analytic framework with which to interpret observations of albedos, phase curves, and polarization of exoplanetary atmospheres. We generally assume cloud-free\\footnote{We do consider the asymmetric scattering phase function (see Section~\\ref{sec:asymm}) which may be used as an approximation for particulate scattering if the scattering asymmetry factor is known. However, such a treatment still does not address the full problem of Mie scattering by particulates.} homogeneous and semi-infinite atmospheres, and analytically solve the multiple scattering problem for several different scattering phase functions. Detailed numerical models solve the general radiative transfer problem in a plane-parallel inhomogeneous atmosphere (Marley et al. 1999; Seager et al. 2000; Sudarsky et al. 2000; Burrows et al. 2008), with the assumption of radiative equilibrium for given chemistry and sources of absorption and scattering. Such models, however, involve significant computation time and convergence monitoring, and may be computationally prohibitive for formal fits to data, where several consecutive model evaluations are typically required. On the other hand, computationally efficient analytic models exist for several forms of scattering, if the atmosphere can be assumed to be homogeneous in the scattering albedo\\footnote{The scattering albedo is given by $\\omega = \\sigma_{scat}/(\\sigma_{scat} + \\sigma_{abs})$, where $\\sigma_{scat}$ is the single-scattering cross section and $\\sigma_{abs}$ is the absorption cross section.} ($\\omega$). While the latter condition may not be exactly satisfied in giant planetary atmospheres across all temperatures, the analytic approach allows one to derive a representative atmosphere-averaged scattering albedo from the observables, and has been successfully used in fitting observations of solar-system planets (Kattawar \\& Adams 1971; Sromovsky 2005; Schmid et al. 2006).  Many researchers have derived analytic\\footnote{By `analytic' we mean solutions that can be expressed in closed form or  which can be evaluated to arbitrary precision. This includes results derived using ordinary (as opposed to partial) differential equations.} expressions for reflection coefficients, albedos, phase curves, and polarization parameters for different forms of scattering in planetary atmospheres (Chandrasekhar 1950; Horak 1950; Abhayankar \\& Fymat 1970; Kattawar \\& Adams 1971;  Van de Hulst 1981; Bhatia \\& Abhayankar 1982). Chandrasekhar (1950) and Horak (1950) reported analytic solutions for reflected specific intensities from homogeneous atmospheres for the cases of isotropic scattering, asymmetric scattering, and conservative Rayleigh scattering ($\\omega = 1.0$), all for finite and semi-infinite atmospheres. Chandrasekhar (1960) and Horak \\& Chandrasekhar (1960) further generalized their solutions to address non-conservative Rayleigh scattering ($\\omega < 1$) in semi-infinite atmospheres. Formulations using the scalar phase function for single-scattering do not take into account the effect of polarization. Chandrasekhar (1960) and Abhayankar \\& Fymat (1970,1971) reported a general solution including conservative and non-conservative Rayleigh scattering in homogeneous semi-infinite atmospheres using the full Rayleigh phase matrix, so that the Stokes parameters (Stokes 1852; Chandrasekhar 1950,1960), and hence the polarization, could also be computed analytically. Subsequent studies have applied analytic approaches to a wide range of problems in planetary and exoplanetary atmospheres (Kattawar \\& Adams 1974;  Van de Hulst 1980; Bhatia \\& Abhayankar 1982; Sromovsky 2005; Schmid et al. 2006; Natraj et al. 2009; Kane \\& Gelino 2010).  In this work, we present a systematic procedure to compute observables of reflected light for the different scattering phenomena under some basic assumptions. We consolidate the formalism, solution techniques, and results from analytic models that are available in the literature, but are often in scattered and obscure forms. Building on the tradition of Chandrasekhar (1950,1960), we follow the H-function approach to represent emergent fluxes for different scattering phenomena. We consider cloud-free homogeneous semi-infinite atmospheres, scattering in accordance with different scattering phase functions. We consider both conservative ($\\omega = 1.0$) and non-conservative ($\\omega <  1.0$) scattering. We compute geometric albedos and phase curves for Rayleigh scattering, with scalar and vectorial phase functions, isotropic scattering, asymmetric scattering, and Lambert scattering. We also compute polarization curves for the case of Rayleigh scattering.  We provide a step-by-step procedure to obtain the disk-integrated emergent fluxes and polarization, with the appropriate angular transformations to the celestial reference frame. The phase curves and polarization curves are also computed as a function of the mean anomaly, or a time coordinate, for given orbital parameters. We also provide analytic fits to the geometric and spherical albedos for Rayleigh scattering as a function of the scattering albedo. Thus, given an observed geometric albedo at a certain wavelength, an average scattering albedo for the atmosphere can be estimated, which can then be used to predict other observables such as the phase curve and the Stokes parameters.  In what follows, we first describe the formalism in Section~\\ref{sec:methods}. We present our results on phase curves and albedos for different scattering phase functions in Section~\\ref{sec:results}. In Section~\\ref{sec:results_inclination}, we study the dependence of polarization curves for Rayleigh scattering atmospheres on the orbital parameters. We summarize our results in Section~\\ref{sec:summary}.  \\begin{figure}[ht] \\centering %\\includegraphics[width = 0.35\\textwidth]{coord.pdf} %\\includegraphics[width = 0.35\\textwidth]{fig1.pdf} \\includegraphics[width = 0.5\\textwidth]{fig1.eps} \\caption{Planetary coordinate system. See Section~\\ref{sec:geometry} for description.} \\label{fig:coords} \\end{figure} ",
        "conclusions": "\\label{sec:summary}  Observations are becoming increasingly capable of detecting geometric albedos, phase curves, and polarization from extrasolar planets in reflected light. We have provided an integrated analytic framework to interpret such observations. Our approach makes accessible a computationally efficient procedure with which to estimate an average scattering albedo ($\\omega$) from an observable in reflected light, such as a geometric albedo or a phase curve. We follow the H-function approach of Chandrasekhar (1950,1960) which has been used in subsequent studies to derive analytic solutions for scattered emergent flux for various scattering phenomena. We consolidate the formalism, solution techniques, and results from analytic models available in the literature, but often scattered in various papers, and present a systematic procedure to compute albedos, phase curves, and polarization of reflected light.  In this work, we considered cloud-free, homogeneous, semi-infinite atmospheres scattering in accordance with different scattering phase functions. We consider both conservative ($\\omega = 1.0$) and non-conservative ($\\omega < 1.0$) scattering. We compute geometric albedos and phase curves for Rayleigh scattering, with scalar and vector phase functions, isotropic scattering, asymmetric scattering, and Lambert scattering. We also compute polarization curves for the case of vector Rayleigh scattering, and provide a step-by-step procedure to obtain the disk-integrated emergent flux and polarization, for a given scattering albedo and phase angle, with the appropriate angular transformations to the celestial reference frame. The phase curves and polarization curves are also computed as a function of the mean anomaly, or a time coordinate, for given orbital parameters.  Rayleigh scattering is a dominant scattering mechanism in cloud-free planetary atmospheres. An accurate calculation of the geometric albedo ($A_g$) for Rayleigh scattering must include the effect of polarization, using the full Rayleigh phase matrix for single-scattering, contrary to the customarily adopted scalar phase function. For, conservative Rayleigh scattering ($\\omega = 1.0$), $A_g = 0.7977$, as opposed to the familiar value of 0.75 obtained using the scalar treatment. The scalar formulation leads to geometric albedos that are lower than the actual values by up to 9\\%, depending on $\\omega$. We compute albedos and phase curves for Rayleigh scattering for a range of $\\omega$, and report analytic fits to the geometric and spherical albedos as a function of $\\omega$. Based on the polarization curves, we also compute the peak polarization due to Rayleigh scattering as a function of $\\omega$. The emergent intensity and degree of polarization of reflected light for Rayleigh scattering depend strongly on $\\omega$. The emergent intensity at full phase, and, hence, the geometric albedo, both increase with $\\omega$. On the other hand, the degree of polarization ($P$) decreases with increasing $\\omega$.  Motivated by recent efforts to measure polarization of reflected light from exoplanets, we compute polarization curves for homogeneous Rayleigh scattering atmospheres over a wide range of orbital parameters using the vector Rayleigh phase matrix. We explore the dependence of the Stokes parameters on the orbital inclination and eccentricity. We demonstrate how the inclination of a planetary orbit can be constrained by monitoring the extrema in the polarization phase curve of the planet, as well as the ellipticity of the orbit in the Q-U plane.  We summarize the results of our work as follows: \\begin{itemize} \\item{We present an analytic framework to interpret observables of reflected light (phase curves, geometric albedos, and polarization parameters) from extrasolar planets. Our analytic models assume cloud-free homogeneous and semi-infinite atmospheres, and we consider several different scattering phase functions.} \\item{We show that observations of phase curves of exoplanets in reflected light can be used to constrain the underlying scattering mechanisms (e.g. Lambert versus Rayleigh, etc.) in their atmospheres.} \\item{We show that observed geometric albedos of exoplanetary atmospheres can be used to constrain their average single scattering albedos, which are indicative of their chemical and thermal properties, assuming different scattering mechanisms.} \\item{We demonstrate, using the vector Rayleigh phase matrix, how polarization curves can be used to constrain orbital parameters (inclinations and eccentricities) of exoplanets.} \\end{itemize}  The simplification of an analytic approach does introduce some natural caveats. While our requirement of semi-infinite atmospheres is reasonable for all gaseous atmospheres, from super-Earths to giant planets, our assumption of homogeneity is meant to represent only an average scattering albedo in the atmosphere in the observed spectral/photometric bandpass. In general, the scattering albedo in a gaseous atmosphere would not remain constant at all altitudes and wavelengths. Our approach is aimed at providing a simple and efficient procedure to interpret reflected light observations of exoplanets which are currently limited by sparse data.  Detailed numerical models exist in the literature which solve the full radiative transfer problem with an inhomogeneous atmosphere in plane-parallel atmospheres. However, given the large number of free parameters and computational complexity of such models, it is generally not feasible to use them for formally fitting and retrieving atmospheric parameters from the limited data in reflected light. The analytic tools provided in our work makes it feasible to derive averaged atmospheric properties, such as an average scattering albedo or a representative scattering phase function, from the data. Our code can be used in conjunction with standard parameter  estimation methods to retrieve average scattering albedos from observed geometric albedos, as well as orbital parameters from observed phase curves and polarization curves of extrasolar planets. Our code is freely available on request by email to either of the authors."
    },
    "1112/1112.4169_arXiv.txt": {
        "abstract": "\\noindent We consider the growth of cosmological perturbations to the energy density of dark matter during matter domination when dark matter is a scalar field that has undergone Bose-Einstein condensation.  We study these inhomogeneities within the framework of both Newtonian gravity, where the calculation and results are more transparent, and General Relativity.  The direction we take is to derive analytical expressions, which can be obtained in the small pressure limit.  Throughout we compare our results to those of the standard cosmology, where dark matter is assumed pressureless, using our analytical expressions to showcase precise differences.  We find, compared to the standard cosmology, that Bose-Einstein condensate dark matter leads to a scale factor, gravitational potential and density contrast that increase at faster rates. ",
        "introduction": "Astrophysical observations indicate that 23\\% of the energy density of the Universe is of an unknown nonbaryonic form, known as dark matter.  While a precise explanation remains elusive, a weakly interacting, nonrelativistic massive particle is favored.  Leading candidates include supersymmetric particles, such as neutralinos, and the axion, originally proposed to solve the strong CP problem in QCD.  Dark matter is often modeled as a pressureless, nonrelativistic particle, known as cold dark matter.  While this model has achieved significant success, in particular with early universe and large-scale cosmology, it meets with difficulty on galactic scales.  Cold dark matter simulations of galactic halo formation predict density profiles with a central cusp \\cite{navarro}, while observations indicate constant density cores \\cite{burkert}.  Scalar field dark matter that has undergone Bose-Einstein condensation \\cite{sin,leekoh,hubak,silmal,wang,mfs,sikyan,sahni} has been considered as a solution to this problem since the resulting density profiles agree with observed rotation curves \\cite{hubak,BH,harkocusp, chavanis2}.  At low temperatures, where wave aspects dominate, a many-body system of bosons exhibits Bose enhancement, whereby bosons favor joining highly populated states.  As a result, bosons can pile up in the same ground state forming a coherent matter wave of macroscopic size known as a Bose-Einstein condensate (BEC).  This will occur when the thermal de Broglie wavelength, $\\lambda = \\sqrt{2\\pi\\hbar^2/2mk_B T}$, begins to exceed the interparticle spacing, $n^{-1/3}$, so that wave functions of individual bosons begin to overlap, where $m$ is the boson mass and $n$ its number density.  Equivalently, Bose-Einstein condensation occurs when the temperature, $T$, drops below the critical temperature, \\begin{equation} T_c = \\frac{1}{m} \\frac{2\\pi \\hbar^2}{k_B [\\zeta(3/2)]^{2/3}} n^{2/3}, \\end{equation} where $\\zeta(x)$ is the Riemann-Zeta function.  The experimental realization of trapped BECs in dilute alkaline atoms in 1995 \\cite{anderson} has led to a renewed interest in BECs, which has been a subject that unifies many disciplines, for example neutron stars \\cite{glendenning}, superconductivity \\cite{mahan} and, what is our interest here, dark matter \\cite{sin,leekoh,hubak,silmal,wang,mfs,sikyan,sahni}.  A great virtue in the study of BECs is that in ultracold atomic systems important system parameters, including the interaction strength between particles and the dimensionality of the system, can be tuned precisely.  This opens up the possibility of using ultracold atomic systems in the laboratory to simulate phenomena on galactic and cosmological scales.  Possible examples include creating a system where the BEC is subject to so-called electromagnetically induced ``gravity\" with a $1/r$ interatomic attractive potential \\cite{odell} and creating a controlled explosion of atoms by suddenly making the s-wave scattering length negative, a phenomenon dubbed ``bosenova\" because of its resemblance, on a vastly lower energy scale, to the core collapse in a supernova \\cite{donley}.  Our focus here is with the cosmological applications of BECs, in particular BEC dark matter.  In addition to the density profile and rotation curves mentioned above, investigations of BEC dark matter include the study of vortex formation \\cite{silmal, yumor, brook, kl, shapiro}, additional aspects of galactic structure \\cite{galactic}, Bose-Einstein condensation in the early universe \\cite{urena,fukuyama,harko2} and axions \\cite{sikivie}.  Recently a study of its cosmology was initiated by Harko \\cite{harko} and Chavanis \\cite{chavanis} (see also \\cite{fukuyama,velten,kamion,khlopov}). They investigated the evolution of inhomogeneities in the dark matter energy density.  Such inhomogeneities eventually become nonlinear and lead to galaxy formation and indirectly affect anisotropies in the cosmic microwave background radiation.  They derived evolution equations for the density contrast in Newtonian gravity \\cite{chavanis} and post-Newtonian gravity \\cite{harko,chavanis} and presented numerical solutions to these equations.  In this work we also analyze the evolution of cosmological inhomogeneities after dark matter has undergone Bose-Einstein condensation.  Our departure from previous results is twofold.  First, we present simple analytical solutions which allow for a precise understanding of how BEC dark matter differs from standard cold dark matter.  Since dark matter is believed to have a small pressure (indeed, it is often modeled as having zero pressure), by taking the small pressure limit such analytical solutions are obtainable.  Second, while we will begin with Newtonian gravity, we derive the evolution equations using the complete theory of General Relativity.  The use of General Relativity is necessary when considering superhorizon perturbations, which are beyond the reach of Newtonian gravity.  It is also necessary for considering anisotropies in the radiation spectrum.  We do not study anisotropies here, as it is outside the scope of our work, but such a study would be interesting and important and would rely on our results.  In the next section we review the Gross-Pitaevskii equation coupled to the Poisson equation and the Thomas-Fermi approximation, which has become the standard framework for describing a gravitating BEC.  We then derive the equation of state for BEC dark matter and use it to analyze the homogeneous, unperturbed cosmology.  In section \\ref{inhomogeneities} we derive the evolution equations for inhomogeneities during matter domination, first in Newtonian gravity and then in General Relativity.  Throughout we compare our results to those of the standard cosmology, where dark matter is pressureless.  When making these comparisons we shall refer to such dark matter as standard cold dark matter (SCDM).  We conclude in section \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  In this work we studied the resulting cosmology when dark matter is a scalar field that has undergone Bose-Einstein condensation.  Such a model of dark matter has been shown to be in better agreement with the density profiles of galactic halos than standard cold dark matter in the $\\Lambda$CDM model. We focused on the growth of inhomogeneities, i.e. perturbations to the dark matter energy density.  Since, as is well known, such perturbations only grow appreciably during matter domination, we considered only this epoch.  Such an analysis had been considered in earlier work \\cite{harko,chavanis} within the context of Newtonian and post-Newtonian gravity.   Our analysis differed from these papers principally in two ways.  First, the direction we took was to derive simple, analytical formulas that clearly showcased the modifications BEC dark matter produces in the standard cold dark matter solutions.  Second, while we did make our analysis first within Newtonian gravity, we then used the complete theory of General Relativity, finding solutions for the scale factor, gravitational potential and density contrast.  We found analytically that each one of these quantities increases at a faster rate compared to when dark matter is in the form of standard cold matter, consistent with the numerical results in \\cite{harko,chavanis}.  Our fully relativistic solutions for the gravitational potential in (\\ref{grphisoleta}) and the density contrast in (\\ref{grsola}) are valid for both subhorizon and superhorizon perturbations, the latter of which is beyond the reach of Newtonian gravity.  These fully relativistic solutions are also necessary  for studying anisotropies in the radiation spectrum.  While we did not include an analysis of anisotropies here, such an analysis we expect to lead to new and interesting physics and would rely on our results."
    },
    "1112/1112.1320_arXiv.txt": {
        "abstract": "Gravitational potentials in the cosmos are deeper than expected from observed visible objects, a phenomenon usually attributed to dark matter, presumably in the form of a new fundamental particle. Until such a particle is observed, the jury remains out on dark matter, and modified gravity models must be considered. The class of models reducing to MOdified Newtonian Dynamics (MOND) in the weak field limit does an excellent job fitting the rotation curves of galaxies, predicting the relation between baryonic mass and velocity in gas-dominated galaxies, and explaining the properties of the local group. Several of the initial challenges facing MOND have been overcome, while others remain. Here I point out the most severe challenge facing MOND. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.3113_arXiv.txt": {
        "abstract": "If Jupiter and the Sun both formed directly from the same well-mixed proto-solar nebula, then their atmospheric compositions should be similar.  However, direct sampling of Jupiter's troposphere indicates that it is enriched in elements such as C, N, S, Ar, Kr, and Xe by 2--6$\\times$ relative to the Sun \\citep{wla08}.  Most existing models to explain this enrichment require an extremely cold proto-solar nebula which allows these heavy elements to condense, and cannot easily explain the observed variations between these species.  We find that Jupiter's atmospheric composition may be explained if the Solar System's disk heterogeneously accretes small amounts of enriched material such as supernova ejecta from the interstellar medium during Jupiter's formation.  Our results are similar to, but substantially larger than, isotopic anomalies in terrestrial material that indicate the Solar System formed from multiple distinct reservoirs of material simultaneously with one or more nearby supernovas \\citep[\\eg][]{tba07}.   Such temporal and spatial heterogeneities could have been common at the time of the Solar System's formation, rather than the cloud having a purely well-mixed `solar nebula' composition. ",
        "introduction": "The Solar System's composition reflects that of the initial cloud core from which it formed.  On small bodies such as the terrestrial planets, subsequent thermal and chemical processes have altered the composition.  However, for the Solar System's two largest bodies -- the Sun and Jupiter (mass 1~\\mjup\\ $\\approx$ 0.001~\\msol) -- the atmospheric composition is relatively stable against change since the time of formation.  In particular, the relative atmospheric abundances of gases carbon (C), nitrogen (N), sulfur (S), argon (Ar), krypton (Kr), and xenon (Xe) are believed to be fixed at the time of formation 4.5 billion years ago, and not affected substantially by subsequent evolution.  The noble gases are stable against loss due to Jupiter's high escape velocity, and stable against sinking by condensation because they don't combine with other atomic species.  Elements C, N, and S combine into larger molecules such as CH$_4$, NH$_3$, and H$_2$S, but are believed to remain well mixed in Jupiter's troposphere without settling to the interior.  Jupiter's composition at 5--20 bars was directly measured by the Galileo descent probe's mass spectrometer in 1995. Initial results \\citep[][hereafter AMN03]{amn03} showed that Jupiter was enhanced by $\\sim3\\pm1\\times$ relative to the then-current solar abundances of \\citet[][hereafter AG89]{ag89}.  More recent work has better determined the Solar abundances, leading to revised Jupiter:Solar abundances.  Somewhat surprisingly, these new results actually increased the scatter between the elements, to the broader range of $4\\pm2\\times$ \\citep[][hereafter WLA08]{wla08}.  The largest change occurred in the abundance of Jovian Ar, which more than doubled to $5.4\\pm1.1\\times$ solar \\citep[][hereafter GAS07]{gas07}.  A subsequent measurement of solar Ar reduces it again to $2.45\\pm0.7\\times$ \\citep{lod08}(hereafter Lod08).  All enrichments in this paper are in terms of number density relative to hydrogen, normalized to the Sun; \\ie, $(n_i/n_H) / (n_i/n_H)_{\\rm{sun}}$. ",
        "conclusions": "The model we propose here is a departure from existing models for the solution to Jupiter's metallicity problem. Although historically most Solar System formation models have assumed a homogeneous `Solar nebula' composition, our model explicitly assumes the opposite.  We incorporate the fact that the composition of the solar nebula can change spatially and temporally.  These changes occur during the Solar System's first 5--10~Myr, as the Sun is traveling through its birth cloud, experiencing the environmental effects of other stars.  The inclusion of this contamination reflects the latest understandings of the environmental processes affecting star formation.     Our model has three general advantages over the existing amorphous ice and clathrate models that we describe in \\S~\\ref{sect:previous}.  First, it relaxes the very strict low-temperature requirements for the formation of Jupiter's solids.  Second, it allows for Jupiter to have formed at its present location without migration.  Third, it explains the fact that different elements are enriched by different amounts (\\ie, not a uniform 3$\\times$).  Our model is not without problems.  Our fits for N are at the edge (or a little beyond) the $1\\sigma$ level.  Nitrogen is poorly measured in the Galileo probe data, but is a major species and an important constraint.  We could increase the contribution from stellar winds (which supply much of the N), but this would increase the amount of C beyond that observed.  Low-mass stars of 4--8~\\msol\\ produce large quantities of N sufficient to solve the problem, but these stars have lifetimes of 30~Myr or more before they enter the RSG phase.  The other specie our models are all low in, Xe, has a particularly high condensation temperature (making it easy to condense), and this may in part explain its abundance.  Second, the amount of SN contamination required is on the high end of what can be supplied by nearby SNs using direct injection of ejecta into a nearby GMC.  Additional work is necessary to understand the mixing and concentration mechanisms of this ejecta.  However, if the `droplet' model of SN ejecta is correct, then concentrated ejecta can be slowed and cooled over kpc distances, which would solve this problem.  Most work suggests that Saturn, Uranus, and Neptune accreted their atmospheres from the disk in much the same way Jupiter did.  The metallicity enrichments of the outer planets exceed Jupiter's -- Uranus and Neptune have 30 times the Solar C:H ratio, for instance -- suggesting that they were formed at least in part by the traditional chemical condensation series \\citep{lew72}.  The colder nebula and slower evolution make volatile condensation easier at greater distances.  However, `accreted pollution' of the disk from the ISM would affect these planets as well.  Because of their different positions in the disk and their different formation times, our model cannot \\textit{predict} the enrichment they may receive from pollution, except that it could be similar to Jupiter's.  (The outer disk, with its larger cross-section, might be more easily contaminated, but this requires more study.)   More detailed analysis of this awaits better knowledge of their bulk atmospheric compositions, which are largely unknown today.  Jupiter's oxygen composition will be probed by Juno in 2017 and will provide a discriminant between many models including our own (which predicts a global O:H ratio of $\\sim$ 2.6$\\times$ solar) and that of Lodders~2004 (which assumes $\\sim 0.35\\times$ solar).  Jupiter's atmosphere may have formed through a variety of processes, of which polluted accretion could be only one.  For instance, N is among the most difficult species to condense in the various ice models, as it requires a temperature $\\leq$~30~K.  Nitrogen, however, is easily supplied in the cool, slow stellar winds that we study here. It is possible that N was supplied by these winds, and other species were delivered in part by icy planetesimals from a relatively warm ($\\geq$ 50~K) disk.  An advantage of the SN model is that it provides a mechanism for matching not only the chemical abundances which we study here, but also isotopic differences.  It has already been shown that SNs may explain the isotopic differences seen in the Solar System's rocky bodies (\\S~\\ref{sect:isotopes}); additional work on the isotopic differences between the Sun and Jupiter would be a natural extension of the present work.  This path is particularly ripe for future study given the recent Solar composition results from the Genesis mission \\citep[\\eg,][]{wbh07,mmh07}.  The general process of accretion from ISM $\\rightarrow$ disk $\\rightarrow$ planet that we describe here has broader applications for the formation of extrasolar planets.  A strong observed correlation exists between the metallicity of the host star and the existence of extrasolar planets \\citep[\\eg,][]{ufq07}.  We predict that the correlation between stellar and planet metallicities will depend on the formation environment.  Stars that form in dense regions, in the presence of massive stars, will have the highest disk metallicities and thus the greatest chance of forming planets."
    },
    "1112/1112.2029_arXiv.txt": {
        "abstract": "Gamma-ray bursts (GRBs) are brief but intense emission of soft $\\gamma-$rays, mostly lasting from a few seconds to a few thousand seconds. For such kind of high energy transients, their isotropic-equivalent-energy ($E_{\\rm iso}$) function may be more scientifically meaningful when compared with GRB isotropic-equivalent-luminosity function ($L_{\\rm iso}$), as the traditional luminosity function refers to steady emission much longer than a few thousand seconds. In this work we for the first time construct the isotropic-equivalent-energy function for a sample of 95 bursts with measured redshifts ($z$) and find an excess of high-$z$ GRBs. Assuming that the excess is caused by a GRB luminosity function evolution in a power-law form, we find a cosmic evolution of $E_{\\rm iso}\\propto(1+z)^{1.80^{+0.36}_{-0.63}}$, which is comparable to that between $L_{\\rm iso}$ and $z$, i.e., $L_{\\rm iso}\\propto(1+z)^{2.30^{+0.56}_{-0.51}}$ (both $1\\sigma$). The evolution-removed isotropic-equivalent-energy function can be reasonably fitted by a broken power-law, in which the dim and bright segments are $\\psi(E_{\\rm iso})\\propto E_{\\rm iso}^{-0.27\\pm0.01}$ and $\\psi(E_{\\rm iso})\\propto E_{\\rm iso}^{-0.87\\pm0.07}$, respectively ($1\\sigma$). For the cosmic GRB formation rate, it increases quickly in the region of $0 \\leq z \\lesssim 1$, and roughly keeps constant for $1\\lesssim z \\lesssim 4$, and finally falls with a power index of $-3.80\\pm2.16$ for $z\\gtrsim 4$, in good agreement with the observed cosmic star formation rate so far. ",
        "introduction": "Gamma-ray bursts (GRBs) are among the brightest cosmological explosions in the universe, mostly lasting from a few seconds to a few thousand seconds in soft $\\gamma$-ray. Thanks to quick follow-up observations in optical band, the redshifts of some GRBs have been measured by detecting the absorbtion lines of their afterglows or the emission lines of their host galaxies. So far the number of {\\it Swift}-detected GRBs with known redshifts has grown up to about one hundred and thus makes a reliable statistical analysis possible. Among various statistical works, the luminosity function as well as the cosmic formation rate of GRBs are particularly interesting. The luminosity function is a measure of the number of bursts per unit luminosity, which sheds light on the energy release and emission mechanism of GRBs. The cosmic formation rate is a measure of the number of events per comoving volume and time, which can help us understand the production of GRBs in various stages of the universe.  The isotropic-equivalent luminosity ($L_{\\rm iso}$) function of GRBs was firstly assumed to be a standard candle \\citep{b26,b31} and later more realistic shapes of the luminosity function were derived \\citep{b32,b33,b34,b39,b40}. The cosmic GRB formation rate has also been extensively investigated \\citep[e.g.,][]{b4,b5,b6,b7,b8,b9,b32,b15,b36}. For example, \\citet{b4} found a correlation between the variability degree of the prompt gamma-ray light curve and the luminosity and then adopted it to estimate the luminosities/redshifts of 220 bright long GRBs detected by {\\it CGRO}/BATSE. \\citet{b7} used this GRB sample to estimate the luminosity function evolution and the cosmic formation rate of GRBs. Wei \\& Gao (2003) find that there is a tight correlation between the peak energy of the prompt emission spectrum ($E_{\\rm peak}$) and the luminosity. Recently, \\citet{b38} find that both short and long GRBs all comply with this correlation. Using this correlation, \\citet{b9} estimated the luminosities/redshifts of the 689 BATSE GRBs and hence derived their luminosity function and the formation rate. Essentially these two works with simulated GRB redshifts reached quite similar results, i.e., the GRB formation rate increases quickly in the region of  $ 0 \\leqslant z \\leqslant 1$ and keeps constant up to $z \\sim 10$, which is inconsistent with the cosmic star formation rates (SFRs) inferred from UV, optical, and infrared observational data so far \\citep{b11,b12,b13,b14}.  The original concept of luminosity function comes from astrophysical objects such as stars and galaxies which are long-lasting and quite stable in releasing their energy. For GRB-like high energy transients, the total isotropic-equivalent energy ($E_{\\rm iso}$) released in the whole duration of one event can be reliably measured, and its function (i.e.,  the number density of bursts per $E_{\\rm iso}$ interval) likely provides an independent or even more representative clue on the underlying physics. That's why in this work we focus on the so-called ``isotropic-equivalent-energy function'' rather than the traditional luminosity function.  This paper is arranged as follows. \\S~2 introduces our sample and data selection. \\S~3 presents the statistical technique while \\S~4 shows the results. We adopt a robust, nonparametric statistical technique to derive the isotropic-equivalent-energy function and the cosmic formation rate of GRBs from a $E_{\\rm iso}-z$ sample. For comparison, the results from a $L_{\\rm iso}-z$ GRB sample are also presented. In \\S 5, we discuss the implication of our results and compare the cosmic GRB formation rate with the observational cosmic star formation rate. Throughout the paper, we use the standard $\\Lambda$ cold dark matter cosmology with the typical parameters $\\Omega_{\\rm m} =0.27$, $\\Omega_{\\Lambda} =0.73$, and ${\\rm h}=0.7$.    \\section[]{Data analysis} In this work, two sets of data are analyzed. The $E_{\\rm iso}-z$ sample comes from \\citet{b1,b2}, containing 95 long GRBs and X-Ray Flashes (XRF, i.e. particularly soft bursts). This sample is made up of two parts. The first part consists of 70 long GRBs from \\citet{b1} and their redshifts range from 0.033 to 6.3. The $E_{\\rm iso}$ values in this work are slightly different from those in \\citet{b1} because of the different cosmological parameters adopted in two works. The second part is from \\citet{b2} without any modification. All the GRB spectra have been extrapolated and corrected to [1,10000] keV in the cosmological restframe.  The $L_{\\rm iso}-z$ sample is from \\citet{b15} ($L_{\\rm iso}$ represents the isotropic peak luminosity, and the time resolution is 1 $s$). Due to the {\\em Swift}/BAT narrow energy band, only a small fraction of bursts have a well determined spectrum. In order to obtain reasonable estimates of $L_{\\rm iso}$ for all bursts, \\citet{b15} considered the characteristic Band function, i.e., \\begin{equation} N(E)= \\left\\{ \\begin{array}{l} A(\\frac{E}{100{\\rm keV}})^\\alpha \\exp(-\\frac{E}{E_0}),~~~~~~~~~~~~~~~~~~~~~~{\\rm for}~E\\leq (\\alpha - \\beta)E_0\\\\ A(\\frac{E}{100{\\rm keV}})^\\beta [\\frac{(\\alpha - \\beta)E_0}{100{\\rm keV}}]^{\\alpha - \\beta} {\\rm exp}(\\beta - \\alpha),~~~{\\rm for}~E\\geq (\\alpha - \\beta)E_0\\\\ \\end{array} \\right. \\end{equation} where the characteristic parameters $(E_{\\rm peak},~\\alpha,~\\beta)$ taken as $(511~{\\rm keV},~-1,~-2.25)$, respectively. To account for the so-called k-correction, again all spectra have been extrapolated and corrected to [1,10000] keV in the cosmological restframe. To test whether the above Band function applicable to all bursts, \\citet{b15} preformed Monte-Carlo simulation in order to compare GRBs having simulated spectraal parameters with GRBs having measured spectral parameters. The result of such a simulation demonstrates robustness of the $L_{\\rm iso}$ sample adopted above. ",
        "conclusions": "GRBs are brief but intense emissions of soft $\\gamma-$ray, lasting from a few seconds to a few thousand seconds. For GRB-like high energy transients, the total isotropic-equivalent energy, $E_{\\rm iso}$, can be reliably measured, and the number density of bursts per $E_{\\rm iso}$ interval may provide an independent or even more representative clue on the underlying physics of GRBs. In this work, using a sample containing 95 bursts with measured redshifts, we for the first time constructed the isotropic-equivalent-energy function and then adopted it to estimate the GRB formation rate. The fluence-truncation effect has been properly addressed by adopting a $\\tau$ statistical technique. We find there exists cosmic evolution between $E_{\\rm iso}$ and $z$, i.e., $E_{\\rm iso}\\propto g_{\\rm k}(z)=(1+z)^{1.80^{+0.36}_{-0.63}}$ (see Fig.~2), which is comparable with that  between $L_{\\rm iso}$ and $z$, i.e., $L_{\\rm iso}\\propto(1+z)^{2.30^{+0.56}_{-0.51}}$. Our finding of the evolution is largely consistent with previous findings using GRB samples with simulated redshifts \\citep{b7,b10,b9}.    We also find that the evolution power index $k$ is {\\em not} sensitive to the chosen value of fluence/flux as long as it is around the instrumental sensitivity. Changing the fluence limit from $ 8\\times 10^{-7}{\\rm erg}\\,{\\rm cm}^{-2}$ to $ 5\\times 10^{-8}{\\rm erg}\\, {\\rm cm}^{-2}$, the best-fit index $k$ of $E_{\\rm iso}$ sample only varies from 1.80 to 2.05, changing the flux limit from $ 5\\times 10^{-8}{\\rm erg}\\, {\\rm s^{-1}} {\\rm cm}^{-2}$ to $ 1\\times 10^{-8}{\\rm erg}\\, {\\rm s^{-1}} {\\rm cm}^{-2}$, the best-fit index $k$ of $L_{\\rm iso}$ sample only varies from 2.30 to 2.65, all well within the 1$\\sigma$ uncertainty.  As can been seen, the evolution of $E_{\\rm iso}$ is comparable with that of $L_{\\rm iso}$. This might indicate that the durations of GRBs do not evolve significantly with redshift and that the existence of the isotropic-equivalent-energy evolution may reflect the evolution of typical physical parameters of GRB progenitors.  After removing the redshift dependence, the isotropic-equivalent-energy function can be reasonably fitted by a broken power-law. For the dim and bright segments, we have $\\psi(E_{\\rm iso})\\propto E_{\\rm iso}^{-0.27\\pm0.01}$ and $\\psi(E_{\\rm iso})\\propto E_{\\rm iso}^{-0.87\\pm0.07}$, respectively (see Fig.~3). The shape of the energy function (see Fig.~3) is similar to that of the luminosity function (see Fig.~6). Moreover, our indices of $(-0.27,~-0.87)$ are comparable to the indices of $(-0.29,~-1.02)$ reported by \\citet{b9}.  The GRB formation rate as a function of redshift is also calculated. The connection between (long-duration) GRBs with broad-lined Type Ic supernovae \\citep[e.g.,][]{b23,b24,Woosley06,Fan11} suggests that GRB progenitors are very massive, short-lived stars, leading to expectation that the cosmic GRB formation rate would nicely follow the cosmic star formation history \\citep{b25,b26,b27,b28,b29}. So far, the observed star formation rate follows $\\rho(z)\\propto (1+z)^{3.4}$ for $z<1$, $\\propto (1+z)^{-0.3}$ for $1<z<4$ and $\\propto (1+z)^{-3.5}$ for $z>4$ \\citep{b20,b21} (see Fig.~10). Our results (see Fig.~10) show that the GRB formation rate increases quickly for $z\\lesssim1$, then roughly keeps constant for $1\\lesssim z \\lesssim4$, and finally decreases at higher redshift {\\it with a power index of -3.8}  (i.e., $\\propto (1+z)^{-3.8\\pm 2.16}$), in good agreement with the star formation rate. In future a larger GRB sample with measured redshifts would better address this topic and make progress in the fields of massive star formation and GRB physics in the early universe.  \\begin{figure} \\centering \\includegraphics[width=9cm]{fig10.eps} \\caption{GRB rate comparing with SFR. The data of SFR come from \\citet{b22,b20,b21}. We normalized the first point of GRB rate and SFR for convenience. \\label{fig10}} \\end{figure}   Our results are based on the assumption of a power-law redshift evolution of the $E_{\\rm iso}$. Besides introducing the evolution of the isotropic-equivalent-energy function, other models can also reproduce the observed data, such as one could enhance high-$z$ GRB rate by introducing a metallicity preference of GRBs or increase the total number of GRBs (e.g., Virgili et al. 2011). These models might be distinguished by the future observations."
    },
    "1112/1112.0730_arXiv.txt": {
        "abstract": "We present a detailed theoretical analysis of a non-thermal escape of molecular hydrogen from Mars induced by collisions with hot atomic oxygen from martian corona. To accurately describe the energy transfer in O + H$_2(v,j)$ collisions, we performed extensive quantum-mechanical calculations of state-to-state elastic, inelastic, and reactive cross sections. The escape flux of H$_2$ molecules was evaluated using a simplified 1D column model of the martian atmosphere with realistic densities of atmospheric gases and hot oxygen production rates for the low solar activity conditions. An average density of the non-thermal escape flux of H$_2$ of $1.9\\times10^5$ cm$^{-2}$s$^{-1}$ was obtained considering energetic O atoms produced in dissociative recombinations of O$_{2}^{+}$ ions. Predicted rovibrational distribution of the escaping H$_2$ was found to contain a significant fraction of higher rotational states. While the non-thermal escape rate was found to be lower than Jeans flux for H$_2$ molecules, the non-thermal escape rates of HD and D$_2$ are significantly higher than their respective Jeans rates. The accurate values of non-thermal escape fluxes of different molecular isotopes of H$_2$ may be important in analyses of evolution of the martian atmosphere. The described molecular ejection mechanism is general and expected to contribute to atmospheric escape of H$_2$ and other light molecules from planets, satellites, and exoplanetary bodies. ",
        "introduction": "The interaction of the martian atmosphere with the solar radiation and interplanetary plasma results in its evaporation due to thermal (Jeans) escape and a number of non-thermal mechanisms. The absence of the intrinsic magnetic field on Mars and its low gravitational potential make the martian atmosphere particularly susceptible to erosion \\citep{1998Sci...279.1676A}. The current low atmospheric pressure can be explained well by extrapolating the escape rates to a geological time frame while accounting for a change of the solar activity in the past \\citep{2004P&SS...52.1039C}. The present day degradation of the martian atmosphere occurs mainly via non-thermal escape processes induced by ion charge-exhange, sputtering and ionospheric outflows driven by solar wind \\citep{2004P&SS...52.1039C,2008SSRv..139..355J}. The dissociative recombination (DR) of O$_{2}^{+}$ is a major source of hot O atoms in the upper atmosphere of Mars, responsible for the escape of oxygen and formation of martian hot corona \\citep{1988Icar...76..135I,1993GeoRL..20.1747F,2005SoSyR..39...22K}.   Nascent hot O atoms collide with thermal constituents of the martian atmosphere and eject them from the planetary gravitational field, if a sufficient kinetic energy transfer occurs. Suprathermal neutral oxygen was shown to be important for analyses of Mars' corona and the non-thermal escape of neutral atoms \\citep{2005SoSyR..39...22K,2006SoSyR..40..384K,2009Icar..204..527F}. Recent calculations of the non-thermal He escape from Mars carried out with accurate energy transfer parameters predicted a significant He escape flux induced by hot O atoms \\citep{2011GeoRL..3802203B}.  In this Letter we explore collisional ejection of molecules from the martian atmosphere. Specifically, we report the results of a quantum-mechanical study of the energy transfer from the hot O to H$_2$ molecules and their subsequent escape. Significant computational difficulties arise from the fact that molecular internal rotational and vibrational degrees of freedom can be excited in collisions. In addition, the reactive pathway leading to the production of OH molecules is energetically permitted. To account for the increased complexity, we have computed the cross sections for O($^3$P) + H$_2$ reactive collision using fully quantum-mechanical approach. Kinetic theory was used to calculate the rate of energy transfer, as well as distributions of excited rotational and vibrational (RV) states of the recoiled H$_2$ molecules. The total escape flux of H$_2$ from Mars and RV distributions of escaping molecules have been evaluated for the low solar activity conditions. Also, we have estimated the non-thermal escape fluxes of HD and D$_2$ and compared them to the corresponding Jeans escape rates. Finally, the dependence of the molecular ejection fluxes on the gravitational escape threshold is analyzed for conditions present on other planets, satellites, and exoplanets. ",
        "conclusions": "We find that the collisionally-induced outflow of H$_2$ molecules and their heavier isotopes contributes to the evolution of the martian atmosphere. Namely, the escape rate of molecular hydrogen induced by collisions with hot oxygen from the martian atmosphere was calculated and found to be about six times smaller than the corresponding Jeans escape rate for the low solar activity. For heavier molecules, the collisional escape will dominate over thermal, as we have illustrated in case of HD and D$_2$ isotope molecules. In fact, the described process of molecular ejection induced by collisions may be one of the most important escape mechanisms of HD and D$_2$ from Mars. Consequently, the calculated escape fluxes of H$_2$ isotopologues could be important in analyses of the H/D ratio on Mars and evolution of water in martian history.  The described mechanism of molecular escape, where collisions provides sufficient translational energy to exceed the escape threshold and simultaneously excite internal molecular degrees of freedom, is rather general. It could be used to evaluate non-thermal escape fluxes and RV distributions of heavier molecules, such as CO, N$_2$, or CH$_4$, from Mars, Solar system bodies, and exoplanets. For the escape flux induced by O atoms produced in DR the upper limit on the mass of the escaping molecule is about $30 \\; u$ on Mars. Similarly, we estimate that the non-thermal escape of H$_2$ from a planetary atmosphere is possible for planetary masses up to about 3.4 Earth masses. A number of solar system bodies as well as the lightest currently confirmed exoplanets belong in that mass range. These limits do not include other non-thermal sources of hot atoms.  The escaping H$_2$ molecules exhibit a characteristic internal energy distribution, with a significant fraction of populated higher rotational states. Since H$_2$ molecules do not have a permanent electric dipole moment, they decay to the ground state mainly via collisions with the background gases present in an extended planetary corona. This is true for all escaping molecules that do not have ther permanent dipole moment. It could be possible to indirectly detect the presence of H$_2$ or other rotationally excited light molecules in the extended martian corona from a careful analysis of the collision rates and abundancies of the excited coronal species. Finally, a significant amount of rovibrationally excited H$_2$ molecules remain in the martian atmosphere after colliding with hot O atoms. The cross sections and energy transfer parameters presented in this study can be used to determine non-thermal translational and rovibrational distributions of hot H$_2$ gas in the upper atmosphere of Mars."
    },
    "1112/1112.5053_arXiv.txt": {
        "abstract": "{Studies of dense molecular-cloud cores at (sub)millimetre wavelengths are needed to understand the early stages of star formation.} {We aim to further constrain the properties and evolutionary stages of dense cores in Orion B9. The prime objective of this study is to examine the dust emission of the cores near the peak of their spectral energy distributions, and to determine the degrees of CO depletion, deuterium fractionation, and ionisation.} {The central part of Orion B9 was mapped at 350 $\\mu$m with APEX/SABOCA. A sample of nine cores in the region were observed in C$^{17}$O$(2-1)$, H$^{13}$CO$^+(4-3)$ (towards 3 sources), DCO$^+(4-3)$, N$_2$H$^+(3-2)$, and N$_2$D$^+(3-2)$ with APEX/SHFI. These data are used in conjunction with our previous APEX/LABOCA 870-$\\mu$m dust continuum data.} {All the LABOCA cores in the region covered by our SABOCA map were detected at 350 $\\mu$m. The strongest 350 $\\mu$m emission is seen towards the Class 0 candidate SMM 3. Many of the LABOCA cores show evidence of substructure in the higher-resolution SABOCA image. In particular, we report on the discovery of multiple very low-mass condensations in the prestellar core SMM 6. Based on the 350-to-870 $\\mu$m flux density ratios, we determine dust temperatures of $T_{\\rm dust}\\simeq7.9-10.8$ K, and dust emissivity indices of $\\beta\\sim0.5-1.8$. The CO depletion factors are in the range $f_{\\rm D}\\sim 1.6-10.8$. The degree of deuteration in N$_2$H$^+$ is $\\simeq 0.04-0.99$, where the highest value (seen towards the prestellar core SMM 1) is, to our knowledge, the most extreme level of N$_2$H$^+$ deuteration reported so far. The level of HCO$^+$ deuteration is about 1--2\\%. The fractional ionisation and cosmic-ray ionisation rate of H$_2$ could be determined only towards two sources with the lower limits of $\\sim2-6\\times10^{-8}$ and $\\sim2.6\\times10^{-17}-4.8\\times10^{-16}$ s$^{-1}$, respectively. We also detected D$_2$CO towards two sources.} {The detected protostellar cores are classified as Class 0 objects, in agreement with our previous SED results. The detection of subcondensations within SMM 6 shows that core fragmentation can already take place during the prestellar phase. The origin of this substructure is likely caused by thermal Jeans fragmentation of the elongated parent core. Varying levels of $f_{\\rm D}$ and deuteration among the cores suggest that they are evolving chemically at different rates. A low $f_{\\rm D}$ value and the presence of gas-phase D$_2$CO in SMM 1 suggest that the core chemistry is affected by the nearby outflow. The very high N$_2$H$^+$ deuteration in SMM 1 is likely to be remnant of the earlier CO-depleted phase.} ",
        "introduction": "New stars in our Galaxy form predominantly in the so-called giant molecular clouds (GMCs). The nearest GMC to the Sun is the Orion molecular cloud complex\\footnote{In this paper we adopt a distance of 450 pc to the Orion GMC (\\cite{genzel1989}). The actual distance may be somewhat smaller as, for example, Menten et al. (2007) determined the distance to the Orion Nebula to be $414\\pm7$ pc.}. This cloud complex contains a number of star-forming regions, one example being the Orion B9 region in the central part of Orion B mole\\-cular cloud. Orion B9 represents a very early stage of star formation which is manifest in the fact that most ``dense cores'' in the cloud are either prestellar or contain Class 0 protostars. These objects are cold, and their study requires observations at (sub)millimetre wavelengths.  Using the LABOCA bolometer array on APEX (Atacama Pathfinder EXperiment), we mapped that region at 870 $\\mu$m (\\cite{miettinen2009}; hereafter Paper I). The dust continuum mapping resulted in the discovery of twelve dense cores of which four were found to be associated with IRAS point sources, and eight of them appeared previously unknown sources. Of the newly disco\\-vered cores, two (we called SMM 3 and 4) were found to be associated with \\textit{Spitzer} 24- and 70-$\\mu$m sources, and were tentatively classified as Class-0 protostellar candidates. Six cores showed no signs of embedded infrared emission, and were thus classified as starless. Figure~\\ref{figure:laboca} shows the 870 $\\mu$m emission as contours overlaid on the \\textit{Spitzer}/MIPS 24-$\\mu$m image of the central part of Orion B9\\footnote{Comparison of our LABOCA map to the new SABOCA map together with \\textit{Spitzer} images showed that the LABOCA observations were not well-pointed. The SABOCA peaks appeared to systematically lie southeast from the LABOCA peak positions. The difference in angular resolution between SABOCA and LABOCA data did not appear to be the cause of the offset in peak positions; by smoothing the SABOCA map to correspond the resolution of our LABOCA data, and regridding the maps onto the same grid, the offsets still remained. We used the brightest SABOCA 350-$\\mu$m source in the region, SMM3, as a reference to adjust the pointing of our LABOCA map. The pointing of the LABOCA map was shifted by $(\\Delta \\alpha,\\, \\Delta \\delta)=(+1\\farcs72, \\, -9\\farcs97)$. Unfortunately, the target positions of our molecular-line observations, which were chosen to be the LABOCA peak positions from our previous map, are now slightly offset from the dust maxima (but still within the beam size for most line observations).}. In Paper I, we also derived the degree of deuteration, or the $N({\\rm N_2D^+})/N({\\rm N_2H^+})$ co\\-lumn density ratio, towards selected positions and found the values in the range 0.03--0.04, comparable to those seen in other low-mass star-forming regions (e.g., \\cite{crapsi2005}; \\cite{emprechtinger2009}; \\cite{friesen2010b}). Taking advantage of the H$_2$D$^+$ data from Harju et al. (2006), the ionisation degree and the cosmic-ray ionisation rate of H$_2$ towards the same target positions were estimated to be $x({\\rm e})\\sim10^{-7}$ and $\\zeta_{\\rm H_2}\\sim1-2\\times10^{-16}$ s$^{-1}$, respectively.  The physical properties of dense cores in Orion B9 were studied further by Miettinen et al. (2010; hereafter Paper II). Using the observations of the NH$_3$ $(1,\\,1)$ and $(2,\\,2)$ inversion transitions performed with the Effelsberg 100-m telescope, we determined the gas kinetic temperature, kinematical properties, and dynamical state of the cores. The gas kinetic temperature of the cores were found to be $T_{\\rm kin}\\sim9.4-13.9$ K, and the internal non-thermal motions in the cores appeared to be subsonic, or at most transonic. The virial-parameter analysis showed that the starless cores in the region are likely to be gravitationally bound, and thus prestellar objects. Interestingly, some of the cores were found to have significantly (by $\\sim5-7.5$ km~s$^{-1}$) lower radial velocity than the systemic velocity of the region (${\\rm v_{\\rm LSR}}\\sim9$ km~s$^{-1}$). This suggests that they belong to the ``low-velocity part'' of Orion B which is likely to originate from the feedback from the massive stars of the nearby Ori OB 1b association.  In this paper, we present the results of our new APEX observations of the Orion B9 region. These include the dust continuum mapping with the SABOCA (Submillimetre APEX Bolometer Camera) bolometer array at 350 $\\mu$m, and observations of the molecular-line transitions C$^{17}$O$(2-1)$, H$^{13}$CO$^+(4-3)$, DCO$^+(4-3)$, N$_2$H$^+(3-2)$, and N$_2$D$^+(3-2)$. The SABOCA 350-$\\mu$m observations allow us to probe the peak of the core's spectral energy distributions (SEDs). We also aim to study the degrees of CO depletion and deuterium fractionation, and the fractional ionisation in the cores. These are useful parameters to further constrain the physical and chemical properties of the sources, and their evolutionary stages.  This paper is structured as follows. The observations and data-reduction procedures are described in Sect.~2. The direct observational results are presented in Sect.~3. In Sect.~4, we des\\-cribe the analysis and present the results of the physical and chemical properties of the cores. Discussion of our results is presented in Sect.~5, and in Sect.~6, we summarise the main conclusions of this study.  \\begin{figure*} \\begin{center} \\includegraphics[width=0.85\\textwidth]{laboca_new.ps} \\caption{\\textit{Spitzer}/MIPS 24-$\\mu$m image of the central part of Orion B9 overlaid with contours showing the LABOCA 870-$\\mu$m dust continuum emission. The contours go from 0.1 ($\\sim3.3\\sigma$) to 1.0 Jy~beam$^{-1}$ in steps of 0.1 Jy~beam$^{-1}$. The 24-$\\mu$m image is shown with a logarithmic scaling to improve the contrast between bright and faint features. The colour bar indicates the 24-$\\mu$m intensity scale in units of MJy~sr$^{-1}$. The green plus signs show the target positions of our molecular-line observations (i.e., the submm peak positions of the LABOCA map before adjusting the pointing; see text). The green line shows the base and northwest tip of the HH 92 jet driven by IRAS 05399-0121 (\\cite{bally2002}). The 0.2-pc scale bar and the effective LABOCA beam HPBW ($\\sim20\\arcsec$) are shown in the bottom left (Paper I).} \\label{figure:laboca} \\end{center} \\end{figure*} ",
        "conclusions": "We have carried out a (sub)mm study of dense cores in Orion B9. We used APEX to map the region at 350 $\\mu$m, and to observe the transitions of C$^{17}$O, H$^{13}$CO$^+$, DCO$^+$, N$_2$H$^+$, and N$_2$D$^+$. These data were compared with our previous LABOCA 870-$\\mu$m data. The principal aim of this study was to investigate the dust emission of the cores near the peak of their SEDs, and the degrees of CO depletion, deuterium fractionation, and ionisation in the sources. Our primary results are summarised as follows:  \\begin{enumerate} \\item All the 870-$\\mu$m cores within the boundaries of our SABOCA map were detected at 350 $\\mu$m. The strongest 350-$\\mu$m source in the region is SMM 3, a candidate Class-0 protostellar core. \\item Four of the 870-$\\mu$m cores, namely SMM 3, 4, 6, and 7, were resolved into at least two condensations at 350 $\\mu$m. In particular, the elongated prestellar core SMM 6 was resolved into three to four very low-mass subcondensations, showing that core fragmentation can take place during the prestellar phase of evolution. In all cases, the origin of substructure can be explained by thermal Jeans-type fragmentation. \\item The dust temperatures derived from the 350-to-870-$\\mu$m flux density ratio are very low, only $T_{\\rm dust}\\approx 7.9_{-1.8}^{+3.0}-10.8_{-2.6}^{+5.7}$ K. The corresponding gas temperatures are typically a few kelvins higher. We also derived the submm dust emissivity spectral indices using the assumption $T_{\\rm dust}=T_{\\rm gas}$, and they are in the range $\\beta\\approx 0.5\\pm0.8-1.8\\pm0.6$. The uncertainties in $\\beta$ are large, and within the errors the values are comparable to the fiducial value $\\beta=2$. \\item We refined some of the protostellar SEDs presented in Paper I by adding the observed 350-$\\mu$m flux densities. No ra\\-dical changes were found, and the sources IRAS 05405-0117 and SMM 3 and 4 are classified as Class 0 objects, in agreement with our previous results. \\item The CO depletion factors were found to lie in the range $f_{\\rm D}\\sim 1.6\\pm0.5-10.8\\pm2.2$. We found no systematic difference in $f_{\\rm D}$ between the starless and protostellar cores. In accordance with previous observations and theoretical predictions the most severe CO depletion is seen towards the core with highest average density, SMM 3. The degree of deuteration in N$_2$H$^+$ was found to be in the range $N({\\rm N_2D^+})/N({\\rm N_2H^+})\\simeq 0.035\\pm0.006-0.992\\pm0.267$. The $N({\\rm DCO^+})/N({\\rm HCO^+})$ ratio was found to be about 1--2\\%, comparable to those seen in other low-mass star-forming regions. \\item The fractional ionisation could only be derived towards IRAS 05399-0121 and SMM 4-LVC with the lower limits of $x({\\rm e})>1.5\\times10^{-8}$ and $>6.1\\times10^{-8}$, respectively. These values are comparable to the fractional ionisations we derived earlier towards a few target positions near IRAS 05405-0117 and SMM 4. The cosmic-ray ionisation rate of H$_2$ implied by the derived lower $x({\\rm e})$ limit is $\\zeta_{\\rm H_2}\\sim 2.6\\times10^{-17}$ s$^{-1}$ towards IRAS 05399-0121, and $\\sim 4.8\\times10^{-16}$ s$^{-1}$ towards SMM 4-LVC. The former value, which does not suffer from line-of-sight contamination, is within a factor of two of the 'standard' value $1.3\\times10^{-17}$ s$^{-1}$. \\item The highest degree of deuteration in N$_2$H$^+$, $\\sim 0.99$, was derived towards the prestellar core SMM 1. To our know\\-ledge, this is the most extreme level of N$_2$H$^+$ deuteration reported so far. We also detected D$_2$CO emission towards SMM 1 with an abundance of several times $10^{-11}$. Because D$_2$CO is expected to be formed through the grain-surface chemistry, its presence in the gas phase in SMM 1 could be due to shocks driven by the jet from IRAS 05399-0121, resulting in the release of D$_2$CO from the grain mantles. This conforms to the low CO depletion factor of 1.9 derived towards SMM 1. The very high N$_2$H$^+$ deuteration could be remnant of the earlier CO-depleted phase, and not yet affected by the destroying effect of gas-phase CO. \\item It seems likely that the elongated clump associated with IRAS 05405-0117 and SMM 4 consists of physically independent objects that form a single clump only in projection along the line of sight. \\end{enumerate}   In Papers I and II we discussed the origin of dense cores in Orion B9. The region could have been affected by the feedback from the nearby Ori OB1b subgroup of the Orion OB1 association [resembling the process termed ``cloud shuffling'' by Elmegreen (1979)]. This feedback process may have played a role in sweeping up the gas into dense shells and filaments out of which the dense cores were later fragmented. This picture is supported by the kinematics of Orion B9 (low-velocity line-of-sight components). The present observational results suggest that the further fragmentation of cores into smaller condensations is caused by the thermal Jeans instability. The observations presented in this paper also reveal the intricate chemistry of the cores, such as heterogeneous depletion and deuteration among the cores. This suggests that the individual Orion-B9 cores, though possibly formed collectively in the parent filaments, are evolving chemically at different rates."
    },
    "1112/1112.2997_arXiv.txt": {
        "abstract": "{We present the results of N-body simulations of planetary systems formation in radiatively-inefficient disc models, where positive corotation torques may counter the rapid inward migration of low mass planets driven by Lindblad torques. The aim of this work is to examine the nature of planetary systems that arise from oligarchic growth in such discs. } {We adapt the commonly-used Mercury-6 symplectic integrator by including simple prescriptions for planetary migration (types I and II), planetary atmospheres that enhance the probability of planetesimal accretion by protoplanets, gas accretion onto forming planetary cores, and gas disc dispersal. We perform a suite of simulations for a variety of disc models with power-law surface density and temperature profiles, with a focus on models in which unsaturated corotation torques can drive outward migration of protoplanets. In some models we account for the quenching of corotation torques that arises when planetary orbits become eccentric.} {Approximately half of our simulations lead to the successful formation of gas giant planets with a broad range of masses and semimajor axes. Giant planet masses range from being approximately equal to that of Saturn, up to approximately twice that of Jupiter. The semimajor axes of these range from being $\\sim 0.2$ AU, up to $\\sim 75$ AU, with disc models that drive stronger outward migration favouring the formation of longer-period giant planets. Out of a total of 20 giant planets being formed in our simulation suite, we obtain 3 systems that contain two giants. No super-Earth or Neptune-mass planets were present in the final stages of our simulations, in contrast to the large abundance of such objects being discovered in observation surveys. This result arises because of rapid inward migration suffered by massive planetary cores that form early in the disc life time (for which the corotation torques saturate), combined with gas accretion onto massive cores that leads them to become gas giants. Numerous low mass planets are formed and survive in the simulations, with masses ranging from a few tenths of an Earth mass, up to $\\sim 3$ Earth masses. Simulations in which the quenching of corotation torques for planets on modestly eccentric orbits was included failed to produce any giant planets, apparently because Lindblad torques induce rapid inward migration of planetary cores in these systems.}   {We conclude that convergent migration induced by corotation torques operating during planet formation can enhance the growth rate of planetary cores, but these often migrate into the central star because corotation torques saturate. Outward migration of planetary cores of modest mass can lead to the formation of gas giant planets at large distances from the central star, similar to those observed recently through direct imaging surveys. The excitation of planetary eccentricities through planet-planet scattering during oligarchic growth may quench the effects of corotation torques, however, such that inward migration is driven by Lindblad torques. } ",
        "introduction": "\\label{section:intro}  Observations of extrasolar planets are providing a picture of a highly diverse population of bodies orbiting main sequence stars. At the extreme ends of the distribution, there exist very short-period low-mass rocky planets such as CoRoT-7b \\citep{Leger2009} and Kepler-10b \\citep{Batalha2011}, and very long-period massive gas giant planets detected in recent years by direct imaging \\citep{Marois, Kalas, Lagrange}. In addition, there have been discoveries of short-period hot-Jupiters such as 51-Pegb \\citep{MayorQueloz}, hot-Neptunes such as Gliese 436b \\citep{Butler2004}, and super-Earths such as Gliese 581c \\citep{Bonfils2005}. Multiple planet systems are common, examples being the five planet system orbiting the star 55 Cancri consisting of gas giants and Neptune-mass bodies \\citep{Fischer2008}, and the recently reported Kepler-11 system, consisting of six nearly coplanar low-mass planets \\citep{Lissauer2011}. A question that needs to be addressed is whether or not a particular model of planet formation, such as the core-accretion model, can explain this broad diversity, appealing to variety in the planet forming environment to explain the range of observed systems. Or is it the case that quite different physical processes are operating on different length and/or time scales within protoplanetary discs to form planets with very different characteristics. For example, core-accretion operating on long-times scales at relatively close locations to the central star to form short-period systems, and disc gravitational fragmentation occurring at large radii on short-time scales to form long-period giant planets.  To begin addressing this question in detail, it is necessary to construct global models of planetary formation that allow for the formation of multiple planet systems with a diversity of masses and orbital elements (semimajor axes, eccentricities etc). In this paper, we present the results from global models, that are based on the oligarchic growth scenario for planet formation, that have been constructed using a symplectic N-body code \\citep{cham99}, in conjunction with simplified prescriptions for the gas disc model, planetary migration, capture of planetesimals, gas-envelope accretion, and disc dispersal on Myr time scales. One of our main objectives in this work is to examine how our current understanding of migration of low-mass planets influences the formation of planetary systems, with particular emphasis on the corotation torques in radiatively-inefficient discs \\citep{PBCK2010, PBK2011}. As we anticipate that this is the first paper in a series that will examine this issue, the prescriptions we have adopted in this initial study for a number of physical processes, such as gas accretion, are necessarily very simplified. They serve the useful purpose, however, of enabling N-body simulations to be performed of planetary system formation that lead to a diversity of outcomes, and these can be used to guide future model developments.  There is a substantial body of previous work that has examined the role of migration in the formation of planets using N-body simulations. \\cite{PapaloizouLarwood2000} examined planetary growth through planet-planet collisions using N-body simulations combined with prescriptions for migration and eccentricity/inclination damping. \\cite{McNeilDuncan2005} and \\cite{daisaka} examined the effect of type I migration on terrestrial planet formation, and \\cite{FoggNelson2007,FoggNelson2009} examined the the influence of type I migration on the formation of terrestrial planets in the presence of migrating Jovian-mass planets. \\cite{TerquemPapaloizou2007} examined the formation of hot-super-Earths and hot-Neptunes using N-body simulations with type I migration. \\cite{McNeilNelson2009, McNeilNelson2010} carried out large-scale simulations of oligarchic growth to explore the formation of systems of  hot-Neptunes and super-Earths, such those around the stars Gliese 581 and HD698433, using a novel symplectic integrator with multiple timesteps. An alternative to these approaches has been planetary population synthesis modelling, as presented by \\cite{IdaLin2008, IdaLin2010}, \\cite{Mordasini2009a}, \\cite{Mordasini2009b}, and \\cite{Miguel2011}. These monte-carlo approaches have significant advantages in being able to cover a very broad range of parameter space, allowing meaningful statistical comparisons with observational data to be undertaken. The computational efficiency also allows complex models of planetary atmospheres and gas accretion to be incorporated, as presented by the Mordasini et al work. Accurate treatment of planet-planet gravitational interactions are difficult to include in these models, however, such that predictions about planetary system multiplicity, orbital eccentricities and inclinations are not a natural outcome of the models (we note that models by \\cite{IdaLin2010} now include a simplified treatment of planet-planet interaction dynamics).  This paper is organised as follows. In Sect.~\\ref{method} we present the numerical methods and the physical model. In Sect.~\\ref{initial} we present the initial conditions for the simulations. Results are described and analysed in Sect.~\\ref{endresults}, comparisons are made to observations in Sect.~\\ref{realobs} and a discussion and concluding remarks are provided in Sect.~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} Although the models presented in this paper include a broad range of physical processes relevant to planet formation (migration; planetary growth through mutual protoplanet collisions; accretion of planetesimals and gas; planet-planet gravitational interactions), we have adopted a number of assumptions and simplifications that inevitably affect the realism of the simulations and their results. These include:  \\begin{enumerate} \\item \\emph{Simulation domain}. Even though we have modelled a relatively wide semimajor axis domain with our initial solid matter disc compared to previous N-body work on planetary formation, it is clear that accurate modelling of discs in which significant corotation torques arise requires as wide a domain as possible. Protoplanets move significantly in the disc with some forming at 2-3 AU and migrating out to 70-80 AU. Similarly, planets migrate inward and ought to traverse the terrestrial planet region which we have not modelled. Planets forming in the terrestrial region might also migrate out into the regions that we have investigated. In short, the migration behaviour observed in the simulations presented in this paper indicates that all regions of the disc are coupled during planet formation, and it is no longer sensible to think in terms of a ``terrestrial planet region'' or a ``giant planet region''. As such, a suitable model would involve a domain ranging from as far in as 0.1 AU out to approximately 50 AU. Such a simulation is beyond current modelling techniques because of the required numbers of protoplanets and planetesimals, even for the method presented by \\cite{McNeilNelson2009, McNeilNelson2010} which utilises multiple time steps in a parallel symplectic integrator. Instead, such global models of planetary formation are probably going to require efficient use of modern GPU technology.  \\item \\emph{Gas disc model}. We currently model the gas disc as having fixed power-laws in surface density and temperature, with the disc mass undergoing exponential decay to mimic its viscous and photoevaporative evolution. In reality, the disc is heated by the central star and through local viscous dissipation, and it cools through radiative emission. A significant improvement to the model that we are in the process of implementing will be to evolve the disc surface density and temperature explicity using a 1+1D numerical solution, as described in \\cite{PapaloizouTerquem1999}, for example. This approach will be similar to that described in \\cite{FoggNelson2009} and references therein, and will allow gap formation and type II migration to be simulated directly, along with UV photoevaporation of the disc during its final stages of dispersal.  \\item \\emph{Planetary atmosphere model and enhanced capture radii.} As described in the preceding sections, rapid planetary growth is assisted by the enhanced accretion of planetesimals through implementation of a model for planetary atmospheres that increases the effective accretion cross section \\citep{inaba03}. Although this model works well when accretion is dominated by planetesimals, it is probably inaccurate when accretion includes giant impacts between protoplanets. In particular, a planetary atmosphere can be completely liberated from a planet when it is impacted by a body whose mass is similar to that of the atmosphere, and our implementation of the atmosphere model does not account for this effect. The atmosphere model would clearly be improved in its accuracy if it responded to giant impacts as well as planetesimal accretion rates.  \\item \\emph{Gas envelope accretion}. Our model for gas accretion during the formation of gas giants is very rudimentary, although it serves the purpose of allowing gas giant planets to form in our simulations. While a full accretion model for each planet similar to those presented in \\cite{pollack96} and \\cite{movshovitz} would be ideal, this is computationally beyond the reaches of an N-body code that can model planetary systems formation over Myr time scales. However, there are improvements that can be made that will allow local conditions in the disc to influence the gas accretion rate onto a planet. Coupled with a more sophisticated disc model that allows explicit modelling of gap formation and gas accretion, such an approach would alleviate the requiement to set an artifical upper limit for the planet mass that can form through runaway gas accretion.  \\end{enumerate}   We have presented the results of simulations that include recent torque formulae for type I migration (including Lindblad and corotation torques), with gas envelope accretion, enhanced capture radii due to gas atmospheres, and planet-planet gravitational dynamics included. We have surveyed a range of disc models which all allow for outward migration driven by corotation torques. The main results that we have obtained may be summarised as follows: \\begin{itemize} \\item Convergent migration of protoplanets, and the rapid accretion of planetesimals, can cause the rapid growth of planetary cores to masses in excess of 10 M$_{\\oplus}$ within 0.5 Myr in most disc models. This leads to rapid inward migration of these bodies, driven by Lindblad torques, when the horseshoe libration time scale becomes significantly shorter than the thermal/viscous diffusion time scale and the corotation torques saturate.  \\item Formation of planetary cores with a few Earth masses $\\ge 0.5$ Myr after the simulations are initiated can lead to their migration into the outer regions of the disc (30 - 50 AU). Steady mass growth through gas accretion onto the planet can counterbalance the slow inward migration that occurs as the gas disc mass reduces, and long-term survival in the outer disc can lead to gas giant planet formation there when runaway gas accretion ensues. This mode of giant planet formation was found to be a common outcome in our simulations, especially those with disc models that generate strong outward migration, leading to numerous gas giant planets being formed between semimajor axes 10 - 60 AU. Models such as these could potentially explain the long-period giant planets discovered recently through direct imaging \\citep{Marois, Kalas, Lagrange}.  \\item Out of 40 simulations that used disc models whose masses were either 3 or 5 times more massive than the MMSN, 19 gas giant planets were formed. Most of these are similar in mass to Jupiter (in part because of the gas accretion prescription that was adopted in the models), and are formed at large distances from the star. Short period Jovian mass planets were also formed, however, along with a pair of Saturn-mass bodies at intermediate ($\\sim 1$ AU) orbit distances. These latter systems were formed in discs that generate weaker corotation torques than those that tend to generate the longer-period giant planets.  \\item Multiplanet systems containing more than one giant planet were found to be a rare outcome (3 out of 40 simulations produced two giant planets each), and this has the additional effect of producing systems with very small eccentricities and inclinations due to the low rate of occurence of planet-planet scattering events. In fact, the only planets to be formed with significant eccentricities were a pair of closely separated Jovian-mass objects that underwent significant dynamical interaction.  \\item Our simulations completely fail to produce super-Earth or Neptune-mass planets. This appears to arise because of very rapid inward migration of planets that grow early in the disc lifetime and undergo rapid inward migration, combined with the switching-on of gas accretion that converts planets of intermediate mass into gas giants at later times. Modification of the planetary atmosphere and gas accretion prescriptions will probably result in more surviving planets with intermediate masses. Numerous planets in the Earth mass range were formed in the simulations, however. The `desert' of super-Earths and Neptunes is similar to that reported in the planetary population synthesis models of \\cite{IdaLin2008}, and occurs for much the same reasons as theirs (rapid gas accretion and migration).  \\item Simulations performed where the corotation torque is attenuated when planet eccentricities grow to become larger than the dimensionless horseshoe width appear to produce results quite different from those in which this effect is neglected (i.e. all the runs described above). In particular the growth and survival of planets is reduced because mutual encounters between protoplanets maintains the typical eccentricities above the critical value for which corotation torques diminish. In these latter simulations, no gas giant planets were formed at all. Further work is required to establish the influence of corotation torques on planet formation {\\it via} oligarchic growth, where planet-planet interactions maintain finite values of the eccentricity.  \\end{itemize}  Our models demonstrate that strong corotation torques can substantially alter the qualitative outcomes of planet formation simulations. Even the simplest model of planet formation that involves non linear gravitational interaction between protoplanets and planetesimals during planetary accumulation is by its nature a chaotic process. Given a well-defined distribution of initial protoplanet/planetesimal masses and orbital elements from which individual formation models are drawn, however, an ensemble of such models should give rise to a distribution of outcomes with well-defined statistical properties. Allowing the set of initial conditions to be drawn from a range of disc models whose life-times and radial profiles of density and temperature also have well-defined distributions serves only to modify the distribution of outcomes, as does including additional physical processes such as type I migration. A corollary of this argument is that increasing the complexity of migration processes, as we have done in this paper, also serves to modify the distribution of outcomes in a quantifiable manner. Corotation torques, however, increase the dependency of formation outcomes on the details of the underlying disc model and microphysical processes such as those that control the opacity of disc material. The dependency of migration on opacity, turbulent viscosity and other disc properties, and the role of migration in shaping planetary system architectures, increases the need for more refined observations of protoplanetary disc properties and improved disc models to allow planetary formation calculations to be compared with data on extrasolar planetary systems in a meaningful way. To summarise: planetary formation is a chaotic process, but is deterministic in a statistical sense. Corotation torques do not change the validity of this statement, however, their dependence on detailed disc properties increases the difficulty of constructing realistic planetary formation models for comparison with observational data.  This is the first in a series of papers to examine the oligarchic growth of planets under the influence of type I migration, including corotation torques. Models that include a more sophisticated treatment of the issues raised in Sect.~\\ref{conclusions} will be presented in a forthcoming publication."
    },
    "1112/1112.2673_arXiv.txt": {
        "abstract": "We examine the prospects for producing Nickel-56 from black hole accretion disks, by examining a range of steady state disk models.  We focus on relatively slowly accreting disks in the range of $\\dot{M} = 0.05$ $M_\\odot / {\\rm s}$ to $\\dot{M} = 1$ $M_\\odot / {\\rm s}$, as are thought to be appropriate for the central engines of long-duration gamma-ray bursts.  We find that significant amounts of Nickel-56 are produced over a wide range of parameter space. We discuss the influence of entropy, outflow timescale and initial disk position on mass fraction of Nickel-56 which is produced.  We keep careful track of the weak interactions to ensure reliable calculations of the electron fraction, and discuss the role of the neutrinos. ",
        "introduction": "Over the last decade and a half there has been a dramatic increase in data available for studying long duration gamma-ray bursts (GRBs). Besides the gamma rays, afterglows in the x-ray, optical and radio are seen, and host galaxies have been identified.  In a handful of cases ``supernova bumps'' are seen in the optical afterglows, e.g.  GRB 980326 \\citep{bloom99}.  Furthermore, spectra for Type Ic supernova have been identified in a few cases \\citep{galama98,hjorth03,kawabata03,matheson03,stanek03, malesani04,thomsen04}. These developments have lead to the theoretical conjecture that all long duration gamma ray bursts are the result of the collapse of a massive star; for a review see \\citet{woosley06}.  The temporal evolution of light curve spectra of traditional core collapse supernovae are driven by the decay of Nickel-56 to Cobalt-56 and the subsequent decay of Cobalt-56 to Iron-56.  The half lives for these decays are approximately six days and seventy-seven days respectively.  After these nuclei beta decay, their daughter nuclei emit gamma rays which then thermalize in the ejecta. The energy from these decays is thus re-emitted in the optical and observed as the supernova light curve.  Thus the intensity of a supernova bump should be directly influenced by the amount of $^{56}$Ni produced in the object.  On the order of tenths of solar masses of Nickel are needed.  For example, it has been estimated that around $0.5 {\\rm M}_\\odot$ of $^{56}$Ni is necessary to explain SN2003dh which accompanied GRB030329 \\citep{woosley03}.  The commonly discussed model for the type of supernova that produces a GRB is the collapsar model \\citep{woosley93,macfadyen99}. In this case the star has too much rotation to produce a standard core collapse supernova.  A hot $\\sim 10$ MeV disk surrounding a black hole is produced at the center.  Nickel-56 can be produced in two different ways.  It can be produced in explosive burning, when a shock wave propagates out through the star and pre-existing nuclei are then burned to $^{56}$Ni \\citep{maeda03,fryer06,maeda09}. This is the same mechanism as in traditional core collapse supernovae, but the elemental yields are different since the explosion energy is larger.  However, Nickel-56 can also be produced in a disk wind. This is a primary process where free neutrons and protons are ejected from the disk and combine to form nuclei \\citep{sur05,pru04,sur06}.  Any assessment of $^{56}$Ni production in GRBs must consider both mechanisms.  The latter mechanism is more difficult to model, since it requires a careful accounting of the neutrinos in the disk. In this paper we extend the preliminary studies of \\citet{sur05,pru04,pru05,sur06}, and examine the production of Nickel-56 from the outflow from a variety of disks.  We discuss not only the importance of the electron fraction in determining the abundance yields, but also the role of entropy and outflow timescale.  Recent studies, e.g. \\citet{maeda09} and references therein, have used estimates of ejected $^{56}$Ni mass to constrain the GRB central engine.  This study provides a reference for anyone considering Nickel-56 production in models of long duration gamma-ray bursts. ",
        "conclusions": "Nickel production in the outflows from black hole accretion disks may be an important component of the total GRB Nickel-56 production.  We have outlined the ways in which $^{56}$Ni can be formed by black hole accretion disks under a variety of conditions, taking careful account of the neutrino interactions.  We find that entropy is a strong diagnostic across a broad range of accretion disk models.  Nickel production appears to be optimal for only modest increases in entropy over the in-disk values; for the disks studied here with $\\dot{m}>0.1$, this can be accounted for via neutrino heating alone.  We note that nickel can also be produced via thermonuclear burning from a shock that goes through a star, if the disk is associated with a hypernova explosion.  However, in models without this additional production mechanism, one could constrain the required conditions for nickel production in the object.  For example, if accretion disk outflows are to be a significant fraction of the overall GRB nickel production, the contribution of other processes, such as magnetic heating, to the entropy increase in the outflow from $\\dot{m}>0.1$ disks should not greatly exceed that of the neutrino interactions.  However, conclusions about the mass fraction of Nickel produced cannot be drawn from the entropy alone.  Nickel formation is optimized when the outflow timescale, in particular the ratio of $\\dot{r}/r$ in the outflow, is fast at the beginning stages of nucleosynthesis and slow later.  Furthermore, in the models considered here, outflows from an intermediate disk radius, $r_{0}\\sim 100$ km, more efficiently produced nickel than material ejected closer to the black hole.  Under the optimal conditions up to around 50\\% of the material ejected from black hole accretion disks will become Nickel-56.  These results, when coupled with future wind model calculations and observations, can ultimately be used to constrain the outflow conditions of GRB disks."
    },
    "1112/1112.3958_arXiv.txt": {
        "abstract": "We use \\emph{Hubble Space Telescope} imaging to show that the nearby dwarf starburst galaxy NGC 4449 has an unusual abundance of luminous red star clusters.  Joint constraints from integrated photometry, low-resolution spectroscopy, dynamical mass-to-light ratios, and resolved color-magnitude diagrams provide evidence that some of these clusters are old globular clusters. Spectroscopic data for two massive clusters suggest intermediate metallicities ([Fe/H] $\\sim -1$) and subsolar Mg enhancement ([Mg/Fe] $\\sim -0.1$ to $-0.2$). One of these clusters may be the nucleus of a tidally disrupting dwarf galaxy; the other is very massive ($\\sim 3\\times10^{6} M_{\\odot}$). We have also identified a population of remote halo globular clusters. NGC 4449 is consistent with an emerging picture of the ubiquity of stellar halos among dwarf galaxies, and study of its globular clusters may help distinguish between accretion and \\emph{in situ} scenarios for such halos. ",
        "introduction": "NGC 4449 is one of the nearest starburst galaxies, at a distance of $\\sim 3.8$ Mpc (Annibali \\etal~2008). Its stellar mass is similar to that of the Large Magellanic Cloud (LMC) but has a much higher star formation rate of $\\sim 1 M_{\\odot}$ yr$^{-1}$ (McQuinn \\etal~2010). Extended complexes of \\ion{H}{1} gas and disturbed kinematics have led to the conclusion that NGC 4449 has recently undergone a substantial interaction with another galaxy, although the interaction partner has not been conclusively identified (Hunter \\etal~1998).  Annibali \\etal~(2008) imaged NGC 4449 with the \\emph{Hubble Space Telescope} Advanced Camera for Surveys ($HST$/ACS) and visually identified a number of resolved star clusters. They found a very luminous cluster on the west side of the galaxy and suggested that it could be the nucleus of a disrupting dwarf galaxy.  A surprising number of other massive clusters are apparent from a quick perusal of the $HST$ images. However, unlike the case for an early-type galaxy, many studies have shown (e.g., Seth \\etal~2004; Georgiev \\etal~2009) that a large fraction of luminous star clusters in star-forming dwarf galaxies are likely to be young ($< 1$ Gyr) or of intermediate-age (1--5 Gyr).  In this paper we present a photometric and spectroscopic study of star clusters in NGC 4449, focusing on the older clusters, showing that the galaxy has an unusual abundance of old massive star clusters. ",
        "conclusions": "\\subsection{The Overabundance of Massive Star Clusters in NGC 4449}  The LMC is the dwarf galaxy whose star cluster system is perhaps the best studied, so it makes a useful counterpoint to NGC 4449. First we compare the stellar masses of NGC 4449 and the LMC. The extinction-corrected integrated absolute magnitude of NGC 4449 is $M_B = -18.9$ and $M_K = -20.7$ (Paturel \\etal~2003; Jarrett \\etal~2003). Using mass-to-light ratios from Bell \\etal~(2003) for $(B-V)_{0} = 0.33$, the average of the $B$ and $K$-band implied total stellar mass is $2.6\\times10^{9} M_{\\odot}$. In a review of the properties of the LMC, van der Marel (2004) reports a stellar mass of $2.7\\times10^{9} M_{\\odot}$, nearly identical to that of NGC 4449. So, despite the obvious differences in recent star formation between NGC 4449 and the LMC, the similarity in their overall stellar masses suggests that a direct comparison of their old star cluster populations is reasonable.S  The LMC has sixteen old star clusters (GCs) with ages $> 5$ Gyr (Mackey \\& Gilmore 2004). Fifteen of these have very old ages ($\\sim 12$ Gyr); the other is the unusual low-mass cluster ESO121-SC03, which is slightly younger at 9 Gyr (Mackey \\etal~2006). Perhaps surprisingly, the integrated luminosities of most of the LMC GCs are known only imprecisely. Here we preferentially use the $V$-band luminosities from McLaughlin \\& van der Marel (2005), which are derived from King model fits to the surface brightness profiles of the clusters. Luminosities are available from this source for twelve of the sixteen GCs, including all of the more massive GCs. For the other four GCs we take values compiled in other literature sources (Mackey \\& Gilmore 2004; Pessev \\etal~2008); these values are less certain.  As discussed in \\S 2.1, at the distance of NGC 4449, $V = 20$ corresponds very closely to $M_V = -8$, which we take as a convenient limit for massive GCs. There are five LMC GCs at or above this luminosity; the most luminous is NGC 1916 ($M_V = -8.9$ or  $\\sim 6\\times10^{5} M_{\\odot}$). We do not yet have a full census of GCs in NGC 4449, but we have found at least 19 objects (15 from ACS, 4 from SDSS) above this luminosity for which all available evidence suggests that they are old GCs, with a handful of additional luminous candidates.  Assuming a GC luminosity function similar to that in the Milky Way, these bright clusters would suggest the total number of GCs in NGC~4449 is $\\ga 69$.  This corresponds to lower limits on the GC specific frequency (a luminosity-normalized count of GCs) of $S_N \\ga 1.5$ and $T \\ga 25$, where $T$ is the number of GCs per 10$^9$~M$_\\odot$ of galaxy stellar mass.  These GC frequencies are higher than typical late-type spirals (Chandar \\etal~2001; 2004), but are not atypical of somewhat fainter dwarf irregular (Seth \\etal~2004; Georgiev \\etal~2010) and elliptical galaxies (see, e.g., Durrell \\etal~1996; Strader \\etal~2006; Miller \\& Lotz 2007; Peng \\etal~2008). The luminosity functions of the old NGC~4449 and LMC GCs are compared in Figure~6.  \\begin{figure} \\epsscale{1.2} \\plotone{lumfunc.ps} \\figcaption[lumfunc.eps]{\\label{fig:fig_s} The luminosity function of old (ages $\\ga 1$~Gyr) star clusters in NGC~4449 (solid line) compared to 16 old ($> 3$ Gyr) GCs in the LMC (dotted line). The solid line shows the luminosity function of good quality old NGC~4449 clusters selected from the HST data (Table~1) plus the confirmed SDSS sources from Table~4 (where $V$ magnitudes have been estimated from SDSS photometry).  The $HST$ sample is complete at $F555W \\lesssim 20$, however, it covers only the inner part of the galaxy, and we have included only SDSS objects that have been confirmed via spectroscopy or $HST$ imaging. Even with this incomplete sample, NGC~4449 shows a clear overabundance of massive, old star clusters with respect to the LMC. }   \\end{figure}  We cannot rule out the possibility that a subset of the NGC~4449 GCs with $V<20$, or some of the less luminous clusters, are in fact  of intermediate age ($\\sim$ 1--5 Gyr). As we state above, depending on their exact age, such objects could have masses 2--4 times lower than old clusters of the same luminosity, and the existence of intermediate-age clusters with masses $\\la 10^5 M_{\\odot}$ would not be unexpected. Through multi-band photometry (from the UV to the mid-IR), low-resolution spectroscopic age estimates, mass-to-light ratio measurements, and color-magnitude diagrams, we have shown that there is no compelling evidence that either of the two most luminous star clusters in NGC 4449 (B13 and B15) is of intermediate age, though the case of B15 is less certain. If NGC~4449 did possess luminous intermediate-age GCs, then it would be appropriate to consider intermediate-age LMC clusters in a comparison of their cluster systems, and the abundance of the NGC~4449 system would be less remarkable.  The metallicities and chemical abundance ratios of clusters in the LMC and NGC 4449 offer another puzzling comparison. The median metallicity of the 15 old GCs in the LMC is [Fe/H] $\\sim -1.9$ (Mackey \\& Gilmore 2004). This is much lower than the median spectroscopic metallicity of [Fe/H] $\\sim -1.1$ to $-1.2$ for GCs in NGC 4449, even considering that we do not have metallicities for some of the more distant, possibly metal-poor GCs in NGC 4449. If these GCs are generally ancient then the high mean metallicity is extremely surprising.  Similarly, [Mg/Fe] for LMC GCs is typically [$\\alpha$/Fe] $\\sim 0.2$--0.4 (Johnson \\etal~2006), similar to Galactic GCs, compared to the subsolar values observed for B13 and B15 in this paper. The higher values of [Fe/H] and lower values of [Mg/Fe] are evocative of Pal 12, a Galactic GC thought to have been accreted from the Sgr dwarf (Cohen 2004). Pal 12 has an age of $\\sim 8$--9 Gyr (Mar{\\'{\\i}}n-Franch \\etal~2009). Therefore, a possible interpretation is that many of the GCs in NGC 4449 could have been made in a burst of cluster formation a few Gyr later than the old (12--13 Gyr) metal-poor GCs observed in both the Milky Way and the LMC.  \\subsection{Stellar Halos and Metal-poor Globular Clusters}  The overwhelming weight of evidence now supports the idea that the metal-poor stellar halo of the Milky Way was built largely through the accretion of less massive galaxies, partially via the direct contribution of stars from these dwarf galaxies, and partially from \\emph{in situ} stars propelled to larger radii by the accretion events (Searle \\& Zinn 1978; Ibata \\etal~1994; Bullock \\& Johnston 2005; Zolotov \\etal~2009; Carollo \\etal~2010). Similar accretion events are observed in M31 (Ibata \\etal~2001) and other nearby disk galaxies (Mart{\\'{\\i}}nez-Delgado \\etal~2010). The metal-poor GC systems of these galaxies trace their stellar halo stars in both chemistry and spatial distribution (Brodie \\& Strader 2006).  Many dwarf galaxies also have extended ``halos\" of stars, if one considers stellar components that extend above the outer extrapolation of the inner surface brightness profile, but the relationship between these halos and those observed in more massive galaxies is unclear. Such a halo has indeed been observed in NGC 4449 (Ry{\\'s} \\etal~2011). Theoretically, accretion is expected to be less effective in producing stellar halos around dwarf galaxies because of the decreasing efficiency of star formation in less massive halos (Purcell \\etal~2007). Alternative methods for the formation of stellar halos in dwarf galaxies, such as by internal feedback, have been proposed (e.g., Stinson \\etal~2009).  Ry{\\'s} \\etal~(2011) do not find evidence for a stellar population gradient in the halo of NGC 4449, with no change in the color of the RGB or the relative number of carbon stars. Thus there is no direct evidence for a metal-poor component in this halo. However, the existence of metal-poor GCs shows that such stars are in fact present. Using the color--metallicity relation of Sinnott \\etal~(2010), the $g-z$ colors of the four confirmed outer SDSS GCs correspond to metallicities of [Fe/H] = $-1.1$ to $-1.6$. These findings add to a growing picture that dwarf galaxies (e.g., NGC 6822; Hwang \\etal~2011) have halo GCs.  As mentioned above, there is separate evidence for ongoing accretion (a luminous stellar stream) in NGC 4449 (Mart{\\'{\\i}}nez-Delgado \\etal~2011). The combination of these two observations---a metal-poor stellar population and a stellar steam---are consistent with the idea that the halo of NGC 4449 could indeed have been built by accretion, just as is the case for more massive galaxies. The discovery of more distant GCs (perhaps requiring $HST$) should enable kinematic studies of the halo GC system of the galaxy, which may help distinguish between the \\emph{in situ} and accretion models for the formation of stellar halos in dwarf galaxies."
    },
    "1112/1112.6050_arXiv.txt": {
        "abstract": "A non-ideal MHD collapse calculation employing the axisymmetric thin-disk approximation is used to resolve cloud core collapse down to the scales of the second (stellar) core. Rotation and a magnetic braking torque are included in the model, and the partial ionization resulting in ambipolar diffusion and Ohmic dissipation is calculated from a detailed chemical network. We find that a centrifugal disk can indeed form in the earliest stage of star formation, due to a shut-off of magnetic braking caused by magnetic field diffusion in the first core region. Thus, there is no catastrophic magnetic braking in a model with realistic non-ideal MHD. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.1500_arXiv.txt": {
        "abstract": "We study the time evolution of a rotating, axisymmetric, viscous accretion flow around black holes using a grid based finite difference method. We use the Shakura-Sunyaev viscosity prescription. However, we compare with the results obtained when all the three independent components of the viscous stress are kept. We show that the centrifugal pressure supported shocks became weaker with the inclusion of viscosity. The shock is formed farther out when the viscosity is increased. When the viscosity is above a critical value, the shock disappears altogether and the flow becomes subsonic and Keplerian everywhere except in a region close to the horizon, where it remains supersonic. We also find that as the viscosity is increased, the amount of outflowing matter in the wind is decreased to less than a percentage of the inflow matter. Since the post-shock region could act as a reservoir of hot electrons or the so-called `Compton cloud', the size of which changes with viscosity, the spectral properties are expected to depend on viscosity strongly: the harder states are dominated by low-angular momentum and the low viscosity flow with a significant outflows while the softer states are dominated by the high viscosity Keplerian flow having very little outflows. ",
        "introduction": "Pringle (1981) pointed out that in present of viscosity, most of the matter of the disc accretes into the black hole while most of the angular momentum is taken farther away by very little matter. Though this basic principle is the key to believe that the standard models of Shakura \\& Sunyaev (1973, hereafter SS73) and Novikov \\& Thorne (1973) are possible, there are very few numerical simulations which actually show how a Keplerian disc is formed, or whether it should always form. There are clearly two possibilities: when the companion supplies high angular matter (as in a low mass X-ray binary), viscous stress must be very large to transport this in order that the accretion may occur. On the other hand, when mainly winds from the companion (as in a high mass X-ray binary) is accreted the flow need not produce a Keplerian disc (Chakrabarti \\& Titarchuk, 1995, hereafter CT95; Smith, Heindl and Swank, 2002) since the viscosity may not be high enough to transport angular momentum rapidly to redistribute the angular momentum. It has been shown by Chakrabarti (1989) that low angular momentum, non dissipative flows produce axisymmetric standing shock waves in tens of Schwarzschild radii, but the presence of a large viscosity (Chakrabarti, 1990, hereafter C90) will remove the shock wave since the Rankine-Hugoniot relation is not be satisfied in a highly dissipative flow. A confirmation of such an assertion, originally made in the context of the isothermal flows, came both through numerical simulations (Chakrabarti \\& Molteni, 1995, hereafter CM95) as well as theoretical studies of flows with more general equation of states (Chakrabarti, 1996a, hereafter C96; Chakrabarti \\& Das, 2004; Das \\& Chakrabarti, 2004). The theoretical works showed that the topology of the flow is modified when viscosity is added and beyond a critical viscosity, the centrifugal pressure supported shocks do not form. The numerical work of CM95, which was based on the smoothed particle hydrodynamics (SPH) in a one dimensional flow, clearly demonstrated the transport of angular momentum by the viscosity. It showed that for low enough viscosity parameter $\\alpha$ (SS73), the standing shock continues to form. However, with the increase in viscosity, the shock weakens and moves outward. Ultimately, for high enough viscosity the shock disappears. Most interestingly, the post-shock subsonic flow acquires a Keplerian distribution. Thus, for high enough viscosity, the entire disc becomes subsonic and Keplerian, the assumed condition for the Shakura-Sunyaev disc. {Igumenshchev et al. (1996) studied two-dimensional flows but concentrated only the inner region of the disk, namely, the region less than $20 r_g$. The main interest was to study the transonic nature of the flow just before the matter enters into the horizon. They find that a torus like structure forms close to the black hole and the angular momentum increases outward. Lanzafame et al. (1998) carried out the SPH simulation in a region with a radial extent of $50 r_g$ in two dimensions and concentrated on the shock formation. In particular, they showed that high viscosity can remove the shock waves from the accretion flows. However, SPH method is often intrinsically dissipative and it was necessary to verify the behaviour with other simulation methods. Igumenshchev et al. (1998) used the finite difference method and allowed the heat generated by viscosity heating to be radiated away or absorbed totally. The computational box was up to $300r_g$, but the outer boundary condition was that of a near Keplerian flow having no radial velocity. The inner boundary was kept at $3r_g$. Thus the possibility of having a shock or the inner sonic point were excluded. The disc was found to be stable for very high viscosity parameter $\\alpha$ and less stable for lower $\\alpha$. Igumenshchev et al. (2000) extended earlier work by studying the dependence on the polytropic index $\\gamma$ which varied from $4/3$ to $5/3$ as well as viscosity parameter $\\alpha$ and again found that the stability of the solutions depend on these parameters. }  So far, however, the fate of the solutions which include shocks have not been studied. The formation of shock required injection of very lower angular and low viscosity flows and the inner boundary must be inside $2r_g$ in order that the inner sonic point is formed (C89). In Giri et al. (2010, hereafter Paper I), the results of standing and oscillating shock formations in inviscid flows was presented. We now extend this study by adding viscosity. The theoretical work discussed in C89 or Chakrabarti (1996a) were carried out for a one dimensional flow since they required the sonic point analysis. For a two dimensional flow, a completely self-consistent theoretical solution is not possible. This is why a numerical simulation is necessary to answer the following questions: (a) Are the conclusions based on theoretical considerations continue to remain valid for a two dimensional flow? (b) Do the shocks survive for higher viscosity? (c) How does the outflow rate depend on viscosity? (d) Are all the components of the viscous stress important in a thick accretion flow? (e) Under what condition does the Keplerian flow form?  In the next Section, we show how we introduce the viscosity into the grid-based finite difference method (Molteni, Ryu \\& Chakrabarti, 1996). We consider both the cases, namely, when only the $r\\phi$ component is present, and when all the components, namely, $r\\phi$, $\\phi z$ and $rz$ components are present. In \\S 3 \\& 4, we present the methodology and the results of our simulation respectively. Finally, in \\S 5, we draw concluding remarks. ",
        "conclusions": "In this paper, we have presented the results of the numerical simulations of a two dimensional, axisymmetric, viscous accretion flows. Three parameters, namely, the specific angular momentum, the specific energy and the viscosity parameter determine the complete solution, although results depend somewhat on the injection process at the outer boundary. While both inviscid (Chakrabarti, 1989) and viscous (Chakrabarti \\& Das, 2004) flows allow solutions with or without centrifugal barrier dominated shock waves, we concentrated mostly on the cases when the shocks are formed. We find that the shocks move outward as the viscosity is enhanced and the post-shock region roughly attains a Keplerian distribution. When the viscosity parameter is very high, the shock moves to a large distance and the whole disc becomes a Keplerian disc. We also found that the condition of standing shock wave formation may be satisfied only in a narrow range of the viscous parameter (keeping other parameters as constants) which is in line with the conclusions drawn in Chakrabarti \\& Das (2004). When the Rankine-Hugoniot conditions are not satisfied, the shocks tend to oscillate (Ryu, Chakrabarti, Molteni, 1997; Paper I) and the frequency of oscillation is decreased and the amplitude is increased as the shock moves out.  In contrast to these simulations, the simulations of Igumenshchev and collaborators (1996, 1998, 2000) concentrated on shock-free solutions. Either the simulation boundary was too small or too big with high viscosity, the shocks did not form in these simulations. Thus a direct comparison is not possible. In some of the models (e.g., CT95) the post-shock region or CENBOL is treated as the Compton cloud which plays a major role in shaping the spectrum of the accretion flows. Similarly, the CENBOL is thought to be responsible for the outflows. In Chakrabarti (1999) it was estimated that the ratio of the outflow to the inflow varies with the strength of the shock. In the present case, we varied the strength of the shock by increasing viscosity and found that the outflow rate, though fluctuating in small time scales, becomes smaller as the viscosity parameter is increased.  In Paper I, we emphasized the formation of a weaker inner shock closer to the black hole. This was not predicted by the theoretical works. In the present paper, we find that this inner shock also becomes weaker and moves outward  when the viscosity is introduced. The critical viscosity parameter for removal of the inner shock is lower than what is needed to remove the outer shock.  In Molteni, Sponholz and Chakrabarti (1996) and  Chakrabarti, Acharyya \\& Molteni (2004) it was shown that the shocks oscillated when the cooling time scale matches with the infall time scale. The simulation did not include viscosity. In Ryu, Chakrabarti and Molteni (1997) no cooling or viscosity was added, but the shocks were shown to oscillate when the Rankine-Hugoniot relation was not satisfied. In the present paper, for the first time, we show that even for a viscous flow, the same conclusion holds true as long as the viscosity is less than the critical value and the shocks are not totally removed.  It is well known that the light curves of the radiation coming from an accretion disc around a black hole exhibit quasi-periodic oscillations (QPOs). It appears that in the absence of a hard surface, the standing shock itself behaves like a hard surface and its oscillation changes the size of the post-shock region significantly. This could be the cause of the low and intermediate frequency quasi-periodic oscillations (Chakrabarti \\& Manickam, 2000; Rodriguez et al. 2004; Remillard et al. 2006; Gliozzi et al. 2010; Qu et al. 2010). When the viscosity is increased, the Keplerian rate is enhanced and at the same time, the shock recedes to a large distance and the time period is increased. Asymptotically, this means that a Keplerian disc should not show QPOs. We also observe that away from the equatorial plane the angular momentum is sub-Keplerian (i.e., smaller compared to the specific angular momentum on the equatorial plane). A picture similar to this has been postulated in the two component advective flow model (CT95) where sub-Keplerian flows flank a Keplerian disc on a equatorial plane. However, a complete picture would emerge only after radiative transfer is added.  The authors thank D. Ryu and I. Chattopadhyay for providing helps to include viscosity in the original code for non-viscous flow used in Paper I. They are also thankful to an anonymous referee for helpful suggestions."
    },
    "1112/1112.3655.txt": {
        "abstract": "We evaluate the effects of environment and stellar mass on galaxy properties at 0.85 $< z <$ 1.20 using a 3.6$\\micron$-selected spectroscopic sample of 797 cluster and field galaxies drawn from the GCLASS survey.   We confirm that for galaxies with LogM$_{*}$/M$_{\\odot}$ $>$ 9.3 the well-known correlations between environment and properties such as star-forming fraction (f$_{SF}$), SFR, SSFR, D$_{n}$(4000), and color are already in place at $z \\sim$ 1. %The stellar mass of the cluster galaxies is also correlated with f$_{SF}$, SFR, SSFR, D$_{n}$(4000), and color in a way that is comparable to the environmental trends. We separate the effects of environment and stellar mass on galaxies by comparing the properties of star-forming and quiescent galaxies at fixed environment, and fixed stellar mass.  The SSFR of star-forming galaxies at fixed environment is correlated with stellar mass; however, at fixed stellar mass it is independent of environment.  The same trend exists for the D$_{n}$(4000) measures of both the star-forming and quiescent galaxies and shows that their properties are determined primarily by their stellar mass, not by their environment.  Instead, it appears that environment's primary role is to control the fraction of star forming galaxies.  %In an attempt to understand how environment alters the f$_{SF}$ without affecting the mean properties of the galaxies Using the spectra we identify candidate poststarburst galaxies and find that those with 9.3 $<$ LogM$_{*}$/M$_{\\odot}$ $<$ 10.7 are 3.1 $\\pm$ 1.1 times more common in high-density regions compared to low-density regions.   The clear association of poststarbursts with high-density regions as well as the lack of a correlation between the SSFRs and D$_{n}$(4000)s of star-forming galaxies with their environment strongly suggests that at $z \\sim$ 1 the environmental-quenching timescale must be rapid. Lastly, we construct a simple quenching model which demonstrates that the lack of a correlation between the D$_{n}$(4000) of quiescent galaxies and their environment results naturally if self quenching dominates over environmental quenching at $z >$ 1, or if the evolution of the self-quenching rate mirrors the evolution of the environmental-quenching rate at $z >$ 1, regardless of which dominates. ",
        "introduction": "It has been known for a long time that the properties of galaxies depend strongly on both their stellar mass and their larger-scale environment.  In the last decade, the SDSS and other spectroscopic surveys have revolutionized this field, providing high-quality data that has quantified these correlations with high precision \\citep[see e.g., the review by][]{blanton09}.  It is now clear that in the local universe practically any property that can be used to characterize a galaxy shows some correlation with both its environment and stellar mass. \\newline\\indent Galaxies in denser environments are redder \\citep[e.g.,][]{Kauffmann2004, Balogh2004a, Baldry2006, vandenBosch2008, Skibba2009, vonderlinden2010}, less star forming \\citep[e.g.,][]{Gomez2003, Kauffmann2004, Balogh2004b, Weinmann2006, vonderlinden2010, Peng2010, Weinmann2010}, older \\citep[e.g.,][]{Smith2006, Cooper2010b}, more metal rich \\citep[e.g.,][]{Smith2006,Cooper2010a}, and more frequently have elliptical/S0 morphologies \\citep[e.g.,][]{Dressler1980, Postman1984, Bamford2009}. \\newline\\indent The same trends hold true for stellar mass, with more massive galaxies being redder \\citep[e.g.,][]{Bower1992, Baldry2006, Thomas2010, Peng2010}, less star forming \\citep[e.g.,][]{Weinmann2006, Baldry2006, Peng2010}, older \\citep[e.g.,][]{Kauffmann2003, Nelan2005, Gallazzi2005, Gallazzi2006, Thomas2005, Thomas2010, Graves2009}, more metal rich \\citep[e.g.,][]{Tremonti2004, Gallazzi2005, Gallazzi2006, Nelan2005, Thomas2005, Thomas2010, Graves2009} and more frequently having elliptical/S0 morphologies \\citep[e.g.,][]{Bamford2009, Nair2010}. \\newline\\indent There is also evidence for a covariance between stellar mass and environment, with more massive galaxies residing in denser environments \\citep[e.g.,][]{Kauffmann2004, Baldry2006}, although this issue has not yet reached complete consensus, see e.g., \\cite{vonderlinden2010} for an alternative result.  Either way, it is clear that any potential covariance between stellar mass and environment needs to be isolated in order to determine their relative influence on the galaxy population. \\newline\\indent Recently, \\cite{Peng2010} performed such an analysis and demonstrated that the star formation quenching effects of stellar mass and environment are independent of each other in the SDSS and zCOSMOS datasets up to $z \\sim$ 0.6.  They used this separability to construct an empirically-based model which implies that \"mass quenching\" (which we will refer to in this paper as \"self quenching\") is the primary mechanism that ends star formation in massive galaxies (LogM$_{*}$/M$_{\\odot}$ $>$ 10.6) in all environments at all redshifts; whereas \"environmental quenching\" is increasingly important for lower-mass galaxies at decreasing redshift. \\newline\\indent Although we continue to acquire higher-quality measurements of the correlations between galaxy properties and their stellar mass and environment -- and even have an empirical framework for interpreting how they quench the overall galaxy population \\citep{Peng2010}, we still have only circumstantial evidence about the physical processes that are responsible for quenching.  An obvious next step is to examine the evolution of galaxy properties as a function of stellar mass and environment over a wide range of epochs to determine the frequency of quenching and constrain the relevant timescale, an important clue for identifying the physical processes involved. \\newline\\indent In the last five years extensive spectroscopic surveys of the $z \\sim$ 1 universe such as GOODS \\citep{Elbaz2007, Popesso2011}, DEEP2 \\citep{Cooper2006, Cooper2007, Cooper2008a, Cooper2010b, Gruztbauch2011a}, zCOSMOS \\citep{Iovino2010, Peng2010, Cucciati2010}, ROLES \\citep{Li2011}, VVDS \\citep{Cucciati2006, Scodeggio2009}, as well as the HiZELS narrow-band survey \\citep{Sobral2011} have begun to quantify the effects of stellar mass and environment on galaxy properties.   Photometric surveys are now doing the same at redshifts up to $z \\sim$ 2-3 \\citep[e.g.,][]{Quadri2007, Grutzbauch2011b, Quadri2011}.  So far, this body of work has been somewhat controversial, with studies disagreeing on whether environment is an important factor in the quenching of star formation, and whether the star-formation-rate- (hereafter, SFR) density relation seen at $z \\sim 0$ still exists at $z \\sim$ 1. \\newline\\indent Although the reason for the disagreements has not been dissected in detail, much of it may be semantic in nature, and not due to errors in data analysis or sample bias. %In particular, at this point perhaps the most important question to ask is, what constitutes the SFR-density relation? %\\newline\\indent This is because some authors have studied the color-density relation, others have looked at the specific-star-formation-rate- (hereafter, SSFR) density relation, while others yet have studied the actual SFR-density relation.  Some studies correlate the mean color, mean SFRs, and mean SSFRs of particular subsamples of galaxies with environment.  Others compare the change in the {\\it fraction} of star-forming galaxies (hereafter f$_{SF}$) with environment, where star forming has been defined using any of those three criteria.  At some level all of these correlations have been analyzed under the moniker of the \"SFR-density relation\", even though they are in fact quite different quantities.  As we will show in this paper, the mean properties of galaxies and the f$_{SF}$ behave quite differently as a function of environment and stellar mass, and these different behaviors, as well as some selection effects, may be at the root of many of the past discrepancies. \\newline\\indent Another major challenge for field galaxy surveys that has been largely ignored is that environment itself is difficult to define, and that the choice of parameterization almost certainly affects the interpretation of any analysis.  Typically 3rd or 5th nearest neighbor distance is used as a proxy for galaxy environment; however, because it counts only a few galaxies this metric is noisy and only accounts for the density of galaxies on small scales.  Using a larger number of nearest neighbor galaxies provides a better measurement of the larger scale environment; however, this comes at the cost of losing small scale information.  Furthermore, increasing the number of nearest neighbors substantially leads to problems with the edges of survey fields and requires throwing out a large fraction of the data.  A comprehensive discussion of the issues associated with defining galaxy environment is presented in \\cite{Cooper2005}. %Of course, the volume probed by a survey is important as well, as nearest neighbor densities are usually defined relative to the mean density, narrower surveys probe a smaller range of environments, and hence may see little trends in the max-min environment simply because of this. \\newline\\indent Rich galaxy clusters offer an interesting alternative for measuring the effect of environment on galaxy evolution as they cover a wide range of environmental densities.  In particular, their cores are the most extreme high-density environments at any epoch, and therefore leave little ambiguity about the measurement of local environment.  Of course, the tradeoff with using clusters is that only a small percentage of galaxies in the local or high-redshift universe live in clusters, and therefore the environmental effects seen there may not be generalizable to the galaxy population at large.  Still, they offer the possibility to define the upper limits of how strong environmental effects can be at any epoch. %Certainly if we see mild effects in clusters, we would expect environment %would matter less in both groups or filaments, where most galaxies live. \\newline\\indent A handful of clusters at $z \\sim$ 1 have been studied with regard to environmental and mass evolution effects.  \\cite{Patel2009, Patel2011, Koyama2008, Koyama2010, Rosati2009, Strazzullo2010, Vulcani2010, Demarco2010b} have all found that the correlation between SFR and SSFR with environment is still preserved in clusters around $z \\sim$ 1.  In addition, \\cite{Gobat2008}, and \\cite{Rettura2010, Rettura2011} have also found that the population of quiescent galaxies in clusters at $z \\sim$ 1 is more evolved than quiescent galaxies in the field.  Interestingly, there is some evidence that at $z >$ 1.4 even the cores of clusters are becoming increasingly active \\citep[e.g.,][]{Tran2010, Hilton2010, Hayashi2010, Kuiper2010, Hatch2011}, although this effect is not seen in all clusters \\citep[e.g.,][]{Tanaka2010, Bauer2011}.  Whether this is truly the epoch where the expected change in the star formation properties of cluster galaxies actually occurs, or is simply a selection effect is still difficult to tell as our current knowledge is based on just a handful of clusters. \\newline\\indent Here we present the most comprehensive study of the relative effects of stellar mass and environment on both cluster and field galaxies at $z \\sim$ 1 to date using data from the Gemini Cluster Astrophysics Spectroscopic Survey (GCLASS, PIs: Wilson + Yee).  GCLASS is a spectroscopic survey of 10 rich clusters at 0.85 $< z < $ 1.34 using the GMOS instruments on Gemini-North and Gemini-South.  The GCLASS cluster sample has been drawn from the Spitzer Adaptation of the Red-sequence Cluster Survey \\citep[SpARCS, see][]{Muzzin2009a, Wilson2009, Demarco2010}.  In this paper we use 9 of the clusters, and examine the properties of cluster galaxies as a function of stellar mass and environment for the first time using a large sample of clusters. \\newline\\indent One of the major improvements the GCLASS data offer over other spectroscopic studies is that the spectroscopic targets have been selected based on their observed IRAC 3.6$\\micron$ flux.  This is different from other $z \\sim$ 1 field galaxy surveys such as DEEP2, which is an R-band selected survey (R $<$ 24.1), and zCOSMOS, which is an I-band selected survey (I $<$ 22.5 for zCOSMOS-bright).  At $z \\sim$ 1 these produce rest-frame U- and B-band selected samples, whereas the 3.6$\\micron$ selection from GCLASS produces a rest-frame H-band selected sample and is therefore much closer to a stellar-mass-selected sample. %Furthermore, the GCLASS sample of clusters is fairly large and follows a clear selection criteria and is representative of the galaxy populations in rich, $z \\sim$ 1 clusters. \\newline\\indent This paper is laid out in the following manner.  In $\\S$ 2 we define the cluster sample, $\\S$ 3 discusses the reduction of the spectroscopic data, and in $\\S$ 4 we explain how physical quantities such as stellar masses and SFRs are extracted from the data.  Readers interested in the main results can skip those sections and refer to $\\S$ 5 where we look at the correlation of galaxy properties for the sample of galaxies with stellar masses (LogM$_{*}$/M$_{\\odot}$) $>$ 9.3 as a function of environment and stellar mass, and $\\S$ 6 where we look at the same properties but holding galaxy type, environment, and stellar mass fixed in order to untangle their interdependence.  In $\\S$ 7 we discuss the population of poststarburst galaxies as a function of environment and stellar mass, and in $\\S$ 8 we present a simple model for the evolution of the quiescent galaxy population.  We conclude with a summary in $\\S$ 9.   Throughout this paper we assume a $\\Lambda$CDM cosmology with $\\Omega_{m}$ = 0.3, $\\Omega_{\\Lambda}$ = 0.7, and H$_{0}$ = 70 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "In this paper we have studied the correlations between the properties of galaxies and their environment and stellar mass.  We first considered the average properties of a mass-limited sample of galaxies with LogM$_{*}$/M$_{\\odot}$ $>$ 9.3 ($\\S$ 5), and then subdivided the galaxies into different types (star-forming or quiescent) and studied the properties as a function of environment at fixed stellar mass, and vice versa.  We also identified the population of recently-quenched galaxies and studied their abundance as a function of environment and stellar mass.  Lastly, we constructed several simple quenching models to explain some of the (lack of) correlations in the data. The main results of the paper are, \\newline 1. For galaxies with LogM$_{*}$/M$_{\\odot}$ $>$ 9.3, the well-known correlations between f$_{SF}$, SSFR, SFR, D$_{n}$(4000), and color with environment seen in the local universe are already clearly in place by $z \\sim$ 1. \\newline 2. For cluster galaxies, the well-known correlations between properties such as f$_{SF}$, SSFR, SFR, D$_{n}$(4000), and color with stellar mass seen in the local universe are also in place at $z \\sim$ 1. \\newline 3. Similar to the local universe, there is a covariance between the mean stellar mass of galaxies and their environment, with more massive galaxies living in higher-density environments. \\newline 4. The relative decline in the f$_{SF}$ with environment at fixed stellar mass is a factor of $\\sim$ 2-4 and the relative decline in the f$_{SF}$ with stellar mass at fixed environment is a factor of $\\sim$ 2-3.  Both stellar mass and environment affect the f$_{SF}$, and their relative effects are separable up to $z \\sim$ 1. \\newline 5. Controlling for the covariance in stellar mass and f$_{SF}$ with environment we find that the SSFRs of star-forming galaxies are correlated with stellar mass at fixed environment, but is nearly independent of environment at fixed stellar mass. \\newline 6. The D$_{n}$(4000)s of both star-forming and quiescent galaxies are correlated with their stellar mass at fixed environment, but is nearly independent of environment at fixed stellar mass. \\newline 7. The data suggest that stellar mass is the primary parameter that determines the stellar populations of galaxies.  Environment plays a different role, controlling the f$_{SF}$, and to a lesser degree, the mean stellar mass of galaxies. \\newline 8. The population of poststarburst galaxies is correlated with both environment and stellar mass, in that poststarbursts are more common amongst lower-mass galaxies and in higher-density environments.  Poststarburst galaxies with 9.3 $<$ LogM$_{*}$/M$_{\\odot}$ $<$ 10.7 are 3.1 $\\pm$ 1.1 times more common in clusters at $z \\sim$ 1 compared to the field. \\newline 9. The lack of a correlation of the SSFRs and D$_{n}$(4000) of star-forming galaxies and their environment is best-explained by a rapid environmental quenching timescale.  The rapid environmental quenching timescale is further supported by the clear association of poststarbursts with high-density environments. \\newline 10. Despite an overabundance of poststarbursts in clusters compared to the field, poststarbursts still make a relatively small fraction of the total cluster population.  Only 10\\% $\\pm$ 0.2\\% of cluster galaxies at $z \\sim$ 1 with 9.3 $<$ LogM$_{*}$/M$_{\\odot}$ $<$ 11.0 are poststarbursts. \\newline 11. Simple quenching models show that the lack of a correlation between the D$_{n}$(4000) of quiescent galaxies and their environment can be explained two ways: Either self quenching dominates over environmental quenching at $z >$ 1, or else the evolution of the self quenching rate mirrors the evolution of the environmental quenching rate. \\newline\\indent Overall, this study of $z \\sim$ 1 clusters shows that the trends in galaxy properties with environment and stellar mass that are seen in the local universe \\cite[e.g.,][]{Kauffmann2004,Balogh2004a,Baldry2006,Weinmann2006,Peng2010} are already well-established at $z \\sim$ 1.   It appears clear from our study as well as both lower- and higher-redshift studies that while the fraction of star-forming galaxies is a strong function of their environment, \\cite[e.g.,][]{Balogh2004a,Kauffmann2004,Peng2010,Patel2011,Quadri2011}, the mean properties of the star-forming and quiescent galaxies at fixed stellar mass are not.  This means that stellar mass is the primary parameter linked to determining the stellar populations of galaxies. \\newline\\indent Our simple modeling shows that the lack of correlation of the properties of star-forming and quiescent galaxies with their environment can be understood if the environmental-quenching timescale is rapid, and that the evolution of the self-quenching and environmental-quenching rates mirrors each other -- regardless of which processes dominates the overall quenching process.  This hypothesis is further supported by the substantial population of poststarburst galaxies that are only found in the high-density cluster environment. \\newline\\indent The next obvious step to better understanding of the process of environmental quenching will be to connect the evolution of the cluster galaxy population over a larger redshift baseline.  Doing so will allow us to differentiate between different quenching models and get a better handle on the physical processes behind environmental quenching.\\"
    },
    "1112/1112.3643.txt": {
        "abstract": "The quantitative spectral analysis of low resolution ($\\sim 5$ \\AA) Keck LRIS spectra of blue supergiants in the disk of the giant spiral galaxy M81 is used to determine stellar effective temperatures, gravities, metallicities, luminosites, interstellar reddening and a new distance using the Flux-weighted Gravity--Luminosity Relationship (FGLR). Substantial reddening and extinction is found with E(B-V) ranging between 0.13 to 0.38 mag and an average value of 0.26 mag. The distance modulus obtained after individual reddening corrections is 27.7 $\\pm$0.1 mag. The result is discussed with regard to recently measured TRGB and Cepheid distances. The metallicities (based on elements such as iron, titanium, magnesium) are supersolar ($\\approx$ 0.2 dex) in the inner disk (R $\\la$ 5 kpc) and slightly subsolar ($\\approx$ -0.05 dex) in the outer disk (R $\\ga$ 10 kpc) with a shallow metallicity gradient of 0.034 dex \\kpc. The comparison with published oxygen abundances of planetary nebulae and metallicities determined through fits of HST color-magnitude diagrams indicates a late metal enrichment and a flattening of the abundance gradient over the last 5 Gyrs. This might be the result of gas infall from metal rich satellite galaxies. Combining these M81 metallicities with published blue supergiant abundance studies in the Local Group and the Sculptor Group a galaxy mass metallicity-relationship based solely on stellar spectroscopic studies is presented and compared with recent studies of SDSS star forming galaxies. ",
        "introduction": "The determination of the chemical composition and distances of galaxies is crucial for constraining the theory of galaxy formation and evolution in a dark energy and cold dark matter dominated universe. Ultimately, these measurements lead to ever stronger constraints on the cosmological parameters and the history of cosmic chemical enrichment, from the primordial metal-free universe to the present-day chemically diversified structure. For instance, the relationship between central metallicity and galactic mass appears to be a Rosetta stone to understand chemical evolution and galaxy formation \\citep{lequeux79, tremonti04, maiolino08}. In a similar way, the observed metallicity gradients in spiral galaxies, apparently large for spirals of lower mass and shallow for high mass galaxies \\citep{garnett97, skillman98, garnett04}, provide crucial insight into galaxy formation and evolution. Both the observed mass-metallicity relationship and the abundance gradients are used to test the theoretical predictions of hierarchical clustering, galaxy formation, merging, infall, galactic winds and variability of star formation activity and IMF obtained in the framework of a $\\Lambda$CDM dominated universe \\citep{prantzos00, naab06, colavitti08, yin09, sanchez09, delucia04, derossi07, finlator08, brooks07, koeppen07, wiersma09, dave11a, dave11b}. Note that this is only a small selection of papers relevant to the subject, others are found in the references therein.  However, as intriguing the observations of the mass-metallicity relationship and the metallicity gradients of galaxies are, the published results are highly uncertain. They rely on observations of \\hii~region emission lines, mostly restricted to oxygen, and the analysis method applied is the so-called ``strong-line method'', which uses the fluxes of the strongest forbidden lines of (most commonly) [\\ion{O}{2}] and [\\ion{O}{3}] relative to  H$_{\\beta}$. Unfortunately, abundances obtained with the strong-line method depend heavily on the calibration used. As a striking example, \\citet{kewley08} have demonstrated that the quantitative shape of the mass-metallicity relationship of galaxies can change from very steep to almost flat depending on the calibration used. In the same way, as shown by \\citet{kud08} and \\citet{bresolin09} in their study of the Sculptor spiral galaxy NGC 300, metallicity gradients of spiral galaxies can change from steep to flat and absolute values of metallicity can shift by as much as 0.6 dex, again as the result of different calibrations of the strong line method. In consequence, galaxy metallicities are uncertain by 0.6 to 0.8 dex because of the systematic uncertainties inherent in the strong line methods used. This major problem requires a fresh approach and is begging for the development of a new and independent method less affected by systematic uncertainties.  An obvious alternative method to constrain metallicity is the detailed quantitative spectroscopic analysis of individual blue supergiant stars (BSGs) in galaxies. BSGs of spectral type A and B are massive stars in the mass range between 12 to 40 \\msun in the short-lived evolutionary phase ($10^{3}$ to $10^{5}$ years) when they leave the hydrogen main sequence and cross the HR-diagram at constant luminosity and almost constant mass to become red supergiants. Because of Wien's law massive stars increase their brightness in visual light dramatically when evolving towards lower temperatures and reach absolute visual magnitudes up to M$_{V} \\approx$ -9.5 mag in the BSG phase \\citep{bresolin03}, rivaling with the integrated light of globular clusters and dwarf galaxies. Because of their extreme brightness they are ideal tools to accurately determine the chemical composition of young stellar populations in galaxies.  BSG spectra are rich in metal absorption lines from several elements (C, N, O, Mg, Al, S, Si, Ti, Fe, among others). As young objects with ages of 10 Myrs they provide important probes of the current composition of the interstellar medium. Based on detailed high resolution, very high signal-to-noise (S/N) studies of blue supergiants, which yield abundances as accurate as 0.05 dex \\citep{przybilla06, schiller08, przybilla08}, \\citet{kud08} developed an efficient new spectral diagnostic technique for low resolution spectra (FWHM  $\\sim 5$ \\AA)  with good S/N ratio (50 or better), which allows for an accurate determination of effective temperature, gravity, metallicity, interstellar reddening and extinction. Metallicities accurate to 0.1 to 0.2 dex for each individual target can be obtained at this lower resolution and S/N. The method has  been applied to irregular and spiral galaxies in the Local Group (WLM -- \\citealt{bresolin06}; \\citealt{urbaneja08}; NGC 3109 -- \\citealt{evans07}; IC 1613 -- \\citealt{bresolin07}; M33 -- \\citealt{u09}) and beyond (NGC 300 -- \\citealt{kud08}).  In this paper we present the spectral analyis of low resolution Keck LRIS spectra of 26 BSGs in the disk of the giant spiral galaxy M81. M81 is one of the most massive spirals in the Local Volume \\citep{mccommas09}. It has low foreground extinction with a galactic luminosity of 2.5 L$_{*}$ (corresponding to M$_{K}$ = - 24 mag and M$_{K*}$ = -23 mag) and is characteristic of disk galaxies seen at redshift surveys out to z $\\sim$ 1 \\citep{williams09}. The star formation history and chemical evolution of this galaxy have been subject to extensive recent photometric studies \\citep{dalcanton09, davidge09, williams09, barker09, durell10}. \\hii~regions and Planetary Nebulae have been studied by \\citet{stanghellini10} extending the classical work by \\citet{garnett87} and \\citet{stauffer84}. With our work we provide for the first time direct quantitative spectroscopic information about stellar metallicity of the young disk population.  An important additional aspect of the quantitative spectroscopy of BSGs is their use as accurate distance indicators through the Flux-weighted Gravity--Luminosity Relationship (FGLR). This new distance determination method has been introduced by \\citet{kud03} and \\citet{kud08}. %%%Kudritzki et al. 2003 and 2008. It uses stellar gravity and effective temperature as a measure of absolute bolometric magnitude and provides a distance estimate which is free of the uncertainties caused by interstellar reddening, since the determination of reddening is a by-product of the quantitative spectral analysis. First distance determinations using this method have been carried out by \\citet[][ WLM]{urbaneja08} and \\citet[][ M33]{u09}.  There has been a long history of attempts to measure the distance to M81 from Hubble (1929) to the present (see \\citealt{mccommas09} for references and a plot of distance modulus as a function of time). The work published over the last decade gives a range between 27.60 to 28.03 mag in distance modulus indicating an uncertainty of 20\\%. Our BSG spectroscopy and the effective temperatures and gravities determined will give us a FGLR distance which we can then compare with most recent HST work on Cepheids and the tip of the red giant branch, TRGB. M81 has been used as one of the calibration galaxies for the Tully-Fisher and the surface fluctuation methods in the HST Key project \\citep{freedman01} and by \\citet{mould08, mould09}. %%%Mould and Sakai (2008, 2009).  In section 2 of this paper we describe the observations and data reduction. Section 3 discusses the quantitative spectroscopic analysis and the determination of extinction, effective temperature, gravity and metallicity. Section 4 and 5 discuss interstellar reddening and compare the spectroscopically determined stellar parameters with evolutionary tracks in order to constrain the evolutionary status of the objects observed. Section 6 compares metallicity and metallicity gradient of the BSGs with published metallicity constraints for the older disk population of M81 and discusses chemical evolution over the last Gyrs. In section 7 we provide a galaxy mass-metallicity relationship based on BSG spectroscopic studies and compare with published work using \\hii~region emission lines. In section 8 we determine a new distance to M81 using the FGLR-method and discuss recent Cepheid and TRGB work. Section 9 sumarizes the results and discusses aspects of future work. ",
        "conclusions": "In this paper we have demonstrated that the quantitative spectroscopy of BSGs is a promising tool to constrain the chemical evolution of galaxies and to determine their distances, which can be applied to galaxies clearly beyond the Local Group. Using the relationship between flux-weighted gravity and luminosity we were able to determine a new distance to M81, which compares well with TRGB distances. While there is also agreement with HST Cepheid distances within the error margins, our results with regard to metallicity and metallicity gradient confirmed previous studies that the systematic differences between distance moduli obtained from inner and outer field Cepheids (found in M33, M81, M101, NGC 4258) are very likely not caused by a metallicity dependence of the period-luminosity relationship of Cepheids. There must be another reason for these systematic differences.  An independent check of distances obtained with either the TRGB or Cepheids is important for future work. We note that besides the importance for characterizing the physics of galaxies in the Local Volume accurate distances and a careful discussion of the systematics of stellar distance determination methods are crucial for constraining the dark energy equation-of-state parameter w = p/($\\rho$c$^{2}$). As is well known \\citep{macri06}, the determination of cosmological parameters from the cosmic microwave background is affected by degeneracies in parameter space and cannot provide strong constraints on the value of H$_{0}$ \\citep{spergel06, tegmark04}. Only if additional assumptions are made, for instance that the universe is flat, H$_{0}$ can be predicted with high precision (i.e. 2\\%) from the observations of the cosmic microwave background, baryonic acoustic oscillations and type I high redshift supernovae. If these assumptions are relaxed, then much larger uncertainties are introduced \\citep{spergel07, komatsu09}. The uncertainty of the determination of w is related to the uncertainty of H$_{0}$ through $\\Delta$w/w $\\approx$ 2$\\Delta$H$_{0}$/H$_{0}$. Thus, an independent determination of H$_{0}$ with an accuracy of 5\\% will allow the uncertainty of w to be reduced to \u00b10.1. While extremely promising steps towards this goal have been made by \\citet{macri06} and \\citet{riess09a,riess09b,riess11} using the maser galaxy NGC 4258 as a new anchor point and HST IR Cepheid photometry of recent SNIa galaxies out to 30 Mpc, it is clear that the complexity of this approach requires additional and independent tests. Crucial contributions which can be made using BSGs besides independent distance determinations are to investigate the role of metallicity and interstellar extinction.  We have also shown that the determination of metallicities for individual supergiant stars beyond the Local Group is possible. In this way, we can determine galaxy metallicities and metallicity gradients avoiding the systematic uncertainties of \\hii~region strong line methods. This can be used as an independent way to directly measure the mass-metallicity relationship of galaxies and to correlate metallicity gradients with galactic properties such as mass, angular momentum and morphological type. But it can also be used to find out about systematic uncertainties of \\hii~region strong line method calibrations and to identify the more reliable ones or to develop a new one tested with BSG metallicities. Moreover, in combination with metallicity information of an older population of stars obtained through the analysis of CMDs or the spectroscopy of PNe the chemical evolution history of galaxies can be investigated. In the case of the disk of M81 we have found an indication of a late enrichment of heavy elements, which is significantly different from the Milky Way. We have also provided the first mass-metallicity relationship for star forming galaxies solely based on stellar spectroscopy.  %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the"
    },
    "1112/1112.1732.txt": {
        "abstract": "Characterizing the conversion factor between $\\CO$ emission and column density of molecular hydrogen, $X_\\CO$, is crucial in studying the gaseous content of galaxies, its evolution, and relation to star formation.  In most cases the conversion factor is assumed to be close to that of giant molecular clouds (GMCs) in the Milky Way, except possibly for mergers and star-bursting galaxies. However, there are physical grounds to expect that it should also depend on the gas metallicity, surface density, and strength of the interstellar radiation field. The $X_\\CO$ factor may also depend on the scale on which $\\CO$ emission is averaged due to effects of limited resolution. We study the dependence of $X_\\CO$ on gas properties and averaging scale using a model that is based on a combination of results of sub-pc scale magneto-hydrodynamic simulations and on the gas distribution from self-consistent cosmological simulations of galaxy formation. Our model predicts $X_\\CO\\approx 2-4\\times 10^{20}\\rm\\ K^{-1}\\,cm^{-2}\\,km^{-1}\\,s$, consistent with the Galactic value, for interstellar medium conditions typical for the Milky Way. For such conditions the predicted $X_\\CO$ varies by only a factor of two for gas surfaced densities in the range $\\Sigma_{\\rm H_2}\\sim 50-500\\ \\rm M_{\\odot}\\,pc^{-2}$. However, the model also predicts that more generally on the scale of GMCs, $X_\\CO$ is a strong function of metallicity, and depends on the column density and the interstellar UV flux.  We show explicitly that neglecting these dependencies in observational estimates can strongly bias the inferred distribution of $\\H2$ column densities of molecular clouds to have a narrower and offset range compared to the true distribution. We find that when averaged on $\\sim$ kpc scales the X-factor depends only weakly on radiation field and column density, but is still a strong function of metallicity. The predicted metallicity dependence can be approximated as $X_\\CO\\propto Z^{-\\gamma}$ with $\\gamma\\approx 0.5-0.8$. ",
        "introduction": "\\label{sect:intro} %---------------------- Molecular hydrogen ($\\H2$), the major constituent of cold clouds in the interstellar medium (ISM), is playing a major role in shaping the visible Universe around us. Cooling by $\\H2$ aided the formation of the first stars (e.g., \\citealt{2002Sci...295...93A,2002ApJ...564...23B}) and  star formation in nearby galaxies takes place in molecular and giant molecular clouds\\footnote{The distinction between a molecular cloud and a giant molecular cloud is somewhat blurry and several non-equivalent definitions are in use in the literature. We will use the term ``molecular cloud'' to denote a molecular region of $\\gtrsim{}1$ pc diameter within galaxies that is embedded in a non-molecular component of the ISM. A ``giant molecular cloud'' is then simply a large molecular cloud with a diameter of $\\gtrsim{}50$ pc.} (GMCs). Furthermore, the properties of molecular gas in galaxies, e.g., its distribution, formation and destruction are intricately linked to many other phenomena occurring in galaxies, such as the formation of stars, various stellar feedback processes, and the physics of dust grains.  Unfortunately, detecting $\\H2$ in emission is difficult because the lowest excited levels in the rotational ladder are hardly populated at the low temperatures ($\\sim{}10$ K) of molecular clouds. Moreover, the $\\H2$ molecule, lacking a permanent electric dipole moment, radiates via the much slower quadrupole transition (e.g., \\citealt{1982ARA&A..20..163S,2005fost.book.....S}). Therefore, a practical solution has been to measure the emission from some other molecule, that is assumed to trace $\\H2$, and convert its line intensity into an $\\H2$ column density. The conversion factor is called the X-factor for the particular emission line.  Cooling lines from the carbon monoxide ($\\CO$) isotope $^{12}\\CO$ are often used as $\\H2$ tracers, because $^{12}\\CO$ is the second most abundant molecule in molecular clouds and can be observed even in extragalactic objects at high redshifts \\citep{1991AJ....102.1956B, 2005ARA&A..43..677S, 2006ApJ...640..228T, 2008ApJ...680..246T, 2010ApJ...713..686D, 2010ApJ...714L.118D, 2010Natur.463..781T, 2010MNRAS.407.2091G, 2011ApJ...734L..25E, 2011arXiv1106.2098G}. In this study we will focus on the $^{12}\\CO$ $J=1\\rightarrow{}0$ transition (and on the corresponding X-factor $X_\\CO$), i.e., the transition from the first excited rotational level to the ground state, that is widely used in the literature, particularly in studies of nearby galaxies or molecular clouds in the Milky Way (\\citealt{1970ApJ...161L..43W,1987ASSL..134...21S,1987ApJ...319..730S, 1991ARA&A..29..581Y, 1995ApJS...98..219Y, 2001ApJ...561..218R, 2003ApJS..145..259H, 2007PASJ...59..117K, 2007prpl.conf...81B}).  $X_\\CO$ has been measured for molecular clouds in the Milky Way and in a few other nearby galaxies with a variety of techniques. These methods compare the directly measured $\\CO$ luminosity (or intensity) with an independently determined $\\H2$ mass (or column density). The most common approaches: (i) estimates of virial masses from line-width and cloud size (e.g., \\citealt{1987ApJ...319..730S}),  (ii) measurement of $\\H2$ masses using a different tracer molecule (e.g., $^{13}\\CO$), or a higher order transition, with its own X-factor (e.g., \\citealt{1975ApJ...202...50D, 1978ApJS...37..407D, 2009ApJ...699.1092H}), (iii) measurement of proton number density from gamma rays produced in cosmic ray - proton collision (e.g., \\citealt{1986A&A...154...25B,1996A&A...308L..21S, 1997ApJ...481..205H, 2010ApJ...710..133A}), (iv) measurement of total hydrogen column density using infrared emission (e.g., \\citealt{2001ApJ...547..792D, 2007ApJ...663..866D}), (v) estimates of visual extinction, $A_V$, and thus total gas density from IR star counts \\citep{1923AN....219..109W} or the near-infrared color excess \\citep{1994ApJ...429..694L}. The latter three approaches have to be supplemented with maps of atomic hydrogen unless most of the hydrogen is molecular.  The consensus is that for GMCs in the Milky Way all these different methods give similar values of $X_\\CO\\sim{}1.5-4\\times10^{20}$ cm$^{-2}$ K$^{-1}$ km$^{-1}$ s. The question why the X-factor is remarkably constant in molecular clouds in the Milky Way is not yet entirely settled, but it has been suggested to be a consequence of the narrow range of $\\H2$ column densities and temperatures of Milky Way clouds \\citep{2011MNRAS.412.1686S}.  On the other hand, there is increasing evidence that the X-factor may be different in other galaxies (e.g., \\citealt{2008ApJ...686..948B, 2011arXiv1102.4618L}). Specifically, there has been a long debate as to whether the X-factor is \\emph{larger} in low metallicity galaxies (e.g., \\citealt{1995ApJ...448L..97W, 1996PASJ...48..275A, 2002A&A...384...33B, 2005A&A...438..855I}, but cf. \\citealt{2007prpl.conf...81B}), and possibly \\emph{smaller} in the central regions of disk galaxies (e.g., \\citealt{1988ApJ...325..389M, 1998ApJ...493..730O}), in local starbursts and (ultra-)luminous infrared galaxies (e.g., \\citealt{1992A&A...265..447W, 1994ApJ...433L...9S, 1996A&A...305..421M, 1997ApJ...478..144S, 1998ApJ...507..615D, 1999AJ....117.2632B, 2010AJ....140.1294M}), or in submillimeter galaxies at higher redshifts (e.g., \\citealt{2008ApJ...680..246T}). In a simple toy model in which the $\\CO$ emission stems from an ensemble of optically thick, virialized clumps, the X-factor should scale as $n_{H}^{1/2}/T$ \\citep{1986ApJ...309..326D}. However, one arrives at a more complex scaling of the X-factor if the complex geometry of the supersonically turbulent molecular gas and effects of radiative transfer are taken into account (e.g., \\citealt{2011MNRAS.tmp..943S}).  A further important aspect is that many extragalactic surveys do not resolve individual GMCs, but measure $\\CO$ emission on $\\sim{}$ kpc or even larger scale. A study of the properties of the X-factor on such large scales, however, requires galaxy models with reasonably realistic density distributions. The large dynamical scale necessary to resolve individual $\\CO$ emitting gas clumps ($\\sim{}0.1-1$ pc) in a self-consistent, preferentially cosmological, simulation poses a serious challenge.  In this paper, we study the behavior of the X-factor as a function of gas metallicity, UV radiation field, and as a function of scale. We analyze the impact of the X-factor on star formation laws in a follow-up paper (Feldmann et al, in prep). In section \\ref{sect:methods} we present our novel numerical approach that aims at circumventing the problem of the large dynamical scale. The basic idea is that we combine cosmological simulations, in post-processing, with a model that is calibrated on results of small-scale magneto-hydrodynamic (MHD) simulations of the ISM. This approach allows us to study $\\CO$ emission at GMC scales and above with an effective resolution of $~\\sim{}0.1$ pc in a cosmological volume. In section \\ref{sect:XfacScaleGMCs} we discuss the scaling of the X-factor with $\\H2$ column density on GMC scales. Next, in section \\ref{sect:envDepXfac}, we study the dependence of $X_\\CO$ on the local environment (metallicity and UV radiation field). Subsequently, in section \\ref{sect:SurfDensGMC}, we discuss the observational evidence of a constant surface density of molecular clouds. Finally, in section \\ref{sect:spatav}, we address the effect of spatial averaging of the X-factor. We then summarize our results and conclude.   %--------------------- ",
        "conclusions": "A proper physical interpretation of observational data based on $\\CO$ emission requires a thorough understanding and modeling of the conversion factor, the X-factor, between molecular hydrogen and the $J=1\\rightarrow{}0$ emission line of carbon monoxide. We presented a novel approach to study the properties of the X-factor based on a combination of high resolution ($\\sim{}0.1$ pc) MHD simulations of the ISM \\citep{2011MNRAS.412..337G} as a ``subgrid'' model and gas distribution of $\\sim 60$~pc scales derived from self-consistent cosmological hydrodynamical simulations. This approach is feasible, because the spatial resolution of our cosmological simulations ($\\sim{}60$ pc) is reasonably well matched to the box sizes of the ISM simulations ($\\sim{}20$ pc). In fact, this scale, which roughly corresponds to the scale of GMCs, is a natural scale to separate physical processes occurring on galactic and cosmological scales from those relevant for star formation.  The main advantage of our approach is that it takes into account the complicated interplay of chemical evolution, heating, and cooling in a turbulent ISM on small scales and is thus a step forward compared to, e.g., one-dimensional steady state PDR models or to ad hoc models of small scale physics. A caveat is, of course, that it also suffers from any inaccuracies or missing physics in the underlying small scale simulations, but we are hopeful that these problems can be alleviated in the future. Our approach extents previous work that study the X-factor on ISM scales (e.g., \\citealt{2011MNRAS.412.1686S, 2011MNRAS.tmp..943S}) to galactic scales and can be easily used in cosmological simulations, where all cell properties (densities, metallicities, radiation fields) are known, and thus provide self-consistent boundary conditions for the sub-grid modeling.  The main predictions of our model and the main results of this paper are as follows. \\begin{enumerate} \\item The X-factor on GMC scales depends sensitively on metallicity, dust extinction, $\\H2$ column density, while two orders of magnitude variations of the UV radiation field lead only to moderate changes of $X_\\CO$ at a fixed hydrogen column. The changes in $X_\\CO$ are even smaller at a fixed $\\H2$ column density, because a higher UV field tends to reduce the $\\H2$ column density at a given hydrogen column, which offsets to some extent the increase of $X_\\CO$ at fixed hydrogen column.  \\item In galaxies with solar metallicity the X-factor on GMC scales is predicted to be $\\sim{}2-4\\times{}10^{20}$ K$^{-1}$ cm$^{-2}$ km$^{-1}$ s (variations due to dependence of $X_\\CO$ on $N_\\H2$), in agreement with the canonically assumed galactic conversion factor, for a wide range of gas column densities ($\\sim{}50-500$ $M_\\odot$) and with relatively small scatter at fixed $N_\\H2$ ($<0.3$ dex).  Our model also predicts an increase of the X-factor with decreasing metallicity that is in quantitative agreement with direct measurements of $X_\\CO$ from IR dust emission.  \\item The simulations predict an $\\H2$ column density distributions with a peak around $\\sim{}30-100$ $M_\\odot$ pc$^{-2}$ for an ISM with solar metallicity, in agreement with GMCs observations in the Milky Way.  \\item We show that GMC column densities inferred from $^{12}\\CO$ ($J=1\\rightarrow{}0$) observations via a constant galactic conversion factor are biased toward higher column densities and show a strong peak in their distribution, see Fig.~\\ref{fig:NH2GMC}. In addition, such an approach makes the spurious prediction of many lines of sight with low $\\H2$ columns in ISM regions that are exposed to a large UV field. These observational biases are a consequence of relying on a constant galactic conversion factor that ignores the scaling of $X_\\CO$ with $\\H2$ column density.  \\item Spatial averaging, from GMC to $\\sim{}$ kpc scales (and beyond), decreases the dependence of $X_\\CO$ on the $\\H2$ (or total gas) column and the scatter in the X-factor. The column density dependence only remains at very low or large $\\H2$ columns, see. Fig~\\ref{fig:XfacLS}. The UV radiation field plays only a small role on $\\sim{}$kpc scales, and hence, to a good approximation, the metallicity is the primary driver of the X-factor on $\\gtrsim{}$kpc scales.  \\item On kpc scales and above the X-factor scales with metallicity approximately as a power law with an exponent in the range $[-0.5,-0.8]$, depending on assumptions about the $\\CO$ line width and the applied $\\CO$ intensity detection threshold. The power law exponent on GMC scales is similar $[-0.5,-0.9]$.  \\end{enumerate}  Recent $\\CO$ emission studies \\citep{2010ApJ...714L.118D, 2010ApJ...713..686D, 2010MNRAS.407.2091G, 2010Natur.463..781T, 2011ApJ...734L..25E} have improved our knowledge of the relation between star formation and its fuel, molecular gas, out to redshift 3. With instruments of the next generation, e.g., the square kilometer array, it will further be possible to detect $\\CO$ bright gas at even higher redshifts. Going back in time these galaxies presumably differ strongly in their properties from their present-day (or even their $z\\sim{}1-3$) counterparts. Since all these studies observe $\\CO$ we need to understand whether this gives unbiased estimates of the $\\H2$ density in the environments probed by the observations. Hence, a clear understanding and precise modeling of the systematic trends of the X-factor with metallicity and $\\H2$ column density (among other factors) will remain a crucial challenge in order to properly interpret these future observations.\\\\"
    },
    "1112/1112.1410_arXiv.txt": {
        "abstract": "Observations of disk galaxies at $z\\sim 2$ have demonstrated that turbulence driven by gravitational instability can dominate the energetics of the disk. We present a 1D simulation code, which we have made publicly available, that economically evolves these galaxies from $z\\sim 2$ to $z\\sim 0$ on a single CPU in a matter of minutes, tracking column density, metallicity, and velocity dispersions of gaseous and multiple stellar components. We include an H$_2$ regulated star formation law and the effects of stellar heating by transient spiral structure. We use this code to demonstrate a possible explanation for the existence of a thin and thick disk stellar population and the age-velocity dispersion correlation of stars in the solar neighborhood: the high velocity dispersion of gas in disks at $z\\sim 2$ decreases along with the cosmological accretion rate, while at lower redshift, the dynamically colder gas forms the low velocity dispersion stars of the thin disk. ",
        "introduction": "In the past decade, observations of galaxies near $z \\sim 2$ have revealed compelling evidence for the importance of gravitational instability in their dynamics and evolution. A whole class of galaxies has been observed whose images are dominated by large luminous clumps of gas \\citep{Elmegreen:2004gs,Elmegreen:2005km,Forster-Schreiber:2009gd}, while measurements of the velocity dispersion of such massive star-forming galaxies have found values near 50 km/s spread across the entire disk \\citep{Cresci:2009ta,Genzel:2011dd}. This is difficult to reproduce with supernova feedback, which is strongest near the centers of galaxies where the star formation rate peaks, and which is only strong enough to drive velocity dispersions of $\\sim 10$ km/s \\citep{Joung:2009ec}. Other forms of stellar feedback may drive turbulence \\citep{Thompson:2005vy,Elmegreen:2010iw}, but we will concentrate on the case where turbulence is driven by gravitational instability in the disk.  To a first approximation, the gravitational stability of a thin disk to axisymmetric perturbations is described by Toomre's Q parameter $Q=\\kappa \\sigma/(\\pi G \\Sigma)$, where $\\kappa$ is the epicyclic frequency, $\\sigma$ is the 1d velocity dispersion, and $\\Sigma$ is the gas surface density. The disk is unstable when $Q\\lesssim 1$. The importance of gravitational instability in high redshift galaxies arises from the high cosmological accretion rates they experience, which drive up the value of $\\Sigma$ \\citep{Dekel:2009xh}. This instability gives rise to clumps of the sort observed at high redshift. The inhomogenous and time-varying gravitational field drives turbulence throughout the disk, regardless of the stellar density or supernova rate. The energy source for these random motions must ultimately be the gravitational potential of the galaxy, so gas is transported inwards.  Cosmological simulations with sufficiently high resolution \\citep{Bournaud:2009jb, Ceverino:2010mq} successfully reproduce disks in which gravitational instability forms clumps and causes the inward migration of material through galactic disks. Simulations of isolated galaxies \\citep{Bournaud:2011fb,2011MNRAS.417.1318D,2011MNRAS.413.2935D} with initial conditions set such that $Q < 1$ provide a higher resolution view of such galaxies over a few outer orbits of the disk. These studies, while illuminating, are expensive, since they must solve the equations of hydrodynamics in three dimensions over cosmological times. The model we present here solves the hydrodynamical equations in the limit of a thin axisymmetric disk. Since quantities vary only in the radial direction, the problem is computationally much cheaper to solve, allowing us to explore parameter space efficiently, while still solving the full 1D equations of fluid dynamics instead of relying entirely on semi-analytic models \\citep{Dekel:2009xh,Cacciato:2011ab}  Past 1D models of gravitational instability in disks have a number of shortcomings. The rate at which mass and angular momentum are transported inwards is often parameterized and fit to the results of hydrodynamical simulations, rather than being derived from first principles. The rotation curves are only allowed to be either Keplerian or flat. Energy is frequently assumed to be instantaneously equilibrated, which neglects the possibility that it might be advected through the disk. The pressure support of the disk is often treated as coming from thermal pressure rather than supersonic turbulence. Few models take into account the stellar component of the disk, which becomes increasingly important as the galaxy evolves towards the present day, and can ultimately provide a variety of observable predictions.  In particular, the age-velocity dispersion-metallicity correlation of stars in the solar neighborhood \\citep{Nordstrom:2004av}, might well be explained by means of gravitational instability in high redshift disks \\citep{2009ApJ...707L...1B}. The high velocity dispersion in these disks means that the population of stars formed in that epoch will start with a high velocity dispersion \\citep{Burkert:1992th}. The gas disk cools as a result of slowing cosmological accretion rates, so younger stars are formed in a thinner, more metal-rich disk. This mechanism of thick and thin disk formation contrasts with the more common story that various secular processes and minor mergers heat thin disk stars into a thick disk \\citep[e.g.][]{1987gady.book.....B}.  \\citet{Krumholz:2010kc} (hereafter KB10) found an analytic steady-state solution to the full equations of fluid dynamics in the thin disk limit under the assumption that the disk self-regulates to maintain $Q=1$. To make the problem tractable analytically, however, they required a handful of simplifying assumptions: they use an analytic approximation to Q, which becomes progressively worse at lower redshift as the ratio of gas to stellar velocity dispersion deviates from unity. They also assume that the velocity dispersion of stars and gas are equal, and the gas fraction at all points in the disk remains constant in radius and time. In this paper, we relax these assumptions and include treatments of stellar migration, metallicity, the non-zero thermal temperature of the gas, and evolution of individual stellar populations. These improvements along with an efficient simulation code allow us to realistically evolve disks from high redshift to the present day at minimal computational expense.  In section 2 we derive the equations governing the evolution of the gas over time. Section 3 presents the derivation of the equations governing the stellar dynamics. In section 4, we derive the evolution of metallicity in the gas and stars. In section 5 we discuss how these differential equations are solved numerically, and in section 6 we present the results for fiducial parameters chosen to be similar to the Milky Way. We conclude in section 7. The code we describe, named GIDGET for Gravitational Instability-Dominated Galaxy Evolution Tool, is available at \\url{http://www.ucolick.org/~jforbes/gidget.html} ",
        "conclusions": "Starting from conservation laws and simple assumptions about the gravitational stability of the disk, we have derived evolution equations for the radial profile of a two-component disk. Compared to semi-analytic models, this approach has the advantage that the vast variation in the state variables as a function of radius is resolved rather than averaged over the whole disk. This improvement comes with additional computational costs; however, these are not severe - even using the full Rafikov $Q$ and multiple stellar populations, the code can evolve a disk from $z=2$ to $z=0$ on a single processor in a few days, and using the \\citet{Romeo:2011re} approximation to $Q$ reduces the computation time to under an hour.  This paper is primarily meant to introduce our methodology. However, the fiducial model demonstrates a key point which is often overlooked in galaxy evolution and studies of the thick disk, namely that thick disks need not be formed from thin disks. An age-velocity dispersion correlation appears in our simulation, not because of external perturbers, mergers, or \\textcolor{black}{gradual heating of a thin disk}, but because $\\sigma$ decreases with time and newly formed stars induce transient instabilities in the disk \\citep[see also][]{Burkert:1992th}. Both of these processes are strongly dependent on the cosmological situation in which the disk finds itself, that is, its accretion history. Simulations of isolated thin disks which are gradually heated are therefore unrealistic, in the sense that they are missing the most important drivers of thick disk formation. \\textcolor{black}{The smooth increase in stellar velocity dispersion with age produced in our simulations agrees qualitatively with recent observations which demonstrate the lack of a distinctive bimodality between thick and thin disk stars \\citep{Bovy:2011rh}.}  This approach has several further applications which we intend to explore in future work. For Milky Way-like galaxies, even modern chemodynamical models with sophisticated treatments of stellar migration and evolution rely on highly parameterized treatments of gas inflow in the disk \\citep{Schonrich:2009dg,Spitoni:2011rf}. If the gas evolves to keep the disk marginally gravitationally unstable, its movement in the disk is not this simple - it depends on the evolution of the full non-linear set of equations we have derived here. By accounting for the diffusion of stars in radius as the result of scattering across corotation resonances \\citep{Sellwood:2002kk}, our model could be extended to model the Milky Way in detail and compare directly \\textcolor{black}{with} observations of the metallicity gradient as a function of height above the disk \\citep{Cheng:2011fe}, the age-velocity dispersion correlation \\citep{Holmberg:2009gc}, the age-metallicity relation or lack thereof \\citep{Edvardsson:1993am},\\textcolor{black}{ and the radial and vertical stellar density distributions \\citep{Bovy:2011rl}}.   Galaxy bimodality - the separation of galaxies into a blue cloud of star-forming galaxies and a red sequence of ellipticals - is often viewed as an evolutionary sequence. Blue cloud galaxies gradually accrete gas and smaller galaxies, which fuel star formation. Some of these galaxies will undergo major mergers, leaving red and dead elliptical galaxies. These early-type galaxies can subsequently undergo dry mergers, which extend the red sequence to include extremely massive galaxies. Beyond this canonical view, \\citet{SanchezAlmeida:2011cl} have noted the existence of a significant population of red spirals.  By taking more realistic accretion histories from cosmological simulations, we expect that a certain fraction of disks in the course of their lifetimes will experience a period of low accretion during which they will exhaust their gas supply and become redder, only to return to the blue cloud with the resumption of higher accretion rates. \\textcolor{black}{} Given a set of realistic non-smooth \\textcolor{black}{but quiescent} accretion histories\\textcolor{black}{, appropriate for a large fraction of sub-$L_*$ galaxies}, we may therefore be able to reproduce aspects of phenomenology from the local universe out to $z=2$ as semi-analytic models do, but with the added benefit of a physical treatment of the disk dynamics."
    },
    "1112/1112.4974_arXiv.txt": {
        "abstract": "{Globular clusters are old stellar systems which exhibit very-high stellar densities in their cores. The globular cluster Terzan 5 is characterized by a high stellar encounter rate and hosts the largest detected population of millisecond pulsars. It also features bright GeV gamma-ray emission and extended X-ray radiation. However, no globular clusters have been detected in very-high-energy gamma rays (VHE, E$>$ 100 GeV) so far. In order to investigate this possibility Terzan 5 has been observed with the H.E.S.S. telescope array in this energy band. The discovery of a source of VHE gamma rays from the direction of this globular cluster will be reported. The results of the VHE analysis and a multi-wavelength view of Terzan 5 will be presented in this contribution. No counterpart or model can fully explain the observed morphology of the detected VHE gamma-ray source.} ",
        "introduction": "Several types of galactic VHE gamma-ray sources like pulsar wind nebulae (PWNe) and supernova remnants (SNRs) have been detected to date but none so far in the vicinity of a globular cluster. Globular clusters are old stellar systems which exhibit very-high stellar densities in their cores. This leads to numerous stellar encounters \\cite{pooley2006} and to the prolific formation of millisecond pulsars (msPSRs) \\cite{ransom2008}. Globular clusters are predicted VHE gamma-ray sources. In these models the gamma-ray emission is produced by inverse Compton (IC) up-scattering of stellar radiation fields and the cosmic microwave background by relativistic electrons originating either from the msPSRs themselves or their PWNe \\cite{bednarek2007,venter2009,cheng2010}.  Here the discovery of VHE gamma-ray emission from the direction of the Galactic globular cluster Terzan~5 is reported. This globular cluster is located at a distance of 5.9~kpc \\cite{ferraro2009} at RA(J2000)~17$^\\mathrm{h}$48$^\\mathrm{m}$04$^\\mathrm{s}$.85 and Dec~$-24^{\\circ}$46$^\\prime$44$^{\\prime\\prime}$.6 (Galactic coordinates: $l = 3.84^{\\circ}$, $b = 1.69^{\\circ}$) and exhibits a core radius $r_\\mathrm{c} = 0^\\prime.15$, a half-mass radius $r_\\mathrm{h} = 0^\\prime.52$ and a tidal radius $r_\\mathrm{t} = 4^\\prime.6$ \\cite{lanzoni2010}. Terzan~5 hosts the largest population of detected msPSRs (33) \\cite{ransom2008} and it has been discovered by \\emph{Fermi-LAT} in the GeV range \\cite{kong2010,abdo2010}.  So far only upper limits have been reported from globular clusters in the VHE gamma-ray range (e.g. \\cite{aharonian2009,anderhub2009,mccutcheon2009}). ",
        "conclusions": "Available multi-wavelength data on the one hand do not show any typical VHE gamma-ray emitter (like PWN or SNR) in the vicinity of the detected source. On the other hand the properties of the source, namely the extension and the off-set ($2 \\sigma$ level) are unexpected for a globular cluster. Therefore the results of the H.E.S.S. observations are difficult to interpret. Here scenarios excluding and including the globular cluster as the origin of the VHE gamma-ray emission are briefly discussed.  \\subsection{Chance coincidence}  The positional concurrence between the globular cluster and the VHE gamma-ray source could in principle be just a chance coincidence of physically unrelated objects. Notably the source parameters (extension, photon spectrum) of HESS~J$1747-248$ are compatible with the properties of VHE gamma ray detected PWNe \\cite{mattana2009}.  The probability of chance coincidence can be estimated from the distribution of VHE gamma-ray sources in the galactic disk. From the H.E.S.S. galactic plane scan the lateral distribution of sources in the longitude range from $-85^\\circ - +60^\\circ$ has been obtained. It can be described by a Gaussian profile (containing 48 sources) centered at b=$-0.26^\\circ$ with a width of $0.4^\\circ$ and additional four outliers with b $<-2^\\circ$ below the galactic disk \\cite{chaves2009}. At a latitude of $1.7^\\circ$ Terzan~5 is almost 5 $\\sigma$ away from the center of the Gaussian distribution and there are no other VHE sources detected in the latitude band of $1.5^\\circ - 2^\\circ$. Hence HESS~J$1747-248$ presents an outlier to the lateral source distribution. If it is assumed that one outlier is located in the latitude band of $1.5^\\circ - 2^\\circ$ then the chance coincidence that it is placed within $0.1^\\circ$ of the center of the globular cluster is about 10$^{-4}$. Thus, its quite unlikely that the proximity of Terzan~5 and HESS~J$1747-248$ is due to chance but this possibility cannot be excluded.   \\subsection{Leptonic scenario}  Globular clusters are predicted VHE gamma-ray emitters \\cite{bednarek2007,venter2009,cheng2010}. All these models rely on IC up-scattering of low energy photons (stellar radiation, cosmic microwave background) by energetic electrons. The electrons are proposed to be accelerated either by the population of msPSRs themselves or by their colliding PWNe. Models \\cite{bednarek2007,venter2009} predict a VHE gamma-ray flux for Terzan~5 of about 1\\% of the flux of the Crab nebula for reasonable input parameters. This is in broad agreement with the flux of HESS~J$1747-248$.  For an IC scenario it is expected that the gamma-ray emission should follow the shape of the up-scattered radiation field which is highly centrally peaked in the case of a globular cluster. Therefore, an extended source with off-set from the globular cluster center is challenging to interpret for such a scenario. IC emission in the VHE gamma-ray regime should be accompanied by synchrotron emission in the X-ray regime. However, since the detected diffuse X-ray emission \\cite{eger2010} is also centered on the globular cluster core a simple model where the X-ray and VHE gamma-rays emission originates from the same population of electrons cannot explain the morphology of the VHE source.  To summarize, the population of msPSRs would energetically be capable of producing the VHE gamma-ray emission but the source morphology is not self-evident for such a scenario.  \\subsection{Hadronic scenario}  Globular clusters are believed to boost the rate of stellar mergers due to the extreme stellar densities in their cores (e.g. \\cite{shara2002,grindlay2006}). These collisions lead in some cases to stellar explosions like SN Ia powered by white~dwarf~-~white~dwarf mergers. Remnants of these SN Ia could be the acceleration site of hadronic cosmic rays (see \\cite{acero2010} for the detection of SN 1006 in VHE gamma-rays).  To explain the VHE emission an energy in hadronic cosmic rays of 10$^{51} (n/0.1\\mathrm{cm}^{-3})$~erg ($n$ density of ambient target material) would be needed if it is assumed that the cosmic ray spectrum follows a power-law with index -2 below the region which can be probed with H.E.S.S. This is somewhat high for a supernova. The source morphology could be explained by a hadronic scenario but there is no multi-wavelength support for a SNR origin \\cite{clapson2011} of the VHE gamma-ray emission."
    },
    "1112/1112.1884_arXiv.txt": {
        "abstract": "As shown in previous works, globular clusters can be used to trace the overall structure of the diffuse stellar populations in early type galaxies if the number of clusters per unit stellar mass depends on metallicity. In this paper we further test this assumption in the galaxy NGC 4486 (M 87), by combining several data sources. The results show that globular clusters allow the mapping of the galaxy in terms of the surface brightness profile, integrated colour gradient, chemical abundance, and  mass to luminosity ratios up to 1000 arcsec (or 80.4 kpc) from its centre (i.e. some 10 effective radii). The analysis indicates the presence of a dominant high metallicity bulge associated with the red globulars, whose ellipticity increases outwards, and of a more flattened low metallicity halo connected with the blue globulars. The chemical abundance gradient of the composite stellar population is remarkably similar to that inferred from X ray observations of hot gas. The mass-metallicity spectrum of the stellar population can, in principle, be understood in terms of inhomogeneous enrichment models. In turn, the distribution of the bluest GCs, and lowest metallicity halo stars, has an intriguing similarity with that of dark matter, a feature  shared with NGC 1399. Also, in these two galaxies, the number of blue GCs per dark mass unit is identical within the errors, $\\approx $ $1.0(\\pm 0.3) \\times 10^{-9 }$. The total stellar mass derived for NGC 4486 is $6.8(\\pm 1.1)\\times10^{11} M_\\odot$ with a baryonic mass fraction $f_{b}=$ 0.08($\\pm$ 0.01). ",
        "introduction": "\\label{INTRO}  The potential of globular clusters (GCs) as tracers of early events in the  life of galaxies has been emphasized along the years. The idea has roots in \\citet{EGG62} and \\citet{SEA78} and a landscape of the different issues within that context is given, for example, in \\citet{BRO06}. The theoretical aspects of the problem have been addressed in an effort to define its cosmological background (e.g. \\citealt{BEA02}; \\citealt{BEK08}; \\citealt{MURG10}).  However, as noted by \\citet{KPA09}, and also see \\citet{SPIT10}, some important aspects remain as open questions. Among them, the striking differences observed between the metallicity distribution function (MDF) of the CGs, usually bimodal, and that of field stars as observed in resolved galaxies. In this last case the MDFs are predominantly broad, asymmetric, and  exhibit a low  metallicity tail (e.g. \\citealt{DUR01}; \\citealt{HAR02}; \\citealt{REJ05}; \\citealt{REJ11}).  Differences and similarities between GCs and the underlying stellar populations have been pointed out in the past. For example, \\citet{FOR81} and \\citet{STR81} found that GCs appear significantly bluer than the stellar haloes at all galactocentric distances in four Virgo ellipticals.  These comparisons should take into account, however, that the CGs statistics are ``number weighted'' while halo colours are, naturally, ``luminosity weighted''. Thus, a coincidence between GC and stellar halo colours would be expected only if each GC formed along a given diffuse stellar mass on a constant luminosity basis. This fact is frequently ignored in the literature.  A possible quantitative connection between the features of  GCs and field stars, assumes that clusters are born during major star forming episodes and that, the number of GCs per diffuse stellar mass unit, increases exponentially when the logarithmic chemical abundance $[Z/H]$ decreases.  This connection is consistent with the variation of the GCs specific frequency $S_{n}$ with chemical abundance (\\citealt{HAR02}), and allows the recovery of both the luminosity profile and colour gradient in a galaxy as the result of a multi-metallicity stellar population.  A first approach along these lines was presented in  \\citet{FOR07} (hereafter FFG07). More consequences of that idea were discussed in \\citet{FOR09} (hereafter FVF09) who performed an analysis of 69 galaxies within the Virgo ACS (see \\citealt{FER04}; \\citealt{PEN08}).  FVF09 also find that globular cluster systems (GCS) in Virgo galaxies with stellar masses below $M_*\\approx 10^{11} M_\\odot$, lie on plane defined (in a 3-D logarithmic space) by GC formation efficiency (number of GCs per  unit stellar mass), total stellar mass and projected stellar mass density. NGC 4486, and other bright Virgo ellipticals, lie well outside this plane, probably reflecting the effects of past merger events.  In this paper we revise the case of NGC 4486 (M 87), taking advantage of a number of photometric works (\\citealt{TAM06}; \\citealt{HAR09}; \\citealt{JOR09}) that, combined, allow an analysis on an angular scale of $\\approx$ 1000 arcsec.  An important assumption of the analysis is that GCs and field stars are mostly coeval or formed  within a narrow period of time. Arguments in support of this hypothesis can be found, for example, in \\citet{SPIT10} and references therein. In particular, a direct estimate of the age of the halo of the galaxy NGC 3923, as well as of its associated GCs, was presented by \\citet{NOR08} using Gemini MOS spectra and Lick indices. These results also indicate and old age (about 12 Gyr) for both stars and GCs (see \\citealt{BEA08} for the case of NGC 5128).  More arguments in favour of old ages and multi-metallicity spectra for the dominant stellar populations in elliptical galaxies, stand on the analysis of their integrated colours as discussed in \\citet{SCH09}.  Several papers, on the other side, deal with the surface brightness profile of NGC 4486. In this work we adopt the blue band (Johnson system) surface brightness profile presented  by \\citet{CAO90}. All these works denote the existence of a flattened cD halo (already detected by \\citealt{ARP71}), traced to very low luminosity levels in \\citet{MIH05}. More  recently, \\citet{RUD10} find that the $(B-V)$ colour gradient of the galaxy has a logarithmic dependence along the semi major axis over an angular range of (at least) $1000$ arcsec (i.e., approximately, ten effective radii). This corresponds to 80.1 kpc for a distance modulus $(Vo-M_v)=31.1$, or 16.5 Mpc, (\\citealt{TON01}), adopted in this work .  The paper is organized as follows: Section 2 presents an improved version of the GC integrated colours vs. metallicity relation given in FFG07 and the connection between chemical abundance Z (in solar units hereafter) and the cluster integrated colours, including an attempt to model the so called ``blue tilt''. Section 3 explains the assumed link  between GCs and field stars. Section 4 describes the fit of the GC colour histograms within the inner $300$ arcsec  along the galaxy semi major axis, aiming at determining the number of clusters belonging to the blue or red sub-population as well as their respective chemical scale-lengths. A discussion of the large angular scale features of the galaxy, including the shape of the surface brightness profile, colour gradient, stellar metallicity function and stellar mass to luminosity ratios is given in Section 5. The inferred halo, bulge and total stellar masses, and a comparison with the dark matter content is presented in Section 6. Finally, the GCs formation efficiencies for both the blue and red GCs and the conclusions of this paper, are included in Section 7 and 8, respectively. ",
        "conclusions": "\\label{CONCLU}  The distribution of the integrated GC colours in NGC 4486 is compatible with the presence of two different (blue and red) populations both following an exponential decay (in number) with increasing chemical abundance Z. The chemical scale lengths, however, are very different, being much smaller for the blue GC family. Both GC chemical scales have negative gradients with increasing galactocentric radius. In turn, the composite mass weighted chemical gradient obtained for the composite stellar population is comparable to that derived through X ray observations of hot gas, indicating a strong connection between stars and the interstellar medium.  The main idea behind the GCs-field stars connection adopted in this work is that the metallicity of a given cluster is a proxy to the mass of the associated diffuse stellar population where it is born and which, through pre-enrichment (see \\citealt{BAI09}), determines the cluster metallicity itself. Furthermore, metallicity may be indicative of the mean stellar density of the environment, a factor possibly governing the GCs formation efficiency.  The approach allows, by combining both the halo and bulge components, the reconstruction of the galaxy profile brightness and of the negative integrated colour gradient over an angular scale of, at least, $1000$ arcsec ($\\approx$ 80 kpc). Given the multi-metallicity nature of the stellar population, SSP models are not adequate for a proper understanding of integrated colours in terms of chemical abundance (for a detailed discussion on this subject see \\citealt{SCH09}).  The chemical abundance distribution inferred in this work for the stellar population is similar to those observed in resolved galaxies and exhibits similarities with the predictions given by heterogeneous enrichment models as discussed by \\citet{OEY00}.  In particular, the halo/bulge systems can be identified as formed in low/high density environments typified by small/large number of stellar generations and different filling factors. A comparison of our Fig. 15 and 16 with figs. 2 and 3 in Oey's work, appeals  for a further, more quantitative exploration.  If GCs appear at the very beginning of the star forming process their own chemical distribution may be reflecting the original ``metal production function'' and this would also imply larger mean ages than that of the field stars. In agreement with this, \\citet{SPIT10} suggests that metal rich GCs should be older by about $1$ Gyr compared with the average age of the bulge.  The GC formation efficiency for each (halo, bulge) system shows that the NGC 4486 system clearly deviates from those corresponding to other Virgo galaxies with masses smaller that $10^{11}  M_\\odot$.  Although both the NGC 4486 stellar halo mass, and the large number of blue globulars, could be reached by accreting low mass galaxies, the presence of a chemical abundance gradient rather suggests that accretion might have coexisted with the in-situ formation of indigenous blue GCs (see \\citealt{HAR09} or \\citealt{FORB11} for a similar conclusion). A description of the  mechanisms that may play a simultaneous role in the formation of galaxy haloes is given, for example, in the model analysis by \\citet{MUR07}.  We also stress that the eventual accretion of an external population of blue clusters, through numerous small mergers, would not affect the detectability of the chemical or colour gradient of the indigenous blue GCs  or of the blue tilt (as shown in Section 7). Alternatively, but not in necessary disagreement, \\citet{KISS98} argue that, their detection of rotation for the blue GCs, is possibly the fingerprint of an important single past merger.  On the other side, pure dry merging cannot explain the red GC-bulge population since the chemical scale parameter of the plausible progenitor galaxies is much smaller than that in the central region of NGC 4486. Both dissipative collapse and/or wet merging leading to high star forming efficiencies and chemical enrichment seem required. The low central surface brightness, as well as the shallow chemical gradient in NGC 4486 are compatible with merging processes (e.g. \\citealt{KIM10}). This is also suggested by the position of the galaxy in the S\\'ersic $n$ index-vs-central brightness relation ``normal'' ellipticals (see \\citealt{GRA03} for a comprehensive discussion). A scenario were GC formation is in fact entirely merger-driven is given in \\citet{MURG10}.  The spatial distribution of the halo stars with the lowest chemical abundance is remarkably similar to that inferred for dark matter. This coincidence was already noticed in the case of NGC 1399 (\\citealt{FFG05}). This last galaxy and NGC 4486 seem to have a very similar baryonic mass (see \\citealt{JON97}) and also a comparable number of red GCs. However, they differ both in the number of blue GCs and dark mass content by the same factor, which is $\\approx 3$ larger in NGC 4486 . Besides, a re discussion of the photometric data presented in FFG07 and of the dark matter content presented by Schuberth et al. (2010) for NGC 1399, also indicates a remarkably similar number of blue GCs per dark mass unit, $\\approx$ $1.0(\\pm 0.3)\\times10^{-9}$ in both galaxies.  All these features add to the discussion of the possible origins of GCs, and in particular of the blue clusters, in connection with primordial dark matter sub-haloes (e.g. \\citealt{SAN05}; \\citealt{MAS05}; \\citealt{BEK08}; \\citealt{BOL09}).  Further arguments about the connection between GCs and dark matter come from the very large scale study by \\citet{LEE10}, who report the detection of ``intracluster'' GCs whose projected spatial distribution resembles those of dwarf galaxies and dark matter on very large angular scales. The change in the slope of the GC areal density profile in their work, at $\\approx$ $40$ arcmin (193 kpc) from the centre of NGC 4486, is close to the ``edge'' of the galaxy halo as  determined by \\citet{DOH09} (but see \\citealt{STR11} for a different conclusion).  The analysis presented in this work strongly support the idea that GCs appear along the main star forming events in early type galaxies and can be used as tracers of the dominant stellar populations. The ``two phases'' process, as suggested by \\citet{FORB97}, seems an adequate scenario to explain the overall features of GCs and associated stars. The meaning of ``phase'', however, may require further clarification in terms of the existence of a temporal sequence (i.e. halo globulars appear prior to to bulge globulars as suggested by the very different chemical abundances ?), environmental differences (underlying in the inhomogeneous enrichment models) and, eventually, about how both factors combine in the overall landscape of galaxy formation."
    },
    "1112/1112.5375_arXiv.txt": {
        "abstract": "Bibliometric studies have become increasingly important in evaluating individual scientists, specific facilities, and entire observatories. In this context, the ESO Library has developed and maintains two tools: FUSE, a full-text search tool, and the Telescope Bibliography (telbib), a content management system that is used to classify and annotate ESO-related scientific papers.  The new public telbib interface provides faceted searches and filtering, autosuggest support for author, bibcode and program ID searches, hit highlighting as well as recommendations for other papers of possible interest. It is available at\\\\ \\url{http://telbib.eso.org}. ",
        "introduction": "Telescope bibliographies are important tools to measure scientific output. Typically, they contain all (or all refereed) papers using observational data from specific facilities that were published in the scholarly literature.  Telescope bibliography databases are therefore the ideal source to derive various kinds of reports and statistics. For instance, management and governing bodies of observatories may be interested in publication and citation statistics to evaluate the performance of observatories and telescopes. Instrument scientists often need reports regarding the scientific impact of specific instruments and research programs. Such reports can also provide guidelines for future telescopes and instruments. For the astronomy community at large, it is important that telescope bibliographies interconnect resources: publications are linked to the observing programs that generated the data and to the actual data in the archive, and in turn scientists will be able to go from the archival data directly to all publications that use these data.  In this context, the ESO Library has developed and maintains two tools: (1) FUSE is a full-text search tool that semi-automatically scans defined sets of journal articles for organizational keyword sets, while providing highlighted results in context; (2) the Telescope Bibliography (telbib) is used to classify ESO-related papers, store additional metadata, and generate statistics and reports. Both tools rely heavily on the NASA ADS Abstract Service for bibliographic metadata.  In this paper, we describe how FUSE and telbib link publications and observations and explain the main features of the new public telbib interface. ",
        "conclusions": ""
    },
    "1112/1112.2249_arXiv.txt": {
        "abstract": "We present a systemic analysis of all of the stellar mass black hole binaries (confirmed \\& candidate) observed by the \\textit{Swift} observatory up to June 2010. The broad \\textit{Swift} bandpass enables a trace of disk evolution over an unprecedented range in flux and temperature. The final data sample consists of 476 X-ray spectra containing greater than 100 counts, in the 0.6 -- 10 keV band. This is the largest sample of high quality CCD spectra of accreting black holes published to date. In addition, strictly simultaneous data at optical/UV wavelengths are available for 255 (54\\%) of these observations. The data are modelled with a combination of an accretion disk and a hard spectral component. For the hard component we consider both a simple power-law and a thermal Comptonization model. An accretion disk is detected at greater than the 5$\\sigma$ confidence level in 61\\% of the observations.  Lightcurves and color-color diagrams are constructed for each system. Hardness luminosity and disk fraction luminosity diagrams are constructed and are observed to be consistent with those typically observed by \\textit{RXTE}, noting the sensitivity below 2 keV provided by \\textit{Swift}. The observed spectra have an average luminosity of $\\sim$ 1\\% Eddington, though we are sensitive to accretion disks down to a luminosity of $\\rm 10^{-3}~L_{Edd}$. Thus, this is also the largest sample of such cool accretion disks studied to date. The accretion disk temperature distribution displays two peaks consistent with the classical hard and soft spectral states, with a smaller number of disks distributed between these.  The distribution of inner disk radii is observed to be continuous regardless of which model is used to fit the hard continua.  There is no evidence for large scale truncation of the accretion disk in the hard state ($\\rm at~least~for~L_x \\gtrsim 10^{-3}~L_{Edd}$), with all of the accretion disks having radii $\\rm \\lesssim 40~R_g$. Plots of the accretion disk inner radius versus hardness ratio reveal the disk radius to be decreasing as the spectrum hardens, i.e., enters the hard state. This is in contrast to expectations from the standard disk truncation paradigm and points towards a contribution from spectral hardening. The availability of simultaneous X-ray and optical/UV data for a subset of observations, facilitates a critical examination of the role of disk irradiation via a modified disk model with a variable emissivity profile (i.e., $\\rm T(r) \\propto r^{-p}$). The broadband spectra (X-ray -- Opt/UV) reveal irradiation of the accretion disk to be an important effect at all luminosities sampled herein, i.e., $\\rm p \\lesssim 0.75~for~luminosities \\gtrsim 10^{-3}~L_{Edd}$. The accretion disk is found to dominate the UV emission irrespective of the assumed hard spectral component. Overall, we find the broadband soft state spectra to be consistent with an irradiated accretion disk plus a corona, but we are unable to make conclusive statements regarding the nature of the hard state accretion flow (e.g., ADAF/corona vs jet). Finally, the \\textit{Swift} data reveal a relation between the flux emitted by the accretion disk and that emitted by the corona, for this sample of stellar mass black holes, that is found to be in broad agreement with the observed disk -- corona relationship in Seyfert galaxies, suggesting a scale invariant coupling between the accretion disk and the corona. ",
        "introduction": "Accretion onto a black hole is a fundamental astrophysical process capable of producing copious amounts of energy from radio -- $\\gamma$-ray frequencies. This radiation in turn will have a profound impact on its surroundings. In the local Universe this is illustrated by the discovery of large jet blown cavities around stellar mass black holes \\citep{gallo04,pakull10}. When integrated over the lifetime of the black hole, this energy will have a significant impact on its environment, e.g., feedback from supermassive black holes (SMBH) in the quiescent state ($\\rm \\lesssim 10^{-6}~L_{Edd}$)\\footnote{$\\rm L_{Edd} = 1.3 \\times 10^{39}~(M_x/10~M_{\\sun})~erg~s^{-1}$} is required to reproduce the observed properties of galaxies \\citep{croton06}.  The observation of tight correlations between the SMBH and properties of the host galaxy are further evidence for the importance of feedback from the black hole, and understanding this will provide insight into the process of Galaxy formation and evolution \\citep{ferraresemerritt00,gebhardt00,gultekin09}. At higher redshifts still, the accretion power of early HMXBs could have an impact on the formation of the earliest stars and galaxies \\citep{mirabel11}.  There are over 500 accreting binaries currently known in the Milky Way, though this sample is dominated by systems containing neutron star accretors. Observations of other galaxies have also revealed large populations of accreting objects \\citep{fabbiano06}, which can be used to reveal information on the star formation history \\citep{grimm03} and stellar content of these galaxies \\citep{gilfanov04}.  The X-ray luminosity function of the Milky Way can be used to constrain the population and evolution of X-ray emitting sources in the Galaxy \\citep{grimm02}, which is found to be dominated by accreting binary systems with compact object accretors (BH, NS \\& WD, \\citealt{revnivtsev11a,revnivtsev11b}).  Detailed study of the accretion flow in Galactic X-ray binaries can provide constraints on the nature of accretion in the presence of a large gravitational potential, and the theory of general relativity itself \\citep{psaltis08}. In particular, they provide a nearby window on black hole spin \\citep{miller07,mcclintock11}, the production mechanism of the ubiquitous relativistic jet emission \\citep{markoff01,markoff05,fender06}, and even the nature of supernovae (SNe) and gamma-ray bursts (GRBs; \\citealt{miller11}).  To date, there are 20 confirmed stellar mass black holes (i.e., with dynamical constraints on the mass of the BH) and greater than 20 strong candidate systems, e.g., see \\citet{mcclintockremillard06}.  BH X-ray binaries (XRBs) are typically transient systems that are observed to brighten from $\\rm \\lesssim 10^{-6}~L_{Edd}~- L_{Edd}$ on timescales of weeks to months before returning to a low luminosity state.  The Galactic black hole binaries have been classified based on their behavior in the soft X-ray band (0.1 -- 10 keV). This resulted in the definition of a number of apparently discrete spectral states \\citep{tanakashibazaki96,mcclintockremillard06}. When a black hole is accreting at an appreciable percentage of its Eddington luminosity ($\\rm L_x \\gtrsim 10^{-3}~L_{Edd}$), 2 primary spectral forms are observed.  The first of these is the low-hard state, characterized by a power-law continuum in the 2 -- 10 keV band with a spectral index $\\rm \\Gamma \\lesssim 1.8$, and a spectral cut-off at energies $\\sim$ 100 keV. The RMS spectral variability is observed to be high in this state. The second state is the so-called high-soft state. The X-ray spectrum is dominated by emission in the soft band ($\\rm \\lesssim 2~keV$) consistent with a thermal blackbody component, a softer low luminosity power-law component at higher energies ($\\Gamma \\gtrsim 2.2$), and low RMS noise \\citep{tanakashibazaki96,grove98}. The launch of the \\textit{RXTE} mission \\citep{bradt93} resulted in a refinement in our understanding of these spectral states, and the transitions between them, which have resulted in the definition of a number of intermediate spectral states based on both the spectral state and the timing properties of the source \\citep{homanbelloni05,mcclintockremillard06,vanderklis06}.  The broadband emission from accreting black holes is correlated with the behavior at X-ray wavelengths, as clearly illustrated by the observed phenomenology of the radio jet \\citep{fender06,gallo10}\\footnote{Jets in X-ray binaries are typically observed in 2 distinct forms: (i) transient ejections occurring close to a state transition, and (ii) steady state jets observed in the low-hard state, see \\citet{fender06,gallo10} for details. Hereafter, when we discuss `jets', we are referring to the steady state form unless specifically stated otherwise.}. Studies of jets from accreting black holes on all mass scales have led to the realization that the accretion disk (inflow) and the jet (outflow) are closely related, i.e., `fundamental plane of black hole activity' \\citep{merloni03,falcke04} and `disk-jet coupling' \\citep{fender04}. In particular it has been noted that (i) the low-hard state ($\\rm L_x \\lesssim 0.05~L_{Edd}$ ) is typically associated with the presence of a compact quasi-steady jet (ii) radio/jet emission appears to cease when the system enters the high-soft state ($\\rm L_x \\gtrsim 0.1~L_{Edd}$). This suggests that the accretion disk inflow is intimately involved in the process of launching a jet, e.g., \\citet{fender06,gallo10}.  At higher energies ($\\rm > 10~keV$), the large archive of high energy observations by \\textit{RXTE} \\citep{bradt93}, observations by \\textit{Integral} through the bulge monitoring program \\citep{kuulkers07}, and pointed observations are helping to elucidate the nature of the high energy X-ray component. \\textit{Suzaku} observations have demonstrated the importance of having broadband spectral coverage, where the broad bandpass (0.6 -- 300 keV) has proven to be crucial to simultaneously constrain the soft accretion disk component, the hard X-ray emission and the reflection features, i.e., Fe K line, Compton hump e.g., see \\citet{makishima08,reynolds10a,reynolds10b} for high energy spectra and \\citet{tomsick09,reis11a,reis11b} for constraints on the reflection component.  \\begin{figure}[t] \\begin{center} \\includegraphics[height=0.27\\textheight]{./fig1.eps} \\caption{Histogram illustrating the observational characteristics of the data sample. The filled histogram plots the exposure time for all of the systems, where they have been grouped in 0.5 ks bins (see x1 axis). The exposure peaks at approximately 1 ks, with $\\sim$ 70\\% of the observations having an exposure time less than 2 ks and 90\\% less than 4 ks. The total exposure time is 940 ks. The hatched histogram shows the total number of counts for the spectra in our sample (see x2 axis). The peak is at $\\sim$ 60k counts, while 70\\% and 90\\% of the spectra had greater than 3.5k and 500 total counts respectively. The number of counts in the background relative to the source is negligible in all cases.} \\label{histogram_expos} \\end{center} \\end{figure}  \\begin{figure*}[t] \\begin{center} \\includegraphics[height=0.40\\textheight]{fig2.eps} \\caption{Absorbed flux versus time for all of the black holes observed by \\textit{Swift} up to May 2010, see the symbols indicated on the right of the plot and Table \\ref{BHC_table}. All fluxes are measured in the 0.6 -- 10 keV band via the \\texttt{flux} command in \\textsc{xspec}. The outbursts of these black hole binaries are observed to exhibit a dynamic range of $\\rm 10^3 - 10^5$ over the 0.6 -- 10 keV energy range.}% \\label{BH_all_systems} \\end{center} \\end{figure*}  \\subsection{Specific issues}\\label{spec_issues} Understanding the relationship between the accretion disk inflow and the jet outflow is one of the outstanding goals of modern astrophysics. The anti-correlation between the accretion state of the black hole and the presence or absence of a relativistic jet indicates that understanding the geometry of the accretion flow in each of the primary accretion states will be key, i.e., the low-hard state and the high-soft state. The current paradigm for the configuration of these accretion states is consistent with a hot geometrically thin optically thick multi-temperature blackbody accretion disk \\citep{ss73} extending to the innermost stable circular orbit (ISCO) in the soft state. As the accretion rate and hence disk temperature drop, the inner edge of the accretion disk recedes from the ISCO and the inner region is filled with a geometrically thick ADAF (advection dominated accretion flow, \\citealt{I17}), which produces the observed hard spectral component that dominates the low-hard state X-ray spectrum \\citep{mcclintockremillard06,done07}.  However, a number of recent observations contradict this picture and reveal evidence for the persistent of the accretion disk at the ISCO in the low-hard state, e.g., \\citet{miller06a,miller06b,rykoff07,tomsick08,wilkinson09,reis09,reynolds10a,reynolds10b,reis10}. These observations suggest that the accretion disk is not the primary driver of the observed low-hard $\\Leftrightarrow$ high-soft state transition. Indeed, the latest observations suggest that the disk does not recede from the ISCO until luminosities of $\\rm \\lesssim 10^{-3}~L_{Edd}$ \\citep{tomsick09}. A number of alternative scenarios are consistent with the accretion disk remaining at the ISCO, e.g., the hard X-ray flux may originate in the relativistic jet \\citep{markoff05} or from a hot corona \\citep{belo99,diskpn_gierlinski99,merloni02}, and it is the turning on/off of the jet or evolution of the corona that drives the observed state transitions.  Questions also persist at longer wavelengths. Recent observations have highlighted a problem in our understanding of the UV emission mechanism in X-ray binaries. Standard theory explains the UV emission as direct emission from the accretion disk and observations of the classical $\\rm \\nu^{1/3}$ spectrum in early IUE observations supported this picture \\citep{cheng92,schrader93}. It was soon realized that the large flux of X-rays present in an X-ray binary should modify the standard disk spectrum, with the X-ray emission irradiating the disk, where it is subsequently reprocessed -- disk reprocessing (e.g. \\citealt{vanparadijs94,kingritter98,hynes98,dubus01,russell06}).  A standard steady state accretion disk will be observed to exhibit a radial temperature dependence of the form $\\rm T(r) \\propto r^{-0.75}$ \\citep{ss73}.  It is inherently difficult to constrain the temperature profile of the accretion disk at X-ray wavelengths as the primary modification to the accretion disk spectrum occurs at wavelengths to the red of the typical X-ray detector, i.e., UV/optical or $\\rm E \\lesssim 0.1~keV$. This is illustrated in \\citet{kubota05} who studied a sample of \\textit{ASCA} observations of black holes in outburst and found the temperature profile of the measured spectra to be consistent with the standard $\\rm T \\propto r^{-0.75}$ accretion disk at all times, though there was evidence, of low statistical significance, for a deviation towards a more irradiated disk in some of the observations. In contrast, observations at optical and UV wavelengths have revealed clear signatures for disk irradiation \\citep{schrader94,vanparadijs94,russell06,rykoff07}. Indeed, the addition of X-ray irradiation of the accretion disk to the disk instability model is required in order to reproduce the observed outburst lightcurves of black holes XRBs \\citep{dubus01,lasota01}.  Over the last decade the picture has been further complicated by the realization that the radio jet may contribute significantly across the entire spectral energy distribution, e.g., \\citet{markoff01,russell06}. Furthermore, simultaneous observations at optical and X-ray wavelengths have revealed complex correlation functions, hinting at the presence of a non-thermal component in addition to the accretion disk emission, e.g., \\citet{kanbach01,durant08,gandhi08}.  \\citet{dunn10a,dunn10b} has recently carried out a comprehensive study of black hole X-ray binaries using over 13 years of archival \\textit{RXTE} observations. In \\citet{dunn10a}, hardness intensity (HID) and disk fraction luminosity diagrams (DFLD) were constructed for 25 black holes.  \\citet{dunn10b} focused on the behavior of the accretion disk component in each of the outbursts studied in the previous paper. In particular, they analyzed the observed distribution of temperature versus luminosity for each outburst and found it to be consistent with $\\rm L_{Disk} \\tilde{\\propto} T_{Disk}^4$, and with previous studies, which found the inner radius of the accretion disk to be constant in disk dominated states. They do note a deviation from this relation at the entry to/exit from the disk dominated state, which they suggest is consistent with a varying color correction factor (e.g., \\citealt{shimura95}, see \\S\\ref{accretion_disk}iii for further details). However, due to the 3 keV lower energy limit of \\textit{RXTE}, this study was necessarily restricted to relatively hot accretion disks ($\\rm T \\gtrsim 0.5~keV$).  Despite the significant progress made over recent years, a number of unanswered questions remain. In this study, we focus our attention on the following issues: \\vspace{-2mm} \\begin{enumerate}[(i)] \\item{Does the accretion disk recede from the ISCO during the soft to hard state transition?} \\vspace{-2mm} \\item{What is the nature of the accretion geometry in the low hard state, i.e., Is the accretion disk truncated? What is the source of the observed hard X-ray flux, a jet or a corona?} \\vspace{-2mm} \\item{What produces the UV emission, does it originate from the accretion disk or is there a non-thermal contribution?} \\end{enumerate} \\vspace{-2mm} The \\textit{Swift} observatory was designed to discover and provide rapid multi-wavelength follow-up of gamma ray bursts (GRBs). These same attributes allow detailed study of many other time variable phenomena, in particular accreting stellar mass black holes. Due to the rapid pointing capability and low pointing overhead \\textit{Swift} provides the first detailed CCD resolution monitoring of black hole binaries in outburst. Combined with the strictly simultaneous observations at optical/UV wavelengths provided by UVOT, \\textit{Swift} is the ideal platform to study the accretion process in detail. In comparison to \\textit{RXTE}, the low energy X-ray cutoff allows us to study the accretion disk across the entire outburst cycle, whereas previously the cool disk would have been outside the RXTE low energy cutoff ($\\sim$ 3 keV).  In this paper, we describe observations of black hole binaries undertaken with the \\textit{Swift} observatory prior to June 2010.  In \\S2, we describe the observations and extraction of source spectra and lightcurves. We proceed to analyze the data in \\S3 and the results are presented in \\S4. In particular, we find that there is no evidence for significant truncation of the accretion disk, at least to luminosities $\\rm \\gtrsim 10^{-3}~L_{Edd}$. In the broadband data, we find evidence for an irradiated accretion disk at all luminosities probed, but increasing as we enter the low-hard state. These results are discussed in the context of models for the accretion flow in \\S5, and finally our conclusions are presented in \\S6. ",
        "conclusions": "We have presented a comprehensive analysis of all of the data on stellar mass black binaries in the \\textit{Swift} public archive as of June 2010. This is the largest collection of CCD spectra of accreting stellar mass black holes published to date, amounting to 476 spectra of 27 binary systems. The excellent sensitivity and broad soft X-ray bandpass of \\textit{Swift} allows us to study the evolution of the accretion disk down to temperatures as low as $\\sim$ 0.2 keV for the first time. Approximately half of the spectra are found to require the addition of an accretion disk component at greater than the 5$\\sigma$.  There is no compelling evidence in the data presented herein, which samples the accretion disk evolution down to luminosities of $\\rm \\sim 10^{-3}~L_{Edd}$, for a large scale change in the inner disk radius during the transition from the hard to the soft state, i.e., $\\rm R_{in} \\lesssim 40~R_g$ at all times. Consideration of the evolution of the inner radius versus spectral hardness points towards spectral hardening being a factor in the observed disk radius evolution.  Clearly, future observations probing luminosities below $\\rm 10^{-3}~L_{Edd}$ are necessary to provide further constraints on the evolution of the accretion flow geometry.  Irradiation is found to be present in the majority of observed spectra where simultaneous optical/UV and X-ray data are available. This irradiated disk component is found to dominate the optical/UV emission in both hard and soft states. The soft state spectrum is found to be consistent with a combination of an irradiated accretion disk plus corona, whereas for the hard state a significant contribution from a power-law component (i.e., a jet) cannot be ruled out at optical/UV energies.  We find the relationship between the flux emitted by the accretion disk and that emitted by the corona observed herein is in broad agreement with the observed behavior in Seyfert galaxies, suggesting a scale invariant coupling between the accretion disk and the corona. Though, further study is warranted.  The data and analysis presented herein clearly demonstrates the value of the broadband observational capabilities of the \\textit{Swift} observatory, in the study of accretion physics.  An analysis of the neutron star XRBs observed by \\textit{Swift}, and a comparison with the black hole binary systems presented herein will be presented in a companion paper (Reynolds et al., 2012 in prep.).  Finally, we would like to re-emphasise the importance of the broadband coverage provided by \\textit{Swift} ($\\rm \\sim 1eV - 10~keV$) in constraining the properties of the observed accretion disk emission, which is key if we are to understand the relationship between the accretion disk inflow and the jet outflow."
    },
    "1112/1112.1695_arXiv.txt": {
        "abstract": "We present an analysis of the ages and star-formation history of the F-type stars in the Upper Scorpius (US), Upper Centaurus-Lupus (UCL) and Lower Centaurus-Crux (LCC) subgroups of Scorpius-Centaurus (Sco-Cen), the nearest OB association. Our parent sample is the kinematically-selected {\\it Hipparcos} sample of \\citet{1999AJ....117..354D}, restricted to the 138 F-type members.  We have obtained classification-resolution optical spectra and have also determined the spectroscopic accretion disk fraction.  With {\\it Hipparcos} and {\\it 2MASS} photometry, we estimate the reddening and extinction for each star and place the candidate members on a theoretical H-R diagram.  For each subgroup we construct empirical isochrones and compare to published evolutionary tracks.  We find that 1) our empirical isochrones are consistent with the previously published age-rank of the Sco-Cen subgroups, 2) subgroups LCC and UCL appear to reach the main sequence turn-on at spectral types $\\sim$F4 and $\\sim$F2, respectively.  An analysis of the A-type stars shows US reaching the main sequence at about spectral type $\\sim$A3. 3) The median ages for the pre-main sequence members of UCL and LCC are 16~Myr and 17~Myr, respectively, in agreement with previous studies, however we find that 4) Upper Sco is much older than previously thought. The luminosities of the F-type stars in US are typically a factor of $\\sim$2.5 less luminous than predicted for a 5~Myr old population for four sets of evolutionary tracks. We re-examine the evolutionary state and isochronal ages for the B-, A-, and G-type Upper Sco members, as well as the evolved M supergiant Antares, and estimate a revised mean age for Upper Sco of 11$\\pm$1$\\pm$2~Myr (statistical, systematic). Using radial velocities and {\\it Hipparcos} parallaxes we calculate a lower limit on the kinematic expansion age for Upper Sco of $>$10.5~Myr (99\\% confidence).  However, the data are statistically consistent with no expansion.  We reevaluate the inferred masses for the known substellar companions in Upper Sco using the revised age and find the inferred masses are typically $\\sim$20--70\\% higher than the original estimates which had assumed a much younger age; specifically, we estimate the mass of 1RXS~J1609-2105b to be 14$^{+2}_{-3}$~M$_\\mathrm{Jup}$, suggesting that it is a brown dwarf rather than a planet. Finally, we find the fraction of F-type stars exhibiting H$\\alpha$ emission and/or a K-band excess consistent with accretion to be 0/17 ($<19$\\%; 95\\% C.L.) in US at $\\sim$11~Myr, while UCL has 1/41 ($2^{+5}_{-1}$\\%; 68\\% C.L.) accretors and LCC has 1/50 ($2^{+4}_{-1}$\\%; 68\\% C.L.) accretors at $\\sim$16~Myr and $\\sim$17~Myr, respectively. ",
        "introduction": "Embedded clusters hosting massive stars dominate the star formation of the galaxy \\citep{2003ARA&A..41...57L}, and OB associations appear to be the unbound ``fossils'' of dissolved embedded clusters.  Subgroups within OB associations are thought to be an approximately ``coeval'' stellar population \\citep[e.g.,][]{2007prpl.conf..345B}, sharing a common age, chemical abundance and velocity. Ages of stellar populations are a critical key to determining the evolutionary timescales of various stages of circumstellar disk evolution and, together with the accretion disk fraction, constrain planet formation timescales \\citep[e.g.,][]{2011ApJ...738..122C,2009AIPC.1158....3M}.  In addition, the estimated masses for very low-mass companions to members of the stellar population depend critically on the assumed age \\cite[e.g.,][]{2008ApJ...689L.153L,2011ApJ...726..113I}. It is therefore important to estimate a consistent age for the stellar population so it can be used with confidence for such calculations, and to periodically test this derived age with constantly improving observational data against modern theoretical evolutionary models.  In this paper, we explore the F-type ($\\sim$1.2--1.7 \\msun\\, for ages $\\sim$10-15 Myr) stars of Scorpius-Centaurus (Sco-Cen), the nearest OB association \\citep{2008hsf2.book..235P}.  Sco-Cen consists of three subgroups: Upper Scorpius (US), Upper Centaurus-Lupus (UCL), and Lower Centaurus-Crux (LCC), with mean distances of 145pc, 140pc and 118pc, respectively \\citep{1999AJ....117..354D}.  As the nearest site of recent massive star formation, the stellar membership of Sco-Cen offers important samples of young stars for understanding circumstellar disk evolution across the mass spectrum.  Ages for the Sco-Cen subgroups have been determined primarily from the main sequence turn-off and the low-mass pre-main sequence population combining H-R diagram positions with theoretical evolutionary tracks.  For the high-mass stars, this has been performed most recently for UCL and LCC by \\cite{2002AJ....124.1670M}, obtaining ages $\\sim$17~Myr and $\\sim$16~Myr, respectively.  The most recent published age-dating of the massive stars of US was performed by \\cite{1989A&A...216...44D}, obtaining an age of $\\sim$5~Myr using the evolutionary tracks of \\cite{1981A&A...102..401M} and corroborated by \\citet{2002AJ....124..404P}.  Our goal in this paper is to explore the star-formation history and accretion disk fraction of the F-type members of three subgroups of Scorpius-Centaurus: Upper Scorpius (US), Upper Centaurus-Lupus (UCL) and Lower Centaurus-Crux (LCC).  Specifically, we have employed the kinematically-selected {\\it Hipparcos} sample from \\cite{1999AJ....117..354D} and have 1) placed them on a theoretical H-R diagram, 2) used published evolutionary tracks to obtain ages and masses, 3) compared the star-formation history of the F-type stars with previous results \\citep{2002AJ....124.1670M,2002AJ....124..404P}, and 4) estimated the spectroscopic accretion disk fraction to constrain the disk dispersal timescale for intermediate mass stars. ",
        "conclusions": "We can summarize our conclusions as follows:  \\begin{enumerate} \\item The pre-main sequence ``turn-on'' for UCL and LCC occur near spectral type $\\sim$F2 and $\\sim$F4, respectively.  The F-type members of US appear to be all pre-main sequence.  Examining the A-type members of US allows us to estimate the pre-main sequence turn-on for US to be near spectral type $\\sim$A3.  \\item The good agreement between the isochronal ages of F-type stars studied here and the G- and K-type stars studied in \\cite{2002AJ....124.1670M} provides greater confidence in those ages, and more firmly establishes that UCL is slightly younger than, or roughly coeval with, LCC.  In order of youngest to oldest: Upper Scorpius (US), Upper Centaurus-Lupus (UCL), Lower Centaurus-Crux (LCC).  \\item The median ages obtained with the pre-main sequence F-type members (those later than ~F5) for UCL and LCC agree with previous results for early B-type stars and pre-main sequence G-type stars. \\citep{2002AJ....124.1670M}, with median ages of 16$\\pm$1~Myr and 17$\\pm$1~Myr, respectively.  \\item The F-type members have a median isochronal age of 13$\\pm$1~Myr (68\\% C.L.) for US, which is much older than previous results.  Using the approximate pre-main sequence turn-on point with the A-type stars allows us to estimate an age of 10$\\pm$3~Myr. Re-examining the early B-type stars and the G-type pre-main sequence stars in US yields median ages of 10$\\pm$2~Myr and 9$\\pm$2~Myr, respectively.  Age estimates for the M1.5Iab-Ib supergiant US member Antares give 12$^{+3}_{-1}$~Myr.  Considering these to be independent estimates, we obtain an overall mean age for Upper Sco of 11$\\pm$1$\\pm$2~Myr (statistical, systematic) with the uncertainty in the mean age dominated by systematic differences between the isochronal ages inferred from different mass ranges.  \\item We find a 99\\% confidence lower limit on the kinematic expansion age for Upper Sco of 10.5~Myr using trigonometric parallax and radial velocity data. However, the astrometry and radial velocities are statistically consistent with no expansion.  \\item Based on H$\\alpha$ emission accretion diagnostics, we estimate a spectroscopic accretion disk fraction of 0/17 ($<$19\\%; 95\\% C.L.) in US, consistent with the {\\it Spitzer} results of \\cite{2006ApJ...651L..49C} of 0/30 ($<$11\\%; 95\\% C.L.) for F- and G-type stars. In UCL we find 1/41 ($2^{+5}_{-1}$\\%; 68\\% C.L.) accretors and in LCC we find 1/50 ($2^{+4}_{-1}$\\%; 68\\% C.L.) accretors, which are both consistent with the spectroscopic accretion disk fraction for the G- and K-type members of UCL and LCC from \\cite{2002AJ....124.1670M} for G- and K-type stars.  \\item The revised age of 11~Myr for Upper Sco may explain the lower than expected disk fraction compared to 5~Myr groups (e.g., $\\lambda$~Ori; \\citealt{2007ApJ...664..481B}). We reevaluate the masses of the known substellar companions in Upper Sco using available published data and our revised age of 11~Myr, and find that inferred masses are typically $\\sim$20--70\\% more massive than previously estimated.  All of the imaged companions thus far appear to be more massive than the deuterium-burning limit and so none of the companions are in the planetary-mass regime. In addition to larger substellar masses, the older age for Upper Sco implies a more realistic, lower progenitor mass of $\\sim$18--20~\\msun\\, for the young neutron star RX~J1856.5-3754.  \\end{enumerate}"
    },
    "1112/1112.1971_arXiv.txt": {
        "abstract": "We present a new method to quantify substructures in clusters of galaxies, based on the analysis of the intensity of structures. This analysis is done in a residual image that is the result of the subtraction of a surface brightness model, obtained by fitting a two-dimensional analytical model ($\\beta$-model or S\\'{e}rsic profile) with elliptical symmetry, from the X-ray image. Our method is applied to 34 clusters observed by the \\textit{Chandra} Space Telescope that are in the redshift range $z \\in[0.02,0.2]$ and have a signal-to-noise ratio greater than 100. We present the calibration of the method and the relations between the substructure level with physical quantities, such as the mass, X-ray luminosity, temperature, and cluster redshift. We use our method to separate the clusters in two sub-samples of high and low substructure levels. We conclude, using Monte Carlo simulations,  that the method  recuperates very well the true amount of substructure for small angular core radii clusters (with respect to the whole image size) and good signal-to-noise observations. We find no evidence of correlation between the substructure level and physical properties of the clusters such as mass, gas temperature, X-ray luminosity and redshift. The scaling relations for the two sub-samples (high and low substructure level clusters) are different (they present an off-set, i.e., given a fixed mass or temperature, low substructure clusters tend to be more X-ray luminous), which is an important result for cosmological tests using the mass-luminosity relation to obtain the cluster mass function, since they rely on the assumption that clusters do not present different scaling relations according to their dynamical state. ",
        "introduction": "Clusters of galaxies are the largest virialized objects in the Universe, the upper limit of collapsed halo mass function. In an Universe dominated by a cosmological constant and cold dark matter ($\\Lambda$CDM), dark matter halos are formed by gravitational instability from primordial quantum fluctuations in the mass density field. The amplitude of those fluctuations increases as they cease expanding with the Hubble flux, collapse and virialize, forming dense and relaxed structures. Smaller structures grow to larger ones through mergers, up to clusters of galaxies in the present time. In this hierarchical scenario of structures formation, clusters are thus dynamically young objects and contain evidence of their recent past merging history \\citep[e.g.][]{1993Kauffmann}. We can relate substructures with the cluster dynamical age \\citep[e.g.][]{1992Richstone,2003Suwa}: the more substructure (their total intensity) a cluster presents, the younger (dynamically speaking) it is.    The hot intra-cluster plasma is a powerful X-ray source and its observation reveals the projected spatial distribution of most of the baryonic mass. X-ray studies of galaxy clusters are thus particularly relevant in this context, as they can give us clues to the dynamical age of clusters (e.g., \\citet{2000Henriksen} -- Abell 3266, \\citet{2003LimaNeto} -- Abell 970, \\citet{2005Ferrari} -- Abell 3921). Analysis of substructure in the intra-cluster plasma spatial distribution should help us determining the dynamical state of galaxy clusters. A very good review about the theory and observational status of the study of substructures based on X-ray data in clusters of galaxies is given by \\citet{2005Jeltema}. Here, we only briefly discuss some of the previous work on cluster substructures.  \\citet{1992Jones} made the first X-ray systematic study of structures in galaxy clusters, visually analyzing 208 objects observed by the \\textit{Einstein} satellite, establishing that merging must be a common phenomenon in clusters. \\citet{1992Richstone} developed in an original theoretical study a relation between substructures and cosmology, where they put constraints on cosmological parameters by the fractional rate of major mergers in clusters.  X-ray surface brightness allows us to perform statistical tests such as centroid and ellipticity variation \\citep{1995Mohr}, relating the dynamical age of clusters with its morphology. \\citet{1995Buote,1996Buote} developed a method to quantify X-ray substructures in clusters of galaxies from the moments of the expansion in Fourier series of the X-ray surface brightness. \\citet{2005Jeltema} used the same method, referred to as the \\textit{power-ratio method}, in a sample of 40 clusters of galaxies observed by \\textit{Chandra}. They showed that clusters in general are less relaxed at $z>0.5$, than at $z\\simeq 0$.  Semi-analytic methods give an indication of the expected evolution of cluster substructure and its dependence on cosmological parameters, however, the best method of constraining cosmological models is probably through the comparison with hydrodynamic cluster simulations. For instance, \\citet{2003Suwa} compared simulated clusters in a $\\Lambda$CDM and an OCDM cosmology, at both $z = 0$ and $z = 0.5$, using several methods for quantifying structure. They restrict themselves to comparing the ability of different statistical indicators in distinguishing different simulated cosmologies, showing that cluster structure can potentially constrain $\\Omega_\\Lambda$ or the dark energy equation of state.  Although a lot of effort has been done in order to advance our understanding about substructure in clusters of galaxies, from theoretical to numerical simulation studies, we propose in the present work a novel method of quantifying substructures that has a simple physical interpretation: the substructure level, the way it is defined, reflects the fraction of the total X-ray luminosity that is emitted by the substructures, serving as a tool to understand the underlying physical processes taking place during the cluster evolution.   This paper is organized as follows. In section 2, we describe the sample selection and in section 3 the data reduction and analysis are discussed. In section 4, the substructure level is defined and our method is described, with its calibrations being discussed in section 5. In section 6 the results are presented and discussed and conclusions are finally presented in section 7. The cosmology assumed in this paper is given by $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_0=70$~km~s$^{-1}$Mpc$^{-1}$. ",
        "conclusions": "We have developed a new method to quantify X-ray substructures in clusters of galaxies based on the ratio between the number of counts in the residual and original X-ray images. We calibrated the method and  then applied it to 34 clusters of galaxies in order to obtain the substructure level dependence with physical parameters, such as mass, temperature, X-ray luminosity and redshift.  The calibration was done using Monte Carlo simulations, which showed that the method recuperates very well the true amount of substructure for small angular core radii clusters (with respect to the whole image size) and good signal-to-noise ratio observations.   The substructure level high scatter (spanning from less than 1 to $\\simeq 16\\%$) in the redshift range $z \\in [0.02,0.2]$ shows that clusters are found in all dynamical states in the local Universe: from those relaxed to completely disturbed.    We have not found any strong evidence of correlation between the substructure level and physical properties of the clusters, gas temperature, X-ray luminosity and redshift. However, there seems to be a correlation between the substructure level and the cluster mass, which is given by its Pearson correlation coeficient of 0.49. For our sample of 34 clusters it represents a probability of not presenting correlation of the order of $0.3\\%$.  The distinction between high and low substructure level clusters has shown to be interesting, since different scaling relations were found with these two sub-samples (they present an off-set of $\\sim 2\\sigma$ -- given a fixed mass or temperature, low substructure clusters tend to be more X-ray luminous), which is an important result for cosmological tests which use the cluster mass-luminosity relation to compute the mass function.  A practical application of our method would be the identification of clusters of very low substructure level. Such relaxed clusters would be ideal laboratories for studies where the equilibrium hypothesis is of paramount importance. As an example, the work of Bertolami et al. (2007), on the interaction between dark matter and dark energy, made use of the Layzer-Irvine equation, which must hold for a system in virial equilibrium when there is no interaction in the dark sector. They applied their method to Abell 586, based on the analysis of Cypriano et al. (2005), which suggests that Abell 586 is indeed a very relaxed cluster.  Finally, it is important to say that the method itself is interesting since it concerns a new way to quantify substructures in clusters of galaxies, with a very simple physical interpretation: it reflects the fraction of the X-ray luminosity provided by substructures."
    },
    "1112/1112.1140_arXiv.txt": {
        "abstract": "We present the baseline multimessenger analysis method for the joint observations of gravitational waves (GW) and high-energy neutrinos (HEN), together with a detailed analysis of the expected science reach of the joint search. The analysis method combines data from GW and HEN detectors, and uses the blue-luminosity-weighted distribution of galaxies. We derive expected GW+HEN source rate upper limits for a wide range of source parameters covering several emission models. Using published sensitivities of externally triggered searches, we derive joint upper limit estimates both for the ongoing analysis with the initial LIGO-Virgo GW detectors with the partial IceCube detector (22 strings) HEN detector and for projected results to advanced LIGO-Virgo detectors with the completed IceCube (86 strings). We discuss the constraints these upper limits impose on some existing GW+HEN emission models. ",
        "introduction": "\\label{section:Introdution}  The observation of gravitational-waves (GWs) and neutrinos is entering a new and promising era with newly constructed detectors. GW observatories such as LIGO \\cite{LIGO0034-4885-72-7-076901}, Virgo \\cite{0264-9381-28-11-114002} and GEO \\cite{2010CQGra..27h4003G} will be upgraded to second generation detectors within the next few years. Another advanced GW detector, LCGT \\cite{LCGT0264-9381-27-8-084004}, is being constructed in Japan, while LIGO is considering plans for a third observatory in India \\cite{LIGOindia}.  Both the emission mechanism and detection method of neutrinos can depend greatly on their energy, therefore it is worthwhile to consider different neutrino subgroups.  So far, only astrophysical MeV neutrinos (in addition to those from the Sun) have been detected, and only in one case, from supernova SN 1987A \\cite{PhysRevLett.58.1490,PhysRevLett.58.1494}. Here we concentrate specifically on high-energy neutrinos (HENs): neutrinos $\\gtrsim 100$ GeV that carry information about the particle acceleration region of astrophysical sources \\cite{2008PhR...458..173B}.  HEN observatories currently in operation are IceCube \\cite{IceCubeAhrens2004507}, a cubic-kilometer detector at the geographic South Pole, and \\textsc{Antares} \\cite{ANTARES} in the Mediterranean sea. \\textsc{Antares} is being upgraded to a cubic-kilometer detector called KM3NeT in the following years \\cite{deJong2010445}. A third HEN detector operating at the lake Baikal is also planned to be upgraded to a km$^3$ volume \\cite{Avrorin2011S13}.  The joint (multimessenger) analysis of GW and HEN observations presents multiple advantages over single messenger analyses. Since both GWs and HENs are weakly interacting, their detection requires exceptionally high sensitivity. Search sensitivity can be greatly increased by combining data from GW and HEN detectors. The multimessenger approach could also significantly add to our understanding of the underlying mechanisms that create the astrophysical sources emitting both signals \\cite{ET}. Aso \\emph{et al.} \\cite{asoicecube0264-9381-25-11-114039} designed a GW+HEN multimessenger search algorithm that uses the two LIGO detectors and IceCube to look for spatially and temporally coincident events. After requiring temporal coincidence between the GW and HEN signals within a given time window, the method reconstructs a ring on the sky based on the measured time delay between the signal arrival times in the LIGO detectors, and requires the direction of the neutrino signal to overlap with this ring. Aso \\emph{et al.} show that such coincidence is extremely unlikely to arise from the background, making the requirement of coincidence very effective in reducing the false alarm rate of the joint measurement.  Pradier \\cite{microquasarHENPradier2009268} considered a joint search with initial Virgo and \\textsc{Antares}, and discussed the feasibility of a time coincident search using these two detectors. Pradier discussed microquasars and flares from soft gamma repeaters as plausible galactic GW+HEN sources. He also considered the detectability of quantum gravity effects by measuring the time delay between the arrival of GW and HEN signals.  We are close to the milestone of finishing the first coincident search for GWs and HENs for the initial LIGO-Virgo (S5/VSR1 science runs) and the partial \\textsc{Antares} detector in its 5-string configuration. The analysis, a simpler version of what is presented below, uses the directional distribution and the time of arrival of HENs to trigger a GW follow-up analysis, similar to the analysis used for GW follow-up searches of GRBs (e.g. \\cite{0004-637X-715-2-1438}). There are $\\sim$200 neutrino triggers from the \\textsc{Antares}, most of which are detected by digital optical modules (DOMs) two strings, while 13 neutrinos are detected by DOMs on three strings. The first scientific results of this search will be published soon.  In this article we introduce a joint GW and HEN analysis algorithm that combines the observations of a network of GW detectors and a HEN detector. Besides looking for astrophysical GW+HEN messengers, the search algorithm can also be used to derive upper limits on the population and flux of GW+HEN sources. We estimate the anticipated science reach for initial and advanced detectors, and discuss some of the existing emission models and how they can be constrained in the event of non-detection.  The distribution of astrophysical GW+HEN sources at detectable distances is not uniform. This can be used in a joint search algorithm to reduce false coincidences and increase sensitivity. One can weigh event candidates based on the expected source density in their direction. The method utilizes the blue-luminosity-weighted galaxy distribution to favor astrophysical sources over background coincidences. Blue luminosity is a good tracer of the rate of star formation and therefore CCSN rate (e.g. \\cite{1999A&A...351..459C}). Long GRB rates in galaxies also correlate with the galaxies' blue luminosity \\cite{2010MNRAS.405...57S} (typically long GRBs are more likely to occur in smaller, bluer galaxies compared to CCSNe. The method weighs source directions with the expected detectable source rate from the given directions, assuming that source distribution follows blue-luminosity distribution. We take the blue-luminosity distribution of galaxies up to 40~Mpc from the Gravitational Wave Galaxy Catalogue (GWGC) \\cite{0264-9381-28-8-085016}.  The article is organized as follows. In section \\ref{section:sources} we discuss some anticipated astrophysical sources of GWs and HENs. Section \\ref{section:sciencereach} discusses the expected science reach of joint GW+HEN searches by presenting the interpretation of expected results for constraining existing emission and population models. Section \\ref{section:gwhendata} describes GW and HEN detectors and data. In section \\ref{section:jointanalysis} we introduce the baseline joint GW+HEN analysis, also discussing its relation to the presented science reach. Section \\ref{section:conclusion} presents a summary of the method and the science reach. ",
        "conclusions": "\\label{section:conclusion}  We presented a baseline method for the joint search for GWs and significances of GW and HEN signals, including the point spread functions of event sky locations, as well as the blue-luminosity-weighted distribution of known galaxies.  We estimated the expected science reach of a joint GW+HEN search based on some of the existing GW and HEN emission models. The expected results indicate that the GW+HEN search with initial and particularly with advanced detectors will constrain the parameter space of some of the existing models.  To interpret the science reach of the expected GW+HEN population upper limits, we considered HEN rate expectations based on the Waxman-Bahcall emission model \\cite{waxmanbachall}, as well as the emission model with mildly relativistic jets by Ando and Beacom \\cite{chokedfromreverseshockPhysRevD.77.063007}.  The baseline GW+HEN analysis is expected to result in better upper limits than independent searches if the increase in exclusion distance compared to an independent GW search is greater than $\\sim f_b^{1/3}$ (where $f_b$ is the HEN beaming factor; see also \\cite{PhysRevLett.107.251101}). Further, the main advantage of a multimessenger search over individual searches is that of detection efficiency, as well as additional scientific information available from the source. Due to the non-Gaussianity of the GW data stream \\cite{LIGO0034-4885-72-7-076901}, high-significance events have relatively high false alarm rate and therefore a coincident messenger can greatly increase our confidence in a detection. Additionally, joint detection would help better understand the underlying physics of the source. For instance a HEN detected prior to a GW from a GRB would indicate precursor activity in the GRB. The time delay of the earliest HEN after the GW signature of a collapsar may be indicative of jet propagation within the stellar envelope."
    },
    "1112/1112.3974_arXiv.txt": {
        "abstract": "{Accurate physical properties of eclipsing stars provide important constraints on models of stellar structure and evolution, especially when combined with spectroscopic information on their chemical composition. Empirical calibrations of the data also lead to accurate mass and radius estimates for exoplanet host stars. Finally, accurate data for unusual stellar subtypes, such as Am stars, also help to unravel the cause(s) of their peculiarities.} {We aim to determine the masses, radii, effective temperatures, detailed chemical composition and rotational speeds for the Am-type eclipsing binaries \\SW\\ (A4-5m) and \\HW\\ (A6m) and compare them with similar normal stars.} {Accurate radial velocities from the Digital Speedometers of the Harvard-Smithsonian Center for Astrophysics were combined with previously published $uvby$ photometry to determine precise physical parameters for the four stars.  A detailed abundance analysis was performed from high-resolution spectra obtained with the Nordic Optical Telescope (La Palma).} {We find the masses of the (relatively evolved) stars in \\SW\\ to be 2.10 and 2.24\\,$M_{\\sun}$, with radii of 2.50 and 3.01\\,$R_{\\sun}$, while the (essentially zero-age) stars in \\HW\\ have masses of 1.72 and 1.78\\,$M_{\\sun}$, radii of 1.64 and 1.66\\,$R_{\\sun}$ -- all with errors well below 2\\%. Detailed atmospheric abundances for one or both components were determined for 14 elements in \\SW\\ ([Fe/H] = +0.49/+0.61 dex) and 16 in \\HW\\ ([Fe/H] = +0.33/+0.32 dex); both abundance patterns are characteristic of metallic-line stars. Both systems are well fit by current stellar evolution models for assumed bulk abundances of [Fe/H] = +0.05 and +0.23, respectively ([$\\alpha$/Fe] = 0.0), and ages of $\\sim$700~Myr and 160~Myr.} {} ",
        "introduction": "\\label{sec:intro}  Among the eclipsing binaries with the best determined fundamental properties approximately two dozen are of spectral type A (see Torres et al.\\ \\cite{Torres:10a}).  Roughly 20--30\\% of main-sequence A stars are known to be metallic-lined (Am). The Am stars are found overwhelmingly to be in binary systems (e.g., Abt \\cite{Abt:61}). They are characterized by peculiar abundance patterns including a deficiency of light elements such as Sc and Ca, enhancement of iron-peak elements, and a strong overabundance of Sr and some rare earths, sometimes by an order of magnitude or more. The enhancement generally increases with increasing atomic number. Their typically slower axial rotation and stronger metal lines compared to normal A stars tend to facilitate the radial-velocity measurements necessary for accurate mass determinations.  It is therefore not surprising that systems with Am stars are over-represented among the best-measured eclipsing binaries: fully half of the A-type systems for which we know the absolute masses and radii with relative errors less than 3\\% contain at least one component showing these abundance abnormalities.  Although significant progress in understanding and modeling these chemically peculiar objects has been made in recent years, important questions regarding the role of rotation and the origin of the detailed patterns of element enhancement and depletion are not yet completely settled (see, e.g., Abt \\cite{Abt:00}, B\\\"ohm-Vitense \\cite{Bohm-Vitense:06}, Vick et al.\\ \\cite{Vick:10}). Further observational constraints are highly desirable, particularly ones in which very accurate mass and radius determinations are accompanied by a detailed abundance study of the components.  \\begin{figure}[b] \\resizebox{\\hsize}{!}{\\includegraphics{fig1.eps}} \\epsfxsize=80mm \\caption{Main sequence eclipsing binaries taken from Torres et al.\\ (\\cite{Torres:10a}) in which both components have masses in the range 1.6--2.4 $M_{\\sun}$, and relative errors in the masses and radii of 3\\% or better. Primary and secondary components are connected with dotted lines, and stars that are confirmed or probable Am stars are marked with open circles.  \\SW\\ and \\HW\\ are shown with filled symbols.  The thick solid lines are drawn for reference, and correspond to solar-metallicity isochrones from the series of Yi et al.\\ (\\cite{Yi:01}) for ages of 50~Myr and 1~Gyr.} \\label{fig:debs} \\end{figure}  Beyond a qualitative assessment of their metallic-lined nature, individual element abundances are known for precious few of the dozen or so Am binaries with well determined properties (see, e.g., Lyubimkov et al.\\ \\cite{Lyubimkov:96}). Thus, little can be said about whether or how the detailed abundance patterns that seem to vary to some degree from star to star depend on global properties such as mass, effective temperature, or surface gravity. In this paper we present extensive new observations of two eclipsing Am systems, \\SW\\ and \\HW, with the goal of determining not only accurate absolute dimensions but also individual abundances for all four stars. It is hoped that these results and others like them will present an opportunity for advancement on some of the issues mentioned above. Figure~\\ref{fig:debs} places the two systems in the context of all other main sequence eclipsing binaries in the same mass regime that have the best-determined masses and radii. Am stars are marked with open circles. \\SW\\ is seen to be among the more evolved cases, while \\HW\\ is the least evolved in this mass range.  We describe below our spectroscopic observations of these systems, and an analysis of those data along with previously-published high-precision four-color Str\\\"omgren photometry in order to obtain accurate masses and radii. For \\SW\\ our determinations substantially improve upon previous estimates; for \\HW\\ they represent the first such measurements. We conclude with a detailed comparison with models of stellar and tidal evolution. ",
        "conclusions": "\\label{sec:discussion}  The following sections present a comparison of the measured properties for \\SW\\ and \\HW\\ against predictions from the theory of stellar evolution, stellar structure, and tidal evolution.   \\subsection{Stellar evolution} \\label{sec:stellarevolution}  Accurate mass, radius, and effective temperature determinations, especially when accompanied by a measurement of the chemical composition, are among the most powerful constraints afforded by eclipsing binary stars to test models of stellar structure and evolution (see, e.g., Andersen \\cite{ja91}, Torres et al.\\ \\cite{Torres:10a}, and references therein). When these four quantities are available for both components in the system, there are essentially no free parameters in the comparison with publicly available models: to the extent that the measurements are accurate, the models will either be successful in reproducing all observations simultaneously within their errors, or they will fail.\\footnote{Additional variables such as the helium abundance, mixing length parameter, convective core overshooting, etc., have generally already been chosen in advance by the modelers in published tables, and are not ``tunable'' by the user. Specific models for a particular system can of course be computed with different values of these additional variables, as we do in one example below, adding more freedom to the fits.}  In the case of Am stars such as \\SW\\ and \\HW, however, the measured abundances do not represent the bulk composition of the stars, which is the relevant quantity when comparing with models. Instead, they are only a reflection of local changes in the surface layers attributed to diffusion processes, which give the spectra of these objects their peculiar appearance. Metallicity therefore remains a free parameter when testing models for these stars. Nevertheless, because detailed abundance patterns in Am stars vary from case to case and are not yet completely understood, a comparison with models is still of considerable interest -- even if less constraining -- to establish the precise evolutionary state of the stars. Except for second-order effects due to limb darkening, absolute masses and radii are essentially independent of the theory being tested. For \\HW, however, our light-curve solutions have had to make additional use of model atmospheres to constrain the geometry of the system, because of the lack of secondary eclipses. We point out, though, that those models were used only in a differential sense (to determine \\emph{ratios} of central surface fluxes or bandpass-specific luminosities), so in effect the dependence on theory is very weak.  \\begin{figure}[t] \\epsfxsize=85mm \\epsfbox{sw_plottrackg.eps} \\caption{\\label{fig:swcma_tg} Measurements for \\SW\\ compared against Yonsei-Yale models by Yi et al.\\ (\\cite{Yi:01}) for [Fe/H] $= +0.05$ and [$\\alpha$/Fe] = 0.0.  Evolutionary tracks for the measured masses are shown with solid lines and shaded areas indicating the uncertainty in the location of each track coming from the mass errors. Isochrones from 500~Myr to 900~Myr in steps of 100~Myr are represented by the dashed lines. The best fitting 700~Myr isochrone is drawn with a thicker line.} \\end{figure}  \\begin{figure} \\epsfxsize=85mm \\epsfbox{sw_plotmrt.eps} \\caption{\\label{fig:swcma_mr} Effective temperatures and radii for \\SW\\ shown as a function of the measured mass, with the observational errors represented by the shaded boxes. They are compared against [Fe/H] $= +0.05$ Yonsei-Yale isochrones by Yi et al.\\ (\\cite{Yi:01}) from 500~Myr to 900~Myr, in steps of 100~Myr, with [$\\alpha$/Fe] = 0.0. The best-fitting model is indicated with a heavier line.} \\end{figure}  In this section we compare the absolute dimensions of \\SW\\ and \\HW\\ with stellar evolution models from the Yonsei-Yale series (Yi et al.\\ \\cite{Yi:01}; Demarque et al.\\ \\cite{Demarque:04}), which have been shown previously to match observations very well. We have chosen these models primarily because they are provided with software to interpolate evolutionary tracks (and isochrones) for any given mass and metallicity, the mass being among the most accurately known properties for these stars.  The $\\log g$ vs.\\ $T_{\\rm eff}$ diagram for \\SW\\ in Fig.~\\ref{fig:swcma_tg} allows a simultaneous comparison of the models against the three key observables: $M$, $R$, and $T_{\\rm eff}$. Evolutionary tracks corresponding to a composition slightly above solar ([Fe/H] $= +0.05$, with no $\\alpha$-element enhancement) are shown for the exact masses we measure (Table~\\ref{tab:swcma_absdim}), with the gray areas representing the uncertainty in the location of the tracks that comes from the mass errors (in this case, a small fraction of the mass difference). Isochrones are also displayed for ages 500--900~Myr, in steps of 100~Myr. The models for this composition provide a very good fit to the observations, within the uncertainties, and the age inferred is approximately 700~Myr (indicated with the thicker isochrone). \\SW\\ is seen to be somewhat evolved, and currently near the mid-point of its main sequence phase. The implied bulk composition inferred from theory, [Fe/H] $= +0.05$, is $\\sim$0.4--0.5~dex lower than the iron abundance measured at the surface ([Fe/H] $= +0.49$/+0.61 for the primary and secondary, respectively), which may be interpreted as giving a rough measure of the \\emph{true} enhancement of the iron-peak elements in the photospheres of these stars. Strictly speaking, however, this factor of three enhancement may only be a lower limit, as the effects of metallicity and $\\alpha$-element enhancement in the models can trade off against each other to some extent, leading to mass tracks and isochrones that fit the observations equally well.  For example, a match as good as seen above is achieved also with [Fe/H] $= -0.10$ and [$\\alpha$/Fe] $= +0.22$, with [Fe/H] $= -0.20$ and [$\\alpha$/Fe] $= +0.36$, etc.  Separate diagrams of the radius and effective temperature as a function of mass are seen in Fig.~\\ref{fig:swcma_mr}, along with the same isochrones from the previous figure. The models provide a virtually exact match to the mass and radius of the components at 700~Myr, and a good fit to the temperatures as well.  \\begin{figure}[t] \\epsfxsize=85mm \\epsfbox{hw_plottrackg.eps} \\caption{\\label{fig:hwcma_tg} Similar to Fig.~\\ref{fig:swcma_tg}, but for \\HW. The best-fitting metallicity is [Fe/H] $= +0.23$. Solar metallicity tracks are shown for reference (dotted lines). The dashed lines represent Yonsei-Yale isochrones from 100~Myr to 1~Gyr in steps of 100~Myr. The best fitting 160~Myr isochrone is drawn with a thick solid line.} \\end{figure}   \\begin{figure}[t] \\epsfxsize=85mm \\epsfbox{hw_plotmrt.eps} \\caption{\\label{fig:hwcma_mr} Similar to Fig.~\\ref{fig:swcma_mr}, but for \\HW. In this case the best-fitting metallicity is [Fe/H] $= +0.23$. The dotted lines represent Yonsei-Yale isochrones from 100~Myr to 1~Gyr in steps of 100~Myr. The best fitting 160~Myr model is drawn with a heavier solid line.} \\end{figure}  Similar diagrams for \\HW\\ can be seen in Fig.~\\ref{fig:hwcma_tg} and Fig.~\\ref{fig:hwcma_mr}. In this case, the model composition that best fits the observations is [Fe/H] $= +0.23$ (with [$\\alpha$/Fe] = 0.0), and the implied age is approximately 160~Myr. Solar composition tracks are shown for reference as dotted lines in Fig.~\\ref{fig:hwcma_tg}. \\HW\\ is essentially on the zero-age main sequence. To the extent that [Fe/H] $= +0.23$ represents the overall composition of this system, it would appear that the true enhancement of the iron-peak elements in the surface layers is a modest 0.1~dex in this case, since the measured photospheric abundances are [Fe/H] $= +0.33$/+0.32. As before, this conclusion is dependent on the assumed [$\\alpha$/Fe] enhancement, which is inaccessible to observation in both of these systems.  The higher-than-solar bulk [Fe/H] $= +0.23$ composition for \\HW\\ suggested by the models is somewhat uncommon for stars in the solar neighborhood.  If one enforces in the models a more typical solar-like iron abundance, the observations are about as well reproduced by increasing the [$\\alpha$/Fe] ratio to +0.40. This combination would imply a true surface-layer iron enhancement due to the Am phenomenon in \\HW\\ of $\\sim$0.3~dex, and the inferred age would be the same as before.  However, a value of [$\\alpha$/Fe] as large as this is on the high side of the observed distribution for stars with solar iron abundance, which makes us somewhat skeptical of this scenario. Another possible way of reproducing the \\HW\\ observations with an iron composition closer to solar is to alter the helium abundance.  Tests with the Granada models of Claret (\\cite{Claret:04}) indicate, however, that the radii and effective temperatures can only be matched simultaneously with a very low helium abundance of $Y = 0.24$, at ages of 440--500\\,Myr that are significantly older than before.  A helium abundance this close to the primordial value seems unlikely.  Comparisons (not shown) of the observations for \\HW\\ and \\SW\\ with the standard ([$\\alpha$/Fe] = 0.0) Victoria-Regina stellar evolution models of VandenBerg et al.\\ (\\cite{vr06}) and with the Granada models of Claret (\\cite{Claret:04}) give results that are quite consistent with those from the Yonsei-Yale models, indicating similar overall abundances and evolutionary ages.    \\subsection{Internal structure} \\label{sec:internalstructure}  The observed apsidal motion for \\SW, $\\dot\\omega = 0.00067 \\pm 0.00021$~deg~cycle$^{-1}$ (Clausen et al.\\ \\cite{Clausen:08}), is a measure of the average degree of central mass concentration of the components. The corresponding average apsidal motion constant for the system is $\\log \\bar k_{2,\\rm obs} = -2.57 \\pm 0.30$. Using the absolute dimensions of the stars from Table~\\ref{tab:swcma_absdim}, the models by Claret (\\cite{Claret:04}) predict a theoretical value of $\\log \\bar k_{2,\\rm theo} = -2.582 \\pm 0.050$, in which the formal uncertainty does not account for possible systematic errors in the models. While this is consistent with the empirical value within the errors, the close agreement may be accidental given the difficulty of the measurement (the apsidal period is long: $U = 14\\,900 \\pm 4700$\\,yr). The General relativistic contribution ($\\dot\\omega_{\\rm GR} = 0.00034$~deg~cycle$^{-1}$) amounts to about 50\\% of the total apsidal motion.  While the eccentric system \\HW\\ is also expected to display apsidal motion, the lack of secondary minima precludes such a measurement based on eclipse timings. Spectroscopic detection of the effect would be quite challenging. The models of Claret (\\cite{Claret:04}) predict a value of $\\dot\\omega = 0.00023$~deg~cycle$^{-1}$, and a corresponding apsidal motion constant of $\\log \\bar k_{2,\\rm theo} = -2.410 \\pm 0.050$. The estimated period of this effect is $U = 88\\,900$\\,yr, six times longer than in \\SW. In the case of \\HW\\ the relativistic contribution completely dominates, accounting for 97\\% of the total apsidal motion.    \\subsection{Tidal evolution} \\label{sec:tidalevolution}  In addition to their binarity, Am stars typically display slow rotation compared to field stars of the same spectral type. The components of \\SW\\ and \\HW\\ are no exception.  We noted earlier that while the primary of \\SW\\ appears to be pseudo-synchronized with the orbital motion, the secondary is not, and rotates at only half the corresponding speed. This assumes the spin axes are parallel to the axis of the orbit, a condition that is almost universally taken for granted in binaries. Here we investigate whether synchronization is expected at all in these systems according to tidal theory, and if so, whether the situation for the secondary of \\SW\\ might be understood in terms of spin-orbit misalignment. We examine also the predicted evolution of the orbital eccentricity.  \\begin{figure*} \\epsfxsize=175mm \\hskip 0.07in\\epsfbox{figtidalland_sw.eps} \\caption{Tidal evolution calculations for \\SW, following the prescription by Hut (\\cite{Hut:81}). The vertical dotted line in each panel represents the current evolutionary age of 680\\,Myr ($\\log\\tau = 8.833$), based on the models by Claret (\\cite{Claret:04}). (a) Eccentricity as a function of time; (b) Orbital period; (c) Normalized angular rotation rate of each star. The dot-dashed line represents the evolution of the pseudo-synchronous value, and the horizontal dotted line is the pseudo-synchronous rate at the present age. (d) Angle between the equator of each star and the plane of the orbit. The four curves correspond to different initial values as given in the next panel. (e) Theoretical projected rotational velocity of the primary as a function of time, for four different initial values of the spin-orbit angle, $\\phi_0$. The measured $v \\sin i_{\\rm rot}$ and its uncertainty are indicated with the horizontal dashed line and shaded area. Theory agrees with observation at the current age of the system. (f) Same as (e), for the secondary. In this case theory is not able to match the observation, regardless of the initial values assumed for $\\Omega/\\Omega_{\\rm orb}$ and $\\phi$.\\label{fig:tidal_sw}} \\end{figure*}  We have integrated the differential equations given by Hut (\\cite{Hut:81}) for the time evolution of the semimajor axis ($da/dt$), the eccentricity ($de/dt$), the angular rotation rates of both components ($d\\Omega_1/dt$, $d\\Omega_2/dt$), and the angle between the plane of the orbit and the equator of each star ($d\\phi_1/dt$, $d\\phi_2/dt$).  These six, coupled equations were integrated simultaneously using a 4th-order Runge-Kutta method, with the stellar properties interpolated at each time step from appropriate evolutionary tracks from Claret (\\cite{Claret:04}). The turbulent dissipation timescale for the late evolutionary phases with convective envelopes was taken to be $(M R^2/L)^{1/3}$, where $M$, $R$, and $L$ are the mass, radius, and luminosity of the star. For phases in which the envelopes are radiative (the most relevant here) the timescales adopted follow closely those in Eq.\\,17 and Eq.\\,18 by Claret \\& Cunha (\\cite{Claret:97}).  The initial conditions are of course not known, so they are in effect free parameters of the model. For the orbital period and eccentricity we have adjusted the starting values so as to match $P$ and $e$ at the current evolutionary age of each system. Their time evolution for \\SW\\ is shown in Fig.~\\ref{fig:tidal_sw}a and Fig.~\\ref{fig:tidal_sw}b, in which the vertical dotted lines mark the current age of about 680\\,Myr ($\\log\\tau = 8.833$), according to the Granada models employed in this section. Based on these calculations the orbit is expected to circularize at an age of $\\tau = 835$\\,Myr ($\\log\\tau = 8.921$) for this system.  The evolution of the angular rotation rate $\\Omega = 2\\pi/P_{\\rm rot}$ of each star in \\SW\\ is seen in Fig.~\\ref{fig:tidal_sw}c, where for convenience we have normalized it to the orbital rate, $\\Omega_{\\rm orb} = 2\\pi/P_{\\rm orb}$.  We have no observational constraint on the rotation rates, so the initial values were arbitrarily set to $\\Omega_1/\\Omega_{\\rm orb} = \\Omega_2/\\Omega_{\\rm orb} = 6.0$, reasonable for A-type stars that usually rotate very rapidly (at least in the field). A dot-dashed line in this figure indicates the pseudo-synchronous rate at each age, and the horizontal dotted line is the current pseudo-synchronous rate.  These calculations suggest that pseudo-synchronization for both stars in this system was reached early-on, at an age of roughly 100\\,Myr ($\\log\\tau = 8.0$). This would indeed appear to be the case for the primary of \\SW, based on its measured $v \\sin i$ value, but not for the secondary.  The degree of spin-orbit alignment is represented by the angle $\\phi$, but once again, we have no constraint on this quantity. Furthermore, the relation between $\\phi$ and the orbital and rotational inclination angles $i_{\\rm orb}$ and $i_{\\rm rot}$, both measured with respect to the line of sight, is given by \\begin{equation} \\cos\\phi = \\cos i_{\\rm orb} \\cos i_{\\rm rot} + \\sin i_{\\rm orb} \\sin i_{\\rm rot} \\cos\\lambda~, \\end{equation} which involves an unknown angle $\\lambda$ between the sky-projected angular momentum vectors of the orbit and the stellar spin.  \\begin{figure*}[t] \\epsfxsize=175mm \\hskip 0.07in\\epsfbox{figtidalland_hw.eps} \\caption{Similar to Fig.~\\ref{fig:tidal_sw}, but for the \\HW\\ system. In this case the results are illustrated for four initial values of $\\Omega/\\Omega_{\\rm orb}$ ($\\Omega_0$ for short) as labeled in panel (a), and a single initial value for $\\phi$ of 30\\degr. With the boundary conditions tuned in this way, theory matches the measured rotational velocities of both components shown in (e) and (f). Neither star is predicted to be synchronized, and the calculations suggest their spin axes have not yet been aligned with the axis of the orbit. \\label{fig:tidal_hw}} \\end{figure*}  The spectroscopically measured projected rotational velocities of the stars, which we refer to more properly now as $v \\sin i_{\\rm rot}$, do provide an indirect constraint on a combination of theoretically predictable quantities, but this still involves the unknown angle $\\lambda$. Given that $i_{\\rm orb}$ is rather close to 90\\degr\\ in both \\SW\\ and \\HW\\, we may make the approximation that $\\cos\\phi \\approx \\sin i_{\\rm rot} \\cos\\lambda$. We then have \\begin{equation} \\label{eq:tidal} v \\sin i_{\\rm rot} \\approx {2\\pi\\over P_{\\rm orb}} {\\Omega\\over \\Omega_{\\rm orb}} {\\cos\\phi\\over \\cos\\lambda} R~. \\end{equation} All quantities on the right-hand side of this equation are either known from stellar evolution calculations ($R$), or can be computed from the solution of the differential equations for tidal evolution, with the exception of the angle $\\lambda$, which depends on the observer's viewpoint. In order to make progress, we ignore this term for the moment (or equivalently, we consider $\\lambda$ to be small), so that $\\cos\\phi \\approx \\sin i_{\\rm rot}$.  The evolution of the alignment angle $\\phi$ for \\SW\\ is displayed in Fig.~\\ref{fig:tidal_sw}d for four different initial values (20\\degr, 40\\degr, 60\\degr, 80\\degr). The curves for the primary and secondary are nearly indistinguishable, so only those for the primary are shown. As seen from the convergence of the curves towards zero, spin-orbit alignment for both stars in \\SW\\ is reached at the age of about 200\\,Myr ($\\log\\tau \\approx 8.3$), which is less than a third of the current age of the system.  The predicted evolution of $v \\sin i_{\\rm rot}$ for each star, computed from Eq.\\,\\ref{eq:tidal}, is shown in Fig.~\\ref{fig:tidal_sw}e and Fig.~\\ref{fig:tidal_sw}f compared against the measured values (dashed line and shaded uncertainty region).  Independently of the initial values of $\\phi$ and $\\Omega/\\Omega_{\\rm orb}$, the predictions from tidal evolution theory cannot be made to match the anomalously slow rotation rate of the secondary of \\SW.  The angle $\\lambda$ in Eq.\\,\\ref{eq:tidal} that we have previously ignored is inconsequential for this discussion, as any value greater than zero only makes the disagreement worse.  Evidently current tidal theory is incomplete.  An effect not considered here is the possible decoupling between the core of the star and the external layers, which suffer more directly the actions of tidal forces. Since the spectroscopically measured $v \\sin i_{\\rm rot}$ corresponds only to the layers that we see, it is quite possible that the larger rotation rate predicted by theory is more representative of a more rapidly rotating core than of the outer envelope, or at the very least is some average of the two, which would be larger than the measured value.  Similar tidal calculations for \\HW\\ are shown in Fig.~\\ref{fig:tidal_hw}. In this case we have chosen to illustrate the effects of a change in the initial values of $\\Omega/\\Omega_{\\rm orb}$, rather than $\\phi$. As before, we have adjusted the boundary conditions for the period and eccentricity to match the measured values at the current age, which is 160\\,Myr ($\\log\\tau = 8.21$) according to the Claret (\\cite{Claret:04}) models. As the evolution of both of these quantities (Fig.~\\ref{fig:tidal_hw}a and Fig.~\\ref{fig:tidal_hw}b) depends on $\\Omega$ (see Hut \\cite{Hut:81}), the starting values for $P$ and $e$ are different in each case. The calculations indicate that orbital circularization will occur at an age of $\\sim$1.75\\,Gyr ($\\log\\tau = 9.244$).  The time dependence of $\\Omega/\\Omega_{\\rm orb}$ appears in Fig.~\\ref{fig:tidal_hw}c, and influences the late evolution of the angle $\\phi$ between the equator and orbital plane, as seen in Fig.~\\ref{fig:tidal_hw}d. These two panels indicate that neither star is expected to be pseudo-synchronized at the current age, nor has spin-orbit alignment been achieved. Pseudo-synchronization is predicted to happen at $\\sim$600\\,Myr ($\\log\\tau = 8.78$), and alignment at an age of approximately 1\\,Gyr. The fact that we do see agreement within the measurement errors between the observed rotational velocities and the predicted pseudo-synchronous velocities at the current epoch as listed in Table~\\ref{tab:hwcma_absdim} suggests that this may be due to a combination of non-synchronized angular rotation rates ($\\Omega_{1,2}$) and slightly misaligned spin axes ($\\phi_{1,2} \\approx 20\\degr$ at the current epoch for our particular choice of the initial values for $\\Omega_{1,2}$; see Fig.~\\ref{fig:tidal_hw}d).  The theoretical projected rotational velocities are plotted in Fig.~\\ref{fig:tidal_hw}e and Fig.~\\ref{fig:tidal_hw}f. The lower curves match the measured $v \\sin i_{\\rm rot}$ of each star at the current age, and correspond to initial $\\Omega/\\Omega_{\\rm orb}$ values of 4.0 for both stars, and initial angles $\\phi_{1,2} = 30$\\degr\\ tuned to produce the agreement.  The asynchronous rotation of the secondary of \\SW\\ is not unusual for Am stars, or for normal A stars for that matter. The compilation of Torres et al.\\ (\\cite{Torres:10a}) includes several such systems with well determined properties that show the same lack of synchronization. An extreme case is V459~Cas (Lacy \\cite{Lacy:04}), a pair of A1m stars both rotating 3.5--4.5 times faster than pseudo-synchronous.  There is no obvious correlation that we can see between asynchronous rotation and other global properties of Am or normal A stars. This issue has also been studied by others (Budaj \\cite{Budaj:96}, \\cite{Budaj:97}; Iliev et al.\\ \\cite{Iliev:06}; Prieur et al.\\ \\cite{Prieur:06}), but no credible pattern has emerged."
    },
    "1112/1112.0977_arXiv.txt": {
        "abstract": "Population synthesis models predict that the majority of sdBs form through stable mass transfer leading to long period binaries. To date, about a hundred close short-period binaries with an sdB component have been found, but not a single long-period system has been established. We show preliminary results from our recent effort in determining orbits for a sample of long-period sdB systems using the {\\sc hermes} spectrograph on the Mercator telescope. ",
        "introduction": "Binarity is a key component both in the formation and in the final evolution of stars, when expansion forces one of the companions to fill its gravitational well, and its envelope to spill over to that of the partner. Binary processes strongly influence the composition of stellar populations, and a diverse variety of evolved objects can only be explained through the direct interaction between the bloated envelope of a giant, and accretion by its companion. Such binary interactions are complex and the models invoke a number of poorly constrained parameters, such as the efficiency of envelope ejection, the physical description of the common-envelope phase, the accretion efficiency on to the companion, and many others. Observations that can constrain these model parameters are essential not just for explaining the particular systems, but for population synthesis in general, and for extrapolating to environments of various densities and metalicities as found in clusters and other galaxies.  \\begin{table}[t] \\caption{Eight targets in our sample of composite hot subdwarf / main sequence systems. The spectroscopic class from various sources is given together with the visual magnitude and the number spectra we have collected so far. }\\smallskip \\begin{center}\\small \\begin{tabular}{llrrl}\\tableline\\noalign{\\smallskip} Target name & Sp.Class & $m_V$ & N.sp & References \\\\ \\noalign{\\smallskip}\\tableline\\noalign{\\smallskip} BD--11$^\\circ$162 & sdO+? & 11.2 & 12 & \\citet{zwicky57} \\\\ & sdO+G &      &    & \\citet{berger80} \\\\ & sdOB+K0 &    &    & \\citet{ulla98} \\\\ PG 1104+243       & sdB+K2 & 11.3& 24 & \\citet{ferguson84} \\\\ & sdB+G8 &     &    & \\citet{orosz97} \\\\ Balloon 82800003  & sdB+K1 & 11.4& 16 & \\citet{bixler91} \\\\ BD+29$^\\circ$3070 & sdOB+F & 10.4& 15 & \\citet{berger80} \\\\ & sdB+K0 &     &    & \\citet{ulla98} \\\\ BD+34$^\\circ$1543 & sdB+F  &  9.4& 13 & \\citet{berger80} \\\\ & sdB+G8 &     &    & \\citet{ulla98} \\\\ BD--7$^\\circ$5977 & sdB+K0IV-III & 10.5 & 17 & \\citet{ulla98} \\\\ &        &     &    & \\citet{heber02} \\\\ Feige 80          & sdO+A  & 11.4& 21 & \\citet{berger80} \\\\ & sdO+G8 &     &    & \\citet{ulla98} \\\\ Feige 87          & sdB+?  & 11.7& 15 & \\citet{jeffery98} \\\\ \\noalign{\\smallskip}\\tableline \\end{tabular} \\end{center} \\end{table}  \\citet{han02,han03} made a thorough binary population synthesis study of the hot subdwarfs, using all three binary formation channels that are thought to contribute significantly to the population. The three are; (1) If the subdwarf progenitor has a low mass companion, then mass transfer on the RGB is unstable, and the orbit will shrink until the envelope is ejected.  The study of \\citet{maxted01} as recently completed by \\citet{copperwheat11} finds that $\\sim$50\\%\\ of all sdB stars reside in short-period binary systems ($P_{\\mathrm{orb}} < 10$\\,d). (2) If the companion is more massive than the subdwarf (at least at the end of mass transfer), the orbit will have expanded substantially. Such orbits are hard to measure, but the companion can be detected spectroscopically or from infra-red excess. \\citet{napiwotzki04} found that more than a third of their sdB sample show the spectroscopic signature of main sequence (MS) companions, while \\citet{reed04}, using 2MASS photometry, inferred that about half of the sdBs in the field have main-sequence companions, and are therefore likely to be of this post-stable-Roche-lobe-overflow (pRLOF) type. (3) The final binary formation channel is the merger of two low-mass white dwarfs, and has a much lower efficiency than the two other channels.  By now, $\\sim$100 sdB stars are known to reside in short period binaries, and they are providing clear constraints for common-envelope ejection models. A recent compilation can be found in Appendix A of \\citet{geier11a}, but new systems are discovered at a high pace. The longest period systems are 15 and 29 days respectively \\citep[from][]{morales-rueda03}. For the longer period systems very little is known. \\citet{green01} mention a mean $\\Delta v\\sin i$ of 11.5\\,km/s from 89 observations of 19 composite binaries, and estimate periods averaging 3--4 years, but provide no details about particular systems. Here we will.  \\begin{figure*}[t] \\centering \\includegraphics[width=\\textwidth]{sp.eps} \\caption{Single spectra for four of the targets in the sample. The resolution has been substantially degraded, and the wavelength region truncated to 3800--6800\\,\\AA\\ for illustration purposes. The slope of each spectrum has been corrected by calibration with observations of the single sdB star Feige\\,66, and offset for clarity. It can be seen that the contribution from the companion varies substantially between the targets. However, even in Feige\\,80, where the spectrum is well dominated by the sdB component, there are sufficient lines from the MS companion for the cross-correlation to succeed in establishing precise radial velocity measurements for the cool star.  } \\end{figure*} ",
        "conclusions": "Our long time-base spectroscopy from the Mercator telescope has revealed orbits of $\\sim$500\\,d and longer for eight sdB+F/G binaries (Figure~2 and 3). While pRLOF systems are expected to be found with a wide range of periods, a strong peak is predicted to be found just above 100\\,d \\citep[see Figure~21 of][]{han03}. All the systems shown in Figure~1 appear to have periods much longer than this peak, indicating that some of the parameters that govern the mass transfer process need to be adjusted.  While more than one hundred hot subdwarfs are known to exist in short-period systems, the long-period systems predicted as the outcome of stable Roche-lobe overflow on the first giant branch require extraordinary efforts to nail down. We have presented our first results from a dedicated survey to reveal the orbital periods of hot subdwarfs with main sequence companions, and the periods detected so far are far longer than the distribution predicted by \\citet{han03}, but in excellent agreement with the estimated average of 3--4 years mentioned by \\citet{green01}."
    },
    "1112/1112.2375_arXiv.txt": {
        "abstract": "Important informations on the origin and the propagation mechanisms of cosmic rays may be provided by the measurement of the anisotropies of their arrival direction. In this paper the observation of anisotropy regions at different angular scales is reported. In particular, the observation of a possible anisotropy on scales between $\\sim$10$^{\\circ}$ and $\\sim$30$^{\\circ}$ may be a key-detection for speculations on the presence of unknown features of the magnetic fields the charged cosmic rays propagate through, as well as to potential contributions of nearby sources to the total flux of cosmic rays. Evidence of new weaker few-degree excesses throughout the sky region $195^{\\circ}\\leq$ R.A. $\\leq 315^{\\circ}$ is also reported. ",
        "introduction": "The ARGO-YBJ experiment, located at the YangBaJing Cosmic Ray Laboratory (Tibet, P.R. China, 4300 m a.s.l., 606 g/cm$^2$), is an air shower array able to detect the cosmic radiation at an energy threshold of a few hundred GeV. The full detector is in stable data taking since November 2007 with a duty cycle greater than 85\\%. The trigger rate at the threshold is 3.6 kHz. The detector characteristics and performance are described in \\cite{moon11}. ",
        "conclusions": "This paper reports the observation of CR anisotropy at different angular scales with ARGO-YBJ, as a function of the primary energy. The large scale CR anisotropy has been clearly observed up to about 25 TeV. Evidence of existence of different few-degree excesses in the Northern sky (the strongest ones positionally coincident with the regions detected by Milagro in 2008) is reported. The time distribution of the phenomenon has been showed too. A discussion of the systematic effects which may come from imprecisions in estimating the reference level is outlined.  In fact, a wrong estimation affect the significance and the relative intensity sky maps, even creating artifacts (i.e. fake excesses or deficit regions). Drifts in detector operating conditions or atmospheric effects on the air shower development are quite hard to be modeled, then difficult to be accounted for to sufficient accuracy. As anisotropies of the order 10$^{-4}$ are looked for, operating conditions must be known down to this level, all across the field of view and during all the acquisition time.  Background methods applied in this analysis demonstrated to be effective for analysis of diffuse from even weaker regions (e.g. see \\cite{mrk421}).  Given the importance of the topic, a joint analysis of concurrent data recorded by different experiments in both hemispheres, as well as a correlation with other observables like the interstellar energetic neutral atoms distribution \\cite{ibex09,ibex11}, should be a high priority to clarify the observations. \\bigskip %"
    },
    "1112/1112.5763_arXiv.txt": {
        "abstract": "name}{Abstract} \\renewcommand{\\figurename}{Fig.}   \\def\\theequation{\\thesection.\\arabic{equation}} \\renewcommand{\\thesection}{\\arabic{section}} \\renewcommand{\\thesubsection}{\\arabic{section}.\\arabic{subsection}} \\renewcommand{\\thesubsubsection}{\\arabic{section}.\\arabic{subsection}. \\arabic{subsubsection}}  \\makeatletter \\renewcommand{\\@seccntformat}[1]{{\\csname the#1\\endcsname}.\\hspace{0.5em}} \\makeatother   \\UDK{}   \\ArticleNAME{Black holes and quasiblack holes: Some history and remarks }  \\ArticleAUTHOR{Jos\\'e P. S. Lemos\\\\ {\\small  Multidisciplinary Center for Astrophysics - CENTRA, Physics Department,  Instituto Superior T\\'ecnico - IST, Universidade T\\'ecnica de Lisboa - UTL, Portugal}}  \\ArticleHEAD{Black holes and quasiblack holes \\ldots}  \\ArticleAUTHORHEAD{Jos\\'e P. S. Lemos.}   \\makeabstitle   \\begin{abstract} A brief reference to the two Schwarzschild solutions and what Petrov had to say about them is given.  Comments on how the Schwarzschild vacuum solution describes a black hole are also provided.  Then we compare the properties, differences and similarities between black holes and quasiblack holes.  Black holes are well known. Quasiblack hole is a new concept.  A quasiblack hole, either nonextremal or extremal, can be broadly defined as the limiting configuration of a body when its boundary approaches the body's own gravitational radius (the quasihorizon).  They are objects that are on the verge of being black holes but actually are distinct from them in many ways.  We display some of their properties: there are infinite redshift whole regions; the curvature invariants remain perfectly regular everywhere, in the quasiblack hole limit; a free-falling observer finds in his own frame infinitely large tidal forces in the whole inner region, showing some form of degeneracy; outer and inner regions become mutually impenetrable and disjoint, although, in contrast to the usual black holes, this separation is of a dynamical nature, rather than purely causal; for external far away observers the spacetime is virtually indistinguishable from that of extremal black holes.  Other important properties, such as the mass formula, and the entropy, are also discussed and compared to the corresponding properties of black holes. \\par \\textbf{Key words:} Schwarzschild solution, Petrov, Black holes, Quasiblack holes. ",
        "introduction": "\\label{S:In}    \\subsection{The Schwarzschild solution.}\\label{sch} \\hskip0.1cm Finding vacuum solutions of Einstein's equation \\vskip -0.7cm \\begin{equation} G_{ab}=0\\,, \\end{equation} \\vskip -0.2cm \\noindent where $G_{ab}$ is the Einstein tensor, is an important branch of General Relativity and known to be a non-trivial task. On the other hand, finding solutions of the field equations with matter is a somewhat different setup.  Given any metric, there is always one stress-energy tensor $T_{ab}$ for which Einstein's equations $(G=1\\,,c=1)$ \\vskip -0.3cm \\begin{equation} G_{ab}=8\\pi\\,T_{ab}\\,, \\end{equation} \\vskip -0.2cm \\noindent are trivially satisfied. Now, arbitrarily chosen metrics usually give rise to unphysical stress-tensors, corresponding to matter which is of no interest. Therefore, the task of finding non-vacuum solutions to the field equations is, in a certain way, twice as hard in comparison to solutions in vacuum, one has to choose physically relevant sources, and then solve for the gravitational field in the equations.  Schwarzschild, in 1916, in two strokes, initiated the field of exact solutions in General Relativity, both in vacuum \\cite{Schw1} and in matter for an incompressible fluid \\cite{Schw2}. These solutions are called the Schwarzschild solution and the interior Schwarzschild solution, respectively. The Schwarzschild solution \\cite{Schw1} is perhaps the most well-known exact solution in General Relativity, and its line element can be written in appropriate spherical coordinates $(t,r,\\theta,\\phi)$ as, \\begin{equation} ds^2=-\\left(1-\\frac{2m}{r}\\right)\\,dt^2+ \\frac{dr^2}{\\left(1-\\frac{2m}{r}\\right)}+ r^2\\left(d\\theta^2+\\sin^2\\theta\\,d\\phi^2\\right)\\,. \\label{leschw} \\end{equation} Here $m$ is the mass of the object, outside which there is vacuum. To interpret the solution as a whole vacuum solution, and the emergence of the notion of a black hole it took some time.   \\subsection{Petrov on the Schwarzschild solution.}\\label{pet} \\hskip0.2cm In a Petrov Symposium it is worth to spend some lines on what Petrov had to say on both Schwarzschild solutions. For this we refer to his book {\\it Einstein spaces}, published in Russian in 1961 and then translated into English in 1969 \\cite{petrov1}.    On p. 141 of the book \\cite{petrov1} one can read the rather remarkable phrase: ``It is clear that Einstein, Hilbert, and their contemporaries had a rather primitive idea of what is meant by `spacetime metric' and of its scope. They possessed only a few of the simplest examples (for example Schwarzschild's solution, the solution of Weyl and Levi-Civita with axial symmetry, and cosmological metrics). They did not realize what a powerful instrument they were forging.''  Then there are several mentions, in passing, of the Schwarzschild solution.  On p. 179 it is stated that the Schwarzschild solution is a particular case of solutions included in $T_1$, i.e., solutions with Segre characteristic $(111)$, referring to his algebraic classification of 1954 of the Riemann and Weyl tensors \\cite{petrov2}, repeated in the book in page 99.  On p. 196 Kotler's solution is mentioned, stating it is a generalization of the Schwarzschild solution by including a cosmological term $\\Lambda$.  On p. 360, in Chapter 9, Einstein's equations for a spherically symmetric vacuum are solved, and the Schwarzschild solution is finally displayed, as a textbook should do.  On p. 362, exercises on Schwarzschild and interior Schwarzschild are given, and the Landau and Lifshitz 1948 book {\\it The Classical Theory of Fields} (and the English translation of 1959) is cited \\cite{landaulifsh}.  On p. 386 the two Schwarzschild's papers of 1916 on the vacuum and the interior solutions are quoted in citations 37 and 37a, respectively.   There is an interesting contribution of Petrov to the field of exact solutions. In the paper {\\it Gravitational field geometry as the geometry of automorphisms} \\cite{petrov3}, among a discussion of many solutions, Petrov finds a Type I (111) solution with metric $ ds^2=e^r\\cos\\sqrt{3}r(-dt^2+d\\phi^2)-2\\sin\\sqrt{3}r\\,d\\phi\\,dt + dr^2+e^{-2r}dz^2 $. It is the only vacuum solution admitting a simply-transitive four-dimensional maximal group of motions.  Bonnor \\cite{bonnoronpetr} showed that it is the vacuum solution exterior to an infinite rotating dust, a particular case of the Lanczos-van Stockum solution.  This is no black hole but has relations to the hoop conjecture, closed timelike curves and so on.      \\subsection{Black holes.}\\label{bhsfinally} \\hskip0.2cm It is clear that the Schwarzschild solution (\\ref{leschw}) presents a problem, in the coordinates used, at $r=2m$.  For a long time $r=2m$ was a mysterious place. Only in the 1960s the ultimate interpretation was given and the problem solved. The radius $r_{\\rm h}=2m$ defines the event horizon, a lightlike surface, of the solution. In its full form it represents a wormhole, with its two phases, the white hole and the black hole, connecting two asympotically flat universes \\cite{kruskal} (a work done under the supervision of Wheeler \\cite{wheelbiog}). If, besides a mass $m$ as in the Schwarzschild solution, one includes electrical charge $q$, the Reissner-Norstr\\\"om solution is obtained \\cite{reiss,nordst} (for the interpretation of its full form see  \\cite{gravesbrill}). The inclusion of angular momentum $J$ gives the Kerr solution \\cite{kerr}, and the inclusion of the three parameters ($m,J,q$) is the Kerr-Newman family \\cite{newman}. For a full account of the Kerr-Newman family within General Relativity see \\cite{mtw}.  As predicted early by Oppenheimer and Snyder \\cite{oppenheimer} black holes can form through  gravitational collapse of a lump of matter. As the matter falls in, an event horizon develops from the center of the matter, and stays put, as a null surface, in the spherical symmetric case at $r_{\\rm h}=2m$, while the matter falls in towards a singularity. A posterior important result is that if the matter is made of perfect fluid (such as the Schwarzschild interior solution \\cite{Schw2}) there is the Buchdahl limit \\cite{buch} which states that when the boundary of the fluid matter approaches quasistatically the value $\\frac98\\,r_{\\rm h}$, then the system ensues in an Oppenheimer-Snyder collapse, presumably into a black hole.    The possibility of existence of black holes came with Quasars in 1963. Salpeter \\cite{salpeter} and Zel'dovich \\cite{zeldovich} were the first to advocate that a massive central black hole should be present in these objects in order to explain the huge amount of energy liberated by them.  Lynden-Bell in 1969 then took a step forward and proposed that a central massive black hole should inhabit every galaxy \\cite{lyndenbell}, a prediction that has been essentially confirmed, almost every galaxy has a central black hole. Then with the discovery of pulsars in 1968 and the reality of neutron stars the possibility of small stellar mass black holes became obvious, confirmed in 1973 with the X-ray binary Cygnus X1 and then with other X-ray binary sources (see, e.g., \\cite{lemosbhsgalelempart}).    It is supposed that black holes can form in many ways.  The traditional manner is the Oppenheimer-Snyder type collapse \\cite{oppenheimer}.  Nowadays, one also admits that black holes can form from the collision of particles, or have a cosmological primordial inbuilt origin (see, e.g., \\cite{lemosbhsgalelempart}). The Reissner-Nordstr\\\"om black hole may be not very useful astrophysically, although all black holes should have a tiny, fluctuating, charge. Notwithstanding it might be important in particle physics, perhaps it is an elementary soliton of gravitation, as proposed by some supergravity ideas. Nowadays there is a profusion of theoretical black holes of all types, in all theories, with all charges, in all dimensions (see, e.g., \\cite{profusion}).   Classically, black holes are well understood from the outside: there is astrophysical evidence and theoretical consistency. Perhaps there will be  phenomenological evidence in the near future from the collision of particles.  Quantically, black holes still pose problems.  For the outside, these problems are related to the Hawking radiation and the Bekenstein-Hawking entropy.  For the inside, the understanding of the inside of a black hole is one of the outstanding problems in gravitational theory, and it certainly is a quantum phenomenon. The horizon harbors a singularity. What is a singularity? The two quantum problems, the outside and the inside, are perhaps related. There are many approaches, some try to solve part of the problems others all of them (see, e.g., \\cite{lemosfundamental}). These approaches are the quantum gravity approach, mass inflation, wormhole, regular black hole, holographic reasoning (see, e.g., \\cite{lemosholographic}) and so on. Here, we advocate the quasiblack hole approach to better understand a black hole, both the outside and the inside stories. We do not claim to solve the problems, we look at it through a different angle and see where it leads us to.   \\subsection{Quasiblack holes.}\\label{qbhsfinally} \\hskip0.2cm Following \\cite{buch}, for matter made of perfect fluid there is the Buchdahl limit. However, putting charge into the matter to bypass the limit opens up a new world. The charge can be electrical, or angular momentum, or many other charges. The simplest case is to have matter with electric charge alone, nothing else.  In Newtonian gravitation, i.e., for a Newton-Coulomb system, the solution is easy. Suppose one has two massive charged particles. Then, the gravitational force exerted on each particle is $F_{\\rm g}=\\frac{Gm^2}{r^2}$, where for a moment we have restored $G$, and the electric force is $F_{\\rm e}=\\frac{e^2}{r^2}$. Thus, when $\\sqrt{G}m=e$ it implies $F_{\\rm g}=F_{\\rm e}$. The system is in equilibrium. Of course, if we put another such particle, any number of particles, a continuous distribution of matter, any symmetry, any configuration, the result still holds. For a continuous distribution the relation $\\sqrt{G}\\rho_{\\rm m}=\\rho_{\\rm e}$ must hold, where $\\rho_{\\rm m}$ and $\\rho_{\\rm e}$ are the mass-energy density and the electric charge density, respectively.  In General Relativity, i.e., for an Einstein-Maxwell system, the history is long. Weyl in 1917 \\cite{weyl1} started with a static solution in the form, \\begin{equation} ds^2=-W^2(x^i)\\,dt^2+g_{ij}(x^k)\\,dx^i\\,dx^j. \\label{weyl1} \\end{equation} Then he sought $W$ such that $W^2=W^2(\\phi)$,  in vacuum with axial symmetry, where $\\phi$ is the electric potential. He found $W^2= \\left(\\,\\sqrt{G}\\,\\phi+b\\right)^2+c$, with $b$ and $c$ constants. In 1947 Majumdar \\cite{maj} showed that Weyl's quadratic function works for any symmetry, not only axial symmetry. It was also shown that the (vacuum) extremal Reissner-Nordstr\\\"om solution obeyed this quadratic relation, and that many such solutions could be put together since, remarkably, equilibrium would be maintained, as in the Newton-Coulomb case. Papapetrou \\cite{papa} also worked along the same lines. Hartle-Hawking in 1973 \\cite{hh73} worked out the maximal extension and other properties of a number of extremal black holes dispersed in spacetime. Furthermore, for a perfect square, $W^2(\\phi)= \\left(\\,\\sqrt{G}\\,\\phi+b\\right)^2$, if now there is matter, Majumdar and Papapetrou found that  $\\sqrt{G}\\rho_{\\rm m} =\\rho_{\\rm e}$ \\cite{maj,papa}, and the matter is in an equilibrium state, bringing into General Relativity the Newtonian result. This type of matter we call extremal charged dust matter. The solutions, vacuum or matter solutions, are generically called Majumdar-Papapetrou solutions.   Now, if one wants to make a star one has to put some boundary on the matter. The interior solution is then Majumdar-Papapetrou and the exterior is extremal Reissner-Nordstr\\\"om. This analysis was started by Bonnor who since 1953 has called attention to them, see, e.g., \\cite{bonnor1999}. Examples of Bonnor stars are: (i) A star of clouds, in which each cloud has 1 proton and $10^{18}\\,$neutrons, so to maintain the relation $\\rho_{\\rm m}=\\rho_{\\rm e}$ ($G=1$). For a spherically symmetric star with radius $R$, the star as a whole has $m=q$, and the exterior is extremal Reissner-Nordstr\\\"om, see Figure 1. \\begin{figure}[htb] \\label{bonnorstar} \\centering \\includegraphics{starofclouds.eps} \\caption{A star of clouds as an example of a Bonnor star: Majumdar-Papapetrou (extremal charged dust) matter inside, extremal Reissner-Nordstr\\\"om outside, and a boundary surface joining the inside and outside at the radius $R$.} \\end{figure} (ii) A star made of supersymmetric stable particles with $m_{\\rm s}=e_{\\rm s}$.  Again, the star has total mass $m$ and total charge $q$ related by $m=q$.  Now comes the important point. For any star radius $R$ the star is in equilibrium. Inclusive for $R=r_{\\rm h}$, where $r_{\\rm h}=m$ is the gravitational, or horizon, radius of the extremal Reissner-Nordstr\\\"om metric. What happens when $R$ shrinks to  $r_{\\rm h}$? Something new: a quasiblack hole forms. ",
        "conclusions": "Black holes are generic and stable. Quasiblack holes perhaps not. Any perturbation would lead them into a black hole, although the inclusion of pressure may stabilize the system.  However, stable or not, the quasiblack hole approach can elucidate many features of black holes such as the mass formula and the entropy. The quasiblack hole approach to the understanding of black hole physics seems somehow like the membrane paradigm \\cite{thorne}.  Indeed, by taking a timelike matter surface into a null horizon, in a limiting process, we are recovering the membrane paradigm.  One big difference is that our membrane is not fictitious like the membrane of the membrane paradigm, it is a real matter membrane.  \\vspace{\\baselineskip}  \\noindent {\\bf Acknowledgments.} I would like to thank Vilson Zanchin (S\\~ao Paulo) and Oleg Zas\\-lavskii (Kharkov) for the many works in collaboration related to quasiblack holes. I thank Alexander Balakin and Asya Aminova for inviting me to the Petrov Anniversary Symposium in Kazan held in November 2010. I thank the interest in my talk by Gennady Bisnovatyi-Kogan who raised the important point connected with the stability of quasiblack holes and suggested a way of accessing the problem. I also thank the interest in my talk by Mikhail Katanaev and Dieter Brill who raised the important point connected with the Penrose diagrams of quasiblack holes. One of my motivations to come to this Petrov Anniversary Symposium was a work on the Petrov classification of tensors with four indices, such as the Levi-Civita tensor, with my student Andr\\'e Moita. For personal reasons, he was not able to participate in the Symposium, so our contribution has not appeared in it. Nevertheless, we will publish the results elsewhere. I also thank for financial support the Funda\\c{c}\\~{a}o para a Ci\\^{e}ncia e Tecnologia (FCT) through projects CERN/FP/109276/2009 and PTDC/FIS/098962/2008 and the grant SFRH/BSAB/987/2010, and the Reitoria da Universidade T\\'ecnica de Lisboa for specific support to the presentation of my talk at the Petrov Symposium.    \\renewcommand{\\refname}{\\small {Literature}}"
    },
    "1112/1112.0236_arXiv.txt": {
        "abstract": "We present a near infrared study focused on the detection and characterization of the youngest stellar component of the NGC\\,604 giant star-forming region, in the Triangulum galaxy (M\\,33). By means of color-color diagrams derived from the photometry of {\\it JHK$_s$} images taken with Gemini-NIRI, we have found 68 candidate massive young stellar objects. The spatial distribution of these sources matches the areas where previous studies suggested that star formation might be taking place, and the high spatial resolution of our deep NIRI imaging allows to pinpoint the star-forming knots. An analysis of the fraction of objects that show infrared excess suggests that the star formation is still active, supporting the presence of a second generation of stars being born, although the evidence for or against sequential star formation does not seem to be conclusive. ",
        "introduction": "In this paper we present a near infrared (NIR) photometric study of the NGC\\,604 star-forming region. This study is focused on the detection and first characterization of the NGC\\,604 newest generation of stars. NGC\\,604 is one of the most outstanding giant H{\\sc ii} regions (GHRs) in the Local Group. With a L$_{H\\alpha}$\\,=\\,2.6$\\times$10$^{39}$ erg\\,s$^{-1}$ \\citep{2002MNRAS.329..481B} it is the second most luminous H{\\sc ii} region after 30 Doradus. The observation and study of these two regions play a fundamental role in our understanding of the astrophysical processes associated with massive stars and, in particular, massive star formation and the environments where it takes place. At a distance of 840 kpc, NGC\\,604 is located far enough to observe global characteristics but still close as to resolve individual objects and small structures (1$\\arcsec$\\,$\\sim$\\,4 parsecs). NGC\\,604 and 30 Dor provide the link between Galactic GHRs, which we are able to study in full detail, and distant GHRs, whose global properties can be traced to cosmological distances. Although they share several common characteristics they also exhibit substantial differences such as their spatial structure and ages distribution.  The ionizing stellar cluster is composed of a population of at least 200 O-type stars \\citep{1996ApJ...456..174H}. Wolf-Rayet (WR) and Of stars, or candidates, were first identified and classified in studies by \\cite{1981ApJ...249..471C}, \\cite{1981ApJ...248.1015D}, \\cite{1982A&A...108..339R}, \\cite{1987MNRAS.226...19D}, and \\cite{1998ApJ...505..793M}. New WR objects were detected by \\cite{1993AJ....105.1400D} by means of Hubble Space Telescope (HST) images, and an accurate classification of known WR stars was made in \\cite{2008MNRAS.389.1033D}. \\cite{1996MNRAS.279.1219T} spectroscopically classified two objects, one of them is a transition star from Luminous Blue Variable (LBV) to WR and the other is a red supergiant (RSG) star. There are forty OB stars with accurate spectral classification by \\cite{2003AJ....125.3082B}, three of them exhibit Of/WR signatures (wind profile UV emission lines of N\\,{\\sc v}, C\\,{\\sc iv}, and Si\\,{\\sc iv} together with  strong emission line of He\\,{\\sc ii} at $\\lambda$1640). \\cite{2011MNRAS.411..235E} studied the known evolved stellar population in the region by means of Spectral Energy Distribution (SED) fitting of HST {\\it UVIJHK} photometry. The authors used the {\\it UVI} photometric measurements from \\cite{1996ApJ...456..174H} and they performed NIR photometry using archive images from HST-NICMOS (NIC2) in the F110W, F160W, and F205W filters. The latter HST data set, was originally used by \\cite{2009Ap&SS.324..309B} to perform a photometric study of the region, in which they found five RSG candidates and a dozen candidate massive young stellar objects (MYSOs). However, the selection criteria for IR-excess sources was based on the maximum errors present in the NICMOS photometry. This suggests that a considerable number of sources may yet be detected with deeper images, using photometric errors appropriate for each object.  The age of the main stellar population was estimated by several authors. \\cite{1996ApJ...456..174H} analyzed the color-magnitude diagrams obtained from {\\it UVI} HST-WFPC2 photometry; the authors fitted isochrones and considered the age constrains imposed by the presence of WR candidate stars. They concluded that the average ages for the stars in NGC\\,604 range from 3 to 5 Myr, and that the presence of RSG could suggest the existence of an older subpopulation. This older subpopulation was later studied by \\cite{2011MNRAS.411..235E} who concluded that the RSGs belong to a formation episode that occurred 12.4$\\pm$2.1 Myrs ago. \\cite{2000MNRAS.317...64G} analyzed stellar wind resonance lines in ultraviolet spectra taken with the International Ultraviolet Explorer (IUE), together with the nebular emission lines and the higher order terms of the Balmer series and He\\,{\\sc i} lines in absorption (both obtained from optical spectra taken with the 4.2-m William Herschel Telescope). The authors applied three different techniques of evolutionary synthesis and photoionization models (optimized for young star-forming regions) and concluded, consistently by the three methods, that the main ionizing cluster has an age of $\\sim$\\,3 Myr. The physical and kinematical properties of NGC\\,604 were studied in \\cite{2000PASP..112.1138M} (and references therein), by means of spectroscopic data obtained with the WHT (in particular, the lines formed in the warm ionized gas) and images from HST-WFPC2. The author concluded that NGC\\,604 suffered one large starburst 3.0-3.5 Myr ago that formed the main ionizing cluster in the region.  Regarding the stellar population spatial distribution in the region, it is worth mentioning that the objects in NGC\\,604 are not centrally concentrated (as is the case for 30 Dor) but spread over a large projected area ($\\sim$ 10000 pc$^2$ for cluster A in \\citealt{1996ApJ...456..174H}), composing a structure called the Scaled OB Association (SOBA) by \\cite{2004AJ....128.1196M}.  According to the studies performed by \\cite{2000ApJ...541..720T} and \\cite{2004AJ....128.1196M}, the interstellar medium exhibits two areas with different excitation state: a high excitation region surrounded by a low-excitation region, but not with a concentric geometry. These structures also show different kinematic behavior and physical conditions. On the west side (high ionization) there can be identified three cavities, which seem to be bubbles bursting into the M\\,33 halo (in our direction). There is a boundary, often cited as `the ridge', which extends along the north-south direction defining a clear east-west division within nebulae. The east side (low ionization) exhibits a quieter kinematic behavior. This dichotomy is also observed in the X-ray emission \\citep{2008ApJ...685..919T}.  Detailed studies on the dynamics were also performed by \\cite{1996AJ....112..146Y}, who found that the width of the integrated velocity profile needs the contributions of virial motion, thermal broadening, stellar winds and SNRs. \\cite{1997ApJ...487..163M} performed a study on the dynamics of the region by means of the H$\\alpha$ emission and they concluded that a few tens of WR stars would be enough to feed the necessary kinetic energy. They also stated that the turbulence is relatively young, giving another element in favor of recent processes of stellar formation taking place a few million years ago. Analyzing the velocity dispersion versus intensity diagram over the whole region, \\cite{1996AJ....112.1636M} found that most velocity dispersion values measured in NGC\\,604 are supersonic. The authors also used the same diagram as a tool to identify the zones corresponding to structures such as shells, loops or bubbles which are generated by the interaction of the massive stars' winds with the interstellar medium.   The NGC\\,604 molecular component was widely studied by \\cite{1992A&A...265..437V}, \\cite{1992ApJ...385..512W}, \\cite{2003ApJS..149..343E}, \\cite{2003ApJ...599..258R}, and \\cite{2007ApJ...664L..27T}. In the latter study the authors proposed a sequential star formation scenario triggered by the expansion of the H{\\sc ii} region. This idea was reinforced in a recent study with higher spatial resolution in CO, HCN, and 89-GHz continuum emissions by \\cite{2010ApJ...724.1120M}.  In this study we present new observations of NGC\\,604 acquired using the Gemini-North telescope with NIRI in the {\\it J}, {\\it H}, and {\\it K$_{short}$} ({\\it K$_s$} hereafter) bands. The high quality data obtained allowed us to perform a photometric study to identify and characterize candidate MYSOs in the ionizing cluster, confirming the candidates found by \\cite{2009Ap&SS.324..309B} and increasing the sample presented in that study by a factor of five. Our data and analysis methods are described in Section \\ref{sec:data}. Results and implications are discussed in Section \\ref{sec:results}. ",
        "conclusions": "\\label{conclusions} We have performed a NIR photometric study on the NGC\\,604 massive star-forming region, using Gemini-NIRI images.  Our study was focused on the detection and individual identification of massive young stellar objects that comprise a new stellar generation in the region.  By means of {\\it JHK$_s$} photometry we have identified 68 candidates MYSOs. These objects were selected by their position in the CC diagram (each of these objects exhibits IR excess greater than their photometric error in  ({\\it H-K$_s$})).  The analysis of the spatial distribution shows that most of the MYSOs candidates are located in areas with strong nebular emission, delineated by the radio continuum contours at 8.44 GHz, where previous studies stated the possibility of star formation processes taking place.  We have calculated the fraction of objects with IR excess throughout the whole region. The result yields that several regions in NGC\\,604 show a high fraction of objects with IR excess for the cluster age. Although it is not possible to perform a direct comparison with star forming regions in our Galaxy, this could be interpreted as a recent burst of star formation. Studies of other GHRs are necessary to establish whether there is a general tendency for regions of massive star formation to present higher IR-excess fractions and to understand the true nature of IR-excesses in MYSOs.  We have made a short review of individual objects in NGC\\,604 included in our sample, which were already studied by other authors.  All these results complete a general picture that exhibits the existence of a new stellar generation in the NGC\\,604 region, which is mainly taking place in the molecular gas that extends to the south-east side of the main SOBA population.  This is the first study of the NGC\\,604 star forming region dedicated to the detection of individual MYSOs candidates. New detailed observations for each object are needed to confirm the nature of the MYSOs candidates we have found and also to provide information on the physical conditions of the circumstellar matter which is causing the observed IR excess.\\\\    Acknowledgments  This research was based on observations obtained at the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the Science and Technology Facilities Council (United Kingdom), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), Minist\\'{e}rio da Ci\\^{e}ncia e Tecnologia (Brazil) and Ministerio de Ciencia, Tecnolog\\'{i}a e Innovaci\\'{o}n Productiva (Argentina). The authors would like to specially thank Andrew Stephens, from Gemini Observatory, for his help regarding the data processing procedures of Gemini-NIRI images and Andrew McWilliam for his final revision of the English text. We have also benefited from fruitful discussions with Roberto Terlevich. We are grateful to the anonymous referee for the thorough review which has led to a much improved version from the original manuscript."
    },
    "1112/1112.0478_arXiv.txt": {
        "abstract": "{ According to the theory of hadronic acceleration, ultra high energy cosmic rays are expected to be accompanied by  gamma-rays and neutrinos from pion decays formed in the interactions of protons with photons. While gamma-rays have been linked to astrophysical sources by many experiments (H.E.S.S, MAGIC, Fermi), no point source of UHECRs or neutrinos have been found so far. In this contribution, we present the results of multimessenger stacking sources analysis developed to investigate the correlation of arrival directions of neutrino candidate events and UHECRs. This analysis has been applied on neutrino candidate events detected during 2007-08 by the ANTARES telescope (ANTARES collaboration 2011), and 69 ultra-high energy cosmic rays observed by the PAO (The Pierre Auger Collaboration 2010). } % ",
        "introduction": "Finding point sources of both UHECRs and neutrinos is a very challenging task, as both messengers have many down-sides. UHECRs are rare and do not point back to their sources, since they are scrambled by galactic and intergalactic magnetic fields \\cite{lab1,lab2,lab3,lab4,lab5,lab6,lab7,lab8,lab9,lab10,lab11,lab12,lab13,lab14,lab15,lab16}. Also, due to the interaction with the cosmic microwave background photons, their range may be limited to the distance of about 100Mpc or less \\cite{lab17,lab18,lab19,lab20,lab21}. However, they are detectable with large shower arrays, like the Pierre Auger Observatory (PAO), which so far reported the observation of few tens of events above 55EeV \\cite{lab21,lab22}. On the other side, cosmic neutrinos, as neutral and weak interacting particles, should point back to their sources and their traveling distances should not be limited, but at the same time they are very difficult to detect. Currently operating neutrino telescopes ANTARES and IceCube have not yet observed excess above the atmospheric neutrino flux coming from air showers \\cite{lab23,lab24}.  Previously, the Pierre Auger Observatory (PAO) reported an anisotropy in the arrival directions of UHECRs \\cite{lab21} and indicated a correlation with Active Galactic Nuclei (AGN) from the VCV catalog \\cite{lab25}. The correlation was the most significant for 27 cosmic rays with energies higher than 57EeV and AGNs at distances less than $~$75Mpc. The suggested correlation with the nearby AGN sources mostly following the location of the supergalactic plane decreased in the subsequent analysis \\cite{lab22} with 69 events at energies above 55EeV, observed until 31st December 2009.  In this paper, we investigate the correlation of arrival directions of 2190 neutrino candidate events detected by 5-12 line ANTARES neutrino telescope, and 69 UHECRs observed by the PAO. This stacking sources analysis is developed for a blinded set of neutrino events. Blinding was performed by scrambling in right ascension. At the end, the analysis was done with the unblinded neutrino dataset. ",
        "conclusions": ""
    },
    "1112/1112.5541_arXiv.txt": {
        "abstract": "Measurements of the low energy spectrum of Galactic Cosmic Rays (GCR) by detectors at or near the Earth are affected by Solar modulation. To overcome this difficulty, we consider nearby molecular clouds as GCR detectors outside the Solar system. Using \\gr\\ observations of the clouds by the Fermi telescope we derive the spectrum of GCRs in the clouds from the observed \\gr\\ emission spectrum. We find  that the GCR spectrum has a low energy break with the spectral slope hardening by $\\Delta \\Gamma= 1.1\\pm 0.3$ at an energy of  $E=9\\pm 3$~GeV.  Detection of a low-energy break enables a measurement of GCR energy density in the interstellar space $U=0.9\\pm 0.3$~eV/cm$^3$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.3960_arXiv.txt": {
        "abstract": "A wealth of astronomical data indicate the presence of mass discrepancies in the Universe.  The motions observed in a variety of classes of extragalactic systems exceed what can be explained by the mass visible in stars and gas.  Either (i) there is a vast amount of unseen mass in some novel form -- dark matter -- or (ii) the data indicate a breakdown of our understanding of dynamics on the relevant scales, or (iii) both. Here, we first review a few outstanding challenges for the dark matter interpretation of mass discrepancies in galaxies, purely based on observations and independently of any alternative theoretical framework. We then show that many of these puzzling observations are predicted by one single relation -- Milgrom's law -- involving an acceleration constant $a_0$ (or a characteristic surface density $\\Sigma_\\dagger = a_0/G$) of the order of the square-root of the cosmological constant in natural units. This relation can at present most easily be interpreted as the effect of a single universal force law resulting from a modification of Newtonian dynamics (MOND) on galactic scales. We exhaustively review the current observational successes and problems of this alternative paradigm at all astrophysical scales, and summarize the various theoretical attempts (TeVeS, GEA, BIMOND, and others) made to effectively embed this modification of Newtonian dynamics within a relativistic theory of gravity. ",
        "introduction": "\\label{section:introduction}  Two of the most tantalizing mysteries of modern astrophysics are known as the \\textit{dark matter} and \\textit{dark energy} problems. These problems come from the discrepancies between, on one side, the observations of galactic and extragalactic systems (as well as the observable Universe itself in the case of dark energy) by astronomical means, and on the other side, the predictions of General Relativity from the observed amount of matter-energy in these systems. In short, what astronomical observations are telling us is that the dynamics of galactic and extragalactic systems as well as the expansion of the Universe itself do not correspond to the observed mass-energy as they should if our understanding of gravity is complete.  This thus indicates \\emph{either} (i) the presence of unseen (and yet unknown) mass-energy, \\emph{or} (ii) a failure of our theory of gravity, \\emph{or} (iii) both.  The third case is \\textit{a priori} the most plausible, as there are good reasons for there being more particles than those of the standard model of particle physics \\cite{supersym} (actually, even in the case of baryons, we suspect that a lot of them have not been seen yet and thus literally make up \\textit{unseen mass}, in the form of ``missing baryons''), and as there is \\textit{a priori} no reason that General Relativity should be valid over a wide range of scales where it has never been tested \\cite{Bertolami1}, and where the need for a dark sector actually prevents the theory from being tested until this sector has been detected by other means than gravity itself\\epubtkFootnote{Up to now, all the dark matter particle candidates still elude both direct and indirect non-gravitational detection.}. However, either of the first two cases could be the \\textit{dominant} explanation of the discrepancies in a given class of astronomical systems (or even in all astronomical systems), and this \\textit{is} actually testable.  For instance, as far as (ii) is concerned, if the mass discrepancies in a class of systems are mostly caused by some subtle change in gravitational physics, then there should be a clear signature of a single, universal force law at work in this whole class of systems. If instead there is a distinct dark matter component in these, the kinematics of any given system should then depend on the particular distribution of both dark and luminous mass.  This distribution would vary from system to system, depending on their environment and past history of formation, and should in principle not result in anything like an apparent universal force law\\epubtkFootnote{However, a way to effectively reproduce an apparent universal force law from an exotic dark component could be to enforce an intimate connection between the distribution of baryons, the dark component, and the gravitational field through, e.g., a fifth force effect. This possibility will be extensively discussed in Section~\\ref{sec:covariant}, notably Section~\\ref{sec:dipolar}}.  Over the years, there have been a large variety of such attempts to alter the theory of gravity in order to remove the need for dark matter and/or dark energy. In the case of dark energy, there is some wiggle room, but in the case of dark matter, most of these alternative gravity attempts fail very quickly, and for a simple reason: once a force law is specified, it must fit all relevant kinematic data in a given class of systems, with the mass distribution specified by the visible matter only. This is a tall order with essentially zero wiggle room: at most one particular force law can work. However, among all these attempts, there is one survivor: the \\textit{Modified Newtonian Dynamics} (MOND) hypothesized by Milgrom almost 30 years ago~\\cite{Milgrom1, 22, original} seems to come close to satisfying the criterion of a universal force law in a whole class of systems, namely \\textit{galaxies}. This success implies a unique relationship between the distribution of baryons and the gravitational field in galaxies and is extremely hard to understand within the present dominant paradigm of the concordance cosmological model, hypothesizing that General Relativity is correct on every relevant scale in cosmology including galactic scales, and that the dark sector in galaxies is made of non-baryonic dissipationless and collisionless particles. Even if such particles are detected directly in the near to far future, the success of MOND on galaxy scales as a phenomenological law, as well as the associated appearance of a universal critical acceleration constant $a_0 \\simeq 10^{-10}\\mathrm{\\ m\\, s}^{-2}$ in various seemingly unrelated aspects of galaxy dynamics, will anyway have to be explained and understood by \\textit{any} successful model of galaxy formation and evolution. Previous reviews of various aspects of MOND, at an observational and theoretical level, can be found in~\\cite{Bek06,Bruneton07,Clifton,Ferreira09,MdB98,Blois08,MDorDM,SM02,Scarpa06,Skordis09}. A blog-website dedicated to this topic is also maintained, with all the relevant literature as well as introductory level articles~\\cite{MONDpages} (see also~\\cite{Scilogs}).  Here, we first review the basics of the dark matter problem (Section~\\ref{sec:missing-mass}) as well as the basic ingredients of the present-day concordance model of cosmology (Section~\\ref{sec:LCDM}). We then point out a few outstanding challenges for this model (Section~\\ref{sec:challenges}), both from the point of view of unobserved predictions of the model, and from the point of view of unpredicted observations (all uncannily involving a common acceleration constant $a_0$). Up to that point, the challenges presented are purely based on observations, and are fully independent of any alternative theoretical framework\\epubtkFootnote{The first four sections provide the observational evidence for the MOND phenomenology through the different appearances of $a_0$ in galactic dynamics, but they are actually independent of any specific theory, while the reader more specifically interested into MOND \\textit{per se} could go directly to Section~\\ref{sec:Kepler}.}. We then show that, surprisingly, many of these puzzling observations can be summarized within one single empirical law, Milgrom's law (Section~\\ref{sec:Kepler}), which can be most easily (although not necessarily uniquely) interpreted as the effect of a single universal force law resulting from a modification of Newtonian dynamics (MOND) in the weak-acceleration regime $a < a_0$, for which we present the current observational successes and problems (Section~\\ref{sec:MOND}). We then summarize the various attempts currently made to embed this modification in a generally covariant relativistic theory of gravity (Section~\\ref{sec:covariant}) and how such theories allow new predictions on gravitational lensing (Section~\\ref{sec:lensing}) and cosmology (Section~\\ref{sec:cosmology}). We finally draw conclusions in Section~\\ref{sec:summary}.   \\newpage ",
        "conclusions": "\\label{sec:summary}  In this review, after briefly presenting the currently favored $\\Lambda$CDM model of cosmology (which clearly works overwhelmingly well on large scales despite its slightly unelegant mixture of currently unknown elements, Sects.~2 and 3), we reviewed the few most outstanding challenges that this model is still facing (Sect.~4), which will have to be addressed one way or the other in the coming years. These include coincidences at $z=0$ between the scale of the energy density in dark energy, dark matter, and baryonic matter, as well as a common natural scale for the behavior of the dark matter and dark energy sectors. What is more, as far as galaxy formation is concerned, many predictions made by the model (keeping in mind that baryon physics could modify these predictions) were ruled out by observations: these include many observations indicating that structure formation should take place earlier than predicted, the low number of observed satellites around the Milky Way (especially the missing satellites at the low and high mass ends of the mass function), the phase-space correlation of satellite galaxies of the Milky Way as opposed to their predicted isotropic distribution, the apparent presence of constant DM density cores in the central parts of galaxies instead of the predicted cuspy dark halos, the over-abundance of large bulgeless thin disk galaxies that are extremely difficult to produce in simulations, or the presence of spiral arms in disks that should be immune to such instabilities. But even more challenging is the appearance (Figure~\\ref{btflongbaseline}) of an acceleration constant $a_0 \\simeq 10^{-10}\\mathrm{\\ m\\, s}^{-2}$ (i.e., the common scale of the dark matter and dark energy sectors as $a_0 \\sim \\Lambda^{1/2}$ in natural units) in many \\textit{a priori} unrelated scaling relations for DM and baryons in galaxies. These scaling relations involve a possibly devastating amount of fine-tuning for all collisionless dark matter models (Sect.~4.3), and can all be summarized by Milgrom's empirical formula (Sect.~5), meaning that the observed gravitational field in galaxies is mimicking a universal force law generated by the baryons alone.  \\epubtkImage{btfdeviationslongbaseline.png}{% \\begin{figure}[htbp] \\centerline{\\includegraphics[width=\\textwidth]{btfdeviationslongbaseline}} \\caption{The acceleration parameter $\\sim V_f^4/(GM_b)$ of extragalactic systems, spanning ten decades in baryonic mass $M_b$. X-ray emitting galaxy groups and clusters are visibly offset from smaller systems, but by a remarkably modest amount over such a long baseline. The characteristic acceleration scale $a_0 \\sim \\sqrt{\\Lambda}$ is in the data, irrespective of the interpretation. And it actually plays various other independent roles in observed galaxy phenomenology. This is natural in MOND (see Sect.~5.2), but not in $\\Lambda$CDM (see Sect.~4.3).} \\label{btflongbaseline} \\end{figure}}  With inert, collisionless and dissipationless DM, making Milgrom's law emerge requires a huge, and perhaps even unreasonable, amount of fine-tuning in the expected feedback from the baryons. Indeed, the relation between the distribution of baryons and DM should \\textit{a priori} depend on the various different histories of formation, intrinsic evolution, and interaction with the environment of the various different galaxies, whereas Milgrom's law provides a sucessful unique and \\textit{history-independent} relation. Given this puzzle, the central idea of Modified Newtonian Dynamics (MOND) is to rather explore the possibility that the force law is indeed effectively modified (Sect.~6). The main motivation for studying MOND is thus a fully \\textit{empiricist} one, as it is driven by the \\textit{observed} phenomenology on galaxy scales, and not by an aesthetic wish of getting rid of DM. The corollary is that it is \\textit{a priori} not a problem for a theory designed to reproduce the uncanny successes of the MOND phenomenology to replace CDM by ``dark fields'' (see Sect.~7) or more exotic forms of DM, different from simple collisionless DM particles, contrary to the common belief that this would be against the spirit of the MOND paradigm (although it is true that it would be more elegant to avoid too many additional degrees of freedom). It is perhaps more important that, if MOND is correct in the sense of the acceleration $a_0$ being a \\textit{truly} fundamental quantity, the strong equivalence principle \\textit{cannot} hold anymore, and local Lorentz invariance could perhaps be spontaneously violated too.  At this juncture, it is worthwhile to summarize the general predictions of MOND, as a paradigm, and their observational tests (Table~\\ref{table:mondtests}). As a mathematical description of the effective force law, MOND works remarkably well in individual galaxies.  As a modified gravity theory (at the classical level), it makes some predictions that are both unique and challenging to reproduce in the context of the $\\Lambda$CDM paradigm.  However, MOND faces sharp challenges, particularly with cosmology and in rich clusters of galaxies, which will not be conclusively addressed without a viable parent theory (Sect.~7), based on first principles and underlying the MOND paradigm (if such a theory exists at all). In any case, in his series of papers introducing the idea in 1983, Milgrom~\\cite{Milgrom1} made a few very \\textit{explicit} predictions, which we quote hereafter, and compare with modern observational data (see also the Kepler-like laws of galactic dynamics in Sect.~5.2):  \\begin{itemize} \\item \\textit{``Velocity curves calculated with the modified dynamics on the basis of the observed mass in galaxies should agree with the observed curves.''}  It is now well established that MOND provides good fits to the rotation curves of galaxies (Figure~\\ref{figure:RCresiduals}~\\cite{SM02,things}), including bumps and wiggles associated to a baryonic counterpart (Figure~\\ref{figure:RCexamples}, Kepler-like law n$^\\circ$10 in Sect.~5.2).  These fits are obtained with a single free parameter per galaxy, the mass-to-light ratio of the stars. What makes them most impressive is that the best-fit mass-to-light ratios, obtained on purely dynamical grounds assuming MOND, vary with galaxy color exactly as one would expect from stellar population synthesis models~\\cite{Bel03}, that are based on astronomers' detailed understanding of stars. Note that the rotation curves of galaxies are predicted to be \\textit{asymptotically} flat, even though this flatness is not always attained at the last observed point (see Kepler-like law n$^\\circ$1 in Sect.~5.2, and last explicit prediction hereafter).  \\item \\textit{``The relation between the asymptotic velocity and the mass of the galaxy is an absolute one.''}  This is the Baryonic Tully--Fisher relation with $M_b = a_0 G V_f^4$ (see Kepler-like laws n$^\\circ$2 in Sect.~5.2).  It appears to hold quite generally~\\cite{M05}, even for galaxies that we would conventionally expect to deviate from it~\\cite{GentileTDG,MdB98,McGgasrich}.  \\item \\textit{``Analysis of the $z$-dynamics in disk galaxies using the modified dynamics should yield surface densities which agree with the observed ones.''}  This states that in addition to the radial force giving the rotation curve, the motions of stars perpendicular to the disk must also follow from the source baryons (see Sect.~6.5.3).  This proves to be a remarkably challenging observation, and such data for external galaxies are dear to obtain~\\cite{bershady}.  To make matters still more difficult, the radial acceleration usually dominates the vertical ($V^2/r \\gg \\sigma_z^2/z$).  This has the consequence that the distinction between MOND and conventional dynamics is not pronounced in regions that are well observed, becoming pronounced only at rather low baryonic surface densities~\\cite{MdB98}.  The vertical velocity dispersions in low surface density regions (see Sect.~6.5.3) is typically $\\sim 8$~km/s~\\cite{BJBB,KdN08}.  This exceeds the nominal Newtonian expectation (typically $\\sim 2$~km/s for $\\Sigma = 1\\,M_{\\odot}\\,\\mathrm{pc}^{-2}$, depending on the thickness of the disk), and is more in accordance with MOND. However, it would require a considerably more detailed analysis to consider this a test, let alone a success, of MOND. The Milky Way (Sect.~6.5.2) may provide an excellent test for this prediction~\\cite{Bienayme,Salcedo} as more precision data become available.  \\item \\textit{``Effects of the modification are predicted to be particularly strong in (low surface brightness) dwarf galaxies.''}  The dwarf spheroidal satellite galaxies of the Milky Way have very low surface densities of stars, so (see Kepler-like law n$^\\circ$8 in Sect.~5.2) are far into the MOND regime. As expected, these systems exhibit large mass discrepancies~\\cite{Walker09,SG}.  Detailed fits to the better observed ``classical'' dwarfs~\\cite{AngusdSph} are satisfactory in most cases (see Sect.~6.6.2). The so-called ``ultrafaint'' dwarfs appear more problematic~\\cite{MWolf}, in the sense that their velocity dispersions are higher than expected.  This might be an indication of the MOND-specific external field effect (see Sect.~6.3 and~\\cite{BMdwarfs}), as the field of the Milky Way dominates the internal fields of the ultrafaint dwarfs.  If so, these objects are not in dynamical equilibrium, which considerably complicates their analysis.  \\item \\textit{Locally measured mass-to-light ratios should show no indication of hidden mass when $V_c^2/R \\gg a_0$, but rise beyond the radius where $V_c^2/R \\approx a_0$.}  We have paraphrased this prediction for brevity (see also  Kepler-like law n$^\\circ$7 in Sect.~5.2). The test of this prediction is shown in Figures~\\ref{figure:MDRA}, \\ref{fig:mdasolsys}, and \\ref{figure:MLRSd}. The predicted effect is obvious in the data with populations synthesis mass-to-light ratios for the stars~\\cite{Bel03}, or with dynamical mass-to-light ratios~\\cite{MdB98} that make no assumption about stellar mass.  In HSB spirals, there is no obvious need for dark matter in the inner regions, with the mass discrepancy only appearing at large radii as the acceleration drops below $a_0$ (Figure~\\ref{figure:MDRA}).  \\item \\textit{``Disk galaxies with low surface brightness provide particularly strong tests.''}  Low surface brightness means low stellar surface density, which in turns means low acceleration. LSB galaxies are thus predicted to be well into the modified regime (see also  Kepler-like law n$^\\circ$8 in Sect.~5.2). This was a strong \\textit{a priori} prediction, because few bona-fide examples of such objects were known at the time. Indeed, in 1983, when these predictions were published, it was widely thought that nearly all disk galaxies shared a common high surface brightness. One specific consequence of MOND for LSB galaxies is that they should lie on the same BTFR, with the same normalization, as high surface brightness spirals.  This was subsequently observed to be the case~\\cite{ZwaanTF,sprayTF}. There is \\textit{no} systematic deviation from the BTFR with surface brightness (Figure~\\ref{figure:btfresiduals}), thus contrary to what is naturally expected in conventional dynamics~\\cite{MdB98,CRix}. Another consequence of low surface density is that the acceleration is low ($< a_0$) everywhere. As a result, the mass discrepancy appears at a smaller radius in low surface brightness galaxies, and is larger in amplitude than in high surface brightness galaxies.  This effect was subsequently observed (Figure~\\ref{figure:MLRSd}~\\cite{MdB98}).  \\item \\textit{``We predict a correlation between the value of the average surface density of a galaxy and the steepness with which the rotational velocity rises to its asymptotic value.''}  MOND does \\textit{not} simply make rotation curves flat.  It predicts that high surface brightness galaxies have rotation curves that rise rapidly before becoming flat, and may even fall towards asymptotic flatness.  In contrast, low surface brightness galaxies should have slowly rising rotation curves that only gradually approach asymptotic flatness (see also  Kepler-like law n$^\\circ$8 in Sect.~5.2).  Both morphologies are observed (Figure~\\ref{figure:RCshape}).  The expected connection between dynamical acceleration and the surface density of the source baryons is illustrated in Figures~\\ref{figure:XiSd} and \\ref{figure:ARSd}.  \\end{itemize}  The original predictions listed above cover many situations, but not all.  Indeed, once one writes a specific force law, its application must be completely general.  Such a hypothesis is readily subject to falsification, provided sufficiently accurate data to test it -- a perpetual challenge for astronomy. Table~\\ref{table:mondtests} summarizes the tests discussed here.  By and large, tests of MOND involving rotationally supported disk galaxies are quite positive, as largely detailed above (see Sect.~6.5).  By construction, there is no cusp problem (solution to the challenge n$^\\circ$6 of Sect.~4.2), and no missing baryons problem (solution to the challenge n$^\\circ$10 of Sect.~4.2), as the way the dynamical mass-to-light ratio systematically varies with the circular velocity is a direct consequence of Milgrom's law (Kepler-like law n$^\\circ$4 of Sect.~5.2). There does appear to be a relation between the quality of the data and the ease with which a MOND fit to the rotation curve is obtained, in the sense that fits are most readily obtained with the best data~\\cite{BBS}.  As the quality of the data decline~\\cite{sanders96}, one begins to notice small disparities.  These are sometimes attributable to external disturbances that invalidate the assumption of equilibrium~\\cite{SV98}.  For targets that are intrinsically difficult to observe, minor problems become more common~\\cite{dBM98,swatersmond}.  These typically have to do with the challenges inherent in combining disparate astronomical data sets (e.g., rotation curves measured independently at optical and radio wavelengths) and constraining the inclinations of low surface brightness galaxies (bear in mind that all velocities require a $\\sin(i)$ correction to project the observed velocity into the plane of the disk, and mass in MOND scales as the fourth power of velocity).  Given the intrinsic difficulties of astronomical observations, it is remarkable that the success rate of MOND fits is as high as it is:  of the 78 galaxies that have been studied in detail (see Sect.~6.5.1), only a few cases (most notably NGC~3198~\\cite{bottema,things}) appear to pose challenges.  Given the predictive and quantitative success of the majority of the fits, it would seem unwise to ignore the forest and focus only on the outlying trees.  One rotationally supported system that is very familiar to us is the Solar system (see Sect.~6.4).  The Solar system is many orders of magnitude removed from the MOND regime (Figure~\\ref{fig:mdasolsys}), so no strong effects are predicted.  However, it is of course possible to obtain exquisitely precise data in the Solar system, so it is conceivable that some subtle effect may be observable~\\cite{Sanderssun}.  Indeed, the lack of such effects on the inner planets already appears to exclude some slowly varying interpolation functions~\\cite{BlanchetNovak}. Other tests may yet prove possible~\\cite{Mag1,Milgromsun}, but, as they are strong-field gravity tests by nature, they all depend strongly on the parent relativistic theory (Sect.~7) and how it converges towards GR~\\cite{galileon}. So in Table~\\ref{table:mondtests}, we list the status of Solar system tests as unclear, depending on the parent relativistic theory.  An important aspect of galactic disks is their stability (see Sect.~6.5.3).  Indeed, the need to stabilize disks was one of the early motivations for invoking dark matter~\\cite{OstPeeb}.  MOND appears able to provide the requisite stability~\\cite{BradaM}.  Indeed, it gives good reason~\\cite{Milstab} for the observed maximum in the distribution of disk galaxy surface densities at $\\sim \\Sigma_{\\dagger} = a_0/G$ (Freeman's limit: Figure~\\ref{fig:freeman} and  Kepler-like law n$^\\circ$6 in Sect.~5.2).  Disks with surface densities below this threshold are in the low acceleration limit and can be stabilized by MOND.  Higher surface density disks would be purely in the Newtonian regime and subject to the usual instabilities. Going beyond the amount of stability required for existence, another positive aspect of MOND is that it does not over-stabilize disks. Features like bars and spiral arms are a natural result of disk self-gravity.  Conventionally, large halo-to-disk mass ratios suppress the growth of such features, especially in low surface brightness galaxies~\\cite{Mihos}.  Yet such features are present\\epubtkFootnote{\\cite{MdB98} utilized this fact to predict that conventional analyses of low surface brightness disks would infer abnormally high mass-to-light ratios for their stellar populations -- a prediction that was subsequently confirmed~\\cite{fuchs,Saburova}.}.  The suppression is not as great in MOND~\\cite{BradaM}, and numerical simulations appear to do a good job of reproducing the range of observed morphologies of spiral galaxies (solution to the challenge n$^\\circ$9 of Sect.~4.2, see~\\cite{Tiretgas}). Bars tend to appear more quickly and are fast, while warps can also be naturally produced (Sect.~6.5.3). There appears to be no reason why this should not extend to thin and bulgeless disks, whose ubiquity poses a challenge to galaxy formation models in $\\Lambda$CDM.  This particular point of creating large bulgeless disks (challenge n$^\\circ$8 of Sect.~4.2) can actually be solved thanks to early structure formation followed by a low galaxy-interaction rate in MONDian cosmology (see Sect.~9.2), but this definitely warrants further investigation, so we mark this case as merely promising in Table~\\ref{table:mondtests}.  Interacting galaxies are by definition non-stationary systems in which the customary assumption of equilibrium does not generally hold.  This renders direct tests of MOND difficult.  However, it is worth investigating whether commonly observed morphologies (e.g., tidal tails) are even possible in MOND.  Initially, this seemed to pose a fundamental difficulty~\\cite{MdB98}, as dark matter halos play a critical role in absorbing the orbital energy and angular momentum that it is necessary to shed if passing galaxies are to not only collide, but stick and merge. Nevertheless, recent numerical simulations appear to do a nice job of reproducing observed morphologies~\\cite{Tiretinteract}.  This is no trivial feat. While it is well established that dark matter models can result in nice tidal tails, it turns out to be difficult to simultaneously match the narrow morphology of many observed tidal tails with rotation curves of the systems from which they come~\\cite{dubinski}. Narrow tidal tails appear to be natural in MOND, as well as more extended resulting galaxies, thanks to the absence of angular momentum transfer to the dark halo (solution to the challenge n$^\\circ$7 of Sect.~4.2).  Additionally, tidal dwarfs that form in these tails clearly have characteristics closer to those observed (see Sect.~6.5.4) than those from dark matter simulations~\\cite{GentileTDG,Miltidal}.  Spheroidal systems also provide tests of MOND (Sect.~6.6).  Unlike the case of disk galaxies, where orbits are coplanar and nearly circular so that the centripetal acceleration can be equated with the gravitational force, the orbits in spheroidal systems are generally eccentric and randomly oriented.  This introduces an unknown geometrical factor usually subsumed into a parameter that characterizes the anisotropy of the orbits.  Accepting this, MOND appears to perform well in the classical dwarf spheroidal galaxies, but implies that the ultrafaint dwarfs are out of equilibrium (see Sect.~6.6.2).  For small systems like the ultrafaint dwarfs and star clusters (Sect.~6.6.3) within the Milky Way, the external field effect (Sect.~6.3) can be quite important.  This means that star clusters generally exhibit Newtonian behavior by virtue of being embedded in the larger Galaxy.  Deviations from purely Newtonian behavior are predicted to be subtle and are fodder for considerable debate~\\cite{Rodrigo2419,Sanders2419}, rendering the present status unclear (Table~\\ref{table:mondtests}).  At the opposite extreme of giant elliptical galaxies (Sect.~6.6.1), the data accord well with MOND~\\cite{dearth}.  Indeed, bright elliptical galaxies are sufficiently dense that their inner regions are well into the Newtonian regime.  In the MONDian context, this is the reason that it has historically been difficult to find clear evidence for mass discrepancies in these systems.  The apparent need for dark matter does not occur until radii where the accelerations become low.  That only spheroidal stellar systems appear to exist at surface densities in excess of $\\Sigma_{\\dagger}$ is the corollary of Freeman's limit:  such dense systems could not exist as stable disks, so must per force become elliptical galaxies, regardless of the formation mechanism that made them so dense.  That populations of elliptical galaxies should obey the Faber-Jackson relation (Kepler-like law n$^\\circ$3 in Sect.~5.2, Figure~\\ref{figure:faberjackson}) is also very natural to MOND~\\cite{Sanders94,sanders10}.  The largest gravitationally bound systems are also spheroidal systems: rich clusters of galaxies.  The situation here is quite problematic for MOND (Sect.~6.6.4).  Applying MOND to ascertain the dynamical mass routinely exceeds the observed baryonic mass by a factor of 2 to 3.  In effect, MOND requires additional \\textit{dark matter} in galaxy clusters. The need to invoke unseen mass is most unpleasant for a theory that otherwise appears to be a viable alternative to the existence of unseen mass.  However, one should remember that the present-day motivation for studying MOND is driven by the observed phenomenology on galaxy scales, summarized above, and not by an aesthetic wish of getting rid of DM. What is more, parent relativistic theories of MOND might well involve additional degrees of freedom in the form of ``dark fields''. But in any case, one must be careful not to conflate the rather limited missing mass problem that MOND suffers in clusters with \\textit{the} non-baryonic collisionless cold dark matter required by cosmology.  There is really nothing about the cluster data that requires the excess mass to be non-baryonic, as long as it behaves in a collisionless way.  There could for instance be baryonic mass in some compact non-luminous form (see Sect.~6.6.4 for an extensive discussion). This might seem to us unlikely, but it does have historical precedent.  When Zwicky~\\cite{zwicky} first identified the dark matter problem in clusters, the mass discrepancy was of order $\\sim 100$.  That is, unseen mass outweighed the visible stars by two orders of magnitude.  It was only decades later that it was recognized that baryons residing in a hot intracluster gas greatly outweighed those in stars.  In effect, there were [at least] \\textit{two} missing mass problems in clusters.  One was the hot gas, which reduces the conventional discrepancy from a factor of $\\sim 100$ to a factor of $\\sim 8$~\\cite{giodini} in Newtonian gravity.  From this perspective, the remaining factor of two in MOND seems modest.  Rich clusters of galaxies are rare objects, so the total required mass density can readily be accommodated within the baryon budget of big bang nucleosynthesis.  Indeed, according to BBN, there must still be a lot of unidentified baryons lurking somewhere in the Universe. But the excess dark mass in clusters need not be baryonic, even in MOND.  Massive ordinary neutrinos~\\cite{sandersclusters,Sandersneut} and light sterile neutrinos~\\cite{angussterilenu,AFD} have been suggested as possible forms of dark matter that might provide an explanation for the missing mass in clusters.  Both are non-baryonic, but as they are \\textit{hot} DM particle candidates, neither can constitute \\textit{the} cosmological non-baryonic \\textit{cold} dark matter. At this juncture, all we can say for certain is that we do not know what the composition of the unseen mass is. It could even just be an evidence for the effect of additional ``dark fields'' in the parent relativistic formulation of MOND, such as massive scalar fields, vector fields, dipolar dark matter, or even subtle non-local effects (see Sect.~7).  There are other aspects of cluster observations that are more in line with MOND's predictions.  Clusters obey a mass--temperature relation that parallels the $M \\propto T^2 \\propto V^4$ prediction of MOND (Figures~\\ref{figure:clustersbtf} and \\ref{btflongbaseline}) more closely than the conventional prediction of $M \\propto T^{3/2}$ expectation in $\\Lambda$CDM, without the need to invoke preheating (a need that may arise as an artifact of the mismatch in slopes).  Indeed, Figure~\\ref{btflongbaseline} shows clearly both the failing of MOND in the offset in characteristic acceleration between clusters and lower mass systems, and its successful prediction of the slope (a horizontal line in this figure).  A further test which may be important is the peculiar and bulk velocity of clusters.  For example, the collision velocity of the bullet cluster is so large\\epubtkFootnote{The observed shock velocity of $\\sim 4700$~km/s is thought to be enhanced by hydrodynamical effects.  The collision velocity is improbable \\textit{after} a substantial ($\\sim 1700$~km/s) correction for this~\\cite{Mastropietro}.} as to be highly improbable in $\\Lambda$CDM (occurring with a probability of $\\sim 10^{-10}$~\\cite{Lee}). In contrast, large collision velocities are natural to MOND~\\cite{Angusbullet}.  Similarly, the large scale peculiar velocity of clusters is observed to be $\\sim 1000$~km/s~\\cite{Kash}, well in excess of the expected $\\sim 200\\;\\mathrm{km}\\,\\mathrm{s}^{-2}$.  Ongoing simulations with MOND~\\cite{anguscosmo} show some promise to produce large peculiar velocities for clusters.  In general, one would expect high speed collisions to be more ubiquitous in MOND than $\\Lambda$CDM.  An important line of evidence for mass discrepancies in the Universe is gravitational lensing in excess of that expected from the observed mass of lens systems.  Lensing is an intrinsically relativistic effect that requires a generally covariant theory to properly address.  This necessarily goes beyond MOND itself into specific hypotheses for its parent theory (Sect.~7), so is somewhat different than the tests discussed above. Broadly speaking, tests involving strong gravitational lensing fare tolerable well (Sect.~8.1), whereas weak lensing tests, that are sensitive to larger scale mass distributions, are more problematic (Sects.~8.2, 8.3, and 8.4) or simply crash into the usual missing mass problem of MOND in clusters. Note that weak lensing in relativistic MOND theories produces the same amount of lensing as required from dynamics, so this is not the problem. The problematic fact is just that some tests seem to require more dark matter than the effect of MOND provides.  On larger (cosmological) scales, MOND, as a modification of classical (non-covariant) dynamics, is simply unsatisfactory or mute.  MOND itself has no cosmology, providing analogs for neither the Friedmann equation for the dynamics of the Universe, nor the Robertson-Walker metric for its geometry. For these, one must appeal to specific hypotheses for the relativistic parent theory of MOND (Sect.~7), which is far from unique, and theoretically not really satisfactory, as none of the present candidates emerges from first principles.  At this juncture, it is not clear whether a compelling candidate cosmology will ever emerge. But on the other hand, there is \\textit{nothing} about MOND as a paradigm that contradicts \\textit{per se} the empirical pillars of the hot big bang:  Hubble expansion, big bang nucleosynthesis, and the relic radiation field (Sect.~9). The formation of large scale structure is one of the strengths of conventional theory, which can be approached with linear perturbation theory.  This leads to good fits of the power spectrum both at early times ($z \\approx 1000$ in the cosmic microwave background) and at late times (the $z = 0$ galaxy power spectrum~\\cite{Tegmark04}).  In contrast, the formation of structure in MOND is intrinsically non-linear.  It is therefore unclear whether MOND-motivated relativistic theories will inevitably match the observed galaxy power spectrum, a possible problem being how to damp the baryon acoustic oscillations~\\cite{Dodprob,SMFB}. At this stage, a unique prediction does not exist. Nevertheless, there are two aspects of structure formation in MOND that appear to be fairly generic and distinct from $\\Lambda$CDM. The stronger effective long range force in MOND speeds the growth rate, but has less mass to operate with as a source. Consequently, radiation domination persists longer and structure formation is initially inhibited (at redshifts of hundreds). Once structure begins to form, the non-linearity of MOND causes it to proceed more rapidly than in GR with CDM.  Three observable consequences would be (i) the earlier emergence of large objects like galaxies and clusters in the cosmic web (as well as the associated low interaction rate at smaller redshifts) providing a possible solution to challenge n$^\\circ$2 of Sect.~4.2~\\cite{anguscosmo}, (ii) the more efficient evacuation of large voids (possible solution to challenge n$^\\circ$3 of Sect.~4.2), and (iii) larger peculiar (and collisional~\\cite{Angusbullet}) velocities of galaxy clusters (solution to challenge n$^\\circ$1 of Sect.~4.2). The potential downside to rapid structure formation in MOND, however, is that it may overproduce structure by redshift zero~\\cite{nusser,Llinares}.  The final entries in Table~\\ref{table:mondtests} regard the cosmic microwave background, discussed in more detail in Sect.~\\ref{subsection:CMB}.  The third peak of the acoustic power spectrum of the CMB poses perhaps the most severe challenge to a MONDian interpretation of cosmology.  The amplitude of the third peak measured by WMAP is larger than expected in a Universe composed solely of baryons~\\cite{WMAP3}.  This implies some substance that does not oscillate with the baryons. Cold dark matter fits this bill nicely.  In the context of MOND, we must invoke some other massive substance (i.e., non-baryonic dark matter such as, e.g., light sterile neutrinos~\\cite{angussterilenu}) that plays the role of CDM, or rely on additional degrees of freedom in the relativistic parent theory of MOND (see Sect.~7) that would have the same net result (see the extensive discussion in Sect.~9.2), or even combine non-baryonic dark matter with these additional degrees of freedom~\\cite{SMFB}.  While these are real possibilities, neither are \\textit{a priori} particularly appealing, any more than it is to invoke CDM with complex fine-tuned feedback to explain rotation curves that apparently require only baryons as a source.  The missing baryon problem that MOND suffers in rich clusters of galaxies and the third peak of the acoustic power spectrum of the CMB are thus the most serious challenges presently facing MOND.  But even so, the interpretation of the acoustic power spectrum is not entirely clear cut.  Though there is no detailed fit to the power spectrum in MOND (unless we invoke $10 \\,$eV-scale sterile neutrinos~\\cite{angussterilenu}), MOND did motivate the \\textit{a priori} prediction~\\cite{McGCMB1} of two aspects of the CMB that were surprising in $\\Lambda$CDM (see Sect.~9.2). The amplitude ratio of the first-to-second peak in the acoustic power spectrum was outside the bounds expected ahead of time by $\\Lambda$CDM for $\\Omega_b$ from big bang nucleosynthesis as it was then known (see Sect.~\\ref{subsection:CMB}). In contrast, the first:second acoustic peak ratio that is now well measured agrees well with the quantitative value predicted in advance for the case of the absence of cold dark matter~\\cite{McGCMB2,McGCMB3}.  Similarly, the rapid formation of structure expected in MOND leads naturally to an earlier epoch of re-ionization than had been anticipated in $\\Lambda$CDM~\\cite{McGCMB1,McGCMB3}.  Thus, while the amplitude of the third peak is clearly problematic and poses a severe challenge for any MOND-inspired theories, the overall interpretation of the CMB is debatable.  While the existence of non-baryonic cold dark matter is \\textit{a priori} the most obvious explanation of the third peak indeed, it is not at all obvious that straightforward CDM -- in the form of rather simple massive inert collisionless particles -- is uniquely required.  Science is in principle about theories or models that are falsifiable, and thus that are presently either falsified or not. But in practice it does not (and cannot) really work that way: if a model that was making good predictions up to a certain point suddenly does not work anymore (i.e., does not fit some new data), one obviously first tries to adjust it to make it fit the observations rather than throwing it away immediately. This is what one calls the requisite ``compensatory adjustments\" of the theory (or of the model): Popper himself drew attention to these limitations of falsification in \\textit{The Logic of Scientific Discovery}~\\cite{Popper}. In the case of the $\\Lambda$CDM model of cosmology, which is mostly valid on large scales, the current main trend is to find the ``compensatory adjustments'' to the model to make it fit in galaxies, mainly by changing (or mixing) the mass(es) of the dark matter particles, and/or through artificially fine-tuned baryonic feedback in order to reproduce the success of MOND. Incidentally, exactly the same is true for MOND, but for the opposite scales: MOND works remarkably well in galaxies but apparently needs compensatory adjustments on larger scales to effectively replace CDM. Now does that mean that falsification is impossible? That all models are equal? Surely not. In the end, a theory or a model is really falsified once there are too many compensatory adjustments (needed in order to fit too many discrepant data), or once these become too twisted (like Tycho Brahe's geocentric model for the Solar system). But there is  obviously no truly quantitative way of ascertaining such global falsification. How one chooses to weigh the evidence presented in this review necessarily informs one's opinion of the relative merits of $\\Lambda$CDM and MOND.  If one is most familiar with cosmology and large scale structure, $\\Lambda$CDM is the obvious choice, and it must seem rather odd that anyone would consider an alternative as peculiar as MOND, needing rather bizarre adjustments to match observations on large scales. But if one is more concerned with precision dynamics and the observed phenomenology in a wide swath of galaxy data, it seems just as strange to invoke non-baryonic cold dark matter together with fine-tuned feedback to explain the appearance of a single effective force law that appears to act with only the observed baryons as a source. Perhaps the most important aspect before one throws away any model is to have a ``simpler'' model at hand, that still reproduces the successes of the earlier favored model but also naturally explains the discrepant data. In that sense, right now, it is absolutely fair to say that there is \\textit{no} alternative which really does better \\textit{overall} than $\\Lambda$CDM, and in favor of which Ockham's razor would be. It would however probably be a mistake to persistently ignore the fine-tuning problems for dark matter and the related uncanny successes of the MOND paradigm on galaxy scales, as they could very plausibly point at a hypothetical better new theory. It is also important to bear in mind that MOND, as a paradigm or as a modification of Newtonian dynamics, is not itself generally covariant. Attempts to construct relativistic theories that contain MOND in the appropriate limit (Sect.~7) are correlated but distinct efforts, and one must be careful not to conflate the two. For example, some theories, like TeVeS (Sect.~7.4), might make predictions that are distinct from GR in the strong-field regime.  Should future tests falsify these distinctive predictions of TeVeS while confirming those of GR, this would perhaps falsify TeVeS as a viable parent theory for MOND, but would have no bearing on the MONDian phenomenology observed in the weak-field regime, nor indeed on the viability of MOND itself. It would perhaps simply indicate the need to continue to search for a deeper theory. It would for instance be extremely alluring if one would manage to find a physical connection between the dark energy sector and the possible breakdown of standard dynamics in the weak-field limit, since both phenomena would then simply reflect discrepancies with  the predictions of GR when $\\Lambda \\sim a_0^2$ is set to zero (see, e.g., Sect.~7.10). It is of course perfectly conceivable that such a deep theory does not exist, and that the apparent MONDian behavior of galaxies will be explained through small compensatory adjustments of the current $\\Lambda$CDM paradigm, but one has yet to demonstrate how this will occur, and it will inevitably involve a substantial amount of fine-tuning that will have to be explained naturally. In any case, the existence of a characteristic acceleration $a_0$ (Figure~\\ref{btflongbaseline}) playing various different roles in  many seemingly independent galactic scaling relations (see Sects.~4.3 and 5.2) is by now an empirically established fact, and it is thus mandatory for \\textit{any} successful model of galaxy formation and evolution to explain it. The future of this field of research might thus still be full of exciting surprises for astronomers, cosmologists, and theoretical physicists.  \\begin{table}[htbp] \\caption{Observational tests of MOND.} \\label{table:mondtests} \\centering {\\small \\begin{tabular}{lcccc} \\toprule Observational Test & Successful & Promising & Unclear & Problematic \\\\ \\midrule \\textbf{Rotating Systems} & & & & \\\\ solar system & & & X & \\\\ galaxy rotation curve shapes & X & & & \\\\ surface brightness $\\propto \\Sigma \\propto a^2$  & X & & & \\\\ galaxy rotation curve fits & X & & & \\\\ fitted M$_*$/L & X & & & \\\\ \\midrule \\textbf{Tully--Fisher Relation} & & & & \\\\ baryon based & X & & & \\\\ slope & X & & & \\\\ normalization & X & & & \\\\ no size nor $\\Sigma$ dependence & X & & & \\\\ no intrinsic scatter & X & & & \\\\ \\midrule \\textbf{Galaxy Disk Stability} & & & & \\\\ maximum surface density  & X & & & \\\\ spiral structure in LSBGs   & X & & & \\\\ thin \\& bulgeless disks & & X & & \\\\ \\midrule \\textbf{Interacting Galaxies} & & & & \\\\ tidal tail morphology  & & X & & \\\\ dynamical friction & & & X & \\\\ tidal dwarfs & X & & & \\\\ \\midrule \\textbf{Spheroidal Systems} & & & & \\\\ star clusters & & & X & \\\\ ultrafaint dwarfs & & & X & \\\\ dwarf Spheroidals & X & & & \\\\ ellipticals & X & & & \\\\ Faber-Jackson relation & X & & & \\\\ \\midrule \\textbf{Clusters of Galaxies} & & & & \\\\ dynamical mass  & & & & X \\\\ mass--temperature slope & X & & & \\\\ velocity (bulk \\& collisional)  & & X & & \\\\ \\midrule \\textbf{Gravitational Lensing} & & & & \\\\ strong lensing & X & & &  \\\\ weak lensing (clusters \\& LSS) & & & X &  \\\\ \\midrule \\textbf{Cosmology} & & & & \\\\ expansion history & & & X & \\\\ geometry  & & & X & \\\\ big bang nucleosynthesis & X & & & \\\\ \\midrule \\textbf{Structure Formation} & & & & \\\\ galaxy power spectrum & & & X & \\\\ empty voids & & X & & \\\\ early structure & & X & & \\\\ \\midrule \\textbf{Background Radiation} & & & & \\\\ first:second acoustic peak & X & & & \\\\ second:third acoustic peak  & & & & X \\\\ detailed fit  & & & & X \\\\ early re-ionization & X & & & \\\\ \\bottomrule \\end{tabular} } \\end{table}   \\newpage"
    },
    "1112/1112.3367_arXiv.txt": {
        "abstract": "We demonstrate a novel technology that combines the power of the multi-object spectrograph with the spatial multiplex advantage of an integral field spectrograph (IFS). The Sydney-AAO Multi-object IFS (SAMI) is a prototype wide-field system at the Anglo-Australian Telescope (AAT) that allows 13 imaging fibre bundles (``hexabundles'') to be deployed over a 1--degree diameter field of view. Each hexabundle comprises 61 lightly--fused  multimode fibres with reduced cladding and yields a 75 percent filling factor.  Each fibre core diameter subtends 1.6 arcseconds on the sky and each hexabundle has a field of view of 15 arcseconds diameter. The fibres are fed to the flexible AAOmega double--beam spectrograph, which can be used at a range of spectral resolutions ($R=\\lambda/\\delta\\lambda\\approx$ 1700--13000) over the optical spectrum (3700--9500\\,\\AA). We present the first spectroscopic results obtained with SAMI for a sample of galaxies at $z \\approx 0.05$. We discuss the prospects of implementing hexabundles at a much higher multiplex over wider fields of view in order to carry out spatially--resolved spectroscopic surveys of $10^4-10^5$ galaxies. ",
        "introduction": "\\label{sec:intro}  Galaxies are intrinsically complex with multiple components and varied formation histories. This complexity is the primary reason that unravelling the physics of galaxy formation and evolution is so challenging. Galaxies are made up of baryons confined to dark matter haloes, and often have multiple distinct kinematic components (e.g.\\ bulge and/or disc). There are complex interactions between the stars, gas, dust, dark matter and super-massive black holes. These can lead to both positive and negative feedback on the formation rate of stars.  Experimental efforts to address galaxy formation have generally taken two directions. First, galaxy imaging and spectroscopic surveys have progressively moved to higher redshift, in an attempt to directly observe galaxy evolution and formation. This approach has had much success, placing strong constraints on the evolution of the global star formation rate \\citep[e.g.][]{2006ApJ...651..142H}, unveiling the strong evolution of black hole accretion over most of cosmic time \\citep[e.g.][]{2004MNRAS.349.1397C,2009MNRAS.399.1755C,2006AJ....131.2766R}, tracing the evolution of galaxy size and morphology \\citep[e.g.][]{1997ApJ...490..577D} and much more.  The second approach has been to expand our view in wavelength rather than cosmic time. The physical processes occurring in galaxies cause emission over the entire range of the electromagnetic spectrum. In order to have a full picture of galaxy formation, a multi-wavelength approach is vital. This has been made more achievable with recent generations of satellites covering the X-ray, ultra-violet, mid- and far-infrared. While the spectral energy distributions of stars tend to peak in the optical or near-infrared, obscuration and reprocessing by dust generates strong mid- and far-infrared emission. Young stars (when not obscured) are most directly traced in the ultraviolet, while black hole accretion can generate radiation from the radio to X-ray and gamma-ray bands. Surveys such as the Galaxy And Mass Assembly (GAMA) survey \\citep{2011MNRAS.413..971D} and the Cosmic Evolution Survey \\citep[COSMOS;][]{2007ApJS..172....1S} show the value of this multi-wavelength approach.  The third route, and the one that we address in this paper, is to focus on spatially resolving galaxies; in particular, obtaining spatially resolved spectroscopy. Optical spectroscopy allows us to measure a wide range of parameters including current star formation rates (e.g.\\ via \\ha), gas phase metallicities, stellar ages, stellar metallicities, black hole accretion rates, ionization structure and extinction due to dust (e.g.\\ via the Balmer decrement). The major spectroscopic surveys to date have used a single fibre \\citep{2001MNRAS.328.1039C,2003AJ....126.2081A} or single slit \\citep{2005A&A...439..845L,2007ApJ...660L...1D} on each object, and so obtain just one measurement of these parameters for each galaxy. Moreover, these measurements may not be representative of the galaxy as a whole, but biased, depending on where the aperture is placed. This fundamental problem is addressed by integral field spectrographs (IFS). In the last decade, projects such as SAURON \\citep{2001MNRAS.326...23B} have demonstrated the power of integral field spectroscopy to capture a range of key observables that are simply not available to single-aperture surveys. As well as studying the properties listed above in a spatially-resolved context, obtaining gas and stellar kinematics over an entire galaxy enables us to separate dynamical components, measure dynamical mass, examine the impact of winds and outflows, and discover merging systems via dynamical disturbances.  Integral field spectroscopy has almost exclusively been limited to single-object instruments, meaning that it is time-consuming to build large samples. The largest current data set, using the SAURON system on the William Herschel Telescope, contain $\\sim260$ objects \\citep[ATLAS-3D;][]{2011MNRAS.413..813C}. The CALIFA project \\citep{2011califa} aims to target 600 objects with the PMAS integral field unit (IFU) on the Calar Alto Telescope using $\\sim$200 nights of telescope time. The only multi-object integral field spectrograph currently available is that on the VLT FLAMES instrument \\citep{2002Msngr.110....1P}, which contains 15 IFUs, each with 20 spatial elements of size 0.52 arcsec and a total field of view (FoV) of $2\\times3$\\,arcsec. This has enabled measurements of the Tully-Fisher relation at $\\sim$0.6 \\citep{2008A&A...477..789Y} as well as a number of other projects. However the small FoV of each IFU, combined with the high spectral resolution ($\\ge$9000) and associated narrow wavelength range limits its applicability for large--scale surveys.  Astrophotonic technology \\citep{2009OExpr..17.1880B} is now opening the way to new instrumentation that can address the need for highly multiplexed integral field spectroscopy. Hexabundles \\citep{2011OExpr..19.2649B,2011MNRAS.415.2173B} are optical fibre bundles where the cladding has been stripped from each fibre to a minimum over a short length ($\\sim$30\\,mm) and the fibres then gently fused together at the input end to provide an IFU ($\\sim$1\\,mm aperture) with high filling factor. These can then be used in conventional multi-fibre spectrographs.  In this paper, we report on the Sydney-AAO Multi-object Integral field spectrograph (SAMI), the first fully operational demonstrator instrument to use hexabundles. SAMI has 13 IFUs that can be positioned anywhere over a 1 degree diameter field of view. In Section~\\ref{sec:rationale}, we discuss in detail the scientific rationale for such an instrument, along with some practical considerations regarding sensitivity. In Section~\\ref{sec:instrument}, we describe the SAMI instrument in detail. In Section~\\ref{sec:comm}, we outline the observations carried out during the commissioning of the instrument, the results from which are discussed in Section~\\ref{sec:results}. Our conclusions, and goals for the future, are laid out in Section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  In this paper we have presented a new instrument, the Sydney-AAO Multi-object IFS (SAMI). SAMI makes use of astrophotonic technology in the form of hexabundles (multi-core fibre bundles) to enable simultaneous IFU observations of 13 objects over a 1--degree diameter accessible field. Each IFU contains 61 elements, each 1.6 arcsec in diameter, giving a FoV for each IFU of 15 arcsec. SAMI has now been commissioned on the Anglo-Australian Telescope (AAT) and we demonstrate its science potential via preliminary observations of galaxies selected from the 6dF Galaxy Survey.  We make the case that multiplexed integral field spectroscopy is the natural next step in galaxy surveys, which to date have been dominated by multiplexed single-aperture observations. The extra information gained by IFU spectroscopy, combined with the statistical power of a large survey, will enable a fundamental step forward in our understanding of galaxy formation and evolution by distinguishing the spectroscopic properties of the major structural components of galaxies. Assuming that a survey could observe 3 fields per night (i.e.\\ nominal exposure times between 2 and 3\\,hours), SAMI will allow 10,000 galaxies to be targeted in 260 clear nights on the AAT.  Much greater gains could be made with an increased FoV and larger numbers of IFUs. A key requirement for making full use of the larger fields of view available on current and future telescopes is to have sufficient spectrographs to handle all the fibres from the IFUs. A 2--degree diameter FoV (e.g.\\ that available with the AAT's 2dF corrector) could accommodate $\\sim$50 or more IFU bundles comparable to those in SAMI; this would, however, require $\\sim$3300 fibres (including sky fibres). Such large numbers of fibres naturally drive designs towards mass-produced fixed-format spectrographs, such as those currently being developed for projects such as MUSE \\citep{2004SPIE.5492.1145B} and VIRUS \\citep{2006NewAR..50..378H}.  More development work on the bundle technology is under way. One possible improvement is to force the fibre cores into a perfect hexagonal configuration to simplify the data analysis, particularly when using dithering to smooth out the discrete sampling. This will not be easy to do with our current core size while at the same time ensuring an AR-coated polished front facet that is not overly stressed by the hexagonal grid. With the successful commissioning on the AAT of fibre-based OH suppression technology \\citep{2004OExpr..12.5902B,2010SPIE.7735E..40E}, we can envisage whole bundles that are fully sky-suppressed at near-infrared wavelengths. Once the near-infrared sky is rendered as dark as the optical sky, exposure times in both regimes are expected to be the same and it makes sense to run optical and infrared spectrographs simultaneously."
    },
    "1112/1112.3017_arXiv.txt": {
        "abstract": "{ The ATLAS and CMS collaborations have reported an excess of events in the $\\gamma\\gamma$, $ZZ^*\\to 4\\ell$ and $WW^*$ search channels at an invariant  mass $m \\simeq 125$~GeV, which could be the first evidence for the long-awaited Higgs boson. We investigate the consequences of requiring $m_h\\simeq 125$~GeV in both the mSUGRA and NUHM2 SUSY models. In mSUGRA, large values of trilinear soft breaking parameter $|A_0|$ are required, and universal scalar $m_0\\agt  0.8$~TeV is favored so that we expect squark and slepton masses typically in the multi-TeV range. This typically gives rise to an ``effective SUSY'' type of sparticle mass spectrum. In this case, we expect gluino pair production as the dominant sparticle creation reaction at LHC. For $m_0\\alt 5$~TeV, the superpotential parameter $\\mu\\agt 2$~TeV and $m_A\\agt 0.8$~TeV, greatly restricting  neutralino annihilation mechanisms. These latter conclusions are softened if $m_0\\sim 10-20$~TeV or if one proceeds to the NUHM2 model. The standard neutralino abundance tends to be far above WMAP-measured values unless the neutralino is higgsino-like. We remark upon possible non-standard (but perhaps more attractive) cosmological scenarios which can bring the predicted dark matter abundance into accord with the measured value, and discuss the implications for direct and indirect detection of neutralino cold dark matter. } ",
        "introduction": "\\label{sec:intro}  Recently, the ATLAS and CMS experiments have performed a combined search~\\cite{combined} for the Standard Model (SM) Higgs boson $H_{SM}$ using 1-2.3~fb$^{-1}$ of integrated luminosity with the result that the region $141$~GeV $<m_{H_{SM}}<$ 476~GeV is now excluded as a possibility at 95\\%CL. Even more recently, using the full data sample in excess of $5$~fb$^{-1}$ per experiment collected in 2011, the ATLAS~\\cite{atlas} and CMS~\\cite{cms} collaborations have reported excesses in the Higgs search $\\gamma\\gamma$, $ZZ^*\\to 4\\ell$ and $WW^*\\to 2\\ell$ channels with reconstucted invariant mass $m(\\gamma\\gamma)\\sim m(4\\ell )\\sim 125$~GeV. The combined statistical significance lies at the $2.5\\sigma$ level. These latest results might be construed as the first emerging direct evidence of the Higgs boson. Indeed, these new Higgs search results are consistent with the combined LEP2~\\cite{lep2higgs}/Tevatron precision electroweak analyses~\\cite{higgsEW} which favor the existence of a Higgs boson with mass not much beyond the LEP2 limit of $m_{H_{SM}}>114.4$~GeV.  While the putative $m_h\\sim 125$~GeV signal is consistent with SM expectations, it is rather stunning that it is also well in accord with expectations from supersymmetric models (SUSY), where the window of possible Higgs masses $m_h$ is far smaller. In the Minimal Supersymmetric Standard Model (MSSM), the Higgs sector consists of two doublet fields $H_u$ and $H_d$, which after the breaking of electroweak symmetry, result in the five physical Higgs bosons: two neutral $CP$-even scalars $h$ and $H$, a neutral $CP$-odd pseudoscalar $A$ and a pair of charged scalars $H^\\pm$~\\cite{wss}. At tree level, the value of $m_h$ is bounded by $M_Z|\\cos 2\\beta|$, where $\\tan\\beta \\equiv v_u/v_d$ is the ratio of Higgs field vacuum expectation values. Including radiative corrections, which depend on various sparticle masses and mixings that enter the $h$-boson self-energy calculation, one finds instead that $m_h\\alt 135$~GeV~\\cite{hmass}. In fact, using $\\sim 1$~fb$^{-1}$ of data in summer 2011, ATLAS~\\cite{atlww} and CMS~\\cite{cmsww} had already reported some excess of $WW^*$ events. In Ref.~\\cite{hww}, such events had been shown to favor a rather high mass light Higgs scalar $h$, with mass in the $m_h\\sim 125-130$~GeV range, and with large scalar masses $m_0$ and large trilinear soft breaking terms $A_0\\sim \\pm 2m_0$~\\cite{hww}.  Over most of the MSSM parameter space, the lighest Higgs boson $h$ is nearly SM-like so that SM Higgs search results can also be directly applied to $h$ (for exceptions, see Ref.~\\cite{belyaev}). A calculation of the light (heavy) scalar Higgs boson mass at 1-loop level using the effective potential method gives \\be m_{h,H}={1\\over 2}\\left[(m_A^2+M_Z^2+\\delta)\\mp\\xi^{1/2}\\right]\\, , \\ee where $m_A$ is the mass of the $CP$-odd pseudoscalar $A$ and \\be \\xi = \\left[ (m_A^2-M_Z^2)\\cos 2\\beta +\\delta\\right]^2+\\sin^2 2\\beta (m_A^2+M_Z^2)^2 \\, . \\ee The radiative corrections can be approximated as follows \\be \\delta =\\frac{3g^2m_t^4}{16\\pi^2M_W^2\\sin^2\\beta}\\log\\left[ \\left(1+{m_{\\tst_L}^2\\over m_t^2}\\right)\\left(1+{m_{\\tst_R}^2\\over m_t^2}\\right)\\right] . \\ee Thus, in order to accommodate a value of $m_h\\sim 125$~GeV, we anticipate rather large values of top squark soft masses $m_{\\tst_{L,R}}$ typically at least into the few-TeV range.  For our calculation of $m_h$, we include the full third generation contribution to the effective potential, including all sparticle mixing effects~\\cite{bisset}. The effective Higgs potential, $V_{eff}$, is evaluated with all running parameters in the $\\overline{DR}$ renormalization scheme evaluated at the scale choice $Q_{SUSY}=\\sqrt{m_{\\tst_1}m_{\\tst_2}}$, {\\it i.e.} the mean top squark mass scale. Of particular importance is that the $t$, $b$ and $\\tau$ Yukawa couplings are evaluated at the scale $Q_{SUSY}$ using 2-loop MSSM RGEs and including full 1-loop MSSM radiative corrections~\\cite{pbmz}. Evaluating $V_{eff}$ at this (optimized) scale choice then includes the most important two-loop effects~\\cite{hh}. This calculational procedure has been embedded in the Isajet mass spectra program Isasugra~\\cite{isasugra}, which we used for the present work. We note that just a few~GeV theory error is expected in our $m_h$ calculation. Also, it should be noted that our value of $m_h$ is typically a couple GeV below the corresponding FeynHiggs~\\cite{feynhiggs} calculation, mainly due to the fact that we are able to extract and use the two-loop $\\overline{DR}$ Yukawa couplings including 1-loop threshold corrections in our calculation of radiative corrections to $m_h$. Our calculation of $m_h$ agrees well with results from SuSpect, SoftSUSY and Spheno codes~\\cite{SuSpect}.  Our goal in this paper is to calculate the implications of a 125~GeV light Higgs scalar $h$ for supersymmetry searches at LHC, and for direct and indirect neutralino dark matter searches. In Sec.~\\ref{sec:msugra}, we examine implications of a 125~GeV light Higgs scalar in the paradigm mSUGRA model~\\cite{msugra}. In Sec.~\\ref{sec:nuhm2}, we examine implications in the more general 2-parameter non-universal Higgs model (NUHM2). In Sec.~\\ref{sec:dm}, we examine implications of a 125~GeV light Higgs scalar for $(g-2)_\\mu$, $BF(b\\to s\\gamma )$, $BF(B_s\\to\\mu^+\\mu^- )$ and for direct and indirect searches for neutralino cold dark matter (CDM). In Sec.~\\ref{sec:conclude}, we present our conclusions. ",
        "conclusions": "\\label{sec:conclude}  Evidence has been presented by ATLAS and CMS at the $\\sim 2.5\\sigma$ level for the existence of a light Higgs scalar with mass $m_h\\simeq 125$~GeV. If this evidence is bolstered by an increased data sample in 2012, then the discovery will have strong implications for supersymmetric models. We have examined both the mSUGRA model and the NUHM2 model under the restriction that $m_h=125\\pm 1$~GeV.  In the case of the mSUGRA (CMSSM) model, we conclude the following. \\bi \\item The common GUT scale scalar mass $m_0\\agt 0.8$~TeV. This tends to imply that squark and slepton masses are $>2$ TeV with $m_{\\tq}>m_{\\tg}$. In fact, the entire low $m_0$, low $m_{1/2}$ region of the mSUGRA plane is ruled out independent of $A_0$ or $\\tan\\beta$ values. \\item The soft breaking trilinear parameter $|A_0|\\alt 1.8m_0$ is excluded for $m_0<5$ TeV, or $|A_0|\\alt 0.3m_0$ is excluded if $m_0$ ranges up to 20 TeV. \\item The superpotential Higgs mass term $\\mu \\agt 2$~TeV for $m_0\\alt 5$ TeV. This strongly restricts mixed higgsino-bino states as a source of thermal neutralino CDM, as would be found in the HB/FP region. This constraint is relaxed if $m_0$ lies in the $5-20$ TeV range. \\item $m_A\\agt 0.8$~TeV, which means $m_{\\tz_1}\\agt 0.4$~TeV if neutralinos annihilate through the  $A$-resonance. \\ei  In the case of NUHM2 model, we find: \\bi \\item $m_0\\agt 0.8$~TeV as in mSUGRA, \\item for $m_0<5$ TeV, then $A_0\\alt -1.8m_0$ or $A_0\\sim +2m_0$, \\item for $m_0\\sim 5-20$ TeV, then just $A_0\\alt 2.5m_0$ is required, \\item unlike mSUGRA, the entire ranges of $\\mu$ and $m_A$ are still allowed, \\item thermally produced neutralinos match the WMAP-measured relic abundance for a mixed higgsino state at $m_{\\tz_1}\\sim 0.7$ TeV. \\ei  In addition, for NUHM2 and mSUGRA models, \\bi \\item A value of $m_h\\simeq 125$~GeV is inconsistent with the $(g-2)_\\mu$ anomaly. If the anomaly turns out to be real, it may imply alternative models such as ``normal scalar mass hierarchy''~\\cite{nmh} where first/second generation GUT scalar masses  $m_0(1,2)$ are much lighter than third generation scalars $m_0(3)$. \\item A value of $m_h\\simeq 125$~GeV is completely consistent with the measured values of $BF(b\\to s\\gamma )$, $BF(B_s\\to\\mu^+\\mu^- )$ and $BF(B_u\\to\\tau^+\\nu_\\tau )$. \\item Neutralino CDM is typically overproduced in the standard MSSM cosmology, unless the neutralino is higgsino-like, in which case its mass is around 0.8~TeV. In non-standard cosmologies, such as those including late decaying moduli fields or mixed axion/LSP CDM, the CDM abundance can be easily brought into accord with measured values. \\item Direct and indirect WIMP detection rates tend to be very low for models with a standard overabundance of CDM. In the case of higgsino-like WIMPs with a standard underabundance, direct and indirect detection prospects are rather bright. \\ei  \\noindent{\\bf Note Added:} After this work was finished, several papers appeared that also investigated implications of the recent LHC Higgs search results on mSUGRA and NUHM models~\\cite{others}. Their results tend to agree with ours although small differences do arise due to differences in the considered ranges of model parameters."
    },
    "1112/1112.1823_arXiv.txt": {
        "abstract": "Radio continuum observations of barred galaxies revealed strong magnetic fields of $\\ge50-100\\,\\mu$G in the circumnuclear starbursts. Such fields are dynamically important and give rise to magnetic stress that causes inflow of gas towards the center at a rate of several solar masses per year, possibly along the spiral field seen in radio polarization and as optical dust lanes. This may solve the long-standing question of how to feed active nuclei, and explain the relation between the bolometric luminosity of AGN nuclei and the star-formation rate of their hosts. The strong magnetic fields generated in young galaxies may serve as the link between star formation and accretion onto supermassive black holes. -- Magnetic fields of $\\ge160\\,\\mu$G strength were measured in the central region of the almost edge-on starburst galaxy NGC~253. Four filaments emerging from the inner disk delineate the boundaries of the central outflow cone of hot gas. Strong Faraday rotation of the polarized emission from the background disk indicates a large-scale helical field in the outflow walls. ",
        "introduction": "Magnetic fields are a major agent in the ISM and also control the density and propagation of cosmic rays. The radio--infrared correlation indicates that turbulent fields are strongest in star-forming regions. Magnetic fields and cosmic rays can provide the pressure to drive galactic outflows. Outflows from starburst galaxies in the early Universe may have magnetized the intergalactic medium. In spite of our increasing knowledge of cosmic magnetic fields, many important questions, especially their origin, strength in intergalactic space, first occurrence in young galaxies and their dynamical importance for galaxy evolution remain unanswered. ",
        "conclusions": "Magnetic fields in the central regions of galaxies are dynamically important. Magnetic stress drives gas inflow from the circumnuclear regions towards the nuclei. This process connects the star formation activity to the accretion rate and could be important for the formation of supermassive black holes. Ordered fields as observed by polarized radio emission may trace the direction of gas in- and outflows.  Results have been obtained so far only for radio-bright and nearby galaxies. A better insight into the interaction between gas and magnetic fields, especially near galactic nuclei, needs higher resolution and higher sensitivity. The {\\it EVLA}\\ has largely improved the sensitivity for radio continuum observations. The {\\it EVLA}\\ project {\\it CHANGES}\\ (PI: Judith Irwin) will search for outflows in a large sample of edge-on galaxies. The goal for the next decade is the {\\it Square Kilometre Array (SKA)}\\ with a collecting area of about 100x that of present-day telescopes. The SKA will open a new era in the observation of cosmic magnetic fields (Beck \\cite{beck10}). Starburst galaxies will be observable out to redshifts of $z\\simeq3$ (Murphy \\cite{murphy09})."
    },
    "1112/1112.3895_arXiv.txt": {
        "abstract": "We present a new scheme for the general computation of cosmic propagators that allow to interpolate between standard  perturbative results at low-$k$ and their expected large-$k$ resummed behavior. This scheme is applicable to any multi-point propagator  and allows the matching of perturbative low-$k$ calculations to any number of loops to their large-$k$ behavior, and can potentially be applied in case of non-standard cosmological scenarios such as those with non-Gaussian initial conditions.  The validity of our proposal is checked against previous prescriptions and measurements in numerical simulations  showing a remarkably good agreement.  Such a generic prescription for multi-point propagators provides the necessary building blocks for the computation of  polyspectra in the context of the so-called $\\Gamma$-expansion  introduced by Bernardeau et al. (2008). As a  concrete application we present a consistent calculation of the matter bispectrum at one-loop order. ",
        "introduction": "The large-scale structure of the universe that we observe today is thought to emerge from gravitational instabilities out of primordial metric perturbations, therefore precise observations of the large-scale structure of the local universe can be used to put constraints on  cosmological models. The connection between cosmological parameters and observations are however only straightforward when they correspond to the linear regime, i.e. when the observables can be computed as a linear transform of the primordial metric perturbations; this connection is less trivial when nonlinearities, due in a large part to the gravitational dynamics itself, are present. Those nonlinearities arise  during the late stage of the gravitational instabilities. This is an epoch during which the universe can be safely assumed to be matter dominated (at least in the context of the standard model of cosmology where the dark energy  component forms an homogeneous fluid.)  According to the cosmological principle  cosmic fields are statically homogeneous and isotropic. In the context we are interested in, two degrees of freedom are then relevant (see~\\cite{2002PhR...367....1B} for details): the local density contrast, $\\delta(\\vx)$ and the velocity divergence, $\\theta(\\vx)=\\partial_{i}\\vu_{i}(\\vx)$.  It is then very fruitful to introduce the Fourier modes of the cosmic fields $\\delta(\\vk)$ and $\\theta(\\vk)$, which evolve independently of one another in the linear regime. One can then introduce the doublet $\\Psi_{a}(\\vk)=\\left\\{\\delta(\\vk),-\\theta(\\vk)/\\mH\\right\\}$, where ${\\cal H}$ is the conformal expansion rate with its time evolution described by the Friedmann equation, to write down the equations of motion in compact form and facilitate the implementation of resummation techniques.  In this context, the goal of the theoretical and numerical calculations is a precise description, beyond the linear regime, of the statistical properties of $\\Psi_{a}$. We are  particularly interested in the multi-component power  spectra $P_{ab}(k)$ defined as, \\begin{equation} \\mg \\Psi_{a}(\\vk)\\Psi_{b}(\\vk')\\md=\\Dirac(\\vk+\\vk')P_{ab}(k) \\end{equation} where the Latin indices $a$ and $b$ vary from 1 to 2, which are implicitly or explicitly measured in observations such as galaxy surveys, and higher-order spectra such as bispectra, \\begin{equation} \\mg \\Psi_{a}(\\vk_{1})\\Psi_{b}(\\vk_{2})\\Psi_{c}(\\vk_{3})\\md=\\Dirac(\\vk_{1}+\\vk_{2}+\\vk_{3})B_{abc}(\\vk_{1},\\vk_{2}). \\end{equation}  This problem is in general very complicated if one wants to solve it from first principles. It can be made slightly easier to address in the context of the standard cosmological model where the primordial metric perturbations are expected to follow Gaussian or near Gaussian statistics. In that case the primordial properties of the fields are entirely determined by the initial power spectra, $P_{ab}^{\\init}(k)$. The question is then to uncover the functional dependence of $P_{ab}(k)$ as time evolves with $P_{ab}^{\\init}(k)$.  In the last few years attempts have been made to present perturbative schemes in the context of the growth of structure. The standard perturbation theory is unambiguously defined but leads to uncontrollable results \\cite{2002PhR...367....1B}. On the other hand alternative approaches have been proposed which produce more robust results, such as the Renormalized Perturbation Theory (RPT) in \\cite{2006PhRvD..73f3519C}, the Time Renormalization Group (TRG) approach in \\cite{2008JCAP...10..036P}, the closure theory in \\cite{2009PhRvD..79j3526H,2008ApJ...674..617T}, or with the help of perturbation theory expansion expressed in Lagrangian variables \\cite{2008PhRvD..77f3530M,2008PhRvD..78h3503B,2011JCAP...08..012O,2011PhRvD..83h3518M},  etc. One of such approaches, advocated in~\\cite{2008PhRvD..78j3521B}, is based on what has been called the $\\Gamma$-expansion.  This method exploits the following relation, \\begin{equation} \\mg \\Psi_{a}(\\vk)\\Psi_{b}(\\vk')\\md =\\sum_{p}\\frac{1}{p!} \\int\\dd^3\\vq_{1}\\dots \\dd^3\\vq_{p} \\left\\langle \\frac{\\dpartial \\Psi_{a}(\\vk)}{\\dpartial\\Phi_{a_{1}}(\\vq_{1})\\dots \\dpartial\\Phi_{a_{p}}(\\vq_{p})}\\right\\rangle P^{\\init}_{a_{1}b_{1}}(q_{1})\\dots P^{\\init}_{a_{p}b_{p}}(q_{p}) \\left\\langle \\frac{\\dpartial \\Psi_{b}(\\vk')}{\\dpartial\\Phi_{b_{1}}(\\vq_{1})\\dots \\dpartial\\Phi_{b_{p}}(\\vq_{p})}\\right\\rangle \\label{PkRecons} \\end{equation} where the standard Einstein convention (repeated indices are summed over) is used and where $\\Phi_{a}(k)$ is the value of $\\Psi_{a}(k)$ at the initial time. This relation expresses the fact that ensemble average over primordial fluctuations can be reorganized in an alternative way to that in standard perturbation theory \\cite{2002PhR...367....1B}. It exhibits the multi-point propagators defined as the ensemble average of the infinitesimal variation of the cosmic fields with respect to the initial conditions. More precisely  the multi-point propagators\\footnote{It is important to note that in this paper we call the $\\Gamma^{(p)}$ functions the $(p+1)$-propagators -- because it connects $p+1$ lines -- but alternative conventions can be found in the literature where $\\Gamma^{(p)}$ is called the $p$-point propagator.} $\\Gamma^{(p)}(\\vq_{1},\\dots,\\vq_{p})$ are defined as \\begin{equation} \\frac{1}{p!}\\left\\langle \\frac{\\dpartial \\Psi_{a}(\\vk)}{\\dpartial\\Phi_{a_{1}}(\\vq_{1})\\dots \\dpartial\\Phi_{a_{p}}(\\vq_{p})}\\right\\rangle \\equiv \\Dirac\\Big(\\vk-\\sum_{r=1}^{p}\\vq_{r}\\Big)\\, \\Gamma^{(p)}(\\vq_{1},\\dots,\\vq_{p}). \\end{equation} The relation (\\ref{PkRecons}) is valid for Gaussian initial conditions but can be extended  for non-Gaussian initial conditions~\\cite{2010PhRvD..82h3507B}. It clearly shows that the propagators are key ingredients for calculating nonlinear power spectra. They are the building blocks of the $\\Gamma-$expansion approach~\\cite{2008PhRvD..78j3521B} and the focus of this paper.  The second reason why these quantities appear to be important ingredients in perturbation theory calculations is that their asymptotic properties, e.g. how they behave for large wave-numbers, can be computed beyond perturbation expansions. This has been pioneered in \\cite{2006PhRvD..73f3519C,2006PhRvD..73f3520C} and recently revised in \\cite{2011arXiv1109.3400B} with the so-called eikonal approximation. More specifically it has been shown that in the high-$k$ limit the multi-point propagators are damped by a function that depends on the displacement field alone, irrespectively of the dynamics responsible of this displacement. Predictions on the behavior of those objects are then robust and can be computed in various approximations.  This motivates their use as the building blocks of a perturbation theory scheme. To achieve this end, one must have a description of multi-point propagators at all scales, matching the perturbative calculations at low-$k$ to the resummed asymptotic behavior at high-$k$.  As resummed propagators also contain information on loop contributions, both regimes can only be matched if a consistent interpolation schemes can be built. Furthermore, it is arguably this interpolation regime the most important for the prediction of loop corrections to (equal-time) correlation functions, as e.g. the power spectrum: while the low-$k$ limit can be safely computed using standard perturbation theory, the high-$k$ limit only adds a (time-dependent) phase-shift to Fourier modes and thus does not contribute to  the power spectrum or bispectrum.  This problem  has been solved for the two-point propagator in \\cite{2006PhRvD..73f3520C} with the help of an exponentiation scheme that interpolates between the one-loop results and its high-$k$ behavior.  However, this  prescription is somewhat ad-hoc in that is specific  to matching one-loop to resummed behavior of the two-point propagator  but it is not clear a priori how it can be extended  to incorporate  higher-loop information at low-$k$ or its generalization to multi-point propagators. The aim of this paper is to revisit this problem and propose a consistent solution which would be valid for any propagators and incorporate perturbative information to any loop order.  This paper is organized as follows. In the section II we recall the basic  equations of motion, define our notation including the diagrammatic description and present the $\\Gamma$-expansion of power spectra and bispectra in terms of multi-point propagators.  In section III we review the expected properties of the propagators. In section IV the proposed interpolation scheme is presented in detail, while a comparison of our theoretical predictions for the tree-point propagator to numerical simulations is presented in section V. The implications of our results are illustrated with a bispectrum computation in the  section VI. Lastly section VII contains our conclusions. ",
        "conclusions": "\\label{sec:conclusions}   In this paper we propose a systematic interpolation scheme that aims at describing the (multi-point) propagators in such a way that their expressions interpolate between their perturbation theory forms - to any loop order - and their large-$k$ behavior obtained from non-perturbative re-summations. This scheme is based on a separation of scales between the long wave-modes, whose effects are fully resumed and lead to the large-$k$ behavior, and  the short wave-modes that are fully taken into account in the perturbation theory treatment. Our prescription is different from the exponentiation scheme proposed in \\cite{2006PhRvD..73f3520C} but departs from it only very weakly. Our new prescription however is very general and can be used for any loop order calculations and for any propagator. Furthermore, this construction is totally unambiguous. We note that it can even be used in the context of primordial non-Gaussianities although new terms arise only at two loop order.  As the construction proposed here has a very general range of application, it should in principle be tested for a large variety of quantities, from two-point and multi-point propagators of the density field to the ones for the velocity field. We proposed here some comparisons with N-body results for the quantities that are of most interest  for the use of the multi-point propagators in the context of the $\\Gamma$-expansion applied to the calculation of the  density power spectrum. Because of the property we mentioned in the previous paragraph, this prescription for the two-point propagator is found to give very accurate results for the two-point propagator. We leave for further studies the impact of two-loop effects. In this work we further check the validity of our prescription against numerical results for the one-loop level of the density three-point propagator. These comparisons are presented in Fig. \\ref{regGamma2Num}. We found that it gives a satisfactory form for a large range of configurations, i.e. interpolating the low-$k$ one-loop result with the high-$k$ exponential decay, even when the signs of tree order and one-loop forms differ.  In the context of the $\\Gamma$-expansion approach this construction therefore provides us with the necessary recipes for constructing poly-spectra incorporating any order of perturbation theory results. In coming papers we will explicitly compare predictions of the $\\Gamma$-expansion for the power spectra with numerical simulations when higher order PT loop corrections are included. These prescriptions provide a good opportunity for giving the explicit form of the bispectra in the context of the $\\Gamma$-expansion. The explicit mathematical forms are given in Section \\ref{sec:bispectra} when bispectra are computed up to one-loop order. Such expressions make use in particular of the three-point propagators at one-loop order whose explicit forms are given in the appendix. We found in particular that the bispectra terms can be separated in a tree order contribution and coupling terms contributions that contribute at different scales, i.e. in subsequent bumps,  in a similar way to the power spectrum. Comparison of the proposed forms for the bispectrum with $N$-body results is however left for further studies.  We finally note that such a construction is restricted here to the case of a single pressure-less fluid. Whether it could be used in cases the fluid content of the universe is richer (with extra degrees of freedom carried by baryons, see \\cite{2010PhRvD..81b3524S,2011arXiv1109.3400B}, massive neutrinos, dark energy components, see \\cite{2011JCAP...06..019B,2011JCAP...03..047S}, etc.) is still largely an open issue. Results obtained in  \\cite{2011arXiv1109.3400B} however suggest that the effects of {\\sl adiabatic} long wavelength modes can be resummed and therefore could be incorporated in a large variety of cases."
    },
    "1112/1112.1542_arXiv.txt": {
        "abstract": "\\noindent{Recent calculations of the hydromagnetic deformation of a stratified, non-barotropic neutron star are generalized to describe objects with superconducting interiors, whose magnetic permeability $\\mu$ is much smaller than the vacuum value $\\mu_0$. It is found that the star remains oblate if the poloidal magnetic field energy is $\\gtrsim 40\\%$ of total magnetic field energy, that the toroidal field is confined to a torus which shrinks as $\\mu$ decreases, and that the deformation is much larger (by a factor $\\sim \\mu_0/\\mu$) than in a non-superconducting object. The results are applied to the latest direct and indirect upper limits on gravitational-wave emission from Laser Interferometer Gravitational Wave Observatory (LIGO) and radio pulse timing (spin-down) observations of 81 millisecond pulsars, to show how one can use these observations to infer the internal field strength. It is found that the indirect spin-down limits already imply astrophysically interesting constraints on the poloidal-toroidal field ratio and diamagnetic shielding factor (by which accretion reduces the observable external magnetic field, e.g. by burial). These constraints will improve following gravitational-wave detections, with implications for accretion-driven magnetic field evolution in recycled pulsars and the hydromagnetic stability of these objects' interiors.} ",
        "introduction": "The external magnetic field of a neutron star is (relatively) easily inferred from its spin-down rate, but its internal magnetic field is not directly observable. The main clue suggesting the existence of strong internal neutron star fields comes from the 1998 August 27 giant flare from the soft gamma-ray repeater (SGR) 1900+14 \\citep{fetal99,hetal99,metal99}. The giant flare, which released $\\sim\\ 10^{37}$ J of energy as X-rays, was accompanied by a 2.3-fold increase in the spin-down rate \\citep{metal99,wetal99,tetal00}. To explain this, \\citet{i01} proposed that the flare and the enhanced spin down were caused by a global reconfiguration of the internal magnetic field of $\\sim 10^{13}$ T, well above the external dipole field of $6.4\\times 10^{10}$ T\\footnote{Throughout this paper, we use SI units: 1 T $= 10^{4}$ G.}. \\citet{co11} generalised the \\citet{i01} calculation (by allowing the toroidal field strength to change, as well as the moment of inertia) and concluded that an internal field strength of $\\sim 10^{12}$ T gave rise to the 1998 August 27 event, lower than the first estimate by \\citet{i01}, but still significantly higher than the observed external field.  Stellar ellipticity can also be used to constrain the strength of a star's internal field \\citep{c02,detal09,aetal10,p11}. It is well known that a strong magnetic field can deform a star \\citep{cf53,f54,g72,k89,pm04,hetal08,metal11}. The ellipticity $\\epsilon$ is roughly proportional to the magnetic energy \\citep{c02,hetal08,detal09}. Neutron stars, with their intense magnetic fields, possess significant ellipticities, making them good candidates for gravitational wave sources \\citep{bg96,mp05,setal05,hetal08,detal09}. Recent data from the fifth Laser Interferometer Gravitational Wave Observatory (LIGO) Science Run set an upper limit of $\\epsilon\\lesssim 1.4\\times 10^{-4}$ on the Crab Pulsar \\citep{aetal08,aetal10}, translating into an internal magnetic field of $\\lesssim 10^{12}$ T under standard assumptions. LIGO non-detections of the central compact object (CCO) in the supernova remnant Cassiopeia A (Cas A) have constrained its ellipticity as well. The Cas A CCO has not been detected electromagnetically, making it impossible to infer its external magnetic field from the spin-down rate. However, \\citet{wetal08} and \\citet{w10} constrained its ellipticity as a function of gravitational wave frequency (e.g., $\\epsilon\\lesssim 3.6 \\times 10^{-4}$ for 100 Hz, $\\epsilon\\lesssim 0.6\\times 10^{-4}$ for 200 Hz, and $\\epsilon\\lesssim 0.38\\times 10^{-4}$ for 300 Hz), implying an internal magnetic field $\\lesssim 10^{14}$ T. Lastly, \\citet{cetal11} showed that it will be possible to use future data from LIGO to set a lower limit of $10^7$ T on the magnetic field ($\\epsilon \\lesssim 10^{-4}$) of the putative 24-year-old neutron star in the supernova remnant SNR 1987A.  Gravitational waves are generated by a rotating star when it is not spherically symmetric and when its `wobble angle' $\\theta$, the angle between its total angular momentum vector and symmetry axis, is nonzero. In general, therefore, the star precesses as it radiates. The magnetic field analysed in this paper is axisymmetric, which deforms the star into an ellipsoid. The most general expression for the gravitational wave signal of a precessing triaxial ellipsoid [given by, e.g., Eqs. (19)--(26) of \\citep{jks98}] depends on $\\theta$ and the angle $\\iota$ between the angular momentum vector and line of sight to the Earth, as well as $\\epsilon$. The signal is strongest when $\\theta=\\pi/2$ (i.e., the rotation axis is perpendicular to the symmetry axis, which is in turn parallel to the magnetic axis) and $\\iota=0$ (i.e., the rotation axis is directed towards Earth). In this paper, to simplify matters, we henceforth assume implicitly that $\\theta=\\pi/2$ and $\\iota=0$, so that gravitational wave emission and signal detection are assumed to be optimal and there is no precession, in order to focus on $\\epsilon$ without geometric complications.   In \\citet{metal11}, we constructed hydromagnetic equilibria for stratified, \\emph{non-barotropic} stars. The commonly adopted barotropic assumption, while simplifying calculations, severely restricts the form of the field that can be `fitted' into the star, e.g., \\citet{hetal08} found that the field must vanish at the surface, contrary to observations, and \\citet{lj09} and \\citet{cfg10} found that only configurations dominated by the poloidal component (poloidal energy $\\gtrsim 90\\%$ of total) are allowed, contrary to the numerical simulations of \\citet{bn06}. By abandoning the barotropic assumption, we are able to construct a simple, self-consistent hydromagnetic equilibrium with an internal field that can be matched to an external dipole. In our configuration, because we do not require the pressure to be a function of density alone, we are less restricted in the choices of poloidal and toroidal components; they need not be of any particular relative strengths and are independently adjustable. Incidentally, this means that the magnetic field configuration is independent of the equation of state chosen, cf. \\citet{lj09} and \\citet{cfg10}.   In this paper, we present one possible astrophysical application of the aforementioned result, namely to constrain the internal magnetic fields of neutron stars in conjunction with gravitational wave observations. As we can match the internal field to an external dipole, we are able to relate one set of observations (external magnetic field strength, from spin period and spin-down rate) with another (gravitational wave upper limits), at least in principle. In Sec. 2, we briefly describe our field structure and summarize the results of \\citet{metal11}. Then, in Sec. 3, we generalize the earlier work to include a superconducting interior. We compare the ellipticity calculated using our superconducting model to the gravitational-wave upper limits of 81 known millisecond pulsars, including also the important effect of accretion-induced diamagnetic shielding. Lastly, in Sec. 4, we summarize our results and discuss the possibility that the internal fields of millisecond pulsars are stronger than currently thought. ",
        "conclusions": "This paper makes a first attempt at combining non-barotropic magnetized stellar models \\citep{metal11} with LIGO non-detections and radio timing data to constrain the ratio $\\Lambda$ and hence the internal magnetic field of recycled pulsars. The models resemble the `twisted torus' of \\citet{cetal09}, \\citet{lj09}, and \\citet{cfg10}, in that they consist of a potentially strong internal toroidal field, which is not observable directly and is confined inside the star, and an external field, which can be measured. The main difference is that we assume the star is stably stratified but non-barotropic \\citep{r09}. Because of this, our models allow the poloidal and toroidal components to be adjusted independently of each other; in barotropic models (such as the twisted torus model), the poloidal and toroidal fields must obey certain relations to ensure that pressure is always a function of density purely [see also, e.g., \\citet{hetal08}]. This means that our models can easily accommodate field configurations with $0.01\\lesssim\\Lambda\\lesssim 0.8$, the range found to be stable by the numerical simulations of \\citet{bn06}.   As seen in Fig. \\ref{eps}, our calculated $\\epsilon$ is much lower than the LIGO and spin-down upper limits when the inferred dipole field strengths are input directly into Eq. (5).\\footnote{Recall that the magnetic axis is assumed to be perpendicular to the rotation axis throughout this paper.} Note that the lower black dashed curve in Fig. \\ref{eps} already assumes a strong toroidal field ($\\Lambda=10^{-3}$); a weaker toroidal field will generate a curve even lower down. Next, we calculate how interior superconductivity ($\\mu_r=10^{-3}$) and accretion-induced screening or burial (red and blue bands in Fig. \\ref{eps}) enhance $\\epsilon$. This latter possibility is further explored in Fig. \\ref{xip}, which shows how, in principle, a gravitational wave detection from a recycled pulsar with a certain surface field (measured from spin down) can be used to infer both $\\Lambda$ and shielding factor $\\xi$. Fig. \\ref{xip} also tells us the sets of parameters that are ruled out by current observations: for a given pulsar with a measured spin-down field strength, non-detection means that $\\xi$ and $\\Lambda$ lie above that particular curve in Fig. \\ref{xip}. For example, the upper solid curve of Fig. \\ref{eps}, corresponding to $\\xi=10^{-4}$ and $\\Lambda=10^{-3}$, lies at the corner of the theoretically allowed shaded region in Fig. \\ref{xip} and is therefore ruled out even for a non-superconducting star with a relatively weak observed field strength of $10^4$ T (because it lies below the dashed blue curve in Fig. \\ref{xip}).   Is it possible that the magnetic field of a recycled pulsar is much stronger just below the surface than the inferred dipole field? \\citet{k91} and \\citet{a93}, motivated by the discrepancy between field strengths inferred from dipole spin down and from cyclotron line measurements of accreting X-ray pulsars, raised such a possibility. They proposed that higher-order multipoles may exist close to the surface. \\citet{a93} then showed that surface field strengths $\\sim 10^4$ times higher than the observed dipole field can account for the aforementioned discrepancy, as well as the location of the millisecond pulsars on the $P$--$\\dot{P}$ diagram. The surface field can also be masked by magnetic field burial, whereby accreted matter compresses the polar magnetic flux into a narrow belt around the equator. It has been shown that accretion of $\\gtrsim 10^{-5} M_\\odot$ is enough to alter the dipole moment significantly \\citep{l99,pm04,mp05,zk06,vm09}, depending on the equation of state \\citep{pmp11}.  \\citet{pg07} raised the possibility that the large-scale magnetic field of a neutron star is supported by long-lived currents in its superconducting core, while small-scale, fast-decaying (decay time $10^5$--$10^7$ yr) magnetic structures also exist near the surface, supported by short-lived currents in the crust. After trying several different initial poloidal and toroidal field strengths, \\citet{pg07} found that all their models eventually (after $\\sim 1$ Myr) reconfigure into a long-term stable configuration comprising a dipolar poloidal component and a quadrupolar/octupolar toroidal component. The decay of the crustal currents is dominated by Hall drift in the first $10^3$--$10^4$ yr if the initial field is $\\gtrsim 10^{10}$ T, i.e. magnetar strength \\citep{gr92,rg02,pg07}. When the external dipole field is of magnetar strength, the internal crustal field can still be $\\sim 10$ times greater.  \\citet{pg07} did not calculate the magnetic deformation of their star, but one can extend their work and that of \\citet{metal11} to investigate the effects of this strong, internal, quadrupolar/octupolar toroidal field on stellar ellipticity. To do so, one must remember that Eq. (5) applies to an internal field that is continuous with an external dipole, as shown in Fig. \\ref{field}. When the field is distorted significantly by accretion [into, e.g., the `equatorial tutu' shape found by \\citet{mp01} and \\citet{pm04}], then the internal toroidal field described by Eq. (2) must be recalculated to ensure that the total field still has the desired properties listed at the start of Sec. 2. A similar adjustment must be carried out to accommodate multipolar external fields. It can be shown, for example, that one needs two toroidal field belts around the new neutral circles at $(0.7R,\\pi/4)$ and $(0.7R,3\\pi/4)$ to match a quadrupolar external field \\citep{m10}.     Lastly, we remind the reader that a strong internal field is not the only possible cause for stellar deformation. Other effects that can deform a star are:  \\begin{itemize}[labelindent=\\parindent,leftmargin=*] \\item rotation; ellipticity induced by centrifugal forces is $\\epsilon_\\Omega\\approx 0.3 (\\nu/\\textrm{kHz})^2$, where $\\nu$ is the spin frequency of the neutron star \\citep{ps72,c02}. Magnetar spin frequencies are $\\sim 0.1$ Hz, giving $\\epsilon_\\Omega\\sim 10^{-9}$. However, the deformation is aligned with the rotation axis [except for a fraction $\\sim 10^{-4}$ proportional to the shear modulus \\citep{g70}], so it does not contribute to gravitational wave emission or precession \\citep{m00}; \\item crustal shear stresses; ellipticity supported by the crust's elasticity \\citep{hja06}, up to a maximum of $\\sim 10^{-7}$ for conventional neutron stars, $\\sim 10^{-5}$ for hybrid quark-baryon or meson-condensate stars, or $\\sim 10^{-4}$ for solid strange quark stars \\citep{o05,hk09}. \\item nuclear processes; a sustained, asymmetric temperature step of $\\sim 5\\%$ at the base of the crust drives electron capture reactions which produce compositional asymmetries \\citep{ucb00}. \\end{itemize}"
    },
    "1112/1112.1068_arXiv.txt": {
        "abstract": "Direct $N$-body calculations are presented of the early evolution of exposed clusters to quantify the influence of gas expulsion on the time-varying surface brightness. By assuming that the embedded OB stars drive out most of the gas after a given time delay, the change of the surface brightness of expanding star clusters is studied. The influence of stellar dynamics and stellar evolution is discussed.\\\\ The growth of the core radii of such models shows a remarkable core re-virialisation. The decrease of the surface mass density during gas expulsion is large and is only truncated by this re-virialisation process. However, the surface brightness within a certain radius does not increase noticeably. Thus, an embedded star cluster cannot reappear in observational surveys after re-virialisation. This finding has a bearing on the observed infant mortality fraction. ",
        "introduction": "Since stars form in dense clumps within molecular clouds, very young clusters are observed to be embedded in gas. Typically $1/3$ of the initial cluster-forming gas mass is converted into stars \\citep{Lada:2003il,Machida:2011uw}. Following star formation, and especially after formation of OB stars, most of the residual gas is driven out, so that there is nearly no gas observed in populations older than a few million years. While the gas ejection process begins, self-gravitation decreases, resulting in a rapid expansion of the young star cluster (\\citealt{1997MNRAS.284..785G}, \\citealt{2006MNRAS.373..752G}, \\citealt{2003MNRAS.338..665B}, \\citealt{2007MNRAS.380.1589B}). In this contribution we investigate cluster formation and evolution during this process, focusing on the surface brightness of star clusters responding to residual star-forming gas expulsion. We want to answer the question of how the surface brightness is influenced by a) stellar dynamics and b) stellar evolution, and how both relate to each other.  In the context of gas expulsion, numerical $N$-body computations show a re-virialisation process taking place after a few million years in the centre of the system. While the cluster outskirts continue to expand, the centre contracts \\citep*{2001MNRAS.321..699K}, leading to an increase of the density in the central region. Since at any given age, the surface mass density has an impact on the surface brightness, this process is of special interest. We ask what role this effect plays in the evolution of the surface brightness and whether its impact is sufficient to let the cluster core disappear under a given detection limit owing to its expansion, then re-appear after re-virialisation.  We have performed direct $N$-body computations, including stellar evolution, with initial conditions similar to \\cite{2001MNRAS.321..699K}, with the emphasis being put on the time-evolution of the cluster surface brightness and surface mass density. The residual gas is modelled by a time-dependent background potential. ",
        "conclusions": "In the present work, we investigate the evolution of the surface brightness of star clusters under the assumption that the residual gas of the cluster-forming region is driven out by evolving OB stars, leading to a strong expansion of the total cluster. Five different models of star clusters with different sizes and masses are numerically evolved, every model with a SFE of $1/3$ and with five different gas expulsion time scales $\\tau$. Thereby we compare the influences of stellar dynamics and stellar evolution on the surface brightness. Based on the expansion histories the question came up whether the impact of the re-virialisation process is sufficient to let the cluster core disappear under a given detection limit owing to its expansion, then re-appear after re-virialisation. From our computations we conclude that the surface brightness is not changed significantly by this process, because it is clearly dominated by stellar evolution.  Independent by this finding, we can draw conclusions of the expansion behaviour. We find that the expansion of the inner parts of a star cluster is more sensitive to the gas expulsion time scale $\\tau$ than the outer parts are. While the outer Lagrange radii are only weakly responsive to changing $\\tau$, the inner Lagrange radii react very sensitively. E.g. the core expands moderately in the case of slow expulsion, while a quick expulsion on the other hand can lead to a rapid disruption of the core."
    },
    "1112/1112.3771_arXiv.txt": {
        "abstract": "By studying the stellar population properties along the radius in 15 nearby S0 galaxies, I have found that the outer stellar disks are mostly old, with the SSP-equivalent ages of 8--15~Gyr, being often older than the bulges. This fact puts into doubt a currently accepted paradigm that S0 galaxies have formed at $z=0.4$ by quenching star formation in spiral galaxies. ",
        "introduction": "Lenticular galaxies have been introduced by Edwin \\cite[Hubble (1936)]{hubble} as an intermediate type between ellipticals and spirals: they have large-scale stellar disks as spirals but lack blue spiral arms and HII-regions, and they look smooth and red as ellipticals. Now lenticulars are thought to be (trans-)formed from spirals by removing gas and quenching star formation in their disks; also dynamical heating is required to make the S0 disks stable against spiral wave perturbations. There are some evidences that this transformation might take place at the redshift of $z\\approx 0.4$, within dense environments, groups or clusters, where and when suddenly the dominance of blue (spiral?) galaxies is replaced by S0 dominance (e.g., \\cite[Fasano et al. 2000]{fasano}, \\cite[Wilman et al. 2009]{wilman}). But if this scenario is valid, the star formation in the outer disks of nearby S0s should proceed only 4 Gyr ago, and we must see intermediate-age stellar populations there. With deep long-slit spectroscopy of nearby S0 galaxies, I study the stellar population properties along the radius beyond several scalelengths of their large-scale disks to check the validity of this paradigm. ",
        "conclusions": "I have measured the Lick indices in 15 S0s along the radius beyond 2--4 exponential scalelengths of their disks, and have estimated the ages and abundances of the stellar populations for the bulges and for the disks. The bulges have solar metallicities or higher, and the disks have solar metallicity or lower. The disks demonstrate very high magnesium-to-iron ratios of their stellar populations so they cannot be descendants of spiral galaxies with prolonged star formation. The SSP-equivalent ages of the disks are mostly old -- 60\\%\\ of the sample demonstrate the ages older than 10~Gyr, -- and almost always (except one case) larger than the SSP-equivalent ages of their corresponding bulges (Fig.~\\ref{fig1}).  The parameters of the stellar populations of the disks do not show any correlation with the disk luminosities or masses. The only correlations found are those with the photometric disk scaleheight taken from the decomposition results by \\cite[Mosenkov et al. (2010)]{mosenkov}. The whole scaleheight range is from 0.3 kpc (found in the disks with the ages of 2--3~Gyr and [Mg/Fe]$< +0.2$) to 0.6--0.9 kpc (found mostly in the disks with the ages $>10$~Gyr and [Mg/Fe]$>+0.3$). All the old disks of our sample galaxies are so thick disks, while the only two young outer disks, those of NGC~4111 and NGC~7332, are certainly thin disks. The disk of NGC~2732, with its SSP-equivalent age of 8~Gyr and the scaleheight of 0.5 kpc, is halfway between thin and thick disks.  Many arguments evidence for S0s and spiral galaxies being relatives, but this time the thing looks like S0s are progenitors of spirals, opposite to what is thought before. Indeed, if we compare stars of the thick disk of our own Galaxy with the thick outer stellar disks of S0s studied here we will see full resemblance: the ages $>10$ Gyr, [Mg/Fe]$>+0.2$, the total metallicity [Z/H] is between 0.0 and --0.7 (such parameters of the thick disk of our Galaxy are found by e.g. \\cite[Bernkopf \\& Fuhrmann 2006]{bernfuhr} or by \\cite[Schuster et al. 2006]{schuster}). So if one provides fresh cold gas accretion into the disks of these S0s, after several Gyrs of star formation we would get a typical spiral galaxy, with the thick old stellar disk and thin younger stellar disk. I propose now the following evolutionary sequence. All disk galaxies were S0s immediately after their birth; later, at $z<1$, some of them acquired cold gas accretion sources -- these became spirals, -- and some of them failed to find such sources -- those remained lenticulars. In large cluster-size and group-size dark haloes, there are little chance to find external sources of cold gas accretion, due to hot gas effect, -- so in nearby clusters the dominant disk-galaxy population is S0s. The open question remains what can be these sources of cold-gas prolonged accretion -- they may be cosmologically motivated filaments (\\cite[Dekel \\& Birnboim 2006]{dekel}) or rich systems of irregular-type dwarf satellites which have been merging with the host one after another."
    },
    "1112/1112.1774_arXiv.txt": {
        "abstract": "The extended Galileon models possess tracker solutions with de Sitter attractors along which the dark energy equation of state is constant during the matter-dominated epoch, i.e. $w_{{\\rm DE}}=-1-s$, where $s$ is a positive constant. Even with this phantom equation of state there are viable parameter spaces in which the ghosts and Laplacian instabilities are absent. Using the observational data of the supernovae type Ia, the cosmic microwave background (CMB), and baryon acoustic oscillations, we place constraints on the tracker solutions at the background level and find that the parameter $s$ is constrained to be $s=0.034_{-0.034}^{+0.327}$ (95 \\% CL) in the flat Universe. In order to break the degeneracy between the models we also study the evolution of cosmological density perturbations relevant to the large-scale structure (LSS) and the Integrated-Sachs-Wolfe (ISW) effect in CMB. We show that, depending on the model parameters, the LSS and the ISW effect is either positively or negatively correlated. It is then possible to constrain viable parameter spaces further from the observational data of the ISW-LSS cross-correlation as well as from the matter power spectrum. ",
        "introduction": "The main target of the dark energy research over the next few years or so is to distinguish between the $\\Lambda$-Cold-Dark-Matter ($\\Lambda$CDM) model and dynamical models with time-varying equations of state $w_{{\\rm DE}}$. {}From the observational data of WMAP7 combined with baryon acoustic oscillations (BAO) \\cite{Percival} and the Hubble constant measurement \\cite{Riess09}, Komatsu \\textit{et al.} \\cite{WMAP7} derived the bound $w_{{\\rm DE}}=-1.10\\pm0.14$ (68 \\% CL) for the constant equation of state. Adding the supernovae type Ia (SN Ia) data provides tighter constraints on $w_{{\\rm DE}}$, but still the phantom equation of state ($w_{{\\rm DE}}<-1$) is allowed by the joint data analysis \\cite{WMAP7}. This property persists for the time-varying dark energy equation of state with the parametrization such as $w_{{\\rm DE}}=w_{0}+w_{a}(1-a)$ \\cite{para}, where $a$ is the scale factor \\cite{Wang}.  In the framework of General Relativity (GR) it is generally difficult to construct theoretically consistent models of dark energy which realize $w_{{\\rm DE}}<-1$. In quintessence \\cite{quin} with a slowly varying scalar-field potential, for example, the field equation of state is always larger than $-1$. A ghost field with a negative kinetic energy leads to $w_{{\\rm DE}}<-1$ \\cite{phantom}, but such a field is plagued by a catastrophic instability of the vacuum associated with the spontaneous creation of ghost and photon pairs \\cite{Trodden}.  In modified gravitational theories it is possible to realize $w_{{\\rm DE}}<-1$ without having ghosts and Laplacian-type instabilities (see Refs.~\\cite{review}). In $f(R)$ gravity, where the Lagrangian $f$ is a function of the Ricci scalar $R$, the dark energy equation of state crosses the cosmological constant boundary ($w_{{\\rm DE}}=-1$) \\cite{Hu,Amen08,Tsuji08,Moto} for the viable models constructed to satisfy cosmological and local gravity constraints \\cite{Hu,Amen08,fRviable,Tsuji08}. This is also the case for the Brans-Dicke theory \\cite{Brans} with a field potential which accommodates the chameleon mechanism \\cite{chame} to suppress the propagation of the fifth force \\cite{Yoko}. In modified gravity models of dark energy based on the chameleon mechanism (including $f(R)$ gravity), the effective potential of a scalar degree of freedom needs to be carefully designed to pass cosmological and local gravity constraints \\cite{Gan}.  There is another class of modified gravity models of dark energy in which a nonlinear self-interaction of a scalar degree of freedom $\\phi$ can lead to the recovery of GR in a local region through the Vainshtein mechanism \\cite{Vain}. The representative models of this class are those based on the Dvali-Gabadadze-Porrati (DGP) braneworld \\cite{DGP} and the Galileon gravity \\cite{Nicolis} (see Refs.~\\cite{DGPnon,Galileonva} for the implementation of the Vainshtein mechanism in these models). The nonlinear interaction of the form $(\\partial\\phi)^{2}\\square\\phi$, which appears from the brane-bending mode in the DGP model \\cite{DGPnon}, gives rise to the field equation invariant under the Galilean shift $\\partial_{\\mu}\\phi\\to\\partial_{\\mu}\\phi+b_{\\mu}$ in the flat spacetime. This was extended to more general field Lagrangians satisfying the Galilean symmetry in the limit of the Minkowski spacetime \\cite{Nicolis,DEV}.  The cosmology based on the covariant Galileon or on its modified versions has been studied by many authors \\cite{Galileondark,Galileoninf}. In Refs.~\\cite{DePRL,DePRD} the dynamics of dark energy was investigated in the presence of the full covariant Galileon Lagrangian. In this model the solutions with different initial conditions converge to a common trajectory (tracker). Along the tracker the dark energy equation of state $w_{{\\rm DE}}$ changes as $-7/3$ (radiation era) $\\to$ $-2$ (matter era) $\\to$ $-1$ (de Sitter era) \\cite{DePRL,DePRD,Vikman10,Kimura}. There exists a viable parameter space in which the ghosts and Laplacian instabilities are absent. However, the joint analysis based on the observational data of SN Ia, CMB, and BAO shows that the tracker solution is disfavored because of the large deviation of $w_{{\\rm DE}}$ from $-1$ during the matter era \\cite{Nesseris,Kimura}. The solutions that approach the tracker only at late times are allowed from the combined data analysis \\cite{Nesseris}.  As an extension of the covariant Galileon model, Deffayet \\textit{et al.} \\cite{Deffayet11} obtained the most general Lagrangian in scalar-tensor theories with second-order equations of motion. In four dimensions the corresponding Lagrangian is of the form (\\ref{Lagsum}) with the four functions (\\ref{eachlag2})-(\\ref{eachlag5}) given below. In fact this is equivalent to the Lagrangian found by Horndeski \\cite{Horndeski} more than 3 decades ago \\cite{Koba11,Cope}. The conditions for the avoidance of ghosts and Laplacian instabilities were recently derived in Ref.~\\cite{Defe2011} in the presence of two perfect fluids (non-relativistic matter and radiation).  The covariant Galileon corresponds to the choice $K=-c_{2}X$, $G_{3}=c_{3}X/M^{3}$, $G_{4}=M_{{\\rm pl}}^{2}/2-c_{4}X^{2}/M^{6}$, $G_{5}=3c_{5}X^{2}/M^{9}$ in Eqs.~(\\ref{eachlag2})-(\\ref{eachlag5}), where $c_{i}$'s are dimensionless constants, $X=-\\partial^{\\mu}\\phi\\partial_{\\mu}\\phi/2$, $M_{{\\rm pl}}$ is the reduced Planck mass, and $M$ is a constant having the dimension of mass. Kimura and Yamamoto \\cite{Kimura} studied the model with the functions $K=-c_{2}X$, $G_{3}=c_{3}M^{1-4n}X^{n}$ ($n\\geq1$), $G_{4}=M_{{\\rm pl}}^{2}/2$, and $G_{5}=0$, in which case the dark energy equation of state during the matter era is given by $w_{{\\rm DE}}=-1-s$ with $s=1/(2n-1)>0$. At the background level this is equivalent to the Dvali-Turner model \\cite{DTurner}, which can be consistent with the observational data for $n$ larger than the order of 1. If we consider the evolution of cosmological perturbations, the LSS tends to be anti-correlated with the late-time ISW effect. This places the tight bound on the power $n$, as $n>4.2\\times10^{3}$ (95\\% CL) \\cite{Kimura2}, in which case the dark energy equation of state is practically indistinguishable from that in the $\\Lambda$CDM model.  In Ref.~\\cite{Defe2011} the present authors proposed more general extended Galileon models with the functions $K=-c_{2}M_{2}^{4(1-p_{2})}X^{p_{2}}$, $G_{3}=c_{3}M_{3}^{1-4p_{3}}X^{p_{3}}$, $G_{4}=M_{{\\rm pl}}^{2}/2-c_{4}M_{4}^{2-4p_{4}}X^{p_{4}}$, and $G_{5}=3c_{5}M_{5}^{-(1+4p_{5})}X^{p_{5}}$, where the masses $M_{i}$'s are fixed by the Hubble parameter at the late-time de Sitter solution with $\\dot{\\phi}=$\\,constant. For the powers $p_{2}=p$, $p_{3}=p+(2q-1)/2$, $p_{4}=p+2q$, $p_{5}=p+(6q-1)/2$, where $p$ and $q$ are positive constants, there exists a tracker solution characterized by $H\\dot{\\phi}^{2q}=$\\,constant. During the matter-dominated epoch one has $w_{{\\rm DE}}=-1-s$, where $s=p/(2q)$, along the tracker. This covers the model of Kimura and Yamamoto \\cite{Kimura} as a specific case ($p=1$, $q=n-1/2$, $c_{4}=0$, $c_{5}=0$). In the presence of the nonlinear field self-interactions in $G_{4}$ and $G_{5}$, the degeneracy of the background tracker solution for given values of $p$ and $q$ is broken by considering the evolution of cosmological perturbations. Hence the ISW-LSS anti-correlation found in Refs.~\\cite{Kimura,Kimura2} for $c_{4}=c_{5}=0$ should not be necessarily present for the models with non-zero values of $c_{4}$ and $c_{5}$.  In this paper we first place constraints on the tracker solution in the extended Galileon models by using the recent observational data of SN Ia, CMB, and BAO. The bound on the value $s=p/(2q)$ is derived from the background cosmic expansion history with/without the cosmic curvature $K$. We then study the evolution of cosmological density perturbations in the presence of non-relativistic matter to break the degeneracy of the tracker solution at the background level. We will show that the LSS and the ISW effect are either positively or negatively correlated, depending on the parameters $c_{4}$ and $c_{5}$. This information should be useful to distinguish between the extended Galileon models with different values of $c_{4}$ and $c_{5}$ from observations. ",
        "conclusions": "In this paper we have studied cosmological constraints on the extended Galileon models of dark energy. For the functions (\\ref{geGali}) with the powers (\\ref{power}) there exist tracker solutions along which the field equation of state $w_{{\\rm DE}}$ changes as $-1-4s/3$ (radiation era) $\\to$ $-1-s$ (matter era) $\\to$ $-1$ (de Sitter era), where $s=p/(2q)>0$. Unlike the case of the covariant Galileon ($s=1$), $w_{{\\rm DE}}$ can be close to $-1$ during the radiation and matter eras for $0\\leq s\\ll1$. Moreover, even for $w_{{\\rm DE}}<-1$, there are viable model parameter spaces in which the ghosts and Laplacian instabilities are absent.  Using the recent data of SN Ia, CMB, and BAO, we placed observational constraints on the background tracker solutions with $s>0$. In the flat Universe we found that the model parameters are constrained to be $s=0.034_{-0.034}^{+0.327}$ and $\\Omega_{m,0}=0.271_{-0.010}^{+0.024}$ (95\\% CL) from the joint data analysis. The chi-square for the best-fit case ($s=0.034$) is slightly smaller than that in $\\Lambda$CDM. However the difference of the AIC information criteria between the two models is $\\Delta({\\rm AIC})=1.85$, so that the extended Galileon is not particularly favored over the $\\Lambda$CDM. We also carried out the likelihood analysis in the presence of the cosmic curvature $K$ and obtained the bounds $s=0.067_{-0.067}^{+0.333}$, $\\Omega_{m,0}=0.277_{-0.022}^{+0.023}$, and $\\Omega_{K,0}=0.0077_{-0.0127}^{+0.0039}$ (95\\,\\%~CL). The tracker for the covariant Galileon ($s=1$) is disfavored from the data, in which case only the late-time tracking solution is allowed observationally \\cite{Nesseris}.  The background quantities for the tracker are independent of the values of $\\alpha$ and $\\beta$. This means that the models with different $\\alpha$ and $\\beta$ cannot be distinguished from the observational constraints derived from the background cosmic expansion history. In order to break this degeneracy we studied the evolution of cosmological perturbations in the presence of non-relativistic matter for the flat FLRW background. As shown in Ref.~\\cite{Kimura}, for $\\alpha=\\beta=0$, the matter density perturbation $\\delta_{m}$ divided by the scale factor $a$ is anti-correlated with the effective gravitational potential $\\Phi_{{\\rm eff}}$ for the modes relevant to the LSS. This leads to the anti-correlation between the LSS and the ISW effect in CMB, so that the parameter $s$ for the tracker is severely constrained to be $s<1.2\\times10^{-4}$ (95\\,\\% CL).  For the models with $\\alpha\\neq0$ and $\\beta\\neq0$, however, the correlation between $\\delta_{m}/a$ and $\\Phi_{{\\rm eff}}$ depends on the values of $\\alpha$ and $\\beta$. If $p=1$ and $q=5/2$ (i.e. $s=0.2$), for example, $\\delta_{m}/a$ and $\\Phi_{{\\rm eff}}$ tend to be positively correlated for the model parameters close to the border (b) in Fig.~\\ref{validfig} ($\\beta=\\alpha/2-1/15$), whereas they show anti-correlations for $\\alpha$ and $\\beta$ close to the border (a) ($\\beta=\\alpha/2$). The typical examples of the positive and negative correlations are plotted in Figs.~\\ref{pofig} and \\ref{anfig}, respectively. The qualitative differences between these two cases can be understood by estimating the effective gravitational coupling $G_{{\\rm eff}}$ and the quantity $\\xi=(G_{{\\rm eff}}/G)(1+\\eta)/2$ derived under the quasi-static approximation on sub-horizon scales. As the model parameters approach the border (a) in Fig.~\\ref{validfig}, $\\xi$ gets larger while $G_{{\\rm eff}}$ decreases, so that the anti-correlation between $\\delta_{m}/a$ and $\\Phi_{{\\rm eff}}$ tends to be stronger. We studied the evolution of perturbations for different values of $s$ in the range $0<s<0.36$ and found that the basic properties for the ISW-LSS correlation are similar to those discussed for $s=0.2$.  We also estimated the growth index $\\gamma$ of the matter perturbation and found that, for $p=1$ and $0<s<0.36$, it typically varies in the range $0.4\\lesssim\\gamma\\lesssim0.6$ at the redshifts for $0\\le z\\lesssim1$. Hence it is also possible to distinguish between the extended Galileon models and the $\\Lambda$CDM model from the galaxy clustering. However, we expect that the tightest observational bounds on the values $\\alpha$ and $\\beta$ should come from the ISW-LSS correlation. We leave such observational constraints for future works."
    },
    "1112/1112.1404_arXiv.txt": {
        "abstract": "The detection of gravitational waves from the inspiral of a neutron star or stellar-mass black hole into an intermediate-mass black hole (IMBH) promises an entirely new look at strong-field gravitational physics.  Gravitational waves from these intermediate-mass-ratio inspirals (IMRIs), systems with mass ratios from $\\sim$10:1 to $\\sim$100:1, may be detectable at rates of up to a few tens per year by Advanced LIGO/Virgo and will encode a signature of the central body's spacetime.  Direct observation of the spacetime will allow us to use the ``no-hair'' theorem of general relativity to determine if the IMBH is a Kerr black hole (or some more exotic object, e.g. a boson star).  Using modified post-Newtonian (pN) waveforms, we explore the prospects for constraining the central body's mass-quadrupole moment in the advanced-detector era.  We use the Fisher information matrix to estimate the accuracy with which the parameters of the central body can be measured.  We find that for favorable mass and spin combinations, the quadrupole moment of a non-Kerr central body can be measured to within a $\\sim15\\%$ fractional error or better using 3.5 pN order waveforms; on the other hand, we find the accuracy decreases to $\\sim 100\\%$ fractional error using 2 pN waveforms, except for a narrow band of values of the best-fit non-Kerr quadrupole moment. ",
        "introduction": "Within the coming decade, the field of gravitational wave astrophysics will become an effective tool for exploring the universe.  As gravitational waves are emitted from any system with a time dependent quadrupole moment (such as the inspiral and merger of two compact objects), this new field will provide us with an insight into the universe independent of conventional electromagnetic observations.  For the first time, we will be able to directly detect the coalescence of compact objects, such as neutron stars (NS) and black holes (BH). Of particular interest is the information about strong-field gravitational physics that these gravitational waves contain.  The analysis of the final moments of binary inspiral will provide us with the first direct tests of general relativity in the strong-field regime \\cite{CutlerThorneReview}.  With the LIGO and Virgo gravitational wave observatories scheduled to reach their advanced sensitives \\cite{AdvLIGO, AdvVirgo} around 2015, a large effort is underway to develop techniques capable of estimating the parameters of a coalescing compact binary composed of neutron stars and/or black holes.  Most of this work has been geared towards the detection of inspirals \\jontext{of}{in which the two components are of} approximately equal mass, with the total mass of the system less than $\\approx 35 M_{\\odot}$ \\cite{LSC4,LSC5}; however, more recent searches in the high-mass regime have \\jontext{probed}{looked} for inspirals with total mass \\jontext{from}{in the range} $25M_{\\odot} \\leq M \\leq 100M_{\\odot}$ and component masses \\jontext{from}{in the range}  $1M_{\\odot} \\leq m_1\\text{, }m_2 \\leq 99M_{\\odot}$, searching for more extreme mass ratios \\cite{HighMass}.  These systems emit gravitational waves in the frequency range best suited for detection by Advanced LIGO, the peak sensitivity of which will lie at $\\sim 100$ Hz.  Both observational evidence and theoretical models suggest the existence of intermediate-mass black holes (IMBHs) with masses ranging from under a hundred to ten thousand solar masses (see \\cite{MillerColbert:2004} for a review).  While the recently discovered ultra-luminous X-ray source in ESO 243-49 may turn out to be the first confident IMBH detection \\cite{ObserveIMBH}, many questions about IMBH formation mechanisms and their prevalence in dense stellar environments remain unanswered.  Gravitational waves emitted during the intermediate-mass-ratio inspirals (IMRIs) \\jontext{}{of stellar-mass compact objects into IMBHs} will be in the observable frequency band of Advanced LIGO \\cite{BrownIMRI} and dynamics studies suggest possible IMRI detection rates of up to a few tens per year \\cite{IlyaRates}.  The gravitational waveforms generated by these inspirals will encode information about the higher-order multipole moments of the central body's spacetime.  Since the ``no-hair'' theorem predicts that all the higher-order multipole moments of a Kerr black hole can be expressed as a function of its mass and spin, a simultaneous measurement of the mass, spin, and mass-quadrupole moment would serve as a test of the null hypothesis that massive compact objects are Kerr black holes and inspirals are governed by general relativity.  In this paper, we describe the potential for constraining the mass-quadrupole moment of the central body in an IMRI.  This idea was first rigorously developed by Ryan, who showed that the inspiral of a nearly circular, nearly equatorial extreme mass ratio binary encodes information about \\textit{all} the higher order mass and spin moments of the spacetime \\cite{RyanMultipole}.  While Ryan studied the detectability of all higher order multipole moments \\cite{RyanDetection}, more recent work has focused on isolating the mass-quadrupole moment (IMRIs in LIGO by Brown et al. \\cite{BrownIMRI}, and extreme-mass-ratio inspirals in LISA by Barack and Culter \\cite{BCEMRI}).  Whereas previous work focused on \\jontext{}{deviations about} the null case of a Kerr spacetime, we offer the first exploration of the quadrupole moment parameter space, looking at non-Kerr objects and comparing their detectability to that of the null case.   For simplicity, we restrict ourselves to the \\carl{quasi-circular, single-spin} case where the orbital and IMBH angular momenta are aligned.  Using the Fisher matrix formalism, we \\jontext{measure the upper limit uncertainties on}{compute the precision of} parameter estimation \\jontext{given the}{that will be possible using} Advanced LIGO \\jontext{design noise}{operating at its design} sensitivity.  By including an additional term in the waveform representing deviations from the theoretically predicted quadrupole moment, we \\jontext{will}{find that we should} be able to constrain any deviation of the quadrupole moment with a fractional error of $\\sim 1-10$ when the best-fit parameters are those of a Kerr central body, and to $\\sim 15\\%-300\\%$ or better for certain non-Kerr central bodies, depending on our choice of pN order and central-body mass and spin.  In some cases, this will make it possible to discern between a Kerr black hole and some more exotic, horizon-less object, such as a boson star \\cite{RyanBoson}.  If the object \\textit{is} a Kerr black hole, we will be able to \\jontext{determine how well}{demonstrate that} it matches the predictions of general relativity \\jontext{}{to the same precision}.  In section~\\ref{IMRIrate}, we review the current theoretical understanding of IMBH formation and IMRIs, with a particular \\jontext{eye to}{focus on} detection rates in Advanced LIGO.    \\jontext{Section~\\ref{IMRIwave} we cover}{In section~\\ref{IMRIwave}  we describe} the frequency domain waveforms and their numerical implementation in the Fisher information matrix (FIM).  Section~\\ref{results}  reports the results of two cases: the null case (a Kerr IMBH), and the anomalous case (in which the best-fit quadrupole moment deviates from the no-hair Kerr IMBH).  We also discuss technical issues arising from the FIM, as well as our use of the post-Newtonian waveforms for IMRI systems.  Unless otherwise stated we use geometrized units, with $G=c=1$. ",
        "conclusions": "We summarize here the key results of this paper.  Using frequency domain post-Newtonian waveforms, we tested the possibility of measuring the quadrupole moment of a massive, spinning compact body as a parameter independent of the object's mass and spin.  This served two distinct purposes:  first, a simultaneous measurement of the mass, spin, and quadrupole moment of a Kerr black hole would serve as a test of the no-hair theorem of general relativity, acting as a check on the internal consistency of the theory.  Secondly, the detection of an off-Kerr quadrupole moment would suggest the existence of some new, exotic type of star or other supermassive object, allowing Advanced LIGO to identify possible new forms of matter in the Universe.  For this analysis, we first explored whether we would be able to rule out significant deviations from a Kerr black hole in the case when the best-fit parameters correspond to $Q_{\\text{anom}} = 0$. For this null test, we found that the ability of Advanced LIGO to measure an anomalous quadrupole moment will depend heavily on the parameters of the system.  The ideal systems are medium mass, highly spinning systems with an ISCO frequency close to $250$Hz, for which the final moments of the inspiral signal lie at frequencies in the highest-sensitivity range of the detector.  Additionally, since systems with more extreme mass ratios spend more orbital cycles in the strong-field of the central body, their inspirals will encode more information about the quadrupole moment of the spacetime.  It is the interplay between these two effects, the ISCO frequency and the mass ratio, that determines the best system for these tests.  The optimal system also depends on the pN order of the waveform used, due to the increased mass precision at higher orders and the correlation between the masses and the quadrupole moment of the central body.  If the central body has a non-zero quadrupole moment, the fractional uncertainty with which $Q_{\\text{anom}}$ can be measured varies as the $\\mathcal{M}_c$-$\\eta$-$Q_{\\text{anom}}$ error ellipse is slowly deformed.  The error has a local minimum at $Q_{\\text{anom}}^{\\text{(min)}}$, representing an error ellipse who principal axes nearly align with the $\\mathcal{M}_c$-$\\eta$-$Q_{\\text{anom}}$ coordinates, and increases linearly with $|Q_{\\text{anom}}-Q_{\\text{anom}}^{\\text{(min)}}|$.  The slope of this linear effect is sufficiently small that certain systems could be easily identified to within 15\\% accuracy for large values of $Q_{\\text{anom}}$.  If one of these near-optimal systems was observed, Advanced LIGO would be able to definitively rule out that it was a Kerr black hole.  \\jon{The space-based gravitational wave detector (e)LISA will also have the capability to measure anomalous quadrupole moments of black holes. Barack and Cutler~\\cite{BCEMRI} showed that a LISA observation of an EMRI with a signal-to-noise ratio of $100$ would be able to detect a deviation in the quadrupole moment of $10^{-4}$--$10^{-2}$, depending on the system parameters, when the central object is a Kerr black hole. If the effect we described in this paper, that the quadrupole moment can be measured more accurately when it is not equal to the Kerr value, also holds for EMRI systems, (e)LISA could potentially do even better. However, that analysis needs to be revisited in light of the recent redesign of the eLISA mission~\\cite{LISAdescope}, which has led to a change in the frequency sensitivity of the mission and has meant that an SNR of $100$ is now rather optimistic for EMRI sources. eLISA has the potential to provide more precise constraints than Advanced LIGO because EMRI systems will be observed for many more cycles in the strong-field regime than IMRI systems. However, even optimistically, eLISA will not be producing data until $5$--$10$ years after Advanced LIGO. The results in this paper indicate that Advanced LIGO could provide the first hints for departures from the Kerr metric, which would then be confirmed with more precision by future detectors including eLISA.}  The primary limitations in the study presented in this paper were the lack of accurate  waveforms covering the IMRI mass regime and the limitations of the FIM approach.  In order to properly study the parameter space, we would require a realistic signal from an IMRI template, studied with a \\jontext{more Bayesian complete}{Bayesian} technique (e.g.\\ Markov Chain Monte Carlo).  The primary limitation on our ability to measure $Q_{\\text{anom}}$ is the correlation in the mass/quadrupole subspace. We would therefore expect to measure $Q_{\\text{anom}}$ more precisely if we were able to measure the masses more precisely. We have used a pN approximant waveform in this study, so we may be missing pieces of the frequency evolution that would help break this degeneracy. A further study using a more realistic waveform model would be required to investigate if we could do better in practice than the results described here. \\carl{Current work is underway in developing such waveforms via the ``effective one-body -- numerical relativity'' technique (see \\cite{PanEOBNR}), though these studies have yet to produce fully spinning IMRI waveforms in the strong-field regime.}  However, even using this simplified waveform model, we find that we would be able to differentiate between black holes, boson stars, and naked singularities in the Advanced LIGO era, which would represent a significant achievement for compact-object physics.  \\carl{Assuming IMRI waveforms will be available in the next few years, the realistic measurements may be even more precise than the results quoted here.}  Future work will be needed to quantify the biases introduced by our use of a pN approximant. While we have explored some of these  errors by comparing different pN orders, there also exists the bias introduced by modeling a waveform with $Q_{\\text{anom}}$ present when it does not exist in nature (or the converse).   This work can also be extended to proposed third generation detectors (e.g.\\ the Einstein telescope \\cite{JonET}), which could offer increased sensitivity at the same SNR due to their lower low-frequency cutoff."
    },
    "1112/1112.6224_arXiv.txt": {
        "abstract": "Hydrophobic silica aerogels with ultra-low densities have been designed and developed as cosmic dust capture media for the Tanpopo mission which is proposed to be carried out on the International Space Station. Glass particles as a simulated cosmic dust with 30 $\\mu $m in diameter and 2.4 g/cm$^3$ in density were successfully captured by the novel aerogel at a velocity of 6 km/s. Background levels of contaminated DNA in the ultra-low density aerogel were lower than the detection limit of a polymerase chain reaction assay. These results show that the manufactured aerogel has good performance as a cosmic dust collector and sufficient quality in respect of DNA contamination. The aerogel is feasible for the biological analyses of captured cosmic dust particles in the astrobiological studies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.0035_arXiv.txt": {
        "abstract": "We present a study of the nebular phase spectra of a sample of Type II-Plateau supernovae with identified progenitors or restrictive limits. The evolution of line fluxes, shapes, and velocities are compared within the sample, and interpreted by the use of a spectral synthesis code. The small diversity within the dataset can be explained by strong mixing occurring during the explosion, and by recognising that most lines have significant contributions from primordial metals in the H envelope, which dominates the total ejecta mass in these type of objects. In particular, when using the \\Oi\\ 6300, 6364 \\AA\\ doublet for estimating the core mass of the star, care has to be taken to account for emission from primordial O in the envelope. Finally, a correlation between the \\ha\\ line width and the mass of \\nick\\ is presented, suggesting that higher energy explosions are associated with higher \\nick\\ production. ",
        "introduction": "\\label{intro}  Supernovae (SNe) are responsible for the production of most of the heavy elements that we observe in the Universe today. In a core-collapse SNe (CC-SN), this nucleosynthesis occurs both by hydrostatic burning during the star's evolution and by explosive burning behind the shock after collapse. \\cite{woo95} analysed the yields of elements that are synthesised in massive stars and concluded that stars of 15-25 \\msun\\ are the typical stars responsible for producing heavy elements. However, the initial mass function shows that low mass stars are more common and therefore, depending on stellar yields, stars in the range 8--15 \\msun\\ could also make a non-negligible contribution to the chemical enrichment of the Universe.  Type II-Plateau (IIP) SNe are a subset of CC-SNe, which result from the core-collapse of massive H-rich stars, and therefore display prominent H features in their spectra, along with an extended plateau of nearly constant luminosity in their light curves. They are the most common SN type making up $\\sim$40 per cent of all SNe per unit volume \\citep{sma09}. Stellar evolutionary models predict that stars in the range 8--30 \\msun\\ should explode as red supergiants \\citep{heg03, lim03, eld04, hir04}. However, the progenitors of some IIP SNe have been identified, and apart from  a few rare cases of blue supergiant progenitors \\citep{wal87, pas05, kle11}, they have been found to be red supergiants in a lower than expected mass range of 8--17 \\msun\\ \\citep{sma09}. \\cite{utr08, utr09} have modelled the light curves and spectral evolution of IIP SNe (some with progenitor identifications) and find larger masses of the exploding stars than those from the direct detections of the progenitors.  This discrepancy between the progenitors obtained from direct imaging and those from modelling has also been analysed and discussed in \\cite{sma09} and \\cite{mag10} using a sample of nearby IIP SNe. One way of studying this discrepancy is to investigate and compare the nucleosynthetic yields from IIP SNe with identified progenitors.  We do this by studying a sample of IIP SNe at epochs when we can see emission from newly synthesised heavy elements. During the plateau phase of the light curve, which generally lasts for $\\sim$80--120 d, the light curve is powered by{remission of the energy deposited by the shock wave, including H recombinations. At these times, the H envelope is optically thick and we can not see emission from the metal core. However, after the H has recombined, the envelope becomes transparent and the core becomes visible. The light curve then enters the radioactive decay phase, which is powered by the decay of \\cob\\ to $^{56}$Fe. Radiative transfer effects are still significant in this phase, but decrease with time as the line optical depths fall. On the other hand, after approximately two years, the SNe generally become too faint to obtain good signal-to-noise spectra. Therefore, the optimal window for studying and interpreting the nucleosynthesis that has occurred is between one and two years post explosion.  Very few IIP SNe have been studied in detail during the nebular phase and few late-time nebular spectra of IIP SNe exist, with the notable exception of SN 1987A, which due to its proximity was very well monitored. This lack of late-time spectral data is partly due to the long exposure times that are required to obtain sufficient signal-to-noise. The best observed IIP SNe at the necessary late-times are SN 1987A \\citep[e.g.,][]{phi90, bou93, mei93,wood93}, SN 1990E \\citep{ben94}, SN 1999em \\citep{leo02, elm03a}, SN 2004et \\citep{sah06, mag10}, SN 2007it \\citep{and11}, and SN 2007od \\citep{and10, ins11}.   It is generally assumed that the emission lines of specific elements are emitted from regions containing predominantly that particular element, i.e.~that the O zone (formed in the late stages of pre-SN evolution or during the SN explosion) would be responsible for the production of the O emission lines. However, the final H envelope mass for an 8--12 \\msun\\ star is $\\sim$6--8 \\msun\\ \\citep[e.g.][]{hir04}, much larger than the mass of synthesised metals of $\\sim$0.5 \\msun\\ and it is not clear if emission by primordial metals in the envelope may dominate that of synthesised ones. This has been shown by \\cite{li93} to be the case in SN 1987A for the \\Caii\\ 7291, 7323 \\AA\\ and the \\CaiinF\\ 8600 \\AA\\ triplet, which are produced by a primordial abundance of Ca in the inner (velocity $<$ 2500 km s$^{-1}$) region of the H and He envelope and not by newly synthesised Ca. Ca emission lines are very efficient cooling lines, and so even small quantities of primordial Ca can result in prominent emission lines in the late-time spectra \\citep{fra89}. To obtain the observed flux of the Ca lines, \\cite{li93} needed only a Ca mass of 2$\\times10^{-4}$ M$_{\\odot}$ in $\\sim$5 \\msun\\ of H in the inner envelope of SN 1987A, a factor of 10 times lower than that typically predicted for newly synthesised Ca in a SN 1987A-type progenitor star.  Using models that included realistic mixing of the nuclear burning zones, \\cite{koz98} showed that for SN 1987A, the \\CaiinF\\ 8600 \\AA\\ triplet was at all times formed by emission from primordial Ca in the envelope, whereas the \\Caii\\ 7291, 7323 \\AA\\ lines were formed by both synthesised and primordial Ca until $\\sim$400 d, but mainly by primordial Ca after that. This does not imply that there is no newly synthesised Ca but instead that if it is present at late times, that the zones in which it is formed are too small and it cannot capture enough of the radioactive energy to create any strong emission, unless it is mixed with a greater mass of other elements such as H and He \\citep{li93}.  Therefore, to calculate the relative emission strengths from primordial and synthesised material, the deposited energy in the different zones, and the dominant processes responsible for the production of these lines must be modelled using a spectral synthesis code. Then the synthesised mass of O, Ca, and Fe can be estimated by first subtracting the emission contribution from primordial material in the envelope. Previous studies of samples of CC-SN nebular phase spectra have mainly focused on stripped envelope SNe (Type Ib (lacking H) and Type Ic (lacking H and He)) \\citep[e.g.][]{mae08, mod08, tau09, mil10}, and have studied the \\Oi\\ 6300, 6364 \\AA\\ profiles to look for possible signs of ejecta asphericity in the line profile shapes and how they evolve with time.  \\cite{elm11} have studied a sample of  CC-SNe including IIP SNe as well Type Ib/c SNe, focusing again on an analysis of the \\Oi\\ 6300, 6364 \\AA\\ doublet. Using the overall luminosity of the \\Oi\\ doublet and the ratio of the two components, they have estimated the synthesised O mass, assuming all the O luminosity is produced by freshly made O and not from primordial O in the envelope. We will in this paper show that this is a questionable assumption.  In this paper, we study a sample of 25 late-time nebular phase spectra of nine IIP SNe with identified progenitors to investigate how their spectral features change as a function of time, and how the flux, profile shape and velocity of the features vary among SNe. This can then be linked to the material distribution and physical conditions in the SNe. We will apply a spectral synthesis model to the nebular phase spectra to study line flux ratios and interpret their evolution with time as well as their internal differences. This paper is complemented by Jerkstrand et al. (in prep.), which derives detailed nucleosynthesis results for one of the SNe in the sample, SN 2004et, by a complete analysis of the spectral evolution from the ultra-violet to far-infrared. ",
        "conclusions": "In this paper we have investigated the nebular phase spectra of a sample of IIP SNe and compared their observational analysis to a detailed radiative transfer model, to try to better understand the formation of their emission lines and the variation within the sample. We have focused on IIP SNe with either progenitor detections or restrictive limits on their progenitor masses.  Despite the large variation in \\nick\\ masses and expansion velocities, the characteristics of the emission lines in the nebular spectra of IIP SNe are found to be relatively homogeneous and independent of progenitor masses. For each SN, the shapes of the different emission line profiles are very similar and suggest a similar spatial distribution of the different elements. This is also seen when the velocities of the prominent emission lines are measured using their HWHM, with only \\ha\\ showing a noticeably higher velocity, which suggests a larger contribution from farther out in the ejecta.  For SN 1987A, as has been previously shown, we have shown that the \\ha\\ velocity was significantly larger than the \\Oi\\ velocity, while for the other SNe in our sample the differences between H and \\Oi\\ were not as pronounced. This suggests that the zones responsible for \\Oi\\ emission are different in the rest of the SN sample compared to those in SN 1987A. From analysis of the line profiles and widths, it appears that all IIP SNe (except SN 1987A) have similar spatial distributions of \\Oi\\ which is closer in velocity to \\ha\\ than for SN 1987A. This implies that the \\Oi\\ flux could be predominantly emitted by primordial O in the H zones and may not be coming from the synthesised O from the inner regions of the ejecta. The line profiles are also shown to be intrinsically peaked in shape, even after taking into account the convolution with the instrumental resolution. This shows that mixing must have occurred for all the elements to be producing emission from both the centre and the outer regions of the ejecta.  These observational results are corroborated by the output of the radiative transfer model of a 12 \\msun\\ star,  where a significant proportion of the flux of the prominent emission lines is found to originate in the H-rich zones.  For SN 1987A, the \\Caii\\ and several \\FeiiF\\ emission lines were also found to be predominantly formed in the H zones, but here we present the new result that for the 12 \\msun\\ progenitor model used here, the \\Oi\\ flux is produced in approximately equal amounts by O in the H zones and O in the inner O-rich zones, and this explains the much smaller difference between the \\ha\\ and \\Oi\\ velocities seen for our sample of SNe. This has major implications for current methods of calculating the synthesised O mass in IIP SNe, which assume that the \\Oi\\ 6300, 6364 \\AA\\ line emission is produced solely by synthesised O. A calculation of the synthesised O mass of SN 2004et, that takes into account these new results will be presented in Jerkstrand et al.~(in prep.).  The flux ratios of the emission lines from the 12 \\msun\\ model were found in general to agree quite well with the observed flux ratios. No trends were found between progenitor mass and the observed flux ratios, in agreement with the result that the flux in these lines has a significant contribution from primordial abundances. We also investigated possible correlations between the spectral properties and other observational properties of the SNe and found a correlation between the ejected \\nick\\ mass and FWHM of the \\ha\\ emission line. An empirical relation was obtained, which can be used to estimate the \\nick\\ mass of SNe with limited bolometric coverage and/or poor constraints on their extinction and distance. The benefit of this method is that it requires only the width of the \\ha\\ emission line during the nebular phase to estimate the \\nick\\ mass. Even more importantly, this provides a strong constraint on explosion models, particularly the amount of fall-back material in them."
    },
    "1112/1112.5426_arXiv.txt": {
        "abstract": "We continue our exploration of whether the flyby anomalies can be explained by scattering of spacecraft nucleons from dark matter gravitationally bound to the earth, with the addition of data from five new flybys to that from the original six. We continue to use our model in which inelastic and elastic scatterers populate shells generated by the precession of circular orbits with normals tilted with respect to the earth's axis. With 11 data points and 8 parameters in the model, a statistically meaningful fit is obtained with a chi-squared of 2.7.  We give plots of the anomalous acceleration along the spacecraft trajectory, and the cumulative velocity change, for the five flybys which exhibit a significant nonzero anomaly.  We also discuss implications of the fit for dark matter-nucleon cross sections, give the prediction of our fit for the anomaly to be expected from the future Juno flyby, and give predictions of our fit for flyby orbit orientation changes. In addition we give formulas for estimating the flyby temperature increase caused by dark matter inelastic scattering, and for the fraction of flyby nucleons undergoing such scatters. Finally, for circular satellite orbits, we give a table of predicted secular changes in orbit radius.  These are much too large to be reasonable -- comparing with data for COBE and GP-B supplied to us by Edward Wright (after the first version of this paper was posted), we find that our model predicts changes in orbit radius that are too large by many orders of magnitude. So the model studied here is ruled out.  We conclude that further modeling of the flyby anomalies must simultaneously attempt to fit constraints coming from satellite orbits. ",
        "introduction": "In this paper we follow up our earlier investigations \\cite{adler1}, \\cite{adler2} of the anomalous geocentric frame orbital energy changes that are observed during earth flybys of various spacecraft, as reported by Anderson et al. \\cite{anderson}.   Some flybys show energy decreases, and others energy increases, with the largest anomalous velocity changes of order 1 part in $10^6$.  While the possibility that these anomalies are artifacts of the orbital fitting method used in \\cite{anderson} is still being actively explored, there is also a chance that they may represent new physics.  In \\cite{adler1} we explored the possibility that the flyby anomalies result from scattering of spacecraft nucleons from dark matter particles in orbit around the earth, with the observed velocity decreases arising from elastic scattering, and the observed velocity increases arising from exothermic inelastic scattering, which can impart an energy impulse to a spacecraft nucleon.  Many constraints on this hypothesis were analyzed in \\cite{adler1}, with the conclusion that the dark matter scenario is not currently ruled out, but requires dark matter to be non-self-annihilating, with the dark matter scattering cross section on nucleons much larger, and the dark matter mass much lighter, than usually assumed.  These constraints on the dark matter scattering cross section and mass will be revisited, and somewhat weakened, in the present paper.  In \\cite{adler2} we constructed a model for the spatial and velocity distribution functions for dark matter particles in earth orbit, based on assuming two populations of dark matter particles, one of which scatters on nucleons elastically, and the other of which scatters inelastically, with each population taken to fill a shell-like distribution of  orbits generated by the precession of a tilted circular orbit around the earth's rotation axis.  Fit of this model to the data were presented, but since the model has 8 parameters, and there were only 6 data points in the original flyby set, the statistical significance of the fits was doubtful.  After \\cite{adler2} was written, Anderson and Campbell \\cite{andersoncampbell}  furnished us with data for 5 additional flybys, all of which show no velocity anomaly.  In the present paper, we add these 5 flybys to the original 6 and  fit the model of \\cite{adler2} to all of the data points.  We find that with 11 data points and 8 parameters, we get a statistically meaningful fit with a chi-squared of 2.7.  In this fit, the dark matter shells lie much closer to earth than they did in the fits exhibited in \\cite{adler2}.  The formulas defining the model were given in Sec. II of \\cite{adler2} (with further detail in the Appendices of the arXiv version), and will not be repeated here. So reading the earlier paper is a necessary prerequisite for fully understanding the model analyzed in this paper. Our focus here is on presenting the orbital parameters of the new flybys and the results of the fit to the full 11 flyby set.  In addition to giving the fits to the total velocity anomaly, we give the contributions of the shells of elastic and inelastic scatterers to this total. We also plot the anomalous acceleration along the spacecraft trajectory and the cumulative velocity change, for the five flybys which show significant nonzero velocity anomalies,  give constraints on dark matter-nucleon cross sections implied by our fit, and give formulas for estimating the flyby temperature change and the fraction of nucleons that undergo exothermic dark matter scatters.  For the future Juno flyby, our 11 flyby fit predicts an anomaly of 11.6 mm/s.  We conclude by giving predictions of our model for flyby orbit orientation anomalies, and for secular changes in orbit radius of satellites in circular orbits. The latter, obtained from a program that can treat non-circular orbits as well, shows that the predicted orbit radius changes for COBE and GP-B are orders of magnitude larger than observational data supplied to us by Edward Wright after the first version of this paper was posted to arXiv: see the Added note to Section VII. Hence the model studied here is ruled out, and any further fitting of more general forms of the dark matter model, or of other types of models, to the flyby anomalies, must simultaneously attempt to fit constraints coming from closed satellite orbits.  This greatly complicates the analysis, and we shall pursue this further only if the future Juno flyby, which has orbital parameters similar to GLL-I, confirms the presence of a kinetic energy anomaly. ",
        "conclusions": "In \\cite{adler2} we gave an estimate for the total mass $M_e$ and $M_i$  in the elastic and inelastic dark matter shells, \\begin{align}\\label{eq:mass} M_e\\simeq &4 \\pi^{5/2} \\rho_eD_e m_1/\\sigma_{\\rm el}~~~,\\cr M_i \\simeq & 4 \\pi^{5/2} \\rho_iD_i m_1/(B_{\\rm inel}\\sqrt{2 \\Delta m/m_2'}~)~~~,\\cr \\end{align} where in $M_i$ we have allowed for the possibility that $\\Delta m$ is much smaller than $m_2'$  (rather than $\\Delta m \\sim m_2'$, as assumed in \\cite{adler2}). Here $\\sigma_{\\rm el}$ is the threshold cross section for the scattering of the elastic dark matter population on nucleons,  $B_{\\rm inel}$ is the coefficient of the $\\cos \\theta$ term in the near threshold cross section for exothermic scattering of the inelastic dark matter population on nucleons (see Eq. (7) of \\cite{adler2}), and $m_1$, $m_2'$, and $\\Delta m$ are as defined in Section VI and Appendix B. Using these estimates, and the upper bound \\cite{adler3}  of $4 \\times 10^{-9} M_{\\oplus} \\sim 1.4 \\times 10^{43} {\\rm GeV}/c^2$ on the mass of dark matter in orbit around the earth between the 12,300 km radius of the LAGEOS satellite orbit and the moon's orbit, one can get lower bounds on $\\sigma_{\\rm el}$ and $B_{\\rm inel}$.  For example, referring to Table V, from  $\\rho_eD_e=4.537\\, {\\rm km}^2$  we get \\begin{equation}\\label{eq:bound1} \\sigma_{\\rm el} \\geq 2.2\\times 10^{-31} {\\rm cm}^2~~~, \\end{equation} while from $\\rho_iD_i=0.0000811\\,{\\rm km}^2$, taking into account the fact that only a fraction \\begin{equation}\\frac{1}{\\surd \\pi} \\int_{L}^{\\infty} dz \\exp(-z^2) \\simeq 0.50 \\times 10^{-3} \\end{equation} \\big(with $L=(12,300-7561)/2038=2.325$\\big) of the inelastic dark matter shell lies above the LAGEOS orbit, we correspondingly get \\begin{equation}\\label{eq:bound2} B_{\\rm inel}\\sqrt{\\Delta m/m_2'}   \\geq 1.4 \\times 10^{-39} {\\rm cm}^2~~~. \\end{equation} These bounds are consistent with the cross section range arrived at from various constraints in \\cite{adler1}.  In comparing these bounds with cross section limits derived from dark matter direct detection experiments, two points should be kept in mind.  The first is that if the dark matter mass is much lower than a GeV, the recoils from elastic scattering on nucleons of galactic halo dark matter will be below the threshold for detection in direct searches, and so no useful constraints on the elastic dark matter nucleon scattering cross section are obtained.  The second is that with two species of dark matter, as in our model, the fraction $f_i$ of the galactic halo consisting of inelastically scattering dark matter can be much smaller than unity.  In this case a direct detection cross section bound $\\sigma_{\\rm MAX}$, obtained assuming a single species of dark matter in the galactic halo, implies a significantly larger bound $\\sigma_{\\rm MAX}/f_i$ for the dark matter-nucleon scattering cross section of the inelastic component.     As opposed to the fits given in \\cite{adler2}, the current fit to 11 flybys places the elastic and inelastic dark matter shells much closer to earth, well within the orbits of geostationary and Global Positioning System satellites.  So the effect on these satellites should be minimal, but it now becomes important to examine the implications of our model for satellites that are in low or medium earth orbit, as initiated  in Section VII and Table IX. (See in this regard the Added note in Section VII.)    It will also be important to evaluate  implications of a low earth orbit exothermic scattering dark matter shell for earth's heat balance.  We pose these as significant questions to be addressed in future investigations."
    },
    "1112/1112.4750_arXiv.txt": {
        "abstract": "A preceding paper  \\cite{guidAsy} demonstrated that explicit asymptotic methods generally work much better for extremely stiff reaction networks than has previously been shown in the literature. There we showed that for systems well removed from equilibrium explicit asymptotic methods can rival standard implicit codes in speed and accuracy for solving extremely stiff differential equations. In this paper we continue the investigation of systems well removed from equilibrium by examining quasi-steady-state (QSS) methods as an alternative to asymptotic methods. We show that for systems well removed from equilibrium, QSS methods also can compete with, or even exceed, standard implicit methods in speed, even for extremely stiff networks, and in many cases give somewhat better integration speed than for asymptotic methods. As for asymptotic methods, we will find that QSS methods give correct results, but with non-competitive integration speed as equilibrium is approached. Thus, we shall find that both asymptotic and QSS methods must be supplemented with partial equilibrium methods as equilibrium is approached to remain competitive with implicit methods. ",
        "introduction": "Stiff networks of differential equations rather uniformly have  been viewed as requiring special implicit or semi-implicit methods for integration in order to maintain stability while taking reasonably efficient timesteps \\cite{gear71,lamb91,press92,oran05,timmes,hix05}. Purely explicit methods are not competitive in speed for most applications because they are limited by stability criteria to integration timesteps that are far too short. Various asymptotic and steady-state schemes have been proposed to stabilize explicit methods by removing some of their stiffness (overviews may be found in Refs.\\ \\cite{oran05,mott99}). These methods have had some success in moderately stiff systems, but it generally has been concluded that in very stiff systems, such as those encountered in astrophysical thermonuclear networks, asymptotic and steady-state schemes do not work \\cite{oran05,mott99}.  In a preceding paper on asymptotic methods \\cite{guidAsy}, the present paper on quasi-steady-state (QSS) methods, and a following paper on partial equilibrium methods \\cite{guidPE}, we challenge these conclusions and present strong evidence that algebraically-stabilized explicit integration may in fact not only compete with, but in some cases may have the potential to outperform traditional implicit methods, even for the stiffest networks. In this paper we deal specifically with the QSS method and show that, for systems well removed from equilibrium, the QSS approximation can give highly-competitive integration of extremely-stiff systems. ",
        "conclusions": "In this paper we have compared quasi-steady-state (QSS) calculations with asymptotic and implicit calculations for extremely stiff networks and concluded that \\begin{enumerate} \\item QSS and asymptotic methods give similar results, but QSS timesteps are generally at least as large as for asymptotic methods, and can be larger by as much as an order of magnitude in some cases. \\item Both QSS methods and asymptotic methods are uniformly capable of stable, accurate solutions, even for extremely stiff thermonuclear networks, with timesteps that are substantially larger than those for standard explicit methods. The only question then is whether such methods can use large enough integration timesteps to be competitive with implicit methods. \\item As for asymptotic methods \\cite{guidAsy}, QSS methods give integration speeds that compete with or even exceed that for implicit methods in extremely stiff networks as long as the system is well removed from (microscopic) equilibrium, but fail to deliver competitive timesteps in the approach to equilibrium. Solution of this problem will require explicit partial equilibrium methods that we shall discuss in Ref.\\ \\cite{guidPE}. \\end{enumerate} Thus, we find compelling evidence that  quasi-steady-state and asymptotic methods may have significant application in the integration of large networks for even the stiffest systems if they are not close to microscopic equilibrium.  Although these conclusions indicate that asymptotic and QSS methods must be supplemented by partial equilibrium methods to make explicit integration viable across a full range of stiff problems, the results of this paper and those of Ref.\\ \\cite{guidAsy} suggest that the practical utility of the asymptotic and QSS methods alone for application in astrophysics and many other fields may be substantial. As we have seen, there are important, extremely stiff problems for which the system never becomes significantly equilibrated. This is most likely to occur in explosive scenarios, where we expect rapid expansion on a hydrodynamical timescale. The expansion will typically lead to reaction freezeout for those reactions that are strongly temperature-dependent, and in rapidly-changing environments this may occur before the system has had time to establish significant microscopic equilibration. The nova calculation of \\S\\ref{novaExplosions} and the tidal supernova calculation of \\S\\ref{tidalSupernova} are examples of realistic situations where this occurs. For such problems we have presented evidence that quasi-steady-state or asymptotic approximations alone (even without partial equilibrium methods) may provide integration speeds that rival or even substantially exceed those for the best current implicit methods, particularly for larger networks.  Even for problems involving large networks coupled to hydrodynamics where the preceding is not true globally, it will usually be that for many hydrodynamical zones over various time ranges the conditions will not favor equilibration. Thus, at each hydrodynamical timestep it may prove most efficient to integrate the reaction networks for all zones not exhibiting significant reaction-network equilibration using explicit QSS or asymptotic methods.  For those zones exhibiting significant equilibration at a given hydrodynamical timestep, more work will be required to determine whether standard implicit methods such as backward Euler, or explicit asymptotic or QSS augmented by partial equilibrium methods are most efficient. If the latter turns out to be true, it likely will be most useful to integrate all zones with an asymptotic or QSS plus partial equilibrium method, but it could turn out that the most efficient approach is a hybrid reaction network algorithm capable of switching among asymptotic plus partial equilibrium, QSS plus partial equilibrium, and implicit methods as conditions dictate."
    },
    "1112/1112.4799_arXiv.txt": {
        "abstract": "Recently the {\\em Fermi}-LAT data have revealed two gamma-ray emitting bubble-shaped structures at the Galactic center. If the observed gamma rays have hadronic origin (collisions of accelerated protons), the bubbles must emit high energy neutrinos as well.  This new, Galactic, neutrino flux should trace the gamma ray emission in spectrum and spatial extent.  Its highest energy part, above 20--50 TeV, is observable at a kilometer scale detector in the northern hemisphere, such as the planned KM3NeT, while interesting constraints on it could be obtained by the IceCube Neutrino Observatory at the South pole. The detection or exclusion of neutrinos from the {\\em Fermi} bubbles will discriminate between hadronic and leptonic models, thus bringing unique information on the still mysterious origin of these objects and on the time scale of their formation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.1033_arXiv.txt": {
        "abstract": "Origin of extragalactic $\\gamma$-ray background (EGRB) is still a matter of debate. EGRB can either have truly diffuse or unresolved discrete point sources origin. Majority of the Fermi and EGRET detected identified sources are blazar. So, they are expected to be a significant contributors to the EGRB. In order to estimate their contribution to the EGRB one needs to construct their luminosity functions. Here, we examine and construct the evolution and luminosity function of blazars using first LAT AGN catalog. We consider both pure luminosity and pure density evolution models for flat spectrum radio quasars (FSRQs). We also describe the methodology to estimate the contribution from misaligned blazars to the EGRB. ",
        "introduction": "Observational $\\gamma$-ray astronomy got a massive boost after the launch of the Fermi Gamma-ray Space Telescope (Fermi) on 11th June, 2008.  The Large Area Telescope (LAT) onboard Fermi provides an increase in sensitivity by more than an order of magnitude over EGRET and AGILE.  The origin of extragalactic $\\gamma$-ray background (EGRB) is one of the fundamental unsolved problem in astrophysics. EGRB can arise from some diffuse processes like black hole evaporation, large scale structure formation, matter-antimatter annihilation, etc (e.g., \\cite{loeb2000,gabici2003,cohenrujulaglashow1998}). Alternatively, due to the limited instrument sensitivity, unresolved $\\gamma$-ray sources could contribute significantly to the observed EGRB.  Recently, the EGRB has been rederived from Fermi Large Area Telescope (LAT) data \\citep{abdo2010b}. EGRB is consistent with a featureless power law of index $(2.41\\pm 0.05)$. The integrated flux above $100$ MeV is $(1.03\\pm 0.17)\\times 10^{-5}$ photon cm$^{-2}$ s$^{-1}$ sr$^{-1}$.     \\par Fermi already detected $\\sim 1500$ sources in 11 months of observation (First Fermi catalog; \\cite{abdo2010c}) and is expected to detect many more. In contrast to the less than $100$ active galactic nuclei (AGNs) in the EGRET catalog, the first Fermi-LAT AGN catalog \\citep{abdo2010a} consist of $\\sim 600$ blazars. Fermi also detected 11 non-blazar AGN \\citep{abdo2010d}, three normal (M $31$, LMC and Milky Way) and two starburst galaxies (M $82$ and NGC $253$). Thus, Fermi provides an excellent opportunity to understand the origin of EGRB.    Since the dominant class of identified EGRET and Fermi sources are blazars, they are expected to contribute to the EGRB. The contribution from blazars to the EGRB is estimated using two different approaches: \\newline (a) Considering a relationship between the blazar $\\gamma$-ray luminosity with their luminosity at some other wavelength, their $\\gamma$-ray luminosity function is assumed to be the scaled luminosity function at that wavelength (e.g., \\cite{stecker1996,stecker1993,narumoto1, narumoto2,abazajian2010,inoue2011,li2011,cavadini2011,stecker2011}). \\newline (b) Constructing the $\\gamma$-ray luminosity function directly from observed $\\gamma$-ray blazars (e.g., \\cite{chiang1995,chiang1998,deb1,deb2}).  In an earlier work \\citep{deb1,deb2} following a similar approach of \\citet{chiang1998} but with almost twice the number of sources, we constructed the $\\gamma$-ray luminosity functions of flat spectrum radio quasars (FSRQs) and BL Lacs from EGRET-detected sources only. We found strong positive evolution for FSRQs, while BL Lacs did not show any significant evolution. The maximum contribution from blazars to EGRB was found to be $\\sim 20$\\%. However, due to the limited number of sources in the EGRET catalog, the luminosity function and hence, the contribution from blazars was not well constrained. \\citet{mucke} estimated the contribution from FSRQs and BL Lacs to the EGRB in the context of AGN unification paradigm, and found that unresolved blazars (both ``aligned'' and ``misaligned'') contribute $20\\%-40\\%$ to the EGRB. In a similar approach, \\citet{dermer2007} estimated that FSRQs and BL Lacs contribute $\\sim 10\\%-15$\\% and $2\\%-3\\%$ at $1$ GeV respectively. From the first three months of Fermi observation, Abdo et al.~\\cite{abdo2009a} found that FSRQs exhibit strong evolution while BL Lacs show no evolution. Using the 1st Fermi catalog, \\citet{abdo2010e} found from the source count distribution that the point source contribution is $\\sim 16 \\%$ of the EGRB.    \\par  Here, we construct the luminosity function of FSRQs and BL Lacs following the methodology described in \\citet{deb1}. Utilizing the Fermi observed sources we construct a complete source list, apply the ${V}/{V_{max}}$ method to search for the presence of evolution, find parameters by modified ${V}/{V_{max}}$ method, and finally construct the luminosity function using ${1}/{V_{max}}$ method and maximum likelihood estimator.    In the next section we discuss the evolution of blazar source class. We discuss in section 3 the luminosity function construction of FSRQs. The methodology to calculate the misaligned blazars contribution to the EGRB is given in section 4. We conclude with our discussion in section 5. ",
        "conclusions": "Using first Fermi catalog, we investigated the luminosity function and nature of evolution of FSRQs and BL Lacs separately in $\\gamma$-rays. The nature of evolution of FSRQs and BL Lacs as found in this work, are in agreement with earlier works (e.g., \\cite{deb1,chiang1998}). We considered both PLE and PDE models for FSRQs. Both power law and exponential evolution functions have been considered. The luminosity function is considered to be of a broken power law form and the parameters are found using likelihood analysis. From this analysis it is not possible to strongly conclude on the true evolution model. Fig \\ref{fig-f3} suggests that PDE model explains observations marginally better than that of PLE model. In contrast to earlier works of determination of luminosity function from $\\gamma$-ray observations \\citep{deb1,chiang1995,chiang1998}, our derived luminosity function is well constrained. In a recent work, Ajello et al.~\\cite{ajello2011} showed that the redshift of the peak in the number density of FSRQs is luminosity dependent. Hence, they considered luminosity dependent density evolution for FSRQs (similar to that of radio-quiet AGN). They considered a much lower limiting flux ($\\sim 10^{-8}$ ph cm$^{-2}$ s$^{-1}$) and hence, a much larger sample ($186$ FSRQs) than us.  For misaligned blazars, we plan for SED fitting of all detected off-axis sources and then construct $\\xi(\\theta)$ from model. This work is under progress and will be reported elsewhere."
    },
    "1112/1112.3807_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\begin{table}[t] \\begin{center} \\caption{Number of abstracts on ADS containing the word ``wavelet'' for given date ranges.} \\begin{tabular}{cc} \\hline Years & Number of Abstracts \\\\ \\hline Until 1995 & 251 \\\\ 1996 - 2000 & 679 \\\\ 2001 - 2005 & 1221 \\\\ Since 2006 & 1797 \\\\ \\hline \\end{tabular} \\end{center} \\label{tab:adsresults} \\end{table}%  The usage of the wavelet transform in astrophysics has become very popular in recent years. Table \\ref{tab:adsresults} compiles the number of publications on ADS\\footnote{NASA Astrophysics Data System, \\protect\\url{http://adswww.harvard.edu/}} in a given range of years that have the word ``wavelet'' contained in their abstract. Clearly, the usage of wavelets has gained popularity fast. Typical applications for wavelet-transform-based methods are morphological separation of sources in images and noise removal. The success of wavelet based methods in astrophysics is in part due to the fact, that astrophysical data often contains information on different angular or spectral scales. For example, an optical image of a galaxy contains compact, bright stars as well as extended and faint emission from the bulge and spiral arms. Multi-scale methods, such as the wavelet transform, allow to investigate the different scales of an image separately \\citep{5299269}.  The most widely used type of wavelet transformation is the so called undecimated or redundant, isotropic wavelet transformation. This is in part due to the algorithmic simplicity of the method but also because undecimated wavelet transforms have proven to be more efficient for noise removal then their decimated counterpart. Apart from that, they also provide a number of computational advantages when reconstructing an image from a subset of its wavelet coefficients \\citep{Starck:2010uf}.  In this paper we review the basics of denoising based on the undecimated wavelet transformations in Section \\ref{sec:wtdenoise} and present an extension of the wavelet transform to three dimensional data as proposed in \\cite{2009A&A...504..641S}. Section \\ref{sec:impl} describes the implementation of the transform in C++ along with a description of where the implementation departs from the original algorithm. A first application of the algorithm is shown in Section \\ref{sec:application}, where we use the algorithm to implement a source finder and test the performance on simulated H\\,{\\sc i} galaxies. We close the paper with a summary and an outlook on future applications and potential improvements to the algorithm. ",
        "conclusions": "We have shown how 2D-1D wavelet denoising schemes can be used for source finding. Even with very simple post-processing of the denoised data, we set up an efficient source finding pipeline. Especially the robustness of the algorithm seems promising that it will work well on real data, which is certainly the next test to be passed.  Even though the splitting of the wavelet transformation in a 2D and 1D part avoids some of the difficulties that arise with anisotropic sources, it is far from perfect. A better denoising would be the usage of a full 3D curvelet transformation \\citep{Candes:2007wf,2005SPIE.5914..351Y}. This transformation is however computationally much more difficult and demanding on the available hardware. We therefore think that our approach is a good compromise between sensitivity and computational complexity. But with more powerful hardware or more optimized algorithms, denoising by usage of the curvelet transform might become feasible, even for the large data sets we expect from the future radio telescopes.  Furthermore we like to stress that even though this algorithm was developed with H\\,{\\sc i} surveys in mind, it will in principle work for every kind of data that is similar to the data product of such a survey, e.g. other spectral line surveys."
    },
    "1112/1112.0279_arXiv.txt": {
        "abstract": "The negative effective magnetic pressure instability discovered recently in direct numerical simulations (DNS) may play a crucial role in the formation of sunspots and active regions in the Sun and stars. This instability is caused by a negative contribution of turbulence to the effective mean Lorentz force (the sum of turbulent and non-turbulent contributions) and results in formation of large-scale inhomogeneous magnetic structures from initial uniform magnetic field. Earlier investigations of this instability in DNS of stably stratified, externally forced, isothermal hydromagnetic turbulence in the regime of large plasma beta are now extended into the regime of larger scale separation ratios where the number of turbulent eddies in the computational domain is about 30. Strong spontaneous formation of large-scale magnetic structures is seen even without performing any spatial averaging. These structures encompass many turbulent eddies. The characteristic time of the instability is comparable to the turbulent diffusion time, $L^2/\\etat$, where $\\etat$ is the turbulent diffusivity and $L$ is the scale of the domain. DNS are used to confirm that the effective magnetic pressure does indeed become negative for magnetic field strengths below the equipartition field. The dependence of the effective magnetic pressure on the field strength is characterized by fit parameters that seem to show convergence for larger values of the magnetic Reynolds number. ",
        "introduction": "The 11-year solar activity cycle manifests itself through the periodic variation of the sunspot number. Sunspots consist of vertical magnetic fields with a strength of up to $3\\kG$ \\citep[see, e.g.][]{P79,O03}. It is generally believed that these fields continue in a similarly concentrated fashion also beneath the surface in the form of magnetic flux tubes or fibers \\citep{P82}. It was thought that fibral magnetic fields constitute a lower energy state and might therefore be preferred. Tube-like magnetic fields were frequently seen in hydromagnetic turbulence simulations \\citep{Nor92,Bra95,Bra96}, but those tubes were similar to the vortex tubes in hydrodynamic turbulence. Those vortex tubes have typical diameters comparable to the viscous length. Similarly, the aforementioned magnetic tubes have a thickness comparable to the resistive length. For the Sun, however, the resulting tube thickness would be too small to be relevant for sunspots.  At the surface, strong magnetic flux concentrations also form in regions of strong flow convergence, but the size of these regions is too small for sunspots, because sunspots are usually much bigger than a single granular convection cell. In fact, a typical sunspot can have a diameter of some 30 pressure scale heights. The tremendous size of sunspots has therefore been used as an argument that they are not made near the surface, but at much greater depth far away from the surface. At the bottom of the convection zone the convection cells are big enough and could in principle be responsible for producing much bigger flux concentrations. Magnetic flux tubes can also form through the action of shear as shown by various simulations \\citep{Cline,GK11}. While shear is likely an important ingredient of the solar dynamo, it remains unclear whether the resulting magnetic tubes are really able to produce sunspots as a result of them piercing through the solar surface and, more importantly, whether one should consider them being tied to the deep shear layers of the Sun.  In this paper we discuss an alternative scenario in which sunspots are shallow  phenomena that are not anchored at the bottom of the convection zone. Various mechanisms have been discussed. Of particular interest here are mechanisms that are based on the suppression of turbulence by magnetic fields. In the mechanism of \\cite{KM00} it is the suppression of the turbulent heat flux while in the mechanism of \\cite{RK07} it is the suppression of the turbulent hydromagnetic pressure. Both mechanisms lead to a linear large-scale instability in a stratified medium. However, these mechanisms may be of different importance in different layers.  In this study we focus mainly on the second mechanism, which has recently been studied in direct numerical simulations (DNS) as well as in mean-field calculations. This mechanism is called the  negative effective magnetic pressure instability (or NEMPI for short). NEMPI is a convective type instability that is similar to the interchange instability in plasmas \\citep{T60,N61,P82} and the magnetic buoyancy instability \\citep{P66}. However, the free energy in interchange and magnetic buoyancy instabilities is due to the gravitational field, while in NEMPI it is provided by the small-scale turbulence. NEMPI is caused by the suppression of turbulent hydromagnetic pressure by the mean magnetic field. When fluid Reynolds numbers is larger than 1 and the mean magnetic field is less than  the equipartition magnetic fields, the negative turbulent contribution to the mean Lorentz force is enough large so that the effective mean magnetic pressure (the sum of turbulent and non-turbulent contributions) become negative \\citep{KRR90,RK07}. This is the main reason for the excitation of the large-scale instability that results in formation of large-scale inhomogeneous magnetic structures.  The effect of the suppression of the turbulent heat flux has not yet been studied as extensively as NEMPI since the original paper of \\cite{KM00}, although \\cite{KO06} have used that model to study the decay of sunspots. Moreover, there is now some evidence that the effects anticipated by \\cite{KM00} may have already been in operation in various simulations of solar convection where the spontaneous formation of pores has been seen \\citep{Stein}. Another example is the large-eddy simulation of \\cite{KKWM10} where an imposed vertical large-scale magnetic field gets concentrated into giant vortices. This result is reminiscent of that of \\cite{Tao98}, who found a segregation into magnetized and nearly unmagnetized regions in stratified convection simulations. However, it has not yet been possible to obtain the large-scale magnetic structures resembling the sunspots, in models with a strong imposed magnetic flux tube structure at the bottom of the domain \\citep[see, e.g.][]{Rem09a,Rem09b,Cheung,Rem11a,Rem11b}.  In the rest of this paper, we focus on NEMPI, which was first found in mean-field calculations of a stratified layer \\citep{KMR96,RK07,BKR10,BKKR11}. However, those results remained unconvincing until NEMPI was also discovered in DNS \\citep[][hereafter BKKMR]{BKKMR}. It it therefore most appropriate to begin our discussion with the latter. ",
        "conclusions": "In the present paper we have performed detailed investigations of NEMPI detected recently by BKKMR. Most notably, we have extended the values of the scale separation ratio, $\\kf/k_1$, from 15 to 30. In this case, the spontaneous formation of magnetic structures becomes particularly evident and can clearly be noticed even without any averaging. Whether or not the particular structures seen in DNS really have a correspondence to phenomena in the Sun, cannot be answered at the moment, because our model is still quite unrealistic in many respects. For example in the Sun, $\\kf$ and $\\urms$ change with depth, which is not currently taken into account in DNS. Also, of course, the stratification is not isothermal, but convectively unstable. However, DNS in turbulent convection by \\cite{KBKMR11} have shown that $\\Peff(\\beta)$ still has a negative minimum in that case.  With regards to the production of sunspots, it is likely that NEMPI will shut off before the magnetic energy density has reached values comparable with the internal energy of the gas, as is the case in sunspots. Thus, some other mechanism is still needed to push the field of flux concentrations into that regime. One likely candidate is the mechanism of KM where the suppression of convective heat flux by the magnetic field is crucial. This impression is further justified by resent calculations of \\cite{Stein} where pores are seen to form spontaneously in a simulation where horizontal magnetic fields are injected at the bottom of the domain.  Pores are small sunspots, so something else is needed to make these structures bigger and to amplify this mechanism further. Again, the answer could be related to scale separation. Thus, it will now be important to undertake detailed investigations of instabilities in strongly stratified layers with finite heat flux and finite magnetic field.  \\begin{acks} We acknowledge the NORDITA dynamo programs of 2009 and 2011 for providing a stimulating scientific atmosphere. Computing resources provided by the Swedish National Allocations Committee at the Center for Parallel Computers at the Royal Institute of Technology in Stockholm and the High Performance Computing Center North in Ume{\\aa}. This work was supported in part by the European Research Council under the AstroDyn Research Project No.\\ 227952. \\end{acks}"
    },
    "1112/1112.3750_arXiv.txt": {
        "abstract": "We present the results of a multi-epoch observational campaign on the mini-broad absorption line quasar (mini-BAL QSO) PG 1126-041 performed with \\textit{\\textit{XMM-Newton}} from 2004 to 2009. Time-resolved X-ray spectroscopy and simultaneous UV and X-ray photometry were performed on the most complete set of observations and on the deepest X-ray exposure of a mini-BAL QSO to date. Complex X-ray spectral variability, found on time scales of both months and hours, is best reproduced by means of variable and massive ionized absorbers along the line of sight. In the highest signal-to-noise observation, highly-ionized X-ray absorbing material outflowing much faster ($\\upsilon_{out}\\sim 16500$ km s$^{-1}$) than the UV absorbing one ($\\upsilon_{out}\\sim 5000$ km s$^{-1}$) is detected. This highly-ionized absorber is found to be variable on very short time scales of a few hours. ",
        "introduction": "Accretion disk winds are among the most promising physical mechanisms able to link the small and the large-scale phenomena in active galactic nuclei (AGN), to shed light on the physics of accretion/ejection around supermassive black holes (SMBHs), and to help understanding the impact of the AGN phase on the host galaxy evolution \\citep[e.g.][]{1998A&A...331L...1S, 2004ApJ...616..688P, 2005Natur.433..604D, 2009MNRAS.398...53B}.  Such winds are currently directly observed, as blueshifted and broadened absorption lines in the UV and X-ray spectra of a substantial fraction of AGN. In the UV band we observe broad absorption lines (BALs, FWHM $> 2000$ km s$^{-1}$), mini-broad absorption lines (mini-BALs, 500 km s$^{-1}$ $<$ FWHM $<$ 2000 km s$^{-1}$), and narrow absorption lines (NALs, FWHM $< 500$ km s$^{-1}$) in about 30\\% of AGN, and their intrinsic fraction is estimated to be as high as 50\\% \\citep{2008MNRAS.386.1426K, 2008ApJ...672..102G, 2011MNRAS.410..860A}. These absorbers can be outflowing with speeds as high as $\\sim 0.2c$. In the X-ray band, we observe lower-velocity (100-1000 km s$^{-1}$) ionized gas (the so-called ``warm absorber'') in 50\\% of AGN \\citep{2005A&A...432...15P, 2007MNRAS.379.1359M}. X-ray high-velocity absorbers outflowing with speeds up to $\\sim 0.3c$ (called ultra-fast outflows, UFOs) are observed in about 30-40\\% of local AGN \\citep{2010A&A...521A..57T}. X-ray BALs, outflowing with speeds up to $0.7c$, have been also observed in a handful of AGN \\citep{2002ApJ...579..169C,2009ApJ...706..644C,2011xru..conf..240L}. These observational results indicate that outflowing matter is a key ingredient of the inner regions of AGN, and that understanding the launching and structure of these winds is fundamental for constructing complete models of the central engine.  Here we present the main results of a multi-epoch \\textit{XMM-Newton} observational campaign on PG 1126-041, a low-redshift ($z=0.06$) AGN that shows UV mini-BALs blueshifted by 2000-5000 km s$^{-1}$ in the C~IV, N~V, and Si~IV species \\citep{1999MNRAS.307..821W}. An early \\textit{ROSAT} observation suggested the presence of ionized absorption as well as possible short-term variability \\citep{1999MNRAS.307..821W,2000A&A...354..411K}. The new \\textit{XMM-Newton} campaign provided the largest dataset (a total of four pointings: one archival in 2004, two in 2008, one in 2009) and the deepest X-ray exposure (133 ks in the 2009 observation) ever on a mini-BAL QSO. Our analysis confirmed the presence of massive ionized absorbers along the line of sight, and revealed complex spectral variability on time scales of both months and hours. We focus here on the variability of the X-ray absorbers; details on the X-ray analysis can be found in \\citet{2011A&A...536A..49G}. ",
        "conclusions": "The \\textit{XMM-Newton} observational campaign on the mini-BAL QSO PG 1126-041 provided new results about the link of the accretion and ejection processes in AGN. Time-resolved X-ray spectroscopy performed on long and short time scales revealed the presence of massive and variable ionized absorbers along the line of sight, as well as a varying intensity of the intrinsic X-ray continuum. In particular, the long 2009 observation provided some of the best evidences of the transient character of high-velocity, highly ionized X-ray absorbing outflows over very short time scales of only a few hours.  The appearance of the X-ray spectrum of the mini-BAL QSO PG 1126-041 is dominated by variable ionized absorption. The variability of the moderately-ionized absorber is qualitatively consistent with the scenario depicted by \\citet{2010ApJ...719.1890M}, where a variable X-ray absorbing screen is invoked to explain the observed variability over (rest frame) week/month time scales of the UV mini-BALs in the QSO HS 1603+3820. Our observational findings are also consistent with UV line-driven accretion disk winds theoretical scenarios, where a thick, variable screen of X-ray absorbing gas (an overionized ``failed wind'') close to the central SMBH shields the more distant portion of the flow from the strong ionizing continuum, and allows for its effective acceleration \\citep[e.g.][]{2004ApJ...616..688P}. Indeed, the observed X-ray spectra of PG 1126-041 are quite similar to those predicted for highly-accreting AGN by the hydrodynamical simulations of \\citet{2004ApJ...616..688P}, as developed by \\citet{2010MNRAS.408.1396S}.  The strong variability of the highly-ionized outflowing absorber suggests that the assumption of constant velocity, usually used to estimate the mass outflow rate associated to the highly-ionized part of AGN winds, is likely to be incorrect: the dynamics of the inner accretion/ejection flow around SMBHs needs to be further investigated and, to this end, time-resolved X-ray spectroscopy is the most powerful tool that we have. Many exciting questions are still open: what is driving the variability of the ionized absorbers over different time scales? Is there any relation with the X-ray continuum flux level? What are the duty cycles of the ionized absorbers? Most importantly, how all this reflects on the mass ouflow rate and on the energy budget associated to the wind? These questions may be tackled by means of long and deep X-ray monitoring campaigns of bright BAL and mini-BAL QSOs."
    },
    "1112/1112.4493_arXiv.txt": {
        "abstract": "With this third paper of a series we present Johnson-Gunn $B,g,V,r,i,z$ multicolour photometry for 79 objects, including a significant fraction of the faintest galaxies around NGC~5044, assessing group membership on the basis of apparent morphology (through accurate S\\'ersic profile fitting) and low-resolution ($R = 500-1000$) optical spectroscopy to estimate redshift for 21 objects.  Early- and late-type systems are found to clearly separate in the S\\'ersic parameter space,  with the well-known luminosity vs. shape relation being mostly traced by different morphological types spanning different ranges in the shape parameter $n$. A significantly blue colour is confirmed for Magellanic irregulars (Sm/Im's), while a drift toward bluer integrated colours is also an issue for dE's. Both features point to a moderate but pervasive star-formation activity even among nominally ``quiescent'' stellar systems. Together, dE's and Im's provide the bulk of the galaxy luminosity function, around $M(g) \\simeq -18.0\\pm 1.5$, while the S0 and dSph components dominate, respectively, the bright and faint-end tails of the distribution. This special mix places the NGC~5044 group just ``midway'' between the high-density cosmic aggregation scale typical of galaxy clusters, and the low-density environment of looser galaxy clumps like our Local Group. The bright mass of the 136 member galaxies with available photometry and morphological classification, as inferred from appropriate M/L model fitting, amounts to a total of $2.3\\,10^{12}$~M$_\\odot$. This is one seventh of the total dynamical mass of the group, according to its X-ray emission. Current SFR within the group turns to be about 23~M$_\\odot$~yr$^{-1}$, a figure that may however be slightly increased facing the evident activity among dwarf ellipticals, as shown by enhanced H$\\beta$ emission in their spectra.  Lick narrow-band indices have been computed for 17 galaxies, probing all the relevant atomic and molecular features in the 4300-5800~\\AA\\ wavelength range. Dwarf ellipticals are found to share a sub-solar metallicity ($-1.0 \\lesssim [\\mathrm{Fe/H}] \\lesssim -0.5$) with a clear decoupling between Iron and $\\alpha$ elements, as already established also for high-mass systems. Both dE's and dS0's are consistent with an old age, about one Hubble time, although a possible bias, toward higher values of age,  may be induced by the gas emission affecting the H$\\beta$ strength. ",
        "introduction": "}  Low-mass and low-surface brightness (LSB) galaxies, including in this definition both dwarf ellipticals and irregulars, have since long surged to a central role in the general debate about galaxy formation and evolution. These objects, typically characterized by an absolute blue magnitude fainter than $M_B \\simeq -16$ and a mean surface brightness $\\langle \\mu_\\mathrm{e}\\rangle \\gtrsim 22$~mag~arcsec$^{-2}$ \\citep{fb94}, display in fact a quite wide range of morphological and spectrophotometric properties \\citep[e.g.][]{jerjen00} such as to rise the question whether they assume an individual role as a sort of pristine building blocks to form more massive systems \\citep[as in the CDM theoretical scheme, e.g.][]{kaufmann93}, or they rather originate as coarse ``by-products'' of the formation of ``standard'' ellipticals and spirals \\citep{knezek99,lisker07}.  The problem of the genetic signature of low-mass systems is not a negligible one as, although in a more silent way, dwarf galaxies display in many cases continuous star formation activity, which makes them possible efficient ``engines'' to enrich the interstellar medium \\citep[e.g.][]{arimoto95}. Actually, a major problem in this framework resides in the apparent dichotomy between active ongoing star formation, widely observed in LSB galaxies, and the exceedingly low metallicity of their present-day stellar populations \\citep{arimoto90}. The most striking case in this context is surely that of IZw~18, as recently discussed by \\citet{aloisi07}. Again, while for some dwarf ellipticals we might be dealing with a somewhat quiescent evolutionary scenario, reminiscent of old globular cluster history, in other situations a ``tuned'' star formation could have proceeded allover the entire galaxy life \\citep[e.g.][]{grebel05,held05} still leaving, nowadays, an important fraction of fresh (primordial?) gas \\citep[e.g.][]{carrasco95,kuzio04}.  As far as galaxy spatial distribution is concerned, the ubiquitous presence of LSB irregulars in the different cosmic environments \\citep{cowie91,fb94,saracco99} is clearly at odds with the morphology-density relation \\citep{dressler80}, which has been proven to hold for ``standard'' galaxies for a wide range of densities and scales \\citep{postman84,maia90,hp03,kelm05}. A systematic observational work to map the LSB galaxy distribution in selected zones of the sky has been carried out by different teams, leading to complete surveys and morphological catalogues of some loose groups of galaxies and nearby clusters \\citep{binggeli85,karachentseva85,ichikawa86,davies88,ferguson89,fs90,jd97,sh97}.  These morphological surveys serve as a basic reference to any further spectrophotometric analysis of individual galaxies \\citep{bm88,bh91,cellone94,held94,secker97,cellone99,smith08}. However, efficient detection and accurate observation of LSB galaxies still remain a formidable task, even for HST observations \\citep[e.g.][]{ferrarese06} especially when moving to distances beyond the Virgo and Fornax Clusters, the two best studied aggregates so far \\citep[e.g.][]{hilker99,drinkwater01,conselice01,deady02,mieske04,gavazzi05,sabatini05}.  A complementary view may be gained through the study of smaller groups. In fact, the wide range of densities which characterizes these aggregates makes them the natural link between galaxy clusters and the field. Detailed mapping of the distribution and nature of low-mass galaxies in small groups will fill the ``void'' between fully relaxed systems \\citep[like actually the case of NGC 5044 group, e.g.][]{forbes07} to very loose and possibly unbound aggregates \\citep{campos04,roberts04}. Besides probing the properties of dwarf galaxies in different environments, these have the advantage that even a moderate-sized sample would be fairly representative of the whole group population. At the same time, depth effects would be minimized allowing a better analysis of distance-dependent quantities.  In this work we want to further follow up a long-range study on the low-mass LSB galaxy population of the NGC~5044 group. In particular, after assessing group morphology and membership in previous contributions \\citep{cb01,cb05}, we would like to tackle here the problem of galaxy evolutionary properties, according to a multiwavelength observational approach that includes integrated and surface-brightness multicolour photometry, further complemented in most cases with a spectroscopic input. Our analysis will rely on original stellar population synthesis models, to which refer galaxy data in order to lead to a convenient and physically self-consistent picture of the galaxy evolutionary status for the NGC~5044 group.  In this framework, we will arrange our discussion by presenting the observed database in Sec.~2. Its morphological characterization will be carried out in Sec.~3 by fitting galaxy surface brightness with a \\citet{S68} profile. These results in most cases refine our previous discussion in \\citet{cb05} leading also to a homogeneous morphological classification of the full galaxy population in the NGC~5044 group. Together with up-to-date redshift measurements, the revised morphological parameters proved to be a useful additional piece of information to constrain group membership and eventually assess the overall physical properties of the galaxy group as a whole. As a central step in this direction, in Sec.~4 we study galaxy distribution in different magnitude and colour domains. Starting from apparent magnitudes and morphological types, a match with the \\citet{buzzoni05} template galaxy models also allows us to lead to a realistic estimate of galaxy stellar mass and a reference birthrate figure for the galaxy group as a whole.  The spectroscopic properties for our sample are finally discussed in some detail in Sec.~5, where Lick narrow-band indices are derived for most of the observed galaxies. The diagnostic diagrams compare $\\alpha$-element features, like Magnesium and Calcium, versus Iron and Balmer line stregth, also including the contribution from Carbon molecules like CH (G band) and  C$_2$. We will discuss our results and summarize our main conclusions in Sec.~6. ",
        "conclusions": "With this third paper of a series \\citep[see also][]{cb01,cb05}, we conclude our long-term project on the study of the galaxy group surrounding the standard elliptical NGC~5044. As widely recognized in the recent literature \\citep{faltenbacher05,sengupta06,brough06,mendel08,mendel09} this group is an outstanding aggregate in the local Universe for its notable population of low-mass (both dwarf and LSB) galaxies, which take part in a massive ($M_{\\rm tot} \\sim 1.6\\,10^{13}~M_\\odot$, \\citealp{betoya06}) dynamically relaxed stucture, as traced by the diffuse X-ray emission of the group \\citep{buote03,brough06}.  Along several observing runs carried out at the {\\sc Casleo} (San Juan, Argentina) and {\\sc Eso} (La Silla, Chile) telescopes we collected multicolour $B,g,V,r,i,z$ photometry for a significant fraction (79 objects) of the faintest galaxies projected on the NGC~5044 sky region assessing cluster membership on the basis of apparent morphological properties (through accurate S\\'ersic profile fitting) and low-resolution ($R = 500-1000$) spectroscopy to estimate redshift for 25 objects in the field.  Both on the basis of the morphological and spectrophotometric properties of the member galaxies, it is evident from our analysis a marked segregation effect within the global population. As far as the galaxy surface-brightness profile is concerned, a clear separation can be made in the S\\'ersic parameter space between early- and late-type systems, the first showing in nearly all cases a ``concave'' (shape parameter $n \\lesssim 1$) radial brightness profile. Magellanic systems, on the contrary, stand out in the distribution for their clean exponential profile ($n \\simeq 1$, see Fig.~\\ref{f3}), while dwarf spheroidals confirm their sort of ``acephalous'' structure, where the ``convex'' profile ($n > 1$) lacks of any nuclear design, and the galaxy body vanishes into a thiny external envelope of extremely faint mean surface brightness ($\\mu(g) \\gtrsim 26$~mag arcsec$^{-2}$, see Fig.~\\ref{f4}). Overall, compared with a de Vaucouleurs $R^{1/4}$ profile, typical for standard ellipticals, the dwarf early-type galaxies in the group (dE, dS0, dSph) characterize for their ``broader'' (i.e.\\ poorly-nucleated) surface-brightness distribution.  The intrinsically bluer colour and the larger colour spread, compared to the other morphological sub-groups (see, in particular, Fig.~\\ref{f5} and Table~\\ref{t3}), marks the population of Magellanic irregulars, pointing to a wide range of star-formation histories. The case of the Sm galaxy N134 is a notable one in this regard, appearing to be the youngest member of the N5044 group with an age well less than 1~Gyr (see Fig.~\\ref{f8}), probably as a result of its ongoing interaction with NGC~5054. On the other hand, a drift toward bluer integrated colours is also an issue for dE's, a feature that may point to some moderate star-formation activity even among these nominally ``quiescent'' stellar systems \\citep[see, e.g.\\ ][for similar results with the Fornax cluster dE's]{cellone96}.  Once rescaling for the group distance, the absolute magnitude distribution indicates that in the NGC~5044 group we are dealing with a prevailing fraction of systems fainter than $1.8\\,10^9~L_\\odot$ (see Sec.~4.2). In particular, dwarf ellipticals and irregulars, together, mark the bulk of the galaxy population around $M(g) \\simeq -18.0\\pm 1.5$, while dSph's characterize the faint-end tail of the galaxy distribution (Fig.~\\ref{f6} and \\ref{f7}). The leading role of disk galaxies may even better impose when considering that, apart from NGC~5044 that leads the group as an elliptical, the other three brightest members (heading then the underlying dS0 line in the c-m diagram) are all spirals, namely NGC~5054 (type Sb), N68 (Sab) and N18 (Sab).  To some extent, this further emphasizes the quite special location of the NGC~5044 group, within the cosmic aggregation scale, being strategically placed just ``midway'' between the high-density environment of galaxy clusters, and the low-density conditions of looser galaxy clumps like our Local Group. Recalling the recent discussion of \\citet{gavazzi10} about the ``nurture'' effects on galaxy evolution, it is interesting to remark that this ``duality'' of the NGC~5044 group reflects also in the prevailing galaxy mix within the aggregate, where the gas-depleted population of dwarf ellipticals (a typical sign of dense cosmic environments) coexists with an important population of dwarf (gas-rich) irregulars (the typical inhabitants of low-density regions in the Universe).  A tentative assessment of the mass budget sampled by the group member galaxies has been carried out in Sec.~4.2 on the basis of the observed colours and the morphology of each object. By relying on the \\citet{buzzoni05} template models of Table~\\ref{t4}, this combined piece of information allowed us to assign a representative M/L such as to convert absolute magnitudes into bright (baryonic) mass. Overall, the 63 member galaxies in our sample complemented by other 50 likely members comprised in the \\citet{fs90} original catalog (which includes NGC~5044 and 5054 themselves), and by 23 additional member galaxies with available photometry and morphological type from the \\citet{mendel08} survey (see Table~\\ref{t5} and \\ref{t6}) collect a total of $M_{\\rm tot}^{\\rm bright} = 2.3\\,10^{12}$~M$_\\odot$, with a mean representative value of $\\langle \\log M^*_{\\rm gal}\\rangle = 9.2 \\pm 1.0$ for the member galaxies (see Fig.~\\ref{f9}). Roughly one fourth of $M_{\\rm tot}^{\\rm bright}$ is stored in NGC~5044 itself, while the three brightest members, alone, represent nearly half the total bright mass of the group.  The derived total bright mass is about one seventh of the total dynamical mass of the group as inferred from the X-ray emission map. This fraction might however be getting even smaller if one accounts for a larger estimate between $M_{\\rm tot}^{\\rm dyn}\\sim 3.7\\,10^{13}~M_\\odot$ \\citep{gastaldello07b} and $9.9\\,10^{13}$~M$_\\odot$, according to \\citet{mendel08}, just on the basis of the galaxy velocity dispersion.  Like in the luminosity function, a clear morphological segregation is also in place for galaxy masses with dE and dS0 systems which surmount the Im and dSph component. The latter, in particular, constrains the low-mass tail of the mass function providing the key connection between galaxies and globular clusters, around the $10^7$-$10^8$~M$_\\odot$ mass range. A more marked presence of later-type systems among the faintest galaxy population naturally complies with the downsizing mechanism \\citep{gavazzi10} making active star formation to better confine among low-mass aggregates, as shown in Fig.~\\ref{f10}.  As discussed in Sec.~4, fresh stars are produced within the NGC~5044 group at a rate of roughly 23~M$_\\odot$~yr$^{-1}$, implying a global birthrate $b \\sim 0.15$. This figure might however be a prudent lower limit, as it does not account for the contribution of the two brightest members of the group and for the possible star-forming activity even among dE and dS0 galaxies. Some evidence in this sense stems, in fact, from the broad colour spread of these galaxies in Fig.~\\ref{f8}, but even clearer signs emerge from the spectroscopic analysis of Sec.~5. In particular, Fig.~\\ref{f17} shows that the general presence of H$\\beta$ emission in the spectra of many dwarf ellipticals (and in NGC~5044 itself) is resilient of a moderate but pervasive activity of fresh star formation throughout in the group environment.  The good spectroscopic material, collected along the same photometric observing runs added a number of useful details to characterize the evolutionary status of the NGC~5044 group. These data allowed a univocal membership classification in most ambiguous cases discriminating between genuine dwarf galaxies belonging to the group and bright background galaxies (see Table~\\ref{t7}); the case of objects B3 and N39 is illustrative in this sense, as both ellipticals eventually resulted to be part of a background structure at $z \\sim 0.095$. In this regard, at least three relevant galaxy aggregates appear beyond our group, the farthest one being a confirmed galaxy cluster at $z \\sim 0.28$ \\citep{gastaldello07}.  The resolution and wavelength coverage for {\\sc Eso} and {\\sc Casleo} spectra also allowed us to self-consistently derive a series of Lick narrow-band indices, as summarized in Table~\\ref{t8}. In particular, eight relevant features of Iron between 4300 and 5800~\\AA\\  have been sampled, while $\\alpha$ elements can be traced by means of the 4455~\\AA\\ Ca feature and the striking Mg feature at 5170~\\AA. In addition, two molecular bands of Carbon (CH at 4300~\\AA, and C$_2$ at 4668~\\AA) complete our diagnostic tools. As far as early-type galaxies are concerned, spectroscopic analysis clearly indicate for them a sub-solar metallicity ($-1.0 \\lesssim [Fe/H] \\lesssim -0.5$, see especially Fig.~\\ref{f18}), and in general an old age, consistent with the Hubble time. In the latter respect, however, one has to be aware of the possible hidden bias (toward higher values) in age diagnostic via H$\\beta$ absorption strength (Fig.~\\ref{f17}) due to the effect of gas emission.  As for the established relationship for high-mass ellipticals \\citep{worthey92}, also the dwarf population of NGC~5044 shows a decoupled trend between Iron and $\\alpha$ elements (as in the uppel panel of Fig.~\\ref{f18}), consistent with the mild trend between $[\\alpha/Fe]$ and stellar mass for $M > 10^{8.5}~M_\\odot$ galaxies shown by \\citet{mendel09}. This feature supports and further extends also to the low-mass framework the canonical picture for galaxy chemical evolution as a result of the two distinct enrichment channels dealing with Type I and Type II SNe \\citep[see, e.g.][]{matteucci86}."
    },
    "1112/1112.1425_arXiv.txt": {
        "abstract": "We present a new distance determination to the Galactic globular cluster 47 Tucanae by fitting the spectral energy distributions of its white dwarfs to pure hydrogen atmosphere white dwarf models.  Our photometric dataset is obtained from a 121 orbit Hubble Space Telescope program using the Wide Field Camera~3 UVIS/IR channels, capturing F390W, F606W, F110W, and F160W images.  These images cover more than $60$ arcmin$^2$ and extend over a radial range of $5-13.7$ arcmin ($6.5-17.9$ pc) within the globular cluster. Using a likelihood analysis, we obtain a best fitting unreddened distance modulus of (m - M)$_{\\rm{o}}=13.36\\pm0.02\\pm0.06$ corresponding to a distance of $4.70\\pm0.04\\pm0.13$ kpc, where the first error is random and the second is systematic.  We also search the white dwarf photometry for infrared excess in the F160W filter, indicative of debris disks or low mass companions, and find no convincing cases within our sample. ",
        "introduction": "\\label{sec:intro}  Many astrophysical analyses involving globular clusters rely on well determined metallicity, reddening, and distance measurements.  In particular, the estimated ages of Galactic globular clusters can depend heavily on these assumed quantities, so their careful determination is necessarily important.  Ages of Galactic globular clusters can place strong constraints on the epoch of formation of our Galaxy, identify when cluster formation occurred with respect to the era of reionization in the Universe and the cosmic star formation peak, constrain the formation history of the Galaxy by dating the halo and bulge, as well as constrain $\\Omega_M$ \\citep{gratton03} when coupled with measurements of $H_{\\rm{o}}$ by WMAP \\citep{spergel03} and the HST Key Project \\citep{freedman01}.  Most of the Galactic globular clusters have ages between 9-14 Gyr \\citep[e.g.][]{marinfranch09,dotter10}, with most absolute ages commonly determined by fitting theoretical isochrones to the main sequence turnoff.  This, however, depends strongly on the assumed distance modulus to the cluster, which is the most significant contributor to the age uncertainty \\citep{bolte95}.  Although less sensitive to distance, even the more recently employed age dating technique via the white dwarf cooling sequence \\citep{hansen04,hansen07} requires a precise distance modulus. \\begin{deluxetable*}{lll} \\tablecolumns{3} \\tabletypesize{\\small} \\tablecaption{Comparative Absolute Distance Estimates to 47 Tuc\\label{tab:distance}} \\tablewidth{0pt} \\tablehead{ \\colhead{$\\rm{(m - M)_o}$\\tablenotemark{a}} & \\colhead{Method}& \\colhead{Reference}  \\\\ \\colhead{ } &\\colhead{ } &\\colhead{ }  \\\\ }\\startdata $13.36\\pm0.02\\pm0.06$\\tablenotemark{b} & WD Spectral Energy Distributions & This Paper \\\\ $13.23\\pm0.08$ & Eclipsing Binary & \\cite{thompson10}\\\\ $13.25$ & Main Sequence Fitting & \\cite{bergbusch09} \\\\ $13.32\\pm0.09$ & Tip of the Red Giant Branch & \\cite{bono08}\\\\ $13.47\\pm0.11$ & RR Lyrae & \\cite{bono08}\\\\ $13.28\\pm0.07$  &  Eclipsing Binary & \\cite{kaluzny07} \\\\ $13.18\\pm0.03\\pm0.04$\\tablenotemark{b} & Horizontal Branch Fitting & \\cite{salaris07}\\\\ $13.02\\pm0.19$& Cluster Kinematics & \\cite{mclaughlin06}\\\\ $13.40\\pm0.08$\\tablenotemark{c,d} & Main Sequence Fitting &  \\cite{gratton03} \\\\ $13.25+0.06-0.07$  & Main Sequence Fitting & \\cite{percival02} \\\\ $13.21\\pm0.04\\pm0.07$\\tablenotemark{b} & Main Sequence Fitting & \\cite{grundahl02} \\\\ $13.15\\pm0.14$ & WD Cooling Sequence Fitting & \\cite{zoccali01} \\\\ $13.45\\pm0.09$\\tablenotemark{c} & Main Sequence Fitting & \\cite{carretta00} \\\\ $13.32\\pm0.20$ & Tip of the Red Giant Branch & \\cite{ferraro00} \\\\ $13.29\\pm0.20$\\tablenotemark{e} & Horizontal Branch Fitting & \\cite{ferraro99} \\\\ $13.57\\pm0.15$ & Main Sequence Fitting & \\cite{reid98} \\\\ $13.32\\pm0.03\\pm0.04$\\tablenotemark{b} & Horizontal Branch Fitting & \\cite{kaluzny98} \\\\ $13.23\\pm0.17$ & Baade-Wesselink of RR Lyrae & \\cite{storm94} \\\\ $13.38\\pm0.05$ & Horizontal Branch Fitting & \\cite{salaris98} \\\\ $13.52\\pm0.08$\\tablenotemark{c} & Main Sequence Fitting & \\cite{gratton97}\\\\ $13.28$ & Horizontal Branch Fitting & \\cite{hesser87} \\\\ $13.18$ & Main Sequence Fitting & \\cite{hesser87} \\\\ \\enddata \\tablenotetext{a}{Dereddened using E(B~-~V)~$=0.04$ and $A_V/$~E(B~-~V)~$=3.1$.} \\tablenotetext{b}{The first uncertainty is random and the second is systematic.} \\tablenotetext{c}{Derived from the quoted distance modulus uncorrected for binary bias.  A binary correction removes 0.02 mag from the distance modulus.} \\tablenotetext{d}{Using their mean distance derived from two colors.} \\tablenotetext{e}{From their global metallicity derived distance modulus.} \\end{deluxetable*}  The Galactic globular cluster 47 Tucanae (47 Tuc, NGC 104) is considered the metal-rich proto-type for globular cluster work \\citep[$\\rm{[Fe/H]}=-0.74\\pm0.003\\pm0.026$, ][]{carretta09}. As the closest and least reddened metal-rich globular cluster, 47 Tuc serves as an important science, as well as a calibration target, for the Hubble Space Telescope  (HST).  It is used as a metal-rich anchor when comparing observations of resolved stars in nearby galaxies to stellar evolutionary models \\citep[][]{cassisi97,zoccali99,ferraro99,dicecco10} and in examining the star formation history of galaxies in the Local Group  \\citep[e.g.][]{monelli10}. As another example, the Wide Field Camera~3 (WFC3) Galactic Bulge Treasury program (GO-11664) has targeted 47 Tuc in their resolved stellar population study  \\citep{brown09}.  With these observations, ranging from the UV to near-IR, 47 Tuc will serve as an empirical population template, helping to calibrate the reddening-free indices and correcting the transformation of theoretical isochrone libraries into the WFC3 photometric system.  With relatively well measured values of metallicity and reddening \\citep[E(B~-~V)~$=0.04\\pm0.02$, ][]{salaris07,harris10}, the distance becomes the crucial input parameter to determine the age of 47 Tuc.  A considerable effort has gone into determining the distance to this cluster.  We list recent work in Table~\\ref{tab:distance}. There has been quite a variety of methodologies used to determine the distance of this cluster, including main sequence fitting \\citep{hesser87,gratton97,reid98,carretta00,grundahl02,percival02,gratton03,bergbusch09}, the tip of the red giant branch \\citep{ferraro00,bono08}, horizontal branch fitting \\citep{hesser87,kaluzny98,salaris98,ferraro99,salaris07}, RR Lyrae \\citep{storm94,bono08}, cluster stellar kinematics \\citep{mclaughlin06}, eclipsing binaries \\citep{kaluzny07,thompson10}, and white dwarf (WD) cooling sequence fitting \\citep{zoccali01}. This set of measurements spans unreddened distance moduli between $13.02-13.57$ for 47 Tuc which has important implications for the derived age. Examining only the main sequence turnoff luminosity for a cluster at the extremes in distance moduli, this yields a difference in age of $\\sim6$ Gyr\\footnote{To first order, we find that when examining the turn off magnitude, the age uncertainty is approximatley equal to (10 Gyr/mag)*(uncertainty in the distance modulus).}, altering any inferred history of the formation of the Milky Way that one might derive from 47 Tuc.  Our goal is to determine the distance to 47 Tuc following the example of \\cite{zoccali01} by using the cluster's WD population.  \\cite{zoccali01} constructed a WD sequence using local calibrators with measured trigonometric parallaxes and magnitudes to compare with the WD cooling sequence of 47 Tuc, yielding an unreddened distance modulus of (m - M)$_{\\rm{o}}=13.15\\pm0.14$ (see Table~\\ref{tab:distance}).  Here, rather than use local calibrators, we determine the distance by comparing our measured photometry to the hydrogen atmosphere (DA) WD models of \\cite{tremblay11} and the helium atmosphere (DB) models of \\cite{bergeron11}.  Previous spectroscopic studies have found solely DA WDs in globular clusters within the temperature range of our sample, including the globular clusters M4 \\citep{davis09}, NGC 6752 and NGC 6397 \\citep{moehler04}, as well as in the open cluster NGC 2099 \\citep{kalirai05}. While this has not been tested in 47 Tuc to date, we make the assumption that our sample consists of DA WDs, but we also compare our results to that obtained considering DB models and a mixture of DA and DB WDs.  WDs within globular clusters are faint objects in crowded fields, making their identification quite challenging.  However, with carefully planned observations with the HST, combined with thorough data reduction and analysis, it becomes possible to obtain high quality multi-band photometry of WDs in large samples within one cluster system \\citep[see the recent observations of WDs in M4 and NGC 6397 from][]{richer02,hansen02,hansen07,richer08}. ",
        "conclusions": "\\label{sec:conclusions}  \\subsection{A Search for Infrared Excess} \\label{sec:irexcess}  Our dataset also provides us the opportunity to search for evidence of extrasolar planetary systems.  A previous search for planets in the core of 47 Tuc was performed by \\cite{gilliland00} using a time-series analysis over $8.3$ days searching for close in Jupiter-sized planets around $\\sim34 000$ main sequence stars.  They concluded that no light curve in their sample strongly suggested the presence of a planet, indicating that the planet frequency in 47 Tuc may be an order of magnitude less than in our solar neighborhood.  This may be expected in the dense stellar environment of the core of 47 Tuc which could lead to orbit disruption and/or hinder planet survival.  In addition, the planet-metallicity correlation indicates that the frequency of planets increases with the host star's metallicity \\citep[e.g.][]{fischer05}, making it even less likely to find a planet around a star with subsolar metallicities in 47 Tuc. \\begin{figure}[H] \\plotone{f12.eps} \\caption{The observed minus the expected model magnitudes as a function of effective temperature (T$_{\\rm{eff}}$) for our WD sample in the four filters F390W ({\\it upper left}), F606W ({\\it upper right}), F110W ({\\it lower left}), and F160W ({\\it lower right}). } \\label{fig:obs_exp} \\end{figure}  Our dataset enables us to instead search for evidence of extrasolar planetary systems with the SEDs of our selected WDs.  WDs have essentially metal-free atmospheres \\citep{paquette86,zuckerman03} and low luminosities, allowing for  the detection of cool low mass counterparts, such as debris or circumstellar disks, as well as planets or brown dwarfs \\citep{probst83,zuckerman87a,farihi05,burleigh06}. One form of evidence for the existence of these features is the presence of IR excess \\citep[][among others]{zuckerman87b,graham90,kilic06,vonhippel07,farihi08}. The likelihood of a WD having an IR excess is correlated with its T$_{\\rm{eff}}$ and/or presence of photospheric metals \\citep{farihi09}, the most well studied of which is calcium.  The pollution of hydrogen-rich photospheres of WDs with metals could be caused by circumstellar accretion.  If a planetesimal or asteroid is tidally disrupted, a circumstellar dust disk could arise \\citep{jura03,jura08} and likely pollute the WD's atmosphere with detectable metals.  With no prior information about the calcium abundance of the WDs in our study, we examine their T$_{\\rm{eff}}$.  Of our 59 WDs, all but 3 have fitted T$_{\\rm{eff}}>10 000$~K, with the remaining three having T$_{\\rm{eff}}=9684\\pm1041$~K, T$_{\\rm{eff}}=9804\\pm921$~K, and $9924\\pm841$~K.  We can estimate how many cases of IR excess we expect to see in our sample using upper and lower bounds provided by various studies in the literature. Recent work by \\cite{zuckerman10} show that $\\sim30\\%$ of the photospheres of white dwarfs with T$_{\\rm{eff}}$ between $13 500-19 500$~K show evidence for being polluted with metals. In the same temperature range, approximately $20\\%$ of DA WDs with metal bearing photospheres also show evidence for warm circumstellar material \\citep{kilic08}.  This indicates a fraction of dusty WDs to be $6\\%$, which we take as an upper limit.  A lower limit to the fraction of warm DA WDs with dusty disks is $\\sim1\\%$ \\citep{farihi09}, obtained from recent large WD surveys \\citep{liebert05,koester05,farihi05,hoard07,mccook99}.  From our 171 (59) WDs, we therefore expect between $\\sim2-10$ ($\\sim0-4$) WDs to have IR excess.  However, our reddest photometric measurement is F160W, a wide-H band filter with a central wavelength of $1536.9$~nm.  This is just where the spectral energy distribution of a typical warm WD would begin to show evidence for an IR excess given a companion or disk with T$_{\\rm{eff}}\\sim800$~K.  In addition, we must also consider our photometric uncertainties. It thus seems unlikely that we would find clear evidence for an IR excess in our best sample of WDs.  However, if we do find indication for IR-excess in the H band, it could be confirmed with additional observations in redder filters.  We have searched the SEDs in Figure~\\ref{fig:sed_prob} of all our WD candidates for evidence of an IR excess, focusing on the examination of the F160W band.  Any evidence that we found for an enhanced F160W measurement appeared to be due to contamination by neighboring stars, diffraction spikes, or cosmic rays.  In addition, for each filter, we plot the observed magnitude minus the expected model magnitude as a function of effective temperature for each of our 59 WDs, shown in Figure~\\ref{fig:obs_exp}. The few outlying points in the F160W filter at lower temperature are objects that are too faint compared to the models, driven by large photometric uncertainties.  If IR excess is present, we would expect to see points with negative values on the y-axis.  Figure~\\ref{fig:obs_exp} indicates that in the F160W band, our WDs are at most in excess of 0.2 magnitudes compared to the model.  We question whether this could indicate a cool brown dwarf companion or debris disk.  Using the evolutionary models for cool brown dwarfs and exosolar giant planets of \\cite{baraffe03}, we combine a typical WD in our sample with the contribution from a 10 Gyr dust free cool brown dwarf. Here, we convert the model IR data from the CIT system to that of 2MASS using the transformations of \\cite{carpenter01}, as 2MASS is a nearly identical photometric system to that in WFC3 and provides the closest comparison possible.   In Figure~\\ref{fig:sed_companion}, we show that for masses of a companion below $0.06$ M$_{\\odot}$ ($\\sim60$ M$_{J}$), we would be unable to detect it. However, a cool brown dwarf with $0.06$ M$_{\\odot}$, may be detectable, as this leads to a difference in magnitude in the F160W band of $>0.23$ magnitudes. \\begin{figure} \\plotone{f13.eps} \\caption{The observed F606W, F110W, and F160W magnitudes for a typical WD in our sample ({\\it solid squares}).  Also shown are the added contribution to these magnitudes from a 10 Gyr dust free cool brown dwarf companion with masses of $0.05$ M$_{\\odot}$ ($\\sim50$ M$_{J}$) as {\\it open blue squares}, $0.06$ M$_{\\odot}$ ($\\sim60$ M$_{J}$) as {\\it solid red triangles}, $0.07$ M$_{\\odot}$ ($\\sim70$ M$_{J}$) as {\\it open green triangles}, and $0.075$ M$_{\\odot}$ ($\\sim75$ M$_{J}$) as {\\it solid cyan circles} from the \\cite{baraffe03} models. Here, we have converted the model magnitudes to that of 2MASS using the transformations in \\cite{carpenter01}. In order for us to detect a cool brown dwarf companion from our sample, the companion must be at least $0.06$ M$_{\\odot}$, which is an excess of $0.23$ magnitudes in the F160W filter. {\\it A color version of this figure is available in the electronic version.}} \\label{fig:sed_companion} \\end{figure} Examining Figure~\\ref{fig:obs_exp}, it is clear that none of our 59 WDs have F160W in excess of 0.2 from the models. We conclude that we do not find any case in our sample that shows strong evidence for a genuine IR excess.   Our search for planets in 47 Tuc continues, however, with this cycle 17 dataset.  We have time series data of two fields outside the core of 47 Tuc, one is the stare field described in Section~\\ref{sec:dataset} consisting of F110W and F160W observations over $\\sim20$ days with $10000$ stars, and the other field has ACS images in F606W and F814W filters spanning $\\sim9$ months with $20000$ stars.  Our analysis, currently underway, will be able to put a constraint on the binary frequency within 47 Tuc and search for exoplanets by analysis of the stellar light curves \\citep{woodley12}. Not only will we examine the light curve of main sequence stars down to $\\sim0.2$ M$_{\\odot}$, but we will also be able to search for planets around the brighter WD population.  A small radius ratio between the WDs and giant gas planets will significantly aid in our planetary search.  \\subsection{Summary} \\label{sec:summary}  With this new HST dataset of the Galactic globular cluster 47 Tuc, we have examined the SEDs of a sample of 59 WDs that are devoid of photometric contamination by neighboring stars or image defects.   From this sample of SEDs, we have applied a novel approach to simultaneously estimate the distance of 47 Tuc from the comparison of theoretical model WD atmospheres \\citep{tremblay11}.  We obtain a true distance modulus of $13.36\\pm0.02$(random)$\\pm0.06$(systematic) corresponding to a distance of $4.70\\pm0.04$(random)$\\pm0.13$(systematic)~kpc."
    },
    "1112/1112.4508_arXiv.txt": {
        "abstract": "We show that the Big Bang singularity of the Friedmann-Lema\\^itre-Robertson-Walker model does not raise major problems to General Relativity. We prove a theorem showing that the Einstein equation can be written in a non-singular form, which allows the extension of the spacetime before the Big Bang.  These results follow from \tour research on singular semi-Riemannian geometry and singular General Relativity \\cite{Sto11a,Sto11b,Sto11d} (which we applied in previous articles to the black hole singularities \\cite{Sto11e,Sto11f,Sto11g,Sto11c}). \\bigskip \\noindent ",
        "introduction": "The solution proposed here is simple: to show that the singularities of the {\\FLRW} model don't break the evolution equation, we show that the equations can be written in an equivalent form which avoids the infinities in a natural way. We consider useful to prepare the reader with some simple mathematical observations, which will clarify our proof. These observations can be easily understood, and combined they help us understanding the Big Bang singularity in the {\\FLRW} spacetime.   \\subsection{Distance separation \\vs topological separation}  To understand the singularities in the {\\FLRW} model, it is useful to make a parallel with another type of singularities, and their standard resolution in mathematics. Let's consider a surface in $\\R^3$. In general it can be defined locally as the image of a map $f:U\\to\\R^3$, where $U\\in\\R^2$ is an open subset of the plane. If the function $f$ is not injective, the surface will have self-intersections. Another way to define the surface is implicitly, as the solution of an equation. In this case it may happen again to have self-intersections. The typical example is the cone, defined as \\begin{equation} \\label{eq_cone} x^2-y^2-z^2=0. \\end{equation} We can desingularize it by making the transformation \\begin{equation} \\label{eq_desing_cone} \\begin{array}{l} \\Bigg\\{ \\begin{array}{ll} x&=u \\\\ y&=uv \\\\ z&=uw \\\\ \\end{array} \\\\ \\end{array} \\end{equation} which maps the cylinder $v^2+w^2=1$ from the space parametrized by $(u,v,w)$, to the cone from equation \\eqref{eq_cone}, in the space parametrized by $(x,y,z)$. This procedure is very used in mathematics, especially in \\textit{algebraic geometry}, and it is was studied starting with Isaac Newton \\cite{newt60}.  \\image{cone2cylinder}{0.8}{The old method of resolution of singularities shows how we can ``untie'' the singularity of a cone and obtain a cylinder. This illustrates the idea that it is not necessary to assume that, at the Big Bang singularity, the entire space was a point, but only that the space metric was $0$.}  The natural metric on the space $(x,y,z)$ induces, by pull-back, a metric on the cylinder $v^2+w^2=1$ from the space $(u,v,w)$. The induced metric on the cylinder is singular: the distance between any pair of points of the circle determined by the equations $u=0$ and $v^2+w^2=1$ is zero. But the points of that circle are distinct.  From the viewpoint of the singularities in General Relativity, the main implication is that just because the distance between two points is $0$, it doesn't mean that the two points coincide. We can see something similar even in Special Relativity: the $4$-distance between two events separated by a lightlike interval is equal to $0$, but those events may be distinct.   \\subsection{Degenerate warped product and singularities}  The mathematics of General Relativity is a branch of differential geometry, called \\textit{{\\semiriem} (or pseudo-Riemannian) geometry} (see \\eg \\cite{ONe83}). It is a generalization of the Riemannian geometry, to the case when the metric tensor is still {\\nondeg}, but its signature is not positive. In this geometric framework are defined notions like contraction, Levi-Civita connection, covariant derivative, Riemann curvature, Ricci tensor, scalar curvature, Einstein tensor. These are the main ingredients of the theory of General Relativity \\cite{HE95,ONe83,Wal84}.  The problem with the singularities is that there, these main ingredients can't be defined, or become infinite. The perfection of {\\semiriem} geometry is broken there, and by this, it is usually concluded that the same happens with General Relativity.  In \\cite{Sto11a} we introduced a way to extend {\\semiriem} geometry to the degenerate case. There is a previous approach \\cite{Kup87b,Kup96}, which works for metric of constant signature. Our need was to have a theory valid for variable signature (because the metric changes from being {\\nondeg} to being degenerate), and which in addition allows us to define the Riemann, Ricci and scalar curvatures in an invariant way, and something like the covariant derivative for the differential forms and tensor fields which are of use in General Relativity. After developing this theory, introduced in \\cite{Sto11a,Sto11d}, we generalized the notion of warped product to the degenerate case, providing by this a way to construct useful examples of singularities of this nice behaved kind \\cite{Sto11b}.  From the mathematics of degenerate warped products it followed that a warped product like that involved in a {\\FLRW} metric (equation \\ref{eq_flrw_metric}) has only singularities which are well behaved, and which allow the extension of General Relativity to those points. At these singularities, the Riemann curvature tensor $R_{abcd}$ is not degenerate, and it is smooth if $a$ is smooth. The Einstein equation can be replaced by a densitized version, which allows the continuation to the singular points and avoids the infinities.   \\subsection{What happens if the density becomes infinite?} \\label{s_infinite_density}  In the Friedmann equations \\eqref{eq_friedmann_density}, \\eqref{eq_acceleration}, and \\eqref{eq_fluid}, the variables are $a$, the mass/energy density $\\rho$ and the pressure density $p$. When $a\\to 0$, $\\rho$ appears to tend to infinity, because a finite amount of matter occupies a volume equal to $0$. Similarly, the pressure density $p$ may become infinite. How can we rewrite the equations to avoid the infinities? As it will turn out, not only there is a solution to do this, but the quantities involved are actually the natural ones. As present in the equations, both $\\rho$ and $p$ are scalar fields. But the adequate, invariant quantities actually involve the \\textit{volume element}, or the \\textit{volume form} \\begin{equation} \\label{eq_dvol} \\vol := \\sqrt{-g}\\de t\\wedge\\de x\\wedge\\de y\\wedge\\de z, \\end{equation} where by the factor $\\sqrt{-g}$ we mean $\\sqrt{-\\det g_{ab}}$. The densities are in fact not the scalars $\\rho$ and $p$, but the quantities $\\rho\\vol$ and $p\\vol$. They are differential $4$-forms on the spacetime, and the components of these forms in a coordinate system are $\\rho\\sqrt{-g}$ and respectively $p\\sqrt{-g}$.  Another hint that the natural quantities are the densitized ones is given by the stress-energy tensor. When it is obtained from the Lagrangian, what we actually get is the tensor density \\begin{equation} \\label{eq_energy_tensor:volume_form} T^{ab}\\sqrt{-g} = -2 \\dsfrac{\\delta(\\mc L_M\\sqrt{-g})}{\\delta g_{ab}} \\end{equation}  The values $\\rho$ and $p$ which appear in the Friedmann equations coincide with the components of the corresponding densities only in an orthonormal frame, where the determinant of the metric equals $-1$, and we can omit $\\sqrt{-g}$. But when $a\\to 0$, an orthonormal frame would become singular, because $\\det g\\to 0$. A coordinate system in which the metric has the determinant $-1$ will necessarily be singular when $a(t)=0$. In a non-singular coordinate system, $\\det g$ has to be variable, as it is in the comoving coordinate system of the {\\FLRW} model. From \\eqref{eq_flrw_metric} , the determinant of the metric in the {\\FLRW} coordinates is \\begin{equation} \\label{eq_det_g_flrw} \\det g = -a^6 \\det{}_3 g_{\\Sigma}, \\end{equation} where by $\\det_3 g_{\\Sigma}$ we denoted the determinant of the metric tensor of the $3$-dimensional typical space $\\Sigma$. Since the typical space is the same for all moments of time $t$, $\\det_3 g_{\\Sigma}$ is constant.  Given that the metric's determinant in the comoving coordinates is \\begin{equation} \\label{eq_sqrt_det_g_flrw} \\sqrt{-g} = a^3 \\sqrt{g_{\\Sigma}}, \\end{equation} which tends to $0$ when $a\\to 0$, we see that it might be possible for $\\sqrt{-g}$ to cancel the singularity of $\\rho$ and $p$ in $\\rho\\vol$, respectively $p\\vol$. ",
        "conclusions": ""
    },
    "1112/1112.4487_arXiv.txt": {
        "abstract": "Cosmic bubble collisions provide an important possible observational window on the dynamics of eternal inflation. In eternal inflation, our observable universe is contained in one of many bubbles formed from an inflating metastable vacuum. The collision between bubbles can leave a detectable imprint on the cosmic microwave background radiation. Although phenomenological models of the observational signature have been proposed, to make the theory fully predictive one must determine the bubble collision spacetime, and thus the cosmological observables, from a scalar field theory giving rise to eternal inflation. Because of the intrinsically non-linear nature of the bubbles and their collision, this requires a numerical treatment incorporating General Relativity. In this paper, we present results from numerical simulations of bubble collisions in full General Relativity. These simulations allow us to accurately determine the outcome of bubble collisions, and examine their effect on the cosmology inside a bubble universe. We confirm the validity of a number of approximations used in previous analytic work, and identify qualitatively new features of bubble collision spacetimes. Both vacuum bubbles and bubbles containing a realistic inflationary cosmology are studied. We identify the constraints on the scalar field potential that must be satisfied in order to obtain collisions that are consistent with our observed cosmology, yet leave detectable signatures. ",
        "introduction": "Accelerated expansion of the universe and spontaneous symmetry breaking, two powerful ideas in modern physics, conspire in many theories to give rise to a phenomenon known as eternal inflation. This phenomenon occurs when a region of the universe is in a metastable vacuum with positive energy density. Classically, these vacua are stable, and drive inflation (exponential expansion). Quantum mechanically, they can decay via the nucleation of expanding bubbles \\cite{Coleman:1977py,Coleman:1980aw} containing a new phase. Unless the rate of bubble formation outpaces the expansion of the inflating background, the original phase is never completely consumed, and inflation becomes eternal. In this picture, many different phases can be seeded from an eternally inflating ``parent\" vacuum, leading to a patchwork-universe with diverse spacetime-dependent physical properties. For a review of eternal inflation, see e.g. Ref.~\\cite{Aguirre:2007gy}.  Surprisingly, there may be directly  observable consequences of living in an eternally-inflating multiverse. In eternal inflation, our observable universe resides inside one member of an ensemble of bubbles. These bubbles necessarily collide, perturbing the homogeneity and isotropy of our own bubble interior, and providing a direct observational test of eternal inflation~\\cite{Aguirre:2007an}. Assessing the likelihood of seeing bubble collisions, determining their effects on cosmology, and predicting their observational signatures has been the subject of a large body of work~\\cite{Gott:1984ps,Garriga:2006hw,Hawking:1982ga,Wu:1984eda,Moss:1994pi,Freivogel:2007fx,Aguirre:2007wm,Aguirre:2008wy,Chang:2007eq,Chang:2008gj,Czech:2010rg,Freivogel:2009it,Easther:2009ft,Giblin:2010bd,Kleban:2011yc,Lim:2011kd,Dahlen:2008rd,Johnson:2010bn}. This body of literature has established the plausibility of models where collisions are likely, compatible with our observed cosmology, and leave observable signatures. However, an existence proof does not yet exist. For a review of much of this work, see Ref.~\\cite{Aguirre:2009ug}.  The primordial inhomogeneities left by bubble collisions are imprinted in the cosmic microwave background (CMB) radiation. The signature of an individual collision is an azimuthally symmetric modulation of the CMB temperature in a localized region. In general, the theory predicts a set of observable collisions, with the properties of each drawn from an in-principle calculable probability distribution. A search for the signatures of bubble collisions in the Wilkinson Microwave Anisotropy Probe~\\cite{Bennett:2003ba} (WMAP) 7-year CMB data~\\cite{Komatsu:2010fb} was performed in Refs.~\\cite{Feeney:2010dd,Feeney:2010jj}, which yielded an upper bound on the average number of observable collision signatures on the CMB sky.  A key element in the analysis of Refs. \\cite{Feeney:2010dd,Feeney:2010jj} is the fact that each bubble collision is expected to leave a fairly generic set of signatures in the CMB~\\cite{Feeney:2010dd,Chang:2008gj,Czech:2010rg,Kleban:2011yc}, described by a phenomenological template with only a few free parameters. However, there is currently no way to connect the parameters in this phenomenological model to the parameters in the fundamental theory. In order to go from the Lagrangian to the signatures of bubble collisions, one must determine the properties of the colliding bubbles, the immediate outcome of the collision, and the cosmological evolution of the perturbed bubble interior. Because of the non-linear nature of the field equations, a full and unambiguous determination of the cosmological signatures of bubble collisions requires numerical relativity. See  e.g. Ref.~\\cite{2001CQGra..18R..25L} for a modern review of this field.  In this paper, we present the results of fully-relativistic simulations of the collision between Coleman-de Luccia (CDL) bubbles containing a realistic inflationary cosmology. Although any given bubble undergoes many collisions, the future domains of influence of each of the collision events typically do not overlap until relatively late times inside the bubbles. It is therefore a good approximation to treat each collision as occurring in isolation, and we restrict ourselves to the study of collisions between two bubbles.  The spacetime resulting from the collision between two CDL vacuum bubbles has SO(2,1) symmetry~\\cite{Hawking:1982ga,Wu:1984eda}, allowing us to treat the problem of bubble collisions in full generality using a 1+1D simulation. This symmetry is expected to be mildly broken by fluctuations on the walls  of the colliding bubbles~\\cite{Adams:1989su,Garriga:1991tb,Garriga:1991ts,Aguirre:2005sv}. However, expanding bubbles become more and more spherical (the perturbations remain much smaller than the overall size of the bubbles), and so our truncation to 1+1 dimensions should capture the most important dynamics. Einstein's equations in 1+1D are fully constrained by the scalar field configuration, allowing for a straightforward and efficient numerical implementation.  The goal of the present paper is to study the outcome of bubble collisions and identify the properties that models must have in order to be consistent with our observed cosmology. This has been addressed to a certain extent in previous work using the Israel Junction Condition formalism~\\cite{Hawking:1982ga,Wu:1984eda,Freivogel:2007fx,Aguirre:2007wm,Chang:2007eq,Johnson:2010bn}, as well as in numerical simulations that neglected gravitational effects~\\cite{Aguirre:2008wy,Easther:2009ft,Giblin:2010bd}. In addition, Ref.~\\cite{Blanco-Pillado:2003hq} studied collisions in full numerical relativity for the collision of vacuum bubbles (see also Ref.~\\cite{Carone:1989nj} for a relativistic treatment of single bubbles). Our fully relativistic treatment allows us to go beyond these previous studies and completely determine the phenomenology of bubble collisions given a particular model of the scalar field theory driving eternal inflation. In addition, we determine how well the approximations used in previous work hold up when compared against the fully relativistic solutions. In a forthcoming paper, we will extend this numerical framework to extract the observational signatures of bubble collisions.  The remainder of the paper is organized as follows. We begin in Sec.~\\ref{sec:summary} by reviewing conclusions obtained from the Israel Junction Condition formalism and summarizing the set of questions we attempt to address with our simulations. In Sec.~\\ref{sec:equationsofmotion}, we perform the truncation to the 1+1D system of equations to be solved numerically, and discuss the generation of initial data. We then briefly outline the numerical implementation and tests of the code in Sec.~\\ref{sec:numerical_implementation}, and present the results of our simulations in Sec.~\\ref{sec:results}. More detailed treatments of a few technical aspects of the setup are presented in a set of appendices. We work in units where $M_{\\rm Pl} = 1$ unless otherwise noted; Newton's constant is given by $G_N = M_\\mathrm{Pl}^{-2}$. ",
        "conclusions": "This work establishes the groundwork for rigorously determining the connection between an underlying potential landscape giving rise to eternal inflation and the cosmological signal of bubble collisions. Such a goal requires the ability to study the coupled Einstein and scalar field equations without approximations. To this end we have developed a robust code that implements the full system of relativistic field equations, thoroughly tested it on a number of cases and applied it to determine the outcome of collisions between bubbles.  Our results highlight the utility of two complementary analytic tools for predicting the outcome of bubble collisions. For the collision of vacuum bubbles, we have confirmed that the thin-wall solutions obtained from the Israel junction condition formalism approximate the numerical solutions to good accuracy. In particular, the junction condition solutions give a good quantitative prediction for the metric functions to the future of the collision, as well as a good qualitative prediction for the late-time behavior of the post-collision domain wall. The limitation is that the junction condition formalism cannot capture the dynamics of the collision event. This is particularly problematic when considering bubbles with an interior cosmology. The dynamics are, however, captured to a certain extent by the free passage approximation~\\cite{Giblin:2010bd} which states that, in the immediate future of the collision, the field profiles making up the two bubbles simply superpose. This illustrates that the structure of the potential can be more important for determining the outcome of a collision than the kinematics. The free passage approximation breaks down shortly after the collision, at which point the full field dynamics must be considered. A full numerical solution is necessary to bridge the gap between the realm of validity of the free passage approximation and the junction condition formalism.  Some of the main conclusions we have reached based on our numerical solutions are as follows: \\begin{itemize} \\item The energy released in the collision of identical vacuum bubbles goes mostly into the formation localized field configurations such as oscillons and pockets of the false vacuum. \\item Classical transitions occur in the presence of gravity, and lasting regions of higher energy density can be created through gravitational effects. \\item The structure of the potential is the dominant factor in determining the immediate outcome of a collision. For the collision between identical bubbles, the field is initially displaced a distance equal to the barrier width towards its minimum. For the collision of non-identical bubbles, when the potential barriers are symmetric, the collision excites a breathing mode of the post-collision domain wall. For asymmetric barriers, the post-collision domain wall forms with some initial velocity in the direction of the thicker barrier, with no breathing modes. \\item We have demonstrated conclusively that slow-roll inflation can occur to the future of a collision. Large field models of inflation are more robust to collisions than small-field models, which are completely disrupted unless the potential barriers are highly asymmetric (with the barrier $T1$ much thicker than the barrier $T2$). \\item The direction in which the post-collision wall accelerates is determined not only by the height of the potential at the exit from the barrier, but also by the slope of the potential as the field tends towards the minimum. Therefore, if there is an inflationary potential in the Observation bubble, for the Collision bubble not to intrude, the interior of the latter is required to be either a vacuum (no cosmology) or an inflationary plateau that is either higher or broader than the one in the Observation bubble. \\item The amplitude of the perturbation to the inflaton is set mostly by the barrier widths in the potential, with only a mild dependence on the kinematics. In addition, for the examples we have studied, there are always fewer $e$-folds to the future of the collision than in the undisturbed portions of the bubble. \\end{itemize}  Based on these observations, we can determine which potentials are most likely to give rise to cosmologies including bubble collisions that are compatible with our current observable universe, yet leave detectable signatures. Potentials resembling L2 in Fig.~\\ref{fig:SRpots} are good candidates. Since this potential includes a large-field model of inflation, the collision does not disrupt cosmological evolution inside the bubble. Adjusting the relative width of the two barriers changes the amplitude of the perturbation to the inflaton. The intrinsic amplitude of the perturbation along with the total duration of inflation inside the Observation bubble determines the amplitude of the signal in the CMB or other cosmological observables.  While the existing code is able to accurately obtain the dynamics of the collision, it cannot yet be used for studying the ensuing cosmology: the uniform grid employed and coordinate conditions adopted are unable to track the increasingly narrow structures that develop around walls in a computationally efficient manner. To address this issue, we will extend our code to adopt adaptive mesh refinement (see e.g. Refs.~\\cite{1989fnr..book..206C,1993PhRvL..70....9C} for a first example of its application in numerical relativity). We will present the derivation of the observational signatures from these simulations in a forthcoming publication."
    },
    "1112/1112.5162_arXiv.txt": {
        "abstract": "We use data from the \\textit{Wide-Field Infrared Survey Explorer} (\\textit{WISE}) to perform a statistical study on  the mid-infrared (IR) properties of a large number ($\\sim$10$^2$) of BL~Lac objects --- low-luminosity Active Galactic Nuclei (AGN) with a jet beamed toward the Earth.  As expected, many BL~Lac objects are so highly beamed that their jet synchrotron emission dominates their IR spectral energy distributions.  In other BL~Lac objects, however, the jet is not strong enough to completely dilute the rest of the AGN emission.  We do not see observational signatures of the dusty torus from these weakly beamed BL~Lac objects.  The lack of observable torus emission is consistent with  suggestions that BL~Lac objects are fed by radiatively inefficient accretion disks.  Implications for the ``nature vs.\\ nurture\" debate for FR~I and FR~II radio galaxies are briefly discussed.  Our study supports the notion that, beyond orientation,  accretion rate plays an important role in AGN unification. ",
        "introduction": "\\label{sec:intro}  The standard orientation-based unification paradigm posits that every Active Galactic Nucleus (AGN) is comprised of the same basic components: an accretion disk around a supermassive black hole, a broad emission line region (BELR), an obscuring torus, and a narrow line region; radio-loud AGN also launch large-scale relativistic jets \\citep[e.g.,][]{antonucci93,urry95}.   The torus is a key  component in this unified model, as it can block the BELR along certain lines of sight, and its dust  also reprocesses UV/X-ray radiation into the infrared (IR).   In simple unification, the only other way to hide the BELR is also via a geometric argument:  emission lines are outshown by a relativistic  jet beamed toward  Earth.  In addition to orientation, intrinsic differences among AGN  also play a role in AGN unification.    For example,  \\citet{urry95} conclude their seminal review  with 10 outstanding questions, including one regarding Fanaroff-Riley galaxies \\citep[FR;][]{fanaroff74}: ``do FR~Is have broad emission line regions?\"  The significance of this question is that if FR~I galaxies lack BELRs, then intrinsic properties might play a prominent role in driving the so-called FR~I/II dichotomy (in which the more powerful FR~II galaxies have edge-brightened radio lobes).  Indications so far are that emission line luminosities from FR~I galaxies are 5--30 times weaker than for FR II galaxies, and FR I galaxies  may also have weaker tori (\\citealt{zirbel95, chiaberge99, donato04}; although also see \\citealt{cao04, leipski09}).  However, the requisite observations and their interpretation are technically challenging \\citep[][]{capetti00_cena, evans06}, and there are still many ongoing ``nature vs.\\ nurture\" debates.  We can  turn to beamed versions of FR~I  and FR~II galaxies for additional guidance --- i.e., BL Lac objects and flat spectrum radio quasars (FSRQs), respectively, which are collectively  called blazars.  Broad emission lines and sometimes dusty tori are routinely detected in the spectral energy distributions (SEDs) of FSRQs, but less often from BL Lac objects \\citep[e.g.][]{abdo10, giommi11}.  However, it is  generally still not observationally clear if BL Lac BELRs and tori appear weak solely due to dilution by jet emission \\citep[e.g.,][]{chen11_spitzer, malmrose11}, or rather if they really are intrinsically weak \\citep[e.g.,][]{ghisellini11, sbarrato11_ph}.   Constraints on blazar BELRs and tori are furthermore important, as both components are potential sources of seed photons for Comptonized gamma-ray emission, especially for FSRQs and low synchrotron peaked (LSP) BL~Lac objects.  Note, gamma-ray emission from high synchrotron peaked  (HSP)  BL~Lac objects is generally consistent with synchrotron self Compton \\citep[e.g.][]{abdo10}.   The preliminary data release of the \\textit{Wide-field Infrared Survey Explorer }\\citep[\\textit{WISE};][]{wright10} recently opened a new multiwavelength window over a large area of the sky.  In this letter, we investigate the tori of  low-luminosity radio galaxies  by examining BL Lac objects detected by \\textit{WISE}.   We describe our BL Lac sample in \\S \\ref{sec:sample}.  In \\S \\ref{sec:irprop} we describe their IR properties,  and results are discussed in \\S \\ref{sec:disc}.     All  spectral indices are defined as $f_{\\nu} \\propto \\nu^{-\\alpha_{\\nu}}$, and we define the broad-band radio to IR spectral index $\\alpha_{ri} = -\\log(L_r/L_i)/4.25$, where $L_r,$ and $L_i$ are monochromatic luminosities at 5~GHz and 3.4~$\\mu$m  rest-frame, respectively.   We adopt the following cosmology:  $H_0=71$~km~s$^{-1}$~Mpc$^{-1}$; $\\Omega_M=0.27$; $\\Omega_\\Lambda=0.73$.   \\begin{figure*} \\centering \\includegraphics[scale=0.5]{f1} \\caption{Sample SDSS optical spectra for four BL Lac objects (left), and corresponding SEDs (right).  For illustrative purposes, we  assume $z=0.2$ for the top two rows.   \\textit{WISE} data points are shown as open squares, and the other SED multiwavelength data are assembled as in \\S 4.3.5 of \\citet{plotkin11_ph}.  For comparison, we include a typical quasar SED \\citep[from][R06]{richards06}, which shows  thermal torus dust emission peaking redward of 1--2$~\\mu$m (3.0--1.5$\\times10^{14}$~Hz). } \\label{fig:sampspec} \\end{figure*}  \\newpage ",
        "conclusions": "\\label{sec:disc} We  conclude from the above that our weakly beamed BL Lac objects lack observational signatures of the dusty torus in the mid-IR.  Thus, BL Lac dusty tori appear weaker than the tori of luminous quasars (i.e., one cannot simply scale a normal quasar torus down to BL Lac luminosities).  Another way to express this conclusion is that BL Lac tori have  different properties than the tori of normal quasars.    It is reasonable to extrapolate this  result to highly beamed BL Lac objects, if beaming is primarily a geometric argument.  If one interprets the torus as an extension of the BELR \\citep[e.g.,][]{elitzur06}, then our study suggests that BL Lac BELRs are also intrinsically weak.   We stress that this is a statistical conclusion.  The simulations including the torus do predict a small population of weakly beamed BL~Lac objects with \\textit{WISE} colors matching the observations, but that population is not large enough to explain the IR colors for the  BL~Lac population as a whole.  However, we cannot exclude the presence of the torus/BELR from every single BL~Lac object.   Indeed,  some BL~Lac objects can show very weak broad emission lines in their optical spectra (especially some LSPs).   One explanation for weaker BL Lac tori is that their physical properties (i.e., torus geometries, dust properties, ionizing SEDs, etc.) are different than for luminous quasars.   Another  interpretation is  the same one invoked for optically dull AGN  and ``naked'' Seyfert galaxies \\citep[e.g.,][]{hawkins04, trump11}, as follows.  Consider a scenario where the BELR and torus are fed by a radiatively driven disk wind \\cite[e.g.,][]{murray98}.  At sufficiently low accretion rate (below a few percent $L/L_{Edd}$),  the inner region of the accretion disk is replaced by a radiatively inefficient accretion flow (RIAF).  The RIAF is then no longer able to support a sufficient  wind to populate the BELR and torus \\citep[e.g.,][]{nicastro00}.  In this scenario, the strength of the BELR and torus is intimately connected to the efficiency of the accretion flow and therefore accretion rate.\\footnote{Note, WLQs and their weak BELRs require a different explanation because WLQs show normal dusty tori and are radiatively efficient \\citep{lane11}.}  Indeed, such an ``accretion rate dichotomy'' has already been suggested to explain the BL Lac/FSRQ divide \\citep[with BL Lac objects fed by RIAFs, e.g.,][]{boettcher02, ghisellini09}.  Such a scheme may also be responsible for the FR~I/II divide \\citep[e.g.,][]{ghisellini01, wold07}.  In fact, an important (supporting) point is that the classification of unbeamed radio galaxies based on high-excitation vs.\\ low-excitation emission lines (which are likely connected to the accretion flow) is more physically meaningful than the older FR~I/II scheme based on radio morphology/jet power \\cite[see, e.g.,][]{jackson99, hardcastle09}.  Furthermore, such a division is naturally expected from  the hysteresis displayed by accreting stellar mass black holes, where changes in X-ray spectral states for individual objects (i.e., at fixed orientations) may  similarly be related to the efficiency of the inner accretion flow.  \\textit{WISE}  offers a new perspective on the strength of the obscuring torus in low-luminosity radio galaxies.  The \\textit{WISE} view of BL Lac objects is  consistent with other (independent) indications of an accretion rate dichotomy among AGN, adding a complementary piece to the puzzle.   Further constraints on such a division in accretion mode   is an important step for better understanding AGN unification."
    },
    "1112/1112.2574_arXiv.txt": {
        "abstract": "{A fully autonomous data reduction pipeline has been developed for FRODOSpec, an optical fibre-fed integral field spectrograph currently in use at the Liverpool Telescope. This paper details the process required for the reduction of data taken using an integral field spectrograph and presents an overview of the computational methods implemented to create the pipeline. Analysis of errors and possible future enhancements are also discussed. } ",
        "introduction": "The Liverpool Telescope \\citep[LT,][]{2004SPIE.5489..679S} is a 2.0 metre robotic telescope that is operating unattended at the Observatorio del Roque de Los Muchachos Observatory on La Palma, Spain. Since robotic operations started in April 2004, the LT has produced data for a variety of science programmes using software and instruments that were designed and developed in-house at the Astrophysics Research Institute of LJMU. Achieving first light in April 2009, the Fibre-fed RObotic Dual-beam Optical Spectrograph \\citep[FRODOSpec,][]{2004AN....325..215M} is the successor to the now decommissioned Meaburn Spectrograph and has been a common user instrument on the telescope since February 2010.  FRODOSpec is a bench mounted spectrograph with two optical paths, known as arms, that are utilised by separating the incident light around 5750\\AA \\ into two bandwidths using a dichroic beam-splitter. The light down each arm is collimated, dispersed and focused onto CCDs, with the elements of each optical chain separately optimised for blue and red light. An optical schematic is shown in Figure \\textbf{\\ref{fig:frodooptics}}.  Two dispersive elements are available for each arm: a conventional diffraction grating and a higher resolution Volume Phase Holographic (VPH) grating. The VPH is bonded to a prism so that the light is dispersed at the same angle as the grating, requiring no parts other than the pneumatic stage they are mounted upon to be moved when selecting between them. Resolving power and wavelength ranges for each arm and dispersive element are shown in Table \\textbf{\\ref{tab:capabilities}}.  Light is transmitted from the focal plane of the telescope to the spectrograph by a bundle of one hundred and forty-four optical fibres. At the telescope focal plane, the fibres are arranged in a regular pattern (see Figure \\textbf{\\ref{fig:frodooptics}}) to form a $12\\times12$ integral field unit (IFU), with each fibre coupled to a microlens to minimise light losses. Each fibre/microlens covers a field of view on sky of $\\sim$ $0.83''\\times0.83''$, corresponding to a total field of view of approximately $10''\\times10''$. At the input of the spectrograph, the fibres are rearranged to form a pseudo-slit which acts as the entrance aperture.  \\begin{table*}[ht] \\begin{center} \\begin{tabular}{l l l l l} \\hline & & & & \\\\ [-1ex] \\textbf{Arm / Dispersive Element} & \\textbf{Wavelength Start} & \\textbf{Wavelength End} & \\textbf{Resolution} & \\textbf{Dispersion} \\\\ [1ex] & (\\AA) & (\\AA) & & (\\AA/px) \\\\ [1ex] \\hline & & & & \\\\ Red Grating & 5800 & 9400 & 2200 & 1.6\\\\ Red VPH & 5900 & 8000 & 5300 & 0.8 \\\\ Blue Grating & 3900 & 5700 & 2600 & 0.8\\\\ Blue VPH & 3900 & 5100 & 5500 & 0.35 \\\\ & & & & \\\\ \\hline \\\\ \\end{tabular} \\caption{Wavelength ranges, resolving powers and dispersions for the different FRODOSpec arms/dispersive elements.} \\label{tab:capabilities} \\end{center} \\end{table*}  Despite the availability of instrument unspecific software packages to reduce data taken using an integral field spectrograph, the constraints of their generic design typically limit the highest achievable degree of automation to either manual like IRAF \\citep{1992ASPC...25..417V} or semi-automatic like R3D \\citep{2006AN....327..850S}, kungifu \\citep{2007NJPh....9..443B} and P3D \\citep{2010A&A...515A..35S}. As these packages do not constitute an end-to-end system whereby data products can be produced without the need for human interaction, development of a bespoke pipeline to reduce FRODOSpec data was necessary in order that the following objectives could be fulfilled:  \\begin{itemize} \\item To autonomously produce a science-ready data product. \\item To autonomously produce a ``quicklook'' data product, allowing quick data quality assessment. \\item To have full quality control over the data products produced. \\item To provide feedback for the automated LT scheduler. \\end{itemize}  This paper details the second version of the pipeline, deployed in May 2011, and is structured as follows. An explanation of the input data and the output data products is given in \\textbf{\\S \\ref{sec:overview}}. \\textbf{\\S \\ref{sec:methods}} outlines the computational methods used to process the data with error analysis. Key pipeline performance indicators are presented in \\textbf{\\S \\ref{sec:performance}}. Concluding remarks and possible future enhancements to the pipeline are discussed in \\textbf{\\S \\ref{sec:conclusions}}. ",
        "conclusions": "\\label{sec:conclusions}  This paper has explained the characteristics of data taken with FRODOSpec, and has presented a brief overview of the computational methodology used to develop a fully autonomous pipeline to reduce it. Throughout the paper, representative data taken using each dispersive element and arm was used to assess the magnitude of errors and how they propagate through the pipeline. The effect of cross-talk has been discussed, and its contribution found to be negligible. The L2 output data products have also been described, along with some key performance indicators.  The L2 is in a state of continued development. Possible future enhancements include the optimal extraction of flux, which has been shown to increase S/N with a corresponding maximum increase in effective exposure time of up to 70\\% \\citep{1986PASP...98..609H}. All enhancements and addenda will be published on the instrument website at \\footnotesize\\burl{http://telescope.livjm.ac.uk/Info/TelInst/Inst/FRODOspec/}\\normalsize."
    },
    "1112/1112.2929_arXiv.txt": {
        "abstract": "It has been suggested that besides stellar companions, substellar objects in close orbits may be able to trigger mass loss in a common envelope phase and form hot subdwarfs. In an ongoing project we search for close substellar companions combining time resolved high resolution spectroscopy with photometry. We determine the fraction of as yet undetected radial velocity variable systems from a sample of 27 apparently single sdB stars to be $\\simeq16\\%$. We discovered low-mass stellar companions to the He-sdB CPD$-$20$^{\\rm \\circ}$\\,1123 and the pulsator KPD\\,0629$-$0016. The brown dwarf reported to orbit the eclipsing binary SDSS\\,J0820+0008 could be confirmed by an analysis of high resolution spectra taken with UVES. Reflection effects have been detected in the light curves of the known sdB binaries CPD\\,$-$64$^{\\rm \\circ}$\\,481 and BPS\\,CS\\,22169$-$0001. The inclinations of these systems must be much higher than expected and the most likely companion masses are in the substellar regime. Finally, we determined the orbit of the sdB binary PHL\\,457, which has a very small radial velocity amplitude and may host the lowest mass substellar companion known. The implications of these new results for the open question of sdB formation are discussed. ",
        "introduction": "About half of the sdB stars reside in close binaries with periods ranging from a few hours to a few days \\citep{maxted01,napiwotzki04}. Because the components' separation in these systems is much less than the size of the subdwarf progenitor during its red-giant phase, these systems must have experienced a common-envelope and spiral-in phase \\citep{han02,han03}.  Although the common-envelope ejection channel is not yet properly understood in detail, it provides a reasonable explanation for the strong mass loss required to form sdB stars. However, for about half of all known subdwarfs there is no evidence for close stellar companions as no radial velocity variations are found. Among other formation scenarios, the merger of two helium white dwarfs has often been suggested to explain the origin of single sdB stars \\citep{han02,han03}.  \\citet{soker98} suggested that substellar objects like brown dwarfs and planets may also be swallowed by their host star and that common envelope ejection could form hot subdwarfs. Substellar objects with masses higher than $\\simeq10\\,M_{\\rm J}$ were predicted to survive the common envelope phase and end up in a close orbit around the stellar remnant, while planets with lower masses would entirely evaporate or merge with the stellar core. The stellar remnant is predicted to lose most of its envelope and evolve towards the extreme horizontal branch (EHB). A similar scenario has been proposed to explain the formation of apparently single low mass white dwarfs \\citep{nelemans98}. The discovery of a brown dwarf with a mass of $0.053\\pm0.006\\,M_{\\rm \\odot}$ in an $0.08\\,{\\rm d}$ orbit around such a white dwarf supports this scenario and shows that substellar companions can influence the outcome of stellar evolution \\citep{maxted06}.  The planet discovered to orbit the sdB pulsator V\\,391\\,Peg with a period of $1\\,170\\,{\\rm d}$ and a separation of $1.7\\,{\\rm AU}$ was the first planet found to have survived the red-giant phase of its host star \\citep{silvotti07}. Serendipitous discoveries of two substellar companions around the eclipsing sdB binary HW\\,Vir \\citep{lee09} and one brown dwarf around the similar system HS\\,0705$+$6700 \\citep{qian09} followed. These substellar companions to hot subdwarfs have rather wide orbits, were not engulfed by the red giant progenitors and therefore could not have influenced the evolution of their host stars. But the fact that substellar companions in wide orbits around sdBs seem to be common suggests that similar objects closer to their host stars might exist as well \\citep[for a review see][]{schuh10}.  Here we present new results from our ongoing search for sdBs with close substellar companions using high-resolution, time-resolved spectroscopy as well as time-resolved photometry. First we determine the fraction of as yet undetected RV variable systems from a sample of apparently single sdB stars and review the known candidate systems. Then we present newly discovered systems with low-mass stellar as well as substellar companions and discuss the implications of these systems for sdB formation. ",
        "conclusions": "In Fig.~\\ref{fig:substellar} the orbital parameters of all known reflection effect sdB binaries are plotted. From the 21 systems known, 11 may have substellar companions. Since only about $100$ close sdB binaries with orbital solutions are known, the fraction of systems with low-mass companions exceeds $20\\%$. The fraction of possible substellar companions is higher than $10\\%$. Due to selection effects, the true fraction must be higher. Reflection effects or eclipses necessary to clarify the nature of the companions are more likely to be found in close and highly inclined binaries. Recent discoveries of low-mass objects around sdBs with periods of more than ten days \\citep{geier11c,barlow11} indicate that such systems do exist as well and that they might be quite numerous.  It can be seen in Fig.~\\ref{fig:substellar} that the minimum companion mass seems to decrease with longer orbital periods. This is contrary to what is expected from common envelope evolution theory: The smaller the mass of the companion, the more the system has to shrink before enough energy and angular momentum is transferred to the common envelope to be ejected. However, this scenario assumes that the companion survives the CE phase. But the fate of substellar objects engulfed by close red giants is unclear. \\citet{soker98} argues that substellar objects might also evaporate during the CE phase or merge with the core of the red giant \\citep[see also][]{politano08}. Most recently, we discovered the rapidly rotating single sdB EC\\,22081$-$1916, the possible outcome of such a common envelope merger event \\citep[][Geier et al. these proceedings]{geier11a}. The destruction of the companions during or shortly after the CE phase might explain the lack of lowest-mass objects with very short periods seen in Fig.~\\ref{fig:substellar}. Those objects might have come too close to the red-giant core to survive.  We conclude that close substellar companions to sdB stars do exist and that they are by no means rare. They play a significant role in the formation of sdB stars. They may either form close binaries and survive, merge with the red-giant core creating rapidly rotating single sdBs or evaporate in the red-giant envelope. In the latter case they create ordinary single sdBs that rotate slowly (see also Geier et al. these proceedings).  \\begin{figure} \\begin{center} \\includegraphics[width=7cm]{periodK_substellar.eps} \\end{center} \\caption{The RV semiamplitudes of all known reflection effect sdB binaries plotted against their orbital periods. Eclipsing systems are plotted as diamonds, non-eclipsing systems as circles. Open symbols mark all known systems from the literature \\citep[see][and references therein]{geier11b}, while the filled symbols mark the newly discovered systems presented here and in Geier et al. (these proceedings). The solid line marks the border between the stellar and the substellar mass regime, the dashed line the border between brown dwarfs and planets.} \\label{fig:substellar} \\end{figure}"
    },
    "1112/1112.5815_arXiv.txt": {
        "abstract": "In terms of energy resolution, temporal response to burst events, and thermal stability, lanthanum bromide doped with Ce is a much better choice than the traditional NaI(Tl) scintillator for hard X-ray astronomy. We present the test results of a phoswich detector with a diameter of 101.6~mm consisting of 6~mm thick LaBr$_3$:Ce and 40~mm thick NaI(Tl), which is the largest one of this type reported so far. The measured energy resolution is 10.6\\% at 60 keV, varying inversely proportional to the square root of the energy, and the energy nonlinearity is found to be less than 1\\%, as good as those of smaller phoswiches. The coupled scintillators and phototube also show excellent uniformity across the detecting surface, with a deviation of 0.7\\% on the pulse amplitude produced by 60~keV gamma-rays. Thanks to the large ratio of light decay times of NaI(Tl) and LaBr$_3$:Ce, 250~ns vs.\\ 16~ns, pulse shape discrimination is much easier for this combination than for NaI(Tl)/CsI(Na). As the light decay time of LaBr$_3$:Ce is about 15 times faster than that of NaI(Tl), this phoswich is more suitable for detection of bright, transient sources such as gamma-ray bursts and soft gamma-ray repeaters. The internal activity of lanthanum produces a count rate of about 6~counts~s$^{-1}$ at 37.5~keV in the detector. This peak could be used for in-flight spectral calibration and gain correction. ",
        "introduction": "The phoswich detector has the advantages of high photopeak efficiency and outstanding background rejection capability over a single scintillator. The most traditional configuration of a phoswich is a combination of NaI(Tl) and CsI(Na) crystals, with NaI being the sensitive volume and CsI for anti-coincidence.  This has been widely used in hard X-ray and soft gamma-ray astronomy over the past decades, especially where large detecting area, fast timing, and/or low background are required, e.g., the Phoswich Detection System (PDS) \\citep{fro97} onboard BeppoSAX and the High Energy X-ray Timing Experiment \\citep{rot98} onboard the Rossi X-ray Timing Explorer. Phoswiches like NaI/CsI are also selected for use in future space missions such as the Hard X-ray Modulation Telescope \\citep{lit07} and the Space-based multi-band astronomical Variable Object Monitor \\citep{don10}. Other types of phoswich are also in use, such as the GSO(Ce)/BGO hard X-ray detector on Suzaku \\citep{kam96,tak07}. Therefore, investigation of new, advanced phoswiches would be interesting in X-ray astronomy.  Multiple crystals with distinct light decay times coupling to a single photomultiplier tube (PMT) provides a simple and reliable readout configuration, in contrast to high-Z semiconductor arrays of the same effective area, and is thus suitable for observations where the need for large area is prior to that for position sensitivity. However, such a large format of detector in front of a single readout has a potential disadvantage that pulses pileup is more likely to occur given the same input flux density, which will limit the maximum count rate that one can measure with the detector \\citep{mee09}. The dead time of scintillators is limited by its light decay time. The best way to overcome this issue without complicating the system is to find new crystals with shorter decay times.  Recently, lanthanum bromide crystals doped with cerium (LaBr$_3$:Ce) has been extensively studied and used due to its excellent physical properties. The energy resolution of LaBr$_3$:Ce is found to be around 3\\% at 662~keV \\citep{van01}, which, versus 7\\% for NaI(Tl), is a great improvement and is even comparable to the resolution of semiconductor (e.g.\\ CdZnTe) detectors typical of 1--2\\%. LaBr$_3$:Ce has a decay time as fast as 16~ns, more than 10 times shorter than that of NaI(Tl). Therefore, LaBr$_3$:Ce should be a better choice than NaI(Tl) for applications in hard X-ray astronomy as mentioned above. Also, its excellent thermal stability \\citep{mos06} and high tolerance to radiation damage \\citep{nor07,owe07} make it suitable and robust for space applications. Balloon flights for X-ray astronomy with a single LaBr$_3$:Ce crystal have been carried out \\citep{man11}.  \\citeauthor{maz10} \\citep{maz10} reported characteristics of a phoswich assembly with LaBr$_3$:Ce and NaI(Tl), with a diameter of 3~inch (= 76.2~mm). \\citeauthor{man11} \\citep{man11} mentioned that a similar phoswich had been designed and under study. In this paper, we report on the test of a 4~inch (= 101.6~mm) diameter LaBr$_3$:Ce/NaI(Tl) phoswich (\\S~\\ref{sec:det}), about its temporal and spectral characteristics, uniformity (\\S~\\ref{sec:char}), and internal radioactivity (\\S~\\ref{sec:int}). A brief discussion and summary of the results is presented in \\S~\\ref{sec:diss}. ",
        "conclusions": "\\label{sec:diss}  This is so far the largest LaBr$_3$:Ce/NaI(Tl) phoswich reported in the literature. The detector presents an energy resolution of 10.6\\% at 60 keV and varies following the $1/\\sqrt{E}$ relation, indicating that the energy uncertainty is dominated by Fano fluctuations. The pulse width spectrum shows two distinct peaks respectively from energy deposits in LaBr$_3$:Ce and NaI(Tl). Compared with the NaI(Tl)/CsI(Na) phoswich, whose decay time ratio is $630/250 \\approx 2.5$, the two peaks are more widely separated due to a larger time ratio of $250/16 \\approx 17$. This makes the pulse shape discrimination much easier, and will lead to a lower probability of misidentification. The lowest detectable energy of the detector is estimated to be around 2--3~keV, limited by the electronics noise and also the Be window; a Ti K$\\alpha$ line at 4.5~keV can be detected distinctly. The variation of the pulse amplitude is only about 0.7\\% across the sensitive surface of the detector; this is much smaller than the energy fluctuation and its contribution to the spectral resolution is negligible in this energy band. The detector performance, based on our tests, is found to be as good as previously measured for smaller, single LaBr$_3$:Ce crystals, indicating that the growth and sealing techniques required for a large phoswich of this type is ready.  In principle, this phoswich can achieve a dead time $\\sim$10 times smaller than the NaI(Tl)/CsI(Na) phoswich can do. The maximal allowed count rate is inversely proportional to the dead time. Thus, the LaBr3:Ce crystal can measure a count rate 10 times higher than NaI(Tl) can without pulses pileup. This would bring interesting applications in detecting burst events such as the gamma-ray bursts and soft gamma-ray repeaters. Despite the excellent performance of LaBr$_3$:Ce/NaI(Tl), its high internal activity seems to be an issue for applications where low background is needed. This does not affect the detection of burst events, which occur at a short duration with high flux and thus the instrumental background can be ignored. For long gamma-ray bursts that last up to $\\sim$100 seconds, the internal background may become significant and should be taken into account. In the 200--300~keV band, the internal background level of our detector estimated from simulations is about 0.3~counts~s$^{-1}$, or 0.0037~counts~s$^{-1}$~cm$^{-2}$. Above 300~keV, the internal background spectrum is almost flat in units of counts per keV. For comparison, the orbital background of BeppoSAX/PDS\\footnote{See Table 6 in section 4.6.3.3 of the BeppoSAX handbook available at ftp://ftp.asdc.asi.it/sax/doc/handbook/} has a count rate of 0.014~counts~s$^{-1}$~cm$^{-2}$ in 200--300~keV. We note that the background of BeppoSAX/PDS is lower than many other space instruments due to its low inclination orbit. Thus, the internal background by $^{138}$La is a minor or negligible component compared with the total orbital background in the hard band above 200~keV, and will not affect the detection of the high energy tail in relatively long astronomical phenomena. Interesting applications may include measuring the peak energy and its temporal evolution of long gamma-ray bursts \\citep{fro00}, or the search of hard X-ray emission in the afterglow \\citep{tka00,mai05}. While in the soft band below 200 keV, the estimated internal background is about 1.6--2 times of the orbital background of PDS, and will reduce the sensitivity by a factor of $\\sim$1.7. On the other hand, the electron capture of $^{138}$La with a pair of 37~keV and 1436~keV events allows coincident measurement of the 37~keV peak, which can be used for in-orbit calibration and gain correction. For that purpose, a proper design of the readout is needed, to precisely measure both the fast, low energy event in LaBr$_3$:Ce, and a slow, high energy event in NaI(Tl). This is under investigation in our lab and will be reported in the future."
    },
    "1112/1112.4602_arXiv.txt": {
        "abstract": "\\src\\ is the Central Compact Object associated with the Puppis A supernova remnant. Previous X-ray observations suggested \\src\\ to be a young neutron star with a weak dipole field and a peculiar surface temperature distribution dominated by two antipodal spots with different temperatures and sizes. An emission line at 0.8 keV was also detected. We performed a very deep (130 ks) observation  with \\xmm, which allowed us to study in detail the phase-resolved properties of \\src. Our new data confirm the existence of a narrow spectral feature, best modelled as an emission line, only seen in the `Soft' phase interval -- when the cooler region is best aligned to the line of sight. Surprisingly, comparison of our recent observations to the older ones yields evidence for a variation in the emission line component, which can be modelled as a decrease in the central energy from $\\sim$$0.80$ keV in 2001 to $\\sim$$0.73$ keV in 2009--2010. The line could be generated via cyclotron scattering of thermal photons in an optically thin layer of gas, or -- alternatively -- it could originate in low-rate accretion by a debris disk. In any case, a variation in energy, pointing to a variation of the magnetic field in the line emitting region, cannot be easily accounted for. ",
        "introduction": "Central Compact Objects (CCOs) in Supernova Remnants (SNRs) are a handful of about 10 point-like, thermally-emitting X-ray sources located close to the geometrical centres of non-plerionic supernova remnants, with no counterparts at any other wavelength. CCOs are supposed to be young, isolated, radio-quiet neutron stars \\citep[see][for a review]{deluca08}.  While the first discovered CCO (the one in the RCW103 SNR) turned out to be a unique object \\citep{deluca06}, results of X-ray timing on a sub-sample of sources with fast periodicities have recently set very useful constraints for a general picture of CCOs as a class. Analysis of multi-epoch \\xmm\\ and \\cxo\\ observations of 1E 1207.4--5209 inside G296.5+10.0 and CXOU J185238.6+004020 at the centre of Kes 79 ($P=424$ ms and $P=105$ ms, respectively) yielded unambiguous evidence for very small period derivatives \\citep{halpern10,halpern11}. This implies, under standard rotating dipole assumptions, characteristic ages exceeding the age of the host supernova remnants by more than 3 orders of magnitude, as well as very small ($\\la$$10^{11}$ G) dipole magnetic fields. This points to an interpretation of such sources as `anti-magnetars' -- weakly magnetised isolated neutron stars, born with a spin period almost identical to the currently observed one.  Such a picture has been recently strengthened by the discovery of 112 ms pulsations from \\src, the CCO in the Puppis A SNR \\citep{gotthelf09}, in two archival \\xmm\\ datasets collected in 2001, previously used by \\citet{hui06} for a comprehensive spectral analysis of the neutron star and of the surrounding SNR. \\citet{gotthelf10}, using a 2010 \\cxo\\ observation, set a $2\\sigma$ upper limit of $3.5\\times10^{-16}$ s s$^{-1}$ to the period derivative, corresponding to a dipole magnetic field $B<2\\times10^{11}$ G, and to a characteristic age $\\tau_\\mathrm{c}>5$ Myr, much larger than the age of the host SNR \\citep[3.7 kyr][]{winkler88}. Thus, \\src\\ can be included in the anti-magnetar family.  The X-ray emission properties of \\src\\ are very peculiar. \\citet{gotthelf09} report on a $180^{\\circ}$ shift in the phase of the pulse peaks occurring abruptly at $\\sim$$1.2$ keV. This has been interpreted as due to the existence of two antipodal hot spots on the star surface, with two different temperatures (`warm' and `hot') -- indeed seen in the emission spectrum as two different blackbody components. The star rotation bringing into view or hiding such regions makes the observed spectrum to change from a soft phase (maximum alignment of the warm spot with the line of sight) to a hard phase (maximum alignment of the hot spot). Even more intriguing, \\cite{gotthelf09} report on the evidence for an emission line at $\\sim$$0.8$ keV, possibly associated with the warm spot.  Here we report on a very deep \\xmm\\ observation of \\src, performed in 2009 and 2010, which  allows us to study in detail the phase-resolved emission properties of the neutron star and to check for any time variability. ",
        "conclusions": "Our multi-epoch \\xmm\\   analysis shows a phase-dependent, {\\em time variable} spectral feature, best modelled as an emission line with a variable central energy, in the X-ray spectrum of the `anti-magnetar' candidate \\src.  To put such a peculiar result in context, we first note that our observations confirm the picture of \\src\\ as a weakly-magnetised neutron star. Indeed, we  improved the upper limit on $\\dot{P}$, bringing the   dipole component of the magnetic field down to  $<2.3\\times10^{11}$ G at $3\\sigma$ level. Based on a larger statistics, we also confirm the geometric model by \\citet{gotthelf10}, explaining the phase-resolved thermal emission with two antipodal spots of different temperature (compare e.g. our Fig.~\\ref{pf} to their Fig.~6). Lack of any measurable time variation in the continuum properties suggests that the warm and hot regions are intrinsic, persistent features in the thermal map of the star. To explain such a large surface anisotropy for \\src\\ (and for CCOs in general) we may consider that a large difference in intensity could exist between the internal and the external magnetic fields of the neutron star, as proposed by \\citet{turolla11} to explain the properties of the low-magnetic field Soft Gamma Repeater SGR 0418+5729 \\citep{rea10}. In our case, an internal (toroidal + poloidal) field of a few $10^{13}$ G would be large enough to effectively channel the heat flux from the core \\citep{geppert04,geppert06}, but would not induce crustal fractures with consequent magnetar-like bursting activity.  The most natural interpretation of the variable emission line is that of a cyclotron feature produced by electrons. If its central energy is associated to the fundamental $e^-$ cyclotron frequency, the magnetic field in the line emitting region would be 6--$7\\times10^{10}(1+z)$ G (where $z\\sim0.25$ is the gravitational redshift). This is quite compatible with the upper limit  from the spin-down rate. The line is very narrow ($\\sigma\\leq40$ eV). If it is a cyclotron line then $\\Delta E / E = \\Delta B / B$ and the relative variation of $B$ over the emitting region (conservatively) needs to be $\\leq 10\\%$. Thus, the line  should be produced in a very compact region. A variation of the central energy of the feature would require  a change either in position of the emitting plasma within a non-variable magnetic field, or in the intensity of the magnetic field itself. The lack of changes in the phase-resolved continuum emission rules out simpler, purely geometric explanations such as precession of the neutron star.  To explain the generation of an emission line in the spectrum of an isolated neutron star, the possibility of cyclotron scattering of surface thermal photons by a geometrically thin, optically thin layer of plasma could be considered. However, under simple assumptions (emission from the entire star surface; plane-parallel geometry; pure, conservative scattering), a scattering layer would produce an absorption line. One might invoke  a spatially-limited scattering medium, possibly some distance away from the star surface. Photons coming from the part of the surface not covered by the layer could be scattered along the line-of-sight giving rise to an emission feature. The value of $B$ derived from the line energy is somehow smaller than the upper limit on the surface field, so the line could indeed form at some height in the magnetosphere. A confined medium seems also to be required by the results of phase-resolved spectroscopy which shows that the emission line is seen mostly when the cooler spot is into view. Still, such a picture seems rather contrived, since the nature of the layer and the mechanism keeping the plasma suspended and confined in a compact blob remains to be understood.   To ease the problem, an energy source unrelated to the surface thermal emission should be invoked to excite $e^-$ to higher Landau levels in the line emitting region. Indeed, \\citet{nelson93} predicted that for neutron stars accreting at a low rate ($L_{\\mathrm{accr}}<10^{34}$ erg s$^{-1}$), and endowed with magnetic fields of $10^{11}$--$10^{13}$ G,   accreting ions may lose energy to atmospheric electrons via magnetic Coulomb collisions. Electrons, excited to high Landau levels, radiatively decay and part of the cyclotron photons are expected to escape producing an emission line. According to \\citet{nelson93}, at $B<10^{12}$ G, the fraction of the accretion-powered flux escaping in the line is expected to be very small ($\\ll$$5\\%$), the largest part being reprocessed and emitted in a thermal continuum. Thus, one should postulate that the bulk of the X-ray luminosity of \\src\\ is accretion-powered, at variance with observations. It would require an accretion rate of   $\\sim2\\times10^{13}$ g s$^{-1}$ implying -- under standard relations for propeller spin-down \\citep{menou99} -- a $\\dot{P}$ value more than 10 times larger than our upper limit. Although the model by \\citet{nelson93} does not fit to our case, low-level accretion of supernova fallback material (which cannot be ruled out, based on X-ray timing, as well as on the optical upper limits set by Mignani et al.2009) could play some role in generating an emission line in a low-$B$ field atmosphere. A detailed investigation of such possibility is beyond the scope of this paper.  We will not go into further speculations about the line emitting mechanism. We stress that the evidence for time evolution of the spectral feature is model-independent and represents to date the first evidence for variability in an ``anti-magnetar'' candidate. Likely, such `activity' is related to a variation of the magnetic field of the  star. A $\\sim$$10\\%$ decrease in $\\sim$$8$ yr seems vastly too steep to be attributed to the large-scale dipole field. Possibly, we are witnessing evolution of a localised multipole component, dominating close to the star surface. This would hint at the presence of a large internal field, as proposed to explain the anisotropic thermal map of the star. Precise X-ray timing, assessing the $\\dot{P}$ and measuring the star dipole field, will add a crucial piece of information. Coupled to further sensitive phase-resolved spectroscopy to monitor spectral variability, this could help to solve the puzzles set by our results on \\src\\, which would have important implications for the understanding of the nature of CCOs and of their relations with other families of neutron stars.  \\vspace{-2mm}"
    },
    "1112/1112.4328_arXiv.txt": {
        "abstract": "{Westerlund I is the richest young cluster currently known in our Galaxy, making it one of the most massive clusters for which we can resolve the individual stars even in the crowded centre. This makes it an ideal target to assess whether massive clusters formed currently will remain bound or will disperse and contribute significantly to the stellar field population.} {Here we report a measurement of the radial velocity dispersion of Westerlund I to explore whether the cluster is currently in virial equilibrium, if it is in the process of collapse or if it is expanding and dispersing into the field.} {We obtained MIKE/Magellan high resolution optical spectra of 22 post main-sequence stars in Westerlund I for 2 or 3 epochs with a maximum baseline of about one year. Radial velocities variations between these spectra have been measured through cross correlation.} {We calculate the velocity dispersion from the cross correlation of five yellow hypergiants and one luminous blue variable, that show little radial velocity variations between epochs and have many spectral features in common. After taking into account the effect of small number statistics and undetected binaries, we estimate the velocity dispersion for the massive stars in Westerlund I to be $2.1^{+3.3}_{-2.1}$ \\, km s$^{-1}$. For several different assumptions concerning possible mass segregation and the elongation of the cluster, we find that Westerlund I is subvirial at the 90\\% confidence level.} {We can rule out that the cluster is significantly supervirial at the 97\\% confidence level, indicating that Westerlund I is currently bound. This implies that Westerlund I has survived past the point where any gas expulsion has taken place and is expected to survive for billions of years.} ",
        "introduction": "Soon after its launch about twenty years ago, the Hubble Space Telescope identified many massive young star clusters in starburst and merging galaxies \\citep{Hol92}. These clusters are massive enough that they appear unlikely to disperse in a Hubble time due to tidal stripping \\citep{Lam05}, which means they could possibly become future globular clusters. However, violent internal or external events could cause the cluster to become unbound, in which case the cluster will dissolve into the field, significantly contributing to the stellar populations of their host galaxy. A likely candidate for such an event would be the expulsion by stellar winds or the first supernovae of the remainder of the gas cloud out of which the cluster formed. This scenario has been extensively studied both analytically \\citep{Hills80, Elmegreen83,Mathieu83} and with N-body simulations \\citep{Lada84,Goo06}. These studies find that the expulsion of gas lowers the potential well in which the stars reside, causing a previously virial cluster to expand. The efficacy of the gas expulsion to disrupt clusters depends on several variables, including the star formation efficiency, the timescale of the expulsion \\citep{Hills80,Mathieu83} and whether the stars have achieved virial equilibrium with the potential of the surrounding gas before the gas is ejected \\citep{Goo09}. Without studying the dynamical state of these clusters, we do not know which fraction of them would survive to become future globular clusters analogues and which ones will disperse into the field.  The only information on the dynamic state of these distant massive clusters, that we can currently obtain is the radial velocity dispersion \\citep[e.g.][]{Bas06}. Under the assumption of virial equilibrium the one-dimensional velocity dispersion can be used to estimate the mass of the cluster, if its size is resolved. This mass estimate is called the dynamical mass. The dynamical mass can be compared to the mass derived from the luminosity of the cluster, assuming a universal IMF. Many clusters have been found to have a significantly higher velocity dispersion, than expected from the luminosity \\citep{Bas06}. Several explanations have been proposed for this, including: (i) the IMF might actually be different under extreme star forming conditions \\citep{Men02}, (ii) the velocity dispersion might be significantly increased due to binaries \\citep{Kou08,Gie10}, or (iii) the cluster might be supervirial, for example due to a recent expulsion of gas from the cluster \\citep{Goo06,Bau07}. These scenarios can be distinguished through observations of a cluster for which we can resolve the individual stars and thus measure the binary properties and the IMF.  Westerlund I is the most massive young cluster currently known in our Galaxy \\citep{Por10}. Photometric observations with the NTT/SOFI in the JHKs bands \\citep{Bra08,Gen11} show an elongated cluster with a mass of $4.9^{+1.8}_{-0.5} \\times 10^4 M_\\odot$, a half mass radius $r_{\\rm hm}=1.0$ pc, and a distance of $4 \\pm 0.2$ kpc. This distance is confirmed by comparing the radial velocity of the H II gas in Westerlund I with the rotation curve of the galaxy \\citep{Kot07}. At this distance from the galactic centre stellar evolution and galactic tidal stripping are expected to disperse Westerlund I over $\\sim 9$ Gyr, if it is currently in virial equilibrium \\citep{Bau03}. Thus Westerlund I might represent a precursor of a future low-mass, metal-rich \"globular cluster\". However, at an age of 4.5-5 Myr \\citep{Cro06,Gen11} we expect that Westerlund I might be out of virial equilibrium due to any recent expulsion of gas, which has not been converted into stars \\citep{Goo06}. In this paper we describe our measurements of the radial velocity dispersion in Westerlund I to explore whether the cluster is likely to be bound or not. If Westerlund I is in virial equilibrium we expect a one-dimensional velocity dispersion of $\\sim 4.6$ km s$^{-1}$ (see Sect. \\ref{sec:disc}) for the mass and radius measured by \\citet{Bra08} and \\citet{Gen11}.  \\citet{Men07,Men09} previously measured the radial velocity dispersion of Westerlund I using K-band spectra. They found a velocity dispersion of $9.2 \\pm 2.5$ km s$^{-1}$ based on a sample of ten stars. This value could be severely inflated due to the presence of binaries \\citep{Kou08,Gie10}. \\citet{Rit09,Rit09b} have also measured the radial velocities for a large sample of members and candidate members of Westerlund I using VLT/FLAMES+GIRAFFE spectra from 8350 to 9000 \\AA{} over multiple epochs to check for variability and investigate the binary properties of these massive stars.  In this paper we present multi-epoch high-resolution spectroscopy to find the velocity dispersion corrected for binaries of Westerlund I and thus constrain the dynamical state of Westerlund I.  In Sect. \\ref{sec:obs} we present the observations and discuss the data reduction. We present the radial velocity variations between epochs due to binaries or atmospheric instabilities in the target stars, and calculate the velocity dispersion using the stars with low variability in Sect. \\ref{sec:res}. The implications of the measured velocity dispersion on the dynamical state of Westerlund I are discussed in Sect. \\ref{sec:disc}. Finally we present our conclusions in Sect. \\ref{sec:conc}. ",
        "conclusions": "} We present multi-epoch high resolution spectra of 22 massive post main-sequence stars in the young, massive, Galactic cluster Westerlund I. For 21 of these 22 stars the velocity variations have been measured between multiple epochs by cross correlation (see Table \\ref{tab:vel} or Fig \\ref{fig:binarity}). Many of these stars seem to show radial velocity variations at the one to ten km s$^{-1}$ level. Some of these variations might be explained by an underestimate of our uncertainties. Alternatively these variations could be caused either by binaries or by instabilities/waves in the complex atmospheres of these super- and hypergiants.  Using the five least variable yellow hypergiants and the LBV, we find a velocity dispersion for the massive stars in Westerlund I of $\\sigma=2.1^{+3.4}_{-0.9}$ km s$^{-1}$ with a confidence limit of 95\\%, assuming no significant errors in measured radial velocity differences between the stars. In reality even these least variable stars show radial velocity variations between epochs of a few km s$^{-1}$, implying that the quoted velocity dispersion is only an upper limit of the true velocity dispersion.  In Figure \\ref{fig:p_vdisp} the probability distribution of the velocity dispersion from our measurements is compared to several estimates of the velocity dispersion which we expect in virial equilibrium. These estimates have been calculated using different estimates of the total mass of Westerlund I, as well as different rough approximations of the effects of possible mass segregation and the effects of a velocity anisotropy, which is suggested by the observed spatial elongation of Westerlund I. For all of these cases the observations favor a subvirial cluster. For the lower photometric mass estimate of \\citet{Gen11} the observations are consistent with the cluster in virial equilibrium. In all cases we exclude that Westerlund I is significantly supervirial, implying that the cluster is bound, has survived any recent gas expulsion. Barring any violent interactions with other clusters or molecular clouds, Westerlund I, which might be the most massive young cluster in our Galaxy, is expected to survive for billions of years."
    },
    "1112/1112.4058_arXiv.txt": {
        "abstract": " ",
        "introduction": "The mass-to-light ($M/L$) ratio is a very important parameter that is determined by galactic stellar population (the distribution of stars in age, mass, and, to a lesser extent, stellar metallicity), dust (through affecting the galaxy's luminosity) and the relative mass of its nonstellar components. The latters include the dark matter, which is usually comparable in mass to the stellar component within the galaxy's optical boundaries (see, e.g., Khoperskov et al. 2010; Zasov et al. 2011; Bizyaev and Mitronova 2009), as well as the gas whose mass fraction can be significant in late-type galaxies. The stellar population models computed for certain accepted initial mass function (IMF) obviously predict $M/L_B$  and $M/L_V$ to be of the order of 1-10 solar units in the visible bands, and  $M/L_R$ or $M/L_K \\sim 0.5-2$ solar units in the near infrared, in dependence of the relative number of young stars (and therefore the galaxy's color) and the adopted IMF. The presence of the dark matter or the gas along with internal extinction can only increase the overal galactic $M/L$. Therefore, those galaxies whose total $M/L$ estimated within the optical boundaries looks lower than that provided by the stellar population modeling using the universal IMF (even ignoring any dark halo) are of especial interest. Such candidate galaxies were selected from the sample of objects with known rotational velocities by Saburova et al. (2009).  Since the reliable luminosity, color, and rotational velocity estimates are available not for all selected objects, the conclusion about the $M/L$ anomaly for each specific galaxy should be verified. An unusually low $M/L$ for a given color index, if confirmed, can point to the existence of stellar population with a non-standard IMF: such a galaxy must have a low relative number of stars with a mass less than the solar one because these stars determine the total mass of stellar population once have a little affect on the photometric features. This would imply the existence of special conditions for the formation of the the bulk of its stars. Note that the questions about the IMF universality and the possibility of a non-typical IMF shape are being actively discussed in the literature (see, e.g., Gilmore 2001; Kroupa 2002; Hoversten and Glazebrook 2008; Meureret al. 2009; Bastian et al. 2010; Dabringhausen et al. 2010).  In this paper we present the results of our photometry for nine galaxies selected from Saburova et al (2009) that were previously identified as presumably having low total $M/L$ ratios within the optical radius $R_{25}$ (the radius of the 25 $^m/\\Box''$ isophote in $B$) and analyze some possible sources of errors in the observed quantities. We also compare the maximum rotation velocity of the disk component $v_{max~disk}$, expected for a ''normal'' stellar population with the maximum observed  rotation velocity $v_{max}$. As a reference, we employ the model values of the mass-to-light ratios for the stellar population with Salpeter IMF following the evolutionary models by Bell \\& de Jong (2001). The condition $v_{max~disk} > v_{ max}$ suggests that the stellar population model with a regular IMFs is inapplicable and that the the real stellar disk is lighter than it is predicted by the stellar population models even if non-stellar components (dark matter and gas) are ignored. To verify the total luminosity estimate as well as the photometric parameters: color and inclination of the disk, we performed photometric observations of nine galaxies in which we suspect the abnormally low $M/L$. ",
        "conclusions": "We present the results of our surface photometry in the B, V and R bands for nine disk galaxies in which the discrepancy between low M/L ratio  and color of stellar population is suspected. We obtain the photometric profiles, the radial profiles of the color indices, the position angle, and the flattening of isophotes. We decomposed the images using the estimation of photometric parameters for individual components (bulge, bar and disk) of the galaxies.  We apply two methods for the dynamical and photometric disk mass estimates: by checking whether the colors agree with the total $M/L$ ratios of the galaxies, which may be considered as an upper limit of $M/L$ for the stellar population, and by comparing the observed rotational velocities of the galaxies with the maximum circular velocity of the disk expected from its photometric parameters. Our results show that there are no obvious contradictions between the dynamical and photometric mass estimates in most considered objects because the abnormally low dynamical $M/L$ estimates can be naturally explained by the uncertainty in determining the rotational velocity or by the errors in photometry. At the same time our photometric data for NGC 4826, NGC 6814, NGC 5347, and with lesser confidence for UGC 03685 and NGC 4214 suggest that these galaxies may actually be too ''light'' for their luminosity, and hence may have the IMF defficient by low-mass stars. The presence of dark matter in the galaxies which was ignored in our study would only increase the discrepancy between the dynamical and photometric $M/L$ estimates. The galaxies we discuss require a more complete study: first of all, more careful measurements of the rotational velocity far from the center are needed.  The low values of the mass-to-light ratio in galaxies, even if this is not a result of anomalous stellar population, suggest the low dark-to-luminous matter mass fraction in them.\\\\"
    },
    "1112/1112.2560.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {Conceptually reconstructing the physical conditions in relativistic jets, given the observed electromagnetic spectrum, is a complex inverse problem. } % aims heading (mandatory) { We aim to improve our understanding of the mechanisms operating in relativistic jets by modeling  their broadband electromagnetic spectrum.} % methods heading (mandatory) {We develop an inhomogeneous jet model where the injection of relativistic primary and secondary particles takes place in a spatially extended region.  We calculate the contribution of all particles species to the jet emissivity by several radiative processes, and assess the effect of gamma-ray absorption in internal and external photon fields. A number of specific models with different parameters are computed to explore the possibilities of this scenario. } % results heading (mandatory) {We obtain a variety of spectral shapes depending on the model parameters, some of them predicting significant gamma-ray emission. The observed broadband spectrum of the low-mass microquasar XTE J1118+480 can be satisfactorily reproduced by the model.} % conclusions heading (optional), leave it empty if necessary {Our results indicate that outbursts similar to those displayed in the past by XTE J1118+480 might be detected with present-day gamma-ray instruments.} ",
        "introduction": "\\label{sec:introduction}  Jets are collimated outflows. Astrophysical objects that launch jets have two common components: a rotating compact object and an accretion disk. Examples of these systems are gamma-ray bursts, active galactic nuclei, microquasars, and young stellar objects.  Microquasars are X-ray binaries (XRBs) that produce relativistic jets. The jets of microquasars have bulk Lorentz factors of \\mbox{$\\sim1-10$}, and a typical power of \\mbox{$10^{35-37}$ erg s$^{-1}$} (e.g. Gallo et al. 2005, Sell et al. 2010, see also Heinz \\& Grimm 2005 and references therein).  Steady jets in microquasars appear in the so-called low-hard spectral state. In this state, the X-ray spectrum is the sum of a low-luminosity black body component and a hard power-law that cuts off at  \\mbox{$\\sim 100-150$} keV. These characteristics are well reproduced by ``disk+corona''  models (see for example Poutanen et al. 1997, Dove et al. 1997, and Malzac et al. 2001). In these models, the thermal emission originates in an accretion disk (truncated at tens or hundreds of gravitational radii from the compact object), and the hard power-law is generated in a ``corona'' of hot plasma that surrounds the black hole.  The main radiative mechanism of the corona is the Comptonization of disk photons by thermal electrons or electron-positron pairs.  The presence of the corona can also explain the detection of the Fe K$\\alpha$ line at $\\sim7$ keV in some XRBs (e.g. Miller et al. 2002, Tomsick et al. 2009, Duro et al. 2011). The transition from an outer thin disk to an inner quasi-spherical corona is self-consistently predicted in some accretion models, such as the advection-dominated accretion flows (ADAFs) studied by Narayan \\& Yi (1994, 1995) and the adiabatic inflow--outflow solutions (ADIOS) of Blandford \\& Begelman (1999).  The ejection of jets  during the low-hard state is revealed by the flat radio spectrum, which is interpreted as synchrotron radiation of non-thermal electrons in an expanding outflow (Blandford \\& K\\\"onigl 1979). Broadband observations of some XRBs indicate that there is a correlation of the form $S_{\\rm{radio}}\\propto S_X^{0.7}$ between the radio and the X-ray flux (Gallo et al. 2003). This suggests that the jets might contribute significantly to, or even dominate, the X-ray emission in the low-hard state. Two confirmed microquasar - Cygnus X-1 (McConnell et al. 2002, Albert et al. 2007, Sabatini et al. 2010, Laurent et al. 2011) and Cygnus X-3 (Abdo et al. 2009, Tavani et al. 2009) - have been detected in gamma rays. The jets - loaded with high-energy particles and moving at relativistic bulk velocities - appear as the probable site of the gamma-ray emission.\\footnote{If a fraction of the particles in the corona are relativistic, the corona might also radiate gamma rays; see Romero et al. (2010) and Vieyro \\& Romero (in preparation).}  Jet radiative models have been typically divided into ``leptonic'' (e.g. Markoff et al. 2001, Kaufman Bernad\\'o et al. 2002, Bosch-Ramon et al. 2006) and ``hadronic'' (e.g. Romero et al. 2003, 2005). In  hadronic (leptonic) models, the bulk of the gamma-ray emission is produced in interactions initiated by relativistic protons (electrons). Although microquasar jets can propagate over distances of hundreds of AU (Mirabel \\& Rodr\\'iguez 1999), most of these  models are based on the ``one-zone'' approximation: the emission region is assumed to be homogeneous and compact.  In this work, we present an inhomogeneous lepto-hadronic model for the radiation of jets. We intend to expand available models by adopting more general assumptions about the contents of relativistic particles and the characteristics of the acceleration/emission region. In our model, relativistic protons and electrons are injected in a spatially extended and inhomogeneous region of the jet. We investigate the consequences of the propagation and cooling of these particles (and also of secondary muons, pions, and pairs) on the jet radiative output. The characteristics of the spectral energy distributions (SEDs) can provide insight into the physical conditions in the jets. As an application, we fit the spectrum of the low-mass XRB XTE J1118+480, a very well-studied black hole candidate in the Galactic halo.  The paper is organized as follows. In Sect. \\ref{sec:jet_model}, we introduce the main features of the model. We present the general assumptions made about the geometry and energetics of the jet, and briefly describe the interaction processes between the relativistic particles and the matter, radiation, and magnetic field. We outline the method of calculation of the particle distributions and the jet radiative spectrum, including absorption corrections. In Sect. \\ref{sec:general_results}, we exhibit and analyze the results of a series of models calculated for different sets of parameters. Section \\ref{sec:XTE} presents the application to XTE J1118+480. We show the best-fit SEDs and discuss the possible detection of the source with present and future gamma-ray telescopes. We close with a short review of the model results and the perspectives for future studies.   %__________________________________________________________________ ",
        "conclusions": "We have developed an inhomogeneous lepto-hadronic jet model, and presented some general results for the relativistic particle distributions, cooling rates, and the jet radiative spectrum. The model can satisfactorily reproduce the observed SED of the low-mass microquasar XTE J1118+480 in low-hard state. The best-fit SEDs have revealed some interesting consequences of the injection of relativistic protons in an inhomogeneous extended region. In particular, we have demonstrated that the gamma-ray emission can escape absorption, and that the source might be detectable when a new outburst of similar characteristics occurs.  With some improvements, the same jet model can be applied to active galactic nuclei and gamma-ray bursts. The main modifications concern the calculation of the particle distributions. Eq. (\\ref{eq:transport_equation}) would have to be modified to a covariant version, which could then be applied to outflows with very high bulk Lorentz factors. In addition, within environments where the radiation field is ultra-dense (as for the jets of long gamma-ray bursts), the coupling among the kinetic equations for particles and photons could not be ignored. A further generalization of the transport equation would be to include time dependence, allowing the study of flares. Other interesting improvements to the model would be in the treatment of the jet kinematics and energetics, to develop a jet model that consistently couples the dynamical and radiative aspects. All these topics are part of our research project, and will be addressed in future works.   %______________________________________________________________"
    },
    "1112/1112.2162_arXiv.txt": {
        "abstract": "The desorption characteristics of molecules on interstellar dust grains are important for modelling the behaviour of molecules in icy mantles and, critically, in describing the solid-gas interface. In this study, a series of laboratory experiments exploring the desorption of three small molecules from three astrophysically relevant surfaces are presented. The desorption of CO, O$_2$ and CO$_2$ at both sub-monolayer and multilayer coverages was investigated from non-porous water, crystalline water and silicate surfaces. Experimental data was modelled using the Polanyi-Wigner equation to produce a mathematical description of the desorption of each molecular species from each type of surface, uniquely describing both the monolayer and multilayer desorption in a single combined model. The implications of desorption behaviour over astrophysically relevant timescales are discussed. ",
        "introduction": "The interaction of molecular species with surfaces is of critical importance in astrophysical environments. Much interstellar chemistry occurs on or in the icy layers which cover dust grains in molecular clouds, the birthplace of stars and planets \\citep{Fraser02}. These dust grains are believed to be composed of silicates and carbonaceous material \\citep{Draine03}. The silicate is at least $\\sim$~95~\\% amorphous, as determined from the breadth of the 9.7~$\\mu$m band \\citep{Li2002}, and although its composition has not been determined precisely, is well approximated by amorphous olivine (MgFeSiO$_4$) \\citep{Sofia01,Draine03}. The silicates are believed to be 'fluffy' in nature, with a large surface area upon which reactions can occur \\citep{Mathis98}.  At the low temperatures in molecular clouds (T~$<$~20~K) atoms and molecules freeze out onto these silicate surfaces, forming the icy mantles. Amorphous solid water (ASW) is by far the largest component of this ice, with abundances of $\\sim$~1~$\\times$~10$^{-4}$ with respect to the total H column density \\citep{Williams02}, equivalent to coverages of up to 100 monolayers (ML). Given that the extinction threshold for H$_2$O mantles is A$_V$~$\\sim$~3.3 mag \\citep{Whittet88}, it is reasonable to assume that between the dense pre-stellar cores (where dust grains are completely coated by this icy mantle) and the cloud edges (where competition between ice formation and photodesorption of H$_2$O yields a population of bare silicate grains \\citep{Smith93}), there must be a region where icy surfaces and bare silicates co-exist.  As interstellar regions evolve, the icy grains are further processed, either by gentle heating or cyclic desorption-deposition events \\citep{Visser09}, producing crystalline H$_2$O. This has been detected in various stellar and pre-planetary environments, including, for example, M giant stars \\citep{Omont90}, Quaoar in the Kuiper Belt \\citep{Jewitt04}, Trans Neptune Objects \\citep{Merlin07}, comets \\citep{Lisse06}, and in the outer disks around T Tauri stars \\citep{Schegerer10}.  As gas-grain modelling of interstellar environments becomes more sophisticated \\citep{Wakelam10,Acharyya11,Cuppen11}, key questions impeding the full implementation of surface chemistry include: what effect does the underlying grain surface have on the desorption characteristics of key molecules, and to what extent is this desorption affected as we move from the multilayer to sub-monolayer coverage regime? A third issue is how to realise the transition between the multilayer and monolayer regimes in a gas-grain model without overloading the model with many more layers of complexity.  CO, O$_2$ and CO$_2$ are all interstellar molecules that potentially could populate bare interstellar grains at sub-monolayer coverages. On ice, only one previous study focussed on sub-monolayer coverages of CO, highlighting the spectroscopic rather than desorption characteristics of the porous ASW:CO system \\citep{Collings05}. To date, no desorption studies of CO, O$_2$ or CO$_2$ have been made on a silicate surface, at either multilayer or sub-monolayer coverages.  CO is the second most abundant interstellar molecule \\citep{Tielens05}, and is known to form in the gas phase, then freeze-out on to H$_2$O covered grains to form overlayers of pure CO ice, the layer thickness critically depending on gas density rather than grain temperature \\citep{Pontoppidan03}.  Consequently, extensive temperature programmed desorption (TPD) studies have been made of multilayer CO coverages on various surfaces, including Au \\citep{Acharyya07,Bisschop06,Collings04,Fuchs06}, H$_2$O \\citep{Collings03a,Collings03b}, a meteorite sample \\citep{Mautner06} and highly oriented pyrolytic graphite (HOPG) \\citep{Ulbricht06}.  It is widely accepted that both CO$_2$ and H$_2$O form on dust grains in molecular clouds, and recent studies suggest that at least some of the CO$_2$ and H$_2$O is formed concurrently \\citep{Goumans08,Noble11,Ioppolo11}. Since CO is the key precursor to CO$_2$ formation, such mechanisms would require CO freeze-out onto bare grain surfaces long before CO ice, or even large quantities of H$_2$O ice, are detected. Although tenuous, spectroscopic observational and experimental evidence exists for such a freeze-out process \\citep{Fraser05}, but an investigation of the desorption behaviour of CO at sub-monolayer coverages on ice and silicate surfaces is vital if we want to be certain that it can reside at an interstellar grain surface long enough to form more complex species.  Similarly, O$_2$ is a potential precursor in H$_2$O formation \\citep{Ioppolo08,Oba09}, and is also a species likely to form in the gas phase then freeze-out onto grain surfaces, rather than forming on the grain itself. However, as a homonuclear diatomic O$_2$ is infrared inactive, so its detection in interstellar ice, though occasionally claimed via a forbidden transition \\citep{Elsila97} remains elusive. It is not clear whether this is related to the weak transition probability, lack of a significant O$_2$ population in the ice, or that the O$_2$ has rapidly reacted (upon adsorption) to exclusively form H$_2$O.  Nevertheless, the multilayer desorption behaviour of O$_2$ has been studied previously alongside CO and N$_2$ \\citep{Acharyya07,Fuchs06}, as well as on Au, porous ASW \\citep{Collings04} and TiO$_2$ \\citep{Dohnalek06}. These studies show that the multilayer desorption characteristics of O$_2$ are very similar to CO, so it is interesting from both a chemical and astrophysical viewpoint to also investigate the desorption behaviour of O$_2$ at sub-monolayer coverages on ice and silicate surfaces.  The desorption characteristics of CO$_2$ have not been extensively studied on any surface, despite it being one of the most abundant solid phase molecular species in the interstellar medium. Recently the desorption characteristics of multilayer CO$_2$ were reported from porous ASW, and have been previously studied on HOPG \\citep{Ulbricht06}, porous ASW and Au \\citep{Collings04}. However, given that the observational \\citep{Pontoppidan06} and experimental \\citep{Noble11} evidence shows that a fraction of CO$_2$ ice must form concurrently with the water ice layer, some CO$_2$ molecules must populate both the bare silicate grains and the ice layers at sub-monolayer coverages.  Here we present an experimental study of the desorption of CO, O$_2$ and CO$_2$ from three different surfaces; non-porous ASW, crystalline ice, and amorphous olivine-type silicate. As these are, for the first time, all undertaken in the same experimental set-up -- FORMOLISM \\citep{Amiaud2006} -- we are able to investigate both the individual effect of each surface on the desorption characteristics, as well as determining whether the molecular composition or morphology of the surface is most relevant in determining the desorption behaviour of molecules. For each of the nine combinations of CO, O$_2$ and CO$_2$ on each surface, our study has encompassed both the multilayer and sub-monolayer regimes. These data are modelled to determine a simple analytical expression which accurately calculates both the sub-monolayer and the multilayer desorption energies. By changing the model to incorporate interstellar, rather than experimental, heating rates, we are able to address the key questions above, namely: what effect does the underlying surface have on the desorption characteristics of the molecules adsorbed there, and how is this desorption modified in the sub-monolayer coverage regime?  \\section[]{Experimental}  The experiments were conducted using the FORMOLISM set-up (FORmation of MOLecules in the InterStellar Medium), described in detail elsewhere \\citep{Amiaud2006,Lattelais11}. Briefly, the set-up consists of an ultra-high vacuum chamber (base pressure $\\sim$~10$^{-10}$ mbar), containing a silicate-coated copper sample surface, operating at temperatures between 18 and 400~K. The system is equipped with a quadrupole mass spectrometer (QMS), which is used for the TPD experiments. A sample of either O$_2$, CO or CO$_2$ was deposited onto the surface at 18~K via the triply differentially pumped beam line; a linear temperature ramp was then applied to the surface, and the QMS used to measure the desorption of each species into the gas phase, as a function of temperature.  \\begin{table} \\caption{Description of experimental exposures.} \\label{tbl:depot} \\begin{tabular}{@{}lll@{}} \\hline Molecule                        & Surface              & Depositions$^a$\\\\ &                          & ML     \\\\ \\hline \\multirow{4}{*}{O$_2$}   & H$_2$O$_{(np)}$ & 0.25, 0.50, 0.75, 1, 1.25$^c$, 1.38, 1.75, \\\\ &                          & 2.25$^c$, 2.75\\\\ & H$_2$O$_{(c)}$   & 0.08, 0.17, 0.50, 0.67, 1, 1.33\\\\ &SiO$_x$              & 0.20, 0.60, 1, 1.20, 1.60, 1.67\\\\ \\hline \\multirow{4}{*}{CO}        & H$_2$O$_{(np)}$ & 0.25, 0.50, 0.75, 0.88, 1, 1.06$^b$, 1.13$^b$, \\\\ &                          & 1.25$^b$, 1.75$^b$, 2.50\\\\ & H$_2$O$_{(c)}$   & 0.10, 0.20, 0.60, 0.80, 1, 1.20$^b$, 1.60\\\\ &SiO$_x$              & 0.50, 1, 1.25$^b$, 1.75, 2.25\\\\ \\hline \\multirow{3}{*}{CO$_2$} & H$_2$O$_{(np)}$ & 1, 2$^b$, 4$^b$, 6$^b$, 10$^b$, 20\\\\ & H$_2$O$_{(c)}$   & 1, 2, 3, 5$^c$, 10, 20\\\\ &SiO$_x$              & 0.5, 1, 2, 3, 5, 10\\\\ \\hline \\end{tabular}  $^a$ defined as $\\frac{exposure}{N_{mono}}$, where N$_{mono}$ is defined in \\S~\\ref{sec:expt}.\\\\ $^b$ Omitted from model (shown in grey in Figure~\\ref{fgr:expt}).\\\\ $^c$ Modelled, but omitted from Figure~\\ref{fgr:model} for clarity. \\end{table}  Three surfaces were investigated in this study: non-porous ASW (H$_2$O$_{(np)}$), crystalline ice (H$_2$O$_{(c)}$) and amorphous silicate (SiO$_x$). The SiO$_x$ surface was recently installed in the FORMOLISM experiment \\citep{Lemaire10}, and mimics bare dust grains in molecular clouds. The silicate is amorphous in nature, as evidenced by infrared spectroscopic studies, while TPD experiments, including those presented in this study, reveal the surface to be non-porous on the molecular scale.  The silicate surface is capable of reaching base temperatures of 18~K, which is how the molecular deposition temperature on all the surfaces was pre-determined. For the ice surfaces, 50~ML films were grown on top of the silicate surface by spraying water vapour from a microchannel array doser located 2~cm in front of the surface. The water vapour was obtained from deionised water which had been purified by several freeze-pump-thaw cycles, carried out under vacuum. H$_2$O$_{(np)}$ mimics the ASW which comprises the bulk of interstellar ice, and H$_2$O$_{(c)}$ mimics the crystalline ice seen in some star-forming regions. To produce H$_2$O$_{(np)}$, water was dosed while the surface was held at a constant temperature of 120~K. To form H$_2$O$_{(c)}$, the surface was held at 120~K during the deposition, then flash heated at 50~Kmin$^{-1}$ to 140~K, and finally at 10~Kmin$^{-1}$ to 142.5~K. For each type of ice surface, the temperature was then held constant until the background pressure in the chamber stabilised, before cooling it back down to 18~K, at which temperature adsorbates were dosed onto the respective surfaces.  \\begin{figure*} \\centering \\includegraphics[width=\\textwidth]{noble_fig1} \\caption{Temperature-programmed desorption spectra of O$_2$, CO and CO$_2$ from H$_2$O$_{(np)}$, H$_2$O$_{(c)}$ and SiO$_x$. For each molecule-surface combination, the main window shows a series of TPD spectra from the lowest sub-monolayer coverage up to 1~ML; inset are TPD spectra from 1~ML up to the highest multilayer coverage, as specified in Table~\\ref{tbl:depot}. Spectra shown in black were included in the model, whilst those in grey were omitted. In all cases, the line profile of the TPD spectra in the multilayer regime differ from those in the sub-monolayer regime and, with the exception of CO, the form of the TPD profiles is very similar in the sub-monolayer regime, independent of surface. The desorption profiles of all three molecules from the SiO$_2$ surface in the sub-monolayer regime strongly resemble each other, unlike the ices. See text for full discussion.} \\label{fgr:expt} \\end{figure*}  For each of the nine combinations of molecule and surface, a series of depositions were made, varying the adsorbate surface coverage, from sub-monolayer to multilayer, as outlined in Table~\\ref{tbl:depot}. After the first adsorbate exposure, the surface was then heated at a rate of 10~Kmin$^{-1}$ until the adsorbate was fully desorbed from the surface (below 100~K on the ice surfaces and 130~K on the silicate surface). The original surface therefore remained intact, and by purposefully keeping the temperature of the ice below 100~K no further thermal annealing of the ice occurred during an experimental sequence. Subsequently, the surface was then re-cooled to 18~K, and a new coverage of the same adsorbate added. In this way we could be certain the same surface was used for a single series of adsorbate exposures, allowing real comparison of the coverage effect for the first time. In the multilayer, the underlying surface is generally not important to the desorption, as adsorbate-adsorbate interactions between molecules determine the kinetics. In the sub-monolayer, the interaction between the adsorbate and the substrate is fundamental in determining the desorption characteristics. Thus, by including both regimes, we hoped to probe the differences between these two coverages for different molecules and different surfaces.  A porous ASW surface was not included in this study for two reasons. First, the presence of pores in the ice surface can produce very complex desorption profiles, due to the competing kinetics which arise after entrapment of adsorbates in the pores \\citep[e.g.][]{Collings03a}. Worse still, pore collapse starts at only 10--15~K, so the experimental approach used here, of exposing a single ice surface to a variety of adsorbate coverages, is simply not possible with the porous ASW. A new ASW sample would have been required for each experiment; each time we form a porous ice in the laboratory its structure will be slightly different, resulting in a massive change in surface area and non-reproducible adsorbate exposures being required to obtain monolayer coverages. Furthermore, one aim of this experiment is to compare surface characteristics using a single model for the whole dataset, and such an approach would not have been possible when including both porous and non-porous ices, again because the monolayer surface area is difficult to define. Secondly, it has recently been shown experimentally that it is very likely that all water ice in interstellar regions will be non-porous ASW \\citep{Accolla11}, even though observations previously suggested that invoking a porous ASW was the only way to explain observed quantities of mixed H$_2$O and CO ices \\citep{Pontoppidan03}. ",
        "conclusions": "A key aim of this work was to address what effect the underlying dust grain surface could have on the desorption characteristics of molecules adsorbed there, and how this desorption is modified in the sub-monolayer coverage regime compared to the multilayer regime.  This work used temperature programmed desorption to measure the desorption characteristics of O$_2$, CO and CO$_2$ from H$_2$O$_{(np)}$, H$_2$O$_{(c)}$ and SiO$_x$ over sub-monolayer to multilayer coverages. The experimental data were modelled using the Polanyi-Wigner equation, combining different approaches to reproduce the sub-monolayer and multilayer TPD traces, in particular by varying E$_{ads}$ as a function of coverage in the sub-monolayer regime. The desorption can be categorised as 'full wetting', 'intermediate' or 'non-wetting' behaviour, and the switching point is well described by N$_{mono}$, an empirical measure of the coverage at which pure monolayer desorption stops. On both a laboratory and an astrophysically relevant timescale, the desorption characteristics of molecules from the amorphous substrates H$_2$O$_{(np)}$ and SiO$_x$ were found to be very similar, while on the crystalline surface H$_2$O$_{(c)}$, molecules desorbed at higher temperatures and on a longer timescale.  It would be relatively trivial to implement the data presented here into astrochemical models; it would require only a subroutine to describe the sub-monolayer, coverage dependent, E$_{ads}$, which can be calculated using the parameters presented in this paper, and the coverage N$_{mono}$, at which the true monolayer desorption ends and beyond which a fixed E$_{ads}$, and zeroth order desorption can be implemented.  Previously, \\citet{Green09} concluded that surface type is relatively unimportant, and rather it is the heating rate and grain size that dominate desorption kinetics; we do not disagree that both these factors are important, but by not considering the sub-monolayer case, this previous work overlooks certain intricacies. Firstly, the results presented here suggest that the surface type is the dominant factor controlling desorption under sub-monolayer coverage conditions. This definition of surface type must include both the surface material and the degree of crystallinity, and not simply the material alone. From the data presented here, it is clear that the desorption characteristics of molecular species from H$_2$O$_{(np)}$ and SiO$_x$, two amorphous surfaces, are much more similar than those from H$_2$O$_{(np)}$ and H$_2$O$_{(c)}$, which are composed of the same underlying material. Secondly, while the size of the grain is critical to desorption characteristics, it is both the surface area and the surface-adsorbate interaction which must be considered. Without including information on whether the surface-adsorbate interaction is wetting or non-wetting, it is not possible to know the coverage at which behaviour switches from sub-monolayer (to intermediate) to the multilayer regime, i.e. the value of N$_{mono}$. As it is notoriously difficult to calculate physisorption surface-adsorbate energies theoretically, this is necessarily an empirically determined value, highlighting the importance of experiments such as those presented here."
    },
    "1112/1112.2448_arXiv.txt": {
        "abstract": "We present estimation of proton-air inelastic interaction cross-section obtained for different energy values are. Results are compared with different hadron interaction models. ",
        "introduction": "For more than 40 years the Yakutsk array continuously records ultra-high energy cosmic rays (UHECR) by measuring electromagnetic, muonic and Cherenkov components of extensive air showers (EAS)~\\cite{Artamonov1994}. Through these measurements a lateral distribution of charged particles and muons with $\\ethr \\ge 1$~GeV~\\cite{Dyakonov1991,Knurenko2006} was studied and longitudinal development of air showers was reconstructed by measuring the lateral distribution of Cherenkov light emitted in EAS~\\cite{Dyakonov1993,Knurenko2011}.  Estimation of the depth of maximum shower development ($\\xmax$) in individual showers with energy $10^{15} - 10^{19}$~eV at the Yakutsk array was performed using readings from integral and differential (tracking) Cherenkov light detectors. Integral detectors provide observation of Cherenkov photons lateral distribution in wide distance range, while differential detectors record time-base of the Cherenkov light pulse at different core distances. It follows from model calculations that characteristics such as $P = \\log_{10}{Q(200)/Q(550)}$ (a relation between the density of Cherenkov light fluxes at $200$ and $550$~m from shower axis) and $\\tauhw$ (the width of the Cherenkov pulse at the half-height) are sensitive to the depth of maximum shower development. Besides, the $P$ parameter depends weakly on conditions of Cherenkov light distribution in the atmosphere ($\\sim 8$\\,\\%), but depends strongly on the accuracy of shower geometry reconstruction ($\\sim 13$\\,\\%). The $\\tauhw$ is measured by equipment that provides high precision of the Cherenkov pulse shape reproduction and thus assures better accuracy in $\\xmax$ measurement ($\\sim 25$~g/cm$^{-2}$). The methodology of Cherenkov light measurement and analysis is presented in a greater detail in papers~\\cite{Dyakonov1991, Ivanov2007}. ",
        "conclusions": "The Yakutsk EAS array is a medium sized array compared to compact and large arrays. It effectively controls the energy range from $10^{15}$~eV to $5 \\times 10^{19}$~eV. In several decades of continuous observations a large amount of experimental data was accumulated. Both irregularities in cosmic ray energy spectrum are manifested in these data and the mass composition of primary particles initiating ultra-high energy air showers has been estimated using one technique. Such simultaneous analysis of the spectrum and mass composition allows one to test various hypotheses on cosmic rays origin and their propagation in the Universe. To increase the precision of this analysis it would take a choice of a single hadron interaction model from the group of models presented today. Studying the energy dependence of inelastic interaction cross-section just helps make such a choice. For example, the choice of the model for EAS development would allow experimenters to improve the precision in energy estimation and obtain the model-independent estimation of cosmic ray mass composition in ultra-high energy domain, more reliable than those adopted nowadays.  The work is supported in part by SB RAS (integral project ``Modernization of the Yakutsk array''), RFBR (grant \\#11--02--00158) and the Russian Ministry of Education and Science (contract \\#02.740.11.0248)."
    },
    "1112/1112.5432_arXiv.txt": {
        "abstract": "In this paper we consider a new mechanism for stopping the  inward migration of a low-mass planet embedded in a gaseous protoplanetary disc. It operates when a low-mass planet (for example a super-Earth), encounters outgoing  density waves excited by another  source in the disc. This source could be a gas giant  in an orbit interior to that of the low-mass planet.  As the super-Earth passes through the wave field, angular momentum is  transferred to the disc material and then   communicated to the planet through coorbital dynamics, with the consequence that its inward migration can be  halted or even reversed.  We illustrate how the mechanism we consider works in a variety of different physical conditions employing global two-dimensional hydrodynamical calculations. We confirm our results  by performing local shearing box simulations in which  the super-Earth interacts with density waves excited by an independent harmonically varying potential.  Finally, we discuss the constraints arising from the process considered here, on  formation scenarios  for  systems containing a giant planet and lower mass planet in an outer orbit with a 2:1 commensurability such as GJ876. ",
        "introduction": "\\label{introduction} The orbital migration of low-mass planets in  protoplanetary discs may play an important role in shaping planetary systems. In our previous studies \\citep{papszusz2005, paperI, paperII, papszusz2010} we have concentrated on the possibility of the occurrence of a resonant structure during early phases of the evolution of the  planetary system. We have considered two-planet systems, containing at least one low-mass planet of several Earth masses,  embedded in a gaseous disc. In \\cite{paperI} we  found that in a system with one low-mass planet (described here as a super-Earth) and a  gas giant, the formation of commensurabilities  readily occurs  when  the  migration of the two planets is convergent with  the gas giant being  on the larger orbit. However, as  shown in \\cite{paperII},  this is no longer true if the migration is convergent but the gas giant is on the smaller  orbit. In this case  the super-Earth never gets close enough to the location of the 2:1 commensurability in order to be caught in it. The super-Earth  is initially found to  migrate towards the central star  but at some point stops or even starts to migrate slowly outward. However, no explanation of this phenomenon was given. The aim of our paper is to explore this behavior more fully using both global and local shearing box simulations with the aim of delineating the mechanism responsible for the outward migration, showing  how it depends on model parameters such as the mass of the giant planet and the disc aspect ratio.  Since it was realized that the fast migration of  low-mass planets \\citep{ward97} could pose a serious challenge to scenarios for  planet formation, possibilities for slowing down such migration have been the subject of intensive  studies. A number of potential mechanisms have been found by focusing on a single low-mass planet embedded in a  protoplanetary disc. These mechanisms involve, among others,  entry into a magnetospheric cavity close to the star \\citep{lin1996},  effects arising from the orbital eccentricity of a protoplanet \\citep{paplar},  effects due to the possible  eccentricity of the protoplanetary disc  \\citep{pap2002}, magnetic fields \\citep{terquem},  MHD turbulence \\citep{laughlin2004, nelpap2004, johnson, adamsbloch},  sudden  jumps in disc state variables \\citep{menougood, matsumura},  corotation torques  \\citep{masset, ppa, ppb}, disc thermodynamics \\citep{PaarMel2006, BarMas2008a, KleyCrida2008, pp2008, kley2009, Paa2010, hasegawa, Paa2011, yamada} and instabilities of partial gap edges in low viscosity discs  \\citep{li2009, li2010}.  However, there may be mechanisms for affecting migration that depend on having a system of planets. For example, \\cite{ogi} have suggested  the possibility of  ``an eccentricity trap'' operating for a resonantly interacting convoy of planets in a way to halt type I migration  near the inner edge of a protoplanetary disc. \\cite{thommes05}  suggested that gas giants may be   efficient at capturing low-mass planets in their exterior mean-motion resonances, forming in this way a barrier in a protoplanetary disc ``a safety net for fast migrators''.  A follow up study based on numerical simulations of the disc planet interactions  has  been been presented by \\cite{pierens}. They found that planets with masses in the range 3.5-20 $M_{\\oplus}$ become trapped at the edge of the gap formed by the giant planet,  while more massive planets are captured into resonance.  They  noted that the positive surface density gradient  resembled the configuration discussed by \\citet{masset}  for their planet trap. Accordingly they raised the possibility that corotation torques operating as in a tidally  undisturbed disc could be responsible. However, the configuration differs significantly from that of \\citet{masset}. The surface density profile in the  latter's   planet trap is maintained by a viscosity that increases rapidly inwards.  In the absence of a viscous angular momentum flux produced by the action of viscous  stresses at an inner boundary,  acceleration of the accretion flow would occur, that in turn causes a decrease in the surface density in a steady state. However, in a disc tidally disturbed  by a giant planet,  the profile is maintained by the outward tidal transport of angular momentum produced by dissipating density waves rather than by either an accretion flow, or effects of comparable magnitude resulting from applied viscous stresses. The  wave transport also occurs in the coorbital region of a migrating planet and thus coorbital torques may be expected to behave differently  from those occurring  in an undisturbed disc.  Notably,  having an  external source of angular momentum for the coorbital region means that  positive corotation torques can be sustained without  either the action of applied viscous stresses, or  the accretion  of exterior  fresh material into the coorbital zone followed by its  loss into the inner regions, which would otherwise be required. Put another way, an external source of angular momentum for the coorbital zone could  prevent corotation torques from saturating  as would be expected for a tidally undisturbed disc with small applied viscosity. Thus, as seen in our simulations,  long term positive torques may be communicated to the planet even when the nominal viscosity is set to zero. In this way, even though there are some similarities, effects seen in a disc that is tidally disturbed by a giant planet   may differ from those   found in a disc with a surface density profile maintained through the application of a variable viscosity.  \\begin{figure*} \\begin{minipage}[!htb]{160mm} \\centering \\includegraphics[width=70mm]{fig1alast.pdf} \\includegraphics[width=70mm]{fig1blast.pdf} \\caption{\\label{2jow}{The left panel shows the evolution of the semi-major axis of a super-Earth in a presence of a gas giant with  one Jupiter mass (black curve) and two Jupiter masses  (grey curve). The right panel shows the variation of the torque acting on the super-Earth: coming from the  disc  matter located  in the vicinity of the planet, in the region between $r=1.4$ and $r=1.7$. The timescale covers many periods of the super-Earth, which is located at $r \\sim 1.545$  which is close to the azimuthally averaged surface density maximum that occurs slightly exterior to the gap edge. Note during this time period, the giant planet is located at $r\\sim 0.9,$ to which it has migrated.}} \\end{minipage} \\end{figure*}    \\begin{figure*} \\begin{minipage}[!htb]{160mm} \\centering \\includegraphics[width=75mm]{fig2alast.pdf} \\includegraphics[width=75mm]{fig2clast.pdf} \\end{minipage} \\begin{minipage}[!htb]{160mm} \\centering \\includegraphics[width=75mm]{fig2blast.pdf} \\includegraphics[width=75mm]{fig2dlast.pdf} \\end{minipage} \\caption{\\label{oneorbit}{({\\it Left panel}) The semi-major axis of the super-Earth (top) and its distance from the central star (bottom) as a function of time. The time interval in each picture corresponds  roughly to one orbital period of the super-Earth. ({\\it Right panel}) The torque acting on the super-Earth: total torque (top), the one coming from the disc matter located in the close vicinity of the planet, namely between $r=1.4$ and $r=1.7$ (bottom). The vertical lines in the right bottom panel indicate particular moments of time for which the planet positions in the disc are shown in Fig. \\ref{rysunek}. }} \\end{figure*}       The mechanism,  we  describe  in this paper, operates  in a similar physical situation where a low mass planet migrates in a disc which contains  trailing density waves  excited by another planet.  A  specific configuration of this type that we consider has  a gas giant in an interior orbit with a low mass planet in an exterior orbit around a central mass \\citep[see][]{paperII}. However, local simulations that we perform indicate that the essential features are a single planet  immersed in a field of propagating density waves produced  by  an external source. The waves are associated with an angular momentum flux which can enter the coorbital region of the low mass planet and transfer angular momentum to the disc material there through the dissipation of shocks. This in turn can transfer angular momentum to the low mass planet through coorbital dynamics. It is important to note that this does not depend on vortensity gradients as do standard corotation torques \\citep[e.g.][]{masset, ppa, ppb}. In addition because the waves act as a source of angular momentum from outside the coorbital zone and the transfer involves viscosity independent shock wave dissipation, the usual saturation considerations relevant to standard corotation torques do not apply.   This paper is organized as follows. In Section \\ref{modeleq} we describe the physical model and give the basic governing equations. Then in Section \\ref{setup} we give a brief review of the  type I migration theory  applicable to an isolated  super-Earth embedded in a protoplanetary disc without propagating density waves. We go on to present simulations of  a super-Earth embedded  in a protoplanetary  disc in which  density waves excited by an interior  giant planet propagate in Section   \\ref{supergiant}.  We show that in the presence of outward propagating trailing density waves,  the inward type I migration of a super-Earth can be reversed  with the consequence that attainment of a 2:1 commensurability by migrating inwards from large radii is prevented. We go on to describe local shearing box simulations  of a  planet  embedded in a disc with density  waves excited by an imposed harmonically varying potential in Section \\ref {shearingbox}.  We find similar results to those for the global simulations. The waves are found to induce angular momentum transport in the  disc material that is then passed on to the planet. In Section \\ref{theory} we consider this  mechanism, which  leads  to the  outward migration of the super-Earth in the global simulations, from  the point of view of the conservation of energy and angular momentum. We then consider some consequences of the wave-planet interaction for the resonance capture  into an outer 2:1 commensurability of a super-Earth by a gas giant in Section \\ref {observations}. In general density waves excited by the giant planet tend to prevent this, indicating that such commensurabilities , as seen for example in GJ876, may have formed through planetesimal migration after the gas disc dispersed. Finally in Section \\ref{conclusions} we give our conclusions. ",
        "conclusions": "\\label{conclusions} In this paper we have studied a new mechanism that can  reverse the type I migration of a low mass planet that would  occur in an unperturbed  locally isothermal gaseous disc. This mechanism operates when there is a source of trailing density waves that propagate through the disc. The most natural example of such source arises from the gravitational perturbation by another planet and for global simulations we have focused on the case when a gas giant is present. We have also carried out local simulations in a shearing box where the forcing was due to an imposed harmonically varying potential with varying amplitude. In all cases the density waves produce shocks which are associated with  angular momentum transfer to the disc material. Coorbital disc material  then transfers angular momentum to the low mass planet in the same direction, when  it is scattered by it. This is analogous to what happens with coorbital material undergoing a slow drag in horseshoe orbits in the case of zero pressure particle dynamics. However, in the cases considered here the dynamics of the coorbital material is not stationary in an appropriately rotating frame and generally more complex. In particular, we stress that although coorbital dynamics plays a role, the situation is unlike that pertaining to standard corotation torques in initially unperturbed discs. For example,  because the dissipation and effective drag are produced by shocks, which are viscosity independent and the propagating density waves provide an external  source of angular momentum for the coorbital zone, issues of torque saturation become  irrelevant.  Furthermore  although there is an interior gap, the mechanism described here depends on the  same source of angular momentum,  rather than  applied viscous stresses or an accelerating inward accretion flow, to maintain the surface  density distribution and  in this way it differs from the trap mechanism of \\citet{masset}.  The  effect of the interaction with the density waves is  that the  migration of a low-mass planet located in an exterior orbit relative to the gas giant can be slowed down and finally reversed. Thus we found that a  planet  with mass in the super-Earth range cannot approach a Jupiter mass  planet close enough in order to form first order mean-motion resonances with it. The migration was found to halt with semi-major axis ranging between $1.6$ and $2$ times that of the giant. Only when the low mass planet exceeded $\\sim 40 M_{\\oplus}$ was it able to attain a 2:1 commensurability.  As the giant planet mass is increased, even larger low mass planets would be required.  Our results indicate that migrating the outermost planet in the GJ876 system to its observed 2:1 commensurability through planet interaction with the gaseous disc alone would be problematic for the reasons outlined above. This may indicate that migration induced by planetesimals after the clearance of the gas disc may have been significant in that case.  \\newpage"
    },
    "1112/1112.0447_arXiv.txt": {
        "abstract": "We measured the angular diameter of the exoplanet host star $\\epsilon$ Eridani using the Navy Optical Interferometer. We determined its physical radius, effective temperature, and mass by combining our measurement with the star's parallax, photometry from the literature, and the Yonsei--Yale isochrones \\citep{2001ApJS..136..417Y}, respectively. We used the resulting stellar mass of 0.82$\\pm$0.05 $M_\\odot$ plus the mass function from \\citet{2006AJ....132.2206B} to calculate the planet's mass, which is 1.53$\\pm$0.22 $M_{\\rm Jupiter}$. Using our new effective temperature, we also estimated the extent of the habitable zone for the system. ",
        "introduction": "\\citet{2000ApJ...544L.145H} discovered a planet orbiting $\\epsilon$ Eri (HD 22049, K2 V) using radial velocity measurements that spanned 20 years. They determined the planet's orbit and a minimum mass of 0.86 $M_{\\rm Jupiter}$. They assumed a stellar mass of 0.85 $M_\\odot$ from \\citet{1993ApJ...412..797D}, who used spectroscopic observations to derive parameters such as effective temperature, surface gravity, metallicity, and luminosity and then used those parameters to calculate the star's mass.  \\citet{2004AandA...426..601D} observed $\\epsilon$ Eri using the Very Large Telescope Interferometer (VLTI) during the commissioning period of the VLT Interferometer Commissioning Instrument (VINCI). They measured a limb--darkened angular diameter of 2.148$\\pm$0.010$\\pm$0.027 mas with a global uncertainty of 0.029 mas, the last of which leads to an error of 1.4$\\%$. The first two errors on the diameter are the statistical and systematic errors, respectively. The statistical error is based on the dispersion of the data while the systematic error is based on uncertainties in the calibrators' diameters. Di Folco et al. then used the CESAM code version 4 evolutionary models by \\citet{1997AandAS..124..597M} to obtain various parameters for the host star, such as luminosity, effective temperature, age, mass, radius, surface gravity, and metallicity. The resulting CESAM mass for $\\epsilon$ Eri was 0.83 $M_\\odot$ while they list their estimated mass as 0.90$\\pm$0.10 $M_\\odot$ and the adopted mass as 0.84 $M_\\odot$.  \\citet{2006AJ....132.2206B} combined Hubble Space Telescope Fine Guidance Sensor observations of $\\epsilon$ Eri with ground--based astrometric and radial velocity data to produce an orbit for the system and the mass of the planetary companion. They assumed a mass for the central star of 0.83 $M_\\odot$ and a young age for the star, $\\sim$800 Myr. Their resulting $M_{\\rm planet}$ was 1.55$\\pm$0.24 $M_{\\rm Jupiter}$.  We used the Navy Optical Interferometer\\footnote{The NOI underwent a recent name change. It was previously known as the Navy Prototype Optical Interferometer.} (NOI) to measure the angular diameter of $\\epsilon$ Eri and then calculated its physical radius and effective temperature using the combination of our new measurement and other observed quantities, such as parallax, interstellar absorption, and bolometric correction. We provide an independent estimate of the star's mass by placing it on an H--R diagram using our values, which then leads to an estimate of the planet's mass. Section 2 describes the instrument and our observing procedure, Section 3 discusses how $\\epsilon$ Eri's angular diameter, physical radius, and $T_{\\rm eff}$ were determined, and Section 4 explores the physical implications of the new measurements including the mass determination and habitable zone calculation. ",
        "conclusions": "As a check for our result we estimated the LD diameter for $\\epsilon$ Eri using two additional methods. First, we created an SED fit for the star as described in Section 2 and Figure \\ref{cal_sed} shows the result. The only free parameter for the SED fit is the angular diameter. Second, we used the relationship described in \\citet{2004AandA...426..297K} between the ($V-K$) color and log $\\theta_{\\rm LD}$. Our measured $\\theta_{\\rm LD}$ is 2.153$\\pm$0.028 mas, the SED fit estimates 2.04$\\pm$0.11 mas, and the color--diameter relationship produces 2.04$\\pm$0.82 mas.  The main sources of errors for the three methods are uncertainties in visibilities for the interferometric measurement, uncertainties in the comparison between observed fluxes and the model fluxes for a given $T_{\\rm eff}$ and log~$g$ for the SED estimate, and uncertainties in the $K$ magnitude and in the parameters of the relation and the spread of stars around that relation for the color--diameter determination. All three diameters agree within their errors but our interferometric measurement provides an error 4 and 29 times smaller than the other methods, respectively.  Our measurement is consistent with that derived using the FLUOR beam combining instrument on the Center for High Angular Resolution Astronomy Array at 2.126$\\pm$0.014 mas \\citep{2007AandA...475..243D}. FLUOR observes in the $K'$--band at $1.94 - 2.34$ $\\mu$m and our measurement in visible wavelengths confirms their value. This was expected because both analyses incorporate limb darkening in their respective wavelengths, which will compensate for any difference due to the wavelength of the observations.  With our newly calculated $R$, $T_{\\rm eff}$, and $L$, we estimated the mass of the central star using the Yonsei--Yale isochrones \\citep[Y$^2$,][]{2001ApJS..136..417Y}. Figure \\ref{yy} shows the tracks for several masses and the best fit is for a 0.82 $M_\\odot$ star. We combined this mass with the mass function ($5.9 \\times 10^{-10}\\pm 1.0 \\times 10^{-10} M_\\odot$) and inclination (30.1$\\pm$3.8$^\\circ$) from \\citet{2006AJ....132.2206B} and calculated the planet's mass to be 1.53$\\pm$0.22 $M_{\\rm Jupiter}$, which is consistent with that listed in Benedict et al. of 1.55$\\pm$0.24 $M_{\\rm Jupiter}$. The Y$^2$ isochrones also provided an approximate age for the system, which is on the order of 1 Gyr. This is a slightly older age than the results from other studies, which are all less than 1 Gyr \\citep[e.g.,][]{1991ApJ...375..722S, 1996AJ....111..439H, 2000ApJ...533L..41S, 2004AandA...426..601D, 2005AandA...443..609S}. We consider our age to be consistent with those measured using other methods because we do not claim high accuracy in that parameter.  We also calculated the size of the system's habitable zone using the following equations from \\citet{2006ApJ...649.1010J}: \\begin{equation} S_{b,i}(T_{\\rm eff}) = (4.190 \\times 10^{-8} \\; T_{\\rm eff}^2) - (2.139 \\times 10^{-4} \\; T_{\\rm eff}) + 1.296 \\end{equation} and \\begin{equation} S_{b,o}(T_{\\rm eff}) = (6.190 \\times 10^{-9} \\; T_{\\rm eff}^2) - (1.319 \\times 10^{-5} \\; T_{\\rm eff}) + 0.2341 \\end{equation} where $S_{b,i}$($T_{\\rm eff}$) and $S_{b,o}$($T_{\\rm eff}$) are the critical fluxes at the inner and outer boundaries in units of the solar constant. The inner and outer physical boundaries $r_{i,o}$ in AU were then calculated using \\begin{equation} r_i = \\sqrt{ \\frac{L/L_\\odot}{S_{b,i}(T_{\\rm eff})} } \\; \\; \\; \\; \\; {\\rm and} \\; \\; \\; \\; \\; r_o = \\sqrt{ \\frac{L/L_\\odot}{S_{b,o}(T_{\\rm eff})} }. \\end{equation} We obtained habitable zone boundaries of 0.50 AU and 1.00 AU. $\\epsilon$ Eri's planet has a semimajor axis of 3.39$\\pm$0.36 AU \\citep{2006AJ....132.2206B}, so there is no chance the planet orbits anywhere near the habitable zone. Our values are slightly smaller than those of Jones et al. (0.572 and 1.123 AU).  $\\epsilon$ Eri features a dust ring that was imaged at a wavelength of 850 $\\mu$m using the Submillimetre Common User Bolometer Array at the James Clerk Maxwell Telescope by \\citet{1998ApJ...506L.133G}. They found the ring has a mass between 0.01 and 0.4 $M_{\\rm Earth}$, a mass comparable to the estimated mass of comets orbiting in our Solar System. The disk's emission peaks at 100 $\\mu$m \\citep{2000MNRAS.314..702D} and its flux compared to the central star is on the order of 10$^{\\rm -5}$, a contribution that puts it well outside the dynamical range of the NOI, particularly at the optical wavelengths that the NOI uses. At a distance of 55 AU from $\\epsilon$ Eri \\citep{2003ApJ...598L..51L}, the ring is also larger than the field of view of the NOI. Additionally, \\citet{2007AandA...475..243D} found no evidence of dust close to the star that would influence our visibility measurements."
    },
    "1112/1112.3358_arXiv.txt": {
        "abstract": "\\bub\\ is a pair of interacting galaxies at redshift $z=0.123$ that hosts a luminous obscured quasar in its northern nucleus.  Here we present two long-slit Magellan LDSS-3 spectra that reveal a pair of symmetric $\\sim$ 10 kpc-size outflows emerging from this nucleus, with observed expansion velocities of $\\sim$250 \\kms\\ in projection. We present a kinematic model of these outflows and argue that the deprojected physical velocities of expansion are likely $\\sim 1000$~\\kms\\ and that the kinetic energy of the expanding shells is likely $10^{44-45}$ erg~s$^{-1}$, with an absolute minimum of $>10^{42}$ erg~s$^{-1}$. Although a radio counterpart is detected at 1.4GHz, it is faint enough that the quasar is considered to be radio-quiet by all standard criteria, and there is no evidence of extended emission due to radio lobes, whether aged or continuously powered by an on-going jet. We argue that the on-going star formation is probably insufficient to power the observed energetic outflow and that \\bub\\ is a good case for radio-quiet quasar feedback.  In further support of this hypothesis, polarimetric observations show that the direction of quasar illumination is coincident with the direction of the outflow. ",
        "introduction": "The discovery of a tight relationship between the masses of black holes (BHs) in nearby galaxies and the velocity dispersions and masses of their stellar populations \\citep{magorrianetal1998, gebhardtetal2000a, ferraresemerritt2000} suggests that the active phase of the supermassive BH evolution has profound effects on the formation of the galaxy. Indeed, there is direct evidence for such `feedback' by radio-loud active galactic nuclei (AGNs) onto the large-scale gaseous environment of their host galaxies, both at low and at high redshifts \\citep{nesvadbaetal2008, fustockton2009}. In these objects, ionized gas is outflowing with velocities above the escape speed from the galaxy, often along the direction of the radio jet. This also appears to be the mode in which AGNs heat up galaxy clusters \\citep{mcnamaranulsen2007}. However, as only a minority ($\\sim$10\\%) of active BHs are radio-loud, this mechanism may not be sufficient to explain the colors and the luminosity function of massive red galaxies (e.g., \\citealt{springeletal2005}), and it is as yet unclear if the remaining $\\sim$90\\% of AGNs are affecting the gas, and thus the star formation, in the host galaxies in a similar fashion.  There is strong evidence that star formation driven winds play an important role in galaxy evolution.  Their imprint is apparent in the mass-metallicity relation at low redshift \\citep[e.g.,][]{tremontietal2004}, and they are required in galaxy simulations to reproduce the observations \\citep[e.g.,][]{springeletal2005,daveetal2008}.  Most importantly, winds driven by star formation are clearly seen in low-redshift dwarf galaxies \\citep[e.g.,][]{martin1998}, in local ultra-luminous infrared galaxies \\citep[ULIRGs; e.g.,][]{rupkeetal2005}, in galaxies at $z\\sim1$ \\citep[e.g.,][]{weineretal2009} and high-redshift star-forming galaxies as well \\citep[e.g.,][]{shapleyetal2003}.  See \\citet{veilleuxetal2005} for a recent review of low redshift observations.  What has been much harder to pin down is the role of radio-quiet BH accretion in driving winds \\citep[e.g.,][]{tremontietal2007,alexanderetal2010,krugetal2010, coiletal2011,hainlineetal2011,rupkeveilleux2011}.  One class of BHs that may be showing radio-quiet feedback processes in action are broad absorption line quasars, in which high-velocity outflows appear to be driven by the radiation of the quasar itself \\citep[e.g.,][]{progaetal2000}. Because the distance of the outflowing material from the BH and its covering factor are difficult to determine, the kinetic energies of these outflows are unknown in most cases.  In some sources, it may be a substantial fraction of the accretion energy \\citep{crenshawetal2003, moeetal2009}, and the detailed photo-ionization models place these outflows at a considerable distance (several kpc) from the center of the galaxy \\citep[but see also][]{faucher-giguereetal2011}.  These observations are promising, but so far have been possible only for a handful of sources.  We seek unambiguous evidence for galaxy scale winds.  A handful of strong cases have been made for radiatively driven AGN outflows on galaxy-wide scales at high redshift \\citep[e.g.,][]{tremontietal2007,alexanderetal2010,hainlineetal2011}.  Locally, new exciting observations of molecular lines, both with ground-based interferometers \\citep[][]{alataloetal2011} and with Herschel \\citep[e.g.,][]{fischeretal2010} are revealing large molecular outflows that are plausibly AGN-driven.  In particular, \\citet{sturmetal2011} make a strong case that the molecular outflow velocity in local ultra-luminous infrared galaxies (ULIRGs) is strongly correlated with the strength of the active nucleus, such that more luminous AGNs drive faster outflows \\citep[although see also][]{krugetal2010}.  Perhaps the best low-redshift case study of large-scale outflows is of Mrk 231.  Mrk 231 is the most luminous local ULIRG \\citep{sandersetal1988}, the nearest broad emission-line AGN, a rare low-ionization broad absorption-line quasar, and contains a radio jet.  This galaxy also hosts a massive young starburst that likely dominates the bolometric output of the galaxy \\citep[e.g.,][]{lonsdaleetal2003}.  The galaxy has been studied in detail at every conceivable wavelength \\citep[see, e.g., ][and references therein]{feruglioetal2010}.  Most relevant to our work are the recent results pointing to a massive molecular \\citep{fischeretal2010,feruglioetal2010} and neutral \\citep{rupkeveilleux2011} outflow in this galaxy.  It is argued strongly in these papers that radiative driving, rather than star formation or jet activity, power the observed outflows.  In this paper we try to make a similar case for radiation-pressure driving, except that we focus on a radio-quiet obscured quasar, for which there is little evidence of massive star formation in the optical spectrum.  By focusing on obscured, radio-quiet systems, we hope to isolate unambiguously the sources of acceleration in the observed outflows.  We also hope to test the hypothesis that AGN feedback occurs most vigorously in the obscured phase \\citep[e.g.,][]{sandersetal1988,hopkinsetal2006}.  Since 2007, we have been conducting a study of the morphologies and kinematics of narrow-line regions in obscured radio-quiet quasars \\citep{zakamskaetal2003, reyesetal2008}, in search of tell-tale signs of the interactions between the supermassive BH in its radiative (rather than jet-launching) phase and its large-scale gaseous environment. In our program conducted with the Magellan telescopes (\\citealt{greeneetal2009}, hereafter Paper I, and \\citealt{greeneetal2011}, hereafter Paper II), we obtained 15 long-slit spectra of obscured quasars. We find that in the quasar luminosity regime probed by our observations (\\loiii$\\approx 10^{42}-10^{43}$~ erg~s$^{-1}$ or inferred \\lbol$\\approx 10^{45}-10^{46}$~ erg~s$^{-1}$), the quasar easily ionizes the galaxy along any direction not obscured by dust.  Furthermore, the gas is disturbed on galaxy-wide scales, with both broad emission-line widths and kinematics clearly inconsistent with galaxy rotation. Among these 15 sources, one object, SDSS J135646.10+102609.0 (hereafter SDSS J1356+1026), shows very extended ($\\ga 50$ kpc) ionized line emission which is highly kinematically organized on $\\sim 20$ kpc scales. In this paper, we focus on this source, and build a case for AGN-driven feedback using an analysis of the kinematics and geometry of the ionized gas. ",
        "conclusions": "\\bub\\ is an on-going galaxy merger at redshift $z=0.123$ that contains at least one obscured quasar in its northern nucleus. The bolometric luminosity of the object, estimated at $1.5\\times 10^{46}$ erg~s$^{-1}$, places it among the more luminous nearby quasars \\citep[e.g.,][]{hopkinsetal2007}. We have conducted a study of the kinematics and spatial distribution of the ionized gas in this source. We find a $\\sim$ 50 kpc tidal feature close to the systemic velocity of the object that is photo-ionized by the quasar along an unobscured direction. High optical linear polarization is observed, also indicating that the emission-line regions are powered by the AGN. In addition, we observe a pair of expanding shells, or bubbles, centered on the northern nucleus and expanding into the intergalactic medium. The remarkable scale of the expanding structure ($\\sim 40$ kpc in total extent), its high expansion velocities (observed velocity range of 460 \\kms, with physical velocities likely reaching 1000 \\kms), and its large opening angle and luminosity imply that the kinetic energy of the outflow likely exceeds $>10^{44-45}$ erg~s$^{-1}$.  We find that while star formation may be occurring concurrently with quasar activity in \\bub, it is energetically insufficient to power the observed outflow, likely by an order of magnitude, as further supported by the polarization observations. The merging quasars are radio quiet, and there is no evidence for extended radio emission; therefore, the outflow is unlikely to be powered by a relativistic jet in a manner similar to outflows seen in some radio galaxies. We propose that \\bub\\ may be an example of on-going feedback from a quasar in its most common radiative phase. Such feedback has been postulated to account for the colors and luminosities of massive galaxies \\citep[e.g.,][]{springeletal2005}, but it has been difficult to assemble direct evidence of interactions between radio-quiet quasars and their large-scale gaseous environments aside from individual cases noted in the introduction. \\bub\\ provides compelling evidence that accretion energy can drive wide-angle outflows of gas beyond the confines of the host galaxy.  The presence of extended ionized emission by itself is not a signature of quasar feedback \\citep{stocktonmackenty1987}. If one could obtain arbitrarily sensitive observations, one would be bound to find small amounts of gas, and if this gas finds itself along an unobscured direction to the quasar it can be photo-ionized out to large distances. For example, \\citet{villarmartinetal2010} find a 180 kpc emission-line feature near an SDSS type 2 quasar, but its spatial distribution and kinematics make it clear that the feature is tidal debris illuminated by the quasar. Images of this source taken with the \\emph{Hubble Space Telescope} \\citep{zakamskaetal2006} demonstrate that the scattering regions (and therefore illumination cones) have the same E-W orientation as the 180 kpc tidal tail. We find a very similar feature in \\bub\\ extending in the N-S direction, which also happens to be the direction of quasar illumination, as indicated by the polarimetric observations. Rather than swept from the galaxy by quasar feedback to 50-100 kpc away, the gas seen in these two cases is most likely tidal debris  which happened to be along an unobscured line of sight to the quasar.  Similarly, \\citet{liuetal2009} find several obscured quasars with companions out to $\\ga 30$ kpc away from the quasar showing high equivalent width emission lines with high-ionization line ratios. Again, rather than instances of quasar feedback, these occurrences are probably small companion galaxies which happen to fall within one of the quasar illumination cones.  What sort of signatures specifically implicate quasar feedback? The absence of stellar continuum is a necessary (although not a sufficient) requirement, as it helps rule out illuminated companion galaxies and tidal features. Gas velocities that exceed the escape velocity of the galaxy are a clear sign, but since the outflows from type 2 quasars are directed in the plane of the sky, the chances of observing such high velocities in these objects are low.  The presence of a kinematically organized and symmetric structure, such as that seen in \\bub, is a very compelling clue, as it completely rules out an instance of a gas inflow illuminated by the quasar (while the nature of radial velocity observations technically makes inflow and outflow scenarios degenerate, it is hard to imagine the gas conspiring to accrete in such symmetric manner).  The high linear-polarization fraction of $P=8\\%$ provides unambiguous evidence that there is a luminous obscured quasar in \\bub\\ and severely limits the fraction of the bolometric luminosity of this source that may be due to star formation. Furthermore, through the measurement of the polarization position angle, polarimetric observations indicate that the outflow direction (N-S) is well aligned with the predominant direction of quasar illumination.  Clearly a narrow-band image or integral-field spectroscopy would provide a more complete picture of the energetics.  We did not find similarly well-organized kinematics in the rest of the Magellan sample \\citep{greeneetal2011}, but the large velocity widths and the lack of rotation seen even in prominent disk galaxies observed along their semi-major axes make quasar feedback the likely explanation in those cases as well.  Our work to study three-dimensional outflows in luminous obscured quasars is ongoing with integral-field spectroscopy, which we hope will provide better statistics on the population as a whole (Zakamska, N. L. et al. in preparation)."
    },
    "1112/1112.4608_arXiv.txt": {
        "abstract": "Gravitational microlensing events provide a potentially powerful tool for the study of stellar populations.  In particular, they can be used to discover and study a variety of binary systems.  A large number of binary lenses have already been found by the microlensing surveys and a few of these systems show strong evidence of orbital motion on the timescale of the lensing event.  We expect that more of such binary lenses will be detected in the future.  For binaries whose orbital period is comparable to the event duration, the orbital motion can cause the lensing signal to deviate drastically from that of a static binary lens.  The most striking property of such lightcurves is the presence of quasi-periodic features, produced as the source traverses the same regions in the rotating lens plane. These repeated features contain information about the orbital period of the lens.  If this period can be extracted, a lot can be learned about the lensing system even without performing a time-consuming detailed lightcurve modeling.  However, the relative transverse motion between the source and the lens significantly complicates the problem of period extraction.  To resolve this difficulty, we present a modification to the standard Lomb-Scargle periodogram analysis.  We test our method for four representative binary lens systems and demonstrate its efficiency in correctly extracting binary orbital periods. ",
        "introduction": "During the past two decades gravitational microlensing surveys have transformed from a bold theoretical idea \\citep{Pac86} into an important tool for studying the stellar population of our Galaxy. Following Paczy\\'{n}ski's proposal, four different monitoring programs, MACHO \\citep{Alc93}, OGLE \\citep{Uda93}, EROS \\citep{Aub93} and MOA \\citep{Yoc98}, have for years collected data on the Galactic Center as well as LMC and SMC, detecting many thousands of microlensing events. Two microlensing surveys that are still currently in operation, OGLE-IV and MOA, continue to add to this number at a rate of roughly 1000 events per year.  While the majority of the detected events are consistent with being produced by a single-mass lens, a sizable fraction shows clear evidence of lens binarity \\citep{Alc00, Jar04, Jar10}. These are particularly interesting, because more complicated lightcurves \\citep[see e.g.][]{MaoPac} produced by binary lenses allow us to learn considerably more about these systems \\citep[see e.g.][]{MDS95, DSP97}.  The signatures of lens binarity are most prominent when the projected orbital separation between the two binary components, $a$, is of the order of the Einstein radius: \\begin{equation} \\re=\\sqrt{\\frac{4 G M}{c^2}\\frac{(D_s-D_l)D_l}{D_s}}, \\end{equation} where $D_s$ and $D_l$ are the distances to the source and the lens, respectively and $M$ is the total mass of the lens.  The timescale for the duration of the lensing event is determined by the Einstein radius crossing time \\begin{equation} \\te=\\re/v, \\end{equation} where $v$ is the relative velocity between the lens and the source as viewed by the observer.  For solar-mass lenses in the Galactic disk $\\te$ is typically on the order of months, while the orbital period of the lensing binary, $\\Torb$, is more likely to be measured in years.  As a result, most binary lensing events can be modeled while completely neglecting orbital motion of the two masses \\citep[see e.g. recent estimates by][]{PMK11}.  Nevertheless, a handful of microlensing lightcurves do show convincing evidence of binary phase change during the event \\citep{Dom98, Alb00, An02, Jar05, Hwa10, Ryu10}. And recently there has been an increase of interest in the possibility of finding lensing events for which the timescale ratio, defined as \\begin{equation} {\\cal R} = \\te/\\Torb, \\end{equation} is of order of or greater than unity \\citep{PKM11, DSE11}.  If the orbital period can be determined for such systems, it will provide a new avenue toward extracting the lens parameters. Specifically, the period relates the total mass to the orbital separation. If the value of $\\alpha$ and $q$ can also be determined, we have some of the key elements needed to specify the binary. In cases when some additional information is available, e.g., because the lens is located nearby and/or because parallax or finite-source-size effects are detected, a full binary solution may be obtained.  Unlike single-lens lightcurves which have a simple analytical form \\citep{Pac86}, binary lenses can produce very complex magnification patterns which cannot be expressed in a closed form.  Even for a stationary binary lens, the process of fitting an observed lightcurve can be complex \\citep{DSM96}. The addition of orbital motion complicates the fitting process even further.  Fortunately, in the limit when ${\\cal R} \\gtrsim 1$ the binarity features observed in the microlensing lightcurves show very strong periodicity which allow $T_{\\rm orb}$ to be determined via simple Lomb-Scargle (LS) periodogram analysis \\citep{DSE11}.  In this paper, however, we are primarily interested in the regime when ${\\cal R} \\lesssim 1$. For these longer-period binaries the interplay between the orbital motion and the projected motion of the source causes the period determined via LS analysis to differ significantly from $T_{\\rm orb}$ \\citep{DSE11}.  Here we describe a modified timing analysis method which allows us to compensate for this relative motion and extract the correct binary period of the lens. ",
        "conclusions": "When the Einstein crossing time is comparable to the orbital period of a binary lens, microlensing lightcurves display quasi-periodic features which can be used to determine $\\Torb$.  In the limit when ${\\cal R} \\gtrsim 1$, the standard LS periodogram analysis can be used to find the period.  However, in this paper we focus on the regime with ${\\cal R} \\lesssim 1$ when the variability timescale in observed lightcurves is set by a combination of binary rotation and relative source-lens motion.  We have shown that this regime includes specific systems with orbital periods in the range from months to years.  For such lightcurves the standard periodogram yields incorrect periods (see Fig. \\ref{fig:oldnewpowerspectrum}).  Here we proposed a modification to the timing analysis method designed to compensate for the source motion and extract the correct orbital period.  We tested our new timing analysis method on simulated lightcurves for four different binary lens systems and calculated period detection rates for a wide range of impact parameter values.  The results are very encouraging. As long as the orbital period is detectable (i.e. not longer than the length of the simulated lightcurve) our proposed method finds and identifies it correctly in a large fraction of cases.  We find that the period detection rates are determined primarily by the photometric precision and three interdependent parameters, $\\alpha$, ${\\cal R}$ and $b$. Perhaps unsurprisingly, the photometric precision is the most important.  It determines the minimum value of $\\alpha$ for which binarity of the lens is detectable \\citep{DSE11} and sets the maximum value of $b$ for which a microlensing event is distinguishable from the noise. It also essentially sets the maximum practical length for a lightcurve which in turn determines the maximum detectable orbital period and therefore places a lower limit on  ${\\cal R}$.  We consider two levels of photometric uncertainty: 1\\%, characteristic of ground-based data,  and 0.001\\%, relevant for dedicated space missions such as Kepler.  In the first case, the reasonable event duration is $6 \\te$ (beyond this, the wings disappear into the noise) and, correspondingly, the detection rate drops off for ${\\cal R} \\lesssim 1/6$, as illustrated in Fig. \\ref{fig:modlscompare} and well as by our results for the wide binary lens in \\S\\ref{sec:wide}. High-precision photometry allows us to significantly expand the parameter space for which the orbital periods can be reliably detected with our method.  The results we discuss in this paper are based on the analysis of lightcurves produced by binary lenses in face-on circular orbits.  We are now in the process of extending our calculations to include elliptical and inclined orbits.  In these more complex situations, the projected binary separation, and with it the shape of the lensing magnification pattern, will vary with time during the microlensing event.  This effect will certainly affect the spacing of the features in the residual lightcurves, possibly affecting the extraction of the orbital period.  Our preliminary results indicate that the orbital period should still be detectable in a large fraction of lightcurves, but the full report will be the subject of a follow-up paper.  Finally, we want to point out that in its present form, our timing analysis method works best for binary lenses with $q$ not very different from unity.  It is predicated on the assumption that the oscillatory features repeat twice per orbit.  The fact that period detectability does not change much from $q=1.0$ to $q=0.5$ (Fig. \\ref{fig:6sys}[a] and [b]) shows that some asymmetry in the magnification pattern can be easily tolerated.  We tested lower values of $q$ for the unequal-mass system and found that the detectability decreases by a factor $\\sim 2$ for $q=0.25$ and drops below $20\\%$ for $q=0.1$.  It is clear that in the planetary regime (i.e., when $q << 1$) a factor of 2 in Eqs. (\\ref{eq:modCj}) and (\\ref{eq:modSj}) does not apply, since the magnification pattern is no longer even remotely symmetrical \\citep{DS11}.  Fortunately, such systems produce visibly different lightcurve morphologies characterized by very spiky rather than sinusoidal features \\citep[see][for some examples of lightcurves produced by low-$q$ binary lenses]{DSE11}, and so can be flagged. We are now working on ways to extend our method to such low-$q$ systems."
    },
    "1112/1112.4925_arXiv.txt": {
        "abstract": "We report the identification of the double degenerate system NLTT~16249 that comprises a normal, hydrogen-rich (DA) white dwarf and a peculiar, carbon-polluted white dwarf (DQ) showing photospheric traces of nitrogen. We disentangled the observed spectra and constrained the properties of both stellar components. In the evolutionary scenario commonly applied to the sequence of DQ white dwarfs, both carbon and nitrogen would be dredged up from the core. The C/N abundance ratio ($\\approx 50$) in the atmosphere of this unique DQ white dwarf suggests the presence of unprocessed material ($^{14}$N) in the core or in the envelope. Helium burning in the DQ progenitor may have terminated early on the red-giant branch after a mass-ejection event leaving unprocessed material in the core although current mass estimates do not favor the presence of a low-mass helium core. Alternatively, some nitrogen in the envelope may have survived an abridged helium-core burning phase prior to climbing the asymptotic giant-branch. Based on available data, we estimate a relatively short orbital period ($P\\la13$ hrs) and on-going spectroscopic observations will help determine precise orbital parameters. ",
        "introduction": "The presence of a large concentration of carbon, detected via C$_2$ molecular bands, in the otherwise helium-rich atmosphere of many old ($\\gtrsim 10^9$ years) white dwarfs has been viewed for many years \\citep[e.g.,][]{vau1979,koe1982} as evidence of the carbon-rich nature of their cores: The material diffuses upward, away from the core, and is dredged-up to the surface by the deep helium convection zone developing in cool white dwarfs \\citep{fon1984,pel1986,mac1998}. These objects, collectively known as DQ white dwarfs, form a homogeneous sequence of stars with surface temperatures ranging from $\\sim11,000$ to $\\sim5,000$ K and with temperature-correlated surface carbon abundances between $10^{-2}$ and $10^{-7}$ relative to helium by number \\citep{duf2005,koe2006}.  Evolutionary models \\citep[e.g.,][]{sch1992} show that stars of low to intermediate masses, i.e., below 8 to 10 \\msun, end their nuclear burning history with a core composed primarily of carbon and oxygen or heavier elements (Ne, Mg) in the upper mass range \\citep{gar1997}. In the past, spectroscopic searches for oxygen (ultraviolet CO bands) or nitrogen (violet CN bands) have yielded abundance limits O/C$\\approx0.1-1$ increasing with temperature \\citep{koe1982} or O/C,N/C$\\approx 10^{-2}-10^{-3}$ \\citep{weg1984}. Although oxygen is a primary product of the triple-alpha (3-$\\alpha$) nuclear burning process, the actual amount of oxygen dredged-up along with carbon from the interior depends on its abundance profile in the core \\citep{mac1998}. On the other hand, nitrogen should be fully exhausted along with helium by the 3-$\\alpha$ process and is not expected in the atmosphere of white dwarfs with a fully-evolved core. Recently, \\citet{gae2010} reported the discovery of two oxygen-rich white dwarfs alongside the carbon-rich DQ sequence. The measured O/C abundance ratios ($\\log{\\rm O/C}\\approx 0.6-1.8$) imply core compositions dominated by oxygen (80 to 99\\% by mass). The stars were suspected by \\citet{gae2010} to be the descendents of the most massive stars to avoid core collapse \\citep[9-10\\msun;][]{gar1997}.  Clearly, the surface composition reflects the white dwarf core composition which, in turn, is determined by the evolutionary path followed by the progenitor. In this context, we present in Section 2 spectroscopic observations of the peculiar object NLTT~16249 that show a strong Balmer line series along with C$_2$ Swan and CN violet band absorptions. The observed radial velocity variations and a successful spectral decomposition and model atmosphere analysis (Sections 3.1-3.3) show that the object is in fact a close double degenerate system comprising a normal DA white dwarf and a peculiar DQ of similar optical luminosities. We examine possible evolutionary scenarios in Section 3.4, and we propose that the photospheric carbon to nitrogen abundance ratio (C/N$\\approx50$) indicates that the core material was left partially unprocessed or that burning ceased before destroying all nitrogen in the envelope. Finally, we summarize in Section 4. ",
        "conclusions": "We report the discovery of a peculiar DQ white dwarf in a close degenerate system ($P\\la13$ hrs). The presence of nitrogen in the atmosphere of the DQ also suggests the presence of unprocessed material in the core of the star, most probably off-center and below the extinct helium-burning shell. Helium burning may have been interrupted following a common-envelope event leaving a nitrogen-rich layer above a normal carbon/oxygen core. Our spectral decomposition suggests a system mass in the 1.3 to 1.8 \\msun\\ range with cooling ages between 1.3 and 2.6 Gyr depending on the mass. The binary will merge within 9 Gyr for a system mass of 1.3 \\msun, or 13 Gyr for 1.8 \\msun\\ following \\citet{rit1986}. New radial velocity measurements will be used to set a precise binary mass ratio and confirm the evolutionary prospects of the binary NLTT~16249."
    },
    "1112/1112.2438_arXiv.txt": {
        "abstract": "{ The globular clusters are probably good targets for dark matter (DM) searches in $\\gamma$-rays due to the possible adiabatic contraction of DM by baryons. In this work we analyse the three-year data collected by {\\it Fermi} Large Area Telescope of globular clusters NGC 6388 and M 15 to search for possible DM signals. For NGC 6388 the detection of $\\gamma$-ray emission was reported by {\\it Fermi} collaboration, which is consistent with the emission of a population of millisecond pulsars. The spectral shape of NGC 6388 is also shown to be consistent with a DM contribution if assuming the annihilation final state is $b\\bar{b}$. No significant $\\gamma$-ray emission from M 15 is observed. We give the upper limits of DM contribution to $\\gamma$-ray emission in both NGC 6388 and M 15, for annihilation final states $b\\bar{b}$, $W^+W^-$, $\\mu^+\\mu^-$, $\\tau^+\\tau^-$ and monochromatic line. The constraints are stronger than that derived from observation of dwarf galaxies by {\\it Fermi}.} ",
        "introduction": "A standard model of cosmology is developed, in which the universe consists of $4\\%$ ordinary baryonic matter, $\\sim 23\\%$ dark matter (DM), $\\sim 73\\%$ dark energy, and a tiny abundance of relic neutrinos \\cite{2011ApJS..192...18K}. The nature of DM particle remains a mystery. One of the leading candidates is the weakly interacting massive particle (WIMP), which is predicted in several models, such as neutralino in supersymmetry model (see the reviews \\cite{1988ARNPS..38..751P, 1996PhR...267..195J,2005PhR...405..279B}). In this kind of models, the mass and interaction strength of DM particles can produce the correct relic density of DM if the WIMPs are thermally ``freeze-out'', which is called ``WIMP miracle''.  If DM particles annihilate or decay into standard model particles, they can be detected indirectly from the cosmic ray (CR) radiation. Among many kinds of CR particles, $\\gamma$-rays are the best probe due to their simple propagation. {\\it Fermi} gamma-ray telescope, which was launched in 2008, has surveyed the $\\gamma$-ray sky with very high resolution and sensitivity for more than three years. Nearly $2000$ sources as well as the diffuse $\\gamma$-ray emission were detected by {\\it Fermi} Large Area Telescope ({\\it Fermi}-LAT) \\cite{2011arXiv1108.1435T,2009ApJ...703.1249A,2009PhRvL.103y1101A, 2010PhRvL.104j1101A}. The analysis of the {\\it Fermi}-LAT data in the Galactic center region did see some excesses with respect to the background model \\cite{2010ApJ...717..825D,2010ApJ...724.1044S}, however, there is no strong indication of signals from DM annihilation or decay\\footnote{See also the argument of possible DM explanation of the $\\gamma$-ray haze/bubble \\cite{2011ApJ...741...25D}.}. The constraints on DM model parameters can be derived according to the non-detection of DM signals from e.g., dwarf galaxies \\cite{2010ApJ...712..147A,2010PhRvD..82l3503E,2011PhRvL.107x1303G, 2011PhRvL.107x1302A,2011arXiv1111.2604C}, galaxy clusters \\cite{2010JCAP...05..025A,2010PhRvD..82b3506Y,2010JCAP...12..015D, 2011PhLB..698...44K,2011arXiv1105.3240P,2011arXiv1110.1529H} and the diffuse $\\gamma$-rays \\cite{2010JCAP...04..014A,2010NuPhB.840..284C,2010JCAP...03..014P, 2010JCAP...06..027Z,2010JCAP...07..008H,2010JCAP...10..023Y, 2010JCAP...11..041A,2010arXiv1011.5090A,2011arXiv1105.4230C, 2011PhRvD..83l3513Z,2011JCAP...04..020Y}.  Due to the very weak interactions of DM particles, it is important to investigate the sites with high DM density when searching for DM annihilation signals. The proposed good candidates include the Galactic center, dwarf galaxies, Galactic subhalos and cluster of galaxies. The Milky Way globular clusters (GCs), defined as spherical ensemble of stars that orbits the Galaxy as satellites, are also potentially good targets for indirect detection of DM. The formation of GCs remains a poorly understood problem. There are generally two scenarios to describe the formation of GCs. The primordial formation scenario suggests that GCs were formed in cosmological DM minihalos before the formation of galaxies \\cite{1984ApJ...277..470P}. The other way to form GCs might be the star-forming events such as the merger of galaxies. There was evidence to show that metal-poor GCs might have a cosmological origin and metal-rich GCs might form in the galaxies \\cite{2006ARA&A..44..193B}. If the GCs were formed in cosmological DM minihalos, they would experience the adiabatic contraction (AC) due to the infall of baryons during the evolution of GCs and leave a high density spike of DM. GCs are not usually discussed for DM detection due to the poor knowledge about their origin and the observational fact that there is in general no significant amount of DM in vicinity of GCs \\cite{2009MNRAS.396.2051B,2010MNRAS.406.2732L,2010arXiv1010.5783C}. However, there is possibility that the high density spike of DM due to the AC process may still play an important role for the annihilation signals. The previous works to search for or constrain DM models with $\\gamma$-rays from GCs include \\cite{2008PhRvD..78b7301Z,2008ApJ...678..594W}.  Recently the atmospheric Cherenkov telescope array High Energy Stereoscopic System (H.E.S.S.) had investigated two GCs NGC 6388 and M 15 to search for possible DM signals \\cite{2011ApJ...735...12A}. No $\\gamma$-ray signal was detected by H.E.S.S. and strong constraints on the DM model parameters were given. In this work, we use the three-year data of {\\it Fermi}-LAT to study the $\\gamma$-ray emission from DM annihilation in these two GCs. Detections of $\\gamma$-ray emission from some GCs with {\\it Fermi}-LAT were reported \\cite{2009Sci...325..845A,2010A&A...524A..75A, 2010ApJ...712L..36K,2011ApJ...729...90T}, including NGC 6388 studied here. For M 15 there is no detection yet. In this work we will focus on the possible DM component of the $\\gamma$-ray emission, if any, from the GCs. The upper limits of DM contribution will be derived and the constraints on DM model parameters will be presented. ",
        "conclusions": "The GCs are thought to form in the cosmological context with AC process at the beginning which pulls DM into the halo center and results in a very high annihilation luminosity. Thus search for $\\gamma$-rays from GCs may be effective to probe the particle nature of DM. In this work we analyze the {\\it Fermi}-LAT three-year data (Pass 7) of GCs NGC 6638 and M 15 and constrain the DM annihilation models. A clear detection of $\\gamma$-ray emission from NGC 6388 is found, with TS value $\\sim600$. The spectrum of NGC 6388 can be well fitted with a power-law + exponential cutoff function, which is expected for the emission of a population of MSPs. We find that a DM scenario with $m_{\\chi}\\sim25$ GeV and $b\\bar{b}$ final state can also fit the SED. For M 15 no significant $\\gamma$-ray emission if found (the spectral fit indicates a potential source with TS $\\approx12$).  Assuming there is an additional spectral component from DM annihilation of these two GCs, we derive the upper limits of the DM component for different DM masses ($10\\,{\\rm GeV}-10\\,{\\rm TeV}$) and annihilation final states ($b\\bar{b},\\,W^+W^-,\\,\\mu^+\\mu^-,\\,\\tau^+\\tau^-,\\, \\gamma\\gamma,\\,\\gamma Z$) (Figures \\ref{fig:ul} and \\ref{fig:ulline}). The constraints on the DM annihilation cross section are given (Figures \\ref{fig:svbw}-\\ref{fig:svline}). Except for the line emissions, the constraints are stronger than that derived according to the {\\it Fermi}-LAT observations of dwarf galaxies. For DM mass smaller than TeV our constraints are also stronger than that given by H.E.S.S. observations of the same GCs. For $b\\bar{b}$ and $W^+W^-$ final states which are generally expected from supersymmetric DM model, the constraints can reach the natural scale with which DM is thermally produced. Especially the leptonic annihilation models to explain the CR $e^{\\pm}$ excesses can be excluded by the current analysis.  However, the uncertainties of the present analysis, for example the properties of the hypothetical DM halo and the origin and evolution of the GCs, are far from clear. The GCs were assumed to be formed in the cosmological context, and the DM density profiles in the GCs are modelled taking into account the most probable astrophysical processes, e.g., the AC by baryons, adiabatic growth of an IMBH and the scattering by stars. Future studies on the observations and modelings of the DM distribution in the GCs are necessary to improve the current work."
    },
    "1112/1112.0860_arXiv.txt": {
        "abstract": "Hypothetical long-lived massive colored particles (MCPs or $Y$s) would be confined in colorless exotic strongly interacting massive particles (SIMPs) at color confinement temperature of $T_{\\rm C}\\sim 200$~MeV.  Two long-lived MCPs form a bound state $(Y\\bar{Y})$ at collisions of two SIMPs.  We study sensitivities of MCP annihilation to decay properties of resonances $(Y\\bar{Y})$, and binding energies or energy levels of exotic SIMPs.  The $(Y\\bar{Y})$ formation is assumed to dominantly proceed through resonances of $(Y\\bar{Y})$ in this paper.  We make a toy model of the effective cross section for $Y\\bar{Y}$ annihilation.  Abundances of SIMPs are then calculated for different sets of parameters specifying properties of $(Y\\bar{Y})$ resonances, binding energies of SIMPs, the initial abundance and the mass of MCP.  Calculated relic abundances for respective SIMP species are $2\\times 10^{-8}$--$3\\times 10^{-4}$ times that of baryon.  They can be much higher but cannot be much smaller than the previous estimate. The abundances can be consistent depending on parameters with the possible scenario that SIMPs bind to nuclei and subsequent exotic nuclear reactions reduce the primordial abundance of $^7$Li or enhance those of $^9$Be and/or B in the early Universe.  A unique information on the quark-hadron phase transition in the early Universe may become available in future by elaborated studies on the annihilation process with light element abundances as observables. ",
        "introduction": "\\label{sec1}  In some particle models beyond the standard model (SM), long-lived massive colored particles (MCPs) can exist.  The models include a long-lived gluino in split supersymmetry~\\cite{ArkaniHamed:2004fb,ArkaniHamed:2004yi}, the next-to-lightest supersymmetric particle (NLSP) gluino~\\cite{Raby:1997bpa,Shirai:2010rr,Covi:2010au}, and NLSP squark~\\cite{Sarid:1999zx} both in a weak scale supersymmetry.  In addition, extended theories with other kinds of colored particles have been considered~\\cite{Hisano:2010bx,Nakayama:2010vs}.  Such models will be tested at weak scale in the Large Hadron Collider (LHC).  If long-lived MCPs exist at the temperature of color deconfinement, i.e., $T_{\\rm C}\\sim 200$~MeV, they would be confined inside exotic heavy hadrons, so-called strongly interacting massive particles (SIMPs, or $X$ particles)~\\cite{Wolfram:1978gp,Dover:1979sn,Starkman:1990nj}.  Dover, Gaisser \\& Steigman \\cite{Dover:1979sn} have originally estimated the relic abundance of SIMPs, and studied the possibility of exotic stable hadrons with limits from terrestrial superheavy element searches.  They adopted the thermal rate of MCP annihilation in collisions of SIMPs which is assumed to be the same as (or less than) that for nucleons, and $\\sigma v\\sim 30$~mb.  The final abundance thus estimated was about $10^{-11}$ times the baryon abundance.  Nardi \\& Roulet \\cite{Nardi1990} have studied the possibility of new stable exotic quarks predicted by the E$_6$ model.  They concluded that the possibility is ruled out unless there is some mechanism for decay of such quarks.  Also, they pointed out that the annihilation cross section used in Ref. \\cite{Dover:1979sn} is significantly overestimated.  The first reason is that the cross section includes extra final states irrelevant to the MCP annihilation.  The second is the de Broglie wave length of a two-MCP system is the maximum radius scale for the total annihilation cross section.  It is never larger than the interaction radius of $\\sim$1~fm until the temperature decreases down to $\\sim$10 MeV~$(m/{\\rm 1~TeV})^{-1}$ with $m$ the mass ($m\\gg 1$~GeV) of the MCPs.  However, the previous work~\\cite{Dover:1979sn} used a cross section larger than $1~{\\rm fm}^{-2}$ above the temperature.  Thermal relic abundances after the freeze-out of annihilations in the early Universe have been studied with various assumptions on annihilation cross sections~\\cite{Baer:1998pg}.  Theoretical estimates, however, predicted various values ranging over more than several order of magnitude depending on adopted cross sections~\\cite{Baer:1998pg}.  Kang, Luty and Nasri~\\cite{Kang:2006yd} studied aspects of annihilation of SIMPs at $T\\sim T_{\\rm C}$.  They showed that the annihilation proceeds through the following steps: the formation of bound states of two MCPs with large orbital angular momentums, their transitions to more bound lower-energy levels, and the annihilation of MCPs.  They suggested that the bound state formation can realize from multiple partial waves of initial relative angular momenta in the two-SIMP system.  They then suppose a geometrical cross section for annihilation at temperatures $T\\sim T_{\\rm C}$.  The final abundance of SIMP has been estimated~\\cite{Kang:2006yd} to be \\begin{equation} \\frac{N_X}{s} \\sim \\sqrt[]{\\mathstrut \\frac{15}{\\pi}} \\frac{g_\\ast^{1/2}}{g_{\\ast s}} \\frac{m^{1/2}}{\\sigma T_B^{3/2} m_{\\rm Pl}}~~, \\label{eq1} \\end{equation} where $N_X$ is the number density of SIMP, $s=2\\pi^2 g_{\\ast s} T^3/45$ is the entropy density with $g_{\\ast s}\\sim 10$ the total number of effective massless degrees of freedom in terms of entropy~\\cite{kolb1990} below the phase transition of quantum chromodynamics (QCD), $g_\\ast$ is the total number of effective massless degrees of freedom in number~\\cite{kolb1990}, $\\sigma$ is the annihilation cross section of the SIMP, $T_B$ is the temperature of the Universe at which the colorless SIMPs are formed, and $m_{\\rm Pl}$ is the Planck mass. The ratio of number abundance of SIMP  to that of baryons is then ${N_X}/{n_b}\\sim 0.5\\times 10^{-8}$, where $n_b$ is the number density of baryons.  Diaz-Cruz et al.~\\cite{DiazCruz:2007fc} have adopted models in which the lighter stop $\\tilde{t_1}$ is the NLSP and the gravitino is the LSP.  They studied cosmological behaviors of the stop, and find that stops are finally hadronized into mesinos or sbaryons.  They used data on energy levels of normal hadron and the estimation of annihilation rate by Kang et al. \\cite{Kang:2006yd}, and showed an abundance evolution of exotic hadrons.  Jacoby \\& Nussinov~\\cite{Jacoby:2007nw,Nussinov:2009hc} have refined the estimate of relic abundance in Ref.~\\cite{Kang:2006yd}.  They assumed that MCPs are quarklike, and considered a reaction between SIMPs and normal hadrons, i.e., mesons and nucleons, which transform constituents or energy levels of exotic hadrons $X$.  They suggested that the dissociation of bound states of two MCPs must be considered.  They roughly took all those effects into account, and concluded that the relic abundance after the freeze-out of the annihilation at $T\\lesssim T_{\\rm C}$ is as small as suggested in Ref.~\\cite{Kang:2006yd}.  Recently the ''Quirk Model'' has been suggested by M. Luty~\\cite{Jacoby:2007nw}.  In the model, colored particles transform under an additional non-Abelian gauge group.  The phenomenology relating to a new exotic quanta of the corresponding gauge field originally named theton had been studied about three decades ago~\\cite{Okun:1980kw,Okun1980b}.  They have discussed results for the situation in which the mass of the lightest quark was heavier than the QCD parameter $\\Lambda_{\\rm had}\\sim 1$~GeV~\\cite{Gupta:1981ve}. Taking account of the recent suggestions~\\cite{Kang:2006yd,Jacoby:2007nw}, an estimate of relic abundances has derived, and collider phenomenology of such particles or the possibility of detection at LHC or Tevatron has been studied~\\cite{Jacoby:2007nw,Kang:2008ea}.  Strassler and Zurek~\\cite{Strassler:2006im} have also studied the phenomenology of a new confining gauge group added to SM as a kind of hidden-valley models.  In this paper, we construct a toy model which roughly includes effects of all possible processes operating in the annihilation at $T\\lesssim T_{\\rm C}$.  Abundances of SIMPs during the annihilation epoch are calculated for four cases of different properties of resonant or barely-bound states of two-MCP systems and binding energies of SIMPs.  We show results for each case with different initial abundances and masses of MCPs.  We then suggest that final abundances can be much higher but cannot be much smaller than the naive estimations of Refs. \\cite{Kang:2006yd,Jacoby:2007nw}.  In Sec.~\\ref{sec2}, we describe input physics and assumptions adopted in this paper.  We show a hypothetical quantum mechanical picture of annihilation of SIMPs (Sec.~\\ref{sec2}A), a toy model of the effective cross section for the annihilation (Sec.~\\ref{sec2}B), a comparison of the formation rate of bound states involving two MCPs with the cosmic expansion rate (Sec.~\\ref{sec2}C), the transformation of mesonlike SIMPs to nucleonlike SIMPs (Sec.~\\ref{sec2}D), and the annihilation cross section at low temperature after the transformation (Sec.~\\ref{sec2}E).  In Sec.~\\ref{sec3}, we describe models of different decay properties of resonances of two MCPs, and different binding energies of SIMPs.  Effective annihilation cross sections for the models are derived.  In Sec.~\\ref{sec4}, we show result of calculations of SIMP abundance and effects of decay properties of resonances and binding energies.  In Sec.~\\ref{sec5}, we briefly mention the case in which heavy hadrons including two MCPs exist stably.  In Sec.~\\ref{sec6}, this work is summarized, and a possible impact of results about relic abundance of SIMPs on big bang nucleosynthesis is discussed.  In this paper, the following units are adopted: the Planck's constant, the Boltzmann's constant, and the light speed are unity, i.e., $\\hbar=1$, $k_{\\rm B}=1$ and $c=1$. ",
        "conclusions": "\\label{sec6}  We have investigated effects of decay properties of resonances made of two long-lived massive colored particles (MCPs), i.e., $(Y\\bar{Y})$, and binding energies or energy levels of exotic strongly interacting massive particles (SIMPs) including one MCP in them, on the annihilation of MCPs at color confinement temperature $T_{\\rm C}\\sim 200$~MeV.  Magnitudes of cross sections of two-SIMP collision for final states including a $(Y\\bar{Y})$ is expressed by a parameter $f_{\\rm w}$.  In addition, binding energies of SIMPs are taken as other parameters.  We studied effects of the parameters on the $Y\\bar{Y}$ annihilation, and resulting relic abundances of SIMPs by calculating with different parameter sets.  It is assumed that the $(Y\\bar{Y})$ bound state formation from initial states of two SIMPs dominantly proceeds through resonances (or barely-bound states) of $(Y\\bar{Y})$.  The cross section for bound state formation is then formulated utilizing a well-known expression for resonant cross section.  We suppose that bound states or resonances of MCPs decay into more bound states mainly through the dipole photon emission rather than hadron emissions, based upon an ansatz from particle data on heavy quarkonia \\cite{PDG2010}.  In addition, we roughly estimate fates of $(Y\\bar{Y})$ states from a comparison of the rate for cascading down energy levels by the spontaneous emission and collisions with background photons to that for being destroyed by collisions.  We then make a toy model of the effective cross section for $Y\\bar{Y}$ annihilation by two-SIMP collisions which accounts for only the reaction which ends with the $Y\\bar{Y}$ annihilation without dissociated into two SIMPs by background photons.  In this model, we roughly take into account the transition of SIMP, i.e, $Y\\bar{q}\\rightarrow Yqq$, where $q$ is a light quark.  Evolutions of SIMP abundances are calculated as a function of the temperature for four cases.  Different parameter sets for three initial abundances and three masses of MCP are considered in the calculation.  In Figs.~\\ref{fig:case1}--\\ref{fig:case4} we show results of abundance evolution at temperature $T\\lesssim T_{\\rm C}$ for some cases.  Final abundances for all cases are listed in Table \\ref{tab6}.   We show that larger $f_{\\rm w}$ values leads to smaller final abundances of SIMP, while smaller mass gains of SIMPs lead to larger final abundances.  We then suggest that informations on exotic heavy hadrons such as their energy levels and partial decay widths are necessary to estimate relic abundances of SIMPs precisely.  The calculated relic abundances (Table \\ref{tab6}) for respective SIMP species are $2\\times 10^{-8}$--$3\\times 10^{-4}$ times that of baryon depending on parameters.  This large uncertainty stems from uncertainties in level structures and decay widths of SIMPs, and the mass and the initial abundance at color confinement of MCP.  We finally comment on possible scenarios in a case that long-lived heavy hadrons composed of two $Y$ particles exist.  We note that the relatively small final abundances of SIMPs with a huge uncertainty can be important since even such small amounts of SIMPs possibly leave observational signatures on primordial light element abundances through bindings of SIMPs to nuclei and nuclear reactions of the bound states \\cite{Kusakabe:2009jt,Kawasaki:2010yh}.  If a long-lived SIMP interacts with nucleon by the strength of the potential between two nucleons, the binding of the SIMP to a nucleon triggers formations of heavy nuclei up to carbon which are bound to the SIMP.  The decay lifetime of the SIMP, therefore, should be $\\tau_X \\lesssim 200$~s to avoid heavy element synthesis conflicting with observations of old stars \\cite{Kusakabe:2009jt}.  If the potential between a long-lived SIMP, i.e., $X$, and a nucleon is suitably weaker than that between two nucleons, primordial abundances of $^7$Be and $^7$Li can reduce by the reactions $X+^7$Be$\\rightarrow ^3$He$+^4$He$_X$ and $X + ^7$Li$\\rightarrow t+^4$He$_X$, respectively \\cite{Kawasaki:2010yh}.  The discrepancy in $^7$Li abundance between predictions in standard big bang nucleosynthesis model and observations of metal-poor stars~\\cite{Spite:1982dd,Ryan:2000zz,Melendez:2004ni,Asplund:2005yt,bon2007,Shi:2006zz,Aoki:2009ce} can be caused through the exotic nuclear reactions enabled by the existence of long-lived relic SIMPs.  The abundance of SIMPs necessary for a reduction of $^7$Li abundance down to the observed level in an example case (figures 5 and 7 of Ref.\\cite{Kawasaki:2010yh}) was $\\sim 1.7\\times 10^{-4}$ times as large as the baryon abundance.  That relatively high SIMP abundance is predicted especially in Case 3 in this study.  The present study indicates that a high relic abundance tends to result from small branching ratios of $(Y\\bar{Y})$ states for decays to lower $(Y\\bar{Y})$ bound states, a large initial abundance of MCP, and large MCP mass. Although we did not treat a case that the colored particle $Y$ is electrically neutral, its relic abundance may be larger than in charged cases as we have mentioned at the end of the Sec.~\\ref{sec3}.  The neutral $Y$ may, therefore, also be a viable candidate for SIMP needed for a $^7$Li reduction.  There are interesting possible predictions to be tested in future astronomical observations adding to that on $^7$Li \\cite{Kusakabe:2009jt}.  One is $^9$Be and/or B abundances larger than the prediction of standard big bang nucleosynthesis.  Another is a high isotopic ratio $^{10}$B/$^{11}$B different from predictions of other models for boron production such as the cosmic ray nucleosynthesis ($^{10}$B/$^{11}$B$\\sim 0.4$~\\cite{Prantzos1993,Ramaty1997,Kusakabe2008}) and the supernova neutrino process ($^{10}$B/$^{11}$B$\\ll 1$~\\cite{Woosley:1995ip,Yoshida:2005uy}).  If a long-lived SIMP affected the nucleosynthesis, primordial light element abundances would have a unique information on the quark-hadron phase transition in the early Universe.  The relic abundance of SIMP significantly depends on hadronic characteristics through annihilation rates as shown in this study.  In view of this, more elaborated studies on the annihilation process are necessary to derive an information on the phase transition in future."
    },
    "1112/1112.2865_arXiv.txt": {
        "abstract": "Since its early phases of operation, the AGILE  mission is successfully observing Gamma Ray Bursts (GRBs) in the hard X-ray band with the SuperAGILE imager and in the MeV range with the Mini-Calorimeter. Up to now, three firm GRB detections were obtained above 25 MeV and some bursts were detected with lower statistical confidence in the same energy band. When a GRB is localized, either by SuperAGILE or Swift/BAT or INTEGRAL/IBIS or Fermi/GBM or IPN, inside the field of view of the Gamma Ray Imager of AGILE, a detection is searched for in the gamma ray band or an upper limit is provided. A promising result of AGILE is the detection of very short gamma ray transients, a few ms in duration and possibly identified with Terrestrial Gamma-ray Flashes. In this paper we show the current status of the observation of Gamma Ray Bursts and Terrestrial Gamma-ray Flashes with AGILE. ",
        "introduction": "The italian small satellite mission AGILE (see \\cite{Tavani_et_al_2008} for further information) has been designed, developed and produced having in mind the Gamma Ray Bursts (GRBs) as one of the most important scientific objectives. For this reason the X-ray monitor SuperAGILE (see \\cite{Feroci_et_al_2007} for a complete description of the instrument) and the non-imaging Mini-Calorimeter (described by \\cite{Labanti_et_al_2009}) are equipped with on-board trigger algorithms and dedicated telemetry data packets (see \\cite{Del_Monte_et_al_2007} and \\cite{Fuschino_et_al_2008} for a complete description). Similar methods are used in the ground based analysis in order to detect GRBs.  At the time of writing, that is in two years of operations including the Commissioning and Science Verification Phase (until November 2007), the whole Cycle One (December 2007 -- November 2008) and more than one half of the Cycle Two (December 2008 -- November 2009), SuperAGILE localized 28 GRBs, about one per month. SuperAGILE is based on 1-D imaging and its field of view is composed of a central region, of $68 \\degree \\times 68 \\degree$ in which both coordinates are encoded with an error radius of 3', and four wings, where only 1-D images are produced as narrow strips of $6' \\times 68 \\degree$. Out of 28 GRBs, 17 have an error radius of 3' while for the remaining 11 the position is given by the 1-D strip described above. In the same interval of time the Mini-Calorimeter (MCAL), a non imaging and all-sky scintillator detector, detected about one GRB per week in the energy band between 300 keV and 100 MeV (see \\cite{Marisaldi_et_al_2008} for the description of the MCAL GRB capabilities). During the same time span the AGILE Gamma Ray Imaging Detector (sensitive in the energy band between 30 MeV and 30 GeV and described by \\cite{Tavani_et_al_2008}) found three firm detections (GRB 080514B, GRB 090401B and GRB 090510) above 30 MeV and two less significant ones (GRB 080721 and GRB 081001) in the same energy band. ",
        "conclusions": "The observation of GRBs by Fermi and AGILE missions is showing that the emission of gamma rays from this class of sources is rather uncommon, with a detection rate of about one event per month. The continuing observation will allow to increase the statistic of the data sample and may help to find correlations between the detection of gamma rays and other features of the GRBs. So far the gamma rays are detected mainly during the prompt phase of the GRBs and usually belong to an extended emission, lasting longer than the prompt emission in hard X-rays. The fraction of gamma rays in the extended emission may vary, depending on the position of the peaks in the lightcurve. Another important property, reported in the case of GRB 080514B, is that the fluence in gamma rays is found on the extrapolation of the spectrum of the prompt emission detected up to the MeV region. A similar feature has been found in the spectrum of GRB 080916C, detected by Fermi/LAT (\\cite{Abdo_et_al_2009}). The analysis of the energetics of other gamma ray emitting bursts is still in progress. Up to now no peculiar feature is found in the X-ray afterglow of gamma ray emitting GRBs but the sample is still small.    After the onset of the timescale shorter than 64 ms on the on-board trigger of the AGILE MCAL, the instrument is detecting a population of short and intense bursts with properties compatible with those of TGFs. Both the geographical and local time distributions as well as the cumulative spectrum of the AGILE TGF sample well match the same distributions for the RHESSI sample, confirming that AGILE is actually detecting TGFs and the goodness of the selection criteria applied. Improvements in the selection strategy and trigger logic are expected to increase the TGF detection rate. Thanks to the almost equatorial orbit, AGILE can provide a continuous monitoring of the equatorial region, especially concerning central Africa and South East Asia, where some of the most severe TGF-producing thunderstorms develop."
    },
    "1112/1112.4444_arXiv.txt": {
        "abstract": "{ Evolution in the mass function of galaxy clusters sensitively traces both the expansion history of the Universe and cosmological structure formation. Robust cluster mass determinations are a key ingredient for a reliable measurement of this evolution, especially at high redshift. Weak gravitational lensing is a promising tool for, on average, unbiased mass estimates. }{ This weak lensing project aims at measuring reliable weak lensing masses for a complete X-ray selected sample of $36$ high redshift ($0.35\\!<\\!z\\!<\\!0.9$) clusters. The goal of this paper is to demonstrate the robustness of the methodology against commonly encountered problems, including pure instrumental effects, the presence of bright ($8$--$9$ mag) stars close to the cluster centre, ground based measurements of high-$z$ ($z\\!\\sim\\!0.8$) clusters, and the presence of massive unrelated structures along the line-sight. }{ We select a subsample of seven clusters observed with MMT/\\textsc{Megacam}. Instrumental effects are checked in detail by cross-comparison with an archival CFHT/\\textsc{MegaCam} observation. We derive mass estimates for seven clusters by modelling the tangential shear with an NFW profile, in two cases with multiple components to account for projected structures in the line-of-sight. }{ We firmly detect lensing signals from all seven clusters at more than $3.5\\sigma$ and determine their masses, ranging from $10^{14}\\,M_{\\odot}$ to $10^{15}\\,M_{\\odot}$, despite the presence of nearby bright stars. We retrieve the lensing signal of more than one cluster in the CL~1701+6414 field, while apparently observing CL~1701+6414 through a massive foreground filament. We also find a multi-peaked shear signal in CL~1641+4001. Shear structures measured in the MMT and CFHT images of CL~1701+6414 are highly correlated. }{ We confirm the capability of MMT/\\textsc{Megacam} to infer weak lensing masses from high-$z$ clusters, demonstrated by the high level of consistency between MMT and CFHT results for CL\\,1701+6414. This shows that, when a sophisticated analysis is applied, instrumental effects are well under control. } ",
        "introduction": "Cosmological observables probing different physics are found to agree within their uncertainties on a $\\Lambda$CDM cosmological model dominated by Dark Energy and Dark Matter \\citep[e.g.,][]{2008ApJ...686..749K,2010A&A...516A..63S,2011ApJS..192...16L}. Investigating the unknown physical nature of Dark Energy ranks among the foremost questions for cosmologists. In particular, the presence or absence of evolution in Dark Energy density is expressed by the equation-of-state parameter $w_{\\mathrm{DE}}$. State-of-the-art measurements \\citep[e.g.,][]{2009ApJ...692.1060V,2010MNRAS.406.1759M,2011ApJS..192...18K} are consistent with $w_{\\mathrm{DE}}\\!=\\!-1$, and hence with Dark Energy being an Einsteinian cosmological constant \\citep[e.g.][]{2010A&ARv..18..595B}. The tightest constraints on $w_{\\mathrm{DE}}$ can be achieved only by combining results from different probes showing complementary dependencies on cosmological parameters.  The evolution of galaxy clusters is understood to be determined by cosmological parameters through both cosmic expansion and hierarchical structure formation. Thus, clusters provide information to constrain cosmological parameters that is complementary to other tests (using e.g., the cosmic microwave background, type Ia supernovae, or baryonic acoustic oscillations). Information about the expansion history of the Universe and structure formation is encoded in the \\emph{cluster mass function} $n(M,z)$; different cosmological models imply a different cluster mass function at high $z$, compared to the local mass function \\citep[e.g.][]{1996MNRAS.282..263E,2002ARA&A..40..539R,2005RvMP...77..207V,2005RvMA...18...76S,2006A&A...453L..39R}. Refining the original analytic model of \\citet{1974ApJ...187..425P}, various fitting formulae have been developed based on numerical simulations \\citep[see][for an overview]{2010MNRAS.402..191P}. The measured cluster mass function not only gives strong evidence for the existence of Dark Energy \\citep[e.g.][]{2009ApJ...692.1060V} but also adds valuable information to the joint constraints of cosmological parameters \\citep[for a recent review, see][]{2011arXiv1103.4829A}.  Galaxy cluster cosmology relies on the accurate determination of cluster masses and a thorough understanding of the different mass proxies in use. Cluster masses are most commonly inferred from a variety of X-ray observables (X-ray luminosity $L_{\\mathrm{X}}$, gas mass $M_{\\mathrm{gas}}$, the quantity $Y_{\\mathrm{X}}\\!=\\!T_{\\mathrm{X}}M_{\\mathrm{gas}}$) or gravitational lensing (both weak and strong), but as well using the Sunyaev-Zel'dovich effect or the motions of member galaxies. For distant galaxy clusters which have low masses owing to the early epoch of structure formation they represent, the small number of available photons prohibits a detailed (spectral) X-ray analyses available for local clusters. Nevertheless, weighing a large number of high-redshift clusters will yield the best constraints on cosmological parameters. This strategy will be adopted by the upcoming generation of cluster surveys, e.g. using \\textsl{eROSITA} \\citep{2010SPIE.7732E..23P,2011arXiv1111.6587P}. Hence, scaling relations connecting quantities like $L_{\\mathrm{X}}$ or $M_{\\mathrm{gas}}$ with the total cluster mass will continue to play important roles and need to be understood and calibrated thoroughly.  Rooted in the thermodynamics of the intracluster medium (ICM), X-ray methods rely on assumptions of hydrostatic equilibrium, elemental composition, and, to a large extent, sphericity of the cluster's gravitational potential well \\citep{1988xrec.book.....S,2010A&ARv..18..127B}. Weak gravitational lensing \\citep[WL; e.g.][]{2006glsw.book.....S} offers an alternative avenue for determining cluster masses, which is completely independent of these assumptions, directly mapping the projected mass distribution of matter, Dark and luminous.  Merging clusters deviate strongly from thermal and hydrostatic equilibrium, with a significant amount of the internal energy being present as kinetic energy of bulk motions or turbulent processes, e.g.\\ merger shocks. Thus, the expectation from numerical simulations is that X-ray masses for merging systems might be biased after a significant merger, with a relaxation timescale of $\\mathcal{O}(1\\,\\mathrm{Gyr})$. Simulations agree with the expectation from the hierarchical structure formation picture that mergers are more frequent at higher redshifts than in the local Universe \\citep[e.g.][]{2005APh....24..316C}. There is no consensus between different simulations yet which of several suggested physical effects dominates after the phase in which a disturbed morphology can be seen \\citep[e.g.,][]{2006ApJ...650..128K,2007ApJ...668....1N,2010ApJ...715.1508S}. Most importantly, bulk motions induce non-thermal pressure, providing support for the gas against gravity, thus possibly leading the hydrostatic mass to underestimate the true mass by $5$--$20$\\% even in \\emph{relaxed} clusters \\citep{2006MNRAS.369.2013R,2007ApJ...655...98N,2010A&A...514A..93M}.  Therefore, studying scaling relations of X-ray observables with weak lensing masses has become an important ingredient in refining cluster masses from X-ray observations \\citep[e.g.,][]{2008A&A...482..451Z,2010ApJ...711.1033Z,2010A&A...514A..93M}. Relative uncertainties of the individual WL cluster masses are higher than those from X-rays, largely due to intrinsic shape noise. But the power of weak lensing comes through the statistical analysis of $M_{\\mathrm{wl}}/M_{\\mathrm{X}}$ for the whole sample, under the assumption that WL mass estimates are, on average, unbiased. This means, while WL mass estimates for \\emph{individual clusters} are subject to an error due to the projection of filaments or voids along the line-of-sight, the stochastic nature of these errors makes them cancel out when averaging over a well-defined cluster sample. Statistical comparisons to X-ray masses \\citep[e.g.,][]{2010A&A...514A..93M} help us to investigate WL systematic uncertainties, i.e.\\ triaxiality \\citep{2009MNRAS.396..315C} and projection of unrelated LSS \\citep{2003MNRAS.339.1155H}, to which X-ray observables are far less sensitive.  This article presents the second part of a series on weak lensing analyses following up the 400 Square Degree Galaxy Cluster (\\emph{400d}) Survey, initiated in \\citet[hereafter Paper I]{2010A&A...520A..58I}. The \\emph{400d} Survey presents a flux-limited sample of galaxy clusters detected serendipitously in a re-analysis of all suitable \\textsl{ROSAT} PSPC pointings \\citep{2007ApJS..172..561B}. From the resulting catalogue, \\citet[V09]{2009ApJ...692.1033V} drew the \\emph{cosmological subsample} of $36$ X-ray luminous and distant clusters, for which high-quality \\textsc{Chandra} X-ray observations were obtained and analysed. The \\textsc{Chandra}-based cluster mass function resulting from the \\emph{Chandra Cluster Cosmology Project} was published by V09, for the complete redshift range of $0.35\\!\\leq\\!z\\!<\\!0.90$ spanned by the clusters in the \\emph{cosmological subsample}, as well as divided into three redshift bins. Building on this mass function, \\citet{2009ApJ...692.1060V} constrained cosmological parameters, in particular $w_{\\mathrm{DE}}$.  Determining accurate weak lensing masses for the distant clusters in the \\emph{400d} cosmological subsample opens the way to observationally test the assumptions \\citet{2009ApJ...692.1033V,2009ApJ...692.1060V} make for the scaling relations and their evolution. Put briefly, the WL follow-up of the \\survey ~cosmological sample clusters provides us with a control experiment for the resulting X-ray mass function. With $36$ clusters, the \\emph{400d} WL sample ranks among the largest \\emph{complete} high-$z$ WL samples. %  In Paper~I, we presented the results of our feasibility study, performing a detailed lensing and multi-method analysis of CL\\,0030+2618. In particular, we showed the \\megacam ~instrument to be well suited for WL studies. As the next step of the project, we investigate seven further clusters from our sample, all of which were also observed with \\megacam ~at MMT. The resulting WL mass determination and the status after $8$ out of $36$ clusters have been analysed are the subjects of this paper.  We consistently assume a $\\Lambda$CDM cosmology specified by the dimensionless Hubble parameter $h\\!=\\!0.72$ and matter and Dark Energy density parameters of $\\Omega_{\\mathrm{m}}\\!=\\!0.30$ and $\\Omega_{\\Lambda}\\!=\\!0.70$.  This paper is organised as follows: After giving a short overview on our data set and its reduction in Sect.~\\ref{sec:sec2}, we give salient details of the WL analysis methods we used in Sect.~\\ref{sec:sec3}. In Sect.~\\ref{sec:sec4}, we take a closer look at two clusters which show a more complicated shear morphology and devise a two-cluster shear model. Comparing our MMT results to a CFHT weak lensing analysis of one of our clusters, we once more prove MMT weak lensing to be reliable and provide an external calibration (Sect.~\\ref{sec:cfhtsec}). In Sect.~\\ref{sec:accuracy}, we provide details of the error analysis for our main results, the cluster masses, which are then discussed in Sect.~\\ref{sec:disc}. Section~\\ref{sec:summary} presents our summary and conclusion. ",
        "conclusions": "\\label{sec:summary}  In this study, the second in the series on the \\emph{400d} survey WL follow-up, we reduced and analysed MMT/\\textsc{Megacam} observations for seven clusters of galaxies. Building on Paper~I, data reduction is performed using \\texttt{THELI}, and WL shear catalogues are extracted using an implementation of the KSB+ algorithm.  In the three cases, where we have MMT observations in $g'r'i'$, we define lensing catalogues based on a refined version of the three-colour method used in Paper~I. By comparing with the colours observed for a \\citet{2006A&A...457..841I} photo-$z$ field, we exclude sources from regions in colour-colour-magnitude space containing a large fraction of foreground galaxies from the analysis. For clusters with only one MMT band, we apply a cut in magnitude as background selection.  We detect all of our $0.39\\!<\\!z\\!<\\!0.80$ clusters using the aperture mass method ($S$-statistics) at the $>\\!3.5\\sigma$ level. Performing a \\citet{2001A&A...374..740S} mass reconstruction, we find the projected mass to follow the $S$-statistics closely. The WL masses of our clusters are determined from NFW modelling of their tangential shear profiles, yielding masses in the $10^{14}\\,\\mbox{M}_{\\sun}\\!\\leq\\!M_{200}^{\\mathrm{wl}}\\!<\\!2\\!\\times\\!\\,10^{15}\\,\\mbox{M}_{\\sun}$ interval.  Two of our clusters are exceptional due to their complicated shear morphology: For CL\\,1701+6414, where several known clusters lie close to one another in projection, we simultaneously fit the shear of the two strongest $S$-peaks, identified with A~2246 and our target. The field of CL\\,1641+4001 also exhibits multiple shear peaks, but we find no evidence for the presence of more than one cluster.  An independent analysis of the CL\\,1701+6414 field using archival CFHT data confirms the superposition of several weak lensing sources. By matching shear catalogues from MMT and CFHT, we find ellipticities measured with both instruments to be consistent with the assumption of noisy, but unbiased measurements of the same quantity. Hence, MMT/\\textsc{Megacam} is proven to be equally good for WL science as the well-established CFHT/\\textsc{Megacam}. We further produced a photo-$z$ catalogue based on CFHT $g'r'i'z'$ data of the field. Despite the shallowness of three bands, we are able to devise a coarse foreground/background selection for CL\\,1701+6414. This experiment again confirms not only shear peaks for several known clusters but also the validity of the magnitude-cut selection. Better photo-$z$ data will be needed to potentially turn this cross-check into a calibration for single-band lensing data.  We find the error budgets for our cluster masses to be dominated by statistical uncertainties (which can be suppressed by using a large cluster sample), but with a significant contribution of systematic uncertainties. Statistical uncertainties are naturally higher for relatively high-$z$ clusters like ours, but the role of data quality (weather and instrumental conditions that led to drastic reductions in data depth) can hardly be overestimated.  As the second paper in the \\emph{400d} WL series, the main results of this study are that: \\begin{itemize} \\item Instrumental effects are well under control. \\item Reliable masses can be obtained in the presence of bright stars close to the cluster centre -- an important finding for a successful follow-up of a \\emph{complete} sample! \\item Ground-based WL works at least till $z\\!\\approx\\!0.8$. \\item We can correct for extreme cases of massive foreground structures. \\end{itemize} Furthermore, we identify areas of possible future improvements of the methods we applied: As observational constraints will forbid complete homogeneity of the data analysis to some degree for each survey. Therefore, better methods to calibrate WL analyses in particular with a different number of available filters need to be developed. This applies specifically to background selection and correction for cluster members."
    },
    "1112/1112.1595.txt": {
        "abstract": "The discovery of ``hot Jupiters'' very close to their parent stars confirmed that Jovian planets migrate inward via several potential mechanisms.  We present empirical constraints on planet migration halting mechanisms.  We compute model density functions of close-in exoplanets in the orbital semi-major axis -- stellar mass plane to represent planet migration that is halted via several mechanisms, including the interior 1:2 resonance with the magnetospheric disk truncation radius, the interior 1:2 resonance with the dust sublimation radius, and several scenarios for tidal halting.  The models differ in the predicted power law dependence of the exoplanet orbital semi-major axis as a function stellar mass, and thus we also include a power law model with the exponent as a free parameter.  We use a Bayesian analysis to assess the model success in reproducing empirical distributions of confirmed exoplanets and Kepler candidates that orbit interior to 0.1 AU.   Our results confirm a correlation of the halting distance with stellar mass.  Tidal halting provides the best fit to the empirical distribution of confirmed Jovian exoplanets at a statistically robust level, consistent with the Kozai mechanism and the spin-orbit misalignment of a substantial fraction of hot Jupiters.  We can rule out migration halting at the interior 1:2 resonances with the magnetospheric disk truncation radius and the interior 1:2 resonance with the dust disk sublimation radius, a uniform random distribution, and a distribution with no dependence on stellar mass.   Note, our results do not rule out Type II migration, but rather eliminate the role of a circumstellar disk in stopping exoplanet migration.  For Kepler candidates, which have a more restricted range in stellar mass compared to confirmed planets, we are unable to discern between the tidal dissipation and magnetospheric disk truncation braking mechanisms at a statistically significant level.  The power law model favors exponents in the range of 0.38--0.9.  This is larger than that predicted for tidal halting (0.23--0.33), which suggests that additional physics may be missing in the tidal halting theory. ",
        "introduction": "Many ``hot Jupiter'' planets have been discovered to orbit very close to their central stars \\citep[e.g.,][]{planetref1,planetref2,keplerref}.   It is well-established that these planets must form further out from their host stars, likely beyond the snow line, and either migrate embedded in a primordial disk or via dynamical interactions \\citep[e.g., ][]{lin96,lubow10,kozai,spinorbitref}. Observations of the Rossiter-McLaughlin effect to identify stellar spin -- planet orbit misalignment show that a significant fraction of ``hot Jupiters'' are aligned, and a significant fraction are also misaligned \\citep[]{morton,spinorbitref,spinorbitref2}.  The mis-aligned planets are likely directed inwards via planet-planet scattering, the Kozai mechanism, secular chaos, or analogous mechanisms \\citep[e.g.,][]{kozai,wukozai,scatterref2,scatterref3,secularchaos}.  Recent work by \\citet[]{dawson1,dawson2} suggests that less than 15\\% of hot Jupiters undergo migration via the Kozai mechanism, instead favoring planet-planet scattering.  For aligned close-in planets, planet migration embedded in a primordial disk is suspected to explain the observed planet location.  Most disk migration models involve similar physical processes, but contain differences in the underlying assumptions about the structure and properties of the primordial disk and the planet \\citep[e.g., viscosity, density, scale height, temperature, dissipation time-scale, toroidal magnetic fields embedded in the disk, etc., ][]{lin96,lubow10,migrateref3,migrateref4,migrateref5,toroidalref}.  Current hypotheses often combine a few models together in an attempt to explain observed exoplanet mass and semi-major axis distributions.  The justification cited is that conditions within the disk and the planet mass and density change over time, resulting in different models being applicable at different times in the planet migration process.  Type I migration assumes that the density structure of the disk is not affected by planets.  Instead, turbulence determines the density structure of the disk.  For this reason, this form of migration is most applicable to small mass planets.  In the case of type II migration, a gap is formed between the disk and a high mass planet.  This gap is the result of the tidal torques from the planet becoming stronger than the viscous torques of the disk.  Initial theoretical models of Type I and Type II migration suggested a rapid planet migration rate that could result in planet destruction by dispersing the accreted material \\citep[]{ward,ward2}.  Various mechanisms are thought to decrease the planet migration rate, including eccentric planet orbits and disk turbulence \\citep[]{lubow10,migrateref4}.  In all cases, planet-disk migration must take place while the primordial disk of gas and dust is still present during the classical T Tauri phase, or first $\\sim$5 Myr, of the host star's life \\citep[]{disklifetimeref,disklifetimeref2}.  Type II migration offers a mechanism to transport gas giants that must form beyond the snow line inward to their host star.   However, it is not well-constrained how planet migration is halted once started, lest the planet be tidally disrupted by the host star.  Possible braking mechanisms include -- tidal circularization \\citep[]{tidalref4,tidalref5,tidalref6, tidalref,tidalref2,tidalref3}, trapping the exoplanet in the 1:2 interior orbital resonance with the magnetospheric truncation radius \\citep[]{eisner05}, or trapping the planet in the 1:2 interior orbital resonance with the dust sublimation radius \\citep[e.g., ][]{kuchner02}.  For an approximately solar mass star, the gas disk truncation radius is comparable to the dust sublimation radius for a typical T-Tauri star magnetic field strength of $\\sim$2 kGauss \\citep[]{eisner05}, but that is not the case for lower and higher mass stars.  For lower-(higher-)mass T Tauri stars, the dust sublimation radius can be interior (exterior) to the estimated magnetospheric truncation radius (${\\S}$3).  In this work we investigate exoplanet distributions as a function of semi-major axis and host stellar mass as a test for migration halting mechanisms. The increasing number of exoplanet discoveries provides sufficiently large ensembles of close-in (e.g. $<$0.1 AU) exoplanets over a range of host stellar masses to discern which mechanism may be responsible for halting exoplanet migration.  In ${\\S}$2, we outline our empirical samples.  In ${\\S}$3, we present each migration halting mechanism model and its corresponding prediction for the exoplanet distribution density function.  In ${\\S}$4, we present our methodology to evaluate the success of each migration halting model at reproducing empirical distributions, and in ${\\S}$5 we present the results of these statistical tests.  In ${\\S}$6, we present our conclusions. ",
        "conclusions": "We use the empirical distribution of confirmed exoplanets and Kepler planet candidates in the host stellar mass -- exoplanet orbital semi-major axis plane as a diagnostic for migration halting mechanisms.  Migration halting from tidal circularization provides the best posterior probabilities for all empirical samples, favored by factors of 2.5--61 for the different sub-samples investigated herein when compared to halting at the 1:2 interior resonance with the magnetospheric disk truncation radius.  We can rule out migration halting at the 1:2 interior resonance with the dust sublimation radius,  a uniform random halting radius, and an exponential decay halting radius as viable models for the majority of observed close-in exoplanets.  Our generalized power law model favors a dependence of the halting distance with stellar mass that is stronger than predicted from tidal dissipation theories, and clearly rules out the independence of the halting distance on stellar mass, with posterior probability ratios ranging from 1.5--10$^4$.  The favoring of a stronger power law dependence than predicted for the tidal halting model suggests future theoretical work may be needed to better reproduce the observed sub-structure in the empirical distribution of exoplanets as a function of semi-major axis and stellar mass.  The authors would like to thank the referee Dan Fabrycky for their patience, corrections and comments to our analysis, which substantially improved this work.    This research has made use of the NASA Exoplanet Archive, which is operated by the California Institute of Technology, under contract with the National Aeronautics and Space Administration under the Exoplanet Exploration Program.  This research has made use of the Exoplanet Orbit Database and the Exoplanet Data Explorer at exoplanets.org.   The authors would like to thank Thayne Currie for his useful (and rapid turnaround) comments on the manuscript.  We would also like to acknowledge utility of the Bayesian statistics class taught by John Johnson at Caltech.  \\clearpage \\begin{landscape} \\begin{deluxetable}{lp{2.2cm}p{2.2cm}p{2.2cm}p{2.2cm}p{2.2cm}p{2.2cm}} \\tabletypesize{\\scriptsize} % \\rotate \\tablewidth{0pc} \\tablecolumns{5} \\tablecaption{Model Best Fit Parameter Values\\tablenotemark{a}} \\tablehead{ \\colhead{Model} & \\colhead{ M$_{pl}<10$M$_{\\oplus}$} & \\colhead{10M$_{\\oplus}<$M$_{pl}<0.3$M$_{J}$} & \\colhead{M$_{pl}>0.3$M$_{J}$} & \\colhead{R$_{pl}$ $<2R_\\oplus$} & \\colhead{$2R_\\oplus \\leq $R$_{pl}$$\\leq6R_\\oplus$} & \\colhead{R$_{pl}$$>6R_\\oplus$}} \\startdata $M^{1/3}$ Tidal Model & $C=0.052$,\\newline $\\sigma=0.34$ & $C=0.072$,\\newline $\\sigma=0.14$ & $C=0.042$,\\newline $\\sigma=0.16$ & $C=0.07$,\\newline $\\sigma=0.3$ & $C=0.102$,\\newline $\\sigma=0.26$ & $C=0.056$,\\newline $\\sigma=0.24$ \\\\ $M^{3/13}$ Tidal Model & $C=0.05$,\\newline $\\sigma=0.36$ & $C=0.07$,\\newline $\\sigma=0.14$ & $C=0.044$,\\newline $\\sigma=0.16$ & $C=0.068$,\\newline $\\sigma=0.3$ & $C=0.1$,\\newline $\\sigma=0.26$ & $C=0.056$,\\newline $\\sigma=0.24$ \\\\ Magnetospheric Truncation & $B_0=1.1547$ kG,\\newline $\\sigma=0.38$ & $B_0=2.4075$ kG,\\newline $\\sigma=0.16$ & $B_0=1.0683$ kG,\\newline $\\sigma=0.16$ & $B_0=2.2884$ kG,\\newline $\\sigma=0.3$ & $B_0=4.338$ kG,\\newline $\\sigma=0.26$ & $B_0=1.6292$ kG,\\newline $\\sigma=0.24$ \\\\ Dust Sublimation & $\\sigma=1.02$ & $\\sigma=1.02$ & $\\sigma=0.86$ & $\\sigma=1.02$ & $\\sigma=1.02$ & $\\sigma=1.02$ \\\\ Power Law & $\\alpha=0.9$,\\newline $C=0.062$,\\newline $\\sigma=0.34$ & $\\alpha=0.58$,\\newline $C=0.078$,\\newline $\\sigma=0.14$ & $\\alpha=0.68$,\\newline $C=0.042$,\\newline $\\sigma=0.16$ & $\\alpha=0.6$,\\newline $C=0.074$,\\newline $\\sigma=0.3$ & $\\alpha=0.38$,\\newline $C=0.104$,\\newline $\\sigma=0.26$ & $\\alpha=0.4$,\\newline $C=0.058$,\\newline $\\sigma=0.24$ \\\\ Uniform Random & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata \\\\ Exponential Decay & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata \\\\ Stellar Mass Independent & $A=0.048$,\\newline $\\sigma=0.5$ & $A=0.07$,\\newline $\\sigma=0.18$ & $A=0.044$,\\newline $\\sigma=0.16$ & $A=0.066$,\\newline $\\sigma=0.3$ & $A=0.098$,\\newline $\\sigma=0.26$ & $A=0.056$,\\newline $\\sigma=0.24$ \\\\ % Roche Radius & $\\sigma=1.02$ & $\\sigma=1.02$ & $\\sigma=1.02$ & $\\sigma=1.02$ & $\\sigma=1.02$ & $\\sigma=1.02$ \\\\  \\enddata \\tablenotetext{a}{$\\sigma$ values are the Gaussian kernel widths in units of log(AU) that are convolved with the analytic 1:1 curves to produce the model density functions as described in ${\\S}$3.  The units of log$(A)$ are also log(AU).  The second through fourth columns are for the confirmed exoplanet sub-samples, and the last three columns correspond to the KOIs.} \\end{deluxetable}  \\begin{deluxetable}{lllllll} \\tabletypesize{\\scriptsize} % \\rotate \\tablewidth{0pc} \\tablecolumns{5} \\tablecaption{Relative Bayesian Posterior Probabilities (Unnormalized)\\tablenotemark{a}} \\tablehead{ \\colhead{Model} & \\colhead{ M$_{pl}<10$M$_{\\oplus}$} & \\colhead{10M$_{\\oplus}<$M$_{pl}<0.3$M$_{J}$} & \\colhead{M$_{pl}>0.3$M$_{J}$} & \\colhead{R$_{pl}$ $<2R_\\oplus$} & \\colhead{$2R_\\oplus \\leq $R$_{pl}$$\\leq6R_\\oplus$} & \\colhead{R$_{pl}$$>6R_\\oplus$}} \\startdata $M^{1/3}$ Tidal Model\t& $1.2190 \\times 10^{-01}$ & $6.9976 \\times 10^{07}$ & $3.8776 \\times 10^{43}$ & $2.3146 \\times 10^{123}$ & $4.3350 \\times 10^{164}$ & $8.7087 \\times 10^{28}$ \\\\ $M^{3/13}$ Tidal Model\t& $7.9400 \\times 10^{-02}$ & $3.6697 \\times 10^{07}$ & $6.1489 \\times 10^{42}$ & $6.7614 \\times 10^{122}$ & $3.4503 \\times 10^{164}$ & $7.7497 \\times 10^{28}$ \\\\ Magnetospheric Truncation\t& $3.1100 \\times 10^{-02}$ & $8.4200 \\times 10^{06}$ & $6.4026 \\times 10^{41}$ & $8.3070 \\times 10^{121}$ & $8.9229 \\times 10^{163}$ & $3.5485 \\times 10^{28}$ \\\\ Dust Sublimation\t& $3.6000 \\times 10^{-03}$ & $7.7860 \\times 10^{-04}$ & $6.1275 \\times 10^{11}$ & $1.2069 \\times 10^{34}$ & $5.6110 \\times 10^{35}$ & $4.7453 \\times 10^{10}$ \\\\ Power Law\t& $2.1550 \\times 10^{-01}$ & $7.2185 \\times 10^{07}$ & $2.4825 \\times 10^{44}$ & $3.5253 \\times 10^{123}$ & $2.2953 \\times 10^{164}$ & $6.3967 \\times 10^{28}$ \\\\ Uniform Random\t& $7.5600 \\times 10^{-02}$ & $3.1060 \\times 10^{05}$ & $1.5013 \\times 10^{19}$ & $2.8514 \\times 10^{110}$ & $1.0866 \\times 10^{139}$ & $2.1531 \\times 10^{25}$ \\\\ Exponential Decay\t& $7.6600 \\times 10^{-02}$ & $3.0870 \\times 10^{-01}$ & $5.4692 \\times 10^{12}$ & $9.1955 \\times 10^{49}$ & $6.3116 \\times 10^{53}$ & $1.3048 \\times 10^{14}$ \\\\ Stellar Mass Independent\t& $3.2300 \\times 10^{-02}$ & $5.9364 \\times 10^{06}$ & $2.5710 \\times 10^{40}$ & $1.0893 \\times 10^{121}$ & $9.7962 \\times 10^{163}$ & $4.2242 \\times 10^{28}$ \\\\ %Roche Radius\t& $7.3107 \\times 10^{-04}$ & $4.8581 \\times 10^{-05}$ & $3.3891 \\times 10^{07}$ & $4.7831 \\times 10^{19}$ & $8.4426 \\times 10^{20}$ & $2.9812 \\times 10^{07}$ \\\\ \\enddata \\tablenotetext{a}{The second through fourth columns are for the confirmed exoplanet sub-samples, and the last three columns correspond to the KOIs.} \\end{deluxetable}   \\begin{deluxetable}{lllllll} \\tabletypesize{\\scriptsize} % \\rotate \\tablewidth{0pc} \\tablecolumns{6} \\tablecaption{Reduced $\\chi^2$ Values\\tablenotemark{a}} \\tablehead{ \\colhead{Model} & \\colhead{ M$_{pl}<10$M$_{\\oplus}$} & \\colhead{10M$_{\\oplus}<$M$_{pl}<0.3$M$_{J}$} & \\colhead{M$_{pl}>0.3$M$_{J}$} & \\colhead{R$_{pl}$ $<2R_\\oplus$} & \\colhead{$2R_\\oplus \\leq $R$_{pl}$$\\leq6R_\\oplus$} & \\colhead{R$_{pl}$$>6R_\\oplus$}} \\startdata $M^{1/3}$ Tidal Model\t&\t0.51739\t&\t0.32685\t&\t0.89054\t&\t1.0987\t&\t1.2792\t&\t0.35351\t\\\\ $M^{3/13}$ Tidal Model\t&\t0.54308\t&\t0.41063\t&\t1.6656\t&\t1.1376\t&\t1.2416\t&\t0.35557\t\\\\ Magnetospheric Truncation\t&\t0.57783\t&\t0.39862\t&\t2.4932\t&\t1.2987\t&\t1.2235\t&\t0.37129\t\\\\ Dust Sublimation\t&\t0.78871\t&\t1.8692\t&\t5.4704\t&\t13.1747\t&\t18.9103\t&\t2.7711\t\\\\ Power Law\t&\t0.5709\t&\t0.39799\t&\t0.55665\t&\t1.1233\t&\t1.3821\t&\t0.37203\t\\\\ Uniform Random\t&\t0.88506\t&\t0.7415\t&\t4.7696\t&\t3.7036\t&\t3.8442\t&\t1.1901\t\\\\ Exponential Decay\t&\t0.72727\t&\t1.6272\t&\t5.5312\t&\t11.078\t&\t15.8296\t&\t2.4884\t\\\\ Stellar Mass Independent\t&\t0.62176\t&\t0.51691\t&\t5.3648\t&\t1.518\t&\t1.4775\t&\t0.42532\t\\\\ % Roche Radius\t&\t0.87309\t&\t2.0973\t&\t6.1121\t&\t15.5875\t&\t22.0302\t&\t3.2525\t\\\\ \\enddata \\tablenotetext{a}{The second through fourth columns are for the confirmed exoplanet sub-samples, and the last three columns correspond to the KOIs.} \\end{deluxetable} \\end{landscape}  \\clearpage \\begin{landscape} \\begin{figure} \\centering \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1a} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1b} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1c} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1d} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1e} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1f} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1g} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1h} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1i} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1j} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1k} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1l} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1m} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1n} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1o} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1p} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1q} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f1r} \\caption{First row: Empirical distributions of exoplanet semi-major axes and host star masses for (from left to right): confirmed exoplanets as of February 2012 (NASA Exoplanet Archive) with M$_{pl}$$<$10 M$_{\\oplus}$, confirmed exoplanets with 10M$_{\\oplus}<$M$_{pl}<0.3$M$_{J}$, confirmed exoplanets with M$_{pl}>0.3$M$_{J}$, Kepler Objects of Interest from \\citet[]{batalha} with R$_{pl}$ $<2R_\\oplus$, Kepler Objects of Interest with $2R_\\oplus \\leq $R$_{pl}$$\\leq6R_\\oplus$, and Kepler Objects of Interest with R$_{pl}$$>6R_\\oplus$.  Second row: Empirical density functions generated for each data set using a method analogous to that in \\citep[]{wasserman}.  Third row: Planet frequency as a function of stellar mass for each data set in the first row, shown with arbitrary normalization.  This is used as a normalization for our migration braking models to remove survey sample selection effects as a function of stellar mass.} \\end{figure} \\end{landscape}  \\clearpage \\begin{landscape} \\begin{figure} \\centering \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2a} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2b} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2c} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2d} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2e} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2f} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2g} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2h} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2i} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2j} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2k} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2l} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2m} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2n} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2o} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2p} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2q} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f2r} \\caption{Top: Empirical density functions as in the second row of Figure 1.  Middle: Model probability density function with optimal parameters in Table 1 for the halting model at the interior 1:2 orbital resonance with the magnetospheric disk truncation radius.  Bottom: Subtracted residuals.} \\end{figure} \\end{landscape}  \\clearpage \\begin{landscape} \\begin{figure} \\centering \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3a} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3b} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3c} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3d} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3e} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3f} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3g} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3h} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3i} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3j} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3k} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3l} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3m} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3n} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3o} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3p} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3q} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f3r} \\caption{Same as Figure 2 for the Power Law halting model, with an optimized exponent free parameter.} \\end{figure} \\end{landscape}  \\clearpage \\begin{landscape} \\begin{figure} \\centering \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4a} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4b} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4c} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4d} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4e} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4f} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4g} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4h} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4i} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4j} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4k} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4l} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4m} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4n} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4o} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4p} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4q} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f4r} \\caption{Same as Figure 2 for the Roche radius a $\\propto$ M$^{1/3}$ tidal halting model.} \\end{figure} \\end{landscape}  \\clearpage \\begin{landscape} \\begin{figure} \\centering \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5a} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5b} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5c} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5d} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5e} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5f} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5g} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5h} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5i} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5j} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5k} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5l} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5m} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5n} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5o} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5p} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5q} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f5r} \\caption{Same as Figure 2 for the a $\\propto$ M$^{3/13}$ tidal halting model.} \\end{figure} \\end{landscape}  \\clearpage \\begin{landscape} \\begin{figure} \\centering \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6a} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6b} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6c} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6d} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6e} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6f} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6g} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6h} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6i} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6j} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6k} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6l} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6m} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6n} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6o} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6p} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6q} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f6r} \\caption{Same as Figure 2 for the stellar-mass independent halting model at a constant semi-major axis.} \\end{figure} \\end{landscape}  \\clearpage \\begin{landscape} \\begin{figure} \\centering \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7a} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7b} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7c} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7d} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7e} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7f} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7g} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7h} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7i} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7j} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7k} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7l} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7m} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7n} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7o} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7p} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7q} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f7r} \\caption{Same as Figure 2 for the uniform random halting model.} \\end{figure} \\end{landscape}  \\clearpage \\begin{landscape} \\begin{figure} \\centering \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8a} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8b} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8c} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8d} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8e} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8f} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8g} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8h} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8i} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8j} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8k} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8l} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8m} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8n} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8o} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8p} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8q} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f8r} \\caption{Same as Figure 2 for the exponential decay with semi-major axis tidal halting model.} \\end{figure} \\end{landscape}  \\clearpage \\begin{landscape} \\begin{figure} \\centering \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9a} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9b} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9c} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9d} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9e} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9f} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9g} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9h} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9i} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9j} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9k} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9l} \\\\ \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9m} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9n} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9o} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9p} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9q} \\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{f9r} \\caption{Same as Figure 2 for the halting model at the interior 1:2 orbital resonance with the dust sublimation radius.} \\end{figure} \\end{landscape}  %\\clearpage %\\begin{landscape} %\\begin{figure} %\\centering %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10a} %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10b} %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10c} %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10d} %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10e} %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10f} \\\\ %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10g} %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10h} %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10i} %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10j} %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10k} %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10l} \\\\ %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10m} %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10n} %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10o} %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10p} %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10q} %\\includegraphics[width=0.20\\textwidth,clip=true,trim=0cm 0cm 0cm 0cm]{figs/f10r} %\\caption{Same as Figure 2 for the Roche radius halting model.} %\\end{figure} %\\end{landscape}"
    },
    "1112/1112.4819_arXiv.txt": {
        "abstract": "The process referred to as ``semi-convection'' in astrophysics and ``double-diffusive convection in the diffusive regime'' in Earth and planetary sciences, occurs in stellar and planetary interiors in regions which are stable according to the Ledoux criterion but unstable according to the Schwarzschild criterion. In this series of papers, we analyze the results of an extensive suite of 3D numerical simulations of the process, and ultimately propose a new 1D prescription for heat and compositional transport in this regime which can be used in stellar or planetary structure and evolution models. In a preliminary study of the phenomenon, Rosenblum et al. (2011) showed that, after saturation of the primary instability, a system can evolve in one of two possible ways: the induced turbulence either remains homogeneous, with very weak transport properties, or transitions into a thermo-compositional staircase where the transport rate is much larger (albeit still smaller than in standard convection).  In this paper, we show that this dichotomous behavior is a robust property of semi-convection across a wide region of parameter space. We propose a simple semi-analytical criterion to determine whether layer formation is expected or not, and at what rate it proceeds, as a function of the background stratification and of the diffusion parameters (viscosity, thermal diffusivity and compositional diffusivity) only. The theoretical criterion matches the outcome of our numerical simulations very adequately in the numerically accessible ``planetary'' parameter regime, and can easily be extrapolated to the stellar parameter regime.  Subsequent papers will address more specifically the question of quantifying transport in the layered case and in the non-layered case. ",
        "introduction": "\\subsection{The physics of semi-convection: double-diffusive convection} \\label{sec:intro-basic}  The concept of ``semi-convection'', first introduced by \\citet{schwartzchildharm1958}, is often invoked in a number of rather different situations \\citep{merryfield1995} which nevertheless have one point in common: they occur in regions which are stable to the Ledoux criterion, but unstable to the Schwarzschild criterion. Mathematically speaking, this condition can be expressed as \\begin{eqnarray} 0 < \\left( \\frac{\\partial \\ln T}{\\partial \\ln p} \\right) - \\left( \\frac{\\partial \\ln T}{\\partial \\ln p} \\right)_{\\rm ad} & < & \\left( \\frac{\\partial \\ln \\mu}{\\partial \\ln p} \\right) \\mbox{   , equivalently    } 0 < \\nabla - \\nabla_{\\rm ad} < \\nabla_\\mu \\label{eq:sccriterion} \\end{eqnarray} where $T$, $p$ and $\\mu$ are the temperature, gas pressure, and mean molecular weight, and where the subscript ``ad'' denotes a derivative at constant entropy. Physically speaking, (\\ref{eq:sccriterion}) describes regions which are thermally unstably stratified, but where standard convection is suppressed by the presence of significant compositional gradients.  The first linear analysis of the stability of ``semi-convective'' regions was presented by \\citet{walin1964} in the oceanographic context, and later by \\citet{kato1966} in the astrophysical context. They both showed that a semi-convective region is hydrodynamically unstable, but to a much more gentle instability than the one associated with standard convection because it relies on doubly-diffusive processes to grow. It is in fact one of two forms of  so-called ``double-diffusive convection'' (the other being fingering convection, otherwise known as ``thermohaline convection'' in astrophysics), and is often referred to as ``double-diffusive convection in the diffusive regime'' in physical oceanography. For the sake of clarity and brevity, we will refer to the phenomenon as ``diffusive convection'' in this series of papers.  As reviewed by \\citet{rosenblum2011}, diffusive convection is principally controlled by three non-dimensional parameters. The first two characterize the nature of the fluid considered, and are the Prandtl number Pr and the diffusivity ratio $\\tau$, \\begin{equation} {\\rm Pr} = \\frac{\\nu}{\\kappa_T} \\mbox{    ,   } \\tau = \\frac{\\kappa_\\mu}{\\kappa_T} \\mbox{    ,   } \\label{eq:Prtaudef} \\end{equation} where $\\nu$, $\\kappa_T$ and $\\kappa_\\mu$ are the viscosity, thermal diffusivity and compositional diffusivity respectively. For reference, note that Pr and $\\tau$ are very roughly of the order of $10^{-2}$ in giant planet interiors and $10^{-6}$ in stellar interiors, and that $\\tau$ is usually somewhat smaller than $\\Pr$.  The third parameter is the inverse density ratio, defined as \\begin{equation} R^{-1}_0 = \\frac{\\nabla_\\mu}{\\nabla - \\nabla_{\\rm ad}}\\mbox{   , } \\label{eq:R0def} \\end{equation} which measures the relative importance of the destabilizing thermal stratification compared with the stabilizing compositional one. A semi-convective region is unstable \\citep{walin1964,kato1966} if \\begin{equation} 1 \\le R_0^{-1} \\le \\frac{{\\rm Pr} + 1}{{\\rm Pr} + \\tau}  = R^{-1}_{\\rm c} \\mbox{   , } \\label{eq:instabrange} \\end{equation} whereas regions with $R_0^{-1} < 1$ are unstable to overturning convection and those with $R_0^{-1} > R^{-1}_{\\rm c}$ are absolutely stable. Within this parameter range, it can be shown that the growth rate of the linear modes is complex. Furthermore, in the low Prandtl number regime characteristic of stellar and planetary interiors, the real part of the growth rate is proportional to the square root of the Prandtl number times the Br\\\"unt-V\\\"ais\\\"al\\\"a frequency (see Appendix A). In the same limit, the typical lengthscale of the unstable modes is a thermal diffusion scale. %  Linear stability is unfortunately of rather limited utility, in particular when it comes to estimating the mixing rates induced by the diffusive convection. Experiments -- laboratory or numerical -- are the only way forward.  Since no terrestrial fluid exists with similar values of Pr and $\\tau$, it is tempting to use experimental measurements of heat and compositional fluxes in different parameter regimes, in particular laboratory experiments in the heat-salt system relevant for physical oceanography \\citep{lindenshirtcliffe1978}, and extrapolate them to the astrophysical case \\citep{stevenson1982,Gu04}. However one must be very cautious in doing so since the Prandtl number of water is typically 4-7, depending on temperature, whereas the Prandtl number in stellar and planetary interiors is much lower than one. Since the typical scale of the instability is a thermal diffusion scale, double-diffusive mixing is a mostly-laminar process at high Pr and a turbulent one at low Pr. There is no reason to expect that the laminar scalings should apply to the turbulent case.  Nevertheless, from a qualitative point of view, one of the most interesting results from laboratory \\citep{turnerstommel1964} and field experiments \\citep{timmermans2008} in salt water is the fact that diffusive convection has a tendency to form thermohaline staircases, i.e. well-defined mixed layers separated by thin, very strongly stratified, and essentially diffusive interfaces. It is often thought that the necessarily weak transport through these interfaces is what controls and limits the efficiency of transport by ``layered'' convection. Such considerations have led \\citet{spruit1992} and \\citet{Chabrier07c} for example to propose theories for heat and compositional transport in astrophysics, which rely on assumptions about the layer heights and the interface thicknesses. However, it is important to remember that, until very recently, layered convection had never been demonstrated to exist at low Prandtl number. As a result, these theories have, by and large, remained un-tested (see \\citet{rosenblum2011} for a review of prior numerical work).  \\subsection{Recent numerical and theoretical results} \\label{sec:intro-rosen}  Recent 3D numerical simulations have finally shed some light on the subject of transport by diffusive convection. \\citet{rosenblum2011} ran an exhaustive numerical study of the phenomenon for fixed Prandtl number and diffusivity ratio Pr$=\\tau=0.3$, and with the inverse density ratio $R_0^{-1}$ spanning the entire instability range (\\ref{eq:instabrange}). While still far from any astrophysically-relevant regime, the selected values of Pr and $\\tau$ in these simulations were below unity and therefore in the right ``region'' of parameter space. \\citet{rosenblum2011} found that the instability grows as expected and that the early behavior of the fluid can be satisfactorily explained by considering the fastest-growing modes of instability only. However, they discovered that two very different regimes of diffusive convection are possible after saturation of what we will refer to as the ``primary'' instability, depending on the value of the inverse density ratio. The various regimes are illustrated in Figure \\ref{fig:layerregime}.  \\begin{figure} \\includegraphics[width=\\textwidth]{f1.jpeg} \\caption{\\small Illustration of the various regimes of diffusive convection. In systems without compositional gradients, the Schwarzschild criterion marks the stability boundary between overturning convection and absolute stability. In the presence of a stable compositional gradient, diffusive convection occurs for $R_0^{-1}$ between 1 (which corresponds to the Ledoux-stability limit) and $R_c^{-1} = (\\Pr+1)/(\\Pr+\\tau)$. Within this range, two possibilities arise: for $R_0^{-1} \\in [1,R_L^{-1}]$, spontaneous transition into layered convection is observed, while for $R_0^{-1} \\in [R_L^{-1}, R_c^{-1}]$, the system remains in a state of weak oscillatory convection. Note that both $R_L^{-1}$ and $R_c^{-1}$ depend on $\\Pr$ and $\\tau$. } \\label{fig:layerregime} \\end{figure}  For large enough $R_0^{-1}$, i.e. for more ``stable'' stratifications, \\citet{rosenblum2011} showed that the system settles into a homogeneous, statistically stationary, weakly turbulent state. The turbulence is dominated by internal gravity waves\\footnote{Recall that the background stratification is stably stratified in terms of the density and therefore supports standard gravity waves which oscillate with the buoyancy frequency.} and mixing occurs principally via wave-breaking. Heat and compositional transport are fairly inefficient, and depend sensitively on the inverse density ratio. For low enough $R_0^{-1}$ on the other hand, which corresponds to systems closer to the Ledoux-stability criterion, they observed the spontaneous emergence of thermo-compositional staircases after a short adjustment period. These staircases take the form of vigorously convective layers, which are thermally and compositionally well-mixed, and separated by fairly sharp interfaces, as observed in the oceanographic case. The interfaces, however, are far from merely diffusive and are instead very dynamic, and often pierced by strong localized updrafts and downdrafts. Later on, the layers are observed to merge, and each merger is accompanied by a significant increase in the overall transport across the staircase. \\citet{rosenblum2011} found that transport in that regime depends sensitively on the mean layer height, and proposed preliminary scalings to quantify it. The latter remain to be verified across a wider region of parameter space.  In fact, insight into the reason for this dichotomous ``layered vs. non-layered'' convection, can be gained from recent oceanographic studies of fingering convection, a related double-diffusive instability of thermally stable fluids which are destabilized by adverse compositional gradients \\citep{stern1960sfa}. Such conditions are found in the tropical thermocline for example, where surface heating and evaporation continually warm up the upper layers of water, and increase its salt concentration. Crucially, thermohaline staircases are also commonly found in fingering regions of the ocean \\citep{schmitt2005edm}. \\citet{radko2003mlf} studied their formation in this regime, and showed that they can naturally emerge as a secondary large-scale ``mean-field'' instability of the system. More precisely, he showed that horizontally invariant but vertically sinusoidal density perturbations grow exponentially out of the homogeneous, small-scale fingering convection, and eventually overturn into a regularly-spaced staircase. His theory was later validated by \\citet{stellmach2011} via  three-dimensional numerical simulations.  A crucial result of Radko's theory is his identification of a necessary and sufficient condition for the layer-forming instability, namely that the turbulent flux ratio $\\gamma_{\\rm turb}$, defined as the ratio of the turbulent buoyancy flux due to heat transport, to the turbulent buoyancy flux due to salt transport, should be a decreasing function of the density ratio $R_0$. He thus named the instability ``the $\\gamma-$instability''. The fact that $\\gamma_{\\rm turb}$ in salt water decreases with $R_0$ for low density ratios, then increases again for higher density ratios, explains why thermohaline staircases in the tropical ocean are only found in regions with low enough $R_0$.  Recently, \\citet{rosenblum2011} showed that Radko's $\\gamma-$instability theory can very easily and naturally be extended to explain the emergence of staircases in their own simulations of diffusive convection. The equivalent condition for instability is that the {\\it total} flux ratio $\\gamma_{\\rm tot}$ (i.e. the ratio of the total buoyancy fluxes, diffusive plus advective, of heat to composition respectively), should be a decreasing function of $R_0$. Since the inverse density ratio $R_0^{-1}$ is a more convenient parameter in diffusive convection, an equivalent sufficient condition for instability is that $\\gamma_{\\rm tot}^{-1}$ should be a decreasing function of $R_0^{-1}$. For completeness, the $\\gamma-$instability theory is rederived and discussed in Section \\ref{sec:newtheory}. \\citet{rosenblum2011} found through their systematic exploration of the instability range that $\\gamma_{\\rm tot}^{-1}$ has a minimum at about $R_L^{-1} = 1.4$ when  $\\Pr=\\tau=0.3$. This explains why layers are seen to form for $R_0^{-1} < 1.4$ in their simulations (at these values of Pr and $\\tau$) but not for larger $R_0^{-1}$. Furthermore, in the simulations which do lead to layering, \\citet{rosenblum2011}  found that theory and numerical experiments agree remarkably well on the growth rate of the $\\gamma-$instability. Their preliminary study thus suggested that, in order to know under which conditions layer formation is possible in stars and giant planets, one simply needs to determine if and when $\\gamma_{\\rm tot}^{-1}$ decreases with $R_0^{-1}$, for a given parameter pair $(\\Pr, \\tau)$.  \\subsection{Work outline}  The findings of \\citet{rosenblum2011} lay a very clear path towards creating a practical model for transport by diffusive convection (semi-convection) in astrophysics: \\begin{enumerate} \\item Model the function $\\gamma_{\\rm tot}^{-1}(R_0^{-1};\\Pr,\\tau)$ (from numerical simulations and/or theoretical calculations) to determine if and when layered convection is expected, and verify whether the $\\gamma$-instability predictions continue to hold at lower $\\Pr$ and $\\tau$. \\item Characterize transport by layered convection, and in particular, its dependence on layer height, $\\Pr$, $\\tau$ and $R_0^{-1}$. \\item Characterize transport by homogeneous diffusive convection (i.e. in the absence of layers). \\end{enumerate} The first step is addressed in this paper, while steps 2 and 3 are deferred to subsequent publications in the series.  In the present paper, we therefore focus our efforts on a precise determination of the region of parameter space where layered convection is expected to occur. We do so using a combination of numerical simulations and theory.  We discuss the numerical model and present typical results in Section \\ref{sec:model}. We review the $\\gamma-$instability theory in Section \\ref{sec:newtheory}. In Section \\ref{sec:gammaextract} we outline the methodology used for extracting the value of the flux ratio $\\gamma^{-1}_{\\rm tot}$ from simulations at numerically-accessible parameters, and present our results. In Section \\ref{sec:gammapred} we present a simple semi-analytical theory which enables us to estimate $\\gamma^{-1}_{\\rm tot}$ for any set of parameters, and compare it with our  numerical results. In Section \\ref{sec:layers}, we show that the predicted growth rates from the $\\gamma-$instability theory match the results of our numerical simulations very well. Finally, we conclude in Section \\ref{sec:ccl}. ",
        "conclusions": "\\label{sec:ccl}  The ultimate goal of this series of papers is to propose a new model for mixing by diffusive convection. In the work presented here, we ran and analyzed a very extensive suite of numerical simulations of the process in a wide range of parameter space. We have shown that in the astrophysically-relevant low Prandtl number ($\\Pr = \\nu/\\kappa_T$), low diffusivity ratio ($\\tau = \\kappa_\\mu/\\kappa_T$) regime, diffusive convection can take one of two forms depending on the local inverse density ratio $R_0^{-1} = \\grad_\\mu/(\\grad - \\grad_{\\rm ad})$: moderately efficient layered convection at lower $R_0^{-1}$, or inefficient wave-dominated ``oscillatory'' convection for higher $R_0^{-1}$ (see Figure \\ref{fig:layerregime}). We have confirmed through numerical and analytical work that a spontaneous transition into layered convection occurs, under predictable circumstances, through a linear mean-field instability of the initial state of oscillatory convection. This instability is called the $\\gamma-$instability. It was originally suggested in the oceanographic context of fingering convection by \\citet{radko2003mlf} and later applied to double-diffusive convection in astrophysics by \\citet{rosenblum2011}.  Since transport in layered convection is much more efficient than in the non-layered case, a crucial element of any new model of diffusive convection will be the availability of a practical criterion for determining, for given $R_0^{-1}$, $\\Pr$ and $\\tau$, whether a system is expected to transition into layers or not, and on what timescale. The original $\\gamma-$instability theory provides such a criterion, but the latter can only be used in practice provided experimental measurements of the turbulent fluxes in homogeneous diffusive convection, for the same parameters, are available. We provide such measurements here for the planetary parameter regime, but similar results are unlikely to ever be available for the much more extreme stellar parameter regime.  Based on these considerations, we then proposed a new empirically motivated model for the turbulent fluxes, which enable us to derive a completely parameter-free semi-analytical criterion to determine, for {\\it any} given fluid in the astrophysical regime ($\\Pr, \\tau \\ll 1$) and any given stratification ($R_0^{-1}$), (a) whether a system is expected to transition into layers or not, and (b) on what timescales the layers are expected to emerge. Our model was found to fit the available numerical data very well, and can therefore be used very reliably within the same region of parameter space (e.g. the planetary parameter regime, for which $\\Pr, \\tau \\sim 10^{-2}-10^{-1}$). We further propose that it should also be used in regions of parameter space for which fully resolved simulations are not available, namely in the stellar parameter regime ($\\Pr, \\tau \\sim 10^{-7}-10^{-5}$). Based on this model, we find that layered convection is theoretically expected in stellar interiors for a fairly wide range of parameter space, with $R_0^{-1}$ between 1 and about 1000.  Our results answer, at least approximately, the first of the three questions we initially raised: (1) Under which conditions do layers form? (2) What is the transport rate in layered convection and (3) What is the transport rate in non-layered ``oscillatory'' convection. In subsequent papers in this series we will continue our investigation by answering questions (2) and (3)."
    },
    "1112/1112.1920_arXiv.txt": {
        "abstract": "We present results from the first generation of multi-dimensional general relativistic neutrino hydrodynamics simulations of core-collapse supernovae. A comparison with models computed using either the purely Newtonian approximation or the ``effective gravitational potential'' approach reveals appreciable quantitative differences in the heating conditions and the gravitational wave spectra. Our results underscore the important role of general relativity in the supernova problem (which appears to be on par with other important factors such as the dimensionality and the equation of state) both for our understanding of the explosion dynamics as well as for predictions of observable signatures. ",
        "introduction": "A core-collapse supernova occurs at the end of the life of a massive star when it has exhausted its nuclear fuel at its centre. The nuclear ashes form an iron core which undergoes a gravitational collapse once it reaches a mass of roughly $1.4 M_\\odot$. When the core reaches supranuclear densities, the collapse is halted due the stiffening of the equation of state, and a shock wave is launched as the core rebounds. The initial kinetic energy of the shock is quickly spent for disintegrating the nuclei in the infalling material material into free nucleons, and the shock is further weakened by the rapid emission of several $10^{51} \\ \\mathrm{erg}$ in neutrinos from the post-shock region during the neutronization burst. The stalled shock typically hovers at a radius of $\\sim 100 \\ldots 200 \\ \\mathrm{km}$ -- possibly for hundreds of milliseconds -- until it is revived, and expels the envelope of the progenitor.  The nature of the explosion mechanism has been a subject of intense research in computational astrophysics for several decades. Several mechanisms have been proposed to explain the revival of the shock, the most prominent one being the ``delayed neutrino-driven mechanism'' \\citep{bethe_85,wilson_85}, which relies on the deposition of energy by neutrinos in the ``gain layer'' behind the shock to re-energize the shock. Except for a special class of progenitors with an O-Ne-Mg core \\citep{kitaura_06}, this mechanism only works in concert with multi-dimensional hydrodynamic effects such as convection and the so-called ``standing accretion shock instability'' (SASI, \\citealp{blondin_03,foglizzo_06,ohnishi_06}), which both increase the efficiency of neutrino heating in the gain region. While some of the most ambitious supernova simulations in axisymmetry \\citep{buras_06_b, marek_09,bruenn_10,suwa_10} indicate that the neutrino-driven mechanism may indeed work, there are also alternative scenarios such a magnetohydrodynamically-driven explosions \\citep{bk_70,akiyama_03,dessart_07_a}, acoustically-powered supernovae (\\citealp{burrows_06}; see however \\citealp{quataert_08} for criticism), and explosions triggered by a QCD phase transition \\citep{sagert_09}.  From a computational point of view, core-collapse supernovae present a number of challenges: Not only are they inherently multi-dimensional phenomena due to the operation of hydrodynamical instabilities, but in order to accurately capture the crucial effects of neutrino cooling and heating in the optically thick and thin regimes, the problem needs to be treated within the framework of kinetic theory. As the direct solution of the full Boltzmann equation for neutrinos is currently only feasible under the assumption of spherical symmetry \\citep{yamada_99,liebendoerfer_04}, a variety of different approximation strategies for the neutrino transport are used in the most sophisticated multi-dimensional supernova simulations, including ``ray-by-ray'' transport combined with variable Eddington factor techniques \\citep{buras_06_a} or flux-limited multi-group diffusion schemes \\citep{bruenn_10}, multi-angle Boltzmann transport without energy-bin coupling \\citep{ott_08_a}, two-moment schemes with an analytic closure \\citep{obergaulinger_11}, and the isotropic diffusion source approximation \\citep{liebendoerfer_09}.  Moreover, \\emph{general relativity} (GR) plays a major role in the supernova problem due to the compactness of the proto-neutron star $(GM/Rc^2 \\sim 0.1 \\ldots 0.2)$ and the occurrence of high velocities ($\\sim 0.3 c$). A general relativistic treatment is also required for precise predictions of the gravitational wave signal from core-collapse supernovae.  However, relativistic supernova simulations with up-to-date neutrino transport have long been limited to the case of spherical symmetry \\citep{bruenn_01,yamada_99,liebendoerfer_04}. In order to study the effects of GR in the context of multi-dimensional models, and in particular its impact on the gravitational wave emission from the supernovae core, we have recently \\citep{mueller_10} introduced a generalization of the ray-by-ray-plus variable Eddington factor method \\citep{rampp_02,buras_06_a}. In this paper, we present results from the first generation of axisymmetric relativistic supernova simulations using our \\textsc{Vertex-CoCoNuT} code and summarize the results of a comparison with the Newtonian case as well as with the pseudo-Newtonian ``effective potential'' approximation. ",
        "conclusions": "Multi-dimensional general relativistic simulations of core-collapse with a sophisticated treatment of the microphysics and the neutrino transport on par with the best currently available Newtonian models have only recently become possible, but the first available results presented here in this paper already serve to underscore the importance of general relativity in the supernova problem. For a $15 M_\\odot$ progenitor, we found that GR somewhat improves the heating conditions compared to models computed in the Newtonian and the effective potential approximations, which, unlike the GR model, fail to explode. The gravitational wave spectra are also considerably changed by GR effects, which shift the typical frequency of the time-integrated signal upward by $\\sim 80\\%$ compared to the purely Newtonian case. We therefore conclude that an accurate treatment of GR effects may be no less relevant for a better understanding of the neutrino-driven explosion mechanism and quantitative predictions of the signals from core-collapse supernovae than other (undoubtedly important) key factors that have recently been discussed such as dimensionality issues \\citep{nordhaus_10,hanke_11,takiwaki_11} and the nuclear equation of state."
    },
    "1112/1112.3100_arXiv.txt": {
        "abstract": "{Baryon Acoustic Oscillations (BAOs) are oscillatory features in the galaxy power spectrum which are used as a standard rod to measure the cosmological expansion. These have been studied in Cartesian space (Fourier or real space) or in spherical harmonic space in thin shells.}{Future wide-field surveys will cover both wide and deep regions of the sky and thus require a simultaneous treatment of the spherical sky and of an extended radial coverage.  The Spherical Fourier-Bessel (SFB) decomposition is a natural basis for the analysis of fields in this geometry and facilitates the combination of BAO surveys with other cosmological probes readily described in this basis. In this paper, we present a new way to analyse BAOs by studying the BAO wiggles from the SFB power spectrum.} {In SFB space, the power spectrum generally has both a radial ($k$) and tangential ($\\ell$) dependence and so do the BAOs. In the deep survey limit and ignoring evolution, the SFB power spectrum is purely radial and reduces to the Cartesian Fourier power spectrum. In the opposite limit of a thin shell, all the  information is contained in the tangential modes described by the 2D spherical harmonic power spectrum.}{We find that the radialisation of the SFB power spectrum is still a good approximation even when considering an evolving and biased galaxy field with a finite selection function.  This effect can be observed by all-sky surveys with depths comparable to current surveys. We also find that the BAOs radialise more rapidly than the full SFB power spectrum. }{Our results suggest the first peak of the BAOs in SFB space becomes radial out to $\\ell \\sim 10$ for all-sky surveys with the same depth as SDSS or 2dF, and out to $\\ell \\sim 70$ for an all-sky stage IV survey. Subsequent BAO peaks will also become radial, but for shallow surveys these may be in the non-linear regime.  For modes that have become radial, measurements at different $\\ell$'s are useful in practice to reduce measurement errors.} ",
        "introduction": "The study of Large Scale Structure (LSS) with galaxy surveys is a promising tool to study the dark universe \\citep{WGFC,DETF}. Baryon Acoustic Oscillations (BAOs) are a special feature in the galaxy power spectrum present on scales $100h^{-1}$Mpc, which are due to oscillations in the coupled baryon-photon fluid before recombination \\citep{Sunyaev:1970Z,Peebles:1970Y,Eisenstein:2005,Seo:2003Eis, Seo:2007Eis}. The BAOs are considered a powerful cosmological tool as the BAO scale acts as a standard ruler with which to probe cosmic expansion both in the radial and tangential directions.  BAOs were first detected by \\cite{Eisenstein:2005} in SDSS data \\citep{SDSS} and later with 2dF galaxies \\citep{2dFdata:2003,Cole:2005} and finally with both surveys \\citep{Percival:2007}, though others suggest current data cannot currently probe the BAO scale sufficiently \\citep{Cabre:Gaz}.  Until now, BAOs have been studied in Cartesian space, either in Fourier space \\citep{Seo:2003Eis,Seo:2007Eis}, or in real space \\citep{Eisenstein:2005,baoconfig1,baoconfig2}, and in 2D spherical harmonic space on thin spherical shells  \\citep{Dolney:2006}. These descriptions use different information and therefore have different constraining power for cosmological parameters \\citep{Rassat:2008bao}.  Future wide-field BAO surveys will, however, cover both large and deep areas of the sky, and thus require a simultaneous treatment of the spherical sky geometry and of extended radial coverage. The Spherical Fourier-Bessel (SFB) decomposition is a natural basis for the analysis of fields in this geometry.  The SFB analysis is powerful as it uses a coordinate system in which the radial selection function and physical effects are naturally described. Moreover, this description facilitates the combination of BAO surveys with other cosmological probes which are readily described in the SFB decomposition such as the smooth power spectrum and redshift space distortions \\citep{Heavens:1995,Fisher:1995,Percival:2004,Erdogdu:2005wi,Erdogdu:2006dv}, weak lensing \\citep{Heavens:2003,CHK:2005,Kitching:2008}, and the Integrated Sachs-Wolfe effect \\cite[ISW,][]{ISW3D}. Studying BAOs from wide-field surveys with an SFB expansion,  is therefore natural both for the geometry considered, and for unifying the treatment with the other probes.  In \\S\\ref{sec:theory} we first review the general decomposition of a 3D field in SFB space and introduce the concept of \\emph{radialisation} when the field is statistically isotropic and homogeneous. In \\S\\ref{sec:3dBAO:1}, we introduce a new way to consider the BAOs, by studying the baryon wiggles in SFB space, in a similar way as in Fourier space, while in \\S\\ref{sec:radialisation} we discuss the radialisation of BAOs in SFB space in the context of existing and future surveys. In \\S\\ref {sec:conclusion} we present our conclusions. ",
        "conclusions": "\\label{sec:conclusion}  In this paper, we have presented a new way to study Baryon Acoustic Oscillations (BAOs) using a Spherical Fourier-Bessel (SFB) decomposition of a wide-field deep galaxy survey. The BAO signal in SFB space can be studied by considering the ratio $R^C_\\ell(k)$ of the SFB power spectrum with wiggles $\\left(C^{\\rm b}_\\ell(k)\\right)$ to the smooth SFB power spectrum $\\left(C_\\ell^{\\rm no b}(k)\\right)$, similarly to what is done in Fourier space. In this decomposition, BAOs can be observed simultaneously in both the radial ($k$) and tangential ($\\ell$) modes.  For a field which is statistically isotropic and homogeneous (SIH), the SFB power spectrum is purely radial, i.e. independent of tangential modes, and reduces to the Cartesian Fourier power spectrum. In the other extreme limit where the field is in a thin shell, we showed that all the SFB information is contained in the tangential modes and is simply related to the 2D spherical harmonic power spectrum.  Considering the practical case where the field is observed with a radial selection function (thus partially violating the SIH condition), we find that the radial limit can still be reached for selection functions covering a large radial range (see Appendix \\ref{app:gauss}), which we refer to as the \\emph{radialisation} of the power spectrum. The radialisation will be limited to a region in ($\\ell,k$) space corresponding to $k \\gg \\sqrt{\\ell (\\ell+1)}/r_0$, where $r_0$ parameterises the radial selection function. We find that  radialisation is a good approximation for these modes even when evolution due to growth and bias are considered, both of which can be considered as a further violation of the SIH condition. To study this limit, we have derived an analytic solution to the window function of a non-evolving, unbiased galaxy field when the radial selection function is a Gaussian in $r$ centered on the observer (Equation \\ref{eq:analytic:window}).  We also find the BAOs considered in SFB space radialise as the survey depth is increased, meaning both the phase and amplitude of the BAOs tend towards the Fourier space ratio $R^P(k)$. This means that BAOs for a wide-field shallow survey have smaller amplitude and are spread across the $(\\ell,k)$ space, while BAOs for a wide-field deep survey have a larger amplitude and are confined to the radial modes (Figures \\ref{fig:clk} and \\ref{fig:slices}).  We study the radial limit analytically (Figure \\ref{fig:radial}) and find that it can in principle be observed (for a limited $\\ell$-range and for small physical scales) with all-sky surveys with current surveys depths or for future stage-IV surveys.  This suggest that the first BAO peak becomes radial up to $\\ell \\sim 10$ for an all-sky survey with similar depth to SDSS or 2dF, and up to $\\ell \\sim 70$ for an all-sky stage IV survey. In practice though, the radialisation might be observable to even higher $\\ell$ since we observe that the BAOs radialise more rapidly than the full SFB power spectrum, especially at large $\\ell$. Subsequent BAO peaks also become radial for large values of $\\ell$ and for shallower surveys, though these may be already be in the non-linear regime. For modes that have become radial, measurements at different $\\ell$'s are useful in practice, to reduce measurement errors due to cosmic variance and shot noise.  We note that we have ignored redshift-space distortions in our analysis though the prescription for these in SFB space is well known, \\citep[see for e.g.,][]{Heavens:1995}. These distortions may affect the radialisation as they will introduce mode-mixing, albeit with a distinct signature and are readily described in the SFB basis.  Further $\\ell$-mode mixing will occur for incomplete sky coverage, which can be corrected for using the mask geometry."
    },
    "1112/1112.3931_arXiv.txt": {
        "abstract": "We study the damping of the gravitational radiation-driven $f$-mode instability in rotating neutron stars by nonlinear bulk viscosity in the so-called supra-thermal regime. In this regime the dissipative action of bulk viscosity is known to be enhanced as a result of nonlinear contributions with respect to the oscillation amplitude. Our analysis of the $f$-mode instability is based on a time-domain code that evolves linear perturbations of rapidly rotating polytropic neutron star models. The extracted mode frequency and eigenfunctions are subsequently used in standard energy integrals for the gravitational wave growth and viscous damping. We find that nonlinear bulk viscosity has a moderate impact on the size of the $f$-mode instability window, becoming an important factor and saturating the mode's growth at a relatively large oscillation amplitude. We show similarly that nonlinear bulk viscosity leads to a rather high saturation amplitude even for the $r$-mode instability. In addition, we show that the action of bulk viscosity can be significantly mitigated by the presence of superfluidity in neutron star matter. Apart from revising the $f$-mode's instability window we provide results on the mode's gravitational wave detectability. Considering an $f$-mode-unstable neutron star located in the Virgo cluster and assuming a mode amplitude at the level allowed by bulk viscosity, we find that the emitted gravitational wave signal could be detectable by advanced ground-based detectors such as Advanced LIGO/Virgo and the Einstein Telescope. ",
        "introduction": "Neutron stars are exciting cosmic laboratories of matter at extreme conditions. However, unveiling the physics of these ultra-dense and relativistic systems is a real challenge even after more than four decades of astronomical observations. The situation is likely to improve significantly with the advent of gravitational wave astronomy. The forthcoming generation of improved gravitational-wave detectors, such as the Advanced LIGO and Virgo interferometers and the planned Einstein Telescope (ET),  should be sensitive to various aspects of neutron star dynamics. The observation of neutron star oscillations (normal modes), usually dubbed neutron star ``asteroseismology'', could become an excellent tool for probing the interior properties of these objects~\\citep{1998MNRAS.299.1059A}.  Of particular importance for neutron star asteroseismology are the oscillation modes that can become unstable (i.e. grow in amplitude) under the emission of gravitational radiation. This is the well-known CFS instability which develops when a retrogade mode (with respect to the stellar rotation) becomes prograde, as a result of rotational dragging, and then grows while emitting gravitational radiation~\\citep{1970PhRvL..24..611C,1978ApJ...222..281F}. Exhaustive work on the subject has shown that among the various stellar modes the ones that could become CFS-unstable in realistic neutrons stars, and likely to be strong sources of gravitational waves, are the inertial $r$-mode and the fundamental $f$-mode (see~\\citet{2003CQGra..20R.105A} for a review).     The study of dissipative mechanisms that could limit or entirely suppress the gravitational wave instability of the $r$- and $f$-modes has become a small industry since the early days of this field~\\citep{2001IJMPD..10..381A}. Given that astrophysical neutron stars are viscous systems, a clear understanding of the various dissipative mechanisms is a necessary ingredient of any realistic theoretical model. In fact, this is another facet of neutron star asteroseismology as the absence of unstable modes can be related to the strength of dissipation, thus providing indirect information about the stellar interior. For instance, this possibility has been recently illustrated by \\citet*{2011PhRvL.107j1101H}, who compared  the latest theoretical parameter space of the $r$-mode instability against the observed population of accreting  neutron stars in low mass X-ray binaries.    The present paper makes another contribution to our understanding of the $f$-mode instability and its damping. We focus on the action of bulk viscosity, which accounts for the energy drained from the oscillating fluid as a result of the induced departure from local beta equilibrium~\\citep{1989PhRvD..39.3804S, 1983bhwd.book.....S}. The novelty of our work is the study of bulk viscosity in the {\\em nonlinear} regime (with respect to the oscillation amplitude), thus extending previous models based on amplitude-independent viscosity~\\citep[e.g.][]{1991ApJ...373..213I,1995ApJ...438..265L, 2011PhRvL.107j1102G}.    The bulk viscosity formalism used in our analysis is based on the work by Alford and collaborators~\\citep{2010JPhG...37l5202A,2011AIPC.1343..580A,2011arXiv1103.3521A}. These authors have recently studied the role of nonlinear bulk viscosity in the context of the $r$-mode instability (we also briefly discuss the $r$-mode below in Sec~\\ref{sec:Rmode}). A key result in these papers is an analytical expression for the bulk viscosity coefficient in the nonlinear regime which can be easily incorporated in the standard machinery for calculating viscous mode damping. This latter procedure consists of solving the inviscid hydrodynamics equations and determining the mode frequency and eigenfunctions, and subsequently using these results in the available expressions for the mode's volume-integrated energy and damping rate (see section~\\ref{sec:Form}).   Our $f$-mode calculation is numerical and is based on a 2D time-domain code developed by~\\citet{2009MNRAS.394..730P}. The code evolves the linearised hydrodynamical equations of rapidly rotating stars in Newtonian theory. The frequency and eigenfunctions of specific modes can be then extracted from the time-evolved perturbations with a method developed by~\\citet{Stergioulas:2003ep} and~\\citet{Dimmelmeier:2005zk}. An overview and a sample test of the code is provided in section~\\ref{sec:NSmodel}. In this work we focus on the sextupole ($l=m=4$) $f$-mode, which is known to be the dominant unstable multipole both in Newtonian and Relativistic stars~\\citep{1983PhRvL..51...11F, 1991ApJ...373..213I,2011PhRvL.107j1102G}. We calculate the resulting instability ``window'' (that is, the unstable portion of the parameter space described by the stellar rotational frequency and temperature) after accounting for the action of shear and bulk viscosity (section~\\ref{sec:InsWin}). Our calculation provides results for the $f$-mode {\\em saturation} amplitude as determined by the balance between nonlinear viscous damping and gravitational radiation-driven growth (section~\\ref{sec:InsWin}). We also calculate the gravitational wave strain associated with an unstable $f$-mode and with its amplitude limited by nonlinear bulk viscosity (section~\\ref{sec:GW}). Our conclusions and a discussion of remaining issues can be found in section~\\ref{sec:concl}. ",
        "conclusions": "} \\label{sec:concl}  We have studied the gravitational wave-driven instability of the $f$-mode in Newtonian polytropic neutron star models. The novelty of our work is the inclusion of bulk viscosity in the supra-thermal regime. We have found that the ensuing nonlinear enhancement of bulk viscosity leads to a moderate change in the $f$-mode instability window. For the particular case of the $l=m=4$ $f$-mode (the most unstable multipole) we have shown that nonlinear bulk viscosity associated with the modified Urca process has a notable effect provided the mode amplitude, as defined in equation~(\\ref{eq:norm}),  is  $\\alpha \\gtrsim 0.1$ (see figure~\\ref{fig3}). This amplitude translates to a mode energy $E \\gtrsim 10^{-5} E_{\\rm rot}$ for $\\Omega \\approx \\Omega_{\\rm K}$, increasing to $E \\gtrsim 10^{-3} E_{\\rm rot}$ for $\\Omega \\approx \\Omega_{c}$. However, even for the extreme case of a mode amplitude as large as $\\alpha \\sim 1$ bulk viscosity cannot entirely suppress the instability.   If the more powerful direct Urca reactions can operate (which would be the case in the most massive neutron stars) then the resulting $f$-mode instability window is significantly reduced already in the sub-thermal bulk viscosity regime. The transition to the supra-thermal regime has again a moderate effect to the shape of the instability window. It turns out that in the temperature range where the change is important the instability is already likely to be suppressed by superfluid mutual friction (see figure~\\ref{fig4}).  The presence of superfluid components in neutron star matter could diminish the impact of bulk viscosity by suppressing the rates of the chemical reactions involved. By considering a temperature regime where the protons are superfluid and the neutrons normal and by treating in an approximate manner the superfluid suppression, we have shown that nonlinear bulk viscosity is significantly weakened  (see figure~\\ref{fig5}).   Apart from calculating the bulk viscosity-modified $f$-mode instability window we have also provided upper limits for the gravitational wave signal from a mode with $\\alpha = 0.1-1$, i.e. the amplitude range where nonlinear bulk viscosity can partially saturate the mode (figure~\\ref{fig6}). Our results suggest that the $l=m=4$ $f$-mode instability (operating in a neutron star at the fiducial distance of 20~Mpc) could be detectable by a third generation instrument like the ET for an amplitude $\\alpha \\ge 0.1$ and an observation time of $\\sim 1$ week. Detectability by Advanced LIGO/Virgo requires a rather high amplitude ($\\alpha \\sim 1$) and/or a longer observation time (e.g. the evolution timescale (\\ref{eq:tausd})). Moreover, the $f$-mode  detectability is moderately sensitive to the stellar polytropic index, with lower indices (stiffer equation of state) being favoured. It should be pointed out that our work provides only `static' upper limits, in the sense that the combined rotational/temperature evolution of an $f$-mode unstable neutron star has not been properly incorporated. A more detailed analysis of this issue will the subject of future work.  Our brief excursion to the study of the damping of the $r$-mode instability by nonlinear bulk viscosity has revealed that a large mode amplitude ($\\alpha > 1$, see equation~(\\ref{eq:aAl})) is required for the effect to be relevant. Thus, most likely, the saturation of the $r$-mode is controlled by nonlinear mode couplings rather than bulk viscosity. A similar statement could also be true for the $f$-mode but this remains to be (dis)proved, given the lack of any concrete results on the $f$-mode nonlinear couplings. Some light on this issue can be shed by a recent study of the nonlinear dynamics of the $l=2$ $f$-mode~\\citep{2010PhRvD..82j4036K}. Extrapolation of  those results to the $l=4$ mode suggests that nonlinear hydrodynamics could outperform bulk viscosity in saturating a growing mode.   Besides the obvious need for a better understanding of the $f$-mode  nonlinear dynamics, several other improvements are required in this research area. For instance, our own analysis can be clearly improved by moving from Newtonian to General Relativistic gravity and from polytropic matter to more realistic equations of state. Indeed, as suggested by recent work on the relativistic $f$-mode, the instability's growth time may shrink as much as a factor $\\sim 10$ compared to the Newtonian result. However, as we pointed out in section~\\ref{sec:GW}, this effect is likely to be counter-balanced by a similar increase in the bulk viscosity strength as a result of using more realistic equations of state. The presence of a solid crust adds one more level of complexity to the problem. Despite its central role in limiting the $r$-mode instability~\\citep{2001IJMPD..10..381A}, the dynamics of the crust has been totally ignored in the context of the $f$-mode instability. Another interesting possibility arises when exotic matter (for example, deconfined quarks) is present in the stellar interior. Previous work on the $r$-mode instability in neutron stars with exotic cores has unveiled several important differences with respect to normal hadronic neutron stars~\\citep[e.g.][]{1992PhRvD..46.3290M,2002MNRAS.337.1224A}. This is likely to be the case also for the $f$-mode instability. Clearly, all the above issues are of great interest and we plan to address some of them in the near future."
    },
    "1112/1112.1619_arXiv.txt": {
        "abstract": "With the evolving technology in CMOS integration, new classes of 2D-imaging detectors have recently become available. In particular, single photon avalanche diode (SPAD) arrays allow detection of single photons at high acquisition rates ($\\ge\\unit[100]{kfps}$), which is about two orders of magnitude higher than with currently available cameras. Here we demonstrate the use of a SPAD array for imaging fluorescence correlation spectroscopy (imFCS), a tool to create 2D maps of the dynamics of fluorescent molecules inside living cells. Time-dependent fluorescence fluctuations, due to fluorophores entering and leaving the observed pixels, are evaluated by means of autocorrelation analysis. The \\multitau correlation algorithm is an appropriate choice, as it does not rely on the full data set to be held in memory. Thus, this algorithm can be efficiently implemented in custom logic. We describe a new implementation for massively parallel \\multitau correlation hardware. Our current implementation can calculate 1024 correlation functions at a resolution of $\\unit[10]{\\mu s}$ in real-time and therefore correlate real-time image streams from high speed single photon cameras with thousands of pixels. ",
        "introduction": "Fluorescence correlation spectroscopy\\cite{MAGDE1974,MAGDE1974a} (FCS) is a powerful experimental technique for measuring the dynamics of fluorescently labeled molecules in solution and also inside living cells. It allows one to determine the particle number, the diffusion coefficient, flow speeds and also photophysical and chemical reaction rates (for an overview, see Ref.~\\onlinecite{KRICHEVS2002}). In FCS the time trace of the fluorescence intensity fluctuations $I(t)$ inside a small observation volume (usually around $\\unit[10^{-15}]{l}=1\\unit[]{\\mu m^3}$) is monitored. The fluctuations originate from particles entering and leaving the focus, or transitions between states having different quantum yields. Faster dynamics of the fluorescing particles also lead to faster fluctuations, which can be quantified by means of a temporal first-order autocorrelation function (ACF): \\begin{equation}\\label{eq:acf1} g(\\tau)=\\frac{\\mean{I(t)\\cdot I(t+\\tau)}_t}{\\mean{I(t)}_t^2},\\ \\ \\ \\mean{I(t)}_t:=\\lim\\limits_{\\tilde{T}\\rightarrow\\infty}\\frac{1}{\\tilde{T}}\\int_0^{\\tilde{T}}I(t)\\;\\mathrm{d}t \\end{equation} The ACF usually contains features that are spread over several orders of magnitude in time (nanoseconds to seconds).  The standard FCS setup uses a confocal microscope in combination with single-photon sensitive detectors to acquire the fluorescence time trace $I(t)$ from one focal volume. Then the data are fed into a ``correlator'' (hardware or software component), which estimates the ACF over a certain dynamic range.  In the accompanying paper by \\citeauthor{mocsar2012}\\cite{mocsar2012}, a field programmable gate array (FPGA) implementation of such a correlator circuit for use with a confocal microscope and up to two single photon avalanche diode (SPAD) detectors is presented.  In recent years, the availability of fast cameras has triggered the development of different imaging modalities for FCS\\cite{WOHLAND2007,DROSS2009a,WOHLAND2010}, which allow spatial mapping of the diffusion coefficient and other dynamic properties by calculating an ACF for each of potentially many pixels. This requires correlation hardware that can process the data stream very fast, ideally in real time, for all pixels of the image sensor. Here we extend the idea of hardware reuse, presented before\\cite{mocsar2012}, for application to imaging FCS.  Our FPGA-based implementation can calculate 1024 autocorrelation functions in parallel and in real time. The dynamic range of the ACFs is $\\tau_{\\text{min}}\\ldots\\tau_{\\text{max}}=\\unit[10]{\\mu s}\\ldots\\unit[1]{s}$. As a data acquisition device we use the single photon avalanche diode array (SPAD array) \\mbox{\\textit{Radhard2}} % \\cite{note1,carrara2009gamma}. From our sensor we read a $32\\times32$ pixel frame every $\\Delta t_{\\text{frame}}=\\unit[10]{\\mu s}$ where each pixel contains $\\unit[1]{bit}$ of information (no photon or at least one photon in the last $\\Delta t_{\\text{frame}}$). The design presented here could easily be adapted to other image sensors such as customized complementary metal oxide semiconductor (CMOS) or electron multiplying charge-coupled device (EMCCD) cameras. ",
        "conclusions": "In this paper we presented the implementation of an FPGA-based \\mtau correlator design that can calculate $1024$ correlation functions in real time at a minimum lag time of $\\unit[10]{\\mu s}$. To our knowledge this is the largest number of real-time \\mtau correlators implemented so far in a single device. The minimum lag time of $\\unit[10]{\\mu s}$ in our design is considerably longer than that of currently available hardware correlators (e.g. from ALV GmbH, Langen, Germany or correlator.com, Bridgewater, USA and \\lcite{MAGATTI2003}). However, those are limited to at most $32$ auto-correlators.  We use our design to correlate the output of a single-photon avalanche diode array used as image sensor. This combination will allow us to perform imaging fluorescence correlation spectroscopy at sufficiently fast time scales to resolve even the motion of small molecules in solution and living cells.  Our design is flexible, so beside estimating ACFs, temporal cross-correlation functions of different pixels\\cite{DERTINGE2007,Colyer2010} or multiple colors\\cite{Bestvater2010} (using spectrally resolved detectors) can also be calculated."
    },
    "1112/1112.1569_arXiv.txt": {
        "abstract": "Using \\chandra High Energy Grating (HEG) observations of 32 AGNs,  we present a systematic study of the X-ray Baldwin effect (XBE, i.e. the anti-correlation between narrow \\fekalfa\\ line EW and X-ray continuum luminosity for AGN samples) with the highest spectral resolution currently available. We have previously reported an anti-correlation with $EW~\\propto~L_{\\rm 2-10~keV}^{-0.22}$ in a HEG sample, and the correlation is much weaker after averaging multiple observations of individual AGNs ($EW~\\propto~L^{-0.13}_{\\rm 2-10~keV}$). This indicates that rapid variation in X-ray continuum plays an important role in producing the XBE, and such an effect should also be visible in individual AGNs. In this paper, by normalizing the line EWs and continuum luminosities to the time-averaged values for each AGN in our sample with multiple HEG observations, we find a strong anti-correlation between EW and L$_X$ (EW/$\\langle EW \\rangle \\propto (L/\\langle L \\rangle)^{-0.82\\pm0.10}$), consistent with the XBE expected in an individual AGN if the narrow line flux remains constant while the continuum varies. This is first observational evidence that the \\fekalfa line flux in a large sample of AGNs lacks of a corresponding response to the continuum variation, supporting that the narrow Fe-K line emission originates from a region far from the nucleus. We then performed Monte-Carlo simulations to address whether the global XBE can be produced by X-ray continuum variation solely, and found that such interpretation of XBE cannot be ruled out statistically. One thus should be very cautious before reaching any scientific conclusion based on an observed XBE. ",
        "introduction": "\\label{intro} The narrow (FWHM~$<10,000 \\ \\rm km \\ s^{-1}$) \\fekalfa fluorescent emission at $\\sim 6.4$~keV is a common feature in the X-ray spectra of active galactic nuclei (AGN) (e.g. Sulentic \\etal 1998; Lubi\\'{n}ski \\& Zdziarski 2001; Weaver, Gelbord, \\& Yaqoob 2001; Perola \\etal 2002; Yaqoob \\& Padmanabhan 2004; Bianchi et al. 2007; Winter \\etal 2009; Shu, Yaqoob \\& Wang 2010 (hereafter Paper I); Fukazawa et al. 2011). Such a narrow \\fekalfa line is believed to be produced in cold, neutral matter far from the nucleus (see Paper I, and references therein). Possible origins of the \\fekalfa line include the outer regions of an accretion disk, the broad-line region (BLR), and a parsec-scale torus (Nandra 2006; Bianchi et al. 2008; Paper I; Liu \\& Wang 2010; Shu, Yaqoob \\& Wang 2011; Jiang, Wang \\& Shu 2011). In some cases, the line has been resolved by the \\chandra high-energy grating (HEG; see Markert et al. 1995) with FWHM typically less than 5000 $\\rm km \\ s^{-1}$ (e.g. Kaspi et al. 2000; Yaqoob et al. 2001; Paper I; Shu et al. 2011). In Paper I, we measured the intrinsic width of the narrow \\fekalfa line core in a large sample of type I AGNs using the \\chandra HEG, and obtained a weighted mean of FWHM=$2060\\pm230$ km s$^{-1}$. Interestingly, this width of the line is fully consistent with that obtained in type II AGNs observed with \\chandra HEG (Shu et al. 2011), indicating a common origin of the narrow \\fekalfa line in both AGN populations.  The equivalent width (EW) of the \\fekalfa line was found to be anti-correlated with the X-ray continuum luminosity (also known as X-ray Baldwin effect) in several AGN samples (e.g., Iwasawa \\& Taniguchi 1993; Nandra et al. 1997b; Page et al. 2004; Jiang, Wang, \\& Wang 2006; Bianchi et al. 2007; Winter et al. 2009; Chaudhary et al 2010). Note the observed X-ray Baldwin effect in some samples could be (at least) partly due to the contamination of radio-loud objects (Jimen$\\acute{\\rm e}$z-Bail$\\acute{\\rm o}$n et al. 2005; Jiang et al. 2006). It is also interesting to note a recent work by Krumpe et al. (2010) who found that the X-ray Baldwin effect flattens above $L_X\\sim10^{44}$ erg s$^{-1}$, claiming that the anti-correlation (if exists) is only relevant in radio-quiet AGNs with $L_X<10^{44}$ erg s$^{-1}$.   There are several physical factors that could produce the X-ray Baldwin effect. One possibility is that  the covering factor and/or the column density of line-emitting material decreases with increasing X-ray continuum luminosity (e.g., Page et al. 2004). Alternatively, the X-ray Baldwin effect could be artificial due to the rapid X-ray continuum variation in AGNs. Jiang et al. (2006) reported that X-ray continuum variation model can naturally produce a global X-ray Baldwin effect in their composite Chandra/XMM sample. In Paper I, with Chandra HEG data, we also found that averaging multiple observations of individual sources could significantly smear the X-ray Baldwin effect from $EW~\\propto~L^{-0.22}$ to $EW~\\propto~L^{-0.13}$. This clearly indicates that the observed X-ray Baldwin effect could at least partially attributed to continuum variation.   Another possibility is contamination from the Fe K$\\alpha$ line emission from the accretion disk,  which could be more and more ionized as the X-ray luminosity increases, leaving less low-ionization material to produce the Fe K$\\alpha$ line at $\\sim$ 6.4 keV (e.g., Nandra et al. 1997b; Guainazzi et al. 2006; Nayakshin 2000a, 2000b). Therefore, it is essential to study the X-ray Baldwin effect for narrow Fe K line emission with the highest spectral resolution available, to isolate the narrow cores from the broad component.  Using a sample that consists only of HEG data\\footnote{\\chandra HEG affords best spectral resolution currently available in the Fe K band, at 6.4 keV is $\\sim$39 eV, or $\\sim$1860 km s$^{-1}$ FWHM.}, we can investigate the X-ray Baldwin effect with a spectral resolution in the Fe K band that is nearly four times better than in previous studies that relied on CCD data (or data with a resolution worse than CCDs), and therefore provide the best isolation of the narrow core that is currently possible. In this paper, we extend the results of Paper I to investigate the origin of the Baldwin effect, in particular taking into consideration of the effect of X-ray continuum variability. Throughout this paper, we adopt a cosmology of $\\Omega_M=0.3, \\Omega_{\\lambda}=0.7$, and H$_0=70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "In this paper, we presented a systematic study on the narrow \\fekalfa line Baldwin effect in AGNs observed with the \\chandra high energy grating. By normalizing the line EWs and continuum luminosities to the time-averaged values for AGNs with multiple HEG observations, we found a strong anti-correlation between both quantities with a best-fitting slope of $-0.82\\pm0.10$, statistically consistent with a slope of $-1$. This result indicates generally that the flux of the \\fekalfa line of an individual AGN in \\chandra grating sample is nearly constant and lacks a corresponding response to the continuum variation. We simulated the X-ray continuum variation to test the possibility whether such observational bias could totally account for the observed X-ray Baldwin effect. We find that the simulations could yield significant anti-correlation between EW and continuum luminosity, and the strength of the simulated Baldwin effect is sensitive to the X-ray variation model and variation amplitude. Based on our HEG sample and current knowledge of AGN's X-ray variation, the possibility that the observed X-ray Baldwin effect is totally an observational bias due to X-ray variation cannot be ruled out. This bias should be taken into account before any scientific conclusion can be made based on an observed X-ray Baldwin effect. One should endeavor to make as many repeated observations of each AGN as possible, covering the relevant timescales, and averaging the results. This could minimize the bias and help us to reveal the underlying intrinsic X-ray Baldwin effect, if there is any.      X.W.S. thanks the support from China postdoctoral foundation. We acknowledge support from Chinese National Science Foundation (Grant No. 10825312, 11103017), and the Fundamental Research Funds for the Central Universities (Grant No. WK2030220004, WK2030220005). This research made use of the HEASARC online data archive services, supported by NASA/GSFC. The authors are grateful to the \\chandra instrument and operations teams for making these observations possible."
    },
    "1112/1112.2659.txt": {
        "abstract": "If the Standard Model is valid up to very high energies it is known that the Higgs potential can develop a local minimum at field values around $10^{15}-10^{17}$ GeV, for a narrow band of values of the top quark and Higgs masses. We show that in a scalar-tensor theory of gravity such Higgs false vacuum can give rise to viable inflation if the potential barrier is very shallow, allowing for tunneling and relaxation into the electroweak scale true vacuum. The amplitude  of cosmological density perturbations from inflation is directly linked to the value of the Higgs potential at the false minimum. Requiring the top quark mass, the amplitude and spectral index of density perturbations to be compatible with observations, selects a narrow range of values for the Higgs mass, $m_H=126.0\\pm 3.5$ GeV, where the error is mostly due to the theoretical uncertainty of the 2-loop RGE. This prediction could be soon tested at the Large Hadron Collider. Our inflationary scenario could also be further checked by better constraining the spectral index and the tensor-to-scalar ratio. ",
        "introduction": "Despite the many experimental successes of the Standard Model (SM) of particle physics, its scalar sector is still to be confirmed by accelerator experiments. Extensions of the SM are intensively explored, motivated by the fact that the SM does not explain neutrino masses, dark matter, the baryon-antibaryon asymmetry, a primordial stage of inflation, etc. From a theoretical  point of view,  useful guides for extending the SM are represented by a possible unification of gauge couplings, by fine-tuning problems (such as the lightness of the Higgs mass) and by the inclusion of gravity.    We consider here the possibility that no major extensions of the SM Higgs sector are required. In this framework, we consider the possibility of realizing inflation by introducing only modifications in the gravitational sector.  The idea that inflation is realized in the SM Higgs sector has been proposed recently ~\\cite{arXiv:0710.3755}, by allowing for a non-minimal coupling of the Higgs to gravity. This is an appealing point of view because it does not introduce new degrees of freedom. It is however controversial whether such scenario is viable at the quantum level, because radiative corrections could modify the potential \\cite{arXiv:0809.2104,arXiv:0812.4946}.  In this paper we propose a slightly less minimalistic scenario, in which we {\\it do} introduce a new scalar degree of freedom, but assume it to be decoupled from the SM: in such a way the SM is not altered at all and only gravitational physics is affected.  We use the fact that the Higgs potential has a local minimum between $10^{15}-10^{17}$ GeV, which is indeed known to exist for a narrow band of the top and Higgs mass values \\cite{CERN-TH-2683,hep-ph/0104016,strumia2}. We will assume that the Universe started with the Higgs in this false vacuum, which leads to a stage of exponential inflation, where density perturbations are produced.  It is already a non-trivial fact that such a local minimum can exist in the Higgs potential, but is even more non-trivial that there exist an allowed parameter range in the Higgs and top masses, for which the energy density of this minimum has the right value (GUT scale) to give rise to the correct amplitude of density perturbations from inflation.   The additional scalar has the  role of providing a {\\it graceful exit} to inflation, which is possible in a  Brans-Dicke scalar-tensor theory of gravity, as shown in ref.~\\cite{astro-ph/0511396,hep-ph/0511207} (and earlier in ~\\cite{johri,extended,hyperextended} for power-law inflation). After the stage of exponential inflation,  the expansion is drastically slowed down by the Brans-Dicke scalar and the Higgs field can tunnel out efficiently, via production and collision of bubbles, and subsequently rolls down to its present vacuum expectation value, $v \\sim 246$ GeV. Moreover quantum fluctuations in the Brans-Dicke scalar generate a spectrum of density perturbations, which can be understood also as a tunneling which happens at slightly different times in different regions of space. In particular, we show that successful inflation can be achieved for a very broad class of couplings of the scalar-tensor theory and that the spectral index of density perturbations is largely independent on the particular choice of couplings.  We use the information from the amplitude and spectral index of cosmological density perturbations generated during inflation to predict the Higgs mass, finding \\be m_H= (126.0 \\pm 3.5) \\,{\\rm GeV} \\,\\,. \\ee where the error is mostly due to the theoretical uncertainty of the 2-loop RGE. This range is within the present experimental window set by direct searches, namely the $115$ GeV lower bound set by LEP \\cite{PDG} and the $140$ GeV  upper bound set recently by LHC \\cite{HCP11}, which restricted the previous $155$ GeV bound set by Tevatron \\cite{arXiv:1107.5518}. We note that the range above is also compatible with the July 2011 global electroweak SM precision fit \\cite{GFitter}, which gave $m_H=125^{+8}_{-10} $ GeV at $1\\sigma$.  After the first release of this paper on the arXiv, LHC updated the Higgs boson exclusion limits \\cite{last}: as of 13 December 2011 ATLAS excludes at the $95 \\%$ CL masses outside $116-130$ GeV and CMS excludes at the $95\\%$ CL masses outside $115-127$ GeV. Remarkably, the allowed band for SM false vacuum inflation includes the region $124-127$ GeV, where ATLAS and CMS recorded excesses of events in the di-photon as well as 4-leptons channels \\cite{last}. At present, the significance of this result is however still low.  Hopefully, our inflationary model prediction will be soon better tested by LHC. In the case of compatibility with LHC results, future more precise cosmological measurements of the scalar spectral index and the tensor-to-scalar ratio could provide further tests for the idea of Higgs false vacuum inflation \\cite{Masina:2011un}.    The paper is organized as follows. In section \\ref{Model} we present our inflationary model, based on a scalar-tensor theory of gravity. First, we derive the expression for the slow roll parameters, the scalar spectral index and the tensor-to-scalar ratio. We then establish the relation between the amplitude of density perturbations and the Higgs potential at the false minimum. In section \\ref{Minimum} we study values of the top and Higgs masses giving a SM false vacuum within the range consistent with inflation. In section \\ref{Results} we show that the set of values derived in this way are only in part consistent with the measurement of $m_t$. This allows to predict a narrow range for the Higgs mass.  In section \\ref{Post} we provide details about the post-inflationary period, while in section \\ref{General} we generalize the model by considering higher order terms. We finally draw our conclusions in section \\ref{Conclusions}.  %%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{Conclusions}  We proposed to exploit the possible presence of a false minimum in the SM Higgs potential at very high-energies, $10^{15}-10^{17}$ GeV, to provide a large amount of potential energy which can drive primordial inflation. In the framework of scalar-tensor theories of gravity, a graceful exit from inflation can be achieved through tunneling of the Higgs field and subsequent relaxation down to its present vacuum expectation value $v\\sim 246$ GeV.  Requiring the amplitude and spectral index of cosmological density perturbations from inflation and the top quark mass to be compatible with observations, we showed that this possibility is realized only within a small region of values for the Higgs mass, see fig.\\ref{fig-mtmh}, leading to the prediction $$ m_H = (126.0 \\pm 3.5)\\,  {\\rm GeV}\\,\\,, $$ where the error is mostly due to the present theoretical uncertainty of the 2-loop RGE. This prediction can be tested soon by LHC, in particular for the decay mode $H \\rightarrow \\gamma \\gamma $. The inflationary model proposed here could thus meet experimental support or be ruled out. It is exciting that preliminary results from LHC \\cite{last} show an excess of events in the range $124-127$ GeV.  If these results will be confirmed, further checks of the model will be offered by better determinations of the scalar spectral index $n_S$, predicted in this model to be within $0.93-0.96$. Using the constraints coming from the top quark mass, the scalar-to-tensor ratio $r$ in this model is constrained to be within $ 10^{-4}\\lesssim r \\lesssim10^{-1}$. Actually, it can be shown \\cite{Masina:2011un} that any model in which the false vacuum is very shallow, the relation between the potential at the false minimum and the amplitude of perturbations, eq. (\\ref{eq-Dr}), implies $r \\gtrsim 10^{-4}$ . With forthcoming more precise cosmological measurements, such as the Planck satellite mission, one can test the region of large values of $r$, while improving the top quark mass measurement can further constrain $r$ from below.  Moreover, we showed that our scenario, and in particular the prediction for $n_S$, can be obtained in a wide class of scalar-tensor theories. In particular, higher dimensional operators can be safely present.  As a completely general remark, we point out that discovering a Higgs with mass close to $126$ GeV is a very suggestive hint in favor of the existence of a {\\it false minimum} in the SM Higgs potential at energies close to $10^{16}$ GeV, which could be {\\it the starting point for inflation in our Universe.} Indeed, it would lead to a period of exponential expansion producing  density perturbations with the right amplitude. We have shown here that in a scalar-tensor theory of gravity this inflationary stage can end, allowing the Higgs field to tunnel out of the false minimum and subsequently to relax down to its present electroweak scale value.     %%%%%%%%%%%%%%%%%%%%%%%%% \\vskip .7cm {\\bf Acknowledgements} We thank Tirthabir Biswas for discussions in the early stages of this project and Pietro Dalpiaz for interesting suggestions. We also thank Jaume Garriga and Giovanni Villadoro for useful comments.     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1112/1112.1796_arXiv.txt": {
        "abstract": "We present the results of a 500\\,ksec long \\textsl{XMM-Newton} observation and a 120\\,ksec long quasi-simultaneous \\textsl{Chandra} observation of the Narrow Line Seyfert 1 galaxy 1H0707$-$495 performed in 2010 September. Consistent with earlier results by \\citet{Fabian2009a} and \\citet{Zoghbi2010a}, the spectrum is found to be dominated by relativistically broadened reflection features from an ionised accretion disc around a maximally rotating black hole. Even though the spectra changed between this observation and earlier \\textsl{XMM-Newton} observations, the physical parameters of the black hole and accretion disc (i.e., spin and inclination) are consistent between both observations. We show that this reflection spectrum is slightly modified by absorption in a mildly relativistic, highly ionised outflow which changed velocity from around $0.11\\,c$ to $0.18\\,c$ between 2008 January and 2010 September. Alternative models, in which the spectral shape is dominated by absorption, lead to spectral fits of similar quality, however, the parameters inferred for the putative absorber are unphysical. ",
        "introduction": "The Narrow Line Seyfert 1 galaxy 1H0707$-$495 ($z=0.04057$, \\citealt{Jones2009aMNRAS}) is most famous for its huge and sharp drop in flux at $\\sim7\\,$keV. Discovered by \\citet{Boller2002a}, this feature was initially interpreted as an absorption edge despite the lack of a fluorescent emission line. Subsequent models for the feature included partial covering by a thick absorber \\citep{Tanaka2004a}, which also explains the temporal changes seen in its location \\citep{Gallo2004a}. In a long \\textsl{XMM-Newton} observation, however, \\citet{Fabian2009a} and \\citet{Zoghbi2010a} showed that relativistically broadened fluorescent iron K$\\alpha$ emission is a more likely explanation, especially given that these authors could also show that the source's soft-excess can be modelled by a relativistic iron L$\\alpha$ line, the first one ever observed. That the position of the two lines and the ratio of the normalisations agree with expectations adds further evidence to the reflection scenario \\citep{Fabian2009a}. In order to produce such a strong feature and an observable Fe L$\\alpha$ line, the iron abundance in 1H0707$-$495 has to be several times the Solar composition. The new observations also ruled out the partial covering explanation for the feature, as the strong absorption lines and edges predicted by these models are not compatible with high resolution X-ray spectra \\citep{Zoghbi2010a}. Recently, \\citet{Fabian2012a} observed 1H0707$-$495 at low flux, confirming the reflection dominated interpretation. The same physical interpretation also successfully describes the spectra of similar AGN, e.g., IRAS13224$-$3809 \\citep{Boller2003a,Ponti2010a}, too.   In this \\emph{paper} we analyse data of 1H0707$-$495 taken by the \\textsl{XMM-Newton} satellite from 2010 September 12 until 2010 September 19. The results are compared to a previous observation \\citep{Fabian2009a,Zoghbi2010a} with similar exposure and flux. Section~\\ref{sec:data-analysis} is dedicated to present the data reduction we applied and general properties of the observation are discussed. The outcome of our spectral modelling is presented in Sect.~\\ref{sec:x-ray-spectrum} and finally the results are discussed in Sect.~\\ref{sec:results}. ",
        "conclusions": "\\label{sec:results}  In this paper we have presented the results from a spectral analysis of 500\\,ksec of new \\textsl{XMM-Newton} and \\textsl{Chandra} data from the AGN 1H0707$-$495 and have compared these results to the previous long observation of the source. In agreement with earlier work \\citep{Fabian2009a,Fabian2012a,Zoghbi2010a}, the simple picture deduced from these observations is that 1H0707$-$495 an accretion disc around a maximally rotating black hole is irradiated by a central and compact source closely above the black hole, producing the underlying complex X-ray continuum by reflection. In both observations, the primary X-ray source which irradiates the accretion disc has a height of $\\le 4\\,r_\\mathrm{g}$ above the black hole. Such a low emitter is in line with timing measurements \\citep{Zoghbi2010a} and implies that the photons are extremely focused onto the inner parts of the disc. Fits with an empirical broken power law emissivity confirm this interpretation ($\\epsilon_{r<r_\\mathrm{br}} = 10.2^{+0.8}_{-0.9}$, $\\epsilon_{r>r_\\mathrm{br}} = 2.3^{+0.5}_{-0.1}$, $r_\\mathrm{br} = 2.9^{+0.2}_{-1.0}\\,r_\\mathrm{g}$ for \\rn). Despite significant changes in the continuum shape between both observations, the black hole in 1H0707$-$495 is consistently well determined to be maximally rotating. This result also adds further credibility to the relativistic line interpretation of the spectrum, as the black hole parameters are not expected to change on such short time scales. Such a behavior would be significantly more difficult to explain in alternative models explaining the line via an ionised and blue-shifted radial wind.  The accretion disc causing the observed reflection features has a complex ionisation structure, which is here approximated by two reflectors of strongly different ionization as modeled by \\citet{Ross2005a} and \\citet{ross2007a}. Even though the highly ionised reflector dominates the soft X-ray spectrum, its normalisation is only 0.1--0.2\\% of that of the weakly ionised reflector, i.e., most of the irradiated X-rays are intercepted by the colder medium. This result fits nicely into the picture of a rather neutral accretion disc with a thin skin of highly ionised material caused by the incident X-rays \\cite[see, e.g.,][]{garcia:10a,garcia:11a}. As a caveat, however, note that our approach implicitly assumes that the ionisation fraction of the disc is independent of distance from the black hole. From basic arguments it is clear that the photoionisation due to the incident X-rays will result in an ionisation structure of the disc surface which strongly depends on radius, as even in the simplest and most conservative models photons are strongly focused towards the inner parts. Additionally, the energy release within the accretion disc itself will give rise to a temperature and ionisation gradient \\citep[e.g.,][and references therein]{hubeny:01a,davis:05a}.  As shown by our modeling of the weaker spectral features, the relativistically blurred continuum is then modified by absorption in an ultra-fast wind ($0.11$--$0.18\\,c$, i.e., 30000--50000\\,$\\mathrm{km}\\,\\mathrm{s}^{-1}$). Due to the change of the minor features at 2--5\\,keV and the constancy of the 7\\,keV drop between the two observations (see \\ref{sec:highly-ioniz-outfl}), the ionised wind could be uniquely identified to exist along with the reflection, which dominates the spectrum. The overall properties of this wind are in line with the relativistic, highly ionized winds that have now been detected in more than 40 radio-quiet AGN, including Narrow Line Seyfert 1 galaxies. These winds have been mainly detected through strongly redshifted K$\\alpha$ and K$\\beta$ absorption lines from H- and He-like iron. See, e.g., \\citet{chartas:02a}, \\citet{pounds:03a}, \\citet{turner:04a}, or \\citet{cappi:09a} for discussions of individual sources and \\citet{Tombesi2010a} for a recent comprehensive study with \\textsl{XMM-Newton}.  Radiation hydrodynamical calculations show that such winds can in principle be formed as line driven winds from an accretion disc \\citep[][and references therein]{proga:00a,kurosawa:09a}. The spectral signature imprinted on an X-ray continuum are in rough agreement with the features seen here \\citep[e.g.,][]{schurch:09a}, although further theoretical work such as proper inclusion of Compton broadening is clearly needed \\citep[e.g.,][]{sim:10a}. Unfortunately, the current wind models available for X-ray spectral modeling also do not yet allow us to self-consistently model absorption with abundances consistent with the significant overabundance in the accreted material inferred from the X-ray reflection. Despite this problem, however, the agreement between our simple wind model and the data is remarkable. In passing, we note that the measured outflow velocity is also consistent with the unified model for quasars proposed by \\citet{Elvis2000a}, where a ``Warm Highly Ionised Medium'' is ejected at speeds of 10000--$60000\\,\\mathrm{km}\\,\\mathrm{s}^{-1}$ at an inclination of around $60^\\circ$. Coincidentally, this value is in agreement with our best-fit inclination angle. Note that in order to explain the apparent shift of minor spectral features in the 2--5\\,keV band, the line of sight velocity of this wind must have changed between both observations. As such a wind is expected to be highly structured \\citep{sim:10a}, only very slight changes in the line of sight would be required to explain the observed change in velocity.  Despite the overall success of the modeling, however, some significant broad residuals remain (Fig.~\\ref{fig:spectra}d). There are several major issues which could explain these discrepancies that could not be treated properly in our analysis due to limitations of our best-fit model (Sect.~\\ref{sec:basic-model}). First of all, while measurements of the reflection show a high Fe abundance in reflection, all other elements are assumed to be of Solar abundance. From basic arguments it is clear that a high Fe abundance likely implies that other elements are over-abundant as well, which could not be taken into account in our approach due to the limitations of the \\texttt{reflionx} model. Such a restriction might therefore underpredict emission from other elements, while at the same time over-predict the Fe abundance. Improving on modeling the ionisation gradient in the accretion disc as discussed above could also slightly reduce the artificially high Fe abundance obtained from fitting the data with simple reflection models \\citep{Reynolds1995a}. As already discussed by \\citet{Zoghbi2010a}, however, NLS1 galaxies such as 1H0707$-$495 exhibit enhanced star formation \\citep{Sani2010a} and therefore are expected to be Fe enriched. This assumption is confirmed in near-IR measurements \\citep[see, e.g., ][]{Shemmer2002}. Moreover, in similar sources like MCG$-$6-30-15 \\citep{Miniutti2007amnras}, 1H0419$-$577 \\citep{Fabian2005a}, or IRAS13224$-$3809 \\citep{Ponti2010a} iron is also required to be over-abundant.  Regarding the ionised wind seen in absorption, we measure a lower wind velocity and $N_\\mathrm{H}$ compared to earlier absorption dominated models \\citep{Gallo2004a,done2007a}. This difference is probably due to the fact that in earlier models the full continuum was seen to be dominated by the wind, while in the present model the soft-excess is mainly described by reflection of a highly ionised accretion disc. This interpretation is in line with our identification of weaker spectral features as absorption lines from H-like ions (Si, S, Ca, and possibly Ar), which are found at the correct energies expected from the inferred wind velocites. As discussed above, since the 7\\,keV feature is non variable, it cannot be due to the wind. This result significantly simplifies the wind modeling, as in order to explain the whole 7\\,keV feature as a wind a very complex absorber is necessary \\citep[see also][note, however, that the ionised and blue-shifted absorption is required to describe the narrow absorption feature right above the drop]{done2007a}.  \\emph{Acknowledgements.} This work is based on observations obtained with \\textsl{XMM-Newton}, an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA. JW and TD acknowledge partial support from the European Commission under contract ITN 215212 ``Black Hole Universe'' and TD by a fellowship from the Elitenetzwerk Bayern. JS acknowledges support from the Grant Agency of Czech Republic (GACR 202/09/0772). We thank John Davis for the development of the \\texttt{SLxfig} module used to prepare the figures in this paper, and Alex Markowitz and Moritz B\\\"ock for their many useful discussions."
    },
    "1112/1112.0335_arXiv.txt": {
        "abstract": "We model neutrino emission from a newly born neutron star subsequent to a supernova explosion to study its sensitivity to the equation of state, neutrino opacities, and convective instabilities at high baryon density.  We find the time period and spatial extent over which convection operates is sensitive to the behavior of the nuclear symmetry energy at and above nuclear density.  When convection ends within the proto-neutron star, there is a break in the predicted neutrino emission that may be clearly observable. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.0502.txt": {
        "abstract": "We first study dark energy models with a minimally--coupled scalar field and exponential potentials, admitting exact solutions for the cosmological equations: actually, it turns out that for this class of potentials the Einstein field equation exhibit alternative Lagrangians, and are completely integrable and separable (i.e. it is possible to integrate the system analytically, at least by quadratures). We analyze such solutions, especially discussing when they are compatible with a late time quintessential expansion of the universe. As a further issue, we discuss how such quintessential scalar fields can be connected to the inflationary phase, building up, for this class of potentials, a quintessential inflationary scenario: actually, it turns out that the transition from inflation toward late--time exponential quintessential tail admits a kination period, which is an indispensable ingredient of this kind of theoretical models. All such considerations have also been done including radiation into the model. ",
        "introduction": "Scalar fields have been used extensively in cosmology since inflationary theories were first conceived in 1981 \\cite{guth}. They are in fact useful to fill in the universe with a negative--pressure content, hence giving the typical inflationary behavior, with its exponentially increasing scale factor $a(t) \\sim \\exp(H_0t)$, where $H_0$ is the present value of the Hubble parameter. Scalar fields have been used both in minimal and non minimal coupling with geometry. (An overview of scalar--tensor theories is given in Refs. \\cite{dam} \\cite{de1}; a more updated review can be found in Ref. \\cite{STreview}.)  When observational results on Type Ia supernovae (SNeIa) have strongly supported the possibility that the universe is now in an accelerated stage of expansion \\cite{per1} \\cite{per2} \\cite{rie} \\cite{gar}, it has become necessary to consider again a negative--pressure component. The late-time inflation so generated is however {\\em soft} with respect to the earlier one and is again based on the dominance of at least one scalar field referred to as {\\em quintessence} \\cite{ost} \\cite{cal} \\cite{zla} \\cite{ste1} \\cite{wan} or, earlier, as {\\em x--field} \\cite{tur1} \\cite{chi}. This gives rise to a sort of energy component which only appears measurable through gravitational effects. There is in fact a wide agreement that today a form of such a {\\em dark energy} \\cite{van} \\cite{tur2} has to be taken into account in any realistic cosmological model. Seen as a constant $\\Lambda$--term \\cite{car1} \\cite{car2} \\cite{sah1}, dark energy density has been usually seen as the vacuum energy density \\cite{rob}, leading therefore to relevant problems due to huge discrepancies between theory and experiments \\cite{wei} \\cite{car2} \\cite{sah1} \\cite{rugh}.  In this context, the scalar field in itself has gained new attention, and speculations about its nature and evolution have got new strength. Papers like Refs. \\cite{rat} and \\cite{pee} have been reconsidered again under a new light, as they show much of the role a scalar field could play when one wants to have an accelerated expansion today. Many models have thereof been constructed, first of all trying to consider more appealing kinds of potentials driving the dynamics of the scalar field. Among others, the potential has in fact to be seen as an important ingredient of all the involved theoretical models, and many are the trials to give prescriptions in order to reconstruct its form according to observational data (see Refs. \\cite{sai} \\cite{rub1}, for instance; but, for more theoretical considerations, see also Refs. \\cite{lid} \\cite{cal} \\cite{zla} \\cite{ste1} \\cite{bra} \\cite{bin}). Here we take into account a specific class of potentials useful for the present evolution of the universe, later investigating all the relative involved cases so derived in the cosmological equations. The exponential--type potential has received much attention since the late 80's \\cite{bar} \\cite{bur} \\cite{rat}, (but see also Refs. \\cite{lid} \\cite{wet1} \\cite{wet2} \\cite{cop} \\cite{fer} \\cite{fab} \\cite{barr} \\cite{sah2} \\cite{bra} \\cite{bat}, \\cite{ack}, just to cite only some other old papers). In Refs. \\cite{rub2} \\cite{rub3} we have found general exact solutions for two classes of exponential potentials for a scalar field minimally coupled to gravity, in presence of another dust component (ordinary non--relativistic pressureless baryonic matter), also showing \\cite{pav} that such solutions seem to fit SNeIa data given in Ref. \\cite{per2}. (See Ref. \\cite{dem} for much more on confrontation of theoretical predictions with observational data.) Furthermore, in Ref. \\cite{keko} other approximate and exact solutions for cosmology with exponential potentials can be found, while up to more recently it is still widely referred to (as, for example, in Ref. \\cite{batsapa}, where the Noether Symmetry Approach \\cite{de2} \\cite{de1} is used to probe the nature of dark energy, well underlining the role of the exponential potential in such a task). This seems to us a good motivation to investigate also on generalizations of this kind of potentials. (As a matter of fact, it is often studied also the case with the hyperbolic sine potential, which can easily be assimilated with some situations studied below).  Moreover, also for such a class of generalized exponential potentials it is possible to exhibit some \\textit{special} properties: it turns out, actually, that the Einstein field equations are integrable (in the Liouville sense) and separable, that is they can be analytically integrated, at least by quadrature. Finally, it is worth noting that our class of exponential potentials can be selected by finding the most general variables transformation which diagonalizes the scalar field kinetic-energy form, leaving the transformed Lagrangian \\emph{simple}.  Therefore, among other considerations, we then generalize the exponential potential to the extent allowed by the particular technique used in Refs. \\cite{rub2} \\cite{rub3} in order to integrate cosmological equations, still bearing in mind that, in order to describe the present universe, solutions should always allow a late--time acceleration. Such a technique is generally known as the Noether Symmetry Approach to cosmology \\cite{de2}, \\cite{de1}. (Note that in Refs. \\cite{de2}, \\cite{de1}, such an application of the Noether Symmetry Approach already leads to prefer the exponential--type potentials.) As in Refs. \\cite{rub2} and \\cite{rub3} we anyway do not completely adopt it here, limiting ourselves to import from it only the change of variables needed to solve equations. As a matter of fact, we have to point out that this transformation (together with the choice of a specific kind of potential) has to be considered as the main drawback of the approach above, since it leads to a point symmetry which probably applies only to the background evolution, having nothing to do with the actual symmetries of the full theory. We analyze the exact solutions we find in this way, discussing when they are compatible with a late time quintessential expansion of the universe.  As a further step, we also consider and illustrate a possible scenario in the framework of the \\textit{quintessential inflation paradigm} with scalar fields, where  an inflationary  potential drives also the quintessential phase of the scalar field evolution, by means of an exponential form of its late time tail. In this connection, we discuss how such an evolution mechanism for the scalar field potential can be compatible with the so-called inflation--kination transition, when the field energy density is dominated by the kinetic energy of the scalar field $\\varphi$. All this clearly leaves apart the greater complexity involved by the more general forms of the potential studied above, and will be the subject of a forthcoming paper trying to discuss more in general the quintessential inflation. Even if in this second paper we should be able to overcome the dichotomy present in this work, where the two different parts of the paper could well live apart, we believe that it is nonetheless already interesting to offer both the issues in a single paper, trying to unify the whole subject a little forcibly.  In Sec. II we introduce the class of potentials and the general cosmological setting for further considerations. In Sec. III, we systematically derive, when possible, general exact solutions case by case. In Sec. IV we discuss connections between exponential potentials quintessence and inflation. Finally, Sec. V is devoted to a conclusive discussion. ",
        "conclusions": "After the historical role of scalar fields in inflationary cosmology, the quite recent reconsideration of their crucial importance relies upon the fact that they can improve a dynamical mechanism for giving rise today to a repulsive component in cosmic energy, the dark energy. The interest for a scalar field in the universe, otherwise usually thought of as a matter--dominated one, is in fact tied to the very recent cosmological history revealed by astrophysical observations.  In this paper, our attention has been generally focused on the role the exponential--like potential has still to play in nowadays cosmology. At first, we have discussed some models generalizing the simple exponential form of the potential, in order to derive general exact solutions. The technique used has been very simple, being it based on a given {\\em ad hoc} change of variables. Such a procedure is generally possible either \\emph{by chance} or because there exists a sort of a method to deduce that useful transformation. In our work we have adopted the second procedure, i.e. the Noether Symmetry Approach to cosmology, borrowing from it here not only the suggestion on the kind of useful transformations to apply to the variables $a$ and $\\varphi$ involved, but also on the \\emph{natural} kind of potential $V(\\varphi)$ to study. On the other hand, the results we have found in this first part of the paper are many and not always easily discernable. Basically, they have been mathematically derived but not yet appropriately discussed on physical ground. This deserves, of course, further investigations in a forthcoming paper, but we can however try here to sum up what has come out as more interesting for nowadays cosmology.  When $\\lambda = 0$ in the potential, first of all, the only case we have discussed here is for $A^2 \\neq 0\\,, B^2 \\neq 0\\,,\\epsilon = - 1$. The universe then displays infinitely many periods of closed evolution, and one should choose which branch to investigate, anyway without succeeding in describing the currently observed behaviour. The situation with $\\lambda > 0$, on the other hand, depends on the sign of $\\epsilon$, even if it presents accelerated asymptotic evolutions both for $\\epsilon = + 1$ and $\\epsilon = - 1$. In the latter case, such an acceleration is however obtained together with an oscillating behavior. This is also true for $\\lambda < 0$ when $\\epsilon = - 1$, while for $\\epsilon = + 1$ we again find a closed universe.  It remains to be seen that the only physically acceptable situation (and solution) is given for $\\lambda > 0$ and $\\epsilon = + 1$. In a forthcoming paper we are going into details of the cosmological evolution, with regard to each specific solution discussed above. In particular, we are going to perform the necessary confrontation of this theoretical output with the observational data sets.  The second part of the paper is still connected with the exponential potential. But, now, we instead focus on considering and illustrating a possible quintessential inflationary scenario, formulated in terms of a multi--branch scalar field, driving both the inflationary and the quintessential phases of the evolution of the universe. The quintessential tail is realized through an exponential potential (choosing, thus, the latter as our working potential for quintessence); on the contrary, for describing the inflationary plateau we do not fix any inflationary potential, but we propose a parametrization of the inflationary scalar field equation of state, and implement the transition from an inflationary stage to a kination evolution, which is characterized by the value $w_\\varphi=1$ of the equation of state, which is needed in any reasonable quintessential inflationary model and corresponds just to the asymptotic in the past value for the equation of state of the exponential potential scalar field. It turns out that the reconstructed potential changes rather roughly from its value during inflation to its final values, and the \\textit{asymptotic in the future} value for the equation of state $w_\\varphi=1$ corresponds to the \\textit{asymptotic in the past} value for the quintessential exponential potential. Thus, we want to stress again that these behaviours can be naturally connected, and the exponential form of the scalar field potential driving the late stage of the universe could indeed be the asymptotic late time behaviour of the inflationary scalar field, which transits from inflation into kination, toward the late quintessential stage. It is worth noting that such a conclusion is somehow independent of the mechanism proposed for the evolution of the scalar field potential, being only based on a parametric description of the very early inflationary dynamics of the universe, and on the property of the equation of state $w_\\varphi=1$, which characterizes our exponential potential. In our cited  forthcoming paper we are also going to select, among all solutions, the cases which \\textit{preserve} such asymptotic behaviour   $w_\\varphi=1$.  Of course, this mechanism for driving the transition from the inflationary evolution toward the late time accelerated expansion has to be considered as mainly exploratory, and some topics need more investigation. For example, the particles production, that is the gravitational production during the reheating or the preheating, in which particles are produced by virtue of the variation of the classical inflaton field, needs to be investigated in a forthcoming paper, considering the scalar field potential given by Eq. (\\ref{gamma4}). (As to this, we have to remind that, in order to study preheating, it is necessary to couple the classical inflaton field $\\varphi$ to a massless quantum scalar field $\\chi$, as for instance in the Lagrangian \\begin{equation} {\\cal L}_{coupled}= \\frac{1}{2} \\partial_\\mu \\chi \\partial^\\mu \\chi - \\frac{1}{2} \\xi R \\chi^2 -\\frac{1}{2} g^2\\varphi^2\\chi^2\\,, \\end{equation} where $R$ is the scalar curvature, $\\xi$ is the gravitational coupling and $g$ is a quartic coupling constant between $\\phi$ and $\\chi$.)  To conclude, let us note that considering radiation in the model only changes the evolution of the $\\Omega$ parameters. Radiation initially dominates on matter, while later on the energy densities of matter and radiation become equal; after that, for some time matter dominates in the universe, while dark energy starts dominating the expansion rate of the universe only afterwards. These results seem to confirm other investigations of effects of radiation on the evolution of the quintessence field, according to which the energy density of the scalar field during the epoch of nucleosynthesis ($z\\sim 10^{9}$) is much smaller than the energy density of radiation. While the kinetic term in the scalar field energy density vanishes and the potential term becomes constant, the dark--energy term behaves like an effective cosmological constant, not affecting the process of primordial nucleosynthesis."
    },
    "1112/1112.6333_arXiv.txt": {
        "abstract": "A fundamental ingredient in wormhole physics is the presence of exotic matter, which involves the violation of the null energy condition. Although a plethora of wormhole solutions have been explored in the literature, it is useful to find geometries that minimize the usage of exotic matter. In the context of modified gravity, it has also been shown that the normal matter can be imposed to satisfy the null energy condition, and it is the higher order curvature terms, interpreted as a gravitational fluid, that sustain these non-standard wormhole geometries, fundamentally different from their counterparts in general relativity. In this paper, we review recent work in wormhole physics in the context of modified theories of gravity. ",
        "introduction": "Traversable wormholes are hypothetical tunnels in spacetime given by the following line element \\begin{equation} ds^{2}=-e^{2\\Phi(r)}dt^{2}+\\left[1-\\frac{b(r)}{r}\\right]^{-1}\\,dr^2+r^{2}(d\\theta ^{2}+\\sin ^{2}\\theta d\\phi ^{2}) \\,, \\label{WHmetric} \\end{equation} where $\\Phi(r)$ and $b(r)$ are arbitrary functions of the radial coordinate, $r$, denoted as the redshift function, and the shape function, respectively \\cite{Morris:1988cz}. The radial coordinate $r$ decreases from infinity to a minimum value $r_0$, representing the location of the wormhole throat, where $b(r_0)=r_0$, and then increases from $r_0$ back to infinity. For the wormhole to be traversable, one must demand that there are no horizons present, which are identified as the surfaces with $e^{2\\Phi}\\rightarrow0$, so that $\\Phi(r)$ must be finite everywhere. Note that in any region where the $t$ coordinate is timelike requires that $r > b(r)$.  Now, a fundamental property of wormholes is the flaring out of the throat, which is  translated by the condition $(b-b^{\\prime}r)/b^{2}>0$ \\cite{Morris:1988cz}. At the throat $b(r_{0})=r=r_{0}$, the condition $b'(r_{0})<1$ is imposed in order to have wormhole solutions. It is precisely these restrictions that impose the violation of the null energy condition (NEC) in classical general relativity. At this stage, it is important to emphasize that the NEC violation involves a stress-energy tensor such that $T_{\\mu\\nu}k^\\mu k^\\nu <0$, where $k^\\mu$ is any null vector \\cite{Morris:1988cz}.  However, recently it has been shown in the context of modified theories of gravity that the matter threading the wormhole may satisfy the energy conditions, and it is precisely the effective stress-energy tensor involving higher order derivatives that is responsible for the violation of the NEC. This approach has been explored in $f(R)$ gravity \\cite{modgravity1}, curvature-matter couplings \\cite{modgravity2a,modgravity2b}, conformal Weyl gravity \\cite{modgravity3} and in braneworlds \\cite{modgravity4}, amongst other contexts. In this paper, we review some of these scenarios. ",
        "conclusions": "In this work, we have explored the possibility that wormholes be supported by modified theories of gravity. We imposed that the matter threading the wormhole satisfies the energy conditions, and it is the higher order curvature derivative terms (in the context of $f(R)$ gravity and its curvature-matter coupling generalization) and the local high-energy bulk effects and nonlocal corrections from the Weyl curvature in the bulk (in braneworld gravity), that support these nonstandard wormhole geometries, fundamentally different from their counterparts in general relativity.     \\subsection*"
    },
    "1112/1112.1725_arXiv.txt": {
        "abstract": "The companion to the G0V star HR7672 directly imaged by \\citet{liu_02} has moved measurably along its orbit since the discovery epoch, making it possible to determine its dynamical properties. Originally targeted with adaptive optics because it showed a long-term radial velocity acceleration (trend), we have monitored this star with precise Doppler measurements and have now established a 24 year time baseline. The radial velocity variations show significant curvature (change in the acceleration) including an inflection point. We have also obtained a recent image of HR7672B with NIRC2 at Keck. The astrometry also shows curvature. In this paper, we use jointly-fitted Doppler and astrometric models to calculate the three-dimensional orbit and dynamical mass of the companion. The mass of the host star is determined using a direct radius measurement from CHARA interferometry in combination with high resolution spectroscopic modeling. We find that HR7672B has a highly eccentric, $e=0.50^{+0.01}_{-0.01}$, near edge-on, $i=97.3^{+0.4}_{-0.5}$ deg, orbit with semimajor axis, $a=18.3^{+0.4}_{-0.5}$ AU. The mass of the companion is $m=68.7^{+2.4}_{-3.1}M_J$ at the 68.2\\% confidence level. HR7672B thus resides near the substellar boundary, just below the hydrogen-fusing limit. These measurements of the companion mass are independent of its brightness and spectrum and establish HR7672B as a rare and precious ``benchmark\" brown dwarf with a well-determined mass, age, and metallicity essential for testing theoretical evolutionary models and synthetic spectral models. It is presently the only directly imaged L,T,Y-dwarf known to produce an RV trend around a solar-type star. ",
        "introduction": "\\label{sec:intro} Brown dwarfs have complicated atmospheres. Unlike stars, for which it is possible to infer bulk physical properties from spectra alone, the emergent radiation from substellar objects is currently not well understood (e.g., \\citealt{cushing_08}). More than a decade following the direct detection of the first L and T dwarfs \\citep{becklin_zuckerman_88,nakajima_95,oppenheimer_95,rebolo_95,basri_96,kirkpatrick_99_lt}, and now faced with the recent detection of the first Y dwarfs \\citep{cushing_11,liu_11,luhman_11}, theoretical spectral models are still undergoing major developments as they presently do not capture all of the relevant physics involved in shaping the spectra of cold bodies \\citep{allard_97,marley_02,baraffe_03,burrows_06,saumon_08}.  Factors that complicate the interpretation of brown dwarf spectra include the necessity to simultaneously model: the opacity of molecular species having millions of absorption lines, which results in an ill-defined continuum; the formation and settling of dust grains; non-equilibrium chemistry resulting from convective mixing; and temporal changes from weather patterns, among other phenomena \\citep{cushing_08,marley_10}. Furthermore, the basic model input parameters, such as mass, radius, age, metallicity, and effective temperature, are often degenerate with one another, particularly for the faintest objects for which only broadband photometry or low resolution spectroscopy are available \\citep{dupuy_09,janson_11,galicher_11}. In order to improve our understanding of low temperature atmospheres and guide the development of more sophisticated models, it is necessary to measure one or more of these physical properties independently of spectra and photometry.  Substellar companions found orbiting nearby stars serve as useful laboratories for calibrating theoretical models. Properties of the host star are more readily measured and may be used to infer properties of the companion, such as chemical composition or possibly age, under the assumption that the star and brown dwarf formed from the same material at the same time \\citep{pinfield_06}. A number of these ``benchmark\" systems have recently been discovered \\citep{lane_01,close_05,liu_07, ireland_08_BD,dupuy_09,bowler_09,biller_10_pztel,wahhaj_11}.  It is also possible to measure the \\emph{mass} of brown dwarfs (and extrasolar planets) using orbital dynamics. This has been accomplished in the past using the transit technique in combination with Doppler radial velocity (RV) observations \\citep{stassun_06,johnson_11_lhs6343}. However, substellar companions with large semimajor axes and proximate distances may be imaged directly (spatially resolved from their host) using adaptive optics (AO) and high-contrast imaging technology \\citep{marois_06,oppenheimer_hinkley_09,absil_mawet_10,crepp_11}. The companion sky-projected orbit can then be traced by measuring its position relative to the star over multiple epochs. This type of astrometry has been used for decades to characterize the orbits of binary stars \\citep{aitken_19} and is significantly easier to realize in practice compared to monitoring the ``wobble\" of a star in the sky relative to an inertial reference grid with $<0.1$ mas accuracy.  If Doppler measurements of the primary star are also obtained over a comparable time baseline, one can construct a three-dimensional orbit. Since the astrometry breaks the $sin(i)$ inclination degeneracy resulting from RV measurements alone, a Keplerian model that self-consistently combines both data sets provides an estimate of the companion true mass (e.g., \\citealt{boden_06}). Knowing the mass of non-hydrogen-fusing objects is crucial because it governs their luminosity evolution in time \\citep{stevenson_91,burrows_97}.   It is commonly thought that the long orbital periods of wide-separation companions prohibits the calculation of dynamical masses. However, those with semimajor axes as large as $a\\approx30$ AU orbiting nearby ($d\\lesssim100$ pc) stars may be characterized in several years to tens of years, because the size of the astrometric uncertainties are $\\sim$1-2 orders of magnitude smaller than the size of the orbit on the sky (thanks in part to the recent miniaturization in the physical size of detector pixels). Orbital solutions converge when both the astrometry and RVs become ``unique\", which generally corresponds to showing curvature or a change in the acceleration (also known as the ``jerk\").\\footnote{See \\citet{lu_09} for analysis of an equivalent problem involving stars orbiting the supermassive black hole in the center of the Galaxy. The orbits are well-characterized even though the periods are of order thousands of years. Likewise, the orbits of comets residing in the outer solar system may be determined with a time baseline of only several days \\citep{bernstein_00}.} With accuracies of $\\approx10$ mas (e.g., the size of a pixel in the NIRC2 camera at Keck -- PI: Keith Matthews), it is possible to calculate the inclination, as well as other parameters, in only several tenths of an orbital phase wrap (tens of degrees change in the true anomaly) depending on the orbit orientation and observing dates \\citep{konopacky_10,currie_11_betapic,dupuy_11}. The constraints become tighter with time as the companion executes its motion along the ellipse.  A natural method to identify promising high-contrast imaging targets is to select stars that show long-term Doppler accelerations (trends). This approach is convenient because RV programs have now established time baselines exceeding a decade. For example, the L$4.5\\pm1.5$ dwarf companion to HR7672 ($=$Gl 779, $=$HD190406) was discovered by \\citet{liu_02} using this technique. This G0V star was originally chosen for AO observations because it showed unambiguous evidence for the existence of a distant companion with substellar minimum mass: an $\\approx \\;$-24 m/s/yr acceleration over a 14 year baseline. HR7672B is currently the only directly imaged L,T,Y-dwarf known to produce an RV trend around a solar-type star.  In this paper, we present an updated RV time series that shows significant orbital motion of the companion since the epoch of the direct imaging discovery. We have also obtained a new astrometric measurement with NIRC2 at Keck. Combining our observations with astrometry from the literature, and performing a joint-fit analysis to the full data set, we construct a three-dimensional orbit and calculate the dynamical mass of HR7672B with a fractional error of only 4\\% using Newtonian dynamics (95\\% confidence level). By explicitly connecting the mass of a brown dwarf to its radiative properties, the results may be used to inform both theoretical evolutionary models and synthetic spectral models. This study is afforded by an equally detailed characterization of HR7672A, for which we have modeled its high resolution spectrum to determine its surface gravity and metallicity, acquired a spectral energy distribution (SED) to determine its luminosity, and measured its radius directly using interferometry. ",
        "conclusions": "We present the first three-dimensional orbit solution and mass determination for a directly imaged L-dwarf companion to a solar-type star. HR7672 is a unique benchmark system because: (i) the primary is a G0V star amenable to ultra-precise Doppler measurements; (ii) the metallicity of the primary is reliably determined from spectroscopy; (iii) the mass and age of the primary are well-constrained; and (iv) the parallax has been measured with a fractional error of only 1.2\\%.  \\begin{figure}[!t] \\begin{center} \\includegraphics[height=2.25in]{f11.eps} \\caption{Cumulative summation of the posterior companion mass distribution. The canonical boundary at $72M_J$ separating stars from brown dwarfs is labeled for reference. HR7672B is a brown dwarf at 90.1\\% confidence. Results from orbital dynamics are more accurate and precise than estimates based on spectrophotometry.} \\end{center}\\label{fig:cumsum} \\end{figure}  We have monitored HR7672 with precise RV measurements for 24 years, covering 33\\% of an orbital cycle. The observational cadence is high compared to other trend stars because this system was also a target of the NASA-UC $\\eta$-Earth program \\citep{howard_10}, which began conveniently around the same time as the Doppler signal started to show curvature. We have also obtained a recent high-contrast image of the system with NIRC2 at Keck, significantly increasing the astrometric time baseline by extending the sky-projected orbital coverage to 13\\% of a full cycle. Both data sets have a high ($\\approx$100) signal-to-noise ratio and exhibit substantial orbital motion. We have performed a joint Doppler-astrometric MCMC analysis to compute the companion orbital elements, true mass, and their uncertainties.  We find that HR7672B has a near edge-on, $i=97.3^{+0.4}_{-0.5}$ deg, highly eccentric, $e=0.50^{+0.01}_{-0.01}$, orbit, with semimajor axis $a=18.3^{+0.4}_{-0.5}$ AU. Presently accelerating on its path back towards the primary star, both from the perspective of the Earth and from the perspective of HR7672A, this companion will gradually becoming more difficult to image directly. We predict that HR7672B will reach periastron in the year 2014 just prior to disappearing behind HR7672A. This companion does not yet have JH spectra. We recommend that high-contrast imaging programs employing integral-field spectrographs observe this target as soon as possible, while the components are reasonably well separated.  The mass of HR7672B is $68.7^{+2.4}_{-3.1}M_J$. This L$4.5\\pm1.5$-dwarf therefore resides near the stellar/substellar boundary, though is likely a brown dwarf if not a transition object: 90.1\\% of the mass probability distribution falls below the $72M_J$ dividing line. Our dynamical measurement is more accurate and precise than that obtained previously from spectrophotometry and the use of theoretical models. With a fractional error of only 4\\%, it is comparable to the results for transiting brown dwarfs \\citep{stassun_06,johnson_11_lhs6343}. Uncertainty in the mass will decrease further as additional RV and astrometric measurements are obtained.  These observations provide a mass determination that is completely independent of the companion spectrum, colors, and luminosity (see Appendix A for an estimate of the companion bolometric luminosity). Presently, few benchmark systems have dynamical mass measurements and provide such comprehensive information. For instance, GD165B has been well-characterized with optical and near-infrared spectroscopy, but its age is uncertain because it orbits a white dwarf \\citep{kirkpatrick_99,jones_94}. HD203030B has also been studied across a wide bandpass, conveniently orbits a G8V star, and has a narrow age range (130-400 Myr), but its projected separation is 487 AU, thus prohibiting a dynamical mass \\citep{metchev_hillenbrand_06}. Likewise, HD3651B and HD114762B orbit solar-type stars but their projected separations are 480 AU and 128 AU respectively \\citep{liu_07,bowler_09}. Many other dynamical measurements have targeted companions to M-stars, thus helping to circumvent issues with contrast \\citep{zapatero_04,ireland_08_BD,dupuy_09_lhs,konopacky_10}, but M-star metallicities are notoriously difficult to determine \\citep{johnson_09}. Low-mass companions to solar-type stars have also been studied using Doppler measurements in combination with HST astrometry \\citep{benedict_10} and Hipparcos intermediate astrometry \\citep{sahlmann_11}, but indirect observations preclude a comparison between spectral models and dynamical mass measurements by definition. One comparable system, HD~130948BC, consists of a double brown dwarf orbiting a G2V primary. Although RVs have not been obtained for the pair, which are nominally separated by $\\approx$100 mas, a near-unity luminosity ratio permits a precise individual mass determination \\citep{dupuy_09,konopacky_10}. We conclude that HR7672B is thus a rare and precious mass, age, and metallicity benchmark brown dwarf companion that may be used to explicitly calibrate theoretical evolutionary models and synthetic spectral models by anchoring input parameters to measured values.  Detailed substellar companion characterization studies are particularly relevant now that the first extrasolar planet spectra are securely in hand \\citep{bowler_10,barman_11_2m1207,barman_11}. As the next generation of high-contrast imaging instruments come online (e.g., \\citealt{hinkley_11_PASP}), and more cold bodies orbiting nearby stars are discovered, HR7672B will serve as a helpful guide for the development of more sophisticated theoretical models by providing a link between our understanding of stars and planets."
    },
    "1112/1112.3666_arXiv.txt": {
        "abstract": "As a part of our ongoing Volume-limited A-Star (VAST) adaptive optics survey, we have obtained observations of 26 binary systems with projected separations $<$100 AU, 13 of which have sufficient historical measurements to allow for refinement of their orbital elements. For each system with an estimated orbit, the dynamical system mass obtained was compared with the system mass estimated from mass-magnitude relations. Discrepancies between the dynamical and theoretical system mass can be explained by the presence of a previously unresolved spectroscopic component, or by a non-solar metallicity of the system. Using this approach to infer the presence of additional companions, a lower limit to the fraction of binaries, triples, and quadruples can be estimated as 39, 46, and 15 per cent, for systems with at least one companion within 100 AU. The fraction of multiple systems with three or more components shows a relative increase compared to the fraction for Solar-type primaries resolved in previous volume-limited surveys. The observations have also revealed a pair of potentially young ($<$100 Myr) M-dwarf companions, which would make an ideal benchmark for the theoretical models during the pre-Main Sequence contraction phase for M-dwarfs. In addition to those systems with orbit fits, we report 13 systems for which further orbital monitoring observations are required, 11 of which are newly resolved as a part of the VAST survey. ",
        "introduction": "The orbits of binary stars offer one of the few techniques to determine stellar masses, or masses and radii in cases with a favorable geometry (e.g. \\citealp{Andersen:1991wd}). Orbits of binaries with young ages are particularly important to test theoretical evolutionary models, and examples include a double-lined eclipsing binary in Orion (e.g.\\citealp{Stassun:2008fp}). Low mass stars and brown dwarfs also represent a regime requiring empirical calibration, and the visual orbits of nearby M-, L-, and T- dwarfs have been used to measure system masses and compare with theoretical mass-luminosity relations \\citep{Dupuy:2009jba}.  Visual orbits also provide a method to search for evidence of additional components and determine higher order multiplicity, by identifying systems with dynamical masses significantly in excess of the theoretical predictions. These visual binaries with an indication of unresolved companions can be monitored with spectroscopy or interferometry to determine the properties such as period and mass ratio of the closer pair and augment the statistics compiled from catalogues \\citep{Tokovinin:2008bl}. The properties of  higher order multiple systems represent tests of formation scenarios including fragmentation of cores (e.g. \\citealp{Pringle:1989vl,Bonnell:2001tb}) and disks (e.g. \\citealp{Stamatellos:2007be,Kratter:2010gf}) and may be influenced by processes such as accretion (e.g. \\citealp{Bate:2000ir}) and dynamical interactions (e.g. \\citealp{McDonald:1995uw,Lodato:2007ja})  In this paper, a subset of the systems resolved by our ongoing volume-limited A-star (VAST) survey are used to determine dynamical system masses from orbit fits and to compare the results with theoretical models and search for additional unresolved stellar companions. The sample of AO-imaged binaries considered in this study is detailed in Section 2, and a short summary of the new observations is given in Section 3. The data analysis, including the AO image processing to determine the relative positions and the subsequent orbit determination from the compilation of position data, is explained in Section 4. Section 5 reports the astrometric results from the new measurements and the orbital elements and masses based on the orbit fits. The discussion in Section 6 covers a comparison with theoretical mass-magnitude relations, an assessment of the higher order multiplicity, and the identification of a set of targets for continued monitoring. Finally, Section 7 provides a summary and future directions for the project. ",
        "conclusions": "\\begin{table*} \\caption{Comparison of dynamical and theoretical system masses for each system within the orbit subsample} \\begin{tabular}{lr@{$\\pm$}lccr@{$\\pm$}lccr@{$\\pm$}lccr@{$\\pm$}l} &\\multicolumn{2}{c}{Dynamical}&$\\cdot$&\\multicolumn{3}{c}{\\citeauthor{Lejeune:2001fq}}&$\\cdot$&\\multicolumn{3}{c}{\\citeauthor{Marigo:2008fy}}&$\\cdot$&\\multicolumn{3}{c}{\\citeauthor{Siess:2000tk}}\\\\ &\\multicolumn{2}{c}{System}&$\\cdot$&Age&\\multicolumn{2}{c}{System}&$\\cdot$&Age&\\multicolumn{2}{c}{System}&$\\cdot$&Age&\\multicolumn{2}{c}{System}\\\\ HIP&\\multicolumn{2}{c}{Mass (M$_{\\odot}$)}&$\\cdot$&(Myr)&\\multicolumn{2}{c}{Mass (M$_{\\odot}$)}&$\\cdot$&(Myr)&\\multicolumn{2}{c}{Mass (M$_{\\odot}$)}&$\\cdot$&(Myr)&\\multicolumn{2}{c}{Mass (M$_{\\odot}$)}\\\\ \\hline \\hline \\multicolumn{15}{c}{Systems identified as doubles (\\S 6.1.1)}\\\\ \\hline 5300&3.16&0.34&$\\cdot$&280&3.56&0.13&$\\cdot$&400&3.47&0.13&$\\cdot$&355&3.69&0.16\\\\ 9480&2.72&0.13&$\\cdot$&250&3.15&0.20&$\\cdot$&280&3.12&0.20&$\\cdot$&160&3.25&0.31\\\\ 11569 AaAb$^{\\dagger}$&2.12&0.25&$\\cdot$&100&2.45&0.07&$\\cdot$&100&2.89&0.15&$\\cdot$&100&2.97&0.25\\\\ 76952&4.19&0.30&$\\cdot$&315&4.38&0.45&$\\cdot$&400&4.10&0.40&$\\cdot$&315&4.38&0.45\\\\ \\hline \\multicolumn{15}{c}{Low-mass binary pair within hierarchical system (\\S 6.1.2, \\S 6.1.3)}\\\\ \\hline 44127 BC&0.68&0.04&$\\cdot$&--&\\multicolumn{2}{c}{--}&$\\cdot$&250&0.88&0.01&$\\cdot$&50&0.83&0.01\\\\ 82321 BC&1.16&0.09&$\\cdot$&400&1.91&0.02&$\\cdot$&445&1.87&0.02&$\\cdot$&500&1.90&0.02\\\\ \\hline \\multicolumn{15}{c}{Systems with one unresolved spectroscopic component (\\S 6.2.1)}\\\\ \\hline 80628&4.99&0.75&$\\cdot$&630&3.03&0.08&$\\cdot$&630&3.06&0.18&$\\cdot$&445&3.15&0.13\\\\ \\hline \\multicolumn{15}{c}{Systems with two unresolved spectroscopic components (\\S 6.2.1)}\\\\ \\hline 28614&6.36&0.62&$\\cdot$&630&3.98&0.32&$\\cdot$&630&3.96&0.28&$\\cdot$&500&4.18&0.34\\\\ 36850&5.42&0.97&$\\cdot$&280&4.26&0.28&$\\cdot$&355&4.15&0.24&$\\cdot$&315&4.34&0.31\\\\ \\hline \\multicolumn{15}{c}{Systems with evidence of unresolved spectroscopic components (\\S 6.2.2)}\\\\ \\hline 17954&4.15&0.39&$\\cdot$&400&3.49&0.11&$\\cdot$&400&3.44&0.12&$\\cdot$&225&3.67&0.16\\\\ 47479&5.83&0.53&$\\cdot$&560&3.99&0.10&$\\cdot$&560&3.96&0.10&$\\cdot$&445&4.17&0.14\\\\ 93506&5.26&0.37&$\\cdot$&710&4.11&0.25&$\\cdot$&710&4.07&0.23&$\\cdot$&500&4.45&0.36\\\\ \\hline \\multicolumn{15}{l}{$^{\\dagger}$ - The ages and theoretical masses may be significantly biased due to the presence of additional companions}\\\\ \\multicolumn{15}{l}{\\;\\;\\;\\; within the resolution limit of the \\textsl{Tycho2} and \\textsl{2MASS} observations.} \\label{tab:comparison} \\end{tabular} \\end{table*} \\begin{figure} \\resizebox{0.99\\hsize}{!}{{\\includegraphics{fig7.eps}}} \\caption{Using the method introduced in \\S 4.4, a comparison can be made between the dynamical mass determined from the orbit and the mass estimated from theoretical mass-magnitude relations for the stars within the orbit subsample (the \\citet{Marigo:2008fy} models are used for this example). The systems with an A-type star primary which are known to consist only of two components are denoted as black points (\\S 6.1.1). Two hierarchical systems were fully resolved with our data, and the lower-mass pair of each system are in green (\\S 6.1.2). The systems which have a significantly discrepant dynamical mass can be explained by the presence of an unresolved spectroscopic companion within our data. The targets with known spectroscopic components are plotted in red (\\S 6.2.1), while those systems with evidence suggesting a previously unknown spectroscopic component are plotted in blue (\\S 6.2.2). The dashed line denotes the equivalence   between the dynamical and theoretical mass.} \\label{fig:comparison} \\end{figure} We have presented high resolution observations obtained for 26 nearby multiple systems with A-type primaries with projected separations within 100 AU. The subset of 12 targets with orbit fits have been further divided into four distinct categories primarily based on a comparison between the dynamical mass and the theoretical mass estimated from the mass-magnitude relations, as shown in Figure \\ref{fig:comparison}. The theoretical mass estimates for each system using the \\citet{Lejeune:2001fq}, \\citet{Marigo:2008fy}, and \\citet{Siess:2000tk} grids are shown in Table \\ref{tab:comparison}. Those systems with only two known components are discussed in \\S 6.1.1, where the importance of metallicity is described. Two hierarchical systems resolved within our data are discussed in \\S 6.1.2 and \\S 6.1.3, which allow for a comparison to the models within the K- to M-type spectral range. For systems with a dynamical mass excess, significantly higher than the mass predicted from the theoretical mass-magnitude relations, the subset with known spectroscopic components is discussed in \\S 6.2.1, and we present three systems with evidence suggestive of an additional unresolved component in \\S 6.2.2. The remaining targets are discussed in the context of continued monitoring of the orbital motion in \\S 6.3, of these 11 are newly resolved as a part of the VAST survey, 2 were resolved in recent multiplicity surveys of nearby Southern A-type stars \\citep{Ivanov:2006jn,Ehrenreich:2010dc}, and 1 resolved within a large speckle interferometry survey \\citep{McAlister:1987ko}.  \\subsection{Comparison to theoretical models} \\subsubsection{A-type binaries} \\begin{figure*} \\resizebox{0.99\\hsize}{!}{{\\includegraphics{fig8.eps}}} \\caption{The system mass as a function of age based on the mass-magnitude relations derived from three of the model grids. Each panel is similar in nature to Figure \\ref{fig:guide}, with the spread in the mass-age relation introduced by uncertainties in both the measured $K$-band magnitude, and the distance determination. The horizontal dot-dashed line indicates the dynamical mass determined from the system orbit, with the uncertainty denoted by the line-shaded region. The vertical dotted line indicates the age of the system derived from the isochrones, with the line-shaded region denoting the range of ages consistent with the uncertainties in the position of the primary on the CMD. The primary of the HIP 11569 system is significantly bluer than expected for a Main Sequence star, and as such its age has been assigned to 100 Myr, with no corresponding uncertainty. The presence of additional components to the HIP 11569 system, within the resolution limit of both the \\textsl{Tycho2} and \\textsl{2MASS} observations, is the likely cause of this bias.} \\label{fig:true} \\end{figure*} \\begin{figure} \\resizebox{0.99\\hsize}{!}{{\\includegraphics{fig9.eps}}} \\caption{\\textit{(top panel)}: The dynamical mass of the HIP 5300 system (horizontal line-shaded region) is significantly below that of the system mass estimated from mass-magnitude relations (solid shaded region) calculated from the Solar metallicity models from \\citet{Lejeune:2001fq}. \\textit{(middle panel)}: Assuming a higher metallicity ($Z=0.04$), the position of the components on the colour-magnitude diagram suggests a younger age (denoted by the vertical line-shaded region), and the mass-magnitude relation derived from the metal-enhanced models suggests a systematically lower system mass, although the dynamical mass is still discrepant. \\textit{(bottom panel)}: With a metallicity of $Z=0.10$, the age estimate is younger still ($<$10 Myr), and the discrepancy between the dynamical mass and the mass estimated from the mass-magnitude relation is removed.} \\label{fig:metallicity} \\end{figure} Four targets within the orbit subsample are systems where the two known components have been resolved within our high resolution data; HIPs 5300, 9480, 11569, and 76952. For each of these four targets, the mass-magnitude relations were used to determine how the system mass changes as a function of our estimate of the system age, shown in Figure \\ref{fig:true}. For one system, HIP 76952, the dynamical system mass is consistent with the theoretical system mass (Figure \\ref{fig:true}, top panel). The ages of the systems estimated from the position of the primary on the CMD are consistent with their position within the Local Interstellar Bubble (LIB). The minimum age of a star within the LIB, excluding those with relatively high space motions (e.g. the $\\beta$ Pic moving group - \\citealp{Ortega:2002je}), has been shown to be 160 Myrs \\citep{Abt:2011ha}.  The dynamical system masses of the three remaining systems are consistently lower than their theoretical system masses. One possible explanation for the apparently low dynamical masses is a non-Solar metallicity. Varying the metallicity has a significant effect on the system age estimate and the mass-magnitude relations. As an example, a 2 M$_{\\odot}$ star with super-Solar metallicity will be more luminous and have a redder $V-K$ colour index than a similar-mass star of Solar metallicity. A super-Solar metallicity star will appear to be significantly older based on its position on a CMD if Solar metallicity models are used. To explore the effect of our assumption of Solar metallicity for the entire sample, the HIP 5300 system (Figure \\ref{fig:true}, second panel) was studied at varying metallicity values. Using the Solar metallicity models, the dynamical system mass of this system is significantly lower than the theoretical system mass (Figure \\ref{fig:metallicity}, top panel). Increasing the metallicity causes the star to appear both younger, and less massive (Figure \\ref{fig:metallicity}, bottom panel). Only eight of the targets included within this study have metallicity measurements, either from spectroscopic analysis (e.g. \\citealp{Erspamer:2003hz}) or derived from Str\\\"{o}mgren photometry (e.g. \\citealp{Song:2001bv}), demonstrating the need for further study in this area. \\subsubsection{K- and M-type binaries - HIP 44127} \\begin{figure*} \\resizebox{0.99\\hsize}{!}{{\\includegraphics{fig10.eps}}} \\caption{\\textit{Top row} - \\textit{(left)}: An AO observation of HIP 44127 triple system. The orbit of the lower-mass pair (BC) taken from Table \\ref{tab:orbits} is plotted as a solid line for reference. The bright point source to the lower-right of the primary PSF is a known filter ghost. The axis labels are given in arcseconds, with each arcsecond equal to approximately 14.5 AU. \\textit{(centre)}: A colour-magnitude diagram highlighting the position of the primary with respect to the theoretical models; \\citet{Marigo:2008fy} - solid line, \\citet{Siess:2000tk} - dotted line, while the \\citet{Baraffe:1998ux} models do not extend to high enough masses. The age of the system was estimated using the isochrone which best fit the position of the primary. \\textit{(right)}: A colour-magnitude diagram highlighting the position of the two lower-mass components (B - blue, C - red). The line styles are the same as the centre panel \\textit{Bottom row} - The theoretical system mass plotted as a function of system age for each of the three models. The vertical dotted line within each panel indicates the age of the system estimated from using each of the models, while the horizontal dot-dashed line indicates the dynamical system mass estimated from the orbital fit. As the \\citet{Baraffe:1998ux} models do not extend to the A-type star mass range, an age estimate is not possible.} \\label{fig7} \\end{figure*}  The detection of two hierarchical systems, with pairs of lower-mass companions in a wide orbit around an A-type primary, allows for a comparison of the theoretical models in the low-mass regime where the models differ in the treatment of the contraction phase of these objects (e.g. Figure \\ref{fig:massmag}). Our observation of the HIP 44127 system, shown in Figure \\ref{fig7}, resolves three components to this hierarchical system, with an A-star primary (A) separated by $\\sim 4\\arcsec$ from two fainter, gravitationally bound, companions (BC).  The orbit presented in the Sixth Orbit Catalog of the BC pair around the A-type primary is in disagreement with our recent observations of this system. Although the phase coverage is insufficient for a robust orbital determination, the high proper motion of the primary ($\\Delta \\alpha = -441.1$ mas yr$^{-1}$, $\\Delta \\delta = -215.2$ mas yr$^{-1}$) suggests that the BC pair is co-moving. In addition, radial velocity variations detected within the spectra of the primary suggest the presence of a spectroscopic component to this system with a period of $\\sim$11 years \\citep{Abt:1965fz}. The separation of this component was anticipated to be between 0$\\farcs$2 and 0$\\farcs$6, reaching maximum separation in the middle of 2007 \\citep{Docobo:2006bh}. We do not resolve any companion in any epoch of our observations which are consistent with these predictions, placing an upper limit to the separation of 0$\\farcs$08, 0$\\farcs$10, and 0$\\farcs$10 in 2008, 2010, and 2011 respectively. The magnitude difference between the primary and the suggested companion, estimated to be $\\Delta m=1.2$, is well within our detection limits at the expected separation \\citep{Docobo:2006bh}.  A refined orbit of HIP 44127 BC is presented in Table \\ref{tab:orbits} and Figure \\ref{fig:orbitb}, with an estimated dynamical system mass of 0.68$\\pm$0.04 M$_{\\odot}$. The small magnitude difference between the two components ($\\Delta K \\sim 0$) suggests a mass ratio close to unity. Assuming individual masses of $\\sim 0.34$ M$_{\\odot}$ the components would be of early- to mid-M spectral type \\citep{Baraffe:1996ip}, a region of particular disagreement between theoretical models. The \\citet{Baraffe:1998ux} and \\citet{Siess:2000tk} models both predict that for a binary consisting of two stars of magnitudes equal to the measured magnitudes of the two components ($M_K=5.87$), the system mass will increase from 0.3 to 0.9 M$_{\\odot}$ between 10 and 100 Myrs (Figure \\ref{fig7}), as these models take into account the phase of contraction onto the Main Sequence for lower-mass stars. No such change is predicted by the \\citet{Marigo:2008fy} models, with the system mass remaining unchanged between 10 Myrs and 1 Gyr (Figure \\ref{fig7}). The theoretical system masses are only consistent with the dynamical mass at $\\sim 30$ Myrs in the \\citet{Baraffe:1998ux} and \\citet{Siess:2000tk} models, and are not consistent with any age in the \\citet{Marigo:2008fy} models. The metallicity of this system has been measured to be almost Solar \\citep{Wu:2010cu}, removing the degeneracy which exists between metallicity and the mass and age of the system (e.g. \\S 6.2.6).  The position of the primary on the CMD suggests a relatively young age of the system between 50 and 250 Myr. The $V-K$ colours of the BC pair also suggest a young age, between 40 and 100 Myr using the \\citet{Siess:2000tk} and \\citet{Baraffe:1998ux} models. The inferred young age of the system is not inconsistent with the minimum age of stars found within the LIB \\citep{Abt:2011ha}, given the relatively high $UVW$ space velocity of the HIP 44127 system \\citep{Palous:1983ty}. The detection of X-ray emission from this system is also of interest. Previous studies have shown that A7 stars such as the primary should not emit X-rays, and that any detection of X-rays from the position of the star can be indicative of a lower-mass companion (e.g. \\citealp{DeRosa:2011ci}). The \\textsl{ROSAT} source J085913.0+480227 is coincident with the optical position of HIP 44127 \\citep{Voges:1999ws}.  This hierarchical triple system warrants further study, primarily in order to refine the orbital fit as the BC pair approaches apastron passage in 2018. This system also makes an ideal candidate for future spectroscopic observations to search for the narrow spectral lines of the two faint companions. With a double-lined spectroscopic orbit fit, model-independent masses can be calculated for the individual components. If the young age suggested by both the position on the CMD and the system mass estimated from the mass-magnitude relations is correct, the BC pair would be an ideal calibrator for the theoretical models in the young, low-mass regime.  \\subsubsection{K- and M-type binaries - HIP 82321} \\begin{figure*} \\resizebox{0.99\\hsize}{!}{{\\includegraphics{fig11.eps}}} \\caption{The theoretical system mass plotted as a function of system age for each of the three models for the HIP 82321 BC system. The vertical dotted line within each panel indicates the age of the system estimated from using each of the models, while the horizontal dot-dashed line indicates the dynamical system mass estimated from the orbital fit. The clear discrepancy between the dynamical and theoretical system masses can be explained by the lack of complete coverage of the orbit presented in Figure \\ref{fig:orbitb}. Further high angular resolution observations of this system will allow for an improved determination of the orbital parameters.} \\label{fig:82321} \\end{figure*}  The hierarchical triple system HIP 82321 was resolved within a single epoch of our AO observations. The A-type primary (A) is $\\sim 2\\arcsec$ from a binary pair of two lower-mass companions (BC) in a wide orbit. The significant proper motion of the primary ($\\Delta  \\alpha = 22.8$ mas yr$^{-1}$, $\\Delta \\delta = -51.4$ mas yr$^{-1}$), and near constant separation of the three stars, suggest that the lower-mass pair is co-moving with the primary. Spectroscopic measurements of the lower-mass components of this system are possible given the small magnitude difference of the pair with respect to the primary ($\\Delta K \\sim 2.5$); such spectra would allow for a determination of the individual masses independent of the distance to the system. The \\textsl{Hipparcos} parallax and \\textsl{2MASS} $K$-band magnitude measurements of this system allow for a tight constraint of the system age to between 300 and 400Myrs, based on the position of the primary on the CMD. The primary is also a possible member of the Ursa Major moving group, an association of Solar metallicity stars with an age between 300 Myrs \\citep{Soderblom:1993kr} and 500 Myrs \\citep{King:2003df}.  A refined orbit of the BC pair is presented in Table \\ref{tab:orbits}, with an estimated dynamical system mass of 1.32$\\pm$0.05 M$_{\\odot}$. The theoretical system mass, assuming the distance to each component is the same as the primary, is significantly higher at $\\sim 2.0$ M$_{\\odot}$ (Figure \\ref{fig:82321}). The sparse coverage of the orbit (Figure \\ref{fig:orbita}) suggests that the orbit fit could be poor, resulting in an incorrect dynamical system mass. The orbit fit would be significantly improved with subsequent observations. An alternative scenario is that the pair are a background binary with a proper motion similar to the primary, which would bring the dynamical and photometrically derived masses into agreement.  \\subsection{Higher-order multiplicity} \\begin{figure} \\resizebox{0.99\\hsize}{!}{{\\includegraphics{fig12.eps}}} \\caption{A schematic of the 7 hierarchical systems within both subsamples, including only components within a projected separation of 100AU to the A-type primary. The schematic diagram does not include the suspected spectroscopic binaries described in \\S 6.2.2. Those systems which were newly resolved as a part of the VAST survey are highlighted in blue. Each filled circle represents an individual component, with the period of the binary pair denoted by its vertical position. Our high resolution observations are only sensitive to systems with orbital periods greater than $\\sim$ 10 years, denoted by the dashed horizontal line. For wide separation systems with indeterminate orbital periods, such as HIP 128 AB--C, the orbital period has been arbitrarily set at 1000 years, denoted by the dotted horizontal line.} \\label{fig:hierarchical_line} \\end{figure} \\subsubsection{Known spectroscopic binaries} \\begin{figure*} \\resizebox{0.99\\hsize}{!}{{\\includegraphics{fig13.eps}}} \\caption{The six systems for which spectroscopic components are either known to be present through previous spectroscopic observations (HIPs 28614, 36850, 80628), or thought to exist due to a significant discrepancy between the theoretical system mass and the dynamical system mass. For clarity only one model has been used to estimate the theoretical system mass \\citep{Lejeune:2001fq}. The horizontal displacement between the intersection of the dynamical system mass (dot-dashed line) and the age estimate of the system (dotted line), and the theoretical system mass curve (solid shaded region), gives an order of magnitude estimate of the mass of the unresolved spectroscopic components.} \\label{fig:spectroscopic} \\end{figure*} The presence of an unresolved spectroscopic binary can have a significant impact on the magnitude assigned to each component of the multiple system, increasing the estimated component magnitude by as much as 0.75 mag for an unresolved equal-mass spectroscopic binary. For each target within this study, the literature was searched for references to additional components resolved through spectroscopic or interferometric observation which would influence any comparison made between the dynamical system mass and the theoretical system mass (e.g. Figure \\ref{fig:spectroscopic}). Of the targets within the orbit subsample, three are known to have additional spectroscopic components (HIPs 28614, 36850, 80628). The spectroscopic component to HIP 9480 resolved by \\cite{Abt:1965fz} is resolved within our adaptive optics observations. Similarly, for the monitoring subsample, two are known spectroscopic binaries (HIPs 128, 116611). The greater frequency of spectroscopic binaries within the orbit sample, relative to the monitoring sample, can be explained by the narrow spectral lines of the former sample due to their relatively low radial velocities, $\\langle v\\sin i\\rangle = 71$ km s$^{-1}$ compared with $\\langle v\\sin i\\rangle = 115$ km s$^{-1}$ for the latter sample \\citep{Abt:1995eo}. The magnitudes listed for the resolved components of these systems (Table \\ref{tab:components}) are the blended magnitudes of the listed spectroscopic components. In addition to these known spectroscopic binaries, three members of the orbit subsample are resolved as hierarchical triples within our high resolution observations (HIPs 11569, 44128, 82321). A schematic representation of the higher-order multiplicity systems is given in Figure \\ref{fig:hierarchical_line}.  For the three systems with known spectroscopic components within the orbit sample, the dynamical system mass includes the mass of each component, regardless of whether it is resolved within our data. This causes a significant discrepancy when the dynamical system mass is compared with the theoretical system mass (e.g. Figure \\ref{fig:comparison}), with the dynamical system mass being systematically higher. The discrepancy between the two values cannot be directly converted into a mass for the unresolved components however, as a blended magnitude would have been used when determining the mass from the mass-magnitude relations.  \\subsubsection{Evidence of spectroscopic components}  Systems with significantly higher dynamical masses than theoretical masses obtained from mass-magnitude relations are strong candidates for multiple systems with unresolved components which have not been detected in spectroscopic observations. We find two such systems with the signature of a possible unresolved component: HIP 17954 and HIP 93506 (Figure \\ref{fig:spectroscopic}), with masses of the order of 0.5 and 1.5 M$_{\\odot}$. Sensitive spectroscopic observations of both systems may lead to the detection of the spectral lines from an unresolved lower-mass component. A similar phenomenon is observed for the HIP 47479 system, although the number of measurements used to determine the orbit is particularly low, making the dynamical mass less certain. The narrow spectral lines of HIP 47479, implied by the low measured stellar rotational velocity \\citep{Royer:2007dj}, make it an ideal candidate for spectroscopic follow up. Previous spectroscopic observations of this system reveal $v\\sin i$ variations with a magnitude of 40 km s$^{-1}$ \\citep{Moore:1932vo}.  From the total sample of 26 systems, and only considering stellar companions within 100 AU, a lower-limit on the higher-order multiplicity of A-type stars can be estimated. Assuming the suspected unresolved companions described earlier in this section are true, there are five double, six triple, and two quadruple systems within the orbit subsample, corresponding to frequencies of 39, 46, and 15 per cent, respectively. This lower-limit shows an enhancement on the higher-order multiplicity of A-type stars when compared with Solar-type primaries (74 per cent double, 20 per cent triple, and 6 per cent quadruple or higher-order - \\citealp{Raghavan:2010gd}), and is more consistent with the fraction reported for more massive O-type primaries (46 per cent double, and 54 per cent triple or higher-order - \\citealp{Mason:2009fm}).  \\subsection{Continued monitoring targets} \\subsubsection{Newly-resolved binaries} \\begin{figure*} \\resizebox{0.9\\hsize}{!}{{\\includegraphics{fig14a.eps}}} \\caption{Observations of twelve of the thirteen systems with projected separations $< 100$AU suggested as future orbital monitoring targets. The companion in each image is highlighted, with a logarithmic image scale between $1$ (white) and $10^{-6}$ (black) relative to the peak intensity of the primary.} \\label{fig:futurea} \\end{figure*} \\addtocounter{figure}{-1} \\begin{figure*} \\resizebox{0.9\\hsize}{!}{{\\includegraphics[trim=0mm 152mm 0mm 0mm]{fig14b.eps}}} \\caption{\\textit{(continued)}: An observation of one of the thirteen systems with projected separations $< 100$AU suggested as future orbital monitoring targets. The companion in each image is highlighted, with a logarithmic image scale between $1$ (white) and $10^{-6}$ (black) relative to the peak intensity of the primary.} \\label{fig:futureb} \\end{figure*} We have identified 13 binary systems with projected separations ranging between 13 AU and 96 AU which would make ideal candidates for future orbital monitoring projects (Table \\ref{tab:futuresample_obs}). The 100 AU projected separation cut-off was applied to select only systems for which orbital motion could be detected with several years of observations. The binaries resolved within the monitoring subsample typically have lower-mass ratios than for the orbit subsample, a demonstration of the effectiveness of AO observations at detecting high-contrast binaries. Based on their position on the CMD, two  members of the monitoring subsample (HIP 5310, HIP 18217) appear to lie on the Zero-Age Main Sequence, and the measured magnitude difference between primary and secondary would correspond to a late K or early M-type companion in each case. These companions are of particular interest as they will allow for further tests of the theoretical models within the young, low-mass regime. A gallery of the observations obtained of each target within the monitoring subsample is shown in Figure\\ref{fig:futurea}.  \\subsubsection{HIP 77660} \\begin{figure} \\resizebox{0.99\\hsize}{!}{{\\includegraphics{fig15.eps}}} \\caption{\\textit{(top panel)}: The historical measurement of the system HIP 77660, plotted with the orbital fit estimated by \\citet{Docobo:2010hf}. This system demonstrates the quadrant ambiguity limitation of the speckle interferometry technique. The binary companion was measured by \\citet{Tokovinin:2010ft} using speckle interferometry to be at a position angle of 248.4$^{\\circ}$ to 246.8$^{\\circ}$ between mid-2007 and mid-2009 (blue filled circles, top right). Our AO images, which do not suffer from the same quadrant ambiguity, shows the binary at a position angle of 69.7$^{\\circ}$ to 67.3$^{\\circ}$ between mid-2006 and mid-2011 (red filled stars, left). \\textit{(bottom panel)}: The relative position of the secondary, with 180$^{\\circ}$ added to the four speckle interferometry measurements which were inconsistent with the observations presented within this study. The solid line connects the first and last observations, showing a deviation away from linear motion. Further observations are required to determine if this deviation is indicative of the true orbital path of the secondary about the primary.} \\label{fig8} \\end{figure} An advantage of monitoring binary systems for orbital motion with AO instrumentation over interferometric techniques is the elimination of the quadrant ambiguity. In some cases, the output of the image processing of speckle data results in a 180$^{\\circ}$ ambiguity on the position angle measurement \\citep{Bagnuolo:1992ck}. This may lead to a scenario where orbital motion is thought to exist for a binary pair, when in fact this motion is the product of such a quadrant uncertainty combined with the true linear motion of the companion. The ambiguity can be resolved by observing the system using AO imaging, where no reconstruction is required to obtain the final science image.  Based on our AO observations of the binary system HIP 77660, it appears that such a quadrant ambiguity has occurred, making linear motion appear as orbital motion, as shown in Figure \\ref{fig8}. The significant proper motion of the system ($\\Delta\\alpha = -91$ mas/yr, $\\Delta\\delta = -28$ mas/yr), as measured by \\textsl{Hipparcos}, is inconsistent with a stationary background object. While future measurements of this system will be able to resolve the presence of orbital motion after a sufficient time baseline, spectroscopy or multi-colour photometry will allow for a rapid characterisation of the properties of both components."
    },
    "1112/1112.2660_arXiv.txt": {
        "abstract": "Gravitational stability of a disc consisting of the gaseous and the stellar components are studied in the linear regime when the gaseous component is turbulent. A phenomenological approach is adopted to describe the turbulence, in which not only the effective surface density but the velocity dispersion of the gaseous component are both scale-dependent as power-law functions of the wavenumber of the perturbations. Also, the stellar component which has gravitational interaction with the gas is considered as a fluid. We calculate growth rate of the perturbations and in the most of the studied cases,  the stability of the disc highly depends on the existence of the stars and the  exponents of the functions for describing the turbulence. Our analysis suggests that the conventional gas and star threshold is not adequate for analyzing stability of the two-component discs when turbulence is considered. ",
        "introduction": "\\label{sec:1}    Various instability mechanisms may operate in an accretion disc depending on the physical properties of the disc itself. An imbalance of the heating and cooling rates in an thin accretion disc may lead to a disc which is thermally unstable \\citep[e.g.,][]{LF74,pringle76,piran78}. Under some conditions, an accretion disc is hot enough to be fully or partially ionized which then implies to expect significant dynamical effects of the magnetic field. In fact, an instability related to the rotational profile and the magnetic field of an accretion disc is known to be the driven mechanism of generating turbulence inside the accretion discs \\citep{balbus91}. In the most of the accreting systems, the mass of the disc is negligible in comparison to the mass of the central object and so, one may safely neglect the self-gravity of the disc. However, gaseous or stellar discs are subject to the gravitational instability if the Toomre parameter drops below unity \\citep{toomre64}  and the cooling time scale is less than a few time of the dynamical time scale \\citep{gam2001}.  Formation of structures like planets in protoplanetary discs \\citep[e.g.,][]{boss98,rafikov2005,rafikov2007} or stars at the Galactic center \\citep[e.g.,][]{tan2004,nayakshin2005,nayakshin2006,levin2007} are described based on the gravitational instability.  Analysis of the gravitational stability of a disc started by the  pioneer works of \\citet{safr60} and \\citet{toomre64}.  The onset of axisymmetric gravitational instability is determined by the following condition \\citep{toomre64} \\begin{equation}\\label{eq:Toomre} Q=\\frac{c_{\\rm s} \\kappa}{\\pi G \\Sigma} < Q_{0}, \\end{equation} where $Q_{0}=1$ and $c_{\\rm s}$, $\\kappa$ and $\\Sigma$ are the sound speed, epicyclic frequency and the surface density of the disc, respectively. The non-dimensional Toomre parameter $Q$ has a vital role in the gravitational stability of an accretion disc. The non-axisymmetric perturbations  grow for somewhat higher $Q_{0}$.  However, the Toomre condition (\\ref{eq:Toomre}) is not the only criteria for the fragmentation of a disc. In fact,  thermodynamics of the disc and its efficiency of cooling dictates another condition for the fragmentation of a disc as has been suggested in a seminal work by \\citet{gam2001} and investigated via extensive numerical simulations \\citep*[e.g.,][]{gammie2003,rice2005,clarke2007,cos2009}.    Although validity of the Toomre condition for the stability of the self-gravitating discs has also been confirmed by direct numerical simulations, in some of the astrophysical systems like galaxies one component description is not adequate to study their stability \\citep[e.g.,][]{Yim2011,bournaud,bruce2011}. A typical galaxy consists of the gas and  the stars. Over recent two decades, many authors studied gravitational stability of such two-components discs \\citep[e.g.,][]{jog84,romeo92,silk94,jog96}. The main outcome is the stellar component has a destabilizing role either a fluid approximation or a collisionless description is used to model the stellar component. Unfortunately, it is not possible to represent the stability condition of a two-component disc is a closed analytical form  like the one component case, though some authors introduced approximate analytical formulae for the onset of the instability in a disc \\citep[e.g.,][]{silk94,romeo2011b}.    Recently, \\citet*{romeo2010} (hereafter; RBA) studied possible effects of the turbulence on the gravitational stability of the accretion discs using a simple, but very illustrative model for the turbulence. They used scale-dependent relations  for the effective surface density and  the velocity dispersion of the gaseous component. Then, they could obtain conditions under which the self-gravitating disc becomes unstable depending on the properties of the turbulence. For a range of the input parameters, they showed the conventional Toomre parameter is modified  to correctly describe the stability of the disc when turbulence of the gas is considered.    However, such kind of approach for studying gravitational stability of the discs has already been applied by a few authors, but for some specific cases (e.g., Elmegreen 1996). In another relevant study, \\citet*{bruce2011} investigated gravitational instabilities in two-component disc galaxies when the gas dissipates on the local crossing time. However, he did the linear analysis by considering the energy equation and prescribing a scale-dependent dissipation rate instead of  following a phenomenogical approach to describe the turbulence. The main finding is an increase to the conventional gas and star threshold by a factor of two or three because of the existence of the dissipation which removes the stabilization from turbulent pressure.  In this paper, we study gravitational stability of a two fluid disc where the turbulence of the gaseous component is considered similar to RBA. However, we will show that different regimes of the (in)stability according to RBA are significantly modified in the presence of the stars. One can also compare our results to \\citet{jog84} to see how stability of a two fluid disc is changed when the turbulence of the gaseous component is considered. In the next section, our basic assumptions and the main dispersion relation are presented. Then, we will analyze stability of the system for a wide range of the input parameters and the results are compared to RBA. ",
        "conclusions": "We studied gravitational stability of a disc consisting of two components (i.e., gaseous and stellar fluids) using a perturbation analysis. We also considered the turbulent nature of the gaseous fluid using  a phenomenological description which is supported by  the observations and the numerical simulations. Although stability of a two-fluid disc  without turbulence of the gas \\citep*[e.g.,][]{jog84,jog96,silk94,rafikov2001} or one component  turbulent gaseous disc \\citep*{romeo2010} have already been studied analytically, but our analysis is a first step towards considering not only the two-fluid nature of such systems but the turbulence of the gas. We found that turbulence of gas highly affects the stability of the two-component stellar and fluid discs. According to our analysis, not only there is not a closed analytical relation for onset of the instability in a two-component disc, but the conventional gas and star threshold does not correctly address the stability of such systems in the presence of turbulence of the gas. We also explored situations without the thickness correction for the gas. In fact, this correction factor has a stabilizing effect, in particular at high wavenumber $k$. For the perturbations much larger than the thickness of the disc, the thickness correction becomes unity. But as the wavenumber increases so that for the scales larger than the thickness of the disc, the correction factor introduces an additional inverse $k$ dependence to the effective surface density. This weakens the self-gravitational effect significantly. In the presence of the turbulence of the gas, we also found this stabilizing effect.   In dwarf galaxies or the far-outer regions of the galaxy discs, the conventional stability criteria suggests a high level of stability \\citep[e.g.,][]{van,hunter}. Although \\citet*{bruce2011} showed that  such regions are more unstable when  turbulence of the gas is considered, our analysis shows that the stability of such systems highly depends on the nature of the turbulence (i.e., values of $a$ and $b$) in the presence of stars. However, detailed numerical simulations of  the galaxy discs (including the turbulence of the gas component) are needed to confirm our linear stability analysis."
    },
    "1112/1112.0179_arXiv.txt": {
        "abstract": "Recent data from the VVV survey have strengthened evidence for a structural change in the Galactic bulge inwards of $\\vert l\\vert\\le4\\dg$.  Here we show with an N-body barred galaxy simulation that a boxy bulge formed through the bar and buckling instabilities effortlessly matches measured bulge longitude profiles for red clump stars.  The same simulation snapshot was earlier used to clarify the apparent boxy bulge - long bar dichotomy, for the same orientation and scaling.  The change in the slope of the model longitude profiles in the inner few degrees is caused by a transition from highly elongated to more nearly axisymmetric isodensity contours in the inner boxy bulge. This transition is confined to a few degrees from the Galactic plane, thus the change of slope is predicted to disappear at higher Galactic latitudes. We also show that the nuclear star count map derived from this simulation snapshot displays a longitudinal asymmetry similar to that observed in the 2MASS data, but is less flattened to the Galactic plane than the 2MASS map. These results support the interpretation that the Galactic bulge originated from disk evolution, and question the evidence advanced from star count data for the existence of a secondary nuclear bar in the Milky Way. ",
        "introduction": "\\label{sec:intro}  The barred nature of the Milky Way is well-established from NIR photometry \\citep{Blitz+Spergel91, Dwek+95, Binney+97, Bissantz+Gerhard02} and star counts \\citep{Stanek+94, LopezCorredoira+05,Benjamin+05,Cabrera-Lavers+07b}, from comparing HI and CO lv-diagrams with hydrodynamic models \\citep{Englmaier+Gerhard99, Fux99, Bissantz+03}, and from dynamical modelling of the kinematics of bulge stars \\citep{Zhao96,Fux97,Bissantz+04,Shen+10}.  One open issue is the structure of the inner Galactic bulge, for $R\\lsim 0.5$ kpc. Is this simply the scaled-down, higher-density parts of the surrounding boxy bulge? Does it become axisymmetric towards the center? Does it contain a ``classical'' bulge component? Or does it even contain a secondary nuclear bar as do about one third of barred galaxies like ours \\citep{Erwin11}?  This issue is important; for example, a nuclear bar misaligned with the main Galactic boxy-bulge bar would substantially modify the morphology of the gas flow in the central Galaxy, as illustrated by the simulations of \\citet[][hereafter RFC08]{Rodriguez-Fernandez+Combes08} and \\citet{Namekata+09}.  \\citet{Nishiyama+05} analyzed red clump (RC) star counts in the inner bulge. Because RC stars are approximate standard candles, their magnitude distribution is a proxy for the stellar density along the line-of-sight \\citep{Stanek+94}.  Determining the maximum of the RC magnitude distribution at various longitudes in a strip at $b=-1\\dg$, \\citet{Nishiyama+05} found a clear change of slope in the RC longitude profile at $\\vert l\\vert=4\\dg$, separating a steeper slope in the outer bulge from a shallow slope in the nuclear regions.  Recently, \\citet{Gonzalez+11b} analyzed new RC star counts from the VVV survey and found excellent agreement with \\citet{Nishiyama+05}. They also determined the RC maximum versus longitude for $b=+1\\dg$ and found good agreement with the $b=-1\\dg$ results, as expected for a triaxially symmetric structure. This also shows that reddening corrections, which differ greatly between these latitudes, are unlikely to cause the observed change of slope in the longitude profiles.  The RC star counts therefore indicate a significant and fairly sudden change in the structure of the bulge at $\\vert l\\vert\\simeq 4\\dg$. Here we analyze the distribution of RC maxima for a suitably oriented boxy bulge - barred galaxy model. Surprisingly, we find that the RC longitude profiles predicted by this model are very similar to the observed profiles.  We show that this can be traced back to the model's density structure becoming more axisymmetric towards the center, and finally we compare the asymmetry of the predicted star count map in the inner bulge region with available analysis of 2MASS data. ",
        "conclusions": "\\label{sec:discuss}  We have shown that the RC maximum longitude profiles measured for $b=\\pm1\\dg$ are well-matched by an N-body model of a boxy bulge and bar which formed from a bar-unstable disk. This model is the same as in MVG11, and we have used the same orientation and scaling to avoid parameter fitting. The change in slope of the model profiles in the inner few degrees is caused by a transition from highly elongated to nearly axisymmetric density contours around this scale. The average FWHM of the simulated RC profiles in the range $l=\\pm4\\dg$ ($\\simeq\\pm560$ pc at $R_0=8$ kpc) is $\\simeq 0.56$ mag, consistent with the VVV data after convolution with estimated stellar population and patchy extinction broadening. This corresponds to a deconvolved bulge FWHM scale of $\\pm 750$ pc along the LOS in this region. At higher latitudes the transition in the model isodensities weakens, predicting that the change of slope in the longitude profile disappears.  The projected star count map of the model has the same asymmetry pattern as inferred previously from 2MASS observations. The central surface density contours in the Galactic bulge are significantly more flattened, however. This could imply the existence of a dense central disk around the location of the ILR.  Our results generally support the idea that the main part of the Galactic bulge formed from the disk through dynamical instabilities. They also show that the evidence from star counts so far does not require a secondary nuclear bar in the Milky Way.  \\noindent"
    },
    "1112/1112.2726_arXiv.txt": {
        "abstract": "The collapse of baryons into extremely massive stars with masses $\\ga$ 10$^4$ M$_{\\odot}$ in a small fraction of protogalaxies at $z \\ga$ 10 is a promising candidate for the origin of supermassive black holes, some of which grow to a billion solar masses by z $\\sim$ 7.  We determine the maximum masses such stars can attain by accreting primordial gas.  We find that at relatively low accretion rates the strong ionizing radiation of these stars limits their masses to $M_{\\rm *}$ $\\sim$ 10$^3$ M$ _{\\odot}$ ($\\dot{M}_{\\rm acc}$/10$^{-3}$ M$_{\\odot}$ yr$^{-1}$)$^{\\frac{8}{7}}$, where $\\dot{M}_{\\rm acc}$ is the rate at which the star gains mass.  However, at the higher central infall rates usually found in numerical simulations of protogalactic collapse ($\\ga$ 0.1 M$_{\\odot}$ yr$^{-1}$), the lifetime of the star instead limits its final mass to $\\sim$ 10$^6$ M$_{\\odot}$.  Furthermore, for the spherical accretion rates at which the star can grow, its ionizing radiation is confined deep within the protogalaxy, so the evolution of the star is decoupled from that of its host galaxy. Ly$\\alpha$ emission from the surrounding H~{\\sc ii} region is trapped in these heavy accretion flows and likely reprocessed into strong Balmer series emission, which may be observable by the {\\it James Webb Space Telescope}.  This, along with strong He~{\\sc ii} $\\lambda$ 1640 and continuum emission, are likely to be the key observational signatures of the progenitors of supermassive black holes at high redshift. ",
        "introduction": "The existence of 10$^8$ - 10$^9$ M$_{\\odot}$ black holes (BH) in massive galaxies by $z \\sim 7$, less than a billion years after the Big Bang \\citep{fan03,wil03,mort11}, remains one of the great mysteries of cosmological structure formation.   In the $\\Lambda$CDM paradigm, early structure formation is hierarchical, with small dark matter halos at early epochs evolving into ever more massive ones by accretion and mergers through cosmic time. Hence, it is generally held that the supermassive black holes (SMBH) of the $z \\sim 7$ \\textit{Sloan Digital Sky Survey} (\\textit{SDSS}) quasars grow from much smaller seeds at high redshifts.  The origin of these seeds, and how they reach such large masses by such early times, remains to be understood. At least four main processes have been proposed for their formation \\citep[see][]{vol10,ah11rev,primrev}: the collapse of Pop III stars into 100 - 300 M$_{\\odot}$ BH at $z \\sim 20$ \\citep[e.g.][]{mad01,awa09,milos09,th09}, the direct collapse of extremely hydrogen molecule-poor primordial gas in $\\sim$ 10$^8$ M$_{\\odot}$ dark matter halos into 10$^4$ - 10$^6$ M$_{ \\odot}$ BH at $z \\sim 10$ \\citep[e.g.][see also Colgate et al. 2003]{bl03,kbd04,begel06,ln06,spaans06,osh08,rh09,sethi10,io11}, the collapse of dense primeval star clusters into 10$^4$ - 10$^6$ M$_{\\odot}$ BH \\citep[see e.g.][]{brmvol08,dv09}, and the collapse of primordial overdensities in the immediate aftermath of the Big Bang \\citep[see e.g.][]{mack07,carrrev}.  The processes by which black holes form at high redshift and evolve into SMBH must account for how they become so large by $z \\sim 7$ and why their numbers at that redshift are so small, about 1 Gpc$^{-3}$.  Pop III seed BH are plentiful at $z \\sim$ 20 - 30 but must grow at the Eddington limit without interruption to reach 10$^8$ - 10$^9$ M$_{\\odot}$ by $z \\sim$ 7.  This is problematic because they and their progenitors either expel all the baryons from the shallow potential wells of the halos that create them, so they are \"born starving\" \\citep{wan04,jb07,pdc07,awa09,jeon11}, or they eject themselves from their halos, and thus their fuel supply, at hundreds of km/s if they are born in core-collapse supernova explosions (Whalen \\& Fryer 2011 in prep).  Also, accretion onto Pop III BH has been found to be inefficient on small scales, typically at most 20\\% Eddington \\citep{milos09, pm11a}, making the constant duty cycles required for sustained growth difficult \\citep[but see][]{Li11}.  If halos can instead congregate into primitive galaxies of $\\sim$ 10$^8$ M$_{\\odot}$ at $z \\sim$ 10 - 15 that are devoid of the coolant molecular hydrogen (H$_{\\rm 2}$), they reach virial temperatures of $\\sim$ 10$^4$ K and begin to atomically cool.  Numerical simulations of this process \\citep{wta08,rh09,sbh10} find that in analogy to Pop III star formation in much smaller halos at higher redshifts \\citep{nu01,bcl02,abn02,on07,y08,turk09, stacy10,clark11,greif11}, baryons rapidly pool at the center of the halo and form a hydrostatic object. But this object is thought to become far more massive than Pop III stars modeled to date because atomic line cooling and the deeper potential well of the halo lead to much higher infall rates at its center, 0.1 - 1 M$_{\\odot}$ yr$^{-1}$ rather than the 10$^{-4}$ M$_{\\odot}$ yr$^{-1}$ typical of primordial star-forming minihalos at $z \\sim$ 20.  Such objects could collapse into black holes that are far more massive than Pop III BH \\citep[e.g.][]{shib02}, with Bondi-Hoyle accretion rates that allow them to grow into SMBH in less time \\citep[e.g. Wyithe \\& Loeb 2011; but see][]{dotan2011}.   One difficulty with this scenario is that it is not yet understood how primitive galaxies can form in the high Lyman-Werner UV backgrounds needed to fully suppress H$_{\\rm 2}$ molecule formation and quench Pop III star formation in its constituent halos prior to assembly \\citep[1 - 10 times those expected at the epoch of reionization at $z \\sim 6$;][see also Agarwal et al. in prep] {dijkstra08,sbh10,whb11}. Furthermore, it is not known if the object at the center of the protogalaxy even becomes a star, because at such high accretion rates it reaches very large masses on timescales that are short in comparison to Kelvin-Helmholtz times and the onset of nuclear burning (although Ohkubo et al. 2009 have modeled the evolution of Pop III stars under accretion at much lower rates).  It could be that the central object reaches such high entropies and densities that it is enveloped by an event horizon before reaching the main sequence \\citep{fwh01}.  If the central object becomes a star, what governs its rate of growth, its final mass, and thus the mass of the SMBH seed it becomes? Numerous authors have studied how stellar radiation regulates accretion onto Pop III protostars, but at much lower inflow rates than those in the centers of collapsing protogalaxies \\citep[e.g.][]{stahler86b,op01,oi02,tm08,stacy11,hoyy11}.  Do the much higher accretion rates in primitive galaxies quench radiative feedback?  We have performed extensive semianalytical calculations of radiative feedback by Pop III supermassive stars in collapsing protogalaxies at $z \\sim$ 10.  In $\\S \\, 2$, we examine the primary forms of radiative feedback on accretion onto the star (deferring the processes that can be ignored to thorough examination in the Appendix).  We also derive the maximum mass that the supermassive star, and hence the SMBH seed, can achieve as a function of accretion rate, accounting for its prodigious ionizing UV flux.  In $\\S \\, 3$, we estimate the maximum mass that the star can reach if it is limited by its finite lifetime rather than by radiative feedback.  We discuss the observational signatures of rapidly accreting supermassive primordial stars in $\\S \\, 4$.  Finally, in $\\S \\, 5$ we review the implications of final supermassive stellar mass for SMBH seed mass and the appearance of the first quasars in the Universe.  \\begin{figure*}[t] \\centering \\includegraphics[width=3.5in]{Figure1_new.ps} \\caption{ Schematic representation of accretion onto a supermassive Pop~III star of mass $M_{\\rm *}$ and radius $r_{\\rm *}$.  Gas accretes at a constant rate $\\dot{M}_{\\rm acc}$, being acted upon only by gravity at $r$ $>$ $r_{\\rm HII}$, the radius of the H~{\\sc ii} region of the star.  Inside $r_{\\rm HII}$ the gas, whose infall velocity is $v(r)$, is decelerated by momentum absorbed during ionizations, which is indicated by the shortening of the velocity vector arrows.  While continuum photons with energies $h\\nu$ $<$ 13.6 eV escape the H~{\\sc ii} region and do not couple strongly to the gas even outside $r_{\\rm HII}$, resonant line photons do couple strongly and so in principle can also impart momentum to the gas.  In particular, a large amount of momentum is emitted by H~{\\sc i} recombinations as Ly$\\alpha$; however, these photons couple so strongly to the gas that within $r_{\\rm Ly\\alpha}$ more than 50 percent of the emission is trapped in the accretion flow and cannot propagate outward and affect the dynamics of the flow.  Within $r_{\\rm HII}$, Ly$\\alpha$ photons are reprocessed into Balmer series photons which escape the H~{\\sc ii} region and may provide an observable signature of accreting supermassive stars (see $\\S \\, 4$).} \\end{figure*} ",
        "conclusions": "We find that the masses of the stellar seeds of SMBH forming from baryon collapse in early protogalaxies are primarily governed by their intense ionizing UV flux and their finite lifetimes.  For spherically-symmetric accretion at constant rates $\\dot{M}_{\\rm acc}$ $\\la$ 0.1 M$_{\\odot}$ yr$^{-1}$, the maximum mass the star can reach is governed by radiative feedback and is $M_{\\rm *}$ $\\simeq$ 10$^{3}$ ($\\dot{M}_{\\rm acc}$/10$^{-3}$ M$_{\\odot}$ yr$^{-1}$)$^{\\frac{8}{7}}$ M$_{\\odot}$.  At higher masses, the H II region breaks out to large radii and terminates accretion.  We have verified that other forms of feedback, such as gas pressure, radiation pressure from trapped line emission, photodissociation of H$^-$, and scattering of Ly$\\alpha$ photons are much less effective at slowing accretion (see $\\S \\, 2.2$ and the Appendix).  At accretion rates above $\\ga$ 0.1 M$_{\\odot}$ yr$^{-1}$ the lifetime of the star, not radiative feedback, determines its final mass by limiting the time for which gas can accumulate on the star. Radiative feedback limits supermassive Pop III stars to final masses of $\\sim$ 3 $\\times$ 10$^5$ M$_{\\odot}$ and time-limited accretion limits them to $\\sim$ 10$^6$ M$_{\\odot}$ at the highest accretion rates ($\\sim$ 1 M$_{\\odot}$ yr$^{-1}$) found in numerical simulations of protogalactic collapse \\citep{wta08,sbh10,jlj11}.  We caution that our analytical calculations do not account for all conceivable effects that could stem the growth of supermassive stars.  For instance, \\citet{tm08} note that rotation of infalling gas leads to lower circumstellar densities and larger H~{\\sc ii} regions with greater radiative feedback \\citep[see also][]{hoyy11,stacy11}, implying lower final stellar masses.  Our spherically-symmetric calculation excludes rotation, so the mass limits we find for a given accretion rate (Fig. 3) are upper limits.  We do, however, find agreement with \\citep{oi02}, who performed a similar calculation, although they only considered the growth of stars up to $\\sim$ 10$^3$ M$_{\\odot}$.  Furthermore, we also note that the spherical symmetry and constant accretion in our models do not address accretion that is episodic or clumpy and self-shielded from ionizing radiation from the star \\citep{wet08b,wet10,krum09}.  These processes could cause accretion to proceed at lower time-averaged rates than those predicted here \\citep[but see][]{kuiper11}.  Another process that could truncate the growth of supermassive stars well before feedback or stellar lifetimes is the onset of a general relativistic instability in the core of the star that causes it to collapse when it becomes sufficiently massive \\citep{chandra64}.  Such instabilities are predicted to set in once the star has grown to $\\sim$ 10$^5$ M$_{\\odot}$ \\citep{iben63, fowler64}, but stellar rotation could stabilize it against collapse up to much larger masses \\citep{fowler66,bk98, baum99}.  While our models ignore rotation, the accreting gas is likely to have some \\citep[e.g.][]{col03}; if so, the supermassive star will inherit the angular momentum of the gas from which it formed and perhaps bypass the general relativistic instability.  Additional insight into the processes that limit the growth of supermassive stars will be gleaned from stellar evolution calculations accounting for the continual accretion of mass at high rates (Heger et al. in prep).  If supermassive stars formed and grew to masses of $\\ga$ 10$^5$ M$_{\\odot}$ in the early universe, recombination emission from their H~{\\sc ii} regions may be bright enough to be detected by future missions such as the \\textit{JWST}.  In particular, Ly$\\alpha$ photons trapped in the accretion flow are reprocessed into Balmer series photons that could escape into the IGM.  Consequently, the formation of these objects is accompanied by distinctive strong H$\\alpha$ emission together with strong continuum and He~{\\sc ii} $\\lambda$1640 emission, the latter arising from the hard spectrum of hot Pop III stars.  However, their black holes may be easier to discover in observational surveys, given the small numbers and brief lifetimes of their progenitors.  Indeed, these black holes may be the very ones that have already been found at the centers of massive galaxies and quasars at high redshifts.  Current numerical simulations of SMBH seed formation in early protogalaxies are now at an impasse because the formation of the hydrostatic supermassive protostar restricts Courant times to values that are too short to evolve central flows for even one dynamical time (but see Johnson et al. 2011 for an alternative approach using the sink particle technique).  Our results suggest that in the early and intermediate stages of the growth of the star its evolution is essentially decoupled from flows on even slightly larger spatial scales. Consequently, it should now be possible to retreat from the extreme spatial resolutions previously applied to the protostar and evolve flows at the center of the galaxy over enough dynamical times to capture its structure at the time of the death of the star and breakout of x-rays from the BH seed into the IGM. Thus, it will soon be possible to witness the births of the first quasars in the universe with supercomputers."
    },
    "1112/1112.1638_arXiv.txt": {
        "abstract": "\\noindent I show that cosmological bubble walls in the thin wall approximation are unstable to the creation of ``barnacles'' -- loci of different wall tension adjacent to regions filled with a third vacuum. Barnacle formation leads to the same observational consequences as the extensively studied bubble collision scenario, but occurs exponentially more often. The process is described by a saddle point of the thin wall action with two negative modes. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.4647_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{S:defineterms} The fine-tuning of the universe for intelligent life has received much attention in recent times. Beginning with the classic papers of \\citet{1974IAUS...63..291C} and \\citet{1979Natur.278..605C}, and the extensive discussion of \\citet{1986acp..book.....B}, a number of authors have noticed that very small changes in the laws, parameters and initial conditions of physics would result in a universe unable to evolve and support intelligent life.  We begin by defining our terms. We will refer to the laws of nature, initial conditions and physical constants of a particular universe as its \\emph{physics} for short. Conversely, we define a `universe' be a connected region of spacetime over which physics is effectively constant\\footnote{We may wish to stipulate that a given observer by definition only observes one universe. Such finer points will not effect our discussion}. The claim that the universe is fine-tuned can be formulated as: \\begin{quote} \\textbf{FT:} In the set of possible physics, the subset that permit the evolution of life is very small. \\end{quote}   FT can be understood as a counterfactual claim, that is, a claim about what would have been. Such claims are not uncommon in everyday life. For example, we can formulate the claim that Roger Federer would almost certainly defeat me in a game of tennis as: ``in the set of possible games of tennis between myself and Roger Federer, the set in which I win is extremely small''. This claim is undoubtedly true, even though none of the infinitely-many possible games has been played.  Our formulation of FT, however, is in obvious need of refinement. What determines the set of possible physics? Where exactly do we draw the line between ``universes''? How is ``smallness'' being measured? Are we considering only cases where the evolution of life is physically impossible or just extremely improbable? What is life? We will press on with the our formulation of FT as it stands, pausing to note its inadequacies when appropriate. As it stands, FT is precise enough to distinguish itself from a number of other claims for which it is often mistaken. FT is not the claim that this universe is optimal for life, that it contains the maximum amount of life per unit volume or per baryon, that carbon-based life is the only possible type of life, or that the only kinds of universes that support life are minor variations on this universe. These claims, true or false, are simply beside the point.  The reason why FT is an interesting claim is that it makes the existence of life in this universe appear to be something remarkable, something in need of explanation. The intuition here is that, if ours were the only universe, and if the causes that established the physics of our universe were indifferent to whether it would evolve life, then the chances of hitting upon a life-permitting universe are very small. As \\citet[][pg. 121]{leslieuniverses} notes, ``[a] chief reason for thinking that something stands in special need of explanation is that we actually glimpse some tidy way in which it might be explained''. Consider the following tidy explanations: \\begin{itemize} \\setlength{\\itemsep}{-2pt} \\item This universe is one of a large number of variegated universes, produced by physical processes that randomly scan through (a subset of) the set of possible physics. Eventually, a universe will be created that is a member of the life-permitting set. Only such universes can be observed, since only such universes contain observers. \\item There exists a transcendent, personal creator of the universe. This entity desires to create a universe in which other minds will be able to form. Thus, the entity chooses from the set of possibilities a universe which is foreseen to evolve intelligent life\\footnote{The counter-argument presented in Stenger's book (page 252), borrowing from a paper by Ikeda and Jeffreys, does not address this possibility. Rather, it argues against a deity which intervenes to sustain life in this universe. I have discussed this elsewhere: ikedajeff.notlong.com}. \\end{itemize} These scenarios are neither mutually exclusive nor exhaustive, but if either or both were true then we would have a tidy explanation of why our universe, against the odds, supports the evolution of life.  Our discussion of the multiverse will touch on the so-called anthropic principle, which we will formulate as follows: \\begin{quote} \\textbf{AP:} If observers observe anything, they will observe conditions that permit the existence of observers. \\end{quote} Tautological? Yes! The anthropic principle is best thought of as a selection effect. Selection effects occur whenever we observe a non-random sample of an underlying population. Such effects are well known to astronomers. An example is Malmquist bias --- in any survey of the distant universe, we will only observe objects that are bright enough to be detected by our telescope. This statement is tautological, but is nevertheless non-trivial. The penalty of ignoring Malmquist bias is a plague of spurious correlations. For example, it will seem that distant galaxies are on average intrinsically brighter than nearby ones.  A selection bias alone cannot explain anything. Consider the case of quasars. When first discovered, quasars were thought to be a strange new kind of star in our galaxy. \\citet{1963Natur.197.1040S} measured their redshift, showing that they were more than a million times further away than previously thought. It follows that they must be incredibly bright. The question that naturally arises is: how are quasars so luminous? The (best) answer is: because quasars are powered by gravitational energy released by matter falling into a super-massive black hole \\citep{1964SPhD....9..195Z,1969Natur.223..690L}. The answer is not: because otherwise we wouldn't see them. Noting that \\emph{if} we observe any object in the very distant universe \\emph{then} it must be very bright does not explain why we observe any distant objects at all. Similarly, AP cannot explain why life and its necessary conditions exist at all.  In anticipation of future sections, Table \\ref{T:symbols} defines some relevant physical quantities. \\begin{center} \\begin{table} \\begin{tabular}{| l | l | l | } \\hline Quantity \t\t\t\t\t& Symbol \t\t\t\t\t&  Value in our universe \\\\ \\hline Speed of light \t\t\t\t\t& $c$\t\t\t\t\t& 299792458 m s$^{-1}$ \\\\ Gravitational constant \t\t\t& $G$\t\t\t\t\t& 6.673 \\ten{-11} m$^3$ kg$^{-1}$ s$^{-2}$ \\\\ (Reduced) Planck constant \t\t& $\\hbar$ \t\t\t\t\t& 1.05457148 \\ten{-34} m$^2$ kg s$^{-2}$\t \\\\ Planck mass-energy\t\t\t\t& $m_\\ro{Pl} = \\sqrt{\\hbar c / G}$\t& 1.2209 \\ten{22} MeV \\\\ Mass of electron; proton; neutron \t& $m_\\ro{e}$; $m_\\ro{p}$; $m_\\ro{n}$ \t& 0.511; 938.3; 939.6 MeV \\\\ Mass of up; down; strange quark \t& $m_\\ro{u}$; $m_\\ro{d}$; $m_\\ro{s}$ \t& (Approx.) 2.4; 4.8; 104 MeV \\\\ Ratio of electron to proton mass\t& $\\beta$\t\t\t\t\t& $(1836.15)^{-1}$ \\\\ Gravitational coupling constant\t\t& $\\alpha_G = m_\\ro{p}^2 / m_\\ro{Pl}^2$ & 5.9 \\ten{-39} \\\\ Hypercharge coupling constant\t& $\\alpha_1$\t\t\t\t& $1/98.4$ \\\\ Weak coupling constant\t\t \t& $\\alpha_2$ \t\t\t\t& $1/29.6$ \\\\ Strong force coupling constant\t\t& $\\alpha_s = \\alpha_3$\t\t& 0.1187  \\\\ Fine structure constant\t\t\t& $\\alpha = \\frac{\\alpha_1 \\alpha_2}{\\alpha_1 + \\alpha_2}$ & 1/127.9 (1/137 at low energy) \\\\ Higgs vacuum expectation value\t& $v$\t\t\t\t\t& 246.2 GeV \\\\ QCD scale\t\t\t\t\t& \\lqcd\t\t\t\t\t& $\\approx 200$ MeV \\\\ Yukawa couplings\t\t\t\t& $\\Gamma_i = \\sqrt{2} m_i/v$\t& Listed in \\citet{2006PhRvD..73b3505T} \\\\ \\hline Hubble constant\t\t\t\t& $H$\t\t\t\t\t& 71 km/s/Mpc (today) \\\\ Cosmological constant (energy density) & $\\Lambda$ $(\\rho_\\Lambda)$\t\t& $\\rho_\\Lambda = (2.3 \\ten{-3} eV)^{-4}$ \\\\ Amplitude of primordial fluctuations & $Q$\t\t\t\t\t& $2 \\ten{-5}$ \\\\ Total matter mass per photon\t\t& $\\xi$\t\t\t\t\t&$ \\approx 4$ eV \\\\ Baryonic mass per photon\t\t& $\\xi_\\ro{baryon}$\t\t\t\t\t& $\\approx 0.61$ eV \\\\ \\hline \\end{tabular} \\caption{Fundamental and derived physical and cosmological parameters, using the definitions in \\citet{Burgess2006}. Many of these quantities are listed in \\citet{2006PhRvD..73b3505T}, \\citet[][Table A.2]{Burgess2006} and \\citet{Nakamura2010}. Unless otherwise noted, standard model coupling constants are evaluated at $m_Z$, the mass of the $Z$ particle, and hereafter we will use Planck units: $G = \\hbar = c = 1$, unless reintroduced for clarity. Note that often in the fine-tuning literature \\citep[e.g.][pg. 354]{1979Natur.278..605C,1986acp..book.....B}, the low energy weak coupling constant is defined as $\\alpha_w \\equiv G_F m^2_\\ro{e}$, where $G_F = 1/\\sqrt{2}v^2 = (292.8 \\ro{ GeV})^{-2}$ is the Fermi constant. Using the definition of the Yukawa coupling above, we can write this as $\\alpha_w = \\Gamma_\\ro{e}^2 / 2\\sqrt{2} \\approx 3 \\ten{-12}$. Note that this means that $\\alpha_w$ is independent of $\\alpha_2$.} \\label{T:symbols} \\end{table} \\end{center} ",
        "conclusions": "We conclude that the universe is fine-tuned for the existence of life. Of all the ways that the laws of nature, constants of physics and initial conditions of the universe could have been, only a very small subset permits the existence of intelligent life.  Will future progress in fundamental physics solve the problem of the fine-tuning of the universe for intelligent life, without the need for a multiverse? There are a few ways that this could happen. We could discover that the set of life-permitting universes is much larger than previously thought. This is unlikely, since the physics relevant to life is low-energy physics, and thus well-understood. Physics at the Planck scale will not rewrite the standard model of particle physics. It is sometimes objected that we do not have an adequate definition of `an observer', and we do not know all possible forms of life. This is reason for caution, but not a fatal flaw of fine-tuning. If the strong force were weaker, the periodic table would consist of only hydrogen. We do not need a rigorous definition of life to reasonably conclude that a universe with one chemical reaction (2H $\\rightarrow$ H$_2$) would not be able to create and sustain the complexity necessary for life.  Alternatively, we could discover that the set of possible universes is much smaller than we thought. This scenario is much more interesting. What if, when we really understand the laws of nature, we will realise that they could not have been different? We must be clear about the claim being made. If the claim is that the laws of nature are fixed by logical and mathematical necessity, then this is demonstrably wrong --- theoretical physicists find it rather easy to describe alternative universes that are free from logical contradiction \\citep[Davies, in][]{Manson2003}. The category of ``physically possible'' isn't much help either, as the laws of nature tell us what is physically possible, but not which laws are possible.  It is not true that fine-tuning \\emph{must} eventually yield to the relentless march of science. Fine-tuning is not a typical scientific problem, that is, a phenomenon in our universe that cannot be explained by our current understanding of physical laws. It is not a gap. Rather, we are concerned with the physical laws themselves. In particular, the anthropic coincidences are not like, say, the coincidence between inertial mass and gravitational mass in Newtonian gravity, which is a coincidence between two seemingly independent physical quantities. Anthropic coincidences, on the other hand, involve a happy consonance between a physical quantity and the requirements of complex, embodied intelligent life. The anthropic coincidences are so arresting because we are accustomed to thinking of physical laws and initial conditions as being unconcerned with how things turn out. Physical laws are material and efficient causes, not final causes. There is, then, no reason to think that future progress in physics will render a life-permitting universe inevitable. When physics is finished, when \\emph{the} equation is written on the blackboard and fundamental physics has gone as deep as it can go, fine-tuning may remain, basic and irreducible.  Perhaps the most optimistic scenario is that we will eventually discover a simple, beautiful physical principle from which we can derive a unique physical theory, whose unique solution describes the universe as we know it, including the standard model, quantum gravity, and (dare we hope) the initial conditions of cosmology. While this has been the dream of physicists for centuries, there is not the slightest bit of evidence that this idea is true. It is almost certainly not true of our best hope for a theory of quantum gravity, string theory, which has ``anthropic principle written all over it'' \\citep{Schellekens2008}. The beauty of its principles has not saved us from the complexity and contingency of the solutions to its equations. Beauty and simplicity are not necessity.  Finally, it would be the ultimate anthropic coincidence if beauty and complexity in the mathematical principles of the fundamental theory of physics produced all the necessary low-energy conditions for intelligent life. This point has been made by a number of authors, e.g. \\citet{1979Natur.278..605C} and \\citet{Aguirre2005}. Here is \\citet{Wilczek2006}: \\begin{quote} ``It is logically possible that parameters determined uniquely by abstract theoretical principles just happen to exhibit all the apparent fine-tunings required to produce, by a lucky coincidence, a universe containing complex structures. But that, I think, really strains credulity.'' \\end{quote}  \\appendix"
    },
    "1112/1112.3067_arXiv.txt": {
        "abstract": "Nighttime visible-light sky brightness and transparency are reported for the Polar Environment Research Laboratory (PEARL), located on a 610-m high ridge near the Eureka research station, on Ellesmere Island, Canada. Photometry of Polaris obtained in $V$ band with the PEARL All Sky Imager (PASI) over two winters is supported by standard meteorological measurements and visual estimates of sky conditions from sea level. These data show that during the period of the study, October through March of 2008/09 and 2009/10, the sky near zenith had a mean surface brightness of $19.7~{\\rm mag}~{\\rm arcsec}^{-2}$ when the sun was more than $12\\arcdeg$ below the horizon, reaching $20.7~{\\rm mag}~{\\rm arcsec}^{-2}$ during astronomical darkness with no moon. Skies were without thick cloud and potentially usable for astronomy 86\\% of the time (extinction~$<2$ mag). Up to 68\\% of the time was spectroscopic ($\\leq 0.5$ mag), attenuated by ice crystals, or clear with stable atmospheric transparency. Those conditions can persist for over 100 hours at a time. Further analysis suggests the sky was entirely free of ice crystals (truly photometric) $48\\pm3$\\% of the time at PEARL in winter, and that a higher elevation location nearby may be better. ",
        "introduction": "Winters in the High Arctic are cold, dry, and dark. In northern Canada, within the Territory of Nunavut, the coasts of the ice-locked eastern arctic archipelago combine these conditions with high terrain, providing locations potentially above much of the surface cloud and turbulence, and leading to the prospect of good observing sites for optical/near-infrared astronomy \\citep[See][]{Steinbring2010}. One such site under study is the location of the Polar Environment Atmospheric Research Laboratory (PEARL) at $80\\arcdeg$ north latitude, on Ellesmere Island. The PEARL facility is at an elevation 610 m on a ridge at the tip of the Fosheim Peninsula. It is accessible via a 15 km long road from the sea-level research base at Eureka, providing reliable logistical support: power, broadband satellite communications, an all-weather airstrip, and sea access in summer. Eureka is operated by the civilian weather service of the Canadian government, Environment Canada (EC), providing hourly meteorological data (air temperature, barometric pressure, relative humidity, wind speed and direction) and visual cloud-cover estimates.  With the onset of continuous polar night in early October, air temperature drops quickly at Eureka. It reaches an average of $-30$C within a few weeks, and by December is typically near $-40$C. A strong thermal inversion layer develops, with a peak $10$C warmer than sea level at a median height near 800 m, although it can be below 600 m elevation (lower quartile) and above 1200 m elevation (upper quartile), based on a 50-year climatology \\citep{Lesins2009a}. It is already well known from visual sky-quality estimates obtained at Eureka that under these conditions skies are infrequently obscured by thick cloud. The mean precipitable water vapor column at sea level drops below 2 mm, freezing out into the form of ice crystals. Most often (just over 50\\% of the time) a layer of this ``diamond dust\" is suspended within the inversion layer, although it can precipitate from otherwise clear skies. As seen from PEARL, under these conditions Eureka appears blanketed in haze. Looking upwards, the opacity of crystals can be thin, producing visible halos for the moon and bright stars.  Observations with millimetre-wave radar combined with Laser Detection and Ranging (LIDAR) from sea level near Eureka provide the best available data for characterizing the size and vertical distributions of these boundary-layer ice crystals, differentiating them from mixed-phase water and ice clouds at higher altitudes by measuring their returned relative backscatter cross-sections and depolarization \\citep{Bourdages2009}. These show ice crystals are typically 15-70 $\\mu$m in cross section, and have vertical density distribution which decays exponentially with height (linearly with air pressure), dropping by two orders of magnitude from sea level to 1200 m \\citep[Figure 7 in][]{Bourdages2009}. This is a reduction by a factor of 10 by 600 m, with a further factor of 2 decrease at 800 m elevation. That also correponds to roughly the elevations of the highest local terrain near Eureka, and events of wind-blown snow being the dominant component have been recorded \\citep{Lesins2009b}. In fact, attaining the highest practical elevation, and rising above much of the ice crystal layer was an important aspect of siting PEARL, as its primary purpose is for the observation of chemical and dynamical signatures in the middle and upper atmosphere by optical means. An issue for astronomy though is the opacity of ice crystals and the fraction of time that the atmosphere above PEARL is subject to thin attenuation by any other contributors - one possibility being so-called ``arctic haze\", aerosol pollutants transported from southern latitudes \\citep[See][and references therein]{Quinn2007}. Quantification of transparency and sky brightness using an all-sky camera is desirable for comparison with other sites worldwide. Such an analysis for Mauna Kea using the CFHT SkyProbe was reported in \\cite{Steinbring2009}.  The PEARL All Sky Imager (PASI) was deployed in 2007, and is suitable for an analysis similar to that of the Mauna Kea SkyProbe study. PASI \\citep[See][for details of the instrument]{Veselinovic2008} was designed for the detection of atmospheric emission lines in airglow; primarily for the purpose of characterizing atmospheric waves (primarily tides and buoyancy waves - typically termed ``gravity waves\" in the atmospheric community) and aurora, which can appear near the horizon, and sometimes overhead. It is an autonomous imager composed of an up-looking $1024\\times1024$ CCD camera and a fisheye lens, yielding approximately square-$0.5\\arcdeg$ pixels near zenith, viewing through a Plexiglass dome-shaped window in the ceiling of a warm room on the PEARL roof. Color information is provided by cycling through a filter wheel containing 5 narrowband (10~\\AA-wide) filters. Although not selected for the purposes of astronomy, one of these has a central wavelength of 5725~\\AA, comparable to $V$ band. This region was chosen to avoid airglow and aurora, providing a background sky measurement once every 15 minutes. This provides a useful cadence and a sufficiently deep exposure to image a few bright stars, including $\\alpha$ Ursae Minoris (Polaris) which is just 10 degrees off zenith when viewed from Eureka.  This paper outlines our study of sky transparency and sky brightness at PEARL based on PASI measurements of extinction estimated from photometry of Polaris, correlated with meteorological and visual sky quality estimates from Eureka. The data and their reduction is outlined in Section~\\ref{observations}, with further analysis of the fraction of time free of ice crystals follows in Section~\\ref{analysis}, followed by a summary and conclusions in Section~\\ref{summary}. ",
        "conclusions": "\\label{summary}  We have presented an analysis of stellar photometry using an all-sky camera at PEARL and supported by meteorological and visual sky-quality estimates from Eureka, over two consecutive winters beginning in 2008. The observed PASI lightcurve of Polaris is of sufficient quality to discern its 28 millimag, 4 day periodicity, and provides two basic outcomes. First, it confirms skies are dark in the visible, $19.7~{\\rm mag}~{\\rm arcsec}^{-2}$ on average for all sun elevations below $-12\\arcdeg$, and reaching $20.7~{\\rm mag}~{\\rm arcsec}^{-2}$ for sun elevations below $-18\\arcdeg$ when there is no moon. Further observations in standard filters would be helpful to characterize sky brightness in the infrared, and due to the contribution of scattered artificial light from the facility itself or the nearby Eureka base. Second, PEARL is at a high enough elevation to be above the bulk of the ice crystal boundary layer in winter. For 86\\% of the time this has an extinction less than 2 mag, suitable for astronomical observations. Up to 68\\% of the time extinctions are less than 0.5 mag. This spectroscopic/clear-sky fraction is superior to higher elevation sites in southern Canada, for example, those reported (by different methods) for Mount Kobau \\citep[60.5\\%;][]{Brosterhus1972} and similar mountains in the southern British Columbia interior \\citep{Hickson1989}. That these conditions can persist for over 100 hours at a time allows for observations that would be impossible from a single mid-latitude site (not part of a global network).  Further analysis of the data show that the truly photometric fraction of time (without any cloud) is $48\\pm3$\\%, which may be comparable to some of the best sites worldwide; Dome A in Antarctica, and Mauna Kea in Hawaii. The attenuation by air only, in instrumental magnitudes, was $0.12\\pm0.03$ mag, which suggests only a small contribution from arctic haze is possible. The uncertainties are limited by photometry with PASI. Best conditions at PEARL are associated with calm winds. A possible trend in extinction with barometric pressure suggests that an increase in elevation may appreciably reduce optical depth of ice crystals. According to this analysis, an elevation increase of 225 m would effectively ensure skies are clear when winds are calm in Eureka. One location nearby Eureka in principle providing such an elevation advantage is the top of Blacktop Ridge, and so this may provide encouragement to further explore high terrain near Eureka for candidate sites."
    },
    "1112/1112.4192_arXiv.txt": {
        "abstract": "{ We present new global fits of the constrained Minimal Supersymmetric Standard Model (cMSSM), including LHC 1/fb integrated luminosity SUSY exclusion limits, recent LHC 5/fb constraints on the mass of the Higgs boson and XENON100 direct detection data. Our analysis fully takes into account astrophysical and hadronic uncertainties that enter the analysis when translating direct detection limits into constraints on the cMSSM parameter space. We provide results for both a Bayesian and a Frequentist statistical analysis. We find that LHC 2011 constraints in combination with XENON100 data can rule out a significant portion of the cMSSM parameter space. Our results further emphasise the complementarity of collider experiments and direct detection searches in constraining extensions of Standard Model physics. The LHC 2011 exclusion limit strongly impacts on low-mass regions of cMSSM parameter space, such as the stau co-annihilation region, while direct detection data can rule out regions of high SUSY masses, such as the Focus-Point region, which is unreachable for the LHC in the near future. We show that, in addition to XENON100 data, the experimental constraint on the anomalous magnetic moment of the muon plays a dominant role in disfavouring large scalar and gaugino masses. We find that, should the LHC 2011 excess hinting towards a Higgs boson at 126 GeV be confirmed, currently favoured regions of the cMSSM parameter space will be robustly ruled out from both a Bayesian and a profile likelihood statistical perspective. } ",
        "introduction": "\\label{secintro}  The CMS and ATLAS collaborations have recently published new results for the search for Supersymmetry (SUSY) \\cite{arXiv:1111.4820,arXiv:1109.2352,arXiv:1109.6572}, based on proton-proton collisions with a center-of-mass energy $\\sqrt{s} = 7$ TeV, and an integrated luminosity of 1 fb$^{-1}$, recorded during 2011. No significant excess above the Standard Model (SM) predictions was measured, so that new constraints on SUSY could be derived, in particular for what concerns the cMSSM (constrained Minimal Supersymmetric Standard Model, see Refs. \\cite{Chamseddine:1982jx,CTP-TAMU-24-92,RAL-92-005,hep-ph/9311269,Kane:1993td} and references therein), a simple construction of minimal SUSY widely studied in the literature, which has become the {\\em de facto} default model to evaluate SUSY searches and, often, prospects for dark matter searches \\cite{Bertone:2010zz,Bergstrom:2000pn,Munoz:2003gx,Bertone:2004pz}. Additionally, an update on the search for the SM Higgs boson based on 5 fb$^{-1}$ of data has recently been announced by the ATLAS and CMS experiments \\cite{Higgs_results}. In particular, the ATLAS experiment has reported a new 95$\\%$ exclusion region constraining the allowed mass range for the SM Higgs boson to $115.5 - 131.0$ GeV, as well as a 2.3$\\sigma$ excess at a mass of $126$ GeV.   In Ref.~\\cite{arXiv:1107.1715}, we presented global fits of the cMSSM including data from precision tests of the Standard Model (SM), constraints from cosmology on the dark matter relic abundance, LEP bounds, and constraints on SUSY from LHC 2010 data. Additionally we studied the impact of recent exclusion limit on the WIMP-proton spin-independent scattering cross-section $\\sigmaSI$ from the XENON100 experiment \\cite{xenon:2011hi}, a two-phase time projection chamber with a liquid xenon target, on the cMSSM parameter space. Similar studies of the impact of LHC  and XENON100 data on global fits of the cMSSM and other SUSY models are found in Refs.~\\cite{arXiv:1106.2529,arXiv:1104.3572,arXiv:1105.5162,arXiv:1111.6098,arXiv:1110.3568}. Our approach differs from other studies, because we provide results for both a Bayesian and frequentist statistical analysis, and included for the first time the relevant uncertainties (astrophysical and related to hadronic matrix elements) for the analysis of direct detection data.  In this work we update the global fits of the cMSSM presented in Ref.~\\cite{arXiv:1107.1715} to include LHC 1 fb$^{-1}$ exclusion limits and LHC 5 fb$^{-1}$ constraints on the Higgs boson (See also Refs.~\\cite{arXiv:1104.3572,arXiv:1111.6098,arXiv:1110.3568} for a discussion of the impact of LHC 2011 SUSY searches on the cMSSM, as well as on more general SUSY models. See Ref.~\\cite{Buchmueller:2011ab} for a study of the impact of a hypothetical Higgs discovery at 125 GeV on the cMSSM parameter space). In our analysis, we apply all the experimental bounds used in Ref.~\\cite{arXiv:1107.1715}, most notably constraints on $\\sigmaSI$ derived from the XENON100 experiment. We use nested sampling as a scanning algorithm \\cite{Feroz:2007kg} and study both the resulting Bayesian posterior distributions and the frequentist best fit regions. The methodology used has been described in detail in previous papers  \\cite{deAustri:2006pe,Trotta:2008bp,Feroz:2011bj,arXiv:1107.1715}, in particular we closely follow the procedure presented in~\\cite{arXiv:1107.1715}. We do not repeat the description of the methodology here and instead focus on the discussion of new physical results from the tighter LHC constraints. For further details on the statistical and theoretical background of our approach we refer the reader to Ref.~\\cite{arXiv:1107.1715} and references therein.  Constraints on the WIMP mass and spin-independent cross-section derived from direct detection data strongly depend on the parameters describing the dark matter halo model, namely, the WIMP velocity distribution and the local density~\\cite{ST}. Additionally, constraints derived on the cMSSM parameters from the XENON100 exclusion limit depend on the hadronic matrix elements, which parameterise the contributions of light quarks to the proton composition. None of these quantities is precisely known, and neglecting the uncertainty on these quantities in the modelling can lead to incorrect inferences for the cMSSM parameters. Our inference procedure consistently takes into account our (lack of) knowledge of these parameters. We account for both astrophysical and nuclear physics uncertainties entering in the calculation of the event rate in direct detection experiments, by treating them as nuisance parameters in the analysis, which in the end are marginalised/maximised over to correctly infer the cMSSM parameters.  This paper is organised as follows. In Section 2 we present the theoretical and statistical framework for the analysis. In Section 3 we present the results, Section 4 is devoted to the discussion of the prospects for discovering SUSY in the cMSSM in different detection channels. We present our conclusions in Section 5. ",
        "conclusions": "\\label{secconclusion}  In this work we have presented new global fits of the cMSSM, including exclusion limits from LHC 2011 SUSY and Higgs searches and XENON100 direct detection data. In addition to uncertainties on Standard Model quantities, our analysis takes into account both astrophysical and hadronic uncertainties that enter the analysis when translating direct detection limits into constraints on the cMSSM parameter space. XENON100 data in combination with LHC 2011 constraints can rule out a significant portion of the cMSSM parameter space. In addition to the FP region being ruled out at the $99\\%$ level by XENON100 data, we have shown that the size of the SC region is significantly reduced by LHC 2011 constraints on $m_0$ and $m_{1/2}$. LHC 2011 data strongly impacts on the cMSSM parameter space, pushing allowed particle masses to higher values and reducing the favoured values of both spin-dependent and  spin-independent neutralino-proton cross-sections.  This has important implications for direct detection of the cMSSM, with LHC constraints cutting into the low $m_{\\neut}$ region of cMSSM parameter space and, in combination with direct detection constraints, pushing the contours to lower $\\sigmaSI$. Our results further emphasise the complementarity of collider experiments and direct detection searches. While LHC data are placing tight constraints on large portions of cMSSM parameter space, direct detection data is crucial to rule out regions of high SUSY masses, such as the FP region, which is unreachable for the LHC in the near future (see also the discussion in Ref. \\cite{Bertone:2011kb}).  We further investigated the importance of the experimental constraint on the anomalous magnetic moment of the muon. We found that this constraint is the dominating reason for large scalar and gaugino masses being disfavoured in our analysis. All other constraints (aside from XENON100) leave this range of masses fairly unconstrained.  Finally, we investigated constraints on the lightest Higgs mass resulting from LHC 2011 and XENON100 data. $m_h$ is pushed towards larger values due to the new LHC 2011 95$\\%$ confidence level. We have qualitatively shown that, should the excess around $m_h \\sim 126$ GeV be confirmed by future data sets, the residual cMSSM parameter space will be strongly disfavoured.  While this work was under completion, similar analyses of the constraints on the cMSSM parameter space resulting from LHC 2011 SUSY searches and XENON100 data were performed in Ref.~\\cite{arXiv:1111.6098,arXiv:1110.3568}. Our conclusions are qualitatively similar to Ref.~\\cite{arXiv:1110.3568}, while in Ref.~\\cite{arXiv:1111.6098} it was claimed that the FP region can not be excluded at $95\\%$ level by XENON100 data and a negative outlook on direct detection prospects of the cMSSM was given. This is a result of a very conservative estimation of the XENON100 exclusion limit in Ref.~\\cite{arXiv:1111.6098}. Our analysis differs from Ref.~\\cite{arXiv:1110.3568}, where astrophysical uncertainties are not included in the analysis and results are presented for the profile likelihood only. In contrast, we provide a detailed comparison between Bayesian and profile likelihood results. Ref.~\\cite{arXiv:1111.6098} provides a similar comparison, however, the scanning parameters chosen are not suitable for an accurate mapping of the profile likelihood function (according to the prescriptions given in Ref.~\\cite{Feroz:2011bj}), so that the extent of the resulting confidence levels may vary strongly from scan to scan. Additionally, in Ref.~\\cite{arXiv:1111.6098} the impact of astrophysical and hadronic uncertainties is taken into account by smearing out the XENON100 exclusion limit, instead of marginalising/maximising over the corresponding nuisance parameters. In contrast to both Ref.~\\cite{arXiv:1111.6098} and Ref.~\\cite{arXiv:1110.3568} we include recent 5 fb$^{-1}$ LHC constraints on the Higgs mass, so that we provide a more sophisticated, up-to-date statistical analysis of the cMSSM parameter space."
    },
    "1112/1112.1853_arXiv.txt": {
        "abstract": "We discuss the properties of neutron stars and their modifications due to the occurrence of hyperons and quarks in the core of the star. More specifically, we consider the general problem of exotic particles inside compact stars in light of the observed two-solar mass pulsar. In addition, we investigate neutron star cooling and a possible explanation of the recently measured cooling curve of the neutron star in the supernova remnant Cas A. ",
        "introduction": "The study of strong interaction physics under extreme conditions is at the center of many experimental and theoretical efforts in nuclear physics. Ultrarelativistic heavy-ion collisions allow for the investigation of the phase structure of QCD at high temperatures. Here one assumes that the conditions in the fireball of the collision zone generate chirally restored and deconfined matter, whose properties can be deduced by analyzing observables like particle multiplicities and the collective flow. At the other end of the QCD phase diagram, the investigation of neutron stars is the main tool to understand extremely dense and  cold matter. As the theoretical approach of using lattice gauge simulations is not applicable in the high density/chemical potential regime one has to rely on model descriptions of the dense hadronic and possibly quark matter. In the following we study such a model that includes hadronic flavor SU(3)  particles, as well as quarks, in an extension of this approach. ",
        "conclusions": ""
    },
    "1112/1112.0314_arXiv.txt": {
        "abstract": "We present a novel, virtual X--ray observatory designed to obtain synthetic observations from hydro--numerical simulations, named \\name{}. In particular, we provide a description of the code constituting the photon simulator and of the new approach implemented. We apply \\name{} to simulated galaxy clusters in order to demonstrate its capabilities. In fact, X--ray observations of clusters of galaxies continue to provide us with an increasingly detailed picture of their structure and of the underlying physical phenomena governing the gaseous component, which dominates their baryonic content. Therefore, it is fundamental to find the most direct and faithful way to compare such observational data with hydrodynamical simulations of cluster--like objects, which can currently include various complex physical processes. Here, we present and analyse synthetic {\\it Suzaku} observations of two cluster--size haloes obtained by processing with \\name{} the hydrodynamical simulation of the large--scale, filament--like region in which they reside. Taking advantage of the simulated data, we test the results inferred from the X--ray analysis of the mock observations against the underlying, known solution. Remarkably, we are able to recover the theoretical temperature distribution of the two haloes by means of the multi--temperature fitting of the synthetic spectra. Moreover, the shapes of the reconstructed distributions allow us to trace the different thermal structure that distinguishes the dynamical state of the two haloes. ",
        "introduction": "\\label{sec:intro}  As optimal laboratories for both cosmology and astrophysics, galaxy clusters have been thoroughly investigated with large surveys as well as via dedicated single--object observations. In this respect, X--ray measurements still represent one of the best ways to fully study their structure, traced by the hot plasma filling their potential wells. In particular, the most important and difficult quantity to infer is the total gravitating mass, which can be determined from measurements of the intra--cluster medium (ICM) temperature and density profiles, assuming hydrostatic equilibrium to hold. So far, studies of the X--ray--emitting ICM have been mainly limited to the innermost region, where the investigation of the baryonic physics is observationally less challenging than in the outskirts. Although very difficult to map in the X rays, a characterization of the outer region of clusters is in fact crucial to understand the formation and evolution of this structures and to use them as cosmological probes, since it encloses a significant fraction of the cluster volume, where the properties of the accreting gas and of the dark matter halo are still uncertain \\cite[e.g., see][]{ettori2011}. Only very recently, X--ray observations performed with the Japanese satellite {\\it Suzaku} have reached the virial radius for a few clusters of galaxies \\cite[e.g.][]{fujita2008,bautz2009,george2009,reiprich2009,hoshino2010,kawaharada2010,simionescu2011,akamatsu2011}, promisingly taking a step forward on the temperature profile debate and indirectly opening new perspectives in controlling the biases on cluster--based cosmological investigations. These results are just one example of the improvements by which X--ray observations continue to provide us with an increasingly detailed picture of galaxy clusters, for which a clear interpretation of the underlying physical processes is very challenging.  On the theoretical side, many interesting issues on cluster formation and evolution are addressed by means of hydro--numerical simulations, which represent a powerful tool to investigate in detail the nature of these complicated astrophysical objects \\cite[e.g.][for a recent review]{borgani2009}. Simulations of large cosmological boxes that include several massive cluster--like haloes can now be performed incorporating not only the dominating dark matter component but also baryonic matter, in the form of gas-- and star--like particles (or cells), governed by hydrodynamical processes. With particular concern for galaxy clusters, current simulations are not only able to account for basic gas hydrodynamics, but implement also more complicated models for star formation from multi--phase medium \\cite[e.g.][]{katz1992,katz1996,springel2003,marri2003} and thermal or kinetic feedback from supernovae--driven winds \\cite[][]{navarro1993,scannapieco2006,dallavecchia2008}, chemical enrichment, metal and molecule cooling \\cite[e.g.][]{mosconi2001,yoshida2003,tornatore2004,tornatore2007,scannapieco2005,maio2007,maio2010}, thermal conduction \\cite[][]{cleary1999,jubelgas2004,ruszkowski2010}, AGN feedback \\cite[e.g.][]{springel2005b,dimatteo2005,sijacki2006,sijacki2007,sijacki2008,puchwein2008,fabjan2010,dubois2010,teyssier2011}, cosmic rays \\cite[][]{pfrommer2007,jubelgas2008} and magnetic fields \\cite[][]{phillips1985,dolag1999,brueggen2005,price2005,dolag2009}, to name the most important effects.  The ideal achievement would be the combination of both these sources of information, by directly comparing simulated clusters to X--ray observations of real objects. Devoted to this goal, sophisticated numerical codes, such as X--MAS/X--MAS2 \\cite[][]{gardini2004,rasia2008} and, more recently, XIM \\cite[][]{xim2009,heinz2010}, have been developed in the last years in order to turn hydro--simulation outputs into mock images with a given X--ray telescope. In particular, the X--MAS virtual telescope is explicitly designed to process outputs obtained from SPH codes like GADGET, whereas XIM is particularly dedicated to grid--based hydrodynamical simulations, e.g. performed with the AMR code FLASH \\cite[][]{fryxell2000}, which makes them fundamentally complementary. Similarly, both X--MAS and XIM use a plasma thermal emission code by which they calculate the projected emission associated to the gas component in the simulation. The computational effort, required to calculate this, is usually reduced by interpolating among the model spectra externally stored in a library of templates. Essentially, the emissivity integrated along the chosen line of sight is calculated in terms of flux, depending on the properties of all the gas at a given projected position. Lastly, the convolution with instrumental response is performed and the final photon event file is generated. Such virtual telescopes represent an important step forward with respect to visualization tools used to produce surface brightness maps from simulation outputs, since they properly calculate the emission associated to the gas component accounting for the structure of the simulated source along the line of sight. Moreover, they are able to include accurate calculations required to obtain realistic mock X--ray observations, for instance convolving the simulated spectra with any given instrument response and telescope PSF. In fact, it is vital that the synthetic observations produced by such simulators are as much as possible similar to the standards commonly used by observers, so that the comparison can be the most faithful. As such, this advanced approach easily leads to a challenging increase of the parameters to handle in order to realistically model the plasma emission and the detailed three--dimensional structure of the simulations.  Analogous in the scope, the virtual telescope presented in this paper, \\name{}, is also dedicated to convert hydro--numerical outputs into mock X--ray observations. However, the novelty of our simulator lies in the method adopted, by which the spectral emission calculated singularly for each gas element in the simulation is immediately converted into a discrete sample of photons, collected and stored before projecting along any line of sight and convolving with any desired instrument. With this strategy we are able to significantly gain in computational effort, since it requires to process the original simulation only once, independently of the specific synthetic observation to be performed afterwards. Furthermore, the guarantee for high spatial and spectral resolution, preserved without dramatically increasing the computational cost, offers the possibility to anticipate the observational achievements of upcoming X--ray missions, such as {\\it IXO/ATHENA}, in which the high--resolution spectroscopy will allow us to explore the intrinsic structure of galaxy clusters through the study of their spectral features. In fact, it has been suggested in the last decade that the diagnostics of the broadening of heavy--ion emission lines in highly--resolved X--ray spectra of galaxy clusters could actually provide valuable information about the underlying structure of the ICM and its velocity field (see \\cite{pawl2005}, for a preliminary, interesting work in this field), useful for the dynamical classification of these objects. The expectations for such line diagnostics are related in particular to the most prominent emission line in X--ray spectra, namely the $\\sim 6.7 \\kev$ line from helium--like iron. In fact, the large atomic mass of the FeXXV ion significantly reduces the thermal line broadening, so that the line width turns out to be definitely more sensitive to turbulent or bulk gas motions \\cite[][]{inogamov2003,sunyaev2003}. This kind of studies can eventually help constraining the detectability of non--thermal motions that are likely to establish in the ICM \\cite[e.g.][]{rebusco2008} and can compromise X--ray mass measurements \\cite[][]{rasia2006,fang2009,lau2009,biffi2011}.  The paper is structured as follows: \\sec\\ref{sec:method} is devoted to the description of the method implemented in the photon simulator and the fundamental units of the code. In \\sec\\ref{sec:application} we apply \\name{} to the hydrodynamical simulation of a filament--like structure, describing in detail each phase, from the photon generation till the mock {\\it Suzaku} observation of two massive galaxy clusters residing in the filament and the spectral analysis of the synthetic spectra. Results about the recovering of the ICM emission measure distribution of the two cluster--like haloes from the multi--temperature fitting are presented in \\sec\\ref{sec:results}. Finally, we discuss the results obtained, summarise and draw our conclusions in \\sec\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} In this paper, we have presented a novel X--ray photon simulator, named \\name{}\\footnote{The source code of the presented version of \\name{} is made available by request to the authors (see http://www.mpa-garching.mpg.de/~kdolag/Phox/). We also plan to make the code available directly for anonymous download in the near future.}, dedicated to obtain synthetic X--ray observations from the output of hydrodynamical numerical simulations.  The development of \\name{} has been strongly motivated by the ultimate aim of comparing the output of hydrodynamical simulations of cluster--like objects to real X--ray observations in the most faithful and reliable way possible. In fact, X--ray observations still provide us with increasingly detailed data that require a deep understanding of the underlying structure of galaxy clusters. Numerical simulations, on the other side, continuously improve in describing not only the dark matter but also the complex physical processes governing the baryonic component. Therefore, a strategy devoted to combine the advances and results in both observations and simulations can be promisingly successful.  In the past, software packages like X--MAS \\cite[][]{gardini2004} or XIM \\cite[][]{xim2009} have been developed to produce synthetic observations of astrophysical sources simulated with particle-- and grid--based codes, respectively. The existence of such tools is definitely complementary to our simulator in order to test the robustness of comparison studies between state--of--the--art hydrodynamical simulations including different physical phenomena and real observations. With respect to other codes, the novelty of our virtual telescope resides in the original method implemented. In particular, the order of the three main units that constitute the code offers a significant gain in computational cost and time, making \\name{} extremely flexible and useful to process hydro--numerical simulations of several astrophysical objects, at various spatial scales, as well as simulated light--cones. Moreover, this method uniquely guarantees extremely high spatial and energy resolution. The more demanding phase is indeed the first one, where the output of a hydrodynamical simulation is processed and the X--ray emission of each gas element (either particle or grid cell) is simulated by using the spectral models available in the X--ray software \\xspec{}. Essentially, we perform the sampling of each spectrum with an expected number of photons, and we immediately store the sample of photons insted of constructing a library of template spectra to be populated afterwards. As such, the spectral model chosen to mock up the emission from the gas elements, regardless of its complexity, can be explicitly specified and suited to represent any desired astrophysical source. Also, the relatively reasonable requirements in terms of computational time and memory, make out of \\name{} an ideal tool to process forthcoming large cosmological, hydrodynamical simulations as well as the usual smaller ones.  Concerning galaxy clusters, \\name{} currently uses two possible models for the X--ray emission of the ICM (namely, \\mekal{} and \\apec{}) and accounts for the particular density and temperature of the simulated gas elements. However, the technique implemented offers the chance to account also for different metallicities, or even metal compositions, across the emitting gas elements, without increasing the memory requirements significantly.  Furthermore, \\name{} postpones the projection along the line of sight and the convolution with a real instrument response till after the generation of the discrete sample of photons, permitting to each of them, in principle, to contribute to the final integrated spectrum. Requiring a high energy resolution in the model spectrum associated to every gas element, we can ideally combine the potential of \\name{} with the characteristics of upcoming X--ray satellites, like IXO/ATHENA or ASTRO--H, in order to predict the detectability of fundamental spectral features. This could be of great help, for instance, in constraining the ICM non--thermal motions (e.g. turbulent, bulk, rotational motions) via the velocity broadening of heavy--ion emission lines. In general, the theoretical modelling of X--ray--emitting sources is a powerful way to predict and optimize the capabilities of future X--ray missions and to help in the interpretation of the data when the instruments are in operation.  As an interesting application of this new virtual telescope, we focused on the hydro--numerical simulation of a large--scale region, originally performed with the TreePM/SPH code GADGET--2 and then re--simulated including cooling, star formation and feedback from supernova winds. This filament--like structure is particularly worthy for presenting the simulator in each part and especially for demonstrating the advantages of our approach. From a purely technical perspective, the greatest gain was achieved in terms of computational cost since the entire, large data cube contained in the simulation output was processed with \\name{} Unit 1 only once, providing an output of relatively compacted photon data (see \\tab\\ref{table:timing}). In fact, this first phase permitted to obtain the X--ray photons virtually emitted by all the gas particles in the filament. At a later stage, we restricted our analysis to two cluster--like haloes with different thermodynamical structure, \\leiaD{} and \\leiaB{}, which were geometrically selected from two different regions of the simulation box. Conveniently, the projection and observation units of \\name{} could be run separately for each halo processing the same output generated with Unit 1 and we finally obtained synthetic {\\it Suzaku} images and spectra (\\fig\\ref{fig:xisimg}). The convolution with the response of the X--ray Imaging Spectrometer (XIS) on board {\\it Suzaku}, via the public tool \\xissim{}, provided data that could be analysed with the standard procedures of X--ray analysis. An important aim of this application was the reconstruction of the intrinsic emission measure distribution by the multi--temperature fitting of the simulated spectra. The goal of setting spectroscopic constraints on the multi--phase ICM of galaxy clusters was pursued in X--ray observational studies reported by \\cite{peterson2003}, \\cite{kaastra2004} and \\cite{fabian2006}, although for real clusters it is intrinsically impossible to know the underlying multi--phase structure unambiguously. Instead, numerical simulations uniquely allow to model and predict the theoretical temperature distribution of the ICM \\cite[e.g., studied in][]{kawahara2007,kawahara2008} and constitute therefore the ideal possibility to test the power of the method used to reconstruct the structure of the mock data. The direct comparison between the X--ray data and the theoretical thermal distribution of the input simulation is found to be overall faithful and robust for both \\leiaD{} and \\leiaB{} (see \\fig\\ref{fig:leiaDBem}). Moreover, the different thermal structure that distinguishes the two selected haloes is remarkably highlighted by the multi--temperature fitting. This result promisingly suggests that the temperature distribution of the ICM can be in principle traced and unveiled by spectral analysis of X--ray observations. From the additional investigation of velocity--broadened fitting models (like the \\bapec{} model) for the {\\it Suzaku} spectra, we can also conclude that a higher spectral resolution is required in order to reliably constrain the ICM velocity structure. In general, it is important to notice that the application discussed here does not include other emission components (e.g. a physical X--ray background) in addition to the ICM one and represents therefore an ideal case study.  With regard to technique, the main feature distinguishing \\name{} from other virtual X-ray telescopes, such as XMAS and XIM, is the explicit creation and storage of a 3--dimensional box of photons from the raw simulation data cube itself. This step is performed {\\it before} and without any reference to the observation process and its result is stored as a separate data product, which represents the ideal X--ray emission from all the gas volumes in the simulated region, allowing us to employ Unit 1 of \\name{} to create a photon database associated to the simulation. In Unit 2 this photon cube is projected onto a particular observational plane and the photons are sampled according to the chosen aperture and observation time, as described in \\eq\\ref{fak}. With the photon box pre--calculated, this step is computationally trivial and can be performed very efficiently. As such, it could in fact be made available through an on line web service. The result of this step is a photon list and is processed in this paper by means of the {\\it Suzaku} simulator \\xissim{}. Similarly, we believe it should be possible for most such virtual telescopes to be adapted to the same format without undue difficulties.  This leads us to an interesting modular design, with a photon simulator producing a list of photons directly from a hydrodynamical simulation, and a virtual telescope receiving this list, and convolving it with the characteristics of a real X--ray instrument to produce a photon event file that can be directly compared in all its aspects to the analogous files resulting from real observations. This design presents a proper separation of concerns between theorists and observers, where each module is clearly independent from the others and can therefore be implemented and used separately. In a larger perspective, the well specified model for the data shared and exchanged between these units could represent a first step towards a standardised, interoperable way to publish simulation results in a way that is suitable for investigation by observers. In particular, this is very suitable for the kind of approach pursued by the International Virtual Observatory Alliance (IVOA\\footnote{See http://www.ivoa.net.}), and we indeed plan to deploy a web service that provides access to the results discussed in this paper in that context\\footnote{Please check http://www.mpa-garching.mpg.de/HydroSims for updates on this web service.}."
    },
    "1112/1112.1704_arXiv.txt": {
        "abstract": "We present numerical simulations of disc-planet interactions where the planet opens a gravitationally unstable gap in an otherwise gravitationally stable disc. In our disc models, where the outer gap edge can be unstable to global spiral modes, we find that as we increase the surface density scale the gap becomes more unstable and the planet migrates outwards more rapidly. We show that the positive torque is provided by material brought into the planet's coorbital region by the spiral arms. This material is expected to execute horseshoe turns upon approaching the planet and hence torque it. Our results suggest that standard type II migration, applicable to giant planets in non-self-gravitating viscous discs, is likely to be significantly modified in massive discs when gravitational instabilities associated with the gap occur. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.1018_arXiv.txt": {
        "abstract": "We report on new high-resolution imaging and spectroscopy on the multiple T Tauri star system V773 Tau over the 2003 -- 2009 period. With these data we derive relative astrometry, photometry between the A and B components, and radial velocity (RV) of the A-subsystem components.  Combining these new data with previously published astrometry and RVs, we update the relative A-B orbit model.  This updated orbit model, the known system distance, and A subsystem parameters yields a dynamical mass for the B component for the first time.  Remarkably the derived B dynamical mass is in the range of 1.7 -- 3.0 M$_\\sun$.  This is much higher than previous estimates, and suggests that like A, B is also a multiple stellar system.  Among these data, spatially-resolved spectroscopy provide new insight into the nature of the B component.  Similar to A, these near-IR spectra indicate that the dominant source in B is of mid-K spectral type.  If B is in fact a multiple star system as suggested by the dynamical mass estimate, the simplest assumption is that B is composed of similar $\\sim$ 1.2 M$_\\sun$ PMS stars in a close ($<$ 1 AU) binary system.  This inference is supported by line-shape changes in near-IR spectroscopy of B, tentatively interpreted as changing RV among components in V773~Tau~B.  Relative photometry indicate that B is highly variable in the near-IR. The most likely explanation for this variability is circum-B material resulting in variable line-of-sight extinction.  The distribution of this material must be significantly affected by both the putative B multiplicity, and the A-B orbit. ",
        "introduction": "\\objectname[V773 Tau]{V773~Tau} (HDE~283447, HBC~367) is among the most remarkable pre-main sequence (PMS) stellar systems presently known.  V773 Tau exhibits a broad variety of observable properties traceable to its PMS status, circumstellar material, and multiplicity. V773~Tau was first identified as a T~Tauri star by \\citet{Rydgren1976} based on H$\\alpha$ and Ca II H and K emission, high lithium abundance, photometric variability, and K2 spectral type.  The object presented an enigmatic mixture of classical (CTTS) and weak-lined T Tauri (WTTS) characteristics until it became clear there are multiple components; V773~Tau was resolved as a visual binary independently by \\citet{Ghez1993} and \\citet{Leinert1993} (with visual components designated here as A and B).  \\citet{Martin1994} first suggested, and \\citet{Welty1995} established A as a short-period (51-d) double-lined spectroscopic binary (SB2).  In filled-aperture high-angular resolution studies, \\citet[herein D2003]{Duchene2003} and \\citet{Woitas2003} independently identified an additional, ``infrared'' component in the system (herein designated C -- note D2003 use an alternate component notation), making V773~Tau at least a compact quadruple system with no fewer than four stars within roughly 100 AU.  The SB2 A subsystem was resolved by near-IR \\citep[herein B2007]{Boden2007} and radio \\citep[B2007]{Phillips1996} interferometry, allowing physical orbit reconstruction and component dynamical mass estimates (B2007).  Further, the variable radio emission and resolved morphology has been studied to infer magnetospheric interaction and complicated emission topology \\citep[and references]{Massi2002,Massi2006,Massi2008}.  Finally in a companion paper to this one, the VLBA astrometric study of \\citet[][herein T2011]{Torres2011} has refined the V773 Tau A orbit and revisited the system distance by both orbital and trigonometric parallax, indicating that the system is at 132.8 $\\pm$ 2.4 pc \\citep[an earlier VLBI trigonometric parallax by][had estimated 148.4 $\\pm$ 5.5 pc, while the B2007 orbital analysis yielded 136.2 $\\pm$ 3.7 pc]{Lestrade1999}.   Since the D2003 results and modeling, near-IR imaging and spectroscopic monitoring of V773 Tau (combined with B2007 results on A) have yielded important additional clues to the nature of the B component.  Here we will report on this continued monitoring, and what data taken over the past few years imply about the B component of this remarkable system. ",
        "conclusions": "\\label{sec:discuss}  The most notable result from this analysis is the large value for the V773~Tau~B dynamical mass, 2.35 $\\pm$ 0.67 M$_\\sun$.  This B-mass value is implied by the A-B orbit modeling described here, and independently supported by the companion VLBA study of T2011.  The B dynamical mass estimate makes B comparable to, but slightly less massive than the A subsystem.  Further, this B mass, along with the estimated luminosity (see below) and signs of double-lined features in our near-IR spectra (\\S~\\ref{sec:specAnalysis}) suggests that like A, B is also a multiple stellar system.  Before we discuss the implications of these results we need to consider the luminosity of V773~Tau~B.   In modeling the V773~Tau~B SED, D2003 found a range of possible solutions as a function of possible line-of-sight extinction.  In particular: \\begin{quote} ``With an additional 1 mag of extinction at $V$ toward V773 Tau [B], we find that its SED is well fitted by a K7 dwarf with a luminosity of 2.3 $\\pm$ 0.3 L$_\\sun$; a significant excess at $L$' remains present although the $H$- and $K$-band fluxes are then consistent with photospheric levels.'' \\end{quote}    Here we model the V773~Tau~B spectral energy distribution (SED) informed by the mid-K spectral type determination from \\S~\\ref{sec:specAnalysis}.  Resolved B photometry \\citep[][D2003]{White2001} are well-fit by a range of mid-K SED templates (T$_{\\rm eff} \\sim$ 4100 -- 4300 K) from \\citet{Pickles1998}, \\citet{Lejeune1997}, and PHEONIX \\citep[e.g.][]{Allard2000}, attenuated by considerable line-of-sight extinction (e.g.~$A_{\\rm V} \\sim$ 3 mag).  To facilitate comparison with B2007 results on V773~Tau~A, Figure~\\ref{fig:v773TauBSED} presents a model derived from the same template family (solar-abundance Kurucz-Lejeune template at T$_{\\rm eff}$ = 4250 K and $\\log$ g = 4.0).  This model clearly fits the available photometry well ($\\chi^2$/DOF = 0.68), and results in a B-luminosity estimate of 2.6 $\\pm$ 0.5 L$_\\sun$ with $A_V$ of 2.9 $\\pm$ 0.3 mag.  This extinction is roughly one magnitude higher in $A_V$ than reported toward A ($A_V$ = 1.8 $\\pm$ 0.2 mag) in B2007; taking the A extinction value as interstellar this suggests a circum-B source for the $\\sim$ one magnitude of additional extinction found here.  Like D2003, our SED modeling also finds a marginally significant $L$' excess, further reinforcing the hypothesis that there is circum-B material reprocessing radiation from B.  However, the interpretation of this B SED modeling is complicated by the fact that B is seen to be variable (Figure~\\ref{fig:relPhotometry}), and the input data \\citep[][D2003]{White2001} are not contemporaneous.    \\begin{figure}[t] \\includegraphics[angle=-90,width=14.5cm]{f7.eps} \\caption{Sample V773 Tau B Spectral Energy Distribution Model.  A sample SED model is shown using a 4250 K Kurucz-Lejunne SED template \\citep{Lejeune1997} with significant line-of-sight extinction ($A_V$ = 2.9 mag).  The implied B luminosity is 2.6 $\\pm$ 0.5 L$_\\sun$.  As first reported in D2003, the $L$' flux point (shown in [red]) indicates there is a marginally significant excess compared with the estimated photospheric level. \\label{fig:v773TauBSED}} \\end{figure}     The combination of our B-component dynamical mass (2.35 M$_\\sun$) and luminosity estimate (2.6 L$_\\sun$) suggests a multiple system interpretation.  For instance, \\citet{DAntona1997} models predict a luminosity of 17 L$_\\sun$ and T$_{\\rm eff}$ $\\sim$ 6200 K for a 3 Myr (B2007) single star at 2.35 M$_\\sun$.  (Note that a 1-$\\sigma$ excursion downward in the mass estimate -- 1.7 Msun -- would predict 2.8 L$_\\sun$ and 5200 K from these same models.)  Such a B-luminosity and temperature would seem to be impossible based on determinations presented here.  However assuming the simplest multiple configuration for B as a binary arrangement with a pair of 1.175 M$_\\sun$ stars, the \\citet{DAntona1997} models predict a pair of such stars to have a total luminosity of 2.6 L$_\\sun$, and T$_{\\rm eff}$ $\\sim$ 4800 K, in better agreement with the findings here.  While mass arguments alone would not exclude a single star interpretation, combined with double-lined indications in the near-IR spectroscopic analysis (\\S~\\ref{sec:specAnalysis}), the circumstantial case that V773~Tau~B is a multiple stellar system seems strong.  However, experience (e.g. T2002, D2003) has shown that one must consider the present A-B orbit model as preliminary pending additional phase coverage, and the B-mass estimate is still significantly uncertain ($\\sim$ 29\\% 1-$\\sigma$).    Photometric variability in V773 Tau has typically been attributed to the B component (D2003; Figure~\\ref{fig:relPhotometry}), and interpreted as changing line-of-sight extinction from circum-B material.  This interpretation is at least complicated both by potential B multiplicity, and by the physical size of the A-B orbit (15.3 AU, Table~\\ref{tab:ABphysics}).  Dynamical studies \\citep[e.g.][]{Artymowicz1994,Pichardo2005} indicate that stable regions for diffuse material are offset from binaries by several times the orbital semi-major axis.  Presumably the outer radius of any stable circum-B orbit should be on the order of 15 / 3 $\\sim$ 5 AU. Similarly, the typical size of a putative B subsystem should be several times smaller than this 5 AU scale -- probably $\\leq$ 1 AU -- making the working model for B a short-period binary.  In this picture, circumbinary material in the approximately 2--5\\,AU range would be responsible for both the near-infrared excess and photometric variability of B.  The large amount of dust needed to explain both extinction and variability (see below) raises the question whether it could remain dynamically stable in the region between B as a putative short-period binary and the larger A-B orbit, and whether such a reservoir of material would show other observable signatures such as detectable thermal IR and millimeter flux.  As a system V773~Tau shows a large mid-IR flux above photspheric levels \\citep[e.g.][D2003, Spitzer Taurus Legacy program, Padgett et al.~2011 in prep]{Prusti1992}, and \\citet{Andrews2005} lists significant flux at 850 $\\mu$m and 1.3 mm.  However the presence of the V773~Tau~C component \\citep[D2003,][]{Woitas2003} significantly complicates the interpretation of these flux measurements.  A spatially-resolved mm study with a sensitive facility such as ALMA would clarify the distribution of circumstellar material in the V773~Tau system.  If the putative B variability is due to changing line-of-sight extinction, this degree of variability is extraordinary.  To account for the amplitude of the $K$-variability shown in Figure~3, a variation in $A_V$ of up to 25\\,mag (assuming ISM-like dust opacity) would be indicated.  Such large $A_V$ excursions seem discrepant with the results of the SED modeling above.  A more plausible alternative might be an optically-thick structure -- perhaps a disk that may or may not be perturbed by complicated dynamics -- obscuring the line-of-sight to B components, allowing only a small amount of scattered light to reach the observer.  Variability on short timescale could be related to one of the components of B temporarily coming into direct view along its orbit, akin to the KH~15D system \\citep[][and references therein]{Herbst2010}.  The current sampling of the light curve for B is too sparse to decisvely test this hypothesis.  In either case, this large attenuation could explain why no clear detection of B is seen in extensive optical spectroscopic monitoring \\citep[B2007, Table~\\ref{tab:RVdata}]{Welty1995}.      If the inference of B multiplicity and circum-B material is correct, then the V773~Tau system seems remarkably similar to the \\objectname[HD 98800]{HD~98800} system, where circumbinary material is found around one but not both binaries in the system \\citep{Prato2001}.  In HD~98800 the inner B and outer A-B orbit significantly complicates the dynamics of B circumbinary material, leading to an inference of disk truncation and warping \\citep{Boden2005,Akeson2007,Verrier2007,Pichardo2008}.  The apparent increasing system complexity raises questions of stability in V773~Tau.  While more information on the B and C components will be necessary to address the specific stability of V773~Tau, it is clear that such complex systems are a common outcome of the star formation process, and can remain stable for Gyr \\citep[e.g.][]{Duquennoy1991,Tokovinin2008,Raghavan2010}."
    },
    "1112/1112.3647_arXiv.txt": {
        "abstract": "{We study phenomenological implications of the ATLAS and CMS hint of a $125\\pm 1$~GeV Higgs boson for the  singlet, and singlet plus doublet non-supersymmetric dark matter models, and for the phenomenology of the CMSSM. We show that in scalar dark matter models the vacuum stability bound on Higgs boson mass is lower than in the standard model and the 125~GeV Higgs boson is consistent with the models being valid up the GUT or Planck scale. We perform a detailed study of the full CMSSM parameter space keeping the Higgs boson mass fixed to $125\\pm 1$~GeV, and study in detail the freeze-out processes that imply the observed amount of dark matter. After imposing all phenomenological constraints except for the muon $(g-2)_\\mu,$ we show that the CMSSM parameter space is divided into well separated regions with distinctive but in general heavy sparticle mass spectra. Imposing the $(g-2)_\\mu$ constraint introduces severe tension between the high SUSY scale and the experimental measurements -- only the slepton co-annihilation region survives with potentially testable sparticle masses at the LHC. In the latter case the spin-independent DM-nucleon scattering cross section is predicted to be below detectable limit at the XENON100 but might be of measurable magnitude in the general case of light dark matter with large bino-higgsino mixing and unobservably large scalar masses.}   \\arxivnumber{1112.3647}    \\begin{document} ",
        "introduction": "In the standard model (SM) of particle interactions the only unknown quantity is the Higgs boson mass~\\cite{Higgs:1964ia,Guralnik:1964eu,Englert:1964et,Higgs:1964pj}. Any assumption that fixes the Higgs boson quartic self-coupling at any scale $\\Lambda$ implies a prediction for the Higgs boson mass. Many models of that sort have been proposed in the past based on different arguments of new physics beyond the SM. In general, the properties of the SM Higgs potential are among the best studied quantities in particle physics~(\\cite{Xing:2011aa}; for a review and references see \\cite{EliasMiro:2011aa}).    Based on  data collected in 2011,  both the ATLAS and CMS experiments have published their results for searches for the SM-like Higgs boson~\\cite{Chatrchyan:2012tx,ATLAS:2012si} confirming and improving their earlier claims~\\cite{ATLAS-CONF-2011-163,CMS-PAS-HIG-11-032} for the inconclusive evidence of a signal of a $M_H=124$~GeV (CMS) or $M_H=126$~GeV (ATLAS) Higgs boson; we will assume that the mass is in this $M_H=125 \\pm 1$~GeV range. The corresponding local significances of the excess in ATLAS and CMS are $3.5\\sigma$ and $3.1\\sigma$, respectively, while the  global significances after taking into account the look-elsewhere-effect  are $2.2\\sigma$ and $2.1\\sigma$. Although definitive confirmation of the observed evidence requires more data,  the LHC result motivates studies of fundamental scalars in particle physics and in cosmology.    If  the present  inconclusive evidence for $M_H \\approx 125$~GeV Higgs boson will be confirmed, this result will have a profound impact on building models beyond the SM and on their phenomenology. In the context of the SM, the Higgs boson mass 125~GeV is below the vacuum stability bound  $M_H>128$~GeV coming from the requirement of the SM validity up to the scale of gauge coupling unification $\\Lambda_{\\rm GUT}.$ Vanishing SM Higgs boson self-coupling $\\lambda (\\Lambda)=0$ below the GUT scale, $\\Lambda<\\Lambda_{\\rm GUT},$ implies that  the fundamental scale of new physics related to electroweak symmetry breaking and, perhaps, to flavour generation, might be lower than the GUT scale. On the other hand, the Higgs boson mass $M_H \\approx 125$~GeV may imply that there is new physics beyond the SM not too far  from the electroweak scale that modifies the Higgs boson mass prediction. The most popular such a framework is low energy supersymmetry (SUSY) that prefers a light Higgs boson. For SUSY scenarios the lightest Higgs boson mass $M_H \\approx 125$~GeV is unusually high, close to the upper bound in popular models, and implies a higher SUSY breaking scale than one expects from naturalness arguments. Clearly those arguments mean that the present hint for the Higgs boson mass requires re-assessment of  several  ``standard\" concepts both in SUSY and in non-SUSY models.  The aim of this work is twofold. First, assuming that the Higgs boson mass is in the range $M_H=125\\pm 1$~GeV, we study the implications of this assumption on the vacuum stability in scalar dark matter (DM)  models. In those models the DM and Higgs sectors are related via the Higgs portal and the scalar potentials are in general rather complicated. Due to many new self-interactions in the scalar sector, the SM Higgs quartic coupling renormalization is modified and  one might expect that the triviality $\\lambda (\\Lambda)=0$ may be achieved for higher values of $\\Lambda.$ We show that this is indeed the case and the SM vacuum stability results will be changed in the non-SUSY scalar DM models compared to the SM prediction. As a new result we show that in those scenarios the 125~GeV Higgs boson is consistent with the vacuum stability  up to $ \\Lambda_{\\rm GUT}$ and, therefore, the scalar DM models do not require new fundamental scales between TeV and  the GUT scales.   Second, a technically much more involved question is what is the implication of  the $M_H=125\\pm 1$~GeV LHC result for SUSY predictions of generating DM relic abundance, DM direct detection and for the LHC phenomenology. Generically such a heavy Higgs boson requires rather heavy stops, \\ie, a large SUSY breaking scale\\footnote{In the context of the 125~GeV Higgs boson this point has already been noted in \\cite{Carena:2011aa,Moroi:2011aa,Moroi:2011ab,Draper:2011aa,Arbey:2011ab,Heinemeyer:2011aa,Li:2011ab,Baer:2011ab,Hall:2011aa,Arbey:2011aa}.}. This, in general,  implies a large fine tuning to obtain the correct electroweak scale,  very fine tuned DM annihilation channels and poor prospects for discovering SUSY at the LHC. We analyze those issues in detail in  the constrained minimal supersymmetric standard model (CMSSM) and show that the requirements of $M_H=125\\pm 1$~GeV and correct DM relic abundance together select out parameter regions with  well defined sparticle spectra. We work out CMSSM predictions for DM direct detection cross sections in those parameter regions. The most important new result of this paper is to predict sharp linear relationship between the gluino, lightest stop and slepton masses  in the stop and slepton co-annihilation regions that are the only ones accessible to the LHC experiments.  If, in addition, also the muon anomalous magnetic moment $(g-2)_\\mu$ constraint is imposed on the CMSSM, only a tiny parameter region is singled out that induces DM via the slepton co-annihilation channel. In this parameter space the LHC has a good chance to observe gluinos and the lightest stop but the DM direct detection experiments like  XENON100 are predicted to obtain null result. In the other DM freeze-out channels that also predict the correct amount of DM  the situation might be an opposite -- only TeV scale DM is observable in DM direct detection experiments while the heavy gluinos and scalars decouple from the spectrum. We classify all  those possibilities and discuss their phenomenology.   In section~\\ref{sec:scalars} we present results for models of the SM extended with scalars: a complex $SU(2)$ singlet, an inert doublet or both. In section~\\ref{sec:cmssm} we give scans for CMSSM with both with and without the $(g-2)_\\mu$ constraint. We conclude in section~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  The recent LHC searches for the SM-like Higgs boson motivate studies of the fundamental scalars  in particle physics models and in cosmology. In this paper we analyzed the implications of  the $M_H\\approx 125$~GeV Higgs boson for the vacuum stability in scalar DM models and for the phenomenology of CMSSM. This value of the Higgs boson mass is interesting in both cases because it does not fit to the standard expectation neither in the SM nor in minimal supersymmetric  models with SUSY breaking scale below 1~TeV.  We have shown that in the case of non-SUSY scalar DM models  the vacuum can be stable up to the GUT scale even for Higgs boson masses much below the corresponding SM bound. Therefore, unlike the SM, the scalar DM models can be valid up to the GUT or Planck scales even for the Higgs boson mass as low as $M_H\\approx 125$~GeV.   In minimal SUSY models, to the contrary, the   $M_H\\approx 125$~GeV Higgs boson is heavier than expected in scenarios that address naturalness of the electroweak scale. In order to generate such a large Higgs boson mass at loop level, the SUSY breaking scale must be rather high and could be unobservable at the LHC. This problem can be overcome with extremely large stop $A$-term so that the lightest stop is light due to large mixing. At the same time the DM neutralino can also be light, either because of dominant slepton co-annihilation processes or because of large bino-higgsino mixing. In the latter case the DM-nucleon scattering cross section can be observable in direct detection experiments like the XENON100.  To quantify those results we studied the CMSSM by scanning over its parameter space  allowing the sparticle mass parameters to be very large. We first considered the case without attempting to explain the $(g-2)_\\mu$ in the context of CMSSM. We confirmed that for very large $A$-terms there exists a stop co-annihilation region where all DM,  stop and gluino are preferably light. Due to the mass degeneracy between stop and DM the stops can also be long lived resulting in non-trivial LHC phenomenology. The second parameter region that is potentially reachable at the LHC is the slepton co-annihilation region. The most important result of this work is to make sharp predictions of gluino, stop and slepton masses, shown in Fig.~\\ref{fig3},  for the CMSSM parameter regions that remain testable at the LHC.  For other channels of generating the correct DM  relic abundance the $M_H\\approx 125$~GeV Higgs boson implies very heavy sparticle masses. The exception is, of course, the DM that can be light due to bino-higgsino mixing even if other sparticles are as heavy as 10~TeV. In this case the CMSSM cannot be tested at the LHC but  the DM spin-independent scattering cross section off nuclei may be large due to the large higgsino component. The latter scenario may be discoverable already in the running XENON100 experiment.  If, in addition, one attempts to explain  also the $(g-2)_\\mu$ in this framework, there is immediate tension between the high SUSY scale and the large value of the needed $(g-2)_\\mu$ contribution. We found that after  imposing the $(g-2)_\\mu$ constraint on the CMSSM, only the slepton co-annihilation region survived at $3\\sigma$ level, see Fig.~\\ref{fig:cmssm:g-2}. This implies that the CMSSM has a definite prediction for the sparticle masses and spectrum to be tested at the LHC experiments.    \\vskip 0.5in \\vbox{ \\noindent{ {\\bf Acknowledgements} } \\\\ \\noindent We thank A. Strumia for several discussions. This work was supported by the ESF grants  8090, 8499, 8943, MTT59, MTT60, MJD140, JD164,  by the recurrent financing SF0690030s09 project and by  the European Union through the European Regional Development Fund. }"
    },
    "1112/1112.0472_arXiv.txt": {
        "abstract": "We present new {\\it Suzaku} and {\\it Fermi} data, and re-analyzed archival hard X-ray data from INTEGRAL and {\\it Swift}--BAT survey, to investigate the physical properties of the luminous, high-redshift, hard X-ray selected blazar IGR J22517+2217, through the modelization of its broad band spectral energy distribution (SED) in two different activity states. Through the analysis of the new {\\it Suzaku} data and the flux selected data from archival hard X-ray  observations, we build the source SED in two different states, one for the newly discovered flare occurred in 2005 and one for the following quiescent period. Both SEDs are strongly dominated by the high energy hump peaked at 10$^{20}$-10$^{22}$ Hz, that is at least two orders of magnitude higher than the low energy (synchrotron) one at 10$^{11}$-10$^{14}$ Hz, and varies by a factor of 10 between the two states. In both states the high energy hump is modeled as inverse Compton emission between relativistic electrons and seed photons produced externally to the jet, while the synchrotron self--Compton component is found to be negligible. In our model the observed variability can be accounted for by a variation of the total number of emitting electrons, and by a dissipation region radius changing from within to outside the broad line region as the luminosity increases. In its flaring activity, IGR J22517+2217 shows one of the most powerful jet among the population of extreme, hard X--ray selected, high redshift blazar observed so far. ",
        "introduction": "In the past years, {\\it Integral}--IBIS (Ubertini et al. 2003) and {\\it Swift}--BAT (Barthelmy et al. 2005) catalogs of hard X--ray selected extragalactic sources have been published (Sambruna et al. 2007; Bird et al. 2010), opening a new window in the study of blazars.  The hard-X-ray selection produces a set of diverse sources, including extremely ``red'' flat spectrum radio quasars (FSRQs), with Inverse Compton (IC) peak at keV-MeV energies, or ``blue'' BL Lacs with synchrotron peak at these frequencies, and provides a way to test the validity of the blazar sequence, looking for outliers not predicted by the sequence itself (Giommi et al. 2007). In particular, by selecting hard X--ray luminous FSRQ at high redshift, i.e. with the highest intrinsic bolometric luminosities, it is possible to collect samples of supermassive black holes (SMBHs) in the early Universe, thus introducing important observational constraint on the number density of heavy black holes at high redshift and hence on their formation processes and timing.  IGR J22517+2217 was first reported by Krivonos et al. (2007) as an unidentified object detected by  {\\it Integral}--IBIS (150 ks of total IBIS exposure). {\\it Swift} follow-up observations were used to associate the source with MG3 J225155+2217 (Bassani et al. 2007). MG3 J225155+2217 has been optically identified as a QSO by Falco et al. (1998) in a redshift survey of 177 FSRQs, on the basis of  S$_{IV}$, Ly$\\alpha$, C$_{II}$ and C$_{IV}$ emission lines. IGR J22517+2217 is the highest redshift ($z=3.668$) blazar detected in the fourth {\\it Integral}--IBIS hard X--rays catalog (Bird et al. 2010). The source has a {\\it Swift}--BAT counterpart (SWIFT J2252.0+2218) in the 3 years BAT survey (Ajello et al. 2009) and is present in the multifrequency ``Roma--BZCAT\" catalog of blazars (Massaro et al. 2009).  Using XRT, IBIS and archival optical/IR data, Bassani et al. (2007) constructed a non simultaneous SED of the source, showing an extremely bright X--ray emission with respect to the optical emission ($\\alpha_{OX}<0.75$), and suggested that IGR J22517+2217 could be a rare FSRQ with synchrotron peak at X--ray frequencies, or a more canonical FSRQ, i.e. with the synchrotron peak at radio--mm frequencies and IC peak at MeV-GeV energies, but with an exceptionally strong Compton dominance.  This ``controversial\" blazar has been studied also by Maraschi et al. (2008). They reanalyzed the existing {\\it Swift} (XRT and BAT) and {\\it Integral}--IBIS data, and propose a ''standard leptonic one-zone emission model'' (Ghisellini \\& Tavecchio 2009, see Sect. 3) with the peak of the synchrotron component at microwave/radio frequencies, and a high luminosity external Compton (EC) component peaking in hard X-rays to reproduce the SED  of the source. This model ruled out both a synchrotron and a synchrotron self--Compton (SSC) interpretation for the X--ray emission.  Ghisellini et al. (2010) included IGR J22517+2217 in their sample of 10 X--ray selected blazar at $z>2$: the intent of the paper was to characterize the physical properties of these powerful sources, and to confirm the capability of the hard X--ray selection in finding extreme blazars with massive SMBH, powerful jets and luminous accretion disks. IGR J22517+2217 is the highest redshift FSRQ in their sample and shows the highest total jet power (P$_{Jet}=1.5\\times10^{48}$ erg s$^{-1}$).  All these previous studies have been performed through the analysis of the ''average'' X-ray spectra obtained with the INTEGRAL and {\\it Swift}-BAT surveys, without taking into account any possible flux variation of the source during the period of monitoring (5 years for INTEGRAL and BAT). In this paper we present the discovery of strong flaring activity, in X--ray (IBIS and BAT) archival data, of this extremely bright and peculiar FSRQ, and the modelization of both its flaring and quiescent SEDs. New {\\it Suzaku} and {\\it Fermi} data are used for characterizing the quiescent state. Our goal is to investigate the evolution of the SED and obtain information on the physical condition of the source in the two different states.  The paper is organized as follows. In \\S2 we report the multiwavelength data analysis of the instruments involved in the SED building. In \\S3 we describe the model adopted to reproduce the broad band SED, while in \\S4 we discuss the SED fitting of both the flaring and quiescent state. The summary of our results is presented in \\S5. Throughout the paper, a $\\Lambda-CDM$ cosmology with $H_0 = 71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_\\Lambda = 0.73$, and $\\Omega_m = 0.27$ is adopted. ",
        "conclusions": "Thanks to a new {\\it Suzaku} observation in the X-ray energy band, the {\\it Fermi} upper limits in the 0.1-100 GeV band, the flux selected spectra obtained through a re-analysis  of the IBIS and BAT hard X-ray data, and other optical and radio archival data sets, we were able to identify a strong flare episode in the high redshift, hard X-ray selected blazar IGR J22517+2217, which occurred in Jan 2005, followed by a period of quiescence, extending up to the present days. To model the overall SEDs of the source in the flare and quiescent states, we adopted a leptonic, one-zone synchrotron and inverse Compton model. The optical/UV emission is interpreted as thermal emission from the accretion disk, plus IC from the corona, and reprocessed emission from the BLR.  The curvature observed in the X-ray spectra was proposed to be due to intrinsic, moderate absorption (N$_H\\sim2\\times10^{22}$ cm$^{-2}$). However, in the context of the broad band SED modelization proposed in this paper, it appears to be inherently accounted for by an intrinsic softening of the EC component around $\\sim10^{18}$ Hz.  In both states a very strong Compton dominance is observed, with the high energy hump (produced by EC), that is at least two orders of magnitudes higher than the low energy (synchrotron) one. The high energy peak flux varies by a factor 10 between the two states, while the high energy peak frequency remain almost constant, between $10^{20}-10^{22}$ Hz. The observed large Compton dominance constrains the value of the magnetic field, and hence the relevance of the self--Compton component, that is found to be negligible in both states. The model can explain the observed variability as a variation of the total number of emitting electrons (a variation of factor $\\sim7$) and as a change in the dissipation radius, moving from within to outside the broad line region as the luminosity increases.  In the flaring state, the jet power lower limit, represented by the radiative component $P_{\\rm r}$, requires a power comparable to the disk luminosity to produce the observed radiation. The total jet power upper limit, dominated by the bulk motion of protons, and estimated assuming one proton per electron, is more than $\\sim30$ times more powerful than the accretion luminosity ($2.6\\times10^{48}$ erg s$^{-1}$). Such extreme values have been derived only recently for a handful of extreme, high redshift, hard-X/soft-$\\gamma$ ray selected FSRQs, showing similar strong Compton dominance, and comparable with the value achieved by 3C 454.3 during the 2009 exceptional flare."
    },
    "1112/1112.0644_arXiv.txt": {
        "abstract": "$B$-modes are special patterns in cosmic microwave background (CMB) polarization. Degree-scale $B$-modes are smoking-gun signatures of primordial gravitational waves. The generic strategy of CMB polarization experiments is to employ a large number of polarimeters for improving the statistics. The Q/U Imaging ExperimenT-II (QUIET-II) has been proposed to detect $B$-modes by using the world's largest coherent polarimeter array (2,000 channels). A unique detection technique using QUIET's polarimeters, which comprises a modulation/demodulation scheme, enables us to directly extract polarization signals. The extracted signals are free from unpolarized components and intrinsic $1/f$ noise. We developed a data readout system for the QUIET-II experiment. We employed a ``master'' clock strategy, on-board demodulation functions, and end-to-end Ethernet connections for logical simplicity and high-density compact electronics for physical compactness. A clock module acts as a single master and guarantees phase matching between the modulation by the polarimeters and the on-board demodulation by ADC modules. Each ADC module has 64 ADC chips in the VME-6U single slot size. Both modules have hardware processors for Ethernet TCP/UDP. All these modules and control computers are connected via end-to-end Ethernet. Physical compactness and logical simplicity enable us to easily handle a large number of polarimeters, while maintaining quality of the $B$-mode experiments. The developed electronics (the clock modules and the ADC modules) fulfill these requirements. Tests with a setup similar to that of the real experiment proved that the system works appropriately. The performance of all system components is validated to be suitable for $B$-mode measurements. ",
        "introduction": "\\PARstart{D}{etection} of primordial gravitational waves could provide a new and unique window on the very early universe~\\cite{dodelson}. Although various approaches for detecting them \\cite{ligo,toba,cassini}, the most promising approach is the measurement of $B$-modes, which are odd-parity patterns in cosmic microwave background (CMB) polarization. Because $B$-modes are very faint ($<100$~nK), it is important to detect a large number of CMB photons for achieving sufficient sensitivity. The Q/U Imaging ExperimenT (QUIET) is a ground-based experiment that aims to detect the $B$-modes in the Atacama Desert in Chile, which is $5,080$~m above sea level. At the initial phase of the QUIET experiment (QUIET-I), we observed CMB polarization in the 95 (43)~GHz bands with 360 (72) channels of coherent polarimeter elements (polarimeter array). The obtained upper bounds for the $B$-modes are one of the most stringent limits to date; however they are still limited by statistical errors \\cite{QUIET}. An upgrade of the experiment (hereafter referred to as QUIET-II) has been proposed to achieve better sensitivity with a larger polarimeter array. Its primary goal is to detect $B$-modes at $r \\approx 0.01$ by using 2,000 channels composed of 500 elements of four-output polarimeters, where $r$ is the tensor-to-scalar ratio that indicates the intensity of the primordial gravitational waves \\cite{r}.  We developed a readout system for the QUIET-II experiments. For handling the QUIET-II polarimeter array, we employed a master clock strategy, on-board demodulation functions, and end-to-end Ethernet connections for logical simplicity and developed high-density electronics to achieve physical compactness. The master clock strategy guarantees phase matching between modulation and demodulation. Since polarization signals are modulated to suppress polarimeter noise and unpolarized components, phase matching and demodulation functions are essential requirements for the system. The on-board demodulation has the advantage of reducing the load of data transfer. Furthermore, physical compactness is also important to attach the system on a telescope mount to avoid picking up unexpected noise on the cables between the polarimeter array and the readout system. In this paper, we describe the system design and prototype tests. ",
        "conclusions": "We developed a readout system for QUIET-II. The conceptual requirements (handling a large number of channels) are fulfilled by the master clock strategy, on-board demodulation functions, and end-to-end Ethernet connections. Physical compactness and logical simplicity enable us to further increase the polarimeters for future experiments. We demonstrated the extraction of the polarization signals from the polarimeter by using a prototype system with a configuration similar to that of the experiment. In long term operation, data loss and time counter mismatch were not observed. We validated that the system works according to the design. Therefore, the system is ready to be used in the QUIET-II experiment.   \\appendix[Formulation for the polarimeter outputs]\\label{app} On the strip-line-coupled monolithic microwave integrated circuit devices, split right and left circular polarizations at the spectrum polarizer are amplified. Noises by HEMT amplifiers are added as follow: $g_1 E_R + n_1$ and $g_2 E_L + n_2$. Each phase switch modulate only the CMB signal components ($E_R$ and $E_L$). After the phase switches, two microwaves can be written as $c_1(t)g_1 E_R + n_1$ and $c_2(t)g_2 E_L + n_2$, respectively, where $c_1(t)$ and $c_2(t)$ are defined as the modulation terms in Sec.~\\ref{sec-demod}. Consider only one diode that is sensitive to the Stokes $Q$ parameter; then the recombined microwave by the coupler is $1/\\sqrt{2} (c_1(t)g_1 E_R + n_1 + c_2(t)g_2 E_L + n_2)$. The received power in the detector diode can be written as \\begin{equation}\\label{eq-app} D(t) =  |\\frac{1}{\\sqrt{2}} (c_1(t)g_1 E_R + n_1 + c_2(t)g_2 E_L + n_2)|^2. \\end{equation} Assuming that polarimeter noises $n_1(t)$ and $n_2(t)$ are stationary Gaussian noises, i.e., $\\langle n_1(t) \\rangle$ = 0, $\\langle n_2(t) \\rangle$ = 0, and $\\langle n_1(t)n_2(t) \\rangle$ = 0, where $\\langle X \\rangle$ is the average of $X$ during each phase state of the primary modulation, Eq.~(\\ref{eq-app}) is simplified as: \\setlength{\\arraycolsep}{0.0em} \\begin{eqnarray}\\label{eq-app1} D(t_i)&{}={}& g_1 g_2 E_L E_R c_1(t_i)c_2(t_i) \\IEEEnonumber \\\\ &&+ \\frac{g_1^2E^2_L}{2} c_1(t_i)^2 + \\frac{g_2^2E^2_R}{2} c_2(t_i)^2  \\IEEEnonumber \\\\ &&+{} \\frac{(n_1(t))^2}{2} |c_1(t_i)|^2 + \\frac{(n_2(t))^2}{2} |c_2(t_i)|^2. \\end{eqnarray} \\setlength{\\arraycolsep}{5pt} When $P = g_1 g_2Q/2$~~(here, $Q=2E_RE_L$), $L = g_1^2E^2_L/2$,  $R = g_2^2E^2_R/2$, $N_1(t_i)=(n_1(t_i))^2/2$, and $N_2(t_i) = (n_2(t_i))^2/2$, Eq.~(\\ref{eq-app1}) will be transformed to Eq.~(\\ref{eq-Dt}). By setting $c_2(t_i) = 1$, we can obtain Eq.~(\\ref{eq-Dt0})."
    },
    "1112/1112.4567.txt": {
        "abstract": "We present K band spectra for 133 nearby (d $<$ 33 parsecs) M dwarfs, including 18 M dwarfs with reliable metallicity estimates (as inferred from an FGK type companion), 11 M dwarf planet hosts, more than 2/3 of the M dwarfs in the Northern 8 pc sample, and several M dwarfs from the LSPM catalog. From these spectra, we measure equivalent widths of the Ca and Na lines, and a spectral index quantifying the absorption due to H$_2$O opacity (the H$_2$O-K2 index).  Using empirical spectral types standards and synthetic models, we calibrate the H$_2$O-K2 index as an indicator of an M dwarf's spectral type and effective temperature.  We also present a revised relationship that estimates the [Fe/H] and [M/H] metallicities of M dwarfs from their Na I, Ca I, and H$_2$O-K2 measurements. Comparisons to model atmosphere provide a qualitative validation of our approach, but also reveal an overall offset between the atomic line strengths predicted by models as compared to actual observations. Our metallicity estimates also reproduce expected correlations with Galactic space motions and H$\\alpha$ emission line strengths, and return statistically identical metallicities for M dwarfs within a common multiple system. Finally, we find systematic residuals between our H$_2$O-based spectral types and those derived from optical spectral features with previously known sensitivity to stellar metallicity, such as TiO, and identify the CaH1 index as a promising optical index for diagnosing the metallicities of near-solar M dwarfs. ",
        "introduction": "Contrary to their dimness, M dwarf stars hold significant promise for illuminating the processes that govern the formation and evolution of stars, planets, and the Milky Way.  M dwarfs possess masses of 0.6 M$_{\\odot} >$ M$_* >$ 0.08 M$_{\\odot}$ \\citep{Delfosse2000}, straddling the peak of the stellar initial mass function and dominating stellar populations by number \\citep{Bastian2010}.  With main sequence lifetimes that exceed a Hubble time, Galactic M dwarfs also represent a complete archeological record of the chemical evolution and star formation history of the Milky Way \\citep[e.g.][]{2007AJ....134.2418B}.  Moreover, M dwarfs are of great interest as potential exoplanet host stars, as all of a planet's observable signals will be significantly easier to detect if it orbits an M dwarf compared to a similar planet in orbit around a G dwarf \\citep[e.g., ][]{2008PASP..120..317N}.  Calibrating the fundamental parameters of M dwarfs, however, is a difficult challenge from both an observational and theoretical perspective.  Empirical measurements of M dwarf masses, luminosities, temperatures, and radii are anchored primarily by the information extracted from the orbits of M dwarf binaries, either the very rare eclipsing systems \\citep[$\\sim$15 systems known; e.g., ][]{Kraus2011,2011ApJ...742..123I} or less rare (but also somewhat less informative) spectroscopic/astrometric systems \\citep[e.g., ][]{Shkolnik2010,Martinache2007}.  Overcoming the intrinsic faintness of these low-luminosity systems, however, makes their characterization a taxing observational challenge, and analyses of known binaries reveal systematic offsets in inferred temperatures and radii that correlate with both orbital period and tracers of magnetic activity \\citep{Ribas2006,Lopez-Morales2007}, suggesting that more wide binaries are needed to accurately infer the parameters typical isolated M dwarfs \\citep{Kraus2011}.  Theoretical constraints on M dwarf parameters have been similarly difficult to achieve: accurately modeling the deep convective zones in M dwarf interiors \\citep{Mullan2001, 2008ApJ...676.1262B} and the mixture of molecules and grains that dominate M dwarf atmospheres \\citep{Tsuji1996,Allard2000} requires significant computational resources, as well as extensive databases of oscillator strengths and opacities measured from laboratory experiments.  Confronting theoretical predictions with empirical measurements have also identified significant offsets even for isolated field stars: empirical effective temperature scales appear $\\sim$200-300 K cooler than would be inferred from theoretical model atmospheres \\citep[e.g.][]{2005MNRAS.358..105J}, for example.  While M dwarf masses, temperatures, and radii have proven challenging to calibrate, those parameters appear strikingly pedestrian when contrasted with the difficulty of inferring M dwarf metallicities.  \\citet{1976A&A....48..443M} was the first to systematically assess the strengths of M dwarf spectral features as a function of metallicity, constructing a grid of synthetic model atmospheres that included molecular opacities and spanned temperatures of 4250--3000K and metallicities of \\feh = -2.0 -- 0.0 dex.  With these models, Mould established several foundations of subsequent investigations of M dwarf metallicities, such as the metallicity sensitivity of TiO absorption, the gravity sensitivity of CaH absorption, and the rough metallicities of Galactic M subdwarfs \\citep[ \\feh $\\sim$ -1.0 dex; ][]{Mould1976b}.  These studies enabled the subsequent photometric and spectroscopic identification of M subdwarfs \\citep{Stauffer1986,Ruiz1993}, the development of increasingly detailed models of M dwarf atmospheres \\citep{Allard1990}, and represent the fundamental origin of the TiO and CaH-based indices used today to identify metal-poor M subdwarfs \\citep{1997AJ....113..806G, 2005MNRAS.356..963W, 2007ApJ...669.1235L}.  While techniques to identify metal-poor M dwarfs have been developing for more than 35 years, methods for identifying metal-$rich$ M dwarfs have received significantly less attention.  Interest in identifying M dwarfs with super-solar metallicities was only recently spurred by the realization that such stars could be highly promising targets for searches for Earth-like planets.  The existence of a correlation between (gas giant) planet frequency and host-star metallicity is well established for FGK-dwarfs \\citep{1997MNRAS.285..403G, 2004A&A...415.1153S, 2005ApJS..159..141V}, but the first investigations to examine if the planet-metallicity correlation extended to M dwarf hosts returned mixed results \\citep{2005A&A...442..635B,2006ApJ...652.1604B}.  Bonfils et al. (2005, B05 hereafter) analyzed V and K photometry for 20 wide M dwarf companions to FGK-dwarfs whose metallicities could be determined reliably via standard spectroscopic techniques. Assuming that both binary components inherit the same metallicity from their parent molecular cloud material, B05 assigned the metallicity measurements for each primary to the M dwarf secondary, and derived iso-metallicity contours in the M$_V$ vs. V-K$_s$ color-magnitude plane.  This photometric metallicity calibration suggested that nearby M dwarfs, including the planet hosts Gl 876 and Gl 436, were slightly metal-poor compared to the mean metallicity of the Galactic disk.  This conclusion was reinforced by the work of \\citet{2006ApJ...652.1604B}, whose spectroscopic analysis indicated sub-solar metallicities for Gl 436, Gl 581, and Gl 876.  The persistence of the planet-metallicity correlation into the M dwarf regime has been suggested, however, by the recent re-calibration of the photometric metallicity calibration by \\citet{2009ApJ...699..933J}.  Using six M dwarfs with wide, metal-rich FGK companions to update the M$_V$ vs. V-K metallicity contours, Johnson \\& Apps (2009, JA09 hereafter) found that the B05 calibration systematically underestimated the compositions of these metal-rich stars.  JA09 used their revised calibration to estimate the metallicities of six M dwarfs with planetary mass companions, concluding that M dwarf planet hosts are indeed preferentially metal rich, just like FGK hosts.  This conclusion was supported by the subsequent work of \\citet{2010A&A...519A.105S}, who updated the B05 and JA09 calibrations by using theoretical models to inform the functional form of the M$_K$ vs. V-K relation, and by inferring the mean metallicity of M dwarfs in the Solar Neighborhood from a volume-limited and kinematically-matched sample of F \\& G dwarfs in the Geneva-Copenhagen Survey \\citep{Nordstrom2004}.  The photometric metallicity calibration developed by B05 and extended by JA09 and SL10 has proven to be a valuable resource, but its dependence on precise V magnitudes and trigonometric parallaxes limits its utility to M dwarfs in the immediate solar neighborhood, at least for the remainder of the pre-Gaia era.  In a recent contribution, we presented an alternative, spectroscopic technique for estimating the metallicities of near-solar metallicity M dwarfs \\citep[][hereafter RA10]{2010ApJ...720L.113R}. This technique is capable of providing metallicity estimates with an accuracy comparable to that of the photometric technique, and requires only a moderate resolution K band spectrum of the M dwarf target, providing a significantly lower observational burden for empirically estimating the metallicities of distant, near-solar or super-solar M dwarfs. An example of the utility of this technique is provided by a recent paper on M dwarf planet hosts identified in the most recent Kepler data release by Muirhead et al. (2011; submitted)  While RA10 provided a concise introduction to the spectroscopic K band metallicity indicator, space limitations prevented a full exploration of the technique.  In this paper, we provide a full analysis of the motivation, calibration, demonstration and application of the technique using spectra we obtained for 133 Solar Neighborhood M dwarfs.  In Section 2, we describe the composition of our sample, and the acquisition and reduction of our spectroscopic observations.  In Section 3, we describe our measurements of the new H2O-K2 (a modified version of the H2O-K index utilized by RA10), \\nai and \\cai spectral features upon which the K band metallicity technique depends. In Section 4, we analyze model atmospheres to demonstrate the H2O-K2 index's insensitivity to metallicity, and use the H2O-K2 measurements from our sample to calibrate the index as a spectral type and T$_{eff}$ indicator.  In Section 5, we use the H2O-K2 index as the basis for a revised K band metallicity calibration; by incorporating an additional M dwarf binary into our sample of metallicity calibrators, and adopting a modified functional form to better normalize the temperature dependence of the Na and Ca features, we derive an updated K band \\feh  metallicity indicator, as well as a new calibration for overall metallicity ([M/H]), along with robust estimates of the uncertainty in each calibration.  In Section 6 we perform several sanity checks of the K band metallicity technique, demonstrating that the K band metallicity estimates preserve expected correlations between metallicity and Galactic kinematics as well as signatures of chromospheric activity.  We conclude in Section 7 by using the K band metallicity estimates, in combination with other stellar parameters such as rotational velocity ($v$ sin $i$), magnetic activity (as diagnosed by H$_{\\alpha}$ emission) and optical brightness, to identify a sample of nearby M dwarfs with particular promise for exoplanet surveys.  We summarize our findings in Section 8, and present our full K band spectral atlas in an appendix (we also make our spectra available to the community online). ",
        "conclusions": "We present K band spectra for a sample of 133 M dwarfs, including 18 M dwarfs with reliable metallicity estimates (as inferred from an FGK type companion), 11 planet hosts, more than 2/3 of the Northern 8 pc sample, and additional M dwarfs from the Lepine Shara Catalog. From these spectra, we measured EWs of the \\cai and \\nai lines, as well as an index quantifying the absorption due to H$_{2}$O opacity, features we have previously shown to be of use for predicting M dwarf temperatures and metallicities.  \\begin{enumerate}  \\item{From the subset of our stars which possess primary or secondary KHM spectral types, we calibrated the H$_2$O-K2 index as a spectral type indicator. We also estimated effective temperatures for the stars in the sample by interpolating their H$_2$O-K2 indices onto a solar metallicity [H$_2$O-K2, T$_{eff}$]  model grid.}  \\item{We revise the functional form adopted to predict the [Fe/H] metallicities of M dwarfs based on measurements of their \\nai, \\cai, and H$_{2}$O features.  We perform a linear regression upon 18 M dwarfs with reliable metallicity estimates from FGK companions to calibrate this updated relationship.  Statistical tests demonstrate that the [Fe/H] estimates produced by this relationship are accurate to $RMSE$=0.141 dex, and confirm that this relation accounts for more of the variance within our calibrant sample than any other existing technique for estimating M dwarf metallicities.}  \\item{For the first time, we derive an expression for an M dwarf's overall metallicity [M/H]; quantitative comparisons to model atmospheres, which are benchmarked according to overall metallicity, provide a qualitative validation of our approach for estimating metallicities from \\nai, \\cai, and H$_{2}$O features, but also demonstrate an overall offset between the atomic line strengths predicted by models as compared to actual observations.}  \\item{We examined previous optical molecular indices sensitive to stellar metallicity using out metallicity estimates. We identify that the CaH1 feature as a potentially valuable optical metallicity discriminator for solar-metallicity early type M dwarfs.}  \\item{We perform several sanity checks of our metallicity estimates, confirming that our metallicity estimates reproduce expected correlations between metallicity and Galactic space motions and H$\\alpha$ emission line strengths, and return statistically identical metallicities for members of M-M multiple systems.}  \\item{We recovered the results from previous metallicity studies that nearby M dwarf planet host exhibit solar to supersolar metallicities, where stars with Jovian-mass planets are more metal-rich than stars with Neptune-likes or super-Earths. A list of the best targets in the sample for planetary searches selected by metallicity, activity and rotation-rates is given.}    \\end{enumerate}"
    },
    "1112/1112.2707_arXiv.txt": {
        "abstract": "NGC~326 is one of the most prominent X- or Z-shaped radio galaxies (XRGs/ZRGs) and has been the subject of several studies attempting to explain its morphology through either fluid motions or reorientation of the jet axis.  We examine a 100~ks \\textit{Chandra X-ray Observatory} exposure and find several features associated with the radio galaxy: a high-temperature front that may indicate a shock, high-temperature knots around the rim of the radio emission, and a cavity associated with the eastern wing of the radio galaxy.  A reasonable interpretation of these features in light of the radio data allows us to reconstruct the history of the AGN outbursts.  The active outburst was likely once a powerful radio source which has since decayed, and circumstantial evidence favors reorientation as the means to produce the wings.  Because of the obvious interaction between the radio galaxy and the ICM and the wide separation between the active lobes and wings, we conclude that XRGs are excellent sources in which to study AGN feedback in galaxy groups by measuring the heating rates associated with both active and passive heating mechanisms. ",
        "introduction": "A small fraction of double-lobed extragalactic radio sources have long, centro-symmetric tails of radio emission extending from the nucleus or primary lobes in a very different direction, thereby producing ``X'' or ``Z''-shaped radio galaxies (XRGs/ZRGs).  The tails, or ``wings'', are faint, do not currently harbor a relativistic jet, and can be substantially longer (in projection) than the primary lobes. They have elicited considerable interest due to the suggestion by \\citet{merritt02} that they result from a recent supermassive black hole (SMBH) merger that reoriented the jet, but their origin is uncertain.  It has also been suggested that the wings result from the deflection of lobe plasma by pressure fronts in the intergalactic or intracluster medium (IGM/ICM) \\citep{leahy84,worrall95,capetti02,gopal03,zier05,hodges-kluck10a}.  Neither of these hypotheses have yet been proven, even in a single source.  B2~0055+26 is a dramatic XRG/ZRG, with wings longer (in projection) than the primary lobes (Figure~\\ref{radio_maps}).  The wings are highly collimated over $\\sim$100~kpc.  Its host galaxy is the northern component of the dumbbell NGC~326 \\citep{colla75} and the entire system resides in the galaxy cluster Zw~0056.9+2636 \\citep{zwicky68}. Hereafter, we adopt the convention in the literature and refer to the whole system as NGC~326.  The cluster atmosphere has been characterized by \\citet[][hereafter, W95]{worrall95} on large scales using ROSAT, and has an average temperature of $kT \\sim 2$~keV.  The cluster atmosphere is substantially asymmetric (Figure~\\ref{rosat}), suggesting that it is not in hydrostatic equilibrium and may be made up of two or more merging subclusters or that it is a composite of two or three poorer clusters seen in projection (W95).  A detailed multifrequency study of the radio emission measured the spectral index $\\alpha$ as a function of position in the lobes and wings, finding spectral steepening in the wings and away from the ends of the jets \\citep[][hereafter M01]{murgia01}.  Owing to its dramatic structure (Figure~\\ref{radio_maps}), NGC~326 is frequently used as the prototypical XRG/ZRG in formation models purporting to explain the X-shaped morphology.  \\citet{ekers78}, who discovered the wings but not their true length, attributed them to ongoing regular precession of the relativistic jet inflating the primary lobes.  Later, \\citet{merritt02} argued that the striking X-shaped structure necessitated rapid secular precession.  Noting that black holes are difficult to reorient via external forces other than accretion disk instabilities, \\citet{merritt02} suggested instead that the X-shaped structure is the result of the final coalescence of a SMBH binary.  Supposing that the components of the binary have misaligned angular momenta, the final spin axis (and jet axis) will be nearly instantaneously reoriented.  \\citet{wirth82} also argued in favor of a single rapid jet reorientation caused by a close encounter between the constituents of the dumbbell galaxy.  \\begin{figure*}[t] \\begin{center} \\includegraphics[width=0.42\\textwidth,angle=-90]{1.46ghz.eps} \\includegraphics[width=0.42\\textwidth,angle=-90]{1.55ghz.eps} \\caption{\\textit{Left:} 1.46 GHz VLA C Array image.  The contour levels are 1, 2, 4, 8, 16, 32, 64, and 128~mJy/beam.  The restoring beamsize is $16.^{\\prime\\prime}7\\times 14.^{\\prime\\prime}4$. \\textit{Right:} 1.55 GHz VLA A Array.  The contour levels are 0.1, 0.2, 0.4, 0.8, 1.6, and 3.2~mJy/beam.  The restoring beamsize is $1.^{\\prime\\prime}35\\times 1.^{\\prime\\prime}15$.  Note that the images are on different size scales, and the full extent of the wings is not seen at right. } \\label{radio_maps} \\end{center} \\end{figure*}  Based on the large-scale, hot, asymmetric cluster atmosphere, W95 found that the buoyancy of the lobe material could explain the wings provided that they were driven by rapid backflows expected near the terminal shocks of the radio jet heads.  This scenario is fundamentally different from those above in that it presumes a stable jet axis.  In support of the backflow hypothesis, we \\citep{hodges-kluck10a} reported a local asymmetric enhancement of the hot gas in a preliminary analysis of an archival 100~ks \\textit{Chandra X-ray Observatory} (CXO) image on the scale of $\\sim$50--75~kpc where the radio beam is aligned with the major axis and the wings with the minor axis.  This is also consistent with the expectation of \\citet{capetti02} that X-shaped sources are aware of the geometry of their environments.  However, it is not clear how the wings would grow so long in these models.  In this paper, we carefully consider the high-resolution CXO data in order to shed light on the history of NGC~326, evalauting possible hydrodynamic hypotheses and characterizing the interaction between the radio galaxy and its environment, which could not be resolved in the W95 study. NGC~326 is located at a redshift $z = 0.0474$ where $1^{\\prime\\prime} = 0.92$~kpc in the \\textit{Wilkinson Microwave Anisotropy Probe} cosmology \\citep[$H_0 = 71$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\Lambda} = 0.73$, and $\\Omega_m = 0.27$;][]{spergel07}.  We use a Galactic photoelectric absorption in all spectral fitting of $N_H = 5.69\\times 10^{20}$~cm$^{-2}$ \\citep{kalberla05}.  The total extent of the cluster atmosphere in the ROSAT image is about 20$^{\\prime}$ on a side (1.1~Mpc).  The remainder of the paper is organized as follows: in Section~2 we describe the CXO observation, data, and basic features, then in Section~3 we use the X-ray and radio data to derive quantities of interest related to the radio galaxy.  In Section~4 we use these data to reconstruct the history of the radio galaxy and argue that NGC~326, as well as other XRGs, are interesting sources with regard to feedback from active galactic nuclei (AGNs).  We summarize our results in Section~5. ",
        "conclusions": "We report on a 100~ks CXO image of NGC~326, an XRG with wings 100~kpc long that resides in a poor cluster.  Owing to its dramatic wings, NGC~326 has frequently been the subject of studies attempting to explain wing origin \\citep[W95;][]{wirth82,merritt02,gopal03,dennett02,zier05,capetti02,lal05}. The CXO image provides an important step forward in evaluating these hypotheses due to the identification of several X-ray features associated with the radio galaxy.  These include a high-temperature front to the southeast of the radio galaxy (Figure~\\ref{tempmap}), which we assume to be a shock front, high-temperature knots of emission around the radio galaxy akin to cavity rims (Figure~\\ref{hardness_compare}), and a cavity associated with the eastern wing (Figure~\\ref{cavity}).  The high angular resolution of CXO compared to ROSAT also allows us to remove point sources and characterize the broader ICM, revealing an X-ray peak associated with the northern lobe of the radio galaxy rather than the ISM and an elliptical enhancement around the radio galaxy \\citep[first identified in][]{hodges-kluck10a}.  By characterizing the immediate environment of the radio galaxy and taking into account spectral aging of the radio emission in the wings, we can construct a reasonable history for the AGN outbursts. The jets, lack of hot spots, and low radio power ($10^{41}$~erg~s$^{-1}$) indicate a weak FR~I-like source today, but the putative shock front, lobe size and morphology, and GHz emission at the wing tips indicate a powerful FR~II source no more than $\\sim$30~Myr ago.  We therefore suggest the kinetic jet power could be described as $P_{\\text{jet}} = 10^{45} e^{-t/t_0}$~erg~s$^{-1}$, where $t_0 = 10$~Myr for the current outburst.  This function would account for the observed features as well as produce about the right amount of cavity enthalpy.  The energy required to inflate cavities the size of the lobes places the current kinetic jet power between $5\\times 10^{42} - 5\\times 10^{43}$~erg~s$^{-1}$.  The radio power suggests that the flow is therefore still radiatively inefficient.  NGC~326 is the best studied XRG, and the variety of wing formation models associated with it is a testament to the difficulty of proving any one scenario. The sound speed in the ICM is only 500~km~s$^{-1}$ ($0.5$~kpc~Myr$^{-1}$), implying that the long (100~kpc) wings must have expanded supersonically.  This disfavors the backflow models, which have yet to convincingly demonstrate that internal lobe flows can generate sustained supersonic expansion; simulations indicate that they become turbulent and dissipate their thrust before reaching the wing edges.  In fact, any model in which the wing growth is related to the sound speed is disfavored in systems with long wings (Figure~\\ref{wingages}).  The most likely alternative is that the wings are fossil relics illuminated by the present outburst, but \\citet{gopal03} note that a simple spin-flip would not explain the Z-shaped morphology of the wings.  However, their model relies on bending a moderately weak jet, which would not be expected to propel hypersonic lobes.  As fossil radio lobes would be subject to hydrodynamic forces, it is possible that the Z-shaped morphology of the wings is produced by buoyant evolution in an anisotropic atmosphere.  It would take less than 50~Myr for the wings to rise buoyantly away from an axis connecting them through the nucleus to their present-day location, which is consistent with the age of the active lobes and the age required for the wings to respond to the anisotropy of the large-scale cluster emission to form the understated C-shaped arc (Figure~\\ref{hardness_compare}).  However, this scenario fails to explain the plume on the eastern wing.  It is also notable that the tendency for XRGs to have jets co-aligned with the major axes of their host galaxies \\citep{capetti02,saripalli09} and hot atmospheres \\citep{hodges-kluck10a} extends to XRGs with very long wings.  This is unexplained by any reorientation model, and remains compelling evidence that the environments of most XRGs help determine their morphology.  XRGs provide an important target for characterizing AGN feedback through the study of ``ghost'' cavities, both as pathfinders (they illuminate the cavities in the GHz bands) as well as targets which are readily distinguished from the active outburst.  The heating rate associated with dead outbursts can then be measured in isolation, and compared to the timescales required to thermalize and isotropize the energy in the ICM.  In cases such as NGC~326, where we believe most of the energy associated with the active outburst has already been deposited, we can use the AGN lifetime to estimate an average heating rate, assuming some duty cycle.  XRGs are therefore interesting beyond the origin of their bizarre morphology."
    },
    "1112/1112.2525_arXiv.txt": {
        "abstract": "We examine here the acceleration of particles and high energy collisions in the the Janis-Newman-Winicour (JNW) spacetime, which is an extension of the Schwarzschild geometry when a massless scalar field is included. We show that while the center of mass energy of collisions of particles near the event horizon of a blackhole is not significantly larger than the rest mass of the interacting particles, in an analogous situation, it could be arbitrarily large in the JNW spacetime near the naked singularity. The high energy collisions are seen to be generic in the presence of a photon sphere in the JNW spacetime, whereas an extreme fine-tuning of the angular momentum of the colliding particles is required when the photon sphere is absent. The center of mass energy of collision near the singularity grows slowly for small and extremely large deviations from the Schwarzschild blackhole, but for intermediate strengths of the scalar field it rises moderately fast. As a possible and potentially interesting application, we point out that the presence of such high energy collisions may help the blackhole configurations to be distinguished from a naked singularity. ",
        "introduction": "Various terrestrial particles accelerators collide particles at the center of mass energies upto TeV scale which is significantly below the Planck scale that corresponds to the quantum gravity regime. Clearly, the physics in a very large energy range remains unexplored today. In this connection, an intriguing possibility would be to make use of various naturally occurring processes in the vicinity of various astrophysical compact massive objects where gravity would be very strong.  The blackholes are expected to provide such an environment. With this motivation it was investigated whether it is possible to have ultrahigh energy collisions in the vicinity of the event horizon of the blackholes. Various compact very massive objects that occur in the universe, for instance at the center of most of the galaxies, are assumed to be Kerr blackholes which are characterized by their spin $a$ and mass $M$. In the geometrical units (where $c=G=1$), The event horizon of such a blackhole is located at $r=b(a)=M\\left(1+\\sqrt{1-a^2}\\right)$ in the Boyer-Lindquist coordinates. The center of mass energy of collision between two particles interacting near the horizon released from infinity at rest was investigated. When $a=0$ the blackhole is a Schwarzschild blackhole and the maximum center of mass energy that can be achieved in the collision at $r=b(a=0)=2M$ was shown to be finite \\cite{Bauchev}. The effect of adding rotation to the blackhole was investigated in \\cite{BSW}. It was shown that nothing interesting happens for small values of spin. However, interestingly for large spins close to the extremality $a=1$, it can be shown that the center of mass energy of collision at $r=b(a=1)=M$ is arbitrarily large. Various aspects of this process in Kerr and other blackhole geometries were investigated in \\cite{BSW2},\\cite{Berti}. We extended this process to super-spinning objects with spin larger than unity containing a naked singularity \\cite{Patil},\\cite{Patil2}.  Here we take a different approach to crank up the center of mass energy of collision. Instead of spinning up a blackhole we \"charge it up\" with a static massless scalar field. The two parameter family solution of Einstein equations in this case is given by \\cite{JNW}. One parameter is mass $M$ and the second parameter is $q$, the scalar charge which is a measure of strength of the scalar field as we discuss in the next section. We investigate the collision at $r=b(q)=2\\sqrt{M^2+q^2}$. When $q=0$, the JNW metric reduces to Schwarzschild metric and the center of mass energy of collision at $r=b(q)=2M$ is small. We investigate here the effect of invoking a scalar field on the center of mass energy of collision at $r=b(q)=2\\sqrt{M^2+q^2}$. We show that there is a divergence of center of mass energy of collision. However, the divergence is very slow for small and large values of $q$ and it is moderately fast for intermediate values of $q$ as $r=b$ is approached.  The JNW spacetime \\cite{JNW}, is an extension of the Schwarzschild geometry when a massless scalar field is included instead of being empty. Naively speaking, the blackhole event horizon then deforms into a naked singularity in the JNW spacetime, when any smallest non-zero value $q$ of the scalar field is included. This metric was studied from the perspective of gravitational lensing near a naked singularity \\cite{Vlens}. If naked singularities which are hypothetical astrophysical objects occur in nature, an important problem would be how to distinguish them from the blackholes. Both these entities which occur generically in gravitational collapse of massive matter clouds within the framework of the Einstein gravity \\cite{collapse}, present ultra-strong-gravity regimes which would be clearly of much interest to study and explore the physical effects of very strong gravity fields.  The JNW solution which always contains a naked singularity has been studied extensively from the perspective of gravitational lensing and distinguishing it from blackholes \\cite{Vlens}. It was shown that for certain range of parameters, in the presence of a photon sphere, the gravitational lensing will be qualitatively similar to that by the Schwarzschild blackhole and it will not be possible to observationally distinguish the blackhole from a naked singularity configuration. But for the remaining range of the parameters in the absence of the photon sphere, the observational signatures of the two in terms of lensing will be quite different, allowing us to differentiate between the two.  We show here that it would be possible to have ultrahigh energy collisions generically in the presence of the photon sphere in the case of a JNW naked singularity, whereas in case of the Schwarzschild blackhole such collisions do not take place. Conversion of dark matter and ordinary matter particles like protons into particles with smaller mass, such as electrons, would be efficient at large center of mass energy of collisions. The flux of outgoing particles produced in collisions thus is expected to be larger in the JNW spacetime as compared to the Schwarzschild blackhole. We argue that this could lead to the observational signature of the JNW naked singularities in the presence of the photon sphere, as distinct from the blackholes. ",
        "conclusions": "In this paper we studied the effect of inclusion of a scalar field onto the center of mass energy of collision between the particles. The main purpose here was to show that it is possible to have ultrahigh energy collisions in the JNW spacetime in the presence of a static massless scalar field around naked singularities, while the high energy collisions are absent in the Schwarzschild blackhole case when the scalar field is zero.  We also speculated that the high energy collisions could allow us to set up a scheme to observationally distinguish the Schwarzschild blackhole from the JNW naked singularities which are surrounded by the photon sphere, in the case when the gravitational lensing is ineffective for such a purpose. It is beyond the scope of this paper to rigorously demonstrate it here and we plan to report that work elsewhere.  Also, in the future work we would like to study the process of infall of dark matter particles and also ordinary particles like protons onto the galactic central dark object, modeled as a blackhole and as a JNW spacetime with a photon sphere surrounding the naked singularity. A good understanding of the density of dark matter and ordinary matter particles in the surroundings of these objects will have to be obtained and studied. We intend to study the collisions and investigate various particle physics processes leading to the annihilation of dark and ordinary matter particles into other particles like neutrinos and such other by-products. Finally, one would like to calculate the flux of the particles produced in the collisions at the large distances for JNW naked singularities, and compare it with the corresponding flux in the blackhole case."
    },
    "1112/1112.2655_arXiv.txt": {
        "abstract": "{The H.E.S.S. collaboration (Abramowski et al. 2011) dicovered a new TeV point-like source HESS J1943+213 in the Galactic plane and suggested three possible low-energy-band counterparts: a $\\gamma$-ray binary, a pulsar wind nebula (PWN), or a BL Lacertae object.} {We measure the distance to the radio counterpart G57.76-1.29 of HESS J1943+213.} {We analyze Very Large Array observations to obtain a reliable HI absorption spectrum.} {The resulting distance limit is $\\ge$ 16 kpc. This distance % strongly supports that HESS J1943+213 is an extragalactic source, consistent with the preferred counterpart of the HESS collaboration}. ",
        "introduction": "Very-high-energy (VHE) $\\gamma$-ray observations shed a light on the question of origin and acceleration of cosmic rays (Caprioli 2011). A multi-wavelength approach helps to solve the question, especially by seeking VHE sources' counterparts in lower energy bands. More than 110 VHE $\\gamma$-ray sources have been detected by ground-based high energy telescopes recently (http://www.mppmu.mpg.de/~rwagner/sources/).  About 30 of them have no counterparts identifed so far. The identification of counterparts plays a key role in understanding the nature of the VHE sources and in distinguishing different emission mechanisms for TeV $\\gamma$-rays (Halpern \\& Gotthelf 2010, Misanovic et al. 2011, Tian et al. 2008).  HESS J1943+213 was recently discovered by the H.E.S.S. collaboration (Abramowski et al. 2011) and was suggested to have three possible low-energy-band counterparts: a $\\gamma$-ray binary, a pulsar wind nebula (PWN), or a BL Lacertae object. HESS J1943+213 has a compact radio counterpart G57.76-1.29 which has been observed in the VGPS (also a NVSS source, NVSS J194356+211826, see Condon et al. 1998).  We analyze 1.4 GHz continuum and HI-line observations of G57.76-1.29 from Very Large Array (VLA) Galactic Plane Survey (VGPS, Stil et al. 2006). Based on our well-tested methods to build HI absorption spectrum against a continuum background radio source, we extract an HI absorption spectrum towards G57.76-1.29/HESS J1943+213. We report our results in this research note. ",
        "conclusions": "\\label{res}   \\begin{figure} \\vspace{50mm} \\begin{picture}(50,50) \\put(-40,230){\\special{psfile=f1.eps hoffset=0 voffset=0 hscale=40 vscale=40 angle=-90}} \\end{picture}  \\caption{1.4-GHz VLA image of HESS J1943+213 (contours 7, 15, 20 K) from the VGPS data.} \\end{figure}   \\begin{figure} \\centering \\includegraphics*[width=8cm]{f2.eps} \\caption{Top panel: HI emission spectrum towards J1943+2118. Bottom panel: HI absorption spectrum. The standard deviation of exp(-Tau) is 0.14, taken from the velocity range of 60 to 100 km/s (the mean value is 1.04)} \\end{figure}  \\begin{figure} \\vspace{50mm} \\begin{picture}(50,50) \\put(-40,230){\\special{psfile=f3.eps hoffset=0 voffset=0 hscale=40 vscale=40 angle=-90}} \\end{picture} \\caption{The HI channel map at -67 km s$^{-1}$ with 1420 MHz continuum contours (the same as in Figure 1).} \\end{figure}  Fig. 1 shows the 1420 MHz continuum map in a 0.1$^\\circ$ field containing G57.76-1.29. The VGPS data has an angular resolution of $1\\arcmin$ and an rms noise inthe HI Line of 2 K (Stil et al. 2006). It is a good database to study Galactic supernova remnants (SNRs) which usually extend from a few arcmin to $\\sim$1 degree in size. We have previously used the data to study some interesting SNRs (see a review paper by Leahy \\& Tian 2010) and have also developed methods to estimate kinemtic distances to them by analyzing 21cm continuum, HI-line and CO-line data (Tian, Leahy \\& Wang 2007; Leahy \\& Tian 2008).  Fig. 2 displays the HI emission spectra of source and background regions (top panel) and the absorption spectrum (bottom panel) towards G57.76-1.29. The background region was taken as directly surrounding the source region. Pixels were assigned to source or background regions depending on whether they have 1420 MHz continuum brightness above or below a specified threshold, respectively. We used the program meanlev in the DRAO (Dominion Radio Astrophysical Observatory) Export Software Package to calculate the source and background spectra. Then we used our methods, which take into account the background continuum radiation, to calculate optical depth in the form $exp(-\\tau)$. The resulting HI absorption spectrum shows absorption at all Galactic velocities. This includes absorption at positive velocities, from the interstellar medium interior to the solar circle, up to the tangient point velocity. Also we see absorption for negative velocities above $\\sim$ -67 km/s, which corresponds approximately to the outer edge of the Galaxy. The lowest velocity absorption feature (Fig. 2) appears at $\\sim$ -67 km s$^{-1}$ and is associated with the emission peak from gas in the outer galaxy. This absorption feature's reality is further supported by the continuum-subtracted HI channel map at the peak of the absorption at -67 km s$^{-1}$ (Fig. 3). This clearly shows the spatial association between the lower HI brightness temperature, caused by absorption, and the peak of the 1420 MHz continuum source. We estimate the uncertainties in our HI absorption spectrum using the standard derivation of exp(-$\\tau$), taken from the velocity range 60 to 100 km/s, where the spectrum is dominated by noise.  The resulting  standard derivation is 0.14. This yields that the absorption feature at $\\sim$ -67 km s$^{-1}$ is $\\sim$6 times the noise level and is real.  The velocity limit gives a lower kinematic distance limit of $\\sim$ 16 kpc to G57.76-1.29, taking a standard Galactic circular rotation model (V$_o$=220 km s$^{-1}$, R$_o$=8.5 kpc). Our absorption spectrum is consistent and gives a picture of absorption out to the most negative velocities for which there is emission, so that G57.76-1.29 is an extragalactic source.  In summary, HESS J1943+213 has three possible low-energy-band counterparts (Abramowski et al. 2011): a $\\gamma$-ray binary, a pulsar wind nebula (PWN), or a BL Lacertae object. Our observations and analysis show that HESS J1943+213, and its confirmed radio counterpart G57.76-1.29, is an extragalactic source."
    },
    "1112/1112.3653_arXiv.txt": {
        "abstract": "We use \\spitzer-IRAC data to identify near-infrared counterparts to submillimeter galaxies detected with \\herschel-SPIRE at 250\\um\\ in the {\\it Herschel} Astrophysical Terahertz Large Area Survey (H-ATLAS).  Using a likelihood ratio analysis  we identify 146 reliable IRAC counterparts to 123 SPIRE sources out of the 159 in the survey area.  We find that, compared to the field population, the SPIRE counterparts occupy a distinct region of 3.6 and 4.5\\um\\ color-magnitude space, and we use this property to identify  further 23 counterparts to 13 SPIRE sources. The IRAC identification rate of 86\\% is significantly higher than those that have been demonstrated with wide-field ground-based optical and near-IR imaging of \\herschel\\ fields.  We estimate a false identification rate of 3.6\\%, corresponding to 4 to 5 sources.    Among the 73 counterparts that are undetected in SDSS, 57 have both 3.6 and 4.5\\um\\ coverage.    Of these,   43 have $[3.6] - [4.5]> 0$ indicating that they are likely to be at $z\\gtrsim$ 1.4. Thus, $\\sim$ 40\\% of identified SPIRE galaxies are likely to be high redshift ($z\\gtrsim$ 1.4) sources. We discuss the statistical properties of the IRAC-identified SPIRE galaxy sample including far-IR luminosities, dust temperatures, star-formation rates, and stellar masses. The majority of our detected galaxies have 10$^{10}$ to 10$^{11}$ $L_{\\sun}$ total IR luminosities and are not intense starbursting galaxies as those found at $z \\sim 2$, but they have a factor of 2 to 3 above average specific star-formation rates compared to near-IR selected galaxy samples. ",
        "introduction": "The extragalactic background at far-infrared (IR) and submillimeter (sub-mm) wavelengths is well-constrained from total intensity measurements (Puget et al.\\,1996; Fixsen et al.\\,1998; Dwek et al.\\,1998). However, the properties of the discrete galaxies that make up this background are still largely unknown. These sub-mm galaxies are expected to capture the dusty star-formation out to redshifts of 4 and beyond and are now understood to be an integral component of galaxy formation and evolution (Hughes et al.\\,1998; Eales et al.\\,1999; Blain et al.\\,2002; Chapman et al.\\,2003; Austermann et al.\\,2010).  A challenge for studies of sub-mm galaxies is that wide-field surveys at these wavelengths have poor spatial resolution, caused by the limited apertures of single dish sub-mm survey telescopes (Smail, Ivison \\& Blain 1997; Scott et al.\\,2002; Coppin et al.\\,2006). Typical resolutions are of the order of tens of arcseconds, making the identification of multiwavelength counterparts challenging, particularly in the optical where sub-mm galaxies are usually faint due to high dust extinction. Furthermore, the surface density of faint optical galaxies is such that several potential counterparts may be situated within the positional error of each sub-mm source. However, the identification of counterparts to sub-mm galaxies is critical for both photometric and spectroscopic studies, which yield redshifts, spectral energy distributions (SEDs) and morphological information. The measurements of these quantities for statistically significant samples of sub-mm galaxies are required to compare their properties with theoretical predictions, and to fully understand the role of the sub-mm bright phase in galaxy evolution.  The most reliable way to pinpoint the positions of sub-mm galaxies is through high-resolution sub-mm interferometry, which directly traces the dust emission at wavelengths comparable to the original selection function (e.g. Downes et al. 1999; Gear et al. 2000; Iono et al. 2006; Younger et al.  2007).  The sensitivity of early generations of sub-mm interferometric instrumentation limited anlayses to samples of no more than a few sources (e.g. Downes et al. 1999; Gear et al. 2000; Iono et al. 2006; Younger et al.  2007, 2008, 2009; Ivison et al. 2008; Cowie et al. 2009; Aravena et al. 2010a; Hatsukade et al. 2010; Tamura et al. 2010; Chen et al. 2011 ), and although rapid progress is being made (e.g. Smolcic 2012a,b) such observations are still unfeasible for large surveys of sources.  In addition, extended sources (e.g. Matsuda et al. 2007), or sources that are a blend of multiple components (e.g. Wang et al. 2011; Smolcic 2012b), can be difficult to detect.  An alternative technique to identify sub-mm galaxies is to take advantage of the far-IR-radio correlation (e.g. Condon et al.\\ 1998; Garrett 2002) and the relatively low surface density of radio sources by searching for counterparts in deep interferometric radio data (e.g. Ivison et al.\\ 1998, 2002; Borys et al.\\ 2004; Dye et al.\\ 2009; Dunlop et al.\\ 2010; Biggs et al.\\ 2011). Similarly, the dusty, active sub-mm galaxies are typically bright at mid-IR wavelengths, a property that has also been utilized to identify counterparts (e.g. Ivison et al.\\ 2004; Pope et al.\\ 2006; Dye et al.\\ 2008, 2009; Clements et al.\\ 2008; Dunlop et al.\\ 2010; Biggs et al.\\ 2011). Deep radio and mid-IR (typically {\\it Spitzer} MIPS 24\\um but also {\\it Spitzer} IRAC 3.6\\um; Biggs et al.\\ 2011) searches typically identify counterparts to 60--80\\% of sub-mm galaxies, and nearest-neighbor positional matching is then used to investigate the properties of these sources at other wavelengths. The remaining 20--40\\% of sources are thought to be at high-redshifts ($z \\ga$ 3) or be dominated by cold dust ($T_D\\la 25$~K at $z\\sim 2$; Chapman et al.\\ 2005).   The {\\it Herschel} Astrophysical Terahertz Large Area Survey (H-ATLAS; Eales et al.\\,2010) is the largest open-time key project being undertaken by the {\\it Herschel} Space Observatory (Pilbratt et al.\\,2010)\\footnote{{\\it Herschel} is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA.}. The total planned survey area is 550 deg$^2$ within which it is expected to detect $>$ 300,000 bright sub-mm galaxies. During the Science Demonstration Phase (SDP), H-ATLAS observed 14.4 deg$^2$ in the Galaxy And Mass Assembly (GAMA; Driver et al.\\,2011) 9~hour field, to 5$\\sigma$ depths of 35 to  132 mJy in five bands from 100 to 500\\um, using PACS (100 and 160\\micron; Poglitsch et al.\\ 2010) and SPIRE (250, 350 and 500\\micron; Griffin et al.\\,2010) in parallel mode.  The lack of deep radio and 24\\um\\ data across the field restricts counterpart identification to optical (e.g.\\ Sloan Digital Sky Survey [SDSS]; York et al.\\,2000) or near-IR data (e.g.\\ VISTA Kilo-Degree Infrared Galaxy Survey [VIKING]; Sutherland et al.\\,in prep, UKIRT Infrared Deep Sky Survey [UKIDSS]; Lawrence et al.\\,2007).  Smith et al.\\ (2011) identified SDSS and GAMA galaxies as the counterparts to 37\\% of the SPIRE sources in the H-ATLAS SDP field.  In the wider area of the whole GAMA-9~hour field Fleuren et al.\\,(2012) increased the fraction of SPIRE sources with reliable identifications to 51\\% by using $K_{\\rm s}$-band imaging from the VIKING survey.  In the GAMA-15~hour field 50\\% of the SPIRE sources have reliable counterparts identified in Wide-Field Infrared Survey Explorer (WISE) data at 3.4\\,\\micron\\ (Bond et al.\\ 2012).  In this paper we cross-identify H-ATLAS SPIRE sources with \\spitzer\\ Infrared Array Camera (IRAC; Fazio et al.\\ 2004a) galaxies, which are selected from 3.6 and 4.5-\\micron\\ observations of 0.4 deg$^2$ of the H-ATLAS SDP field.  The IRAC data are advantageous for counterpart identification because of their mid-IR wavelengths and depth ($\\sim$3 magnitudes deeper than WISE). Thus, counterpart identification is less biased towards galaxies with relatively low redshifts or dust obscurations than the existing analyses.  The paper is organized as follows: the data analysis and sample selection are presented in Section~\\ref{sec:data}. In Section~\\ref{sec:ids} we describe the counterpart identification method and in Section~\\ref{sec:discussion} we discuss our results, including the identification rate and the colors  and properties of SPIRE counterparts. Our conclusions are presented in Section~\\ref{sec:conclusion}. A table of the identified counterparts, including magnitudes and fluxes is presented in the Appendix.  We use J2000 coordinates and $\\Lambda$CDM cosmology with $\\Omega_m$= 0.27, $\\Omega_\\Lambda$= 0.73 and $H_0$ = 70\\,${\\rm km\\,s^{-1}Mpc^{-1}}$ throughout. All photometry is on the AB magnitude system where $1\\,\\mu{\\rm Jy}$= 23.9 mag; IRAC 3.6 and 4.5\\um\\ AB magnitudes are designated $[3.6]$ and $[4.5]$, respectively. ",
        "conclusions": "\\begin{tabular}{lccccccc} \\hline \\hline Method & $N_{\\rm SPIRE}$$^a$ & $N_{\\rm ID}$$^b$ & ID rate$^c$ & $N_{\\rm IRAC}$$^d$& $N_{\\rm multiple}$$^e$  & Contamination rate  \\\\ \\hline LR     &   159             &   123           &   77\\%      &   123       &   0   &   1.9\\%                             \\\\ Color-magnitude&   19      &   13            &   68\\%      &   23        &   7   &   12.6\\%                            \\\\ \\hline Total  &   159             &   136           &    86\\%     &   146       &   7   &   3.6\\%                             \\\\ \\hline\\hline \\end{tabular} \\end{center} \\tablecomments{ $^a$ Number of SPIRE sources considered. $^b$ Number of SPIRE sources with at least one counterpart. $^c$ Fraction of SPIRE sources with at least one identified counterpart. $^d$ Number of IRAC counterparts identified. $^e$ Number of SPIRE sources with multiple IRAC counterparts.} \\end{table*}   \\subsection{LR counterparts} \\label{sec:lrresult}  We apply the LR technique outlined in section~\\ref{sec:lr} and ~\\ref{sec:q0} at 3.6 and 4.5\\um, to identify IRAC counterparts to 123 of 159 \\atlas\\ SPIRE sources (Tables~\\ref{tab:summary} and \\ref{tab:photom}). $q(m)$ and $n(m)$ are calculated in a radius of 3$\\sigma_{\\rm pos}=$ 7.2\\arcsec\\ (although $Q_0$ is radius independent; section~\\ref{sec:q0}) and only IRAC galaxies within 7.2\\arcsec\\ of SPIRE sources are considered to be potential counterparts. This limit includes 99.7\\% of true counterparts, whilst minimizing the potential contamination from unassociated sources.  The analysis is performed separately at 3.6 and 4.5\\um\\ and any IRAC galaxy with $R\\geq$ 0.8 at either wavelength is considered to be a SPIRE counterpart.  In addition, where both 3.6 and 4.5\\um\\ data is available we combine the probabilities ($R$) from each wavelength, and include five counterparts that have a combined probability of being considered to be reliable identifications $\\geq$ 0.8, but have $R<$ 0.8 at 3.6 and 4.5\\,\\micron\\ individually.  All 123 counterparts are presented in table~\\ref{tab:photom}, including 81 galaxies with both 3.6 and 4.5\\um\\ photometry. The false identification rate is 1.9\\% or approximately two sources (equation~\\ref{eq:n}); in the case of counterparts that are identified at both 3.6 and 4.5\\um\\ we use the highest of the two $R_j$ in this calculation.  The identification rate of 77\\% is significantly higher than that from optical (37\\% in SDSS; Smith et al.\\,2011) or near-IR analyses (51\\% in VIKING; Fleuren et al.\\,2012) of \\atlas\\ sources.  This difference is likely to be primarily driven by the wavelengths of the study, which are less sensitive to $K$-correction and dust absorption than optical data.  The counterpart identification rate is also significantly higher than the 50\\% obtained with WISE 3.4-\\micron\\ data (Bond et al. 2012); this is true even if we only consider the 3.6\\,\\micron\\ data, where 97 of the 129 SPIRE sources with 3.6\\,\\micron\\ coverage (75\\%) are identified.  WISE is shallower than our IRAC observations -- 19.7~mag at 3.4\\,\\micron\\ compared to 22.5~mag at 3.6\\,\\micron\\ in IRAC. However, 79\\% of the 3.6-\\micron\\ identified counterparts have $[3.6]\\le$ 19.7 mag, suggesting that IRAC data to this depth would identify counterparts to $\\sim$ 60\\% of SPIRE sources. The remaining difference in the IRAC and WISE identification rates is likely to be due to the resolutions of the instruments -- the WISE 3.4\\,\\micron\\ beam is $\\sim3$ times larger than IRAC -- although Cosmic Variance may also contribute. We note that if instead the data were limited to $[3.6] <$ 20.5, corresponding to the 5$\\sigma$ point-source detection limit reached in a 120 sec integration with IRAC at 3.6\\um, the expected identification rate drops from 77\\% to $\\sim$ 70\\%.  The results presented here are insensitive to the exact value of $Q_0$. When using a 1$\\sigma$ lower value for $Q_0$ (i.e. $Q_0=$ 0.84 at both 3.6 and 4.5\\,\\micron), and comparing to the results presented in section~\\ref{sec:lrresult} and table~\\ref{tab:photom}, the counterparts to 116 sources (95\\%) are unchanged. In this case there are four (3\\%) previously identified sources that no longer have reliable counterparts, and three (2\\%) previously unidentified SPIRE sources that now have $R\\geq$ 0.8 counterparts.       \\subsection{Counterparts identified in IRAC color-magnitude space} \\label{sec:color}   While we reliably identify 77\\% of the SPIRE sources with the LR method, the values of $Q_0$ indicate that $\\sim$ 90\\% of SPIRE counterparts are detected in IRAC.  Figure~\\ref{fig:color} shows the IRAC color-magnitude diagram for galaxies in the \\atlas\\ SDP field, with SPIRE counterparts identified with the LR method highlighted. These galaxies occupy a distinct region of IRAC color-magnitude space -- they typically have $[3.6]-[4.5]\\ga -$ 0.4 and are brighter than the background population.  We use this property to identify counterparts to the SPIRE sources that have both 3.6 and 4.5-\\um\\ IRAC coverage but no $R\\geq$  0.8 counterparts from the LR method. A similar analysis was performed by Biggs et al.\\,(2011) at 3.6 and 5.8-\\micron\\ to identify to LABOCA 870-\\um\\ sources.  We begin by determining the region of color-magnitude space in which there is a minimal chance of contamination by background sources (the ID region). A 20\\% contamination limit is used because this equivalent to $R\\geq$ 0.8 for the LR method. We demarcate the ID region with the simplest reasonable function -- a single diagonal line. The placement of the line is determined by calculating the gradients and intercepts that would yield 20\\% contamination from background sources within 7.2\\arcsec\\ radius of the LR counterparts. The number of LR counterparts returned is maximized to yield a best-fit gradient of $0.515$ for a line that intercepts $[3.6]=$ 16.0 mag at $[3.6]-[4.5]= -$2.142. This division is shown as a solid line in figure~\\ref{fig:color}.  IRAC galaxies that lie above this line and are within 7.2\\arcsec\\ of SPIRE sources that are otherwise unidentified, have $>$ 80\\% probability of being physically associated with the SPIRE source, i.e., $<$ 20\\% chance of being an unrelated foreground or background galaxies, and are considered counterparts.  There are 36 SPIRE sources that do not have LR counterparts. Of these 19 have both 3.6 and 4.5\\,\\micron\\ data, and 23 IRAC counterparts to 13 of these sources are identified with the color-magnitude method (seven have multiple counterparts). These sources are presented in table~\\ref{tab:photom}. The contamination rate for this method is 12.6\\% (approximately two IRAC galaxies), calculated by summing the probabilities of finding a background galaxy within 7.2\\arcsec\\ of a SPIRE source {occupying a distinct region of   color-magnitude space. }    In figure~\\ref{fig:color} we show the $S_{250}/S_{500}$ color-flux plot for \\atlas\\ SDP sources and consider whether sources with identified IRAC counterparts are representative of the whole SPIRE population.  A 2D Kolmogorov-Smirnov (KS) test between the background \\atlas\\ SPIRE population and the sources with identified counterparts has $p=$ 0.172, suggesting that the two samples are drawn from similar, but not necessarily identical, parent populations. Considering the $S_{250}/S_{500}$ color and 250\\,\\micron\\ fluxes separately yields $p=0.710$ and $p=0.289$, respectively. Thus we conclude that the sources with identified IRAC counterparts are very similar to the whole SPIRE population, but may have a slightly different distribution of 250\\,\\micron\\ flux.  \\subsection{Astrometric offset} \\label{sec:offset} To check whether there is any noticeable astrometric shift among the input catalogs for LR analysis,  in figure~\\ref{fig:offset} we plot the positional differences between the SPIRE sources and the corresponding IRAC counterparts. The  median  separation between SPIRE  and  IRAC  sources is  0.10\\arcsec. We also investigate the overall astrometric reference frame difference between IRAC and SDSS catalogs, using $\\sim$ 40\\% of the IRAC sample that is also identified with SDSS galaxies within in 2$''$. The positional difference  between IRAC and SDSS is  0.19 $\\pm$ 0.17\\arcsec and we found no systematic offset either in RA or Dec.  Our IRAC data are complemented with VIKING $Z,  Y, J, H$ and $K_{\\rm s}$-band photometry, which is measured in 2\\arcsec\\ diameter apertures (Sutherland et al.\\,in prep). We cross-match the IRAC (FWHM $\\sim$ 1.6\\arcsec) to the nearest neighbor  VIKING (FWHM $\\sim$ 0.9\\arcsec) source within 2\\arcsec\\ search radius.  Indeed, the mean offset between the IRAC and VIKING positions is 0.09 $\\pm$ 0.50\\arcsec\\ and 0.28 $\\pm$ 0.28\\arcsec\\ in RA and Dec, respectively. We find no statistically significant astrometric error between  two catalogs.  \\begin{figure} \\includegraphics[width=8cm]{offset.ps} \\caption{ The distribution of $\\Delta$RA  and  $\\Delta$Dec offsets between SPIRE sources and the corresponding IRAC counterparts identified in this work. The  astrometric median  separation between  SPIRE sources  and  IRAC  is  0.10\\arcsec\\  which is consistent  with the overall IRAC astrometric uncertainty tied to the SDSS astrometry. The three concentric circles  have radii of  1, 2 and 3$\\sigma_{\\rm pos}$ (7.2\\arcsec) from the SPIRE centroid. The overall offsets in RA and Dec show that there is no systematic offset between input  SPIRE and IRAC sources. } \\label{fig:offset} \\end{figure}       \\begin{figure*} \\begin{minipage}{17.5cm} \\begin{center} \\subfigure{ \\includegraphics[width=4.5cm]{2244sdss.ps} }                \\quad \\subfigure{ \\includegraphics[width=4.5cm]{ID-2244-384-ch1.ps} } \\quad \\subfigure{ \\includegraphics[width=4.5cm]{ID-2244-384-ch2.ps} } \\end{center} \\end{minipage} \\caption{ H-ATLAS ID-2244: an example of source that is identified using the LR method, but where the IRAC counterpart ($R=$ 0.99 at both 3.6 and 4.5\\,\\micron) is different from the SDSS counterpart (SDP.2244, $R=$ 0.97; Smith et al.\\ 2011).  From left-to-right we show a three-color SDSS image ($g$, $r$, and $i$ bands), the 3.6\\,\\micron, and the 4.5\\,\\micron\\ data.  The galaxies identified as the SDSS (SDP.2244) and IRAC (H-ATLAS ID-2244) counterparts to the SPIRE source are marked with an `X' and a small circle, respectively.  The large circle has 7.2\\arcsec\\ radius and encompasses the SPIRE 3$\\sigma_{\\rm pos}$ area in which counterparts are identified. Contours show the SPIRE 250\\,\\micron\\ emission at 5, 7, 9, 11$\\sigma$ levels.  In the absence of high-resolution sub-mm imaging we cannot determine whether this SPIRE source is the results of blended emission from the two identified galaxies, or whether one of those counterparts is a chance association.} \\label{fig:postage1} \\end{figure*}    \\begin{figure*} \\begin{minipage}{17.5cm} \\begin{center} \\subfigure{ \\includegraphics[width=4.5cm]{891sdss.ps} }                \\quad \\subfigure{ \\includegraphics[width=4.5cm]{ID-891-384-ch1.ps} } \\quad \\subfigure{ \\includegraphics[width=4.5cm]{ID-891-384-ch2.ps} } \\end{center} \\end{minipage} \\caption{H-ATLAS ID-891: an example of counterpart identification using the color-magnitude method.  Symbols are as figure~\\ref{fig:postage1}. There is no LR counterpart in the SDSS $r$-band (Smith et al.\\ 2011) or with IRAC at 3.6 or 4.5\\um. However, there are two galaxies that are both identified as SPIRE counterparts on the basis of their IRAC colors and magnitudes. These two galaxies have $R=$ 0.41 and  0.47  at 3.6\\,\\micron, and $R=$ 0.44 and  0.50 at 4.5\\,\\micron\\ and neither is detected in SDSS.  It is probable that the SPIRE source is a blend of the emission from these two IRAC galaxies although, in the absence of spectroscopic redshifts it is unclear whether the configuration is caused by a line-of-sight alignment or an interaction between the two IRAC galaxies.  } \\label{fig:postage2} \\end{figure*}  \\begin{figure*} \\begin{minipage}{17.5cm} \\begin{center} \\subfigure{ \\includegraphics[width=4.5cm]{827sdss.ps} }                \\quad \\subfigure{ \\includegraphics[width=4.5cm]{ID-827-384-ch1.ps} } \\quad \\subfigure{ \\includegraphics[width=4.5cm]{ID-827-384-ch2.ps} } \\end{center} \\end{minipage} \\caption{H-ATLAS ID-827 has four counterparts identified with the IRAC color-magnitude method and is the most complex system in our sample. The westernmost and easternmost galaxies are detected by SDSS (but are not identified as counterparts in these data; Smith et al.\\ 2011), and have photometric redshifts of $z=$ 0.07 $\\pm$ 0.01 and $z=$ 0.20 $\\pm$ 0.03, respectively. The two fainter counterparts are not detected in SDSS but have $[3.6]-[4.5]>$ 0, indicating that they may be at $z\\gtrsim$ 1.4 (Papovich 2008). We conclude that the SPIRE source is most likely to be comprised of blended emission from the four galaxies. Symbols are as figure~\\ref{fig:postage1}.} \\label{fig:postage3} \\end{figure*}       As described in sections~\\ref{sec:lrresult} and \\ref{sec:color} we identified IRAC counterparts to 136 (86\\%) of the SPIRE sources. 123 of these are identified with the LR method and 23 counterparts to 13 SPIRE sources are identified with the color-magnitude method (Table~\\ref{tab:summary}).  The contamination rate of 1.9\\% for the LR method is smaller than that from the color-magnitude method (12.6\\%) but corresponds to a similar number of IRAC galaxies -- two to three. Combining the results from the two identification methods, we expect that the total false identification rate of the catalog presented in Table~\\ref{tab:photom} is 3.6\\%, such that five to six of the IRAC counterparts presented in table~\\ref{tab:photom} are false.    \\begin{figure} \\includegraphics[width=9cm]{stat_72.ps} \\caption{The distribution of IRAC $[3.6]-[4.5]$ color for SPIRE counterparts with photometry at both wavelengths.  For the 47 sources with SDSS spectroscopic or photometric redshifts we also plot the $[3.6]-[4.5]$ color against redshift. The 57 SPIRE-IRAC identifications without reliable redshifts are either faint or undetected in SDSS and 43 of them have $[3.6] - [4.5] >$ 0, an indicator of high redshift source ($z\\gtrsim 1.4$; Papovich 2008).  We expect that more than half of the sub-mm bright sources reside at $z \\gtrsim$ 1.4.  For reference the color-redshift tracks of Arp~220 (Silva et al.\\,1998), Mrk231 (Berta\\, 2006) and the average sub-mm galaxy  SED (Wardlow et al.\\, 2011) are shown.  } \\label{fig:stat} \\end{figure}        \\subsection{Statistics of Identified Sources}   We next compare the counterparts identified in the IRAC data with results based on SDSS (Smith et al.\\ 2011).  Of the 146 IRAC counterparts, 102 (70\\%) are undetected in SDSS, 52 (36\\%) agree with the SDSS identification, and two (1\\%) are different galaxies to the SDSS counterparts. On the basis of the false-identification rates of the two studies (4.2\\% for SDSS and 3.6\\% for IRAC) approximately four counterparts are expected to disagree between the two surveys, which is consistent with that observed here.  The two sources with different counterparts identified in IRAC and SDSS are H-ATLAS ID-2244 and 6962. We show one of these sources, H-ATLAS ID-2244, as an example in figure~\\ref{fig:postage1}. In this case the SDSS (SDP.2244) and IRAC (H-ATLAS ID-2244) counterparts are separated by 3.2\\arcsec\\ and both have a $\\le$ 3\\% probability of being false identifications: $R=$ 0.97 in SDSS and $R=$  0.99 at both 3.6 and 4.5\\,\\micron. In the absence of additional data, such as sub-mm or radio interferometry, the true nature of this source is unclear. It is possible that one of the counterparts is a chance association, although we cannot say which one. It is also possible that the SPIRE source is comprised of a blend of emission from both counterparts in a single SPIRE beam.  Blending is also an important consideration for counterparts that are identified with the color-magnitude method.  Seven of the 13 SPIRE sources identified with this method have multiple IRAC counterparts; conversely the LR method has none. This apparent disparity is not surprising because the LR method implicitly assumes that each SPIRE source has a single IRAC counterpart. Where there are $N$ multiple counterparts contributing equally to the sub-mm flux the average $R$ cannot exceed $1/N$, and will typically be $\\sim 1/N$.  Due to the large SPIRE beam (FWHM = 18.1\\arcsec\\ at 250\\um) it is likely than in at least a few cases a single SPIRE source may be composed of blended emission from multiple galaxies or from galaxy interactions. The color-magnitude method does not consider the presence of other IRAC sources and is therefore not biased against multiple counterparts. In addition, beyond requiring that counterparts to lie within the SPIRE counterpart search radius, the color-magnitude method does not consider the separation between the SPIRE and IRAC centroids and therefore it is not biased against wide-separation counterparts. However, it does require both 3.6 and 4.5\\um\\ fluxes and it assumes that all SPIRE sources have a similar color-magnitude distribution at 3.6 and 4.5\\um.  The seven SPIRE sources in the sample that have multiple IRAC counterparts are likely to be a combination of blended SPIRE sources and individual sub-mm galaxies  associated with multiple interacting IRAC sources. \\atlas\\ ID-891 is shown as an example in figure~\\ref{fig:postage2}. In this case all the detected SDSS galaxies are $\\ga 3.5\\arcsec$ away from the SPIRE centroid and all have $R<$ 0.8 (Smith et al.\\,2011).  There are two potential counterparts in the IRAC data that are not detected in SDSS.  Using the LR method alone these two nearby sources have reliability of $R=$ 0.41 and $0.44$ at 3.6\\um\\ and $R=$ 0.47  and  0.50  at 4.5\\um.  However, both galaxies have $[3.6]$ and $[4.5]$ that place them in the region of color-magnitude space with $<$ 20\\% probability of being random associations with the SPIRE emission and as such they are both identified as counterparts.  The two galaxies have similar colors, with $([3.6]-[4.5])=$ 0.38 and 0.43 mag, consistent with a redshift of $z\\gtrsim$ 1.4 (Papovich 2008).  The two IRAC sources are separated by 2.9\\arcsec\\ (24.9 kpc at $z=$ 2, 23.0 kpc at $z=$ 3.  Furthermore, the SPIRE source is bright ($S_{500}=$ 84 $\\pm$ 9mJy) and red in the sub-mm bands, with a rising spectrum from 250 to 500\\um, indicating a sub-mm photometric redshift of $z\\ga$ 3 (e.g.  Lapi et al.\\,2011), although we caution that the dust temperature and redshift are degenerate in the sub-mm bands.  It is likely that \\atlas\\ ID-891 is a blend of emission from the two identified counterparts and it is possible that the sub-mm emission is the result of a merger or interaction of two gas-rich galaxies (e.g.\\ Aravena et al.\\ 2010b). Spectroscopic data are required to verify this scenario. and we identify ID-891 as a target for additional follow-up, especially to identify the exact nature of the sub-mm emission.  H-ATLAS ID-827, shown in figure~\\ref{fig:postage3}, is the most complex system in our sample. In this case there are four IRAC counterparts identified on the basis of their colors and magnitudes. The four galaxies have a wide range of $[3.6]$ with values from 17.8~mag to 20.2~mag. The counterparts do not appeared to be clustered, as may be expected for gravitational lensing or a multi-component interaction, and each is separated from its nearest neighbor by $\\sim$ 4\\arcsec. The two brightest galaxies are detected by SDSS, but were not identified as SPIRE counterparts by Smith et al.\\ (2011); they have photometric redshifts of $z=$ 0.07 $\\pm$ 0.01 and $z=$ 0.20 $\\pm$ 0.03 for the westernmost and easternmost, respectively. The two fainter counterparts are not detected in SDSS and both have $[3.6]-[4.5]>$ 0, indicating that they are likely to be at $z\\gtrsim$ 1.4 (Papovich 2008). The morphology, astrometry and available redshift information suggests that gravitational lensing in unlikely and that \\atlas\\ ID-827 is most likely to be a blended source, which is comprised of two low-redshift and two high-redshift components.   \\begin{figure} \\includegraphics[width=9cm]{IRAC_Ks_ch1_ch12.ps} \\includegraphics[width=9cm]{IRAC_z_K_K_ch1.ps} \\caption{IRAC and VIKING color-color diagrams. {\\bf Top:} [$K_{\\rm s}$], $[3.6]$ and $[4.5]$. {\\bf Bottom:} [$z^{\\prime}$], [$K_{\\rm s}$] and $[3.6]$.  We highlight SPIRE counterparts and background sources that are within 7.2\\arcsec\\ of SPIRE centroids.  Tracks for Arp~220 (Silva et al.\\,1998), Mrk~231 (Berta\\, 2006) and the average 870-\\um\\ selected sub-mm galaxies (Wardlow et al.\\,2011) are shown, and an $A_{V} = 1$~mag reddening vector and the median error bar are plotted in the lower left-hand corner.  The SPIRE counterparts are slightly bluer than both these tracks suggesting that they may have less dust reddening and $A_V$ values that are 0.5 to 1 magnitudes smaller than the SED templates.  } \\label{fig:viking} \\end{figure}     \\begin{figure} \\includegraphics[width=9cm]{HATLAS-IRAC-Lfir-T.ps} \\includegraphics[width=9cm]{HATLAS-IRAC-z-Lfir.ps} \\caption{ {\\bf Top:} Dust temperature as a function of the infrared luminosity   $L_{\\rm IR}$ (8 - 1100\\um) for H-ATLAS/IRAC galaxies with SDSS redshifts.   $L_{\\rm IR}$ is obtained by fitting the SPIRE flux with $\\beta$ = 1.5 in isothermal modified black-body model (equation~\\ref{eq:modified}). Error bars are median 1$\\sigma$ standard deviation   from   the best-fit models with dust temperature ($T_{\\rm d}$) and total IR luminosities  $L_{\\rm IR}$ as free parameters. {\\bf Bottom:}   Dust temperature as a function of the redshift for H-ATLAS/IRAC galaxies. See text for the comparison galaxy samples plotted here. } \\label{fig:lfir} \\end{figure}  \\subsection{Redshifts of Identified Sources}   Knowledge of the redshift distribution of SPIRE sources is central to understanding their role in the Universe. However, only 50\\% of the IRAC counterparts are detected in SDSS and 74\\% have VIKING photometry. Therefore, any redshift distribution derived from the optical data will be biased to low redshifts where the detection rates in these surveys are high.  Instead, in figure~\\ref{fig:stat} we plot the $[3.6]-[4.5]$ color distribution, and $[3.6]-[4.5]$ against redshift for the counterparts with spectroscopic or photometric redshift from SDSS.  This plot includes the 104 IRAC counterparts that have both 3.6 and 4.5\\um\\ coverage, 47 of which have an optical redshift. All of the optical redshift are $z<$ 0.8; 32 (68\\%) of these galaxies have $([3.6]-[4.5]) <$ 0, while 15 (32\\%) have $([3.6]-[4.5]) >$ 0.  The 57 identifications without reliable redshifts from the optical are undetected in SDSS and 43 of them have $([3.6]-[4.5]) >$ 0, a crude indicator that $z\\gtrsim1.4$ (e.g. Papovich 2008). Thus, up to $\\sim$ 40\\% of the identified SPIRE population may lie at high redshifts. Deep spectroscopy and photometry, particularly at near-IR wavelengths, is required to determine more precise redshift information.  We also investigate the redshift distribution of the SPIRE sources by calculating sub-mm photometric redshifts from the fluxes of the three SPIRE bands. We employ a template fitting method which uses $\\chi^2$ minimization, comparing the observed fluxes to the SEDs of Arp~220 (Silva et al.\\, 1998), Mrk~231 (Berta\\, 2006). We caution that the sub-mm photometric redshifts are limited by the absence of longer wavelength data, and the degeneracy between redshift and dust temperature. Furthermore, flux boosting is an important consideration and will add significant errors to the fluxes of the faintest sources.  Due to these uncertainties, we do not consider the sub-mm photometric redshifts on a source-by-source basis, but instead examine the overall distribution.  \\begin{figure} \\includegraphics[width=9cm]{ssfr_z.ps} \\includegraphics[width=9cm]{sfr_mass_MS.ps} \\caption{ {\\bf  Top:} Specific star formation rate (sSFR) for our IRAC-identified SPIRE galaxies as a function of the redshift. Filled circles  and triangles  represent  H-ATLAS galaxies in stellar mass bins  10$^{9.8}$-  10$^{10.8}$ $M_\\odot$ and 10$^{10.8}$-  10$^{11.5}$ $M_\\odot$, respectively.   Our data are compared to  P\\'erez-Goz\\'alez et al.\\,(2008)  and  Damen et al.\\,(2009). Upper (lower)  two lines  are for the  smaller (larger)  mass bin.  The error bar on the top left shows the average   68\\% confidence range on  $log$(sSFR). {\\bf Bottom:}  Star formation rate as a function of the stellar mass. We plot results from a previous study for comparison:  PACS-COSMOS   and PACS-GOODS South (Rodighiero et al. 2011). We also show the mean  SFR and stellar mass correlation  at $z=2$ (Daddi et al. 2007),  $1$  and $0$ (Elbaz et al. 2007). The majority of the H-ATLAS/IRAC galaxies  fall  between $z = 0$ $\\sim 2$  in SFR-$M_\\star$ relation showing a slight excess of  star formation rate compared with populations at similar redshifts, but selected at optical and near-IR wavelengths. } \\label{fig:ssfr_z} \\end{figure}  This sub-mm photometric redshift distribution peaks at $z\\sim 2$ with a tail out to $z\\sim$ 4, and suggests that the majority of the IRAC counterparts without SDSS redshifts lie at $z\\sim 2$. This result is consistent with the photometric redshifts of the whole SPIRE population (e.g. Amblard et al.\\,2010; Lapi et al.\\,2011) and the redshift distribution of 850\\um\\ sources (e.g.  Chapman et al.\\,2005; Wardlow et al.\\,2011). Now that the counterparts to a significant majority of SPIRE sources are known,  follow-up spectroscopic campaigns to establish the redshifts of all the sources are feasible.   \\subsection{Properties of Identified Sources}   {\\it Near-IR Colors}: In figure~\\ref{fig:viking} we show IRAC and VIKING color-color diagrams, highlighting identified SPIRE counterparts in comparison with background sources that are within 7.2\\arcsec\\ of the SPIRE centroids. As we show in figure~\\ref{fig:color} the SPIRE counterparts have distinct $[3.6]-[4.5]$ color compared to  the background sources, but the two populations are indistinguishable in $[K_{\\rm s}]-[3.6]$ and $[{\\rm z}]-[K_{\\rm s}]$. We also compare to the redshift tracks of Arp~220 (Silva et al.\\,1998), Mrk~231 (Berta\\, 2006) and the average 870-\\um\\ selected sub-mm galaxy (Wardlow et al.\\,2011).  The SPIRE counterparts are slightly bluer than both these tracks suggesting that they may have less dust reddening and $A_V$ values that are 0.5 to 1 magnitudes smaller than the SED templates.   {\\it IR Luminosities and Dust Temperatures}: To understand  properties of the identified H-ATLAS/IRAC galaxies associated with SDSS redshifts, we first consider a SED analysis of the SPIRE data making use of the redshifts based on the IRAC identification. First we consider the far-IR/sub-mm portion of the SED with 3-band SPIRE fluxes at 250, 350, and 500\\um\\ and reproduce the analysis of Amblard et al. (2010) where the dust temperature ($T_{\\rm d}$) and IR luminosities $L_{\\rm IR}$ were studied. We make use of a modified black-body with isothermal dust temperature to describe the dust emission, in which flux can be written as \\begin{equation} f_\\mu \\propto \\mu^{3+\\beta}e^{(h\\mu/kT_{\\rm d}-1)} \\,. \\label{eq:modified} \\end{equation} To be consistent with previous other estimates from the literature we fix $\\beta=1.5$ and fit to the SPIRE fluxes to establish the dust temperature and the overall normalization to the SED.  The IR luminosity is estimated from the best-fit SED over the wavelength range of 8 to 1100\\um. Since we only fit to the dust emission at wavelengths greater than 250\\um\\ our luminosities are likely low by a factor of 2 to 3 if there are any warm dust, heated by AGNs, present in these galaxies. However, our overall uncertainty on the IR luminosity estimates is at least a factor of 5. For the sample as a whole we find an average dust temperature of  26 $\\pm$ 9 K, which compares well with the average dust temperature of 27 $\\pm$ 8 K for $z<  $ 0.1 SPIRE galaxies based on SDSS identifications only (Amblard et al. 2010).  We summarize our results related to the FIR SED analysis in figure~\\ref{fig:lfir} where we show the  $L_{\\rm IR}$ - $T_{\\rm d}$ relation and the $L_{\\rm IR}$ as a function of the redshift, determined from the reliable H-ATLAS IDs with either SDSS spectroscopic or photometric redshift. We also compare the H-ATLAS dust temperatures and IR luminosities with sub-mm bright galaxy samples in the literature. Samples include  the sources in BLAST with COMBO-17 (Wolf et al. 2004) or a SWIRE photometric redshift (Rowan-Robinson et al. 2008) and {\\it Spitzer}-MIPS 70 and 160\\um\\ fluxes (Dye et al. 2009),   local ULIRGS observed with SCUBA at 450 and 850\\um\\  (Clements et al. 2010), and local IRAS-selected galaxies with 60 and 100\\um\\ fluxes along with SCUBA 850\\um\\ (Dunne et al. 2000).  In figure~\\ref{fig:lfir}  we find that the SPIRE-selected galaxies in H-ATLAS with IRAC-based identifications for redshifts span the luminosity range of 10$^{10}$ to 10$^{12}$  $L_{\\sun}$ from sub-LIRG luminosities to ULIRG conditions. These luminosities and the ranges are comparable to BLAST-detected sources, while they span lower than the local IRAS-selected galaxies, which tend to be at lower redshifts $z<0.05$ and have higher luminosities; such a difference is understandable since IRAS is an all-sky shallow survey and detects primarily the rare, bright galaxies in the near-by universe.  {\\it Stellar Masses and Star-Formation Rates}: With the cross-identification  we can expand the SED analysis over 3 orders of magnitude in wavelength from optical to sub-mm. For optical and near-IR data we make use of SDSS, UKIDSS (Lawrence et al.\\,2007), VIKING (Sutherland et al.\\ in prep), and IRAC 3.6 and 4.5\\um\\ fluxes. The analysis of full optical to sub-mm SEDs is similar in spirit to Dunne et al. (2011) where the SEDs of SPIRE sources with reliable identifications using SDSS data were analyzed. We make use of the {\\sc MAGPHYS} SED modelling code (da Cunha, Charlot \\& Elbaz 2008) for this work. The SED modelling in the code is done in a energy-balance manner such that the absorbed light in shorter UV and optical wavelengths is accounted for by the thermal re-radiation at far-infrared and sub-mm wavelengths.  For the H-ATLAS/IRAC galaxy sample we have 5 optical fluxes from SDSS ($u, g, r, i, z$), 5 near-IR fluxes from UKIDSS+VIKING ($Z, Y, J, H, K_{\\rm s}$), and 2 IRAC channels (3.6, 4.5\\um), in addition to 3 SPIRE bands. Redshift information of the sources are   from SDSS spectroscopic observations ($z= $0 - 0.8).   For comparative work with other sub-mm galaxy samples we focus on results related to stellar mass $M_\\star$ and star-formation rate (SFR), as given by {\\sc MAGPHYS} best-fit models. In figure~\\ref{fig:ssfr_z} (top figure) we plot the specific star-formation rate (sSFR) given as SFR/$M_{\\star}$ vs. redshift with our sample sub-divided to two stellar mass bins. We find that the sSFR of  IRAC IDs  marginally show an anticorrelation with galaxy  stellar mass. Such a behavior may be related to the {\\it downsizing}  scenario (Cowie et al. 1996).   However, when compared  to Damen et al.\\,(2009) and P\\'erez-Goz\\'alez et al.\\,(2008), IRAC IDs shows  0.4 $\\sim $ 0.5 dex higher sSFRs. The Damen et al. (2009) galaxy sample is an IRAC-selected sample of galaxies out to $z \\sim $ 1.8 in ECDFS with photometry supplemented with optical, near-IR, and MIPS 24\\um\\ fluxes. The P\\'erez-Goz\\'alez et al.\\,(2008) sample is selected with IRAC at 3.6 and 4.5\\um\\ in a variety of fields. While our SPIRE-selected sample has IRAC identifications,  they are not likely to be the majority of the galaxies in both these studies. This is also clear from  figure~\\ref{fig:color}  where we show that  most of the identified SPIRE sources are distributed  differently from the background population in the color-magnitude space. They occupy the top-end of the star-formation in galaxies selected at near-IR wavelengths.  In Fig.~12 (bottom) we plot the SFR vs. the stellar mass of our sample. The SFR of   {\\it normal} local  star-forming galaxies at $z \\sim 1$  is known to correlate strongly with the stellar mass and form the so-called {\\it star-formation main-sequence} (Brinchmann et al. 2004; Salim et al. 2007; Peng et al. 2010). Such a correlation is also observed at higher redshifts, albeit with a different normalization (e.g. Daddi et al. 2007; Elbaz et al. 2007; Noeske et al.  2007; Pannella et al. 2009; Daddi et al. 2009; Gonzalez et al. 2010; Rodighiero et al. 2010; Karim et al. 2011).   However,  the  typical dispersion  of correlation is known to be  $\\sim$ 0.3 dex  over a wide range of redshifts $z=$ 0 $\\sim$ 2.  Local (U)LIRGs  and high-redshift SMGs  have SFRs in excess of the main-sequence and are defined as starburst galaxies (e.g. Elbaz et al. 2007, 2011).  These galaxies generally occupy a region that is $\\sim$ $\\times$10 above the  main-sequence. However, there is also   evidence that some high-redshift (U)LIRGs and SMGs do not have enhanced star-formation rates relative to their stellar mass and that the starburst fraction may decrease at high redshifts (e.g. Tacconi et al. 2008; Rodighiero et al. 2011). This suggests changes to the mode of star-formation in (U)LIRGs between $z=  $ 0  and $z= $ 2 away from the starburst mechanism (e.g. Daddi et al. 2010b; Genzel et al. 2010; Elbaz et al. 2011; Krumholz et al. 2012; Melbourne et al. 2012). In our work, $\\sim$ 50\\% of the subsample of IRAC-identified SPIRE galaxies  ($z < $ 1) lie near the starburst region. Thus, it is possible that half of the IRAC-identified SPIRE galaxies with SDSS redshifts are in fact starbursts that lie off the main-sequence."
    },
    "1112/1112.1710_arXiv.txt": {
        "abstract": "Current Adaptive Mesh Refinement (AMR) simulations require algorithms that are highly parallelized and manage memory efficiently.  As compute engines grow larger, AMR simulations will require algorithms that achieve new levels of efficient parallelization and memory management.  We have attempted to employ new techniques to achieve both of these goals.  Patch or grid based AMR often employs ghost cells to decouple the hyperbolic advances of each grid on a given refinement level.  This decoupling allows each grid to be advanced independently.  In AstroBEAR we utilize this independence by threading the grid advances on each level with preference going to the finer level grids.  This allows for global load balancing instead of level by level load balancing and allows for greater parallelization across both physical space and AMR level.  Threading of level advances can also improve performance by interleaving communication with computation, especially in deep simulations with many levels of refinement.  While we see improvements of up to $30 \\%$ on deep simulations run on a few cores, the speedup is typically more modest ($5-20 \\%$) for larger scale simulations.  To improve memory management we have employed a distributed tree algorithm that requires processors to only store and communicate local sections of the AMR tree structure with neighboring processors.  Using this distributed approach we are able to get reasonable scaling efficiency ($ > 80 \\%$) out to 12288 cores and up to 8 levels of AMR - independent of the use of threading. ",
        "introduction": "The development of AMR \\citep{Berger1984484, Berger198964, Balsara2001} was meant to provide high resolution simulations for much lower computational cost than fixed grid methods would allow.  The use of highly parallel systems and the algorithms that go with them were also meant to allow higher resolution simulations to be run faster (relative to wall clock time).  The parallelization of AMR algorithms, which should combine the cost/time savings of both methods is not straight forward however and there have been many different approaches \\citep{Paramesh2000, Nirvana2008, Enzo2004, Chombo, FTT} .  While parallelization of a uniform mesh demands little communication between processors, AMR methods can demand considerable communication to maintain data consistency across the unstructured mesh as well as shuffling new grids from one processor to another to balance workload.  In this paper we report the development and implementation of new algorithms for the efficient parallelization of AMR designed to scale to very large simulations. The new algorithms are part of the AstroBEAR package for simulation of astrophysical fluid multi-physics problems \\citep{Cunningham2009ApJS}.  The new algorithmic structure described in this paper constitutes the development of version 2.0 of the AstroBEAR code.  AMR methods come in many varieties.  Meshes can either be unstructured or semi-structured.  Semi-structured methods can be further divided into those which allow grids to be of arbitrary size (patch based) and those which require grids to be of a fixed size (block based or cell-based if the block size is 1).  With block (or cell) based AMR, the additional constraints imposed on the structure of the mesh allow for a simpler type of connectivity within a tree.  For example in block based AMR, any given block will have exactly 8 children or none (if it is a leaf) and will have at most 6 face sharing neighbors.  With patch based AMR, there is no limit to the number of children or neighbors.  In addition, the operation of regridding in block based AMR is much simpler.  As the grid changes a given block will either persist if the physical region continues to require refinement or be destroyed. In patch based AMR, a given region may subsequently be better covered by patches of a different shape requiring transfer of data between physically overlapping previous patches and new patches.  This adds an additional dimension to the tree structure and increases the complexity of maintaining a distributed tree.  For both block (or cell) and patch based AMR, the actual grid data (fluid variables etc.) are always distributed across the various processors.  Usually some overlap in grid data (guard/ghost cells) is desired to allow for frequent access to neighboring values without the need for additional communication.  But the metadata that describes the shape and distribution of the grid data is usually stored on every processor.  For 100's or 1000's of cores, this global tree typically requires less memory than that required for the local grid data and it allows for easy access to data from any part of the domain.  For instance after regridding, the entire tree can be updated and stored locally and then finding new neighbors for local cells/blocks/patches can be done without the need for any further communication.  For patch-based AMR it also allows for 'better' load balancing as each processor can determine from the entire tree which section of data it should be responsible for and can use various knap-sack type algorithms to optimize the degree of interprocessor communication.  Cell based AMR engines typically use an Octree data structure to handle the tree metadata and implementations have been developed that support a memory-distributed tree \\citep{DistributedOctTree, Octor}, or various ways of compressing the global tree \\citep{CompressedOctTrees, FTT} that rely on the simple structure of the tree.  For patch based AMR, algorithms for implementing a distributed tree have not yet been published though implemented in the package SAMRAI \\citep{Chombo} as well as AstroBEAR.  In addition the Chombo library \\citep{Chombo} recently has developed a method of compressing the metadata to avoid having to distribute the tree.   Here we document the distributed tree algorithm used in AstroBEAR 2.0 in which no processor has access to the entire tree but rather each processor is only aware of the AMR structure it needs to manage in order to carry out its computations and perform the necessary communications.  While currently, this additional memory is small compared to the resources typically available to a CPU, future clusters will likely have much less memory per processor similar to what is already seen in GPU's.  Additionally each processor only sends and receives the portions of the tree necessary to carry out its communication as opposed to a pruning approach in which every processor receives every new patch, but only keeps those necessary to maintain the local tree.  AstroBEAR 2.0 also uses extended ghost cells to decouple advances on various levels of refinement.  As we show below this allows for each level's advance to be computed independently on separate threads.  Such inter-level threading allows for total load balancing across all refinement levels instead of balancing each level independently.   Independent load balancing becomes especially important for deep simulations (simulations with low filling fractions but many levels of AMR) as opposed to shallow simulations (high filling fractions and only a few levels of AMR).  Processors with coarse grids can advance their grids simultaneously while processors with finer grids advance theirs.  Without such a capability, each level would need to have enough cells to be able to be distributed across all of the processors.  Variations in the filling fractions from level to level can make the number of cells on each level very different.  If there are enough cells on the level with the fewest to be adequately distributed, there will likely be far too many cells on the highest level to allow the computation to be completed in a reasonable wall clock time.  This often restricts the number of levels of AMR that can be practically used.  With inter-level threading this restriction is lifted.  Inter-level threading also allows processors to remain busy while waiting for messages from other processors.  In what follows we provide descriptions of the new code and its structure as well as providing tests which demonstrate its effective scaling.  In section~\\ref{amr_alg} we review patch based AMR.  In section~\\ref{distributedtree} we describe in detailt the distributed tree algorithm for patch-based AMR, in section~\\ref{threading} we will discuss the inter-level threading of the advance, in section~\\ref{loadbalancing} we will discuss the load balancing algorithm, and in section~\\ref{results} we will present our scaling results and we will conclude in section~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion} As discussed in the introduction Adaptive Mesh Refinement methods were designed to allow high resolution simulations to be carried out at low computational cost.  Highly parallel systems and their algorithms carried the same promise.  The parallelization of AMR algorithms is, however, not straight forward however and there have been a number of different approaches \\citep{Paramesh2000, Nirvana2008, Enzo2004, Chombo, FTT}, to solve the problem.  While parallelization of a uniform mesh demands little communication between processors, AMR methods can demand considerable communication to maintain data consistency across the unstructured mesh as well as shuffling new grids from one processor to another to balance workload.  In this paper we have described our attempt to design and implement a new strategy for AMR parallelization with an eye towards running codes on large machines with $>10^4$ processors. We have found that a threaded approach to the AMR algorithm significantly improves performance by allowing processors to remain busy while waiting for messages as well as to dynamically adjust distributions based on the progress of ongoing coarser grid updates.  We have also shown that a distributed tree algorithm significantly reduces the amount of memory required (and corresponding communication) for simulations run on many processors with modest grid sizes.  We believe this threaded, distributed tree approach described in this paper holds considerable promise as a methodology for implementing AMR on ever-larger numbers of processors.  Future work will entail elaborating refinements to the methods as we experiment with its implementation on different classes of problems and machines."
    },
    "1112/1112.2704.txt": {
        "abstract": "\\vspace{0.2cm} \\noindent We explore models in which weakly interacting massive particle (WIMP) dark matter annihilation is directly responsible for baryogenesis, thereby connecting dark matter with baryogenesis. We call this process ``WIMPy baryogenesis''. The dark matter relic density in these models, as with conventional WIMP models, is obtained with only order one couplings and TeV-scale masses according to the WIMP miracle. Thus, WIMPy baryogenesis models naturally accommodate weak-scale dark matter. Furthermore, an extension of the WIMP miracle simultaneously explains the observed baryon asymmetry and the correct dark matter abundance. The models we present have the further feature that they  create the baryon number asymmetry at the weak scale, thereby avoiding the problems in some models of baryogenesis associated with high reheat temperatures in supersymmetric theories. Some of these models yield observable consequences in ongoing and future experiments. ",
        "introduction": "Generally, baryogenesis and the establishment of the dark matter number density are treated as independent processes.   With the notable exception of one class of models \\cite{McDonald:2011zz},  it has largely been overlooked that models with symmetric, weakly interacting massive particle (WIMP) dark matter can connect dark matter physics with baryogenesis. We present a new mechanism that creates such a link and is based on a simple premise: if WIMP annihilation satisfies the Sakharov conditions, a non-zero baryon number asymmetry can be generated from dark matter annihilation, and in some instances, can account for the entire observed baryon asymmetry. We call this process WIMPy baryogenesis. Our models are distinct from models of asymmetric dark matter, which propose that dark matter and baryons have their origins in a common asymmetry. In our models, the energy densities of baryons and dark matter are more loosely linked but can accommodate the observed dark-matter-to-baryon ratio.  \\cite{Nussinov:1985xr,Kaplan:2009ag,Cohen:2009fz,Buckley:2010ui,Cui:2011qe,Buckley:2011ye}  We list below the  Sakharov conditions and how they are satisfied in WIMPy baryogenesis: % \\begin{enumerate} \\item \\emph{Baryon number violation:} WIMP annihilations violate  baryon or lepton number. A preserved $\\mathrm U(1)$  symmetry is allowed if the  baryon asymmetry is balanced by a negative asymmetry in a decoupled sector  that restores the net global symmetry. We have such a $\\rm U(1)$ symmetry in most models we present. \\item \\emph{$CP$ violation:} WIMP couplings to Standard Model fields violate $CP$. \\item \\emph{Departure from thermal equilibrium:} The cooling of the universe provides the necessary departure from thermal equilibrium. Net dark matter annihilation begins around temperatures $T\\lesssim m_{\\rm DM}$, resulting in a small deviation of the dark matter number density from its equilibrium value. The annihilations can generate a baryon asymmetry that depends on the amount of dark matter annihilation occurring during washout freeze-out, which is comparable to the dark matter density at that time.  \\end{enumerate} % We present models that satisfy all three Sakharov conditions, and that simultaneously generate the observed baryon asymmetry and  WIMP relic density. In particular, we find that there exist successful models of WIMPy baryogenesis with $\\mathcal O(1)$ couplings and CP phases, and weak-scale masses for all new fields. This is an extension of the WIMP miracle to also include baryogenesis, although we show that the size of the generated asymmetry is sensitive to the parameters in the theory and can vary by several orders of magnitude from the observed asymmetry.  %%%%%%%%%%%%%%%%% \\begin{figure} \\begin{center} \\includegraphics[width=7cm]{img/FlowSymmetry}\\hspace{1cm} \\includegraphics[width=6cm]{img/FlowViolated} \\caption{Schematic diagrams showing the evolution of the asymmetry created by dark matter annihilation. {\\bf (left)} Model where asymmetry created in exotic antibaryons is sequestered in a sterile sector through baryon-number-conserving decays. {\\bf (right)} Model where asymmetry created in exotic antibaryons is converted into a Standard Model baryon asymmetry through baryon-number-violating decays. }\\label{fig:flow} \\end{center} \\end{figure} %%%%%%%%%%%%%%%%%   Although WIMP annihilation can  generate a baryon asymmetry, there are other  processes that have the potential to \\emph{wash out} the asymmetry, and their freeze-out is crucial to create the observed baryon asymmetry. In our models, the two leading sources of washout are inverse annihilations of baryons into dark matter and baryon-to-antibaryon processes. Washout scatterings must be suppressed to generate a sizeable baryon asymmetry because, as we show in Section \\ref{sec:wimpybaryogenesis}, any asymmetry generated prior to washout freeze-out\\footnote{The time of washout  freeze-out is defined as when the rate of washout processes falls below the Hubble expansion rate. This is analogous to the freeze-out of WIMP annihilation.} is rapidly damped away. After washout processes freeze out, dark matter annihilations can efficiently create a baryon asymmetry, and the final asymmetry depends on how much dark matter remains when washout scatterings freeze out. Washout freeze-out must occur before that of WIMP annihilation, at which point dark matter annihilation is no longer efficient and no sizeable asymmetry can be created.  Thus, we find our central result: \\emph{if washout processes freeze out  before WIMP freeze-out, then a large baryon asymmetry may accumulate, and its final value is proportional to the WIMP abundance at the time that washout becomes inefficient}.   The early freeze-out of washout processes can occur for kinematic reasons. Inverse annihilations will be Boltzmann-suppressed for $T<m_{\\rm DM}$ because the thermal baryon fields are no longer energetic enough to annihilate back into dark matter. Baryon-antibaryon scatterings, however, can remain rapid at temperatures well below $m_{\\rm DM}$. The only way to suppress baryon-to-antibaryon washout is if all washout processes involve a heavy exotic baryon field in the initial state. We illustrate this scenario in Figure \\ref{fig:flow}, showing how dark matter annihilates to Standard Model baryons plus an exotic baryon, as well as the possible decays of the exotic baryon (either through baryon-preserving or baryon-violating interactions). If this exotic field has a mass $\\gtrsim m_{\\rm DM}$, its abundance is Boltzmann-suppressed at $T<m_{\\rm DM}$ and suppresses the washout rate. Meanwhile, dark matter annihilations are not kinematically allowed if the heavy baryon field has mass $\\gtrsim2m_{\\rm DM}$, so the mass condition  $m_{\\rm DM}\\lesssim m_{\\rm exotic\\,baryon}\\lesssim 2m_{\\rm DM}$ is essential to generate a large baryon asymmetry through WIMPy baryogenesis.  Dark matter annihilations  generate a positive baryon asymmetry stored in Standard Model quarks along with an equal negative  asymmetry stored in the exotic baryon field. It is important that the decays of the exotic baryon do not eliminate the Standard Model baryon asymmetry. In models of WIMPy baryogenesis with a preserved  $\\rm U(1)$ baryon symmetry, the exotic baryon-number-carrying field is charged under an additional discrete symmetry, while Standard Model fields are uncharged, preventing the exotic baryon from decaying into Standard Model baryons and destroying the asymmetry. The heavy baryon-number-carrying field decays instead into light gauge singlet fields that are  charged under the discrete symmetry and decoupled from Standard Model fields at temperatures below the scale of WIMPy baryogenesis. We also present a model where the exotic baryon decays to Standard Model quarks through baryon-number-violating couplings, and such models manifestly satisfy the first Sakharov condition.   For many years, the WIMP miracle -- the fact that dark matter fields with weak-scale masses and annihilation cross sections give the correct dark matter thermal relic density -- has been a compelling paradigm for dark matter model-building. WIMPy baryogenesis preserves the WIMP miracle while also offering an explanation for the observed baryon asymmetry. While WIMPy baryogenesis models do not predict the precise relationship between the dark matter and baryon number densities, natural models do restrict the baryon asymmetry to a range of about seven orders of magnitude (see Section \\ref{sec:wimpybaryogenesis}), and the observed asymmetry is within this range. Since baryogenesis arises from WIMP annihilations, WIMPy baryogenesis is also necessarily connected to weak-scale physics. While we do not discuss an embedding of WIMPy baryogenesis in a particular solution of the  hierarchy problem, we assume that whatever new physics lies at the weak scale stabilizes any scalar potentials in our theory and gives a natural explanation for their weak-scale masses. A consequence of this is that, with weak-scale masses, some of the new fields necessary for baryogenesis may give signals at future experiments. Additionally, the present-day dark matter is symmetric, leading to the possibility of the indirect detection of dark matter annihilations as in conventional WIMP models. This is in contrast with generic asymmetric dark matter models, in which the majority of dark matter annihilations ceased long before the present day, although it is also noteworthy that there do exist scenarios in which the symmetric component of dark matter is regenerated at late times, giving indirect detection signals for some asymmetric dark matter models \\cite{Buckley:2011ye}.  A further advantage of this scenario is that bounds on the reheat temperature in supersymmetric models do not constrain WIMPy baryogenesis. Typical reheat temperature constraints come from overproduction of gravitinos and are in the range $T_{\\rm RH}\\lesssim10^6-10^9$ GeV \\cite{Khlopov:1984pf}, and $T_{\\rm RH}$ is consequently below the scale required for conventional leptogenesis through the decay of heavy, Majorana right-handed neutrinos. Although low-scale mechanisms for baryogenesis are known, such as electroweak baryogenesis, WIMPy baryogenesis is a new way of generating the baryon asymmetry at $T\\sim$ TeV while satisfying the reheat bound.  We discuss the general conditions  for successful WIMPy baryogenesis in Section \\ref{sec:wimpybaryogenesis}, finding that interactions washing out the baryon asymmetry must become ineffective prior to WIMP freeze-out in order to generate the observed asymmetry. In Section \\ref{sec:model}, we focus on a particular model where dark matter annihilates through a lepton-number-violating interaction and the asymmetry is subsequently transferred to baryons by sphalerons. Because sphaleron processes are only rapid in the unbroken electroweak phase, such baryogenesis must occur before the electroweak phase transition. We compute the dark matter relic density and baryon asymmetry, and find the range of masses and couplings that agrees with the observed densities of both. We also consider the implications of additional lepton-number-conserving dark matter annihilation channels. In Section \\ref{sec:alternatives}, we consider  models with WIMPs annihilating directly to quarks, where baryogenesis can occur over a wider range of temperatures because sphalerons are no longer needed to establish the baryon asymmetry. We discuss experimental constraints and possible signals for  models of WIMPy baryogenesis in Section \\ref{sec:signals}. Finally, we summarize in Section \\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  In this paper, we explored a novel scenario called WIMPy baryogenesis that extends the WIMP miracle by generating the observed baryon asymmetry through annihilations of weak-scale dark matter and provides a dynamical connection between the dark matter and baryon abundances. We found that natural couplings and weak-scale masses for the new fields can lead to the correct baryon asymmetry. As a by-product of linking baryogenesis with dark matter annihilation, we introduced a new mechanism for weak-scale baryogenesis, avoiding any conflicts with reheat bounds in supersymmetric theories.  The key observation is that, if dark matter annihilation proceeds via $CP-$ and (Standard Model) $B-$ or $L-$violating operators, all Sakharov conditions for baryogenesis are satisfied. Successful models  also suppress washout prior to dark matter freeze-out. In our models, such suppression results from the heaviness of the field $\\psi$ carrying Standard Model gauge charges and $B-$ or $L-$number that is one of the final states in dark matter annihilation. Additional discrete symmetries forbid such  exotic fields from decaying back to Standard Model fields. We presented models where dark matter annihilates to either quarks or leptons, and  found viable parameter spaces with natural couplings and TeV-scale masses in both scenarios.  In models where dark matter annihilates to leptons, the lepton asymmetry must be generated before the electroweak phase transition so that the asymmetry can be transferred to baryons via sphalerons. As a result, dark matter and $\\psi$ masses must be $\\mathcal O(\\mathrm{TeV})$. Because the new states are heavy and dark matter does not couple directly to quarks, dark matter in this set-up is not in reach of near-future direct and indirect detection experiments. In this scenario, $\\psi$ is  charged only under  weak interactions, making LHC searches challenging, although targeted searches at high integrated luminosity may allow discovery.  If dark matter annihilates to quarks, baryogenesis can occur after the electroweak phase transition, allowing smaller dark matter and $\\psi$ masses. With lighter new states and colored objects, dark matter in these models can be within reach of direct detection experiments and indirect detection searches for antideuterons. LHC searches for $\\psi$  are similar to gluino searches and can exclude $m_\\psi\\lesssim1.44$ TeV at $100\\,\\,\\mathrm{fb}^{-1}$ and $\\sqrt s=14$ TeV.  In both scenarios, WIMPy baryogenesis models can generate both the correct dark matter relic density and the baryon asymmetry at the weak scale. Such models predict new weak-scale particles that can lead to signals in dark matter direct and indirect detection experiments, and that may be accessible at the LHC."
    },
    "1112/1112.3842_arXiv.txt": {
        "abstract": "The neutrino-driven wind, which occurs after the onset of a core-collapse supernova explosion, has long been considered as the possible site for the synthesis of heavy $r$-process elements in the Universe. Only recently, it has been possible to simulate supernova explosions up to $\\sim10$~seconds, based on three-flavor Boltzmann neutrino transport. These simulations show that the neutrino luminosities and spectra of all flavors are very similar and their difference even decreases during the deleptonization of the proto-neutron star. As a consequence, the ejecta are always proton rich which rules out the possible production of heavy $r$-process elements ($Z>56$).  We perform a detailed analysis of the different weak processes that determine the neutrino spectra. Non-electron flavor (anti)neutrinos are produced and interact only via neutral-current processes, while electron (anti)neutrinos have additional contributions from charge-current processes. The latter are dominated by $\\nu_e$-absorption on neutrons and $\\bar\\nu_e$-absorption on protons. At early times, charge-current processes are responsible for spectral differences between $\\nu_e$, $\\bar\\nu_e$ and $\\nu_{\\mu/\\tau}$. However, as the region of neutrino decoupling moves to higher densities during deleptonization, charge-current reactions are suppressed by final state Pauli-blocking. $\\bar\\nu_e$ absorption on protons is suppressed due to the continuously increasing chemical potential of the neutrons. $\\nu_e$ absorption on neutrons is blocked by the increasing degeneracy of the electrons. These effects result in negligible contributions from charge-current reactions on timescales on the order of tens of seconds, depending on the progenitor star. Hence, the neutrino spectra are mainly determined from neutral-current processes which do not distinguish between the different flavors and results in the convergence of the spectra. These findings are independent of the charge-current reaction rates used. It rules out the possibility of neutron-rich ejecta at late times and the production of heavy $r$-process elements from non-rotating and not magnetized proto-neutron stars. ",
        "introduction": "Core collapse supernova explosions of stars more massive than $8$~M$_\\odot$ is an active subject of research in theoretical astrophysics~\\cite{Janka:2007}. They are related to the revival of the stalled shock, which forms when the collapsing core bounces back at nuclear densities. Right after the formation, the shock looses energy due to the dissociation of heavy nuclei and to the release of the $\\nu_e$-deleptonization burst produced by electron captures on protons. At this moment the shock becomes an accretion front and stalls at a radius between 100 and 200 km. Different mechanism have been suggested to revive the shock and finally explode the star. These are the magneto-rotational mechanism~\\cite{LeBlanc:1970kg,Moiseenko:2007zz,Takiwaki:2007sf}, the acoustic mechanism due to the dumping of sound waves~\\cite{Burrows:2005dv} and the delayed neutrino-heating mechanism~\\cite{Bethe:1985ux}.  Recently, it has been shown that a phase transition to deconfined quark matter can lead to an additional collapse during which a second hydrodynamic shock wave forms that triggers the explosion~\\cite{Sagert:2008ka,Fischer:2011}.  Independently of the explosion mechanism, a proto-neutron star (PNS) forms at the supernova core immediately after core bounce. The initially hot and lepton-rich PNS cools via continuous emission of neutrinos of all flavors in a period of several tens of seconds after the explosion has been launched. Neutrino-absorption processes at the PNS surface deposit energy, which drives a matter outflow known as the neutrino-driven wind. It has been investigated as a possible site for the production of heavy $r$-process elements~\\cite{Woosley:1994ux,Takahashi:1994yz}. These pioneering works where followed by analytic~\\cite{Qian:1996xt}, parametric~\\cite{Hoffman:1996aj} and steady state wind models~\\cite{Otsuki:1999kb,Thompson:2001ys}. They showed that the production of heavy $r$-process elements requires a mass outflow on a short dynamical time scale (a few milliseconds), high entropies per baryon (above 150~k$_\\text{B}$) and a low proton-to-baryon ratio, i.e. electron fraction $Y_e < 0.5$. Recent hydrodynamics simulations~\\cite{Arcones:2006uq,Arcones:2011a} showed that short dynamical timescales can in fact be achieved but fail however to obtain the necessary entropies at times relevant for $r$-process nucleosynthesis~\\cite{Kuroda.Wanajo.Nomoto:2008}.  The nucleosynthesis outcome is very sensitive to the electron fraction of the ejected matter~\\cite{Arcones:2011b} that is determined by the competition between electron neutrino absorption on neutrons and antineutrino absorption on protons and their inverse reactions. Deep in the PNS interior, neutrinos are in chemical equilibrium with matter. However, at the PNS surface where temperature and density drop, neutrinos decouple from matter. Since $\\mu$ and $\\tau$ neutrinos interact only via neutral-current reactions, they are the first to decouple at the highest temperature. For the very neutron-rich conditions found at the PNS surface, $\\bar\\nu_e$ decouple before $\\nu_e$. The neutrino spectra reflect local properties of matter at the position where they decouple. Hence, one expects the following hierarchy of neutrino energies: $\\varepsilon_{\\nu_{\\mu,\\tau}} > \\varepsilon_{\\bar{\\nu}_e} > \\varepsilon_{\\nu_e}$~\\cite{Keil:2003}, where $\\varepsilon = \\langle E^2\\rangle/\\langle E \\rangle$. $\\langle E \\rangle$ is the mean neutrino energy and $\\langle E^2 \\rangle$ is the square value of the root-mean-square (rms) energy.. The fact that $\\bar\\nu_e$ have larger average energies than $\\nu_e$ suggest that the ejecta will be neutron rich. However, one also has to consider the other energy scale in the problem, i.e. the neutron-to-proton mass difference, $\\triangle=m_n-m_p=1.2935$~MeV. It turns out that neutron-rich ejecta are only obtained when $\\varepsilon_{\\bar{\\nu}_e} - \\varepsilon_{\\nu_e} > 4\\,\\triangle$~\\cite{Qian:1996xt,Frohlich:2006} in the presence of similar electron neutrino and antineutrino luminosities.  The treatment of neutrino transport and neutrino matter interactions improved over the years significantly. Simulations of core-collapse supernovae, that include three-flavor Boltzmann neutrino transport, showed rather small neutrino energy differences between $\\bar\\nu_e$ and $\\nu_e$ and hence generally proton-rich conditions were obtained during the early evolution after the onset of explosion~\\cite{Liebendoerfer:2001a,Buras:2005rp}. More recently, it has been possible to simulate core-collapse supernovae consistently through all phases up to 20~seconds after the onset of explosion~\\cite{Fischer:2009af,Huedepohl:2010}, using general relativistic radiation hydrodynamics that employs three-flavor Boltzmann neutrino transport and a sophisticated equation of state. It has been shown that the energy difference between $\\bar\\nu_e$ and $\\nu_e$ decreases continuously after the explosion has been launched and hence generally proton-rich conditions remain for more than 10~seconds. This rules out initial expectations that the ejecta can become neutron-rich at late times. It leaves the $\\nu p$-process~\\cite{Frohlich:2005ys,Pruet:2006,Wanajo:2006ec} as the only nucleosynthesis process that can occur in neutrino-driven winds from non-rotating and not magnetized PNSs.  In this article, we analyze the neutrino spectra obtained in the current simulations focussing in understanding the reasons behind the reduction of the energy difference between $\\bar\\nu_e$ and $\\nu_e$ during the deleptonization of the PNS. We perform a detailed analysis of the main reactions determining the opacities or inverse mean free paths for the different neutrino flavors. Special attention is devoted to neutrino decoupling from matter, in particular the determination of the position of the neutrinospheres which evolves during the PNS deleptonization to continuously higher densities. We also present an alternative analysis, based on the formalism introduced in ref.~\\cite{Liebendoerfer:2004}, that determines the contributions to the neutrino luminosity from different processes and regions of the PNS. Both approaches show that at late times during the deleptonization, the spectra of all neutrino species forms at very similar radii and is determined by neutral-current processes that are insensitive to the neutrino flavor.  The paper is organizes as follows. In section~II we will introduce the procedure used to calculate the inverse mean free paths for the relevant weak interaction processes considered, based on the neutrino distribution functions obtained in three-flavor Boltzmann neutrino transport.  We will furthermore discuss neutrino decoupling and introduce the concept of a scattering atmosphere following the work of ref.~\\cite{Raffelt:2001}.  In section~III, we will apply this analysis and discuss the evolutionary behavior of the neutrino spectra obtained in core collapse supernova simulations of the low-mass 8.8~M$_\\odot$ O-Ne-Mg-core~\\cite{Fischer:2009af}.  In section~IV we extend the analysis from section~III to the 18~M$_\\odot$ iron-core progenitor from ref~\\cite{Fischer:2009af} and find qualitatively similar results. In section~V is show the evolution of the neutrino spectra at different times during the deleptonization. While we explored opacities in section~III and IV, in section~VI we analyze the neutrino emission, focusing on the location at which different contributions to the luminosity at infinity appear. We close the manuscript with a summary in section~VII. ",
        "conclusions": "We have performed a detailed analysis of the different processes that determine the neutrino spectra of all flavors during the deleptonization phase after the onset of supernova explosion. We have explored the 8.8~M$_\\odot$ O-Ne-Mg-core and the 18~M$_\\odot$ iron-core progenitors from ref.~\\cite{Fischer:2009af}, in order to cover a broad range of stellar models.  Using neutrino opacities (inverse mean free paths) for the different processes, we computed energy and transport spheres for each neutrino flavor. Our results confirm the finding of ref.~\\cite{Raffelt:2001,Keil:2003} for $\\nu_{\\mu/\\tau}$, i.e. the presence of a scattering atmosphere between the $\\nu_{\\mu/\\tau}$ energy and transport spheres. In addition we also find a scattering atmosphere for $\\bar\\nu_e$ already early after the onset of an explosion, which has not been observed before. We find that at early times after the onset of an explosion, due to the contribution from charge-current reactions to the total opacity for $\\bar\\nu_e$, the energy spheres of $\\bar\\nu_e$ and $\\nu_{\\mu/\\tau}$ are located at different positions. However, due to Pauli blocking of final state neutrons the $\\bar\\nu_e$-opacity from charge current-reactions decreases continuously and the location of the energy spheres for $\\bar\\nu_e$ and $\\nu_{\\mu/\\tau}$ becomes increasingly similar. At late times, the dominating contribution to the total opacity of $\\bar\\nu_e$ and  $\\nu_{\\mu/\\tau}$ comes from neutral-current reactions which are independent of the neutrino flavor. It explains the increasing similarity of the $\\bar\\nu_e$ and $\\nu_{\\mu/\\tau}$ spectra.  On the other hand, the total opacity for $\\nu_e$ is dominated by charge-current neutrino absorption on neutrons at early times and hence no scattering atmospheres is present yet. As the charged current opacity for $\\nu_e$ is larger than for $\\bar\\nu_e$, due to the much larger abundance of neutrons compared to protons, neutral-current processes become dominating only at later times than for $\\bar\\nu_e$. The reduction in the charged current opacity for $\\nu_e$ is related to Pauli blocking of the final state electrons, which continuously increase their Fermi energy as the neutrinospheres move to higher densities. At the same time, the average neutrino energies decrease. At these late times, a scattering atmosphere also develops for $\\nu_e$, which has also not been observed before, and the spectra of $\\nu_e$ and $\\bar\\nu_e$ become increasingly similar.  In addition to the opacities, we have also explored emission processes (inverse opacities) from which we find qualitatively a similar behavior. Furthermore, we find that charge-current and neutral-current processes contribute to the $\\nu_e$ luminosity at clearly separated density domains. Charge-current contributions decrease during the deleptonization and originate from low densities around $10^{10}$--$10^{13}$~g~cm$^{-3}$. On the other hand, neutral-current contributions originate at high densities between $10^{13}$ to several times $10^{14}$~g~cm$^{-3}$ (depending on the progenitor and the state of deleptonization).  A clear extension of this investigation is the inclusion of improved neutrino interactions, in particular corrections from weak magnetism and nucleon recoil~\\cite{Horowitz:2001xf}, the emission of $(\\mu,\\tau)$-neutrino pairs via the annihilation of trapped electron neutrino pairs and the reverse process~\\cite{Buras:2002wt}. These were included in the study of the O-Ne-Mg-core in ref.~\\cite{Huedepohl:2010}, which was in qualitative agreement with our findings. In addition, the treatment of nuclear correlations~\\cite{Schwenk:2009} and weak processes with light nuclei~\\cite{Arcones:2008,Sumiyoshi:2008} should also be improved. Nevertheless, we expect that none of these additions will change the findings discussed in this paper, i.e. on timescales on the order of tens of seconds charge-current reactions are suppressed due to Pauli blocking of final states and the neutrino spectra are dominated by neutral-current processes. These processes do not distinguish between different neutrino flavors and hence the neutrino spectra of all flavors become very similar at these late times. Moreover, we expect that including the improved phase-space treatment of ref.~\\cite{Reddy:1998} in long-term simulations of supernova explosion, will even enhance the suppression of charge-current processes.  We expect that the neutrino-driven ejecta will always be proton rich. This excludes neutrino-driven winds from non-rotating and not magnetic proto-neutron stars as possible site for the production of heavy $r$-process elements ($Z>56$). Furthermore, the role of neutrino-flavor oscillations at late times for neutrino detection and nucleosynthesis becomes negligible in the presence of very similar neutrino spectra of all flavors. However at early times, when the spectra of the different flavors are still different, it may impact nucleosynthesis under proton-rich conditions~\\cite{MartinezPinedo:2011,Wu:2011} and neutrino detection."
    },
    "1112/1112.1076_arXiv.txt": {
        "abstract": "We review the neutrino flux from gamma-ray bursts, which is estimated from gamma-ray observations and used for the interpretation of recent IceCube data, from a particle physics perspective. We numerically calculate the neutrino flux for the same astrophysical assumptions as the analytical fireball neutrino model, including the dominant pion and kaon production modes, flavor mixing, and magnetic field effects on the secondary muons, pions, and kaons. We demonstrate that taking into account the full energy dependencies of all spectra, the normalization of the expected neutrino flux reduces by about one order of magnitude and the spectrum shifts to higher energies, where we can pin down the exact origin of the discrepancies by the re-computation of the analytical models. We also reproduce the IceCube-40 analysis for exactly the same bursts and same assumptions and illustrate the impact of uncertainties. We conclude that the baryonic loading of the fireballs, which is an important control parameter for the emission of cosmic rays, can be constrained significantly with the full-scale experiment after about ten years. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.4186_arXiv.txt": {
        "abstract": "IDENTIKIT was originally developed as a fast approximate scheme for modeling the tidal morphology and kinematics of disk galaxy encounters and mergers.  In this form, it was first used to implement an interactive modeling tool for galaxy collisions; tests with artificial data showed that the morphology \\textit{and} kinematics of merging galaxies strongly constrain their initial conditions.  This tool is now being applied to real galaxies.  More recently, IDENTIKIT has been used to develop a mapping from the present state of a tidal encounter back to the initial conditions; this offers a way to partly automate the search for dynamical models of galaxy encounters.  Finally, IDENTIKIT's theoretical applications include a comprehensive way to evaluate the mass and extent of tidal features as functions of halo structure. ",
        "introduction": "The tidal theory of galaxy encounters explains the morphological and kinematic features of ``peculiar'' galaxies \\citep{A66} as consequences of gravitational interactions between previously normal disk galaxies \\citep[hereafter TT72]{TT72}.  While this theory has strong theoretical foundations, it gained considerable credibility from TT72's plausible simulations of four well-known interacting galaxies. In the decades since TT72, the number of systems with detailed dynamical models has gradually increased.  There are several reasons, beyond testing the tidal theory, to create models of specific systems: such models (1) help interpret complex three-dimensional morphology, (2) provide access to the time domain, and (3) may be used to test sub-grid models of star formation, feedback, AGN fueling, etc.  Thus, a dynamical merger model can provide a framework unifying disparate observations into a coherent picture.  Finding initial conditions which reproduce a specific pair of interacting galaxies is a time-consuming business.  The parameter set is large, and parameters interact in complex ways.  Hours may be required to run a single self-consistent calculation, and weeks of trial and error may be needed to obtain an acceptable match.  Sixteen parameters specify an encounter between two galaxies; these fall into three distinct groups as follows.  (1) The initial orbit is specified by the pericentric separation $p$, mass ratio $\\mu$, and eccentricity $e$. (2) Disk orientations are specified by inclination angles $i$ and azimuthal angles $\\phi$.  (3) The mapping from a simulation to observables is specified by the time since pericenter $t$, three viewing angles $\\theta_\\alpha$, length and velocity scale factors $\\mathcal{L}$ and $\\mathcal{V}$, position zero-point $\\mathbf{r}_{0}$ and velocity zero-point $v_{0}$.  Together, groups (1) and (2) specify the initial conditions.  The criteria used to decide if a model matches the observations are subtle.  As a rule, it's not enough to reproduce the tidal morphology of an interacting pair of galaxies, since different encounter geometries may produce identical morphologies \\citep[e.g.][]{B11}. Kinematic information on tidal bridges and tails provides much stronger constraints.  Interferometric H{\\small I} data, which traces both tidal morphology \\textit{and} kinematics, is often used to constrain merger models; H$\\alpha$ may also be useful if tidal features contain emission-line regions.  Molecular and/or stellar absorption lines can also provide constraints, but mapping extended structures in these lines is expensive.  Models and observations are typically compared visually, for example by overplotting particles on various projections of a H{\\small I} data cube \\citep[e.g.][]{HM95}. Since interstellar material converts between molecular, atomic, and ionized phases by processes outside the scope of purely dynamical models, visual inspection may be more reliable than quantitative measures derived by differencing model and observational data cubes. However, visual inspection is tedious and inherently subjective; robust quantitative methods would certainly be welcome.  Some progress has been made using genetic algorithms to automate the job of modeling interacting galaxies \\citep[e.g.][]{W98, TK01, S+10}. Interesting results have been obtained, but most tests to date have searched only limited subsets of the relevant parameters and made little use of kinematic information. ",
        "conclusions": ""
    },
    "1112/1112.1240_arXiv.txt": {
        "abstract": "{ \\maxi\\  is a transient black-hole candidate discovered in outburst on 30 August 2011. We report on the available \\inte, \\swift, and \\rxte\\  observations performed in the direction of the source during this event before 55\\,864\\,MJD. Combining the broad band (0.6--200\\,keV) spectral and timing information obtained from these data with the results of radio observations, we show that the event displayed by \\maxi\\  is another example of  ``failed'' outburst. During the first $\\sim$20 days after the onset of the event, the source underwent a transition from the canonical low/hard to the hard intermediate state, while reaching the highest X-ray flux. In the $\\sim$40\\,days following the peak of the outburst, the source displayed a progressive spectral hardening and a decrease in the X--ray flux, thus it entered again the low/hard state and began its return to quiescence. } ",
        "introduction": "\\label{sec:intro}  \\maxi\\ was discovered in outburst by MAXI/GSC \\citep{mihara11} on 30 August 2011 (55\\,803~MJD), and detected simultaneously also with the BAT \\citep{barthelmy05} on-board \\swift\\ \\citep{gehrels04}. At discovery, the flux of the source was 25~mCrab and 40~mCrab in the energy ranges 4--10\\,keV and 15--50\\,keV, respectively \\citep{negoro11}. The analysis of previous MAXI/GSC data suggested that the outburst might have already started on 55\\,802\\,MJD. Follow-up observations with \\swift\\,/XRT and \\swift\\,/UVOT on 55\\,803.7\\,MJD provided the best estimated position so far (RA=18h\\,35m\\,43.43s; Dec=-19d\\,19m\\,12.1s, J2000; associated uncertainty 1.8~arcsec at 90\\% c.l.) and led to the identification of the optical counterpart \\citep{kennea11}. An \\rxte\\,/PCA observation was performed on 55\\,804.5\\,MJD \\citep{strohmayer11}, and permitted a classification of the source as a new black-hole candidate (BHC) possibly undergoing a transition from the low-hard (LHS) to the hard intermediate state \\citep[HIMS; see, e.g.][for a recent review]{belloni09}. Relatively strong radio and infrared emissions from the source were detected from 55\\,806\\,MJD to 55\\,827\\,MJD \\citep{miller11,trushkin11} and are likely associated with the presence of a jet.  In this paper we report on all available \\inte,\\ \\swift\\,/XRT, and \\rxte\\,/PCA data collected during the outburst of the source from 55\\,804 to 55\\,864\\,MJD. ",
        "conclusions": "hid} also supports these conclusions: the source was first observed during the transition from the LHS to HIMS (blue diamonds). It remained in the HIMS for a few days (black dots). It then slowly moved back to the LHS (green stars) and then faded out (red triangles). Based on the timing and spectral properties of  the \\rxte\\ , \\swift\\ , and \\inte\\ data, we conclude that \\maxi\\ is one of the few black-hole candidates, together with  H\\,1743$-$322 in 2008 and SAX\\,J1711.6$-$3808 in 2001 \\citep[see][and references therein]{fiamma09}, which showed a transition to the HIMS, but did not enter a soft state."
    },
    "1112/1112.4613.txt": {
        "abstract": " ",
        "introduction": "The idea that Quantum Gravity effects can be important at astrophysical and cosmological distances has recently attracted much attention.  In particular the framework of Exact Renormalization Group (ERG) approach for quantum gravity \\cite{1998PhRvD..57..971R} has opened the possibility of investigating both the ultraviolet (UV) and the infrared (IR) sector of gravity in a systematic manner.  The essential ingredient of this tool is the Effective Average Action $\\Gamma_k[g_{\\mu\\nu}]$, a Wilsonian coarse grained free energy dependent on an infrared momentum scale $k$ which defines an effective field theory appropriate for the scale $k$. By construction, when evaluated at ${\\it tree}$ level, $\\Gamma_k$ correctly describes all gravitational phenomena, {\\it including all loop effects}, if the typical momentum involved are of the order of $k$. When applied to the Einstein-Hilbert action the ERG yields renormalization group flow equations \\cite{1998CQGra..15.3449D} which have made possible detailed investigations of the scaling behavior of Newton\u00d5s constant at high energies \\cite{1999PThPh.102..181S,2002PhRvD..65b5013L,2002PhRvD..66b5026L,2002CQGra..19..483L, 2002PhRvD..65f5016R, 2002PhRvD..66l5001R, 2004PhRvL..92t1301L, 2005JHEP...02..035B, 2006PhRvL..97v1301C, 2009GReGr..41..983R, 2009PhRvD..79j5005R, 2009arXiv0904.2510M, 2009arXiv0905.4220M}. The  scenario emerging from these studies, first demonstrated by Weinberg \\cite{wein2} in $d=2+\\epsilon$ dimensions, suggests that the theory could be consistently defined  in $d=4$ at a nontrivial UV fixed point where the dimensionless Newton constant, $g(k)=G(k) k^2$, does not vanish in the $k\\rightarrow\\infty$ limit, i.e. $g(k\\rightarrow\\infty)=g^\\ast$. As a consequence the dimensionful Newton  constant $G(k)$ is {\\em antiscreened} at high energies, very much as one would expect based on the intuitive picture that the larger is the cloud of virtual particles, the greater is the effective mass seen by a distant observer \\cite{1993hep.th....4146P}.  Recent works have included matter fields \\cite{2003PhRvD..67h1503P} and have also  considered a growing number of purely gravitational operators in the action. In particular, truncations involving quadratic terms in the curvature have been considered in \\cite{2006PhRvL..97v1301C,niedermaier_asymptotic_2006,2009AIPC.1196...44B,2010NuPhB.824..168B}, while higher powers of the Ricci scalar have been studied in \\cite{2009AnPhy.324..414C,2008IJMPA..23..143C}. In all the investigations the UV critical surface has turned out to be finite dimensional ($d_{\\rm UV}=3$), implying that the theory is nonperturbatively renormalizable. At the UV fixed point the theory has a behaviour very similar to QCD, being weakly coupled at high energies  but the running of the dimensionful Newton constant $G(k)=g(k)/k^2$  in  the deep ultraviolet region is  a power-law, at variance with the logarithmic scaling of QCD.  A weakly coupled gravity at high energies is expected to generate important consequences in several astrophysical and cosmological contexts \\cite{2009arXiv0911.2727B}, and in fact the RG flow of the effective average action, obtained by different truncations of theory space, has been the basis of various investigations of ``RG improved\" black hole spacetimes, \\cite{1999PhRvD..60h4011B,2000PhRvD..62d3008B,2006PhRvD..73h3005B} and Early Universe models \\cite{2002PhRvD..65d3508B,2003GReGr..35.1899B,2004CQGra..21.5005B,2005IJMPA..20.2358B,2006CQGra..23.3103B,2006rdgp.conf..461B}.  However, the behavior of the theory is more complicated at low energies, corresponding at cosmological scales, because the $\\boldsymbol{\\beta}$--functions of any local operator of the type $\\sqrt{g}R$, $\\sqrt{g}R^2$, \\ldots, $\\sqrt{g}R^n$ are singular in the IR due to the presence of a pole at $\\Lambda(k)/k^2=1/2$ being $\\Lambda$ the cosmological constant. The presence of this pole is signaling that the Einstein-Hilbert truncation is no longer a consistent approximation to the full flow equation, and most probably a new set of IR-relevant operators is emerging in the $k\\rightarrow 0$ limit. This singular behavior  in fact appears in the ERG for  nearly all cutoff threshold functions in the Einstein-Hilbert truncation, and it is caused by the presence of negative eigenvalues in the spectrum of the fluctuations determinant of the gravitational degrees of freedom. As discussed in \\cite{2004JCAP...12..001R} the dynamical origin of these strong IR effect is due to the ``instability driven renormalization'', a phenomenon well known from many other physical systems \\cite{1999PhLB..445..351A,2000PhRvD..62l5021L,2004NuPhB.693...36B}. We shall see that the low energy domain of the theory is regulated by an IR fixed point which drives the cosmological constant to zero at very late times $\\Lambda(t \\rightarrow \\infty)=0$. This new-type of ``decaying $\\Lambda$ cosmologies'' is therefore quite different from previous models where the time-dependent cosmological constant encodes the effect of the matter creation process \\cite{1996PhRvD..54.2571L}.  Astrophysical consequences of the possible presence of an IR fixed point for quantum gravity  appeared in \\cite{2002PhLB..527....9B,2004IJMPD..13..107B,2004JCAP...01..001B,2006CQGra..23.3103B,2004ForPh..52..650R} where it was shown that a solution of the ``cosmic coincidence problem'' arises naturally without the introduction of a  ``dark energy field''. In particular in the fixed point regime the vacuum energy density $\\rho_\\Lambda\\equiv\\Lambda/8\\pi G$ is automatically adjusted so as to equal the matter energy density, {\\it i.e.} $\\OL=\\OM=1/2$, and the deceleration parameter approaches $q = -1/4$. Moreover, an analysis of the high-redshift SNe Ia data leads to the conclusion that this {\\em infrared fixed point cosmology} is in good  agreement with the observations \\cite{2004JCAP...01..001B}. Recent works have also considered the possibility that the ``basin of attraction'' of the IR fixed point can act already at galactic scale, thus providing an explanation for the galaxy rotation curve without dark matter \\cite{2004JCAP...12..001R,2004PhRvD..69j4022R,2004PhRvD..70l4028R,2007CQGra..24.6255E}, although a detailed analysis based on available experimental data is still missing. Cosmologies discussing the complete evolution from the early Universe to the present time have been considered in \\cite{2005JCAP...09..012R} and in \\cite{2007JCAP...08..024B}. In particular in \\cite{2007JCAP...08..024B} it has been shown that the ``RG improved'' Einstein equations admit (power-law or exponential) inflationary solutions and that the running of the cosmological constant can account for the entire entropy of the present universe in the massless sector (see also \\cite{2008JPhCS.140a2008B} for a review).  The main purpose of this paper is to explore the idea that quantum effects could dynamically drive the cosmological constant to zero at late times so that $\\Lambda(t\\rightarrow\\infty)=0$ as a result of an explicit dynamical mapping $k\\rightarrow k(t)$ of the renormalization group trajectories generated by the unstable infrared modes of the gravitational sector. %In particular, at variance with previous studies %\\cite{2004JCAP...01..001B}, the infrared fixed point will arise as a result %of the of instability induced  {This idea will then be studied using the so-called Dynamical Systems Approach (DSA), a technique  already used in cosmology by Bogoyavlensky \\cite{oleg} and further developed by Collins, and  Ellis and Wainwright to analyze non trivial cosmologies (i.e. Bianchi models) in the context of pure General Relativity \\cite{ellisbook}. Some work has also been done in the case of (minimally coupled) scalar fields in cosmology \\cite{Copeland:2009be,Copeland:1997et} and, more recently, of scalar tensor theories of gravity \\cite{Carloni:2007eu},  $f(R)$ theories of gravity \\cite{Carloni:2004kp,Leach:2006br,Abdelwahab:2007jp,Carloni:2007eu,Carloni:2007br} and  Ho\\v{r}ava-Lifschits gravity \\cite{Carloni:2009jc}. Studying cosmologies using the DSA has the advantage of offering a relatively simple method to obtain particular exact solutions and to obtain a (qualitative) description of the global dynamics of these models.}  %\\textcolor{red}{In our specific context the DSA allows us to prove that}, in {In our specific context the DSA allows us to prove that} the presence of  a singular behavior of the RG evolution for the cosmological constant in the infrared put strong constraints on the possible RG trajectories. In particular we will show that it is possible to realize a scenario where the Universe has a transition from an early unstable phase of power-law evolution of the type $a\\propto t^\\beta$ with $\\beta \\in ]0,1[$ to a de-Sitter phase.  The structure of the paper is the following. In Section 2 the basic mechanism of the dynamical suppression of the cosmological constant  by the unstable low energy modes is discussed. In Section 3 the dynamical system analysis of the resulting cosmologies is presented. Section 4 is devoted to the conclusions. ",
        "conclusions": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% In this paper we have investigated the effect on cosmological scale of a quantum gravity related decaying of the cosmological constant. This investigation was performed  using the Dynamical System Approach that allows a semi-quantitative treatment of the model via the construction of a phase space directly connected to the behavior of the equations. The phase space is two dimensional and it contains a total of nine fixed point, of which three are finite and six are asymptotical. The stability of these points is found to depend on the parameters $A$ and $\\alpha$, but not on $B$ which has a marginal position in the entire analysis. The determination of the fixed points and their stability allows us to infer the shape of the phase space orbits and to deduce the qualitative features of the cosmic histories possible in this models.  Among these, one is particularly relevant in terms of the problem of  dark energy domination i.e. the presence of  a cosmic history characterized by a first phase in which the scale factor grows as a power law with exponent included in $]0,1[$ followed by a second phase with a faster growth. In terms of the  features of the phase space that  such scenario can be realized if one has an unstable fixed point representing the first phase and an attractor representing the second\\footnote{The second point could  be unstable as well, but then the set of the initial condition able to realize such scenario would be smaller and more difficult to calculate.}. Looking at  Tables \\ref{tbl:stabilit punti fissi1} - \\ref{tavola stasbilita punti fissi asintotici EF 2} one can easily see that there is only is one set of the parameters $(A, \\alpha)$ for which this can be realized: $(A<\\frac{1+3w}{1+w}, \\; \\alpha<2)$, for which $\\mathcal{A}$ is a repeller. It is reassuring to notice that in this case  $A$ is of the order  unity  and, as $A=\\xi^2\\lambda_\\ast$, we deduce that also $\\xi$ is  of the order unity being $\\lambda_\\ast < 1/2$.  When these conditions are satisfied there can exist, depending on initial conditions, one orbit that starts at $\\mathcal{A}$ and  either go directly to $\\mathcal{C}$ or bounce  at $\\mathcal{E}_{\\infty}$. The first case can be seen as a ``classical'' Friedmann-de Sitter transition. In the second case the cosmological evolution could be richer because  the solution  $\\mathcal{E}_{\\infty}$ can indicate, depending on the sign of $\\alpha$, a growth whose rate saturates or an expansion phase comparable with a de Sitter one. One could interpret this as a ''two regimes`` dark energy phase.  Note that, as we have seen, in the fixed points $\\mathcal{A}$ and $\\mathcal{B}$ the parameter $B$ is zero these points represent states for the cosmology that are indistinguishable from standard general relativity. This means that our model in the neighborhood of these points is indistinguishable from the standard cosmology. An example of the evolution of some key cosmological quantities compared with the standard de Sitter cosmology is given in Figure \\ref{PlotCosm1}, and in Figure \\ref{PlotCosm3} we compare the additional RG term in (\\ref{useful fried}) with the standard cosmological constant. %{\\it It is interesting to notice that the Newton constant increases with time: this behavior does not poses any %{\\bf It is important to stress that the running of $G$ is only significant on %cosmological distance scales, let's say after the homogeneity scale, %because  the dynamical renomalization process induced by the instability takes over %only near $k\\rightarrow 0$  and $k~H$ because (\\ref{cui}). % %On small (solar system for instance)  scales our $G(k)$ %coincides with the Newton constant determined %by standard laboratory experiments (see also the discussion in ).}  {It is important to emphasize that the standard experimental value of Newton's constant, $G_{\\rm exp}$, does not coincide with the value $G(k=\\xi H_0)$ which is relevant for cosmology today. $G_{\\rm exp}$ is measured (today) at $k_{\\rm exp}\\propto \\ell^{-1}$ where the length $\\ell \\equiv \\ell_{\\rm sol}$ is a typical solar system length scale, say $10^{12}$ m. Thus, in terms of the running Newton constant, $G_{\\rm exp} = G(k=\\xi'/\\ell_{\\rm sol})$, since $\\ell_{\\rm sol}\\ll H_0$. %, and since in presence of several scales the relevant %cutoff is always the larger one %\\footnote{See ref.[14] for a detailed discussion of this point.}. It is only the cosmological quantity $G(k=\\xi H)$ which dynamically evolves in the fixed point regime, %which grows $\\propto t^2$ in the fixed not $G_{\\rm exp}$. This remark entails that the dynamical evolution of the cosmological Newton constant in the recent past does not ruin the predictions about primordial nucleosynthesis which requires that $G(k=\\xi H_{\\rm nucl})$ coincides with $G_{\\rm exp}$ rather precisely. In fact, at the time $t=t_{\\rm nucl}$ of nucleosynthesis for which $H=H(t=t_{\\rm nucl})$, the cosmological Newton constant was indeed $G(k=\\xi H_{\\rm nucl})\\approx G_{\\rm exp}$ since $ct_{\\rm nucl} \\approx H_{\\rm nucl}^{-1}$ and $\\ell_{\\rm sol}$ are of the same order of magnitude.}   The phase space structure also sheds light on aspects of the type of RG cosmologies not directly related to the problem of dark energy. For example, it is clear that it is possible to interpret the transition to a dark era as the beginning of a primordial inflationary phase. {Our results then imply that there is only one set of values that allows a ``graceful exit'' i.e. the transition from inflation to a Friedmann cosmology. Specifically in the case $(A<\\frac{1+3w}{1+w}, \\; \\alpha>2)$ we can have this kind of scenario. However, since there is only one point which represents accelerating expansion one can see that the RG model can be used to represent {\\em either} the inflationary era {\\em or} the dark energy era, but not both.}  Finally the analysis of the asymptotic fixed points gives information on the possibility of changes in the sign of expansion rates in this type of cosmologies. In general relativity such phases are normally associated  either to the so-called \"bounces\" and are associated to cyclic Universes (when $\\dot{H}>0$) or to recollapsing universes (when $\\dot{H}<0$). In modified theories of gravity changes in the sign of $H$ can occur also when the size of the Universe is large because the scale factor does not need to be monotonic \\cite{Carloni:2005ii}. In terms of the phase space  this kind of behavior is associated to the presence of asymptotic attractors and the sign of the quantity $-(1+q)$  as expressed in (\\ref{deceleration par}) in the fixed point. In particular we have that if  $-(1+q)>0$ we have a deceleration followed by an acceleration and if  $-(1+q)<0$ the opposite situation. A quick analysis shows, for example, that only some of the asymptotic attractors for very specific values of the parameters can give origin to a bounce.  The previous results point towards a cosmology with interesting features that, in our opinion, deserves some more study. In particular one could try to test the transition to Dark Age we have found against the SnIa data to obtain some more constraints on the free parameters. This, and other issues, will be discussed in a following work.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %\\begin{figure}[ht] %\\centering %\\includegraphics[scale=0.70]{PlotCosm1.eps} % %\\caption %{Evolution of $ \\ln H$ (solid), $\\ln \\rho$ (dashed), $\\ln G$ (dotted)  as a function of $\\ln a$ in RG cosmologies with $w=0$ $A=1/2$, $\\alpha=1$.}%\\label{PlotCosm1} %\\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %\\begin{figure}[ht] %\\centering %\\includegraphics[scale=0.60]{PlotCosm2.eps} %\\caption %{Behaviour of $ \\ln H$ as a function of $\\ln a$. It is clear that in the RG case (dashed) the cosmological constant assumes a higher value with respect to the standard relativity case (solid).}\\label{PlotCosm2} %\\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %\\begin{figure}[ht] %\\centering %\\includegraphics[scale=0.60]{PlotCosm3.eps} %\\caption %{Behaviour of $ \\ln \\rho$ as a function of $\\ln a$. It is clear that in the RG case (dashed) the energy density evolves at the same rate as in standard relativity (solid). This implies that different orbits of this type correspond to different behaviors of the gravitational constant.}\\label{PlotCosm3} %\\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %\\begin{figure}[ht] %\\centering %\\includegraphics[scale=0.60]{PlotCosm4.eps} %\\caption %{Behaviour of $ \\ln G$ as a function of $\\ln a$. It is clear that in the RG case (dashed) the gravitational cosntant incerases in time.}\\label{PlotCosm4} %\\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %\\begin{figure}[ht] %\\centering % %\\subfigure[Behaviour of $ \\ln H$ as a function of $\\ln a$. It is clear that in the RG case (dashed) the cosmological constant assumes a higher value with respect to the standard relativity case (solid).]{ %\\includegraphics[scale=0.32]{PlotCosm2.eps} %\\label{fig:subfig2} %} % %\\subfigure[Behaviour of $ \\ln \\rho$ as a function of $\\ln a$. It is clear that in the RG case (dashed) the energy density evolves at the same rate as in standard relativity (solid). This implies that different orbits of this type correspond to different behaviors of the gravitational constant.]{ %\\includegraphics[scale=0.32]{PlotCosm3.eps} %\\label{fig:subfig3} %} %\\subfigure[Behaviour of $ \\ln G$ as a function of $\\ln a$. It is clear that in the RG case (dashed) the gravitational cosntant incerases in time.]{ %\\includegraphics[scale=0.32]{PlotCosm4.eps} %\\label{fig:subfig4} %} %\\caption[] %{Comparison of $ \\ln H$, $\\ln \\rho$, $\\ln G$ as a function of $\\ln a$ in RG cosmologies with $w=0$ $A=1/2$, $\\alpha=1$  and the standard Relativity with the cosmological constant $\\Lambda$ and  gravitational constant $G$ normalized to 1.}\\label{PlotCosm2} %\\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1112/1112.3480_arXiv.txt": {
        "abstract": "We present a study of the formation of clustered, massive galaxies at large look--back times, via high resolution spectroscopic imaging of CO in the unique GN20 proto-cluster. The data reveal evidence for rich structure and gas dynamics in unprecedented detail, allowing us to image the molecular gas with a resolution of only 1 kpc just 1.5 Gyr after the Big Bang. These state-of-the-art data give new insight into the detailed physical processes involved in early massive galaxy formation, and they provide a first glimpse of the morphological studies that will become feasible on a regular basis with ALMA. ",
        "introduction": "\\label{Intro}   The majority of submillimeter galaxies \\citep[SMGs;][]{2002PhR...369..111B} are believed to be starburst--dominated major mergers \\citep[e.g.,][]{2003ApJ...599...92C}. Indeed, there is direct evidence for multiple CO components and/or disturbed kinematics in some SMGs, supporting the merger picture \\citep[e.g.,][]{2008ApJ...680..246T, 2010ApJ...724..233E, 2011MNRAS.412.1913I, 2011ApJ...739L..31R}. However, it has recently been suggested that other mechanisms that drive extreme star formation rates may also be at play. In particular, the cold mode accretion phenomenon \\citep[CMA; e.g.][]{2009ApJ...703..785D, 2009Natur.457..451D} has been extended to SMGs by a number of authors \\citep{2001astro.ph..7290F, 2006ApJ...639..672F, 2010MNRAS.404.1355D}. In this scenario, SMGs are massive galaxies sitting at the centers of deep potential wells and fed by smooth accretion.  GN20 is an interesting SMG in this regard. Situated at z$=$4.05, CO observations hint at a large spatial extent and ordered rotation \\citep{2010ApJ...714.1407C}. However, the previous VLA observations had some severe spectral limitations; in particular, the limited bandwidth truncated the line profiles, and the reliance on continuum mode precluded any information on kinematics. We therefore obtained 124 hours on the Expanded Very Large Array (EVLA) to image the CO(2--1) emission in the GN20 field in the B-- and D--configurations. We centered the two 128 MHz IFs at 45.592 GHz and 45.720 GHz, for a total bandwidth of 246 MHz or 1600 km s$^{-1}$.  Each of the two IFs had 64 channels, corresponding to $\\sim$13 km s$^{-1}$ per channel. The $\\sim$20 hours of lower--resolution D--array data were presented in the EVLA Special Issue of ApJ by \\citet{2011arXiv1105.2451C}. Here, we present the full dataset (B+D--configurations) on GN20, providing greatly improved spatial resolution and image fidelity. ",
        "conclusions": "\\label{discussion}   The high--quality, high--resolution data on GN20 presented in this work reveal an extended gas reservoir, 14 $\\pm$ 4 kpc in diameter, which is consistent with a rotating disk. The data further resolve the reservoir into multiple, kpc--sized gas clumps which each constitute a few percent of the total gas mass and together comprise only 30\\% of the total gas mass. These observations may favor a scenario where the star formation is fueled by a process other than a major merger, e.g., cold mode accretion. To determine if this is in fact the case, a crucial question must be addressed: is it possible to have a highly--obscured starburst without a major merger? This is a question that, due to the faintness of this particular source, may have to await the launch of JWST."
    },
    "1112/1112.4269_arXiv.txt": {
        "abstract": "\\altaffiltext{6}{Physics Department, Xiangtan University, 411105 Xiangtan, Funan Province, P. R. China}  We report here the tentative discovery of a Jovian planet in orbit around the rapidly pulsating subdwarf B-type (sdB-type) eclipsing binary NY Vir. By using new determined eclipse times together with those collected from the literature, we detect that the observed-calculated (O-C) curve of NY Vir shows a small-amplitude cyclic variation with a period of 7.9\\,years and a semiamplitude of 6.1\\,s, while it undergoes a downward parabolic change (revealing a period decrease at a rate of $\\dot{P}=-9.2\\times{10^{-12}}$). The periodic variation was analyzed for the light-travel time effect via the presence of a third body. The mass of the tertiary companion was determined to be $M_3\\sin{i^{\\prime}}=2.3(\\pm0.3)$\\,$M_{Jupiter}$ when a total mass of $0.60$\\,$M_{\\odot}$ for NY Vir is adopted. This suggests that it is most probably a giant circumbinary planet orbiting NY Vir at a distance of about 3.3 astronomical units (AU). Since the rate of period decrease can not be explained by true angular momentum loss caused by gravitational radiation or/and magnetic braking, the observed downward parabolic change in the O-C diagram may be only a part of a long-period (longer than 15 years) cyclic variation, which may reveal the presence of another Jovian planet ($\\sim2.5$$M_{Jupiter}$) in the system. ",
        "introduction": "HW Vir-like eclipsing binaries are a group of detached binary systems that consists a very hot subdwarf B-type (sdB) primary and a fully convective M-type secondary with periods shorter than 4 hours (e.g., Menzies \\& Marang 1986; Kilkenny et al. 1998; Drechsel et al. 2001; Ostensen et al. 2007; Polubek et al. 2007; Wils et al. 2007; For et al. 2010; Ostensen et al. 2010). The hot sdB components in this group of systems are on the extreme horizontal branch (EHB) of the Hertzsprung-Russell diagram burning helium in their cores and and having very thin hydrogen envelopes (e.g., Heber 2009). Theoretical investigations have shown that they are formed through a common-envelope evolution (e.g., Han et al. 2003), and will evolve into normal cataclysmic variables (CV) (e.g., Shimansky et al. 2006). The discovery of circumbinary substellar objects orbiting these peculiar binaries has very important implications of several outstanding problems in astrophysics, e.g., the formation of sdB stars and the fates of low-mass companion systems. To date, three substellar companions were found in HW Vir and HS0705+6700 systems as circumbinary brown dwarfs and planets (Lee et al. 2009; Qian et al. 2009).  As for substellar objects in orbits around single sdB-type stars, Silvotti et al. (2007) discovered a giant planet with a mass of $3.2\\,M_{Jupiter}$ and an orbital separation of 1.7\\,AU, around the hot sdB star V391 Peg. Geier et al. (2009) later announced the detection of a close substellar companion to the hot sdB star HD\\,149382 with a short period of 2.391\\,days. However, Jacobs et al. (2011) found no evidence for the presence of the claimed substellar object by analyzing He lines. Recently, Geier et al. (2011) detected a brown dwarf companion to the hot subdwarf SDSS J083053.53+000843.4 with a very short period of 0.096\\,days and a mass of 0.045-0.067\\,$M_{\\odot}$. Here we report the tentative discovery of a Jovian planet with a mass of $\\sim2.3\\,M_{Jupiter}$ around the HW Vir-like binary star NY Vir. The planet has the lowest mass among the substellar objects discovered before orbiting sdB stars. Moreover, evidence may show that there is another circumbinary planet companion in the system. ",
        "conclusions": "Investigations of several authors (e.g., Qian et al. 2008; Brinkworth et al. 2005) showed that the magnetic activity cycles of a fully convective component star (i. e., the Applegate mechanism) (Applegate 1992) can not explain the observed cyclic variations in the O-C diagrams of close binaries because of the problem of energy. Moreover, a recently statistical study by Liao and Qian (2010) suggests that light travel-time effect is the most probable mechanism to cause cyclic changes in the O-C diagrams of close binary stars. Therefore, we analyzed NY Vir for the light travel-time effect that arises from the gravitational influence of a tertiary companion. The possible presence of a companion object produces the relative distance changes of the eclipsing pair as it orbits the barycenter of the triple system.  By considering a typical mass of $M_{1}=0.46M_{\\odot}$ for sdB-type stars and using the mass ratio of $q=0.3$ determined by Kilkenny et al. (1998), the mass of the secondary is estimated as $M_{2}=0.14M_{\\odot}$. Then, the mass function and the mass of the possible tertiary companion were derived by using the following equation, \\begin{equation} f(m)=4\\pi^{2}(a_{12}\\sin{i^{\\prime}})^{3}/GT^{2}=(M_{3}\\sin{i^{\\prime}})^{3}/(M_{1}+M_{2}+M_{3})^{2}, \\end{equation} where $M_{3}$ is the mass of the third body, G is the gravitational constant, and T is the period of the $(O-C)_2$ oscillation. $a_{12}\\sin{i^{\\prime}}$ can be determined by, \\begin{equation} a_{12}\\sin{i^{\\prime}}=A\\times{c} \\end{equation} where A is the semi-amplitude of the $(O-C)_2$ oscillation and $c$ the speed of light. The results are: $f(m)=3.0(\\pm0.4)\\times{10^{-8}}\\,M_{\\odot}$ and $M_3\\sin{i^{\\prime}}=2.3(\\pm0.3)$\\,$M_{Jupiter}$. When the orbital inclination of the third body is larger than $9.6^{\\circ}$, the mass of the tertiary component corresponds to $2.3\\,M_{Jupiter}\\le{M_3}\\le14\\,M_{Jupiter}$. In this situation, as a few good cases have been made for circumbinary planets (e.g., Lee et al. 2009; Beuermann et al., 2010; Qian et al. 2010, 2011), it should be a planetary object. The orbital separation $d_3$ of the Jupiter-like planet candidate is about $3.3(\\pm0.8)$\\,AU.  The downward parabolic change in the O-C diagram (the dashed line in Fig. 2) implies a period decrease at a rate of $\\dot{P}=-9.2\\times{10^{-12}}$ (or 1s in about 3,447 years). This corresponds to $\\dot{P}=-3.36\\times{10^{-9}}$\\,days/year, which is close to the value obtained in literature (Kilkenny et al. 2011). The orbital separation of the two component stars in NY Vir is estimated as $0.77\\,R_{\\odot}$. By using the equation given by Kraft et al. (1961) and Faulkner (1971), \\begin{equation} \\frac{\\dot{P}_{GR}}{P}=-3\\frac{32G^3}{5c^5}\\frac{M_1M_2(M_1+M_2)}{d^4}, \\end{equation} where $P$ is the orbital period and $d$ the distance between both components, the contribution of gravitational radiation (GR) to the period decrease can be computed as $\\dot{P}_{GR}=-0.026\\times{10^{-9}}$\\,days/year. This is about two orders smaller than the observed value indicating that it can not be caused by GR. This decrease can be explained by magnetic braking (MB) as in the case of HW Vir (Qian et al. 2008; Lee et al. 2009). However, it is more widely accepted that MB is stopped for fully convective stars (Rappaport et al. 1983; Spruit \\& Ritter 1983).  Therefore, as in the case of HU Aqr, the more plausible reason for the observed long-term period decrease is that it is only a part of a long-period cyclic variation (longer than 15 years), revealing the possible presence of another planet in the planetary system. This is the same as that of HU Aqr where a planetary system is orbiting the eclipsing polar (Qian et al. 2011). We estimate that the mass of the second planet is about $2.5$\\,$M_{Jupiter}$. Some investigators (e.g., Horner et al. 2011) think that the serious problem for the presence of the fourth body is the dynamical stability of the circumbinary planetary system. However, the 3D dynamic-aware analysis of the stability of circumbinary orbits by Doolin et al. (2011) reveals that these orbits are surprisingly stable throughout binary mass fraction -- eccentricity parameter space.  If the cyclic variations in the O-C diagram are caused by the light-travel time effects via the presence of planetary objects, they should be strictly periodic. To check the existence of the planetary system, more eclipse times are needed in the future. Moreover, At the maximum point of the O-C curve, the central eclipsing binary NY Vir is at the farthest position of the orbit, while the circumbinary planet is closest to the observer. If the tertiary companion is coplanar to the central binary, the circumbinary planet should be in the light of the eclipsing binary and transit the binary component stars. Therefore, searching for the transits of the binary components by the circumbinary planets at the O-C maxima can ascertain the planetary system. More recently, a Saturn-like planet transits an M-type eclipsing binary was reported by Doyle et al. (2011), which provides the first direct evidence for the presence of a circumbinary planet.  As in the cases of other HW Vir-like binary stars, NY Vir was evolved through a common envelope (CE) after the more massive component star in the original system evolves into a red giant. The ejection of CE removed a large amount of the angular momentum, and the present EHB star with a low-mass stellar companion in a short-period binary was formed. The distance between the central EHB binary and the circumbinary planet is about 3.3\\,AU. By assuming the mass of the main-sequence (MS) progenitor of the EHB star is $M_{MS}=1.0$\\,$M_{\\odot}$, the orbital separation of the planet around the MS progenitor is estimated as (e.g., Bear \\& Soker 2011): $a_0\\cong{(M_{EH}+M_2)/(M_{MS}+M_2)}\\approx1.74$\\,AU (by ignoring tidal interaction). When the primary star evolved to a red giant branch star before the system entered a CE, it has a reached radius of $\\sim{100}$\\,$R_{\\odot}\\sim{0.5}$\\,AU. Therefore, the secondary should be at a distance of a little large, say 0.8 AU. Can a system of two stars ($M_1=1.0\\,M_{\\odot}$ and $M_2=0.14\\,M_{\\odot}$) at an orbital separation of 0.8\\,AU have a planet at 1.74\\,AU in a stable orbit? The 3D numerical simulations by Doolin et al. (2011) indicates that the orbit is stable. A system with two M-type stars ($M_1=0.69\\,M_{\\odot}$ and $M_2=0.20\\,M_{\\odot}$) at an orbital separation of 0.22\\,AU was reported to have a Saturn-like planetary companion at a distance of 0.7\\,AU (Doyle et al. 2011). If the detection of a planetary companion to NY Vir at orbital separation of 3.3\\,AU is true, it can give some constraints on the stellar evolution and the interaction between red giants and their planetary companions. Moreover, it is interesting to pointed out that the orbital separations between the central binaries and the inner substellar objects in the three systems, i.e., HW Vir (Lee et al. 2009), HS\\,0705+6700 (Qian et al. 2009), and NY Vir, are all in a small range of 3.3-3.6\\,AU. It is unclear whether these agreements are fortuitous or there are physical reasons."
    },
    "1112/1112.2165_arXiv.txt": {
        "abstract": "We present {\\it Kepler} observations of the bright (V=8.3), oscillating star HD 179070. The observations show transit-like events which reveal that the star is orbited every 2.8 days by a small, 1.6 $R_{\\rm Earth}$ object. Seismic studies of HD 179070 using short cadence {\\it Kepler} observations show that HD 179070 has a frequency-power spectrum consistent with solar-like oscillations that are acoustic p-modes. Asteroseismic analysis provides robust values for the mass and radius of HD 179070, 1.34$\\pm$0.06 M$_{\\odot}$  and 1.86$\\pm$0.04 R$_{\\odot}$ respectively, as well as yielding an age of 2.84$\\pm$0.34 Gyr for this F5 subgiant. Together with ground-based follow-up observations, analysis of the {\\it Kepler} light curves and image data, and blend scenario models, we conservatively show at the $>$99.7\\% confidence level (3$\\sigma$) that the transit event is caused by a 1.64$\\pm$0.04 R$_{\\rm Earth}$ exoplanet in a 2.785755$\\pm$0.000032 day orbit. The exoplanet is only 0.04 AU away from the star and our spectroscopic observations provide an upper limit to its mass of $\\sim$10 $M_{\\rm Earth}$ (2-$\\sigma$). HD 179070 is the brightest exoplanet host star yet discovered by {\\it Kepler}. ",
        "introduction": "The NASA {\\it Kepler} mission, described in  Borucki et al. (2010a), surveys a large area of the sky spanning the boundary of the constellations Cygnus and Lyra. The prime mission goal is to detect transits by exoplanets as their orbits allow them to pass in front of their parent star as viewed from the Earth. Exoplanets, both large and small and orbiting bright and faint stars, have already been detected by the {\\it Kepler} mission (see Borucki et al. 2010b; Koch et al. 2010a; Batalha et al. 2010; Latham et al. 2010; Dunham et al. 2010; Jenkins et al. 2010a; Holman et al. 2010; Torres et al. 2010; Howell et al. 2010; Batalha et al. 2010; Lissauer et al. 2011) and many more exoplanet candidates are known (Borucki et al. 2011).  We report herein on one of the brightest targets in the Kepler field, HD 179070 (KIC 3632418, KOI 975). This V=8.3 F6IV star has been observed prior to the Kepler mission as a matter of course for essentially all of the Henry Draper catalogue stars. Catalogue information\\footnote{http://simbad.u-strasbg.fr/simbad/sim-fid} for HD179070 provides a T$_{\\rm eff}$ of 6137K (spectral type F6 IV), an Hipparcos distance of 108$\\pm$10 pc, [Fe/H]=-0.15, an age of 2.8 GYr, a radial velocity of -28 km/sec and E(b-y)=0.011 mag. The age, metallicity, and color excess reported in the catalogue come from the photometric study by Nordstr\\\"{o}m,  et al. (2004). nothing special to distinguish it from any other stars in the HD catalogue.  HD 179070 is strongly saturated in the {\\it Kepler} observations. Techniques to obtain good photometry from saturated stars are employed by the {\\it Kepler} project and have been used to analyze light curves of other {\\it Kepler} saturated stars (e.g., Welsh et al. 2011) While a star will saturate the CCD detectors on {\\it Kepler} if brighter than V$\\sim$R$\\sim$11.5, collecting all the pixels which contain the starlight, including the bleed trail, allow a photometric study to be performed (see Gilliland et al. 2010a). {\\it Kepler} light curves of HD179070 from the early data as well as subsequent quarters of observation have shown a repeatable transit-like event in the light curve. We will examine the nature of this periodic signal using Quarter 0 to Quarter 5 {\\it Kepler} observations and will show that it is caused by a small exoplanet transiting HD 179070 with a period near 2.8 days. We will refer to the host star as HD179070 and the exoplanet as Kepler-21b throughout this paper.  Our analysis for this exoplanet follows the very complete, tortuous validation path fully described in Batalha et al. (2010). The interested reader is referred to that paper for the detailed procedures which we do not completely repeat herein. Section 2 will discuss the {\\it Kepler} observations including both the light curves and the pixel image data. \\S3 \\& \\S4 discuss ground-based speckle and spectroscopic follow-up observations we have obtained and Section 5 presents the asteroseismic results. Using a full analysis of the observations, we produce a transit model fit for the exoplanet, determine many observational properties for it, and end with a discussion of our results. ",
        "conclusions": "{\\it Kepler} photometry of the bright star HD 179070 reveals a small periodic transit-like signal consistent with a 1.6 R$_{\\rm Earth}$ exoplanet. The transit signal repeats every 2.8 days and the complete phased light curve shows all of the events to be consistent in phase, amplitude, and duration. Analysis of the {\\it Kepler} image data and difference images are well matched by model fits. Detailed point response function (PRF) models conclude that the source of the transit event is centered on or near to the center of HD 179070 itself. Furthermore, these models show that many faint background eclipsing binary scenarios, capable of blending light with that from HD 179070 to produce the transit-like event, can be eliminated.  High resolution ground-based optical speckle imaging reveals no nearby companion star to within 5 magnitudes of HD179070 itself. Near-IR AO observation, however, reveal a faint companion star 0.75\" away and $\\sim$4 magnitudes fainter in $K$. Using a color transformation, this star  is expected to be R$\\sim$14.2, just below the detection limit of the speckle results. Spectroscopic observations also confirm that no bright star is present near HD 179070 (within 0.5\"). Asteroseismology was performed for HD 179070 using the {\\it Kepler} light curves. Adopting the spectroscopically determined values for T$_{\\rm eff}$, log g, and [Fe/H], the mass,  radius, and age of HD 179070 were well determined.  Putting all of the above observations and models together, we conclude that the cause of the periodic transit event is indeed a small 1.6 R$_{\\rm Earth}$ exoplanet orbiting the subgiant star HD 179070. Transit models were fit to the highly precise {\\it Kepler} light curve data revealing that the exoplanet orbits every 2.78 days at an inclination of 82.5 degrees. The exoplanet has an equilibrium temperature near 1900K and is located 0.04 AU from its host star. Kepler-21b has been validated by detailed modeling of blend scenarios as a true exoplanet at greater than 99.7\\% confidence. We can only determine an upper mass limit for the exoplanet, $\\sim$10 M$_{\\rm Earth}$, resulting in a upper limit to the mean density of $\\sim$13 g cm$^{-3}$.  {\\it Kepler} continues to monitor HD 179070 and will eventually build up higher S/N phased transit light curves. These long term observations may allow other planets within this same system may be directly detected or detected via transit timing variations. Given the brightness of HD 179070, it is likely that continued radial velocity monitoring will take place with Keck or other current or planned radial velocity instruments. Given a consistent level of instrumental precision, the observed stellar jitter will slowly be averaged out and velocity signals from this or other planets orbiting the host star may be detected. Finally, using a technique such as described in Schuler et al. (2011) high-resolution, high signal-to-noise echelle spectroscopy will provide detailed metal abundance values of the host star's atmosphere which may hold clues as to the formation, or not, of planetary bodies."
    },
    "1112/1112.0026_arXiv.txt": {
        "abstract": "Observations of stars and pulsars orbiting the black hole in the center of the Milky Way offer the potential of measuring not only the mass of the black hole but also its spin and quadrupole moment, thereby providing observational verification of the no-hair theorem. The relativistic effects that will allow us to measure these higher moments of the gravitational field, however, are very small and may be masked by drag forces that stars and pulsars experience orbiting within the hot, tenuous plasma that surrounds the black hole. The properties of this plasma at large distances from the central object have been measured using observations of the extended X-ray emission that surrounds the point source. At distances comparable to the black-hole event horizon, the properties of the accretion flow have been constrained using observations of its long-wavelength emission and polarization, as well as of the size of the emitting region at 1.3~mm. I use models of the plasma density and temperature at various distances from the black hole to investigate the effect of hydrodynamic drag forces on future measurements of the higher moments of its gravitational field.  I find that hydrodynamic drag does not preclude measurements of the black hole spin and quadrupole moment using high-resolution observations of stars and pulsars that orbit within a few thousand gravitational radii from its horizon. ",
        "introduction": "The black hole in the center of the Milky Way, Sgr~A$^*$ has a number of characteristics that make it an ideal candidate for testing general relativity in the strong field regime. It is surrounded by a swarm of young stars at very close orbits, the tracking of which has already led to a direct measurement of its mass (see Ghez et al.\\ 2008; Gillessen et al.\\ 2009 for recent work). It is accreting from the surrounding medium at rates high enough to make it a detectable source both in the X-rays (e.g., Baganoff et al.\\ 2003) and at long wavelengths (e.g., Falcke et al.\\ 2008; Yusef-Zadeh et al.\\ 2009). It is also the black hole with a horizon that subtends the largest angle in the sky, making it possible to directly image the region very close to its horizon (e.g., Doeleman et al.\\ 2008).  Three distinct technological advances in the next decade offer the possibility of not only improving the accuracy of our knowledge of the black hole's mass but also of measuring directly its spin and quadrupole moment. This last quantity of the gravitational field is not an independent quantity in the Kerr metric; instead it depends in a very particular way on the mass and spin of the black hole, as required by the no-hair theorem (see, e.g., Ryan 1995). Measuring all three moments of the spacetime of Sgr~A$^*$ will, therefore, allow us to test the black-hole nature of the source, verify the validity of the no-hair theorem, and consequently test general relativity in the strong-field regime (see Will 1998; Psaltis \\& Johannsen 2011).  Interferometric observations in large telescopes equipped with adaptive optics (such as GRAVITY on the VLT; Gillessen et al.\\ 2010) will detect and track stars in very close orbits around Sgr~A$^*$. The orbits of these stars will precess because of the relativistic corrections to the black-hole's Newtonian field at rates that depend directly on the compact object's spin and quadrupole moments (Will 2008).  The discovery of a single radio pulsar in orbit around Sgr~A$^*$ will also allow for a direct measurement of the black hole's spin and quadrupole moment, by modeling relativistic effects in the residuals of its orbital solution (Wex \\& Kopeikin 1999). The presence of a cluster of young, massive stars in orbit around Sgr~A$^*$ makes it highly likely that such pulsars exist (e.g., Pfahl \\& Loeb 2004) and current searches place a conservative upper limit of $\\simeq 90$~pulsars within the central parsec of the galaxy (e.g., Macquart et al.\\ 2010).  The difficulty in detecting these pulsars lies in the fact that interstellar scattering heavily disperses and broadens the emission from each pulsar.  Finally, current mm-VLBI observations of Sgr~A$^*$ with only three baselines have already confirmed the expectation that the accretion flow is optically thin at such wavelengths (Doeleman et al.\\ 2008) and has led to the first constraints on the orientation and properties of the black hole (Broderick et al.\\ 2009). Future observations with 5-10 baselines will provide the first direct image of horizon-scale structures in the vicinity of a black hole (Fish \\& Doeleman 2009) and allow for a model-independent measurement of the black-hole's spin and quadrupole moment (Johannsen \\& Psaltis 2010).  Measuring the mass, spin, and quadrupole moment of the black hole in the center of the Milky Way with three independent techniques is crucial in distinguishing relativistic effects from other astrophysical complications. Indeed, the presence of gas and stars in the vicinity of the black hole introduce perturbations to the orbits of stars and pulsars, that may mask or even bias the measurements. The cluster of stars (known and anticipated) in orbit around Sgr~A$^*$ introduces multipole components to the gravitational field, which in turn cause precession of their orbits. In order for this classical effect to be negligible compared to the relativistic precession that depends on the spin and quadrupole of the black-hole spacetime, the orbital separations of the stars from the black hole have to be less than a milliparsec (Merritt et al.\\ 2010).  Stars and pulsars orbiting very close to the black hole, however, will interact with the hot surrounding plasma, which is provided by the stellar winds of the more distant stars and powers the accretion luminosity of the black hole. This interaction will cause them to leave their geodesic orbits and very slowly spiral towards the black hole. If any of the various components of hydrodynamic drag is comparable to or larger than the perturbation to the Newtonian gravitational acceleration due to relativistic effects, then tracking the orbits of stars or pulsars around the black hole will not lead to a clean measurement of the moments of its spacetime.  In this article, I first infer the density and temperature of the hot plasma at different distances from Sgr~A$^*$ using models of the accretion flow and of its feeding region that are consistent with current X-ray and millimeter observations. I then estimate the effect of hydrodynamic drag on the orbits of objects in this plasma and show that it is negligible for stars and pulsars orbiting within a few thousand gravitational radii from the black-hole horizon. Through this work, I will assume that the black hole has a mass of $4.5\\times 10^6 M_\\odot$ and lies at a distance of $8.4$~kpc, as inferred from the constrained fit of Ghez et al.\\ (2008; see also Yelda et al.\\ 2010). ",
        "conclusions": "In this paper, I used the most recent models of the X-ray and millimeter observations of the black hole in the center of the Milky Way in order to estimate the density and temperature of the plasma in a wide range of distances from the central black hole. I then calculated the effect of hydrodynamic drag on the orbits of star and pulsars that can be used to probe relativistic effects and test the no-hair theorem.  I found that in both the case of orbiting stars and of pulsars, the hydrodynamic drag dominates over the gravitational interaction of the object with its wake. For the case of 10~$M_\\odot$ stars or 2~$M_\\odot$ pulsars around Sgr~A$^*$, the hydrodynamic drag is negligible compared to the quadrupole terms of the gravitational acceleration, as long as their orbits are within $\\simeq 10^5$ and $\\simeq 10^4$ gravitational radii, respectively, from the central black hole. Albeit tight, such orbits are also required for the perturbations in the gravitational field due to the presence of the stellar cluster to be negligible compared to relativistic effects (Merritt et al.\\ 2010). In other words, for those stars and pulsars for which the interactions with the other objects of the stellar cluster are negligible, hydrodynamic drag will not preclude measuring the spin and quadrupole moment of the black hole spacetime."
    },
    "1112/1112.3928_arXiv.txt": {
        "abstract": "Gravitational wave observations will probe non-linear gravitational interactions and thus enable strong tests of Einstein's theory of general relativity. We present a numerical relativity study of the late inspiral and merger of binary black holes in scalar-tensor theories of gravity. We consider black hole binaries in an inhomogeneous scalar field, specifically binaries inside a scalar field bubble, in some cases with a potential. We calculate the emission of dipole radiation. We also show how these configurations trigger detectable differences between gravitational waves in scalar-tensor gravity and the corresponding waves in general relativity. We conclude that, barring an external mechanism to induce dynamics in the scalar field, scalar-tensor gravity binary black holes alone are not capable of awaking a dormant scalar field, and are thus observationally indistinguishable from their general relativistic counterparts. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.4480_arXiv.txt": {
        "abstract": "\\noindent We study the properties of the population of optically dark events present in a carefully selected complete sample of bright {\\it Swift} long gamma-ray bursts. The high level of completeness in redshift of our sample (52 objects out of 58) allow us to establish the existence of a genuine dark population and we are able to estimate the maximum fraction of dark burst events ($\\sim 30 \\%$) expected for the whole class of long gamma-ray burst. The redshift distribution of this population of dark bursts is similar to the one of the whole sample. Interestingly, the rest-frame X-ray luminosity (and the de-absorbed X-ray flux) of the sub-class of dark bursts is slightly higher than the average luminosity of the non-dark events. At the same time the prompt properties do not differ and the optical flux of dark events is at the lower tail of the optical flux distribution, corrected for Galactic absorption. All these properties suggest that dark bursts events generate in much denser environments with respect to normal bright events. We can therefore exclude the high-{\\it z} and the low-density scenarios and conclude that the major cause of the origin of optically dark events is the dust extinction. ",
        "introduction": "Gamma Ray Bursts (GRBs) are brief and intense flashes of high energy gamma rays, originating at cosmological distances and often associated with radiation emitted at longer wavelengths for longer periods, identified as the afterglow. The afterglow is almost always detected in the X-ray band (for $\\sim 95 \\%$ of the events detected by {\\it Swift}-XRT; Evans et al. 2009) while the optical afterglow is not. The GRBs that have no optical afterglow or a very low optical-to-X-flux ratio are classified as ``dark burst\". Jakobsson et al. (2004) proposed that a GRB~should be classified as ``dark\" if the slope of the spectral energy distribution between the optical and the X-ray band ($\\beta_{\\rm OX}$) is $< 0.5$. This working definition is a direct implication of the simplest fireball model. In fact, the spectral index $\\beta$ ($F_{\\nu} \\propto \\nu^{-\\beta})$ is related to the power-law index of the electron energy distribution ($p$) and the location of the cooling frequency ($\\nu_{\\rm c}$), independently of the nature of circumburst environment (homogeneous or wind-like) and the collimation of the outflow (see Sari et al. 1998). The average value of $\\beta_{\\rm OX}$ is then expected to be between 0.5 (if $p$=2 and  $\\nu_{\\rm c}$ lies above the X-ray band) and 1.25 (if $p$=2.5 and  $\\nu_{\\rm c}$ lies below the optical frequency). Therefore any optically sub-luminous burst should be located below a constant line $\\beta_{\\rm OX} = 0.5$ in an optical vs. X-ray flux ($f_{\\rm O}-f_{\\rm X}$) diagram, providing that the fluxes are estimated at a common time, chosen to be $t_{\\rm obs} = 11$~hr post-burst (Jakobsson et al. 2004). A slightly more elaborated method was presented by Rol et al. (2005) that compared the optical and X-rays fluxes at a given time extrapolating the latter to the optical band using not only the spectral index but also the temporal power-law index, in the context of the standard fireball model. Both Jakobsson et al. (2004) and Rol et al. (2005) found similar results ($\\sim 10-20\\%$ of possible dark bursts) on samples of pre-{\\it Swift} GRBs and these methods are still used as immediate diagnostic tools to discriminate between optically bright and dark bursts.  Recently, van der Horst et al. (2009) proposed a new method for the optical classification of dark GRBs. Their method is less affected by assumptions about the emitting region with respect to the one done in Jakobsson et al. (2004). They improved the previous method by defining the region of optically (sub-luminous) dark bursts in the $\\beta_{\\rm OX}-\\beta_{\\rm X}$ plane: they are located below the dividing line of $\\beta_{\\rm OX} = \\beta_{\\rm X}-0.5$. In the same diagram, optically bright bursts are placed above the line $\\beta_{\\rm OX} = \\beta_{\\rm X}$, while all the GRBs that are still consistent with the fireball model will lie in the region defined by the relations $\\beta_{\\rm OX} = \\beta_{\\rm X}$ and $\\beta_{\\rm OX} = \\beta_{\\rm X}-0.5$. Optically bright events, for which the optical luminosity is too high if compared to the X-ray luminosity, are pretty rare, while there is a sizable fraction of events for which the X-ray emission seems to be in excess with respect to the observed optical one.  Based on these definitions, studies on GRBs samples in the {\\it Swift}-era showed that, despite an advancement on the GRB~detection quality both in the prompt response and position accuracy, the fraction of genuinely dark GRBs remains significant. Melandri et al. (2008), Cenko et al. (2009), Zheng et al. (2009), Gehrels et al. (2008), Fynbo et al. (2009) and Greiner et al. (2011) found a fraction of dark GRB~in their samples of about 50\\%, 50\\%, 20\\%, 20\\%, 30\\% and 40\\% respectively. A higher fraction of dark GRBs is found in samples based on observations done by single ground based telescopes (i.e. the 2-m Liverpool and Faulkes Telescopes, Melandri et al. 2008; the 60-inch Palomar telescope, Cenko et al. 2009; the 2.2-m GROND telescope, Greiner et al. 2011) while a smaller fraction is detected if the whole {\\it Swift} sample is considered (Gehrels et al. 2008 for GRBs up to end of 2007). This difference is certainly related to the different properties of the GRBs samples considered. In any case the population of dark bursts seems to be $\\geq 20\\%$ of the entire GRB~class.  In the era of rapid follow-ups the darkness of these events could not be ascribed to lack of sensitivity, late observational times or rapid temporal decays (Roming et al. 2006, Melandri et al. 2008). Different scenarios have been proposed to explain it: \\begin{itemize} \\item {\\it low-density scenario}: if the relativistic ejecta decelerate in a uniform low-density medium then the optical afterglow can be intrinsically faint with respect to the X-ray emission; \\item {\\it dust scenario}: if dark bursts are exploding in galaxies with local thick and dusty (i.e., high metallicity) environments (with possibly some intervening systems along the line of sight) their optical afterglows could be suppressed by extinction, without thereby affecting their higher energy radiation. The extinction law characterising the bursts environment might be similar to the one observed in the local Universe or biased toward large dust grains; \\item {\\it high redshift scenario}: if the burst is occurring at very large distances, its visible light could be completely extinguished as a result of the absorption of the Ly-$\\alpha$ forest and Ly-$\\alpha$ dropout redshifted into the optical bands. \\end{itemize} The latter explanation seems to be responsible for the non-detection only for a small fraction of the population of ``dark burst\" (Greiner et al. 2011), while a combination of the first two effects seems to be a more realistic scenario (Perley et al. 2009).  In this paper, we will investigate the properties of the population of ``dark burst\" present in a complete sub-sample of {\\it Swift} long GRBs, with a high percentage of redshift determination (Salvaterra et al. 2011, Nava et al. 2011). The use of a complete sub-sample of GRBs allowed us to draw more firm conclusions about the properties of this class of events with no bias in the selection criteria. The text is organised as follows: in Section 2 we will describe the general properties of the dark bursts that belong to our selected sample. We then discuss the results on the dark bursts population, their redshift distribution and their luminosity in Section 3 and finally we draw our conclusions in Section 4. Throughout the paper we assume a standard cosmology with with $H_{0} = 70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{m} = 0.3$ and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "We studied the properties of the sub-class of optically dark bursts detected in the complete sample of bright {\\it Swift} long GRBs presented in Salvaterra et al. (2011) and Nava et al. (2011). From our analysis we find that this population has the same redshift distribution of the whole sample. Thanks to the high completeness in redshift of this sample we have been able to confirm the existence of a genuine fraction ($\\sim 25-35 \\%$) of optically dark events.  The majority of those do not show any darkness evolution, being optically dark from very early time. Those events cannot be explained in the context of to the high-{\\it z} scenario and we confidently exclude that their darkness is due to the Ly-$\\alpha$ absorption in the optical bands.  The dark bursts do not have different prompt properties compared to the normal events (see Table \\ref{KStDB} and Fig. \\ref{figlumx}). However the former display lower optical flux and relatively higher X-ray flux with respect to the latter, as also noted by van der Horst et al. (2009). In particular they are always located in the high tail of the X-ray luminosity distribution, showing that they are, at the same time, not only fainter in the optical but also brighter in the X-ray. Because we are dealing with bright prompt events, it is unlikely that the darkness of the events in our sample could be ascribed only to their intrinsic faintness or to rapid temporal decay. These bursts are indeed faint events in the optical band but they also have a slight excess of emission in the X-ray band (although the excess is not statistically significant as the optical deficiency). The higher X-ray column densities observed for the dark bursts in our sample (Campana et al. 2011) clearly indicate that they formed in metal-rich environments where a fair amount of dust must be present. This disfavour the low-density scenario.  Therefore the most plausible explanation is left to be found in the context of the dust scenario. Since we take into account the Galactic absorption, if this scenario is correct then the observed darkness is due to high local extinction in a dense environment or to absorption from intervening material. Unfortunately we do not have complete information regarding all the possible intervening systems that the light of these events may encounter along their lines of sight and probably this effect becomes relevant only for GRBs at high redshift (Campana et al. 2006, Campana et al. 2010). Therefore it can play a significant role only for a small fraction of dark bursts in our sample. The former possibility instead plays for sure an important role for a couple of events in our sample: GRB~060210 (A$_{\\rm V} \\geq 4$; Curran et al. 2007) and GRB~080607 (A$_{\\rm V} \\geq 3.5$; Perley et al. 2009).  High values of X-ray column densities are hint of high local absorption; however, we do not know what happens to the dust environment in the vicinity of the GRB. Their prompt emission could, for example, alter the local dust composition, destroying small dust grains in favour of the bigger ones. This may change the extinction law that will become flatter, nearly constant and independent from the observed wavelength, in the UV-optical band, with respect to the one commonly observed in the local Universe. This effect, known as gray dust, even if it is not easy to recognise was successfully invoked to explain few GRB spectral energy distributions (Stratta et al. 2005). We do not have clear hints of this scenario in our sub-sample of dark bursts but we cannot exclude that it plays a significant role. However the investigation of the ``dust scenario\" and the ``gray dust scenario\" needs a more detailed analysis of the spectral energy distribution of each single event in our sample; this is beyond the aim of this work and will be exhaustively treated in a dedicated forthcoming work.  On the basis of our results we were able to understand more about the nature of dark bursts when compared to bright events: \\begin{itemize} \\item they have similar prompt properties; \\item they have a higher X-ray flux and X-ray luminosity and, at the same time, lower optical flux; \\item they are located in different (denser) environments; \\item they cannot be explained in the context of the high-{\\it z} or low-density scenarios; \\item their darkness must be mainly related to circum-burst dust absorption. \\end{itemize}"
    },
    "1112/1112.5954_arXiv.txt": {
        "abstract": "We describe mixing scalar particles and Majorana fermions using Closed-Time-Path methods. From the Kadanoff-Baym equations, we obtain the charge asymmetry, that is generated from decays and inverse decays of the mixing particles. Within one single formalism, we thereby treat Leptogenesis from oscillations and recover as well the standard results for the asymmetry in Resonant Leptogenesis, which apply when the oscillation frequency is much larger than the decay rate. Analytic solutions for two mixing neutral particles in a constant-temperature background illustrate our results qualitatively. We also perform the modification of the kinetic equations that is necessary in order to take account of the expansion of the Universe and the washout of the asymmetry. ",
        "introduction": "There are several different methods of calculating the lepton asymmetry, that may be generated in the Early Universe. This variety of approaches is a feature that Leptogenesis~\\cite{Fukugita:1986hr} has in common with many other topics in Particle Theory.  The approach that is perhaps most commonly applied is to calculate the decay asymmetry of individual right handed, singlet Majorana neutrinos $N$ from $S$-matrix elements~\\cite{Fukugita:1986hr,Luty:1992un,Covi:1996wh,Buchmuller:1997yu}. These asymmetries, obtained by Quantum Field Theoretical methods, are subsequently substituted into classical Boltzmann equations, that describe the macroscopic non-equilibrium dynamics and the evolution of the asymmetry. By LSZ-reduction, $S$-matrix elements can be obtained from time ordered $n$-point Green functions. We then enjoy the usual simplifications of time-ordered perturbation theory, {\\it i.e.} we can calculate the $n$-point functions using the standard diagrammatic Feynman rules. The price to pay for this simplification is cutting the direct link between the non-equilibrium dynamics and the generation of the $CP$-asymmetry at the quantum level. A particular problem turns out to be due to the instability of the singlet neutrinos, which causes that the $S$-matrix, that involves the $CP$-asymmetric decays and inverse decays of these singlets, is not unitary, or rather, that it does not correspond to an $S$-matrix in the proper sense. As a consequence, the unitary evolution of the non-equilibrium system must be carefully re-implemented, which is performed in practise through the non-trivial method of real intermediate state subtraction~\\cite{Kolb:1979qa}. Without real intermediate state subtraction, one would predict from substituting the rates from the $S$-matrix elements into the network of Boltzmann equations the presence of a lepton asymmetry even in thermal equilibrium, which would be an unacceptable violation of the $CPT$ theorem.   For Resonant Leptogenesis~\\cite{Flanz:1996fb,Covi:1996wh,Pilaftsis:1997dr,Pilaftsis:1997jf,Pilaftsis:2003gt,Pilaftsis:2005rv}, it is assumed that the mass difference $\\Delta M$ of two of the singlet neutrinos is small compared to their average mass $\\bar M$, or, more precisely, the energy difference is small compared to the average energy. Decays and inverse decays, that induce off-diagonal correlations of the singlet neutrinos in their mass eigenbasis, then lead to a relatively large mixing, which enhances $CP$-violation. In the $S$-matrix approach, the mixing appears as a wave function correction to the singlet neutrino propagator in the decay diagrams~\\cite{Covi:1996wh,IgnatievKuzminShaposhnikov,Botella:1990vf,Liu:1993tg}. The mixing that is calculated in this way is time-independent, which excludes the possibility of incorporating the effect of oscillations.  Alternatively, one can avoid the subtleties concerning unitarity by computing the real-time evolution of the quantum state from the outset. This can either be achieved by using the canonical Hamiltonian approach or functional, Closed-Time-Path (CTP)~\\cite{Schwinger:1960qe,Keldysh:1964ud,Chou:1984es,Calzetta:1986cq} methods. This may appear as a complicated task, but one should keep in mind that the $CP$-asymmetry arises in the $S$-matrix language at the one-loop level. In a real time approach, it turns out that in a diagrammatic representation, we have to go to two-loop level. However, the leading $CP$-violating effects can be extracted by imposing some of the internal particles to be on-shell, which may be considered as a generalisation of the Cutkosky rules~\\cite{Cutkosky:1960sp}. In the end, the integrals that are relevant for the leading order prediction of the $CP$ asymmetry are identical to those encountered in the $S$-matrix approach, but with the advantage that an overcounting of particular processes in the Boltzmann equations, that would violate unitarity, is avoided by construction.  The canonical approach makes use of the property of the Hamiltonian as the time evolution operator. A blueprint of this canonical approach, is the theory of neutrino oscillations in Refs.~\\cite{Dolgov:1980cq,Barbieri:1990vx,Enqvist:1990ad,Enqvist:1991qj,Sigl:1992fn}. It has been noticed that when supplemented by $CP$-violation, the oscillations may lead to the creation of a lepton asymmetry~\\cite{Akhmedov:1998qx,Asaka:2005pn,Gagnon:2010kt}. The possibility of describing mixing oscillations is one key advantage of the canonical approach when compared to using $S$-matrix elements. On the other hand, without a set of diagrammatic rules, the perturbation expansion is rather cumbersome. This is why the remaining important $CP$-violating contribution, that is sometimes referred to as direct $CP$-violation ({\\it e.g.} in Ref.~\\cite{Flanz:1996fb}), has not been calculated in the Hamiltonian approach. In the $S$-matrix approach, this contribution arises from the penguin-shaped vertex diagram. Besides, the decays of the singlet neutrinos into Standard Model leptons are described in practise by appealing again to the $S$-matrix formalism. Therefore, one may state that within the Hamiltonian approach, no entirely closed picture for the emergence of the lepton asymmetry has yet been developed.  In principle, it should be possible to derive a set of diagrammatic rules that simplifies the calculations within the Hamiltonian formalism. However, such a set of rules is readily available. It arises in an elegant and rather intuitive manner from the functional formulation of the CTP formalism~\\cite{Schwinger:1960qe,Keldysh:1964ud,Chou:1984es,Calzetta:1986cq}. In the CTP approach, the Hamiltonian time evolution of the density matrix is replaced by Schwinger-Dyson equations for the two-point Green functions. These equations have four matrix components from the two indices that denote the two branches of the CTP. The Kadanoff-Baym equations are those of the Schwinger-Dyson equations, that encompass the macroscopic evolution of the system. Recently, this approach has proved successful in order to describe Leptogenesis within a single theoretical framework~\\cite{Buchmuller:2000nd,De Simone:2007rw,Garny:2009rv,Garny:2009qn,Anisimov:2010aq,Garny:2010nj,Beneke:2010wd,Beneke:2010dz,Garny:2010nz,Garbrecht:2010sz,Anisimov:2010dk}. Both, $CP$-violation from wave-function and vertex corrections are incorporated. Unitarity issues are resolved and an accurate account of all quantum-statistical effects on the asymmetry is made. Moreover, the formulation in terms of Green functions bears the potential of incorporating corrections from Thermal Field Theory within this formalism. For the $CP$-conserving production rate of singlet neutrinos, several phenomenologically relevant thermal corrections have recently been calculated in Refs.~\\cite{Asaka:2006rw,Kiessig:2010pr,Anisimov:2010gy,Salvio:2011sf,Kiessig:2011fw,Kiessig:2011ga,Laine:2011pq}. When it comes to the rate of $CP$-violation, an aspect that is not dealt with in Refs.~\\cite{Buchmuller:2000nd,De Simone:2007rw,Garny:2009rv,Garny:2009qn,Anisimov:2010aq,Garny:2010nj,Beneke:2010wd,Beneke:2010dz,Garny:2010nz,Garbrecht:2010sz,Anisimov:2010dk} is the possibility of flavour oscillations between the singlet neutrinos. The present work addresses this shortcoming.  While in Ref.~\\cite{Beneke:2010wd}, the wave function correction is calculated using the perturbative loop expansion, the present work captures the same effect by solving the Kadanoff-Baym equations for the singlet neutrino propagator directly. In the limit where the mass difference is much larger than the decay rate $\\Gamma_{\\rm D}$ of the singlets, $\\Delta M\\gg\\Gamma_{\\rm D}$, the standard perturbative result is reproduced. When $\\Delta M\\sim\\Gamma_{\\rm D}$ or smaller, it becomes important that the solution to the Kadanoff-Baym equations corresponds to a resummation of all wave-function insertions. This is a generalisation of the familiar procedure for time-independent situations, which can be performed by the summation of a geometric series. The result obtained in the present work from the Kadanoff-Baym equations encompasses both, time-independent and oscillating contributions to the mixing of the singlet neutrinos. We note that for a model of the asymmetric decay of a scalar field, it was clarified in Refs.~\\cite{Liu:1993ds,Covi:1996fm} that the Hamiltonian approach incorporates indeed the wave-function contribution to the $CP$-asymmetry from decays in the vacuum as it is usually calculated from $S$-matrix elements. The present work goes beyond this proof of principle not only because it also treats the decays of Majorana fermions, but also because it formulates the evolution of the asymmetry in an interacting, finite-density system within the single framework of the CTP formalism. Our results may therefore be employed for phenomenological studies of Resonant Leptogenesis in the Early Universe.  In order to further describe the context of this present work, we note that for computations of the asymmetry from $S$-matrix elements, it has been realised that the resonant enhancement factor $1/\\Delta M$ should be regulated by the finite width of the singlet neutrinos~\\cite{Pilaftsis:1997dr,Pilaftsis:1997jf,Buchmuller:1997yu,Pilaftsis:2003gt,Pilaftsis:2005rv,Anisimov:2005hr}. For solutions to the Kadanoff-Baym equations in Wigner space,\\footnote{The Wigner transform is the Fourier transform of a two-point function with respect to the relative coordinate, while retaining the average coordinate, see Appendix~\\ref{appendix:CTP:Wigner}.} it turns out that the non-equilibrium distribution functions are not multiplied by a finite-width, Breit-Wigner distribution, but rather by a zero width, Dirac $\\delta$-function. The observation of this somewhat curious property has lead to the speculation that these zero-width solutions arise due to an incorrect treatment of pinch singularities that must be addressed by a non-perturbative resummation~\\cite{hep-ph/9802312,Greiner:1998ri}, that is yet to be specified. This may be perceived as an obstacle for the formulation of an effective theory of Resonant Leptogenesis within the CTP framework. However, it was recently demonstrated analytically and numerically, that the vanishing width is a correct property of the out-of-equilibrium Wigner functions. Within loop diagrams, the zero width is corrected for by a resummation of derivative operators, which effectively shifts the argument of the $\\delta$-function from the real axis into the imaginary direction of the complex plane~\\cite{arXiv:1108.3688}. In the Kadanoff-Baym equations, this effect may be of relevance when the tree-level scattering processes are kinematically forbidden. As this is not the situation that we assume for the present purposes, the corrections that are explained in Ref.~\\cite{arXiv:1108.3688} are subdominant in the final result for the asymmetry and therefore are neglected here. The results of Ref.~\\cite{arXiv:1108.3688} imply however, that a straightforward solution of the Kadanoff-Baym equations should indeed regulate the asymmetry when $\\Delta M\\to 0$. Indeed, this is what we explicitly find in the present work.  After these comments, we present the outline of this paper: In Section~\\ref{section:scalar}, we use the Schwinger-Dyson equations on the CTP, in particular the Kadanoff-Baym equations, in order to calculate the asymmetry that emerges from the decays of mixing neutral scalar particles. The Kadanoff-Baym equations decompose into constraint and kinetic equations. When $\\Delta M\\gg \\Gamma_{\\rm D}$, one may use either of these equations in order to derive the asymmetry. If this is not the case, one should use an expansion based on $\\Delta M\\ll \\bar M$. In this limit, the kinetic equations take the form that is familiar from the time-evolution of a density matrix in the Hamiltonian approach. In order to gain a qualitative understanding, we derive analytical solutions for the mixing propagator in the case of two mixing flavours in a constant-temperature background.  In Section~\\ref{section:fermions}, we apply the same approximation strategies as for scalars to mixing Majorana neutrinos. As an additional complication, we need to solve for the various spinor components of the singlet neutrino propagator. However, it turns out that the flavour dynamics shares the same essential features with the scalar model. Even though our main interest is the decay of Majorana fermions, a reader who is less interested in these technical details can skip this Section and continue with Section~\\ref{section:illustrative}. There, we discuss some features of the analytical solutions for the asymmetry in a model in constant temperature background, for both, fermions and scalars. In particular, we see that when $\\Delta M\\gg \\Gamma_{\\rm D}$, the effect of oscillations of the singlet neutrinos averages out and the standard results for the decay asymmetry from $S$-matrix elements apply. However, when the condition $\\Delta M\\gg \\Gamma_{\\rm D}$ does not apply, the oscillations can have a significant effect on the asymmetry.  In Section~\\ref{section:effective:theory}, we present the modifications of the kinetic equations that arise in an expanding Universe. We therefore eventually arrive at an effective theory for Resonant Leptogenesis, that is formulated within the CTP framework. In Section~\\ref{section:conclusions}, there are a short summary of this paper as well as concluding remarks. ",
        "conclusions": "\\label{section:conclusions}  In this paper, we have developed the theory of Leptogenesis from mixing Majorana fermions and from mixing scalars in the CTP formalism. Our calculations rely on the expansion parameters $\\Gamma_{\\rm D}\\ll \\Delta M$ or $\\Delta M\\ll \\bar M$. Due to the smallness of the couplings, $\\Gamma_{\\rm D}\\ll\\bar M$, such that the parametric regimes where the two approximations apply overlap. When $\\Gamma_{\\rm D}\\ll \\Delta M$, the usual result for the $CP$-asymmetry from the wave-function correction applies~\\cite{Flanz:1996fb,Covi:1996wh,Pilaftsis:1997dr,Pilaftsis:1997jf,Pilaftsis:2003gt,Pilaftsis:2005rv,Beneke:2010wd}. Otherwise, oscillations are relevant for the $CP$ asymmetry, and it is necessary to solve the kinetic equations for the singlet neutrino propagator. In Section~\\ref{section:effective:theory}, we explain how such a solution may be obtained when taking the expansion of the Universe into account and how this solution can be used to calculate the final value of the lepton asymmetry of the Universe.  Within our results, we recover many elements that have previously been derived using Boltzmann equations and $S$-matrix elements or the Hamiltonian approach. However, there are a few features that differ from the earlier approaches: \\begin{itemize} \\item For Leptogenesis from the decays of Majorana fermions, we show that the kinetic evolution equations~(\\ref{flavour:dynamic:general}) can accurately reproduce the classic result~\\cite{Flanz:1996fb,Covi:1996wh,Pilaftsis:1997dr,Pilaftsis:1997jf,Pilaftsis:2003gt,Pilaftsis:2005rv,Beneke:2010wd} for the asymmetry in Resonant Leptogenesis~(\\ref{source:well-known}). This requires a detailed analytical computation of the evolution of the two helicity states of the singlet neutrinos, and has not been provided in earlier work based on the Hamiltonian approach~\\cite{Akhmedov:1998qx,Asaka:2005pn,Gagnon:2010kt}. \\item Our results consistently include the quantum statistical (Fermi-Dirac and Bose-Einstein) corrections that appear for all on-shell particles in the finite-density background. Note that these also apply to the on-shell particles within the loops of the diagrams that describe $CP$-violation. \\item When $\\Delta M \\sim \\Gamma_{\\rm D}$ or smaller, it is important to take account of the lepton and Higgs chemical potentials, because in this situation, the lepton number violation by the singlet neutrinos proceeds so slowly, that we have to take account of its time-dependence. This may be crucial in order to treat the washout of the asymmetry in a quantitatively correct manner. \\item Our effective theory is formulated within one single framework, the CTP formalism. In contrast, earlier approaches typically combine elements from $S$-matrices, classical Boltzmann equations or Hamiltonian transition amplitudes. In particular, in combination with Ref.~\\cite{Beneke:2010dz}, the methods presented here cover both, direct $CP$-violation and $CP$-violation from mixing, even when oscillations are important. Working within a single formalism may have the advantage, that it is easier to implement controlled approximations. The CTP approach combines the advantage of the $S$-matrix approach, that perturbative expansions can be efficiently formulated in terms of Feynman diagrams, with the advantage of the Hamiltonian real-time evolution, that avoids unitarity violation from the overcounting of real intermediate states. \\item Since the CTP approach is formulated in terms of Green functions and Feynman diagrams, this method for calculating the lepton asymmetry may be particularly suitable to be combined with corrections due to scattering processes in the thermal bath~\\cite{Asaka:2006rw,Kiessig:2010pr,Anisimov:2010gy,Salvio:2011sf,Kiessig:2011fw,Kiessig:2011ga,Laine:2011pq}. \\end{itemize}  The CTP approach appears as a well-suited method for a consistent description of macroscopic statistical systems that consist of High Energy Particle Physics degrees of freedom. In particular, it may provide an intuitive and accurate framework in order to describe the emergence of the baryon asymmetry of the Universe from the quantum effect of $CP$-violation. Further significant progress on the theory of Leptogenesis appears to be a realistic prospect. For Resonant Leptogenesis, the present work is intended to pave the way to improve the quantitative predictions and the qualitative understanding.  \\subsubsection*{Additional Notes} Another successful calculation of the asymmetry for Resonant Leptogenesis in the CTP approach, using partly different calculational strategies than those in the present work, was reported by Garny et al.~\\cite{Garny:2011}. During the final stage of the preparation of this manuscript, Ref.~\\cite{Asaka:2011wq} appeared. Based on the Hamiltonian approach, it is also pointed out there, that the lepton chemical potential is relevant for computing the correlations of the singlet neutrino flavours.  \\subsubsection*{Acknowledgements}  \\noindent This work is supported by the Gottfried Wilhelm Leibniz programme of the Deutsche Forschungsgemeinschaft and by the Alexander von Humboldt Foundation.  \\begin{appendix}"
    },
    "1112/1112.0606_arXiv.txt": {
        "abstract": "{We present a multi-object optical spectroscopy follow-up study of X-ray sources in a field along the Galactic Plane ($l=327.42^{\\circ}$, $b=2.26^{\\circ}$) which is part of the \\chandra\\ Multi-wavelength Plane survey (ChaMPlane). We obtained spectra for 46 stars, including 15 likely counterparts to X-ray sources, and sources showing an H$\\alpha$ color excess. This has led to the identification of a new cataclysmic variable (CV), CXOPS J154305.5-522709, also named ChaMPlane Bright Source 7 (CBS~7), and we identified 8 X-ray sources in the field as active late-type stars. CBS~7 was previously studied in X-rays and showed a hard spectrum and two periods: $1.22\\pm0.08$~h and $2.43\\pm0.26$~h. We present here clear evidence that the source is a CV through the detection of H, HeI and HeII emission lines in its optical spectrum. The hard X-ray spectrum and the presence of the HeII $\\lambda$4686 in emission with a large equivalent width suggest a magnetic CV. The near-infrared counterpart is significantly variable, and we found a period consistent with the longest X-ray period at $2.39\\pm0.05$~h but not the shortest X-ray period. If this period is the orbital period, this would place the system in the CV period gap. The possible orbital period suggests a dM4$\\pm$1 companion star. The distance is then estimated to be $\\sim$1~kpc. The system could be a relatively hard and X-ray luminous polar or an intermediate polar, possibly nearly-synchronous. } ",
        "introduction": "Cataclysmic variables (CVs) are semi-detached binary stars with orbital periods typically of the order of hours, consisting of a white dwarf (WD) primary accreting via Roche lobe overflow from a companion star which is usually a late-type, main-sequence star. Subtypes of CVs include non-magnetic CVs where an accretion disk forms around the primary, and magnetic systems (MCVs) where the white dwarf magnetic field is strong enough to truncate (intermediate polars --- IPs --- with $10^5<B<10^7$~G) or disrupt (polars, $B>10^7$~G) the accretion disk. The mass flow is then channeled along the magnetic field lines and creates a shock at the magnetic pole of the WD \\citep[e.g.][]{Cropper:1990p816,Patterson:1994p1017}. {Polars are identified through their polarized optical emission, due to cyclotron emission processes near the surface of the WD.} Polars are synchronous, locked systems, but the magnetic field in IPs is generally not strong enough to tidally lock the WD to the companion star and IPs are asynchronous rotators. IPs are thus often characterized by two distinct periods in their X-ray lightcurves. The first is associated with the binary orbital period ($P_{orb}$), and the second, shorter, with the spin period of the WD ($P_{spin} \\ll P_{orb}$). IPs tend to have larger orbital periods than polars, and the latter could be the end-products of the evolution of part of the IPs \\citep[e.g.][]{Norton:2008p4760}.  The sample of known field CVs suffers from various selection effects \\citep[e.g.][]{Pretorius:2007p4659,Servillat:2011p4345}, and is likely biased towards bright and highly variable systems. {Population synthesis models indicate that a fainter population showing fewer outbursts may dominate the population of DNe (up to 70\\%, \\citealt{Kolb:1993p1211}; \\citealt{Howell:1997p1209}). Recent detections in the SDSS of faint CVs with short periods confirmed the existence of such a dominant population \\citep{Gansicke:2009p1204}, but its proportion in absolute numbers is not fully constrained. The overall Galactic CV population, and the demographics of the various subclasses of CVs, are thus not fully understood, and this is especially true in the obscured regions of the Galactic Plane.}  In order to obtain an unbiased sample, we conduct the \\chandra\\ Multi-wavelength Plane Survey (ChaMPlane) which allows the identification of serendipitous X-ray sources discovered by the \\chandra\\ X-ray Observatory along the Galactic Plane \\citep{Grindlay:2005p738}. Based on a multi-wavelength dataset, we aim to identify rare Galactic populations of objects such as accreting white dwarfs, neutron stars, and black holes in order to study the nature, distribution and evolution of accreting compact objects in the Galaxy.  ChaMPlane has focused on the dense population of thousands of unidentified X-ray sources detected in the Galactic Bulge \\citep[e.g.][]{Muno:2009p875,Hong:2009p2767} with deep observations of three low extinction windows (Baade's window at latitude $b=-3.8^{\\circ}$, Stanek's window at $-2.2^{\\circ}$ and Limiting Window at $-1.4^{\\circ}$, South of the Galactic Center). This led to the discovery of a significant population of candidate CVs \\citep{vandenBerg:2006p5431,VanDenBerg:2009p1014,Hong:2009p2775,Koenig:2008p5597}, suggesting the presence of a large number of MCVs in the Galactic Bulge. We recently identified 10 periodic sources in the Limiting Window \\citep{Hong:2011p4354}. These sources seem consistent with IPs based on their X-ray luminosity and spectra, but they are more likely to be polars based on their detected periods (1 to 3~h). {These puzzling properties could be explained if the sources are polars with hard X-ray spectra, or members of a rare sub-class of MCVs, nearly synchronous IPs ($P_{spin}/P_{orb}>0.3$, 6 known sources, \\citealt{Ritter:2003p2684,Ritter:2010p2694}), some of which are considered as the missing link in the evolution of the MCVs from IPs to polars.}  We selected a subset of the brightest unidentified X-ray sources in the ChaMPlane database in order to characterize the Galactic Plane population of X-ray sources (\\citealt{vandenBerg:2011p7449}; \\citealt{Penner:2008p376}). One of these sources, CXOPS J154305.5-522709 or ChaMPlane Bright Source 7 (CBS~7) was tentatively classified as a CV based on its X-ray properties and H$\\alpha$ color excess. CBS~7 was found in the field of the \\chandra\\ archived observation of the high mass X-ray binary pulsar 4U 1538--52 (ObsId 90, $l=327.42^{\\circ}$, $b=2.26^{\\circ}$) which is part of ChaMPlane.  We report here new multi-object optical spectroscopy observations of selected X-ray sources in this field and near-infrared (NIR) images centered on CBS~7 in Section~\\ref{data}. We present results for CBS~7 confirming the source as a CV in Section~\\ref{results} and discuss in more detail its nature in Section~\\ref{discuss}. ",
        "conclusions": "We presented multi-wavelength follow-up observations of an archived \\chandra\\ observation from the ChaMPlane survey. Among the X-ray sources, one is identified as a CV (CBS~7), 8 as active late-type stars, 5 may be chance alignments. The stars with an H$\\alpha$ excess from photometry not corresponding to Chandra sources are all identified as late type stars.  {The source CBS~7, selected from its X-ray emission and an H$\\alpha$ excess is clearly identified as a CV through optical spectroscopy. It is possibly an IP, a near-synchronous IP or a polar. Both properties (X-ray emission and H$\\alpha$ excess) thus appear to be essential in the selection of accreting compact binary candidates, but are not sufficient to clearly identify them (one CV among 4 candidates in this work). The detection of variability as a third property seems more conclusive (see also \\citealt{Hong:2009p2775}). The use of those criteria in the ChaMPlane survey thus led to a majority of targets identified as foreground, late-type, active stars. However, and more importantly, we identify a few rare accreting compact binaries at intermediate distance in the Galactic Plane such as CBS~7 ($d \\gtrsim 1.0$~kpc), which are crucial in order to better understand the Galactic accreting compact binary population.}"
    },
    "1112/1112.5940_arXiv.txt": {
        "abstract": "We discuss the prospects for the detection of gamma-ray bursts (GRBs) by the Cherenkov Telescope Array (CTA), the next generation, ground-based facility of imaging atmospheric Cherenkov telescopes (IACTs) operating above a few tens of GeV. By virtue of its fast slewing capabilities, the lower energy threshold compared to current IACTs, and the much larger effective area compared to satellite instruments, CTA can measure the spectra and variability of GRBs with excellent photon statistics at multi-GeV energies, which would revolutionize our understanding of the physics of GRBs, test their validity as the origin of ultra-high-energy cosmic rays, and provide powerful probes of the extragalactic background light as well as Lorentz-invariance violation. Employing a model of the GRB population whose properties are broadly consistent with observations by the Gamma-ray Burst Monitor (GBM) and Large Area Telescope (LAT) onboard {\\it Fermi}, we simulate follow-up observations of GRBs with the Large Size Telescopes (LSTs), the component of CTA with the fastest slew speed and the best sensitivity at energies below a few hundred GeV. For our fiducial assumptions, we foresee that the LSTs can detect $\\sim$~0.1 GRBs per year during the prompt phase and $\\sim 0.5$ per year in the afterglow phase, considering only one array site and both GBM and the Space-based multi-band astronomical Variable Object Monitor ({\\it SVOM}) as the alert instruments. The detection rates can be enhanced by a factor of about 5 and 6 for the prompt emission and the afterglow, respectively, assuming two array sites with the same sensitivity and that the GBM localization error can be reduced to less than 1$^\\circ$. The expected distribution of redshift and photon counts are presented, showing that despite the modest event rate, hundreds or more multi-GeV photons can be anticipated from a single burst once they are detected. We also study how the detection rate depends on the intrinsic GRB properties and the delay time between the burst trigger and the follow-up observation. ",
        "introduction": "Gamma-ray bursts (GRBs) are the most violent explosive phenomena in the Universe. The prompt emission of GRBs typically has luminosity of $L_{\\gamma} \\sim 10^{51}$--$10^{52}~{\\rm erg}~{\\rm s}^{-1}$ at $\\sim$~0.1--1~MeV, and the following long-lasting afterglow emission is observed in the radio to X-ray bands~\\citep[see reviews,][]{Mes06,Zha07}. High-energy gamma-ray emission above $\\sim 0.1$~GeV was first seen by the Energetic Gamma-Ray Experiment Telescope (EGRET) detector on-board the Compton Gamma Ray Observatory~\\citep[e.g.,][]{Hur+94,Gon+03}. The Large Area Telescope (LAT) on-board the \\textit{Fermi} satellite has recently detected high-energy gamma rays from a much larger sample of GRBs~\\citep{Abd+09a,Abd+09b,Abd+09c,Abd+10,Abd+11,Ack+10,Ackermann2011}. The LAT GRBs exhibit the following features~\\citep[see recent reviews,][]{Granot2010,BG11,Pee11,Ino+12}. (a) Most of the LAT GRBs do not show significant suppression at the high-energy end of their spectra, though observed LAT limits on the GeV fluence for the GRBs detected by the Gamma-ray Burst Monitor (GBM) onboard the {\\it Fermi} satellite may suggest a steeping or cutoff in the high-energy spectrum~\\citep{Beniamini2011}; (b) Some of the LAT GRBs have an anomalous extra component in the $>0.1$~GeV range~\\citep{Abd+09c,Ack+10,Ackermann2011}, while others show a high-energy spectrum consistent with the Band function~\\citep{Band1993}; (c) In many cases, the emission onset in the $>0.1$~GeV energy range is delayed relative to that in the $<1$~MeV energy range. (d) The LAT GRBs often show the long-lived high-energy emission lasting longer than the duration of the sub-MeV component; (e) Not only long GRBs but also short GRBs seem to have the above features~\\citep{Ack+10}. Understanding these features is likely to give us important clues to GRB mechanisms and related astrophysics.  The Cherenkov Telescope Array (CTA), under plan as the next generation international gamma-ray observatory~\\citep{CTA10,CTA11}, will provide a great step forward in studying these issues. CTA will be constructed at two sites, one each in the northern and southern hemispheres, and comprises three types of telescopes: the Large Size Telescopes (LSTs) with 23~m diameter and 4.6$^\\circ$ field of view (FOV); the Medium Size Telescopes (MSTs) of 12~m and 8$^\\circ$ FOV; and the Small Size Telescopes (SSTs) of 7~m and 10$^\\circ$ FOV. With CTA, the sensitivity will improve by a factor of 5--10 in the 0.1--10~TeV range compared to existing Cherenkov telescopes. It will cover about 5 decades in energy, including energies below a few tens of GeV and above $100$~TeV, and the angular and energy resolution will be appreciably increased. In the following, we summarize open issues that can be unraveled by CTA. We refer the reader to \\citet{Ino+12} for an extensive overview on the science prospects for GRB observations with CTA. \\begin{enumerate} \\item The bulk Lorentz factor $\\Gamma$ is one of the key quantities in understanding the properties of relativistic jets making GRBs.  It is thought to be limited to order of $\\sim 1000$ in the classical fireball model with baryons, whereas higher Lorentz factors may be achieved by Poynting-dominated jets~\\citep[e.g.,][]{SDD01} or radiation-dominated jets with dissipation via e.g., jet-confinement~\\citep{Iok+11} or baryon-entrainment~\\citep{Iok10}.  The bulk Lorentz factor can be constrained from observations of the high-energy end of the spectra~\\citep{LS01}, and absence of high-energy cutoffs for the LAT GRBs have indeed given us lower limits on $\\Gamma$~\\citep[e.g.,][]{Abd+09a,Abd+09b,Abd+09d}. In the case of GRB 090926A~\\citep{Ackermann2011}, a sharp softening of the spectrum is observed at $\\sim 1.4$~GeV, and interpreting this as the pair-creation cutoff leads to $\\Gamma \\approx 720$. However, such estimates can be affected by the finite extent of the emission region~\\citep{Bar06}, time-dependence of the photon field~\\citep{Gra+08,AM11} and multiple emission regions~\\citep{Li10,Aoi+10,Has+11}, so that the resulting high-energy spectra can be more complicated.  For more detailed studies, observations with much better statistics are required. CTA will be ideal for these purposes. \\item The mechanism of the prompt emission is one of the most critical issues in the theory of GRBs. In the classical scenario, $\\sim$~MeV emission is explained by the optically-thin synchrotron radiation from electrons accelerated at internal shocks~\\citep{RM94}. However, this classical scenario has several problems in explaining observations~\\citep[see reviews,][]{Mes06,Zha07}, and many alternative scenarios such as the photospheric dissipation scenario~\\citep[e.g.,][]{Tho94,MR00,RM05,Iok+07,Bel10,Mur+12} and magnetic dissipation scenarios~\\citep[e.g.,][]{Lyu06,ZY11,MU11} have been suggested.  For the mechanisms of the $>0.1$~GeV emission, both leptonic and hadronic emission have been have been proposed.  Synchrotron emission~\\citep{Wan+09,Iok10,Dai+11}, synchrotron self-Compton (SSC) emission~\\citep{CGP10a,CGP10b}, external inverse-Compton (EIC) emission~\\citep{TWM09,TWM11}, and proton synchrotron or proton-induced cascade emissions~\\citep{AGM09,AIM10,RDF10,Mur+12} are currently viable explanations.  However, once the prompt emission is detected by CTA, it may be possible to discriminate among the emission mechanisms by analyzing the very-high-energy (VHE) spectrum and variability with high photon statistics. \\item It is believed that the long-lived high-energy component is related to the afterglow emission.  In the simplest scenario, the high-energy emission is explained by synchrotron emission from the non-radiative external forward shock with extreme parameters~\\citep{KB11} or radiative external forward shock~\\citep{Ghisellini10}.  The Klein-Nishina effect may play a role in the initial rapid decay~\\citep{Wan+10}, but the more natural explanation is the gradual turn-off of the prompt emission~\\citep{LW11,He+11,Max+11}. One of the important tests for synchrotron external shock scenarios is to see the maximum synchrotron cutoff that decreases as the Lorentz factor declines~\\citep{PN10}, and CTA may eventually see the SSC emission component.  Observations by CTA can also provide a clue to the origin of X-ray shallow-decay emission.  The shallow-decay behavior is observed in most X-ray afterglows of \\textit{Swift} GRBs~\\citep{Nou+06,Obr+06,Zha+06}, though it was not predicted by the standard afterglow model~\\citep{MR97,SPN98}. Its origin has been debated for many years, and possible models include energy injection \\citep[e.g.,][]{DL98,RM98,Zha+06}, the long-lasting internal activity \\citep{Ghisellini07,Kum+08,Mur+11}, the long-lasting reverse shock~\\citep{GDM07,UB07}, time-dependent microphysics~\\citep{ITYN06}, effect of cosmic ray escape~\\citep{Der07}, bulk Compton emission~\\citep{Pan08a}, prior emission~\\citep{Yam09}, and multi-component jets~\\citep[e.g.,][]{EG06,TIYN06}. These models should predict different high-energy emissions~\\citep{Fan+08,Pan08b,Mur+10,Mur+11}, so that the simultaneous observations by {\\it Swift} and the Space-based multi-band astronomical Variable Object Monitor ({\\it SVOM})~\\citep{Paul2011} as well as CTA would be useful for discriminating among the models. On the other hand, high-energy emission has been observed concurrently with X-ray flares, which are often seen in the early afterglow phase~\\citep{Abd+11}.  Although they are considered to originate from some internal dissipation processes rather than external shocks~\\citep{IKZ05}, their mechanism and the origin of the high-energy emission is still unclear.  CTA may provide a breakthrough for revealing these problems. \\item VHE signals from GRBs can be useful for other purposes. GRBs may be the main sources of the observed ultrahigh-energy cosmic rays (UHECRs)~\\citep{Wax95,Vie95}.  If this hypothesis is true, VHE gamma rays provide one of the crucial probes, and detections of characteristic signals, including hadronic cascade radiation~\\citep{BD98,PW05,DA06,AIM09,Mur+12}, ion synchrotron radiation~\\citep{BD98,Tot98a,Tot98b}, synchrotron pair-echo emission induced by ultra-high-energy photons~\\citep{Mur11}, and emission via nuclear photodisintegration and Bethe-Heitler pair production~\\citep{MB10}, are relevant to test this GRB-UHECR hypothesis. Also, VHE gamma rays from GRBs are useful as a probe of the Universe.  Since they are attenuated by the extragalactic background light (EBL), one can constrain the EBL by measuring the attenuation as in the case of blazars~\\citep[e.g.,][]{Aha+06,Alb+07}.  VHE gamma rays also lead to secondary emission as a result of interactions with the EBL, and detections or non-detections of this pair-echo emission enable us to probe the uncertain intergalactic magnetic fields in voids as well as the EBL~\\citep[e.g.,][]{RMZ04,Tak+08,Mur+09}. GRBs may occur even at the very distant universe, $z > 5$~\\citep[e.g.,][]{Kaw+06,Sal+09,Tan+09}, and VHE gamma rays might be generated by GRBs originating from population III stars~\\citep{TSM11}.  If we detect the attenuated VHE gamma rays from such distant GRBs, it would be possible to constrain the EBL in the high-redshift universe~\\citep{Gil+09,Ino+10} and/or cosmic magnetic fields at high redshifts~\\citep{Tak+11}. In addition, one can obtain insight into fundamental physics.  For example, using the temporal (and spectral) information, one can test the Lorentz-invariance violation predicted in some quantum gravity theories~\\citep[e.g.,][]{Ame+98,Abd+09d,SM10}. \\end{enumerate}  As described above, CTA should be a powerful tool to understand GRB physics and test other various ideas and hypotheses related to GRBs.  VHE gamma-ray astronomy has now been firmly established by state-of-the-art IACTs such as H.E.S.S., MAGIC, and VERITAS.  A firm detection of GRBs has not been reported so far~\\citep{Alb+07,Aha+09,Ale+10}, but CTA may eventually achieve the goal with much lower energy threshold and much better sensitivity, which will also allow us to obtain time-resolved spectra of much better quality than ever before in the energy range greater than a few tens of GeV. In this work, we report the prospects for detecting long GRBs (durations longer than 2~sec) with CTA, and how the results depend on some of the array performance and GRB properties. The situation considered here is that GRBs are detected by some satellites and followed up by the LSTs, which will play a principal role in GRB observations. In our simulation, we mainly assume {\\it Fermi}/GBM as the burst trigger, which has already provided data with high statistics and may be active in the CTA era. We also make a rough estimate for the alerts from the {\\it SVOM} satellite, which is planned to be launched before CTA operation. Some complementary aspects of this work are also presented in \\citet{Ino+12}.  The organization of the paper is as follows. In section~2, we describe the intrinsic GRB properties of both prompt and afterglow emissions in our simulations. In section~3, we model the performance of the LSTs and the follow-up observation of GRBs. Performing Monte Carlo simulation, we estimate the detection rate of GRBs with the LSTs in section~4.  Finally, section~5 is devoted to summary and discussion. We use cosmological parameters $H_0 = 70$~km~s$^{-1}$Mpc$^{-1}$,  $\\Omega_{m} = 0.3$ and  $\\Omega_\\Lambda = 0.7$. ",
        "conclusions": "\\label{sec:summary}  In this paper, we have presented the prospects for detection of GRBs with the LSTs of CTA, the most vital component for GRB observations with their fast slewing capability and the best sensitivity at the lowest energies. We have modeled and simulated the follow-up observation of GRBs with $T_{90}>2$~sec and $z<5$ alerted by {\\it Fermi}/GBM. Our GRB population is modelled according to a given luminosity function that is consistent with Swift observations together with well-known spectral correlations. We note the following strengths of our model: (1)~It reproduces the fluence and duration distributions of GBM bursts (Figure~\\ref{fig:S_vs_T90,S-pdf}). (2)~The differential sensitivity of our model LSTs is consistent with the official one given by the \\citet{CTA10} in the energy range less than a few 100~GeV (Figure~\\ref{fig:toy_model,0.5h}). (3)~The {\\it Fermi}/LAT detection rate is predicted to be 12~yr$^{-1}$ for the bursts satisfying both $T_{90}>2$~sec and $z<5$, which is roughly consistent with the actual rate of 7--8~yr$^{-1}$.  Assuming {\\it Fermi}/GBM alerts alone, our fiducial parameter set predicts the GRB detection rate with LSTs for one array site (i.e., north or south) as 0.03~yr$^{-1}$ for the prompt emission and 0.13~yr$^{-1}$ for the afterglow emission (Table~\\ref{table:GRBrate_theory1}). The expected event rates become larger when two array sites of CTA and additional alerts from {\\it SVOM}/ECLAIRs are taken into account --- if the array performance for the two sites are the same in the energy range less than a few 100~GeV, the detection rates go up to about 0.3~yr$^{-1}$ and 1~yr$^{-1}$ for the prompt and afterglow emissions, respectively, where we assume no overlap of FOVs of GBM and ECLAIRs. Note that in the above estimates, we optimistically estimate the onboard-localization ability of GBM in the prompt phase, expecting improvement by the CTA era (see Section \\ref{subsec:GBMlocalization}). For the afterglow emission, our treatment of the GBM localization ability hardly affects the rate estimation.  For our fiducial assumptions,  once CTA succeeds in detecting the prompt emissions for GBM alerts, the total photon counts $N_\\gamma$ are expected to be $> 10^2$ for 60~\\% of CTA detected events. Our simulation also shows that $N_\\gamma$ is 10--100 times larger than the number of GeV photons expected by LAT for the same bursts. This suggests that CTA can obtain the time resolved light curve in the energy range greater than a few tens of GeV with higher statistics than ever before. We expect 90~\\% of the prompt burst detected by CTA to have redshifts less than 3.5, and 90~\\% of the afterglows to have less than 2.9. Because of the follow-up observation after the GBM alerts, the redshift distribution is more affected by the GBM sensitivity rather than that of CTA. Hence, more frequent detections of high-$z$ prompt emission are expected for {\\it SVOM} alerts.  Studying the dependence of the detection rate on our model parameters, we found the following results: \\begin{description} \\item[(i)] For the prompt emission, if all GRB samples have the extra power-law component with luminosity 10~\\% of the Band component, the detection rate with CTA increases by a factor of 1.9. \\item[(ii)] For the afterglow, the spectral index has a relatively large effect on the detection rate. It decreases by a factor of $3$ with the extreme assumption that all bursts have the softer spectral energy index of $-1.5$, while in the harder case of $-0.5$, the rate increases by a factor of 1.5. \\item[(iii)] If the typical value of the delay time between the start of observations and the GRB trigger time is 1~min, the detection rate of the prompt emission increases by a factor of about 1.3 compared to our fiducial case. On the other hand, the rate of the afterglow depends slightly on the delay time. Although the shorter delay time does not have much effect on the enhancement of the detection rate itself, the ability of fast slewing ($\\lesssim 100$~sec) is necessary for catching the prompt emission. \\item[(iv)] It seems better to make follow-up observations for the alerts with the error radius of up to $\\sim 5^\\circ$, as long as all 4 LSTs simply point toward the best position localized by GBM, although the detection efficiency is decreased. \\item[(v)] The detection rate varies by less than 30~\\% for variation of the other parameters including those of the luminosity function. \\end{description}  Note that the LAT detection rate is estimated to be 14~yr$^{-1}$ in the case of (i), which is about 1.8--2.0 times as large as the observed LAT event rate. This discrepancy may be due to the existence of the spectral softening below the LAT band, suggested for subset of bright events \\citep{Beniamini2011}, our crude (but reasonable) LAT detection conditions, etc. Even if the spectral softening effect is a dominant cause, the detection rate with CTA is not necessarily expected to decrease from our results by a similar factor, 1--2, due to the following possible reasons. First, the spectral softening below the LAT band have been only suggested for some of bright bursts~\\citep{Beniamini2011}.  CTA will be sensitive even for typical (less bright) bursts, which might not have the spectral softening. Second, the bursts with the steep Band component below the LAT band may be hardly detected by LAT, whereas they might have an extra hard component at $>10$~GeV, which can be relevant for CTA. We have not taken into account such possibilities in our simulations, which remain as future work. In any case, CTA may clarify the properties of the extra hard component and spectral softening of the Band component, which would strongly constrain the prompt emission mechanism.  It is crucial to increase the alert frequency with good localization in order to achieve higher detection rates. Hence the improvement of the localization accuracy by future GRB alert facilities is important. Let us consider the alerting detector which has the same characteristics (e.g., sensitivity) as GBM except for the localization ability, and simply assume that the alerts from the detector have a constant error radius $\\sigma$ for all the alerted bursts. In this case, if we can use the localization error $\\sigma$ less than the threshold $\\sigma_{\\rm th}$, the probability of localization in the FOV gets higher, while the frequency of follow-up observations do not vary ({\\it CTAobs}~$\\sim 5$~yr$^{-1}$). The increase of the probability saturates at $\\sigma \\sim 1^\\circ$, and then the detection rate of {\\it Pdet} and {\\it Adet} for this detector are enhanced to $\\simeq 0.2$~yr$^{-1}$ and $\\simeq 1.2$~yr$^{-1}$, respectively. It implies that future alerting detectors with the better localization ability are more desirable.  As countermeasures against the alerts with large localization error, two observing strategies are proposed: one is to point each LST at different directions in the GBM error circle, while the other is to scan the region instead of pointing one location. The optimal strategy for each observation can be determined by the estimate of the detection probability, which requires the information on the trigger (such as the degree of localization error, the brightness, the expected delay time, and the zenith angle). In order to perform this and to increase the chance of detection, we need more detailed study on the sensitivity for each strategy as a function of some important parameters (such as zenith angle, the extended FOV, scan speed, etc.), which remains as a future work.  In the VHE region, the future High Altitude Water Cherenkov (HAWC) Observatory mission \\citep{AbeysekaraHAWC2011} may also detect GRBs with its large FOV ($\\sim15$~\\% of all the sky) and high duty cycle ($\\sim100$~\\%), while the GRB spectrum may be difficult to determine by HAWC, because of its small effective area and low energy resolution at $\\sim30$~GeV compared to CTA \\citep{goodmanHAWC2011}.  Hence, CTA and HAWC are complementary and both types of observations are important.  Very recently, \\citet{Gilmore2011a,Gilmore2011b} and \\citet{Bouvier2011arXiv} independently studied on expectations for GRB detection rates by CTA, and we find their results are broadly in agreement with our conclusions here."
    },
    "1112/1112.0562_arXiv.txt": {
        "abstract": "In 2000 July/August a microlensing event occurred at a distance of 2\\farcm33 from the center of the globular cluster M22 (NGC6656), observed against the dense stellar field of the Milky Way bulge. We have used the adaptive optics system NACO at the ESO Very Large Telescope to resolve the two objects that participated in the event: the lens and the source. The position of the objects measured in 2011 July is in agreement with the observed relative proper motion of M22 with respect to the background bulge stars. Based on the brightness of the microlens components we find that the source is a solar-type star located at a distance of $6.0\\pm1.5$~kpc in the bulge, while the lens is a $0.18\\pm0.01\\msun$ dwarf member of the globular cluster located at the known distance of $3.2\\pm0.2$~kpc from the Sun. ",
        "introduction": "The effect of gravitational microlensing of background stars by compact objects located in globular clusters was analyzed for the first time by \\cite{1994AcA....44..235P}. He showed that thanks to the usually well-known distances to the source and the lens, and transverse velocity between the populations to which the objects belong, it is possible to derive the lens mass when the event time scale is measured. Paczy\\'nski suggested to monitor globular clusters like M22 or 47~Tuc in front of the rich background of either the Galactic bulge or the Small Magellanic Cloud. According to his calculations one should detect up to a few microlensing events in one year of continuous monitoring of M22 with a 1~m class ground-based telescope. Some microlensing events detected so far toward the bulge in microlensing surveys such as OGLE \\citep{2000AcA....50....1U}, MACHO \\citep{2000ApJ...542..281A}, and MOA \\citep{2001MNRAS.327..868B} might be associated with globular clusters \\citep{1998A&A...336..411J,2008IJMPD..17.2305D}.  \\cite{2005AcA....55..261P} presented the results of a search for erupting objects in the field of the globular cluster M22. The cluster was observed over 10 weeks in 2000-2001 with the 1.0~m Swope telescope at Las Campanas Observatory as one of the targets of the Cluster AgeS Experiment \\citep[CASE;][]{2005AIPC..752...70K}. Besides two erupting dwarf novae they found a probable microlensing event located at $\\alpha_{2000.0}=$18:36:22.40, $\\delta_{2000.0}=-$23:56:29.4, i.e., only 2\\farcm33 $=1.75~r_{\\rm c}=0.69~r_{\\rm h}$ from the cluster center, where $r_{\\rm c}=1\\farcm33$ and $r_{\\rm h}=3\\farcm36$ are the core radius and half-mass radius of the cluster, respectively (taken from the 2010 version of \\cite{1996AJ....112.1487H} catalog). The brightness of the object increased by about 0.8~mag in $V$ over 20~days. Around 2000 August~5 it reached a maximum of $V=19.1$~mag and then faded to a constant level of $V=19.9$~mag. Based on its color at maximum brightness the authors excluded the possibility that the object could be a dwarf nova. They fitted a single lens model to the light curve and found that the most likely geometry of the event places the source in the Galactic bulge and the lens in the cluster. The fitted parameters are: the epoch of maximum $t_0=2451759.70^{+0.33}_{-0.34}$, characteristic (Einstein) time $t_{\\rm E}=15.9\\pm1.1$~days, impact parameter $u_0=0.54^{+0.02}_{-0.18}$ in units of Einstein radius $r_{\\rm E}$, $V_{\\rm S}=19.92^{+0.62}_{-0.02}$~mag, and $V_{\\rm L}=24.8^{+\\infty}_{-4.0}$~mag, where $V_{\\rm S}$ and $V_{\\rm L}$ are $V$-band magnitudes of the source and lens, respectively. The authors assessed the mass of the lens to $M=0.14^{+0.10}_{-0.02}\\msun$. Large uncertainties of the above values result from the faintness of the object and partial coverage of the event.  After many years, in some special cases it is possible to directly detect the lens, measuring its mass and the geometry of the microlensing event \\citep[e.g.,][]{2001Nature...414..617,2007ApJ...671..420K}. In this Letter we resolve the microlensing system components based on new near-IR high-resolution images, measuring the complete geometry of the event, and the parameters of the source and lens stars. The event reported here is the first confirmed microlensing event in a globular cluster. We note that brightening episodes detected in {\\it Hubble Space Telescope} ({\\it HST}) images of M22 by \\cite{2001Natur.411.1022S} were later reexamined and interpreted either as a dwarf nova type outburst \\citep{2003ApJ...597L.137A} or as a result of cosmic ray double hits \\citep{2002ApJ...565L..21S}. ",
        "conclusions": "In the Letter we have shown that the microlensing event which occurred $2\\farcm33$ from the center of the globular cluster M22 in 2000 July/August involved a $0.18\\pm0.01\\msun$ dwarf of the cluster and a background solar-like star located in the Galactic bulge. Almost 11 years after the event, using high-resolution near-IR image we resolved the two microlensing components. The observed position of the source and lens stars as well as their brightness are consistent with the proposed earlier geometry of the event. Additional evidence comes from the constant brightness of the target object in the last two years (2010-2011)."
    },
    "1112/1112.4388_arXiv.txt": {
        "abstract": "Gravitational wave detectors from the advanced generation onwards are expected to be limited in sensitivity by thermal noise of the optics, making the reduction of this noise a key factor in the success of such detectors. A proposed method for reducing the impact of this noise is to use higher-order Laguerre-Gauss (LG) modes for the readout beam, as opposed to the currently used fundamental mode. We present here a synopsis of the research program undertaken by the University of Birmingham into the suitability of LG mode technology for future gravitational wave detectors. This will cover our previous and current work on this topic, from initial simulations and table-top LG mode experiments up to implementation in a prototype scale suspended cavity and high-power laser bench. ",
        "introduction": "Advanced gravitational wave detectors such as Advanced LIGO \\cite{Harry10} and Advanced Virgo \\cite{VIR-027A-09} are being constructed with the goal of a tenfold increase in sensitivity over the first generation of detectors \\cite{Sigg04}\\cite{virgo2006short}. This is expected to be sufficient to make the first conclusive gravitational wave detection. The projected sensitivity of both Advanced LIGO and Advanced Virgo will be limited over the observation frequency band by several fundamental noise sources: seismic noise, radiation pressure noise, suspension thermal noise, shot noise, and thermal noise of the test masses themselves \\cite{Rowan05}. The dominant contribution to the total thermal noise of the test masses is expected to be from the Brownian motion of the atoms that make up the reflective coatings. This motion causes an uncertainty in the phase of the light reflected from the test mass, and will therefore appear as noise in the detector output. This thermal noise has until now not been a limiting noise source in detectors, but as progress within the field pushes other noise sources ever lower, its mitigation becomes a more urgent priority.  One way to reduce the effects of the coating Brownian thermal noise is to use a larger beam spot on the test mass surfaces, since a larger beam averages more effectively over the random motions of the surface. The coating Brownian thermal noise power spectral density has been shown to scale with the inverse square of the Gaussian beam size parameter $w(z)$ \\cite{Vinet09}. However, as the beam spots are made larger, a larger fraction of the light power is lost over the edge of the reflective surface. This fraction is known as the \\emph{clipping loss}. If one uses a beam with a more homogeneous radial intensity distribution than the commonly used LG$_{00}$ beam, the Brownian thermal noise level can be reduced without increasing the clipping loss. Such beams are often referred to as \\emph{flat} beams. Candidate flat beams for gravitational wave detectors include mesa beams \\cite{mesa}, optimised conical beams \\cite{conical} and higher-order LG beams \\cite{mours}.  Higher-order LG modes have the advantage over other flat beams of being theoretically compatible with the spherical mirror surfaces that are currently used with the LG$_{00}$ mode. Equation \\ref{eqn:lghxamp} shows the complex amplitude function for the helical higher-order mode set, \\begin{equation} u_{p,l}^{\\textrm{hel}}(r,\\phi,z) =\\frac{1}{w(z)}\\sqrt{\\frac{2p!}{\\pi(|l|+p)!}}\\quad e^{ \\left(2p+|l|+1\\right)\\Psi(z)}\\\\ \\times \\,\\left(\\frac{\\sqrt{2}r}{w(z)}\\right)^{|l|}L^{(|l|)}_{p}\\left(\\frac{2r^2}{w(z)^2}\\right) e^{-\\textrm{i}k\\frac{r^2}{2q(z)}+\\textrm{i}l\\phi},\\\\ \\label{eqn:lghxamp} \\end{equation} where $p$ is the radial mode index, $l$ is the azimuthal mode index, $w$ is the beam radius, $k$ is the wave-number, $q$ is the complex Gaussian beam parameter, $L^{(|l|)}_{p}$ are the associated Laguerre polynomials, and $\\Psi$ is the Gouy phase, given by \\begin{equation} \\Psi(z)=\\textrm{arctan}\\left(\\frac{\\lambda z}{\\pi w_0^2}\\right). \\label{eqn:Gouyphase} \\end{equation} The order of a higher-order LG mode is given by $2p+|l|$. An alternative sinusoidal form of LG modes with a sinusoidal amplitude dependence in azimuthal angle can be used equally well. In this article we will discuss the research performed by the Birmingham group towards evaluating the feasibility of LG mode technology for implementation in detectors of the advanced generation and beyond to reduce the test mass thermal noise limit. ",
        "conclusions": "We have investigated the suitability of higher-order LG mode technology for use in future gravitational wave interferometers for reducing the test mass thermal noise, from many different angles. An initial simulation study showed that the longitudinal and angular control performance of the LG$_{33}$ mode in an Advanced Virgo-like interferometer was significantly better than the LG$_{00}$ mode with the same clipping losses at the cavity mirrors. It was also shown that increases in the observable event rates of Advanced Virgo by up to a factor of 2 could be gained by using the LG$_{33}$ mode in place of the LG$_{00}$ mode. Table-top experiments verified the equivalence of the LG$_{33}$ mode and LG$_{00}$ mode for PDH longitudinal control signal generation in an optical cavity, and demonstrated the mode cleaning of SLM generated LG$_{33}$ beams to reach a mode purity of over 99\\,\\%. More detailed theoretical investigations into the effects of mirror surface distortions on LG mode purity within optical cavities led to the derivation of an analytical formula for the coupling between LG modes caused by specific mirror distortions. Using this formula, mirror specifications required for compatibility with the LG$_{33}$ mode were subsequently derived. We are currently engaged in a collaborative effort with the University of Glasgow to verify the results of these theoretical investigations, by means of experiment within the 10\\,m prototype facility in Glasgow. We are also currently investigating the feasibility of generating high-power and high-mode purity LG$_{33}$ beams at the Hannover high-power laser facility. Higher-order LG modes can potentially offer a significant improvement in the sensitivity of future gravitational wave interferometers: our research program will evaluate how much of that potential is likely to be realisable in practice. \\ack We thank B. Barr, A. Bell, S. Huttner, J. Macarthur, B. Sorazu and K. Strain for providing the opportunity and helping to perform the LG mode experiments at the Glasgow 10\\,m prototype, and we thank C. Bogan, P. Kwee and B. Willke for the same assistance in the high-power LG modes experiment in Hannover. This document has been assigned the LIGO document number LIGO-P1100146."
    },
    "1112/1112.3561_arXiv.txt": {
        "abstract": "We study  coupled axial and polar axisymmetric oscillations of a neutron star endowed with a strong magnetic field, having both poloidal and toroidal components.  The toroidal component of the magnetic field is driving the coupling between the polar and axial oscillations. The star is composed of  a fluid core as well as  a solid crust. Using a two dimensional general relativistic simulation and a magnetic field $B=10^{16}$G, we study the change in the polar and axial spectrum caused by the coupling. \\textit{We find that the axial spectrum suffers a dramatic change in its nature, losing its continuum character}. In fact, we find that only the \\textquoteleft edges' of the continua survive, generating a discrete spectrum. As a consequence the crustal frequencies, that in our previous simulation could be absorbed by the continua, if they were embedded inside it, are now long living oscillations.  They may lose their energy only in the very special case that they are in resonance with the  \\textquoteleft edges' of the continua. ",
        "introduction": "In the last few years, a great attention was devoted to the interpretation of the quasi periodic oscillations (QPOs) discovered in the tail of giant flares in Soft Gamma Repeaters (SGRs).  Those oscillations cover a wide range of frequencies going from a few Hz up to kHz. Until now only in two SGRs it was possible to observe clearly QPOs: in the SGR 1806-20 and in the SGR 1900+14 (see \\cite{I2005}, \\cite{Ws2006a} for a  review). In the SGR 1900+14  four frequencies have been detected: 28, 53, 84, and 155 Hz while in the SGR 1806-20 several QPOs have been discovered: 18, 26, 30, 92, 150, 625 and 1840 Hz. A recent reanalysis of the data for the SGR 1806-20  carried on by \\cite{2011A&A...528A..45H} with a different method, shows the presence of additional QPOs at 16.9, 21.4, 36.4, 59.0 and 116.3 Hz.  Understanding the nature of the QPOs will offer  an unique opportunity to investigate the properties of neutron stars and  to constrain the nuclear matter in neutron stars as well as the topology and strength of their magnetic field. For those reasons, since their discoveries, QPOs have been the subject of an intensive study. The first hypothesis, in order to explain QPOs, was to consider  pure shear modes of the crust, excited after the burst generated by giant flares. This idea proposed by \\cite{1998ApJ...498L..45D} was investigate in several analytical and numerical works (see for example \\cite{2007MNRAS.375..261S}, \\cite{SA2007}). Those works evidenced that not all the lower QPOs could be explained as pure crustal frequencies. An alternative scenario has been then proposed by \\cite{2006MNRAS.368L..35L} and \\cite{2006MNRAS.371L..74G}, involving  global crust-core oscillations of the star. In particular,  \\cite{2007MNRAS.377..159L} showed with a toy model that the QPOs spectrum should  be continuum, with turning points that were called \\textquoteleft edges' of the continuum. Following this idea, a significant amount of  work has been devoted  in understanding global Alfv\\'en oscillations in neutron star, using general relativistic models (\\cite{2008MNRAS.385L...5S}, \\cite{2009MNRAS.396.1441C}, \\cite{2009MNRAS.397.1607C}) or Newtonian model (\\cite{2008MNRAS.385.2069L}).  All those works made clear that global Alfv\\'en modes in a pure fluid neutron star could not   explain the small gap between the lower frequencies observed in QPOs.  A further step was then needed: to add a solid crust and to study the effect of its presence on the oscillations of the core ( see  \\cite{2010JPhCS.222a2031K}). \\cite{2011MNRAS.410.1036H}, using a non-relativistic model, showed that the presence of the crust strongly influences the continuum spectrum. In particular, they found that near the edges of the continua other discrete Alfv\\'en modes appear. This results was partially confirmed by \\cite{2011MNRAS.410L..37G} and \\cite{2012MNRAS.421.2054G}, using a non-linear relativistic method: the authors could find some frequencies in the gaps between two near continua, however they stated that those oscillations were quickly damped. In a linear relativistic simulation, \\cite{2011MNRAS.414.3014C} confirmed the presence of those discrete Alfv\\'en modes in the gap between two  continua and showed that the crustal modes could live long if they are located in those gaps.  The discrete Alfv\\'en modes as well as the crustal modes in the gap between the first two continua could explain the somehow dense spectrum of the   observed lower  QPOs. However as pointed out in \\cite{2012MNRAS.420.3035V}, it was still difficult to explain the higher frequencies as the 625Hz because at higher frequencies the continuum dominates, absorbing all the discrete frequencies, and in addition the continua  overlap to each other: in this way it is not possible to isolate the edges of the continuum and to have a unique interpretation for the 625Hz QPO.  A possibility to explain the high frequency QPOs came from the study of polar oscillations in magnetars. The work by   \\cite{2009MNRAS.395.1163S} showed that the polar Alfv\\'en spectrum has a discrete nature and, in addition, the lower polar Alfv\\'en modes have frequencies  around few hundred Hz. However, it was not clear how this result could affect the axial spectrum.  In this work, we make an attempt to investigate the effect of the coupling between axial (torsional)  and polar oscillations  on the  spectrum. Our magnetar model is composed of a fluid core and a solid crust and is permeated by a magnetic field with strength $B=10^{16}$G which acquires both poloidal and toroidal components. The toroidal component is the coupling \\textquoteleft pipeline' between the axial and polar part of the spectrum. By using a two dimensional relativistic, linear code, we find that the axial spectrum is strongly modified by the presence  of the coupling: its continuum feature is destroyed, leaving behind only its \\textquoteleft edges'. In this new picture, both Alfv\\'en  and crustal modes have a discrete nature and the latter ones, if out of resonance with the   \\textquoteleft edges' of the continuum, could live long enough to be seen.  The paper is organized as follows: in Section 2 we present our background model, while in Section 3 we derive the perturbation equations and  test our code by comparing it with earlier results. In Section 4 we study a configuration where the coupling between axial and polar oscillations is active only on the axial part (i.e. we do not take into account the coupling to  the polar part). In Section 5 the full coupling between axial and polar oscillations is presented, including boundary conditions and initial data. The Section 6 and the Section 7 are devoted to comments and conclusions. ",
        "conclusions": "In this paper we study the coupling between the axial and the polar type of oscillations in a strongly magnetised neutron star, composed of a fluid core and a solid crust and permeated by a magnetic field with both a toroidal and a poloidal component. This mixed magnetic field configuration  is though to be stable (see \\cite{2006A&A...450.1077B}, \\cite{2011ApJ...735L..20L}, \\cite{2011ApJ...736L...6C} and references therein). Following the work by  \\citep{2009MNRAS.395.1163S}  (for the polar part) and \\cite{2011MNRAS.414.3014C} (for the axial part), we derive a coupled system of equations that involves both the polar perturbation functions $W$ and $V$ and the axial function $\\cal {Y}$. The equations are coupled through the toroidal component of the magnetic field. After solving those equations with the appropriate boundary conditions and  various sets of initial data, we obtain the axial and polar spectrum.  However, we found that both the spectra are strongly influenced by the coupling: the axial spectrum, that in absence of coupling had a discrete nature with the lower frequency located around a few hundred Hertz, preserves its discrete nature but some very low frequencies appear in the 20-300 Hz range. Those lower frequencies seems to be strictly related to the energy transfer between the axial and the polar oscillations via toroidal magnetic field: the lower frequencies found in the axial spectrum are the same found in the polar one, with a tiny shift of 1 or 2 Hz. However, the most relevant change is the one that concerns the axial spectrum: examining the FFT amplitude at various points in the star by varying the phase, \\textit{we observe that the axial spectrum has a discrete nature}. In fact, we can observe only the edges of the continuum that we found in \\cite{2011MNRAS.414.3014C}: those edges are still Alfv\\'en modes, in fact they scale with the strength of the magnetic field and their restore force is the magnetic force. \\textit{Due to the presence of the crust in our model, also \\textquoteleft crustal modes' appear in the spectrum}. In the axial case, they are not anymore absorbed by the continuum so they are long-living modes. The only case in which a damping of the crustal modes  can be observed is if they are in resonance with the edges of the continuum. In this case the energy leaks from the crust to the core with a velocity that depends on the strength of the magnetic field. The efficiency of the  process depends on the  magnetic field strength. Finally, we made an attempt  to resolve the puzzle of the 625 Hz QPO: we find that in our model, although the spectrum is discrete, at higher frequencies it becomes  too dense to allow an unambiguous interpretation of this high frequency mode.  In the future we would like to investigate the effect of superconductivity on the oscillations spectrum as well as non-axisymmetric oscillations. Note that a first study of non-axisymmetric  oscillations in a magnetised neutron star was made by \\cite{2010MNRAS.405..318L}, \\cite{2011MNRAS.412.1730L} (with a Newtonian model) and by \\cite{2011JPhCS.314a2081S} (with a general  relativistic model): they both found that the spectrum of non-axisymmetric oscillations has also a discrete nature."
    },
    "1112/1112.1087_arXiv.txt": {
        "abstract": "We discuss a hydrodynamical model for the dispersal of protoplanetary discs around young, low mass ($<1.5$ \\msun) stars by photoevaporation from the central object's energetic radiation, which considers the far-ultraviolet as well as the X-ray component of the radiation field. We present analytical scaling relations and derive estimates for the total mass-loss rates, as well as discussing the existence of similarity solutions for flows from primordial discs and discs with inner holes. Furthermore, we perform numerical calculations, which span a wide range of parameter space and allow us to provide accurate scalings of the mass-loss rates with the physical parameters of the systems (X-ray and FUV luminosity, stellar mass, disc mass, disc temperature and inner hole radius).  The model suggest that the X-ray component dominates the photoevaporative mass-loss rates from the inner disc. The mass-loss rates have values in the range from 10$^{-11}$ to 10$^{-7}$ \\msunyr and scale linearly with X-ray luminosity, with only a weak dependence on the other parameters explored. However, in the case of high FUV to X-ray ($L_{\\rm FUV}/L_X>100$) luminosity ratios, the FUV constricts the X-ray flow and may dominate the mass-loss.  Simulations of low mass discs with inner holes demonstrate a further stage of disc clearing, which we call `thermal sweeping'.  This process occurs when the mid-plane pressure drops to sufficiently low values. At this stage a bound, warm, X-ray heated region becomes sufficiently large and unstable, such that the remaining disc material is cleared on approximately dynamical time-scales. This process significantly reduces the time taken to clear the outer regions of the disc, resulting in an expected transition disc population that will be dominated by accreting objects, as indicated by recent observations. ",
        "introduction": "The evolution and dispersal of protoplanetary discs is an important step in the star and planet formation process. Protoplanetary discs provide the material from which planets form, implying that the initial evolution of newly formed planets is coupled to the properties of the parent disc. Moreover the disc dispersal time-scale ultimately sets the time in which (gas) planets must form.  It is observationally well established that at an age of 1~Myr most stars are surrounded by discs that are optically thick at infra-red (IR) wavelengths, which have in most cases disappeared by typical ages of 10~Myr (e.g. Haisch et al. 2001, Mamajek 2009). Optically thick discs around young stars evolve in a manner that is consistent with  `standard' viscous evolution theory, where the disc diagnostics decline with time in an approximately power law fashion (Hartmann et al. 1998). However, the final evolution from disc-bearing (primordial) to disc-less status appears to be much more rapid than viscous theory predicts (e.g Luhman et al. 2010; Muzerolle et al. 2010; Ercolano et al. 2011) and to occur over the entire range of disc radii probed by IR observations. Furthermore, given that the sub-millimetre observations show that most non-accreting stars (WTTs) are devoid of emission out to several hundreds of AU (Duvert et al. 2000, Andrews \\& Williams 2005, 2007), the clearing must be correlated across all disc radii from 1AU to 100AU.  Further clues regarding the mechanism for clearing discs are provided by objects which have been identified as `transition' discs. These are discs which show a deficit in opacity at near-IR (NIR) wavelengths compared to a standard optically thick `primordial'  disc, but are consistent with standard optically thick discs at Mid-IR/Far-IR (MIR/FIR) wavelengths; furthermore some of these objects show small NIR excesses above the photosphere (not consistent with a primordial disc) have been labelled as `pre-transition' discs (e.g. Espaillat et al. 2010; Furlan et al. 2011). Such spectral characteristics are most readily explained in terms of a hole or gap in the dust in the inner regions of the disc. However, there is no clear consensus as to the origin of such structures and the debate is complicated by the fact that they display a  wide variety of other observational characteristics: for example some show evidence for accretion (e.g. Calvet et al. 2005, Najita et al. 2007,Hughes et al. 2009) while others do not (Cieza et al. 2010; Merin et al. 2010); there is also evidence in some systems that the deficit in opacity at NIR wavelengths does not necessarily preclude significant quantities of gas and small amounts of dust in the inner disc (Espaillat et al. 2010) although this is not always the case (Calvet et al. 2002). The frequency of these { `transition/pre-transition'} discs compared to those of primordial discs suggest that approximately 10-20\\% of discs are in this stage (Strom et al. 1989, Skrutskie et al. 1990, Luhman et al. 2010, Ercolano et al. 2011a; Furlan et al. 2011). Assuming that these objects are in transition from a primordial to a fully cleared state, this observation provides  two important constraints: the transition time from primordial to cleared is approximately 10\\% of the discs total lifetime, and that the clearing process proceeds from the inside out.  While many theories have been developed in order to explain the observations of `transition' discs: grain growth (Dullemond \\& Dominiki, 2005); photophoresis (Krasuss et al. 2007); MRI-driven winds (Suzuki \\& Inutsuka, 2009), { only two mechanism have been proposed that can explain the observations of `transition' discs in the context of total (gas \\& dust) disc clearing (assuming that indeed all discs pass through a `transition' disc phase). These are} photoevaporation (e.g. Clarke et al. 2001) and planet formation (e.g. Armitage \\& Hansen, 1999). Photoevaporation and planet formation are processes which compete to remove gas from the disc, and it is still a matter for debate whether one process dominates over the other. Ercolano \\& Clarke (2010) discuss the implications of metallicity on both processes and Yasui et al. (2009) provide very preliminary observations of smaller disc frequencies in low metallicity regions in the extreme outer Galaxy, consistent with the predictions from X-ray photoevaporation models and arguing against planet formation as a main dispersal mechanism.  Like any theory, the photoevaporation model has progressed through several stages of evolution, { with the original theories only including photoevaporation without viscous disc evolution (Shu et al. 1993; Hollenbach et al. 1994; Yorke \\& Welz 1996; Richling \\& Yorke 1997), or photoevaporation from external sources (Johnstone et al. 1998; Richling \\& Yorke 1998; Richling \\& Yorke 2000; Adams et al. 2004). It was} Clarke et al. (2001) who realised that mass-loss from the surface of a disc due to heating (photoevaporation), coupled with viscous evolution results in a gap opening in the inner disc when the photoevaporation rate and accretion rate become comparable. The inner disc then drains onto the central star on its own much faster local viscous time-scale. While the basic ideas have remained identical (photoevaporation and viscous process compete until photoevaporation eventually opens a gap in the inner disc which then rapidly drains on to the central star), the actual calculations of the photoevaporation rates have progressed greatly. The original model used semi-analytic mass-loss rates due to EUV radiation calculated by Hollenbach et al. (1994), while later the photoevaporative flow was modelled hydrodynamically  by Font et al. (2004).  In these models the disc is mainly irradiated by a diffuse field of recombination photons from the atmosphere of the inner disc. Subsequently Alexander et al. (2006a) realised that, following the draining of the inner disc,   the inner-edge of the hole is instead directly exposed to the radiation field from the star; this results in higher mass-loss rates which rapidly clear the remaining disc out to large radii (Alexander et al. 2006b).  Recently, a lot of attention has been paid to other high energy radiation fields; X-rays (Alexander et al. 2004, Ercolano et al. 2008,2009, Owen et al. 2010,2011b) and FUV (Gorti \\& Hollenbach 2008,2009; Gorti et al. 2009). These calculations have yielded much higher photoevaporation rates than the original EUV rate. {These higher mass-loss rates primarily  arise from significant mass-loss at larger radii in the disc ($R>$10 AU), due to  the ability of the FUV and X-rays to heat much larger columns of gas to of order the escape temperature}. While the central star certainly produce significant fluxes of FUV, EUV and X-rays (e.g. Alexander et al. 2005, Ingleby et al. 2011) it is important to determine which, or which combination of high energy rates are responsible for setting the photoevaporation rate.  { Some of the main observational tests of photoevaporation are based on observational statistics of disc populations as a function of time}. Owen et al. (2011b) showed that the X-ray photoevaporation model is consistent with the current diagnostics of disc evolution, and made some predictions about how these diagnostics should correlate with X-ray luminosity. Furthermore, Drake et al. (2009) reported that accretion rate measurements and X-ray observations in  Orion, point to a process of X-ray driven `photoevaporation-starved accretion', where higher mass-loss rates prevent material accreting onto the central star, leading to lower accretion rates. This process was theoretically demonstrated by Owen et al. (2011b) and is also likely to  explain the systematic discrepancy between the X-ray luminosities of  accreting and non-accreting young stars. Unlike the previous EUV models (e.g. Alexander et al. 2006b; Alexander \\& Armitage, 2009), the X-ray photoevaporation model can account for a large fraction of the observed `transition' discs. In particular, photoevaporation can account for the population of small inner holes $R<20$AU, those with low accretion rates $\\dot{M}_*<10^{-8}$\\msunyr and those with gas inside the inner-hole. However, `transition' discs with large inner holes and large accretion rates are likely to be created through a different mechanism (e.g. Espaillat et al. 2010).  { The first direct evidence and best test of individual photoevaporation models is the detection of a blue-shifted 12.8$\\mu$m NeII line from the surface of the disc around TW Hya (Pascucci \\& Sterzik 2009), which is consistent with either  an EUV or X-ray driven photoevaporative wind (Pascucci et al. 2011)}. Moreover, Ercolano \\& Owen (2010) recently showed that  the previously unexplained low velocity component of the OI 6300\\AA, - emission  blue-shifted by $\\sim5$km s$^{-1}$ detected by Hartigan et al. 1995 around most primordial discs - is naturally produced in an X-ray driven photoevaporative flow. The observed luminosity is too high to be consistent with an EUV wind and the blue-shift is too large to be consistent with an FUV wind. { Currently, the only well resolved OI 6300\\AA~ line is detected from the accreting `transition' disc around TW Hya, where the line presents with no blue-shift (Pascucci et al. 2011).  Unfortunately, hydrodynamic models of the inner regions of photoevaporating transition discs during the accreting phase do not exist for any of the photoevaporation models (only transition discs with cleared inner gas discs were considered by Ercolano \\& Owen 2010) and useful comparisons must wait until these calculations are completed. } Therefore, pending these calculations, and the availability of high resolution spectroscopic observations of many objects, direct comparisons of predicted to observed OI line profiles are not possible.  While observations are providing a fruitful avenue for direct comparisons of photoevaporation models, in this work we develop a theoretical framework in which we can analyse the combined effects off all three heating mechanisms (EUV,FUV and X-rays). The most relevant scale in any photoevaporation model is the radius at which a gas of a given temperature ($T_{\\textrm{gas}}$) becomes unbound from the central star and is given by\\footnote{Throughout this work we use $\\{R,\\phi,z\\}$ to represent the cylindrical co-ordinate system and $\\{r,\\theta,\\varphi\\}$ to represent the spherical co-ordinate system, in all cases centred on the star.}: \\begin{equation} r_g=8.9 \\textrm{AU} \\left(\\frac{T_\\textrm{gas}}{10^{4} \\textrm{ K}}\\right)^{-1}\\left(\\frac{M_*}{1\\textrm{ \\msun}}\\right) \\end{equation} For the EUV photoevaporation model this becomes a fixed radius (for a given disc), since EUV heating gives rise to essentially isothermal gas at $10^4$K. For the X-ray and FUV model it is better to think of this radius as an escape temperature scaled length, allowing the consideration of different radial scales around different mass stars. However, in order to be consistent with the previous EUV only literature { (Hollenbach et al. 1994; Clarke et al. 2001; Font et al. 2004; Alexander et al. 2006a)}, we separately fix $R_g=8.9(1/1\\textrm{\\msun})$AU. Furthermore, as we will discuss in detail  in Section 2, this radius has an important physical significance as it is a measure of the  radius at which a flow at temperature $T_{\\textrm{gas}}$ can go through the sonic surface. For X-ray heating in the range $T_{\\textrm{gas}}\\sim1000-10000$K, $r_g\\sim 1-10$AU for a very low mass star (0.1\\msun) and $r_g\\sim10-100$AU for a solar type star. For FUV heating,  $T_\\textrm{gas}\\apprle 1000$K, $r_g\\apprge10$AU for a very low mass star and $r_g\\apprge100$AU for solar type stars.  In this work,  we lay out the theoretical foundation for photoevaporation, deriving the scaling with X-ray luminosity and stellar mass in Sections~2. In Section~3 we discuss the likely interplay between the effects of combined EUV, FUV and X-ray heating on the calculation of photoevaporation rates. Section~4 presents the results of radiation-hydrodynamic calculations that test the theoretical foundations laid out in the  Section 2, while Section ~5 contains some simple numerical experiments that address the discussion of combined FUV and X-ray heating contained in Section 3.  Section~6 describes the final clearing of the outer disc by `thermal sweeping'. In Section~7 we describe the implications of photoevaporation  on disc evolution and dispersal. In Section~8 we discuss the presented results and draw our conclusions in Section~9.   In Appendix~A we discuss the possibility of self-similarity in X-ray heated flows and in Appendix~B we provide numerical fits to our photoevaporation calculations for use by the community.     \\section {The theory of X-ray photoevaporation}  The problem of thermally driven winds from discs has been studied analytically in the past (Begelman et al. 1983; Liffman 2003; Adams et al. 2004). While a full analytic solution to the problem has not been obtained in even the most simplistic thermodynamic constructions, considerable progress can still be made. In the Appendix we discuss the existence of similarity solutions with regard to the problem of X-ray photoevaporation. The results from the analysis are as follows: \\begin{itemize} \\item For primordial discs, the total mass-loss rates scale linearly with X-ray luminosity and are independent of stellar mass. \\item For primordial discs, the density at a given value of $\\mathbf{r/R_g}$ scales with X-ray luminosity and stellar mass as $n(L_X,M_*)\\propto L_X\\,M_*^{-2}$ for primordial discs.  \\item For discs with inner holes of fixed radius, the total mass-loss scales linearly with the X-ray luminosity.  \\item Discs with inner holes with a range of radii are not strictly mutually self-similar: see section 4.3.1 for an approximate self-similarity argument that can be used to understand the numerical results in this case.  \\end{itemize}  In this section, we derive the scalings listed above using a series of simple analytic arguments. Since the mass-flux is conserved in any steady-state photoevaporative flow we are free to evaluate this quantity at any point in the flow. Therefore, choosing to do this at the sonic surface the total mass-loss rate is given by:  \\begin{equation} \\dot{M}_w=\\int_Sn_s(\\mathbf{r_s})c_s(\\mathbf{r_s})\\mathbf{l}(\\mathbf{r_s})\\cdot\\dd\\mathbf{S} \\end{equation} where $S$ is the sonic surface, $n_s$ \\& $c_s$ are the density and sound speed at the sonic surface respectfully, and $\\mathbf{l}$ is the unit vector along a streamline. In the case of a flow from a disc, this can be approximately written as: \\begin{equation} \\dot{M}_w=2\\int_0^\\infty 2\\pi R n_s(R)c_s(R)\\dd R\\label{eqn:m_tot} \\end{equation} where we note that the true mass loss flux will differ from this by a factor of order unity due to the fact that the flow is not purely vertical at the sonic surface. Therefore, provided we can determine the flow properties at the sonic surface we can estimate the total mass-loss rates.  \\subsection{The Sonic Surface in X-ray Heated Flows} The sonic point is a critical point of the steady-state momentum equation when $u=c_s$, and physically arises as a consequence of wave-steeping requirements, namely that no strong shock forms (Landau \\& Lifshitz, 1987).    Several authors have derived the conditions on the sound speed at the sonic surface for a flow emerging from a disc (e.g. Begelman et al. 1983; Adams et al. 2004; also shown in Appendix A for completeness). Under the conditions expected for disc winds (sub-sonic launching, flow time-scale $>$ thermal time-scale, height of the sonic surface ($z_s$) $\\ngg\\,\\, R$) the sound speed at the sonic surface should be to order unity the Parker value at each cylindrical radius, namely:  \\begin{equation} c_{s}^2 \\approx {{GM_*}\\over{2R}}\\label{eqn:cs} \\end{equation} Therefore, we can {\\it a priori} determine the sound speed, and hence the temperature at the sonic surface.  \\subsubsection{Density-Sound-Speed Coupling in X-ray Heated Flows} If we can specify both the density and sound speed at the sonic surface then we have determined the mass-flux in the entire flow as mass-flux is conserved. We argued above that in the case of disc winds, the sound speed at the sonic surface is roughly given by the Parker value (Equation~\\ref{eqn:cs}). This means that the gas temperature at the sonic surface is also simply given as a function of cylindrical radii.  It is well known that the temperature of X-ray heated gas can be described in terms of the local ionization parameter (e.g. Tarter et al. 1969), given by: \\begin{equation} \\xi=\\frac{4\\pi F_X}{n} \\end{equation} where $F_X$ is the X-ray flux ($L_X/4\\pi r^2$). Ercolano et al. (2009ab) \\& Owen et al. (2010) found that the temperature of gas within a column density to the central star of $10^{22}$ cm$^{-2}$ (this being the penetration depth of 1~KeV photons\\footnote{see discussion in Section 4.2.2 for implications of attenuation which is neglected here}) to the central star can be described in terms of optically thin X-ray heating. Thus, adopting the assumption of optically thin X-ray heating  the ionization parameter becomes $\\xi=L_X/n r^2$, and the gas temperature is specified in terms of local variables only, where the form of this profile $T=f(\\xi)$ is shown in Figure~\\ref{fig:ionT}. \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{ionT} \\caption{Fit to temperature-ionization parameter calculated by Owen et al. (2010).}\\label{fig:ionT} \\end{figure}  At the sonic surface (provided $z_s\\ngg R$ so $R\\sim r$) we may write: \\begin{equation} c_s^2=\\frac{k_Bf\\left(L_X/n_s R^2\\right)}{\\mu m_h}=\\frac{GM_*}{2R}\\label{eqn:den1} \\end{equation} Therefore, we may re-arrange Equation~\\ref{eqn:den1} to give an expression for the density at the sonic surface\\footnote{where $f^{-1}\\circ f =1$ and $f^{-1}\\equiv \\hspace{-0.28cm} /\\hspace{0.2cm} 1/f$}: \\begin{equation} n_s=\\frac{L_X}{R^2}\\left[f^{-1}\\left(\\frac{GM_*\\mu m_h}{2k_BR}\\right)\\right]^{-1}\\label{eqn:den_s} \\end{equation} One may also write this at fixed $R/R_g$ to obtain the mass-scaling of the density, namely $n\\propto M_*^{-2}$. Thus, provided we know the form of $f(\\xi)$, the density at the sonic surface is also fixed {\\it a priori} through the temperature-ionization parameter relation, which provides a density-sound-speed coupling. \\subsubsection{Consequences of Density-Sound-Speed Coupling} The density-sound-speed coupling found at the sonic surface in optically thin  heated X-ray heated flows provides some very powerful constraints on the flow properties of X-ray photoevaporative winds, and implies that the structure and temperature of the underlying disc has little effect on the resulting flow.  This arises since in any steady state flow mass-flux must be conserved; therefore, in this case since we are able to described the mass-flux at the sonic surface {\\it a priori} with reference to the X-ray physics alone, we have  determined the mass-flux in the {\\it entire} flow without reference to the disc properties.  This picture is very different to pure EUV (Hollenbach et al.  1994) or FUV (Gorti \\& Hollenbach 2008,2009) flows where the underlying density structure in the launch region is crucial to the mass-loading of the flow. Unlike EUV/FUV flows where the flow will adjust to changes in the underlying disc's structure, an X-ray flow will adjust the underlying disc structure to `feed' the flow at the required rate, { provided the sonic surface remains optically thin to the X-rays responsible for the heating}.  \\subsection{Estimates of the total mass-loss rate} Armed with the knowledge that it is the X-ray physics, and not the disc structure that sets the overall mass-loss rates we are free to integrate Equation~\\ref{eqn:m_tot} over the sonic surface to estimate the total mass-loss rates. We can use Equations~\\ref{eqn:cs} \\& \\ref{eqn:den_s} to substitute for $n_s$ and $c_s$ in the integrand. Furthermore, if we change the variable of integration in Equation~\\ref{eqn:m_tot} from cylindrical radius $R$ to escape temperature we find: \\begin{eqnarray} \\dot{M}_w(M_*,L_X)\\!\\!\\!\\!\\!&=&\\!\\!\\!\\!L_X\\sqrt{\\frac{8k_B\\pi^2}{\\mu m_h}}\\int_0^\\infty\\!\\!\\!\\dd\\Tesc\\, \\frac{\\Tesc^{-1/2}}{ f^{-1}(\\Tesc/2)}\\label{eqn:inttesc}\\\\ &\\approx&8\\times 10^{-9}\\left(\\frac{L_X}{1\\times 10^{30}\\,\\textrm{erg s}^{-1}}\\right) \\textrm{M}_\\odot\\,\\textrm{yr}^{-1}\\label{eqn:tot} \\end{eqnarray} (where the second result comes from evaluating the integral out to an escape temperature of 1000K, or a radius of $\\OO (100)$AU for a solar mass star). This mass-loss rate  is in good agreement with the numerical result of $6.7\\times 10^{-9}$M$_\\odot\\,$yr$^{-1}$ obtained from the hydrodynamic simulations of Owen et al. (2011b) and is also consistent with the approximately linear scaling between mass-loss rate and X-ray luminosity found in the simulations. Furthermore, the results indicates there is no explicit dependence between stellar mass and mass-loss rate (although the $M_*-L_X$ scaling introduces a strong positive implicit scaling). We also derive this scalings in more detail through the existence of similarity solutions  presented\\footnote{ Note that the scalings presented in Appendix~A do not rely on the approximate expression for the temperature at the sonic surface contained in Equation\\ref{eqn:cs}.} in Appendix A.  \\subsection{Importance of Numerical Results}\\label{sec:Theory}  In Sections 4 and 5 we present hydrodynamical simulations that have been designed to test and refine the ideas set out above. However, the role of the numerical simulations is not merely one of  validation: although Equation~\\ref{eqn:tot} is of great utility in predicting the integrated mass-loss from the system (and can be used to derive the mass loss rate per unit area {\\it at the sonic surface}), it does not  - without knowledge of the detailed streamline topology below the sonic surface - allow one to calculate the mass loss rate per unit area from the disc itself ($\\dot \\Sigma_w(R)$). { For example the form of the integrand in Equation~\\ref{eqn:tot} suggests that the mass-loss is dominated in the lower temperature regime (2000-5000K), but these stream-lines connect the sonic surface to the disc surface in { a non-uniform way, such that there is not a one to one correspondence between the mass-flux profile at the sonic surface and the mass-flux profile from the disc.} This results in a mass-loss profile that peaks at smaller radii in the disc, rather than the large radius suggested by the calculations at the sonic surface.} Although this may appear to be a detail, it turns out that, when photoevaporation is combined with secular disc evolution, the pattern of disc clearing is rather sensitive to the form of $\\dot \\Sigma_w(R)$. { In particular, it is the form of $\\dot{\\Sigma}_w(R)$ that determines } where the gap first opens in the disc, the extent to which the system undergoes a period of `photoevaporation-starved accretion' prior to gap opening and hence the surface density profile of the disc outside the gap. This latter then determines the rate at  which the disc is evaporated once the inner disc is completely cleared out. This is of obvious importance if one wants to compare the predictions of photoevaporation theory with the observed incidence of transition discs with holes of different sizes (see Section 8.1). ",
        "conclusions": "In this work we have outlined the basic properties of the  photoevaporation of discs around low mass stars by the X-rays from  the central  star. In this regime, the photoevaporation rates are set by the properties in the X-ray flow and are fairly insensitive to variations in disc structure and extra heating sources e.g. FUV radiation. At the most basic level this indicates factors that determine discs' lifetimes and evolution may stem from the stars' intrinsic X-ray luminosities. Our main conclusions are summarised below:   \\begin{enumerate} \\item The photoevaporative mass-loss rates are primarily set by the X-ray heating alone, as the sonic surface always occurs within the X-ray heated region. Extra UV heating provides little effect on the determined structure of the X-ray wind and its associated sonic surface geometry  and  mass-loss rate  if $L_X/L_{\\rm FUV}>0.01$. \\item We have derived from first principles the scaling of photoevaporative mass loss from primordial discs with X-ray luminosity and stellar mass and have also derived analytical estimates of the total mass loss rates. The mass-loss rates scale approximately linearly with X-ray luminosity and have no explicit dependence on stellar mass. These scalings and absolute values have been confirmed through a large suite of numerical simulations.  \\item We do not predict a strong dependence of disc lifetime on stellar mass for stars less massive than $1 M_\\odot$, in agreement with observations. Higher mass stars have higher X-ray luminosities but also have more massive discs to be dispersed. The resulting time-scale for disc dispersal would be a strong function of stellar mass only in the unlikely case that the viscous properties of discs were an extremely strong function of stellar mass.  \\item We have explained why, in the case of discs with inner holes, the photoevaporative mass-loss rates scale linearly with X-ray luminosity and is independent of inner hole size and have calculated numerical models to calibrate the mass-loss rates.  \\item We find that once the mid-plane pressure of a disc with an inner hole drops below a critical value (given by Equations~15 \\& 16) the remaining  disc is cleared rapidly (on a roughly dynamical  time-scale). This process (which we have termed `thermal sweeping') sets in once the layer of bound X-ray gas on the inside of the disc rim becomes thick (i.e. comparable with its vertical scale height) since at this stage the heated gas escapes preferentially normal to the disc plane and increases the exposure of the disc rim. Such rapid dispersal removes the possibility of  the long lived `relic' discs hypothesised in Owen et al. (2011b). We find that thermal sweeping sets in at around $10-15$ A.U. in the case of low mass stars ($0.1 M_\\odot$) and we roughly estimate that the corresponding radius is $\\apprle 50$ AU for higher mass T Tauri stars ($\\sim 0.7 M_\\odot$).  \\item We find that when we take into account the spread in masses and X-ray luminosities of T Tauri stars, we can produce transition discs that span a wide range of properties (hole radii and accretion rate onto the star). In this way we can account for a large fraction ($\\apprge50\\%$) of the observed transition discs. We cannot, however, account for large holes with high accretion rates (see Figure \\ref{fig:trans}).  \\item Photoevaporation and `thermal sweeping' ultimately destroy the final disc, but it is still accretion onto the central star that removes most of the original disc mass ($\\sim 10 \\%$ M$_*$): in most cases only $\\sim 10-20\\%$ of the original disc mass lost through in the processes of photoevaporation or thermal sweeping.  \\end{enumerate}"
    },
    "1112/1112.1614_arXiv.txt": {
        "abstract": "We present $BV$ photometry of four Sagittarius dwarf spheroidal galaxy globular clusters: Arp~2, NGC~5634, Palomar~12, and Terzan~8, obtained with the Danish telescope at ESO-La Silla.  We measure the structural parameters of the clusters using a King profile fitting, obtaining the first reliable measurements of the tidal radius of Arp~2 and Terzan~8. These two clusters are remarkably extended and with low concentrations; with a concentration of only $c=0.41 \\pm 0.02$, Terzan~8 is less concentrated than any cluster in our Galaxy.  Blue stragglers are identified in the four clusters, and their spatial distribution is compared to those of horizontal branch and red giant branch stars. The blue straggler properties do not provide evidence of mass segregation in Terzan~8, while Arp~2 probably shares the same status, although with less confidence. In the case of NGC~5634 and Palomar~12, blue stragglers are significantly less populous, and their analysis suggests that the two clusters have probably undergone mass segregation. ",
        "introduction": "\\label{sec:intro}  Stars known as blue stragglers (BSs) occupy a position on the color-magnitude diagram (CMD) of stellar clusters at a higher luminosity and temperature than the main-sequence (MS) turnoff point, appearing as a sparse prolongation of the MS \\citep[e.g.,][]{sandage53,burbidge58}. This position implies that these stars have a larger mass than the one expected from their parent cluster evolution. Two different, but non-exclusive \\citep[e.g.,][]{ferraro09}, scenarios have been postulated for BS formation, namely mass exchange or the merger of close binaries \\citep[e.g.,][]{mccrea64,carney01} and the direct collision of stars \\citep[e.g.,][]{hills76,leonard89,davies94,leonard95}, which is naturally enhanced in crowded environments.  The formation of BSs is then intimately related to the structure and dynamical status of their surroundings. In this respect the Sagittarius dwarf spheroidal galaxy (hereafter Sgr), due to its ongoing disruption into our Galaxy, offers different environments where the BS phenomenon can be studied.  Since the discovery of Sgr \\citep{ibata94}, a plethora of globular clusters have been putatively associated with it. Four of them, M54 (NGC~6715), Arp~2, Terzan 7 and Terzan~8, lying within the body of Sgr, are typically considered as bona fide members \\citep{dacosta95}. Using several arguments, including dynamical modeling, the position over the stream of stars from the disrupting Sgr and abundance ratios, several other globular clusters have been considered as (former) members of the Sgr cluster system that have been stripped away due to tidal interactions with the Galaxy \\citep{dinescu00,palma02,bellazzini03,cohen04,carraro07,law10,forbes10}.  In the present paper we study the BS population in four of the Sgr globular clusters: the bona fide members Arp~2 and Terzan~8, and the stripped NGC~5634 and Palomar~12. These clusters cover a range of ages, metallicities, galactocentric distances and structural properties, which however have heretofore been scarcely studied with appropriate data.  Arp~2 (${\\rm [Fe/H]}=-1.83$; \\citealt{mottini08}) is a particularly remarkable cluster, since despite several attempts, there is still no agreement whether it has a somewhat younger age \\citep{buonanno95,marin09} or the same as the bulk of the Galactic globular clusters \\citep{layden00,dotter10}. A young age would clash with its mostly blue horizontal branch \\citep[HB;][]{buonanno95}. Terzan~8, also lying in the Sgr body, is the most metal-poor (${\\rm [Fe/H]}=-2.34$; \\citealt{mottini08}), and also perhaps the oldest, cluster in Sgr \\citep{montegriffo98,dotter10}. NGC~5634 is much like Terzan~8 in terms of age and metallicity \\citep{bellazzini02} and has been associated to Sgr given its position and radial velocity \\citep{bellazzini02,law10}. The final cluster in this study, Palomar~12, has been associated with Sgr due to its proper motion \\citep{dinescu00} and anomalous abundances \\citep{cohen04}. In addition, traces of the Sgr stream have been found in its surroundings \\citep{martinez02}. Palomar~12 is a relatively young \\citep[$\\sim$\\,8--9\\,Gyr;][]{stetson89,rosenberg98} and metal-rich (${\\rm [Fe/H]}=-1.0$, \\citealt{brown97}) cluster.  The paper is organized as follows: we present our observations and reduction procedures in Sect.~\\ref{sec:obs}; the measurement of the structural parameters of the clusters is given in Sect.~\\ref{sec:structure}. We present the results on the BS spatial distribution in Sect.~\\ref{sec:stragglers}. Finally, our main results concerning the BS spatial distribution and dynamical status of the clusters in our sample are discussed in Sect.~\\ref{sec:discussion}, and a brief summary is presented in Sect.~\\ref{sec:summary}.  \\section {Observations and data reduction} \\label{sec:obs}  \\begin{figure} \\includegraphics[width=0.49\\textwidth]{sgr_fig1.ps} \\caption{Medianed images produced with MONTAGE2. These images are used to establish the initial positions of the stars that are introduced to ALLFRAME. All images have approximately $14\\arcmin\\times14\\arcmin$. North is up and East to the left.} \\label{fig:clusters} \\end{figure}   \\begin{figure} \\includegraphics[width=0.49\\textwidth]{sgr_fig2.ps} \\caption{Level of completeness for the $V$ photometry in each cluster. Stars were counted in 0.5\\,magnitude bins. The dashed blue line indicates the completeness inside the inner 2\\arcmin of each cluster, while the red solid line indicates the completeness outside this radius. NGC~5634 is divided into three bins with $30\\arcsec<R<60\\arcsec$ (dotted line), $60\\arcsec<R<120\\arcsec$ (dashed line), and $R>120\\arcsec$ (solid line).} \\label{fig:completeness} \\end{figure}   Bessel $B$ and $V$ images were obtained using the Danish Faint Object Spectrograph and Camera (DFOSC) mounted on the 1.54\\,m Danish Telescope at ESO La Silla between June 27 and 30, 2003. A series of short exposures were obtained ranging from 200 to 600 seconds in the $B$ band and 100 to 300\\,seconds using the $V$ filter. The seeing quality was good during the four nights, with FWHM in the range 0.8\\arcsec--1.3\\arcsec. The DFOSC field of view is 13.5\\arcmin$\\times$13.5\\arcmin\\, with a scale of $0\\farcs39$ pixel$^{-1}$.  \\begin{table*} \\centering \\caption{Exposure times and seeing conditions for the images used.} \\label{table:seeing} \\begin{tabular}{@{}lllc@{}} \\hline Cluster name & $B$ exp. time &      $V$ exp. time &  Seeing range\\\\ &[seconds]&[seconds]&[$\\arcsec$]\\\\ \\hline NGC~5634& 6$\\times$250+1$\\times$300+4$\\times$400+&6$\\times$100+1$\\times$120+1$\\times$150+ & $B$: 0.9--1.1\\\\ & 1$\\times$450+2$\\times$500+1$\\times$600&4$\\times$200+2$\\times$250+1$\\times$300& $V$: 0.8--0.9\\\\ Arp~2 & 6$\\times$250+9$\\times$400&5$\\times$120+1$\\times$100+9$\\times$200& $B$: 0.9--1.0\\\\ & & &$V$: 0.8--0.9\\\\ Terzan~8 & 5$\\times$200+4$\\times$250+6$\\times$400&5$\\times$80+4$\\times$120+6$\\times$200&$B$: 0.9--1.0\\\\ & & &$V$: 0.8--1.0\\\\ Palomar~12& 2$\\times$250+1$\\times$300+7$\\times$400+&1$\\times$120+1$\\times$150+9$\\times$200+&$B$: 0.9--1.0\\\\ & 1$\\times$450+4$\\times$500& 2$\\times$220+2$\\times$250&$V$: 0.9--1.0\\\\ \\hline \\end{tabular} \\end{table*}  Reduction of the images, including bias and overscan subtraction and flat-fielding, was done using standard tasks within IRAF.\\footnote{IRAF is distributed by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy (AURA) under cooperative agreement with the National Science Foundation.}  Point-spread function (PSF) photometry for all the images was done using DAOPHOT/ALLSTAR \\citep{stetson87}, where typically $\\sim$50 bright and isolated stars were selected in the NGC~5634 and Pal~12 fields and $\\sim$100 stars in the more populated Arp~2 and Terzan~8 fields (see Figure~\\ref{fig:clusters}) in order to derive the PSF.  For the cluster photometry, the best 15 images in each band for each cluster were selected. Details of exposure time and image quality can be seen in Table \\ref{table:seeing}. The entire set of images has been used for a reassessment of the clusters' stellar variability, as will be presented in a forthcoming paper, extending the results from \\citet{salinas05}.  Transformation equations in pixel space for the positions of the stars between the images were derived using DAOMATCH/DAOMASTER \\citep{stetson93}. This step requires the selection of a ``master image'' to which transformation equations are anchored to. The master image for each cluster was selected to be the one with the best seeing and longer exposure. The subset of images for each cluster from both filters was combined together using MONTAGE2 \\citep{stetson94}, which weighs each image according to its seeing and total flux (after sky subtraction), hence making sharper and deeper images more preponderant in the final combined frame. From these combined images, a master list of stars was generated by applying iteratively DAOPHOT/ALLSTAR. This list of stars together with the images were then fed into ALLFRAME \\citep{stetson94}, which fits simultaneously the PSF to all stars in all the available images. Mean magnitudes and errors for each band from the output ALLFRAME photometry were derived again using DAOMATCH/DAOMASTER, which match the star coordinates in the different frames and apply a zero-point correction to the magnitudes respect to the master image.  \\begin{figure*} \\includegraphics[width=\\textwidth]{sgr_fig3.eps} \\caption{Color-magnitude diagrams of the four clusters. Field contamination can be seen as most severe in the case of Arp~2.} \\label{fig:cmds} \\end{figure*}  \\citet{landolt92} standard stars were observed in the fields Mark~A, G26 and PG~1525 during the first, second and fourth nights of observations. These fields were observed at a range of airmasses, $1\\!<\\!X\\!<\\!1.4$, bracketing the science exposures. Transformation equations to the standard system were derived using IRAF/PHOTCAL. Aperture corrections, also calculated within the IRAF/PHOTCAL environment, were found to be always below $-0.02$\\,mag for both filters in all clusters.  Globular clusters studies using ALLFRAME photometry usually deal with large image sets comprising hundreds if not thousands of images \\citep[e.g.,][]{walker11}. Artificial stars experiments hence become prohibitive and completeness has to be estimated using the behaviour of the faint end of the observed luminosity function. In our case, the modest amount of images allows us to study the completeness directly with artificial stars experiments. In our procedure we included 1000 artificial stars in random positions for each cluster, except for Palomar~12 where 600 stars were added. The magnitudes of these artificial stars were between $19.5<V<25$, i.e. roughly from the BS magnitude level to the level of the faintest stars detected with ALLFRAME. These artificial stars were initially included in each cluster master image and then into the rest of the image set by using the inverse of the transformation equations derived with DAOMASTER. In this way, artificial stars are placed in the exact same relative position inside each frame. For each cluster 10 new image sets were generated in this way, each having 15 $B$ and 15 $V$ images like the original set. Each new image set was photometered with ALLFRAME in the same way as the original image set, i.e. the one with no additional stars added.  The level of completeness for each cluster can be seen in Figure~\\ref{fig:completeness}. In Arp~2, Palomar~12 and Terzan~8, which have uncrowded fields, we analyzed the completeness in two areas. Beyond 2\\arcmin\\, the photometry of these three clusters can be considered as 100\\% complete to around $V\\!\\sim\\!23$ (solid line), while inside this radius, 100\\% is reached down to $V\\!\\sim\\!22.5$. For NGC~5634, which provides us with the most crowded central fields (see Figure~\\ref{fig:clusters}), we made the count of artificial stars in three concentric annuli, with $30\\arcsec<R<60\\arcsec$, $60\\arcsec<R<120\\arcsec$, and $R>120\\arcsec$; the level of completeness in these areas is depicted in Figure~\\ref{fig:completeness} as the dotted, dashed and solid lines, respectively. The incompleteness in the inner parts is especially severe, with 100\\% being reached only down to $V\\sim20.5$) but in any case not affecting the BS magnitude level.  In conclusion, the photometry can be considered as complete in the entire magnitude range where BSs are present in the four clusters, with the exception of the inner 30\\arcsec of NGC~5634, where at $V\\sim19.5$ the completeness level is only about 50\\%. This central area is then excluded from our analysis. Final CMDs of the four clusters can be seen in Fig.~\\ref{fig:cmds}.  \\begin{figure*} \\centering \\includegraphics[width=0.9\\textwidth,angle=270]{sgr_fig4.eps} \\caption{Surface stellar density profile for each of our four clusters. The solid lines show the best-fitting King profiles, while the lower panels show the fit residuals.} \\label{fig:structure} \\end{figure*} ",
        "conclusions": "\\label{sec:discussion}  Our analysis of the BS distribution shows that Terzan~8, and probably Arp~2, have not relaxed yet. This is not altogether surprising, since relaxation time depends on the size of the dynamical system and, as seen in Sect.~\\ref{sec:structure}, these two clusters have large cores. We calculate the central relaxation times of the four clusters using Eq. 10 from \\citet{djorgovski93}, assuming a mean stellar mass of $1/3\\,M_{\\odot}$, adopting cluster distances from \\citet{harris96} and taking the core radii values calculated in Sect.~\\ref{sec:structure}. Results can be seen in the last column of Table~\\ref{table:king}. Terzan~8 has a central relaxation time of $t_\\mathrm{crt}=16.6$~Gyr~-- significantly larger than the cluster's age. Arp~2 has $t_\\mathrm{crt}=7.9$~Gyr. This is a factor $\\sim1.7$ shorter than the age of the oldest clusters in the Galaxy, and indicates that if Arp~2 were as old as them, mass segregation should already be visible. Since according to the A-D test the BS distribution and the specific frequency profile show at most weak evidence of segregation, this could imply that Arp~2 is indeed a slightly young ($\\sim 9$~Gyr) cluster, as first postulated by \\citet{buonanno95}.  Arp~2 and Terzan~8 are among the globular clusters with the lowest concentration measured. Since in these environments direct collisions and binary disruption are expected to be negligible, the BS properties should be more similar to the field than to more massive and concentrated clusters. Nevertheless, the specific frequency of BSs in the local halo is found to be significantly higher, with $S_{\\mathrm{BS}}\\sim 4$ \\citep{preston00}, much higher than in the lowest-concentration clusters, suggesting that the specific frequency is primarily driven by the total luminosity of the cluster \\citep{sandquist05}, instead of concentration. Notably, the BSs specific frequency in these clusters is within the range estimated for the Galactic bulge, namely $0.31<S_{\\mathrm{BS}}<1.23$ \\citep{clarkson11}.     We have presented $BV$ photometry of four globular clusters associated with the Sagittarius dwarf: Arp~2, NGC~5634, Palomar~12 and Terzan~8. Arp~2 and Terzan~8 have remarkably low concentrations and with half-light radii of $\\sim$15 pc, not unlike the ``extended star clusters'' now commonly found in Local Group galaxies \\citep[e.g.][and references therein]{hwang11}. Terzan~8 presents clear evidence of a non-segregated BS population, indicating its dynamical youth and joining in this condition much more massive clusters as $\\omega$~Cen \\citep{ferraro06} and NGC~2419 \\citep{dalessandro08b}, along with the seemingly dissolving Palomar~14 \\citep{beccari11,sollima11}. Although not as strongly as Terzan~8, Arp 2 also shows at best weak evidence of mass segregation, supporting the idea that it is a relatively young globular cluster, in spite of its predominantly blue HB morphology."
    },
    "1112/1112.0518_arXiv.txt": {
        "abstract": "We present infrared spectra of the binary \\XX\\ obtained with the Infrared Spectrograph on the \\sirtf. The data show some evidence for the presence of solid C$_{60}$  -- the first detection of C$_{60}$ in the solid phase -- together with the well-known ``Unidentified Infrared'' emission features. We suggest that, in the case of \\XX, the C$_{60}$ is located close to the hot component, and that in general it is preferentially excited by stars having effective temperatures in the range $15\\,000-30\\,000$\\,K. C$_{60}$ may be common in circumstellar environments, but un-noticed in the absence of a suitable exciting source. ",
        "introduction": "The possible existence of buckminsterfullerene (C$_{60}$) in astrophysical environments has long been suggested \\citep{kroto}, but only recently has observational evidence for emission from C$_{60}$ in the gas phase been forthcoming. Gas phase C$_{60}$ has now been detected in the environments of young planetary nebulae \\citep{cami,zhang}, RCB stars \\citep*{garcia} and in the reflection nebulae NGC\\,2023 and NGC\\,7023 that are illuminated by B~stars \\citep{sellgren}. In the case of the low-excitation planetary nebula Tc~1, \\cite{cami} argue that the C$_{60}$ (and C$_{70}$) molecules are attached to the surfaces of cooler carbonaceous grains. Many of the objects displaying C$_{60}$ also have strong ``Unidentified Infrared'' (UIR) features.  The formation of C$_{60}$ and other fullerenes in terrestrial laboratories usually requires a hydrogen-deficient environment, and this seems to be consistent with their presence in the environments of (evolved) H-deficient carbon stars \\citep{cami,garcia}. However the detection of C$_{60}$ in the reflection nebulae NGC\\,2023 and NGC\\,7023 \\citep{sellgren} indicates that fullerene formation is possible in young (H-rich) environments. UIR features, as well as ``Extended Red Emission'' attributed, among other hypotheses, to small -- possibly ionised -- hydrocarbon molecules, are seen in the environment of NGC\\,7023 \\citep{berne,sellgren}.  We report here the possible detection of solid phase C$_{60}$, in the environment of the peculiar binary \\XX, observed with the \\sirtf\\ \\citep{wernera,gehrz}. ",
        "conclusions": "We have reported the possible detection of solid-phase C$_{60}$ in the environment of the peculiar binary \\XX. Contrary to previous work we conclude that the hot star is a B subdwarf that is surrounded by an ionised shell and a C$_{60}$-bearing shell, most likely in the form of a disc. Variations in the optical depth of the latter results in variations in the excitation of H lines.  We will present a detailed discussion of the \\XX\\ system and its environment in a forthcoming paper."
    },
    "1112/1112.5337.txt": {
        "abstract": "According to the standard $\\Lambda$CDM model, the accelerated expansion of the Universe will go on forever. Motivated by recent observational results, we explore the possibility of a finite phase of acceleration which asymptotically approaches another period of decelerated expansion. Extending an earlier study on a corresponding homogeneous and isotropic dynamics, in which interactions between dark matter and dark energy are crucial, the present paper also investigates the dynamics of the matter perturbations both on the Newtonian and General Relativistic (GR)  levels and quantifies the potential relevance of perturbations of the dark-energy component. In the background, the model is tested against the Supernova type Ia (SNIa) data of the Constitution set and on the perturbative level against growth rate data, among them those of the WiggleZ survey, and the data of the 2dFGRS project. Our results indicate that a transient phase of accelerated expansion is not excluded by current observations. ",
        "introduction": "It is generally believed that our presently observable Universe is dynamically dominated by a dark sector which is composed of a dark-energy component with a large negative pressure and pressureless dark matter. The physical nature of both these components remains a mystery, notwithstanding the intense research activity in the field since the discovery of the accelerated expansion of the Universe in \\cite{riessperl}. Direct and indirect support for this result has been accumulating over the past years. This comprises further results from advanced data sets for supernovas of type Ia (SNIa) as well as results from studies of the large scale structure \\cite{lss}, cosmic microwave background \\cite{cmb}, the integrated Sachs--Wolfe effect \\cite{isw}, baryonic acoustic oscillations \\cite{eisenstein} and  gravitational lensing \\cite{weakl}. By now, there exists a standard model, the $\\Lambda$CDM model, which, \\textit{grosso modo}, is compatible with the cosmological data. (Notice, however, that there is an ongoing discussions, see, e.g., \\cite{nieuwenh}, on apparent shortcomings of this model). Nevertheless, because of the cosmological constant problem in its different facets, including the coincidence problem, a still growing number of competing models has been developed over the last years, most of them ``dynamizing\" the cosmological constant or even generalizing Einstein's theory. Observations force these models to have  a dynamics that is very similar to that of the $\\Lambda$CDM model around the present epoch. Moreover, the past evolution is restricted by the necessity of a matter-dominated epoch to guarantee cosmic structure formation. The future cosmological evolution within alternative models, however, may be very different from a de Sitter phase, which is the final fate of a $\\Lambda$CDM universe and also of several other approaches like Chaplygin-gas scenarios. Phantom-type dark energy with a constant equation-of-state parameter, e.g., will end in a big-rip singularity after a finite time \\cite{caldwell,caldwell2}. More recently, a still different scenario, called ``little rip\", was proposed \\cite{frampton}. Already at the beginning of the past decade several authors discussed the possibility that the currently observed accelerated expansion might be a transient phenomenon, i.e. that there might occur a transition back to decelerated expansion \\cite{albrecht,barrow,bertolami}. Some recent observations seem to back up this idea. Evidence was found for a slowing-down of the expansion rate of the Universe, equivalent to an increase of the deceleration parameter $q(z)$ for decreasing redshifts $z$ close to the present epoch $z=0$ \\cite{sastaro,antonio,liwuyu,caituo}. This could indicate a scenario, according to which the observed accelerated expansion of the Universe is a transient phenomenon, implying a transition back to a decelerated expansion either for the future evolution with $z\\lesssim 0$ or even around the current epoch $z\\gtrsim 0$. Scenarios of transient acceleration were previously discussed in \\cite{alcaniz} and  \\cite{alcaniztr}. The model on which the present paper relies was developed in \\cite{transient}. It describes transient cosmological acceleration as the consequence of an interaction between dark matter and dark energy. Such a dynamics cannot be obtained if the interaction represents a small correction to the standard $\\Lambda$CDM model. For models of such type  the long-time cosmological dynamics will always be determined by the cosmological term and result in accelerated expansion. To achieve transient accelerated expansion, a twofold role of the interaction is necessary. At first, it has to cancel the ``bare\" cosmological constant and at second it has to generate a phase of accelerated expansion by itself. Acceleration has to be an interaction phenomenon. As it was shown in \\cite{transient}, these requirements can be fulfilled by interaction terms that combine powers and exponentials of the cosmic scale factor. Even though this specific choice was made for mathematical convenience, we expect that the mentioned two features will be crucial for a broader class of models. While the study in \\cite{transient} was restricted to the homogeneous and isotropic background, the present paper investigates the corresponding perturbation dynamics as well. In particular, we calculate the growth rate of the matter perturbations and the matter power spectrum. We compare our results with the growth-rate data collected in \\cite{gong} as well as with those of the  WiggleZ survey \\cite{blake} and with the data from the 2dFGRS program \\cite{cole}. Emphasis is also put on the potential relevance of perturbations of the dark-energy component which are neglected by many studies of the matter perturbation behavior (for exceptions see, e.g. \\cite{sapone,Park-Hwang,vernizzi}). In concordance with parallel investigations for other models \\cite{saulo}, we find that they are small indeed on scales that are relevant for structure formation. However, there are indication that their role is increasing with increasing scale. The background dynamics is reconsidered on the basis of the SNIa observations by the Constitution set \\cite{Hicken}.  The paper is organized as follows. In section \\ref{model}  we reanalyze  the basic features and the homogeneous and isotropic background dynamics of the transient acceleration model. Section \\ref{newton} is devoted to a Newtonian treatment of the perturbation dynamics. A fully relativistic, gauge-invariant investigation and a calculation of the matter power spectrum are the subjects of section \\ref{rel}. In section \\ref{summary} we summarize and discuss our results. ",
        "conclusions": "\\label{summary}  We have tested a phenomenological model of transient accelerated expansion in which an interaction in the dark sector is constitutive for the cosmological dynamics. The interaction has both to cancel a ``bare\" cosmological constant and, at the same time, to generate a phase of accelerated expansion. While the detailed structure of the interaction was chosen for mathematical convenience, it admits an analytic solution of the background dynamics, we think that it can be used to discuss general features of transient acceleration models. We reconsidered the background dynamics of this model and performed a statistical analysis based on the SNIa data of the Constitution sample. The model predicts a future minimum of the deceleration parameter $q$ which afterwards switches to positive values again (Fig.~\\ref{figq}). The dark-energy density becomes negative for $a\\gtrsim 6$ but this contribution is exponentially suppressed. The direction of the energy transfer is from dark energy to dark matter for $a\\lesssim 8$. It is reversed for  $a\\gtrsim 8$ but the entire interaction term becomes exponentially suppressed as well.  The perturbation dynamics of the model was investigated both on the Newtonian and on the GR levels. Using a simple parametrization for the perturbed interaction term in the Newtonian setting  and including perturbations of the dark-energy component within a simple ansatz, we carried out a statistical analysis, using the growth-rate data in \\cite{gong} and \\cite{blake}. This allowed us to quantify, although with a large dispersion, the role of the perturbed interaction term and the contribution of the dark-energy perturbations. For $a\\lesssim 2$ the fractional matter perturbations deviate stronger from the Einstein-de Sitter behavior than for the $\\Lambda$CDM model (Fig.~\\ref{figdelta1}). This corresponds to a larger difference for the growth rate $f(z)$ as well. For $a\\gtrsim 2$, however, the matter perturbations continue to grow while they approach a constant value for the  $\\Lambda$CDM model. The effective gravitational constant $G_{eff}$ deviates considerably from $G$ for $a>1$ but the ratio $G_{eff}/G$ approaches unity again in the far-future limit (Fig.~\\ref{figGeff}).  The relativistic perturbation analysis for the interacting two-component system was performed in terms of gauge-invariant quantities with physical interpretation in the comoving gauge. Perturbations in such type of systems are intrinsically non-adiabatic.  To decouple the relevant perturbation equation we assumed for simplicity that the fractional perturbations of the dark energy are proportional to the fractional matter perturbations. The statistical analysis, based on the data from the 2dFGRS project reveals that the factor of proportionality is very small. In other words, dark-energy perturbations are small on scales that are relevant for structure formation. However, our analysis also shows that the considered data range is not adequately described by a constant factor. For the smallest scales we have much more data than for larger scales. Consequently, the small-scale date have a higher statistical weight than the fewer data on larger scales. On the smallest scales, the dark-energy fluctuations are irrelevant indeed. On the other hand it is obvious that on larger scales the overall best-fit curve does not provide a good description of the observations. A considerably larger (but still small) value of the mentioned factor shows a much better performance (Fig.~\\ref{Spectrum32}). This indicates an increasing role of the dark-energy perturbations with increasing scale, a subject that deserves attention in future research."
    },
    "1112/1112.5084_arXiv.txt": {
        "abstract": "We have succeeded in developing hydrophobic silica aerogels over a wide range of densities (i.e. refractive indices). A pinhole drying method was invented to make possible producing highly transparent aerogels with entirely new region of refractive indices of 1.06-1.26. Obtained aerogels are more transparent than conventional ones, and the refractive index is well controlled in the pinhole drying process. A test beam experiment was carried out in order to evaluate the performance of the pinhole-dried aerogels as a Cherenkov radiator. A clear Cherenkov ring was successfully observed by a ring imaging Cherenkov counter. We also developed monolithic and hydrophobic aerogels with a density of 0.01 g/cm$^3$ (a low refractive index of 1.0026) as a cosmic dust capturer for the first time. Consequently, aerogels with any refractive indices between 1.0026 and 1.26 can be produced freely. ",
        "introduction": "Silica aerogel is widely used as a Cherenkov radiator in high energy and nuclear experiments. We developed aerogel production technique for the Belle aerogel Cherenkov counter in 1990s, and have studied for its upgrade program since 2002 \\cite{example1}. In our conventional production method \\cite{example2}, we are able to produce hydrophobic aerogels for a refractive index range of 1.005-1.07 by using alcohol solvent for the alcogel (wet gel) synthesis. Introduction of new solvent, N,N-dimethylformamide (DMF), extended the upper limit of refractive index up to 1.11. Furthermore, the range of refractive index was extended again by a new production method in a recent study \\cite{example3}. The production method also generates more transparent aerogels than conventional ones for a refractive index range of 1.06-1.26. The aerogels with high refractive index and high transparency will expand future possibilities of particle identification.  In addition, aerogel with a refractive index under 1.005 was also developed in recent years. It is a material with an extremely low density under 0.02 g/cm$^3$ because refractive index ($n$) of aerogel is expressed as a function of its density ($\\rho $): $n = 1+\\alpha \\rho $, where $\\alpha $ is a constant. It is expected that the low density aerogels will be utilized in cosmic dust sample return experiments, for example, the TANPOPO \\cite{example4}, which is an astrobiology mission planned on the Japanese Experiment Module (JEM) in the International Space Station (ISS).  In this article, the development of extremely low density aerogels is described in the next section. Section 3 is devoted to the introduction of the new production method for high refractive index aerogels. In Section 4, the transparency of new aerogels is explained. Then, conclusion is given after a description of a result of a beam test in Section 5. ",
        "conclusions": "We have succeeded in developing hydrophobic silica aerogel with refractive indices over a range of 1.0026-1.26. Especially, the pinhole drying method enables to obtain highly transparent aerogels with high refractive indices up to $n$ = 1.26. It was confirmed that pinhole-dried aerogels can be used as a Cherenkov radiator by the beam test. On the other hand, extremely low density aerogel ($\\rho $ = 0.01 g/cm$^3$) will be utilized as a cosmic dust capturer in the sample return experiment on the JEM/ISS."
    },
    "1112/1112.3341_arXiv.txt": {
        "abstract": "{We find ourselves in an extended era of entropy production.  Unlike most other observations, the arrow of time is usually regarded as a constraint on initial conditions.  I argue, however, that it primarily constrains the vacuum structure of the theory.  I exhibit simple scalar field potentials in which low-entropy initial conditions are not necessary, or not sufficient, for an arrow of time to arise.  I argue that the string theory landscape gives rise to an arrow of time independently of the initial entropy, assuming a plausible condition on the lifetime of its metastable vacua. The dynamical resolution of the arrow of time problem arises from the same structural properties of the string landscape that allow it to solve the cosmological constant problem without producing an empty universe, particularly its high dimensionality and the large difference in vacuum energy between neighboring vacua.}   \\begin{document} ",
        "introduction": "\\label{sec-intro}  The entropy $S$ of a closed macroscopic system either (a) increases or (b) is constant under unitary evolution. This is easily understood on statistical grounds: under reasonable assumptions, (a) is the most probable evolution of a system that has not yet reached its maximum entropy, $S_{\\rm max}=\\ln \\dim \\cal H$, where $\\cal H$ is the Hilbert space of the system; while (b) is the most probable evolution of a system with $S=S_{\\rm max}$.  In theory, this statement is invariant under time reversal: Given a state at time $t$, the entropy will probably increase or stay constant, no matter whether we evolve in the past or future direction.  In practice, however, we observe that the entropy of matter systems increases only towards the future, never towards the past.  This asymmetry is called the (thermodynamic) arrow of time.  It can be explained by noting that the early universe was in a special state of relatively low entropy.  But why was this the case?  This question is usually referred to the arrow-of-time ``problem'', but the arrow of time is an observation like any other: it can become a problem only in the context of a particular cosmological theory.  A theory is ruled out if its predictions conflict with the observed arrow of time.  And a theory is incomplete if it is compatible with but does not predict the observed arrow of time.  The purpose of this paper is twofold.  First, I will show that the conditions required for an arrow of time depend on the theoretical framework and cannot be equated with low initial entropy.   I will exhibit simple models that demonstrate that a low-entropy initial state is neither necessary nor sufficient for explaining the observed arrow of time.  Secondly, I will argue that the landscape of string theory furnishes a powerful example in which the observed arrow of time can arise from arbitrarily high-entropy initial conditions.  This property rests on some of the same key features that allow the string landscape to solve the cosmological constant problem without fine-tuning.  It constitutes a second, independent success of the framework.    \\paragraph{Outline and Summary}  In Sec.~\\ref{sec-obs}, I characterize the observed arrow of time and distinguish it from other observations.  The arrow of time is defined as the entropy produced in our causal past since the time of big bang nucleosynthesis, $\\Delta S\\sim 10^{103}$.  In Sec.~\\ref{sec-one}, I consider cosmological models with a unique vacuum.  If the cosmological constant, $\\Lambda$, is negative (Sec.~\\ref{sec-neg}), one finds that the requirement of large maximum entropy leads to multiple challenges, in addition to the smallness of the initial entropy.  Most importantly, the magnitude of the cosmological constant must be small.  This requirement is shared by all other models considered later: because of the covariant entropy bound, the observed arrow-of-time alone requires the existence of a vacuum with unnaturally small cosmological constant.  In this sense, the cosmological constant problem is part of the arrow of time problem.  Of course, even if we set $-\\Lambda\\ll 1$, this class of models is ruled out by the observed positive sign of the cosmological constant~\\cite{Rie98,Per98}.  In the remainder of the paper, I consider only models that contain at least one vacuum with positive vacuum energy.  In Sec.~\\ref{sec-pos}, I still consider models with only one vacuum, but now the sign of its cosmological constant is taken to be positive.  Generic initial conditions lead to eternal inflation; and because of the absence of vacua with negative energy, all possible states recur infinitely many times, even in a single causal patch.  As was first pointed out by Dyson, Kleban, and Susskind~\\cite{DysKle02}, this implies that the overwhelming majority of observers arise from the smallest possible fluctuations relative to the maximum entropy state, empty de~Sitter space.  Such observers are known as Boltzmann brains.  They observe no arrow of time other than the small arrow associated with their own existence.  This conclusion obtains independently of initial conditions, so these models conflict violently with the observed arrow of time.  They, too, are ruled out.  In Sec.~\\ref{sec-two}, I consider small landscapes that contain both de~Sitter ($\\Lambda>0$) and terminal ($\\Lambda<0$) vacua, with initial conditions in one of the de~Sitter vacua.  This leads to eternal inflation.  In Sec.~\\ref{sec-cp}, I review the associated measure problem and the causal patch measure.  In Sec.~\\ref{sec-twovac}, a model with only two vacua is analyzed.  An arrow of time arises only if two conditions hold: the initial state must have low entropy, and the de~Sitter vacuum decay must proceed faster than the production of Boltzmann brains.  In Sec.~\\ref{sec-domhis}, a fast method for reproducing these results is developed, based on the branching tree implementation of the causal patch measure.  This method is applied to two models, each with four vacua, in Sec.~\\ref{sec-four}.  The two models differ only through the ordering of the vacua; the values of the cosmological constant are the same.  Yet, in the first model, low-entropy initial conditions are necessary for an arrow of time; whereas in the second, an arrow of time is predicted even with high-entropy initial conditions.  This constrast provides an instructive reference point for the analysis of the string landscape.  The key feature of the second model is that observers of any type---Boltzmann or ordinary---can be produced only after passing through a vacuum with large cosmological constant, and hence, very small entropy.  This feature is shared by the string landscape.  In Sec.~\\ref{sec-big}, I claim that the string landscape predicts an arrow of time, independently of the initial entropy, if and only if all de~Sitter vacua decay faster than they produce Boltzmann brains.  Only four key properties of the string landscape, listed in Sec.~\\ref{sec-properties}, are used in the proof of the claim, in Sec.~\\ref{sec-proof}.  They are the high dimensionality of the landscape, the large but not double-exponentially large number of vacua, the lack of tuning, and the fact that neighboring vacua tend to have vastly different vacuum energy.  It is interesting that some of these properties also underly the solution of the cosmological constant problem.  \\paragraph{Discussion}  The approach followed in this paper is based on semiclassical gravity.  The universe is assumed to have a beginning, and time is taken to end at spacelike singularities such as the big crunch in a vacuum with negative cosmological constant.  Whether these features survive in a full quantum gravity theory remains a major open question.  They are crucial here because they allow some of the models considered to evade ergodicity.  This is necessary for a large arrow of time to arise from arbitrarily high-entropy initial conditions.  All models considered except for the single vacuum in Sec.~\\ref{sec-neg} exhibit the dynamics known as eternal inflation, and thus require a regulator or ``measure''.  Here the causal patch measure~\\cite{Bou06,BouFre06a} is used, which is equivalent to the light-cone time cutoff~\\cite{Bou09,BouYan09,BouFre10b}.  Results would be similar with closely related proposals such as the fat geodesic~\\cite{BouFre08b,LarNom11}, scale factor~\\cite{DesGut08a}, and CAH+ cutoffs~\\cite{Vil11}.  In some of the literature on the arrow of time problem, it is taken for granted that the initial conditions of the universe should somehow be selected at random, from a Hilbert space that is imagined to be large enough to describe the presently visible universe.  To some~\\cite{Wal05} (including myself, until recently), this seemed like a rather implausible assumption, since it was unclear how the observed arrow of time could arise dynamically from high entropy initial conditions.  In this paper, however, I demonstrate explicitly how this {\\em can\\/} happen.  Therefore, high entropy initial conditions are {\\em not\\/} automatically in conflict with observation.\\footnote{This does not mean, of course, that high initial entropy should be taken for granted.  Since we have no clue as to the correct theory of initial conditions (if there is one), we have no right to exclude the possibility that the universe simply began small and/or with low entropy.  See, e.g., Refs.~\\cite{Lin84b,Vil83b}.} Whether they are depends on the vacuum structure of the theory.  For example, the Hartle-Hawking no boundary proposal~\\cite{HarHaw83} exponentially favors initial conditions corresponding to empty de~Sitter space with the smallest positive vacuum energy available in the theory.  In a theory with a single vacuum, this rules out the proposal.  But in certain potential landscapes such as the second model in Sec.~\\ref{sec-four}, the Hartle-Hawking proposal is not excluded as a theory of initial conditions.  (In the landscape of string theory, the Hartle-Hawking proposal is not viable~\\cite{BouZukTA}.\\footnote{I would like to thank D.~Page for discussions of this point.})   \\paragraph{Related work}  This paper builds on several results of earlier work, particularly on the seminal insights of Dyson, Kleban, and Susskind about recurrences in de Sitter vacua~\\cite{DysKle02} (Sec.~\\ref{sec-pos}), and on an analysis of the abundance of Boltzmann brains in the multiverse~\\cite{BouFre08b,DesGut08b}.  Some of the assumptions and the proof in Sec.~\\ref{sec-proof} differ slightly from those given in Ref.~\\cite{BouFre08b}; in the context of the present paper, assumptions were desirable that distill out the key properties of the string landscape that lead to a dynamical explanation of the arrow of time.  The main points of the present paper are (i) the arrow-of-time problem is ill-defined outside a specific theoretical context; (ii) the vacuum structure of a theory is central to the question of whether and how an arrow of time arises; and (iii) the vacuum structure of the string landscape does lead to the prediction of an arrow of time, if a plausible condition on the lifetime of de~Sitter vacua is satisfied.  For a clear and concise introduction to the arrow of time problem, see Ref.~\\cite{Wal05}.\\footnote{The present paper provides counterexamples to two propositions discussed in Ref.~\\cite{Wal05}: the claim that the arrow of time cannot arise dynamically but must come from initial conditions (it arises dynamically from high-entropy initial conditions in some models I exhibit, including the string landscape); and the claim that conditioning on the existence of observers cannot explain the arrow of time because it would be equivalent to assuming the arrow of time (we will encounter models in which such conditioning leads to the prediction of an empty universe except for a single Boltzmann brain, in conflict with the observed arrow of time).}  Interesting recent papers that relate to the arrow of time problem in eternal inflation or the string multiverse include Refs.~\\cite{AlbSor04,CarChe04,Pag06,BouFre06b,Ban07,% Alb09,CarTam10,AguCar11,Teg11,Nom11b}.  The present analysis is not directly related to the {\\em Entropic Principle}~\\cite{Bou06,BouHar07}.  Arguably, observation requires free energy.  Although this is not sufficient, matter entropy production turns out to be a successful proxy for observers in many settings~\\cite{Bou06,BouHar07,CliFre07,BozAlb09,% PhiAlb09,BouLei09,BouHar10}, and it has the advantage that allows for conditioning on observers without making specific anthropic assumptions such as the existence of galaxies, or stable nuclei, etc.  Employing the Entropic Principle for this purpose in no way prejudices the question of the arrow of time, since both Boltzmann brains and ordinary observers operate by increasing the entropy.  The question addressed here is why the entropy produced in the universe we see is much larger than necessary for our own operation: why is the arrow of time so large? ",
        "conclusions": ""
    },
    "1112/1112.3388_arXiv.txt": {
        "abstract": "A bright prominence associated with a coronal mass ejection (CME) was seen erupting from the Sun on 9~April 2008.  This prominence was tracked by both the {\\it Solar Terrestrial Relations Observatory} (STEREO) EUVI and COR1 telescopes, and was seen to rotate about the line of sight as it erupted; therefore, the event has been nicknamed the ``Cartwheel CME.'' The threads of the prominence in the core of the CME quite clearly indicate the structure of a weakly to moderately twisted flux rope throughout the field of view, up to heliocentric heights of 4~solar radii. Although the STEREO separation was 48{\\degrees}, it was possible to match some sharp features in the later part of the eruption as seen in the 304~{\\AA} line in EUVI and in the H$\\alpha$-sensitive bandpass of COR1 by both STEREO {\\it Ahead} and {\\it Behind}.  These features could then be traced out in three-dimensional space, and reprojected into a view in which the eruption is directed towards the observer.  The reconstructed view shows that the alignment of the prominence to the vertical axis rotates as it rises up to a leading-edge height of $\\approx\\!2.5$ solar radii, and then remains approximately constant.  The alignment at 2.5 solar radii differs by about 115{\\degrees} from the original filament orientation inferred from H$\\alpha$ and EUV data, and the height profile of the rotation, obtained here for the first time, shows that two thirds of the total rotation is reached within $\\approx\\!0.5$~solar radii above the photosphere. These features are well reproduced by numerical simulations of an unstable moderately twisted flux rope embedded in external flux with a relatively strong shear field component. ",
        "introduction": "The bandpass of the inner coronagraph (COR1) telescopes on the two {\\it Solar Terrestrial Relations Observatory} (STEREO) spacecraft runs from 650--670~nm \\cite{howard:2008}. This range was selected to include the hydrogen H{$\\alpha$} line at 656~nm, so as to be sensitive to erupting prominences associated with coronal mass ejections (CMEs).  A particularly bright prominence eruption was observed by both COR1 telescopes on 9~April 2008.  Although the relative contributions of H{$\\alpha$} emissions and Thomson-scattered light cannot be derived from the COR1 observations alone, it is presumed that H{$\\alpha$} is a significant contributor for the extremely bright parts of the prominence early in the eruption.  This CME and prominence eruption was also observed by the STEREO {\\it Extreme Ultraviolet Imager} (EUVI) telescopes \\cite{howard:2008} before entering the COR1 field of view, as well as by TRACE at 171~{\\AA} \\cite{handy:1999}, and by the {\\it Hinode} XRT telescope \\cite{golub:2007} and EIS spectrometer \\cite{culhane:2007}. \\citeauthor{Landi&al2010} (2010) analyze this event, combining {\\it Hinode}/XRT and EIS data, together with {\\it Solar and Heliospheric Observatory} (SOHO)/UVCS/EIT/LASCO and STEREO/EUVI/COR1/COR2 images, to characterize the thermal properties of the ejected plasma, and constrain the heating rate. The post-CME current sheet following this event is analyzed by \\inlinecite{savage:2010} and \\inlinecite{ko:2010}. \\inlinecite{Patsourakos&Vourlidas2011} determine the geometrical parameters of the CME and post-CME current sheet in the field of view of the outer coronagraph (COR2) on STEREO. This event has become known as the ``Cartwheel CME'' because of the highly visible rotation around the line of sight seen during the initial stages of the eruption.  Rotation about the vertical direction can also be observed through triangulation as the prominence material rises through the EUVI and COR1 height ranges. The plane of sky motions which give this event its ``Cartwheel'' name is actually a combination of the rotation about the vertical axis together with a deflection of the prominence in latitude and longitude during the initial rise phase. A rotation around the axis of the prominence, known as ``roll effect'' \\cite{martin:2003}, may have contributed as well. Further rotation about the vertical direction is indicated by the CME orientation in the COR2 data.  This paper presents an exploration of the rotation about the vertical direction in the EUVI and COR1 height ranges, using triangulation from the two STEREO viewpoints. The height profile of the rotation angle for the erupting prominence is obtained, which could not be done prior to the STEREO mission.  The separation of the two STEREO spacecraft at this time was 48{\\degrees}. Co-identification of features in both views is quite difficult with such a large separation.  However, it was possible to match some sharp features in the later part of the eruption as seen by EUVI in the He~{\\sc ii} line at 304~{\\AA}.  Features were also tracked in white light (with a possible H{$\\alpha$} component) as seen by COR1, although not necessarily the same features seen at 304~{\\AA}.  The derived {height profile of prominence rotation is compared with the orientation of the  polarity inversion line of the radial field component (PIL) at various heights, to check whether the rotation is simply the result of alignment with the PIL. The radial field component is obtained from a potential-field source-surface (PFSS) extrapolation. Additionally, the rotation} and rise profiles of the prominence, as well as the STEREO images, are compared with the corresponding data obtained in a series of CME simulations in a companion paper (\\opencite{Kliem&al2011}; in the following Paper~II), in order to {explore a broader range of possible origins} of the large total rotation observed. The {comparisons suggest} that the major part of the rotation was due to the presence of a shear field component of considerable strength in the source region \\cite{isenberg:2007} and that the weak helical kink instability of a moderately twisted flux rope \\cite{torok:2004} contributed as well. Field line plots of the best matching case are included here to demonstrate the correspondence with the observed overall shape and helical threads of the rising prominence. ",
        "conclusions": "The images of the STEREO EUVI and COR1 telescopes resolve the basic structure of the erupting prominence in the core of the ``Cartwheel CME'' on 9~April 2008 out to heliocentric distances of $4~R_\\odot$. A single flux rope with indications of weak to moderate twist is revealed. The flux rope expands approximately self-similarly in the COR1 field of view of (1.5--4)~$R_\\odot$.  The true path of the prominence could be reconstructed from stereoscopic observations even when the angular separation between the STEREO spacecraft was as much as 48{\\degrees}.  The reconstruction reveals that the {dextral} prominence rotated counter-clockwise by a large angle of $\\approx\\!115\\degrees$ up to a heliocentric height of $2.5~R_\\odot$, where the rotation leveled off. A gentle backward rotation by $\\approx\\!15${\\degrees} may have followed in the height range up to $3.3~R_\\odot$. Two thirds of the rotation were acquired within $0.5~R_\\odot$ from the photosphere. The coronal height profile of the rotation angle in a CME is thus derived for the first time. To our knowledge, this ranks as one of the largest rotations so far measured in the corona.  Taken jointly with recent results about the orientation of the CME at a heliocentric distance of $\\approx\\!13~R_\\odot$, obtained from fitting a croissant-shaped flux rope model to stereoscopic images from COR2 \\cite{Patsourakos&Vourlidas2011}, a further counter-clockwise rotation by $\\approx\\!35\\degrees$ from the value at $2.5~R_\\odot$ is indicated. This aligned the erupted flux closely with the heliospheric current sheet above the active region. {However, the structure of the field within the first $\\sim\\!0.5~R_\\odot$ above the original location of the prominence, obtained from a PFSS extrapolation, precludes alignment with the PIL as mechanism for the initial rotation.}   The rotation and rise profiles of the prominence in the EUVI-COR1 height range are reproduced by a numerical model that follows the evolution of an unstable force-free flux rope. A parametric study of flux rope rotation in this model, detailed in Paper~II, suggests that the main part of the rotation in {this height range} was caused by the shear field in the source volume. While the rotation and rise profiles can be modeled nearly equally well by a weakly kink-unstable flux rope of $3.5\\pi$ initial twist and by a kink-stable flux rope of $2.5\\pi$ initial twist, the indications of twist in the COR1 images and the absence of a strong initial perturbation in the EUVI-{\\it Ahead} data favor the $3.5\\pi$ model, suggesting that the helical kink instability did occur and contributed the remaining part of the {rotation}.  {In summary, the large total rotation of the erupting flux rope in the Cartwheel CME is due to a combination of three effects. The internal Lorentz forces associated with the shear field and with the tension of twisted field lines caused a strong rotation in the inner corona. The resulting orientation relative to the heliospheric current sheet was favorable for the subsequent alignment with the sheet in the outer corona and inner solar wind to proceed with the same sense of rotation. Since the heliospheric current sheet had an orientation relatively near the original direction of the erupting flux, a very large final rotation angle was reached.}"
    },
    "1112/1112.3207_arXiv.txt": {
        "abstract": "In preparation for the {\\it HERMES} chemical tagging survey of about a million Galactic FGK stars, we estimate the number of independent dimensions of the space defined by the stellar chemical element abundances [X/Fe{]}. This leads to a way to study the origin of elements from observed chemical abundances using Principal Component Analysis. We explore abundances in several environments, including solar neighbourhood thin/thick disk stars, halo metal-poor stars, globular clusters, open clusters, the Large Magellanic Cloud and the Fornax dwarf spheroidal galaxy. By studying solar-neighbourhood stars, we confirm the universality of the $r$-process that tends to produce [$n$-capture elements/Fe{]} in a constant ratio. We find that, especially at low metallicity, the production of $r$-process elements is likely to be associated with the production of $\\alpha$-elements. This may support the core-collapse supernovae as the $r$-process site. We also verify the over-abundances of light $s$-process elements at low metallicity, and find that the relative contribution decreases at higher metallicity, which suggests that this lighter elements primary process may be associated with massive stars. We also verify the contribution from the $s$-process in low-mass AGB stars at high metallicity. Our analysis reveals two types of core-collapse supernovae: one produces mainly $\\alpha$-elements, the other produces both $\\alpha$-elements and Fe-peak elements with a large enhancement of heavy Fe-peak elements which may be the contribution from hypernovae. Excluding light elements that may be subject to internal mixing, K and Cu, we find that the [X/Fe{]} chemical abundance space in the solar neighbourhood has about 6 independent dimensions both at low metallicity ($-3.5 \\la \\mbox{[Fe/H{]}} \\la -2$) and high metallicity ($\\mbox{[Fe/H{]}} \\ga -1$). However the dimensions come from very different origins in these two cases. The extra contribution from low mass AGB stars at high metallicity compensates the dimension loss due to the homogenization of the core-collapse supernovae ejecta. Including the extra dimensions from [Fe/H{]}, K, Cu and the light elements, the number of independent dimensions of the [X/Fe]+[Fe/H] chemical space in the solar neighbourhood for {\\it HERMES} is about 8 to 9. Comparing fainter galaxies and the solar neighbourhood, we find that the chemical space for fainter galaxies such as Fornax and the Large Magellanic Cloud has a higher dimensionality. This is consistent with the slower star formation history of fainter galaxies. We find that open clusters have more chemical space dimensions than the nearby metal-rich field stars. This suggests that a survey of stars in a larger Galactic volume than the solar neighbourhood may show about $1$ more dimension in its chemical abundance space. ",
        "introduction": " ",
        "conclusions": "\\subsection[]{The $\\bmath{n}$-capture elements subspace} \\subsubsection[]{The $r$-process contribution} Several authors including \\citet{hon04a,iva06} and \\citet{hay09} noted that the $r$-process production ratio for the $n$-capture elements in [X/Fe] is almost constant despite the large scatter in the abundance of each element. Therefore, in the $\\mathcal{C}$-space, if the major contribution is from the $r$-process, one would expect the data cloud to form a tight straight line in the $n$-capture element space.  It is pleasing to see that, at low metallicity where the $r$-process should be dominant \\citep{tru81}, the PCA shows a strong first component made up of all elements positively correlated. This $r$-process component contributed  $\\ga 80\\%$ of the data cloud variance for the metal-poor stars. This agrees with other conclusions from the literature.  \\subsubsection[]{The overabundance of light $s$-process elements} Observations have also reported more scatter in the ratio of light $s$-process elements to heavy $s$-process elements such as [Y/Ba] \\citep{bur00,nor01,iva03,aok05,hon07} at [Fe/H] $<-3$. As shown in Fig.~\\ref{fig:light-s-over-abundances}, the light $s$-process elements seem to be overproduced at lower metallicity relative to the heavier elements. Various mechanisms have been proposed to explain this observation including the weak $r$-process \\citep[e.g.][]{izu09} and the CPR \\citep{qia07}. Our PCA method clearly show the existence of this unknown source. For the low metallicity sample (Fig.~\\ref{fig:Barklem-neutron}), the second principal component is made up of light $s$-process elements anticorrelated with heavier $n$-capture elements. This suggests a second mechanism which either overproduces light $s$-process elements or overproduces heavier $n$-capture elements. From Fig.~\\ref{fig:light-s-over-abundances}, the former interpretation is preferred.  Also, as shown in Table~\\ref{tab:LEPP}, the second primary source for light $s$-process elements contributes more as the metallicity decreases. In Fig.~\\ref{fig:ubiquitous-r}, the First Stars Survey sample shows slightly more dimensions than the Barklem's sample. This could be interpreted as a larger contribution from the LEPP since the First Stars Survey sample has a lower mean metallicity. This decreasing trend of LEPP with increasing metallicity suggests that the culprit for LEPP is restricted to low metallicity and/or massive progenitors. If this is due to the mass dependence, this is consistent with the proposed LEPP candidates-- the weak $r$-process, CPR, collapsars or the truncated $r$-process. The relative contribution of the LEPP becomes smaller at higher metallicity and the production of $n$-capture elements is dominated by the main $r$-process and the $s$-process. The results in Table~\\ref{tab:LEPP} also illustrates one of the strengths of the PCA method compared with the traditional [X/Fe]--[Fe/H] plots-- one can quantify the contributions of different mechanisms.  \\subsubsection[]{Low-mass AGB contribution} At high metallicity, one would expect a contribution from low-mass AGB stars with initial mass range of $1.5$ -- $3~\\mbox{M}_\\odot$. Unlike LEPP, the $s$-process in low-mass AGB stars produces both light $s$-process elements and heavy $s$-process elements such as Ba \\citep[see][]{her05,kap11}, but not so much of the mostly $r$-process elements such as Nd and Eu relative to the $r$-process contribution \\citep{arl99}.  In our PCA analysis, we find that the relative contribution from the first component (i.e. the $r$-process) decreases at higher metallicity (Fig.~\\ref{fig:ubiquitous-r} and Fig.~\\ref{fig:cumul-Reddy-neutron}). We also find that the second component is made up of both light and heavy $s$-process elements (Fig.~\\ref{fig:Reddy-neutron}). Both are consistent with the $s$-process in low-mass AGB stars, which is dominant at high metallicity.  To show that this transition between high metallicity and low metallicity is real, we have also analyzed the intermediate metallicity sample taken from \\citet{bur00}. We separate the the sample into two extreme metallicity bin: $-2.7<$ [Fe/H] $<-2$ and $-1.5<$ [Fe/H] $<-1$. The former shows results that resembles Fig.~\\ref{fig:Barklem-neutron} and the latter shows results that is similar to Fig.~\\ref{fig:Reddy-neutron}. This further confirms that the transition is indeed happening and occurs around $-2 \\la$ [Fe/H] $\\la -1.5$.  \\subsection[]{Satellite galaxies} In Section~\\ref{subsec:dwarf-galaxies-results}, we presented results for the dwarf galaxies samples in comparison with the solar neighbourhood stars in the same metallicity range. Recall that in this metallicity range, many stars in the solar neighbourhood are likely to be thick disk stars, and the belief is that the thick disk had a brief and intense star formation history (SFH) \\citep*[see also][]{chi01,kob06}. On the other hand, the SFH is believed to be slower for dwarf galaxies \\citep[e.g.][]{wei11}.  In our study, our PCA method results are consistent with the speculation that dwarf galaxies have slower SFH. We showed that dwarf galaxies seem to have more independent $\\mathcal{C}$-space dimensions than the solar neighbourhood. This implies that the ISM in dwarf galaxies is impacted by a higher level of stochastic enrichment: the less vigorous history of enrichment events in the dwarf galaxies may allow more chemical traces to persist.  Now we will discuss the principal components for $n$-capture elements. Our PCA analysis shows two peculiarities of the $n$-capture elements in the dwarf galaxies. Firstly, the second principal component gives a significant contribution to the data cloud variance. We interpret the second principal component as the LEPP contribution (because it has a large contribution from light $s$-process elements).  In contrast, for the solar neighbourhood stars in the same metallicity range, the LEPP contribution is hardly visible and only contributions from the $r$-process and the $s$-process in low-mass AGB stars are seen. This is consistent with the view that the SFH in dwarf galaxies is slower and therefore retains more chemical substructure from the previous generation of ejecta.  Secondly, the third principal component in this analysis was not seen before. It has heavy $s$-process elements (Ba, La) anticorrelated with the other $n$-capture elements, including the light $s$-process elements. Recall that previously at high metallicity, for $n$-capture elements in the Milky Way's solar neighbourhood, the $s$-process in low-mass AGB stars tends to produce both light and heavy $s$-process elements. If the third component represents the AGB contribution, one would expect it to have the same sign for the light and heavy $s$-process elements. However, this is not the case for the dwarf galaxies. This result is not totally unexpected. As shown in Fig.~\\ref{fig:dwarf-light-s} and Fig.~\\ref{fig:dwarf-heavy-s}, observations for dwarf galaxies have shown over-abundances of [{\\it hs}/\\mbox{Fe}] at high metallicity but similar abundances or even under-abundances for the  [{\\it ls}/\\mbox{Fe}]. \\begin{figure} \\begin{minipage}{80mm} $\\begin{array}{c} \\hspace{-0.7cm}\\includegraphics[width=3.4in]{Fig22a.ps}\\\\ \\hspace{-0.7cm}\\includegraphics[width=3.4in]{Fig22b.ps} \\end{array}$ \\caption{The [{\\it hs}/Fe]--[Fe/H] of the solar neighbourhood, Fornax and the LMC stars. The symbols in {\\it grey} represent the solar neighbourhood compilation from \\citet{bur00} ({\\it plus}), \\citet{ful00,ful02} ({\\it crosses}) and \\citet{red03,red06} ({\\it open circles}). The {\\it blue filled circles} are Fornax stars from \\citet{let10}. The {\\it red filled triangles} are the LMC stars from \\citet{pom08}. The two dwarf galaxies are overabundant in heavy $s$-process elements, compared to solar neighbourhood stars.}\\label{fig:dwarf-light-s} \\end{minipage} \\end{figure} \\begin{figure} \\begin{minipage}{80mm} $\\begin{array}{c} \\hspace{-0.7cm}\\includegraphics[width=3.4in]{Fig23a.ps}\\\\ \\hspace{-0.7cm}\\includegraphics[width=3.4in]{Fig23b.ps} \\end{array}$ \\vspace{-0.3cm} \\caption{The [{\\it ls}/Fe]--[Fe/H] of the solar neighbourhood, Fornax and the LMC stars. The symbols are the same as in Fig.~\\ref{fig:dwarf-light-s}. The two dwarf galaxies have similar or lower abundances of light $s$-process elements compared to solar neighbourhood stars.}\\label{fig:dwarf-heavy-s} \\end{minipage} \\end{figure}  This could be explained by the argument that the slower SFH in these galaxies allows low-mass AGB stars to contribute at lower metallicity:  AGB stars working in a metal-poor environment have higher ratio of neutrons to Fe-peak seeds and will therefore preferentially produce heavier $s$-process elements (cf. Section~\\ref{sec:chemev}).  If this is the case, we should see some similarities between solar neighbourhood CEMP-s stars and dwarf galaxies. It has been proposed that CEMP-s stars have suffered from AGB binary mass transfer at low metallicity. AGB stars which operate at low metallicity should, by the same argument, preferentially produce {\\it hs} than {\\it ls}. As CEMP-s stars are rare objects, we use compilation from the SAGA database \\citep{sud08}, and select stars with [C/Fe] $>1$ and [Ba/Fe] $>1$ \\citep[criteria from][]{bee05}. For those stars with multiple measurements, we take the measurement from the study with the highest resolution and signal-to-noise ratio \\citep[for more details, see][]{sud11}. In Fig.~\\ref{fig:metal-poor-AGB}, we show the [Y/Ba] for the solar neighbourhood not-CEMP stars, CEMP-s stars and dwarf galaxies stars. This figure is consistent with our speculation. Both solar neighbourhood CEMP-s stars and dwarf galaxies stars show lower [Y/Ba] than the solar neighbourhood not-CEMP stars. This suggests that $n$-capture elements for both CEMP-s stars and dwarf galaxies could have been produced by the $s$-process in AGB stars operating in a metal-poor environment. \\begin{figure} \\hspace{-0.3cm}\\includegraphics[scale=0.5, angle=0]{Fig24.ps} \\vspace{-0.5cm} \\caption{The Y ({\\it ls}) to Ba ({\\it hs}) ratios as a function of [Fe/H] for solar neighbourhood not-CEMP stars, CEMP-s stars and dwarf galaxies stars. The solar neighbourhood not-CEMP stars symbols are the same as in Fig.~\\ref{fig:light-s-over-abundances}. CEMP-s stars are plotted with {\\it black open diamonds}. Fornax stars are plotted with {\\it black crosses} and the LMC stars are plotted with {\\it black open squares}.}\\label{fig:metal-poor-AGB} \\end{figure}  Alternatively, the `normal' abundances of light $s$-process element for dwarf galaxies can be explained if dwarf galaxies do not form many massive stars during their formation. This leads to a smaller contribution from LEPP and smaller production of light $s$-process elements in the formation phase of these galaxies. The under-production of light $s$-process elements from LEPP is then compensated by the onset of low-mass AGB star at low metallicities. The smaller LEPP contribution at low metallicity is not necessarily inconsistent with the PCA analysis results showing that LEPP contribution is still visible at the Fornax metallicity range ($-1.2<$ [Fe/H] $<-0.8$), as shown in Fig.~\\ref{fig:dsph-neutron}, because there are also less $r$-process events to homogenize the ejecta.  \\subsection[]{All elements} \\subsubsection[]{Low metallicity} The interpretation of the PCA analysis of all elements at low metallicity as shown in Fig.~\\ref{fig:all-Barklem} can be summarized as follows:  The first principal component could be the contribution of the $r$-process since it involves substantial production of all of our $n$-capture elements. There is also substantial $\\alpha$-element production together with the $r$-process production. This supports the view that the $r$-process occurs in core-collapse supernovae rather than neutron star mergers. It is generally agreed that the $\\alpha$-elements at low metallicity are produced in core-collapse supernovae. If this interpretation is correct, among $\\alpha$-elements, the composition of the first principal component as shown in Fig.~\\ref{fig:all-Barklem} suggests that core-collapse supernovae that involve the $r$-process produce more heavier $\\alpha$-elements (Ca and Ti) than light $\\alpha$-elements (Mg). In the yields from \\citet{kob06}, heavier $\\alpha$-elements are produced more in massive progenitors and/or with higher explosion energy. Our PCA results may suggest that the $r$-process is associated with more massive progenitors and/or high explosion energy.  The second and third principal components should be discussed together. The second component shows anticorrelation between $\\alpha$-elements and Fe-peak elements whereas the third component shows the production of {\\it both} $\\alpha$-elements and Fe-peak elements. We suggest that this could be tentatively explained by the hypernovae-supernovae model \\citep{ume02,ume05,kob06}. The `normal' supernovae produce more $\\alpha$-elements than iron \\citep[see also][]{woo95,chi02}. However, the amounts are constrained from Galactic chemical evolution models \\citep[e.g.][]{tim95, kob06}. Theoretically, since the formation of neutron stars or blackholes is uncertain, the remnant mass is determined by a mass-cut. For low-mass supernovae ($\\la 20~\\mbox{M}_\\odot$), the ejected Fe mass is constrained to be about $0.07 M_\\odot$ from the light-curves and spectra of individual supernovae. For massive supernovae ($\\ga 20~\\mbox{M}_\\odot$), the observational data is limited, but the significant production of Fe (together with $\\alpha$-elements) is supported by the light-curve and spectra. With higher explosion energy ($\\ga 10 \\times 10^{51}$ erg), it is possible to produce enough amounts of heavy Fe-peak elements including Zn and Co to fit the observed abundances ratios \\citep{kob06}. Our PCA analysis seems to well-match this scenario: the third component could be interpreted as the contribution of hypernovae and the second component could be interpreted as the contribution from `normal' supernovae.  We examine this idea by comparing the First Stars Survey sample which has a lower mean metallicity with the Barklem's sample. Since hypernovae are preferentially produced by more massive progenitors and therefore slightly shorter life span, \\citet{kob11a} introduced metallicity dependence of the hypernovae fraction and argued that the contribution of hypernovae is larger at lower metallicity. However, the fraction of `normal' supernovae and hypernovae is uncertain. To examine these ideas, we analyze the $\\mathcal{C}$-subspace of Mg, Ca, Ti, Co, Ni, Zn. In this subspace, we interpret the component that has positive contribution from all elements to represent the hypernovae contribution and the component that shows anticorrelation between the $\\alpha$-elements (Mg, Ca, Ti) and heavy Fe-peak elements (Co, Ni, Zn) to be the usual supernovae contribution.  With these assumptions,  we found that for the First Stars Survey sample, the hypernovae component makes up $58.5 \\%$ of the two contributions, whereas for the higher-metallicity Barklem's sample, the ratio decreases to $52.8\\%$.  The $\\alpha$-elements contribution appears in all of the first three components. However the second and third components do not show a visible contribution from all $n$-capture elements. This might imply that, although $r$-process production is accompanied by $\\alpha$-elements yield, the reverse is not always true. If the $r$-process occurs in core collapse supernovae, then perhaps it can be triggered only if certain conditions or progenitor mass range are satisfied, as suggested by previous work \\citep{wan06}.  The Al abundance was corrected by $+0.6$ dex for NLTE effects. If this correction is not made, the results remains largely unaltered, except that Al is contributing to the second component.  Al is an light odd-Z element which comes from $^{22}$Ne, and thus the yields depends on the metallicity of progenitor stars \\citep[e.g][]{kob06}; see also \\citet{arn71}. Fig.~\\ref{fig:odd-even-effect} shows that, relative to light even-Z elements such as Mg, light odd-Z elements such as Na and Al are indeed metallicity-dependent and are suppressed at low metallicity. This shows that the production of light odd-Z elements is indeed metallicity-dependent and is suppressed at low metallcity. Therefore, although our NLTE-correction for Al is crude and does not depend on the stellar parameters, the correction is in the correct direction and is required to give a sensible PCA result-- i.e. Al should not be contributing to the core-collapse supernovae component at low metallicity. \\begin{figure} \\begin{minipage}{80mm} $\\begin{array}{c} \\hspace{-0.65cm}\\includegraphics[width=3.4in]{Fig25a.ps} \\vspace{-0.4cm}\\\\ \\hspace{-0.65cm}\\includegraphics[width=3.4in]{Fig25b.ps} \\end{array}$ \\vspace{-0.4cm} \\caption{These plot show the ratio of light odd-Z elements (Al and Na) to light even-Z elements/$\\alpha$-elements (Mg) as a function of [Fe/H] of solar neighbourhood halo and disk stars. The symbols are the same as in Fig.~\\ref{fig:light-s-over-abundances}. Al, Na and Mg abundances for First Stars Survey are NLTE-corrected according to \\citet{and07,and08,and10} in these two plots. [Al/Fe] for Barklem's sample are NLTE-corrected by $+0.6$ dex. Both plots show an odd--even effect, i.e. light odd-Z elements are suppressed at low metallicity.}\\label{fig:odd-even-effect} \\end{minipage} \\end{figure}  The fourth Cr--Mn component could be interpreted as the contribution from incomplete Si-burning, favoring the idea that the production of Cr and Mn relies on a different mass cut region as proposed by \\citet{nak99}. If this interpretation is correct, this might suggest that the production ratio of elements synthesized in the incomplete Si-burning region are strongly correlated among themselves, but their production ratios are decoupled from the production of other elements such as Zn, Co, and Ni that are synthesized in the complete Si-burning region. This component should however be viewed with caution. Only \\mbox{Cr\\,{\\sc i}} is detected in Barklem's sample, and both \\mbox{Cr\\,{\\sc i}} and Mn are suspected to require significant NLTE correction \\citep{ber08,ber10b}. Therefore this dimension could be spurious. Further discussion of this component requires the availability of a large homogeneous NLTE-corrected data set.  \\subsubsection[]{High metallicity} Stars in this metallicity range have ages up to about $10$ Gyrs, corresponding to more than a thousand generations of core-collapse supernovae \\citep[refer to fig.~8 in][]{kob09}. The PCA results for the high metallicity sample shown in Fig.~\\ref{fig:all-Reddy} can be summarized as follows:  The first component is a strong component ($40 \\%$ of the total variance) which consists of all light odd-Z elements, $\\alpha$-elements and Fe-peak elements. This could be interpreted as the loss of chemical dimensions after many generations of short-time-scale core-collapse supernovae as shown in \\citet{karl01} and \\citet{bla10}. After many generations of enrichment, elements that can be formed by short-time-scale core-collapse supernovae would have been well mixed in the ISM. In short, the first component is just a measure of the metal-richness of the gas reservoirs from which the stars formed. Much information was lost after the first few generations of supernovae.  It is also interesting that Al contributes to the first component. This is different from the low metallicity sample for which the Al does not contribute. This is not unexpected. As discussed,  Al is expected to be suppressed at low metallicity, but starts to contribute as metallicity increases.  The second component can be interpreted as an $r$-process contribution, since all $n$-capture elements contribute to this component in the same sign. Unlike the low-metallicity case, $\\alpha$-elements do not seem to contribute to this component. As discussed above, this may be due to the small parameter space of core-collapse supernovae that involve the $r$-process. As massive stars become fewer at high metallicity, most of the core-collapse supernovae will preferentially only produce $\\alpha$-elements but not $r$-process elements. These $\\alpha$-elements, as they continue to be produced and homogenized within the ISM, will only contribute to the hyperplane orthogonal to the $n$-capture elements. If this is the case, the PCA results will show a strong first component that includes supernovae ejecta and leaves the occasionally-produced $n$-captured elements to form a second component. This is also consistent with the results that the $s$-process in low-mass AGB stars starts to dominate $n$-capture elements production at high metallicity. Alternatively, the PCA results at high metallicity may imply that there is another $r$-process site that does not produce $\\alpha$-elements, such as neutron star mergers \\citep[e.g.][]{fre99b,ros99}  The third component includes the contribution from light and heavy $s$-process elements, and as discussed in the $n$-capture elements section,  could be interpreted as the low-mass AGB contribution. The appearance of $\\alpha$-elements constribution is unclear, but could be due to the envelope mass loss releasing the natal gas with heavier elements during the AGB phase.  \\subsection[]{Wider region of survey} As shown in Fig.~\\ref{fig:OC-MW-compare-all}, among the 12 elements that we investigated, compared to solar neighbourhood disk stars, we find about $1$ more dimension for open clusters which were sampled in a larger Galactic volume. The extra dimension for a larger Galactic volume is not surprising. Although stars from the solar neighbourhood could come from different regions of the galaxies via radial mixing, stars from open clusters far from the solar radius are unlikely to have undergone the same mixing processes as the solar neighbourhood stars. The larger Galactic volume ($6 \\la r_G \\la 20$ kpc for open clusters, compared to $7.5 \\la r_G \\la 8.5$ kpc for the solar neighbourhood) is likely to retain more variation in $\\mathcal{C}$-space.  \\subsection[]{K and Cu; \\textit{APOGEE};  the Ca-triplet region} In this section, we briefly discuss the effects of adding K and Cu to the list of elements, the dimensionality of the {\\it APOGEE} $\\mathcal{C}$-space, and the dimensionality of the $\\mathcal{C}$-space accessible from observations of the Ca-triplet region of the spectrum.  We did not study the contribution from K and Cu, because these elements are absent from most of our samples except for Reddy's high metallicity sample. The {\\it HERMES} bands are designed to include these two elements.  We found that adding K and Cu to the high metallicity sample only slightly increases the effective dimensionality-- from 6 (for 17 elements) to about 6.5 (for 19 elements)-- with K and Cu contributing equally. However, we suspect that K and Cu may add more dimensions at low metallicity because we can still distinguish contributions from different types of core-collapse supernovae. Cu is an interesting element; it is an odd Z elements preferably produced with hypernovae \\citep{kob06} and may have significant contribution from the weak $s$-process \\citep[like Ga \\& Ge,][]{pig10}. K may be affected by aspherical explosions \\citep{mae03} and also by the neutrino luminosity \\citep{kob11d}. The neutrino luminosity and weak $s$-process contribution could add promising extra dimensions at low metallicity.  In this study, we focused on the {\\it HERMES} $\\mathcal{C}$-space.  The {\\it APOGEE} high resolution near-infrared survey will measure the elements C, N, O, Na, Al, K, S, Mg, Si, Ca, Ti, V, Cr, Mn, Co, Ni but not $n$-capture elements, and we now estimate the dimensionality of this $\\mathcal{C}$-space.  We exclude C, N, O, Na as before because they may be subject to internal mixing. First, we considered the 9-dimensional $\\mathcal{C}$-space (Al, Mg, Ca, Ti, V, Cr, Mn, Co, Ni) measured by both Barklem's low metallicity sample and Reddy's high metallicity sample. We found that the low metallicity sample has about $4.5$ independent dimensions, whereas the high metallicity sample has $4$ independent dimensions. This is consistent with our previous observation that $n$-capture elements contribute 2 independent dimensions at high metallicity-- from $r$-process and $s$-process in low-mass AGB stars, but at low metallicity, $n$-capture elements only contribute $1$ -- $1.5$ dimensions. \\begin{table*} \\begin{center} \\caption{Summary of effective dimensions in various environments and $\\mathcal{C}$-subspace. \\label{tab:summary}} \\begin{tabular}{llll} \\hline Environment & Elements considered & Eff. dims. & No. of elements \\\\ \\hline & \\underline{{\\it Neutron-Capture Elements}} & \\\\ Solar neighbourhood halo stars & Y,Zr,Ba,La,Nd,Eu\t& 1 -- 2  & 6 \\\\ Solar neighbourhood disk stars  & Y,Zr,Ba,La,Nd,Eu\t& 2 & 6 \\\\ Fornax dSph galaxy & Y,Ba,La,Nd,Eu\t& 3 & 5 \\\\ \\\\ & \\underline{{\\it All Elements (HERMES)}} & \\\\ Solar neighbourhood halo stars & Al,Sc,Mg,Ca,Ti,V,Cr,Mn,Co,Ni,Zn,Y,Zr,Ba,La,Nd,Eu\t& 6 & 17 \\\\ Solar neighbourhood disk stars & Al,Sc,Mg,Si,Ca,Ti,V,Cr,Mn,Co,Ni,Zn,Y,Zr,Ba,Nd,Eu\t& 6 & 17 \\\\ Solar neighbourhood disk stars & Al,K,Sc,Mg,Si,Ca,Ti,V,Cr,Mn,Co,Ni,Cu,Zn,Y,Zr,Ba,Nd,Eu\t& 6 -- 7 & 19 \\\\ MW open clusters & Al,Sc,Mg,Si,Ca,Ti,V,Cr,Co,Ni,Y,Ba,Nd\t& 6 -- 7 & 13 \\\\ MW globular clusters & Mg,Si,Ca,Ti,Y,Ba,Eu & 4 & 7 \\\\ Fornax dSph galaxy & Mg,Si,Ca,Ti,Cr,Ni,Y,Ba,Nd,Eu\t& 6 & 10 \\\\ The Large Magellanic Cloud & Sc,Mg,Si,Ca,Ti,V,Cr,Co,Ni,Y,Zr,Ba & 6 -- 7 & 12 \\\\ \\\\ & \\underline{{\\it All Elements (APOGEE)}} & \\\\ Solar neighbourhood halo stars &  Al,Mg,Ca,Ti,V,Cr,Mn,Co,Ni & 4 -- 5 & 9 \\\\ Solar neighbourhood disk stars &  Al,Mg,Ca,Ti,V,Cr,Mn,Co,Ni & 4 & 9 \\\\ Solar neighbourhood disk stars &  Al,K,S,Mg,Si,Ca,Ti,V,Cr,Mn,Co,Ni & 5 & 12 \\\\ \\hline \\end{tabular} \\end{center} \\flushleft $^*$ We assume measurement uncertainty of $0.1$ dex and adopt $85 \\%$ cut-off for the ranked-eigenvalues cumulative percentages to derive the effective dimensions. This might not hold for open/globular clusters and dwarf galaxies as they have higher measurement/systematic uncertainty. Therefore the dimensions for open/globular clusters and dwarf galaxies in this table should be regarded as upper limit \\end{table*}  In addition to these 9 dimensions, Reddy's sample also included Si, S and K. We found that adding Si alone does not add to the effective dimensionality, as Si is strongly correlated with the other $\\alpha$-elements. However, we found that including S and K will add around $1$ additional dimension to the $\\mathcal{C}$-space. Finally we summarize the effective dimensions of the $\\mathcal{C}$-space (and subspaces) of various environments in Table~\\ref{tab:summary}.  The Ca-triplet region is widely used for large-scale spectroscopic surveys, so we also estimated the dimensionality of the $\\mathcal{C}$-space defined by this region of the spectrum.  For the metal-rich stars, the elements Al, Mg, Si, Ca, Ti and Ni define a space with 2 independent dimensions.  For the metal-poor stars, the elements Al, Mg, Ca, Ti and Ni define a space with 3 independent dimensions. Note that this estimate excludes possible extra dimensions from the light elements C, N and O."
    },
    "1112/1112.5220_arXiv.txt": {
        "abstract": "Some asteroids eject dust, unexpectedly producing transient, comet-like comae and tails.  First ascribed to the sublimation of near-surface water ice,  mass losing asteroids (also called ``main-belt comets'') can in fact be driven by a surprising diversity  of mechanisms.  In this paper, we consider eleven dynamical asteroids losing mass, in nine of which the ejected material is spatially resolved.  We address  mechanisms for producing mass loss including rotational instability, impact ejection, electrostatic repulsion, radiation pressure sweeping, dehydration stresses and thermal fracture, in addition to the sublimation of ice.  In two objects (133P and 238P) the repetitive nature of the observed activity leaves ice sublimation as the only reasonable explanation while, in a third ((596) Scheila), a recent impact is the cause. Another impact may account for activity in P/2010 A2 but this tiny object can also be explained as having shed mass after reaching rotational instability.  Mass loss from (3200) Phaethon  is probably due to cracking or dehydration at extreme ($\\sim$1000 K) perihelion temperatures, perhaps aided by radiation pressure sweeping.  For the other bodies, the mass loss mechanisms remain unidentified, pending the acquisition of more and better data.   While the active asteroid sample size remains small, the evidence for an astonishing diversity of mass loss processes in these bodies is clear. ",
        "introduction": "The classification of small bodies in the inner solar system as either asteroids or comets has historically been attempted by different scientists using different techniques and employing different criteria.  Observational astronomers classify small bodies having transient, unbound atmospheres (usually made visible by the scattering of sunlight from entrained micron-sized dust particles) as ``comets''.  Bodies having instead a constant geometric cross-section are called ``asteroids''.  To planetary scientists, comets and asteroids are distinguished by their ice content or perhaps by their formation location.  Comets are icy (because they formed beyond the ``snow-line'') while asteroids are not (supposedly because they formed at higher mean temperatures inside it).    Lastly, to dynamicists, comets and asteroids are broadly distinguished by a dynamical parameter, most usually the Tisserand parameter measured with respect to Jupiter (Kresak 1982, Kosai 1992).   It is defined by  \\begin{equation} T_J = \\frac{a_J}{a} + 2\\left[(1-e^2)\\frac{a}{a_J}\\right]^{1/2}\\cos(i) \\label{tisserand} \\end{equation}  \\noindent where $a$, $e$ and $i$ are the semimajor axis, eccentricity and inclination of the orbit while $a_J$ = 5.2 AU is the semimajor axis of the orbit of Jupiter.  This parameter, which is conserved in the circular, restricted 3-body problem, provides a measure of the close-approach speed to Jupiter. Jupiter itself has $T_J$ = 3. Main belt asteroids have $a \\le a_J$ and  $T_J >$ 3 while dynamical comets have $T_J < 3$.    The three systems of classification (observational, compositional and dynamical) are independent but imperfect.  For example, whether a coma or tail is detected on a given object depends strongly on the parameters of the observing system used.  A puny telescope may not reveal a coma that is easily rendered visible by  a more powerful one.  There is neither an agreed quantitative ice fraction with which to divide comets from asteroids nor, more importantly, any reliable way to measure the ice fraction in a small body.   The formation locations and dynamical histories of small bodies are rendered uncertain by hard-to-model non-gravitational forces from both electromagnetic radiation (the Yarkovsky effect) and mass loss itself (the rocket effect of Whipple 1950) and also by the chaotic nature of solar system dynamics.  The Tisserand criterion is an imperfect classifier because it is based on an idealized representation of the Solar system (e.g.~Jupiter's orbit is not a circle, the gravity of other planets is not entirely negligible, and so on).  Therefore, the utility of Equation (\\ref{tisserand}) as a dynamical discriminant is limited for objects with $T_J$ very close to 3.  As an example, the quasi-Hilda comets at $a \\sim$ 4.0 AU have $T_J \\sim$ 2.9 - 3.04 but are clearly interacting with Jupiter through the  3:2 mean-motion resonance (Toth 2006).  Some recognized Jupiter family comets (e.g. 2P/Encke with $T_J$ = 3.03) also fall in this category (Fernandez et al.~2002, Levison et al.~2006).    Given these and other imperfections it is remarkable that, for a majority of objects, the observational, compositional and dynamical definitions of asteroids and comets lie in close agreement.   For the most part, objects classified as asteroids (comets) based on their orbits, have the physical properties expected of asteroids (comets) as far as can be observed.  Exceptions have given rise to a somewhat confusing and evolving system of nomenclature, used to describe small solar system bodies by a combination of their orbital properties and physical appearances.  To clarify this we show, in Figure (\\ref{classification}), a two-parameter classification based on morphology, on the one hand, and the Tisserand parameter on the other.    Traditional comets lose mass and have  $T_J <$ 3.  Those with 2 $\\le T_J <$ 3 are called Jupiter family comets and are thought to originate in the Kuiper belt.  Comets with $T_J <$ 2 are long-period or Halley family comets, with a source in the Oort cloud.  Inactive counterparts to the Jupiter family comets are called, variously, extinct, dead or dormant comets (Hartmann et al.~1987).  They are presumably former comets in which past heating by the Sun has removed all near surface ice, although buried ice might remain and these objects could, in principle, reactivate.  Inactive counterparts to the long-period and Halley family comets are called Damocloids (Jewitt 2005).  Again, these are likely objects in which near-surface ice has been lost.  In this paper, we focus attention on the sub-set of  the small-bodies which are dynamically asteroid-like  ($a < a_J, T_J >$ 3) but which lose mass, like comets.   These were called ``main-belt comets'' by Hsieh and Jewitt (2006) but here we use the term  ``active asteroids'', since some of the examples to be considered, while dynamically asteroid-like and showing comet-like properties, are not in the main-belt.   Numerical integrations show that these are not recently captured comets from the Kuiper belt (Fernandez et al.~2002, Levison et al.~2006).  Scientific interest in these objects lies in the possibility that primordial water ice could have survived in asteroids despite early heating from embedded radioactive nuclei (Grimm and McSween 1989) and heating by the sun. Even greater interest is added by the possibility that the outer asteroid belt may have supplied part of the volatile inventory of the Earth (Morbidelli et al.~2000).  Additionally, active asteroids are a source of dust for the Zodiacal cloud, while unseen counterpart bodies may supply dust to  the debris disks of other stars (e.g. Shannon and Wu 2011).  After briefly summarizing the current observational evidence concerning these active asteroids, we discuss the surprisingly varied mechanisms through which a body is capable of losing mass.  A recent and complementary discussion focused on observational properties has been offered by Bertini (2011). ",
        "conclusions": "\\label{discussion}  Collisions are implicated in active asteroids both directly, as in the case of (596) Scheila and, perhaps, P/2010 A2, and indirectly as a trigger for activity (for example, to expose buried ice), as in 133P and 238P.  Here, we briefly examine the expected rate of collision between asteroids.  The typical collision probability per unit area in the asteroid belt is $P_c \\sim$ 3$\\times$10$^{-18}$ km$^{-2}$ yr$^{-1}$, with variations by a factor of several reflecting a collisional environment that varies with location in the belt  (Bottke et al.~1994).  The interval between impacts onto an asteroid of radius $r$ is  \\begin{equation} \\tau_c = \\frac{1}{\\pi r^2 P_c N_p (\\ge r_p)} \\label{tauc} \\end{equation}  \\noindent where $N( \\ge r_p)$ is the number of impactors larger than $r_p$.  Estimates of the size distribution of the asteroids are many and varied, with significant uncertainties resulting from the unmeasured albedos of most asteroids, as well as from severe observational bias effects (Jedicke et al.~2002).  The uncertainties are particularly acute for sub-kilometer asteroids because such objects are faint and remain largely unobserved.  For radii $r >$ 1 km, the best-fitting differential power law index is about -2, albeit with significant, size-dependent deviations from this value.  For radii $r  \\le$ 1 km, perhaps the best constraints on the distribution come  from the  impact crater size distribution on asteroid Gaspra.  There, craters from 0.4 km to 1.5 km in diameter (caused by projectiles perhaps 10 to 20 times smaller) are distributed as a power-law with a differential size index -3.7$\\pm$0.5 (Belton et al.~1992; note that Chapman et al.~(1996) report that \\textit{fresh} craters on Gaspra follow an even steeper distribution, with differential power law index -4.3$\\pm$0.3).  We assume that the total main-belt  population is $\\sim$1.4$\\times$10$^{6}$ asteroids with diameters $>$1 km.  Combining these results and integrating over the size distribution we take the number of projectiles with radius $\\ge r_p$ as  \\begin{equation} N_p (\\ge r_p) = 2.6\\times10^{13} \\left(\\frac{r_p}{1 m}\\right)^{-2.7} \\label{N} \\end{equation}  \\noindent where $r_p$ is expressed in meters.  Equations (\\ref{tauc}) and (\\ref{N}) together give   \\begin{equation} \\tau_c = 4200 \\left(\\frac{r_p}{1 m}\\right)^{2.7} \\left(\\frac{r}{1 km} \\right)^{-2}  ~[yr] \\label{tauc2} \\end{equation}  \\noindent  with $r$ in km.  Lines of constant collision time from Equation (\\ref{tauc2}) are plotted in Figure (\\ref{tauc_plot}). Impacted asteroid (596) Scheila is marked on Figure (\\ref{tauc_plot}) with an error bar indicating different estimates of the projectile radius from Jewitt et al.~(2011; $a \\sim$ 17 m) and Ishiguro et al.~(2011; $a \\sim$ 40 m), respectively. The corresponding collision times are  3$\\times$ 10$^{3}$ $\\le \\tau_c \\le$3$\\times$ 10$^{4}$ yr.  Given that there are $\\sim$250 known asteroids as large as or larger than (596) Scheila, the mean interval between similar events is $\\sim$(10 - 100) yr, statistically consistent with the detection of Scheila within the first decade of efficient, nearly real-time sky monitoring.  Disrupted asteroid P/2010 A2 (interpreted as a $\\sim$60 m radius object impacted by a projectile of characteristic size 2 to 4 meters; Jewitt et al.~2010) is also marked in Figure (\\ref{tauc_plot}), indicating a collision time $\\tau_c \\sim$ 7 to 50$\\times$10$^6$ yr.  With $N (\\ge 60) \\sim$ 4$\\times$10$^8$ (Equation \\ref{N}), the expected rate of similar events is $\\sim$8 to 60 yr$^{-1}$.   Jewitt et al.~(2011) estimated the detection probability of P/2010 A2 clones as $\\lesssim$6\\%, so that of 8 to 60 similar events per year we would currently detect only $\\sim$0.5 to 4. This is still higher than the actual rate of detection, counted as one object in perhaps a decade of efficient, nearly real-time sky monitoring by LINEAR and other survey telescopes.  The high rate could simply indicate the extreme uncertainty in using Equation (\\ref{N}) to estimate the number of meter-sized projectiles. For example, the equation gives $N_p (\\ge 5)$ = 3$\\times$10$^{11}$, while published estimates of this number range from $\\sim$10$^{9.5}$ to $\\sim$10$^{12}$ (Davis et al.~2002), indicating a large uncertainty in $\\tau_c$ at small projectile sizes.     Or it could indicate that activity in P/2010 A2 is caused by another process, possibly YORP spin-up, as remarked in Section (\\ref{p2010a2}).  Further progress in understanding the impact rate in the asteroid belt hinges strongly upon better measurements of the sub-kilometer size distribution.  Future wide-area sky surveys may help elucidate the mechanisms operating to cause asteroidal mass loss.  For example, we expect the spatial distribution of collisionally produced or triggered objects to be correlated with the regions of the asteroid belt in which the collision probability per unit time is highest.  Unfortunately, the published surveys for active asteroids so far are either biased (Hsieh 2009) or, if unbiased, detected no objects and so provide insufficient information (Gilbert and Wiegert 2009, 2010, Sonnett et al.~2011).  All but one of the known active asteroids were discovered serendipitously by a variety of surveys and methods, most too poorly quantified in terms of reported areal coverage and limiting magnitude to be used to interpret the spatial distribution.   This dismal situation deserves to be soon corrected by well-characterized, all-sky surveys having sensitivity sufficient to discover substantial numbers of active asteroids.  \\clearpage"
    },
    "1112/1112.0319_arXiv.txt": {
        "abstract": "We construct axisymmetric Schwarzschild models to measure the mass profile of the local group dwarf galaxy Fornax.  These models require no assumptions to be made about the orbital anisotropy of the stars, as is the case for commonly used Jeans models.  We test a variety of parameterizations of dark matter density profiles and find cored models with uniform density $\\rho_c=(1.6 \\pm 0.1) \\times 10^{-2} \\, M_{\\odot} \\text{ pc}^{-3}$ fit significantly better than the cuspy halos predicted by cold dark matter simulations. We also construct models with an intermediate-mass black hole, but are unable to make a detection.  We place a $1$-$\\sigma$ upper limit on the mass of a potential intermediate-mass black hole at \\mbh$\\leq 3.2 \\times 10^4 \\, M_{\\odot}$. ",
        "introduction": "Low-mass galaxies provide a unique testing ground for predictions of the cold dark matter (CDM) paradigm for structure formation, since they generally have a lower fraction of baryons than massive galaxies.  These galaxies allow for a more direct measurement of the underlying dark matter potential, as the complicated effects of baryons on the dark matter are less pronounced.  A particularly testable prediction of CDM is that all galaxies share a universal dark matter density profile, characterized by a cuspy inner power law $\\rho \\propto r^{-\\alpha}$ where $\\alpha=1$ (\\citealt{nav96}, hereafter \\citetalias{nav96}).  Many authors have investigated low-mass spirals and found, in contrast to the predictions of CDM, dark matter density profiles with a flat inner core of slope $\\alpha=0$ \\citep{bur95,per96,deb01,bla01,sim05}.  This has launched the debate known as the core/cusp controversy.  A number of other studies have investigated the mass content of dwarf spheroidal galaxies (dSphs). \\citet{gil07} give a comprehensive review of recent attempts to constrain the inner slope of their dark matter profiles with Jeans modeling (\\citealt{jea19}; \\citealt{bt87}, chapter 4). When significant, cored profiles are preferred for all dSphs modeled (\\citealt{gil07}, and references therein).  These results, however, are subject to a major caveat of Jeans modeling; it is complicated by the effect of stellar velocity anisotropy. Models fit to the line-of-sight component of the velocity dispersion, but anisotropy can severely affect the modeling of enclosed mass.  Therefore, additional assumptions must be made.  The studies presented in \\citet{gil07} assume spherical symmetry and isotropy.  \\citet{eva09} show that a weakness of Jeans modeling is that given these assumptions combined with the cored light profiles observed in dSphs, the Jeans equations do not allow solutions with anything other than a cored dark matter profile.  \\citet{wal09b} construct more sophisticated models and attempt to parameterize and fit for the anisotropy.  As a result, preference for cored profiles becomes model-dependent.  They therefore are unable to put significant constraints on the slope of the dark matter profile.  This highlights the main problem with Jeans modeling---it is highly dependent on the assumptions made.  Distribution function models are more general than Jeans models, and progress has been made applying them to a number of dSph systems \\citep{kle02,wu07,amo11}. Nevertheless these models still make strong assumptions such as spherical symmetry or isotropy, and models that do fit for anisotropy do so without using the information about the stellar orbits contained in the line-of-sight velocity distributions (LOSVDs).  We employ a fundamentally different modeling technique, known as Schwarzschild modeling, that allows us to use this information to self-consistently calculate both the enclosed mass and orbital anisotropy. Schwarzschild modeling is a mature industry, but one that has seldom been applied to the study of dSph galaxies (see \\citealt{val05}).  In addition to being well-suited for measuring dark matter profiles, Schwarzschild modeling has often been used to search for black holes at the centers of galaxies. Another unresolved issue relevant to the study of dSphs is whether they host an intermediate-mass black hole (IMBH).  In a hierarchical merging scenario, smaller galaxies are thought to be the building blocks of larger galaxies.  It is thought that all massive galaxies host a supermassive black hole (SMBH) at their center, therefore it is logical to believe that their building blocks host smaller IMBHs.  Evidence for these IMBHs is scarce, however, and dynamical detections are even scarcer. The closest and lowest mass example of a dynamical measurement is an upper limit on the local group dSph NGC~205 of \\mbh$\\leq 2.2 \\times 10^4 \\, M_{\\odot}$ \\citep{val05}.  Black holes in this mass range can provide constraints on theories of black hole growth and formation.  The two most prominent competing theories of nuclear black hole formation are direct collapse of primordial gas \\citep{ume93,eis95,beg06} or accretion onto and mergers of seed black holes resulting from the collapse of the first stars \\citep{vol05}.  In this paper we present axisymmetric, three-integral Schwarzschild models in an effort to determine the inner slope of the dark matter density profile as well as the orbit structure of the Fornax dSph.  We also investigate the possibility of a central IMBH.  We assume a distance of 135 kpc to Fornax \\citep{ber00}. ",
        "conclusions": "\\subsection{Cores and Cusps}  Our analysis shows that for the Fornax dwarf an NFW dark matter halo with inner slope $\\alpha = 1$ is rejected with high confidence.  We have kinematics from $30$ pc-$1.6$ kpc, and over this range the models prefer an $\\alpha = 0$ uniform density core with $\\rho_c = 1.6 \\times 10^{-2} \\, M_{\\odot} \\, \\mathrm{pc}^{-3}$.  We do not attempt to fit for models with an intermediate value of the slope $0 \\le \\alpha \\le 1$.  Further investigation is necessary before we can conclude that the best fitting dark matter profile is the logarithmic model.  The steep $\\alpha=1$ cusp of the NFW profile is, however, robustly ruled out.  The models, in general, seem to prefer less mass in the areas over which we have kinematic constraints.  In NFW models, the concentration $c$ sets the normalization (or y-intercept) of the density profile.  Because $c$ cannot be lowered below an astrophysically reasonable limit, NFW models enclose more mass than cored models.  This difference is reflected in the $\\chi^2$ difference between cored and NFW models, as the kinematics are best fit by models with less mass.  Figure \\ref{moments} hints at this as the best fit NFW model (dashed line) typically has higher values for $\\sigma$ than either the data or best-fitting cored model (solid line). %   Several groups have approached the core/cusp issue in dSphs by taking advantage of the fact that some dSphs host multiple populations of tracer stars that are chemically and dynamically distinct. By fitting models to each component, the underlying dark matter profile can be modeled more accurately.  \\citet{amo11b} fit two-component distribution function models to Sculptor, while \\citet{wal11} apply a convenient mass estimator (discussed below) to each stellar component in Sculptor and Fornax.  It is believed that this mass estimator is unaffected by orbital anisotropy, thus their method yields a robust determination of the dynamical mass at two locations in the galaxy---allowing for the slope of the dark matter profile to be measured. Each of these studies finds models with a cored dark matter halo preferable to the predicted cuspy NFW profile.  It must be noted, however, that we are not observing the pristine initial dark matter distribution in this galaxy.  Rather, it has likely been modified by complex baryonic processes over the lifetime of the galaxy.  These processes may include: adiabatic compression \\citep{blu86}, halo rebounding following baryonic mass loss from supernovae \\citep{nav96b}, or possibly dynamical friction acting on clumps of baryons (\\citealt{elz01}; but see also \\citealt{jar09}).  Although we chose this galaxy because these effects were likely to be small, they are nevertheless not well understood and our result must be taken in that context.  \\subsection{Central IMBH}  We are unable to place a significant constraint on the mass of a central IMBH.  Figure \\ref{chi2res} (upper right) shows the marginalized $\\chi^2$ curve against IMBH mass for cored dark matter density profiles.  The curve asymptotes to low values of IMBH, thus we are only capable of placing an upper limit on the mass of any potential IMBH.  Furthermore, our best-fit cored model with and without an IMBH have the same $\\chi^2$.  We therefore impose a $1$-$\\sigma$ upper limit on $M_{\\bullet} \\leq 3.2 \\times 10^4 \\, M_{\\odot}$. It is unfortunate that we are not able to place a lower limit on \\mbhs because measurements of black holes in the range \\mbh$\\lsim 10^4 \\, M_{\\odot}$ place direct constraints on SMBH formation mechanisms (van Wassenhove et al. 2010 ).  Our models, however, do robustly rule out a black hole of larger mass.  \\begin{figure}[t] \\includegraphics[width=9cm]{vplot.eps} \\caption{Anisotropy in various components of the velocity dispersion tensor. Shaded regions correspond to the 68\\% confidence regions, solid lines plot the best fit model.  Left and right hand panels plot stars near the major and minor axes, respectively. \\label{aniso}} \\end{figure}   In massive galaxies it is thought that the radius of influence, $R_{\\mathrm{inf}} \\sim G M_{\\bullet}/\\sigma^2$ must be resolved in order to detect and precisely measure a black hole \\citep{geb03, kor04, fer05,gul09}. Using our upper limit on $M_{\\bullet}$ we can calculate the maximum radius of influence of a potential black hole.  Estimating the central velocity dispersion at $\\sigma \\sim 10 \\text{ km s}^{-1}$ gives an upper limit for $R_{\\mathrm{inf}} \\lsim 14$ pc.  Our kinematics start at $R = 26$ pc, so it is not surprising that the minimum black hole mass we were able to detect has $R_{\\inf}$ close to $ 26$ pc. To detect smaller black holes, we require kinematics of stars closer to the center of the galaxy.  We are able to detect the dynamical influence of a black hole with a similar mass as \\citet{val05} detect in NGC~205, however with kinematics of much lower resolution.  Our innermost model bin is centered around $30 \\mathrm{~pc}$ whereas they use high-resolution kinematics from the \\emph{Hubble Space Telescope} to resolve spatial scales less than $1 \\mathrm{~pc}$.  The advantage we have is that the central velocity dispersion is much smaller in Fornax, which makes $R_{\\mathrm{inf}}$ larger for fixed \\mbh.  NGC~205 is also more than five times as distant as Fornax.  \\subsection{Mass Estimators}  Several authors have come up with convenient estimators of total mass within a given radius for local group dSphs.  \\citet{stri08} use the mass enclosed within $300$ pc while \\citet{wal09b} and \\citet{wol10} find a similar expression for the mass contained within the projected and un-projected half-light radii, respectively. These estimators bear striking resemblance to a result obtained by \\citet{cap06} derived from integral field kinematics of massive elliptical galaxies, and they all hint at an easy way to determine dynamical masses without expensive modeling.  They are believed to be insensitive to velocity anisotropy based on the derivation in \\citet{wol10}, and we compare their estimates to our models as a check on this.  For the mass contained within $300$ pc we measure $M_{300}=3.5^{+0.77}_{-0.11} \\times 10^6 M_{\\odot}$, roughly a factor of three smaller than \\citet{stri08} who measure $M_{300} = 1.14^{+0.09}_{-0.12} \\times 10^7 M_{\\odot}$ using Jeans models with parameterized anisotropy.  The \\citet{cap06}, \\citet{wal09b}, and \\citet{wol10} mass estimators are all of the form:  \\begin{equation} M(r_{\\mathrm{est}}) = k \\langle\\sigma^2_{LOS}\\rangle R_e \\label{esteq} \\end{equation}   \\ni where $r_{\\mathrm{est}}$ is the radius at which the estimator is valid. For \\citet{cap06} and \\citet{wal09b} $r_{\\mathrm{est}} = R_e$ (the projected half-light radius), while for \\citet{wol10} $r_{\\mathrm{est}} = r_e$ (the un-projected half-light radius). Other than the projected/un-projected difference, each estimator differes only by the value of the constant $k$.  In order to more fairly compare between these estimators and our models, we use the values for the luminosity-weighted line-of-sight velocity dispersion $\\langle\\sigma^2_{LOS}\\rangle=11.3^{+1.0}_{-1.8} \\, \\mathrm{ km~s}^{-1}$ projected half-light radius $R_e=689 \\, \\mathrm{pc}$, and un-projected half-light radius $r_e=900 \\, \\mathrm{pc}$ that we calculate from the data used in our models.   Our best-fitting model has $M(R_e)=3.9^{+0.46}_{-0.11} \\times 10^7 \\, M_{\\odot}$ and $M(r_e)=5.8^{+1.0}_{-0.2} \\times 10^7 \\, M_{\\odot}$. With each group's value for $k$ and our kinematics, the mass estimates are: $M(R_e) \\approx 5.1^{+1.0}_{-1.5} \\times 10^7 \\, M_{\\odot}$ \\citep{wal09b}, $M(r_e)\\approx 8.1^{+1.6}_{-2.4} \\times 10^7 \\, M_{\\odot}$ \\citep{wol10}, and $M(R_e)\\approx 1.0^{+0.3}_{-0.2} \\times 10^8 \\, M_{\\odot}$ \\citep{cap06}. Our model is broadly consistent with both the \\citet{wal09b} and \\citet{wol10} estimators.  The evidence that mass estimators are anisotropy-independent comes largely from comparison to spherical Jeans models (except \\citealt{cap06}). The weakness of these models is that the anisotropy must be parameterized and is restricted to be a function of radius only.  Our models are not subject to these constraints since the anisotropy is calculated non-parametrically and is free to vary with position angle.  We suggest that the best way to prove the accuracy of mass estimators is to compare with models that can self-consistently calculate both mass and anisotropy for realistic potentials.  For bright elliptical galaxies, \\citet{cap06} and \\citet{tho11} have done just that.  In these cases, the mass estimates are checked against masses derived from axisymmetric Schwarzschild modeling and good agreement is found.  Ours is the first study to perform a similar test with dSphs, and there is no reason to assume that success with bright ellipticals guarantees accuracy in the dSph regime.  The results from our comparison above are nevertheless reassuring.   \\subsection{Tidal Effects}  The principle of orbit superposition, and hence our entire modeling procedure, relies on the assumption that the galaxy is bound and in a steady state. The amount of tidal stripping in Fornax due to the effect of its orbit through the Milky Way's halo is not well-known.  For reasonable values of Fornax total mass $m$, Milky Way mass $M$, and Galactocentric radius $R_0$, the tidal radius of Fornax is $r_t \\sim (m/3M)^{1/3} R_0 \\sim 13.5 \\text{ kpc}$.  This estimate of $r_t$ is sufficiently larger than our model grid that we would not expect tidal effects to be important over the radial range of our models.  If Fornax is on an eccentric orbit about the Milky Way, however, the above equation for $r_t$ is not valid and estimation of the tidal radius is not as straightforward.  Fortunately, studies investigating its transverse motion suggest the orbit of Fornax is roughly circular \\citep{pia07,wal08}."
    },
    "1112/1112.2770_arXiv.txt": {
        "abstract": "An emerging theme in modern astrophysics is the connection between astronomical observations and the underlying physical phenomena that drive our cosmos.  Both the mechanisms responsible for the observed astrophysical phenomena and the tools used to probe such phenomena -- the radiation and particle spectra we observe -- have their roots in atomic, molecular, condensed matter, plasma, nuclear and particle physics.  Chemistry is implicitly included in both molecular and condensed matter physics.  This connection is the theme of the present report, which provides a broad, though non-exhaustive, overview of progress in our understanding of the cosmos resulting from recent theoretical and experimental advances in what is commonly called laboratory astrophysics.  This work, carried out by a diverse community of laboratory astrophysicists, is increasingly important as astrophysics transitions into an era of precise measurement and high fidelity modeling. ",
        "introduction": "Laboratory astrophysics and complementary theoretical calculations are the foundations of astronomy and astrophysics and will remain so into the foreseeable future.  The impact of laboratory astrophysics ranges from the scientific conception for ground-based, airborne and space-based observatories, all the way through to the scientific return of these projects and missions.  It is our understanding of the underlying physical processes and the measurement or calculation of critical physical parameters that allows us to address fundamental questions in astronomy and astrophysics.  The field of laboratory astrophysics comprises both theoretical and experimental studies of the underlying physics that produce the observed astrophysical processes.  We have identified six areas of physics as relevant to astronomy and astrophysics.\\footnote{The authors comprise past and current members of the American Astronomical Society Working Group on Laboratory Astrophysics.} Astronomy is an observational science focused primarily on detecting photons generated by atomic, molecular and condensed matter physics.  Chemistry is implicitly included here as part of molecular and condensed matter physics.  Our understanding of the universe also relies on knowledge of the evolution of matter (nuclear and particle physics) and of the dynamical processes shaping it (plasma physics).  Planetary science, involving {\\it in-situ} measurements of solar system bodies, requires knowledge from atomic, molecular, condensed matter and plasma physics. Hence, our quest to understand the cosmos rests firmly on scientific knowledge in six areas: atomic, molecular, condensed matter, plasma, nuclear and particle physics.  Here we review recent advances in our astrophysical understanding of the cosmos arising from work in laboratory astrophysics.  We focus primarily on the past decade.  Our work complements that of previous reviews on laboratory astrophysics in atomic physics (Beiersdorfer 2003; Kallman and Palmeri 2007; International Astronomical Union [IAU] Commission 14), molecular physics (Salama 1999; Tielens 2005; Herbst and van Dishoeck 2009; IAU Commision 14), condensed matter physics (Draine 2003; Whittet 2003), plasma physics (Drake 1999; Remington \\etal 2006; Zweibel and Yamada 2009; Yamada \\etal 2010), nuclear physics (K\\\"appeler \\etal 2011; Wiescher \\etal 2010; Adelberger \\etal 2011) and particle physics (Grupen 2005; Aprile and Profumo 2009).  Because laboratory astrophysics, as implied by its name, is astrophysically motivated, we have structured our report into five broad categories which blanket the field of astronomy and astrophysics.  This helps to bring out the synergy between the various subareas of laboratory astrophysics.  The specific categories are: planetary systems and star formation (section~\\ref{sec:PSF}), stars and stellar evolution (section~\\ref{sec:SSE}), the galactic neighborhood (section~\\ref{sec:GAN}), galaxies across time (section~\\ref{sec:GAT}) and cosmology and fundamental physics (section~\\ref{sec:CFP}).  This structure parallels the scientific divisions used by the rececent U.S.\\ National Research Council Astro 2010 Survey on Astronomy and Astrophysics (Blandford \\etal 2010a). These five sections are further subdivided into relevant subareas of laboratory astrophysics.  Space limitations prevent these subsections from being exhaustive.  Rather they are aimed at giving the reader an overview of recent successes in the field and appropriate citations to provide entry into the relevant research.  We conclude with a brief discussion and outlook for the future in section~\\ref{sec:end}. ",
        "conclusions": "\\label{sec:end}  Our astrophysical understanding of the cosmos continues to be propelled forward by advances in laboratory astrophysics.  This review has touched on many, but far from all, of the achievements of the past decade.  The coming decade promises to be equally, if not more, fruitful.  The Astro 2010 Survey Report and Panel Reports (Blandford \\etal 2010a,b) have laid out a series of exciting scientific objectives, the achievement of which they point out are going to require numerous advances in laboratory astrophysics.  We direct the reader to those reports for a detailed discussion.  Additional in depth discussions about the laboratory astrophysics needs and opportunities for the coming decade can be found in a number of White Papers written over the past few years.  These include those submitted by the Working Group on Laboratory Astrophysics (WGLA) to the Astro 2010 Survey (Brickhouse \\etal 2009a,b,c,d,e) as well as community input to Astro2010 through the Science White Papers ({\\tt http://sites.nationalacademies.org/BPA/BPA\\_050603}) and the Laboratory Astrophysics White Papers ({\\tt http://sites.nationalacademies.org/BPA/BPA\\_051118}).  Another White Paper is that submitted by the WGLA to the U.S. National Research Council Planetary Science Decadal Survey: 2013-2022 (Gudipati \\etal 2009).  In plasma laboratory astrophysics, there have been a couple of reports recently released by the community (Prager \\etal 2010; Rosner and Hammer 2010).  And most recently there is the White Paper from the 2010 Laboratory Astrophysics Workshop sponsored by the Astrophysics Division of the Science Mission Directorate which covered atomic, molecular, condensed matter and plasma laboratory astrophysics (Savin \\etal 2011).  These all point the way to the future and the richness of discovery which we can only just begin to guess.  \\ack  The authors thank their many colleagues including J.\\ E.\\ Bailey, P.\\ Beiersdorfer, G.\\ V.\\ Brown, J.\\ R.\\ Crespo L\\'opez-Urrutia H.\\ Ji, H.\\ Kreckel, J.\\ E.\\ Lawler, M.\\ Medvedev, T.\\ Plewa, D.\\ Sasselov, R.\\ K.\\ Smith, C.\\ Sneden, B.\\ J.\\ Wargelin and S.\\ Widicus Weaver for stimulating conversations. NSB was supported in part by NASA contract NAS8-03060 to the Smithsonian Astrophysical Observatory for the Chandra X-ray Center. JJC is supported in part by the National Science Foundation through grant AST 0707447. RPD acknowledges support from DOE/NNSA Defense Sciences and Advanced Scientific Computing, from DOE/Science Office of Fusion Energy Sciences and from the Defense Threat Reduction Agency. GJF acknowledges support by NSF (0908877), NASA (07-ATFP07-0124, 10-ATP10-0053, and 10-ADAP10-0073) and STScI (HST-AR-12125.01 and HST-GO-12309). MSG acknowledges funding from NASA Astrobiology Institute ``Icy Worlds'' and support from the Jet Propulsion Laboratory, California Institute of Technology, under a contract with the National Aeronautics and Space Administration. WCH was supported in part by the US Department of Energy under grant DE-SC00046548 to the University of California at Berkeley. EH acknowledges the support of NASA through its program in laboratory astrophysics and through the {\\it Herschel} program. SP is partly supported by the US Department of Energy with an Outstanding Junior Investigator Award and by Contract DE-FG02-04ER41268 and by NSF Grant PHY-0757911. FS acknowledges the support of the Astrophysics Research and Analysis Program of NASA Science Mission Directorate. DWS is supported in part by the NASA Astronomy and Physics Research and Analysis program, the NASA Solar Heliospheric Physics program, and the NSF Division of Astronomical Sciences Astronomy and Astrophysics Grants program. EGZ was supported in part by the NSF grant PHY-0821899 to the Uiversity of Wisconsin.  \\pagebreak  \\appendix"
    },
    "1112/1112.5961_arXiv.txt": {
        "abstract": "{ Adiabatic modeling of solar-like oscillations cannot exceed a certain level of precision for fitting individual frequencies. This is known as the problem of near-surface effects on the mode physics. We present a theoretical study which addresses the problem of frequency precision in non-adiabatic models using a time-dependent convection treatment.  We find that the number of acceptable model solutions is significantly reduced and more precise constraints can be imposed on the models.  Results obtained for a specific star ($\\beta$ Hydri) lead to very good agreement with both global and local seismic observables. This indicates that the accuracy of model fitting to seismic data is greatly improved when a more complete description of the interaction between convection and pulsation is taken into account.}  \\abstract*{ Adiabatic modeling of solar-like oscillations cannot exceed a certain level of precision for fitting individual frequencies. This is known as the problem of near-surface effects on the mode physics. We present a theoretical study which addresses the problem of frequency precision in non-adiabatic models using a time-dependent convection treatment.  We find that the number of acceptable model solutions is significantly reduced and more precise constraints can be imposed on the models.  Results obtained for a specific star ($\\beta$ Hydri) lead to very good agreement with both global and local seismic observables. This indicates that the accuracy of model fitting to seismic data is greatly improved when a more complete description of the interaction between convection and pulsation is taken into account.} ",
        "introduction": "Two conditions are necessary for precise probing of the physics of stellar interiors with asteroseismology.  One condition concerns the quality of the observed data.  Space missions have resulted in data of unprecedented quality and improved data analysis.  The other condition is that models should be able to reproduce the complex physics at work in these stars, allowing for a precise interpretation of the observations.  The treatment of convection is one of the largest sources of uncertainty in stellar astrophysics.  For cool pulsating stars, the convective envelope introduces severe complications for modelling the oscillations. In these stars, the interaction between pulsation and convection is very important for proper modeling of the oscillations. For example, this interaction has been found to be a major source of damping in $\\delta$~Scuti stars \\cite{2005AA...435..927D} and in driving of oscillations in $\\gamma$~Doradus stars \\cite{2004AA...414L..17D}.  For stars with solar-like pulsations, it turns out that the transition region where the thermal relaxation time is of the same order as the pulsation period is located inside the superficial convective zone. Since these near-surface layers have not been properly modeled, they give rise to discrepancies between the computed and observed frequencies \\cite{1988AA...200..213D}.  Gabriel's formalism of time-dependent convection (TDC) \\cite{1996BASI...24..233G, 2005AA...434.1055G}, as implemented in the MAD code \\cite{2001AA...366..166D}, leads to very important results for models of $\\delta$~Sct and $\\gamma$~Dor stars \\cite{2004AA...414L..17D, 2005AA...435..927D} and also for the solar damping rates \\cite{2006ESASP.624E..97D}.  In this work, we present TDC results for models of $\\beta$ Hydri, a G2\\,IV sub-giant showing solar-like oscillations. A recent list of observed pulsation frequencies in this star is given in Brand{\\~a}o et al. \\cite{2011AA...527A..37B}. We compare these frequencies with calculations from adiabatic models with and without the near-surface corrections of Kjeldsen et al. \\cite{2008ApJ...683L.175K}. ",
        "conclusions": "The main results of this work are as follows.  (i) TDC leads to a significant improvement in the agreement between calculated and observed frequencies  of solar-like oscillations as compared to the pure adiabatic frequencies. (ii) A physically more robust modeling of the mode physics near the surface provides better constraints on the global and free stellar parameters. (iii) Near-surface corrections are probably only valid for stars most closely resembling the Sun."
    },
    "1112/1112.3295_arXiv.txt": {
        "abstract": "{Solar-like oscillations have been observed by {{\\it Kepler}} and CoRoT in several solar-type stars.} {We study the variations of stellar  p-mode linewidth as a function of effective temperature.} {Time series of 9 months of {{\\it Kepler}} data have been used.  The power spectra of 42 cool main-sequence stars and subgiants have been analysed using both Maximum Likelihood Estimators and Bayesian estimators, providing individual mode characteristics such as frequencies, linewidths and mode heights.} {Here we report on the mode linewidth at maximum power and at maximum mode height for these 42 stars as a function of effective temperature.} {We show that the mode linewidth at either maximum mode height or maximum amplitude follows a scaling relation with effective temperature, which is a combination of a power law plus a lower bound.  The typical power law index is about 13 for the linewidth derived from the maximum mode height, and about 16 for the linewidth derived from the maximum amplitude while the lower bound is about 0.3 $\\mu$Hz and 0.7~$\\mu$Hz, respectively.  We stress that this scaling relation is only valid for the cool main-sequence stars and subgiants, and does not have predictive power outside the temperature range of these stars.} ",
        "introduction": "\\label{sec:intro}  Stellar physics faces a revolution following the great wealth of asteroseismic data made available by space missions such as CoRoT \\citep{Baglin2006} and {\\it Kepler} \\citep{Borucki2009}.   Long observations of solar-like pulsators corresponding to main sequence stars, subgiants and red giants have been performed during more than 6 months by CoRoT \\citep[][and references therein]{Baudin2011,Baudin2011C}.  The {\\it Kepler} mission now provides a larger sample of stars observed for longer durations \\citep{Chaplin2011}.  The study of oscillation mode physics (mode height, linewidth and amplitude) provides information on excitation and damping mechanisms related to the physics of convection and of stellar atmospheres \\citep{Samadi2009}. \\citet{GH99} theoretically derived stellar mode linewidths as a function of stellar mass and age.  They found that stellar mode linewidths would present a depression or {\\it plateau} close to the maximum of mode height.  This depression was first detected in the solar p-mode linewidths by \\citet{CFJR95}.  Such a depression is due to a resonance between the thermal adjustment time of the superadiabatic boundary layer and the mode frequency \\citep{B92a}.   % The frequency location of the maximum of mode height is in turn related to the Mach number (${\\cal M}_a$), the ratio of convective velocity to the sound speed \\citep{Belkacem2011}.  The convective flux giving the maximum mode amplitude is also related to ${\\cal M}_a$ to the power of 3 \\citep[e.g.][]{Belkacem2011,GH99}.  It is therefore interesting to study how the mode linewidth is related to the frequency of maximum amplitude / mode height for several different stars.  Statistical studies over a large number of stars have been performed in order to  validate the scaling relation derived for the amplitude of stellar oscillations by \\citet{HK1995} and recently revised by \\citet{KB2011}.    Scaling relations for mode linewidth have been proposed by \\citet{Chaplin2009} and \\citet{Baudin2011} based upon the stellar effective temperature.   \\citet{Chaplin2009} proposed a scaling relation with linewidth proportional to $T_{\\rm eff}^4$ based upon several ground-based observations.  Using CoRoT observations, \\citet{Baudin2011C} measured linewidths for a sample of solar-like pulsators and provided a scaling relation proportional to $T_{\\rm eff}^{16}$.  With the ability to perform longer observations of stars with {\\it Kepler}, the measurement of mode linewidth becomes easier and more reliable. In this paper, based upon {\\it Kepler} observations and using a larger stellar sample than of \\citet{Baudin2011}, we derive a new relation between mode linewidth and $T_{\\rm eff}$ . ",
        "conclusions": "We studied the dependence of linewidth at maximum mode height and amplitude on $T_{\\rm eff}$ for two sources effective temperature and for two ways of deriving the linewidth.  We showed using 9 months of {\\it Kepler} observations of 42 stars that the mode linewidth at both maximum mode height or maximum amplitude follows a scaling relation based on effective temperature, which is a combination of a power law plus a lower bound.  We stress that this scaling relation is only valid for the cool main-sequence and subgiant stars, and does not have predictive power outside the temperature range of these stars."
    },
    "1112/1112.2220_arXiv.txt": {
        "abstract": "The major morphological features of a galaxy are thought to be determined by the assembly history and net spin of its surrounding dark halo. In the simplest scenario, disk galaxies form predominantly in halos with high angular momentum and quiet recent assembly history, whereas spheroids are the slowly-rotating remnants of repeated merging events. We explore these assumptions using one hundred systems with halo masses similar to that of the Milky Way, identified in a series of cosmological gasdynamical simulations: the Galaxies - Intergalactic Medium Calculation ({\\sc gimic}). At $z=0$, the simulated galaxies exhibit a wide variety of morphologies, from dispersion-dominated spheroids to pure disk galaxies. Surprisingly, these morphological features are very poorly correlated with their halo properties: disks form in halos with high and low net spin, and mergers play a negligible role in the formation of spheroids, whose stars form primarily in-situ. With hindsight, this weak correlation between halo and galaxy properties is unsurprising given that a minority of the available baryons ($\\sim 40\\%$) end up in galaxies. More important to morphology is the {\\it coherent alignment of the angular momentum} of baryons that accrete over time to form a galaxy. Spheroids tend to form when the spin of newly-accreted gas is misaligned with that of the extant galaxy, leading to the episodic formation of stars with different kinematics that cancel out the net rotation of the system. Disks, on the other hand, form out of gas that flows in with similar angular momentum to that of earlier-accreted material. Gas accretion from a hot corona thus favours disk formation, whereas gas that flows ``cold'', often along separate, misaligned filaments, favours the formation of spheroids. In this scenario, many spheroids consist of the superposition of stellar components with distinct kinematics, age, and metallicity, an arrangement that might survive to the present day given the paucity of major mergers. Since angular momentum is acquired largely at turnaround, morphology depends on the early interplay between the tidal field and the shape of the material destined to form a galaxy. ",
        "introduction": "\\label{SecIntro}  Galaxies exhibit a wide variety of morphologies, from spheroids to disks to bars to peculiar galaxies of irregular shape. Many physical properties, such as gas content, average stellar age, and the rate of current star formation, are known to correlate with morphology. Of such properties, the one that seems most tractable from a theoretical perspective is the relative importance of organized rotation in the structure of a galaxy. This is commonly referred to as the disk-to-spheroid ratio, since stellar disks are predominantly rotationally-flattened structures whereas spheroids have shapes largely supported by velocity dispersion.  Since \\citet{Hubble1926} published his original morphological classification scheme, our understanding of the provenance of these two defining features of galaxy morphology has been constantly evolving. Spheroids were once thought to originate in the swift transformation of an early-collapsing, non-rotating cloud of gas  into stars \\citep{ELS1962,Partridge1967,Larson1974}, whereas disks were envisioned to result from the collapse of clouds with high angular momentum and inefficient star formation \\citep{ELS1962,Larson1976}. The role of mergers as a possible transformational mechanism was championed by \\citet{Toomre1977} and gained momentum as the hierarchical nature of structure (and hence, galaxy) formation became accepted \\citep{White_Rees1978,Frenk1985}.  Further development of these ideas led to a broad consensus where disks are thought to form at the center of dark matter halos as a consequence of angular momentum conservation during the dissipational collapse of gas \\citep{Fall1980,Mo1998}, whereas spheroids result predominantly from merger events \\citep[see, e.g.,][and references therein]{Cole2000}. Morphology is thus a transient feature of the hierarchical formation of a galaxy: a disk galaxy may be transformed into a spheroidal one after a major merger, but could then re-form a disk through further gas accretion only to be later disrupted again by another merger. Early galaxy formation simulations gave a visually-compelling demonstration of this scenario, galvanizing support for it \\citep[see, e.g.,][]{SteinmetzNavarro2002}.  This consensus view has been broadly implemented in semi-analytic models of galaxy formation, where the properties of galaxies are deduced directly from the physical properties and assembly history of their surrounding halos \\citep[see, e.g.,][]{Croton2006,Bower2006,Somerville2008}. For example, most models assume that the specific angular momentum of galaxies and halos are similar, and that the merger history of the halos dictates that of the central galaxy.  Recent developments, however, have led to revisiting some of the assumptions of the simple scenario outlined above. For example, it has become clear that major mergers are rare, and therefore probably not the primary formation mechanism of bulges and ellipticals. Instead, ``disk instabilities'' \\citep{Efstathiou1982,Christodoulou1995,Mo1998}, as well as repeated minor encounters, are now claimed to be the main formation path of spheroids \\citep[e.g.,][]{Parry2009,Hopkins2010,DeLucia2011,Bournaud2011}. This has helped to alleviate some tension between the observed evolution of the early-type galaxy population and the major-merger rates predicted by theory \\citep{Bundy2007,Oesch2010}. However, questions might remain open, as the estimation of merger times from observations is non-trivial \\citep{Lotz2011}.  Further scrutiny has come from direct simulation of hierarchical galaxy formation. Conserving enough angular momentum during the hierarchical assembly of a galaxy to form a realistic stellar disk has been challenging \\citep[see, e.g.,][]{Navarro1995,NavarroSteinmetz1997}, as has been pinning down the effect on morphology of repeated merging, especially between gas-rich galaxies \\citep[see, e.g.,][]{Robertson2006,Governato2009}.  The inclusion of energetic feedback, needed to prevent the formation of too many faint or overly massive galaxies, has added an extra level of complexity to the problem, with a number of studies showing that morphologies can be radically altered when even modest changes in the strength of feedback or its implementation are introduced \\citep{Okamoto2005,Scannapieco2008,Ceverino2009,Sales2010,Agertz2011,Brook2011,Piontek2011}. Moreover, the density threshold assumed for star formation also changes the coupling between the stellar winds and the surrounding gas, playing also a major role on the properties of simulated galaxies\\citep{Guedes2011}.   More recently, the {\\it mode} of gas accretion has been recognized as playing a potentially crucial role in galaxy morphology. Gas can flow to galaxies largely unimpeded by shocks \\citep{White1991} and may be collimated by the filamentary structure of the cosmic web, especially in low-mass systems and at high redshift \\citep{Keres2005,Dekel2006,vandeVoort2011}. This complex accretion geometry has been hypothesized to promote the formation of disks by feeding high angular momentum material directly to forming galaxies \\citep[see, e.g.,][]{Dekel2009,Brooks2009}.   Further theoretical progress demands increased sophistication in numerical and semi-analytic modeling. From the simulation perspective, most studies have focussed on individual systems picked according to what the authors believe would facilitate the formation of a galaxy of predetermined morphology; for example, a recent major merger to study ellipticals \\citep[e.g.,][]{Meza2003} or a quiet, rapidly-rotating halo to study spirals \\citep[e.g.,][]{Abadi2003a,Governato2007}. Note that this presupposes the morphology of the resulting galaxy, and often results in the tuning of star formation and feedback parameters until, unhelpfully, results match prejudice.  Statistically-significant samples of galaxies selected in an unbiased way and simulated at high resolution are needed for new insights, a goal that, despite valiant efforts \\citep{Croft2009, Sales2009c, Sales2010}, has so far proved beyond reach of even the fastest computers and best algorithms.  The situation, however, is starting to change, with the advent of simulations of volumes large enough to include dozens of well-resolved $\\sim L_\\star$ galaxies \\citep{Crain2009,Schaye2010,Hahn2010,Cen2011,Vogelsberger2011}.  We explore these issues here using the {\\sc gimic} gasdynamical simulation series \\citep{Crain2009}.  {\\sc gimic} targeted several carefully selected regions from the Millennium Simulation \\citep{Springel2005a} in an attempt to maximize the resolution of individual galaxy systems while at the same time sampling a cosmologically-significant volume. The first analyses of $z=0$  {\\sc gimic} galaxies \\citep{Crain2010,Font2011b,McCarthy2011} show that they are fairly realistic, so we feel confident that we can use them to gain insight into the origin of galaxy morphology.  This paper is organized as follows. In Sec.~\\ref{sec:numexp} we present briefly the numerical method and simulations. We analyze the morphologies of simulated galaxies and their origin in Secs.~\\ref{SecResults} and \\ref{SecOrGxMorph}, respectively. We summarize our main conclusions in Sec.~\\ref{sec:concl} ",
        "conclusions": "\\label{sec:concl}  We use gasdynamical cosmological simulations of galaxy formation to study the origin of different galaxy morphologies in the $\\Lambda$CDM cosmogony at redshift $z=0$.  The {\\sc gimic} simulation series covers a large volume and has a resolution high enough to study the structure and kinematics of the stellar components of $100$ central galaxies in Milky Way-sized halos. We focus our analysis on the origin of galaxy morphology, somewhat narrowly defined as the relative importance of rotational support vs velocity-dispersion support (the disk-to-spheroid ratio) in the structure of the galaxy. Our main results may be summarized as follows:  \\begin{itemize}  \\item The simulated galaxies span a wide range of morphological types, from rotation-free spheroids to almost pure disk galaxies where fewer than $5\\%$ of all stars are in counterrotating orbits. Disks have roughly exponential stellar surface density profiles and flat rotation curves, whereas spheroids are dense stellar systems that can be approximated by de Vaucouleurs' $R^{1/4}$ profiles. The resemblance with real galaxies suggests that it should be possible to gain insight into the origin of galaxy morphology by studying the mechanisms responsible for the relative importance of disks and spheroids in {\\sc gimic} galaxies.  \\item The morphology of simulated galaxies seems mostly unrelated to the spin or assembly history of their surrounding dark matter halos. Most stars form {\\it in-situ} and comprise on average about $\\sim 40\\%$ of all available baryons in the halo. Most baryons in a halo therefore end up not making part of the central galaxy, which helps to explain the weak correlation between the properties of halos and those of central galaxies.  Contrary to simple expectations, disks form in halos with low and high angular momenta, and spheroids form even in galaxies where most stars form {\\it in-situ}, suggesting a formation path for spheroids that does not rely on merging.  \\item The star formation history provides an interesting clue to the origin of morphology. Disks tend to have young stars, and to form gradually over long periods of time. This is because gradual cooling from a hot corona delays the accretion of gas and promotes late star formation.  \\item Star formation in spheroids proceeds episodically, leaving behind populations of stars of similar age but distinct kinematics. These populations originate from the accretion of gas whose angular momentum is misaligned relative to that of earlier-accreted material. The misalignment destabilizes any pre-existing disk, prompts the rapid transformation of gas into stars, and reduces the net rotational support of the system.  \\item Since angular momentum is largely acquired at the time of maximum expansion of the material destined to form a galaxy, a good indicator of morphology at $z=0$ is the coherence in the alignment of the net spin of various parts of the system at the time of turnaround. Spheroid-dominated galaxies form in systems where misalignments are substantial whereas disks form in systems where the angular momentum of all mass shells is roughly aligned.  \\item The final morphology of a galaxy results from the combined effects of spin alignment and of hot/cold gas accretion. Disk-dominated objects are made of stars formed predominantly {\\it in situ}, and avoid systems where most baryons were accreted cold, or those where spin misalignments are extreme. On the other hand, direct filamentary accretion of cold gas, especially when accompanied by substantial spin misalignments, favours the formation of slowly-rotating spheroids, which may thus form even in the absence of mergers.  \\end{itemize}  Our results suggest a new scenario for the origin of $\\sim L_*$ stellar spheroids that does not rely on merging. This scenario, once developed more thoroughly, should offer a number of predictions falsifiable by observation. For example, the episodic nature of star formation in spheroids envisioned here is expected to leave behind overlapping populations of stars of distinct age, kinematics, and, possibly, metallicity) that survive to the present because of the paucity of mergers. We plan to explore the observational signatures of these populations in future work.  The scenario we propose here also offers clues to the origin of pure disk galaxies. A number of our simulated galaxies have virtually no ``classical'' spheroid, with fewer than $5\\%$ of their stars in counterrotating orbits. These galaxies form either in systems where spin alignment is extraordinarily coherent or where most of the baryons in the galaxy get accreted late from a hot corona. Although at this time limited numerical resolution precludes a more detailed study, we plan to use these clues to resimulate some of these systems at higher resolution with the goal of shedding light into the origin of bulgeless galaxies.  We emphasize that, although coherent spin alignment at early times is clearly an important clue, it should be considered as one ingredient of the complex process that determines the morphology of a galaxy. Strong feedback, for example, may expel baryons from galaxies and cycle them through a hot corona before they get re-accreted and turned into stars, potentially erasing the spin alignment dependence at turnaround we report here. Furthermore, aligned spins in the accreting gas might not be enough to ensure the survival of a stellar disk, especially if the dark matter halo is strongly triaxial and its principal axes are not coincident with the disk. Finally, although mergers are rare in the mass range we explore here, they likely play a more important role in the formation of more massive spheroids. Until simulations can reproduce not only the properties of individual systems, but the full statistical distribution of galaxy morphologies and their dependence on mass and environment, it is likely that a full understanding of the origin of galaxy morphology will remain beyond reach."
    },
    "1112/1112.2698.txt": {
        "abstract": "We have calculated improved photonic passbands for the $UBVRI$, Hipparcos and Tycho $H$$_{p}$, $B$$_{T}$, $V$$_{T}$ standard systems using the extensive spectrophotometric libraries of NGSL and MILES. Using the $H_{p}$ passband, we adjusted the absolute flux levels of stars in the spectrophotometric libraries so their synthetic H$_{p}$ magnitudes matched the precise Hipparcos catalog value. Synthetic photometry based on the renormalized fluxes were compared to the standard $UBVRI$ and $B_{T}$, $V_{T}$ magnitudes and revised synthetic zero-points  were determined. The Hipparcos and Tycho photometry system zero-points were also compared to the $V$ magnitude zero-points of the SAAO $UBVRI$ system, the homogenized $UBV$ system and the Walraven $VB$ system.  The confusion in the literature concerning broadband magnitudes, fluxes, passbands and the choice of appropriate mean wavelengths is detailed and discussed in an appendix. ",
        "introduction": "The Hipparcos Catalogue \\citep{Perr97} is a high precision photometric (plus parallax and proper-motion) catalog of more than 100,000 stars measured with the $H_{p}$ band; the Tycho2 catalog \\citep{Hog00} contains 2.5 million stars measured (mostly) with lower precision in the $B_{T}$ and $V_{T}$ bands. The remarkable collection of data was obtained during the 4 year (1989--1993) mission of the Hipparcos satellite. The Hipparcos and Tycho photometric systems and their measured median precisions were discussed by \\citet{vLee97b} and the passbands were given in \\citet{vLee97a}. However, the detectors suffered degradation throughout the mission as a result of being launched into an incorrect orbit and this degradation invalidated the measured pre-launch passbands.  \\citet{Bess00} devised self-consistent Hipparcos and Tycho passbands by comparing regressions of $V-H_{p}$, $V-B_{T}$ and $V-V_{T}$ versus  $V-I$ for a sample of precise E-region $UBVRI$ standard stars with synthetic photometry computed from the $R\\sim$100 Vilnius averaged spectra \\citep{Stra72}. This indicated the necessity of a significant redward shift of the blue edge of the published $H_{p}$ band but only small changes for the Tycho bands. However, the passbands may not have been definitive because of the small number of averaged Vilnius spectra used and their low resolution.  In the last few years, two libraries of accurate higher resolution ($R\\sim$1000--2000) spectrophotometric data have become available -  the Next Generation Spectral Library (NGSL) \\citep[][\\url{http://archive.stsci.edu/prepds/stisngsl/index.html}]{Heap07} and the Medium Resolution INT Library of Empirical Spectra (MILES) \\citep[][\\url{http://www.iac.es/proyecto/miles/}]{Sanc06}. Many of the stars in these spectral libraries also have Hipparcos and Tycho magnitudes - providing the opportunity to reexamine the passbands of the Hipparcos and Tycho systems.  Furthermore, the high precision of the Hipparcos magnitudes \\citep[see section 1.3.1][]{Perr97} enables them to be used to adjust the flux levels of the data in the NGSL libraries and make the stars extremely valuable for whole sky spectrophotometric calibration of imaging surveys such as SkyMapper \\citep{Kell07}.  Being space based, a unique property of the Hipparcos photometric systems is the absence of any seasonal or hemisphere-related effects seen in some ground based photometric systems due to variations in temperature, atmospheric extinction and instrumental orientation. The Hipparcos photometry database can therefore be compared with databases of ground-based photometric systems to examine their magnitude zero-points and to look for any systematic offsets in the photometry, as discussed by \\citet{vLee97a} and \\citet{Pell07}.  In this paper, we will outline the derivation of improved $UBVRI$, Hipparcos and Tycho passbands by using synthetic photometry from spectrophotometric atlases and comparing it with broad-band photometry. We will also adjust the absolute levels of the spectrophotometric fluxes by comparing the synthesized $H_{p}$ magnitudes with the Hipparcos catalogue magnitudes. In addition, we will use the mean differences between the synthetic and the observed photometry to determine zero-point corrections for the $UBVRI$, $B$$_{T}$ and $V$$_{T}$ bands.  We will also inter-compare the zero-points of the SAAO $UBVRI$, the homogenized $UBV$ and the Walraven $VB$ systems. Finally, in an appendix we discuss confusion and inexactness concerning the derivation of mean fluxes, response functions and the plethora of expressions for mean wavelengths and frequencies associated with broadband photometry. ",
        "conclusions": "Excellent spectrophotometric catalogs are now available from NGSL \\citep[][\\url{http://archive.stsci.edu/prepds/stisngsl/index.html}]{Heap07}  and MILES \\citep[][\\url{http://www.iac.es/proyecto/miles/}]{Sanc06}. In addition to their intrinsic worth, such stars are very useful to use to calibrate all-sky surveys, such as SkyMapper \\citep{Kell07}. However, the published absolute flux levels are imprecise or non-existent, so we have renormalized the spectra to their precise $H_{p}$ magnitudes. In order to do this it was necessary to determine the best $H_{p}$ passband. We also decided to reexamine  the passbands representing the $UBVRI$ and Tycho $B_{T}$ and $V_{T}$ standard photometric systems using the renormalized NGSL and MILES spectra. We used the CALSPEC stis005 spectrum of Vega to derive the nominal ZP corrections to the AB mag fluxes and synthesized the various magnitudes and slightly adjusted the photonic passbands, achieving better agreement between the synthetic and standard magnitudes than possible with the previous  passbands. In Table 1 and 2 we presented the adopted photonic passbands for $UBVRI$ and $H_{p}$, $B_{T}$ and $V_{T}$, respectively. Table 3 lists the ZP magnitude corrections based on the stis005 Vega spectrum and $V$ = 0.03.  We intercompared the ZPs of the $V$ magnitude scale of the SAAO Cousins-Landolt $UBVRI$ system, the homogenized \\citet{Merm06} $UBV$ system and the Walraven \\citet{Pell07} system and the Hipparcos $H_{p}$ and Tycho $B_{T}$ and $V_{T}$ systems. We found small differences of less than 0.01 mag between them.  The $H_{p}$ magnitude ZP differs by 0.008 mag from the ZP of the $UBVRI$ system.  We analysed the mean magnitude and color differences between the synthetic photometry and the standard photometry and proposed small additional ZP corrections to place the synthetic photometry computed using the AB mag ZPs in Table 3 onto the same ZPs as the standard system photometry. These additional ZP corrections are given in Table 5 together with the passband parameters that should be used to characterize the standard systems.  There was good agreement between the mean differences from the NGSL and MILES catalogs, although the mean level of the MILES blue fluxes deviated slightly, but systematically, from those of the NGSL spectra. The synthetic colors of the fainter CALSPEC spectra also supports the proposed additional ZP corrections except for an unexplained difference in the relative $V$ and $H_{p}$ magnitudes.  Finally, in  the Appendix, we present an extensive discussion on the confusion in the literature concerning measured magnitudes, fluxes and response functions when broad-band photometry is involved and provided equations setting out clearly the derivation of photometric  quantities.  We also cross reference parameters and definitions used in the HST photometric packages \\textit{synphot} and its successor \\textit{pysynphot}. (Table A1)."
    },
    "1112/1112.0822_arXiv.txt": {
        "abstract": "A systematic study of the synchrotron X-ray emission from supernova remnants (SNRs) has been conducted. We selected a total of 12 SNRs whose synchrotron X-ray spectral parameters are available in the literature with reasonable accuracy, and studied how their luminosities change as a function of radius. It is found that the synchrotron X-ray luminosity tends to drop especially when the SNRs become larger than $\\sim$5~pc, despite large scatter. This may be explained by the change of spectral shape caused by the decrease of the synchrotron roll-off energy. A simple evolutionary model of the X-ray luminosity is proposed and is found to reproduce the observed data approximately, with reasonable model parameters. According to the model, the total energy of accelerated electrons is estimated to be $10^{\\rm 47-48}$~ergs, which is well below the supernova explosion energy. The maximum energies of accelerated electrons and protons are also discussed. ",
        "introduction": "Young supernova remnants (SNRs) are widely believed to be the main source of Galactic cosmic rays. \\citet{koyama1995} discovered synchrotron X-rays from shells of SN~1006, which was the first observational clue of cosmic-ray electrons being accelerated up to the TeV energy range. Later, several young SNRs were found to have synchrotron X-ray shells \\citep[c.f.,][]{koyama1997,slane1999,bamba2000,vink2003}. Another piece of evidence for cosmic-ray acceleration in SNRs was obtained from gamma-ray observations. Very-high-energy (VHE) and GeV gamma-rays have been detected from several SNRs, although, it is still unclear whether their origin is hadronic or leptonic \\citep[e.g.,][]{aharonian2004,aharonian2007, abdo2010b,abdo2010c,abdo2011}.  Despite these various pieces of evidence in favor of acceleration, it is still unclear how the acceleration process evolves in SNRs, in particular how the accelerated particles cool down and how they escape from SNRs. One reason for this lack of understanding is that past observational studies have concentrated on individual sources alone. In this paper, we investigate, for the first time, the time evolution of SNR synchrotron X-rays, using all available data in literature for sources in which this emission component is bright. In section~\\ref{sec:sample}, we describe the evolution of synchrotron X-ray luminosity in our sample. An interpretation and a simple model to support it are in section~\\ref{sec:model}. Finally, section~\\ref{sec:discuss} is devoted to discussion of the results. ",
        "conclusions": "\\label{sec:discuss}  We have reproduced the observed fast drop-off in synchrotron X-ray luminosity with SNR radius using a quite simple model. Recently, \\citet{patnaude2010} discovered a decline of a few percent in the synchrotron X-ray power from one source, Cas~A, and reached a similar conclusion.  The total and maximum energy of accelerated particles can be derived from our model. With the adopted values of $A$ and Eq.~(\\ref{eq:N_e}), we can calculate the total energy of accelerated electrons with the energy above $m_ec^2$ to be  $10^{47}$--$10^{48}$~ergs, which is reasonable since this range is much smaller than $E_{\\rm ej}$. Note that these estimated values could increase by 1 order of magnitude if we increase $\\xi$ by up to 10. Even if we alter the value of $p$, the results remain similar if we appropriately reset the normalization $A$. This is expected from equation~(\\ref{eq:Luminoosity}). For example, the results for $p=2.2$ are almost perfectly in agreement with the case of $p=2.0$ if we choose $A=1.7\\times 10^{46}(n_0/1~{\\rm cm}^{-3})$ $n_0=5.0$ and 1.0~cm$^{-3}$ and $A=1.7\\times 10^{46}$ for $n_0=0.1$~cm$^{-3}$. In these cases, the total energy of accelerated electrons with energy above $m_ec^2$ is again $10^{47}$--$10^{48}$~ergs. \\citet{bamba2003} estimated the energy of injected accelerated electrons in SN~1006 for the first time. This was found to be $\\sim 10^{44}$--$10^{45}$~ergs per small segments of filaments, although with rather large uncertainty. The total energy of accelerated electrons is approximately $10^{46}$--$10^{48}$~ergs considering the size of these segments, which is consistent with our result. \\citet{bamba2005} show that the enegy injected into accelerated electrons is always similar in young SNRs, which is also consistent with our results. Theoretically, \\citet{berezhko2006} estimated the energy injected into accelerated protons to be $10^{50}$~ergs in the case of RX~J1713$-$3946. This is also consistent with our results when we consider the electron and proton ratio of $10^{-4}$ in cosmic rays.   At present, no direct information on the accelerated protons is obtained by X-ray observations. One might consider whether we could estimate the total energy of accelerated protons under the assumption of a similar $A$ as for electrons. However, this is not possible because the injection efficiency is different in each case. On the other hand, the maximum energy of accelerated protons, $E_{p}{}^{max}$, can be estimated according to our  model. The synchrotron and $\\pi^0$-decay loss timescales are negligible for protons, and thus $E_{p}{}^{max}$ is determined by the condition $t_{\\rm acc}=t_{\\rm age}$, and we obtain \\begin{equation} E_{p}{}^{max}(t_{\\rm age}) = \\frac{4.8\\times 10^2}{\\xi} \\left(\\frac{v_s}{10^9 {\\rm cm~s^{-1}}}\\right)^2 \\left(\\frac{B_d}{10~\\mu{\\rm G}}\\right) \\left(\\frac{t_{\\rm age}}{10^3~{\\rm yr}}\\right)\\ \\ {\\rm TeV}\\ . \\end{equation} Although the effect of the wave damping and/or escape may be significant in older systems \\citep[e.g.,][]{zirakashvilli2007,ohira2010}, we ignore it here for simplicity. Figure~\\ref{fig:Emax_evolution} shows $E_{p}{}^{max}$ as a function of $t_{\\rm age}$. In the free expansion phase, $E_{p}{}^{max}$ increases in proportion to the age, and protons are accelerated quickly up to around the knee energy in $\\sim$100~yr. When the shock velocity decreases, $E_{p}{}^{max}$ becomes smaller. Hence $E_{p}{}^{max}$ peaks at $t_{\\rm age}=t_1$. In contrast to the case of electron acceleration, $E_{p}{}^{max}(t_1)$ (that is the maximum of $E_{p}{}^{max}$) is larger for higher upstream density $n_0$, and it changes only within a factor of $\\sim2$ for $n_0$ = 0.1--5.0~cm$^{-3}$. Since  $B_d \\propto n_0{}^{1/2}$ and $t_1 \\propto n_0{}^{-1/3}$, one finds $E_{p}{}^{max}(t_1)\\propto B_dt_1\\propto n_0{}^{1/6}$. Protons are accelerated up to $E_{p}{}^{max}(t_1)\\sim 10^{15-16}$~eV in any SNR, with little density dependence. This result might be important for explaining the break at $\\sim 10^{15.5}$~eV in the observed cosmic-ray spectrum with the shock acceleration on SNRs. On the other hand, one can find $E_{e}{}^{max}\\propto B_d{}^{-1/2}v_s\\propto v_s{}^{1/2}$ if $B_d=B_{\\rm amp}\\propto v_s$. Hence the maximum energy of electrons scales as $E_{e}{}^{max}\\propto t^0$ and $\\propto t^{-3/10}$ for the free expansion and Sedov phases, respectively, so that it depends only weakly on the SNR age.  Recent {\\it Fermi} observations of middle-aged SNRs have shown the presence of a gamma-ray spectral break around energies of a few GeV \\citep[e.g.,][]{abdo2009,abdo2010}. Several interpretations have been given \\citep[e.g.,][]{inoue2010,li2010,li2011,ohira2011}, one of which is the escape of particles from acceleration sites. In this scenario, accelerated particles with energies of more than  about 10-100~GeV have already escaped from the acceleration region \\citep{ohira2011}. In contrast, our present model has shown that even for $t_{\\rm age}\\gtrsim10^4$~yr, $\\sim10$~TeV electrons still exist around the shock front (see Figure~\\ref{fig:Emax_evolution}). This apparent discrepancy may come from the fact that most SNRs observed by {\\it Fermi} are interacting with molecular clouds, whereas in our sample there is only one (W~28) undergoing such a collision. Hence molecular clouds may play an important role in dispersing cosmic rays from the SNR shock into the interstellar matter."
    },
    "1112/1112.2685_arXiv.txt": {
        "abstract": "We have undertaken a near-infrared spectral survey of stars associated with compact mid-IR shells recently revealed by the MIPSGAL (24 $\\rm \\mu$m) and GLIMPSE (8 $\\rm \\mu$m) {\\it Spitzer} surveys, whose morphologies are typical of circumstellar shells produced by massive evolved stars. Through spectral similarity with known Luminous Blue Variable (LBV) and Wolf-Rayet (WR) stars, a large population of candidate LBVs (cLBVs) and a smaller number of new WR stars are being discovered. This significantly increases the Galactic cLBV population and confirms that nebulae are inherent to most (if not all) objects of this class. ",
        "introduction": "Despite intensive search efforts over the last several decades the Galactic Luminous Blue Variable (LBV) population has remained sparse. This paucity is difficult to reconcile with our understanding of stellar evolution of the most massive stars. Until the last year or two there were only 12 confirmed Galactic LBVs known, and 23 candidate-LBVs (\\cite[Clark et al. 2005]{Clark2005}). LBVs display rather unique rich infrared emission line spectra, including contributions from H, He, Mg\\,{\\sc ii}, Na\\,{\\sc i}, and Fe\\,{\\sc ii}. Visual inspection of the {\\it Spitzer} GLIMPSE (\\cite[Benjamin et al. 2003]{Benjamin2003}) and MIPSGAL (\\cite[Carey et al. 2009]{Carey2009}) Galactic plane surveys have produced catalogues of previously unknown 8 $\\rm \\mu$m and 24 $\\rm \\mu$m nebulae with concentric point sources that can be traced back to 2MASS $K$-band or even optical sources as possible progenitors of the associated nebulae (\\cite[Gvaramadze et al. 2010a]{Gvar2010a}, \\cite[Wachter et al. 2010]{Wachter2010}), especially with new imaging (\\cite[Stringfellow et al. 2012]{String2012}). We are conducting a near-IR spectral survey to identify new cLBVs and WRs that produced these shells. \\begin{figure}[t] \\begin{center} \\includegraphics[width=5.0in]{Fig1.eps} \\caption{Normalized $K$-band spectra of newly identified cLBVs and WRs. See Figure 1 of \\cite[Stringfellow et al. 2012]{String2012} for images identifying MN 96, including optical recovery in the $I$-band.} \\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "1112/1112.2366_arXiv.txt": {
        "abstract": "Searches for planets via gravitational lensing have focused on cases in which the projected separation, $a$,  between planet and star is comparable to the Einstein radius, $R_E.$ % This paper considers smaller orbital separations and demonstrates that evidence of close-orbit planets can be found in the low-magnification portion of the light curves generated by the central star. We develop a protocol to discover hot Jupiters as well as Neptune-mass and Earth-mass planets in the stellar habitable zone. When planets are not discovered, our method can be used to quantify the probability that the lens star does not have planets within specified ranges of the orbital separation and mass ratio. Nearby close-orbit planets discovered by lensing can be subject to follow-up observations to study the newly-discovered planets or  to discover other planets orbiting the same star. Careful study of the low-magnification portions of lensing light curves should produce, in addition to the discoveries of close-orbit planets, definite detections of wide-orbit planets through the discovery of ``repeating'' lensing events. We show that events exhibiting extremely high magnification can effectively be probed for planets in close, intermediate, and wide distance regimes simply by adding several-time-per-night monitoring in the low-magnification wings, possibly leading to gravitational lensing discoveries of multiple planets occupying a broad range of orbits, from close to wide, in a single planetary system. ",
        "introduction": "Microlensing planet searches have been directed toward discovering planets in orbits whose size is comparable to the size of the Einstein radius, $R_E$, of the central star. Here we study the detectability of planets in much closer orbits. This is necessary, because we now know that many planets are in such close orbits\\footnote{http://exoplanet.eu/catalog.php}. We demonstrate that ground based surveys for lensing events can detect a wide range of close-orbit planets, including ``hot Jupiters'' orbiting sun-like stars, and even Earth-mass planets in the habitable zones of M dwarfs. We also discuss the role of space missions.  We were led to the study of close-orbit planets by work to determine whether or not planets in the habitable zones of nearby stars could produce detectable lensing signatures (\\rd\\ \\& Night 2008). The standard planetary-lensing scenario is most sensitive to planets with $\\alpha=a/R_E$ in the range $0.5-2,$ where $a$ is the orbital separation (Mao \\& Paczy\\'nski 1991; Gould \\& Loeb 1992). In many cases, this corresponds to the region beyond the snowline (e.g., Sumi et al.\\, 2010), while the habitable zone is closer to the star. We found that lensing observations {\\sl would} be able to detect planets orbiting within the habitable zones of nearby M dwarfs (\\rd\\ \\& Night 2008). Here, we demonstrate the importance of lensing by close-orbit planets for a broader range of systems: the planets may be in the habitable zone, closer in, or farther out. The planetary systems may be within a few tens of pc or much farther away. The host stars may be M dwarfs, brown dwarfs, or much more massive stars.  In \\S 2 we study the incidence of close planets. In \\S 3 we study the expected signatures. We sketch the steps needed to mount successful programs to discover these planets in \\S 4. This paper provides a guide to the discovery and analysis of lensing light curves generated by close-orbit planets. ",
        "conclusions": ""
    },
    "1112/1112.4540_arXiv.txt": {
        "abstract": "We address the nature of EUV waves through direct observations of the formation of a diffuse wave driven by the expansion of a coronal mass ejection (CME) and its subsequent separation from the CME front. The wave and the CME on 2011 June 7 were well observed by Atmospheric Imaging Assembly onboard Solar Dynamic Observatory. Following the solar eruption onset, marked by the beginning of the rapid increasing of the CME velocity and the X-ray flux of accompanying flare, the CME exhibits a strong lateral expansion. During this impulsive expansion phase, the expansion speed of the CME bubble increases from 100 km s$^{-1}$ to 450 km s$^{-1}$ in only six minutes. An important finding is that a diffuse wave front starts to separate from the front of the expanding bubble shortly after the lateral expansion slows down. Also a type-II burst is formed near the time of the separation. After the separation, two distinct fronts propagate with different kinematic properties. The diffuse front travels across the entire solar disk; while the sharp front rises up, forming the CME ejecta with the diffuse front ahead of it. These observations suggest that the previously termed EUV wave is a composite phenomenon and driven by the CME expansion. While the CME expansion is accelerating, the wave front is cospatial with the CME front, thus the two fronts are indiscernible. Following the end of the acceleration phase, the wave moves away from the CME front with gradually an increasing distance between them. ",
        "introduction": "Coronal mass ejections (CMEs) are large-scale eruptive phenomena from the Sun. They can carry large amounts of plasma and magnetic field energy into the interplanetary space, which may have severe effects on space environment and human technological systems around the Earth \\citep{gosling93,webb94}. A typical CME has a velocity of several hundred km s$^{-1}$, while the fastest one recorded is over 3000 km s$^{-1}$ \\citep{Yashiro04}. The main acceleration of a CME usually occurs in the inner corona (e.g., $\\leq$ 3.0 $R_\\odot$) \\citep{zhang01}. With the advent of the Atmospheric Imaging Assembly \\citep[AIA;][]{lemen11} onboard Solar Dynamics Observatory (\\emph{SDO}), details on the initiation and structural formation of CME starts to emerge. Recently, it was found that a twisted hot channel ($\\sim$10 MK) as seen in AIA 131 {\\AA} passband starts to form before the flare onset, and its subsequent rise results in the CME and the accompanying flare; this hot channel has been argued to be the existence of magnetic flux rope prior the eruption \\citep{zhang11}. Following the eruption onset, a CME rises up impulsively with a strong acceleration, and also expands quickly along the lateral direction, producing a plasma bubble \\citep{pat09b,cheng11}. The CME bubble is capable of driving a shock, which may generate the observable metric type II bursts \\citep[e.g.,][]{liu09,ma11}.   One interesting phenomenon closely associated with CMEs is the globally-propagating bright feature in the corona dubbed the EIT or EUV wave since its discovery by \\emph{SOHO} spacecraft \\citep{thompson98,thompson99}. The physical nature of the moving front remains somewhat unclear. The front has been interpreted as a fast-mode magnetohydrodynamic (MHD) wave \\citep{thompson99,wang00,warmuth01,veronig08,kienreich09,gopa09,pat09a,pat09b,liu11,olmedo11}, a soliton wave \\citep{wills07}, or slow-mode wave \\citep{wang09}. Whereas, others believe that it is not at all a true wave. \\citet{chen02,chen05}, \\citet{chen09}, and \\citet{chen11} argued that the bright front results from the compression front driven by the successive stretching of the magnetic field overlying the erupting CME flux rope, although there should be a fast-mode wave ahead of the CME front. \\citet{attrill07,attrill09} and \\citet{dai10} claimed that the bright front is related to the magnetic reconnection between the outmost magnetic field of the CME and the magnetic loops from the quiet region. \\citet{delannee08} suggested that the current shell in their numerical simulation can also form the bright front. In order to bring together these opposing views, some authors recently tend to appeal for the hybrid model combining both wave and non-wave explanations \\citep[e.g.,][]{cohen09,liu10,downs11}. Details of various EUV wave models can be found in recent review papers by \\citet{warmuth10}, \\citet{gallagher11}, and \\citet{zhukov11}.  In this Letter, we investigate the physical nature of the EUV wave. From the observations, we find that the wave front has two distinct evolution stages, i.e., a compression front forming the CME front in the early stage and a stand-alone wave front separating from the CME front in the later stage. The data used are mainly from AIA onboard \\emph{SDO} and Sun Earth Connection Coronal and Heliospheric Investigation \\citep[SECCHI;][]{howard08} onboard Solar TErrestrial RElations Observatory (\\emph{STEREO}). Observations and results are presented in Section 2, followed by a summary and discussion in Section 3. ",
        "conclusions": "In this Letter, we focus on studying the separation process of two distinct fronts associated with a solar eruption that occurred on 2011 June 7. Following the eruption onset, the magnetic field of the source region quickly expands and forms a circular bubble. In the early stage of the eruption, the bubble expands strongly with an accelerating velocity. Afterwards, the apparent expansion velocity of the bubble close to the solar surface quickly decelerates to almost zero. In the meantime, a diffuse wave front starts to separate from the sharp bubble front. We conclude that the wave originates from the compression of the surrounding plasma by the impulsively expanding CME bubble. Due to a small standoff distance between the compression front and the driver front, the two fronts can not be distinguished during the early stage of the evolution when the driver is still undergoing acceleration. Through examining the radio data from CALLISTO radio spectrometer, we find that a type II radio burst started at $\\sim$06:26 UT, almost at the same time (slightly earlier) as the wave started to separate from the driver. The radio observation suggests that a shock was generated at that time, indicating that the compression wave had just turned itself into a shock wave. Recent studies by \\citet{pat09a} and \\citet{pat10} also noticed two fronts: the wave front and the bubble front. However, due to the relative low cadence of the \\emph{STEREO}-EUVI observations (5 minutes), the two distinct fronts were seen in only a few frames. Here, the AIA observations in every 12 seconds not only reveal the existence of two fronts but also the detailed separation process of the diffuse front from the sharp bubble front after the expansion of the bubble slows down. These observations clearly demonstrate that the separated diffuse front is a true MHD wave, driven by the early accelerating expansion of the CME bubble.  In conclusion, our observations help understand the physical nature of usually termed EUV waves. The EUV wave is actually a composite phenomenon, consisting of two distinct fronts, a wave front and a CME (compressed plasma) front. We further find that its evolution can be divided into two stages. The first stage takes place in the accelerating expansion phase of the CME bubble, which acts as the piston-driver of the MHD wave. During this stage, the wave front is coupled together with the compression front of the CME bubble. In the second stage, when the expanding velocity of the CME bubble slows down, the MHD wave front decouples from the compression bubble front, forms a distinct front, and propagates across the solar disk. The observational result for the coexistence of both wave and non-wave fronts are in general consistent with the models and numerical simulations for EUV waves by \\citet{chen02,chen05}. We believe that the previous dispute about the nature of EUV waves resides in, at least partly, different parts of the composite phenomenon that the authors may have observed. In fact, the duration of the wave and CME compression front coupling is different for each event, depending on the dynamics of the CME and the surrounding environment."
    },
    "1112/1112.3360_arXiv.txt": {
        "abstract": "Eternal inflation is a seemingly generic consequence of theories that give rise to accelerated expansion of the universe and possess multiple vacuum states. Making predictions in an eternally inflating universe is notoriously difficult because one must compare infinite quantities, and a wide variety of regulating procedures yield radically different results. This is the measure problem of eternal inflation. In this paper, we analyze models of eternal inflation which allow for the possibility of cyclic bubble universes: in each bubble, standard cosmological evolution is re-played over and over again. Eternal inflation can generically arise in cyclic models that include a dark energy dominated phase. In such models, several problematic consequences of standard regulating procedures, such as the youngness and Boltzmann Brain problems, are substantially alleviated. We discuss the implications for making predictions in cyclic models, as well as some general implications for understanding the measure problem in eternal inflation. ",
        "introduction": "In the standard cosmological model, accelerated expansion is invoked at least twice to explain our observations of the cosmos: dark energy in the late universe, and cosmic inflation in the early universe. This theory allows multiple phases, characterized by vastly different rates of accelerated expansion; inflation gave way to a dark energy dominated universe with an expansion rate some $50$ orders of magnitude smaller. In fact, the possibilities may be much more varied, since modern theories of particle physics relying on spontaneous symmetry breaking generically do not predict a unique vacuum state, let alone a unique cosmological history.  Many such theories that allow multiple phases (or vacua) and accelerated expansion give rise to a phenomenon known as eternal inflation. In the simplest example of eternal inflation, different vacua are accessed by a first order phase transition. In this picture, localized pockets containing a new phase are formed from a ``parent\" vacuum undergoing accelerated expansion. These pockets subsequently expand, but if they are formed at a rate slower than the rate of expansion, the phase transition cannot complete, and inflation becomes eternal. The original inflating vacuum continues to grow and spawn regions of all the possible phases (for a review of eternal inflation, see e.g.~\\cite{Aguirre:2007gy}). In this way, eternal inflation makes all possibilities a reality.  There are many conceptual problems that must be confronted in order to make sense of eternal inflation. Since the many possibilities allowed by an underlying theory are {\\em actually} realized somewhere in spacetime, predictions for what we might observe become intrinsically statistical. In a finite universe with a finite number of possibilities, this would not be a problem. However, in an eternally inflating universe each possibility is realized an infinite number of times, and so a frequentist's definition of probabilities becomes ill-defined without some prescription to regulate the infinities. The most straightforward  procedure is to use a geometrical cutoff: count the frequency of some occurrence vs others at a finite time (say, by counting the volume occupied by different phases), and then take the time at which the ratio is evaluated to infinity. However, this procedure is not unique, and choosing different timeslicings yields different sets of probabilities. This is an example of the ``measure problem\" of eternal inflation.  To the extent that they are understood, the predictions made by various measures (cutoff procedures) range from being absurdly in conflict to fairly consistent with data (at least for a few variables such as the cosmological constant). For example, the generic prediction of measures that weight by physical volume is that we should be as ``young\" as possible (i.e. everything else constant, we should exist at the earliest possible stage in the evolution of the universe consistent with our existence). This is known as the Youngness Problem. If there is a cosmological constant in each pocket (as there is evidence for in our universe), then because of the thermodynamic properties of de Sitter space, many measures predict that we should be ``freak observers'' (i.e. that we didn't arise from our observed cosmological evolution) formed by a thermal fluctuation out of the vacuum at some very late time. This is known as the Boltzmann Brain Problem. We will discuss both of these issues at length below.  One strategy for making progress on the measure problem is to rule out those measures which make predictions that do not agree with data, or are not internally consistent (see e.g. the discussion in Ref.~\\cite{Bousso:2008hz}). For example, one would rule out those measures with a Youngness or Boltzmann Brain Problem. Another approach is to look for guidance from fundamental theory, which perhaps would single out the correct prescription~\\cite{Bousso:2010id,Vilenkin:2011yx,Garriga:2008ks,Sekino:2009kv,Freivogel:2006xu}. However, this is a complex problem which is not yet resolved.  An alternative to the inflationary epoch in the standard cosmological model is the ekpyrotic/cyclic universe scenario~\\cite{Khoury:2001wf,Steinhardt:2001st,Steinhardt:2002ih,Khoury:2001bz,Steinhardt:2004gk}, for a review see~\\cite{Lehners:2008vx}. This scenario purports to address both the homogeneity and isotropy of our observed universe as well as seeding structure with appropriate density fluctuations. This is accomplished by invoking an ``ekpyrotic'' phase in the early universe, in which there is slow contraction accompanied by a rapid change in the energy density. This contraction gives way to expansion after a transition from big-crunch to big-bang, after which the standard big-bang cosmology follows. Far in the future, the universe will contract again, and the cycle repeats itself. This picture is motivated by Heterotic M-theory, as we describe in greater detail below.  There are at least two reasons why one must address the question of measure in the cyclic universe. The first arises when there can be differences between cycles, as there are in the so-called Phoenix Universe~\\cite{Lehners:2008qe,Lehners:2010ug,Lehners:2011ig}. In this setup, the amplitude of density fluctuations can vary from place to place in each cycle. However, only those parts of the universe that are sufficiently empty and flat make it from the contracting to the expanding phase, leading to a selection factor for the amplitude of density perturbations. Different possibilities are sampled in different spatiotemporal regions, requiring one to define a statistical measure for predictions (and in an infinite universe, some cutoff procedure). Second, the cyclic universe scenario invokes accelerated expansion to smooth the universe prior to ekpyrotic contraction. As argued above, once accelerated expansion is added to any cosmological model, the existence of eternal inflation becomes a fairly generic possibility. This is even more pronounced given that the cyclic universe is motivated by M-theory, where the many ways to compactify extra dimensions give rise to many possible four-dimensional theories. The picture of the multiverse might then contain both cyclic and inflationary cosmologies, as shown in Fig.~\\ref{Fig-Multiverse}.  \\begin{figure} \\begin{center} \\includegraphics[width=0.75\\textwidth]{Multiverse7.eps} \\caption{The Cyclic/Inflationary Multiverse. \\label{Fig-Multiverse} } \\end{center} \\end{figure}  Can the cyclic universe arise in local pockets spawned from an eternally inflating universe? What are the implications for making predictions in such models? In this paper, we address these questions, finding that using existing measures the youngness and Boltzmann brain problems need not (under a set of conditions we outline in detail) arise in a theory where many cyclic universes are seeded by eternal inflation. In addition, we construct a model where both cyclic and standard inflationary bubbles coexist, and find under which circumstances these problems reappear. One should keep in mind our largest assumptions: 1) that there are multiple phases (including cyclic phases) which can be seeded by eternal inflation and 2) that the full theory allows a predictive transition from a big crunch to a big bang following the ekpyrotic phase. Both assumptions are plausible, however unproven; we offer no new critiques of either the eternal inflation or cyclic scenarios. Rather, we wish to illustrate how the inclusion of cyclic universes can lead to drastically different predictions for where we might find ourselves in the multiverse. Our results also highlight how different assumptions about the physics underlying the multiverse can render previously discarded measure prescriptions consistent with our observed universe.  The plan of the paper is as follows. In Sec.~\\ref{Section-Setup}, we describe how eternal inflation might arise within the cyclic universe scenario. In Sec.~\\ref{Section-RateEqs}, we outline how predictions are extracted, and how the eternally inflating cyclic universe can avoid the youngness and Boltzmann brain problems. We add inflationary bubbles to the setup in Sec.~\\ref{sec:inflationarybubbles} and show under which circumstances these problems reappear, and then conclude in Sec.~\\ref{sec:conclusions}. Finally, we include a discussion of tunneling to the cyclic universe in Appendix~\\ref{sec:twofieldbubbles} and a calculation for worldline based measures in Appendix~\\ref{sec:Boussorates}. ",
        "conclusions": "\\label{sec:conclusions}  In the standard picture of eternal inflation, different phases are seeded from a rapidly inflating parent vacuum through the formation of bubbles, each of which can contain an inflationary cosmology, and possibly describe our observable universe. As we have described in detail above, there exist many ambiguities and subtleties in defining a probability measure suitable for making predictions in an eternally inflating universe. However, many of the problems that arise for various measures can be traced back to a simple fact: during eternal inflation it is much easier to fluctuate to a lower energy state than a higher energy state. The youngness problem arises because space is inflating much more rapidly outside of pocket universes than inside. The Boltzmann Brain problem arises because it is much easier to make observers via a small fluctuation than to re-play the entirety of cosmological evolution.  The model we have introduced in this paper, which incorporates cyclic universes into the picture of eternal inflation, at face value would seem to avoid these problems. In a cyclic bubble universe, cosmological evolution leads periodically to states with higher energy density. Thus, in forming a cyclic bubble, a chain of events is set into motion by which the energy inside the bubble can temporarily exceed that of the background space it is embedded in. The volume inside the bubble can then  grow faster than on the outside, relieving the youngness problem. Cyclic bubbles continually produce observers by normal cosmological evolution, and the nucleation rate for cyclic bubbles can be fast,  alleviating the Boltzman Brain problem. A similar conclusion would be reached in models where the production of false vacuum (higher energy density) bubbles is not highly suppressed, as in the ``de Sitter equilibrium cosmology\" of Albrecht~\\cite{Albrecht:2004ke,Albrecht:2009vr}. These two examples are illustrative of what is generally necessary to avoid a Youngness and BB problem, and suggest that something like a rather drastic violation of the null energy condition is required. Whether or not this is reasonable is of course subject to debate.  The analysis we have performed in this paper confirms the intuition described above. However, the conclusions are somewhat sensitive to the details of the physics connecting the big crunch of one cycle to the big bang of the next. When there are only cyclic and terminal vacua, the volume fraction of the universe is always dominated by regions containing matter and radiation at late times. Thus, the youngness paradox is non-existent. Again, when there are only cyclic and terminal vacua, the existence of a BB problem is dependent on whether there is any loss of volume at the crunch/bang transition. When there is no loss of volume, the BB problem is non-existent. When there is volume loss, the false vacuum must produce cyclic bubbles faster than it produces BBs to avoid a problem. Let us emphasize that there is still a measure problem for the eternal inflation/cyclic universe hybrid; there are many possible cutoff procedures, and no {\\it a priori} way to choose which is correct. However, our results suggest that many commonly considered choices lead to predictions consistent with observations.  We have also analyzed a model that contains both cyclic and inflationary bubbles formed from the same false vacuum. This might be expected more generally if both inflationary and cyclic cosmologies are realized in a fundamental theory. We found that when weighting vacua by physical volume the youngness paradox is re-introduced if the false vacuum energy scale is large and inflationary bubbles are produced more rapidly than cyclic bubbles. This is the regime where the volume fraction inside inflationary bubbles is larger than the volume in the radiation/matter dominated portions of cyclic bubbles. In the causal diamond measure, the cyclic bubbles generically dominate the measure as long as they are produced more frequently than the inflationary bubbles. In the case where there is no loss of volume at the crunch/bang transition, cyclic bubbles in fact always dominate. In cases where the cyclic bubbles dominate the measure, the BB problem is largely absent, although subject to some conditions on the rate of decay of the false vacuum.  Beyond the particular examples we have considered, our results motivate the construction of a more general set of conditions necessary for particular measures to yield results consistent with our observable universe. We hope to explore this question in future work."
    },
    "1112/1112.1237_arXiv.txt": {
        "abstract": "A central quantity in mean-field magnetohydrodynamics is the mean electromotive force $\\meanEMF$, which in general depends on the mean magnetic field. It may however also have a part independent of the mean magnetic field. Here we study an example of a rotating conducting body of turbulent fluid with non-zero cross-helicity, in which a contribution to $\\meanEMF$ proportional to the angular velocity occurs (Yoshizawa 1990). If the forcing is helical, it also leads to an $\\alpha$ effect, and large-scale magnetic fields can be generated. For not too rapid rotation, the field configuration is such that Yoshizawa's contribution to $\\meanEMF$ is considerably reduced compared to the case without $\\alpha$ effect. In that case, large-scale flows are also found to be generated. ",
        "introduction": "Many studies of the large-scale magnetic fields in turbulent astrophysical bodies such as the Sun or the Galaxy are carried out in the framework of mean-field electrodynamics \\citep[see the textbooks by][]{Mof78,Par79,KR80,ZRS83}. It is based on the induction equation governing the magnetic field $\\BB$, \\EQ {\\partial\\BB\\over\\partial t}=\\nab\\times\\left(\\UU\\times\\BB-\\eta\\mu_0\\JJ\\right) \\, , \\label{eq01} \\EN where $\\UU$ is the fluid velocity, $\\JJ=\\nab\\times\\BB/\\mu_0$ the current density, where $\\UU$ is the fluid velocity, $\\JJ=\\nab\\times\\BB/\\mu_0$ is the current density, $\\eta$ the magnetic diffusivity, and $\\mu_0$ the vacuum permeability. Both the magnetic field $\\BB$ and the velocity field $\\UU$ are considered sums of mean parts, $\\meanBB$ and $\\meanUU$, defined as proper averages of the original fields, and fluctuations. The averages are assumed to satisfy the Reynolds averaging rules. The mean magnetic field $\\meanBB$ then obeys the mean-field induction equation \\EQ {\\partial\\meanBB\\over\\partial t}=\\nab\\times\\left( \\meanUU\\times\\meanBB+\\meanEMF-\\eta\\mu_0\\meanJJ\\right) \\, . \\label{eq03} \\EN Here $\\meanEMF=\\overline{\\uu\\times\\bb}$ is the mean electromotive force resulting from the fluctuations of velocity and magnetic field, $\\uu=\\UU-\\meanUU$ and $\\bb=\\BB-\\meanBB$. Generally, $\\meanEMF$ can be represented as a sum \\EQ \\meanEMF = \\meanEMF^{(0)} + \\meanEMF^{(B)} \\label{eq05} \\EN of a part $\\meanEMF^{(0)}$, which is independent of $\\meanBB$, and a part $\\meanEMF^{(B)}$ vanishing with $\\meanBB$. In many representations and applications of mean-field electrodynamics the part $\\meanEMF^{(0)}$ of $\\meanEMF$ is ignored. Only the part $\\meanEMF^{(B)}$, which is of crucial importance for dynamo action, is taken into account.  Here we focus our attention on the part $\\meanEMF^{(0)}$ of $\\meanEMF$. It may depend on non-magnetic quantities influencing the turbulence, in general also on $\\meanUU$. If the magnitude of $\\meanUU$ is small, and if $\\meanUU$ varies only weakly in space and time, we may write \\EQ \\meanemf^{(0)}_i = \\meanemf^{(00)}_i + \\Xi_{ij} \\meanU_j + \\Upsilon_{ijk} \\meanU_{j,k} \\label{eq07} \\EN with $\\meanemf^{(00)}_i$ as well as $\\Xi_{ij}$ and $\\Upsilon_{ijk}$ being independent of $\\meanUU$. Of course, the contribution $\\meanEMF^{(00)}$ to $\\meanEMF^{(0)}$ can only be non-zero if the turbulence allows us to define a direction. For example, turbulence in a rotating body shows in general an anisotropy determined by the angular velocity $\\OO$, and $\\meanEMF^{(00)}$ might then be proportional to $\\OO$, say equal to $c_\\varOmega \\OO$. The $\\Xi_{ij}$ term in \\eq{eq07} can only be unequal to zero if the turbulence lacks Galilean invariance. In the case of isotropic turbulence it describes a contribution to $\\meanEMF^{(0)}$ proportional to $\\meanUU$, say equal to $c_U \\meanUU$. Note that in forced turbulence Galilean invariance can be broken if, independent of the flow, the forcing is fixed in space and shows a finite correlation time \\citep[for an example see][]{RB10}. The $\\Upsilon_{ijk}$ term, if restricted to isotropic turbulence, corresponds to a contribution to $\\meanEMF^{(0)}$ proportional to $\\nab \\times \\meanUU$, say equal to $c_W \\nab \\times \\meanUU$. The coefficients $c_\\varOmega$ and $c_W$ are, in contrast to $c_U$, pseudoscalars. The contributions $c_\\varOmega \\OO$ and $c_W \\nab \\times \\meanUU$ to the mean electromotive force were first considered by \\cite{Yos90}. He found that both $c_\\varOmega$ and $c_W$ are closely connected with the cross helicity $\\overline{\\uu {\\bm\\cdot} \\bb}$. In what follows the occurrence of the contributions $c_\\varOmega \\OO$ and $c_W \\nab \\times \\meanUU$ to the mean electromotive force $\\meanEMF$ is called ``Yoshizawa effect\". This effect has been invoked to explain magnetic fields in accretion discs \\citep{YY93} and spiral galaxies \\citep{Yok96}. It has also been used to explain the surprisingly high level of magnetic fields in young galaxies \\citep{BU98}, because the amplification of the mean field by this effect is independent of any seed magnetic field. The equivalence of a rotation of the frame of reference with a rotation of the fluid body might suggest an equality of $c_\\varOmega$ and $2 c_W$. However, this equivalence exists only in pure hydrodynamics, which is governed by the momentum equation, but no longer in magnetohydrodynamics, where both the momentum equation and the induction equation are important. As a consequence, $c_\\varOmega$ is in general different from $2 c_W$, see \\cite{RB10}, in particular the discussion at the end of Section 3.1.   As for the part $\\meanEMF^{(B)}$ of $\\meanEMF$, we recall here the traditional ansatz \\EQ \\meanEMF^{(B)} = \\alpha_{ij} \\meanB_j + \\eta_{ijk} \\meanB_{j,k} \\, . \\label{eq09} \\EN It can be justified for cases in which $\\meanBB$ varies only slowly in space and time. In the simple case of isotropic turbulence it takes the form $\\meanEMF^{(B)} = \\alpha \\meanBB - \\etat \\nab \\times \\meanBB$, which describes the $\\alpha$ effect and the occurrence of a turbulent magnetic diffusivity \\citep{KR80}.  In this article, we report on numerical simulations of magnetohydrodynamic turbulence in a rotating body, that is, under the influence of the Coriolis force. We present results for the mean electromotive force and discuss them in the light of the above remarks, focussing particular attention on the Yoshizawa effect. ",
        "conclusions": "The mean electromotive force in a turbulent fluid may have a part that is independent of the mean magnetic field and also independent of the mean flow. As an example we have studied forced hydromagnetic turbulence in a rotating body. In this case the Yoshizawa effect occurs, that is, a contribution $c_\\varOmega \\OO$ to $\\bra{\\uu \\times \\bb}$. We have confirmed that $c_\\varOmega$ is determined by the mean cross-helicity $\\bra{\\uu {\\bm\\cdot} \\bb}$. We have also seen that, if an $\\alpha$ effect is present, the Yoshizawa effect can to a large extent be compensated by the action of magnetic fields maintained by this $\\alpha$ effect.  In astrophysics, the occurrence of non-zero cross-helicity is not a very common phenomenon. We give here a few examples in which the findings of this paper could be of interest. In the solar wind the systematic radial flow together with the Sun's large-scale magnetic field give rise to cross helicity of opposite sign in the two hemispheres. Although this primarily implies cross helicity associated with mean flow and mean magnetic field, it also results in cross helicity associated with the fluctuations. Together with the Sun's rotation, the latter should then produce a component of the mean electromotive force that is distinct from that related to the $\\alpha$ effect. Note, however, that the cross-helicity associated with the fluctuations is directly a consequence of the cross helicity from the large-scale field.  Another example where small-scale cross helicity can be generated is in a stratified layer with a vertical magnetic field \\citep{RKB11}. Again, the sign of $\\bra{\\uu {\\bm\\cdot} \\bb}$ is linked to the orientation of the large-scale field relative to the direction of gravity.  Finally, cross helicity can be generated spontaneously and can then be of either sign, such as in the Archontis dynamo \\citep{Arc00}; for kinematic simulations see \\cite{Arc03} as well as \\cite{CG06}. \\cite{SB09} have analyzed this dynamo with respect to the Yoshizawa effect. In this example too, large-scale and small-scale fields are intimately related. This interrelation means that whenever we expect the $\\meanEMF^{(0)}$ term to be present in an astrophysical system, there should also be a mean magnetic field. Such an effect that is odd in the mean magnetic field might therefore instead just as well be associated with an $\\alpha$ effect. As it turns out, this is also the case in the present simulations, where a large-scale magnetic field has been produced. In the present case, we have gone a step further by including also kinetic helicity also, in addition to just cross helicity. This produces an $\\alpha$ effect and, as a consequence of this, a large-scale magnetic field. This field is particularly important when rotation is weak, because then the Yoshizawa effect is strongly quenched by this field."
    },
    "1112/1112.1830_arXiv.txt": {
        "abstract": "We briefly review an important shortcoming --unearthed in previous works-- of the standard version of the inflationary model for the emergence of the seeds of cosmic structure.  We consider here some consequences emerging from a proposal inspired on ideas of Penrose and Di\\'osi \\cite{Penrose,Penrose01,Diosi87,Diosi89} about a quantum-gravity induced reduction of the wave function, which has been put forward to address the shortcomings, arguing that its effect on the inflaton field is what can lead to the emergence of the seeds of cosmic structure \\cite{Perez2006}.  The proposal leads to a deviation of the primordial spectrum from the scale-invariant Harrison-Zel'dovich one, and consequently, to a different CMB power spectrum. We perform statistical analyses to test two quantum collapse schemes with recent data from the CMB, including the 7-yr release of WMAP and the matter power spectrum measured using LRGs by the Sloan Digital Sky Survey. Results from the statistical analyses indicate that several collapse models are compatible with CMB and LRG data, and establish constraints on the free parameters of the models. The data put no restriction on the timescale for the collapse of the scalar field modes. ",
        "introduction": "The great advances made in physical cosmology over the past few years open a window for the consideration of issues long dismissed as philosophical speculations. The agreement between theory and observations of the spectrum of the cosmic microwave background (CMB) anisotropies has strengthened the theoretical status of inflationary scenarios among cosmologists. The inflationary paradigm is said to account for the origin of all cosmic structure: the fluctuations around homogeneity and isotropy that have a fundamentally quantum nature are, according to these ideas, the seeds of galaxies, stars, life and humans.  Hence, all structures in our universe emerge from a featureless stage described by a background Friedmann-Robertson-Walker (FRW) cosmology with a nearly exponential expansion driven by the potential of a single scalar field\\footnote{In the simplest models of inflation: $\\phi$, the inflaton.}, and from its quantum fluctuations characterized by a simple vacuum state.  In particular, the quantum fluctuations transmute into the classical statistical fluctuations that imprint their signature in the CMB photons, and eventually grow into the structures we find in our universe today.   However, when this picture is considered thoroughly, an unavoidable issue arises. According to these ideas, a completely homogeneous and isotropic stage, somehow evolves, after some time, into an inhomogeneous and anisotropic situation. Obviously, this is not simply the result of quantum unitary evolution, since, in this case, the dynamics does not break the initial symmetries of the system.  As discussed in Ref.~\\cite{Shortcomings}, and despite multiple claims to the contrary, there is no satisfactory solution to this problem within the standard physical paradigms.  While much of the focus of the research in inflationary cosmology has been directed towards elucidation of the details of the inflationary model, very little attention has been given by the community to an issue of fundamental principles such as the aforementioned one.   One proposal to handle this shortcoming has been developed in Refs. \\cite{Perez2006,Sudarsky06b,Sudarsky07}. That approach attempts to deal with the problem by introducing a new ingredient into the inflationary account of the origin of the seeds of cosmic structure: \\emph{the self induced collapse hypothesis}, i.e. a scheme in which an internally induced collapse of the wave function of the inflaton field\\footnote{As shown in \\cite{MukCollapso}, one can implement the collapse hypothesis at the level of the Mukhanov Sasaki variable.} is the mechanism by which inhomogeneities and anisotropies arise at each particular length scale. That work was inspired in early ideas by R. Penrose and L.  Di\\'osi \\cite{Penrose,Penrose01,Diosi87,Diosi89} which regarded the collapse of the wave function as an actual physical process (instead of just an artifact of the description of Physics) and it is assumed to be caused by quantum aspects of gravitation. We will not recapitulate the motivations and discussion of the original proposal but instead refer the reader to the above mentioned papers.   The way we treat this process is by assuming that at a certain stage in cosmic evolution there is an induced jump in a state describing a particular mode of the quantum field, in a manner that is similar to the quantum mechanical reduction of the wave function associated with a measurement, but with the difference that in our scheme no external measuring device or observer is called upon as triggering such collapse (as there is nothing in the situation we are considering that could be called upon to play such a role).  The issue that then arises concerns the characteristics of the state into which such jump occurs. In particular, what determines the expectation values of the field and momentum conjugate variables for the after-collapse state.  Previous works by people in our group have extensively discussed both the conceptual and formal aspects of that problem, and the present manuscript will not dwell further into those aspects, except for a very short review.  In Ref.~\\cite{Perez2006} two schemes were considered; one in which, after the collapse, both expectation values are randomly distributed within their respective ranges of uncertainties in the precollapsed state, and another one in which it is only the conjugate momentum that changes its expectation value from zero to a value in its corresponding range, as a result of the collapse.  We will discuss the motivations and detailed characterizations of the two processes in the following sections.  As reported in \\cite{Perez2006}, the different collapse schemes give rise to different characteristic departures from the conventional Harrison-Zel'dovich (HZ) flat primordial spectrum.  These aspects have been preliminarily analyzed in Ref.~\\cite{Unanue2008} using a simple approach which ignores the late-time physics effects and simply considers what can be a reasonable expectation for the allowed deviation from the conventional flat spectrum.  The main objective of this article is to consider in detail the shape of the primordial spectrum emerging from such a scheme, and in particular to explore with precision the deviations that should be expected once one takes into account the physics associated with plasma dynamics, which is responsible for the generation of the acoustic peaks, and other modifications to the spectrum associated with well established physics. All of these must be taken into account so that one can compare directly the predictions corresponding to a particular collapse model to recent data from the CMB fluctuation spectrum and the matter power spectrum from recent galaxy surveys, and thus be able to put bounds on the model's parameters.   The paper is organized as follows. In Sect. \\ref{models}, we describe the theoretical framework in which to study the models to be tested.  Section \\ref{Description_schemes} gives some specifications about the proposed collapse schemes together with a brief physical motivation for them, while Sect. \\ref{Comparing_data} is dedicated to general predictions of these models, and a broad comparison to observational data. We define a fiducial model which is determined by the best fit values obtained by the WMAP collaboration with the same data used in this paper and without assuming any collapse of the wave function. It should be noted that the fiducial model has a value of the spectral scalar index $n_s$ different than 1, while calculations we have carried out here of the primordial fluctuation spectrum for models including collapse are restricted to the case $n_s=1$.  In Sect. \\ref{resultados}, we use data from WMAP 7-year release and the power spectrum of the Sloan Digital Sky Survery DR7 LRG, to put bounds on the free parameters of the collapse models, allowing also other cosmological parameters to vary. In Section \\ref{discusion}, we discuss our results and show our conclusions. ",
        "conclusions": "\\label{discusion}  In this paper, we have studied the phenomenological predictions of the collapse models developed  in Refs. \\cite{Perez2006,Sudarsky08,Unanue2008} and incorporated  the effect of late-time plasma physics. We have calculated the prediction for the CMB fluctuation spectrum in the case where the collapse of the $k$ modes of the scalar field during the inflationary period is included. Furthermore, we have performed a statistical analysis in order to compare the prediction of the collapse models with recent data from the CMB fluctuation spectrum and the matter power spectrum obtained by the SDSS collaboration, setting bounds on the model parameters $A$ and $B$ that characterize the collapse times of the scalar field modes. Results from the statistical analyses were discussed in Sect. \\ref{resultados} and can be summarized as follows:  \\begin{itemize} \\item{For $|A|>20$ in Model I, any value of $B$ provides a  good fit to the data within the range studied in this paper ($B= -10^9 \\cdots 10^9$).} \\item{For $|A| < 20$ in Model I, there is a nontrivial range of values of $B$ that provide a good fit to the data. However, values of $A$ lying in the range: $10^{-4}<A<20$ are not viable since the collapse of the modes occurs after the inflationary period. Furthermore, the data indicate that the values in the range $-20<A<-1$ are preferred.} \\item{The best fit values of $B$ obtained for $|A|<20$ in Model I lie in the range   $B < 1.88 \\,\\ {\\rm Mpc}$.} \\item{In Model II, there is a nontrivial range of values of $B$ that provide a good fit to the data for all values of $A$ tested ($A= -10^{9} \\cdots 10^9$). However, values of $A$ lying in the range: $10^{-4}<A<10^9$ are not viable since the collapse of the modes occurs after the inflationary period. Furthermore, the data indicate that the values in  the range  $|A|>1$ are preferred.} \\item{The best fit values of $B$ obtained for all values of $A$ studied in this paper for Model II  lie in the range   $B < 1.88 \\,\\ {\\rm Mpc}$.} \\item{The value obtained for the cosmological parameters is consistent within $3 \\sigma$ with those obtained by the WMAP collaboration assuming a standard inflationary scenario and also with bounds established by BBN and the value of $H_0$ obtained with the HST.} \\end{itemize}    This analysis allows us to  compare the value of the scale factor at the collapse time $a(\\eta_k^c)$, with the traditional value of the scale factor at ``horizon crossing'' which is often set to  mark the  {\\it quantum to classical transition} in the standard explanation of inflation: $a_k^H$. The ``horizon crossing'' occurs when the length corresponding to the mode $k$ has the same size that the Hubble radius $H_I^{-1}$ (in comoving modes, $k=aH_I$) therefore, $a_k^H \\equiv a(\\eta_k^H) = \\frac{k}{H_I} = \\frac{3k}{8\\pi G V}$. Thus the ratio of the value of the scale factor at horizon crossing for mode $k$ and its value at collapse time is \\begin{equation} \\frac{a^H_k}{a^c_k} = {k\\eta_k^c(k)} = A + B  k. \\end{equation}  The results discussed in Sect. \\ref{resultados} show that two cases are consistent with the data: the one for which the collapse time of inflaton field modes is previous to the time of ``horizon crossing'' of all modes of the inflaton field, and the opposite case. However, the data seem to favor the case in which the collapse happens before the modes cross the horizon.   As we found that there are good fits (for suitable ranges of $B$) for all values of the parameter $A$, we must conclude that the existing data do not lead to a preferred scale for the times of collapse of the wave function for the relevant modes.  However, for each value of $A$, interesting bounds can be placed on the value of the parameter $B$ which in this approach characterizes the modifications of the primordial spectral shape with respect to the HZ conventional flat scale-free spectrum."
    },
    "1112/1112.1692_arXiv.txt": {
        "abstract": "We produce three-dimensional Monte-Carlo radiative transfer models of the edge-on spiral galaxy NGC 891, a fast-rotating galaxy thought to be an analogue to the Milky Way. The models contain realistic spiral arms and a fractal distribution of clumpy dust. We fit our models to {\\it Hubble Space Telescope} images corresponding to the B and I bands, using shapelet analysis and a genetic algorithm to generate 30 statistically best-fitting models. These models have a strong preference for spirality and clumpiness, with average face-on attenuation decreasing from 0.24(0.16) to 0.03(0.03) mag in the B(I) band between 0.5 and 2 radial scale-lengths. Most of the attenuation comes from small high-density clumps with low ($\\lesssim$10\\%) filling factors. The fraction of dust in clumps is broadly consistent with results from fitting NGC 891's spectral energy distribution. Because of scattering effects and the intermixed nature of the dust and starlight, attenuation is smaller and less wavelength-dependent than the integrated dust column-density. Our clumpy models typically have higher attenuation at low inclinations than previous radiative transfer models using smooth distributions of stars and dust, but similar attenuation at inclinations above 70$^\\circ$. At all inclinations most clumpy models have less attenuation than expected from previous estimates based on minimizing scatter in the Tully-Fisher relation. Mass-to-light ratios are higher and the intrinsic scatter in the Tully-Fisher relation is larger than previously expected for galaxies similar to NGC 891.  The attenuation curve changes as a function of inclination, with $R_{B,B-I}=\\frac{A_{B}}{E(B-I)}$ increasing by $\\sim$0.75 from face-on to near-edge-on orientations. ",
        "introduction": "Understanding the 3D structure of dust is crucial for studies of spiral galaxies. In addition to detailing the complex structure of the ISM an accurate representation of dust is necessary to correct measurements of starlight. This is important for a variety of studies including e.g., the Tully-Fisher (TF) relation, where the behavior of the dust as a function of inclination is required to correct integrated photometry \\citep{Verheijen01b}; constructing rotation curves of highly inclined galaxies, since attenuation can censor the rotation curve's inner rise, leading to erroneous mass models \\citep{Matthews01}; correcting measurements of disk kinematics for the asymmetric amount of dust extinction above and below the disk of face-on spirals \\citep{Bershady10a,Bershady10b}; and producing accurate radial surface-brightness profiles corrected for spurious broken-exponential morphologies caused by dust absorption \\citep{deJong96}.  Spirals with inclinations at or near 90$^{\\circ}$ (known as edge-ons) provide the best information about the vertical structure of the dust, both because the corrections for viewing angle are small but also because the dust is projected into easy to see lanes that can attenuate the midplane by more than 10 mag in the optical \\citep{Kylafis87}. As a result edge-ons are an important component to our knowledge of dust structure. Separating the individual components of a dust-light mixture of unknown composition is a difficult endeavor, and is frequently done via radiative-transfer (RT) modeling.  Monte Carlo (MC) RT models, used in astrophysics since the 1970s (for a discussion of early models see \\citealt{Witt77}), have been extensively employed to probe the structure of edge-on spirals. Recently, these models have come into prominent usage (e.g. \\citealt{Bianchi96,Kuchinski98,Matthews01,Baes03,Baes11}), spurred by the need to precisely track multiple scatterings off dust grains. While costlier than the direct analytical approach of \\cite{Kylafis87} in terms of CPU time, increases in processor speed and RAM limits coupled with rapidly dropping costs now allow detailed Monte Carlo models to be run on relatively short timescales. In addition, the current trend in high-performance computing is towards increased numbers of CPUs at slower clock speeds; this is also a benefit to Monte Carlo methods, as they are exceedingly well suited to parallelization and thrive in a distributed environment.  NGC 891 is the most studied nearby (d$\\approx$9.5 Mpc) edge-on spiral galaxy in the universe. It is almost exactly edge-on ($i\\approx 89.7^{\\circ}$ \\citep{Xilouris98}), with a rotation speed (212 km/s) similar to that of the Milky Way. Because the dust morphology of edge-ons appears to change significantly between galaxies with rotation speeds above and below 130 km sec$^{-1}$ \\citep{Dalcanton04}, comparing the Milky Way with other fast-rotating systems is especially relevant. NGC 891 is also one of the first galaxies discovered to have high-latitude HI, a discovery prompted by observations of high-latitude HI in our own galaxy \\citep{Gerard73,Strong78}. Indeed, NGC 891's potential as an analogue of the Milky Way is what aroused much of the early interest in this extragalactic system \\citep{Bahcall83,vanderKruit84}. More recently, evidence for a strong two-armed spiral pattern like that seen in grand design spirals such as M51 has emerged based on asymmetry in H$\\alpha$ and B band emission \\citep{Kamphuis07,Xilouris98}. NGC 891 appears to have a bar \\citep{GarciaBurillo95} as well as a nearby companion (UGC 1807, \\citealt{Mapelli08}), both of which dramatically increase the likelihood of the aforementioned grand design pattern being real \\citep{Elmegreen82}.  NGC 891 is also known for its abundance of high-latitude dust and complex dust substructure. Given the significant high-latitude HI found in NGC 891, it is not surprising to also find dust away from the midplane. \\cite{Howk97} used unsharp masks to study individual high-latitude dust clumps in NGC 891, and later found evidence that these extended substructures may be common to most spirals \\citep{Howk99}. If high-latitude dust caused by star-formation induced outflows is indeed abundant in spiral galaxies, it would act as a foreground screen preferentially over the star-forming spiral arms, increasing the apparent optical depth in the arms without increasing total dust content. While unsharp masks are very useful for enhancing clumpy substructure the masking process destroys quantitative information about the clumps. Therefore, a new procedure for highlighting the dust that preserves this information is needed.  Because there are some indicators of spirality and significant extraplanar dust component, NGC 891 is a clear example of the need for advanced 3D modeling which takes into account these features. Indeed, studies attempting to use smooth, axisymmetric models on NGC 891 have had significant difficulty in fitting the data due to the asymmetry in blue light. \\cite{Xilouris98} split the galaxy in half and used an infinitely long thin disk of stellar emission to reconcile the left and right hand sides of NGC 891 in the B and V bands, while additional dust and light components have been required by smooth models to fit NGC 891's mid-infrared emission \\citep{Popescu00,Bianchi08}. Some groups have added clumpy/non-axisymmetric structure to RT models of edge-on spirals (e.g. \\citealt{Kuchinski98,Mihos99,Matthews01,Misiriotis02,Bianchi08}); however they restrict their analysis to a small range of clumpy models, with dust modeled by hand or based on results for the Milky Way.  In this work we quantitatively fit the {\\it first} RT models to include both dust clumping and realistic spirality to F450W and F814W Hubble images of NGC 891. These RT models resemble real spiral galaxies from any inclination angle, allowing us to make accurate inferences about the dust properties of galaxies like NGC 891 as seen at all inclinations. In Section \\ref{sec:obsconsts} we present the archival data used in this work and our methods of analysis. In Section \\ref{sec:model} we discuss the parameter space of our model, our fitting algorithm, and our fitness metric. We present our results, including a brief {\\it post-facto} analysis of our shapelet-based fitness metric, in Section \\ref{sec:results}. In Section \\ref{sec:discussion} we discuss the photometric properties of our models and compare them to the literature. Finally, in Section \\ref{sec:conclusion}, we present our concluding remarks. ",
        "conclusions": "\\label{sec:conclusion}  We have optimized a 3D Monte Carlo scattered light RT model to HST F450W and F814W images (essentially B and I band) of NGC 891. Our model is the first RT model involving realistic dust clumping and spirality to be quantitatively fit to imaging data. Dust is treated using a fractal algorithm, while spirality is included using a logarithmic spiral parameterization with a coarsely adjustable arm-interarm width ratio.  In order to enhance the effects of non-axisymmetric dust we constructed $\\Delta \\aeff$ maps using smooth models from the literature, then used the coefficients of a high-order shapelet image reconstruction to compare our clumpy models with the data. Computing Monte Carlo models at a high resolution is computationally expensive, so we used a genetic algorithm to find 'good' solutions in our many-dimensional parameter space in a reasonable amount of time. The efficiency of the genetic algorithm also allowed us to include a wide range of parameter values in order to fully probe our parameter space.  We computed $30 \\times 50$-generation runs of the genetic algorithm and find a fairly wide variety of parameter values in the best-fitting models. Some of this variety is likely due to complex interdependencies of the free parameters, (e.g. central disk surface brightness, scale length, and scale height for dust and stars), while the rest is the result of degeneracies from choosing such a large parameter space.  Despite these degeneracies, the best-fitting models from our 30 genetic algorithm runs do a credible job at mimicking the general features of the edge on profile of NGC 891. Perhaps more importantly, our models appear to be realistic simulations of \\textit{face-on} galaxies as well. While some of this is due to our more advanced treatment of spirality, which automatically creates more realistic-looking spiral features than previous work, it is important to note that the amount of spirality and clumping were free parameters; in 29 out of 30 runs our algorithm chose minima with clear spirality (despite not constraining the direction of the spiral structure) and the median fraction of dust placed in clumps is 58\\%. This is the first quantitative measurement of the clumpiness of dust in edge-on spiral galaxies that is not dependent on empirically matching IR SEDs but rather on observable dust morphology, and is in general agreement with clumpy fractions obtained through the SED fitting method.  The preference for visible spiral features and clumpy dust has significant implications for the face-on attenuation in spiral galaxies. We find that while overall our models prefer low (A$_B$ and $A_I < 1$) face-on attenuation, clumping and spirality produce localized regions of enhanced dust absorption, creating Type II surface brightness profiles as well as very small ($\\lesssim 10\\%$) high attenuation filling factors. These clumps result in larger total attenuation than a smooth model at low inclinations while having comparable attenuation when viewed edge-on, where the clumps overlap to more closely mimic a smooth dust distribution.  Additionally, we probed the effect of the dust when viewed as a foreground screen on background light sources. While still highly concentrated in spiral arms, we find significantly higher optical depths than would be suspected when only considering the attenuation --- on average 1.2 mag larger in the B band and 0.5 mag larger in the I band at one scale-length. The discrepancy between the optical depths and the attenuation comes from both the admixture of dust and starlight as well as the effect of scattering. The optical depths of our models are broadly in agreement with some measurements from the literature but not others, which may be a result of a large range in the optical depths in real galaxies.  We find that our models predict a blue bulge and red disk for NGC 891 when viewed face-on, regardless of the attenuation. The smooth, axisymmetric fiducial model from the literature which we use to construct our $\\Delta \\aeff$ maps also has this unusual morphology, and while we use bulge parameters from the the literature the red disk is present in our models even though we allowed the central disk emissivities and scale-lengths and heights to vary. While perhaps not common, an inspection of surface brightness profiles of a sample of face-on galaxies shows that blue, star-forming cores and red disks do exist in intermediate-type spirals. However, we observe neither of these features in the Milky Way, despite the frequently claimed similarity between NGC 891 and our own galaxy.  Even with the increased face-on attenuation in the clumpy models relative to smooth dust models, virtually all of our simulations have smaller attenuations at all inclinations than would be expected given corrections based on simple radiative transfer models and empirical formulae in common use in the literature.  This result indicates that TF studies need reduced attenuation corrections by typically 0.5 mag in B band at i=60$^\\circ$, and 0.25 mag in I at the same inclination. Inclination corrections based on more advanced RT models of edge-on spiral galaxies, however, do fit some of our dustier simulations very well and point to the need to include star-forming regions as well as spirality in future efforts.  Our models also predict a significantly grayer attenuation curve than found for both a simple foreground screen of dust and an attenuation curve derived from UV-bright star forming galaxies. We also find that the attenuation curve becomes $\\sim$40\\% grayer with inclination, something not considered in currently favored attenuation models. Both the large amount of edge-on variability of the models and the gradual systematic increase in R$_{B,B-I}$ with inclination are a strong caution for extrapolating empirically determined attenuation curves to objects with high inclinations.  We have not been able to produce clumpy dust models with the realistic high-latitude dust ``chimneys'' seen in the data, although it is unlikely including them would significantly change our results due to their low spatial frequency and small size. It is unclear whether the lack of high-latitude dust in our models is due to a poor choice of fitness function or the fractal geometry, although the fractal algorithm is certainly not optimized to produce extended dust ``fingers''. We find that the shapelet coefficients of the data have fewer extreme (large or small) coefficient values as well as a better balance of power between odd and even coefficients than most of the clumpy models; these observations will help us to improve our fitness metric.  Finding a method for computing model fitnesses more sensitive to small-scale structures would help improve the reproduction of dust substructure regardless of the clumping formalism. Multi-stage fitting processes may be useful to reduce the size of parameter space; global parameters (like central brightness, scale-length, and spirality) could be determined from smooth, lower resolution models covering the whole galaxy then fed into a much higher resolution model where the only free parameters relate to the dust density and clumping.  In producing quantitative RT modeling which includes spirality and clumping we have made the first foray into a new generation of highly detailed simulations that produce images capable of resembling real galaxy morphologies at any inclination. Extending clumpy, non-axisymmetric models based on our prototype to larger fields of view and large wavelength baselines, we expect to be able to constrain the 3D structure of spiral disks using physically realistic simulations as well as more accurately predict thermal dust re-emission in the mid-infrared. Additionally, the increase in resolution of far-IR images provided by the {\\it Herschel Space Observatory} will enable the next generation of SED models to fit SEDs to multiple positions of individual galaxies, where the signatures of non-axisymmetric structures will be more difficult to average out. By creating a realistic form of spiral structure and dust clumping we provide a blueprint for future work in this area."
    },
    "1112/1112.3374_arXiv.txt": {
        "abstract": "In this paper, we study the heating of the magnetized solar chromosphere induced by the large fraction of neutral atoms present in this layer. The presence of neutrals, together with the decrease with height of the collisional coupling, leads to deviations from the classical MHD behavior of the chromospheric plasma. A relative net motion appears between the neutral and ionized components, usually referred to as ambipolar diffusion. The dissipation of currents in the chromosphere is enhanced orders of magnitude due to the action of ambipolar diffusion, as compared to the standard ohmic diffusion. We propose that a significant amount of magnetic energy can be released to the chromosphere just by existing force-free 10--40 G magnetic fields there. As a consequence, we conclude that ambipolar diffusion is an important process that should be included in chromospheric heating models, as it has the potential to rapidly heat the chromosphere. We perform analytical estimations and numerical simulations to prove this idea. ",
        "introduction": "The lower solar atmosphere $-$ photosphere and chromosphere $-$ is usually treated using the approximation of magnetohydrodynamics (MHD). Different dynamical processes, like, e.g., wave propagation, formation of magnetic structures, flux emergence, etc., have been quite successfully described with this formalism. Nevertheless, because of the rather cool temperatures of the photosphere and low chromosphere, the degree of ionization there is very small, reaching values as low as $10^{-4}$ in the temperature minimum, and remaining always well below unity even at larger heights \\citep[see, e.g.][]{valc}. This fact, together with the decrease of collisional coupling with height, leads to a break of the assumptions underlying MHD.  In the upper photosphere and chromosphere, collisions are unable to couple completely the neutral and ionized components. Due to that, several new effects appear, such as non-ideal Hall effect and ambipolar diffusion.  In plasma physics, the term ``ambipolar diffusion'' refers to the diffusion of positive and negative particles at the same rate due to their Coulomb interaction, which maintains the charge neutrality at scales larger than the Debye length. In astrophysics, however, ambipolar diffusion usually refers to the decoupling of neutral and charged components. Ambipolar diffusion causes the magnetic field to diffuse though neutral gas due to collisions between neutrals and charged particles, the latter being frozen-in into the magnetic field. The use of the same terminology applied to these two different phenomena may lead to confusion. In this paper we follow the astrophysical definition of the ambipolar diffusion\\footnote{Although the term ``ambipolar diffusion'' is widely used in the literature, we consider more appropriate the term ``diffusion by neutrals'', or ``neutral diffusion''. In the rest of the paper we will use both terminologies interchangeably.}. As for the Hall effect, it appears as a result of the different drift velocities of electrons and ions, which are not equally affected by the presence of neutrals. The spatiotemporal scale over which the Hall effect operates in a partially ionized plasma is very different from the case of a fully ionized plasma \\citep{Pandey+Wardle2008}.  Recently, there has been an increasing number of works in the literature showing the importance of the deviations from MHD in different situations. The presence of neutral atoms in partially ionized plasmas significantly affects wave excitation and propagation \\citep{Kumar+Roberts2003, Khodachenko2004, Khodachenko2006, Forteza2007, Pandey2008, Vranjes2008, Soler2009, Soler2010,Zaqarashvili2011}. The relative motion between the neutral and charged species causes an increase of the collisional damping of MHD waves in the photosphere, chromosphere and prominence plasmas. This can be an important way to leak energy from Alfv\\'en and fast mode waves, causing the damping of coronal loop oscillations \\citep{Khodachenko2004, Khodachenko2006} and prominence oscillations \\citep{Forteza2007}. Under certain conditions, the leakage of Alfv\\'en wave energy can become the source of instabilities in the medium due to the non-ideal Hall effect \\citep{Pandey2008}. The excitation rate of Alfv\\'en waves by foot-point motions of magnetic structures significantly decreases if the neutral component is taken into account \\citep{Vranjes2008}. Another effect is the possible appearance of cut-off frequencies for Alfv\\'en and kink waves, as mentioned by \\citet{Soler2009} \\citep[see, however,][]{Zaqarashvili2011}.   The plasma partial ionization is also important for magnetic reconnection, for the same reason as for waves. \\citet{Zweibel1989} has shown that the reconnection rate depends strongly on the collisional coupling between ionized and neutral species and \\citet{Brandenburg+Zweibel1994, Brandenburg+Zweibel1995} concluded that, due to the action of ambipolar diffusion, oppositely oriented magnetic field lines can be brought sufficiently close to facilitate the reconnection. In a series of papers, \\citet{Sakai2006}, \\citet{Smith+Sakai2008}, and \\citet{Sakai+Smith2009} applied a two-fluid approximation to study the reconnection between two current loops. The two-fluid approach allowed these authors to find different temperatures of the ionized and neutral species, and give rise to proton heating and jet-like phenomena that, possibly, can be the reason for footpoint heating of coronal loops, explosive events and jets in the transition region and sunspot penumbrae.   All these non-ideal plasma effects can also modify the equilibrium balance of photospheric and chromospheric structures. An interesting mechanism to concentrate kilo-Gauss magnetic flux tubes at photospheric level, due to the action of the Hall current, was proposed by \\citet{Khodachenko+Zaitsev2002}. At the chromosphere, the order of magnitude stronger dissipation of currents perpendicular to the magnetic field, compared to that of longitudinal currents, was found to facilitate the creation of potential force-free field structures \\citep{Arber2009}. In prominences, the frictional force generated due to collisions between neutrals and ions was found to play a role in supporting their structure against the gravitational force \\citep{Gilbert2002}.  Another phenomenon affected by non-ideal plasma effects is magnetic flux emergence \\citep{Leake+Arber2006, Arber2007}. These authors showed that the amount of emerged flux is greatly increased by the presence of a diffusive layer of partially ionized plasma. Additionally, including neutrals removes the non-realistic lifting of low-temperature plasma into the lower corona and also avoids the subsequent Rayleigh-Taylor instability.   All the examples given above indicate that incorporating non-ideal plasma effects is essential and can lead to new insights into the physics of many phenomena taking place in the solar photosphere and chromosphere. Despite this increasing evidence, we are far from a complete understanding of the influence of these effects. The aim of the present paper is to investigate the consequences of the Joule dissipation of electric currents into the heating of the magnetized chromosphere. In the presence of neutrals, the ambipolar (or neutral) diffusion is expected to be orders of magnitude larger than the classical ohmic diffusion, leading to Reynolds number around unity \\citep[see, e.g.,][]{Khodachenko2004, Arber2007}. Diffusion by neutrals creates, in addition, anisotropy in the properties of the plasma since currents perpendicular to the magnetic field dissipate much quicker \\citep{Arber2009}. Due to that, one might expect important plasma heating. This work was inspired by the recent research by \\citet{DePontieu1998}, \\citet{Judge2008} and \\citet{Krasnoselskikh2010}, pointing in the direction that current dissipation, enhanced by the presence of neutrals in a plasma not entirely coupled by collisions, can play an important role in the energy balance of the chromosphere and above. In the rest of the paper, we will show analytical calculations and numerical simulations in different models (from purely academic to more realistic) of the amount of heating of chromospheric flux tubes due to the ambipolar diffusion effect. ",
        "conclusions": "In this work, we have investigated the heating of the solar magnetized chromospheric plasma produced by current dissipation (Joule heating). The current dissipation is enhanced by orders of magnitude due to the presence of a significant amount of neutral atoms in the chromospheric plasma not entirely coupled by collisions. The net relative motion between the ionized and neutral plasma components results in an additional diffusion term (ambipolar or neutral diffusion). The value of the ambipolar diffusion coefficient depends on the fraction of neutral atoms and on the magnetic field strength and can be as much as five orders of magnitude larger in the chromosphere than the classical ohmic diffusion coefficient. This implies evolution time scales of about 10--100 seconds. Thus, all the processes related to the ambipolar diffusion, including the plasma heating, happen very quickly.  We have performed analytical estimates and numerical simulations in different models of the amount of plasma heating due to diffusion by neutral particles. The simulations covered a variety of scenarios, from a simple atmosphere without gravity to a more complex flux tube model, resembling a chromospheric quiet network area. All our calculations show that the lower chromosphere can be  heated on temporal scales of $10^2$ seconds, giving rise to a temperature increase of several kK. The only requirements for the heating is the existence of a non force-free magnetic field. We obtain that the temperature increase is significant above a height of 1 Mm, depending on the strength of the magnetic field.  From those results, we conclude that current dissipation enhanced by the action of ambipolar diffusion is an important process that is able to provide a significant energy input into the chromosphere. We suggest that this mechanism should be included into the future models of chromospheric heating. Here we would like to stress that, apart from possible dissipation mechanisms, the action of the ambipolar term in the energy equation would stop in two cases. Firstly, this will take place if the temperature increase is so significant that all atoms become ionized leading to $\\eta_A=0$. Obviously, this is not the case of the chromosphere which is observed to be, at least, partially neutral. Note that our calculations indicate that the time scales of heating may become very long after a fast initial temperature increase (see Figures \\ref{fig:ionization300} and \\ref{fig:ft2ionization}). This would mean that the actual complete ionization is never reached in a reasonable amount of time. Secondly, the heating due to ambipolar diffusion will stop if the magnetic field configuration becomes force-free, meaning $\\vec{J} \\parallel \\vec{B}$, i.e. $\\vec{J}_{\\perp}=0$. This might lead to different temperatures and ionization fractions for different starting magnetic field configurations and will need further detailed investigation. Finally, the action of the ambipolar diffusion term has to be balanced by a dissipation mechanism. It is well known since the 60's that the chromosphere is cooled by radiative losses \\citep{Athay1966}. Whether the proposed mechanism is able to balance these losses, and whether it is able to provide sufficient energy to support the chromospheric high temperature, has to be investigated in future realistic 3D MHD chromospheric simulations, including a detailed radiative transfer and all complex chromospheric physics.  Here, we would just like to give hints that the temporal scales of the heating are defined by the  magnetic field strength via the ambipolar diffusion coefficient $\\eta_A$ (Eq.~\\ref{eq:etaa}). For the stronger fields considered by us ($B=300-900$ G), the heating happens very quickly, achieving a temperature increase of 2--3.5 kK in a 1 minute time. For the weaker field strengths (30--40 G), expected in the quiet chromosphere, the plasma is heated by 2 kK in a 10 minutes time. These time scales are comparable to the scales of radiative cooling of the plasma \\citep[a simple Newton formula gives characteristic cooling times at this height of the order of $10^2-10^3$ seconds, see][figure 3]{Mihalas+Toomre1982}. This argument may be used to suggest that one process might, in fact, balance the other.  The radiative damping would tend to decrease the temperature in the chromosphere. This, in turn, would produce a decrease of the ionization fraction, a subsequent increase of $\\eta_A$ and a decrease of the heating time scale (see Figs.~\\ref{fig:ionization300} and \\ref{fig:ft2ionization}). The neutral diffusion term would be put in action again, leading to a new temperature enhancement. The detailed temperature balance would depend on the details of non-LTE chromospheric radiative transfer and on the scales of the non-LTE ionization equilibrium.  The present study has several limitations. The calculation of the ionization balance is done in LTE conditions, assuming instantaneous ionization $-$ recombination. Several recent studies have shown that the {ionization-recombination} of (at least) hydrogen and calcium in the solar chromosphere does not follow LTE relations or instantaneous statistical equilibrium \\citep{Leenaarts2006, Wedemeyer2011}. In addition, we have discussed the energy losses by radiative damping, but have not included them in our simulations. Thus, in the simulations shown in this paper, the temperature would ever increase until the chromosphere is (unrealistically) completely ionized or until the magnetic field configuration becomes force-free. This, however, was done on purpose since we wish here to investigate the action of a newly proposed mechanism in simple situations, to better understand its implications into the general picture. In our future studies, we plan to include the effects of radiative damping in our calculations."
    },
    "1112/1112.4762_arXiv.txt": {
        "abstract": "A dependence of the neutrino masses on the dark energy scalar field could provide a solution to the {\\it why now} problem of dark energy. The dynamics of the resulting cosmological model, {\\it growing neutrino quintessence}, include an attractive force between neutrinos substantially stronger than gravity.  We present a comprehensive approach towards an understanding of the full cosmological evolution including the formation of large--scale neutrino structures. Important effects we account for are local variations in the dark energy and the backreaction on the background evolution, as well as relativistic neutrino velocities. For this aim, we develop a relativistic $N$--body treatment of the neutrinos combined with an explicit computation of the local quintessence field.  At its current stage, the simulation method is successful until $z \\approx 1$ and reveals a rich phenomenology. We obtain a detailed picture of the formation of large--scale neutrino structures and their influence on the evolution of matter, dark energy, and the late--time expansion of the Universe. ",
        "introduction": "\\label{sec:introduction}  In today's cosmological standard scenario, a cosmological constant $\\Lambda$ is assumed to explain the observed accelerated expansion of the Universe. Although so far consistent with all major observational probes \\cite{WoodVasey07, Reid09, Komatsu10, Sullivan11}, it faces two fundamental and unresolved problems.  Besides its disturbingly tiny value (the {\\it cosmological constant problem}), it remains miraculous why $\\Lambda$ has become important just recently (the {\\it why now} or {\\it coincidence problem}).  It is thus tempting to think of alternative cosmological scenarios in which these two problems are alleviated. In this work, we study growing neutrino quintessence, which addresses both of the problems. Growing neutrino quintessence \\cite{Amendola07, Wetterich07} relies on two assumptions. First, the dark energy component is described by a dynamical scalar field.  Second, the neutrino masses depend on this field. Studying the background evolution, it is found that the coincidence problem can indeed be solved for a neutrino--cosmon coupling somewhat larger than gravity.  In order to explore the implications of this model, its evolution has to be understood also on the perturbation level. Linear perturbation theory, however, breaks down even at large scales \\cite{Mota08}. This is due to large overdensities in the neutrino fluid becoming non--linear at $z \\lesssim 2$ on supercluster scales.  An understanding of the non--linear evolution, although of utmost importance for confronting the model with observational data, is still lacking. In first attempts, some aspects of the model have been studied with hydrodynamical and $N$--body methods \\cite{Wintergerst09, Baldi11}. Yet, three crucial aspects of the model have not been included so far. First, as shown by an analytical study of single non--linear neutrino overdensities, local variations of the neutrino mass can become very large \\cite{Nunes11}. Second, numerical results \\cite{Baldi11} show that neutrino particles are likely to reach highly relativistic velocities and thereby leave the Newtonian limit. Third, the local mass variations as well as the relativistic corrections change the average energy--momentum tensor of neutrinos; this {\\it backreaction} effect alters the background evolution of the dark energy scalar field.  It is the purpose of this work to form the basis for an adequate $N$--body based simulation method capable of including all major effects of growing neutrino quintessence. For the first time, we implement an explicit computation of the local dark energy scalar field for every time step.  This allows us to account for local mass variations and backreaction effects. We present a relativistic treatment of the neutrino dynamics as appropriate for the high velocities of neutrinos occurring during the structure formation process. The results of linear perturbation theory are used for the initial conditions. All these aspects are not merely small corrections, but decisive for every quantitative prediction.  The organization of this paper is as follows. We present growing neutrino quintessence and derive the equations of motion required for our simulation in \\sec{growing}. In \\sec{strategy}, we discuss strategies how to account for the peculiar features of growing neutrino quintessence in a numerical simulation. The numerical results are presented in Secs.~\\ref{sec:neutrino} and \\ref{sec:impact}. At first, we have a look at the formation of neutrino structures (\\sec{neutrino}), and then explore the effects on the evolution of dark energy and matter (\\sec{impact}). We summarize in \\sec{summary}. ",
        "conclusions": "\\label{sec:summary}  The cosmological dynamics of growing neutrino quintessence show a much higher complexity than the standard $\\Lambda$CDM cosmology or models of uncoupled dark energy. This work has made a crucial step towards a consistent simulation of growing neutrino quintessence including all its relevant effects.  The standard methods of linear perturbation theory and (Newtonian) $N$--body simulations, while very successful in the $\\Lambda$CDM case, rely on approximations in conflict with the effects of growing neutrino quintessence. Linear perturbation theory breaks down on all scales \\cite{Mota08}, and in Newtonian $N$--body simulations, the particles' velocities exceed the speed of light \\cite{Baldi11}. Furthermore, the local variation of the average neutrino mass due to perturbations in the dark energy scalar field could so far only be studied in idealized configurations \\cite{Nunes11}. None of the applied methods has shed light on the expected backreaction \\cite{Schrempp09} of the highly non--linear neutrino perturbations on the evolution of the cosmological background.  We have thus decided to develop a new method coping with these challenges right from the start. This effort has led to an $N$--body based simulation, designed from scratch and adjusted to growing neutrino quintessence. In particular, we incorporate local variations in the dark energy scalar field and respect relativistic dynamics for the neutrino component. We run the simulation for an exemplary choice of model parameters. Our results show that all the effects that had to be neglected in previous approaches are indeed important.  The simulation confirms the formation of large--scale neutrino structures (already predicted in \\cite{Brouzakis07}), and we give a detailed description of their evolution (\\sec{growth}).  As a main result, we observe a significant neutrino mass suppression of an order of magnitude inside concentrated structures (\\sec{individual}) agreeing with \\cite{Nunes11}. This is the consequence of significant inhomogeneities in the dark energy scalar field, the cosmon $\\varphi$. Once these become too large, our method encounters numerical difficulties (\\sec{numerical}), which is why we do not follow the cosmological evolution for redshifts $z<1$.  We find that the velocities of neutrinos are accelerated to relativistic values during the process of structure formation. The average neutrino equation of state reaches values $w_\\nu \\approx 0.1$ at $z\\approx 1$ (\\sec{relativistic}). Further, we have shown that the relativistic motion induces a difference between the two gravitational potentials. In the vicinity of large neutrino structures, the difference $\\Phi-\\Psi$ typically amounts to $\\sim 10^{-5}$.  Local mass variations and relativistic velocities of the neutrinos lead to strong backreaction effects (\\sec{backreaction}). We have demonstrated that the backreaction can notably modify the late--time expansion of the Universe (\\sec{backreactionandquintessence}). It is likely that the onset of the accelerated expansion occurs later as compared to a homogeneous approximation.  The evolution of matter perturbations is affected by the neutrino--induced gravitational potential, which dominates on large scales (\\sec{matter}). As a consequence, large--scale bulk flows of matter are amplified by a factor close to two at $z=1$. The density field, however, is much less affected. On large scales, we have found an effect of about $10 \\%$.  \\newline  The current stage of the simulation method presented in this work is already promising, and the efforts to extend its range of applicability will continue. When the cosmological evolution can be followed until $z = 0$ for a collection of different model parameters, a confrontation of growing neutrino quintessence with observational data becomes possible.  So far, we are only able to identify several effects that might be interesting regarding observations. The large--scale neutrino--induced gravitational potentials could be observable directly via gravitational lensing or indirectly via the resulting large--scale bulk flows of matter and the enhanced density power spectrum. The time evolution of the gravitational potential, on the other hand, leaves imprints in the integrated Sachs--Wolfe effect. In contrast to $\\Lambda$CDM, we have observed increasing large--scale potentials in growing neutrino quintessence during the formation of neutrino structures.  The main motivation to search for alternatives to the cosmological constant scenario are the difficulties to understand the tiny value of $\\Lambda$ and the {\\it why now} problem. Any competing model should avoid them in the first place. Growing neutrino quintessence offers a mechanism to solve the {\\it why now} problem of dark energy. With a rich phenomenology beyond the standard scenario, the prospects are promising that the model eventually can be put to stringent tests once its non--linear evolution is completely understood.  \\newline"
    },
    "1112/1112.0141_arXiv.txt": {
        "abstract": "{The relativley large spread in the derived metallicities ([Fe/H]) of M dwarfs shows that various approaches have not yet converged to consistency. The presence of strong molecular features, and incomplete line lists for the corresponding molecules have made metallicity determinations of M dwarfs difficult. Furthermore, the faint M dwarfs require long exposure times for a signal-to-noise ratio sufficient for a detailed spectroscopic abundance analysis. } {We present a high-resolution (R$\\sim$50,000) spectroscopic study of a sample of eight single M dwarfs and three wide-binary systems observed in the infrared J-band. } { The absence of large molecular contributions allow for a precise continuum placement. We derive metallicities based on the best fit synthetic spectra to the observed spectra. To verify the accuracy of the applied atmospheric models and test our synthetic spectrum approach, three binary systems with a K-dwarf primary and an M-dwarf companion were observed and analysed along with the single M dwarfs. } {We obtain a good agreement between the metallicities derived for the primaries and secondaries of our test binaries and thereby confirm the reliability of our method of analysing M dwarfs. Our metallicities agree well with certain earlier determinations, and deviate from others. } {We conclude that spectroscopic abundance analysis in the J band is a reliable method for establishing the metallicity scale for M dwarfs. We recommend its application to a larger sample covering lower as well as higher metallicities. Further prospects of the method include abundance determinations for individual elements. } ",
        "introduction": "Although the M dwarfs constitute a large fraction of the detectable baryonic matter, we still lack a great deal of knowledge about our low-mass ($<0.6M_{\\odot}$) hydrogen-burning neighbours. Studies \\citep{1998ASPC..134...28H,2003PASP..115..763C} suggest that as much as 70\\% of all stars in the solar vicinity ($\\sim$10\\,pc) are M dwarfs which makes these objects essential when deriving quantities such as the initial mass function \\citep[IMF,][]{1955ApJ...121..161S} and the present day mass function. These frequently used functions are derived using the luminosity. The transformation to mass, based upon stellar evolution theories, is sensitive to chemical composition. Moreover, a study of the possible time dependence of the IMF function for the low-mass part, needs good stellar metallicity criteria for M dwarfs. Thus, if we are to create a realistic model of the galactic evolution and present day status, detailed studies of these faint but numerous objects are of great interest. The M dwarfs are needed in the understanding of main-sequence stellar evolution and to define a limit between stellar and substellar objects. Finally, a well defined metallicity scale for M-dwarfs is essential to determine whether or not the general trend towards supersolar metallicities among FGK-stars planet hosts \\citep[e.g.][]{2005ApJ...622.1102F} holds also for cooler objects.  Spectroscopic studies of M dwarfs at high resolution have proven to be a difficult task. In the low-temperature regime occupied by these targets (2000$\\lesssim$T$_{\\rm eff}\\lesssim$4100\\,K), the optical spectrum is covered by a forest of molecular lines, hiding or blending most of the atomic lines used in spectral analysis. However, models of low-mass late-type stars have undergone continuous improvements, from the early work by \\citet{1969lls..symp..457T}, \\citet{1969ApJ...157..799A} and \\citet{1975A&A....38..283M,1976A&A....48..443M}, to the work by \\citet{1993PASAu..10..250B}, \\citet{1995A&A...295..736B,1995A&AS..109..263B}, \\citet{1995ApJ...445..433A}, \\citet{1999ApJ...512..377H} with their extensive grid of NextGen models, and the improved MARCS models \\citep{2008A&A...486..951G}. New molecular and atomic line data are collected and organised in large databases such as VALD \\citep{1999A&AS..138..119K} and VAMDC \\citep{2010JQSRT.111.2151D} and will improve the situation further.  The faintness characterising the M dwarfs has limited the number of high-resolution, high signal-to-noise studies. Non-sufficient resolution and dominant molecular features make it difficult to derive accurate atomic line strengths needed for a reliable metallicity determination. Metallicities have been derived via photometric calibrations \\citep{2005A_A...442..635B}, studies using molecular indices \\citep{2006PASP..118..218W}, as well as spectrum synthesis \\citep{2006ApJ...653L..65B,2006ApJ...652.1604B} and spectroscopic calibrations in the K band \\citep{2010ApJ...720L.113R}.  In this paper we present a detailed spectroscopic study in the J-band ($1100-1400$\\,nm), a spectral region relatively free from molecular lines. In the near infrared, atomic lines can be isolated and the lack of molecular lines allows a precise continuum placement. This spectral region was exploited for abundance analysis in the pioneering study of Betelgeuse, based on FTS spectra, by \\citet{VieiraThesis}, but has not been used much for late-type dwarfs until the last decade due to lack of efficient IR spectrometers at large telescopes. The present generation of IR echelles at large telescopes have, however, opened new possibilities, see, e.g. \\citet{2000ApJ...533L..45M,2007ApJ...658.1217M}.  Similarly to previous studies \\citep{2006ApJ...652.1604B,2005A_A...442..635B}, we select a number of binary systems with a solar-type primary and an M-dwarf companion to assess the accuracy of the atmospheric models used in the analysis and to verify the atomic line treatment. We then apply this result to a number of non-binary M dwarfs. Our choice of spectral region makes a careful analysis possible, for a sample of stars that are thought to have a high metal content ($-$0.35 to 0.5~dex), avoiding the large contribution of molecules such as TiO to the spectrum at optical wavelengths.  The paper is organized as follows. In Section~\\ref{sec:prevstud} we briefly summarize previous metallicity studies of M dwarfs. In Section~\\ref{sect:observations}, we describe the programme stars, the observations, and some aspects of the data reduction. In Section~\\ref{sect:analysis}, the ingredients of the spectrum analysis are presented -- compilation and derivation of spectral line data, stellar atmospheric parameters, and the procedure for metallicity estimation. In Section~\\ref{sect:results} we discuss the results of the analysis, and Section~\\ref{sect:conclusions} concludes the paper. ",
        "conclusions": "\\label{sect:conclusions} \\begin{figure}[htbp] \\centering \\includegraphics[scale=0.37]{feh.png} \\caption{Derived metallicities for our sample in comparison with previous abundance determinations (see Table \\ref{tab:feh} for references) as a function of spectral type (see Table \\ref{tab:targets}). For an easier presentation the targets with equal spectral types have been slightly shifted horizontally. In the upper right corner we show an error bar indicating the mean error of the metallicities derived in this paper.} \\label{fig:feh} \\end{figure}  \\begin{figure}[htbp] \\centering \\includegraphics[scale=0.37]{fehvsothers.png} \\caption{Differences in metallicities for our sample in comparison with previous abundance determinations (see Table \\ref{tab:feh} for references) as a function of metallicities from this study. A zero-line (dotted) to guide the eye is also plotted.} \\label{fig:fehvsothers} \\end{figure}  We have derived a new metallicity scale based on a careful spectroscopic analysis of high-resolution spectra (R$\\sim$50,000) observed in the J-band. Previous abundance studies are based on full spectroscopic analyses in the optical regime, low resolution spectroscopic infrared indices as well as purely photometric calibrations. Observations in the infrared J-band have the advantage of few and weak molecular features (FeH) which allows for a precise continuum placement as compared to the optical wavelength regions where the continuum is heavily depressed due to the many and strong TiO lines. We find that we can correct for lines introduced by the Earth's atmosphere quite succesfully by using a rapidly rotating calibration star, and by applying a carefully determined continuum placement. We verify the accuracy of the atmospheric models involved in the metallicity determination by observing three binary systems that are expected to have the same [Fe/H]. In two of the systems GJ105\\,AB and GJ250\\,AB we find that the metallicities agree within 0.01 and 0.02\\,dex, respectively. In the third binary system, HD101930\\,AB, we find a discrepancy of 0.11\\,dex, which is consistent with the derived errors.  We test the convergence of the procedure by starting from different initial guesses and find consistent solutions although some of the cooler objects show a greater spread in the determined metallicity.  Our sample covers a restricted range in \\teff\\ (between $\\approx$3200 and 3400\\,K, and a single object at 3900\\,K), and the extreme metallicities ($<-0.1$ and $>+0.2$) are not well covered. To explore the metallicity scale further, targets with a larger spread in the atmospheric parameters need to be observed, preferrably with an instrument that can cover a larger wavelength range efficiently. Improving the completeness of molecular line lists will improve the accuracy of the metallicity determinations, since there are still a number of unknown molecular blends present.  We conclude that a high-resolution spectroscopic analysis in the near infrared is a reliable method for metallicity determinations in this \\teff\\ regime. It is also the only method which will enable the determination of abundances of individual elements in M dwarfs."
    },
    "1112/1112.5552_arXiv.txt": {
        "abstract": "The ratio of the Fe abundance in the photosphere to that in coronal flare plasmas is determined by X-ray lines within the complex at 6.7~keV (1.9~\\AA) emitted during flares. The line complex includes the He-like Fe (\\fexxv) resonance line $w$ (6.70~keV) and Fe K$\\alpha$ lines (6.39, 6.40~keV), the latter being primarily formed by the fluorescence of photospheric material by X-rays from the hot flare plasma. The ratio of the Fe K$\\alpha$ lines to the \\fexxv\\ $w$ depends on the ratio of the photospheric-to-flare Fe abundance, heliocentric angle $\\theta$ of the flare, and the temperature $T_e$ of the flaring plasma. Using high-resolution spectra from X-ray spectrometers on the {\\em P78-1} and {\\em Solar Maximum Mission} spacecraft, the Fe abundance in flares is estimated to be $1.6\\pm 0.5$ and $2.0 \\pm 0.3$ times the photospheric Fe abundance, the {\\em P78-1} value being preferred as it is more directly determined. This enhancement is consistent with results from X-ray spectra from the {\\em RHESSI} spacecraft,  but is significantly less than a factor 4 as in previous work. ",
        "introduction": "The solar iron abundance remains a subject for discussion in the literature, both the value obtained from photospheric spectra and that from coronal spectra, generally in the X-ray or extreme ultraviolet spectral regions. Recent photospheric abundance determinations \\citep{asp09} using three-dimensional photospheric models give $A({\\rm Fe}) = 7.50 \\pm 0.04$ \\citep{asp09} and $7.52 \\pm 0.06$ \\citep{caf11} (abundances are here expressed logarithmically, on a scale where $A({\\rm H}) = 12)$, less by up to 60\\% (0.2~dex) than values current as recently as the mid-1990s \\citep{bla95}. Since the 1980s, it has become widely recognized that solar coronal element abundances differ from photospheric by factors that depend on the first ionization potential (FIP) of the element, with the coronal abundances of low-FIP (i.e. $\\lesssim 10$~eV) being enhanced but with the coronal abundances of high-FIP elements that are similar to their photospheric abundances. The reviews by \\cite{fel92} and \\cite{fel00} give enhancement factors for some low-FIP elements of up to 4. However, low-altitude coronal flare plasmas and other energetic events are stated to have photospheric abundances, while \\cite{fel90} find that the corona directly above a sunspot also has photospheric abundances but neighbouring coronal regions have low-FIP elements with enhanced abundances. Iron, with a FIP of 7.9~eV, is expected to have a low-FIP element behaviour. A recent analysis \\citep{phi11} of solar flare spectra observed by the {\\em Reuven Ramaty High Energy Solar Spectroscopic Imager} ({\\em RHESSI}) gives a rather different picture. {\\em RHESSI} spectra cover an energy range $\\sim 3$~keV to 17~MeV, and have a spectral resolution of $\\Delta E \\sim 1$~keV; in the soft X-ray range ($\\sim 3-10$~keV, $1.2-4$~\\AA), a strong complex of Fe lines, forming an emission line feature at 6.7~keV ($1.85$~\\AA), is evident on a continuum made up of free--free and free--bound emission. Under certain circumstances -- late in the flare development, with the {\\em RHESSI} thin attenuators in place -- an isothermal assumption for the flare's emitting plasma appears to be valid, and with electron temperature $T_e$ and emission measure ($EM= N_e^2 V$, where $N_e = $ electron density, $V = $ the emitting volume) estimated from the continuum, the Fe abundance can be determined. Analysis of nearly 2000 spectra observed during 20 flares led to the value $A({\\rm Fe}) = 7.91 \\pm 0.10$, with very little variation from flare to flare. This abundance is a factor of $2.6 \\pm 0.6$ more than recent photospheric abundance estimates, significantly less than 4.  \\cite{phi94} describe a means of obtaining the photospheric-to-flare Fe abundance ratio directly using the \\fexxv\\ resonance line $w$ (transition $1s^2\\,^1S_0 - 1s2p\\,^1P_1$), which is frequently the strongest line making up the 6.7~keV line complex, and a feature at 7.06~keV (1.76~\\AA) made up of Fe K$\\beta$ lines, with inner-shell transition ($1s-3p$) and formed by fluorescence of the photosphere by the coronal hot flare plasma. Spectra from the Bragg Crystal Spectrometer on the Japanese {\\em Yohkoh} spacecraft were used to find that the the flare Fe abundance was no more than a factor 2 higher than the photospheric abundance. A preliminary analysis \\citep{phi95} of \\fexxv\\ and Fe K$\\alpha$ line spectra observed by the {\\sc solflex} instrument on {\\em P78-1} spacecraft and the Bent Crystal Spectrometer (BCS) on {\\em Solar Maximum Mission} ({\\em SMM}) points to the photospheric and coronal Fe abundances being equal.  Here the analysis of \\cite{phi95} is revisited, with the theory and observations of the ratio of the Fe K$\\alpha$ lines to the \\fexxv\\ $w$ re-examined. The original observations from the SOLFLEX instrument on the {\\em P78-1} spacecraft are used, together with newly analyzed observations from the {\\em Solar Maximum Mission} Bent Crystal Spectrometer. The photospheric-to-flare abundance ratios are now consistent with the flare Fe abundance estimates from {\\em RHESSI}. ",
        "conclusions": "The possibility of obtaining the flare-to-photospheric Fe abundance ratio using the fluorescence-formed Fe K$\\alpha$ and K$\\beta$ lines and the \\fexxv\\ $w$ line, expressed in our earlier work \\citep{phi94,phi95}, is carried further here in the case of {\\em SMM} Bent Crystal Spectrometer and {\\em P78-1} SOLFLEX spectra in the 6.39--6.70~keV (1.85--1.94~\\AA) range. The six {\\em P78-1} measurements of the Fe K$\\alpha$-to-\\fexxv\\ $w$ line ratio indicate a flare-to-photospheric Fe ratio of $2.0 \\pm 0.3$, while the 19 {\\em SMM} measurements indicate a flare-to-photospheric Fe abundance ratio of $1.6 \\pm 0.5$. The slight difference is likely to be due to $\\gtrsim 20$\\% errors in the relative effective areas of BCS channels~2 and 7. The flare-to-photospheric Fe ratio from the {\\em P78-1} measurements alone is consistent, to within uncertainties, with recent estimates of the flare Fe abundance from {\\em RHESSI} spectra \\citep{phi11} which are a factor $2.6 \\pm 0.6$ higher than the photospheric Fe abundance \\citep{asp09,caf11}. This ratio is also consistent with {\\em Yohkoh} Bragg Crystal Spectrometer measurements of the Fe K$\\beta$ line \\citep{phi95}. These measurements of the Fe abundance enhancement may be relevant to the recent result of \\cite{woo10} suggesting that FIP enhancements in active stars are greatest for Sun-like stars (spectral type G), decrease to zero for cool stars of type K0, while for later spectral types the coronal abundances of low-FIP elements are diminished relative to photospheric. The zero enhancement of Fe in the corona above a sunspot \\citep{fel90} with effective temperature approximately equal to a K0 star and our own result for solar flares may fit such a picture."
    },
    "1112/1112.2691_arXiv.txt": {
        "abstract": "We report the first results from the $Z-$FOURGE survey: the discovery of a candidate galaxy cluster at $z=2.2$ consisting of two compact overdensities with red galaxies detected at $\\ga20\\sigma$ above the mean surface density.  The discovery was made possible by a new deep ($K_s\\la24.8$ AB $5\\sigma$) Magellan/\\fourstar\\ near-IR imaging survey with 5 custom medium-bandwidth filters. The filters pinpoint the location of the Balmer/4000\\AA\\ break in evolved stellar populations at $1.5<z<3.5$, yielding significantly more accurate photometric redshifts than possible with broadband imaging alone.  The overdensities are within $1\\arcmin$ of each other in the COSMOS field and appear to be embedded in a larger structure that contains at least one additional overdensity ($\\sim10\\sigma$).  Considering the global properties of the overdensities, the $z=2.2$ system appears to be the most distant example of a galaxy cluster with a population of red galaxies.  A comparison to a large $\\Lambda$CDM simulation suggests that the system may consist of merging subclusters, with properties in between those of $z>2$ protoclusters with more diffuse distributions of blue galaxies and the lower-redshift galaxy clusters with prominent red sequences.  The structure is completely absent in public optical catalogs in COSMOS and only weakly visible in a shallower near-IR survey.  The discovery showcases the potential of deep near-IR surveys with medium-band filters to advance the understanding of environment and galaxy evolution at $z>1.5$. ",
        "introduction": "Galaxy clusters are the most overdense cosmological regions in the Universe and are unique astronomical tools: their abundance constrains fundamental cosmological parameters and they provide an extreme laboratory for elucidating the role of local environment in the evolution of the massive galaxies.  However, despite extensive multi-wavelength searches, only a few evolved galaxy clusters with red galaxies have been found at $z\\ga1.5$ \\citep[e.g][]{mccarthy_compact_2007,andreon_jkcs_2009,papovich_spitzer-selected_2010,tanaka_spectroscopically_2010,fassbender_x-ray_2011,gobat_mature_2011,santos_discovery_2011}.  More diffuse protoclusters of star-forming galaxies have been found to higher redshift \\citep[e.g.][]{steidel_ly_2000,venemans_protoclusters_2007,capak_massive_2011}, although it is unclear how they relate to lower redshift massive clusters.  Ideally we want to find massive distant structures using spectroscopy. The problem is that most cluster galaxies at $z>1.5$ are too faint for spectroscopy, while photometric redshifts derived from broadband photometry are generally not accurate enough for secure identification.  A novel approach is to use near-infrared imaging with {\\it medium-bandwidth} filters, which are narrower than traditional broadband filters and provide significantly more accurate photometric redshifts \\citep[e.g., $\\delta z/(1+z)\\sim1-2\\%$ at $z\\sim2$;][]{van_dokkum_newfirm_2009,whitaker_newfirm_2011} for thousands of sources simultaneously over large contiguous fields of view \\citep[e.g., ][]{wolf_combo-17_2003}.  Using the newly commissioned \\fourstar\\ near-IR camera \\citep{persson_fourstar_2008} on the 6.5m Magellan Baade Telescope, we have initiated a major survey to obtain deep medium-bandwidth near-IR imaging over several fields. Our custom filters span $1.0\\micron-1.8\\micron$ and hence trace the Balmer/4000\\AA\\ break in galaxies at $1.5<z<3.5$.  The \\fourstar\\ Galaxy Evolution Survey ($Z-$FOURGE\\footnote{{\\url http://z-fourge.strw.leidenuniv.nl/}}) will be described in detail by Labb\\'e et al.,(in prep).  This Letter demonstrates that the improved depth and redshift accuracies allow us to search for massive galaxy overdensities at redshifts $z\\ga1.5$. ",
        "conclusions": "\\label{discussion}  Using the first data from the Magellan/\\fourstar\\ Galaxy Evolution Survey ($Z-$FOURGE), we have detected in the COSMOS field the most distant example of strong galaxy overdensities with red galaxies.  The system at $z=2.2$ extends $r\\sim1.5\\arcmin$ on the sky and appears to be made up of multiple subcomponents, including two overdensities detected at $\\ga20\\sigma$ above the mean and another at $\\sim10\\sigma$.  The two strongest overdensities, A \\& B, each resemble spectroscopically-confirmed galaxy clusters at $z\\ga1.6$ \\citep{papovich_spitzer-selected_2010,tanaka_spectroscopically_2010,gobat_mature_2011}: they contain significant populations of evolved galaxies (mean stellar age $\\sim1$~Gyr or $z_{form}\\sim3.3$) in a compact spatial distribution ($<250$kpc) embedded in massive halos of \\mh$\\sim1-6\\times10^{13}$\\msun.  Perhaps the most outstanding feature of the aggregate system is that it consists of multiple distinct galaxy overdensities in close spatial configuration.  Whether the three $z=2.2$ overdensities are a single gravitationally bound structure is not clear.  Ultimately, high precision spectroscopic redshifts of this structure are needed to confirm its existence and measure its velocity dispersion. Still, even if each overdensity is confirmed with spectroscopy, they could still be unbound and coupled to the Hubble flow. In this case, they possibly trace a filament in the dark matter density field and may evolve into a large-scale structure at $z=0$, e.g., a supercluster \\citep{geller_mapping_1989}.  To help us interpret the $z=2.2$ system as a whole, we searched for analogs in the $1$~Gpc$^3$ GiggleZ cosmological simulation (Poole et al., in prep.; see \\$\\ref{highz}) by running 2~million random line of sight (each a cylinder with $0.7$ Mpc radius and $50$ Mpc length to match the $2\\sigma$ redshift uncertainty, $2\\sigma_{z}\\approx0.06$, of the overdensities).  If a 3-halo system matching the set of overdensity \\mh\\ estimates is found, then 98\\% percent of the time 2 or 3 of these halos will merge into a $z=0$ cluster with \\mh\\ $\\sim10^{14}-10^{15}$ \\msun.  At $z=2.2$ in the simulation, these 2- and 3-halo systems show typical maximum separations of $\\sim0.9-1.8$ Mpc 3D from their mean positions.  Although this comparison is limited by the preliminary halo masses of the $z=2.2$ overdensities, these results may suggest there is high probability that two or more of the $z=2.2$ overdensities will merge by $z=0$.   If two or more of the overdensities are currently gravitationally bound, we may be viewing merging sub-clusters that are each in different evolutionary stages, including some whose galaxies are rapidly evolving and only just forming their red galaxy population (e.g., overdensity~C). Perhaps this signifies the system is in a transitional phase between the known $z\\ga2$ ``protoclusters'' \\citep[e.g.][]{steidel_spectroscopic_2005} with more diffuse distributions of blue galaxies, and the lower-redshift galaxy clusters with prominent red sequences.  Of course, protocluster observations have largely been optically-based, thus deep near-IR observations of such structures and deep optical spectroscopy of our candidate cluster are needed to understand the relationship between the various structures at $z\\ga1.5$.\\\\  The discovery of the $z=2.2$ system demonstrates the powerful combination of near-IR medium-bandwidth filters and deep imaging of the $Z-$FOURGE survey. These systems were undetectable in earlier optical and near-IR catalogs. For example, in the optically-selected ($i<25$) catalog of \\citet{ilbert_cosmos_2009}, the overdensity is completely absent in our redshift slice of interest ($z=2.1-2.3$).  The relatively bright optical limit of this catalog means evolved galaxies at these redshifts are missed entirely (only 2 of our red overdensity galaxies have $i_{total}<25$, see Fig.~\\ref{fig_cmd}).  Even with the $K$-band selected NMBS catalog of \\citet{whitaker_newfirm_2011}, only a weak overdensity is seen in Fig.~\\ref{fig_density}, as the majority of the $z=2.2$ overdensity galaxies are too faint to be detected in the 1.3 mag. shallower NMBS.  The first results of the $Z-$FOURGE survey suggest we are reaching a critical threshold in our ability to study galaxy evolution as a function of local environment at $z\\ga1.5$. By combining the spatial distribution with improved redshift information from deep medium-band filters in the near-IR, we can start to correlate the properties of $K_s-$selected galaxies with their environment, something that has so far only been done in detail for Lyman-break galaxies \\citep[e.g., ][]{steidel_ly_2000}."
    },
    "1112/1112.1417_arXiv.txt": {
        "abstract": "We present ground-based high-spatial resolution mid-infrared (MIR) observations of 20 nearby low-luminosity AGN (LLAGN) with VLT/VISIR and the preliminary analysis of a new sample of 10 low-luminosity Seyferts observed with Gemini/Michelle. LLAGN are of great interest because these objects are the most common among active galaxies, especially in the nearby universe. Studying them in great detail makes it possible to investigate the AGN evolution over cosmic timescale. Indeed, many LLAGN likely represent the final stage of an AGN's lifetime. We show that even at low luminosities and accretion rates nuclear unresolved MIR emission is present in most objects. Compared to lower spatial resolution \\textit{Spitzer}/IRS spectra, the high-resolution MIR photometry exhibits significantly lower fluxes and different PAH emission feature properties in many cases. By using scaled \\textit{Spitzer}/IRS spectra of typical starburst galaxies, we show that the star formation contribution to the 12\\mic emission is minor in the central parsecs of most LLAGN. Therefore, the observed MIR emission in the VISIR and Michelle data is most likely emitted by the AGN itself, which, for higher luminosity AGN, is interpreted as thermal emission from a dusty torus. Furthermore, the 12\\mic emission of the LLAGN is strongly correlated with the absorption corrected 2-10\\,keV luminosity and the MIR--X-ray correlation found previously for AGN is extended to a range from $10^{40}$ to $10^{45}$\\,erg/s. This correlation is independent of the object type, and in particular the low-luminosity Seyferts observed with Michelle fall exactly on the power-law fit valid for brighter AGN. In addition, no dependency of the MIR--X-ray ratio on the accretion rate is found. These results are consistent with the unification model being applicable even in the probed low-luminosity regime. ",
        "introduction": "One of the major discoveries of the last two decades in extragalactic astronomy was that the structure and evolution of the central kiloparsec region of galactic nuclei are strongly connected to the properties and evolution of the central super-massive black hole (SMBH) that presumably resides in the center of each galaxy. In particular, the growth of the SMBH seems to be related to major galaxy mergers leading to various phases of nuclear activity (see contribution by Duschl et al.). These active galactic nuclei (AGN) can be described by the unification model of AGN \\cite{antonucci_unified_1993} that contains the SMBH in the nucleus, surrounded by an accretion disk and, further outward, a dusty obscuring torus-like structure. This torus presumably blocks the direct view onto the central region depending on the viewing angle of the observer, therefore explaining the different AGN sub-classes observed. However, for a more accurate description, various modifications to this standard model may arguably be necessary (e.g., as suggested in the contributions by Elvis et al. and Ramos Almeida et al.). In addition, in the low-luminosity regime various changes in the structure of the AGN might occur. This is indicated by a different spectral energy distribution (SED) observed in these low-luminosity AGN (LLAGN)\\cite{ho_nuclear_2008}, theoretically explained by a change in the accretion disk structure from an optically-thick and geometrically-thin accretion disk \\cite{shakura_black_1973} to an optically-thin and geometrically-thick accretion flow (e.g., advection-dominated accretion flow; ADAF \\cite{narayan_advection-dominated_1994}) at low accretion rates (Eddington ratio $\\eta_\\mathrm{Edd} \\lesssim 10^{-2}$). At even lower rates ($\\eta_\\mathrm{Edd} \\lesssim 10^{-5}$), a jet of outflowing plasma along the rotation axis of the accretion disk might even dominate the whole SED \\cite{falcke_jet_2000}. Finally, theoretical work predicts that the dusty obscuring torus will disappear at low luminosities/accretion rates \\cite{hoenig_active_2007, elitzur_agn-obscuring_2006}. Observational indications for its disappearance have been found mainly for radio-loud LLAGN \\cite{donato_obscuration_2004}, while remaining unproven for the LLAGN in general (but see also  contribution by Fern\\'andez Ontiveros et al.).  Searches for the putative torus are most promising in the mid-infrared (MIR), where the thermal emission of the dust is peaking. In particular, ground-based high-spatial resolutions MIR observations yield valuable information about this probably compact ($\\lesssim 1$\\,pc) structure. It was found that the observed MIR emission of local AGN can be best explained by a torus consisting of many clouds \\cite{hoenig_dusty_2010-1,alonso-herrero_torus_2011} (and also contribution by Ramos Almeida et al.). Although MIR interferometry provides an excellent possibility to investigate and resolve the torus in nearby AGN (see contributions by Burtscher et al., Kishimoto et al., and  Tristram et al.), current its sensitivity is not sufficient to detect most LLAGN.  Even though lacking the high-spatial resolution, early MIR observations led to the discovery of a strong correlation between the nuclear MIR and X-ray emission \\cite{krabbe_n-band_2001, lutz_relation_2004}. This MIR--X-ray correlation can in principle be understood by a close connection between the ultraviolet and X-ray emission originating in the inner accretion disk region. While the hard X-ray emission can penetrate the obscuring torus almost unabsorbed as long as the column density remains in the Compton-thin regime, the ultraviolet is absorbed by the dust and re-emitted in the MIR as thermal radiation. Note that this explanation is independent of the torus orientation, because the MIR emission is nearly isotropic. Therefore hard X-rays are also a good direct tracer of the accretion activity and the current intrinsic power output of an AGN. By using high-spatial resolution MIR VLT/VISIR imaging, we could recently demonstrate for a sample of $\\sim 40$ local AGN that the MIR--X-ray correlation is indeed very strong, in particular when the hard 2-10\\,keV X-ray emission is corrected for absorption \\cite{gandhi_resolving_2009}. The present work represents a direct extension of our previous work into the low-luminosity regime and is driven by the main goal of probing the existence of a torus-like obscuring structure in LLAGN.  However, because of the large number of possible MIR emitters in this luminosity regime (Fig.\\ref{fig:1}), only a first step will be presented: high-spatial resolution Gemini/Michelle and VLT/VISIR MIR imaging of LLAGN. We can investigate the validity of the MIR--X-ray correlation for LLAGN and constrain possible contamination by circum-nuclear star formation affecting the MIR photometry. Star formation can significantly contribute or even dominate the emission at low luminosities. For a detailed description of the VLT/VISIR data we refer to \\cite{asmus_mid-infrared_2011-1}, while comparable observations are also presented in the contribution by Mason et al.  \\begin{figure}[th]  \\includegraphics[width=0.7\\linewidth]{Asmus_Fig1.png} \\hspace{0.05\\linewidth}% \\begin{minipage}[b]{0.25\\linewidth} \\caption{\\label{fig:1} Schematic view of the central structure of an LLAGN. The red arrows indicate the MIR emission of the various components and its nature. } \\end{minipage} \\end{figure} ",
        "conclusions": "The high-angular resolution MIR imaging of $\\sim 30$ nearby LLAGN showed point-like nuclear emission at $\\sim 50$\\,pc scale in all detected objects. No correlation to larger aperture data like Spitzer is evident from a comparison of VISIR photometry and IRS spectra. Therefore one has to be careful when deriving AGN properties from these larger aperture data. In many cases the observed fluxes are significantly lower and spectral features, as the PAH 11.3\\mic emission feature are absent on smaller scales. Furthermore, by using a starburst template SED, we showed that only $\\lesssim 30\\%$ of the high-resolution data may still be dominated by MIR emission from star formation occurring on $\\lesssim 50$\\,pc scale. The majority ($\\gtrsim 70\\%$) of LLAGN are, however, dominated by AGN emission on this scale, while its exact origin -- torus, jet or accretion disk -- remains undetermined at this point. Furthermore, the MIR--X-ray correlation is valid for all AGN observed at high-angular resolution in the MIR over the whole probed luminosity range from $10^{40}$ to $\\sim 10^{45}$\\,erg/s. Owing to the fact that no dependence of either optical classification, or radio-loudness, or host morphology was found, this correlation might put strong constraints on the universal AGN structure. Interestingly, none of the theoretically predicted changes in the accretion structure with decreasing accretion rate, e.g., onset of ADAF or the jet-dominance, seem to have any imprint on the MIR--X-ray luminosity ratio. However, the sample is likely incomplete (selection based on availability of 2-10\\,keV data) and possibly biased against highly obscured (Compton-thick) objects. This along with the small total sample size prevents us from reaching any final conclusions on the AGN structure and nature of the MIR emitter in the low-luminosity regime. At this point we can only conclude that the results agree with those for brighter AGN, for which the MIR emission is well explained by a dusty clumpy torus. Therefore, such a torus might also exist in LLAGN, i.e. the standard unification model of AGN might still hold in this regime, in agreement with \\cite{panessa_x-ray_2006,maoz_low-luminosity_2007}. Furthermore, these results imply that in LINERs, the emission is AGN-dominated in at least central tens of parsecs (compare \\cite{yan_nature_2011}). Finally, the MIR--X-ray correlation represents a powerful observational tool to convert between both wavelength bands for any AGN. In order to settle the question about the nature of the MIR emission, SED modeling with particular focus on the high-resolution infrared data seems most promising (see also contributions from Prieto et al. and Fern\\'andez Ontiveros et al.)."
    },
    "1112/1112.3248_arXiv.txt": {
        "abstract": "{Molecular counterparts to atomic jets have recently been detected within 1000 AU of young stars at early evolutionary stages. Reproducing these counterparts is an important new challenge for proposed ejection models.} {We explore whether molecules may survive in the magneto-hydrodynamic (MHD) disk wind solution currently invoked to reproduce the kinematics and tentative rotation signatures of atomic jets in T Tauri stars.} {The coupled ionization, chemical, and thermal evolution along dusty flow streamlines is computed for the prescribed MHD disk wind solution, using a method developed for magnetized shocks in the interstellar medium. Irradiation by (wind-attenuated) coronal X-rays and far-ultraviolet photons from accretion hot spots is included, with an approximate self-shielding of H$_2$ and CO. Disk accretion rates of $5\\times10^{-6}$, $10^{-6}$ and $10^{-7}\\Msun\\yr$ are considered, representative of low-mass young protostars (so-called `Class~0'), evolved protostars (`Class~I') and very active T~Tauri stars (`Class~II') respectively. } {The disk wind has an \"onion-like\" thermo-chemical structure, with streamlines launched from larger radii having lower temperature and ionisation, and higher H$_2$ abundance.The coupling between charged and neutral fluids is sufficient to eject molecules from the disk out to at least 9~AU. The launch radius beyond which most H$_2$ survives moves outward with evolutionary stage, from $\\simeq0.2$~AU (sublimation radius) in the Class~0 disk wind, to $\\simeq1$~AU in the Class~I, and $>1$~AU in the Class~II. In this molecular wind region, CO survives in the Class~0 but is significantly photodissociated in the Class~I/II. Balance between ambipolar heating and molecular cooling establishes a moderate asymptotic temperature $\\simeq700-3000$~K, with cooler jets at earlier protostellar stages. As a result, endothermic formation of H$_2$O is efficient, with abundances up to $\\simeq10^{-4}$, while CH$^+$ and SH$^+$ can reach $\\geq10^{-6}$ in the hotter and more ionised Class~I/II winds.} {A centrifugal MHD disk wind launched from beyond $0.2-1$~AU can produce molecular jets/winds up to speeds $\\simeq100~\\kms$ in young low-mass stars ranging from Class~0 to active Class~II. The models predicts a high ratio of H$_2$ to CO and an increase of molecular launch radius, temperature, and flow width as the source evolves, in promising agreement with current observed trends. Calculations of synthetic maps and line profiles in H$_2$, CO and H$_2$O will allow detailed tests of the model against observations.} ",
        "introduction": "Powerful supersonic bipolar jets are ubiquitous in young accreting stars, suggesting that jets may play a key role in extracting excess angular momentum from the accretion flow. The high ejection efficiency, and the universal jet collimation properties independently of circumstellar envelope density, both strongly favor a magneto-hydrodynamical (MHD) ejection and collimation process (see e.g.~\\citealt{KoPu00,shu00,cabrit-villard}). However, the main jet origin is still the subject of intense debate. Proposed options include a self-collimated, steady centrifugally-driven disk wind --- either from the inner disk edge near corotation (\"X-wind\") or from a more extended range of radii --- a pressure-driven stellar wind, and unsteady ejections by reconnexions in the sheared stellar magnetosphere, all three processes probably coexisting to some degree in young stars (see \\citealt{Ferr06} and refs.~therein).  Of the wide variety of models published, stationary solutions for \"X-winds\" and extended MHD disk-winds are of particular interest as they provide synthetic predictions in atomic lines that reproduce several aspects of jet collimation and acceleration in microjets from T Tauri stars \\citep{Shang98,SGSL02,Cab99,Garcia01b,Pes04,pyo-DG,pyo-HL-RW,Cabrit-IAU243}. Extended MHD disk winds have received further attention since the tentative detections of jet rotation signatures, that suggest an outer launch radius for the atomic jet% \\footnote{We will use the term \"atomic jet\" rather than \"optical jet\" to refer to the (mostly neutral) fast jet component emitting in atomic/ionic lines, as infrared detections e.g.~in [\\ion{Fe}{ii}] are increasingly frequent \\citep[see][for a review]{Bally-PPV}. Atomic jets differ by their high speed and collimation from the \"neutral atomic outflows\" detected in \\ion{H}{i} in a few young protostars, whose bulk mass is moving at $\\leq20~\\kms$ \\citep{lizano88,svs13}.} in T Tauri stars $\\simeq0.2-3$~AU, well beyond the disk inner edge \\citep{bac02,anderson03,Pes04,cof04,cof07}. Evidence for depletion of refractory species in atomic jets also suggests ejection from relatively large disk radii, outside the dust sublimation radius \\citep{Nis05, Pod06, Pod09, Amb11}. If they remove most of the angular momentum from the accretion flow, such extended MHD disk winds would have a strong impact on the conditions for planet formation and migration in disks \\citep{CombetFerr}. Therefore it is important to find further observational tests of their presence and radial extent.  An important new test for proposed jet engine(s) is whether they can explain the {\\it molecular} jet/wind counterparts recently detected within 1000~AU of young stars. Molecular jets were first discovered emanating from the so-called `Class~0' protostars at the earliest stage of star formation, in CO, SiO, H$_2$, and H$_2$O masers \\citep{Bach90, gui92,gg99,hh211-h2,claussen}. These narrow molecular jets are clearly distinct from the \"standard\" swept-up molecular outflows in that they are much faster, with deprojected speeds of $60-150~\\kms$ as opposed to $\\le20~\\kms$ in the swept-up gas, much more highly collimated on-axis, and with specific chemical enhancements \\citep{Bach90,gg99,tafalla09,tafalla10}. Their collimation scale, ejection/accretion ratio, and variation of physical parameters with velocity are strikingly similar to those of atomic T~Tauri jets, pointing to a similar origin \\citep{Cab07,Lee07,Nis-SiO}. Counterparts in H$_2$, and more rarely in CO, were then identified towards more evolved infrared protostars of Class~I \\citep{CR96,davis01,davis02,davis-AOsvs13,Bally-PPV} and optically visible T Tauri stars of Class~II \\citep{tak-dgtau,tak-hltau,herczeg06,hh30-flow,beck08}. They consist of high-speed H$_2$ in the jet beam itself, and/or of slow wind/cavities at $\\simeq30~\\kms$ encompassing the fast ionic jet, on scales of a few 100~AU. The molecular emission tends to peak at lower velocities than atomic lines \\citep{davis03}.  Three processes are generally invoked to produce molecular jets/winds in young stars: (i) entrainment of ambient molecular material, (ii) gas-phase reformation of H$_2$ in a dust-free atomic wind (e.g.~stellar or magnetospheric); (iii) ejection/reformation of molecules in a dusty wind from the disk surface.  The entrainment hypothesis is intuitively appealing, but appears quite challenging. Turbulent mixing of ambient gas along the walls of a supersonic jet beam predicts too low columns of H$_2$ \\citep{taylor}. Shock entrainment in narrow bowshock wings along the jet is potentially more efficient, but would require that ambient molecular gas can quickly mix-in and refill the cocoon behind preceding bowshocks \\citep{davis-hh111,raga93}. Swept-up molecular cavities driven into ambient gas by an atomic wide-angle wind have also been proposed in Class~II sources where molecular counterparts are broader than the atomic jet \\citep[see e.g.][]{tak-hltau}. However, the nature of such a wide-angle component remains unclear; the fast $\\simeq100~\\kms$ and powerful wide-angle wind predicted by the \"X-wind\" model, propagating into a flattened envelope, would sweep up a cavity of width to length $>1$ at 200~AU from the source in less than 1000 years \\citep{lee01,Shang06}. This is difficult to reconcile with the moderate width/length $<1$ of molecular \"cavities\" in Class~II sources of age $\\simeq10^6$ yrs, unless the wide-angle wind is much slower/weaker (e.g.~a disk wind launched from large radii). In any case, an ambient swept-up cavity would not easily explain the narrow, high-velocity molecular jets seen in Class~0/I sources. The presence of blueshifted H$_2$ or CO up to $50-100~\\kms$ within 10~AU of Class~II / Class~I sources (\\citealt{herczeg06} and Herczeg, private communication) also seems difficult to explain with entrainment.  Alternatively, it has been argued that molecular counterparts may trace the primary ejected material itself \\citep{CR96,tak-dgtau,hh30-flow,hirano}. Studying the formation/survival of H$_2$ in protostellar winds is a difficult problem as it requires the combination of state of the art MHD wind models with chemical networks and realistic radiation fields and heating/cooling terms, few authors having addressed it self-consistently in the past.  Reformation of H$_2$ in {\\it dust-free stellar winds} by ion reactions has been the most extensively studied so far. Spherical models with prescribed temperature laws showed that H$_2$ forms and survives only at high mass-flux rate and if temperature drops rapidly below 2000~K \\citep{Raw88,Glass89,Raw93}. \\citet{Ruden90} solved in parallel for the ionisation, H$_2$ abundance, and thermal evolution, including ion-neutral \"drag\" heating. They found that less dense winds are hotter and more collisionally dissociated, so that the final H$_2$ abundance exceeds a few percents only at very high wind mass-fluxes $>10^{-5}\\Msun~\\yr$. Such a model does not seem able to explain the presence of detectable H$_2$ in low-luminosity Class~0/I/II jet sources, where total mass-flux rates are $1-3$ orders of magnitude smaller \\citep{Lee07,antoniucci,HEG95}. Far-ultraviolet excess from accretion further reduces the predicted H$_2$ (and SiO) abundance at low wind mass-flux, even though CO is efficiently formed \\citep{Glass91}. Alternatively, \\citet{raga-h2form} show that H$_2$ abundances of up to 10\\% could be reached behind dense internal shocks in a dust-free jet. However, the jet velocity variability amplitude should be less than $15~\\kms$, which is quite restrictive. The effect of magnetic cushioning and of far-ultraviolet (FUV) photons remains to be investigated.  Here we investigate the third possibility, namely whether molecular jet counterparts could be tracing {\\it dusty} disk winds. We focus on self-similar, centrifugally-driven MHD disk winds. Pioneering calculations by \\citet{Safi93a} indicated that H$_2$ could survive collisional dissociation in the outer regions of such disk winds, beyond a minimum launch radius $\\sim0.5-3$~AU depending on the adopted MHD solution. However, the rate of collisional dissociation may have been underestimated, as the ionisation fraction was obtained assuming atomic gas. The faster dissociative recombinations occuring in molecular gas would enhance ion-neutral drag, increasing the wind temperature and the dissociation rate. The calculations also did not consider H$_2$ destruction by stellar FUV photons, coronal X-rays, or endothermic reactions with O and OH.  In the present paper, we explore this problem one step further by integrating {\\it in parallel} for the detailed chemistry, ionization, and thermal state along MHD disk wind streamlines. We use a model developed for magnetized molecular interstellar shocks with an extended chemical network of 134 species \\citep{FlPi03b}, adding irradiation by (attenuated) coronal X-ray photons and FUV photons from the accretion shock. Another difference with respect to \\citet{Safi93a} is the use of a slower and denser MHD accretion-ejection solution that better matches the rotation signatures, poloidal velocities, and ejection-accretion ratio in atomic T Tauri jets \\citep{Pes04,Ferr06}. The flow dynamics, thermo-chemistry, and treatment of FUV and X-ray radiation, are described in detail in Section~\\ref{s:model}, which may be skipped in a first reading. Our results are presented in Section~\\ref{s:results} for a range of accretion rates and stellar masses representative of the Class~0, Class~I, and most active Class~II phases of solar-type stars. In Section~\\ref{s:discussion}, the results are compared to static irradiated disk models and to observational trends, and model limitations are discussed. Conclusions are summarized in Section~\\ref{s:concl}. ",
        "conclusions": "\\label{s:concl} We have investigated the non-equilibrium thermal-chemical structure of dusty streamlines in a \"slow\" self-similar MHD disk wind compatible with current observational constraints in atomic T Tauri microjets. We considered a range of accretion rates and stellar masses representative of the Class~0, Class~I, and early Class~II stages of low-mass star formation, and probing a decrease of 2 orders of magnitude in wind density. A detailed chemical network, a complete set of heating/cooling terms, and the effect of irradiation by stellar X-rays and FUV photons, were considered. Our main conclusions are the following:  \\begin{itemize} \\item The MHD disk wind has an \"onion-like\" thermal-chemical structure, with temperature, ionization, and radiation field all decreasing as the launch radius of the streamline increases. The wind is sufficiently ionized to accelerate neutrals out to disk radii of $\\simeq9$~AU, but is sufficiently self-screened and cool to remain molecular on streamlines launched beyond some minimum radius $r_0$. For the MHD solution explored here, this radius is $r_0\\simeq0.2$ AU (sublimation radius) for Class~0 parameters, $r_0\\simeq1$ AU for the Class~I, and $r_0>1$~AU for the Class~II.  \\item Key elements for the survival of H$_2$ in MHD disk winds, as opposed to static irradiated disk atmospheres, are: the short flow crossing timescales ($50-100$ yrs on the 1 AU streamline), the efficient shielding provided by inner wind streamlines against stellar FUV and X-ray photons, and the strong adiabatic cooling which delays gas heating and limits collisional and chemical destruction.  \\item Balance between ion-neutral drag and molecular line cooling establishes asymptotic temperatures $\\simeq700/2\\,000/3\\,000$~K on the 1~AU streamline, increasing with protostellar stage. Temperatures at the wind slow magnetosonic point are within a factor 2 of the surface heating prescribed in the MHD solution, suggesting that the assumed high wind mass-loading might be naturally achievable (although thermal calculations deeper into the disk are needed to test this conjecture).  \\item The chemistry of dusty irradiated MHD disk winds is of a complex and hybrid nature, combining the out-of-equilibrium \"warm\" chemistry of C-shocks heated by ambipolar diffusion, and the photoprocesses present in PDRs and XDRs. Illustrations are the efficient gas-phase synthesis of H$_2$O at all stages, and the substantial abundances reached by SH$^+$, CH$^+$, and $\\Hr^+$ in the hotter and more ionized Class~I/II. In these latter cases, OH, O, and C also reach levels similar to H$_2$O while CO is strongly photodissociated.  \\item The model predictions appear promising to explain several observational trends in molecular wind counterparts, namely: the presence in low-mass Class~0 sources of highly collimated molecular jets with typical speed up to $100~\\kms$, low temperature $\\simeq700$~K, and estimated magneto-centrifugal launch radii $\\simeq0.2-0.6$~AU; as well as the trends for higher H$_2$ temperatures $\\simeq2\\,500$~K, larger magneto-centrifugal launch radii, and lower collimation and speed as the source evolves from Class~0 to Class~I to Class~II.  \\item More discriminant tests of the applicability of MHD disk winds to molecular jets/winds in young stars require synthetic maps and line profiles in H$_2$, CO, and H$_2$O, including a non-local treatment of self-shielding and proper account for beam dilution and NLTE excitation. Such calculations will be presented in following papers in this series, and compared with infrared and submm observations. The effect of internal shocks, in particular the enhancement of species believed to be released from grains, such as SiO and CH$_3$OH, will also be explored. \\end{itemize}"
    },
    "1112/1112.1767_arXiv.txt": {
        "abstract": "We critically analyze the Bell et al. findings on ``anomalous'' widths of high-order Hydrogen radio recombination lines in the Orion Nebula at 6 GHz. We review their method of modified frequency switching and show that the way this method is used for large $\\Delta n$ is not optimal and can lead to misinterpretation of measured spectral line parameters. Using a model of the Orion Nebula, conventional broadening theory, and Monte Carlo simulation, we determine a transition zone $n=224,\\,\\ldots,\\,241$ ($\\Delta n=11,\\,\\ldots,\\,14$), where measurement errors grow quickly with $n$ and become comparable with the measurement values themselves. When system noise and spectrum channelization are accounted for, our simulation predicts ``processed'' line narrowing in the transition zone similar to that reported by Bell et al. We find good agreement between our simulation results and their findings, both in line temperatures and widths. We conclude, therefore, that Bell et al.'s findings do not indicate a need to revise Stark broadening theory. ",
        "introduction": "In the 1990s, Morley Bell and coauthors developed a technique for measuring weak spectral lines by reducing broad baseline variations (\\citet{Bell1997}). This technique was referred to as ``modified frequency switching'' and was used to detect of weak atomic and molecular/maser lines in spectra of gaseous nebulae, circumstellar envelopes, and star formation regions (see, e.g., \\citet{Bell&Feldman1991,Belletal1992,Belletal1993}).  Bell and coauthors used this technique to measure radio recombination line (RRL) widths and temperatures (\\citet{Belletal2000}, hereafter BASV). At 6 GHz, they found that the ``...{[}processed H{]} \\emph{lines at large n are both narrower and stronger than expected from theory}...'' and suggested, ``\\emph{This behavior is ... inconsistent with Griem's }theory ...''. This publication, with its subtitle: ``Confrontation with theory at high principal quantum numbers'' induced a wave of publications where BASV's finding was called an ``anomaly'', ``puzzle'' and even ``mystery''. For example, \\citet{Oks2004} paper is titled ``On the puzzle of the observed narrowing of radio recombination lines,'' \\citet{Griem2005} concludes his paper by writing that the result ``...remains a mystery...,'' \\citet{Gavrilenko&Oks2007} calls the result a ``dramatic discrepancy,'' and \\citet{Watson2006} concludes his abstract with ``Thus this mystery is not resolved by the present calculations''.  While some authors have sought an explanation of these findings in the revision of Stark broadening theory (e.g., \\citet{Oks2004,Watson2006}), some remain skeptical about the modified frequency switching technique or suggest that it requires verification. For example, \\citet{vonProchetal2010} writes ``...it is easy to distort the RRL line shape using data reduction techniques (Bell et al. 2000).'' \\citet{Griem2005} referred to private communication with Bell stressing that ``...measurement errors related to large reduction factors from unprocessed to processed Voigt profile widths can probably not be excluded (M. B. Bell 2004, private communication)''. In their book, \\citet{Gordon&Sorochenko2009} concluded: ``The Bell et al. (2000) results are so different from what had been expected, and the observing technique is so new, that we suggest waiting for an independent confirmation of the observations before accepting a fault in the present theory of RRL Stark broadening.''  BASV's findings are based on frequency switched or ``processed'' observational data, which were recently presented in \\citet{Belletal2011} and re-interpreted in \\citet{Bell2011}. Their findings result from recursive frequency switching, in software, of observational data that were initially recorded at the telescope using hardware frequency switching.  In this paper, we investigate the frequency switching technique and demonstrate that, if applied correctly, it has a number of advantages. Frequency switching removes many gain variations and does not require subjective estimates of the zero level of a spectrum. Though the technique helped BASV to detect RRLs with $\\Delta n$ greater than $\\Delta n=6$ (\\citet{Smirnovetal1984}), the way the method was used was not optimal and, as such, cannot be used to test the theory of spectral line broadening. We present simulations based on the \\citet{Lockman&Brown1975} model of the Orion Nebula and conventional theory of spectral line broadening. We apply observational specifications from BASV and \\citet{Belletal2011}, including frequency range, channel width, frequency switching offset, number of frequency switching overlaps, and noise temperature rms. Results of our simulation demonstrate good agreement with BASV's findings, both in line width and temperature. The computed ``processed'' widths exhibit narrowing similar to that reported in BASV. We show that BASV's spectral line ``narrowing'' is the result of the way the observational data were processed and that BASV's findings do not contradict the existing Stark broadening theory. ",
        "conclusions": "The multiple frequency switching technique reduces amplitudes of broad baseline variations and therefore can be used to systematically ``flatten'' baselines (\\citet{Bell&Feldman1991,Belletal1992,Belletal1993}). Compared to standard spline/polynomial baseline fitting, which allows for subjective decisions, the FS technique reduces baseline variations objectively. For spectral features that are much wider than the offset, $h$, FS can be likened to differentiation, which acts as a high-pass filter in the frequency domain (see, e.g., \\citet{Hamming1998,Owen1995}). In this case, six overlaps approximates the sixth derivative with respect to frequency, which makes FS an efficient high-pass filter. This not only reduces broad variations in the spectrum, but in BASV's case (six overlaps, $h<w$ ), it also reduces broad spectral line wings larger than the FS-offset, which manifests as line narrowing (see, e.g., Figure \\ref{fig:lorentz_2_4_6overlaps}).  We use the \\citet{Lockman&Brown1975} model of the Orion Nebula and simulate the conditions of BASV's observations at 6 GHz: telescope beam size, frequency range, channel width, offset, number of overlaps, and noise temperature. When modeling, we can avoid some difficulties that BASV faced when they processed the real spectrum. The real spectrum is abundant with spectral lines; after overlapping six times, the resulting spectrum is crowded with lines and reference lines. This situation is further complicated when significant noise fluctuations are present, which is BASV's case when $\\Delta n>11$, so cleaning becomes a non-trivial and subjective procedure. In our case cleaning is trivial as we can model and clean each RRL individually, with and without noise.  For all $\\textrm{H}n\\Delta n$ lines studied in BASV, the unprocessed width (i.e., before FS) is greater than the FS-offset used. We show that the processed width weakly depends on $n$ and $\\Delta n$ (see Figure \\ref{fig:width_processed_versus_n_and_overlaps_orion_6ghz}, the six FS-overlaps case). In fact, it remains about $0.8\\, h$ for a wide range of $n$ and $\\Delta n$. This property of the FS technique makes it insensitive to changes in RRL widths when $h<w$. (It is for this reason that we justify the inclusion of BASV's processed widths for W51 at 6 GHz in Figure \\ref{fig:width_processed_versus_n_and_offset_orion_6ghz}.) The FS technique can be used to systematically eliminate broad baseline variations to detect weak lines, however, in BASV's case (six overlaps, $h<w$), this technique is not well suited to study RRL broadening. If one's purpose is to test the theory of spectral line broadening, the choice of the offset made in BASV ($h<w$) is not optimal.  It is shown in the previous sections that FS reduces the S/N (see Equation (\\ref{SNRSNR})) and spectral line amplitudes (Figure \\ref{fig:amps_widths_6_overlaps}). As a result of amplitude reduction, the processed line temperature drops rapidly to the noise temperature $T_{\\textrm{rms}}=1$ mK as $n$ and $\\Delta n$ grow (the lower solid curve in Figure \\ref{fig:T_L_processed_versus_n_orion_6ghz}). The model processed line amplitude reaches 3 mK ($3\\sigma$ level) at $n=224$ ($\\Delta n=11$) and 1 mK ($1\\sigma$ level) at $n=241$ ($\\Delta n=14$). The interval between $n=224$ and 241 (between $\\Delta n=11$ and 14) is a transition zone between observable spectral lines and noise. Figure \\ref{fig:width_processed_versus_n_and_offset_orion_6ghz} shows that the processed width behavior in the transition zone\\emph{ }imitates ``narrowing'' of the processed\\emph{ }spectral lines observed by BASV in spectra of the Orion Nebula and W51 at 6 GHz. This ``narrowing'' is the result of ``forcing'' a program to fit Gaussians to features dominated by noise fluctuations---the procedure used by BASV.  Spectra presented in \\citet{Bell1997}, BASV and \\citet{Belletal2011} result from the overlapping and cleaning steps explained in Section 3. Though BASV applied overlapping to the entire spectrum, cleaning was applied selectively to features located near the expected spectral line frequencies, even if the processed line temperatures were below $3\\,\\sigma$. This subjective and selective approach could create ``spectral lines'' from random fluctuations at the $\\lesssim3\\,\\sigma$ level, which are abundant in the spectrum.  For $\\Delta n\\geq14$ the width errors predicted by our Monte Carlo simulation are comparable to the widths themselves. Therefore, BASV's data points for $\\Delta n\\geq14$ shown in Figures \\ref{fig:T_L_processed_versus_n_orion_6ghz} and \\ref{fig:width_processed_versus_n_and_offset_orion_6ghz} are likely to be misinterpreted noise features.  In conclusion, FS is a useful technique for detecting and measuring weak spectral features, if required corrections are minimal, that is if the FS offset is greater than the line width: $h\\gtrsim w$. Based on our simulation results, we argue that BASV's $T_{\\textrm{rms}}\\approx1$ mK observations of hydrogen RRLs from the Orion Nebula at 6 GHz are limited to $\\Delta n<14$. Within this limit, we find good agreement between their results and a Monte Carlo simulation based on conventional Stark broadening theory and, therefore, we argue that BASV's findings do not necessitate a revision of RRL Stark broadening theory. We suggest that further tests of line broadening theory above this limit will require observations with sub-mK sensitivity and improved baseline stability."
    },
    "1112/1112.0885_arXiv.txt": {
        "abstract": "Recalculating the Bogoliubov coefficients from the solutions in Phys. Rev. D ${\\bf 78}$, 103517 (2008), we obtain the mean number of boson pairs in a uniform electric field in the global coordinates $dS_2$ and $AdS_2$, which have the correct zero-field and zero-curvature limits, and study the vacuum persistence at one-loop. The mean number in $AdS_2$ gives the lowest limit to the Breitenloher-Freedman bound in the uniform electric field, and the mean numbers in $dS_2$ and $AdS_2$ satisfy the reciprocal relation ${\\cal N}_{\\rm dS} (R, E) {\\cal N}_{\\rm AdS} (R, E) = 1$ under the analytical continuation of the scalar curvature $R$. ",
        "introduction": "The gauge-gravity interaction has recently attracted attention partly because of the running gauge coupling due to gravity \\cite{Robinson-Wilczek,Toms} and partly because of charged black holes and AdS/CFT. It is a challenging task even to compute all one-loop diagrams for the gauge-gravity interaction. The quantum states of charged black holes have been studied mostly near the horizon and at the spatial infinity. The complete information of quantum states in a background gauge field and/or curved spacetime may help one understanding the vacuum polarization and vacuum persistence, the Heisenberg-Euler and Schwinger effective action \\cite{Heisenberg-Euler,Schwinger} being the most well known in quantum electrodynamics (QED).   The quantum states of a charged particle coupled to an electromagnetic field may be found in a special class of spacetime, which has a certain symmetry respected by the electromagnetic field. The solutions by Kim and Page \\cite{Kim-Page08} for a charged boson in a uniform electric field in the global coordinates of two-dimensional de Sitter (dS) and anti-de Sitter (AdS) spaces are such an exactly solvable model. There the spacetime has constant curvature, and the uniform electric field respects the symmetry of the spacetime. The symmetry of the spacetime and electric field suggests that the mean number of spontaneously produced pairs should depend on the curvature and the electric field strength. To probe the vacuum polarization and vacuum persistence (twice of the imaginary part of the effective action), one may employ the in-out formalism \\cite{DeWitt,DeWitt-book} using the Bogoliubov coefficients  between the in-vacuum and the out-vacuum. Thus the Bogoliubov coefficients are important in calculating the effective action as well as the Schwinger pair production.   In this paper we revisit the Schwinger pair production and the vacuum persistence under the uniform electric field in $dS_2$ and $AdS_2$ by recalculating the Bogoliubov coefficients using the solutions in Ref. \\cite{Kim-Page08}. For that purpose, we employ the in-vacuum and the out-vacuum in the $dS$ space as defined by Mottola, in which the vacua are the no-particle state with respect to the positive frequency solution in each asymptotic region \\cite{Mottola,BMS,JMP}. The procedure is similar to the Sauter-type time-dependent electric field in the Minkowski spacetime \\cite{KLY08}. For the $AdS$ space we use the vacuum in the St\\\"{u}ckelberg-Feynman picture that resolves the Klein paradox from a tunneling barrier \\cite{Nikishov,Damour,Hansen-Ravndal}, which has an analog of the Sauter-type space-dependent electric field in the Minkowski spacetime \\cite{Kim-Page06,KLY10}. The Schwinger pair production in $dS_2$ reduces to the particle production in the global coordinates in the zero-field limit \\cite{Mottola,BMS,JMP}, and both pair productions in $dS_2$ and $AdS_2$ reduce to the Schwinger formula in the zero-curvature limit \\cite{Schwinger}. It is shown that the pair productions in Ref. \\cite{Kim-Page08} are the leading terms of the exact formula, consistent with the worldline instanton calculation and with the planar coordinates. Further, we find the lowest limit to the Breitenlohner-Freedman bound in the electric field in $AdS_2$ and, surprisingly, discover the reciprocal relation between the mean numbers of produced pairs in $dS_2$ and $AdS_2$.   The organization of this paper is as follows. In Sec. II we study the periodic boundary conditions of the charged bosonic solutions. In Sec. III we recalculate the Bogoliubov coefficients in $dS_2$, find the Schwinger pair production and discuss the zero-field and the zero-curvature limits. In Sec. IV we show that the pair production gives the lowest limit to the Breitenlohner-Freedman bound in the electric field in $AdS_2$. In Sec. V we discuss the vacuum persistence and find the reciprocal relation between the mean numbers of produced pairs in $dS_2$ and $AdS_2$. ",
        "conclusions": "We have recalculated the Bogoliubov coefficients for the solutions in Ref. \\cite{Kim-Page08} for a charged boson in a uniform electric field in $dS_2$ and $AdS_2$. The coefficients are important not only in calculating the Schwinger pair production but also finding the effective action and the vacuum persistence in the in-out formalism \\cite{DeWitt-book}. In QED in the Minkowski spacetime the effective actions and vacuum persistence have been found in a time-dependent electric field \\cite{KLY08} and in a space-dependent electric field of Sauter-type \\cite{KLY10}. The vacuum persistence in this paper is an important tool to check the consistency of QED effective actions both in the uniform electric field and in $dS$ and $AdS$, which will be addressed in the future."
    },
    "1112/1112.5975_arXiv.txt": {
        "abstract": "We study the influence of cluster environment on the chemical evolution of spiral galaxies in the Pegasus I cluster.   We determine the gas-phase heavy element abundances  of six galaxies in Pegasus derived from H~II region spectra  obtained from  integral-field spectroscopy.  These abundances are analyzed in the context of Virgo, whose spirals are known to show increasing interstellar metallicity as a function of H I deficiency.  The galaxies in the Pegasus cluster, despite its lower density and velocity dispersion, also display gas loss due to ISM-ICM interaction, albeit to a lesser degree.  Based on the abundances of 3 H I deficient spirals and 2 H I normal spirals, we observe a heavy element abundance offset of $+0.13 \\pm 0.07$ dex for the H I deficient galaxies.  This abundance differential is consistent with the differential observed in Virgo for galaxies with a similar H~I deficiency, and we observe a correlation between log(O/H) and the H I deficiency parameter DEF for the two clusters analyzed together.  Our results suggest that similar environmental mechanisms are driving the heavy element enhancement in both clusters. ",
        "introduction": "The effect of environment on galaxy evolution has long been a subject of intense research and debate \\citep[see review articles by][for a comprehensive discussion]{haynes84,boselli06}.  Known environmental effects include tidal encounters and mergers, altered morphologies, and stripping of gas from disks.  A natural question is whether the cluster environment has significant effects of the chemical evolution of galaxies.  This topic has received increasing attention in the last few years.  \\citet{skillman96} explore the effect of environment on chemical evolution for spirals in the Virgo cluster.  Examining H II region spectra from 9 Virgo spirals, they find the three most H I deficient objects to have O/H abundances 0.3-0.5 dex higher than their gas-normal counterparts.  They suggest that the abundance differential results in part from a lack of infall of metal-poor gas into the spirals in the cluster core.  \\citet{dors06} fit photoionization models to the Virgo data, confirming the abundance excess for O/H and N/O.  Other studies, involving large-scale spectroscopic surveys \\citep{cooper08,ellison09,zhang09}, see a qualitatively similar galactic metallicity dependence on local galaxy density or gas fraction.  \\citet{zhang09} analyze a sample of 800 galaxies in the HyperLeda catalog, concluding gas-poor galaxies display higher heavy-element content for a given stellar mass.  \\citet{petropoulo11} examine abundances of dwarf and spiral galaxies in the Hercules cluster, finding higher abundances in dwarf galaxies located in relatively dense environments.  While these results suggest a significant impact of cluster environment on chemical evolution, the number of detailed studies for individual clusters is small. This paper presents, for the Pegasus I cluster, a study analogous to the \\citet{skillman96} analysis of the Virgo cluster.  Pegasus I (hereafter referred to as Pegasus) is a low-density, low velocity dispersion cluster of redshift $\\sim$ 3900 km/s in the foreground of the Pisces-Perseus supercluster.  It displays weak X-ray emission, primarily concentrated around the two central ellipticals, without additional substructure \\citep[][Figure 1]{canizares86}.  Full details of the cluster's location and member galaxies can be found in \\citet{levy07}, including a comparison to Virgo in their Table 4.  A comprehensive comparison of the properties of the Pegasus galaxies to other nearby clusters is presented in \\citet{solanes01}.  Because the density of Pegasus is so low, the classical ram pressure stripping effect should not cause significant gas loss in the disks of member spirals.  Furthermore, \\citet{vigroux89} conclude that the cluster is in the early stages of gravitational collapse, which suggests that any environment-driven evolution should be in an early phase as well.  Nevertheless, \\citet{levy07} demonstrate that the Pegasus spirals are in fact experiencing H I loss as they fall into the cluster, suggesting mechanisms other than the classic ram pressure effect \\citep{gunn72} may be at work.  In a follow-up study, \\citet{rose10} find that star formation in Pegasus galaxies is suppressed with higher H I deficiency.   Given that the cluster environment has already caused noticeable changes in the gas content and star formation of these galaxies, one might expect changes in nebular abundance as well.  Using the VIRUS-P integral-field spectrograph \\citep{hill08} on the 2.7m Harlan J Smith Telescope at McDonald Observatory, we have obtained spectra for six Pegasus spirals.  From the H II region spectra, we calculate radial O/H profiles, and examine the extent to which a metallicity offset can be seen between the gas-poor and gas-normal spirals.  We compare our results to those of \\citet{skillman96}, noting the difference in environments of Virgo and Pegasus.  We discuss the possible causes of the offset for Pegasus, and the potential implications for environment-driven chemical evolution in Virgo and other clusters. ",
        "conclusions": "The physical cause of the observed metallicity offset for cluster galaxies has been the subject of some debate.  The typical model for gas-phase chemical evolution in an isolated galaxy is a variant of the ``simple model'' \\citep[e.g.][]{pagel75}.  The simple model treats a galaxy as a ``closed box'' of hydrogen gas, which grows progressively more metal-rich through stellar recycling according to $Z = -y$ ln$\\mu$, where $Z$ is the total metallicity, $\\mu$ is the gas fraction, and $y$ is the yield of metals from massive stars.  This model is easily modified to account for the infall of primordial gas.  As mentioned before,  the inflow of unprocessed gas should be terminated when a galaxy falls into the hot ICM, resulting in a metal enrichment pattern that more closely parallels the pure simple model \\citep[as in][for example]{henry94}.  \\citet{skillman96} explore the possibility that infall cutoff explains the enhanced metal abundance of the core Virgo spirals.  Their chemical evolution models with and without infall demonstrate that infall suppression can create a metallicity increase of ~0.15 dex--a significant fraction of the Virgo offset, but not sufficient to account for the entire effect.  In the case of Pegasus, though, 0.15 dex is adequate to produce the observed enrichment.  One possibility, then, is that as a galaxy first falls into a cluster environment, infall cutoff is the primary external driver of chemical evolution, producing a metallicity transition akin to the H I normal-intermediate shift seen for Virgo.  As galaxies continue to fall through the ICM and experience increasingly more H I stripping, additional processes associated with the more extreme H I deficiencies observed become important, resulting in the elevated abundances seen for the ``deficient'' spirals.  A second mechanism proposed to drive metal enhancement requires the elimination of radial inflows of metal-poor gas.  As described in \\citet{shields91}, if the galactic H I disk is truncated at a radius interior to the stellar disk, then the inward transport of low-metallicity gas from the exterior of the galaxy is inhibited, preventing the dilution of metals in the galactic interior.  Without this dilution, the abundance gradients typically seen for field galaxies tend to flatten out, leading to high characteristic metallicities.  The H $\\alpha$ emission of Pegasus spirals is clearly truncated inside the stellar disks for H I deficient members \\citep{rose10}.  Furthermore, the flattened metallicity gradients predicted for inflow cutoff by \\citet{shields91} are precisely what we observe for our targets.  If the radial inflow effect is comparably important as infall cutoff for Virgo, the same may be true for the Pegasus cluster.  Our results add to a growing number of studies demonstrating evidence for environmental influences on galactic chemical evolution.  \\citet{tremonti04} use strong-line calibrations for integrated spectra \\citep[which are reflective of a galaxy's mean metallicity, see][]{moustakas06} of SDSS galaxies to determine 12 + log(O/H).  \\citet{cooper08} analyze these abundances, finding a strong dependence on local galaxy density (as defined by nearest-neighbor analysis), particularly for objects in group/cluster environments.  Similarly, \\citet{ellison09} use nebular abundances from SDSS galaxies in \\citet{kewley02} to demonstrate a metallicity-galaxy density relationship, although they claim the effect depends only on local density, not cluster membership.  However, using the local galaxy density alone has its limits, as very close galaxy pairs induce inflow of metal-poor gas through tidal interactions, leading to \\emph{lower} characteristic abundances than isolated field galaxies \\citep{kewley06,ellison08}.  These results bear particular significance for NGC 7643, which is within $\\sim 200$ kpc of UGC 12562 \\citep{levy07}.  While this separation is too large by an order of magnitude for tidal interactions to lower the metallicity of NGC 7643 \\citep{kewley06}, it is possible that ``galaxy harassment'' has contributed to that galaxy's high DEF, and therefore elevated (O/H) abundance.  The metallicity dependence on gas content is confirmed by \\citet{zhang09}, who cross-correlate the galaxies in \\citet{tremonti04} with H I data in the HyperLeda catalog, and find that gas-poor galaxies tend to be metal rich.  While our study rejects a dependence on H I content for the metallicities of field galaxies, we note that the \\citet{zhang09} galaxy selection does not account for cluster membership, so their observed correlation may be driven by galaxies in clusters.  The importance of environment for galactic metallicity has not been unanimously accepted.  \\citet{mouhcine07} examine the \\citet{tremonti04} data and find a metallicity correlation with local density that is much weaker than those claimed in \\citet{cooper08} and \\citet{ellison09}.  \\citet{mouhcine07} concede, however, that local environment can drive chemical evolution at least in the cores of dense clusters.  In a study more similar to our own, \\citet{petropoulo11} compute oxygen content from spatially-resolved long-slit spectroscopy of star-forming galaxies in the Hercules cluster.  While they do observe a local density dependence for the metallicity of dwarf galaxies, they claim no environmental effect for major spiral galaxies.  We note two caveats here, though.  First, they do not quantify hydrogen deficiency, making comparisons to Pegasus or Virgo difficult.  Additionaly, they have much lower spatial resolution than we have for Pegasus, and they treat the cores and disks of giant spirals as separate objects, making them insensitive to the kind of metallicity offset we extract from the galactic abundance gradients.  We therefore do not consider the \\citet{petropoulo11} study to be contradictory to ours.  Overall, the evidence indicates that dense environments like those found in galaxy clusters have a significant influence on the abundance patterns of their member galaxies.  As with Virgo, the heavy element abundance differential for Pegasus is considerably larger than the aforementioned statistical samples of SDSS galaxies.  The combined set of data for these two clusters appears to explain why.  It seems that, while considering cluster/group membership alone will reveal a small metallicity offset, the primary importance of the cluster is to drive H I stripping and infall cutoff, which ultimately leads to proportionally higher heavy element content.  To date, \\citet{skillman96} and this study are the only analyses to scrutinize the complementary processes of H I removal and metallicity enhancement as codependent effects of cluster-driven galaxy evolution.  The additions of NGC 7518 and IC 5309 to the metallicity-DEF relationship (Figure \\ref{avgvdef_comp}) are particularly interesting, as they fill a previously unsampled intermediate range in DEF, showing the correlation to be continuous rather than bimodal.  We emphasize that future metallicity surveys of additional clusters should pay careful attention to H I content, as we believe it to be essential to understanding the full influence of a cluster on the properties of its member galaxies."
    },
    "1112/1112.5691_arXiv.txt": {
        "abstract": "We study the effect of local stellar radiation and ultraviolet background radiation (UVB) on the physical properties of damped Lyman-$\\alpha$ systems (DLAs) and Lyman limit systems (LLSs) at $z=3$ using cosmological SPH simulations. We post-process our simulations with the Authentic Ray Tracing (ART) code for radiative transfer of local stellar radiation and UVB. We find that the DLA and LLS cross sections are significantly reduced by the UVB, whereas the local stellar radiation does not affect them very much except in the low-mass halos. This is because the clumpy high-density clouds near young star clusters effectively absorb most of the ionizing photons from young stars. We also find that the UVB model with a simple density threshold for self-shielding effect can reproduce the observed column density distribution function of DLAs and LLSs very well, and we validate this model by direct radiative transfer calculations of stellar radiation and UVB with high angular resolution. We show that, with a self-shielding treatment, the DLAs have an extended distribution around star-forming regions typically on $\\sim 10-30$ kpc scales, and LLSs are surrounding DLAs on $\\sim 30-60$ kpc scales. The DLA gas is less extended than the virial radius of the halo, and LLSs are distributed over the similar scale to the virial radius of the host halo. Our simulations suggest that the median properties of DLA host haloes are: $\\Mh = 2.4 \\times 10^{10} ~\\Msun$, SFR$ = 0.3 ~ \\Msunyr$, $M_{\\star} = 2.4 \\times 10^{8}~ \\Msun$, and $Z / Z_{\\rm \\odot} =  0.1$. About 30 per cent of DLAs are hosted by haloes having $SFR = 1 - 20~\\Msunyr$, which is the typical SFR range for LBGs. More than half of DLAs are hosted by the LBGs that are fainter than the current observational limit. Our results suggest that fractional contribution to LLSs from lower mass haloes is greater than for DLAs. Therefore the median values of LLS host haloes are somewhat lower with $\\Mh = 9.6 \\times 10^{9} ~\\Msun$, SFR$ = 0.06 ~ \\Msunyr$, $M_{\\star} = 6.5 \\times 10^{7}~ \\Msun$ and $Z / Z_{\\rm \\odot} =  0.08$. About 80 per cent of total LLS cross section are hosted by haloes with ${\\rm SFR} \\lesssim 1~\\Msunyr$, hence most LLSs are associated with low-mass halos with faint LBGs below the current detection limit. ",
        "introduction": "\\label{sec:intro} Interplay between neutral hydrogen gas and star formation in galaxies is a crucial issue for understanding galaxy formation and evolution. Absorption systems in the spectrum of QSOs are good tracers of neutral hydrogen, and large surveys such as the Sloan Digital Sky Survey (SDSS) provide abundant data on \\HI\\  gas at high redshift \\citep{Prochaska05-1, Prochaska09}. The absorption systems with $\\NHI \\geq 2\\times 10^{20}\\,{\\rm cm}^{-2}$ show the damping wing in the absorption profile, and are called Damped Ly$\\alpha$ systems (DLAs) \\citep{Wolfe86}. In addition, recently the DLAs has been studied even by GRB aftergrow \\citep[e.g.,][]{Prochaska07b, Ledoux09, Rau10}. It is known that DLAs dominate the neutral gas content of Universe between $z \\sim 0 - 5$ \\citep{Wolfe86, Storrie96, Storrie00}. The global neutral gas content measured from DLAs declines with cosmic epoch in concert with the growth in cosmic stellar mass \\citep{Storrie00}. Therefore, it is generally thought that DLAs are strongly correlated with cosmic star formation. However, in spite of numerous observational sample of DLAs, a deep understanding of DLAs and host galaxies remains elusive.  The nature of DLAs has been revealed gradually by recent progress in observations \\citep[e.g.,][]{Wolfe95, Storrie00, Rao00, Peroux03, Chen03, Prochaska05, Prochaska09, Prochaska10a, Wolfe08, Noterdaeme09} and simulations \\citep[e.g.,][]{Katz96, Gardner97a, Gardner97b, Haehnelt98, Maller01, Gardner01, Razoumov06a, Nagamine04f-1, Nagamine04g-1, Nagamine07, Pontzen08,  Barnes09, Tescari09, Cen10}. Thanks to these works, our understanding on the relationship between DLAs and host galaxies have improved. For example, it has been suggested that a significant fraction of DLAs are distributed in the same dark matter halos as the Lyman break galaxies \\citep[LBGs; e.g.,][]{Adelberger05, Cooke06, Rafelski11}. Theoretically, \\citet{Cen10} examined the number fraction of DLAs related to LBGs using cosmological AMR simulations. \\citet{Lee11} estimated the DLA--LBG cross-correlation function using cosmological SPH simulations, and showed that their result agrees well with the observational estimate of \\citet{Cooke06}.  A connection between DLAs and Ly$\\alpha$ emitters (LAEs) has also been suggested. \\citet{Rahmani10} investigated Ly$\\alpha$ emission from 341 DLAs at $z\\sim 2.86$, and concluded that the overall DLA population seems to originate from the low luminosity end of LAEs at high redshift. \\citet{Barnes11} simulated three dimensional Ly$\\alpha$ radiative transfer (RT) for high-redshift galaxies in cosmological SPH simulations. They showed that DLA host galaxies show extended Ly$\\alpha$ emission around high-column density region. The Ly$\\alpha$ emission is extended over several arc seconds, and has the spectral width of several hundred km \\,s$^{-1}$. A combination of these observational and theoretical works seem to suggest that DLAs are associated with a wide variety of sources, such as LBGs and LAEs.  For a better understanding of DLAs, it is quite important to investigate various feedback effects in galaxies which can drastically change the state of gas. Multiple supernovae can cause galactic wind, and hence change gas distribution largely in galaxies. \\citet{Nagamine04f} investigated the effect of galactic wind on the \\HI\\ column density distribution function $\\fn$, and found that strong galactic winds with a speed of $\\sim$500\\,km\\,s$^{-1}$ is needed to reproduce the observed $\\fn$ at high $\\NHI$ values. \\citet{Tescari09} also found that strong galactic winds with speed of $\\sim$600\\,km\\,s$^{-1}$ are needed in order to reproduce the observed $\\fn$ and the evolution of \\HI\\ mass density with redshift. Several authors have suggested that galactic outflows can produce the large velocity dispersion as observed in DLAs \\citep[e.g.,][]{Razoumov09, Hong10, Cen10}.  Gas inflows from surrounding IGM can also have strong impact on the absorption systems in high-$z$ galaxies. Recent simulations have revealed a new theoretical paradigm for galaxy evolution that a large amount of gas penetrate deep inside dark matter haloes as cold, filamentary streams \\citep{Katz03, Keres05, Keres09, Birnboim03, Dekel06, Ocvirk08, Brooks09, Dekel09}. \\citet{Faucher11} studied the neutral hydrogen content in the cold accretion to LBGs using SPH simulations with UVB RT. They showed that a part of gas could be DLAs, however the covering fraction was quite small. Similarly, \\citet{Fumagalli11} showed the neutral hydrogen gas content in cold stream by cosmological AMR simulations and RT calculation of local stellar source and UVB. They find that most of the gas in cold streams is ionized, however it can survive as Lyman-limit systems (LLSs), which is defined by column densities of $10^{17} < \\NHI < 10^{20.3}$ atoms cm$^{-2}$.  The observations of DLA metallicity can also give important constraints on feedback processes, as the interstellar gas can be enriched by supernovae in the early Universe. \\citet{Moller04} presented the metallicity distribution of DLAs and suggested that the DLAs with Ly$\\alpha$ emission tend to have higher metallicities of $Z \\gtrsim 10^{-1.5}~\\Zsun$. \\citet{Ledoux06} presented the velocity width $\\Delta v$--metallicity relation of DLAs from the line width of  low- ionization metals, and showed that the metallicity is linearly proportional to $\\Delta v$ with a large dispersion. They also suggested that the slope of fitting function for $\\Delta v$--metallicity relation is the same for the sample at  $z > 2.43$ and $z \\le 2.43$, but the median metallicity decreases with increasing redshift. Moreover, \\citet{Prochaska07a} presented the metallicities of a larger DLA sample with a large dispersion similar to that of \\citet{Ledoux06}. Theoretically, recent numerical simulations which include metal enrichment by supernovae have successfully produced the observed metal distribution of DLAs \\citep{Pontzen08, Tescari09, Cen10}.  Another important feedback on DLAs is radiation from UVB and local stellar sources. The UV radiation can ionize hydrogen gas, and largely reduce the neutral fraction of ISM. Earlier \\citet{Katz96} and \\citet{Haehnelt98} included the self-shielding correction for UVB radiation to reproduce DLA statistics in there numerical simulations.  Thereafter, recent simulations have reproduced various observation data of DLAs reasonably well by considering the effect of UV RT \\citep{Razoumov06a, Razoumov08, Kohler07, Pontzen08, Nagamine10-1, Altay11, Faucher11, Fumagalli11, McQuinn11}. Most of them have focused on the RT of UVB, and found that the self-shielding effect is crucial for matching the observed $\\fn$ of DLAs. \\citet{Nagamine10} showed that the simple self-shielding model without RT calculations could reproduce the observed $\\fn$ very well except for the very high $\\NHI$ values. \\citet{Altay11} also reproduced the observed $\\fn$ in the range of $10^{12} < \\NHI  < 10^{22}$\\,atoms\\;cm$^{-2}$ by post-processing cosmological SPH simulations for the UVB RT.  On the other hand, the flux of stellar radiation near the star cluster can be much higher than the local UVB, hence stellar RT can be important for DLAs and LLSs. However, it has been ignored in many simulations. Some authors have suggested that the UV radiation from local stellar sources can be important for absorption systems \\citep{Miralda05, Schaye06}, however, there have not been extensive studies of direct RT calculations that examined the effect of local stellar sources. \\citet{Nagamine10} showed that the $\\fn$ at $z=3$ does not change very much, based on a direct RT calculation from young stellar sources. \\citet{Fumagalli11} carried out the RT calculations of local stellar sources and UVB for seven high-resolution simulated galaxies, and found that the local stellar radiation changed $\\fn$ at $\\NHI \\sim 10^{18}-10^{21}$\\;cm$^{-2}$ by $\\sim$0.5 dex to the other model without local stellar radiation. Their simulation result is in good agreement with observation at $\\log \\NHI > 21$, however it underestimates $\\fn$ at lower $\\NHI$ range of LLS.  The stellar radiation flux onto a gas cloud varies depending on the distance between gas and stars and the type of SED. Therefore in a more realistic inhomogeneous medium, precise RT calculations are needed to study the effects of stellar radiation on absorption systems. Furthermore, since the distribution of stars and gas varies among different galaxies \\citep[e.g.,][]{Yajima11}, we need to carry out the RT calculation for many galaxies to discuss the statistical data such as $\\fn$. In the present work, we perform precise RT calculations with $\\sim 10^{4}$ angular bins for each star particle in cosmological SPH simulations, and examine the distribution of neutral hydrogen gas for a few hundred galaxies. Our RT code is based on the ray-tracing method, which can estimate the ionization structure more accurately than the Monte Carlo method that \\citet{Fumagalli11} used \\citep[see][for a comparison of RT methods]{Iliev06}. In addition, the number of galaxy sample processed with RT in our work is much larger ($N \\gtrsim 100$) than that of \\citet[][$N\\sim 8$]{Fumagalli11}. Thus, our calculation allows us to investigate the effect of RT more reliably on the statistical quantities of DLAs and LLSs, such as $\\fn$ and the mass dependence of cross sections. Moreover, by direct RT calculations of UVB, we verify the self-shielding model of UVB derived by \\citet{Nagamine10}. We also discuss the physical quantities of host galaxies of DLAs and LLSs, such as the halo mass, star formation rate, stellar mass and metallicity.  Our paper is organized as follows. We describe our simulations and our approach for RT calculations in Section~\\ref{sec:model}. In Section~\\ref{sec:result}, we present our results, and show the effect of stellar and UVB radiation on cross sections of DLAs and LLSs, and on $\\fn$. In Section~\\ref{sec:abundance}, the cumulative abundance of DLAs is shown. In addition, we show the typical physical quantities of DLAs and LLSs host galaxies. Finally, we summarize our main conclusions in Section~\\ref{sec:summary}. We focus on redshift $z=3$ in this paper.  \\section[]{MODEL AND METHOD} \\label{sec:model}  \\subsection{Simulations}  We use an updated and modified version of the Tree-particle-mesh (TreePM) smoothed particle hydrodynamics (SPH) code GADGET-3 \\citep[originally described in][]{Springel05e}.  The SPH calculation is performed based on the entropy conservative formulation \\citep{Springel02}.  Our fiducial code includes radiative cooling by H, He, and metals \\citep{Choi09a}, star formation, supernova feedback, a phenomenological model for galactic winds and a sub-resolution model of multiphase ISM \\citep{Springel03a}.  We also include the heating by a uniform UVB, which we will discuss more in Section~\\ref{sec:UVB}.  In this multiphase ISM model, high-density ISM is pictured to be a two-phase fluid consisting of cold clouds in pressure equilibrium with a hot ambient phase.  Cold clouds grow by radiative cooling out of the hot medium, and this material forms the reservoir of baryons available for star formation. The star formation rate (SFR) is estimated for each gas particle that have densities above the threshold density, and the star particles are spawned statistically based on the SFR. For the star formation model, the ``Pressure SF model'' is being used \\citep{Schaye08, Choi09b}. This model estimates the SFR based on the local gas pressure rather than the gas density, and prevents artificial fragmentation of the gas.  The simulations used in this paper also uses the ``Multicomponent Variable Wind\" model developed by \\citet{Choi10}. This wind model is based on both the energy-driven wind and the momentum-driven wind \\citet{Murray05},  and the wind speed depends on the galaxy stellar mass and SFR.  It gives more favorable results when compared to the observations of cosmic C\\,{\\sc iv} mass density and IGM temperature than the previous model with a constant wind speed.  To enable this new wind model, \\citet{Choi10} implemented an on-the-fly group-finder into GADGET-3 to compute galaxy masses and SFRs while the simulation is running.  The group-finder, which is a simplified variant of the \\textup{SUBFIND} algorithm developed by  \\citet{Springel01}, identifies the isolated groups of star and gas particles (i.e., galaxies) based on the baryonic density field. The detailed procedure of this galaxy grouping is described in \\citet{Nagamine04e}.  The simulations used in this paper are performed with $2\\times144^{3}$ particles. The DM and gas particles masses are $1.97\\times10^{7}$ and $4.06\\times10^{6}\\;h^{-1} \\Msun$, respectively. The comoving gravitational softening length is $\\epsilon = 2.78 h^{-1}$\\,kpc in a comoving $(10 h^{-1}\\rm Mpc)^{3}$ calculation box.  The adopted cosmological parameters are consistent with the WMAP results \\citep{Komatsu11}: $H_0 = 72$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ ($h=0.72$), $\\Omega_M=0.26$, $\\Omega_{\\Lambda}=0.74$, $\\Omega_b=0.044$, $\\sigma_8 = 0.80$, and $n_s = 0.96$.   \\subsection{Radiation Transfer}  \\subsubsection{Stellar Radiation} \\label{sec:radiation}  To study the effect of stellar radiation on DLAs and LLSs, we compute the stellar radiation transfer and the ionization structure of gas in each dark matter halo by post-processing the simulation output. The RT scheme used in this paper is the Authentic Radiation Transfer (ART) method, and the treatment is basically the same as in \\citet{Yajima09, Yajima11-1, Yajima11b}. Here we briefly summarize the procedure of RT calculation.  First we set up a uniform grid around each dark matter halo with a grid cell size equal to the gravitational softening length of the simulation, and translate the SPH gas information into a gridded data. As for the scattering of photons, we employ the on-the-spot approximation \\citep{Osterbrock89}, in which the scattered photons are assumed to be absorbed immediately on the spot. In this work, the RT equation is solved along $n_{\\rm g}^{2}$ rays ($\\sim$10$^{4-5}$) with a uniform angular resolution from each source, where $n_{\\rm g}$ is the grid number on a side of the calculation box of each galaxy. We shoot the radiation rays in a radial fashion from each star particles. The number of ionizing photons emitted from the source stars is computed based on the theoretical spectral energy distribution (SED) given by the population synthesis code P\\'{E}GASE v2.0 \\citep{Fioc97}. We consider the effect of age and metallicity of the stellar population by interpolating the table generated from the result of P\\'{E}GASE. We assume the \\citet{Salpeter55} initial mass function with the mass range of $0.1 - 50 \\Msun$.  In addition, we consider the dust extinction. The dust model is similar to \\citet{Yajima11}, e.g., $m_{\\rm d} = 0.01 \\times m_{\\rm g} \\frac{Z}{Z_{\\odot}}$ with a silicate-like mass density of 3\\,g\\,cm$^{-3}$, where $m_{\\rm d}$, $m_{\\rm g}$ and $Z$ are dust mass, gas mass and metallicity. The size distribution is $n_{\\rm d}(a_{\\rm d}) \\propto a^{-3.5}_{\\rm d}$ \\citep{Mathis77} with the dust radius of $0.1 - 1.0~\\rm \\micron$. Since the assumed range of dust size is larger than the wavelength of Lyman limit, we assume the absorption efficiency $Q(\\nu)=1$ for the ionizing photons \\citep{DL84}.   \\subsubsection{UV Background Radiation} \\label{sec:UVB}  The UVB can ionize the baryonic gas in galaxies, and then gas can be heated up to $\\sim 10^4$\\,K. It would be ideal to compute the RT of both UVB and stellar radiation as the simulation runs, however in practice it is a very expensive calculation.  In general, the calculation amount of UVB transfer is larger than that of stellar radiation by $2-4$ orders of magnitude \\citep{Yajima11}. Therefore we first use the simple, four different models of UVB, similarly to \\citet{Nagamine10}. We then compute the actual UVB transfer for some galaxies to check the validity of the UVB models.  Our fiducial simulations include a uniform UVB with a modified \\citet{Haardt96} (hereafter HM96) spectrum \\citep[see][]{Dave99}, where the reionization takes place at $z\\simeq 6$ as suggested by the quasar observations \\citep[e.g.,][]{Becker01} and stellar radiative transfer calculations \\citep[e.g.,][]{Sokasian03}.  We use the following four simulations with different treatment of UVB (but with the same initial condition) to examine the effects of UVB: \\begin{enumerate} \\item {\\bf Optically-Thin Model}: A uniform UVB radiation field with an optically thin approximation. \\\\ \\item {\\bf HM0.5 Model} (HM96 $\\times\\; 0.5$): The ISM is optically thin to the same UVB, however the intensity of UVB is reduced to the half of the Optically-Thin model. \\\\ \\item {\\bf OTUV Model} (Optically Thick UV): The ISM is optically thin to the UVB in the lower density regions with $n_{\\rm H} < 0.01 \\rhoth = 6.34 \\times 10^{-3}$\\,cm$^{-3}$, but completely optically thick in higher density regions of $n_{\\rm H} \\ge 0.01 \\rhoth$, where $ \\rhoth$ is the threshold density, above which star formation is allowed. The value of $\\rhoth$ was determined by \\cite{Choi09b} based on the observed SF cut-off column density of the Kennicutt law.  The OTUV model implicitly assumes that the UVB cannot penetrate into the high density regions by self-shielding. This type of self-shielding model was first proposed by \\citet{Katz96} and \\citet{Haehnelt98} to examine $\\fn$ and DLA kinematics. Later \\citet{Nagamine10} successfully reproduced the observed $\\fn$ by incorporating the OTUV model into cosmological simulations.   See also \\citet{Pontzen08, Fumagalli11} and \\citet{Altay11}. \\\\ \\item {\\bf No-UVB Model}: UVB does not exist at all. \\end{enumerate}  Since we use the same initial condition for these runs, we can compare the effect of radiation halo by halo basis. After discussing the above four models in the beginning of Section~\\ref{sec:fn}, we then move on to discuss the effects of radiative transfer in Sections~ \\ref{sec:UVBtransfer} and \\ref{sec:stellarRT}.  We compute the transfer of local stellar radiation in each halo as post-processing, and estimate the ionization structure in the equilibrium state by evaluating \\begin{equation} (\\Gamma_{\\rm UVB}^{\\gamma}+ \\Gamma_{\\rm star}^{\\gamma}) \\; n_{\\rm H_I} + \\Gamma^{\\rm C}\\; n_{\\rm H_I} \\; n_{\\rm e} = \\alpha_{\\rm B}\\; n_{\\rm H_{II}} \\; n_{\\rm e}, \\end{equation} where $\\Gamma_{\\rm UVB}^{\\gamma}$ and $\\Gamma_{\\rm star}^{\\gamma}$ are the photoionization rate by UVB and stellar radiation respectively; $\\Gamma^{\\rm C}$ is the collisional ionization rate; $n_{\\rm e}, n_{\\rm HI}$ and $n_{\\rm HII}$ are the number density of free-electron, neutral and ionized hydrogen, respectively; $\\alpha_{\\rm B}$ is the total recombination coefficient to all bound excitation levels. The value of $\\Gamma_{\\rm star}^{\\gamma}$ is estimated by the full RT calculation, but $\\Gamma_{\\rm UVB}^{\\gamma}$ is computed with the above UVB model.     \\section[]{RESULTS} \\label{sec:result}  We first compare our simulation results to the recent observations of \\HI\\ column density function \\citep[e.g.,][]{Prochaska09, Noterdaeme09, Prochaska10b}, and examine the effect of UVB and stellar radiation on $\\fn$. We then examine the \\HI\\ cross section of DLAs and LLSs, and study how they change with the UVB models and the RT calculation.  Finally, we study the typical physical properties of DLA and LLS gas, such as SFR and metallicity.   \\subsection{Column density distribution function} \\label{sec:fn}  The column density distribution function $f(\\NHI, X(z))$ is defined such that $f(\\NHI, X) \\dd \\NHI \\dd X$ is the number of absorbers per sight line with \\HI\\ column density in the interval $[\\NHI, \\NHI + \\dd \\NHI]$, and absorption distance in the interval $[X, X + \\dd X]$. The absorption distance $X(z)$ is given by \\begin{equation} X(z) = \\int_{\\rm 0}^{z} (1 + z')^{2} \\frac{H_{\\rm 0}}{H(z')} ~\\dd z'. \\end{equation} In practice, if the comoving box-size of the simulation is $\\Delta L$, then the corresponding absorption distance per sight-line is \\begin{equation} \\Delta X = \\left( \\frac{H_{\\rm 0}}{c} \\right) (1 + z)^{2} \\Delta L. \\end{equation} For example, for $\\Delta L = 10 ~h^{-1} \\rm~Mpc$ and $z=3$, we have $\\Delta X = 0.0534$.  Since the simulation box size is very small, we can safely assume that the DLA clouds do not overlap along a sight- line, and we can compute the $\\NHI$ distribution function by counting the number of projected grid-cells with column densities in the range $[\\NHI, \\NHI + \\dd \\NHI]$ for all halos in the simulation.  We use the same grid that we set around each dark matter halo for the RT calculation, with the grid cell size equal to the comoving gravitational softening length.   \\begin{figure} \\begin{center} \\includegraphics[scale=0.45]{fig1.eps} \\caption{ Column density distribution function at $z=3$. The purple dash-long dashed line indicates the $\\fn$ considering only star-forming galaxies that have young stellar sources ($t_{\\rm age} \\le 10^{7}~\\rm yr$). The orange dot-dash line indicates the $\\fn$ considering only massive galaxies of $M_h \\ge 5 \\times 10^{10}~\\Msun$. The observation data points are from \\citet[][blue filled circles,]{Peroux01}, \\citet[][black filled squares]{Peroux05}, \\citet[][magenta open circles]{OMeara07}, \\citet[][green triangles]{Prochaska09}, \\citet[][cyan bars]{Noterdaeme09} and \\citet[][yellow shade]{Prochaska10b}. The figure is similar to Figure 3 in \\citet{Nagamine10}, but the $\\NHI$ range is extended to a lower value of $\\NHI = 10^{17}$\\,cm$^{-2}$ to discuss LLS. } \\label{fig:cdfall} \\end{center} \\end{figure}   We calculate the \\HI\\ density in  each grid-cell of volume $\\left(\\epsilon/[1+z]\\right)^{\\rm 3}$ after the RT calculation, and estimate $\\NHI$ by projecting the \\HI\\ mass distribution along the $z$-axis direction as \\begin{equation} N_{\\rm HI} = \\sum_{i} \\frac{\\epsilon\\, \\rho_{{\\rm HI}, i}} {m_{\\rm p} (1+z)} \\end{equation} where $\\epsilon$ is the comoving gravitational softening length, $m_{\\rm p}$ is the proton mass, and $z$ is the redshift.  The resulting $\\fn$ is insensitive to the direction of projection. By this method, we are treating each projected grid-cell element as one line-of-sight. We have also checked that we are not missing any LLS columns in the outskirts of the halo by doubling the size of the grid for some of the halos.   Figure~\\ref{fig:cdfall} shows $\\fn$ in the range of $17 < \\log \\NHI < 22$ at $z=3$ together with some observational data. The no-UVB run overpredicts the observed data points at all $\\NHI$ range. In particular, the difference becomes as large as an order of magnitude at $\\NHI \\sim 10^{19}$\\,cm$^{-2}$. This shows that the ionization effect by the UVB is crucial in reproducing the observed LLS and DLA column density distribution functions. On the other hand, the $\\fn$ of the Optically-Thin and the HM0.5 runs are close to the observed data at $\\NHI \\lesssim 10^{19.5}~\\rm cm^{-2}$, however, falls below the observation data at $10^{19.5} ~{\\rm cm^{-2}} \\lesssim \\NHI \\lesssim 10^{21} ~{\\rm cm^{-2}}$ by a significant amount.  The OTUV run agrees with the observed $\\fn$ much better than the other models in a wide range of $10^{17} ~{\\rm cm^{-2}} \\lesssim \\NHI \\lesssim 10^{\\rm 21}~{\\rm cm^{-2}}$. At lower column densities of $\\NHI \\lesssim 10^{18}$\\,cm$^{-2}$, the OTUV run is lower than the Opticall-Thin run, but it still fits with the observed range indicated by the yellow box. Thus, as far as the column density range of LLS and DLA is concerned, we conclude that the UVB treatment of the OTUV run successfully reproduces the observational data in a wide range of $10^{17} ~{\\rm cm^{-2}} \\lesssim \\NHI \\lesssim 10^{\\rm 21}~{\\rm cm^{-2}}$. We note that the result of the OTUV run in Figure~\\ref{fig:cdfall} is slightly different from that in \\citet{Nagamine10}, but this is because the galactic wind model was slightly different between the two works. We have also checked the results with the updated UVB model of \\citet{Faucher09}, and found that the results are not very different.  One notable feature in Figure~\\ref{fig:cdfall} is that the OTUV run is higher than the observed $\\fn$ at $\\NHI \\gtrsim 10^{21.5}$. These high column density sight-lines mainly come from haloes of higher dust-to-gas ratio. In such halos, the molecular hydrogen H$_{2}$ is created efficiently, and the H$_2$ mass fraction can be $\\gtrsim 0.1$ \\citep[e.g.,][]{Hirashita05}. Because our current simulations do not include H$_{2}$ formation, we may overpredict $\\fn$ at high column densities of $\\NHI \\gtrsim 10^{21.5}$\\,cm$^{-2}$.  \\begin{figure} \\begin{center} \\includegraphics[scale=0.43]{fig2.eps} \\caption{ Photoionization rate by UVB radiation ($\\Gamma_{\\rm UVB}$) as a function of hydrogen number density for four different halos in the OTUV run (Halo A : $\\Mh = 9.7\\times 10^{10} \\Msun$, B : $\\Mh = 2.5 \\times 10^{10}\\Msun$, C : $\\Mh = 1.5 \\times 10^{10} \\Msun$, and D : $\\Mh = 9.4 \\times 10^{9}\\Msun$). The halos are taken from the OTUV run, but we are solving the transfer of UVB radiation directly in each cell as shown by the red points.  The black dashed lines and the blue solid lines indicate the ionization rate in the Optically-Thin model and the OTUV model, respectively. } \\label{fig:gammauvb} \\end{center} \\end{figure}    We find that the stellar radiation transfer does not affect the shape of $\\fn$ (overlapping blue dashed line and red solid line in Figure~\\ref{fig:cdfall}), whereas the UVB model changes the $\\fn$ significantly. We will discuss the effect of stellar radiation on DLAs and LLSs in more detail later in the paper.  The change of column density by UVB in the OTUV run is similar to those found by previous works \\citep[e.g., Fig.~A1 in][]{Fumagalli11}.  On the other hand, the effect by stellar radiation in our simulation is much smaller than that of \\citet{Fumagalli11}. There could be several possible reasons for this difference as follows. In our simulations, the high-density gas clumps near young stars can effectively absorb stellar radiation,  and block the propagation of stellar radiation over extended regions in a galaxy \\citep{Yajima11}. Another possible reason could be different numerical methods: our RT method is based on a ray-tracing method, whereas Fumagalli's method is Monte Carlo.  The ray-tracing method can estimate the ionisation structure by point sources more accurately, while the Monte Carlo method can cause some artificial inhomogeneity \\citep[see the comparison work by][]{Iliev06}. Finally, the spatial resolution is always a concern, and we note that the resolution of Fumagalli's AMR simulation in high-density regions is better than ours.  Unfortunately it is very difficult to tell the exact resolution effects unless we compare the two codes closely side by side.  \\begin{figure*} \\begin{center} \\includegraphics[scale=0.4]{fig3a.eps} \\vspace{0.2cm} \\includegraphics[scale=0.4]{fig3b.eps} \\caption{ {\\it Panel (a)} : Neutral fraction of hydrogen gas as a function of number density for the same halos as in Figure~\\ref{fig:gammauvb}. The black dashed lines and blue solid lines indicate the equilibrium solution of ionization balance between the UVB ionization and recombination in the Optically-Thin model and the OTUV model, respectively.  The red filled circles are obtained by including direct radiative transfer calculations of UVB, stellar radiation and collisional ionization. The blue filled circles of Halo C are the neutral fraction with UVB and collisional ionization only (i.e. without stellar radiation transfer).  The green filled circles are the neutral fraction with only UVB (i.e. without stellar radiation and collisional ionization). {\\it Panel (b)} : Neutral fraction of gas clumps irradiated by the radiation from young star clusters at the indicated distance $\\ell$ as a function of hydrogen number density. See text for the details of the calculation.  Different colored lines indicate different conditions as indicated in the legend, where $\\dot{N}_{\\rm ph}$ [s$^{-1}$] is the emissivity of ionizing photons from star clusters, and $\\ell$ [kpc] is the distance between the star clusters and the gas clump. } \\label{fig:uvbcom} \\end{center} \\end{figure*}   Another major source of different results might be the different mass range of simulated galaxies in each simulation. \\citet{Fumagalli11} focused on only seven star-forming galaxies with $\\Mh > 5\\times 10^{10} \\Msun$. On the other hand, our simulations cover a wider range of galaxy masses, and there are numerous low-mass, faint galaxies without young star clusters that contribute to the DLA cross section.  In those low-mass haloes, the local stellar radiation does not impact the ionization of gas, and Fumagalli et al.'s simulation did not include these low-mass halos. In our simulations, about half of the total DLA cross section is contributed by faint low-mass haloes where there is no young stars in our simulations. The DLAs in these faint haloes are unaffected by the stellar radiation. This could explain the little effect of local stellar radiation on $\\fn$ in our simulations, and more effect in Fumagalli's. In addition, Fumagalli et al.'s $\\fn$ showed good agreement with observations at $\\NHI \\gtrsim 10^{20.3}$\\,cm$^{-2}$, but underpredicted at $10^{18} \\lesssim \\NHI \\lesssim 10^{20}$\\, cm$^{-2}$.  They suggested that the reason for the underestimate might be their limited galaxy sample with $\\Mh \\gtrsim 5\\times10^{10} \\Msun$. For comparison, here we investigate the effect of galaxy mass limit in our simulation in Figure~\\ref{fig:cdfall}. % When we consider only the massive galaxies in haloes with $\\Mh > 5\\times 10^{10} \\Msun$, then $\\fn$ decreases for DLAs with a similar slope to the all-halo case, and $\\fn$ is significantly lower and more flattened for the LLS. Integrating $\\fn$ over the column density range of LLS, we find that the contribution from the massive galaxies to the line density of LLSs is only $\\sim$20\\% of the total.  Hence, we confirm that the LLSs mainly come from lower-mass galaxies, and the discrepancy of $\\fn$ with observation at LLS range in \\citet{Fumagalli11} is probably due to the mass limit of their galaxy sample.  In addition, we compute the $\\fn$ using only the star-forming galaxies in our simulations, and find that the result is in-between the all-halo case and only massive galaxy case. This result suggest that LLSs mainly originate from lower mass galaxies in haloes with $\\Mh < 5\\times 10^{10} \\Msun$ and lower star formation rate. We will discuss this point more quantitatively in Section~\\ref{sec:typical}.   \\begin{figure*} \\begin{center} \\includegraphics[width=0.95\\columnwidth,angle=0] {fig4a.eps} \\includegraphics[width=0.95\\columnwidth,angle=0] {fig4b.eps} \\caption{ Physical cross section of DLAs ({\\it panel [a]}) and LLSs ({\\it panel [b]}) as a function of halo mass at $z=3$. Each point represents a halo with star-forming galaxies, including the stellar RT.  Filled triangles indicate the median values in each mass bin. In the bottom right panel, the results from different runs are compared: Optically-Thin (red), OTUV (blue), HM0.5 (green) and No-UVB (magenta) runs, respectively. } \\label{fig:sigmauv} \\end{center} \\end{figure*}  \\subsubsection{Radiation field by UVB transfer} \\label{sec:UVBtransfer}  Next we perform direct radiation transfer calculation of UVB to validate the results of the OTUV run. We perform this calculations by irradiating UVB from the boundary surfaces of calculation boxes with $n_{\\rm g}^{4}\\; (\\approx 10^{8-10})$ rays, and also consider the stellar radiation transfer simultaneously.  Figure~\\ref{fig:gammauvb} shows the photoionization rate versus gas density in each grid cell from direct UVB transfer calculation.  Here we pick four halos randomly from the halo sample of $\\Mh < 10^{11} \\; \\Msun$ in the OTUV run, and recompute the transfer of UVB directly in each cell as shown by the red points. The higher mass haloes are excluded from our sample, because the UVB RT calculation takes too long to process such a large grid ($>$100$^3$) with current computers. The red points in the figure are derived from the RT calculations including both the UVB and stellar radiation. The black dashed and blue solid lines are the photoionization rates by the Optically-Thin model and the OTUV model,  respectively, obtained from the equilibrium calculation. Note that the photoionization rate in this figure is of UVB radiation, $\\Gamma_{\\rm UVB}^{\\gamma}$ in equation (1), not including stellar radiation $\\Gamma_{\\rm star}^{\\gamma}$.  We find that the $\\Gamma_{\\rm UVB}$ has a relatively sharp transition from a high plateau at low densities to a much lower value at $\\log \\nh \\approx -2.2$, and the self-shielding threshold density ($\\nuv=6\\times 10^{-3}$\\,cm$^{-3}$, indicated by the blue solid line) adopted in the OTUV model correctly captures this sharp transition.  The value of $\\nuv$ was determined by \\citet{Nagamine10} by examining the agreement with the observed data on $\\fn$, and one of the main purposes of this paper is to validate this choice of $\\nuv$ through direct RT calculations as demonstrated here. \\citet{Nagamine10} also gave some justification for this particular choice of $\\nuv$, referring to prior radiative transfer simulations. In addition, the adopted threshold density is close to that of \\citet{Haehnelt98}, who suggested the density threshold model based on the tight relation between column density and density. The region with $n_{\\rm H} \\ge 10^{-2}~\\rm cm^{-3}$ mainly contributes to the column density of $N_{\\rm HI} > 10^{17}~\\rm cm^{-2}$, which can self-shield from UVB.  Some limited number of cells at $n_{\\rm HI} > 10^{-2}~\\rm cm^{-3}$ for Halos A \\& C show relatively high values of $\\Gamma^{\\gamma}$, deviating from the main photoionized branch of gas.  These are cells that are probably ionized by the local stellar radiation, and therefore allows UVB to penetrate into deeper densities. In some regions near star-clusters, the stellar radiation can increase the photo-ionization rate than UVB alone by orders of magnitudes, and ionize star-forming regions. However, the impact is not so significant for a statistical quantity such as $\\fn$ (Figure~\\ref{fig:cdfall}).  Figure~\\ref{fig:uvbcom}a shows the neutral fraction of hydrogen gas ($X_{\\rm HI}$) in each cell as a function of number density for the same four halos shown in Figure~\\ref{fig:gammauvb}. In Halos A and C, the stellar radiation boosts the ionization degree than the optically-thin equilibrium values, populating the red points below the black dashed curve. In Halo B and D, the young star clusters are distributed in high-density peaks, hence most ionizing photons are absorbed on the spot.  In addition, the neutral fraction $X_{\\rm HI}$ in low density region is approximately fitted by $X_{\\rm HI} \\sim \\alpha_{\\rm B} \\nh / \\Gamma^{\\gamma}_{\\rm UVB}$, therefore the neutral fraction is simply proportional to the number density of hydrogen. In addtion, we calculate the neutral fraction with different ionization processes (green : UVB only, magenta : UVB + collisional ionization, and red: UVB + stellar radiation + collisional ionization). It shows that the collision boosts the ionization degree at $n_{\\rm H_I} \\sim 10^{-4}~\\rm cm^{-3}$ (magenta points extending below the dashed curve), and the stellar radiation boosts the ionization degree even at higher density regions, as signified by the wide distribution of red points below the dashed curve.  We can further understand the transition from ionized to neutral state by solving for the balance between photoionization by stars and recombination in the one-zone approximation, picturing the gas clumps near young star clusters. In Figure~\\ref{fig:uvbcom}b, we estimate the photoionization rate by $\\Gamma^{\\gamma}_{\\rm star} = \\dot{N}_{\\rm ph} \\sigma_{0} {\\rm exp}(-\\tau) / 4\\pi l^{2}$, where $\\dot{N}_{\\rm ph}$ is the emissivity of ionizing photons from star clusters, $ \\sigma_{0}$ is the ionization cross section at the Lyman limit, $l$ is the distance between star clusters and gas clumps. The optical depth $\\tau$ is roughly estimated by $\\tau = \\sigma_{0} n_{\\rm HI} l$, where $n_{\\rm HI}$ is the number density of neutral hydrogen of the gas clump. The $\\dot{N}_{\\rm ph}$ of a star cluster in our simulations can be $\\sim 10^{50-52}$\\,s$^{-1}$, and the distance $l$ can be proper $\\sim 0.5 - 1.0$\\,kpc, corresponding to the cell size of RT calculation in our simulation. In Figure~\\ref{fig:uvbcom}b,  the neutral fraction steeply rises at the density where $\\tau$ becomes $\\sim 1$. With increasing photon emissivity, the threshold density moves to a higher density as the radiation is able to sink into higher densities, and the ionization degree at lower densities become stronger.  With the above parameter values, we find that the star clusters cannot ionize gas clumps with $n_{\\rm HI} \\gg 1\\; \\rm cm^{-3}$.  In our simulations, the gas density near star clusters can be $> 1$\\,cm$^{-3}$ in massive halos and block the ionizing radiation from stars. From these analyses of UVB radiation transfer, we confirm that the stellar radiation does not strongly affect the neutral fraction of high-density gas at $\\nh > \\nuv$, and the UVB treatment of the OTUV run is reasonable.   \\subsection{Cross Sections of DLAs and LLSs} \\label{sec:sigma}  Once the \\HI\\ column density of each cell in the projected plane is obtained, we can estimate the DLA cross-section of each halo by counting the number of grid-cells that exceed $\\NHI = 2 \\times 10^{20}$\\,cm$^{-2}$ and multiplying by the unit area $\\left(\\epsilon/(1+z)\\right)^{\\rm 2}$ of the grid-cell.  In this paper, we discuss the cross section in physical units.  First, we examine the effect of different UVB models on the cross section. Figure~\\ref{fig:sigmauv}a,b show $\\sigmadla$ and $\\sigmalls$ of each halo as a function of halo mass for different UVB runs. Each point in the figure represent a halo with star-forming galaxies, including the effect of stellar RT. Each run has $\\sim 2000$ galaxies in the simulation box at $z = 3$, and about 200 of them have young star clusters ($t_{\\rm age} \\le 10^{7}~\\rm yr$) that can affect the hydrogen ionization (hereafter we call these galaxies ``star-forming galaxies'').  Here, we select the galaxies with young stars from each run, and compute the RT of local stellar sources.  We find that the cross section increases with increasing halo mass for both DLAs and LLSs in all UVB models. For a quantitative discussion, we plot the median value in each bin with a triangle symbol, and then fit a power law $\\sigmadla  \\propto \\Mh^{\\alpha}$, assuming a functional form of \\begin{equation} \\log \\sigma = \\alpha (\\log \\Mh - 12) + \\beta, \\label{eq:sigma} \\end{equation} where $\\Mh$ is the dark matter halo mass in units of solar masses, and the values of slope '$\\alpha$' and normalization '$\\beta$' are determined by least-square fitting.  We use the cross section at $\\Mh = 10^{12} \\Msun$ as the anchor point of $\\beta$.  The best-fit values of $\\alpha$ and $\\beta$ of all runs are summarized in Table~\\ref{table:alpha}.  The bottom right panels of Figure~\\ref{fig:sigmauv}a,b compare the results of different UVB models. For both DLAs and LLSs, the cross section is the greatest for the No-UVB run, and it drops dramatically once the UVB is turned on. The difference between the Optically-Thin run and the HM0.5 run is small, suggesting that the gas is ionized to relatively high densities under the optically-thin approximation, irrespective of the specific strength of the UVB. For DLAs, the OTUV run has much higher cross sections than in the HM0.5 run, because the high-density clouds with $n>\\nuv$ are the main contributor to $\\sigmadla$, and the OTUV model increases the neutral fraction of such gas over the optically-thin approximation. In fact, the value of $\\beta$ for the OTUV run is much higher than that of HM0.5 run, but the slope is similar in all runs with $\\alpha \\approx 0.65-0.72$ for DLAs. This suggests that $\\sigmadla$ is affected proportionally at all halo mass range by the different UVB models. Obviously the result of the Optically-Thin run is very similar to that of \\citet{Nagamine04f}, as they both used the optically-thin approximation in similar simulations (the main difference is the galactic wind model).  \\begin{table} \\begin{center} \\begin{tabular}{ccccccc} \\hline Run &  $\\alpha$(DLA) & $\\alpha$ (LLS) & $\\beta$ (DLA) & $\\beta$ (LLS)  \\\\ \\hline Optically-Thin     & 0.65   & 0.61 & 2.12 & 3.59  \\\\ HM0.5   & 0.67 & 0.56  & 2.23 & 3.66  \\cr OTUV   & 0.70 & 0.53  & 3.06 & 3.41  \\cr OTUV (no-star)       & 0.72   & 0.55 & 3.11 & 3.46  \\\\ no-UVB   & 0.66 & 0.30  & 3.48 & 3.94  \\cr \\hline \\end{tabular} \\caption{ Best-fit parameters of the fitting function (Eq.\\,[\\ref{eq:sigma}]) for $\\sigmadla$ and $\\sigmalls$ for different UVB models. The value of $\\alpha$ is the slope of the power law, and $\\beta$ is the normalization point at $\\Mh = 10^{12}\\Msun$. } \\label{table:alpha} \\end{center} \\end{table}  The biggest difference between $\\sigmadla$ and $\\sigmalls$ is that the OTUV run (blue line) has almost the same $\\sigmalls$ as the Optically-Thin run (red line), whereas $\\sigmadla$ are very different in the two runs.  This implies that the high-density gas with $n>\\nuv$ is not responsible for the LLS.  The HM0.5 run has a slightly higher $\\sigmalls$ than the Optically-Thin run as expected. We also find that the slope $\\alpha$ is systematically shallower for the LLSs compared to that of DLAs, and the normalization $\\beta$ is higher for the LLSs.  This is expected, because LLSs are more extended in the outskirts of halos.  If we look at the distribution of data points in Figure~\\ref{fig:sigmauv}a,b more carefully, one can see that there is a slight curvature to the median points, and a straight power-law does not necessarily provide a perfect fit. At low halo masses, the median points fall below the power-law fit. At the intermediate to high masses, the median points are slightly higher than the power-law fit, and the cross section seems to flatten out at the highest masses at $\\Mh \\sim 10^{12} \\Msun$.  A similar flattening has been observed in other works \\citep{Pontzen08, Fumagalli11}, as we will discuss in the next section in more detail.   \\begin{figure} \\begin{center} \\includegraphics[scale=0.42] {fig5.eps} \\caption{ Comparison of DLA cross section with other theoretical works. The black dotted and dashed lines are the results of P3 and Q3 models, respectively, by \\citet{Nagamine04g}. The black long-dashed and dot-dashed lines are the strong wind (SW) and weak wind (WW) runs, respectively, in the GADGET SPH simulations of \\citet{Tescari09}. The black filled circles are the cross sections of seven simulated galaxies in the AMR simulation of \\citet{Fumagalli11}. The cyan and orange shade enclose the data points from the simulations of  \\citet[][Gasoline SPH simulation]{Pontzen08} and \\citet[][Enzo AMR simulation]{Cen10}, respectively. } \\label{fig:sigma_com_model} \\end{center} \\end{figure}   \\begin{figure} \\begin{center} \\includegraphics[scale=0.4]{fig6.eps} \\caption{ {\\it Upper panels} : The ratio of cross sections of DLAs and LLSs in each halo with ($\\sigma$) and without ($\\sigma^{0}$) stellar RT: i.e., $\\sigma / \\sigma^{0}$. {\\it Lower panel} : The mass ratio of neutral hydrogen in each halo with ($M_{\\rm HI}$) and without ($M^{0}_{\\rm HI}$) stellar radiation, i.e., $M_{\\rm HI} / M^{0}_{\\rm HI}$. The red points are the mean value in each bin with the bin size of 0.5 dex. } \\label{fig:dlacom} \\end{center} \\end{figure}  \\subsubsection{Comparison with Other Simulation Results}  In Figure~\\ref{fig:sigma_com_model}, we compare the results of DLA cross section from several authors using different simulations \\citep{Nagamine04f, Pontzen08, Tescari09, Cen10, Fumagalli11}. Although there are small differences in the slope and normalization, most of the results are bracketed by our no-UVB run and Optically-Thin run. There are significant overlap between the results of our OTUV run, \\citet{Pontzen08}, \\citet{Fumagalli11}, and the WW (weak wind) model of \\citet{Tescari09}, which is very encouraging.  \\citet{Pontzen08} considered a crude self-shielding model, so it is natural that their result is close to the OTUV run. Although \\citet{Tescari09} also set the self-shielding threshold density criteria above which UVB cannot penetrate, their $\\nuv$ is about one magnitude higher than our value. Hence, their SW (strong wind) model should give lower $\\sigmadla$ than our OTUV run. The SW and WW model have constant wind particle velocities of $600$ and 100\\,km\\,s$^{-1}$ as a stellar feedback, respectively. The difference between the SW and WW model of \\citet{Tescari09} is consistent with the results of \\citet{Nagamine04f}.  Using an Enzo AMR simulation, \\citet{Cen10} reported a much higher $\\sigmadla$ than other simulations for halos with $\\Mh=10^{10}-10^{11} \\Msun$, and the reason for this  difference is not very clear.  Some of the possibilities are: 1) their star formation model is considerably inefficient in converting the neutral gas into stars, 2) their self-shielding treatment is allowing too much neutral gas in massive halos, and/or 3) their stellar feedback is inefficient and does not heat up the ambient neutral gas.  They used only two zoom-in simulations (a high-density cluster region and a low-density void region) to bracket the mean density result, therefore they only bracketed the observed $\\fn$ without reproducing the normalization of $\\fn$.  It appears that his result is more consistent with that of our No-UV run rather than the OTUV run.  We find that the Fumagalli's fit to the $\\sigmadla$ for low-mass halos is much steeper than other simulations, which could be due to a small galaxy sample in their zoom-in simulations (only seven galaxies).  We discussed the effect of limited galaxy sample already in \\S~\\ref{sec:fn} and Figure~\\ref{fig:cdfall}.  \\begin{figure*} \\begin{center} \\includegraphics[scale=0.8]{fig7.eps} \\caption{ Two dimensional map of $\\NHI$ for four halos in the OTUV run at $z=3$. Each column shows different treatment of radiation: column ($a$): no-ionization;  ($b$): UVB $+$ collisional ionization, ($c$): UVB $+$ star $+$ collisional ionization. The column ($d$) is the same as column ($c$), but for the same halo in the Optically-Thin run. Each row corresponds to one halo with a panel size $\\ell_{\\rm phys}$ in physical units: Halo (1): $\\Mh = 6.7 \\times 10^{11} ~{\\rm \\Msun}, \\ell_{\\rm phys} = 300~ {\\rm kpc}$; Halo (2) $\\Mh = 1 \\times 10^{11} ~{\\rm \\Msun}, \\ell_{\\rm phys} = 159~ {\\rm kpc}$; Halo (3) $\\Mh = 2.6\\times 10^{10} ~{\\rm \\Msun}, \\ell_{\\rm phys} = 80 ~{\\rm kpc}$; and Halo (4) $\\Mh = 3.3 \\times 10^{9} ~{\\rm \\Msun}, \\ell_{\\rm phys} = 47~ {\\rm kpc}$. The white vertical tick mark in column ($d$) indicates a scale of proper 30 kpc. } \\label{fig:column} \\end{center} \\end{figure*}   \\subsubsection{Effect of Stellar RT on \\HI\\ Cross Sections} \\label{sec:stellarRT}  The upper panels of Figure~\\ref{fig:dlacom} compares the ratio of DLA and LLS cross section with and without stellar radiation. The data points are only for those halos that contain star-forming galaxies in the OTUV run at $z=3$. We see that both $\\sigmadla$ and $\\sigmalls$ of lower mass haloes are reduced by the stellar RT, but the higher mass haloes are not affected very much. The slope $\\alpha$ and normalization $\\beta$ for the power-law fit are very similar in the two cases of with and without the stellar RT. Since young stars distribute near high-density gas clumps, most of ionizing photons from stars are blocked by the gas clumps, and the region behind high-density gas clumps are not irradiated by the stellar radiation. As a result, the stellar radiation cannot propagate over long distances in the halo, and the stellar radiation does not affect the \\HI\\ cross section very much in high-mass halos. In addition, although relative mass of neutral hydrogen gas to stellar mass somewhat decreases as increasing halo mass, the amount of dense gas clouds around star-forming regions increases with halo mass by frequent minor merging processes \\citep[e.g., Figure~1~in][]{Yajima11}. The photon absorption process and the distribution of stars and gas are determined by complicated nonlinear process, which can be traced by cosmological hydrodynamics and radiative transfer simulations.  The lower panel of Figure~\\ref{fig:dlacom} shows the ratio of \\HI\\ mass in each halo with ($M_{\\rm HI}$) and without ($M^{0}_{\\rm HI}$) the stellar RT, i.e., $M_{\\rm HI} / M^{0}_{\\rm HI}$. It shows that the lower mass halos lose more \\HI\\ masses than the higher mass halos, consistently with Figure~\\ref{fig:dlacom}. The \\HI\\ mass is reduced by about 30 per cent by the stellar RT for halos with $\\Mh < 10^{10}\\;\\Msun$, while it is only about 7 per cent reduction for higher mass halos with $\\Mh > 10^{11}\\;\\Msun$.  Some analytical works discussed the relative strength of local stellar radiation to UVB under the assumption of simple gas distribution \\citep{Miralda05, Schaye06}. \\citet{Miralda05} showed that the effect of stellar radiation is negligible compared to UVB for LLSs, and that it can be comparable for DLAs as an upper limit. On the other hand, the model by \\citet{Schaye06} indicated that the local stellar radiation can be dominant even for LLSs, as well as for DLAs. However, for inhomogeneous gas distribution, the effect of stellar radiation strongly depends on the structure of gas around stars, and it requires RT simulations for proper assessment.  The above analytical works did not discuss the direct impact of local stellar radiation on DLA/LLS cross sections and $\\fn$, and focused only on the {\\it relative} strength between local stellar radiation and UVB. In our simulations, the ionizing flux from local stars can be much higher than the UVB similarly to the previous analytic estimates. However, the neutral clouds survive from ionization by self-shielding and shadowing effect \\citep[See Fig.1 of][]{Yajima11}. As a result, the local stellar radiation cannot change the ionization structure in galaxies very much (particularly in massive galaxies),  and hence the \\HI\\ cross section does not change largely. On the other hand, the UVB enter from outside the halo into low-density region, therefore they can irradiate over greater volume.    \\subsection{Structure of DLAs and LLSs}  As we see in Figure~\\ref{fig:cdfall}, the UVB can drastically change the rate-of-incidence of DLAs and LLSs. In this subsection, we investigate how the UVB and stellar radiation change the spatial distribution of $\\NHI$. Figure~\\ref{fig:column} shows the two dimensional map of $\\NHI$ distribution for four halos in the OTUV run at $z=3$. The column ($a$) is without any ionization by UVB and stellar radiation, corresponding to the No-UV run. The column ($b$) is with the UVB RT and collisional ionization, but no stellar RT. The column ($c$) is with the stellar RT in addition. The column ($d$) is the same as column ($c$), but for the same halo in the Optically-Thin run.  Comparing columns ($a$) and ($b$), we see that the UVB changes the $\\NHI$ distribution drastically. Much of the LLSs in the outer region are ionized by the UVB and becomes lower than $\\NHI = 10^{16}$\\,cm$^{-2}$. The higher columns with $\\NHI \\gtrsim 10^{20.5}$ survive from photo-ionization by UVB because of the high recombination rate in high density regions, but the DLA cross section is significantly reduced.  The DLAs correspond to the yellow and red regions in Figure~\\ref{fig:column}. The DLA distribution does not change very much from column ($b$) to ($c$), and it is not affected by the stellar RT very much as we discussed in \\S~\\ref{sec:stellarRT}.  Comparison of columns ($c$) and ($d$) reveals the effect of optically-thin approximation of UVB in the Optically-Thin run relative to the self-shielding effect in the OTUV run. It shows that the UVB photo-ionizes the gas too deeply in the optically-thin approximation, and the self-shielding model increases the DLA cross section significantly. From the discussion in \\S~\\ref{sec:fn}, we argue that column ($c$) is more realistic than column ($d$) based on the comparison of $\\fn$.  In the OTUV run (column [$c$]), DLAs are extended over $\\sim 10$\\,kpc scales in all halos, which is a stark contrast to the very compact appearance in the Optically-Thin run (column [$d$]), and the LLSs are surrounding DLAs on $\\sim 30-60$\\,kpc scales. \\citet{Nagamine04f} showed the DLA distribution similar to the column ($d$) in their SPH simulations with optically-thin approximation and argued that DLA distribution is very clumpy.  However based on the present work, it seems that DLAs are more broadly and smoothly distributed around LBGs at $z=3$ on $\\sim 10$\\,kpc scales. The highest column density systems with $\\NHI > 10^{22}$\\,cm$^{-2}$ are preserved well even in the Optically-Thin run, as we saw in Figure~\\ref{fig:cdfall}.   \\begin{figure} \\begin{center} \\includegraphics[scale=0.45]{fig8.eps} \\caption{ Cumulative probability distribution function (PDF) of DLAs and LLSs as a function of separation distance (i.e. impact parameter) between DLAs (LLSs) and center-of-mass of the nearest galaxy in projected 2D plane in physical unit. The colored lines are our simulation results in this work. The black dash and dot-dash lines are PDFs with the constant-velocity (CW) wind model + optically-thin UVB approximation \\citep[P3 run in][]{Nagamine07} and the CW+OTUV run \\citep{Nagamine10}, respectively. } \\label{fig:PDF} \\end{center} \\end{figure}   To quantify the relative distribution of DLAs, LLSs and galaxies in different UVB models, we compute the impact parameter distribution of DLAs/LLSs relative to galaxies in each run. Figure~\\ref{fig:PDF} shows the cumulative probability distribution function (PDF) of DLAs and LLSs as a function of separation distance (i.e. impact parameter) between DLAs (LLSs) and the center-of-mass of nearest galaxy in projected 2D plane in physical unit. The colored lines are from the simulations used in this paper. The OTUV run shows a much shallower PDF than the Optically-Thin run. About fifty per cent of DLAs in the OTUV run distribute within $\\sim 7$~kpc from galaxy center, while those in the Optically-Thin run are concentrated within $\\sim 3$~kpc. This is because the UVB cannot penetrate deeper into the galaxy center due to self-shielding, and hydrogen gas near galaxy center can survive as DLAs. Of course, no-UVB run has the widest PDF due to the overproduction of DLAs as noted earlier in this paper.  In addition, we compare the results with different galactic wind models. As described in Section 2.1, we use the \"Multi-component Variable Velocity\" (MVV) wind model of \\citet{Choi10} in the current simulations. Some previous works used the constant-velocity wind (CW) model. The black dashed and solid lines are the PDFs from the CW model + optically-thin UVB approximation \\citep[P3 run in][]{Nagamine07} and CW+OTUV run \\citep{Nagamine10}, respectively. We find that the DLAs in the CW model are concentrated near the galaxy center more than the MVV model. In the CW model, the lower mass galaxies have higher wind velocities than in the MVV run \\citep{Choi10}, therefore the DLAs abundance is somewhat suppressed. In such a situation, only high-density gas around galaxy center can survive as DLAs, resulting in a steeper PDF. In the case of LLSs, the PDFs are in general much wider than those of DLAs'.   \\subsection{Stellar Mass and Metallicity Distribution}  We show the projected maps of stellar mass, $\\NHI$, and metallicity in Figure~\\ref{fig:mtl}. We find a strong correlation between high-$\\NHI$ systems and stellar mass distribution as expected. The metallicity map is obtained by summing up all metal mass along each line-of-sight and dividing by the total gas mass in the same sight-line. The high density region is effectively enriched by the metals created by SNe. The massive halos (1) and (2) show high metallicities around the high density regions and DLAs.  Galactic wind can also carry high-metallicity gas outside the galaxies. All bottom panels of Figure~\\ref{fig:mtl} show that the near-solar metallicity gas is carried well outside of the stellar distribution, up to $\\sim 30-50$\\,kpc away from stars. The effect of wind appears stronger in halos (3) and (4), but this is simply due to the smaller physical scale of the panels for halos (3) \\& (4) compared to (1) \\& (2).  Careful examination shows that halos (1) \\& (2) also show a similar degree of circumgalactic enrichment as halos (3) \\& (4). For example, for halo (1) there is a bubble of sub-$\\Zsun$ gas extending toward north-east direction, about 60\\,kpc away from the center of galaxy on the right-hand-side.  For halo (2), a similar feature can be seen for a near-$\\Zsun$ gas extending out towards north-east direction, again about 60\\,kpc away from the center of galaxy on the right-hand-side.  For halo (3), there is a large bubble of near-$\\Zsun$ gas extending towards south-west direction, about 30\\,kpc away from the stars. For halo (4), a similar feature can be seen towards north direction. For halos (1)$-$(3), notice that the direction of wind propagation is nearly perpendicular to the chain of galaxy distribution.  Our MVV wind model reproduces the observed scaling between wind velocity and galaxy SFR: $v_w \\propto SFR^{1/3} \\propto M_\\star^{1/3}$, which allows greater wind velocity for more massive galaxies.  This is compensated by the deeper potential well of more massive galaxies, therefore the enriched metals reach similar distances for all halos as we discussed above.  The bottom two rows of Figure~\\ref{fig:mtl} also show the virial radius of each halo in white circles. One can see that the DLAs are about the same or slightly smaller than the virial radius, whereas the metal enrichment clearly goes beyond the virial radius.  For halos (1) \\& (2), there are two major galaxies in these massive halos, and each of the galaxies is a host of DLAs that extends to the scale of virial radius.  \\begin{figure*} \\begin{center} \\includegraphics[scale=0.8]{fig9.eps} \\caption{ Two dimensional map of projected stellar mass (top row), $\\NHI$ (middle row), and metallicity (bottom row) for the same four halos shown in Figure~\\ref{fig:column} in the OTUV run. Each column with the number index corresponds to the same halo in each row of Figure~\\ref{fig:column}, and the panel sizes are the same as in Figure~\\ref{fig:column}. The color scheme of stellar mass is normalized in each panel. The middle row is the same as the column (c) in Figure~\\ref{fig:column}, but we show it again for comparison. The white bar in the middle row shows the proper 30 kpc.  Note the different panel sizes for each column. The metallicity map is obtained by summing up all metal mass along each line-of-sight and dividing by the total gas mass in the same sight-line.  The white circles indicate the virial radius of each halo: physical 68, 37, 23 and 11 kpc for halos (1)-(4), respectively. The virial radius is estimated by $R_{\\rm vir} \\simeq 144 (1 + z_{\\rm vir})^{-1} (\\Mh / 10^{11} \\Msun)^{1/3}$\\,kpc \\citep[e.g.,][]{Mo02}, and here we assumed $z_{\\rm vir} = 3$. } \\label{fig:mtl} \\end{center} \\end{figure*}    \\subsection{Cumulative DLA Rate of Incidence} \\label{sec:abundance}  Using the fitting function of DLA cross section derived in \\S~\\ref{sec:sigma}, we are now able to estimate the cumulative number of DLAs per unit redshift. However, there is the mass-resolution limit in a simulation, and we cannot resolve the lower mass haloes that may be DLAs host. It may cause the underestimation of the DLA abundance. In addition, it is hard to produce very massive halos in a limited simulation volume. To overcome this resolution limit, the DLA abundance has been estimated by combining a theoretical fit to the dark matter halo mass function and the derived relationship between DLA cross-section and halo mass \\citep{Gardner97a, Gardner97b, Gardner01, Nagamine04a, Nagamine07, Tescari09}. Here, we follow the same method. The cumulative number of DLAs is estimated by \\begin{equation} \\frac{dN_{\\rm DLA}}{dz} (> M, z) = \\frac{dr}{dz} \\int^{\\infty}_{M} n_{\\rm h}(M^{'}, z) \\sigma_{\\rm DLA}(M^{'}, z) dM^{'}, \\end{equation} where $ n_{\\rm h}(M^{'}, z)$ is the dark matter halo mass function, for which we use the function derived in \\citet{Tinker08}, and $dr / dz = c / H_{0} \\sqrt{\\Omega_{\\rm m}(1 + z)^{3} + \\Omega_{\\rm \\Lambda}}$.   Figure~\\ref{fig:dndz} shows the cumulative DLA number as a function halo mass for different UVB models. The yellow shaded region is the observational estimate $\\log (dn/dz) = - 0.6 \\pm 0.1$ at $z=3$ from SDSS data of \\citet{Prochaska05}. We find that the cumulative abundance largely depends on the UVB models. The OTUV run agrees well with the observed data , if the cross section steeply drops off at $\\Mh < 10^{9}\\,\\Msun$. In this work, we set the lower limit of halo mass to $10^{9}\\,\\Msun$ for the above integration, because the DLA cross section steeply declines at halo masses of $10^{8.5}-10^9 \\Msun$ \\citep{Nagamine04g, Pontzen08}. However, there is some ambiguity in the DLA formation in lower mass haloes, and we would overpredict the abundance if we were to integrate down to $10^{8.5}\\,\\Msun$ for the above integration.  On the other hand, the No-UVB run clearly overproduces DLAs. The Optically-Thin and HM0.5 runs underpredict the observation by some factor. The result of the Optically-Thin run is slightly lower than that of \\citet{Nagamine04f, Nagamine07}. The new MVV wind model used in this work shows the smaller $\\sigmadla$ at higher halo mass. In addition, the halo number density by \\citet{Tinker08} with WMAP 7-year parameters is slightly lower than the mass function used in their paper. Hence, it results in smaller DLA abundance even with the same optically-thin UVB prescription. Although our $\\sigmadla$ in the OTUV run is somewhat larger than that of SW run in \\citet{Tescari09}, our result still agrees with the observation data well.   \\begin{figure} \\begin{center} \\includegraphics[scale=0.4]{fig10.eps} \\caption{ Cumulative abundance of DLAs per unit redshift as a function of halo mass for the Optically-Thin (dotted line), HM0.5 (dot-dash line), OTUV (red solid line) and No-UVB (dash line) runs at $z=3$. The yellow shaded region shows the observed cumulative DLA abundance of \\citet{Prochaska05} from SDSS data. } \\label{fig:dndz} \\end{center} \\end{figure}            \\subsection{Typical physical quantities of DLA host galaxies} \\label{sec:typical}  In this section, we examine the typical physical quantities of DLA host galaxies in the OTUV run. Figure~\\ref{fig:crossdistdla} shows the $\\sigmadla$-weighted probability distribution function (PDF) as a function of certain physical quantity.  For example, PDF\\,$\\propto d\\sigmadla / d\\log M_h$ for the top left panel. The histograms are normalized such that the area under the histogram is equal to unity. The dashed and dotted lines indicate the median and mean values weighted by the cross section, respectively. We find the median halo mass by rank ordering the halo mass according to the $x$-axis quantity, and determining the halo at which the stacked cross section being half of the total. We see that the mean values are biased towards higher values and not representative of the entire distribution, except for the bottom right panel where the distribution is strongly peaked at a certain metallicity.  Therefore median value is the better indicator of typical property of DLA hosts.  The upper left panel of Figure~\\ref{fig:crossdistdla} shows that the median halo mass of DLA hosts is $\\Mh = 2.4 \\times 10^{10} \\Msun$, and about half of the total $\\sigmadla$ is contributed by the halos with $\\Mh = 6.6 \\times 10^{9} - 1.1 \\times 10^{11}~\\Msun$, centered around the median ($\\pm$25\\% around median). Although $\\sigmadla$ increases with increasing $M_h$, the number density of massive halos decreases at the same time. This effect pushes the median DLA halo mass to be in the medium mass range. \\citet{Nagamine07} also presented the median DLA halo mass for various wind models, and our results are similar to their P3 or Q3 model. The cyan shade indicates the mass range of DLAs host galaxies suggested by clustering analysis \\citep{Cooke06}. \\citet{Cooke06} showed that mass of galaxies hosting DLAs were $10^{9.7} < M < 10^{12}~\\Msun$ from DLA-LBG cross-correlation. Our simulation results agree well with Cooke's result, and the simulation work \\citep{Lee11}. In addition, we derive the bias parameter from the median mass. Here we use an ellipsoidal collapse model \\citep{Sheth01, Mo02}, which can reproduce results of $N$-body simulations well. We derive the rms fluctuation of mass density at a mass scale M, $\\sigma(M)$ by using the median mass of DLAs, and assign it in the equation (8) of \\citet{Sheth01}. As a result, we obtain the bias of  $b=1.69$ (Table~\\ref{table:median}). This bias is also in the observed range of $1.3 < b_{\\rm DLA} < 4$ by \\citet{Cooke06}.  The distribution is relatively flat as a function of SFR. The typical LBGs have SFR $\\approx 1-20~ \\Msunyr$, and the fraction of $\\sigmadla$ contributed by such LBGs is $\\sim 30$ per cent based on our simulation. The median SFR in the top right panel of Figure~\\ref{fig:crossdistdla} is $0.3\\,\\Msunyr$. Using the Equation~(1)  in \\citet{kennicutt98} for the Salpeter IMF, this median SFR corresponds to $\\sim 0.02\\,L_{\\star}$ of LBGs at $z=3$, which is much fainter than the observational flux limit ($0.1\\,L_{\\star}$) in the recent LBG survey of \\citet{Reddy09}. Therefore we suggest that more than half of total $\\sigmadla$ are hosted by the fainter LBGs than the current observational limit, which is consistent with the observation \\citep[e.g.,][]{Fynbo99}.  The green shaded region is an approximate range of SFR for DLAs derived by the $\\rm C_{[II]}$ cooling technique \\citep{Wolfe03}. More than eighty per cent of DLA cross section in our simulations overlaps with the green shaded range. There are some observed DLA galaxies with higher SFR of a few tens $\\Msunyr$, which are related to LBGs \\citep[e.g.,][]{Moller02}. The massive haloes with $\\Mh \\gtrsim 10^{11.5}~\\Msun$ in our simulation would correspond to those DLA host galaxies with relatively high SFR.  The bottom two panels of Figure~\\ref{fig:crossdistdla} show the PDF as functions of stellar mass and gas metallicity. The median stellar mass is $\\sim 2.4 \\times 10^{8} \\Msun$, and the median metallicity is $\\langle Z/ \\Zsun \\rangle = 0.1$. Fifty per cent of total $\\sigmadla$ are hosted by galaxies with $3.9 \\times 10^{7} < M_\\star / \\Msun < 1.8 \\times 10^{9}$ and $0.06 <  Z/ \\Zsun < -0.2$, both centered around the median values (i.e. $\\pm$25\\% around median). The distribution in terms of metallicity is strongly peaked around $\\langle Z/ \\Zsun \\rangle = 0.1$, and has an extended tail to lower metallicities. Our mean/median metallicity is not so far from the peak of observational data,  $Z/ \\Zsun \\approx  0.1$ \\citep{Prochaska07a}. However, our PDF is somewhat shifted to higher metallicity than the observed DLA's at $2 \\le z \\le 4$ \\citep{Prochaska07a} and $1.5 \\le z \\le 5$ \\citep{Rafelski12}. Our simulation can successfully reproduce $\\fn$, but it seems to overpredict DLA metallicity distribution. The metallicity--$\\NHI$ relation is not sensitive to the wind models \\citep[e.g., Figure 9 in ][]{Nagamine04f}. Hence, our current wind model does not seem to cause the higher metallicity. The higher metallicity may be caused by the star formation model. If our star formation model is too efficient at high redshift, the gas is quickly enriched by metals, and it leads to higher metallicity of DLAs. In addition,  the resolution of numerical simulations may change the metallicity distribution of DLAs. In general, higher resolution simulations can resolve smaller dwarf galaxies, and many gas clumps in them. If the gas in the dwarf galaxies which have little metals can contribute to the DLA abundance, the typical metallicity of DLAs can be reduced. These topics will be addressed in future works.  Another possibility is the selection bias of DLAs in observation due to the dust reddening \\citep[e.g.,][]{Ellison01, Fynbo10, Fynbo11, Khare12}. Recently \\citet{Fynbo10, Fynbo11} found some metal-rich DLAs using the detection technique based on Si\\,{\\sc{ii}} and Fe\\,{\\sc{ii}} absorption lines. These objects were not detected in the automatic DLA searches in the SDSS using Ly$\\alpha$ \\citep[e.g.,][]{Prochaska05}. They suggested that dust in such metal-rich DLAs could obscure the QSOs, and cause the non-detection of DLAs using Ly$\\alpha$.  This would lead to a distribution of DLAs biased to low metallicities. Hence, the difference between our simulations and observation may come from this selection bias. \\citet{Khare12} concluded that $\\le 10$ per cent of the DLAs in SDSS DR7 cause significant reddening. It is unclear if this reddened population can explain the discrepancy we see in the PDF, but it will bring the two distributions closer.  The dashed histograms are the PDFs in the optically thin run. Although $\\fn$ significantly changes with the UVB model, the typical physical quantities do not change very much. As we see in Figure~\\ref{fig:sigmauv}, $\\beta$ in the fitting function (Equation 5) changes with the UVB model, while the power-law slope $\\alpha$ does not change largely. In our simulations, the physical quantities are roughly proportional to halo mass. Therefore, the typical physical quantities do not depend on the UVB model very much.  Figure~\\ref{fig:crossdistlls} shows the PDF of $\\sigmalls$. The shape of PDF is basically similar to that of $\\sigmadla$. However, in the case of LLSs, the contribution from lower mass (and hence lower SFR and lower $\\Mstar$) haloes becomes greater than for DLAs. The median values are:  $M_h = 9.6 \\times 10^{9}\\,\\Msun$, $\\SFR = 0.06\\,\\Msunyr$, $\\Mstar = 6.5 \\times 10^{7} \\,\\Msun$, and $Z/\\Zsun = 0.08$. About 80 per cent of total $\\sigmalls$ come from haloes with ${\\rm SFR} \\lesssim 1~\\Msunyr$. Hence the majority of LLSs are hosted by the low-mass halos with LBGs that are fainter than the current observational limit. Fifty per cent of total $\\sigmalls$ are hosted by haloes with $2.5 \\times 10^{7} < M_\\star / \\Msun < 5.3 \\times 10^{8}$ and $0.05 < Z/ \\Zsun <  0.15$. We summarize the various median values in Table~\\ref{table:median}.   \\begin{figure} \\begin{center} \\includegraphics[scale=0.43]{fig11.eps} \\caption{ The PDF of DLA cross section in the OTUV run at $z=3$ as a function of various physical quantities: halo mass, SFR, stellar mass, and metallicity, from top-left to lower right, respectively. The dashed histograms are PDFs in the Optically-Thin run. The vertical dashed and dotted lines are the median and mean of PDF, respectively. The green shaded region is the halo mass range of DLAs host galaxies suggested by the clustering analysis \\citep{Cooke06}. The cyan shaded region is an approximate range of SFR for DLAs derived by the $\\rm C_{[II]}$ cooling technique \\citep{Wolfe03}. The vertical, blue dot-dashed line in the upper right panel is the SFR of $L^*$--LBGs at $z=3$ \\citep{Reddy09}. The purple histogram in the bottom right panel shows the PDF of observed DLA metallicities at $2 \\le z \\le 4$ by  \\citet{Prochaska07a}. The blue dashed curve is the Gaussian fit to the metallicity of 207 DLAs at $1.5 \\le z \\le 5$ \\citep{Rafelski12}. } \\label{fig:crossdistdla} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[scale=0.43]{fig12.eps} \\caption{ Same as Fig.~\\ref{fig:crossdistdla}, but for LLS cross section. } \\label{fig:crossdistlls} \\end{center} \\end{figure}  \\begin{table} \\begin{center} \\begin{tabular}{ccccccc} \\hline &  $\\langle \\Mh \\rangle$ & $\\langle {\\rm bias} \\rangle$ &$\\langle {\\rm SFR}  \\rangle$ & $\\langle M_{*}  \\rangle$ & $\\langle Z / \\Zsun \\rangle$  \\\\ &  $(10^{10}~\\Msun)$ && $(\\Msunyr)$ & $(10^{8}~\\Msun)$ &   \\\\ \\hline DLA     & 2.39   &1.69& 0.33 & 2.43 & 0.10  \\\\ LLS   & 0.96 & 1.50&0.06  & 0.65 & 0.08  \\\\ \\hline \\end{tabular} \\caption{ Median values of various physical quantities for the PDF of DLA/LLS cross section. The bias is derived from the median halo mass by using an ellipsoidal collapse model \\citep{Sheth01}. } \\label{table:median} \\end{center} \\end{table} ",
        "conclusions": "\\label{sec:summary}  We examined the physical properties of DLAs and LLSs at $z=3$ using cosmological SPH simulations. In particular, we studied the impact of radiative transfer of UVB and local stellar radiation on DLAs and LLSs. Our major findings are as follows:  \\begin{itemize} \\item The OTUV simulation with a self-shielding density threshold of $\\nuv = 6\\times 10^{-3}$\\,cm$^{-3}$ for the UVB can reproduce the observed $\\fn$ very well (Figure~\\ref{fig:cdfall}), compared to the run with optically-thin approximation (the Optically-Thin run).  We have presented this result earlier \\citep{Nagamine10}, and in this paper we further validated the adopted value of $\\nuv$ by the direct RT calculations. Similar results have been obtained various authors to a varying degree \\citep{Pontzen08, Tescari09, Fumagalli11, McQuinn11}.  As we showed in Figs.~\\ref{fig:gammauvb}, \\& \\ref{fig:uvbcom},  the above value of $\\nuv$ captures the rapid change in the ionization fraction in halo gas very well.  In the Optically-Thin run, the UVB sinks too deeply into the halo gas and ionizes gas too much, resulting in the underestimate of $\\fn$. Given that it is still difficult to perform RT calculations simultaneously with hydrodynamics in current cosmological simulations, our OTUV model provides a useful prescription for any cosmological simulations that cannot resolve the self-shielding by molecular gas on sub-kpc scales.  \\item We find that the local stellar radiation does not strongly affect the distribution of DLAs/LLSs, and $\\fn$ does not change very much by the stellar RT (Figure~\\ref{fig:cdfall}).  This is because clumpy high-density clouds near young star clusters effectively absorb most of the ionizing photons from young stars \\citep{Yajima11}. The effect of stellar RT is stronger on the lower mass halos, and the reduction in $\\sigmadla$ and $\\sigmalls$ were shown in Figs.~\\ref{fig:dlacom}.  \\item We compared the DLA/LLS cross sections in simulations with different UVB models (Figure~\\ref{fig:sigmauv}).  In our simulations, the No-UVB run gives the highest DLA/LLS cross sections as expected, and it overpredicts the rate-of-incidence and the column density distribution.  For the DLAs, the OTUV run gives intermediate $\\sigmadla$ values, and the HM0.5 \\& Optically-Thin run give the lowest $\\sigmadla$.  The similarity of the latter two runs tells us that the gas becomes highly ionized under the optically thin run, no matter what the exact UVB amplitude is. For the LLSs, the latter three runs (OTUV, Opticall-Thin, and HM0.5) all give similar $\\sigmalls$, and the self-shielding model does not matter for LLSs, as they are in the outskirts of galaxies.  \\item We compared the results of DLA cross section from various authors and simulations (Figure~\\ref{fig:sigma_com_model}).   There are some general agreement in the sense that most results are bracketed by our Opticall-Thin run and the No-UVB run, and that both $\\sigmadla$ and $\\sigmalls$ increases with increasing halo mass.   However, there is still significant scatter between different authors.   The origin of the differences could be a combination of many things, such as different models of star formation, galactic wind, UVB and self-shielding treatment. For example, the stronger wind model can sweep high-density gas clouds out of star-forming regions, and decrease DLA cross section. The self-shielding model of UVB can increase the DLA cross section compared to the optically-thin method, and rate of increase can be different according to the threshold density for self-shielding. We have chosen the appropriate value of $\\nuv$ based on the good match to the observed $\\fn$ and through the validation by direct radiative transfer calculation.  \\item The cumulative number of DLAs per unit redshift strongly depends on the UVB models (Figure~\\ref{fig:dndz}). The OTUV run agrees well with the observed data. On the other hand, the No-UVB run overproduces DLAs. The Optically-Thin and HM0.5 runs underpredict the observation.  \\item The visual appearance of DLAs and LLSs is quite different in the OTUV and Optically-Thin runs (Figure~\\ref{fig:column} \\& \\ref{fig:mtl}). In the OTUV run, the DLAs have an extended distribution around star-forming regions on typically $\\sim 10-30$\\,kpc scales, and LLSs are surrounding DLAs on $\\sim 30-60$\\,kpc scales. As shown in the middle row of Figure~\\ref{fig:mtl}, the extent of DLAs are typically smaller than the virial radius of the halo, and the LLSs are about the same size as the virial radius.  Given the good agreement between the OTUV run and the observations of $\\fn$, we consider that these are the most realistic distribution of DLAs and LLSs in our simulation set.  On the other hand, in the Optically-Thin run, the UVB penetrates deeply, ionizes the hydrogen gas effectively, and reduces the column density significantly. As a result, most of the region around star-forming galaxies turn into LLSs, and the DLAs shrink and become compact.  This striking effect of optically-thin approximation is highlighted by the comparison of columns ($c$) and ($d$) in Figure~\\ref{fig:column}, suggesting the importance of self-shielding effect.  The present work supersedes the previous work of \\citet{Nagamine04f}, which  reported clumpy distribution of DLA gas in haloes at $z=3$.  With the proper treatment of self-shielding as in the OTUV run, the DLAs have more extended distribution around star-forming galaxies on $10-30$\\,kpc scales.  We quantified the above features in the PDFs of impact parameters between DLAs/LLSs and the nearest galaxies. As expected, we find a much wider distribution for the OTUV run than for the Optically-Thin run.   \\item We examined the PDF of DLA/LLS cross sections as functions of halo mass, SFR, stellar mass, and metallicity (Figures~\\ref{fig:crossdistdla} and \\ref{fig:crossdistlls}). The median values of distributions are summarized in Table~\\ref{table:median}. About 30 per cent of DLAs are hosted by galaxies having $SFR = 1 - 20~\\Msunyr$, which is the typical SFR range for LBGs. More than half of DLAs ($\\sigmadla$, to be more precise) are hosted by the LBGs that are fainter than the current observational limit. About 80 per cent of total $\\sigmalls$ are hosted by haloes with ${\\rm SFR} \\lesssim 1~\\Msunyr$, hence most LLSs are associated with low-mass halos with faint LBGs below the current detection limit.  \\end{itemize}"
    },
    "1112/1112.2764_arXiv.txt": {
        "abstract": "Objects in the Kuiper belt are small and far away thus difficult to study in detail even with the best telescopes available at earth. For much of the early history of the Kuiper belt, studies of the compositions of these objects were relegated to collections of moderate quality spectral and photometric data that remained difficult to interpret.  Much early effort was put into simple correlations of surface colors and identifications of spectral features, but it was difficult to connect the observations to a larger understanding of the region. The last decade, however, has seen a blossoming in our understanding of the compositions of objects in the Kuiper belt. This blossoming is a product of the discoveries of larger -- and thus easier to study -- objects, continued dedication to the collection of a now quite large collection of high quality photometric and spectroscopic observations, and continued work at the laboratory and theoretical level. Today we now know of many processes which affect the surface compositions of objects in the Kuiper belt, including atmospheric loss, differentiation and cryovolcanism, radiation processing, the effects of giant impacts, and the early dynamical excitation of the Kuiper belt. We review the large quantity of data now available and attempt to build a comprehensive framework for understanding the surface compositions and their causes. In contrast to surface compositions, the bulk compositions of objects in the Kuiper belt remain poorly measured and even more poorly understood, but prospects for a deeper understanding of the formation of the the outer solar are even greater from this subject. ",
        "introduction": "The region of the solar system beyond the planets is populated by a vast swarm of objects in what is called the Kuiper belt. The objects in this swarm -- Kuiper belt objects (KBOs) -- formed on the outer edge of the planetary system where nebular densities were too low and accretional timescale too high to allow objects to accumulate into a single dominant planet, as happened throughout the planetary system. Instead, in the region beyond the planets, accretion proceeded up to objects at least as large as the 2700 km diameter of Triton but then stalled as dynamical excitation of the region caused collision velocities to be too high to allow continued accretion. Today, encounters between objects in the outer solar system lead to erosion to smaller diameters.  Much of we have learned from the Kuiper belt has come from study of the dynamics of these bodies \\citep[i.e.][]{2008ssbn.book..275M}. Understanding the dynamical state of a KBO is straightforward: after an object is discovered, a few observations spaced over several years are usually sufficient to define a precise orbit for the object \\citep{2000AJ....120.3323B}. Since their formation KBOs have been moved around by the gravitational tugs and perturbations of the giant planets, and their current locations preserve a record of these billions of years of interactions. Our current understanding of the rearrangement and migration of the giant planets is based largely on this record.  In principle, the objects in the Kuiper belt also contain a wealth of information on the chemical conditions in the earliest solar system. The materials in these objects have, for example, avoided much of the high temperature and pressure processing to which materials in the inner solar system have been subjected. In practice, however, understanding the detailed composition of KBOs in their native environment has proved difficult. Most objects are small, far away, and thus difficult to study even with the largest telescopes available. For the small number of objects large enough for detailed compositional study to be feasible, only the surface chemistry, which is clearly not representative of the entire body, is accessible to observation. For detailed understanding of ices and minerals incorporated into bodies in the early outer solar system, short-period comets, which are generally derived from the Kuiper belt, are much better targets for study.  For the first decade of the study of compositions of KBOs, most effort was put into simple cataloging and classifying the surfaces, with little insight into underlying causes. The past decade, however, has seen a maturation of both the data sets available and the physical and chemical insights into their meaning. The chemical composition of KBOs serves a tracer of the physical and chemical processes to which these bodies were subjected  at formation and in the time period since. Using studies of surface composition we have gained insight into atmospheric processes, differentiation and cryovolcanism, radiation processing, the effects of giant impacts, and the early dynamical excitation of the Kuiper belt.  In this review we attempt to formulate a coherent picture explaining the compositions of Kuiper belt objects. We show that KBO surface composition divides most strongly by size, with large, medium, and small objects having distinct surface types, and we propose physically and chemically plausible hypotheses for all of these surfaces. We attempt to formulate similar hypotheses for the bulk compositions of KBOs, but we find no satisfactory hypothesis that fits the current data. We end with some thoughts on important open questions that would best allow progress to be made in these areas. ",
        "conclusions": ""
    },
    "1112/1112.0761_arXiv.txt": {
        "abstract": "Dynamics and collapse of collisionless self-gravitating  systems  is described by the coupled collisionless Boltzmann and Poisson equations derived from $f(R)$ gravity in the weak field approximation. Specifically, we describe a system at equilibrium by a time-independent distribution function $f_0(x,v)$ and two potentials $\\Phi_0(x)$ and $\\Psi_0(x)$ solutions of the modified Poisson and collisionless Boltzmann equations. Considering a small perturbation from the  equilibrium and linearizing the field equations, it can be obtained a dispersion relation. A dispersion equation is achieved for neutral dust-particle systems  where a generalized Jeans wave-number is obtained.  This analysis gives rise to unstable modes not present in the standard Jeans analysis (derived assuming Newtonian gravity as weak filed limit of  $f(R)=R$). In this perspective, we discuss several self-gravitating astrophysical systems whose dynamics could be fully addressed in the framework of  $f(R)$-gravity. ",
        "introduction": "\\label{uno}  One of the fundamental goals of  modern cosmology is to probe  Einstein's General Relativity (GR) at any scale beyond the classical tests that confirmed such a theory in the weak field limit and at Solar System level. GR is assumed as the standard  theory of gravity  describing   astrophysical structures up to the whole observed Universe; however there are some inconsistencies at ultraviolet  scales (e.g. the  initial singularity, the quantum gravity issue) and infrared scales (e.g. cosmic acceleration,  concordance problem, flatness problem, galaxy rotation curves, large scale structure, massive stars formation) that strongly suggest that Einstein's approach should be revised or at least extended. Furthermore, astrophysical observations of the last decades suggest that new (dark)  ingredients are necessary to achieve a self-consistent cosmological model. In particular, the observations suggest the Hubble flow is currently accelerating, and the simplest way to explain the cosmic acceleration is to insert a cosmological constant ($\\Lambda$) in the Friedmann-Robertson-Walker Cosmology \\cite{Padmanabhan,Peebles,Sahni}, representing about $70\\%$ of the total amount of energy. On the other hand, the galaxy rotation curves and the large scale structure could be dynamically addressed by introducing huge amounts  of dark matter (about the 25$\\%$ of the total matter). Only  5$\\%$ of the cosmic budget is constituted by  standard matter as stars, neutrinos, radiation, heavy elements and free cosmological hydrogen ad helium. Alternative approaches to GR could be pursued with the aim to  explain  the observed acceleration and missing matter without introducing new ingredients  up to now not observed at fundamental scales. The  so-called $f(R)$-gravity is considered as a possible, straightforward mechanism to explain the cosmic acceleration without inserting  unknown elements as dark energy and dark matter, but extending the geometric part of the field equations by relaxing the strict hypothesis that the gravitational action has to be restricted to $f(R)=R$ as in the Hilbert-Einstein one \\cite{physrep,Capozziello2,felix1,Nojiri, review, Capozziello10}.  These theories have been investigated both at cosmological scales and in the weak field limit \\cite{Capozziello3,f(R)test,Cognola1, Cognola2,f(R)test2}. It has been shown that a late accelerating behavior can be easily recovered \\cite{saffari} and it can be related to an early inflationary expansion \\cite{Nojiri2}. Furthermore,  modifying the gravity action by assuming   non-linear Lagrangians, one obtain corrections to the gravitational potential which can be useful for astrophysical phenomenology  at galactic scales. In particular, without the introduction of dark matter, the rotation curves of spiral galaxies and the haloes of galactic clusters  can be dynamically  addressed \\cite{mnras,salzano,Boehmer, Mendoza, Sobouti}. Several of these extended models  reproduce Solar System tests so they are not in conflict with GR  experimental results  but simply extend them \\cite{Capozziello8,Capozziello9, pogosian}.  It is important to stress that $f(R)$ gravity  has interesting applications also in stellar astrophysics and could contribute to solve several puzzles related to observed peculiar objects (e.g. magnetars, stars in the instability strips, protostars, etc. \\cite{Cooney, Hu}), structure and star formation \\cite{poly, Chang}.  Here we analyze the Jeans instability for self-gravitating systems in $f(R)$-gravity coupled with perfect-fluid matter. The aim is to show that several self-gravitating systems, in particular those involved in star formation (e.g. large molecular clouds or Bok globules), can be exactly addressed in this framework by considering the corrections  to the Newtonian potential coming out from $f(R)$-gravity. This fact could constitute a remarkable signature to retain or rule out these theories at astrophysical level.  The paper is organized as follows. In Section \\ref{due} the classical theory  of gravitational collapse for dust-dominated systems is summarized.  In Section \\ref{tre}, we  discuss the weak field   limit of $f(R)$-gravity obtaining corrections to the standard Newtonian potential that can be figured out as two Newtonian potentials contributing to the dynamics. In Section \\ref{quattro} we recover the dispersion relation and Jeans mass limit while, in  Section \\ref{cinque},  some self-gravitating dust  system are considered in this approach. The difference between GR and $f(R)$-gravity are put in evidence, in particular the Jeans mass profiles with respect to the temperature. We report a catalogue of observed molecular clouds in order to compare the classical Jeans mass  with the  $f(R)$ one. Finally, in Section \\ref{sei}, results are discussed. ",
        "conclusions": "\\label{sei}  $f(R)$-gravity is an approach aimed to address some shortcomings of modern cosmology just assuming extensions of GR without invoking the presence of dark ingredients. In other words, dark energy and dark matter could be effects related to curvature further degrees of freedom instead of new fundamental particles.  Here we have analyzed the Jeans instability mechanism, adopted for star formation,  considering the Newtonian approximation of  $f(R)$-gravity. The related Boltzmann-Vlasov system leads to  modified Poisson equations depending on the $f(R)$-model. In particular, considering  Eqs.\\eqref{EQP} and \\eqref{EQP2},  it is possible to get a new dispersion relation \\eqref{eq:Results} where instability criterion results modified (see also \\cite{poly}). The leading parameter is $\\alpha$, i.e.  the second derivative of the specific $f(R)$-model. Standard Newtonian Jeans instability is immediately recovered for $\\alpha=0$ corresponding to the Hilbert-Einstein Lagrangian of GR. In Fig. \\ref{Fig:Jeans}, dispersion relations for Newtonian and a specific $f(R)$-model are numerically compared. The modified characteristic length van be given in terms of the classical one.  Both in the classical  and  in $f(R)$ analysis, the system damps the perturbation. This damping is not associated to the collisions because we neglect them in our treatment, but it is linked to the so called Landau damping \\cite{BT}.  A new condition for the gravitational instability is derived, showing  unstable modes with faster growth rates. Finally we can observe the instability decrease in $f(R)$-gravity: such decrease is related to a larger Jeans length and then to a lower Jeans mass. We have also compared the behavior with the temperature of the Jeans mass for various types of interstellar molecular clouds (Fig. \\ref{Fig:mass}). In Tables \\ref{table1} and \\ref{table2} we show the results given by this new limit of the Jeans mass for a sample of giant molecular clouds. In our model the limit (in unit of mass) to start the collapse of an interstellar cloud is lower than the classical one advantaging the structure formation.  Real solutions for the Jean mass can be achieved only for $\\alpha<0$  and this result is in agreement with cosmology \\cite{Capozziello10}. In particular,  the condition $\\alpha<0$ is essentials to have a  well-formulated and well-posed Cauchy problem in $f(R)$-gravity  \\cite{Capozziello10}. Finally, it is worth noticing that  the Newtonian value is an upper limit for the Jean mass coinciding with $f(R)=R$.  This work is intended to indicate the possibility to deal with ISM collapsing clouds  under different assumptions about gravity. It is important to stress that we  fully recover the standard  collapse mechanisms but we could also describe proto-stellar systems that escape the standard collapse model. On the other hand,  this is the first step to study star formation in alternative theories of gravity (see also \\cite{Cooney, Hu, poly, Chang,agon}). From an observational point of view, reliable  constraints can be achieved from  a careful analysis of the proto-stellar phase taking into account magnetic fields, turbulence and collisions. Finally, addressing  stellar systems by this approach could be an extremely important to test observationally $f(R)$-gravity.  Moreover, the approach developed in this work admits direct generalizations for other modified gravities, like non-local gravity, modified Gauss-Bonnet theory, string-inspired gravity, etc. In these cases, the constrained Poisson equation may be even more complicated due to the presence of extra scalar(s) in non-local or string-inspired gravity. Developing further this approach gives, in general, the possibility to confront the observable  dynamics of astrophysical objects (like stars) with predictions of alternative gravities."
    },
    "1112/1112.2622_arXiv.txt": {
        "abstract": "We study the gravitational collapse of a magnetized neutron star using a novel numerical approach able to capture both the dynamics of the star and the behavior of the surrounding plasma. In this approach, a fully general relativistic magnetohydrodynamics implementation models the collapse of the star and provides appropriate boundary conditions to a force-free model which describes the stellar exterior. We validate this strategy by comparing with known results for the rotating monopole and aligned rotator solutions and then apply it to study both rotating and non-rotating stellar collapse scenarios, and contrast the behavior with what is obtained when employing the electrovacuum approximation outside the star. The non-rotating electrovacuum collapse is shown to agree qualitatively with a Newtonian model of the electromagnetic field outside a collapsing star.  We illustrate and discuss a fundamental difference between the force-free and electrovacuum solutions, involving the appearance of large zones of electric-dominated field in the vacuum case.  This provides a clear demonstration of how dissipative singularities appear generically in the non-linear time-evolution of force-free fluids. In both the rotating and non-rotating cases, our simulations indicate that the collapse induces a strong electromagnetic transient, which leaves behind an uncharged, unmagnetized Kerr black hole.  In the case of sub-millisecond rotation, the magnetic field experiences strong winding and the transient carries much more energy.  This result has important implications for models of gamma-ray bursts.  Even when the neutron star is surrounded by an accretion torus (as in binary merger and collapsar scenarios), a magnetosphere may emerge through a dynamo process operating in a surface shear layer.  When this rapidly rotating magnetar collapses to a black hole, the electromagnetic energy released can compete with the later output in a Blandford-Znajek jet. Much less electromagnetic energy is released by a massive magnetar that is (initially) gravitationally stable: its rotational energy is dissipated mainly by internal torques.  A distinct plasmoid structure is seen in our non-rotating simulations, which will generate a radio transient with subluminal expansion, and greater synchrotron efficiency than is expected in shock models. Closely related phenomena appear to be at work in the giant flares of Galactic magnetars. ",
        "introduction": "Understanding the gravitational collapse of a massive neutron star is of central importance for its connection to some of the most energetic astrophysical phenomena. Such an event may take place within a core-collapse supernova \\cite{woosley:1993}, or in the late stage of a binary neutron star merger \\cite{eichler:1989,narayan:1992}, and is widely believed to power some types of gamma-ray bursts.  In recent years it has been realized that newly formed stellar mass black holes may be prodigious sources of electromagnetic energy, in addition to driving strong kinetic outflows~\\cite{woosley:1993,1994MNRAS.270..480T,meszaros:1997,mckinney:2006}.  We are faced with the exciting possibility of probing the most extreme forms of gravitational collapse using coordinated measurements of electromagnetic transients and gravitational waves.  Refining models that make observational predictions for joint gravitational and electromagnetic radiation is critical in order to establish efficient observation campaigns for both traditional astronomers and gravitational-wave astronomers. It is also a key step toward predicting the delay between peak electromagnetic and gravitational wave emission, the electromagnetic emission pattern (i.e. the beaming angle) and the electromagnetic and gravitational wave spectrum.  Our focus here is on the evolution of the neutron star's magnetic field during its collapse to a black hole.  We employ fully self-consistent relativistic calculations that follow the dense stellar material as well as the strong electromagnetic and gravitational fields.  A particularly intense electromagnetic transient is generated if the initial magnetic field is very strong ($\\sim 10^{15}-10^{16}$ G), that is, if the star is a magnetar. Rapid rotation will enhance the energy of the transient, to a degree that can only be derived by a full time-evolution of the electromagnetic and gravitational fields.  Rapid rotation also provides a context for generating strong magnetic fields:  when the neutron star is accreting, the shear layer at its surface is a promising site for dynamo action.  A distinct magnetosphere will emerge and hold off the accretion flow if even $\\sim 10^{-4}-10^{-3}$ of the energy that is dissipated in the shear layer is converted to a poloidal magnetic field.  Understanding the evolution of the magnetic field around an {\\it isolated} star is therefore potentially of key relevance for the collapsar and binary merger scenarios of gamma-ray bursts.  Recent success in studying the behavior of plasmas around magnetized, spinning, stable neutron stars within flat spacetime has been presented in Refs.~\\cite{2006ApJ...643.1139C,Spitkovsky:2006np,McKinney:2006sc}. To study the collapse problem, one needs general relativity -- to account for the role of spacetime curvature; and general relativistic magnetohydrodynamics (GRMHD) -- to determine the internal evolution of the star. The electromagnetic (EM) phenomena outside the star can be approached in a variety of ways:  through a full GRMHD calculation, which generally is very expensive given the low matter density; or more approximately using vacuum EM and force-free equations.  In this paper, we compare all three approaches (restricting to ideal MHD in the case of GRMHD).  Within general relativity, studies paying attention to neutron star collapse have been presented in the context of isolated stellar collapse (e.g.~\\cite{Stephens:2006cn,Stephens:2007kb,O'Connor:2010tk,Reisswig:2010cd,Giacomazzo:2011cv}) and binary neutron star mergers (e.g.~\\cite{Anderson:2007kz,Baiotti:2008ra,Liu:2008xy,Anderson:2008zp,Hotokezaka:2011dh,Shibata:2011fj}) with different degrees of realism.  The role of electromagnetic fields are typically examined within general relativistic {\\em ideal} MHD and studies have been presented for single~\\cite{Duez:2005cj,Shibata:2005mz,Duez:2006qe,Kiuchi:2008ss,Liebling:2010bn,Liebling:2010qv} and binary star systems~\\cite{Anderson:2008zp,Liu:2008xy,Giacomazzo:2009mp}. Such simulations have illustrated important details of the extreme dynamical behavior induced by the system which could trigger tremendously energetic phenomena.    Newtonian and general relativistic simulations have shed light on how a binary neutron star merger can amplify pulsar-strength magnetic fields by several orders of magnitude~\\cite{price:2006,Anderson:2008zp,2010A&A...515A..30O,Giacomazzo:2009mp}. The resulting hypermassive star is generically unstable to black hole formation, which opens up the tantalizing possibility that the increasing rotation rate and magnetic field strength would drive an intense electromagnetic outflow. To study such a scenario, our goal here is to combine the GRMHD approach with a suitable description of the magnetically-dominated region outside a collapsing star, by applying the force-free approximation. Our approach is related to that of Ref.~\\cite{Baumgarte:2002b}, which matched a simplified, analytic, relativistic solution for the interior of a collapsing star (dust-ball) with a numerical solution of the coupled Einstein-Maxwell equations in its exterior. Boundary conditions at the star's surface are provided by the analytic, ideal MHD solution.  We go beyond this simplified scenario in three significant ways. First, our stars are evolved through collapse consistently by evolving the ideal GRMHD equations, and we can therefore consider in principle any kind of compact star. In particular, we have studied both rigidly rotating and non-rotating stars, whereas only the non-rotating case was considered in~\\cite{Baumgarte:2002b}. Second, the magnetosphere is described within the force-free approximation, which, as argued in Ref.~\\cite{Goldreich:1969sb}, is a much more realistic model that, for instance, has been instrumental in understanding pulsar spin down~\\cite{Spitkovsky:2006np}. Third and last, our matching of the exterior solution with the interior one is dynamical. Thus the force-free solution can adapt to time-dependent fields sourced by the star.   Our approach thus allows us to examine many interesting scenarios and we apply it here to study the behavior of collapsing, magnetized, compact stars (either rotating or not), the behavior of surrounding plasma (as described within the force-free approximation) and possible electromagnetic radiation induced by the system. We compare our results with recent estimates in Ref.~\\cite{Lyutikov:2011tq} for the non-rotating case which predict that the collapse process is smooth and that the magnetic field remains anchored to the star as a black hole forms, leaving a final black hole with a split-monopole field configuration. Our work, which follows the dynamics of the system, indicates that in both cases the stars radiate significant electromagnetic energy in which reconnection plays a crucial role, and that this radiation ceases shortly after the formation of a black hole. With force-free, such a black hole loses its electromagnetic hair within a dynamical time scale.  We describe our hybrid approximation and the evolution equations in Sec.~\\ref{sec:approach}, followed by a summary of the numerical techniques in Sec.~\\ref{sec:numerical_techniques}. Our choice of initial data is described in Sec.~\\ref{sec:id}, and several tests of the hybrid approach are presented in Sec.~\\ref{sec:testing}. The new results for collapsing, magnetized stars are explained in detail in Sec.~\\ref{sec:physics}, while the astrophysical consequences are described in Secs.~\\ref{sec:astrophysicsI} and ~\\ref{sec:magnetar}. We conclude in Sec.~\\ref{sec:concluding_remarks} with some final comments. ",
        "conclusions": "\\label{sec:concluding_remarks} Understanding the global behavior of strongly gravitating, dynamical systems, containing dense matter coupled to ultra-strong magnetic fields, is of key importance for a thorough understanding of possible signals from them. In this work, we have presented a new approach to this end by combining the ideal MHD and force-free approximations in a suitable manner within general relativity.  We expect that this hybrid scheme represents real progress towards greater realism. The stellar interior utilizes ideal MHD to faithfully model the neutron star without the disadvantages incumbent in the low density exterior. Likewise, the exterior solution uses the force-free approach and so captures the dynamics of the tenuous plasma. The entire domain is described by a fully nonlinear and fully dynamic general relativity solution necessary for strong-field gravity. Indeed, a key aspect of this approach is its generality. It does not require a prescribed stationary stellar boundary, and can be applied, for example, to dynamical systems such as collapsing stars and non-vacuum compact binaries.   We have exploited this approach to study stellar collapse in both rotating and non-rotating cases. Our studies reveal a rich phenomenology in the magnetosphere as the collapse proceeds. In particular, magnetic reconnection plays an important role by inducing strong electromagnetic emission as well as the infall of electromagnetic energy into the black hole, which in a short time loses all its `hair'. When the star starts off rapidly rotating, the energy of the magnetospheric plasma grows significantly during the collapse. As the star rotates faster, its magnetic field lines do not have time to re-adjust to the increased rotation, and are strongly wound up out to the initial light cylinder. It is worth emphasizing that this conversion of dynamical energy into electromagnetic energy does not depend on resistive effects.  Two issues of principle have arisen in performing these calculations.  First, we have shown conclusively that the force-free approximation to the evolution of a dilute, relativistic plasma inevitably leads to singularities.  These singularities are avoided in our calculations by continuously pruning the electric field.  In the simulations that we have run, this procedure appears to cause limited energy dissipation, but its necessity should be kept in mind.  The second issue of principle regards the dependence of the reconnection geometry on the resistivity model.  Fluid pressure is responsible for slowing down the rate of reconnection unless an x-point geometry is able to form. By neglecting fluid pressure, the force-free approximation clearly facilitates the formation of x-points. In spite of this, we observe very limited numerical dissipation in current sheets.  Although the plasma that is represented by our model magnetospheres is, in reality, strongly collisional, it should be kept in mind that fast x-point reconnection still occurs in collisional plasmas in the presence of hydromagnetic turbulence -- as has been demonstrated so far in weakly magnetized plasmas \\cite{kowal:2009}.  Some further exploration of the reconnection geometry is possible in resistive MHD calculations of stellar collapse by varying the spatial dependence of the resistivity.   We have also discussed how the electromagnetic outbursts from collapsing magnetars may have interesting observational effects. In particular, our results are relevant to binary neutron star merger scenarios, in which a hypermassive neutron star forms and collapses to a black hole.  Collapse to a black hole can happen either promptly or after many dynamical times, depending on the masses involved and the equation of state describing the stars (see, for instance,~\\cite{Hotokezaka:2011dh}). Although Kelvin-Helmholtz and magnetorotational instabilities will create ultrastrong magnetic fields ($\\agt 10^{14-16}$ G), we have suggested that the global field (in particular, the amount of magnetic flux threading the merger remnant and eventual black hole) could depend strongly on the lifetime of the remnant.  The formation of a distinct magnetosphere would require tapping only $\\sim 10^{-3}-10^{-4}$ of the energy dissipated in shear layer at the remnant surface.  As such a star collapsed, the magnetosphere would qualitatively follow the behavior outlined here.   The electromagnetic output of the collapse could compete with the later emission from a Blandford-Znajek jet emanating from the black hole horizon, and source a powerful electromagnetic counterpart to the gravity wave signal (e.g.~\\cite{Metzger:2011bv}). In addition, the magnetic field dynamics that is revealed in our simulations has a number of interesting implications for gamma-ray bursts and magnetar flares, as discussed in Secs.~\\ref{sec:astrophysicsI} and~\\ref{sec:magnetar}.  Beyond the work analyzed here, our approach is readily applicable to other relevant systems and will be applied, in particular, to study binary neutron star systems~\\cite{PalenzuelaBNSpassive}, and black hole-neutron star systems (see e.g.~\\cite{McWilliams:2011zi}). It is important to stress, however, that our approach is not free of ambiguities, in particular with respect to how and where the matching between the force-free and ideal MHD regions is implemented. To the extent possible, we have tested the robustness of our results versus known solutions, which give us confidence in this approach.  \\noindent{\\bf{\\em Acknowledgments:}} It is a pleasure to thank  P. Goldreich for insights and discussions during the course of this work.  We also thank E. Berti, P. Cerda-Duran, J. McKinney, R. O'Shaughnessy, E. Ramirez-Ruiz, E. Quataert, A. Spitkovsky and D. Tsang as well as our long time collaborators M. Anderson, E. Hirschmann, P. Motl, D. Neilsen and O. Reula, for discussions and comments. This work was supported by the NSF (PHY-0969827 to Long Island University) and NSERC through Discovery Grants (to CT and LL). Research at Perimeter Institute is supported through Industry Canada and by the Province of Ontario through the Ministry of Research \\& Innovation. Computations were performed at Teragrid and Scinet.  \\appendix"
    },
    "1112/1112.3641_arXiv.txt": {
        "abstract": "Surveying eighteen \\twc-bright galaxies from the \\atlas early-type galaxy sample with the Institut de Radio Astronomie Millim\\'etrique (IRAM) 30m telescope, we detect \\tho and \\tht in all eighteen galaxies, \\hcno in 12/18 and \\hcopo in 10/18. We find that the line ratios \\two/\\tho and \\two/\\hcno are clearly correlated with several galaxy properties: total stellar mass, luminosity-weighted mean stellar age, molecular to atomic gas ratio, dust temperature and dust morphology. We suggest that these correlations are primarily governed by the optical depth in the \\twc lines; interacting, accreting and/or starbursting early-type galaxies have more optically thin molecular gas while those with settled dust and gas discs host optically thick molecular gas. The ranges of the integrated line intensity ratios generally overlap with those of spirals, although we note some outliers in the \\two/\\tho, \\twt/\\tht and HCN/\\hcopo ratios.  In particular, three galaxies are found to have very low \\two/\\tho and \\twt/\\tht ratios. Such low ratios may signal particularly stable molecular gas which creates stars less efficiently than `normal' (i.e. below Schmidt-Kennicutt prediction), consistent with the low dust temperatures seen in these galaxies. ",
        "introduction": "Molecular gas is an essential ingredient for star formation found in many, but not all, galaxies. Early-type galaxies (ellipticals and lenticulars) were classically thought to completely lack molecular gas and to be passively-evolving `red and dead' galaxies. However, we have long known that not all early-types are empty of cold gas; molecular gas was first detected in early-type galaxies by \\citet{wiklind86} and \\citet{phillips87}.  Shortly after these first detections, surveys detected 10-20 galaxies, but were biased towards early-types with particular properties, such as those bright in the far infra-red \\citep{wiklind89,sage89} or with optically obscuring dust \\citep{wang92}. These selection effects made it easy to dismiss any early-type galaxies with molecular gas as peculiar systems. The next generation of surveys increased numbers (30-50 galaxies) and were less biased, if still not complete \\citep{welch03,sage07,combes07}. In \\citet[][hereafter Paper IV]{young11}, we have recently completed an extensive (259 galaxies) molecular gas detection campaign for the volume-limited \\atlas sample of early-type galaxies \\citep[][hereafter Paper I]{cappellari11a}. We find a 22\\% detection rate, down to a typical detection threshold of $6 \\times 10^7$ M$_{\\sun}$ of H$_{2}$, robustly establishing that many early-type galaxies host a substantial amount of molecular gas.  The majority of early-type galaxies with molecular gas are obviously star-forming, based on ultraviolet (UV), optical or infrared data. Indeed, the detection rate for star formation seen in UV nearly matches the molecular detection rate ($\\approx30\\%$, \\citealp{kaviraj07}). According to the current observations, the star formation efficiencies of early-type galaxies broadly follow the Kennicutt-Schmidt law \\citep{shapiro10, crocker11}. A subset shows no obvious sign of ongoing star formation \\citep{crocker08,crocker11}, although these determinations are difficult due to their very low specific star formation rates. Recently, \\citet{saintonge11b} have determined that star formation efficiencies are reduced for more massive, more concentrated and higher stellar surface density galaxies, all properties that positively correlate with early-type morphology.  In a spiral galaxy, the presence of a stellar disc renders the gas disc more locally unstable to axisymmetric perturbations, likely boosting its star formation efficiency \\citep[e.g.][]{jog84}. In \\citet[][hereafter Paper II]{krajnovic11} and \\citet[][hereafter Paper III]{emsellem11}, we found that the vast majority of early-type galaxies (the fast rotators) are consistent with being a family of disc-like systems resembling spiral galaxies with the gas and dust removed \\citep[][hereafter Paper VII]{cappellari11b}. However, the fast rotators are generally characterized by larger spheroids than spiral galaxies. This increase in the depth of the potential well is expected to make their gas discs more stable against fragmentation \\citep{kawata07}.  Indeed, simulations with a centrally-concentrated stellar mass distribution (as found in spheroids) and no stellar disc show that the cool gas is more stable than in spiral galaxies \\citep{martig09}. This stability (termed `morphological quenching')  should lower the efficiency of star formation and produce a cool interstellar medium (ISM) with properties (velocity dispersion, density distribution, etc.) different from those of galaxies with stellar discs. In this paper, we present the first major attempt at constraining the empirical properties of the molecular gas in early-type galaxies, especially looking for any divergence from the properties found for spiral galaxies.  The surveys mentioned above have used the bright \\two emission line to measure the total molecular content of early-type galaxies, but little work has been done to constrain the molecular gas properties using other species and transitions.  Several other molecular species are bright enough to measure, including \\thc, HCN and HCO$^{+}$. These species have been widely observed in starburst and Seyfert galaxies and also in some local spiral galaxies.  The observed \\twc/\\thc ratio is widely used to indicate the average optical depth of the molecular gas\\footnote{Described by: $(\\frac{^{12}\\mathrm{CO}}{^{13}\\mathrm{CO}})_{\\mathrm {obs}}=\\frac{e^{-\\tau(12)}}{e^{-\\tau(13)}}\\frac{^{12}\\mathrm{CO}}{^{13}\\mathrm{CO}}$}, although it may also be influenced by chemical processes. The \\twc isotope is far more abundant and becomes optically thick at lower H$_{2}$ column densities. Its use as a measure of the total molecular hydrogen content (via one of the \\twc-to-H$_{2}$ conversion factors X$_{\\mathrm {CO}}$ or $\\alpha_{\\mathrm {CO}}$) relies on the assumption that it is optically thick, and essentially counts the number of virialized molecular clouds of a similar temperature and density \\citep[e.g.][]{young91}. The less abundant \\thc isotope is optically thin until higher column densities and thus the \\twc/\\thc ratio (often denoted $\\mathfrak{R}$) should reflect differences in average optical depth in \\twc, either within the molecular clouds themselves or because of the additional contribution of a diffuse molecular component. Variations in the $^{12}$C to $^{13}$C abundance ratio and \\twc$\\leftrightarrow$\\thc fractionation (due to charge-ion reactions or selective photodissociation) must also be considered, although they do not seem to drive this ratio in spiral galaxies \\citep{paglione01}.  The \\hcno and \\hcopo lines have higher critical densities (approximately $10^{6}$ and $10^{5}$ cm$^{-3}$, respectively) than CO(1-0) ($n_{\\mathrm crit} \\approx 10^{3}$)\\footnote{Note that effective critical densities, which take into account radiative trapping, are about an order of magnitude lower \\citep[e.g.][]{scoville74}}. The critical density is simply the density at which collisions are more frequent than radiative decays, not a strict limit with no emission at lower densities. Thus depending on the density distribution of the molecular gas, the typical gas density probed by the \\hcno and \\hcopo transitions will vary from system to system, although in all cases it will be higher than the densities probed by CO(1-0).  Higher HCN/\\twc and \\hcop/\\twc ratios in luminous infrared galaxies (LIRGs) and ultra-luminous infrared galaxies (ULIRGs) are thus taken to indicate higher dense gas fractions in these galaxies \\citep{gao04a, gracia-carpio06}. The more moderate increase of \\hcop/\\twc compared to HCN/\\twc in these systems may be explained by the order of magnitude difference in the critical densities of \\hcno and \\hcopo \\citep{juneau09}. In both the Milky Way and M31, the HCN/\\twc ratio declines with radius, signaling the decline of the dense gas content \\citep{helfer97, brouillet05}. Indeed, \\citet{helfer97} find that the \\hcno/\\two ratio correlates with the hydrostatic pressure in the disc: \\hcno/\\two $\\propto P^{\\,0.19\\pm0.04}$.  Molecular chemistry can also influence the HCN and \\hcop emission seen in galaxies. X-ray dominated regions (XDRs) around active galactic nuclei (AGN) may enhance the HCN abundance relative to CO \\citep[e.g.][]{lepp96}. This effect has been observed in Seyfert galaxies and led to the recommendation to prefer \\hcop as a dense gas tracer \\citep{gracia-carpio06, krips07}. However, \\citet{papadopoulos07} warn about possible effects of free electrons in cosmic-ray dominated regions (CRDRs) and XDRs, or highly turbulent molecular clouds destroying \\hcop. Nevertheless, most of these chemical effects rely on the conditions found in starbursts or AGN, and hence HCN and \\hcop are likely to remain good tracers of dense gas in more quiescent regimes.  Until recently, the only early-type galaxy to be studied in molecular species other than \\twc was Centaurus A (Cen A), with its striking lane of dust and gas. The $\\approx5\\times10^{8}$ M$_{\\sun}$ of molecular gas in its disc extends to about 2.6~kpc \\citep{phillips87,eckart90a}, while 4$\\times10^{8}$ M$_{\\sun}$ of atomic gas is seen in a warped disc out to about 6~kpc \\citep{vangorkom90,struve10}. While often considered a galaxy completely in its own class, its molecular and atomic gas content is not very different to that of some of our sample galaxies. In terms of molecular emission line studies, \\citet{wild97} found a \\two/\\tho ratio of 14 towards the center of Cen A and 11 towards two positions in its disc, similar values as found in spirals. As for HCN, \\citet{wild00} find a higher ratio of \\hcno/\\two in the center (0.067) than in offset positions along the disc (0.02-0.04).  In \\citet{krips10}, the \\tho,\\tht, \\hcno and \\hcopo emission lines of four early-type galaxies were measured as a pilot project for this work. Two of these galaxies have particularly low \\two/\\tho ratios of around 3 and two have more typical ratios of around 10. The \\hcno/\\tho and \\hcno/\\two ratios are in the range observed for spiral and Seyfert galaxies, but lower than for starbursts. Most curiously, none were detected in \\hcopo, despite 3/4 being detected in \\hcno, which usually has a similar integrated intensity. In one case, the \\hcno/\\hcopo ratio is constrained to be larger than 2, indicating some significant difference in the chemistry or physical properties of this galaxy's molecular gas compared to those of spiral, Seyfert and starburst galaxies.  In this paper, we extend the original sample of 4 early-type galaxies from \\citet{krips10} to a sample of 18 galaxies from the \\atlas sample. Section~2 describes the sample selection and Section~3 describes our observations and data reduction. We compare the profiles of the different molecular lines within each galaxy and discuss the derived line ratios in Section~4.  In Section~4, we also investigate the variations of the molecular line ratios with other galaxy properties. Our conclusions are presented in Section~5.    \\begin{figure} \\begin{center} \\includegraphics[height=4.5cm]{figures/samplebias.ps}  \\includegraphics[height=4.5cm]{figures/samplebias2.ps} \\caption{Distributions of the 56 galaxies selected for follow-up (solid lines) out of the 59 \\two detections from \\atlas (dashed lines). Top panel: our subsample is clearly biased towards the highest molecular gas masses and mass fractions and slightly biased towards warmer dust temperatures, as measured by the Infrared Astronomical Satellite (IRAS) 60 to 100 \\micron~flux ratio ($f_{60}/f_{100}$). The bottom panel shows no bias in absolute $K$-band magnitude, local environment ($\\Sigma_{3}$), or specific stellar angular momentum ($\\lambda_{\\mathrm{R}}/\\sqrt{\\epsilon_{\\mathrm e}}$).} \\label{fig:sample} \\end{center} \\end{figure} ",
        "conclusions": "Using the IRAM 30m telescope, we have surveyed the \\thc, HCN and \\hcop emission of the 18 galaxies with the strongest \\twc emission in the \\atlas sample of nearby early-type galaxies. We detect all 18 galaxies in both \\tho and \\tht transitions, 12/18 in \\hcno and 10/18 in \\hcopo.  Fitting the \\two lines with a double-peak or Gaussian profile and evaluating the reduced $\\chi^{2}$ value, we find that one third of the galaxies are best fit by a Gaussian, one third with a clear double-peak profile and the remaining one-third with a very flat-topped `double-peak' profile. Two galaxies (NGC~1266 and NGC~5866) have clearly different profiles in the various molecular lines, indicating spatial variations of the  molecular gas conditions within the galaxies. Discrepant profiles in a few other galaxies are consistent with pointing offsets. We hypothesize that at higher signal-to-noise ratios, the majority of the line profiles would be found to be intrinsically inconsistent due to changing molecular gas conditions. Higher signal-to-noise single-dish data or interferometric maps in the various lines will allow us to determine if this is indeed the case.  The molecular gas line ratios of early-type galaxies are generally consistent with the ranges observed in spiral galaxies, although we also identify some outliers. In terms of the HCN/\\two ratio, often used as a dense gas tracer, we extend the work of \\citet{gao04b} to lower FIR luminosities and find an increased scatter in the ratio, compared to late-type galaxies of similar FIR-luminosities. Because of the important effects of the optical depth of \\two seen in the \\Rone ratio, we suggest that optical depth may be undermining the use of HCN/\\two as a dense gas tracer and instead recommend the use of HCN/\\tho. Unfortunately, few spirals have literature HCN/\\tho ratios. Determining whether the early-types have significantly different dense gas fractions thus requires a larger sample of spirals to be observed in both \\hcno and \\tho.  We compare the molecular gas ratios \\Rone, \\two/HCN and \\two/\\hcop to several ISM and stellar galaxy properties. \\Rone is found to correlate with the molecular-to-atomic gas ratio, dust temperature, dust morphology, absolute $K$-band magnitude and SSP age. \\two/HCN similarly correlates with the molecular-to-atomic gas ratio, dust temperature, absolute $K$-band magnitude and SSP age. The persistence of most \\Rone correlations with \\two/HCN rules out abundance variations of \\thc as the driving factor of the relations. No correlation with any of these properties is seen in the \\tho/HCN ratio, further ruling out a dominant contribution from ion-exchange fractionation towards \\thc at low temperatures.  With these options removed, a change in optical depth of the \\two line is the best explanation for the correlations we see. Based on the lower molecular-to-atomic gas ratios, unsettled dust and occasionally high dust temperatures, we find that high \\Rone and \\two/HCN ratios in early-type galaxies are linked to recent or ongoing interactions and/or starbursts. Both of these processes can naturally account for optically thinner \\two line emission, as interactions are known to be more turbulent and star formation feedback can also produce more turbulence or heat up the molecular gas. Early-type galaxies with settled gas and dust discs tend to have less optically thin molecular gas, along with lower dust temperatures and higher molecular to atomic gas mass ratios. The correlations with $K$-band luminosity and SSP age can plausibly be explained by sample selection effects, further work including less massive early-type galaxies with settled gas and dust discs is required before conclusions are reached.  A set of outliers with low \\Rone and \\Rtwo, high HCN/\\two and only upper limits to their HCN/\\hcop ratios are identified. These galaxies may have particularly stable molecular gas which remains very optically thick in \\twc  or have their molecular gas confined to a thin layer by hot gas pressure. NGC~1266, a galaxy with a molecular outflow, also stands apart from both spirals and the rest of the early-type sample. Its HCN/\\tho and \\tho/\\hcop ratios suggest that it has a very high dense gas fraction, but its high \\Rone and \\Rtwo values simultaneously suggest that much of its \\twc emission is optically thin. This combination of ratios is seen in some other extremely active starburst or Seyfert galaxies."
    },
    "1112/1112.6252_arXiv.txt": {
        "abstract": "In a recent article by Benson et al., 2011, the authors show the latest measurements from the South Pole Telescope (SPT) Sunyaev Zel'dovich (SZ) cluster survey to better constrain some cosmological parameters. In particular, the authors found that adding the SPT cluster data significantly improves the constraints on equation of state of dark energy, $w$, beyond those found when using measurements of the CMB, supernovae, BAO and the Hubble constant. The main aim of the present research note is to give a further quantitative estimation of the above better constraints, through the computation of the Figure of Merit (FoM) applied to $\\Omega_m$ and $w$ plots for the $68\\%$ and $95\\%$ confidence regions. This allows a better evaluation and a better comparison of the continuous improvements on the cosmological constraints, obtained using new different cosmological probes and different surveys. ",
        "introduction": "We recall here some brief considerations of the latest SPT survey, mainly devoted to better constrain some cosmological parameters and, in particular, the equation of state of dark energy, as reported in the reference article by Benson et al., 2011. Large scale structures in the Universe, like clusters of galaxies, give important information on different cosmological parameters, in particular, the matter density, the amplitude of the matter power spectrum, and the dark energy equation of state, mainly through its effect on the growth of structures. The cluster abundance measurements are an important tool for testing the standard dark energy model, because they are affected by dark energy properties in a different way than distance-redshift based tests, such as from type Ia supernovae and baryon acoustic oscillations (Benson et al., 2011). Recently there has been significant theoretical and experimental progress in the use of galaxies clusters as efficient cosmological probes. In particular, measurements of the cluster abundance using optical, X-ray, and SZ selection methods have been used to obtain stringent constraints on cosmology and dark energy parameters Currently, the most precise dark energy constraints from galaxies clusters are obtained from X-ray selected samples which use the X-ray emission from the hot intra-cluster gas as an index of the total mass in the cluster. We recall here that hot intra-cluster gas also causes a spectral distortion in the cosmic microwave background (CMB), a phenomenon known as the Sunyaev-Zel'dovich (SZ) effect (Sunyaev and Zel'dovich, 1972). The surface brightness of the SZ effect is redshift-independent and largest at mm-wavelengths. Thus, a mm-wavelength SZ survey with high angular resolution is expected to provide clean, mass-limited catalogs out to high redshift, probing the regime where the cluster abundance is most sensitive to dark energy's effect on the growth rate of structure (Carlstrom et al. 2002). Recently, the first SZ cluster catalogs from three surveys have been released: the South Pole Telescope (SPT, Staniszewski et al. 2009; Vanderlinde et al. 2010; Williamson et al. 2011), the Atacama Cosmology Telescope (ACT, Marriage et al. 2011, Sehgal et al., 2011), and the Planck satellite (Planck Collaboration et al. 2011). We recall that the South Pole Telescope is a 10-meter diameter, mm-wavelength telescope, designed to conduct a large-area survey with low noise and $\\sim 1$ arcminute angular resolution and it is sensitive in three bands, at 95, 150, and 220 GHz. The primary goal of the SPT survey is to search for clusters of galaxies via the SZ effect in a 2500 $deg^2$ survey that was completed in November 2011. In the article by Benson, et al., 2011, the authors use the SPT-SZ survey with the cluster sample from Vanderlinde et al. (2010), by incorporating an externally calibrated X-ray observable mass relation and X-ray measurements of the sample in order to present improved cosmological constraints. More precisely, the authors used a sub-sample of a SZ-selected catalog from the SPT that was described by Vanderlinde et al. 2010 and consisted of 18 clusters at $z > 0.3$ for the cosmological results (for more details see: Benson et al., 2011). ",
        "conclusions": "Using the latest measurements from the South Pole Telescope (SPT) Sunyaev Zel'dovich (SZ) cluster survey, Benson et al., 2011, found that adding the SPT cluster data significantly improves the constraints on equation of state of dark energy, $w$, beyond those found when using measurements of the CMB, supernovae, BAO and the Hubble constant. In this research note we performed a further quantitative estimation of these better constraints, through the computation of the Figure of Merit (FoM) applied to $\\Omega_m$ and $w$ plots for the $68\\%$ and $95\\%$ confidence regions. Our analysis confirms the significant improvements on constraints for $\\Omega_m$, and $w$ parameters gained adding the SPT clusters data to current CMB+BAO+SNe data. The improving factors are $1.56$ and $1.43$ for confidence regions at $68\\%$ and $95\\%$, respectively. Similar conclusions are obtained considering the constraints on the plane: $(\\sigma_8, w)$. This allows us to better evaluate and to better compare the current and future improvements obtained on the cosmological constraints, using both different cosmological probes and different surveys."
    },
    "1112/1112.3194_arXiv.txt": {
        "abstract": "{Supernova remnants (SNRs) are thought to be the main source of Galactic cosmic rays (CRs) up to the ''knee'' in CR spectrum. During the evolution of a SNR, the bulk of the CRs are confined inside the SNR shell. The highest-energy particles leave the system continuously, while the remaining adiabatically cooled particles are released when the SNR has expanded sufficiently and decelerated so that the magnetic field at the shock is no longer able to confine them. Particles escaping from the parent system may interact with nearby molecular clouds, producing $\\gamma$-rays in the process via pion decay. The soft gamma-ray spectra observed for a number of SNRs interacting with molecular clouds, however, challenge current theories of non-linear particle acceleration that predict harder spectra.} {We study how the spectrum of escaped particles depends on the time-dependent acceleration history in both Type Ia and core-collapse SNRs, as well as on different assumptions about the diffusion coefficient in the vicinity of the SNR.}  {We solve the CR transport equation in a test-particle approach combined with numerical simulations of SNR evolution.} {We extend our method for calculating the cosmic-ray acceleration in SNRs to trace the escaped particles in a large volume around SNRs. We calculate the evolution of the spectra of CRs that have escaped from a SNR into a molecular cloud or dense shell for two diffusion models. We find a strong confinement of CRs in a close region around the SNR, and a strong dilution effect for CRs that were able to propagate out as far as a few SNR radii.} {} ",
        "introduction": "Supernova remnants (SNRs) are now widely considered to be sources of Galactic cosmic rays (CRs). Diffusive shock acceleration (DSA) \\citep{Axfetal77, Kry77, Bel78, BlaOst78} and its non-linear modification (NDSA) \\citep{MalDru01} predict a power-law distribution ($N(E)\\propto E^{-s}$) of relativistic particles with index $s=2$ ($s<2$ for the high-energy tail in NDSA). Despite recent advances in DSA and NDSA, it is still unclear why we observe soft ($s>2$) CR spectra at Earth \\citep{Aveetal09} and soft gamma-ray spectra from a number of SNRs: RX~J0852.0-4622 \\citep{Ahaetal07b}, RCW~86 \\citep{Ahaetal09}, SN~1006 \\citep{Aceetal2010_1006}, Cas~A \\citep{2010ApJ...714..163A, Abdetal10_CasA}, and Tycho's SNR \\citep{Accetal11,Gioetal11}.  Understanding gamma-ray SNR spectra in various situations requires the study of the acceleration of the particles in the system (consisting of one or two shock waves), followed by the escape of these particles, their diffusion into the Galactic medium \\citep{Ptuetal10, Capetal10}, and computation of the $\\gamma$-ray emission by various processes.  In an earlier paper \\citep{Teletal11}, we studied the acceleration of particles at SNR shock fronts. In this paper, we follow up this work by studying the escape of these particles from the system, and their interaction with dense media.  A useful probe of the escape of particles from SNRs may be high-energy gamma-ray emission from molecular clouds located close to an efficiently accelerating SNR, the so-called MC-SNR systems. A number of these have been observed by Cherenkov telescopes (CTs) and \\textit{Fermi}. Most sources, including W28 \\citep{Abdetal10_w28}, W44 \\citep{Abdetal10_w44}, IC~443 \\citep{2009ApJ...698L.133A, Abdetal10_ic443}, G349.7+0.2, CTB 37A, 3C391, and G8.7-0.1 \\citep{CasSla10} are located close to dense regions or molecular clouds (MCs), thus one may use them to infer the diffusion parameters in their vicinity. Triggered by these observations, a number of studies \\citep{Gabetal09, Fujetal09} have attempted to explain the soft gamma-ray emission from MC-SNR systems in terms of escaped CRs. These studies have been based on analytical models \\citep{Atoetal95, AhaAto96} and several assumptions: (i) the SNR evolution is either stationary or Sedov-like, (ii) often that particle acceleration is quasi-instantaneous compared with the CR diffusion time, and (iii) the escaping-particle distribution is a power-law or monoenergetic at a given maximum momentum.  The source of CRs is considered to be point-like, i.e., the SNR radius is much smaller than the distance to the MC. \\citet{LiChe10}, \\citet{LiChe11}, and \\citet{Ohietal11}, extended the analytical models to include finite-size sources and finite-size target MCs.  Despite their simplicity, the analytical studies demonstrate that to explain observations, the diffusion coefficient in the SNR vicinity must be roughly two orders of magnitude smaller than the average Galactic value, supporting earlier claims \\citep{Wen74} that the diffusion coefficient might be smaller because of plasma waves that scatter particles. Using a simple model of SNR evolution and Monte Carlo simulations of CR diffusion, \\citet{Fujetal10} showed that particles may be trapped around the SNR for a significant period of time. Subsequently, \\citet{Fujetal11} confirmed this result using a transport equation to describe particle acceleration at the forward shock (FS), and then a simplified transport equation to follow the propagation of particles with $p/m_0c \\geq 10^{2.4}$ from the precursor boundary (set at four Bohm diffusion length-scales) to distances far away from the shock. These findings are consistent with the conclusions of \\citet{Revetal09}, who argue that beyond some critical value, $L$, which they call the free-escape boundary, the number density of CRs falls significantly. The position of $L$ is dictated by the level of excited magnetohydrodynamics (MHD) turbulence, which may be significant out to one SNR radius from the FS \\citep{ZirPtu08a}.  The detailed variation of the turbulence amplitudes and the CR diffusion coefficient depends on the micro-physical balance of streaming instabilities and turbulence damping \\citep{Yanetal12}.  Besides the diffusion coefficient in the SNR vicinity, other important aspects of the study of escaped CRs are their spectral and spatial distributions. These distributions affect the radiative properties of the SNR itself as escaping CRs modify the high-energy part of the CR spectrum, and, if present, the radiative properties of the nearby dense matter consisting of a MC or a dense shell swept up by the winds of a high-mass SN progenitor). Much of the research dealing with stationary and (semi-)analytical solutions cannot recover the spectral and spatial shapes of the CR distribution, but only gives the integrated escaping CR energy flux. Kinetic models provide the shape of the distribution, which depends significantly on the assumed diffusion models.  \\citet{EllByk11} parametrized the spectral distribution of escaping CRs, and then propagated the escaping CRs in the upstream region of the SNR obtaining spatial distributions. They used a spherically symmetric code for the hydro simulations of the SNR evolution that is identical to the code used in this work, but without the grid expansion that significantly increases the resolution. However, the particle acceleration was treated in a plane-parallel steady-state approximation ignoring dilution effects. \\citet{EllByk11} focused on the core-collapse explosion of a massive star in a low density medium surrounded by a dense shell, a case similar to our model for a core-collapse SNR (CC-SNR) in a wind bubble.  In this paper, we investigate how both the spectrum of particles escaped from the SNR into a nearby MC or shell and the particle emission due to interaction with dense matter depends on the acceleration history of young ejecta-dominated Type-Ia and Type-Ic (core-collapse) SNRs, using given different assumptions about CR diffusion in the vicinity of the SNR. For this purpose, we combine a test-particle treatment of CR acceleration with an account of escaping particles. We solve the CR transport equation in a spherically symmetric geometry. Our calculations are based on realistic high-resolution hydrodynamic simulations of SNR evolution. We explore two SN types, focusing on the complexities in the circumstellar environment. In the case of Type-Ia SNRs, we assume that evolution proceeds in a uniform ISM. We also consider a core-collapse SNR arising from a Wolf-Rayet (WR) progenitor, expanding in the wind-blown bubble created by the progenitor star.  The SNR shock wave first evolves in the freely expanding wind of the progenitor, and then, beyond the wind-termination shock, in the shocked-wind region. Subsequently, we would expect the shock to impact the dense shell bordering the wind-blown bubble. However, our intention here is to study the interaction of the accelerated particles produced by the SNR with this dense shell, hence we terminated the SNR evolution before the shock-shell impact took place. We considered particle acceleration at both the forward and reverse shock. In addition, in the CC-SNR case we consider particle re-acceleration at secondary shocks formed when the SNR forward shock collides with the termination shock of the wind zone. The spatial domain considered in our particle simulations extends to about 100 times the SNR radius to facilitate the study of particle escape. The ``absorbing'' boundary located that far from the SNR does not affect the result of our simulations. It is the spatial profile of the diffusion coefficient in the vicinity of the SNR that determines the degree to which particles can escape from the CR precursor.  To estimate the observational consequences of CR escape, we consider the interaction of the escaped particles with a target MC close to a Type-Ia SNR, or in the core-collapse case we consider the interaction of CRs with the dense shell. The MC/shell are located well inside the simulation domain, hence we know the cosmic-ray number density at their location at any time. Using two different models of the CR diffusion upstream of the FS, we calculated the evolutions of both the cosmic-ray spectra at the location of the MC/shell and the hadronic gamma-ray emission as the SNR shock approaches the MC/shell. If not stated otherwise, we refer to particles as escaped when they have escaped from the precursor region of the FS and are found at the location of the MC or shell. The respective particle spectra are averaged over the volume of the MC or shell. ",
        "conclusions": "We have modified our modeling technique developed in \\cite{Teletal11} to examine the distribution of escaped CRs at distances far from SNRs. Our approach combines a realistic treatment of SNR evolution using hydrodynamic simulations with test-particle calculations of CR acceleration, by solving the transport equation of cosmic rays in a spherically symmetric geometry. Both the reverse and forward shocks are included in these simulations. We also account for the re-acceleration at reflected shocks in core-collapse SNR. Our approach inherently includes the effects of dilution due to the expansion of the system, as well as the finite size of the SNR and any gas target illuminated by escaping CRs. It permits the reconstruction of the spectra of escaped CRs at any given distance from the source.  We have shown that the peak energy and intensity of escaped particles strongly depends on the efficiency of the diffusion in the vicinity of the SNR. If diffusion is Bohmian out to the boundary $L = 2 R_{\\rm SNR}$, CRs are very efficiently confined to the SNR, and only the highest-energy particles are able to diffuse out to some distance from the SNR. Even though CRs at $E_{\\rm max}$ leave the SNR, they are still trapped at distances far less than one SNR radius from the FS. If the diffusion coefficient in the FS upstream region is much larger than Bohmian, but not at the level of average Galactic diffusion, then CRs of lower energy are able to escape from the SNR. Consequently, one observes the broader spectra of escaped CRs with lower $E_{\\rm max}$ and the cut off in the spectra of confined CRs is slower than exponential.  It is possible but difficult to constrain the diffusion coefficient in the vicinity of the SNR using the emission from nearby dense material. The particles cannot generally propagate far from the shock. Thus, to study CR acceleration in SNRs via CR escape into a nearby MC or shell, one needs to find MCs or shells located very close, unless the diffusion coefficient in the immediate environment of the SNR is a significant fraction of the average Galactic one, in which case $E_{\\rm max}$ is lower. We note that in all diffusion models, and especially in the ``far`` scenarios, one can clearly observe the striking effect of dilution, i.e. the decrease in the CR intensity upstream of the FS, introduced by the spherical geometry of the SNR.  Ignoring the spherical geometry of the system is a serious mistake when one models the gamma-ray emission from MCs placed at significant distances from the SNRs.  We find that Type-Ia SNR in our simulation has a higher CR energy density in its vicinity than Type-Ic core-collapse SNR. Therefore, Type-Ia SNR may have brighter pion-decay gamma-ray emission than Type-Ic SNR that is still evolving in the wind zones of its progenitor star. The emission of Type-Ia SNR is normally brighter than the MC emission until the SNR shock is very close to the MC. Only then is the gamma-ray flux from the MC comparable to that from the SNR. The emission from the CC-SNR and the wind-blown shell is initially near the background level. At a later stage, the shell emission can significantly dominate over the SNR emission. Although the exact scenario depends on the MC/shell parameters, most importantly the target mass, the described trend may serve to differentiate between these SNR types.  Finally, we find that the commonly used approximation for the $E_{\\rm max}$ evolution based on Sedov scaling is not reproduced when the evolution of the SNR shocks is computed accurately. Moreover, the $E_{\\rm max}$ behavior critically depends on the diffusion coefficient in the vicinity of the SNR."
    },
    "1112/1112.0018_arXiv.txt": {
        "abstract": "The tidal disruption of a solar mass star around a supermassive black hole has been extensively studied analytically$^{1,2}$ and numerically$^3$. In these events the star develops into an elongated banana-shaped structure. After completing an eccentric orbit, the bound debris fall onto the black hole, forming an accretion disk and emitting radiation$^{4,5,6}$. The same process may occur on planetary scales, if a minor body orbits too close to its star. In the Solar System, comets fall directly onto our Sun$^7$ or onto planets$^8$. If the star is a compact object, the minor body can become tidally disrupted. Indeed, one of the first mechanisms invoked to produce strong $\\gamma-$ray emission involved accretion of comets onto neutron stars in our Galaxy$^9$. Here we report that the peculiarities of the `Christmas' $\\gamma-$ray burst  (GRB 101225A$^{10}$) can be explained by a tidal disruption event of a minor body around a Galactic isolated neutron star. This result would indicate either that minor bodies can be captured by compact stellar remnants more frequently than it occurs in our Solar System or that minor body formation is relatively easy around millisecond radio pulsars. A peculiar Supernova associated to a GRB may provide an alternative explanation$^{11}$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.2701_arXiv.txt": {
        "abstract": "We carry out high-resolution adaptive mesh refinement simulations of a cool core cluster, resolving the flow from Mpc scales down to pc scales.  We do not (yet) include any AGN heating, focusing instead on cooling in order to understand how gas gets to the supermassive black hole (SMBH) at the center of the cluster.  We find that, as the gas cools, the cluster develops a very flat temperature profile, undergoing a cooling catastrophe only in the central 10-100 pc of the cluster.  Outside of this region, the flow is smooth, with no local cooling instabilities, and naturally produces very little low-temperature gas (below a few keV), in agreement with observations.  The gas cooling in the center of the cluster rapidly forms a thin accretion disk.  The amount of cold gas produced at the very center grows rapidly until a reasonable estimate of the resulting AGN heating rate (assuming even a moderate accretion efficiency) would overwhelm cooling.  We argue that this naturally produces a thermostat which links the cooling of gas out to 100 kpc with the cold gas accretion in the central 100 pc, potentially closing the loop between cooling and heating.   Isotropic heat conduction does not affect the result significantly, but we show that including the potential well of the brightest cluster galaxy is necessary to obtain the correct result.  Also, we found that the outcome is sensitive to resolution, requiring very high mass resolution to correctly reproduce the small transition radius. ",
        "introduction": "The majority of baryonic matter in galaxy clusters resides in the form of virialized hot gas that emits in the X-ray band \\citep[e.g.][]{ Lin}. Observations of the X-ray surface brightness, along with imaging spectroscopy, allow the measurement of the density and temperature of the intracluster medium (ICM). Many, if not most, relaxed clusters possess a cool core (CC), where the temperature drops by a factor of 2 to 3 compared to the cluster outskirts and the radiative cooling time is much shorter than the Hubble time \\citep[e.g.][]{Cavagnolo08, Sanderson09_2}.  In a steady state, a ``cooling flow'' of $100s$ M$_{\\odot}/$yr would develop in a typical CC cluster \\citep[see review by][]{Fabian94}. However, XMM grating spectra have revealed a dearth of gas at temperatures lower than about one-third of the cluster virial temperature, despite the short cooling time of the gas \\citep{Tam01, Peterson03, Sanders08}. In addition, the observed star-formation rate \\citep[e.g.][]{Cardiel} is typically a factor of $\\sim 10-100$ times lower than the mass deposition rate predicted by classic cooling flow models. This discrepancy implies that there is some process that heats the gas, preventing it from cooling. Many mechanisms have been proposed, including thermal conduction \\citep[e.g.][]{ZN03}, active galactic nuclei (AGN) feedback \\citep[e.g.][]{Bohringer93, McNamara05, BM06}, MHD wave-mediated cosmic rays combined with conduction \\citep[e.g.][]{Loewenstein, GuoOh08}, and turbulence combined with conduction \\citep[e.g.][]{DC05}.  AGN feedback is considered to be the most plausible heating source for a number of reasons \\citep{McNamara07}.  First, there is an observed correlation between the presence of cool cores and signs of AGN activity \\citep[e.g.][]{DF06}.  In addition, in nearby CC clusters, X-ray cavities or bubbles are observed in the core \\citep[e.g.][]{Fabian00} and are thought to be inflated by the interactions between powerful jets and the surrounding gas.  Finally, the estimated energy injected into the surrounding gas by the bubbles can balance cooling in most of these systems \\citep[e.g.][]{Birzan04}. However, how the gas feeds the AGN, i.e. whether the gas accretion is dominated by the ``hot'' mode described in the Bondi-Hoyle accretion model \\citep{Bondi,Allen06} or by the ``cool'' mode assuming that the hot gas first fragments within the sonic radius and falls onto the black hole as cold clouds \\citep[e.g.][]{BS89}, and how AGN deposits the energy to the ICM (whether kinetic or thermal feedback dominates) are still open questions -- see \\citet{Croton} and references therein. To address these issues, we need to first understand exactly how gas cools and eventually accretes onto the black hole.  Theoretical work has been carried out using linear perturbation theory to investigate the local cooling instability \\citep[e.g.][]{Binney, BS89}, deriving the global dynamics of the cooling flow in a steady state \\citep[e.g.][]{QN2000}, or simulating the global cooling flow numerically \\citep[e.g.][]{CT07, Gaspari10, Guo10}. However, due to its complex and non-linear nature, an accurate picture of the cooling flow cannot be derived from analytical models, and the previous simulations usually suffer from a low resolution with finest grid sizes typically around a kpc.  In this paper, we carry out 3D simulations using the adaptive-mesh hydro code Enzo to investigate the onset of the cooling flow. Our primary objective is to explore the cooling catastrophe in detail and examine how gas cools and flows inwards in the cluster center.  We do not attempt to model the accretion on to the SMBH itself and the resulting feedback, and so our model is not a complete description of thermally-balanced feedback; however, by examining the detailed evolution of a cluster as it undergoes a cooling catastrophe, we can get a better idea of the evolution and production of cold gas in the cluster core.  Key issues that this study aims to address include: (i) confirmation of linear calculations that local cooling instabilities do not grow significantly before a cooling catastrophe occurs, (ii) the location and amount of cold gas produced by the global thermal instability, (iii) the rate of gas accretion on to a central SMBH, (iv) observational signatures during the cooling catastrophe (in particular, the lack of cool gas observed in X-rays). (v) the impact of other processes (such as isotropic thermal conduction and Type Ia SN heating rates) on the cooling catastrophe.  This paper is organized as follows, in Section~\\ref{sec:method}, we describe our code, initial conditions and (non-standard) refinement method; in Section~\\ref{sec:results}, we present the main results from our standard simulation; in Section~\\ref{sec:discussion}, we discuss the related physical and numerical issues, and compare our results with observations and previous work. We conclude in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We performed high resolution simulations of an idealized, cool-core X-ray cluster loosely based on Perseus, in order to better understand how gas cools and accretes on to the SMBH.  The use of an AMR code was essential in order to resolve the important scales, ranging from the Mpc virial radius of the cluster, down to the pc scale, which is the size of the cold gas clumps.  We do not include feedback from the AGN and so the current work is not a complete description of cool-core cluster evolution; however, it does shed light on where cold gas condenses out of the flow, where heating needs to occur, and what observational signatures we would predict for such an object.  We simulate an idealized cluster, but include small-scale velocity perturbations, as well as some net rotation, as implied by cosmological simulations.  We include radiative cooling, but no heating term.  The simulation resolves regions of high density and short cooling time.   We present our primary results below.  \\begin{itemize}  \\item We find that as the cluster core cools, the temperature profile remains remarkably flat, from a few hundred pc out to at least 20 kpc, with a change of only a factor of $\\sim 2$ in temperature over this range.  We argue that this occurs because the gas flow adjusts such that adiabatic compression balances cooling at each radius over this range.  The gas that does cool does so via a global cooling catastrophe that occurs first at a `transition' radius of about 10 pc from the SMBH in our best resolved simulation.  This occurs about 300 Myr after the start of the simulation, which is about the cooling time of the gas in the initial configuration.  \\item We stop the simulation about 20 Myr after the cooling catastrophe when the accretion rate onto the black hole (which we measure at a radius of 400 pc) becomes sufficiently large that a reasonable feedback efficiency could balance radiative cooling in the flow.  At this point, when we examine the amount of cold gas, there is a distinct lack of gas below a few keV (because of the flat temperature profile mentioned above), and is in agreement with observational constraints on Perseus.  \\item  Although the simulation is intrinsically three-dimensional, the solution outside of the transition radius is well-described by spherically symmetric flow.  In particular, we find that local thermal instabilities do not grow in the cool-core cluster within 300 Myr (i.e. before the global cooling catastrophe), at least outside of the transition radius around $\\sim 100$ pc. This is despite the fact that we seed the initial flow with substantial perturbations.  \\item Isotropic thermal heat conduction does not significantly affect this result.  It does not lead to a large delay, nor does it significantly affect the density and temperature profile in the center.  Heating from Type Ia SN is similarly unimportant.  \\item The final result is not very sensitive to the gas initial conditions, but is sensitive to the presence of the BCG.  Without a BCG, the gas temperature drops much more rapidly with radius, in conflict with observations.  \\item We find that some aspects of the solution are quite sensitive to the resolution of the simulation.  In particular, we find that very high mass resolution is required in the gas near the cooling flow, and that lower resolution runs resulted in both a steeper slope in the temperature profile in the intermediate range from the transition radius to 20 kpc, as well a larger initial transition radius.  \\end{itemize}  These results are intriguing for a number of reasons.  First, they imply that the real ``cooling-flow\" problem arises very close to the center of the cluster (i.e. within a few hundred pc), in the sense that it is only in this region that gas cools to temperatures below those observed.  Therefore, the lack of cooler gas seen in X-ray spectroscopy may arise not because of the way that heating operates, but because of the temperature plateau caused (we argue) by the BCG potential well.   Mass drop-out occurs only at the very center of the cluster.  The fact that mass drop-out occurs first very close to the SMBH also provides a mechanism for thermal regulation to operate.   One of the unanswered questions for AGN regulation is how the properties of the gas in the entire cooling region, which stretches out to 100 kpc, can control what happens in the center hundred pc.  These simulations help to answer this by showing that this naturally occurs -- the lack of local thermal instabilities insures that the flow is coherent, and, as we show, the result is cold gas condensing out only inside the transition radius, where it can almost immediately form an accretion disk and feed the SMBH.  Of course, this doesn't relieve us of the need for energetic feedback, as we show that the mass inflow rate quickly grows and after a few hundred Myr, reaches hundreds of solar masses per year, in agreement with steady-state models, and in disagreement with observations.  We argue that, before this occurs, the accretion disk that we see feeds gas to the SMBH, generating `radio-mode' feedback that heats up the gas.  How exactly this feedback occurs is not clear and not addressed in this paper.  We will examine this point in more detail in future work.  Another aspect which is not explained by these results is the generation of H$\\alpha$ filaments seen in most cool core clusters.  Here we simply note that they do not naturally occur at large radii via a local cooling instability before the flow becomes globally unstable."
    },
    "1112/1112.0368_arXiv.txt": {
        "abstract": "The Kepler Mission recently identified 997 systems hosting candidate extrasolar planets, many of which are super-Earths. Realizing these planetary systems are candidates to host extrasolar asteroid belts, we use mid-infrared data from the Wide-field Infrared Survey Explorer (WISE) to search for emission from dust in these systems. We find excesses around eight stars, indicating the presence of warm to hot dust ($\\sim$100-500~K), corresponding to orbital distances of 0.1-10~AU for these solar-type stars. The strongest detection, KOI~1099, demands $\\sim$500~K dust interior to the orbit of its exoplanet candidate. One star, KOI~904, may host very hot dust ($\\sim$1200~K, corresponding to 0.02~AU). Although the fraction of these exoplanet-bearing stars with detectable warm excesses ($\\sim$3\\%) is similar to that found by Spitzer surveys of solar-type field stars, the excesses detectable in the WISE data have much higher fractional luminosities (L$_{\\rm dust}$/L$_{*}$) than most known debris disks, implying that the fraction with debris disks of comparable luminosity may actually be significantly higher. It is difficult to explain the presence of dust so close to the host stars, generally corresponding to dust rings at radii $<$0.3~AU; both the collisional and Poynting-Robertson drag timescales to remove dust from the system are hundreds of years or less at these distances. Assuming a steady-state for these systems implies large mass consumption rates with these short removal timescales, meaning that the dust production mechanism in these systems must almost certainly be episodic in nature. ",
        "introduction": "The evolution of solid circumstellar material is a competition between the accretional processes that form planets (when speeds are low) and the violent collisional processes which in post-formation systems inexorably grind down the material into particles small enough that radiation forces can remove them from the system. The disappearance of infrared (IR) excesses around forming stars (on 3-10~Myr time scales) occurs because the accretional processes win and sequester the huge mass of dust into a (relatively) tiny number of objects with very low surface area/mass ratios. In contrast, the observed debris disks (known around hundreds of solar-type stars) are thought to be from collisional dust production as velocities rise during the final stages of planet assembly (for relatively young stars), or triggered around older stars as dust is liberated in higher-speed collisions \\citep{Wyatt2008}. It is these more mature systems where our understanding of the physical picture is arguably incomplete. Although only $\\sim$2\\% of nearby main-sequence stars have massive detectable warm disks \\citep[$>$200K;][]{Lawleretal2009}, several of these disks require dramatic hypotheses to produce the large dust mass estimated. That is, estimates of the collisional grinding of small-body belts indicate that the systems in question should not still possess dust at anywhere near the observed level \\citep{Wyattetal2007}. In fact, \\citet{Rheeetal2008} went so far as to liken the dust mass around BD~+20~307 as a `miracle' given collisional-model predictions.  Debris disks are detected and studied via the excess IR flux they produce in the system, caused by the dust re-radiating the absorbed stellar light. This dust emission can then dominate the stellar photosphere in the IR, where the dust's blackbody peak is set by the equilibrium temperature at the orbital distance (for a narrow ring) or at an extended disk's inner edge, because the warmest dust usually dominates the systems dust emission. For `realistic' (although uncertain) dust grain size distributions, cross-sectional area and thus thermal emission is dominated by small grains, but the system mass is dominated by the unconstrained larger bodies. The true observable is thus the mass in grains in the decade larger than the peak wavelength.  Known disk systems consist of two types. The most common are massive and cold, analogous to the Solar System's Kuiper belt. A few hotter (T$>$200K) systems are known, which are massive cousins of our asteroid belt. An example is HD~69830, hosting 3 RV-discovered Neptune-mass planets inside 0.6 AU \\citep{Lovisetal2006} and a warm dust ring at $\\sim$1~AU, with an estimated 4x10$^{-7}$ M$_{\\oplus}$ in small grains alone \\citep{Lisseetal2007}. Production of this dust is problematic \\citep{Beichman69830}, with hypotheses ranging from comet swarms to super-comets to massive planetary collisions; it is difficult to judge the plausibility of these mechanisms given the rarity of these outcomes.  Dust grains in orbit around a star are quickly destroyed, and must be replenished by a fragmentation cascade, where asteroids within a belt collide and break into progressively smaller pieces \\citep{Wyatt2008}. The lifetime of a debris disk will be the same as the lifetime of the largest bodies in the source population that are available to be pulverized into dust grains. Very small grains ($\\lesssim$1~$\\mu$m) are quickly blown away by stellar radiation pressure. Larger grains suffer Poynting-Robertson (PR) drag and spiral into the star in much less than the system's age; for dust very close to the star ($\\sim$0.1~AU), orbits collapse in only hundreds of years.  In this paper we search for excess infrared emission, indicating the presence of a dusty debris disk, in systems found by the Kepler mission to possess transiting exoplanets. In Section~\\ref{sec:WISE} we define our sample and discuss our methodology for identifying excess candidates, and in Section~\\ref{sec:DDs} we provide additional details on the eight stars that have excesses. Section~\\ref{sec:dust} contains a discussion of the problems associated with having this much dust so close to the host star, and finally Section~\\ref{sec:conc} contains a summary of recent work and this study's contribution relating debris disks to the presence of exoplanets. ",
        "conclusions": "\\label{sec:conc}  We find IR excesses around eight stars that also host planet candidates. The masses and luminosities of these excesses are all too large to be explained by a steady-state collisional cascade. Given that the WISE field coverage was not uniform in depth, and so many of these relatively faint stars suffered from contamination by nearby bright stars, it is challenging to say anything robust about the fraction of these planet-hosting stars that have excesses in our sample. Out of the 186 stars with W3 data of SNR$>$3, five have significant $>$5~$\\delta$ excesses, or $\\sim$3\\%. At first glance, this fraction appears similar to the 2\\% of solar-type stars with warm excesses \\citep{Wyattetal2007,Lawleretal2009}.  However, the fractional luminosities of these Kepler systems with excesses are much higher than most known debris disks \\citep[Figure~\\ref{fig:dustlocations} and][]{Mooretal2006}. Because of the larger distances and thus lower apparent brightness of these Kepler systems as compared to most known debris disk systems, only the brightest debris disks would be detected by WISE. Presumably the eight systems presented here are the very highest luminosity debris disks in the sample, and many more fainter disks are present around these planet-hosting stars. This would indicate that the fraction of stars possessing debris disks and Kepler-detected exoplanets is actually much higher than the fraction of field stars with debris disks.  Surveys with Spitzer \\citep{Brydenetal2009,DodsonRobinsonetal2011} have found that hosting a radial velocity-detected exoplanet does not make a star more likely to host a debris disk, and a recent WISE survey of exoplanets discovered by ground-based transit searches \\citep{Krivovetal2011} has found that the fraction of these stars ($\\sim$2\\%) to host warm debris disks (excesses in W3 and W4) is similar to that found by \\citet{Wyattetal2007}.  The Kepler systems studied here are different from other disk-bearing exoplanet systems, whose planets are primarily discovered by the radial-velocity (RV) technique, and by ground-based transit searches. Exoplanets from the latter two techniques are mostly large and very close to their stars, with an average mass of $\\sim$2 Jupiter masses and $<$20\\% known to be in multiplanet systems \\citep{Wrightetal2011}. The recently announced 1200 Kepler exoplanet candidates \\citep{Boruckietal2011} are in a different regime, with about 1/3 in multiplanet systems, and having many super-Earths (typically 2-6 Earth radii) well inside 1~AU. Among our excess systems, all but two of the planets are Neptune mass or smaller; only one is massive enough to be a hot Jupiter. The resulting correlation we find between a star hosting both small planets and a debris disk inside a few AU agrees with the planetary formation and dynamical evolution models presented by \\citet{Raymondetal2012}. This makes sense in a picture where migrating Jupiters clean out the terrestrial systems they pass through.  Explaining the large dust masses required to produce these excesses remains a huge challenge. Steady-state theories appear to fail completely, as the mass consumption rates required would involve destroying thousands of Earth masses of planetary material over the lifetime of these stars. Episodic events such as catastrophic collisions between bodies may provide an explanation, but these events must last long enough and occur frequently enough to result in $\\sim$3\\% of solar-type stars in this state at any given moment. Much theoretical work needs to be done here.  The upcoming WISE full-sky data release will access additional systems, decrease the error bars on existing measurements, and provide longer-wavelength measurements for some of these systems to see if there is dust at larger orbital distances. Spectral data on these candidates would be invaluable, but the host stars are faint. Future Kepler data releases will contain planets on more distant orbits, perhaps yielding planets in closer proximity to these new debris disks."
    },
    "1112/1112.6191_arXiv.txt": {
        "abstract": "We present an N-body model that reproduces the morphology and kinematics of the Magellanic Stream (MS), a vast neutral hydrogen (HI) structure that trails behind the Large and Small Magellanic Clouds (LMC and SMC, respectively) in their orbit about the Milky Way.  After investigating $8\\times10^6$ possible orbits consistent with the latest proper motions, we adopt an orbital history in which the LMC and SMC have only recently become a strongly interacting binary pair.  We find that their first close encounter $\\sim$2 Gyr ago provides the necessary tidal forces to disrupt the disk of the SMC and thereby create the MS.  The model also reproduces the on-sky bifurcation of the two filaments of the MS, and we suggest that a bound association with the Milky Way is required to reproduce the bifurcation.  Additional HI structures are created during the tidal evolution of the SMC disk, including the Magellanic Bridge, the ``Counter-Bridge,\" and two branches of leading material.  Insights into the chemical evolution of the LMC are also provided, as a substantial fraction of the material stripped away from the SMC is engulfed by the LMC.  Lastly we compare three different N-body realizations of the stellar component of the SMC, which we model as a pressure-supported spheroid motivated by recent kinematical observations.  We find that an extended spheroid is better able to explain the stellar periphery of the SMC, and the tidal evolution of the spheroid may imply the existence of a stellar stream akin to the gaseous MS. ",
        "introduction": "\\label{sec:int}  Disentangling the interaction history of the Large and Small Magellanic Clouds (LMC and SMC, respectively; ``MCs\" collectively) has been facilitated by modeling the formation of the Magellanic Stream (MS), a prominent neutral hydrogen (HI) structure that trails behind the MCs in their orbit about the Milky Way (MW).  Theoretical models have assigned one of two physical mechanisms for creating the MS: tidal stripping (Lin \\& Lynden-Bell 1977; Murai \\& Fujimoto 1980, MF80; Gardiner, Sawa \\& Fujimoto 1994, GSF94; Gardiner \\& Noguchi 1996,  GN96; Yoshizawa \\& Noguchi 2003; Connors et al. 2006, C06) and ram pressure stripping (Meurer et al. 1985; Heller \\& Rohlfs 1994; Moore \\& Davis 1994; Mastropietro et al. 2005, M05).  The common feature of these MS formation models is their reliance on multiple strong interactions between the MCs and the Milky Way.  The plausibility of such an interaction history for the MCs has been challenged, however, by the high-precision proper motion measurements of Kallivayalil et al. (2006a,b; K06).  These proper motion estimates imply such large orbital velocities for the MCs that a ``first passage\" scenario has been proposed in which the MCs are not bound to the Milky Way but are instead passing by for the first time (Besla et al. 2007).  Because the past orbits of the LMC and SMC have been called into question, the origin of the MS must also be revisited.  { For example, the ``blowout hypothesis\" has emerged recently as an alternate scenario for the formation of the MS, which relies on internal mechanisms to expel gas from the MCs (Olano 2004, Nidever et al. 2008).}  Attempts to reconcile the K06 proper motions with the formation of the MS have not yet produced a compelling dynamical model.  { For example, even though the first passage model of Besla et al. (2010, B10) reproduces a tidal tail resembling the MS, the model does not adequately recover the known positions and velocities of the MCs.  In particular, a number of important orbital constraints for the SMC are not satisfied, including on-sky position, line-of-sight velocity, and proper motion (see section~\\ref{sec:dx-compare} for a clarifying discussion as well as Besla et al. 2012).  Ruzicka et al. (2010) searched a large parameter space centered on the K06 proper motions, but they were unable to find a test-particle model that could simultaneously explain the on-sky location and kinematics of the MS.  In our previous work, we were able to identify a good MS model, but we had to adopt a massive isothermal halo for the Milky Way in order to combat the large velocities of the MCs (Diaz \\& Bekki 2011a,  DB11a).}  { Nevertheless, one may salvage some insight from the above three studies.  Ruzicka et al. (2010) point out that their most promising models exhibit a common interaction history: two close encounters between the LMC and SMC at $<$2.5 Gyr and 150 Myr ago, each of which triggers an epoch of tidal stripping from the SMC disk.  Similar conclusions are reached in DB11a and B10, in which the MS forms during the first strong binary interaction between the MCs rather than during an interaction with the Milky Way.  Whereas previous tidal models relied on a combination of tides from both the Milky Way and LMC (e.g., GN96, C06), these recent models indicate that the onset of violent LMC tidal interactions could have been sufficient to disrupt the SMC disk and thereby create the MS.}  This new scenario would imply that the binary state of the MCs is a recent phenomenon and that the MS is an artifact from their dynamical coupling.  { Furthermore, observations appear to corroborate this scenario, particularly because the star formation histories of the LMC and SMC exhibit two correlated bursts of star formation at $\\sim2$ Gyr ago and $\\sim$500 Myr ago (Harris \\& Zaritsky 2009).  These two epochs may correspond to two strong tidal interactions suffered mutually by the LMC and SMC, which in turn would create the physical conditions necessary for episodic bursts of cluster formation (Bekki et al. 2004; Piatti et al. 2005; Piatti 2011).}  Regarding the tension between the K06 proper motions and MS formation models, a possible compromise has been offered by the recent proper motion measurements of Vieira et al. (2010, V10) and Costa et al. (2009): perhaps the velocities of the MCs are indeed larger than traditionally assumed (e.g., C06, M05) but not as large as what K06 would imply.  The V10 measurements are not nearly as precise as those of K06, but Bekki (2011a) has suggested that the V10 proper motions may nevertheless be more \\emph{accurate} than those of K06 due to the larger sample size (3822 stars versus 810 stars, respectively, for the LMC) and the longer baseline (40 years versus 2 years, respectively).  Bekki (2011a) argues that random stellar motions would create unknown systematic errors that could undermine the accuracy of the K06 measurement but would be largely suppressed for V10.  Another unique advantage of V10 is that they measure a precise \\emph{relative} proper motion between the LMC and SMC, achieved by tracking the MCs in the same wide field images ($\\sim$450 square degrees).  Considering that many properties of the MCs can be understood in terms of their recent activity as a binary pair (star formation histories, Harris \\& Zaritsky 2009; bursts of star cluster formation, Bekki et al. 2004, Piatti et al. 2005; formation of MS, DB11a, B10), the V10 constraint on the relative motions of the LMC and SMC are an asset to theoretical models.  In the present work we use orbital models and N-body simulations to address the formation of the MS and other HI structures of the Magellanic system, including the Magellanic Bridge, which extends between the SMC and LMC, and the Leading Arm (LA), which stretches ahead of the MC orbits.  The framework of the present model provides an improvement over previous MS models (e.g., C06, M05, B10) in three important ways.  First, we explore a wide range of orbital histories with velocities constrained by the observational results of V10.  Second, we adopt a multi-component potential for the Milky Way composed of a disk, bulge, and NFW halo, which is a more realistic choice than, for instance, the one-component Milky Way potentials adopted by C06, B10, and DB11a.  And third, we represent the SMC as having both a rotating disk and a non-rotating spheroid, whereas previous MS models have only represented the SMC as a ``pure disk\" system.  We are thus able to correctly reproduce the observed kinematics of the SMC, both its rotating HI component (Stanimirovi\\'c et al. 2004) and its non-rotating stellar component (Harris \\& Zaritsky 2006, HZ06).  Our adopted spheroid model is also consistent with the extended stellar halo of the SMC (e.g., De Propris et al. 2010; Nidever et al. 2011).  In our best model, we find that the SMC disk remains intact until the dynamical coupling of the MCs $\\sim$2 Gyr ago, at which point the MS and LA are violently torn away by the strong tidal forces of the LMC.  Additional tidal debris is subsequently engulfed by the LMC, creating a transfer of mass from the SMC that may provide insight into the chemical enrichment history of the LMC.  A second tidal encounter occurring $\\sim$250 Myr ago is responsible for pulling the Bridge from the SMC disk as well as a complementary structure that we call the ``Counter-Bridge\".  Our model exhibits strong agreement with the HI kinematics and morphology of the MS, particularly its well-known bifurcation into two distinct filaments (Putman et al. 2003a).  For the first time, we are able to provide a structural interpretation of the MS filaments within the context of a dynamical model (however, see Nidever et al. 2008), and we find that a bound association with the Milky Way is required for at least 2 Gyr to explain the bifurcation.  We accordingly suggest that the morphology of the MS provides a strong argument in favor of bound orbits for the MCs, and in particular that a first passage orbit, in which the MCs interact with the Milky Way only recently, may not be able to explain the MS bifurcation.  Using ``disk plus spheroid\" N-body models, we describe for the first time the evolution of the SMC spheroid under the same tidal forces that pulled the MS from the SMC disk.  In order to provide a comparative analysis of important physical parameters, three different spheroid models are presented.  We demonstrate that an extended, low-density spheroid is better able to reproduce observations, including the stellar kinematics of the SMC (HZ06), the recent discovery of a break population of red giants (Nidever et al. 2011), and the observation of kinematically and chemically peculiar stars within the LMC (Olsen et al. 2011).  As a consequence of reproducing these observations, we find that the tidal evolution of the extended spheroid necessarily predicts that a coherent stellar stream is stripped away under the same forces that remove the MS and LA from the disk.  The on-sky location of the predicted stellar stream is slightly offset from the observed MS, particularly at its tip, which possibly explains why previous attempts to observe a stellar counterpart to the MS have failed (e.g., Guhathakurta \\& Reitzel 1998).  The plan of the paper is as follows: in the next section, we outline our numerical model including a brief discussion of adopted parameters.  We represent the SMC as a multicomponent system (disk, spheroid, and dark matter halo), and we describe the evolution of its disk and its spheroid in separate sections for the sake of clarity.  In section~\\ref{sec:disk}, we present our results concerning the tidal evolution of the SMC disk, including a discussion of the MS bifurcation.  In section~\\ref{sec:sph} we compare three different spheroid models which enables us to discuss the possibility of a tidal stream of stars.  In section~\\ref{sec:dx} we provide a general discussion including a comparison of recent MS formation models.  In section~\\ref{sec:con} we summarize and conclude.   \\begin{figure*}[t] \\centering \\includegraphics[angle=0, scale=0.55, trim=0 100 0 100]{f1.pdf} \\caption{The proper motion of the LMC (left) and SMC (right).  Observational ellipses outline the 68.3\\% confidence regions and 1$\\sigma$ error bars for Vieira et al. (2010; V10, solid) and Kallivayalil et al. (2006a,b; K06, dashed).  The parameter space of possible values explored in the present work are shaded by the grey boxes, and the adopted values are indicated by the star.  The values chosen by Gardiner \\& Noguchi (1996; plus sign) and Diaz \\& Bekki (2011a; cross) are also shown for comparison.  Models that fail to reproduce the SMC orbit, e.g., Mastropietro et al. 2005 and Besla et al. 2010, are not included in the plot (see text).  The Vieira et. al (2010) proper motion for the SMC is measured \\emph{relative} to the LMC, and the corresponding ellipse therefore roams over the large grey box (right panel) as the LMC proper motion is varied.  For the purposes of the figure, we have chosen to fix the ellipse by the LMC proper motion adopted in the present model (star, left panel). \\label{fig:pm}} \\end{figure*}  \\begin{deluxetable*}{lcc}[t] \\footnotesize \\tablecolumns{3} \\tablewidth{0pc} \\tablecaption{Mass and Orbital Parameters \\label{tab:param1}} \\tablehead{ \\colhead{  Property } &  \\colhead{Adopted Value}  }  \\startdata  & \\multicolumn{2}{c}{\\bf Milky Way } \\\\ \\\\[-2ex] \\hline \\\\[-2ex] \\tablenotemark{$\\dagger$}Total Mass ($r<$300 kpc)\\tablenotemark{1,2} &  \\multicolumn{2}{c}{1.73 $\\times$ $10^{12}$ M$_{\\odot}$\\tablenotemark{$\\ast$}}  \\\\ Disk Mass ($M_{\\rm d}$)\\tablenotemark{1} &  \\multicolumn{2}{c}{5.0 $\\times$ $10^{10}$ M$_{\\odot}$} \\\\ Bulge Mass ($M_{\\rm b}$)\\tablenotemark{1}  &  \\multicolumn{2}{c}{0.5 $\\times$ $10^{10}$ M$_{\\odot}$} \\\\ NFW Virial Mass ($M_{\\rm vir}$)\\tablenotemark{2} & \\multicolumn{2}{c}{1.30 $\\times$ $10^{12}$ M$_{\\odot}$\\tablenotemark{$\\ast$}} \\\\ NFW Virial Radius ($R_{\\rm vir}$) &  \\multicolumn{2}{c}{175 kpc\\tablenotemark{$\\ast$}} \\\\ Circular Velocity ($V_{ \\rm cir })\\tablenotemark{3}$ &  \\multicolumn{2}{c}{240 km s$^{-1}$\\tablenotemark{$\\ast$}}  \\\\ \\tablenotemark{$\\ddagger$}Distance to Sun ($R_{ \\odot }$)\\tablenotemark{3,4} &  \\multicolumn{2}{c}{8.5 kpc} \\\\ \\tablenotemark{$\\ddagger$}Velocity of Sun\\tablenotemark{5} &  \\multicolumn{2}{c}{(10.0, 5.2 + $V_{ \\rm cir }$, 7.2) km s$^{-1}$}  \\\\  \\\\[-2ex] \\hline \\\\[-2ex] & {\\bf LMC } & {\\bf SMC } \\\\ \\\\[-2ex] \\hline \\\\[-2ex] Mass ($M_{\\rm mc}$)\\tablenotemark{6,7} & $10^{10}$ M$_{\\odot}$ & 3 $\\times$ $10^{9}$ M$_{\\odot}$ \\\\ Scale radius ($a_{\\rm mc}$) & 3 kpc & 2 kpc \\\\ Proper Motion ($\\mu_\\alpha \\cos\\delta $, $\\mu_\\delta$)\\tablenotemark{8} & (1.89, 0.30) mas yr$^{-1}$ &  (0.92, -1.37) mas yr$^{-1}$ \\\\ Line-of-Sight Velocity\\tablenotemark{9,10} & 262 km s$^{-1}$ & 146 km s$^{-1}$ \\\\ On-Sky Position ($\\alpha$, $\\delta$)\\tablenotemark{9,11} & (81.9$^{\\circ}$, -69.9$^{\\circ}$)  & (13.2$^{\\circ}$, -72.5$^{\\circ}$) \\\\ Distance Modulus\\tablenotemark{12,13} & 18.50 & 18.95 \\\\ \\tablenotemark{$\\ddagger$}Position Vector ($X$,$Y$,$Z$) & (-0.8, -41.6, -27.0) kpc & (15.3 -36.9 -43.3) kpc \\\\ \\tablenotemark{$\\ddagger$}Space Velocities ($U$,$V$,$W$) & (-50.7, -226.1, 229.3) km s$^{-1}$\\tablenotemark{$\\ast$} & (-4.2, -223.5, 191.0) km s$^{-1}$\\tablenotemark{$\\ast$} \\\\  \\enddata \\tablenotetext{1}{Binney \\& Tremaine (2008);  $^2$Gnedin et al. (2010); $^3$Reid et al. (2009); $^4$Gillessen et al. (2009); $^5$Dehnen \\& Binney (1998); $^6$Kim et al. (1998); $^7$Stanimirovi\\'c et al. (2004); $^8$Vieira et al. (2010); $^9$van der Marel et al. (2002); $^{10}$Harris \\& Zaritsky (2006); $^{11}$Piatek et al. (2008); $^{12}$Freedman et al. (2001); $^{13}$Cioni et al. (2000)} \\tablenotetext{$\\dagger$}{Includes bulge mass, disk mass, and NFW halo mass within $r<300$ kpc.} \\tablenotetext{$\\ddagger$}{Measured in a galactocentric frame of rest.} \\tablenotetext{$\\ast$}{We find that the tidal evolution of the SMC is nearly identical to the present model when the following parameters are substituted for the NFW halo: M$_{\\rm vir}$=1.90  $\\times$ $10^{12}$ M$_{\\odot}$, R$_{\\rm vir}$=269 kpc.  In this ``alternate model\" (see text), the total mass of the Milky Way within 300 kpc is 2.06 $\\times$ $10^{12}$ M$_{\\odot}$ and the circular velocity at the solar radius is 220 km s$^{-1}$.  The parameters for the LMC and SMC are same as above, expect for the $V$ component of the space velocities which decrease by 20 km s$^{-1}$ owing to the change in circular velocity.} \\end{deluxetable*} ",
        "conclusions": "\\label{sec:con}  We have constructed new N-body models for the tidal evolution of the SMC in a $\\sim3$ Gyr interaction history with the LMC and Milky Way, and we have thereby investigated the formation processes of the MS, LA, and Bridge.  The framework of the present investigation has a number of advantages over previous tidal models of the SMC (e.g., C06, B10), including (i) our adopted orbital model was selected after conducting a large parameter space search centered on the measured proper motion values of Vieira et al. (2010), (ii) the Milky Way is represented by a realistic potential having a bulge, disk, and NFW halo, and (iii) we investigated the stellar kinematics and outer stellar structure of the SMC by adopting ``disk plus spheroid\" models in good agreement with observations (e.g., Harris \\& Zartisky 2006; Nidever et al. 2011).  Our main results are summarized as follows.  (1) The LMC and SMC become dynamically coupled only recently, suffering two strong tidal interactions $\\sim2$ Gyr ago and $\\sim250$ Myr ago.  Portions of the SMC disk are stripped away during each of these strong encounters, with the MS and LA being created during the first encounter, and the Bridge during the second.  The timing of these two tidal interactions is consistent with the correlated star formation epochs determined by Harris \\& Zaritsky (2009).  (2) The morphology of the MS is well reproduced, including its bifurcation into dual filaments (Putman et al. 2003a) and its long extent (Nidever et al. 2010).  The predicted on-sky location of the filaments provides a very strong match to observations, as does their crossing point.  The kinematics of the MS are also well reproduced, particularly the velocity inflection at the MS tip (Nidever et al. 2010).  (3) We interpret the observed MS bifurcation as a wrapping of the original tidal arm (``tidal arm A\") that was stripped from the SMC disk.  The dynamical process responsible for this ``wrapping\" can be traced to a complex three-body interplay between the SMC, LMC, and MW over the past $\\sim2$ Gyr.  We accordingly suggest that the HI morphology of the MS places strong constraints on the orbital history of the MCs, namely that the MCs have executed at least two pericentric passages about the MW during a $\\sim2$ Gyr bound association.  (4) The formation of the Bridge $\\sim250$ Myr ago was accompanied by the formation of a complementary structure called the ``Counter-Bridge\" that extends $\\sim20$ kpc along the line of sight behind the SMC.  Because of this alignment, the Counter-Bridge does not appear as an identifiable structure in column density maps, but it is likely traced by the large line-of-sight depth of the SMC (e.g., Crowl et al. 2001; Groenewegen 2000).  (5) Two branches of leading material are stripped from the SMC disk, but the on-sky location and kinematics do not correspond well with the observed properties of the LA. Even though we do not explicitly account for hydrodynamical effects in our simulation, we have argued that the evolution of the leading material over the past $\\sim300$ Myr should be strongly influenced by ram pressure from the hot halo of the Milky Way (Section~\\ref{sec:dx-la} and Fig~\\ref{fig:rvt}).  Accordingly, we suggest that gravity alone cannot explain the present location and kinematics of the LA, even if this structure is tidal in origin.  (6) The LMC engulfs a large fraction of the material that is tidally stripped from the SMC following their strong tidal interactions $\\sim2$ Gyr ago and $\\sim250$ Myr ago.  The transfer of SMC disk material to the LMC would correspond to the infall of metal-poor gas and would be imprinted into the chemical enrichment history of the LMC (e.g., { Fig 1 of van Loon et al. 2005}; Fig 20 of Harris \\& Zaritsky 2009; Rubele et al. 2012) as well as the nitrogen-poor stellar populations of the LMC (e.g., Bekki \\& Tsujimoto 2010).  The mass transfer from the SMC spheroid would correspond to the accretion of old-type stars onto the LMC and is suggestive of various observed populations of kinematically peculiar stars (e.g., Kunkel et al. 1997a, Olsen et al. 2011).  (7) Of our three models for the SMC spheroid, we find that the extended spheroid is best able to reproduce the relevant observations, including the stellar kinematics of the SMC (Harris \\& Zaritsky 2006), the recent discovery of a break population of red giants (Nidever et al. 2011), and the observation of kinematically and chemically peculiar stars within the LMC (Olsen et al. 2011).  Considering these results, future observations may need to readdress the possibility that SMC stars were tidally stripped into the Bridge region.  (8) The tidal evolution of the extended spheroid predicts that a stellar stream should be stripped away $\\sim2$ Gyr ago { at the same time that the MS is stripped from the SMC disk.  Whereas the gaseous MS originates in the rotating SMC disk, the stellar stream originates in the pressure-supported spheroidal component of the SMC.  This distinct origin suggests that the stripped stars should not correlate with the peak HI regions of the MS, as observed (e.g., Recillas-Cruz 1982; Bruck \\& Hawkins 1983; Guhathakurta \\& Reitzel 1998).  Future wide-field surveys such as the Southern Sky Survey by the SkyMapper telescope (Keller et al. 2007) would be ideal for detecting this possible tidal stream and establishing its connection to the stellar halo of the SMC.}"
    },
    "1112/1112.1913_arXiv.txt": {
        "abstract": "The quest for the supernova explosion mechanism has been one of the outstanding challenges in computational astrophysics for several decades. Simulations have now progressed to a stage at which the solution appears close and neutrino and gravitational wave signals from self-consistent explosion models are becoming available.  Here we focus one of the recent advances in supernova modeling, the inclusion of general relativity in multi-dimensional neutrino hydrodynamics simulations, and present the latest simulation results for an $11.2 M_\\odot$ and a $15 M_\\odot$ progenitor. We also mention 3D effects as another aspect in supernova physics awaiting further, more thorough investigation. ",
        "introduction": "Massive stars end their lives as a core-collapse supernova (SN), a violent event that involves the collapse of the iron core of the progenitor to a proto-neutron star and the subsequent expulsion of the outer layers of the star with a kinetic energy on the order of $10^{51} \\ \\mathrm{erg}$, which is associated with a spectacularly bright optical display. Currently, there is still no final consensus on the supernova explosion mechanism that operates in the optically obscured supernova core, and a number of competing ideas are under discussion.  The delayed neutrino-driven mechanism \\cite{wilson_85,bethe_85}, which relies on neutrino energy deposition in the gain region to revive the stalled shock, remains the most promising candidate, provided that the efficiency of neutrino heating can be sufficiently enhanced by multi-dimensional hydrodynamical instabilities such as convection and the so-called standing accretion shock instability SASI \\cite{blondin_03,foglizzo_06}. This mechanism has worked successfully in several recent 2D simulations \\cite{buras_06_b,marek_09,bruenn_10} (some of which appeared to be only marginally successful \\cite{marek_09}), but has failed in others \\cite{burrows_06,burrows_07}. Alternatives to the neutrino-driven mechanism have also been proposed, such as the acoustic mechanism \\cite{burrows_06,burrows_07} (whose viability has been called into question by \\cite{weinberg_08}, however), magnetohydrodynamically driven supernovae \\cite{akiyama_03,dessart_07_a}, and explosions triggered by a QCD phase transition \\cite{sagert_09}.  As the ``engine'' driving this explosion is not directly accessible by classical, photon-based astronomical observations, our understanding of the supernova explosion mechanism has largely rested on numerical simulations in the past ever since the pioneering work of \\cite{colgate_66}. Over the years, a variety of ambitious numerical approaches has been developed to cope with the challenging interplay of neutrino transport, multidimensional hydrodynamics, general relativity (GR), neutrino physics, and nuclear physics in the supernova problem. The currently most advanced models rely on sophisticated multi-group neutrino transport schemes (e.g. ray-by-ray variable Eddington factor transport \\cite{rampp_02, buras_06_a}, ray-by-ray-diffusion \\cite{bruenn_10}, 2D multi-angle transport without energy bin coupling \\cite{ott_08_a}, or the isotropic diffusion source approximation \\cite{liebendoerfer_09}) with different strengths and weaknesses, and there have only been very tentative attempts to venture forth to 3D models with these methods \\cite{bruenn_10,takiwaki_11}. Elements missing or only partially included in current state-of-the-art-models, such as 3D effects (whose potential is presently being debated \\cite{nordhaus_10,takiwaki_11,hanke_11}) or general relativity may hold the key to a better understanding of the explosion mechanism. Moreover, an improved treatment of such as yet poorly explored aspects is also indispensable for accurate predictions of the neutrino and gravitational wave signal -- the only \\emph{observables} that directly probe the dynamics in the supernova core.  Among the aspects that have not yet been thoroughly investigated in self-consistent multi-D neutrino hydrodynamics simulations of core-collapse supernovae, our group has recently begun to study the influence of GR in more detail. Although the importance of relativistic effects in core-collapse supernovae (due to the compactness of the proto-neutron star and the occurrence of high velocities) has long been recognized an demonstrated \\cite{bruenn_01}, the combination of GR hydrodynamics and multi-group neutrino transport has long been feasible only in spherical symmetry \\cite{bruenn_01,yamada_99,liebendoerfer_04}.  With the relativistic generalization of the ray-by-ray variable Eddington factor method \\cite{mueller_10} used in our neutrino hydrodynamics code \\textsc{Vertex}, we are now able to present first results about the impact of GR on the explosion dynamics and, in particular, the neutrino and gravitational wave emission in axisymmetric (2D) supernova models. ",
        "conclusions": ""
    },
    "1112/1112.1408_arXiv.txt": {
        "abstract": "We present \\Hersc~observations of 62 early-type galaxies (ETGs), including 39 galaxies morphologically classified as \\SO and 23 galaxies classified as ellipticals using SPIRE at 250, 350 and 500\\,$\\mu$m as part of the volume-limited \\Hersc~Reference Survey (HRS).  We detect dust emission in 24\\% of the ellipticals and 62\\% of the S0s. The mean temperature of the dust is $\\langle T_d \\rangle =23.9 \\pm 0.8\\rm \\, K$, warmer than that found for late-type galaxies in the Virgo Cluster. The mean dust mass for the entire detected early-type sample is ${\\rm log}M_d = 6.1 \\pm 0.1\\, M_{\\odot}$ with mean dust-to-stellar mass ratio of ${\\rm log}(M_d/M_*)= -4.3 \\pm 0.1$.  Including the non-detections, these parameters are ${\\rm log}M_d = 5.6 \\pm 0.1$ and ${\\rm log}(M_d/M_*)= -5.1 \\pm 0.1$ respectively. The average dust-to-stellar mass ratio for the early-type sample is fifty times lower, with larger dispersion, than the spiral galaxies observed as part of the HRS, and there is an order of magnitude decline in $M_d/M_*$ between the S0s and ellipticals. We use UV and optical photometry to show that virtually all the galaxies lie close to the red sequence yet the large number of detections of cool dust, the gas-to-dust ratios and the ratios of far-infrared to radio emission all suggest that many ETGs contain a cool interstellar medium similar to that in late-type galaxies.  We show that the sizes of the dust sources in S0s are much smaller than those in early-type spirals and the decrease in the dust-to-stellar mass ratio from early-type spirals to S0s cannot simply be explained by an increase in the bulge-to-disk ratio. These results suggest that the disks in S0s contain much less dust (and presumably gas) than the disks of early-type spirals and this cannot be explained simply by current environmental effects, such as ram-pressure stripping.  The wide range in the dust-to-stellar mass ratio for ETGs and the lack of a correlation between dust mass and optical luminosity suggest that much of the dust in the ETGs detected by \\Hersc~has been acquired as the result of interactions, although we show these are unlikely to have had a major effect on the stellar masses of the ETGs.  The \\Hersc~observations tentatively suggest that in the most massive systems, the mass of interstellar medium is unconnected to the evolution of the stellar populations in these galaxies. ",
        "introduction": "\\label{sec:intro}  In the standard theoretical view of galactic evolution, based on the hierarchical paradigm, galaxies may move both ways along the Hubble sequence during their evolution: ellipticals may form as the result of the merging of late-type galaxies (e.g., Cole et al. 2000; De Lucia et al. 2006); disks form as the result of the accretion of gas on to bulges (Cole et al. 2000). There are also plenty of environmental processes that might transform the morphology of a galaxy, just one example among many being ram-pressure stripping in clusters, which would move a galaxy toward the early-type sequence by removing the interstellar medium (ISM; e.g., Corbelli et al. 2011) and thus quenching star formation (Boselli \\& Gavazzi 2006).  The observational evidence that galaxy transformation is occurring is mixed, especially for galaxies at the current epoch.  The fact that the morphology--density relation is not just the result of early-type galaxies (ETGs)---ellipticals and S0s---being more common in clusters, but extends over a wide range of environmental density, is evidence that galaxy transformation is not a significant process at the current epoch (Dressler 1980).  However, there is also plenty of observational evidence that merging is still important today.  In particular, there is evidence of recent merging or accretion in nearby ellipticals, primarily from obscuration in the optical (Goudfrooij \\& de Jong 1995), where fossils of mergers in the form of dust ripples and dust lanes have been detected (Schweizer \\& Seitzer 1992). Furthermore, the discovery with the {\\it Infrared Astronomical Satellite} ({\\it IRAS}) of the ultraluminous infrared galaxies (ULIRGs), which are often late-stage mergers with optical profiles similar to the profiles of ellipticals (Wright et al. 1990), was persuasive evidence that ellipticals are still being formed by mergers at the current epoch.  Nevertheless, it is not clear whether all the properties of the eventual merged systems would be the same as present-day ellipticals (Naab \\& Ostriker 2009), and many of the distinctive properties of ellipticals seem more naturally explained if ellipticals form in a relatively short early-period in cosmic history (Peebles \\& Nusser 2010).  The properties of the ISM in ETGs potentially have something to tell us about the evolution of these systems. Since late-type galaxies are known to contain much more gas and dust than ETGs, any process that moves a galaxy along the Hubble sequence has to change simultaneously the morphology and the ISM of the galaxy.  In this paper, we present \\Hersc~observations of continuum emission from the dust in a sample of the most massive ETGs in the nearby universe, which we show is the most sensitive method currently available for estimating the mass of the cool ISM in these galaxies.  Although once thought to be devoid of a cool ISM, optical absorption studies suggested that more than 50\\% of ellipticals contain some dust (Goudfrooij et al.\\ 1994; van Dokkum \\& Franx 1995; Tran et al.\\ 2001; Ferrarese et al. 2006) and therefore molecular gas, though the mass of dust is uncertain from these works. In the last decade, previous far-infrared observations with {\\it IRAS}, the {\\it Infrared Space Observatory} ({\\it  ISO}), {\\it Spitzer}, and {\\it AKARI}\\ have shown that dust in ETGs is rather common, with dust masses estimated at some 10--100 times larger than those derived from the optical extinction measurements (e.g., Bregman et al.\\ 1998; Ferrari et al.\\ 2002; Kennicutt et al.\\ 2003; Pahre et al.\\ 2004; Xilouris et al.\\ 2004; Temi et al.\\ 2004, 2007a; Panuzzo et al.\\ 2007; Kaneda et al.\\ 2008; Young et al. 2009). Although these observations have probed the peak of the spectral energy distributions (SEDs) of the brightest giant ellipticals and S0s, they are hampered by poor resolution and/or lack of long wavelength coverage.  Our current ideas about the life cycle of dust in ellipticals suggest that it may be possible for us to test the origin of the dust, particularly whether it is provided by internal (via stellar mass loss) or external processes (e.g., fueled by mergers). As long as dust is not introduced from elsewhere, we essentially have a stellar system in equilibrium: dust is produced in the atmospheres of evolved stars (Athey et al.\\ 2002) and destroyed via sputtering in the hot gas (indeed Bressan et al. 2006 detect silicate features in elliptical galaxies known to be produced in circumstellar envelopes). The dust life cycle in ellipticals may therefore be simpler than that in late-types for the internal origin scenario (Tsai \\& Mathews 1996), and the mass of dust predicted to exist in the steady state through production in evolved stars and destruction through sputtering is approximately $10^5\\,M_\\odot$ (Forbes 1991; Goudfrooij \\& de Jong 1995).  An alternative origin for dust in ETGs is accreted material either due to a merger or from a tidally-interacting companion: Centaurus A has long been known to harbor a dusty disk thought to be formed by a merger (recently revealed in the submillimeter (submm)---Leeuw et al. 2002; Auld et al.\\ 2011; T. Parkin et al.\\ 2011, submitted). Forbes (1991) and Temi et al. (2004, 2007a) found no correlation between the dust emission and optical starlight using {\\it ISO} observations of massive ETGs, suggesting that the dust has an external rather than an internal origin. In 2010, Gomez et al.\\ used \\Hersc~observations of the elliptical galaxy M86 to reveal $10^6\\,M_{\\sun}$ of cold dust coincident with material stripped from the nearby spiral NGC4438 (Cortese et al.\\ 2010a). A similar process has been seen with the ``displaced ISM'' of the elliptical NGC3077 (Walter et al.\\ 2011).  The unprecedented resolution and sensitivity of the recently launched \\Hersc~{\\it Space Observatory} (Pilbratt et al.\\ 2010), combined with the wavelength coverage of the instruments PACS (Poglitsch et al.\\ 2010) and SPIRE (Griffin et al.\\ 2010) from 70 to 500\\,$\\mu$m, allow us to address long-standing issues such as the origin and quantity of dust in ETGs. \\Hersc~provides us with an unbiased view of the interstellar dust (and therefore the interstellar medium); we are now sensitive to the {\\it total} dust mass in galaxies rather than only to those galaxies which had enough warm dust to be detected by {\\it IRAS} and {\\it Spitzer} (e.g., Dunne et al.\\ 2011).  Here, we use observations from the guaranteed time project, the \\Hersc~Reference Survey (HRS; Boselli et al.\\ 2010b) of 62 ETGs including 23 ellipticals. The HRS is a study of dust in nearby galaxies, crucial for calibrating blind surveys of galaxies with \\Hersc~at high redshifts, e.g., HerMES (Oliver et al.\\ 2010) and the \\Hersc-ATLAS (H-ATLAS; Eales et al.\\ 2010).  Cortese et al.\\ (2011b, hereafter C12) recently used the HRS sample of $\\sim$300 nearby galaxies to obtain dust scaling relations, finding that the dust is tightly coupled to the atomic mass and therefore to the cold ISM. C12 measured the dust-to-stellar-mass ratio of galaxies with different morphological types and environment, finding that the ETGs contain less dust mass per unit of stellar mass than the late-type galaxies.  In this work, we further explore the dust properties of the ETGs in the HRS and measure more precisely the variation of dust content with Hubble type and the origin of this variation.  In Section \\ref{sec:red} we introduce the sample and describe the data reduction techniques.  The results are given in Section \\ref{sec:res} and are used to discuss our understanding of the dust content of early- and late-types in Section\\,\\ref{sec:morp} and the link between dust and the hot interstellar medium in Section \\ref{sec:hot}. We discuss the implications for the evolutionary history of ETGs in Section \\ref{sec:disc}.  The conclusions are presented in Section \\ref{sec:conc}.   \\begin{deluxetable*}{lcllllrrrc} \\tabletypesize{\\scriptsize} \\tablecaption{The Sample \\label{tab:sample}} \\tablewidth{0pt} \\tablehead{ \\colhead{HRS} & \\colhead{Other Name} & \\colhead{R.A.} & \\colhead{Decl.} & \\colhead{Type} & \\colhead{$D$} & \\colhead{$D(25)$} & \\colhead{$L_B$} & \\colhead{$L_K$} & \\colhead{Membership} \\\\ \\colhead{} & \\colhead{}  & \\colhead{($\\rm h~~m~~s$)} & \\colhead{($^{\\circ} ~~\\arcmin~~\\arcsec$)} & \\colhead{} &\\colhead{(Mpc)} & \\colhead{($\\arcmin$)} & \\colhead{($L_{\\odot}$)}& \\colhead{($L_{\\odot}$)}&\\colhead{} } \\startdata 3 & NGC 3226 & 10 23 27.01$^r$ & +19 53 54.7$^r$ & E2:pec;LINER;Sy3 & 16.7 &  3.16& 10.12 & 10.59 & Leo Cl. \\\\ 7 & NGC 3245 & 10 27 18.40$^r$ & +28 30 26.3$^r$ & SA(r)0:?;H\\sc{ii};LINER & 18.8 &  3.24 &10.08 & 10.47&  Leo Cl. \\\\ 14 & NGC 3301 & 10 36 56.04& +21 52 55.7 & (R')SB(rs)0/a & 19.2 & 3.55& 9.84 & 10.47  &Leo Cl. \\\\ 22 & NGC 3414 &10 51 16.19$^r$& +27 58 30.2$^r$  & S0 pec;LINER & 20.2 & 3.55 & 9.98 & 10.73&Leo Cl. \\\\ 43 & NGC 3608 &11 16 58.96& +18 08 54.9  & E2;LINER: & 15.8 & 3.16 &10.11 &10.70 & Leo Cl. \\\\ 45 & NGC 3619 & 11 19 21.51$^r$ & +57 45 28.3$^r$ &  (R)SA(s)0+: & 22.1 &  2.69 & 9.87 & 10.57 &Ursa Major Cl.\\\\ 46 & NGC 3626 &11 20 03.80$^r$ & +18 21 24.3$^r$  & (R)SA(rs)0+ & 21.3 &  2.69 &10.12 & 10.70& Leo Cl. \\\\ 49 & NGC 3640 & 11 21 06.85 & +03 14 05.4  & E3  & 17.9 & 3.98 &10.43 &11.06 & Leo Cl. \\\\ 71 & NGC 3945 & 11 53 13.61$^r$ & +60 40 32.3$^r$ & SB(rs)0+;LINER & 18.0 &  5.25 & 10.01 & 11.06 &Ursa Major Cl. \\\\ 87 & NGC 4124 & 12 08 09.64 & +10 22 43.4 & SA(r)0+ & 17.0 & 4.10 &9.70 & 10.38 &Virgo Out. \\\\ 90 & NGC 4179 & 12 12 52.11 &\t+01 17 58.9 & S0\\tablenotemark{a} & 17.0 & 3.80 & 9.89 & 10.60 & Virgo Out.\\\\ 93 & NGC 4203 & 12 15 05.06$^r$ & +33 11 50.2$^r$ & SAB0--:;LINER;Sy3 & 15.6 & 3.39 & 9.89 & 10.73& Coma I Cl. \\\\ 101 & NGC 4251 & 12 18 08.31 & +28 10 31.1 &  SB0? sp & 15.3 & 3.63 & 13.86 & 10.54&Coma I Cl. \\\\ 105 & NGC 4262 & 12 19 30.58 & +14 52 39.8 & SB(s)0--? & 17.0 & 1.87 & 9.69 & 10.43& Virgo A \\\\ 123 & NGC 4324 & 12 23 06.18 & +05 15 01.5 &  SA(r)0+ & 17.0 &  3.52 &9.65 & 10.38& Virgo S Cl.\\\\ 125 & NGC 4339 & 12 23 34.94 & +06 04 54.2 & E0;Sy2 & 23.0 &  2.31 & 9.71 & 10.34&Virgo B \\\\ 126 & NGC 4340 & 12 23 35.31 & +16 43 19.9 &SB(r)0+ & 17.0 &  3.60 & 9.81 &10.44 &Virgo A \\\\ 129 & NGC 4350 & 12 23 57.81 & +16 41 36.1 & SA0;Abs. line & 17.0 &  3.20 &  9.91 & 10.64&Virgo A \\\\ 135 & NGC 4365 & 12 24 28.23 & +07 19 03.1 & E3 & 23.0 & 8.73 & 10.34 & 11.26 &Virgo B \\\\ 137 & NGC 4371 & 12 24 55.43 & +11 42 15.4 & SB(r)0+ & 17.0 & 5.10 &9.92 & 10.68& Virgo A \\\\ 138 & NGC 4374, M84 & 12 25 03.78 & +12 53 13.1 & E1;LERG;Sy2 & 17.0 & 10.07 &10.57 & 11.26 & Virgo A \\\\ 150 & NGC 4406, M86 & 12 26 11.74 & +12 56 46.4 & S0(3)/E3 & 17.0 &  11.37 &10.66 & 11.31& Virgo A \\\\ 155 & NGC 4417 & 12 26 50.62 & +09 35 03.0 & SB0: s & 23.0 & 3.60 & 10.08 & 10.77 & Virgo B \\\\ 161 & NGC 4429 & 12 27 26.56 & +11 06 27.1 & SA(r)0+;LINER;H\\sc{ii} & 17.0 &  8.12 & 10.23 & 11.06& Virgo A \\\\ 162 & NGC 4435 & 12 27 40.50$^r$ & +13 04 44.5$^r$ & SB(s)0;LINER;H\\sc{ii} & 17.0  & 2.92 & 10.05 &10.83 &  Virgo A \\\\ 166 & NGC 4442 & 12 28 03.89 & +09 48 13.0 & SB(s)0 & 23.0 & 5.05 &10.36 & 11.12&  Virgo B \\\\ 174 & NGC 4459 & 12 29 00.04$^r$ & +13 58 42.2$^r$ & SA(r)0+;H\\sc{ii};LINER & 17.0 & 3.36   &10.07 & 10.91 &Virgo A \\\\ 175 & NGC 4461 & 12 29 03.01 & +13 11 01.5 & SB(s)0+: & 17.0 &  3.52 & 9.90 & 10.57&Virgo A \\\\ 176 & NGC 4469 & 12 29 28.03 & +08 44 59.7 & SB(s)0/a? sp & 23.0 & 4.33  &9.97 & 10.82 &Virgo B \\\\ 178 & NGC 4472, M49 & 12 29 46.76 & +08 00 01.7 & E2/S0;Sy2 & 17.0 &  10.25  & 10.90 & 11.59 &Virgo S Cl. \\\\ 179 & NGC 4473 & 12 29 48.87 & +13 25 45.7 & E5 & 17.0 &  4.04 & 10.15 & 10.90& Virgo A \\\\ 180 & NGC 4477 & 12 30 02.17 & +13 38 11.2 & SB(s)0:?;Sy2 & 17.0 & 3.60 &10.12 & 13.77& Virgo A \\\\ 181 & NGC 4478 & 12 30 17.42 & +12 19 42.8 & E2 & 17.0 &  1.89 &9.89 & 10.41 & Virgo A \\\\ 183 & NGC 4486, M87 & 12 30 49.42 & +12 23 28.0 & E+0-1 pec;NLRG;Sy & 17.0 & 11.00 & 10.85 & 11.43 &Virgo A  \\\\ 186 & NGC 4494 & 12 31 24.03 & +25 46 29.9 & E1-2;Sy & 18.7 & 4.79 &10.62 & 11.20 & Coma I Cl. \\\\ 200 & NGC 4526 & 12 34 03.03$^r$ & +07 41 57.3$^r$ & SAB(s)0: & 17.0 &  7.00 & 10.41 & 11.18&Virgo S Cl. \\\\ 202 & IC 3510 & 12 34 19.33& +11 04 17.7 & dE\\tablenotemark{a} & 17.0 &1.10 &8.70 & 9.20 & Virgo A  \\\\ 209 & NGC 4546 & 12 35 29.51 & --03 47 35.5 & SB(s)0--: & 15.0 & 3.31 &10.05 & 10.82& Virgo Out. \\\\ 210 & NGC 4550 & 12 35 30.61 & +12 13 15.4 & SB0: Sy,LINER & 17.0 & 3.95 &9.66 & 10.30& Virgo A \\\\ 211 & NGC 4552, M89 & 12 35 39.88 & +12 33 21.7 & E;LINER;HII;Sy2 & 17.0 &  7.23  &10.29 & 11.06 &Virgo A \\\\ 214 & NGC 4564 & 12 36 26.99 & +11 26 21.5 & E6 & 17.0 &  4.33 & 9.86 & 10.58& Virgo A \\\\ 218 & NGC 4570 & 12 36 53.40 & +07 14 48.0 & S0(7)/E7 & 17.0 & 3.52 & 9.96\t&10.70 &Virgo S Cl.\\\\ 219 & NGC 4578 & 12 37 30.55 & +09 33 18.4 & SA(r)0: & 17.0 & 3.77 &9.71 & 10.41& Virgo E Cl.\\\\ 231 & NGC 4596 & 12 39 55.94 & +10 10 33.9 & SB(r)0+;LINER: & 17.0 & 4.76 &10.08 & 10.79& Virgo E Cl.\\\\ 234 & NGC 4608 & 12 41 13.29 & +10 09 20.9 & SB(r)0 & 17.0 & 4.30 & 9.83 & 10.51& Virgo E Cl.\\\\ 235 & NGC 4612 & 12 41 32.76 & +07 18 53.2 & (R)SAB0 & 17.0 & 2.16 & 9.82 & 10.35& Virgo S Cl.\\\\ 236 & NGC 4621, M59 & 12 42 02.32 & +11 38 48.9 & E5 & 17.0 & 7.67  &10.32& 11.05& Virgo E Cl.\\\\ 240 & NGC 4638 & 12 42 47.43 & +11 26 32.9 & S0- & 17.0 & 2.01 &9.82 & 10.49 &Virgo E Cl. \\\\ 241 & NGC 4636 & 12 42 49.87 & +02 41 16.0 & E/S0/1;LINER;Sy3 & 17.0 & 9.63 &10.51 & 11.18& Virgo S Cl.\\\\ 243 & NGC 4643 & 12 43 20.14 & +01 58 42.1 & SB(rs)0/a;LINER/H\\sc{ii} & 17.0 & 3.00 & 9.98 &10.81& Virgo Out.\\\\ 245 & NGC 4649, M60 & 12 43 40.01 & +11 33 09.4 &  E2 & 17.0 & 5.10 & 10.73 & 11.46 &Virgo E Cl.\\\\ 248 & NGC 4660 & 12 44 31.97 & +11 11 25.9 & E5  & 17.0 & 1.89 & 9.70 & 10.47& Virgo E Cl.\\\\ 250 & NGC 4665 & 12 45 05.96 & +03 03 20.5 & SB(s)0/a & 17.0 & 4.50 & 10.04 & 10.80 &Virgo Out.\\\\ 253 & NGC 4684 & 12 47 17.52 & --02 43 38.6 & SB(r)0+;H\\sc{ii}  & 21.3 & 2.88 &9.82 & 10.42 &Virgo Out.\\\\ 258 & NGC 4697 & 12 48 35.91 & --05 48 03.1 & E6;AGN  & 17.7 & 7.24 & 10.55 & 11.16 &Virgo Out.\\\\ 260 & NGC 4710 & 12 49 38.93$^r$ & +15 09 59.1$^r$ & SA(r)0+? sp;H\\sc{ii} & 17.0 & 4.30 & 9.91 & 10.74 &Virgo Out.\\\\ 269 & NGC 4754 & 12 52 17.56 & +11 18 49.2 & SB(r)0--: & 17.0 & 5.03 &10.05 & 10.81& Virgo E Cl.\\\\ 272 & NGC 4762 & 12 52 56.05 & +11 13 50.9 & SB(r)0 sp;LINER & 17.0 & 8.70  &10.21 & 10.85 &Virgo E Cl.\\\\ 286 & NGC 4866 & 12 59 27.14 & +14 10 15.8 & SA(r)0+;LINER & 17.0 & 6.00 &9.84 & 10.60& Virgo Out.\\\\ 296 & NGC 5273 & 13 42 08.36$^r$ & +35 39 15.1$^r$ & SA(s)0;Sy1.5 & 15.2 &  2.75 &9.54 & 10.21 &Canes Ven. Spur \\\\ 312 & NGC 5576 & 14 21 03.68 & +03 16 15.6 & E3  & 21.2 & 3.55 &10.16 & 10.89&Virgo-Libra Cl.\\\\ 316 & NGC 5638 & 14 29 40.39 & +03 14 00.2 & E1 & 23.9 & 2.69 &10.09 & 10.72& Virgo-Libra Cl. \\enddata \\tablecomments{The columns are as follows. Column 1: the number of the source in the \\Hersc~Reference Survey (Boselli et al.\\ 2010b). Column 2: other common names for the galaxy. Columns 3 and 4: the right ascension and declination of the galaxy (J2000). This is taken from the NASA Extragalactic Database (NED), except positions with a superscript $r$ are radio positions taken from Wrobel (1991) and Filho et al. (2006). Column 5: morphological type taken from Virgo Cluster Catalogue (Binggeli et al. 1985) or from our own classification. The AGN/Seyfert/LINER classifications were taken from Ho et al. (1997), Schmitt (2001), and  Veron-Cetty \\& Veron (2006).  Column 6: distance in $\\rm Mpc$. We assume all objects in Virgo, including those in the outskirts of Virgo, are at a distance of $17 \\,\\rm Mpc$, except for those in Cloud B, which we assume are at a distance of $23 \\,\\rm Mpc$. For the other galaxies we have calculated distances from the heliocentric velocity using a Hubble constant of $70\\rm km\\ s^{-1}\\ Mpc^{-1}$. Column 7: optical isophotal diameter (25\\,mag\\,arcsec$^{-2}$). Column 8: total $K$-band luminosity measured using $K_{S_{\\rm tot}}$ from 2MASS (Skrutskie et al. 2006).  Column 9: cluster or cloud membership from Gavazzi \\& Boselli (1999) for Virgo, otherwise from Tully (1988) or Nolthenius (1993) wherever available, or failing that our own estimate. } \\tablenotetext{a}{These sources were morphologically classified by hand using optical images.} \\end{deluxetable*} ",
        "conclusions": "\\label{sec:conc}  We present the FIR and submm observations with \\Hersc~of 62 ETGs, including 23 ellipticals and 39 \\SO galaxies. We find the following results:  \\begin{enumerate} \\item We detect 24\\% of the ellipticals and 62\\% of the S0s.  The optical and X-ray luminosities of the detected sources show we are detecting the most massive ellipticals, though the S0 sample is representative of the general population. Of the ten `pure' ellipticals in our sample only one is detected by \\Hersc\\ (M86), while six of the eleven `impure' ellipticals are detected. The detection rate for the ETGs outside Virgo appears to be higher than for those inside the cluster but this is not statistically significant.  \\item The mean dust masses for the detected galaxies are ${\\rm log}M_d = 6.3 \\pm 0.1$ and $5.5 \\pm 0.3\\,M_{\\sun}$ for the \\SO and E population, with average dust-to-stellar-mass ratios of ${\\rm log}(M_d/M_*)= -4.1 \\pm 0.1$ and $-5.3 \\pm 0.2$.     Including the upper limits, the average dust masses for the detected {\\em and} undetected \\SO and elliptical sources are ${\\rm log}M_d (M_{\\odot}) =  5.9 \\pm 0.1$ and $5.2 \\pm 0.1$, with dust-to-stellar-mass ratios of ${\\rm log}(M_d/M_*)= -4.4 \\pm 0.1$ and $-5.8 \\pm 0.1$.   \\item  The mean dust temperature for the ETGs {\\em detected} by \\Hersc~is $\\sim 24\\,\\rm K$, warmer than the dust in late-type galaxies.    The {\\em entire} early-type sample, including non-detections, has a dust mass of ${\\rm log}M_d = 5.6 \\pm 0.1$ and ${\\rm log}(M_d/M_*)= -5.1 \\pm 0.1$.  In comparison, the average dust-to-stellar mass ratio for spiral galaxies in the HRS is $-2.59 \\pm 0.03$. The dispersion in the dust-to-stellar mass ratio is much greater for ETGs than spirals.  \\item The NUV--$r$ colors show that virtually all the ETGs lie close to the red sequence, but there is evidence from UV and optical colors, radio continuum observations, and literature searches that a $\\sim 20\\%$ of the sample have had a recent star-forming epoch or have significant residual star formation.  However, the NUV colors have not been corrected for extinction or contamination from the old stellar population. Our ETGs are redder in NUV--$r$ with lower dust-to-stellar mass ratios than the ETGs detected in the H-ATLAS survey.  \\item We show that the detection of cold dust, the ratio of far-infrared to radio emission, and the gas-to-dust ratios all indicate that many ETGs contain a cool ISM in which stars are forming, similar to that seen in spiral galaxies.  \\item We find no evidence for a correlation between the FIR luminosity or dust mass of ETGs with the optical luminosity, suggesting that the main source of the dust in the galaxies detected by \\Hersc~is not mass loss from evolved stars.  This, together with the large spread in the dust-to-stellar-mass ratio of ETGs, suggests a significant fraction of the dust in these galaxies is acquired externally via mergers or tidal interactions.  \\item We use the results from the ATLAS$^{\\rm 3D}$ survey to divide our sample into fast and slow rotators.  We show that the difference in dust properties between S0s and ellipticals is more obvious than the difference between slow and fast rotators. We suggest that this may be additional evidence that the dust in ETGs has been acquired by gravitational encounters and may not tell us much about the evolutionary state of the galaxy.  \\item The low dust-to-stellar-mass ratios of S0s compared with early-type spirals cannot be explained by the larger bulge-to-disk ratios of S0s. The relative sizes of the dust sources in S0s are also smaller than seen in early-type spirals. These suggest that the disks in S0s contain much less dust (and presumably gas) than spiral disks with a similar size.  This effect is probably not being caused by current environmental processes adding to the weight of evidence that early-type spirals are not being transformed into S0s in significant numbers at the current epoch.  \\item If the cool ISM in ETGs is acquired as the result of a tidal interaction or merger, an upper limit on the increase in the stellar mass of the ETG due to the interaction is 1\\% for the ellipticals and 10\\% for the S0s, suggesting that the interactions are not significant galaxy-building events.  \\end{enumerate}  Our sample of ETGs has lower dust-to-stellar-mass ratios and dust masses compared to previously published \\Hersc~samples.  The results from our sample are important for interpreting chemical evolution models which model the production and destruction of dust in ellipticals as well as providing a low-redshift benchmark for understanding the evolution of dust and gas in high redshift surveys (e.g., Rowlands et al.\\ 2012, Dunne et al.\\ 2011). Finally, we note that to obtain a better understanding of some of these issues it will be important to follow-up our statistical study with detailed observational studies of individual ETGs, for example with ALMA."
    },
    "1112/1112.2121_arXiv.txt": {
        "abstract": "We describe the Munich Difference Imaging Analysis pipeline \\linebreak that we developed and implemented in the framework of the Astro-WISE\\footnote{Astronomical Wide-field Imaging System for Europe} package to automatically measure high precision light curves of a large number of stellar objects using the difference imaging approach. Combined with programs to detect time variability, this software can be used to search for planetary systems or binary stars with the transit method and for variable stars of different kinds. As a first scientific application, we discuss the data reduction and analysis performed with Astro-WISE on the pre-OmegaTranS data set, that we collected during a monitoring campaign of a dense stellar field with the Wide Field Imager at the ESO 2.2\\,m telescope. ",
        "introduction": "\\label{sec:Intro} The Munich Difference Imaging Analysis Pipeline (MDia) is a software tool to perform high precision photometric measurements in crowded fields such as the Milky Way disk or the bulge of M31. MDia is based on the difference imaging technique proposed by \\cite{Tomaney96} and \\cite{Alard98}. Common applications are searches for transiting extra-solar planets, microlensing surveys, stellar rotation and oscillation studies as well as supernova projects.\\\\ The MDia tool convolves (degrades) a reference image (constructed from the images collected during the periods of best seeing of the observing campaign, see section \\ref{sec:MDIA_ReferenceFrameImplementation}) to match the seeing of each single image in the data set. Subtracting the convolved reference image from a single image, one obtains a difference image in which every constant source disappears, while variable objects are visible as positive or negative residuals. PSF-photometry performed on the reference image and on each difference image at the position of the detected sources allows to reconstruct the light curves of each object.\\\\ ",
        "conclusions": "\\label{sec:Conclusions} We described the Munich Difference Imaging Analysis pipeline. After explaining the algorithm we presented the Astro-WISE implementation of MDia. Since a complete description with code examples is beyond the scope of this article we recommend any potential user to further read the MDia manual).\\\\ In the last section we presented a subset of the results we obtained with the MDia tool on the pre-OmegaTranS data set. We showed some example light curves of variable stars that demonstrate the high precision that can be achieved with the difference imaging technique.\\\\"
    },
    "1112/1112.0312_arXiv.txt": {
        "abstract": "We present the Arizona CDFS Environment Survey (ACES), a recently-completed spectroscopic redshift survey of the Chandra Deep Field South (CDFS) conducted using IMACS on the Magellan-Baade telescope. In total, the survey targeted $7277$ unique sources down to a limiting magnitude of $R_{\\rm AB} = 24.1$, yielding $5080$ secure redshifts across the $\\sim \\! 30^{\\prime} \\times 30^{\\prime}$ extended CDFS region. The ACES dataset delivers a significant increase to both the spatial coverage and the sampling density of the spectroscopic observations in the field. Combined with previously-published, spectroscopic redshifts, ACES now creates a highly-complete survey of the galaxy population at $R < 23$, enabling the local galaxy density (or environment) on relatively small scales ($\\sim \\! 1$ Mpc) to be measured at $z < 1$ in one of the most heavily-studied and data-rich fields in the sky. Here, we describe the motivation, design, and implementation of the survey and present a preliminary redshift and environment catalog. In addition, we utilize the ACES spectroscopic redshift catalog to assess the quality of photometric redshifts from both the COMBO-17 and MUSYC imaging surveys of the CDFS. ",
        "introduction": "\\label{sec_intro}  \\setcounter{footnote}{3}   Building upon the initial X-ray observations of \\citet{giacconi01, giacconi02}, the Chandra Deep Field South (CDFS, $\\alpha = 03^{\\rm h}32^{\\rm m}25^{\\rm s}$, $\\delta = -27^{\\circ}49^{\\rm m}58^{\\rm s}$) has quickly become one of the most well-studied extragalactic fields in the sky with existing observations among the deepest at a broad range of wavelengths (e.g., \\citealt{giavalisco04}; \\citealt{rix04}; \\citealt{lehmer05}; \\citealt{quadri07}; \\citealt{miller08}; \\citealt{padovani09}; \\citealt{cardamone10}; \\citealt{xue11}; \\citealt{damen11}; Elbaz et al.\\ in prep). In the coming years, this status as one of the very deepest multiwavelength survey fields will be further cemented by the ongoing and upcoming extremely-deep observations with {\\it Spitzer}/IRAC and {\\it HST}/WFC3-IR as part of the Spitzer Extended Deep Survey (SEDS, PI G.\\ Fazio) and the Cosmic Assembly Near-IR Deep Extragalactic Legacy Survey (CANDELS, \\citealt{grogin11, koekemoer11}), respectively.    Despite the large commitment of telescope time from both space- and ground-based facilities devoted to imaging the CDFS, spectroscopic observations in the field have generally lagged those in other, comparably-deep extragalactic survey fields. For example, in the Extended Groth Strip, another deep field targeted by the SEDS and CANDELS programs, the DEEP2 and DEEP3 Galaxy Redshift Surveys (\\citealt{davis03, davis07, cooper11, cooper12, newman12}; Cooper et al.\\ in prep; see also \\citealt{weiner06}) have created a vast spectroscopic database, achieving an impressive $\\sim \\!  60\\%$ redshift completeness down to $R_{\\rm AB} = 24.1$ across more than $0.2$ square degrees and $\\gtrsim \\! 40\\%$ completeness covering a broader area of $\\sim \\! 0.5$ square degrees down to the same magnitude limit. Similarly, there have been a variety of spectroscopic efforts in the GOODS-N field including the Team Keck Redshift Survey \\citep[TKRS,][see also \\citealt{cooper11}]{wirth04} in addition to the independent work of various groups \\citep[e.g.,][]{lowenthal97, phillips97, cohen00, dawson01, cowie04, treu05, reddy06, barger08}. Together, these datasets provide spectroscopic redshifts for $> \\! 90\\%$ of sources down to $z_{\\rm F850LP} = 23.3$ \\citep{barger08}.   These large spectroscopic surveys add significant scientific utility to the associated imaging datasets, making the photometric constraints much more powerful.  For example, spectroscopic redshifts allow critical rest-frame quantities to be derived with increased precision. Furthermore, only through spectroscopy can assorted spectral and dynamical properties (such as the strengths and velocity widths of emission and absorption lines) be measured --- in this regard, the spectral databases provided by surveys such as DEEP2, DEEP3, and TKRS are especially powerful due to their uniform spectral range and resolution. Finally, only with the combination of high (and relatively uniform) sampling density, spatial coverage across a modestly-sized field (e.g., $\\gtrsim \\! 0.05$ square degrees), and moderately high-precision spectroscopic redshifts (i.e., $\\sigma_{z} < 0.01$) can the local galaxy density (or ``environment'') be measured across a broad and continuous range of environments \\citep{cooper05, cooper07}.   In contrast to the EGS and GOODS-N fields, the spectroscopic redshift completeness across the extended $30^{\\prime} \\times 30^{\\prime}$ area of the CDFS is modest, $\\lesssim \\! 25\\%$ down to a limiting magnitude of $R_{\\rm AB} = 23$ and $\\lesssim \\! 20\\%$ at $R_{\\rm AB} < 24$, despite some significant spectroscopic efforts in the field \\citep[e.g.,][]{lefevre04, szokoly04, vanzella05, vanzella06, vanzella08, mignoli05, ravikumar07, popesso09, balestra10, silverman10}.\\footnote{The PRIMUS program \\citep{coil11} has collected spectra for a substantial number of sources in a larger area surrounding the CDFS. However, the relatively low resolution ($R = \\lambda/\\Delta \\lambda \\sim 30$) of the prism-based spectroscopy limits the utility of the PRIMUS dataset for detailed studies of spectral properties (e.g., emission-line equivalent widths) and small-scale environment (e.g., on group scales).} Notably, many of these existing spectroscopic programs have focused their efforts on the smaller GOODS-S region and/or targeted optically-faint, higher-redshift $(z > 1.5)$ sources \\citep[e.g.,][]{dickinson04, doherty05, roche06, vanzella09}.      With the goal of creating a highly-complete redshift survey at $z < 1$ in the CDFS, the Arizona CDFS Environment Survey (ACES) utilized the Inamori-Magellan Areal Camera and Spectrograph \\citep[IMACS,][]{dressler11} on the Magellan-Baade telescope to collect spectra of more than $7000$ unique sources across a $\\sim \\! 30^{\\prime} \\times 30^{\\prime}$ region centered on the CDFS. In Sections \\ref{sec_design} and \\ref{sec_redux}, we describe the design, execution, and reduction of the ACES observations, with a preliminary redshift and environment catalog presented in Sections \\ref{sec_data} and \\ref{sec_environ}. Finally, in Section \\ref{sec_future}, we conclude with a discussion of future work related to ACES. Throughout, we employ a Lambda cold dark matter ($\\Lambda$CDM) cosmology with $w = -1$, $\\Omega_{\\rm m}$ = 0.3, $\\Omega_{\\Lambda} = 0.7$, and a Hubble parameter of $H_{0} = 100\\ h$ km s$^{-1}$ Mpc$^{-1}$. All magnitudes are on the AB system \\citep{oke83}, unless otherwise noted.      \\begin{figure} \\centering \\plotone{f1.eps} \\caption{The ACES target sampling rate as a function of $R$-band magnitude across the entire $\\sim \\! 30^{\\prime} \\times 30^{\\prime}$ COMBO-17/CDFS footprint. The target sampling rate is defined as the percentage of objects at a given $R$-band magnitude in the COMBO-17 photometric catalog that were observed by ACES. The dotted and dashed lines show the sampling rate when only considering sources classified as non-stellar in the COMBO-17 catalog$^{5}$ (dotted) and when accounting for sources with a published spectroscopic redshift in the literature (dashed -- see \\S \\ref{sec_design}). At bright magnitudes, $R_{\\rm AB} < 23$, ACES brings the target sampling rate in the CDFS to $> \\! 80\\%$.} \\label{fig_sampling_rate} \\end{figure}  \\stepcounter{footnote}\\footnotetext{To select the galaxy population, all sources classified as ``Star'' or ``WDwarf'' in the COMBO-17 photometric catalog of \\citet{wolf04} are excluded.}   \\vspace*{0.2in} ",
        "conclusions": "\\label{sec_future}  We present a spectroscopic survey of the Chandra Deep Field South (CDFS), conducted using Magellan/IMACS and aptly named the Arizona CDFS Environment Survey (ACES). The survey dataset includes $7277$ unique spectroscopic targets, yielding $5080$ secure redshifts, within the extended CDFS region. We describe in detail the design and implementation of the survey and present preliminary redshift and environment catalogs.  While this work marks a significant increase in both the spatial coverage and the sampling density of the spectroscopic observations in the CDFS, there remains much analysis of the ACES data to be completed in the future. In particular, work is presently underway to produce a relative throughput correction for the spectra, using observations of F stars that were targeted on many of the IMACS slitmasks. The F stars were observed with the same instrumental set-up as the science targets (i.e., same slitwidth, slitlength, etc.) and were included on multiple slitmasks, such that the spectra fell on each of the $8$ IMACS CCDs, allowing chip-to-chip variations in the throughput to be estimated.  As discussed in \\S \\ref{sec_design}, fainter targets ($R \\gtrsim 22.5$) were observed on multiple ($\\sim \\! 2$--$4$) slitmasks, with the goal of accumulating longer integration times. In the future, these data spanning different slitmasks will be combined, which will likely improve the survey's redshift success rate at fainter magnitudes. In parallel to this work, efforts to improve the spectral reduction procedures are currently underway, which should likewise improve the redshift completeness at all magnitudes. Likewise, future work will include extracting spectra and measuring redshifts for serendipitous detections of objects, an effort that could be helped greatly by the addition of data from different slitmasks. The completion of this ongoing work is expected to coincide with a final data release, including updated redshift and environment catalogs as well as all of the reduced IMACS spectra. Finally, analysis is underway to utilize the ACES data to study the correlations between star-formation history, morphology, and environment at $z < 1$, using the rich multiwavelength data in the CDFS and the increased sample size to improve upon previous efforts \\citep[e.g.,][]{elbaz07, capak07, cooper08a, cooper10a}.    \\vspace*{0.25in}"
    },
    "1112/1112.5192_arXiv.txt": {
        "abstract": "The International Virtual Observatory Alliance (IVOA) developed numerous interoperability standards during the last several years. Most of them are quite simple to implement from the technical point of view and even contain ``SIMPLE'' in the title. Does it mean that it is also simple to build a working VO resource using those standards? Yes and no. ``Yes'' because the standards are indeed simple, and ``no'' because usually one needs to implement a lot more than it was thought in the beginning of the project so the time management of the team becomes difficult. In our presentation we will start with a basic case of a simple spectral data collection. Then we will describe several examples of small'' technologically advanced VO resources built in CDS and VO-Paris and will show that many standards are hidden from managers' eyes at the initial stage of the project development. The projects will be: (1) the GalMer database providing access to the results of numerical simulations of galaxy interactions; (2) the full spectrum fitting service that allows one to extract internal kinematics and stellar populations from spectra of galaxies available in the VO. We conclude that: (a) with the existing set of IVOA standards one can already build very advanced VO-enabled archives and tools useful for scientists; (b) managers have to be very careful when estimating the project development timelines for VO-enabled resources. ",
        "introduction": "On the side of IVOA we now have a comprehensive set of standards\\citep{AGftITCG11} including: data formats (VOTable, \\citealp{VOTable}), VO resource description (Resource Metadata), data models for 1D spectra (Spectrum Data Model) and simulations (SimDM/SimDB) and much more complex and general Characterisation Data Model \\citep{CharDM}, Astronomical Data Query Language, protocols to access tabular data (TAP, \\citealp{TAP}), images (SIAP), and spectra (SSAP, \\citealp{SSAP}), a Simple Application Messaging Protocol \\citep{SAMP} that allows different VO tools and services to talk to each other, authorisation and authentication mechanisms, and others. Some other standards are still at different phases of development. Now it became possible to handle even very complex astronomical datasets in the Virtual Observatory, such as 3D spectroscopy \\citep{CBLM06,ChilVO08} and results of N-body simulations. Some of the standards even carry the adjective ``simple'' in their titles.  Is it really simple and fast to implement operational VO services using these standards? Unfortunately, many (mostly technical) aspects are often hidden from the view. And the implementation of these details may take significant amount of time. ",
        "conclusions": "Our conclusion is that in many cases with the existing set of IVOA standards one can already build very advanced VO-enabled archives and tools useful for scientists. However, project managers have to be very careful when estimating the project development timelines for VO-enabled resources as a significant manpower overhead may be needed."
    },
    "1112/1112.2584_arXiv.txt": {
        "abstract": "*{New approaches for data provenance and data management (DPDM) are required for mega science projects like the Square Kilometer Array, characterized by extremely large data volume and intense data rates, therefore demanding innovative and highly efficient computational paradigms. In this context, we explore a stream-computing approach with the emphasis on the use of accelerators. In particular, we make use of a new generation of high performance stream-based parallelization middleware known as InfoSphere Streams. Its viability for managing and ensuring interoperability and integrity of signal processing data pipelines is demonstrated in radio astronomy. \\newline\\indent IBM InfoSphere Streams embraces the stream-computing paradigm. It is a shift from conventional data mining techniques (involving analysis of existing data from databases) towards real-time analytic processing. We discuss using InfoSphere Streams for effective DPDM in radio astronomy and propose a way in which  InfoSphere Streams can be utilized for large antennae arrays. We present a case-study: the InfoSphere Streams implementation of an autocorrelating spectrometer, and using this example we discuss the advantages of the stream-computing approach and the utilization of hardware accelerators.}  \\abstract{New approaches for data provenance and data management (DPDM) are required for mega science projects like the Square Kilometer Array, characterized by extremely large data volume and intense data rates, therefore demanding innovative and highly efficient computational paradigms. In this context, we explore a stream-computing approach with the emphasis on the use of accelerators. In particular, we make use of a new generation of high performance stream-based parallelization middleware known as InfoSphere Streams. Its viability for managing and ensuring interoperability and integrity of signal processing data pipelines is demonstrated in radio astronomy. \\newline\\indent IBM InfoSphere Streams embraces the stream-computing paradigm. It is a shift from conventional data mining techniques (involving analysis of existing data from databases) towards real-time analytic processing. We discuss using InfoSphere Streams for effective DPDM in radio astronomy and propose a way in which  InfoSphere Streams can be utilized for large antennae arrays. We present a case-study: the InfoSphere Streams implementation of an autocorrelating spectrometer, and using this example we discuss the advantages of the stream-computing approach and the utilization of hardware accelerators.} ",
        "introduction": "\\label{sec:Introduction} Started in the 1930s, radio astronomy has produced some of the greatest discoveries and technology innovations of the 20th century.  One of these innovations -- radio interferometry and aperture synthesis -- was awarded a Nobel Prize for Physics in 1974 (Martin Ryle and Antony Hewish, 1974). An aperture synthesis radio telescope consists of multiple receiving elements in an array that observe the same radiating source(s) simultaneously. Essentially, an array of radio telescopes is used to emulate a much large telescope with size that of the diameter of the array, enabling a better angular resolution of the radio source(s) to be obtained. While angular resolution is determined by the array diameter, another important characteristic is telescope sensitivity, which is determined by its collecting area. The Square Kilometer Array (SKA) will be an aperture synthesis radio telescope, scheduled for completion in the 2020s, that will combine both factors, resolution and sensitivity. The total SKA collecting area of one square kilometer ($10^6$ m$^2$) will provide sensitivity that is 50-100 times higher than that of the best current radio telescope arrays. Its high angular resolution will be provided by distributing the square kilometer of collecting area into many stations that are spread out on a continental scale (with the baseline between some antennae over $3000$ km).  A radio telescope antenna element detects electromagnetic waves by a current induced in an antenna receiver system. This can be measured as a voltage $s_i(t)$ at receiver $i$ that is sampled and digitized at regular times $t$. Whereas a single receiver can measure the source brightness $I(\\vec{d})$ in a specific direction $\\vec{d}$, a pair of receivers $i,j$ separated by a baseline vector $\\vec{B}_{ij}$ can be used as an interferometer to measure the difference in phase between the signals $s_i$ and $s_j$ due to the time delay $\\tau_{ij}=\\vec{B}_{ij}\\cdot\\vec{d}/c$ between the received signals as illustrated in Figure~\\ref{interferometer}.  \\begin{figure}[ht] \\centering \\includegraphics[scale=1.0]{interferometer.pdf} \\caption{Two-element interferometer} \\label{interferometer} \\end{figure}  The time delay $\\tau_{ij}$ can be roughly approximated by the geometry of the antennae relative to the source direction (provided by an \\emph{Ephemerides service}), and more precisely determined by the resulting interference pattern in the cross-correlation between the signals: \\begin{displaymath} (s_i\\star s_j)(\\tau) = \\int_{-\\infty}^\\infty \\overline{s_i(t)}s_j(t+\\tau)\\,dt. \\end{displaymath} The value $V_{ij}=(s_i\\star s_j)(\\tau_{ij})$ is termed a \\emph{visibility} and gives a source brightness measurement at a $(u,v)$ point in the Fourier domain determined by the baseline $\\vec{B}_{ij}$ for that pair of antennae. An array of $n$ antennae has $\\frac{n(n-1)}{2}$ baselines (one per pair of antennae) and so $\\frac{n(n-1)}{2}$ visibilities can be obtained. However, if readings are taken over an interval of hours then each baseline changes over time due to the rotation of the Earth. Hence, over time the baselines sweep out elliptical arcs in the Fourier plane, illustrated in Figure~\\ref{EarthRotationSynthesis}.  \\begin{figure}[ht] \\centering \\includegraphics[scale=1.0]{uv_plot3.pdf} \\caption{A uv-plot shown for the VLBI array in Australia and New Zealand, including Australian Long Baseline Array (five radio telescopes), Warkworth (New Zealand), ASKAP (Western Australia) and AuScope antennas in Katherine and Yarragadee. The uv-plot results in $36\\times 2$ baseline tracks for a $4$ hour observation, $21$-cm wavelength and common source declination of $-70^{\\circ}$ \\cite{sweston}.} \\label{EarthRotationSynthesis} \\end{figure}  The Fourier domain coverage of an array is the combination of the $(u,v)$ tracks from all baselines provided by the array. It shows where the array samples the Fourier transform of the source image. For high quality imaging, it is desirable to have the best possible coverage of the Fourier domain, which is effectively the telescope aperture. A perfect source brightness distribution (the image of the radio source) could be obtained simply by taking the inverse Fourier transform if all $(u,v)$ points in the Fourier domain were able to be measured, but this is never the case. A deconvolution process such as \\emph{Clean} or \\emph{Maximum Entropy} used in any modern interferometer imaging can be thought of as a scheme for interpolating or extrapolating from the measured $(u,v)$ points to all other points in the $(u,v)$ plane \\cite{walker}.  \\begin{figure}[ht] \\centering \\includegraphics[scale=1.0]{pipeline.pdf} \\caption{Radio astronomy array pipeline} \\label{pipeline} \\end{figure}  Measuring the signals from all $n$ antennae over a period of hours results in an enormous dataset for a large array and its processing is a very compute intensive problem. Figure~\\ref{pipeline} shows the operations that are performed in an array on the digitized signals in a simplified pipeline from raw data through to analyzed data products. Computational power required for these operations can be very significant, particularly for the Correlation operation which calculates the visibility $V_{ij}$ for each of the $\\frac{n(n-1)}{2}$ baselines at each time $t$ via cross-correlations. It also performs an autocorrelation of each signal as discussed in Section \\ref{sec:AutocorrelationPipeline}, which together with the visibilities forms the \\emph{datacube} for the array at time $t$.  It is estimated that LOFAR with its 36 antennae stations can produce over 100 TB/day \\cite{schaaf}. For the SKA which will eventually have about 3000 antennae dishes, the data will increase by at least 5 orders of magnitude \\cite{varbanescu}. Such a huge amount of data places very high processing demands and requires a special approach to the overall organization of how data are processed and stored. It is only feasible to store the digitized raw signals required for calculating data cubes for small arrays and is limited to measurements taken over short time periods; in all other cases the data storage requirements are too large to be practical.        In the next section we introduce the stream computing paradigm and how IBM's InfoSphere Streams data management middleware utilizes this paradigm. In Section \\ref{sec:InfoSphereStreams} we describe how InfoSphere Streams can be applied to the operational facets of large radio astronomy telescope arrays to handle the enormous data volumes and compute intensive operations. In Section \\ref{sec:AutocorrelationPipeline} we consider an actual InfoSphere Streams application that performs streaming autocorrelations of actual radio astronomical observations. ",
        "conclusions": "\\label{sec:Conclusion}  This application successfully demonstrated the viability of implementing a real-time  PSD entirely in software using InfoSphere Streams. The SPADE application showed good data throughput without being specifically tailored to a specific accelerator, and allowed dynamic reconfiguration to allow more accelerators to be utilized as necessary or alternative types of accelerators included. Due to the use of standard SPADE operations the management of the data through the pipeline was transparent and the application could be easily extended to provide further analysis or provenance features.  The operations of an FX style autocorrelating spectrometer pipeline are dominated by the algorithmic complexity $n\\log_2(n)$ of the FFT operation. A single dual-PowerCell QS22 blade measured 30 GFlops for the FFT operation. In comparison measuring the performance of the FFT operation utilizing all four cores of a single x86 HS12 blade achieved almost 10 times less performance, and was measured at 3.5 GFlops. Essentially this shows that using various architectures for parallel computing by utilizing suitable accelerator hardware for specific compute intense operations can yield significant speed ups. In our case for an FX style autocorrelation pipeline a speed up of almost 10 times was achieved per QS22 blade.  IBM Info Streams proved it's flexibility to operate using various architectures in unison. Despite the I/O bound links Streams was capable of maximizing the link bandwidth as well as manage the data flow without information loss. Implementing the PSD pipeline in SPADE allowed parameters such as the integration time and FFT point size to be changed in real-time without compromising the flow of data.  Streams facilitates both implicit and explicit parallelization. Implicit parallelization is achieved by fusing operators into processing elements, and explicit parallelization by deployment of processing elements to many physical nodes. Furthermore Streams goes beyond conventional parallelization middleware and frameworks such as MPI (Message Passage Interface) and OpenMP by allowing dynamic operator fusing and processing element deployment to physical nodes during runtime. This degree of dynamic operation enables Streams to provide on demand scalability to increasing data loads and computations.  In this work we mainly focus on reviewing InfoSphere Streams and its potential use for Radio Astronomy. In our opinion the Streams approach has shown positive results to warrant further research and serious consideration for managing DPDM aspects of large antennae arrays. Further research is required to conduct more formal and specific comparative analysis between Streams and other middleware such as MPI and ICE (Internet Connection Engine). Another interesting area that requires more rigorous investigation is Streams scaling capabilities using a larger cluster of x86 nodes as well as combining other accelerators such as GPUs, Intel MIC (Many Integrated Core) architecture and FPGAs.   \\begin{acknowledgement} The first author would like to thank the New Zealand government's Tertiary Education Commission Build IT fund for funding this research and the IBM cooperation for providing the computing hardware under the shared university research grant. \\end{acknowledgement}"
    },
    "1112/1112.5186_arXiv.txt": {
        "abstract": "The number of cosmic strings in the observable universe is relevant in determining the probability to detect the presence of cosmic strings through their gravitational signatures. In particular, we refer to the observation of  gravitational lensing events and anisotropy in the CMB radiation induced by cosmic strings. In this paper, a simple method is adopted to obtain an approximate estimate of the number of segments of cosmic strings, crossing the particle horizon, which fall inside the visible (i. e. observed) part of the universe. We show that a different choice of the CDM cosmological model, which affects volumes, does actually have only a small weight on the expected number of cosmic string segments. Of this number, the fraction realistically detectable may be significantly smaller. ",
        "introduction": "\\label{intro} Cosmic strings are line-like topological defects which may have formed during a symmetry breaking phase transition in the early universe. The existence of topological defects in spacetime was first theorized by Kibble in 1976 \\cite{1}. Formation and evolution of cosmic defects were extensively studied in the subsequent decades \\cite{2,3,4,5} for their cosmological implications. Due to the compatibility of their existence with  cosmological observations and their topological stability, a special interest has always been focused on  cosmic strings \\cite{6,7,8}. This class of linear defects is a generic prediction in quantum field theory and grand unified theories (GUTs) as well as in string and M-theories. The predicted cosmic strings are extremely thin but very massive objects, characterized by a huge energy and interactions essentially of gravitational nature, whose intensity is measured by the dimensionless tension parameter $G \\mu / c^{2}$. The value of this quantity is model dependent; for a string produced at a GUT transition $G \\mu / c^{2}$ is of order $10^{-6}$ or $10^{-7}$. Although observations rule out the idea of topological defects as an alternative theory to inflation, the coexistence between cosmological inflation and cosmic strings is apparent in those inflationary scenarios which can themselves give rise to strings in a phase transition which occurs at the last stages of inflation. This is the case of hybrid inflation models in standard cosmology as well as brane/antibrane inflation models in string cosmology which are of hybrid inflation type. In this framework, a composite family of stable strings may exist in certain classes of models of warped compactifications or large compact extra dimensions which yield string tensions decoupled from Planck scale in the range $10^{-12} \\leq G\\mu / c^{2}\\leq 10^{-6}$ and hence compatible with observational limits \\cite{9,10,11}. Differently from solitonic strings, cosmic superstrings can exist as F and D strings leading to more complicated FD-networks involving the formation of junctions between strings of different tensions \\cite{12,13}. Although they can have different properties related to the different context in which they arise, cosmic superstring are expected to share with the ordinary solitonic strings the same scale-invariance property of the characteristic length scales of the network that they can form. Since their energy density may influence the dynamics of the universe, consistency with present day observations imposes constraints on the number density of cosmic strings. To date, cosmic (super)strings are still theoretical objects, nevertheless theoretical advances together with upcoming observational data might provide stronger constraints on their characteristic parameters and make their eventual detection a test of generic predictions of standard or alternative theories. Observations of such strings would provide direct information on fundamental physics and evolution of our universe and maybe the first experimental evidence of a string theory cosmological model underlying the structure of spacetime.  Cosmic (super)strings may have a number of cosmological effects. The most interesting observational signatures stem from their gravitational interactions. While specific signatures are expected in models where strings couple to other forces, gravitational effects are common to all cosmic (super)strings and are controlled by the dimensionless tension parameter. A special class of gravitational effects originates from the string peculiar way to deform the spacetime around it. The spacetime around a cosmic string is conical, namely, it is flat everywhere except for a missing wedge where points on the opposite edges are identified.Thus, spacetime looks like a cone with a singularity located at the apex of the cone. Any circular path at constant distance from the apex around the string is less than a circumference by the deficit angle $\\Delta = 8\\pi G\\mu /c^{2}$. This geometry leads to observable effects, such as the Kaiser-Stebbins effect (KS) \\cite{14} and gravitational lensing. The former appears as line discontinuities in the CMB temperature caused by a segment of cosmic string moving between the observer and the cosmological photosphere as photons passing by the string move perpendicularly to the plane that contains the string segment. Photons from the last scattering surface streaming by either side of the string will be observed with a redshift or blueshift (ahead of or behind the moving string respectively) proportional to the string tension and the string velocity transverse to the line of sight. Gravitational lensing by a long cosmic string occurs as light from a distant source beyond the string may reach the observer by two different paths around either side of the string thus producing a double image of the source with angular separation proportional to the conical deficit angle $8 \\pi G \\mu /c^{2}$.  The observational signatures of cosmic strings depend on the characteristic parameters and on the details of the evolution of the string network which in turn depend on what happens when two strings intersect. A fundamental feature which differentiates a network of cosmic superstrings from a network of ordinary field-theory cosmic strings is their reduced intercommutation probability $P$ which determines the evolution of a string network towards scaling regime. For FD-networks, $P$ can be significantly less than unity, depending on the type of strings and compactification and on the string coupling constant $g_{s}$. As a consequence of the smaller intercommutation probabilities associated with strings of different tensions, in a cosmic superstrings network each type of string may have a different number density. This means that the same fraction of CMB anisotropy can be sourced either by many light strings or by a few heavy ones but it also means that the total amount of strings in the network must be increased by some factor related to the intercommutation probabilities, which implies stronger constraints on tensions. The eventual presence of Y-shaped junctions of two strings joined by a third after intercommutation processes could be detected by cosmic string gravitational lensing or by observation of the KS effect. In the latter case, a different temperature in each of the three patches of sky is expected to be observed, while in the former case the relativistic motion of the binding strings could also lead to an enhancement of the cosmic string  lensing angle by some factor which would be moderate for a typical network motion \\cite{15,16,17,18}.  In the present paper we consider the case of classical cosmic strings for which the string tension is directly related to the energy scale of the symmetry breaking and the intercommutation probability is essentially $P = 1$. In fact, even though ultrarelativistic Abelian Higgs strings can pass through each other due to a double intercommutation with the net result that no ends are exchanged \\cite{19}, the typical velocities found in simulations for scaling networks of cosmic strings are below the ultrarelativistic threshold, so that the assumption $P = 1$ generally holds. On the other hand, the $rms$ velocities comparable to the speed of light found in numerical simulations  \\cite{20,21,22,23,24,25,26} are sufficiently high as to guarantee the detectability of their signature in the CMB radiation, if such strings do exist. The number of cosmic strings contained in the volume of the observable universe is also a relevant parameter to determine the probability of observing cosmic strings in gravitational lensing events as well as the density of step-like temperature discontinuities related to the KS effect \\cite{27,27bis,28}. In Section 2 we describe a simple method to obtain an approximate estimate of the number of segments of cosmic strings crossing the observable universe. In particular, for their observational relevance, we consider the volume inside the last scattering surface and inside the sphere of optical sources respectively for KS effect and gravitational lensing observations. We extend the previous approach to string networks embedded in cosmological backgrounds, variants of the $\\Lambda$CDM, where the constant parameter $\\Lambda$ is replaced with a dynamical scalar field. For a comparison of results, we consider, for simplicity, spatially flat cosmological models where the cosmic fluid is a mixture of matter and quintessence or a phantom constituent (a time varying component of the energy density of the universe with negative or super-negative pressure, $p < -\\rho$) characterized by the ratio of pressure to energy density in the range $-1 \\leq w < 0$ and $w < -1$, respectively, whereas $w = -1$ coincides with the cosmological constant case. In the above intervals we choose in particular $-1 \\leq w \\leq -0.5$ and $w = -1.5$ respectively for the quintessence and phantom components as they are compatible with constraints from large scale structure and CMB observations combined with SNe Ia data. In Section 3 we outline some features connected with numerical simulations and their underlying models in relation to the scaling property of the string network upon which relies the estimate of the number of string segments. We discuss our results with particular reference to the range of redshifts in the visible universe where the calculation of the number of string segments can be trusted, taking into account the informations about the scaling regimes provided by numerical simulations.  Without a loss of generality, in what follows we shall use the cosmological parameters of the $\\Lambda$CDM concordance model, which fit the majority of observations \\cite{29,30}:  \\begin{equation} \\label{eqn:1} \\Omega_{m} \\simeq 0.27   \\qquad   \\Omega_{\\Lambda} \\simeq 0.73  \\qquad  H_{0} \\simeq 71 km /s /Mpc. \\end{equation} \\noindent ",
        "conclusions": "The number of string segments $N(z)$ has been determined in the range of redshifts of interest for observations by using the constant quantity $A_{m}$ associated with the scaling law, thus assuming that the network has already reached the scaling regime in the matter-dominated era by the time $t_{LSS}$ which is the upper limit of the range of redshifts (\\ref{eqn:11}). As a matter of fact there is a transition period between the radiation-dominated and matter-dominated eras where the string network will be far from the type of equilibrium that characterizes a scaling solution for a time dependent on how fast is the network's response to the change in expansion rate. Thus, the range of redshifts in the matter-dominated era within which we can trust a ``string counting\" based on the scaling property could be well below $z_{LSS}$ and then below the range of radio surveys for searching of cosmic string signatures in the CMB radiation. A similar limit, this time at very small redshifts, namely $z \\approx 0.5$, is also to be taken into account since there is a further transition between the matter-dominated era and the dark-energy-dominated era \\cite{36,37}. The effect of this transition to a cosmological era where the universe expands exponentially will be a reduction of the string velocity, which prevents reconnection, affecting the scaling behavior and, at the same time, a dilution which might lead to a small number density of the leftover network \\cite{38,39}.  The scaling property, which implies it makes sense to consider the possibility of the  string existence in our universe and search for their observable effects, is the first fundamental information about a network of cosmic strings yielded by numerical simulations. These are essentially based on the field theory Abelian Higgs (AH) model and the Nambu-Goto (NG) model. The latter one relies on the assumption that as the scale of the string diameter is much smaller than any cosmological scale, such strings can be studied in the zero width approximation.  In particular, in NG simulations the time required for the long string component relaxation to the scaling regime in the matter-dominated era can appear decidedly longer than the transient taken to reach scaling in the radiation-dominated era so that the transition period might last far beyond the end of the decoupling time. Some difference can be found in the NG simulations results related to a different choice of the long component of the network, which can include all loops but those evolving on non-intersecting trajectories (stable loops that cannot rejoin the network) \\cite{26} or only super-horizon loops. In the latter case the long string component could reach the scaling regime in a relatively short time showing little time delay between the change in the expansion law and the response of the network. However, for what it concerns the loop component, its energy density also reaches a scaling evolution but the loop distribution takes more time to reach the scaling regime and the transition period is longer for the smaller loops \\cite{20,23}.  In AH simulations the effects of matter-radiation transition are not apparent in the network evolution. The first simulations yield approximately even the same value $A_{r} \\simeq A_{m} \\approx 0.3$ \\cite{24}, although this result is admittedly a consequence of a simulation size about the minimum required to study a scaling network of strings and recent improved simulations performed by the same authors yield a different result for the scaling parameters in matter- and radiation-dominated era, with $A_{r} \\simeq 0.26$ and $A_{m} \\simeq 0.29$ \\cite{40}.   The reason for differences between simulation results lies mainly in the model at the basis of the simulations. The AH model includes the small-scale physics near the string width, which has a significant role in the string dynamics. While in AH simulations energy from the strings is converted into massive gauge and Higgs radiation, in NG simulations this decay channel is not available and the string length density is significantly higher, with the long strings first converted into small loops and subsequently decaying via gravitational radiation. This additional decay channel in the field theoretic simulations by direct massive radiation is the primary energy-loss mechanism for long strings and causes AH strings to scale without gravitational radiation. In the NG scenario, particles are produced only near cusps, which determines much weaker constraints on the production of extremely energetic Standard Model particles (the Ultra-High Energy Cosmic Rays) than in the AH model. On the other hand, in the AH model there are no constraints from gravitational radiation whereas in the NG case the strongest limits come from millisecond pulsar timing which depend sensitively on details of the loop distribution and dynamics. The number of loops in the simulation volume is substantially less in field theory simulations than in NG where a population of non-intersecting loops remaining stable is found while in AH, where massive radiation is available, loops radiate and shrink. As regards the relevant component for CMB: the long (super-horizon) strings, the results yielded by the two types of simulations are quite different in most respects, such as the inter-string separation and therefore the string density which results significantly lower in AH than in NG simulations. Similarly, the $rms$ velocities probably due to backreaction from the massive radiation are lower in AH than in NG simulations. In fact, AH simulations give $\\approx 0.5 c$ \\cite{8} in both the radiation and matter eras whereas NG simulations give distinct values in the two cosmological eras both generally above $0.6 c$.  A specific problem related to field theory simulations is that computational constraints require the string width to be artificially increased in order to keep it above the simulation resolution. Although this does not significantly affect the CMB power spectra results a much relevant problem associated with the string width is that the 3D boxes required imply a reduction of the dynamical range. The fact that in AH simulations most of the network energy is emitted through field radiation, until the point that loop production is almost absent, has been supposed to be a consequence of a low resolution in AH simulation compared to NG ones. On the other hand, the accuracy of the NG simulations with regard to small-scale structure and loop production has also been questioned. In fact, the NG approximation fails close to the string width, that is, at cusps and kinks which are expected to form in an intersecting string network. Furthermore, once energy is transferred via intersection events from long strings to small loops, these are assumed to decay into gravitational radiation and are removed from the simulation since the gravitational waves are not included in any simulation for obvious reasons of complexity connected to backreaction on the string network.  \\begin{table} \\caption{Cosmic string parameters (where $v_{r}$ and $v_{m}$ are respectively the $rms$ velocities in the radiation- and matter-dominated era) from high resolution simulations performed from '90s to date. The results obtained by Battye-Moss (BM) refer to both USM modeling NG and AH strings. All the other results but those obtained by Bevis and coauthors (BHKU) refer to NG simulations.}\\label{tab:table_1} \\centering \\resizebox*{1.07\\textwidth}{!}{ \\begin{tabular}{|ccccccccc|} \\hline \\hline &&&&&&&& \\\\ Model & BB & AS & MS & RSB & BHKU (AH) & BM (NG/AH)& BPOS & $\\Delta z$ \\\\ par&\\cite{20}&\\cite{6,21,41}& \\cite{22}& \\cite{23} & \\cite{40}& \\cite{25}& \\cite{26} &\\\\ \\hline \\hline &&&&&&&&\\\\ $\\frac{G\\mu}{c^{2}} $ & $4 \\times 10^{-6}$  & $ 1.5 \\times  10^{-6}$  &  & $7 \\times 10^{-7}$ &$1.8 \\times 10^{-6}$ & $\\left(2.6/6.4\\right) \\times 10^{-7}$ & &\\\\ \\hline &&&&&&&&\\\\ $v_{r}$ & $0.66 c$  &  $0.62 c$  &  $0.63 c$ &          &  $0.5 c$  & $0.65 c$ / $0.4 c$ & $0.63 c$ &\\\\ \\hline &&&&&&&&\\\\ $v_{m}$ & $0.61 c$  &  $0.58 c$  &  $0.57 c$ &          &  $0.5 c$  & $0.60 c$ / $0.4 c$ & $0.59 c$ &\\\\ \\hline &&&&&&&&\\\\ $A_{r}$ & 0.14      &  0.13      &   0.13    & 0.16     &   0.26   & 0.13/0.35           & 0.15 &\\\\ \\hline &&&&&&&&\\\\ $A_{m}$ & 0.18      &  0.17      &   0.20    & 0.19     &   0.29   & 0.21/0.35           & 0.17 &\\\\ \\hline &&&&&&&&\\\\ $N(z)$  &5 -- 320   & 6 -- 380   & 4 -- 230 & 4 -- 270  &  1 -- 76  & 3 -- 200/1 -- 43   &6 -- 380& $ (0.5,7]$\\\\ \\hline &&&&&&&&\\\\ $N(z)$  &360 -- 1000 & 420 -- 1200 & 260 -- 740 & 300 -- 860&85 -- 240& 220 -- 640/48 -- 140&420 -- 1200&  $[8,100]$ \\\\ \\hline &&&&&&&&\\\\ $N(z)$ & (1300,2800) &  (1500,3400)   &  (950,3400) &  (1100,1800) & (310,430)  & (820,3400)/180 & (1500,2200) & 1100   \\\\ \\hline \\end{tabular} } \\end{table}   An alternative approach consists in using simulations of the NG or AH type and the \\textit{Unconnected Segment Model} (USM) \\cite{25}. The USM model represents the string network as a stochastic ensemble of unconnected moving segments, of length $A d_{PH}(t)$ and $rms$ velocity $v$, which are randomly removed at an appropriate rate so as to find sub-horizon decay and the proper string scaling density. The USM model with parameters measured in simulations has been used in the context of CMB anisotropies to reproduce the string power spectra coming from different simulation techniques and derive upper bounds on the string tension parameter. This approach also includes the radiation-matter transition and allows, in principle,  to study the effects of deviations from scaling at the onset of a $\\Lambda$-dominated phase.  As regards the effects on CMB of the network evolution during radiation-matter transition, which sets un upper bound to the applicability of eq. (\\ref{eqn:12}) at redshifts presumably even well below $z_{LSS} = 1100$, until about $z \\sim 100$, it has been shown \\cite{41} that the large-scale anisotropies are primarily produced at $z \\lesssim 100$. Since the coherence length of the string network grows with time, anisotropies on small angular scales are expected to be seeded at early times, and the large angle anisotropies to be seeded at late times. Thus, the contribution of cosmic strings in the redshift range $100 < z < 1100$ should only slightly affect the large angle CMB anisotropy so that the range of interest (\\ref{eqn:11}) in this case can be restricted to $z_{0} < z \\lesssim 100$. At small angular scales the relevant contribution to temperature anisotropy induced by cosmic strings comes from the last scattering epoch where one may expect the strings' signature in the CMB temperature fluctuations to be dominated by the \\textit{integrated Sachs-Wolfe} effect from the last scattering surface. It has been shown that at high multipoles the mean angular power spectrum of string-induced CMB temperature anisotropies can be described by a power law slowly decaying as $\\ell^{-p}$, with $p = 0.889^{+0.001}_{-0.090}$ \\cite{42}. This implies that the power spectrum at small angular scales decays much slower than the exponential Silk diffusion damping of the CMB primordial anisotropies. Consequently, the cosmic string contribution to the angular power spectrum can be subdominant at low $\\ell$ while it dominates the primary anisotropies for large values of $\\ell$. The exponential suppression of the inflationary contribution at high $\\ell$ also means that the string component may dominate for a realistic contribution (i. e. with normalization such that string contribution $f_{10}$ at multipole $\\ell = 10$ is $10\\%$ of the observed power at this multipole) at $\\ell \\gtrsim 3500$. The fraction of the total theoretical spectrum due to strings $f_{\\ell}$ is found to increase \\cite{40} from $f_{1500} \\approx 0.1$ to $f_{3500} = 0.5$,  which points out that a non negligible cosmic strings contribution to the CMB anisotropies is to be searched at small angular scales, corresponding to $\\ell \\gtrsim 2000$, as soon as they  will be accessible. Even though these are the same angular scales for which the \\textit{Sunyaev-Zel'dovich} effect begins to make a significant contribution to the temperature power spectrum, this occurs at two observational frequency bands, namely, 100 and 150 GHz, whereas the string contribution is at frequencies in observational bands near 220 GHz, where the Sunyaev-Zel'dovich effect is suppressed.  Although, as already mentioned, the relaxation time for the cosmic string loops to reach a self-similar evolution with respect to the horizon size appears to be larger for smaller loops, the considerable change in scale factor between the GUT redshift and the last scattering surface might well lead the observable length scales of the infinite component of the network to be in scaling at decoupling, $z = 1089$. Such scaling long strings may contribute to the CMB anisotropies down to fairly small angles. On the other hand, including the non-scaling loops, numerical simulations show some extra power at very small scales, which suggests that non-scaling structures start to have significant effects at very small scales, for $\\ell \\stackrel{<}{\\sim}10^{4}$. As regards the relevant contribution due to the infinite component to the CMB anisotropies at moderately high $\\ell$, the number of string segments in the corresponding range of redshifts close to $z_{LSS}$, where the network is not expected to have reached the matter era scaling regime, can be estimated as:  \\begin{center} \\begin{table} \\caption{The approximate number of string segments for different values of the equation-of-state parameter.}\\label{tab:table_2} \\resizebox*{1.07\\textwidth}{!}{ \\begin{tabular}{|ccccccccc|} \\hline \\hline &&&&&&&& \\\\ Model & BB & AS & MS & RSB & BHKU (AH) & BM (NG/AH)& BPOS & $\\Delta z$ \\\\ par& \\cite{20}& \\cite{6,21,41}& \\cite{22}& \\cite{23}& \\cite{40}& \\cite{25} & \\cite{26} &\\\\ \\hline \\hline &&&&&&&&\\\\ $\\frac{G\\mu}{ c^{2}}$ & $4 \\times 10^{-6}$  & $ 1.5 \\times  10^{-6}$  &  & $7 \\times 10^{-7}$ &$1.8 \\times 10^{-6}$ & $\\left(2.6 /6.4\\right) \\times 10^{-7}$ & &\\\\ \\hline &&&&&&&&\\\\ $v_{r}$ & $0.66 c$  &  $0.62 c$  &  $0.63 c$ &          &  $0.5 c$  & $0.65 c$ / $0.4 c$ & $0.63 c$ &\\\\ \\hline &&&&&&&&\\\\ $v_{m}$ & $0.61 c$  &  $0.58 c$  &  $0.57 c$ &          &  $0.5 c$  & $0.60 c$ / $0.4 c$ & $0.59 c$ &\\\\ \\hline &&&&&&&&\\\\ $A_{r}$ & 0.14      &  0.13      &   0.13    & 0.16     &   0.26   & 0.13/0.35           & 0.15 &\\\\ \\hline &&&&&&&&\\\\ $A_{m}$ & 0.18      &  0.17      &   0.20    & 0.19     &   0.29   & 0.21/0.35           & 0.17 &\\\\ \\hline w = -- 1.5& &&&&&&&\\\\ $N(z)$  &4 -- 330   & 4 -- 390   & 3 -- 240 & 3 -- 280  &  1 -- 78  & 3 -- 210/1 -- 45   & 4 -- 390 & $ (0.5,7]$ \\\\ \\hline w = -- 1& &&&&&&&\\\\ $N(z)$  &3 -- 310   & 4 -- 360   & 3 -- 220 & 3 -- 260  &  1 -- 73  & 2 -- 190/1 -- 42   &4 -- 360 & $ (0.5,7]$ \\\\ \\hline w = -- 0.5& &&&&&&&\\\\ $N(z)$  &3 -- 270   & 4 -- 310   & 3 -- 190 & 3 -- 230  &  1 -- 64  & 2 -- 170/1 -- 36   & 4 -- 310 & $ (0.5,7]$ \\\\ \\hline w = --1.5&&&&&&&&\\\\ $N(z)$  &360 -- 980 & 430 -- 1200 & 260 -- 710 & 310 -- 830 & 86 -- 230 & 230 -- 620/49 -- 130 & 430 -- 1200 &  $[8,100]$ \\\\ \\hline w = --1&&&&&&&&\\\\ $N(z)$  &340 -- 960 & 400 -- 1100 & 250 -- 700 & 290 -- 820 & 83 -- 230& 220 -- 610/47 -- 130 & 400 -- 1100&  $[8,100]$ \\\\ \\hline w = --0.5&&&&&&&&\\\\ $N(z)$  &300 -- 930 & 350 -- 1100 & 220 -- 680 & 250 -- 790 & 72 -- 220 & 190 -- 590/41 -- 130 & 350 -- 1100 &  $[8,100]$ \\\\ \\hline \\end{tabular} } \\end{table} \\end{center}  \\begin{equation} \\label{eqn:25} \\frac{8 \\ B^{3}(z)}{A^{3}_{m}} <  N(z) < \\frac{8 \\ B^{3}(z)}{A^{3}_{r}}. \\end{equation} \\noindent  Let us now consider the deviation from scaling during the matter-dark energy transition, which sets a lower limit at a redshift $z_{0} \\approx 0.5$ to the applicability of eq. (\\ref{eqn:12}). The number of cosmic string segments at $z \\lesssim 1$ could be extremely small so that it would become important to understand how deviation from the scaling behavior around the present cosmological era may affect the network since an eventual network no longer scaling with matter, exponentially slowed down, could be quite difficult to detect. As previously mentioned, the string velocity plays an important role for detecting signatures in the CMB since it determines the amplitude of the line-like discontinuities in the CMB temperature. As for the gravitational lensing there is no such restriction to relativistic velocities for observation to succeed but the number of string segments in the range of optical sources could be a very small fraction of the initial number and approaching $z \\lesssim 1$ it could become a small fraction even of the expected number in the effective observable (visible) universe. This can be seen by averaging on the values presented in Table 1 and taking into account that at $z = 1100$ because of the deviation from scaling $N^{av}(z)$ is expected to be less than the actual number. Thus, assuming that the most distant object we are able to observe by optical methods are located at a redshift $z = 7$, the probability to observe strings crossing the effective observable universe in the optical range is proportional to the volume of the optical sphere, having radius $d(z = 7) = B(z = 7) d_{PH}$. In horizon units then $p \\sim B^{3}(7) = 0.24$ ($B(7) = 0.62$) and the corresponding angular size of a string eventually observed is $\\sim 100^{o}$ $\\left[27\\right]$. For smaller values of the redshift the number of string segments inside the sphere becomes increasingly smaller and the angular size larger. For instance, at a redshift $z = 2$ the probability is reduced to $p \\sim B^{3}(2) = 0.05$ ($B(2) = 0.37$) and the angular size of the string is around $\\sim 134^{o}$. Below $z \\lesssim 1$ the number of string segments inside the sphere of radius $d(z \\lesssim 1)$ is expected to be so small (see Table 1) as to render detection by optical methods in this range a challenging problem unless the deviation from scaling at transition to the dark energy era is such that the decrease in reconnection events, whose effect is to increase the density of strings and the time needed to reach equilibrium, may somewhat compensate the eventual dilution by exponential expansion maintaining the number of strings to a detectable level.  As regards the dark energy variants of the standard cosmology, a different choice of the equation-of-state parameter implies increasing or decreasing distances and hence volumes defined through such distances. These volume effects are due to a different expansion rate: phantom models have always a larger volume than the $\\Lambda$CDM model since the expansion rate is larger, while quintessence models have a smaller volume than the $\\Lambda$CDM model being the expansion rate lower. As a consequence, one can only find an increased or decreased number of string segments, being the characteristic length scale of the network unaltered. The number of string segments for $w \\in \\left\\{-1.5, -1, -0.5\\right\\}$ is presented in Table 2. As we can see from Figure 1, $B(z)$ is slightly dependent of $w$ and a comparison of the results obtained in the two ranges of $z$ (Table 2) points out that at very small redshifts it makes no substantial difference whether we use the $\\Lambda$CDM or some other CDM model. Therefore for the considerations that follow, which apply to small values of redshift, for definiteness, we refer to the $\\Lambda$CDM  cosmology, but they can be easily extended to the above mentioned variants.  Taking into account both the constraints on the string tension parameter set by gravitational lensing ($G\\mu /c^{2} < 3 \\times 10^{-7}$) and by CMB power spectrum which, considering the BM result for NG strings $G\\mu /c^{2} < 2.6  \\times 10^{-7}$ and the stricter constraint coming from the Atacama Cosmology Telescope (ACT) \\cite{43}, $G\\mu/c^{2} < 1.6 \\times 10^{-7}$, may be suitably chosen in the middle of this range, we can get some further insight on the probability of observing cosmic strings through gravitational lensing events for static straight strings. Since the angular separation between the lensed images,  \\begin{equation} \\label{eqn:26} \\delta \\phi = \\Delta \\frac{D_{ls}}{D_{os}} sin \\theta    \\qquad   \\theta \\in \\left[0,\\pi \\right], \\end{equation} \\noindent depends on $\\Delta$ as well as on the ratio between the angular diameter distances from source to string (lens), $D_{ls}$, and from source to observer, $D_{os}$, and on the angle $\\theta$ between the string direction and the line of sight, not all the strings contained in the sphere of the optical sources can be observed through their gravitational lensing effects. For a straight string segment lying at $z$ close to the upper limit of the optical range the corresponding lensing angle would be on the border of observability ($< 0.07^{\\prime \\prime}$), being $D_{ls} \\ll D_{os}$, even in the ideal case $sin \\theta = 1$, although $N(z)$ is still quite large (on average about 270). For the most distant sources at $z \\sim 6-7$ therefore the string-lens is required to be at $z \\lesssim 5$ in order to yield observable effects. For non ideal strings an additional factor less than unity for $\\theta \\neq \\pi/2$ restricts the strings that have more probability to be observed to those located at redshifts around $z \\lesssim 2$ (where the average number of string segments is expected to be $< 60$), which implies a cut on the highest values of $z$ in the optical sources range at redshifts where the expected number of string segments is quite large. Setting a lower bound to the possible redshifts where one can reasonably suppose to find more than one string across the visible universe at the redshift of the transition matter-dark energy results in a further reduction of the probability of observing gravitational lensing events for the most distant and rare sources.  If we consider that the most part of the observed quasars and galaxies are distributed at redshifts $1 \\lesssim z \\lesssim 2$, the range of redshifts at which the probability of observing cosmic strings by their gravitational lensing effects is greatest will be $ 0.6 \\leq z < 1$ (where N(z) is on average $< 15$). For a source at $z_{s} \\in [1,2]$ and a segment of string at a redshift $z_{l} \\in [0.6,1)$ the corresponding lensing angle is in principle always sufficiently large for ideal strings to produce two distinguishable images of the source when this one is a quasar. As regards galaxies, only the points interior to the strip defined by the deficit angle will appear duplicated the other side of the string. The outer points will be cut away resulting in visible sharp edges in the isophotes of the source images \\cite{44}. In the case of non ideal strings even for strings close to the lower limit of the matter dominated era the corresponding lensing angle can be too small to be observable for some values of $\\theta$. Let us consider, for instance, randomly oriented static strings located at redshifts $0.6 \\leq z_{l} < 1$ which could be gravitational lenses for any sources in the above mentioned redshifts range. Outside of a certain interval of values of $\\theta$, $\\delta \\phi$ becomes smaller than the angular resolution even of space-based telescopes ($0.05^{\\prime \\prime}$ to $0.1^{\\prime \\prime}$). Since this interval is almost different for any value of $z_{l}$, in order to have an idea of the number of observable strings for any value of the range of $z_{l} \\in [0.6, 1)$ and for any value of the source redshift $z_{s} \\in [1,2]$ we have to calculate the probability of observing a cosmic string in a given volume which form  with the line of sight an angle $\\theta$ contained in the corresponding interval for each interval separately. By adding the results we find that the total amount of segments of static straight string which can produce observable gravitational lensing effects for any sources located at $z_{s} \\in [1,2]$ is reduced of about a half with respect to the case of ideal strings. For $1\\leq z_{l} < 2$, as $z_{l}$ increases, the ratio between the angular distances decreases so that even small deviations from the ideal case can yield undetectable lensing angles.   $\\delta \\phi$ is no longer detectable for $z_{l} \\sim 1.8$ even for ideal cosmic strings. Of all the strings contained in the spherical shell determined by $z_{l} \\in [1,2)$ only a small fraction (about one over four) may have a $\\delta \\phi$ potentially observable for sources  located at  $z_{s} \\in (1,2]$.  An exhaustive treatment on the topic involving survey fields, distribution of sources at different redshifts and their brightness and the string velocity is beyond the scope of this paper.  The outcome of the above reasoning focused on the most favorable case of the redshifts range of sources where the observed objects are more numerous and brighter is of course also a consequence of the small value chosen for $\\Delta$. Thus, if we apply the same considerations to cosmic superstrings we will have significant differences due to the fact that, while the $D$ and $FD$ components of the network have tensions comparable to those of ordinary cosmic strings, the $F$ component is expected to have smaller tension, $G\\mu_{F}/c^{2} \\sim 10^{-8}$, for a small value of the string coupling constant ($g_{s}\\sim 10^{-2}$). This leads to excessively small values of the lensing angle, $\\delta \\phi \\leq 0.05^{\\prime \\prime}$, even in the ideal case of string segments perpendicular to the line of sight. As a consequence, although a small value of the string coupling constant yields an exceedingly large number of string segments ($\\sim 10^{3}$) for the $F$ component even at small redshifts (down to the value $z_{l} \\sim 0.6$) against the small number found for ordinary cosmic strings, models supporting a small value of the string coupling constant cannot be ruled out on the grounds of the predicted large number of the $F$ component as long as the range of small values of tension involved are still unexplored by means of current technology. The case of cosmic superstrings will be discussed in detail in a forthcoming paper. Here we merely point out that a small value of $\\Delta \\sim 0.09^{\\prime \\prime}$ also leads to temperature jumps in the CMB radiation (KS effect) no longer appreciable so that the presence of a large number of $F$ strings does not conflict with observations."
    },
    "1112/1112.5523_arXiv.txt": {
        "abstract": "We present spectra and mass composition of cosmic rays incoming to the Earth in the energy range $(0.5\\!-\\!2)\\!\\cdot\\!10^{20}$\\,eV. As their sources we consider Seyfert galaxies located at distances $\\lesssim\\!40$~Mpc, following an acceleration model for such moderate--power objects. Mass composition of the particles at sources is assumed to be mixed. Generation spectra are described by a function $E^{-\\gamma_0}$, where $\\gamma_0$ is an arbitrary parameter. It is shown that the assumptions adopted make it possible to describe experimental data provided by HiRes and Pierre Auger Observatory, using different values of $\\gamma_0$. ",
        "introduction": "\\label{intr} Since the Greisen--Zatsepin--Kuzmin (GZK) cutoff for protons with energies $\\gtrsim\\!5\\!\\cdot\\!10^{19}$\\,eV was predicted in 1966 \\cite{GZKI,GZKII}, issues regarding ultrahigh--energy cosmic rays (UHECRs) have remained one of the main topics in high--energy astrophysics. Because of vanishingly small flux of such energetic particles, arrays enormous in area should be used over a~long period of~time. At present a~sufficient number of experimental events have been collected to provide us with an~idea on the behavior of the differential energy spectrum of UHECRs. Now the existence of the GZK cutoff is generally considered to be proved by measurements of the largest UHECR facilities, HiRes \\cite{HiRes_GZK} and Pierre Auger Observatory (Auger) \\cite{PAO_GZK}. It is appropriate, however, to recall that the experimental data reported by AGASA \\cite{AGASA_GZK} indicated the~opposite effect.  Although the suppression of the CR flux at ultrahigh energies has been reported by both HiRes and Auger, their spectra differ significantly in shape. For instance, while the Hires collaboration clearly locates the start point of the GZK cutoff \\cite{Sokolsky_2009}, one of two fits of the Auger spectrum has a~smooth form \\cite{PAO_2009}. Below we discuss the approximations suggested by both teams. Another, more surprising, contradiction between their observations is that the HiRes results demonstrate the proton--dominated CR composition at energies from $2\\!\\cdot\\!10^{18}$\\,eV up to $5\\!\\cdot\\!10^{19}$\\,eV \\cite{HiRes_xmax}, whereas Auger points out a~transition from protons to heavy nuclei, starting with energy of a~few times $10^{18}$\\,eV \\cite{Auger_xmax}. This problem was discussed, for example, in~\\cite{PhysTod}.  To describe experimental data and therefore to determine likely sources of UHECRs, various models of initial mass composition and acceleration conditions were considered (see, for instance, \\cite{Harari_2006}\\,--\\,\\cite{Aloisio_2009}). Nevertheless the origin of ultrahigh--energy particles remains an~open question. In this paper we present spectra and mass composition obtained within a~model \\cite{Uryson_2001,Uryson_2004} in which nuclei of moderate--power Seyfert galaxies are put forward as UHECR sources. Comparison of calculation results with measurements allows us to reveal appropriate acceleration regimes and besides to define somehow distances to the most probable sources. It is well known that particles generated above the GZK threshold shift to the region of lower energies in propagating through the space, shaping peculiarities in the cosmic ray spectrum, the so-called dip and bump \\cite{Hill_1985}\\,--\\,\\cite{Berezinsky_2006}. Study of these spectral features provides insight into UHECR sources. But a~more natural approach for investigation of UHECR particles is to exploit data near their maximum acceleration energy, so here we restrict our consideration to the region above the GZK cutoff. ",
        "conclusions": "In this paper we have analyzed the propagation of UHECR particles in the intergalactic medium from sources to the Earth and obtained their spectra and mass composition in the energy range of $(0.5\\!-\\!2)\\!\\cdot\\!10^{20}$\\,eV, based on the following assumptions: the UHECR sources are considered to be Seyfert galaxies located within the radius of $\\approx\\!40$~Mpc; mass composition of the particles at the sources resembles space content; relations between maximum energies of different nuclei are retained according to the acceleration model suggested for moderate--power sources; generation spectra obey an inverse power law with arbitrary index $\\gamma_0$. The higher is the value of~$\\gamma_0$, the heavier is the composition of nuclei accelerated. Over a~hundred isotopes arising due to interactions with background photons are taken into account.  Our calculations have shown that the~conjecture of nearby sources located at distances of less than 40~Mpc is compatible with the HiRes and Auger experimental data. Although these data are contradictory, it has been possible to represent them in the framework of a~unifying approach with varying injection index.  It is worth noting that our results depend crucially on the model parameters. We hope that a~larger number of UHECR events, which should be available in the near future, will allow to examine generation spectra at sources in more detail."
    },
    "1112/1112.3902_arXiv.txt": {
        "abstract": "The Stockholm Educational Air Shower Array (SEASA) project has established a network of GPS time-synchronised scintillator detector stations at high-schools in the Stockholm region. The primary aim of this project is outreach. A part of the network comprises a dense cluster of detector stations located at AlbaNova University Centre. This cluster is being used to study the cosmic ray anisotropy around the knee. Each station consists of three scintillator detectors in a triangular geometry which allows multiple timing measurements as the shower front sweeps over the station. The timing resolution of the system has been determined and the angular resolution has been studied using Monte Carlo simulations and is compared to data. The potential of this system to study small and large scale cosmic ray anisotropies is discussed. ",
        "introduction": "The 'Stockholm Educational Air Shower Array' (SEASA)~\\cite{seasa}  project has established seven detector stations in the Stockholm area, distributed according to Fig.~\\ref{pic:array}. Each detector station consists of three scintillator detectors, separated by approximately 15~m, which are read out by large area photomultipliers placed directly on top of the scintillators. The design of the detectors is explained in detail in~\\cite{mylic}. Four of the stations have been built by high school students and are located on the attics or roofs at their high schools. The separation between these stations are in the order of kilometers and the energy threshold to trigger multiple stations is therefore above $10^{18}$~eV. To be able to study lower energetic cosmic rays a more dense cluster of three detector stations has been deployed at the AlbaNova University Area (the \\emph{AlbaNova array}). This cluster detects air showers with energies above $~10^{16}$~eV and is being used to study the cosmic ray arrival distribution in this energy regime.  To study the cosmic ray anisotropy the direction of the primary cosmic ray must be possible to reconstruct. This can be done as the air shower front travels in essentially the same direction as the primary particle. Assuming a flat shower front, the direction can be determined by measuring the time difference between hit detector stations. By fitting the geometry of the detectors and the trigger times to the shower plane the incident angles of the shower can be reconstructed. With a fixed geometry the accuracy of the reconstruction  ultimately depends on the trigger time resolution, set by the GPS system. The time resolution of the GPS systems is therefore investigated in section~\\ref{sec:timeres}.  The pointing accuracy, or angular resolution, of the AlbaNova array has been assessed by simulations and this is described in section~\\ref{sec:angres}. The dependence of the angular resolution on the timing accuracy is also investigated here, as well as the trigger efficiency of the array. The final section~\\ref{sec:anisotropy} describes the methods used to study cosmic ray anisotropies. Finally, the hypothesis of a uniform flux of cosmic rays are tested using data taken during approximately six months of operation of the AlbaNova array.   \\begin{figure}[th] \\centering \\includegraphics[width=3.5in]{array} % \\caption{The seven stations constituting the SEASA air shower array. Four stations are located at high schools and are labeled after the corresponding schools. The three stations at the AlbaNova area are shown in greater detail in the inset picture in the upper right corner. The dotted station in this picture is a fourth station planned for installation.} \\label{pic:array} \\end{figure} ",
        "conclusions": "A measurement of the cosmic ray anisotropy has been made using the AlbaNova array which forms a sub array of the SEASA outreach project. The result favors a scenario with a uniform flux of cosmic rays in the energy regime above $10^{16}$~eV, and therefore agree with previous measurements, for example KASCADE~\\cite{kascade}. It is a crude measurement of the cosmic ray anisotropy but it shows that the array is stable and in particular that the GPS timing is reliable. The measurement will be improved in the near future by adding a fourth station at the AlbaNova area which increases the exposure of the array and also improves the shower angle reconstruction. A future improvement is also to enhance the read out electronics with a pulse sampling device. This would make it possible to determine the shower core and curvature and thereby improving the angular resolution of the array. This would allow anisotropy searches on smaller scales than what has been possible up to now."
    },
    "1112/1112.1300_arXiv.txt": {
        "abstract": "While the Low Frequency Array (LOFAR) is still in its commissioning phase, early science results are starting to emerge. Two nearby galaxies, M51 and NGC4631, have been observed as part of the Magnetism Key Science Project's (MKSP) effort to increase our understanding of the nature of weak magnetic fields in galaxies. LOFAR and the complexity of its calibration as well as the aims \\& goals of the MKSP are presented. ",
        "introduction": "LOFAR is a next generation radio telescope currently in its construction \\& commissioning phase. Its core is located in the Netherlands and the international stations are located in Germany, France, Sweden and the UK, with plans to expand to additional European countries. The LOFAR telescope consists of two distinct types of antennas, namely the Low Band Antenna (LBA) which operate between 10 and 90\\,MHz and the High Band Antenna (HBA) which operates between 120\\,MHz and 240\\,MHz. A single LBA antenna is able to see the entire sky and the telescope is pointed by forming beams at each station. Single HBA tiles (each of which contain 16 dipole pairs) have a large field of view but are not able to observe the entire sky (Stappers et al. 2011 \\& van Haarlem, Wise, et al. in prep). The data from all stations are then correlated at a central supercomputer.  LOFAR, due to its bandwidth and low frequency, can measure Faraday rotation measures (RM) with a very high precision. We will use ``RM Synthesis'' (Brentjens \\& de Bruyn 2005) and thereby detect weak magnetic fields and low electron densities in galaxies which are unobservable at higher frequencies. The high RM precision of LOFAR will also help to study RMs from pulsars, stellar and AGN jets, and the lobes of radio galaxies with unprecedented accuracy. ",
        "conclusions": "From these two observations, it is clear that LOFAR has started to give reliable images of extended sources. The extended synchrotron disk and halo of M51 and NGC4631 are easily seen in our images which were expected from observations in other low wavelengths seen in for M51 at 610\\,MHz and for NGC4631 at 327\\,MHz. It is important to note that the presented images are from a single subband. When the whole bandwidth of 24\\,MHz (122 subbands) is combined, a significant increase in sensitivity is expected. This should reveal the very outer regions of these galaxies.  A search for polarization with RM Synthesis is planned in the coming months."
    },
    "1112/1112.3305_arXiv.txt": {
        "abstract": " ",
        "introduction": "}  Planetary rings come in a diverse array of shapes and sizes.  They may be broad or narrow, dense or tenuous, dusty or not, and they may contain various kinds of structures including arcs, wavy edges, embedded moonlets, and radial variations.  Rings share the defining characteristic of a swarm of objects orbiting a central planet with vertical motions that are small compared to their motions within a common plane.  The latter arises because planets in our solar system (with the exceptions of Mercury and Venus, which have no known natural material in orbit) are fast-enough rotaters that their shapes are dominated by an equatorial bulge that adds a strong quadrupole moment ($J_2$) to their gravity fields (see Section~\\ref{OrbElems}).  This is a major contrast between rings and other astrophysical disks, which are not defined by asymmetry in an external gravity field but by the average angular momentum of the disk itself (in both cases, once a preferred plane is established, collisions among particles damp out the motions perpendicular to it).  However, rings do have a number of similarities with other astrophysical disks, which add to the motivation for studying them.  Unlike other known disk systems that are either many light-years away or (like the early stages of our solar system) far back in time, planetary rings can be studied up close and in real time.  Thus, it is worthwhile to consider the parallels that can be drawn between planetary rings and the study of other disks (see Section~\\ref{OtherDisks}).  In this chapter, after some further introductory notes on important concepts (Sections~\\ref{OrbElems} through~\\ref{OpticalDepth}), we will give an overview of the known ring systems, as well as systems where rings are unconfirmed but plausible, in Section~\\ref{RingsByPlanet}.  More detailed descriptions can be found in Section~\\ref{RingsByType}, which contains a discussion of various ring structures organized by type, with a focus on finding commonalities among rings in different locations that share certain qualities.  In Section~\\ref{Experiments} we will discuss experimental methods of learning about rings, and in Section~\\ref{AgeOrigin} we will discuss the age and origin of ring systems.  Finally, in Section~\\ref{OtherDisks}, we will discuss ways that planetary rings can illuminate the study of other astrophysical disks.  \\subsection{Orbital elements \\label{OrbElems}}  Rings are fundamentally populations of orbiting material.  Therefore, to discuss structure within rings, we will occasionally refer to one or more of the six \\textit{orbital elements} that describe the orbit of any object around another object.  These six parameters are simply a transformation of the six Cartesian parameters for position and velocity ($x$, $y$, $z$, $\\dot{x}$, $\\dot{y}$, and $\\dot{z}$) under the assumption that the object moves in the gravity field of a point mass (hereafter the ``planet'').  The derivation can be found in any textbook on orbital mechanics \\citep[e.g., Section~2.8 in][]{MD99}.  \\begin{figure*}[!t] \\begin{center} \\includegraphics[height=5cm]{orbelems_diagram1.pdf} \\includegraphics[height=5cm]{orbelems_diagram2.pdf} \\caption{The geometry of (left) an elliptical orbit within the orbit plane and (right) the orbit plane within 3-D space. \\label{OrbElemsFig}} \\end{center} \\end{figure*}  A diagram of the orbit in space is found in \\Fig{}~\\ref{OrbElemsFig}.  The size of the orbit, and its gravitational potential energy, is described by the semimajor axis $a$, which is the mean distance between the orbiting particle and the planet.  The shape of the orbit is described by the eccentricity $e$; for a circular orbit $e=0$, but real orbits that remain bound to the planet take elliptical shapes with $0 < e < 1$, with most ring particles having $e \\ll 1$.  Unbound orbits, either parabolic or hyperbolic, have $e \\geq 1$.  The orbit plane may be inclined with respect to the reference plane (for ring applications, this is often the planet's equatorial plane), by an angle known as the inclination $I$.  A non-zero inclination requires an account of the orbit plane's orientation, and thus its line of intersection with the reference plane (the ``line of nodes'') is described by the longitude of the ascending node $\\Omega$, measured with respect to a reference axis.  Similarly, a non-zero eccentricity requires an account of the orientation of the ellipse within the orbit plane, and thus the line connecting the planet to the location of the particle's closest approach (its ``periapse'') is described by the argument of periapse $\\omega$.  Finally, once the orbit has been defined by the five parameters already mentioned, the particle's position along the orbit can be given by its actual position (the true anomaly $f$) or its time-averaged position (the mean anomaly $M$) relative to periapse.  Also commonly used are the longitude of periapse $\\varpi = \\Omega + \\omega$ and the mean longitude $\\lambda = \\Omega + \\omega + M$, which are not physical angles since they are the sums of angles not necessarily in the same plane, but they have the virtue of being reckoned from a stationary reference axis rather than a moving line and are useful as long as $I$ is not too large.  The osculating orbital elements, which are most simply calculated and most often used, assume that the planet's gravity field is that of a point mass.  But for ring applications, the known planets are oblate (or bulged at the equators) due to their fast rotation.  This is adequately described by adding to the account of the gravity field a positive quadrupole moment $J_2$ \\citep[for details see, e.g., Section 4.5 of][]{MD99}, though it may be necessary to further include higher moments for applications requiring great precision.  The presence of a non-zero $J_2$, in addition to defining the Laplace plane\\fn{The Laplace plane is the plane about which orbits precess.  When the vertical motions of objects are damped by mutual collisions, material will settle into a ring centered on the Laplace plane.} for orbits near the planet, causes orbits to precess, in the prograde direction for apses ($\\dot{\\varpi} > 0$) and in the retrograde direction for nodes ($\\dot{\\Omega} < 0$).  A non-zero $J_2$ also compromises the physical meaningfulness of the osculating elements, especially for low-eccentricity orbits, introducing fast (i.e., orbit-frequency) variations in all six elements.  The physical meaningfulness of orbital elements can be restored using a revised system of \\textit{epicyclic orbital elements} \\citep{BL87,LB91,BRL94}, which are based on the geometrical shape of streamlines.  These put the orbit-frequency variations back into an analogue of $\\lambda$, leaving the other five elements to again describe a static (or at least slowly-varying) orbit.  A useful algorithm for converting Cartesian coordinates into epicyclic orbital elements was devised by \\citet{RS06}.  \\subsection{Roche limits, Roche lobes, and Roche critical densities \\label{Roche}}  The ``Roche limit'' is the distance from a planet within which its tides can pull apart a compact object.  Simply speaking, a ring would be expected to reside inside its planet's Roche limit, while any disk of material beyond that distance would be expected to accrete into one or more moons.  However, the Roche limit does not actually have a single value, but depends particularly on the density and internal material strength of the moon that may or may not get pulled apart \\citep{CE95}.  A simple value for the Roche limit can be calculated from a balance between the tidal force (i.e., the difference between the planet's gravitational pull on one side of the moon and its pull on the other side) that would tend to pull a moon apart, and the moon's own gravity that would tend to hold it together.  This works out to \\citep[e.g., \\Eqn{}~4.131 in][]{MD99} \\begin{equation} \\label{RocheLimitEqn} a_{\\mathrm{Roche}} = R_{\\mathrm{p}} \\left( \\frac{4 \\pi \\rho_\\mathrm{p}}{\\gamma \\rho} \\right)^{1/3} , \\end{equation} where $R$ is radius and $\\rho$ is internal density, and the subscript ``p'' denotes the central planet.  The dimensionless geometrical parameter $\\gamma = 4 \\pi / 3 \\approx 4.2$ for a sphere, but is smaller for an object that takes a non-spherical shape with its long axis pointing toward Saturn, as one would expect for an actively accreting body and as at least several of Saturn's ring-moons appear to do \\citep{PorcoSci07,Charnoz07}.  Simply distributing the moon's material into the shape of its Roche lobe, with uniform density, yields $\\gamma \\approx 1.6$ \\citep{PorcoSci07}.  However, fully accounting for the feedback between the moon's distorted shape and its (now non-point-mass) gravity field \\citep{Chandra69,MD99} leads to a rather smaller value, $\\gamma \\approx 0.85$; on the other hand, some central mass concentration and the failure of a rubble pile to exactly take its equilibrium shape will likely prevent $\\gamma$ from becoming quite this low.  Note that the moon's internal density $\\rho$ appears in \\Eqn{}~\\ref{RocheLimitEqn}.  Thus the Roche limit is variable; a denser object can venture closer to the planet without danger than can an object that is less dense.  The moon's diameter, on the other hand, does not appear in \\Eqn{}~\\ref{RocheLimitEqn}.  Why then do we commonly imagine a large object getting pulled into smaller pieces when it ventures inside the Roche limit?  This is because the Roche limit has been defined here as the distance within which an object can no longer be held together \\textit{by its own gravity}.  Size becomes relevant for objects small enough to be held together by their internal material strength (which is not considered in \\Eqn{}~\\ref{RocheLimitEqn}) in spite of the tidal forces that enter into the Roche calculation.  \\begin{figure*}[!t] \\begin{center} \\includegraphics[width=8cm]{rochedens.pdf} \\caption{Roche critical density $\\rho_{\\mathrm{Roche}}$ (\\Eqn{}~\\ref{RocheDensEqn}, with $\\gamma = 1.6$) plotted against planetary radii for Jupiter (red), Saturn (cyan), Uranus (green), and Neptune (blue).  An object must have density higher than $\\rho_{\\mathrm{Roche}}$ to be held together by its own gravity; conversely, in the presence of abundant disk material, an embedded object will actively accrete as long as its density remains higher than $\\rho_{\\mathrm{Roche}}$.  The colored bars along the bottom and along the left-hand side show the extent of each planet's main ring system.  For each, a solid circle indicates the outermost extent, and the corresponding minimum $\\rho_\\mathrm{Roche}$, of the main rings. \\label{rochedens}} \\end{center} \\end{figure*}  In fact, it is often more useful in the context of rings to consider the limit from planetary tides not as a critical distance but as a critical density.  At any given distance $a$ from the planet, there is a Roche critical density $\\rho_{\\mathrm{Roche}}$ at which the moon's size entirely fills its region of gravitational dominance (its ``Roche lobe'' or ``Hill sphere'' of characteristic radius $r_\\mathrm{Hill}$).  We can rearrange \\Eqn{}~\\ref{RocheLimitEqn} to obtain \\begin{equation} \\label{HillRadiusEqn} r_{\\mathrm{Hill}} = a \\left( \\frac{m}{3M_{\\mathrm{p}}} \\right)^{1/3} \\end{equation} and \\begin{equation} \\label{RocheDensEqn} \\rho_{\\mathrm{Roche}} = \\frac{4 \\pi \\rho_\\mathrm{p}}{\\gamma (a/R_{\\mathrm{p}})^3}= \\frac{3 M_{\\mathrm{p}}}{\\gamma a^3} , \\end{equation} where $m$ and $M_\\mathrm{p}$ are the masses of the moon and planet, respectively.  The first of these two expressions is the most useful for interpreting \\Fig{}~\\ref{rochedens}.  Within a ring, where material for accretion is plentiful, any pre-existing solid chunk with internal density greater than $\\rho_{\\mathrm{Roche}}$ should accrete a mantle of porous ring material until its density decreases to match $\\rho_{\\mathrm{Roche}}$.  This process should govern the size and density of the largest disk-embedded objects \\citep{PorcoSci07,Charnoz07}.  On the other hand, the density naturally achieved by transient clumps should be compared to $\\rho_{\\mathrm{Roche}}$ in order to predict whether disruption (rings) or accretion (discrete moons) will dominate in a particular location, and the persistent existence of a ring implies that the densities of transient clumps do not exceed $\\rho_{\\mathrm{Roche}}$ (that is, we expect $\\rho \\lesssim \\rho_{\\mathrm{Roche}}$).  As seen in \\Fig{}~\\ref{rochedens}, Saturn's rings extend outward to significantly lower values of $\\rho_{\\mathrm{Roche}}$, approaching 0.4~g~cm$^{-3}$, than are seen in any of the other three known ring systems, probably reflecting their much lower rock fraction (and higher fraction of water ice) as already known from spectroscopy and photometry \\citep{CuzziChapter09}.  That $\\rho_{\\mathrm{Roche}}$ for Saturn's rings reaches values much lower even than the density of solid water ice indicates a high degree of porosity, which is not surprising for a system in balance between disruption and accretion.  If the outer edge of a ring system is taken to be the transition between disruption-dominated and accretion-dominated regions, which is probably true at least for Saturn and Uranus given the large number of moons immediately outward of their main rings, and if the porosity of accreting objects is relatively constant among the different systems, then differences in $\\rho_{\\mathrm{Roche}}$ at the transition location probably reflect differences in bulk composition.  Since Uranus has a transitional $\\rho_{\\mathrm{Roche}}$ three times that of Saturn (\\Fig{}~\\ref{rochedens}), we may well infer that its rings are made of material with a higher grain density, i.e. a significantly higher rock/ice ratio.  Neptune's transitional $\\rho_{\\mathrm{Roche}}$ is intermediate between Saturn's and Uranus', possibly indicating an intermediate rock/ice ratio.  Our inference, from the Roche critical density at the ring/moon transition, that the Uranus system is rockier overall than the Saturn system is consistent with the fact that the average density of Saturn's mid-size moons \\citep{MatsonChapter09} is 1.2~g~cm$^{-3}$, while that of Uranus' major moons \\citep{Jake92} is 1.6~g~cm$^{-3}$.  We cannot test our inference that Neptune's rock/ice ratio is intermediate in this way, as Neptune has no indigenous major moons due to the cataclysm of Triton's capture \\citep{GoldreichNeptune89}.  The extent of Jupiter's Main ring, in contrast to the other three ring systems, is clearly limited by the availability of material (which originates at source moons Metis and Adrastea and evolves inward, and which is not abundant) rather than by a disruption/accretion balance.  However, its high value of $\\rho_{\\mathrm{Roche}} \\sim 1.7$~g~cm$^{-3}$ places the only known limit on the densities (and thus masses) of the source moons.  However, it may not be valid to assume that Metis and Adrastea are held together by gravity, as accreting masses must be, given the large gap in particle size between the $\\sim 10$-km moons and other Main ring particles, which observationally cannot be larger than 1~km \\citep{Show07}.  This large gap in particle size might be explained if Metis and Adrastea are solid bodies originating further out, now held together by material strength, while no bodies of similar size are now able to form through \\textit{in~situ} accretion.  \\subsection{Optical depth \\label{OpticalDepth}}  The amount of material in a system with general disk morphology can be measured in several ways.  The most straight-forward is the surface density, the mass per unit surface area of the disk, though it must be borne in mind that a disk with greater vertical thickness will have proportionately lower volume density than a vertically-thin disk, even if both have the same surface density.  However, surface density can be difficult to measure directly.  A much more common observable is the optical depth $\\tau$, which can be thought of as the attenuation of a beam of light passing through the disk, measured in $e$-folding terms.  That is, \\begin{equation} \\tau = - \\ln \\left( \\frac{I}{I_*} \\right) \\equiv - \\ln T , \\label{TauDef} \\end{equation} \\noindent where $I$ is the observed intensity, $I_*$ is the unocculted intensity (e.g., from a background star), and the ratio between them is defined as the ring's transparency $T$.  The optical depth is sensitive to both the number density and the size of ring particles, which can be obtained when both the optical depth and the surface density are known \\citep[e.g.,][]{Colwell09}.  For a given value of the surface mass density, the optical depth scales inversely with particle size, that is, with the ratio of volume to surface area.  For a vertically-thick homogeneous disk, the optical depth is proportional to the path length through the disk, which in turn is proportional to $\\mu \\equiv \\sin B$, where the elevation angle $B$ is the angle between the line-of-sight and the ring-plane (\\Fig{}~\\ref{TauSlant}a).  In order to compare observations taken over a range of elevation angles, the normal optical depth, $\\tau_\\perp \\equiv \\mu \\tau$, is often used.  However, this parameter must be used with caution for disks that lack homogeneity and/or are close to a single layer thick.  \\citet{Colwell07} found that $\\tau_\\perp$ varies strongly with $\\mu$ in the B~ring, and that the uncorrected $\\tau$ is a more robust parameter in that case, indicating that the B~ring is composed of vertically-thin nearly-opaque clumps with nearly-transparent gaps between them (\\Fig{}~\\ref{TauSlant}b), and that the optical depth is controlled by the relative abundance of clumps and gaps (see Section~\\ref{SGWs}).  \\begin{figure*}[!t] \\begin{center} \\includegraphics[width=16cm]{TauSlant.pdf} \\caption{Schematic showing a slanted path through (a) a homogeneous slab and (b) flattened SGWs.  The measured optical depth $\\tau$ is proportional to $\\sin B$ in the first case, but is relatively insensitive to elevation angle in the second. \\label{TauSlant}} \\end{center} \\end{figure*}  In numerical simulations (see Section~\\ref{NumSims}), the photometric optical depth $\\tau$ (\\Eqn{}~\\ref{TauDef}) is cumbersome to calculate, but a useful proxy known as the dynamical optical depth $\\tau_\\mathrm{dyn}$ can be found by summing the total cross-section area of all simulated particles and dividing by the area of the simulation patch.  This quantity turns out to be equal to the photometric optical depth as long as particles are randomly distributed, as the Gaussian probability of particle overlap plays the same role when using $\\tau_\\mathrm{dyn}$ to calculate the total transparency that the exponential plays when using \\Eqn{}~\\ref{TauDef}.  However, for high values of $\\tau$, when the distance between particles becomes comparable to the particle size, particles become constrained as to the locations in space they can occupy and the two quantities diverge.  Specifically, \\begin{equation} \\tau / \\tau_\\mathrm{dyn} \\simeq 1 + kD , \\end{equation} \\noindent where the volume filling factor $D$ is calculated from particle radius, disk scale height, and $\\tau_\\mathrm{dyn}$, and $k$ is a scalar of order unity \\citep{SK03,Anparsgw10}.  Furthermore, the existence of microstructure such as self-gravity wakes causes the distribution of particles to be strongly non-random, and can cause $\\tau_\\mathrm{dyn}$ to diverge strongly from the photometrically observed $\\tau$. ",
        "conclusions": ""
    },
    "1112/1112.3769_arXiv.txt": {
        "abstract": "Large surveys using the Sunyaev-Zel'dovich (SZ) effect to find clusters of galaxies are now starting to yield large numbers of systems out to high redshift, many of which are new discoveries. In order to provide theoretical interpretation for the release of the full SZ cluster samples over the next few years, we have exploited the large-volume Millennium Gas cosmological $N$-body hydrodynamics simulations to study the SZ cluster population at low and high redshift, for three models with varying gas physics. We confirm previous results using smaller samples that the intrinsic (spherical) $Y_{500}-M_{500}$ relation has very little scatter ($\\sigma_{\\log_{10}Y}\\simeq 0.04$), is insensitive to cluster gas physics and evolves to redshift one in accord with self-similar expectations. Our {\\it pre-heating} and {\\it feedback} models predict scaling relations that are in excellent agreement with the recent analysis from combined {\\it Planck} and {\\it XMM-Newton} data by the Planck Collaboration. This agreement is largely preserved when $r_{500}$ and $M_{500}$ are derived using the hydrostatic mass proxy, $Y_{\\rm X,500}$, albeit with significantly reduced scatter ($\\sigma_{\\log_{10}Y}\\simeq 0.02$), a result that is due to the tight correlation between $Y_{500}$ and $Y_{\\rm X,500}$. Interestingly, this assumption also hides any bias in the relation due to dynamical activity. We also assess the importance of projection effects from large-scale structure along the line-of-sight, by extracting cluster  $Y_{500}$ values from fifty simulated $5^{\\circ} \\times 5^{\\circ}$ sky  maps. Once the (model-dependent) mean signal is subtracted from the maps we find that the integrated SZ signal is unbiased with respect to the underlying clusters, although the scatter in the (cylindrical) $Y_{500}-M_{500}$ relation increases in the {\\it pre-heating} case, where a significant amount of energy was injected into the intergalactic medium at high redshift. Finally, we study the hot gas pressure profiles to investigate the origin of the SZ signal and find that the largest contribution comes from radii close to $r_{500}$ in all cases. The profiles themselves are well described by generalised Navarro, Frenk \\& White profiles but there is significant cluster-to-cluster scatter. In conclusion, our results support the notion that $Y_{500}$ is a robust mass proxy for use in cosmological analyses with clusters. ",
        "introduction": "The Sunyaev-Zel'dovich (SZ) effect \\citep{Sunyaev72} is a powerful method for discovering new clusters of galaxies. It arises generically due to the scattering of Cosmic Microwave Background (CMB) photons off free electrons, leading to a predictable spectral distortion in the CMB that is, in the non-relativistic limit, linearly dependent on the line integral of the electron pressure \\citep{Birkinshaw99}. In modern theories of structure formation, the dominant contribution to the SZ signal comes from the intracluster medium (ICM), a diffuse plasma within clusters that is approximately in hydrostatic equilibrium within the dark-matter dominated potential (see \\citealt{Voit05,Allen11} for recent reviews). The key SZ observable is the $Y$ parameter, defined as \\begin{equation} Y = \\int \\, y \\, {\\rm d}\\Omega, \\end{equation} where the integral is performed over the solid angle subtended by the cluster. The Compton-$y$ parameter is determined by the thermal structure of the ICM \\begin{equation} y = {\\sigma_{\\rm T} k \\over m_{\\rm e} c^2} \\, \\int \\, n_{\\rm e}T_{\\rm e} \\, {\\rm d}l, \\end{equation} where $n_{\\rm e}$ and $T_{\\rm e}$ are the density and temperature of the free electrons respectively and ${\\rm d}l$ is the differential line element along the line-of-sight. Since $Y$ can be expressed as a volume integral of the pressure (when the redshift and cosmological parameters are specified), it measures the total thermal energy of the gas, a property that ought to be strongly correlated with the cluster's mass through the virial theorem. This means that $Y$ ought to be relatively insensitive to the complex micro-physics taking place in the cluster core, unlike other global properties such as X-ray luminosity.   Early observational studies confirmed the detection of an SZ signal towards known massive clusters of galaxies and used this to estimate the Hubble constant (e.g. \\citealt{Jones93,Birkinshaw94}). Over the past decade, SZ observations of known bright X-ray bright clusters have become routine, allowing the investigation of cluster scaling relations to be performed (e.g. \\citealt{McCarthy03b,Benson04,Morandi07,Bonamente08,Huang10,Shimwell11,Lancaster11}). One potential shortcoming of this approach is that the samples are X-ray selected and therefore biased towards luminous, cool-core systems at low redshift.  In the past few years, SZ science has entered the exciting new phase of blind surveys, where detections of new clusters have become possible \\citep{Staniszewski09}. Indeed, SZ surveys are now yielding large numbers of SZ-selected clusters, many of them new detections, especially from the South Pole Telescope (SPT; \\citealt{Vanderlinde10,Andersson11,Williamson11}), the Atacama Cosmology Telescope (ACT; \\citealt{Marriage11,Sehgal11}) and the {\\it Planck} satellite \\citep{Planck11a,Planck11b,Planck11c}. Since the SZ effect is effectively independent of redshift, the SZ selection function tends to favour higher redshift systems than the X-ray counterpart, assuming similar angular resolution. As a result, the new blind SZ surveys are starting to find new massive systems at $z \\sim 1$ \\citep{Planck11d,Foley11,Menanteau11}. In the near future, we should expect to see these numbers increase substantially as the full survey results are published, nicely complementing X-ray surveys such as the MAssive Cluster Survey (MACS; \\citealt{Ebeling01}) and the {\\it XMM} Cluster Survey (XCS; \\citealt{Romer01,Mehrtens11}). Such complementarity will be further exploited with the next generation of X-ray surveys (e.g. with eROSITA) and millimetric telescopes (e.g. CCAT; see \\citealt{Golwala09}).   One of the main goals of SZ surveys is to measure cosmological parameters (e.g. \\citealt{Barbosa96,Carlstrom02,Battye03}). Central to the cosmological application of SZ surveys is the scaling relation between the observables ($Y$ and redshift, $z$) and cluster mass, $M$. Under the assumption that clusters form a self-similar population \\citep{Kaiser86} the SZ flux should scale as $Y \\propto M^{5/3} H(z)^{2/3}$, when measured within a radius enclosing a mean density that is a constant multiple of the critical density of the Universe. Early theoretical studies combined such simple scaling relations with the Press-Schechter formalism \\citep{Press74} to predict the SZ evolution of the cluster population in a variety of cosmological models (e.g. \\citealt{Cole88,Bartlett94,Barbosa96,Eke96,Aghanim97,Kay01,Battye03}). More recently, attention has turned to more detailed studies of how cluster gas physics impacts upon SZ scaling relations, both using semi-analytic models (e.g. \\citealt{McCarthy03a,McCarthy03b,Shaw08}) and full cosmological $N$-body/hydrodynamic simulations \\citep{daSilva00,White02,daSilva04,Motl05,Nagai06,Bonaldi07,Hallman07,Aghanim09,Battaglia11}. Simulations are now also being used to investigate the effects of mergers on SZ scaling relations \\citep{Poole07,Wik08,Yang10,Krause12}. A generic result from these studies is that the self-similar description appears to be approximately valid on cluster scales ($M>10^{14}\\hMsol$) but in detail, differences are seen between the models that are due to the effects of non-gravitational physics (cooling and heating processes), especially at low mass where the gas fraction is depleted.   Two of the main shortcomings in previous simulation studies are the relatively small samples (that are sometimes restricted to lower-mass clusters) and a limited range of (uncertain) cluster gas physics models, often not calibrated to match X-ray data. Some studies may satisfy one of these criteria but usually not both. A new generation of simulations are now starting to overcome both shortcomings. \\citet{Stanek10} recently presented results from two of the {\\it Millennium Gas} simulations \\citep{Hartley08}, large-volume runs based on the Millennium Simulation \\citep{Springel05a} with varying gas physics. These simulations are sufficiently large to enable the full range of cluster masses ($10^{14}-10^{15} \\hMsol$) to be studied and one of the runs, where the gas was pre-heated at high redshift, is able to match the mean X-ray luminosity-temperature relation at $z=0$ \\citep{Hartley08}. Although the work of \\citet{Stanek10} was focused on the more general issue of multi-variate scaling relations, they presented results for the SZ $Y-M$ relation measured within a radius corresponding to a mean internal density equal to 200 times the critical density, $r_{200}$.  The aim of this paper is to use these Millennium Gas simulations to focus in more detail on predictions of the SZ effect and, in particular, the $Y-M$ relation for clusters. Our paper builds on the \\citet{Stanek10} work in three important ways. Firstly, we add a third model that includes a more realistic treatment of feedback, both from supernovae and active galactic nuclei. This model has already been shown to successfully match many of the X-ray properties of non-cool core clusters \\citep{Short09,Short10}. Secondly, we include in our analysis simulated maps of the full SZ effect along the line-of-sight, to assess the projection effects of large-scale structure. Finally, we attempt to produce results for our $Y-M$ scaling relations using methods that are more closely matched with observations. In particular, we present our results for the smaller $r_{500}$ and investigate the impact of assuming hydrostatic equilibrium and a mass proxy ($Y_{\\rm X}$, \\citealt{Kravtsov06}) on the $Y-M$ relation.   We organise the remainder of the paper as follows. In Section~\\ref{sec:simdetails} we outline the simulation details and our methods used to define cluster properties. We also present some basic properties of the sample and SZ maps. Sections~\\ref{sec:pressureprofiles}, \\ref{sec:scalingrelations} and \\ref{sec:hse} contain our main results: in Section~\\ref{sec:pressureprofiles} we present an analysis of the hot gas pressure profiles, before going on to study SZ scaling relations in Section~\\ref{sec:scalingrelations} and the impact of hydrostatic bias in Section~\\ref{sec:hse}. Finally, in Section~\\ref{sec:summary} we summarise our main conclusions and outline future work. ",
        "conclusions": "\\label{sec:summary}  Large surveys are now being performed at millimetre wavelengths exploiting the Sunyaev-Zel'dovich (SZ) effect to detect large samples of galaxy clusters out to high redshift. Such samples will then be used to produce competitive constraints on cosmological parameters, as well as to study the variation in physical properties of the intracluster gas (especially the gas pressure) with mass and redshift. The cosmological application relies on the statistical estimation of cluster mass through the SZ $Y-M$ relation. In recent work (e.g. \\citealt{Andersson11,Planck11c}) the first SZ-selected samples of clusters have already been used to estimate the cluster $Y-M$ relation and full results of cosmological analyses are expected over the next few years.  In this paper, we have analysed some of the largest $N$-body/hydrodynamic simulations of structure formation (the Millennium Gas Simulations) to study the dependence of SZ cluster properties on gas physics, at both low ($z=0$) and high ($z=1$) redshift. The large volume used in these simulations produces significant (hundreds to thousands) samples of clusters over the interesting range of cluster masses ($10^{14}-10^{15} \\hMsol$). We considered three cluster gas physics models: a non-radiative (gravitational heating only) simulation that ought to produce an approximately self-similar cluster population; and two simulations that incorporate additional non-gravitational heating (a model that uniformly pre-heats the gas at high redshift and a model that includes feedback from stars and active galactic nuclei in galaxies). The feedback model is our most realistic, in that it has already been shown to reproduce many of the scaling properties of X-ray clusters, especially those with non-cool cores \\citep{Short10}.  We started by investigating the hot gas pressure profiles of our simulated clusters and how they compare to the pressure profile advocated by \\citet{Arnaud10}. We then compared our derived SZ scaling relations (between $Y_{500}$ and total mass, hot gas mass, X-ray temperature and the X-ray analogue to the SZ $Y$ parameter, $Y_{\\rm X}$) with the recent observational results, in particular those obtained from a combined SZ+X-ray analysis performed by the Planck Collaboration. We also tested two of the key assumptions used in the observed analysis, namely that the mean $Y_{500}-M_{500}$ relation is unaffected by the assumption that the gas is hydrostatic and by the presence of any other hot gas along the line-of-sight. Our main conclusions can be summarised as follows:  \\begin{itemize} \\item In accord with previous studies, our simulation with non-radiative hydrodynamics produces a (spherical) $Y_{500}-M_{500}$ relation that has a self-similar slope (5/3) and also evolves with redshift according to the self-similar expectation, $E(z)^{2/3}$. Simulations with non-gravitational heating (both pre-heating and feedback cases) create slightly steeper $Y_{500}-M_{500}$ relations (with a slope of $1.7-1.8$, when clusters across the mass range, $10^{14}\\hMsol < M_{500} < 10^{15}\\hMsol$, are considered) but the evolution with redshift is still close to self-similar.  \\item The simulations were compared with the {\\it Planck+XMM} results at $z<0.5$ \\citep{Planck11c} and very good agreement was found for a number of scaling relations ($Y_{500}$ versus $M_{500}, M_{\\rm gas,500}$ and $kT_{\\rm X}$) for the pre-heating and feedback models. The scatter in the $Y_{500}-M_{500}$ relation is smaller than observed, however, with $\\sigma_{\\log_{10}Y}\\simeq 0.04$.  \\item Intracluster gas in the non-radiative simulation contains a significant unthermalised component, due to the presence of low-entropy, clumpy gas. This causes an offset in the $Y_{500}-Y_{\\rm X,500}$ relation, which tests the difference between the mass-weighted and X-ray temperatures. As a result, hydrostatic mass estimates are biased low by 20-30 per cent. The pre-heating and feedback simulations on the other hand predict smoother gas distributions, with $Y_{500} \\simeq Y_{\\rm X,500}$ and much smaller hydrostatic bias (estimated masses are only $\\sim 10$ per cent lower).  \\item The estimated $M_{500}-Y_{\\rm X,500}$ relations (assuming the gas is hydrostatic) are in good agreement with the recent observational determination by \\citet{Arnaud10}. When $Y_{\\rm X,500}$ is used as a mass proxy to predict the SZ $Y_{500}-M_{500}$ relation, only a small ($\\sim 20$ per cent) offset in normalisation from the true relation (and thus the observed relation from {\\it Planck+XMM} data) is found. The scatter in the recovered relation is very small ($\\sigma_{\\log_{10}Y}\\simeq 0.02$) due to the strong correlation between $Y_{500}$ and $Y_{\\rm X,500}$. Clusters that are undergoing major mergers are not significantly offset from the mean relation.  \\item Hot gas pressure profiles are well described by generalised NFW profiles, as suggested by \\citet{Nagai07b} and show that the majority of the contribution to the SZ $Y$ parameter (where $r^3 P(r)$ is maximal) comes from radii close to $r_{500}$. Splitting the cluster samples into low and high-mass sub-samples, we find little difference between the two in the run with non-radiative hydrodynamics, as expected. The runs with non-gravitational heating predict that low-mass clusters have lower core pressures and higher pressures in the cluster outskirts, when scaled according to the self-similar expectation. This non-self-similar-behaviour can be attributed to the heating that is more effective in low-mass clusters and acts to push the gas out to large radii. There is also significant cluster-cluster scatter, especially in the core region and in the outskirts, where individiual pressure profiles can be 50 per cent higher than the median profile.  \\item We also compared our median pressure profiles with the Arnaud et al. profile and found good agreement (within 10 per cent) for our high-mass clusters at $r>0.5r_{500}$, in the pre-heating and feedback models. Low-mass clusters are especially discrepant in the core regions, likely due to the absence of radiative cooling in our models. Using the X-ray temperature (rather than hot gas mass-weighted temperature) in the pressure calculation, as well as using hydrostatic estimates of $r_{500}$ and $M_{500}$, only makes a significant ($>10$ per cent) difference to the non-radiative simulation for the reasons already mentioned.  \\item Finally, we considered the effects of projection due to large-scale structure along the line-of-sight, by analysing 50 $5^{\\circ} \\times 5^{\\circ}$ maps of the thermal SZ effect. By measuring the cylindrical SZ flux associated with each cluster and comparing to the flux from the cluster region alone, we were able to discern the contribution from additional structures, in the non-radiative and pre-heating simulations. The pre-heating model showed the largest bias, where low-mass clusters ($M_{500}\\simeq 10^{14}\\hMsol$) had cylindrical $Y_{500}$ values that were around 2-3 times higher than the value from the cluster region. This is due to the large amount of thermal energy injected into the gas at high redshift, as evidenced by the three-fold increase in the mean-$y$ parameter. Subtracting the contribution from an assumed mean background we find the recovered $Y_{500}-M_{500}$ relation to be unbiased with respect to the cluster relation, with some additional scatter that is model-dependent. \\end{itemize}  In summary, we can conclude that when our more realistic models for the intracluster gas are employed (namely those that raise the entropy of the gas to match global X-ray scaling relations), the SZ $Y-M$ relation is in good agreement with the observations (Fig.~\\ref{fig:ymrelz10}) and is largely unaffected by two of the main sources of systematic uncertainty: hydrostatic bias (Figs.~\\ref{fig:y500m500hse} and \\ref{fig:y500m500myx}) and projection effects from large-scale structure (Fig.~\\ref{fig:y500m500cylbsub}).  While our analysis has been one of the most comprehensive to date and used some of the largest and most sophisticated simulation models, there are some significant shortcomings that still need to be addressed. Firstly, the effects of radiative cooling were not included in our most realistic (feedback) model, so the model cannot yet match the full X-ray cluster population (namely the brightest objects with cool cores, \\citealt{Short10}). As we argued, this omission is likely not a significant problem for the $Y-M$ relation but will affect the hot gas pressure profile so it should be addressed in future work. Secondly, we were unable to test projection effects for the feedback model as we only have a sample of clusters rather than the full cosmological volume. Finally, the cosmological model adopted for the simulations (identical to that used in the original Millennium Simulation) is no longer favoured; in particular the value of $\\sigma_8$ is higher than the current best estimate ($\\sigma_8=0.9$ in the simulations, c.f. $\\sigma_8\\simeq 0.8$ from the {\\it WMAP} 7-year data; \\citealt{Komatsu11}). Using the presently-favoured cosmological model is likely to reduce the scale of projection effects, however, as in it structure formation will be less advanced.  We are currently preparing a new generation of Millennium Gas simulations that will rectify all of these problems, starting with a new version of our existing feedback model that will deal with the second and third issues. This new simulation, which is also being run at higher resolution and with an updated semi-analytic galaxy formation model \\citep{Guo10}, will additionally allow cosmologically-dependent statistical predictions for the SZ signal to be performed, namely the SZ power spectrum."
    },
    "1112/1112.5559_arXiv.txt": {
        "abstract": "In the system of a gravitating Q-ball, there is a maximum charge $Q_{{\\rm max}}$ inevitably, while in flat spacetime there is no upper bound on $Q$ in typical models such as the Affleck-Dine model. Theoretically the charge $Q$ is a free parameter, and phenomenologically it could increase by charge accumulation. We address a question of what happens to Q-balls if $Q$ is close to $Q_{{\\rm max}}$. First, without specifying a model, we show analytically that inflation cannot take place in the core of a Q-ball, contrary to the claim of previous work. Next, for the Affleck-Dine model, we analyze perturbation of equilibrium solutions with $Q\\approx Q_{{\\rm max}}$ by numerical analysis of dynamical field equations. We find that the extremal solution with $Q=Q_{{\\rm max}}$ and unstable solutions around it are ``critical solutions\", which means the threshold of black-hole formation. ",
        "introduction": "In a pioneering work by Friedberg {\\it et al}.\\ in 1976 \\cite{FLS76}, nontopological solitons were introduced in a model with a U(1)-symmetric complex scalar field coupled to a real scalar field. In contrast with topological defects, they are stabilized by a global U(1) charge, and their energy density is localized in a finite space region without gauge fields. In 1985 Coleman showed such solitons exist in a simpler model with an SO(2) [viz.\\ U(1)] symmetric scalar field only, and called them Q-balls \\cite{Col85}.  Q-balls have attracted much attention in particle cosmology since Kusenko pointed out that they can exist in all supersymmetric extensions of the standard model \\cite{Kus97b-98}. Specifically, Q-balls can be produced efficiently in the Affleck-Dine (AD) mechanism \\cite{AD} and could be responsible for baryon asymmetry \\cite{SUSY} and dark matter \\cite{SUSY-DM}. Q-balls can also influence the fate of neutron stars \\cite{Kus98}. Since Q-balls are supposed to be microscopic objects, equilibrium solutions and their stability have been intensively studied in flat spacetime \\cite{stability}. It was shown that catastrophe theory is a useful tool for stability analysis of Q-balls \\cite{SS}.  If Q-balls are so large or so massive, on the other hand, their size becomes astronomical and their gravitational effects are remarkable \\cite{Grav-Q,multamaki}. Such gravitating Q-balls, or Q-stars, are analogous to boson stars \\cite{boson-review}. While Q-balls exist even in flat spacetime, boson stars are supported by gravity and nonexistent in flat spacetime. Multamaki and Vilja showed that the size of Q-balls is bounded above due to gravity \\cite{multamaki}. Becerril {\\it et al}.\\ studied evolution of unstable solutions by numerical analysis of dynamical field equations \\cite{BBGN}.  In our previous work \\cite{TS1,TS3,TS24} we studied equilibrium solutions and their stability for several models using catastrophe theory to understand a comprehensive picture of flat Q-balls, gravitating Q-balls and boson stars. One of the main results is summarized in TABLE I. \\begin{itemize} \\item In flat spacetime, as long as the absolute minimum of the potential $V(\\phi)$ is located at $\\phi=0$, there is no upper bound on charge. If we take self-gravity into account, however, there is maximum charge $Q_{{\\rm max}}$ (or maximum  energy $E_{{\\rm max}}$),  at which stability changes, regardless of models. \\item In flat spacetime, in some models such as the AD gauge-mediated model \\cite{AD}, there is nonzero minimum charge $ Q_{{\\rm min}}$, where Q-balls with $Q<Q_{{\\rm min}}$ are nonexistent. If we take self-gravity into account, however, there exist stable Q-balls with arbitrarily small mass and charge. \\end{itemize} The above properties of gravitating solitons hold for general models of Q-balls and boson stars as long as the leading order term of the potential is a positive mass term in its Maclaurin series.  \\begin{table} \\begin{tabular}{|c|c|c|}\\hline & $Q_{{\\rm max}}$ & $Q_{{\\rm min}}$ \\\\\\hline flat Q-balls & $\\infty$ or finite & finite or 0 \\\\\\hline gravitating Q-balls & finite & 0 \\\\\\hline boson stars & finite & 0 \\\\\\hline \\end{tabular} \\caption{Comparison of maximum and minimum charge (or gravitational mass) for flat Q-balls, gravitating Q-balls and boson stars.} \\end{table}  Theoretically the charge $Q$ is a free parameter, and phenomenologically it could increase by charge accumulation. Therefore, a natural question could arise: what happens to Q-balls if $Q$ is close to $Q_{{\\rm max}}$? In this paper we shall investigate what happens if we give perturbations to equilibrium Q-balls with $Q\\approx Q_{{\\rm max}}$.  In connection with this question, Matsuda \\cite{Mat} claimed that inflation occurs in the core of a Q-ball if $Q$ is large enough. He assumed a kind of hybrid potential. In his scenario inflation takes place when the gravity-mediated term, \\beq\\label{AD} V(\\phi)\\equiv \\frac{m^2}{2}\\phi^2\\left[1+K\\ln \\left(\\frac{\\phi}{M}\\right)^2\\right], \\eeq dominates, where $m$ is the gravitino mass, $K$ term a one-loop correction, and $M$ the renormalization scale. If this scenario is true, it would have an important implication for inflationary models. Therefore, the present study is also important as a close examination of this scenario.  This paper is organized as follows. In Sec. II we review previous results of gravitating Q-balls in the AD gravity-mediated model. In Sec. III we make analytic discussions on general properties of Q-ball equilibrium solutions. In Sec. IV we discuss whether Q-ball inflation can take place. In Sec. V we present dynamical field equations and our computing method. In Sec. VI, by analyzing the dynamical field equations, we investigate what happens if $Q$ is so large in the AD gravity-mediated model. Sec. VII is devoted to concluding remarks. ",
        "conclusions": "We have addressed a question of what happens to Q-balls if $Q$ is close to $Q_{{\\rm max}}$. First, without specifying a model, we have shown analytically that the core of an equilibrium Q-ball has attractive nature and inflation cannot take place there. Next, for the Affleck-Dine model, we have analyzed  perturbation of equilibrium solutions with $Q\\approx Q_{{\\rm max}}$ by numerical analysis of dynamical field equations. We have found that the extremal solution with $Q=Q_{{\\rm max}}$ and unstable solutions around it are ``critical solutions\", which means the threshold of black-hole formation.  Specifically, for initial data (\\ref{initial}) with (\\ref{perturbation}), a black-hole is formed if $\\delta\\phi>0$. If $\\delta\\phi<0$, a black hole is not formed, and there are two types of evolutions. If the initial configuration is very close to the extremal Q-ball with $Q=Q_{{\\rm max}}$, the Q-ball continues to oscillate without diffusing mass or charge. In other cases,  the Q-ball diffuses most of mass and charge, and becomes a thick-wall Q-ball with much smaller charge.  Further study is necessary to understand detailed behavior of these critical solutions. It is also interesting to investigate how such behavior of critical solutions depends on models $V(\\phi)$."
    },
    "1112/1112.3525_arXiv.txt": {
        "abstract": "We study the gas accretion onto a supermassive black hole (SMBH) using the 3D SPH code GADGET-3 on scales of $0.1-200$\\,pc. First we test our code with spherically symmetric, adiabatic Bondi accretion problem. We find that our simulation can reproduce the expected Bondi accretion flow very well for a limited amount of time until the effect of outer boundary starts to be visible.  We also find artificial heating of gas near the inner accretion boundary due to the artificial viscosity of SPH. Second, we implement radiative cooling and heating due to X-rays, and examine the impact of thermal feedback by the central X-ray source.  The accretion flow roughly follows the Bondi solution for low central X-ray luminosities, however, the flow starts to exhibit non-spherical fragmentation due to thermal instability for a certain range of central $L_X$, and a strong overall outflow develops for greater $L_X$. The cold gas develops filamentary structures that fall into the central SMBH, whereas the hot gas tries to escape through the channels in-between the cold filaments. Such fragmentation of accreting gas can assist in the formation of clouds around AGN, induce star-formation, and contribute to the observed variability of narrow-line regions. ",
        "introduction": "Feedback from active galactic nuclei (AGN) is considered to be one of the important processes in regulating galaxy growth over cosmic time. Several research groups have implemented AGN feedback in cosmological hydrodynamic simulations \\citep[e.g.][]{DiMatteo05, Booth09}, however they all have to assume a subgrid model for the accretion onto SMBH and feedback with uncertain parameters due to limited resolution in cosmological simulations.  Here we take a different approach and perform small-scale simulations of gas accretion onto SMBH using the GADGET-3 SPH code \\citep[originally described by][]{Springel05e} on small-scales of $r<200$\\,pc.  Our long-term goal is to develop a better model of AGN feedback for large-scale cosmological simulations based on the results of small-scale BH accretion simulations. In this article, we report our initial results of adiabatic Bondi accretion simulations, as well as those with radiative cooling and heating by X-rays from the central SMBH  \\citep{Barai11}. We run a large set of simulations of spherically symmetric Bondi accretion with different resolution, volume, and initial conditions. The Bondi problem is ideal for testing any hydrodynamic code, since it has a semi-analytic solution. ",
        "conclusions": "We find that GADGET-3 SPH code can reproduce the spherically symmetric Bondi accretion flow properly with mainly two limitations: 1) the gas particles escape from the outer boundary due to adiabatic expansion; 2) spurious heating is observed near the inner boundary around SMBH due to the artificial viscosity of SPH.  We also examined the impact of radiative cooling and heating due to X-rays.  The accretion flow roughly follows the Bondi solution when the central X-ray luminosity is relatively low, but the mass accretion rate at $r_{\\rm in}$ is enhanced by a factor of a few over the Bondi accretion rate due to cooling. The simulation starts to exhibit non-spherical fragmentation due to thermal instability once $L_X$ exceeds $\\simeq 0.01 L_{\\rm Edd}$. At $L_X=0.02 L_{\\rm Edd}$, the hot gas escapes in-between the cold filaments. When $L_X$ is further increased to $0.05 L_{\\rm Edd}$, the outflow completely dominates over the inflow, and most of the gas escapes from the computational volume.  A jet-like outflow feature is also observed at this $L_X$, which selects a preferential direction for hot gas to escape rapidly. Such non-spherical features of accreting gas can assist in the formation of clouds around AGN, induce star-formation, and contribute to the observed variability of narrow-line regions.  In the future, we plan to implement rotation, radiation pressure, and different initial geometry of gas distribution. We will measure the efficiencies of thermal, radiative, and kinetic feedback, and compare them with those measured by \\citet{Kurosawa09b} and those used in the cosmological simulations."
    },
    "1112/1112.3713_arXiv.txt": {
        "abstract": "{ Spectroscopic observations of the chemical abundances in metal-poor H~II regions provide an independent method for estimating the primordial helium abundance. H~II regions are described by several physical parameters such as electron density, electron temperature, and reddening, in addition to y, the ratio of helium to hydrogen. It had been customary to estimate or determine self-consistently these parameters to calculate y. Frequentist analyses of the parameter space have been shown to be successful in these parameter determinations, and Markov Chain Monte Carlo (MCMC) techniques have proven to be very efficient in sampling this parameter space.  Nevertheless, accurate determination of the primordial helium abundance from observations of H~II regions is constrained by both systematic and statistical uncertainties. In an attempt to better reduce the latter, and continue to better characterize the former, we apply MCMC methods to the large dataset recently compiled by Izotov, Thuan, \\& Stasi\\'nska (2007). To improve the reliability of the determination, a high quality dataset is needed.  In pursuit of this, a variety of cuts are explored. The efficacy of the He~I $\\lambda$4026 emission line as a constraint on the solutions is first examined, revealing the introduction of systematic bias through its absence. As a clear measure of the quality of the physical solution, a $\\chi^2$ analysis proves instrumental in the selection of data compatible with the theoretical model. In addition, the method also allows us to exclude systems for which parameter estimations are statistical outliers.  As a result, the final selected dataset gains in reliability and exhibits improved consistency.  Regression to zero metallicity yields Y$_{p}$ $=$ 0.2534 $\\pm$ 0.0083, in broad agreement with the WMAP result.  The inclusion of more observations shows promise for further reducing the uncertainty, but more high quality spectra are required. } ",
        "introduction": "Next to the cosmic microwave background radiation, standard big bang nucleosynthesis (SBBN) is the most robust probe of the early universe available  \\citep{wssok,osw,fs}. Furthermore, using the precise baryon density as determined by WMAP  \\citep{wmap,wmap10}, SBBN has effectively become a parameter free theory \\cite{cfo2}. As such, one can use SBBN to make relatively precise predictions of the initial light element abundances of D, $^{3}$He, $^{4}$He, and $^{7}$Li \\citep{cfo,coc,cfo3,coc2,cyburt,coc3,cuoco,serp,cfo5,coc4}. Therefore, an observational determination of these abundances becomes a test of the concordance between SBBN theory and the analyses of microwave background anisotropies. To test these predictions, the observed abundances must be determined with relatively high precision.  Unfortunately, there is a logarithmic relationship between the baryon to photon ratio, $\\eta$, and the primordial helium abundance, Y$_{p}$. Thus, any meaningful test of the theory requires a determination of Y$_{p}$ to an accuracy of $\\lesssim 1\\%$. The 7-year WMAP value for $\\eta$ is $(6.19 \\pm 0.15) \\times 10^{-10}$ \\citep{wmap10}.  For comparison, the SBBN calculation of \\citet{cfo5}, assuming the WMAP $\\eta$ and a neutron mean life of $885.7 \\pm 0.8$ s \\citep{rpp}, yields $Y_p = 0.2487 \\pm 0.0002$, a relative uncertainty of only 0.08\\%.  Here, we discuss the determination of Y$_{p}$ using observations of low metallicity H~II regions in dwarf galaxies.  By fitting the helium abundance versus metallicity, one can extrapolate back to very low metallicity, corresponding to the primordial helium abundance \\citep{ptp74}.  The oxygen to hydrogen ratio, O/H, commonly serves as a proxy for metallicity.  The difficulties in calculating an accurate and precise measure of the primordial helium abundance are well established \\citep{os01,os04,its07}. Previously, we introduced a new method based on Markov Chain Monte Carlo (MCMC) techniques \\citep[][AOS2]{AOS2}, and here, we systematically apply this technique to the data compiled in \\citep[][ITS07]{its07}.  Observations of the helium to hydrogen emission line ratios from extragalactic H~II regions provide a measure of the helium to hydrogen ratio, y$^{+}={n(He~II) \\over n(H~II)}$.  The uncertainty in y$^{+}$ follows from the statistical measurement errors of the helium and hydrogen emission line fluxes, in addition to a myriad of systematic effects.  Interstellar reddening, underlying stellar absorption, radiative transfer, and collisional corrections alter the observed flux, complicating the determination of y$^{+}$, and amplifying the uncertainty.  First, dust along the line of sight scatters the emitted photons (interstellar reddening).  Second, the stellar continuum juxtaposes absorption features under the emission lines (helium and hydrogen underlying absorption).  Third, before leaving the H~II region the emitted photons can be absorbed and re-emitted (radiative transfer).  Finally, in addition to the dominant recombination emission, collisional excitation also contributes to the emission (collisional corrections for helium and hydrogen).  Because none of these processes is directly measured, they cannot be removed independent of the observed emission lines and theoretical models.  As a result, the uncertainty on  y$^{+}$ must reflect the presence of these systematic effects.  Therefore, high fidelity spectra are required to accurately determine y$^{+}$, while simultaneously estimating the model parameters needed to correct for the listed systematic effects.  This desired confidence weighs against the need for larger sample sizes to decrease the uncertainty on Y$_p$ (and $dY/dZ$).  Motivated by the work of \\citet{itl94} and \\citet{ppr00}, the importance of Monte Carlo techniques was demonstrated in a ``self-consistent'' analysis method for determining the nebular helium abundance based upon six helium and four hydrogen lines \\cite{os01,os04}.   In preceding work, \\citet[][AOS]{AOS} updated and extended the physical model and integrated the helium and hydrogen calculations with the goals of increasing accuracy and removing assumptions.  Improving the statistical technique of AOS, \\citet[][AOS2]{AOS2} used a Markov Chain Monte Carlo (MCMC) analysis to explore a global likelihood function for all model parameters including the helium abundance.  The small sample size employed in AOS \\& AOS2 implied a clear avenue for improvement.  The small sample had limited capacity to constrain the slope of He vs. O/H and contributed to the large uncertainty on the determination of Y$_p$.  Since the publication of \\citet{os04}, Izotov \\etal have continued making new observations of H~II regions, significantly increasing their dataset \\citep[][]{it04, its07}.  This work makes use of that expanded dataset, aided by the MCMC analysis, which is used in the process of making quality cuts in the dataset.  Included in this process is the use of $\\chi^{2}$ as a measure of goodness-of-fit between the model and data.  A new, more rigorously selected dataset holds the promise of an improved determination of the primordial helium abundance.  Section \\ref{Model} briefly summarizes the model of AOS \\& AOS2.  A broad overview of the ITS07 dataset is given in \\S \\ref{Sample}, and in \\S \\ref{4026}, the subsets of the observations with and without the detection of He~I $\\lambda$4026 are compared.  Subsequently, further cuts on the dataset are investigated.  First, providing an indication of the quality of the fit, the $\\chi^{2}$ of the best-fit solution is scrutinized in \\S \\ref{Chi}.  Second,  in \\S \\ref{IllBehaved}, ill-behaved points with unreliable determinations are exhibited and discussed.  Third, \\S \\ref {Outliers} examines the distribution of the dataset in terms of model parameters used to correct for systematic effects.  After making cuts on the dataset, Y$_p$ is determined in \\S \\ref{Results}.  Finally, \\S \\ref{Conclusion} offers a discussion of the results, their exploration, and of further improvements in the determination of the primordial helium abundance. ",
        "conclusions": "\\label{Conclusion}  Because of its importance as a probe of big bang cosmology, the predictions of BBN are under constant scrutiny from new observations.  With the baryon-to-photon ratio fixed by the WMAP determination of the baryon density, BBN leads to a distinct set of predictions for the light element abundances. The prediction of the primordial D/H abundance is a major success when compared to determinations of D/H from quasar absorption system observations. To further test the theory, accurate \\he4 abundances are necessary as the BBN prediction for \\he4 is at the 0.1\\% level. Selecting an improved dataset from existing observations was a primary aim of this work.  Unfortunately, \\he4 abundance determinations from emission line fluxes carry significant systematic uncertainties.  In AOS2, we showed that one's ability to accurately estimate the set of physical parameters which are used to determine the \\he4 abundance is greatly improved with MCMC methods.  An eight-dimensional parameter space (which includes the helium abundance) is sampled to calculate a set of line fluxes which can be compared directly with observed fluxes.  Here, we have applied the MCMC method to a large sample of 93 objects, and then carefully screened these objects to remove systematic bias, poor quality model fits, and spurious physical results.  In this process, the importance of He~I $\\lambda$4026 was once again demonstrated.  Data lacking this line was found to be susceptible to a systematic increase in underlying Helium absorption and a corresponding bias to larger abundances. 70 of 93 objects in the initial sample contain the $\\lambda$4026 line, and were the subject of our further analysis.  One of the benefits of the MCMC method is our ability to directly test the goodness-of-fit for a particular set of solutions.  While we are able to find a best fit solution almost every object, there is no guarantee that that solution is actually a good (in a statistical sense) description of the data. As the best fit solution (fitting nine observables with eight parameters) is found by minimizing $\\chi^2$, the resulting minimum $\\chi^2$ is a direct measure of the goodness-of-fit. Unfortunately, 45 of the 70 objects resulted in a $\\chi^2$ per degree of freedom $> 4$. In our quest for high quality and reliable \\he4 data, these data were excluded from further analysis. On a positive note, the practical impact of the 95\\% confidence level cut preferentially removed outlier parameter solutions. Indeed, only a handful of objects surviving the $\\chi^2$ were flagged as outliers.  The cumulative effect of the careful selection led to a dataset with 14 objects.  However, the reliability of this Final Dataset has been methodically evaluated and is a distinct improvement over previous efforts. The precision of the resulting primordial helium abundance is increased, even if not to the level we hope to someday achieve.  The concordance of the various regression datasets incorporating flagged points underscores the robustness of the result and the improved quality of the dataset.  As we have emphasized before, the most promising avenue for significant future improvement in the primordial helium abundance determination lies with higher quality, high resolution spectra.  Decreased measurement uncertainties will clearly better constrain the solution and decrease systematic uncertainty.  Beyond that straightforward benefit, there are a couple of promising, related paths.  First, as discussed in AOS, high resolution spectra afford the chance to measure the absorption underlying the Balmer lines directly.  Removing underlying hydrogen absorption as a solution model parameter would reduce degeneracies and produce better constrained parameters.  Similarly, high signal to noise spectra open up the possibility of adding a weaker Helium or Hydrogen line to our analysis.  Again, increasing the number of lines constraining the parameters would reduce degeneracies and decrease systematic uncertainty.  That both of these improvements would directly benefit the neutral hydrogen fraction, the least well constrained parameter, provides additional impetus.  Finally, we also plan to quantify the correlations between the parameters for each object.  This will yield further insight into the goodness-of-fit.  In summary, we have demonstrated the rigor and transparency of MCMC methods in selecting the best available data to ultimately extract the primordial He abundance.  It supports a stringent screening of candidate spectra, yielding a robust sample.  The fruit of these labors is an improved determination of the primordial helium abundance, both in increased confidence in its accuracy and in a modest increase in its precision.  These uncertainties are still relatively large, however, and the case for needing higher quality spectra is further strengthened."
    },
    "1112/1112.3463_arXiv.txt": {
        "abstract": "Accurate calculations for the ground state of the molecular ions He$_2^{3+}$ and HeH$^{2+}$ placed in a strong magnetic field $B\\gtrsim 10^{2}$~a.u. ($\\approx 2.35 \\times 10^{11}$G) using the Lagrange-mesh method are presented. The Born-Oppenheimer approximation of zero order (infinitely massive centers) and the parallel configuration (molecular axis parallel to the magnetic field) are considered. Total energies are found with 9-10 s.d. The obtained results show that the molecular ions He$_2^{3+}$ and HeH$^{2+}$ exist at $B > 100$\\,a.u. and $B > 1000$\\,a.u., respectively, as  predicted in \\cite{Tu:2007} while a saddle point in the potential curve appears for the first time at $B \\sim 80$~a.u. and $B \\sim 740$~a.u., respectively. ",
        "introduction": "The neutron star surface is characterized by enormous magnetic fields, typically of $B\\sim 10^{12}\\,$G ($\\sim 10^3\\,$a.u.)\\footnote{$1$ a.u. $ \\approx 2.35\\times10^{9}$~G$ = 2.35\\times10^5$~T.}. Even more, for some young neutron stars - magnetars, the magnetic field achieves values up to $B\\sim 10^{16}\\,$G ($\\sim 10^7\\,$a.u.). It raises a general question about the behavior of matter under such extreme magnetic fields (see for a review e.g. \\cite{Lai:2001}). From a practical point of view, because of the intense magnetic fields, the content of a neutron star atmosphere became an issue. What kind of chemical species are formed there? This question got major importance after 2002 when the Chandra X-ray observatory discovered in the spectra of radiation of the isolated neutron star 1E1207.4-5209 two  absorption features at $\\sim 0.7$\\,keV and $\\sim 1.4$\\,keV.  Extensive theoretical studies of traditional and exotic molecular systems in the presence of a strong magnetic field have been carried on (for a review see \\cite{Tu:2007} and \\cite{phrep}). It was found that for magnetic fields $B\\leq 4.414\\times10^{13}\\,$G the most bound one-electron systems are the exotic molecular ions He$_2^{3+}$ and HeH$^{2+}$ besides the traditional molecular ion H$_2^+$. With the hypothesis of a neutron star atmosphere composed of mixed hydrogen-helium molecules, a model was proposed~\\cite{Tur:2005} that could explain the two absorption features discovered by the Chandra X-ray observatory. Even more, if the magnetic field strength is increased, new exotic finite molecular chains start to be present. In particular, for $B = 2.1\\times10^{4}$\\,a.u. the potential curve of the exotic system Li$_2^{5+}$ start to develop a minimum~\\cite{horop:2010}.  The aim of the present work is to make an accurate study of the two molecular ions ${\\rm He}_2^{3+}$ and  ${\\rm HeH}^{2+}$ in the presence of a strong magnetic field $B \\le 4.414\\times 10^{13}$\\,G where the relativistic corrections seem to be unimportant. The molecular axis is assumed to be aligned parallel to the magnetic field line because this configuration is optimal if a magnetic field is sufficiently strong \\cite{phrep, BJS09}. The Born-Oppenheimer approximation is considered where the positively-charged centers are supposed to be infinitely massive.  Following~\\cite{horop:2010} in the case of one-electron systems the {\\it variational method}  and the {\\it Lagrange-mesh method} are two methods which complement each other. The first one has been proven to be very efficient \\cite{phrep} with physically motivated trial functions. On the other hand, the {\\it Lagrange-mesh method} \\cite{h2malla} nowadays provides the most accurate results for some simple one-electron molecular systems in the presence of a strong magnetic field.  In practice, the implementation of the Lagrange-mesh method calculations is made easier and often feasible with {\\it a priori} knowledge of the equilibrium distance and the corresponding total energy. Such a knowledge is provided by a variational calculus. Thus, following the results presented in~\\cite{Tu:2007} the {\\it Lagrange-mesh method} is applied to check and further improve the variational results.  Atomic units $m_{e}= e = \\hbar = 1$ are used throughout, although the energy is given in Rydbergs. ",
        "conclusions": "It is presented a study in the non-relativistic frame and the Born-Oppenheimer approximation of zero order of the ground state of the molecular ions ${\\rm He}_2^{3+}$ and ${\\rm HeH}^{2+}$ in the presence of a strong magnetic field $B\\le 4.414\\times10^{13}$~G in parallel configuration. Although  the {\\it Lagrange-mesh method} is applied, the {\\it variational} results presented in~\\cite{Tu:2007} are taken into account as a reference point for the mesh calculations (see~\\cite{horop:2010}).  The  results give clear indications that the exotic molecular ions ${\\rm He}_2^{3+}$ and ${\\rm HeH}^{2+}$ begin to exist as metastable states starting at the threshold magnetic fields $B_{th} \\sim 10^2$ and  $\\sim 10^3$\\ a.u., respectively.  As the magnetic field increases the potential wells of both systems become sufficiently deep to keep more than one vibrational state. Eventually, for $B= 10^{13}$\\,G the ion  ${\\rm He}_{2}^{3+}$ becomes stable. The ${\\rm HeH}^{2+}$ ion remains unstable towards decay ${\\rm HeH}^{2+}\\rightarrow{\\rm He}^{+} + p$ in the whole domain of magnetic fields considered $B\\leq 4.414\\times10^{13}$\\,G. For both systems and all studied magnetic fields the obtained results are essentially more accurate then those presented in~\\cite{Tu:2007}.  These molecular ions have a similar behavior: when the magnetic field increases, each molecular ion becomes more compact and more bound. For all magnetic fields  the most bound helium-containing molecular system is ${\\rm He}_{2}^{3+}$, $i.e.$, $E_{b}^{{\\rm He}_{2}^{3+}} > E_{b}^{{\\rm HeH}^{2+}}$. From Tables \\ref{Thehe} and \\ref{Theh} it can be seen that the H$_2^+$ molecular system  placed in a magnetic field is the less bound two-center one-electron molecular system at $B \\geq 10^{13}$\\,G.  \\ack{The author wish to thank A.V. Turbiner for the interest to work, valuable discussions and a careful reading of the manuscript and D. Baye for reading the manuscript and for comments. The present work is supported in part by CONACyT (M\\'exico) through a postdoctoral grant.}"
    },
    "1112/1112.4596_arXiv.txt": {
        "abstract": "We present sub-arcsecond resolution HCN (4--3) and CO (3--2) observations made with the Submillimeter Array (SMA \\footnote{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica.}), toward an extremely young intermediate-mass protostellar core, MMS 6-main, located in the Orion Molecular Cloud 3 region (OMC-3). We have successfully imaged a compact molecular outflow lobe ($\\approx$1500 AU) associated with MMS6-main, which is also the smallest molecular outflow ever found in the intermediate-mass protostellar cores. The dynamical time scale of this outflow is estimated to be ${\\leq}$100 yr. The line width dramatically increases downstream at the end of the molecular outflow (${\\Delta}v{\\sim}$25 km s$^{-1}$), and clearly shows the bow-shock type velocity structure. The estimated outflow mass ($\\approx$10$^{-4}$ M$_{\\odot}$) and outflow size are approximately 2--4 orders and 1--3 orders of magnitude smaller, while the outflow force ($\\approx$10$^{-4}$ M$_{\\odot}$ km s$^{-1}$ yr$^{-1}$) is similar, as compared to the other molecular outflows studied in OMC-2/3. These results show that MMS 6-main is a protostellar core at the earliest evolutionary stage, most likely shortly after the 2nd core formation. ",
        "introduction": "Molecular outflows are ubiquitously observed in low- to high-  mass star forming regions \\citep[e.g.,][]{1999osps.conf..227B, 2007prpl.conf..245A}. They have a variety of characteristics such as size (0.01 pc--a few pc), mass (10$^{-4}$--10$^3$ M$_{\\odot}$), and outflow force (10$^{-7}$--1 M$_{\\odot}$ km s$^{-1}$ yr$^{-1}$) \\citep[e.g.,][]{2004A&A...426..503W}. These outflows are known to result from the entrainment of circumstellar gas, swept-up by the primary jet.  It is well accepted that primary jets are launched by magnetohydrodynamic (MHD) processes in the accretion disks surrounding a protostar \\citep[e.g.,][]{2007prpl.conf..277P, 2007prpl.conf..261S}. They play an important role to remove angular momentum from the circumstellar disk and thereby promote mass accretion onto the central star. Furthermore, recent  3D-MHD simulations by \\citet{2008ApJ...676.1088M} focus on the outflows/jets at the earliest evolutionary phase of star formation, between the prestellar phase and Class 0 phase such as outflows from the first adiabatic core and right after 2nd core formation. Theoretical works are able to resolve the jet launching points numerically. However, observational studies with current telescopes are not able to spatially resolve the region. The rarity of suitable targets (short life time at the earliest evolutionary phase) is a persistent problem.  MMS 6 \\citep[414 pc; the distance to Orion;][]{2007A&A...474..515M}, is located in the OMC-3 region \\citep[][; Hereafter, Takahashi et al. 2011a is denoted as Paper II]{1997ApJ...474L.135C, 1998ApJ...509..299L, 1999ApJ...510L..49J, 2005ApJ...626..959M, 2009ApJ...704.1459T, 2011ApJa}. MMS 6-main is the brightest and the most compact submillimeter continuum source in the MMS 6 region \\citep[][]{2009ApJ...704.1459T} as well as in the OMC-3 region \\citep{2011ApJc}. The sub-arcsecond 850 $\\micron$ continuum image (Paper II), spatially resolved a massive envelope (0.29 M$_{\\odot}$) and hot gas ($\\geq$50 K) in the central 120 AU. Comparisons with core models clearly showed the presence of a self-luminous source, which implies that MMS 6-main is an intermediate-mass protostellar core at the earliest evolutionary stage. However, neither molecular outflow, nor radio jet, nor infrared source has ever been detected at the source center in the previous studies \\citep[][Hereafter, Takahashi et al. 2008 is denoted as Paper I]{2005ApJ...626..959M, 2008ApJ...688..344T}.  In order to search for molecular outflows and study their characteristics, we have performed sub-arcsecond angular resolution observations toward MMS 6-main in the CO (3--2) and HCN (4--3) lines with the SMA. Our observations have achieved a factor of seven improvement in terms of beam size ($\\sim$50 times better in terms of the beam surface area) than the previous molecular line studies by \\citet{2009ApJ...704.1459T}. ",
        "conclusions": "An extremely compact molecular outflow was discovered toward an intermediate-mass protostellar core, MMS 6-main in the OMC-3 region. The velocity broadening at outflow tips, (${\\Delta}v{\\sim}25$ km s$^{-1}$), indicates the bow-shock type of velocity structure. The estimated outflow lobe size ($\\approx$1500 AU), mass($\\approx$10$^{-4}$ M$_{\\odot}$), and dynamical time ($\\leq$100 yr), clearly suggest that the outflow is compact, less massive, and young, while the estimated outflow force ($\\approx$10$^{-4}$ M$_{\\odot}$ km s$^{-1}$ yr$^{-1}$) is similar to those derived in other OMC-2/3 outflows. These results most likely suggest that the outflow is originated from a protostar (2nd core), shortly after its formation. These results are consistent with our recent SMA continuum observations at 850 $\\micron$, which imply that MMS 6-main is an intermediate mass protostellar core."
    },
    "1112/1112.6246_arXiv.txt": {
        "abstract": "We describe a major upgrade of a Monte Carlo code which has previously been used for many studies of dense star clusters. We outline the steps needed in order to calibrate the results of the new Monte Carlo code against $N$-body simulations for large $N$ systems, up to $N=200000$. The new version of the Monte Carlo code (called MOCCA), in addition to the old version, incorporates direct FewBody integrator for three- and four-body interactions, and new treatment of the escape process based on \\citet{FH2000}. Now stars which fulfil the escape criterion are not removed immediately, but can stay in the system for a certain time which depends on the excess of the energy of a star above the critical energy. They are called potential escapers. FewBody integrator allows to follow all interaction channels, which are important for the rate of creation of various types of objects observed in star clusters, and assures that the energy generation by binaries is treated in a meaner similar to the N-body model.  There are at most three parameters which have to be adjusted against N-body simulations for large $N$. Two (or one, depends on the chosen approach) connected with the escape process and one responsible for determination of the interaction probabilities. The adopted free parameters are independent on $N$. They allow MOCCA code to reproduce N-body results, in a reasonably precision, not only for the rate of cluster evolution and the cluster mass distribution, but also for the detailed distributions of mass and binding energy of binaries. Additionally, the code can follow the rate of formation of blue stragglers and black hole - black hole binaries. The code computes interactions between binaries and single stars up to a maximum separation $r_{pmax}$, and it is found that the MOCCA code needs rather large value of $r_{pmax}$ to get agreement with the N-body simulations.  The MOCCA code is at present the most advanced code for simulations of real star clusters. It can follow the cluster evolution in details comparable to N-body code, but orders of magnitude faster. ",
        "introduction": "\\label{sec:int}  This is the second paper in a new series of papers in which we attempt to describe the development of the MOCCA (MOnte Carlo Cluster simulAtor) code and its application to the simulations of star cluster evolution. The first in the series \\citep{HyG2011} described in detail recent developments of the previous version of the Monte Carlo code (\\citet{GHH2008}, and references therein) and the first results of simulations concerning blue stragglers (BSS) in an evolving star cluster environment. In this paper, we further develop the code and perform a very detailed comparison with N-body simulations of large $N$ stellar systems up to $N = 2 \\times 10^5$.  The MOCCA code \\citep{HyG2011} is at present the most advanced numerical code for stellar dynamical simulations, capable to follow evolution of real star clusters in detail comparable to N-body simulations, but orders of magnitude faster (several hours for $N = 2 \\times 10^6$). The dynamical ingredients of the Monte Carlo code are essentially the same as those described in \\citet{Gi2006} and \\citet{GHH2008}, whose code embodies several features introduced by \\citet{St1986}, whose code was in turn based on that originally devised by \\citet{He1971}.  Two main features distinguish the MOCCA code from the previous version of the Monte Carlo code: (i) it now incorporates dynamical interactions between binary and single stars and between pairs of binaries based on the FewBody integrator developed by \\citet{FCZR2004}; (ii) it replaces the treatment of the escape process in the static tidal field based on \\citet{Ba2001} by one in accordance the theory proposed by \\citet{FH2000}. The escape process is not instantaneous any more, an object needs time to find its way around the Lagrangian points $L_1$ and $L_2$ to escape. The MOCCA code incorporates most of the processes which are important during stellar system evolution e.g.: the relaxation process, the main engine of dynamical cluster evolution; stellar evolution according to \\citet{Hu2000} for the evolution of single stars, supplemented by the methods of \\citet{Hu2002} for internal evolution of binary stars (BSE code) and also simple approach for colliding stars according to McScatter interface to BSE by \\citet{HPH2006}; the escape process in the static tidal field of the parent galaxy; direct few body integrator to follow interactions between binaries and single stars and other binaries; mass segregated initial cluster configuration according to \\citet{BDK2008} and \\citet{SKB2008}.  There are several factors which motivate this work. Star clusters are the focus of many intensive observational campaigns \\citep[e.g.][and references therein] {bedinetal2001,bedinetal2003,grindlayetal2001,piottoetal2002, kaliarietal2003,kafkaetal2004,richeretal2004,andersonetal2006,miloneetal2011}, which are now turning to the examination of the parameters of their populations of different kinds of binaries, BSS and other \"peculiar\" objects. Dynamical models are needed for the design and interpretation of observational programmes: how is the period distribution and the spatial distribution of binaries affected by dynamical evolution?  What is the influence of environment and dynamical evolution on the formation of \"peculiar\" objects? Understanding the abundance, spatial distribution and channels of formation of BSS can only be attempted by a technique which follows simultaneously both BSS dynamics and internal evolution. While $N$-body technique may ultimately be the method of choice for such studies, systems of the size of most globular clusters are likely to remain beyond reach for some years, simply because of the number of stars and the size of the binary population. After all, it is only recently that the ``hardest\" open clusters M67 \\citep{Hurleyetal2005} and ``easiest\", loosely bound and distant globular cluster Palomar 14 \\citep{Zonoozietal2011} have been modelled at the necessary level of sophistication. To efficiently compute detailed models of large star clusters and to investigate the influence of initial parameters on a cluster's global and local observational properties we need a technique which is much faster than the N-body code and at the same time can give the same level of information about every object in the cluster as the N-body code does. The MOCCA code is such a technique.  One of the drawbacks of the non direct techniques (also the Monte Carlo one) compared to the N-body model is a necessity of use of free parameters which try to describe the complexity of physical processes naturally covered in the direct code. The most important free parameters (from the point of view of MOCCA) are connected with the relaxation process ($\\gamma$ coefficient in the Coulomb logarithm), escape process in the static tidal field and dynamical interactions between different objects (where the parameter, $r_{pmax}$, is the maximum pericentre distance between interacting objects for which few-body interactions are calculated explicitly). The usual method to determine the free parameters is a comparison with the results of N-body simulations. For the previous version of the Monte Carlo code the comparison was done only for small $N$ systems (up to $N = 24000$). The code was successfully used to simulate evolution of real star clusters: M67 \\citep{GHH2008}, M4 \\citep{HG2008}, NGC6397 \\citep[][]{GH2009,HG2009} and 47Tuc \\citep{GH2011}. Despite those successes there were some doubts connected with the $N$ scaling of the escape process, which implementation was based on \\citet{Ba2001}. To fully trust the MOCCA code it has to be tested for larger N, and not only for the global parameters like evolution of the total cluster mass or Lagrangian radii, but also against the properties and spatial distributions of binaries and BSS. That kind of comparison will show how far we can trust results of Monte Carlo simulations and which processes cannot be properly described in the framework of the MOCCA code.  This paper begins in Sec.2 with a summary of the features which have been added to the Monte Carlo scheme during the construction of the new version of the code, the MOCCA code. We also show there how we calibrate the free parameters of the MOCCA code with results of $N$-body simulations.  Next (Sec. 3) we describe the similarities and differences between MOCCA and N-body simulation and discuss the possible reasons for that. The final section summarises our conclusions, and discusses some limitations and future developments of the MOCCA code. ",
        "conclusions": "\\label{sec:conc}  In this paper we have presented an advanced Monte Carlo code (MOCCA) for the evolution of rich star clusters, including most aspects of dynamical interactions involving binary and single stars, internal evolution of single and binary stars and complicated process of escape in the static tidal field. The direct integration of a few body problem was introduced on the base of FewBody code developed by \\citet{FCZR2004}. The stellar and binary internal evolution was done according to BSE code \\citep[]{Hu2000,Hu2002}. The description of the escape process was based on the theory presented by \\citet{FH2000}. So, the MOCCA code is able to follow all channels of interaction up to binary-binary hierarchy including merging of stars, and the escape is not immediate any more, stars need time to find their way to escape from the $L_1$ and $L_2$ Lagrangian points. The probability of escape and probability for interaction are characterized by some free parameters which were adjusted by comparison of MOCCA and N-body simulation results for large $N$ systems, up to $N=200000$.  It was shown that the free parameters of the MOCCA code can be successfully calibrated against N-body simulations and that the free parameters do not depend on $N$. MOCCA code is not only able to follow evolution of the cluster total mass, Lagrangian radii and the core radius, but also is able to reproduce in an reasonably accuracy distributions of binary parameters and number of BH-BH binaries and BSS. It also reproduces well the results obtained by \\citet{Ba2001} for single mass tidally limited systems for the half-mass time and evolution of potential escapers. The code is able to cope with very diverse systems, from single mass, isolated system without primordial binaries to multi mass, tidally limited system with large fraction of binaries. It was shown that simplified and faster version of the MOCCA code (without direct FewBody integrator - MOCCA-NoFB) is a method of choice for projects which aim is to investigate the evolution of star clusters from the point of view their global properties. For other purposes, particularly when properties of \"peculiar\" objects and their distributions are in area of interests one should use slower MOCCA code. It is worth to note that MOCCA and MOCCA-NoFB simulations presented in this paper needs only about three and two hours, respectively, to be completed on a standard Opetron 2.4Ghz CPU.  Despite these successes the MOCA code has still some known shortcomings, which we summarise here. \\begin{enumerate} \\item{\\sl Higher-order multiples}:  It is widely argued that primordial triples and higher multiples should be incorporated into simulations along with primordial binaries. In any case, hierarchical triples form abundantly in binary-binary interactions \\citep{mikkola1984}. Such higher-order multiples are ignored in the present Monte-Carlo code, they are counted and artificially disrupted \\citep{HyG2011}. Hierarchical triples and higher-order multiples are planed to introduce in the new version of the code. \\item{\\sl Rotation}: the Monte Carlo code is based on spherical symmetry, and would require rather fundamental and very difficult reconstruction in order to cope with cluster rotation. Rotation accelerates the rates of core collapse and mass segregation \\citep[e.g.][and references therein]{FSK2006,KYLS2008}. In our models the absence of these rotational effects can be compensated by a modest alteration of the initial conditions. \\item{\\sl Static tide}: the effect of tidal shocks have been extensively studied  and it would be possible to add the effects as another process altering the energies and angular momenta of the stars in the simulations. The addition of tidal shocks will be more important when modelling Galactic globular clusters than open clusters, which usually are confined inside the Galactic disk. \\end{enumerate}  Despite these limitations, some of which are difficult to cure, the MOCCA code presented in this paper shows its potential power in simulations of real star clusters, from open clusters to rich globular clusters. Monte Carlo models are feasible in a reasonable time (a few days) for globular clusters, which are, still for some time, too large for direct N-body simulations. The MOCA code is able to provided data as detailed as N-body code can do. Only those two codes can provides such comprehensive information.  Even when $N$-body simulations eventually become possible, Monte Carlo models will remain as a quicker way of exploring the parameter space for the large scale $N$-body simulations."
    },
    "1112/1112.0782_arXiv.txt": {
        "abstract": "We examine the intrinsic scatter in the correlation between black hole masses and their host bulge masses, and find that it cannot be accounted for by mergers alone. A simple merger scenario of small galaxies leads to a proportionality relation between the late-time black hole and bulge masses, with intrinsic scatter (in linear scale) increasing along the ridge line of the relation as the square root of the mass. By examining a sample of 86 galaxies with well measured black hole masses, we find that the intrinsic scatter increases with mass more rapidly than expected from the merger-only scenario. We discuss the possibility that the feedback mechanism that operated during galaxy formation involved the presence of a cooling flow. ",
        "introduction": "\\label{sec:intro}  Relations between a supermassive black hole (SMBH) mass, $M_{\\rm BH}$, and properties of its host galaxy have been studied intensively for more than a decade. Two of the galactic properties most commonly correlated with $M_{\\rm BH}$ are the stellar mass of the spheroidal component (which we refer to as the bulge), $M_{\\rm G}$ (e.g. \\citealt{KormendyRichstone1995, Magorrian_etal1998, Laor2001, Hu2009, Graham2009}), and its stellar velocity dispersion, $\\sigma$ (e.g. \\citealt{Gebhardt_etal2000, MerrittFerrarese2001, Graham2008a, Graham2008b, Hu2008, Shen_etal2008, Gultekin_etal2009}). These relations are often assumed to have the form of a power law (i.e. linear when plotted on a log-log scale).  However, due to the scarcity of data, large measurement error and intrinsic scatter, there is still no consensus on which galactic parameter has the best correlation with $M_{\\rm BH}$. Some studies examine different combinations of $\\sigma$ and $M_{\\rm G}$; \\citet{Feoli_etal2010}, for example, argued that $M_{\\rm BH}$ is better correlated with the energy parameter, $M_{\\rm G} \\sigma^2$, than with $M_{\\rm G}$ or with $\\sigma$ alone. \\citet{Soker2011} found that a momentum-like parameter $\\mu \\equiv M_{\\rm G} \\sigma /c$ is well correlated with $M_{\\rm BH}$, as predicted by the penetrating-jet feedback model (\\citealt{Soker2009}; although this mechanism is based on energy balance).  Some studies consider the $M_{\\rm BH}$--$M_{\\rm G}$ relationship to be the fundamental one, and further argue that mergers between many low mass galaxies from an initially uncorrelated sample can lead to the observed correlation \\citep{Peng2007, Jahnke2011, Gaskell2011}. \\citet{Peng2007} claimed that the scatter decreases in logarithmic scale toward higher masses; motivated by his view that `fine tuning' is required in feedback, he claimed that AGN feedback is neither necessary nor desirable to produce the observed correlation. \\citet{Jahnke2011} added star formation to the growth of galaxies by mergers, and found the decrease in the relative scatter (i.e. in $M_{\\rm BH}/M_{\\rm G}$) to be slower than when only mergers are considered; however, the decrease in the scatter they obtained was still faster than the observed behaviour.   Following these arguments, we examine the scatter increase based on mergers. We expect, as will be further explained in Section \\ref{sec:scater}, that if only mergers cause BH growth, the intrinsic scatter of BH masses will increase as the square root of number of mergers $n^{1/2}$ while the mass increases as $n$ (thus, the relative scatter will decrease as $n^{-1/2}$). Note that this refers to the intrinsic scatter on {\\it linear scale} rather than logarithmic scale; the reason for discussing the intrinsic scatter in mass-mass relationships in linear scale will also be made clear in Section \\ref{sec:scater}. We also note that in this work, $n$ of a particular galaxy is the number of building blocks that merged over time to form it (cf. `generation number', which is $\\log_2 n$).  In the current paper, we examine the behaviour of the scatter in the $M_{\\rm BH}$--$M_{\\rm G}$ relationship with mass in order to asses the importance of mergers. The sample of galaxies is discussed in Section \\ref{sec:sample}, and the results in Section \\ref{sec:scater}. In Section \\ref{sec:summary} we discuss the implications of the results to the feedback mechanism and give a brief summary.  \\begin{figure*} {\\includegraphics{carmit-Msig.eps}} {\\includegraphics{carmit-MM.eps}} \\caption{The $M_{\\rm BH}$--$\\sigma^4$ (left) and $M_{\\rm BH}$--$M_{\\rm G}$ (right) relations for our sample of 86 galaxies. Here $M_{\\rm BH}$, $M_{\\rm G}$, and $\\sigma$ are the black hole mass, bulge mass, and stellar velocity dispersion of the galaxies, respectively, and the logarithm is base 10; the data are from Table \\ref{tab:Galaxytabel}. The solid and dashed lines are the best-fitting power law relations using the $x$- and $y$-scatter models respectively (see text). The dotted line in each panel has a slope $\\beta=1$.} \\label{fig:corsigmamg} \\end{figure*} ",
        "conclusions": "\\label{sec:summary}  Our objective was to examine the claim that the SMBH mass to bulge mass ($M_{\\rm BH}$--$M_{\\rm G}$) correlation is predominantly a result of mergers of low mass galaxies. The merging process leads to an intrinsic scatter in the $M_{\\rm BH}^{1/2}$--$M_{\\rm G}^{1/2}$ relation which is uniform as given by equation (\\ref{eq:Delta1}). We checked this prediction with a sample of 86 galaxies, listed in Table \\ref{tab:Galaxytabel}. Some properties of the sample are presented in Figs. \\ref{fig:corsigmamg} and \\ref{fig:cor1}, and in Table \\ref{tab:slopes}. In Fig. \\ref{fig:cor2} we show the $M_{\\rm BH}^{1/2}$--$M_{\\rm G}^{1/2}$ correlation and the residuals.  Our main result is that the intrinsic scatter increases with mass more rapidly than expected in a merger-only scenario. Since mergers occur between all types of galaxies, we include all types of galaxies in our study. However, we checked our analysis for ellipticals, spirals, barred, unbarred, AGNs and inactive galaxies and found the main conclusion to hold for each subgroup separately.  While biasing is always a worry when considering these correlations (especially given the scarcity of data), it does not seem like a selection effect can artificially produce the result of our Fig. \\ref{fig:cor2}: the pure merger scenario is only true when the residuals are scattered uniformly; this uniform scatter cannot be smaller than what we see at the high mass end, which is $2\\times 10^4$ in units of ${\\rm M}_\\odot^{1/2}$. Missing galaxies at the high end (e.g. small $M_{\\rm BH}$) can only increase the scatter. At the low mass end, it is only reasonable that we `miss' black holes with masses below the ridge line, but even if they exist, they cannot be scattered over $2\\times 10^4~{\\rm M}_\\odot^{1/2}$ since at the low mass end, $M_{\\rm BH}$ is smaller than that. Also, note the recent work by \\citet{Gultekin_etal2011}, who investigated the possibility that selection effects (namely missing very low mass SMBHs) bias the measurements, and reached the conclusion that it is not likely that the published relations are biased. Our sample was not selected in any special way and largely overlaps with theirs.  Another argument why the pure merger scenario is discrepant with the observations is that the $M_{\\rm BH}$--$M_{\\rm G}$ relationship might not be a true proportionality relation, but have higher order terms. This is indicated by the fact that in two out of the three fitting methods (scatter models) for this relation in logarithmic scale, a slope of 1 (indicating a proportionality relation rather than an arbitrary power law) was more than three standard deviations from the best fitting slope (see Table \\ref{tab:slopes}).  We do not dispute the claim that mergers influence the $M_{\\rm BH}$--$P_{\\rm G}$ relations, where $P_{\\rm G}$ represents any property of the host galaxy. We only argue that based on our results, mergers cannot be the dominant cause of correlation. Rather, we expect this correlation to be determined mainly by a feedback mechanism that operates on all scales, from small galaxies to galaxy clusters.  Let us demonstrate a simple plausible implication of the hypothesis that the same feedback mechanism operates on all scales. It is hard to observe the processes that took place during galaxy formation at the high redshift universe, and many models exist, based on radiation and/or jets from AGNs and/or star formation, e.g. \\cite{SilkRees1998}, \\cite{King2003}, \\cite{Soker2009}, and \\cite{SilkNusser2010}. The situation is much better with cooling flow clusters, that show both star formation and AGN activity. The central galaxy of the Virgo cluster, M87 (NGC 4486), and the central galaxy of Fornax, NGC 1399, are two cooling flow clusters in our sample.  We take the energy in jets from AGN outbursts to be due to accretion on to the SMBH with an efficiency $\\eta_j$: $E_{\\rm jets} = \\eta_j M_{\\rm acc} c^2$. The jets energy is inferred from the energy required to inflate X-ray deficient cavities (bubbles). We here assume that the jet's energy source is the accreted mass (but see \\citealt{McNamara2011}). Under this assumption, \\cite{McNamara2011} find the median value of the molecular mass in the central galaxy to mass accreted during AGN outburst ratio to be $M_{\\rm mol}/ M_{\\rm acc} \\simeq 700  (\\eta_j /0.04)^{-1}$, where the scaling of $\\eta_j=0.04$ is based on the results of \\cite{Soker2011}. This is very close to the ratio $M_{\\rm G}/M_{\\rm BH} = 590 \\pm 70$ found in our sample. The feedback process in cooling flow clusters seems to be able to produce the $M_{\\rm BH}$--$M_{\\rm G}$ correlation. This might hint that the feedback mechanism that operated during galaxy formation involved the presence of a cooling flow, as suggested in an earlier paper \\citep{Soker2010}.  \\medskip  This research was supported by the Asher Fund for Space Research at the Technion, and the Israel Science Foundation.  \\newcommand{\\mnras}{MNRAS} \\newcommand{\\apj}{ApJ} \\newcommand{\\apjl}{ApJ}"
    },
    "1112/1112.2787_arXiv.txt": {
        "abstract": "The proper scientific analysis of a large amount of stellar spectra requires certain capabilities of the analysing tool (e.g. precise semi-automatic fitting for normalisation of the continuum or line-list assisted measurement of spectral lines) as well as  flexible list-driven datafile handling.  While most astronomical legacy packages comprise powerful analysing and data management features allowing rapid processing of quite complex data, current VO tools are lacking support of  even basic capabilities commonly used in stellar spectroscopy.  Especially the high resolution optical spectroscopy of stars with rapid line profile variations or of those with complicated emission profiles benefits from a number of specific methods unsupported in todays Virtual Observatory (VO) tools and even lacking definition in VO spectral protocols.  In our contribution we  identify these techniques, describe their possible implementations and finally give a short overview of several VO-compatible tools emphasising their deficiencies and comparing with the capabilities of common legacy packages. ",
        "introduction": "There is a wealth of calibrated astronomical spectra accessible in current data archives, however most of them are not suitable for direct physical analysis.  Being prereduced by some automatic pipeline or individually by manual reduction, they  are mostly stored in archives in form of  FITS file or ASCII table as the relation of intensity in arbitrary numbers (instrumental counts, digital numbers etc.) and wavelength (or frequency).  The  scientific analysis of such  spectra requires further processing by the variety of different methods. In certain studies a huge number of spectra has to be collected from different servers (e.g. in different spectral regions ) and transformed into  common units. Often the transformed spectra or variables measured in them (e.g. equivalent width) are visually inspected as the time series ordered by time of observation or by the orbital phase.  The tools with built-in VO protocols may considerably increase the efficiency of astronomical analysis by eliminating the data collection, aggregation and  conversion and by allowing the astronomer to fully  concentrate on the visual properties of data (e.g. searching the variability of line profiles or constructing the radial velocity curve). To be accepted quickly by the wide astronomical community a number of methods and recipes has to be implemented in a manner very similar to  the  current legacy applications. We attempted to identify the most useful methods and give  hints for their implementation in VO clients. ",
        "conclusions": "Astronomical spectroscopy uses a wide range of techniques with different level of complexity to achieve its final goal --- to  estimate the most precise and reliable  information about celestial objects. The  large part of spectroscopic analysis today  has been accomplished by several independent non VO-compatible legacy packages, where  each works  with different local files in its own data format. Analysis of large number of spectra is thus very tedious work requiring good data bookkeeping.  Accomplishing the analysis in VO infrastructure may benefit from automatic aggregation of distributed archive resources (e.g. the multispectral research), seamless on-the-fly data conversion, common interoperability of all tools  (using PLASTIC protocol) and powerful graphical visualisation of measured and derived quantities (e.g. in VOPlot or Mirage).  By introduction of  modern  VO-aware tools into the  astronomical spectral analysis  a remarkable increase of effectivity of astronomical research can be achieved."
    },
    "1112/1112.5819_arXiv.txt": {
        "abstract": "The extension of the minimal standard model by three right-handed sterile neutrinos with masses smaller than the electroweak scale ($\\nu$MSM) is discussed in a $Q_6$ flavor symmetry framework. The lightness of the $\\kev$ sterile neutrino and the near mass degeneracy of two heavier sterile neutrinos are naturally explained by exploiting group properties of $Q_6$. A normal hierarchical mass spectrum and an approximately $\\mu$-$\\tau$ symmetric mass matrix are predicted for three active neutrinos. Nonzero $\\theta_{13}$ can be obtained together with a deviation of $\\theta_{23}$ from the maximality, where both mixing angles are consistent with the latest global data including T2K and MINOS results. Furthermore, the tiny active-sterile mixing is related to the mass ratio between the lightest active and lightest sterile neutrinos. ",
        "introduction": "Both the establishment of neutrino oscillation phenomena and the evidence of nonluminous dark matter (DM) demand physics beyond the standard model (SM). On the one hand, compelling evidences from current solar, atmospheric, reactor and accelerator neutrino experiments have told us that neutrinos are massive and lepton flavors are mixed \\cite{PDG}. On the other hand, various cosmological observations have revealed that DM is five times more abundant than normal matter and accounts for about one quarter of the Universe (i.e., $\\Omega_{\\rm DM}\\simeq0.227\\pm0.014$ \\cite{WMAP10}). To explain both problems above, an extension of the SM by three right-handed sterile neutrinos with masses smaller than the electroweak scale ($\\nu$MSM) was first proposed by Asaka {\\it et al.} \\cite{NuMSM1}. The smallness of active neutrino masses is described by the canonical seesaw mechanism \\cite{type1} and one light right-handed sterile neutrino at the $\\rm keV$ scale acts as a candidate of warm dark matter (WDM). Moreover, the model can explain the baryon asymmetry in the Universe through the oscillations \\cite{BANO,BAU} of two heavier right-handed sterile neutrinos with masses at the $\\gev$ scale.  Despite the above phenomenological successes, one unsatisfactory point in the $\\nu$MSM is the lack of a natural explanation for large mass splitting between the $\\kev$ sterile neutrino and the heavier ones. Moreover, the oscillation mechanism for the baryon asymmetry demands strong mass degeneracy between the heavier sterile neutrinos, but the $\\nu$MSM is impotent to be able to explain its origin. In order to resolve theses issues, some interesting ideas have been proposed in the context of a $\\rm U(1)$ flavor symmetry \\cite{U1FS}, the Froggatt-Nielsen mechanism \\cite{FN}, the split seesaw mechanism \\cite{Split} and grand unified theories \\cite{CR}. In the present work, we introduce a non-Abelian discrete flavor symmetry and try to understand the mass splitting and degeneracy due to group properties of the flavor symmetry. The mass degeneracy of two heaver sterile neutrinos could be interpreted as a sign that they constitute a doublet representation of the flavor symmetry, while it may be natural to assign a singlet representation to the $\\kev$ sterile neutrino. Furthermore, if the singlet representation is a complex one, we can prohibit a bare mass term of the $\\kev$ sterile neutrino because of the Majorana nature and may be able to generate a suppressed mass term from higher-dimensional operators. Inspired by these clues, we employ the $Q_6$ group as our flavor symmetry\\footnote{Previous studies about $Q_6$ can be found in Ref. \\cite{kubo}.} since it is the smallest finite group\\footnote{See Ref. \\cite{Frampton} for the classification of discrete groups upto $g=31$, and Ref. \\cite{FSrev} for the recent reviews of non-Abelian discrete flavor symmetries.} which contains both a complex singlet and a real doublet representations. In addition to $Q_6$, we introduce two auxiliary $Z_{N}$ symmetries in order to handle the order of magnitude of some small parameters as well as to forbid unwanted terms. We find that our model predicts a normal hierarchical mass spectrum and an approximately $\\mu$-$\\tau$ symmetric mass matrix for three active neutrinos.  The remaining part of this paper is organized as follows. In section II we present the framework of the $\\nu$MSM with a $Q_6 \\times Z_3 \\times Z_2$ flavor symmetry. Section III is devoted to the realization of the seesaw mechanism and section IV to the diagonalization and the resulting neutrino masses and mixing matrix. A numerical analysis with the focus on nonzero $\\theta_{13}$ is also illustrated in section IV. Finally, we give a summary in section V.% ",
        "conclusions": "In this work, we have proposed a $Q_6$ flavor symmetry realization for the $\\nu$MSM in the presence of two auxiliary $Z_N$ symmetries and succeeded in naturally explaining the lightness of the $\\kev$ sterile neutrino and the mass degeneracy of the two heavier sterile neutrinos. A normal hierarchical mass spectrum and an approximately $\\mu$-$\\tau$ symmetric mass matrix are predicted for three active neutrinos. Nonzero $\\theta_{13}$ can be obtained together with a deviation of $\\theta_{23}$ from the maximality, where both mixing angles are consistent with the latest global data including T2K and MINOS results. Finally, we have derived a tiny active-sterile mixing related to the mass ratio between the lightest active and lightest sterile neutrinos.  The $\\nu$MSM can explain the active neutrino masses, the candidate of dark matter and the baryon asymmetry in the Universe in a unified and elegant way. Although there are already some models in which the $\\nu$MSM is extended by flavor symmetries, our model has more direct connections with the masses and mixing patterns of three active neutrinos. Our realization can also be modified to accommodate the $\\ev$ scale sterile neutrinos \\cite{eVmodel}, which are more or less hinted at by current experimental \\cite{Schwetz4} and cosmological \\cite{Raffelt} data. We shall examine this case with a specific flavor model elsewhere \\cite{TA}."
    },
    "1112/1112.2925_arXiv.txt": {
        "abstract": "We measured projected rotational velocities of more than a hundred apparently single sdBs. A comparison is made with sdB stars in binary systems with orbits so wide, that tidal interaction becomes negligible. All of these stars are slow rotators ($v_{\\rm rot}\\sin{i}<10\\,{\\rm km\\,s^{-1}}$) with EC\\,22081$-$1916 being the only exception. This single star has the highest projected rotational velocity ever measured for an sdB ($v_{\\rm rot}\\sin{i}=163\\,{\\rm km\\,s^{-1}}$) and might have been formed by a merger event. The merger of a red-giant core and a low-mass, main-sequence star or substellar object during a common envelope phase fits particularly well with observations. The implications of our results for hot subdwarf formation are briefly discussed. ",
        "introduction": "The formation of hot subdwarf stars (sdBs) is still unclear. Several different single star and close binary formation scenarios are discussed to explain the enhanced mass loss in the red-giant phase necessary to form sdBs. As an alternative, the merger of compact objects (e.g. helium white dwarfs, He-WDs) has been proposed as a possible origin for single sdBs (for the most recent review see Han et al. these proceedings). Whether and how the progenitor star loses or gains angular momentum should affect the rotational properties of the resulting sdB. Determining the $v_{\\rm rot}\\sin{i}$-distribution of sdBs should therefore be useful to shed light on their formation.  While the rotational properties of horizontal branch stars both in globular clusters and in the field are thoroughly examined \\citep[e.g.][]{behr03}, there is no systematic study for extreme horizontal branch (EHB) stars yet. Most of the sdB stars where $v_{\\rm rot}\\sin{i}$-measurements are available, are slow rotators \\citep[$<10\\,{\\rm km\\,s^{-1}}$, e.g.][]{heber00, edelmann01}. Here we determine the projected rotational velocities of more than a hundred sdB stars by measuring the broadening of metal lines. This study complements the investigation of the rotational properties of close binary sdBs \\citep{geier10}.  \\begin{figure} \\begin{center} \\includegraphics[width=8cm]{distrib_single.eps} \\includegraphics[width=8cm]{distrib_comp.eps} \\end{center} \\caption{Upper panel: Distribution of ${v_{\\rm rot}\\sin\\,i}$ for 72 single stars from our sample. Lower panel: Distribution of ${v_{\\rm rot}\\sin\\,i}$ for 15 sdBs with visible companions.} \\label{fig:rotsingle} \\end{figure} ",
        "conclusions": "The main conclusion to be drawn from our results is that the rotational properties of single sdBs, and both wide and close binary sdBs with orbital periods of the order of a few days are very similar. All of these objects have very small $v_{\\rm rot}\\sin{i}$. However, according to the standard binary formation scenario for sdB stars, these stars should have been formed in quite different ways. In binary systems, mass transfer is proposed to be responsible for the loss of the hydrogen envelope in the red-giant phase. Depending on the initial conditions, this mass loss can be either stable or unstable. Since envelope mass and angular momentum are lost by the sdB progenitor in both cases, it is not surprising that the remnant rotates slowly.  Single sdB stars on the other hand have been proposed to be the remnants of the merger of two He-WDs. Although the angular momentum evolution of the merger product is unknown in detail, it is surprising to find single sdBs stars rotating with basically the same $v_{\\rm rot}\\sin{i}$ as in binaries despite the totally different formation channel. Whether this intriguing behaviour can be explained in the framework of the merger channel remains to be seen (see Podsiadlowski et al. these proceedings).  EC\\,22081$-$1916 is unique among $\\simeq100$ slowly rotating sdB stars and may be the rare outcome of a common envelope merger event."
    },
    "1112/1112.0327_arXiv.txt": {
        "abstract": "We present forecasts on the capability of future wide-area high-sensitivity X-ray surveys of galaxy clusters to yield constraints on the parameters defining the Dark Energy (DE) equation of state (EoS). Our analysis is carried out for future X-ray surveys which have enough sensitivity to provide accurate measurements of X-ray mass proxies and Fe-line based redshifts for about $2\\times 10^4$ \\numero clusters. We base our analysis on the Fisher Matrix formalism, by combining information on the cluster number counts and power spectrum, also including, for the first time in the analysis of the large scale cluster distribution, the effect of linear redshift space distortions. This study is performed with the main purpose of dissecting the cosmological information provided by geometrical and growth tests, which are both included in the analysis of number counts and clustering of galaxy clusters.  We compare cosmological constraints obtained by assuming different levels of prior knowledge of the parameters which define the relation between cluster mass and X-ray observables. This comparison further demonstrates the fundamental importance of having a well calibrated observable-mass relation and, most importantly, its redshift evolution. Such a calibration can be achieved only by having at least $\\sim 10^3$ \\numero net photon counts for each cluster included in the survey, with sufficient angular resolution. We show that redshift space distortions in the power spectrum analysis carry important cosmological information also when traced with galaxy clusters. We find that the DE FoM increases by a factor of 8 \\numero when including the effect of such distortions. Besides confirming the potential that large cluster surveys have in constraining the nature of Dark Energy, our analysis emphasizes that a full exploitation of the cosmological information carried by such surveys requires not only a large statistic but also a robust measurement of the mass proxies and redshift for a significant fraction of the cluster sample, which ought to be derived from the {\\sl same} X-ray survey data. This will be possible with future X-ray surveys, such as those envisioned with the Wide Field X-ray Telescope, with an adequate combination of survey area, sensitivity and angular resolution. ",
        "introduction": "\\label{s:intro}  A number of independent cosmological observations, ranging from type-Ia supernovae (SNIa) \\citep[e.g.][]{riess07,perlmutter99}, to Cosmic Microwave Background (CMB) anisotropies \\citep[e.g.][]{komatsu11,larson11} and Large Scale Structure (LSS) \\citep[e.g.][]{vikhlinin09c,reid10}, convincingly show that the Universe is undergoing a phase of accelerated expansion. One of the main challenges of modern cosmology is in fact understanding the source of such acceleration. To this purpose, a number of models have been proposed that modify the two pillars of modern cosmology, general relativity and the standard model of fundamental interactions \\citep[e.g.][and references therein]{silvestri09}. Models that modify the latter and, therefore, the energy-momentum tensor in the Einstein equation are, for example, the scalar field models like quintessence \\citep[e.g.][]{caldwell09,doran03,zlatev99,ratra88}, k-essence \\citep[e.g.][]{tsujikawa10,mukohyama04}, coupled Dark Energy \\citep[e.g.][]{amendola00} and Chapligin gas \\citep[e.g.][]{bento02}. Models that modify general relativity can produce the cosmic acceleration without including a Dark Energy (DE) component, but they should also satisfy stringent constraint from local (e.g. solar system) tests of gravity.  Examples of such models are the braneworld models, like DGP \\citep[e.g.][]{dgp,movahed09}, $f(R)$ theories \\citep[e.g.][]{hu_sawicki07,brax08,appleby10,sotiriou10} and scalar-tensor theories \\citep[e.g.][]{skordis09}.  The first model proposed to explain cosmic acceleration was based on the introduction of a cosmological constant $\\Lambda$ that can be thought as a fluid with negative pressure and equation of state (EoS) which is constant in space and time, $p=w\\rho c^2$ with $ w =- 1$. However, a cosmological constant able to drive the accelerated expansion leads to the well known problem that $\\Lambda$ needs to be so tiny with respect to any natural energy scale, that there is no theoretical justification for it. Thus, a plethora of models, characterized by different parametrizations of the DE EoS evolution, have been proposed \\citep[e.g.][]{wetterich04}. In principle, it is possible to distinguish among different DE models by combining different cosmological probes, both based on the geometry of the universal background and the growth of density perturbations (e.g., \\citealt{albrecht06}).  In this context, clusters of galaxies have long been recognized as potentially powerful probes of the nature of DE and cosmological models in general \\citep[e.g.][and references therein]{allen11,lombriser10,manera06}. Clusters of galaxies provide cosmological information in a number of different ways. The evolution of the cluster space density depends on cosmological parameters through both the linear growth rate of density perturbations and  the redshift dependence of the volume element. The large-scale clustering of galaxy clusters is also sensitive to cosmological parameters, through the growth rate of perturbations, which affects both the bias parameter and the redshift-space distortions, as well as by sampling the shape of the underlying Dark Matter (DM) power spectrum over a broad range of wavenumbers. As of today, relatively small samples of $\\sim 100$ X-ray selected clusters, originally identified out to $z\\simeq 0.8$ by the ROSAT satellite and then followed-up by \\chandra\\ to obtain robust mass estimates, have provided interesting DE constraints \\citep[e.g.,][]{vikhlinin09c,mantz10a}, which  complement and agree with those from CMB and SNIa observations \\cite[see also][]{rapetti05}.  More recently, by using a complete sample of $\\mincir 1000$ nearby clusters at $z\\mincir 0.2$, identified in the ROSAT All-Sky Survey \\citep[RASS;][]{truemper93},  the large-scale power spectrum has been constructed over a fairly wide scale range \\citep[e.g.,][]{antolinez11}.  While the first Sunyaev-Zeldovich (SZ) surveys have now started producing cluster samples of similar sizes, \\citep[e.g.,][]{ACT11,SPT11,PlanckESZ_sample}, the next generation of X-ray (e.g., eROSITA\\footnote{http://www.mpe.mpg.de/heg/www/Projects/erosita/index.php}, WFXT\\footnote{http://www.wfxt.eu/}) and optical (e.g., DES\\footnote{http://www.darkenergysurvey.org/}, EUCLID\\footnote{http://sci.esa.int/science-e/www/area/index.cfm?fareaid=102}) surveys are expected to increase by orders of magnitude the number of galaxy clusters, further extending the redshift range over which they trace the growth of cosmic structures. Such large cluster surveys have the potential of placing very tight constraints on different classes of DE models, possibly finding signatures of departures from the standard $\\Lambda$CDM predictions. Several studies have dealt with constraints of DE models from future cluster surveys focusing on the impact of uncertainties in cluster mass estimates \\citep[e.g.][]{battye03,majumdar04,lima05,lima07,cunha09,cunha10, basilakos10}. All these analyses generally assume that, when a cluster is detected and included in a survey, the observable (i.e., X-ray luminosity, optical richness) on which the detection is based can be related to the actual cluster mass through a suitable relation, whose functional form is assumed to be known and depends on a set of additional parameters.  A more conservative approach would instead require that for all clusters included in a survey detailed follow-up observations are carried out to calibrate suitable and robust mass proxies. Examples of such mass proxies in X-ray surveys include the total gas mass \\citep[e.g.][]{mantz09II} or the product of gas mass and temperature, the so-called $Y_X$ parameter \\citep{kravtsov06}, which can be computed when a relatively large number of photons is available. Clearly, while measuring flux for an X-ray extended source requires only $\\sim 50$ photons, or less for missions with low background, the measurement of robust mass proxies requires at least $\\sim 10^3$ photons.  In this paper, we will derive forecasts for constraints on the EoS of DE models from future X-ray surveys. We envision that these surveys will be carried out with a telescope with high-enough sensitivity to readily provide robust measurements of mass proxies and Fe-line based redshifts for $\\magcir 10^4$ clusters, which are all characterized by ``\\chandra-quality'' data, thus avoiding the need of external and time-consuming follow-up observations. For example, the Wide Field X-ray Telescope (WFXT) concept was developed to meet such a requirement. This telescope combines large collecting area with a large field-of-view and sharp point spread function (PSF) approximately constant over the entire field of view \\citep[e.g.,][]{giacconi09,murray10,rosati10}.  We adopt the specifications of the WFXT surveys, as an example of next generation X-ray cluster surveys,  and we compute cosmological forecasts using the well-established Fisher Matrix approach \\citep[e.g.,][]{dodelson03}, to combine information from cluster number counts and large-scale clustering. We will quantify the constraints expected on DE models and their dependence on the knowledge of the relation between X-ray observables and cluster mass, for a range of survey strategies (i.e. depth vs. sky coverage). In the course of this study, we will discuss how number counts and the power spectrum of the large-scale distribution of clusters convey cosmological information.  As a novel contribution in this paper, we will show how the detection of Baryonic Acoustic Oscillations (BAOs) and redshift space distortions (RSDs) on cluster scales can significantly contribute to constrain cosmological parameters, similarly to a number of previous studies based on the large scale distribution of galaxies \\citep[e.g.,][]{guzzo08,rassat08,stril10,wang10}.    This paper is structured as follows. In Section 2, we briefly describe the parametrizations of the DE EoS that we use, and describe our Fisher-Matrix approach for cluster number counts and power spectrum.  In Section 3, we first describe the characteristics of the cluster survey, and derive the forecasts on the constraints on DE EoS parameters. In this section, we will also quantify the impact that uncertainties in the scaling relation between X-ray observables and cluster mass have on such constraints.  We discuss our results and present our conclusions in Section 4. ",
        "conclusions": "In this paper, we presented forecasts on the capability of future wide-area high-sensitivity X-ray surveys of galaxy clusters to yield constraints on the parameters defining Dark Energy (DE) equation of states. We considered the standard equation of state (EoS) provided by equation (\\ref{eq:eos}) and the class of Early DE models of equation (\\ref{eq:eos_ede}). Our analysis was carried out for future X-ray surveys which have enough sensitivity to provide accurate measurements of X-ray mass proxies and Fe-line based redshifts for approximately $2\\times 10^4$ \\numero clusters, thus extending by more than two \\numero orders of magnitude the size of the cluster samples presently used to derive cosmological constraints \\citep[e.g.,][]{allen11}. We used the Wide Field X-ray Telescope (WFXT) \\citep[e.g.,][]{giacconi09,murray10,rosati10} as a reference mission concept along with the Wide (20000 sq.deg.), Medium (3000 sq.deg.) and Deep (100 sq.deg.) survey configurations (see Table \\ref{t:sur}). We based our analysis on the Fisher Matrix formalism, by combining information on the cluster number counts and power spectrum, also including the effect of linear redshift-space distortions (RSDs). This analysis has been carried out with the main purpose of dissecting the cosmological information provided by geometrical and growth tests, which are both included in the analysis of number counts and clustering of galaxy clusters. The main results of this study can be summarized as follows.  \\begin{description} \\item[(a)] When constraining the parameters of the DE EoS of equation (\\ref{eq:eos}), we further demonstrate the fundamental importance of having a well calibrated X-ray observable-mass relation and, most importantly, its redshift evolution.  We verify that the Figure of Merit (FoM) of the DE EoS increases up to 106 \\numero when we assume a strong prior on the mass parameters, as resulting from a precise and robust calibration of the mass-observable relation, with respect to the case in which no such prior is available (FoM = 61\\numero) (see Table (\\ref{t:nuis})). Such an internal calibration can be readily achieved from the same X-ray data by having at least $\\sim 10^3$ \\numero net photon counts for each cluster included in the survey. \\item[(b)] We find that the Medium survey is the one carrying most of the constraining power (Table (\\ref{t:de})), since it is expected to yield the largest number of clusters out to redshift $z\\sim 1$. As such, the Medium survey shows the tightest constraints on the evolution of the DE EoS ($\\sigma_{w_0}=0.097$ \\numero and $\\sigma_{w_a}=0.54$\\numero) and the corresponding highest Figure of Merit (FoM $= 60$\\numero).  The Deep survey, although covering a much smaller area than the Wide survey, adds an important contribution to constraining DE parameters (FoM $=20$\\numero). \\item[(c)] We quantify the increase of the constraining power from the three surveys separately and from their combination, by adding progressively information from the cluster number counts, the mean cluster power spectrum analysis and the CMB prior from Planck experiment. We summarize in Fig. \\ref{fig:fom} the resulting improvements on FoM. The slightly different directions of degeneracy of the constraints in the $(w_0,w_a)$ parameter space from cluster number counts and power spectrum explain why the constraints substantially improve when we consider the two contribution together rather then separately. We verified that adding the CMB information improve the corresponding constraints on the DE EoS, mostly as a consequence of the constraint provided by CMB data on the geometry of the Universe (right panel of Fig. \\ref{fig:contr}).  \\item[(d)] We find that RSDs carry important cosmological information through the linear growth of perturbations, also in the case of cluster surveys. Indeed, the DE FoM from the power spectrum analysis of the Wide survey increases by a factor 35 when including RSDs, while increasing by a factor 7.7 and 6.8 \\numero for the Medium and the Deep surveys, respectively. \\item[(e)] As for the information carried by the shape of the power spectrum, a smaller increase in the FoM is instead measured when including BAOs. In this case the FoM from the power spectrum analysis of the Wide survey increases by a factor of 2, while no significant information on BAOs is provided by the Medium and the Deep surveys. Furthermore, relaxing the assumption of CDM and treating the shape of the power spectrum as a free parameter reduces the FoM by a factor of 1.7 \\numero in the analysis of the Wide survey. \\item[(f)] The results obtained for the EDE EoS analysis confirm that the Medium survey is the one carrying most of the information on the DE EoS. \\referee{This emphasizes once more the importance of finding a good balance in the definition of a survey strategy, between the redshift range needed to trace cosmic growth and the survey area.}  By extending the redshift range of the sample and with the ability to internally calibrate the observable-mass relation, we expect to measure the EDE EoS parameter (\\omegade) with an uncertainty of $\\sigma_{\\Omega_{e,de}} = 6.6 \\times 10^{-4}$ \\numero.  \\end{description}  In order to compare the constraints from the WFXT cluster surveys to those expected from other cosmological experiments, we also compare in Fig. \\ref{fig:fom} the FoM expected from WFXT to those presented by \\citet{albrecht06} for different large--scale structure probes. In the DETF report, they showed that Stage II cluster projects (ongoing surveys) provide FoM $\\sim 2$ \\numero (Fig. \\ref{fig:fom}) when combined with Planck priors. This analysis was carried out for a generic cluster count survey covering $200\\, \\rm{deg^2}$ up to $z_{max} = 2$, with the simple assumption of a constant mass selection function. According to Stage IV future experiments, by extending the survey area to $20000\\, \\rm{deg^2}$, the FoM rises in the optimistic configuration \\footnote{ As for the clusters analysis, according to the optimistic configuration, the mean of the mass-observable relation and its variance per redshift interval of $\\Delta z = 0.1$ is assumed in the DETF report to be determined up to a level of $1.6\\%$.}  up to $\\sim 40$ with the contribution of the CMB Planck priors. We point out that the DETF analysis did not include the constraints expected from the power spectrum analysis. A similar value of the FoM was also obtained from an optimistic version of Stage IV project for BAO analysis based on galaxy redshift surveys, again including Planck priors. The constraining power of these optimistic \\footnote{ As for the BAOs analysis, the DETF report introduces the $\\sigma_F$ parameter, which describes the scatter in the relation between the true and the photometric redshifts, $\\sigma_F^2=Var(z-z_{phot})/(1+z)^2$. In the the optimistic configuration $\\sigma_F= 0.01$.  } Stage IV experiments is somewhat weaker than that of the WFXT surveys, with the latter having a FoM larger by a factor of $\\magcir 2$. In Fig. \\ref{fig:fom}, we also show for reference the FoM expected for an optimistic \\footnote{ As for the WL analysis in the DETF report, the r.m.s. bias $\\sigma_{\\ln(1+z)}$ between the mean $z$ and photometric redshift for galaxies in $\\ln(1+z)$ for each bin of width $0.15$ is assumed to be determined with a precision of $0.001$.  Moreover, the shear measurement is assumed to be miscalibrated by a factor $(1+f_{cal})$ that varies independently for each redshift bin. It is assumed that the calibration factor of each redshift bin has a Gaussian prior of width $\\sigma(f_{cal})$. In the optimistic scenario this parameter was fixed to $\\sigma(f_{cal})=0.001$.}  Stage IV Weak Lensing experiment, which should reach a value of about 300.  We stress that the above forecasts from the WFXT surveys are obtained by considering a subsample of clusters with at least 1500 net photon counts. With this restriction robust mass and Fe-line based redshift measurements can be readily available from the same survey data, without resorting to external follow-up calibrations or observations. As such, the derived constraints should be considered as rather conservative since they do not include possible information carried by clusters detected with a smaller number of photons or any other information to constrain mass from external observations (e.g., Sunyaev-Zeldvich fluxes, weak lensing masses and optical richness from future surveys). Lowering the flux limits of the WFXT surveys by a factor of 30 would still guarantee detection of clusters as extended sources, without however allowing a measurement of redshifts and robust mass proxies. Fig. \\ref{fig:fom} shows that by including all the detectable clusters, the FoM increases by about one order of magnitude, even by assuming no prior on the mass parameters to compensate for the lack of robust mass measurements. \\referee{We note that the Wide survey provides the largest constraining power for the DE parameters when we include all clusters down to the detection limit. In fact, in this case the Wide survey dominates the statistics of clusters counts out to redshift 1.5 (see Figure 3 in \\citet{sartoris10}).}  Clearly, the results obtained from all the detected clusters must be considered as optimistic, since they rely on the possibility of confirming all these extended sources and measuring their redshifts with the aid of large follow-up observations.  In general, our analysis emphasizes that for large cluster surveys to be really useful for cosmological applications, not only large samples are needed but also a robust measurement of mass proxies is required for a significant fraction of the cluster sample. This will be possible with future X-ray surveys only with an adequate combination of survey area, sensitivity and angular resolution. Furthermore, our results also indicates that the optimization of the survey strategy depends on the class of cosmological models that one wants to constrain. For instance, in our previous analysis presented in \\cite{sartoris10}, we showed that the Wide survey is best suited to constrain deviations from non--Gaussian initial conditions, due to its ability to sample the long wavelength modes thus detecting a possible scale-dependence of the bias. This differs from the conclusion reached in the analysis presented here, where instead we conclude that the Medium survey is best suited to trace the growth history of perturbation over a large redshift baseline, as required to follow the redshift dependence of the DE EoS. As already mentioned in the introduction, models of modified gravity and of clustered DE represent an other broad class of models for which clusters can provide important constraints \\citep[e.g.,][]{rapetti09,Schmidt_etal09,lombriser10}. General predictions of these models are the scale-dependence of the growth factor of perturbations and of the bias function \\citep[e.g.,][]{parfrey11}. Large cluster surveys, such as those considered in this paper, have the potential of placing important constraints on such signatures of deviations from standard quintessence variants of $\\Lambda$CDM, especially once a suitable survey strategy is chosen.  The upcoming results from the Planck mission will much improved CMB priors based on experimental data. At the same time, the eROSITA satellite will provide in few years a wealth of fresh X--ray view of the large scale distribution of galaxy clusters and the evolution of their population out to $z\\sim\\! 1$. While waiting for a first X--ray satellite optimized for surveys, such as WFXT, there is no doubt that an extensive follow-up campaign will be required to provide a robust mass calibration with independent methods for a significant fraction of eROSITA clusters, so as to fully exploit the cosmological information contained in such surveys.  \\begin{figure} \\includegraphics[width=0.47\\textwidth]{figure/fom_total.eps} \\caption{The histograms represent the Figures of Merit for the $(w_0,w_a)$ parameters from WFXT surveys, as derived in the following configurations: by including in the Fisher Matrix the cluster number counts only (NC), the cluster mean power spectrum only (PS), the sum of the two (NC+PS), and by adding the prior from the Planck experiment (NC+PS+PL). All these FoM are obtained by assuming \\textit{strong prior} on mass parameters. The last group of histograms shows FoMs as obtained in the configuration NC+PS+PL by considering all clusters that can be detected ($\\approx 30$ source counts) in a given survey and by assuming \\textit{no prior} on mass parameters. The FoM for the Deep the Medium and the Wide cluster surveys are shown with the cyan, blue, green histograms respectively. The yellow histogram represent the FoM obtained from the combination of the three surveys. The horizontal lines show the FoM as reported in the DETF \\citep{albrecht06} for Stage II Cluster projects (CL-II; dot--dashed), for optimistic Stage IV BAO and Cluster projects (BAO IVS-o and CL IVS-o, respectively; solid line) and for optimistic Stage IV Weak Lensing project (WL IV-o; dotted line), by combining each probe with CMB Planck priors.} \\label{fig:fom} \\end{figure}"
    },
    "1112/1112.1321_arXiv.txt": {
        "abstract": "{Stellar magnetic activity of solar like stars is thought to be due to an internal dynamo. While the Sun has been the subject of intense research for refining dynamo models, observations of magnetic cyclic activity in solar type stars have become more and more available, opening a new path to understand the underlying physics behind stellar cycles. For instance, it is key to understand how stellar rotation rate influences magnetic cycle period $\\Pcyc$. Recent numerical simulations of advection-dominated Babcock Leighton models have demonstrated that it is difficult to explain this observed trend given a) the strong influence of the cycle period to the meridional circulation amplitude and b) the fact that 3D models indicate that meridional flows become weaker as the rotation rate increases. In this paper, we introduce the turbulent pumping mechanism as another advective process capable also of transporting the magnetic fields. We found that this model is now able to reproduce the observations under the assumption that this effect increases as $\\Omega^2$.  The turbulent pumping becomes indeed another major player able to circumvent the meridional circulation. However, for high rotation rates ($\\Omega \\simeq 5 \\Omega_\\odot$), its effects dominate those of the meridional circulation, entering a new class of regime dominated by the advection of turbulent pumping and thus leading to a cyclic activity qualitatively different from that of the sun.} ",
        "introduction": "Magnetic pumping refers to transport of magnetic fields in convective layers that does not result from bulk motion. One particular case is turbulent pumping. In inhomogeneous convection due to density stratification, convection cells take the form of broad hot upflows surrounded by a network of downflow lanes \\citep{1991MNRAS.253..479C,2008ApJ...673..557M}. In such radially asymmetric convection, numerical simulations show that the magnetic field is preferentially dragged downward \\citep{2001ApJ...549.1183T}. This effect has been demonstrated to operate in the bulk of the solar convection zone. A significant equatorward latitudinal component also arise when rotation becomes important, \\ie when the Rossby number is less than unity. Turbulent pumping speeds of a few$\\ms$ can be reached according to the numerical simulations of \\citet{2006A&amp;A...455..401K}. Therefore, its effects are expected to be comparable to those of meridional circulation.  In spite of those results, the effects of turbulent pumping rarely have been considered in mean field models. As the latter were able to reproduce rather well the large scale magnetic field using only the alpha effect and differential rotation to drive the dynamo, the pumping effect was thought to be an unnecessary complication.  A first approach showing the importance of pumping in the solar cycle was made by \\citet{1992ASPC...27..536B}. Since then, turbulent pumping has been a useful approach to tackle the problem of storage of magnetic fields. Indeed, for magnetic flux ropes to be buoyantly unstable and to emerge at the surface to form bipolar magnetic region with the appropriate tilt, numerical simulations have shown that their strength must be as high as  $10^4 - 10^5 \\rm ~G$ \\citep{2007AN....328.1104J,2004ASPC..325...47F,1987ApJ...316..788C}. One important limitation of this scenario is that $10^5 \\rm ~G$ represents a magnetic energy density higher by an order of magnitude than the kinetic energy density, \\ie one needs to create super equipartition magnetic structures. Therefore, a stable layer is required to store and amplify the magnetic fields. For this process to occur, differential rotation must be able to develop intense toroidal magnetic fields either within the tachocline or in the convection envelope. Pumping could be the transport needed to get the poloidal field down to the tachocline and could also maintain strongly buoyant structures from rising, thus helping them to become even stronger.   More interestingly in the regard of our work, turbulent pumping have shown remarkable properties regarding stellar cycles. \\citet{2008A&A...485..267G} (hereafter GdG2008) have demonstrated that the magnetic cycle period $\\Pcyc$ is no longer dominated by the meridional flow speed $v_0$ but instead by the radial turbulent pumping $\\gamma_r$ following the relation :  \\begin{equation} P_{\\rm cyc} \\propto v_0^{-0.12} \\gamma_r^{-0.51} \\gamma_\\theta^{-0.05} \\end{equation}  \\noindent available in the range $v_0=[500;3000]$, $\\gamma_{r0}=[20;120]$, $\\gamma_{\\theta 0}=[60;140]$ in $\\rm cm ~s^{-1}$. Another interesting feature of this model is that the surface magnetic field no longer shows the strong concentration in the polar region that usually characterizes Babcock Leighton dynamo solutions operating in the advection dominated regime. This can be traced primarily to the efficient downward turbulent pumping that subducts the poloidal field as it is carried poleward by the meridional flow.  With the STELEM code, we aim to study the effect of turbulent pumping on $\\Pcyc$ and to see under which conditions the $\\Pcyc-\\Prot$ relationship can be recovered. The paper is organised in the following manner. In Section \\ref{section_model}, we describe the equations, the initial and boundary conditions, the ingredients of the model and the standard case in which the turbulent pumping is \\emph{not} included. Section \\ref{section_pumping} shows how the dynamo model behaves in the presence of turbulent pumping and we conclude in Section \\ref{section_conclusion}. ",
        "conclusions": "The aim of this study was to address the current issue encountered by 2D mean field models in which the magnetic period decreases with the rotation rate, contrary to the observations \\citep{2009ASPC..416..375S,2011arXiv1109.4634W}. Earlier work of JBB2010 showed that this behaviour was due to the large influence of the MC. More specifically, they have used the results of recent 3D simulations \\citet{brown2008} in which meridional circulation amplitude decrease with rotation rate, and came to the conclusion that the observed $\\Pcyc - \\Prot$ relationship cannot be reproduced, unless a multicellular MC is considered.  The idea of this work was to tackle this problem under another angle introducing the turbulent pumping mechanism which has been shown to have interesting properties on the magnetic period (GdG2008). We have then performed 2D BL flux transport simulations with the STELEM code in the advection dominated regime. We first computed a \\emph{standard} model without turbulent pumping but producing solar characteristics. We found that in such models the magnetic period is indeed very sensitive to the MC amplitude ($\\Pcyc \\propto \\Omega_0^{0.05} s_0^{0.07} v_0^{-0.83}$), and is therefore not able to reproduce the observations, confirming the results of JBB2010.  In the presence of turbulent pumping however, the MC is no longer the only process capable of influencing $\\Pcyc$. We found that the latitudinal speed at the tachocline drives the magnetic cycle period. The weaker the MC, the stronger the influence of pumping on the magnetic cycle period. A reasonable equatorward flow at the BCZ ($\\sim 0.2 \\ms$) with a velocity contrast between the surface and the tachocline of $\\sim 109$ gives the scaling $ P_{\\rm cyc} \\propto v_0^{-0.40} \\gamma_{r0}^{-0.30} \\gamma_{\\theta 0}^{-0.15} $. The observational trend can be thus recovered only if $\\gamma \\propto \\Omega^2$ which has not been reported yet in previous Cartesian simulations \\citep{2001ApJ...549.1183T,2009A&amp;A...500..633K}. However 3D MHD simulations in full spherical shells are currently under development to test this assumption \\citep{matt2011}.  Another successful feature resulting from the presence of turbulent pumping is their ability to reduce the strong concentration at the poles. On one hand, the strong radial component drags the surface field down to the tachocline. On the other hand, the latitudinal component is equatorward everywhere (and therefore at the surface) expanding the strong poloidal field to lower latitudes. As we increase the rotation rates, the pumping becomes stronger and this dilution is enhanced until the surface field becomes completely homogeneous. However, at this level of rotation, the butterfly diagram is strongly affected by the pumping which has become the dominant advective process with a greater amplitude than the MC. For instance, it does not show strong equatorward nor poleward branches. We enter here in a different dynamo regime, in which we are dominated , not by MC, but by turbulent pumping, with characteristics different than the Sun.    Although this new ingredient looks promising with respect to our study, turbulent pumping amplitude must be as high as few $\\sim 10 \\ms$ for $\\sim 5 \\Omega_\\odot$. Such high values have not been reported yet in direct numerical simulations.  A way out would be to mix the two proposed approaches by considering in the same time both pumping and multicellular MC. The latter naturally arised in 3D MHD global simulations \\citep{2002ApJ...570..865B,2007ApJ...669.1190B,brown2008} in which they last for several years and are supported by observational evidences via local helioseismology techniques \\citep{2002AAS...200.0414H}.  Finally, long term modulations have not been considered in this work since no nonlinear effects beside quenching has been used. We intend to take into account those effects such as Malkus Proctor term in a future work since they are potentially important \\citep{2000MNRAS.315..521M}. We also intend to extend the resolution domain by including an atmosphere which would give a prediction for the external toroidal field, but also a bottom boundary condition for stellar wind models \\citep{2011ApJ...737...72P}.  Studying stellar magnetism with mean field models is instructive as we can extract information on the sensitivity of magnetic cycles to parameter change, such a study being very computationally expensive and delicate with 3D MHD simulations \\citep{2011ApJ...731...69B}. However, in the light of this work, a subtle combination of different processes must be considered to account for the many aspects of the observations. More accurate data would be of great help to understand the complex underlying physics of stellar dynamo."
    },
    "1112/1112.4777.txt": {
        "abstract": "Direct numerical simulations of isotropically forced homogeneous stationary turbulence with an imposed passive scalar concentration gradient are compared with an analytical closure model which provides evolution equations for the mean passive scalar flux and variance. Triple correlations of fluctuations appearing in these equations are described in terms of relaxation terms proportional to the quadratic correlations. Three methods are used to extract the relaxation timescales $\\tau_i$ from direct numerical simulations. Firstly, we insert the closure ansatz into our equations, assume stationarity, and solve for $\\tau_i$. Secondly, we use only the closure ansatz itself and obtain $\\tau_i$ from the ratio of quadratic and triple correlations. Thirdly we remove the imposed passive scalar gradient and fit an exponential decay law to the solution.  We vary the Reynolds ($\\Rey$) and P\\'eclet ($\\Pe$) numbers while keeping their ratio at unity and the degree of scale separation and find for large $\\Rey$ fair correspondence between the different methods. The ratio of the turbulent relaxation time of passive scalar flux to the turnover time of turbulent eddies is of the order of three, which is in remarkable agreement with earlier work. Finally we make an effort to extract the relaxation timescales relevant for the viscous and diffusive effects. We find two regimes which are valid for small and large $\\Rey$, respectively, but the dependence of the parameters on scale separation suggests that they are not universal. ",
        "introduction": "Fluid flows in astrophysical bodies are most often highly turbulent. Direct modeling of such high-Reynolds-number flows is currently impossible. Consequently, greatly enhanced diffusivities or modified diffusion operators are often applied in simulations \\cite{BN11}. Such models are challenging in terms of the required computational resources, so wide-ranging parameter studies cannot be performed.  An alternative approach is to separate the large and small scales, and derive equations for the former in which correlations of small-scale quantities are parameterized. This is usually referred to as mean-field theory, see e.g.\\ \\cite{M78,KR80,R89,RH04}. Various schemes have been introduced to close the equations for the correlations of small-scale quantities. In astrophysics, the second-order correlation approximation (SOCA), and the minimal $\\tau$ approximation (MTA, see e.g.\\ \\cite{BF02,BF03}) are widely used. The relevant relaxation time in MTA has been determined numerically for passive scalar transport \\cite{BKM04} as well as for the $\\alpha$ effect in mean-field electrodynamics \\cite{BS05,BS07}. Another approach, related to the MTA, where a relaxation term is invoked to describe the higher order correlations has been introduced in \\cite{O03}. In this `Ogilvie approach' several nondimensional coefficients are invoked to describe physically motivated parameterizations of the higher order correlations in terms of relaxation and isotropization terms. This model has been applied to different physical setups in order to calibrate the coefficients \\cite{GO05,MG07,LKKBL09,GOMS10}.  The validity of the various approximations can in principle be tested by comparing with direct simulations in the same parameter regime. In practise this is often not easy due to the limited parameter range accessible by the simulations. The starting point for such studies has been isotropically forced homogeneous turbulence under the influence of rotation \\cite{KB08} and/or shear \\cite{SKKL09}. In a recent work, the timescales related to diffusion and isotropization that appear in the Ogilvie approach have been studied \\cite{SBKM11}. In the present study we extend the work of \\cite{SBKM11} to the passive scalar case. ",
        "conclusions": "The main conclusion to be drawn from the present work is that the time scales used to model closure terms in the equations for the mean flux $\\overline{\\uu c}$ and the mean square concentration $\\overline{c^2}$ are nearly independent of $\\Rey$ for $\\Rey\\geq10$ and also nearly independent of the scale separation ratio for $\\kf/k_1\\ge3$. Expressed in terms of dynamical times, the resulting non-dimensional time scales can be referred to as Strouhal numbers whose values are around 3 for the $\\overline{\\uu c}$ closure term and around 7 for the $\\overline{c^2}$ closure term. The former value is in good agreement with earlier work using the $\\tau$ approximation \\cite{BKM04}.  Equipped with this knowledge, we may now be better justified in using the closure hypotheses discussed here for the quantities $\\overline{\\uu c}$ and $\\overline{c^2}$. On the other hand, as explained in the present paper, it is quite clear that these closure hypotheses lack thorough justification \\cite{RR07}. One should therefore in future strive to find systematic discrepancies from the anticipated scalings. One example that we alluded to in the present paper is the inhomogeneous case in which the $\\tau$ approach my break down. Future work in that direction seems now to be highly desirable, because in virtually all astrophysical applications the turbulence is inhomogeneous or at least anisotropic."
    },
    "1112/1112.3781_arXiv.txt": {
        "abstract": "\\noindent We have for the first time calculated the population characteristics of the Earth's irregular natural satellites (NES) that are temporarily captured from the near-Earth-object (NEO) population. The steady-state NES size-frequency and residence-time distributions were determined under the dynamical influence of all the massive bodies in the solar system (but mainly the Sun, Earth, and Moon) for NEOs of negligible mass. To this end, we compute the NES capture probability from the NEO population as a function of the latter's heliocentric orbital elements and combine those results with the current best estimates for the NEO size-frequency and orbital distribution. At any given time there should be at least one NES of 1-meter diameter orbiting the Earth. The average temporarily-captured orbiter (TCO; an object that makes at least one revolution around the Earth in a co-rotating coordinate system) completes $(2.88\\pm0.82)\\rev$ around the Earth during a capture event that lasts $(286\\pm18)\\days$. We find a small preference for capture events starting in either January or July. Our results are consistent with the single known natural TCO, 2006 RH$_{120}$, a few-meter diameter object that was captured for about a year starting in June 2006. We estimate that about 0.1\\% of all meteors impacting the Earth were TCOs. ",
        "introduction": "\\label{sec:introduction}  In this work we provide the first study and characterization of the population of temporarily-captured natural Earth satellites (NES; Table \\ref{table:acronyms} contains a list of all acronyms used in the paper) including their steady-state size-frequency distribution (SFD), capture probability as a function of the near-Earth-object (NEO) source population's orbital elements, and their geocentric orbit residence-time distributions.  The NES population provides a test of the NEO population statistics in a meteoroid size range that is not well-sampled by contemporary asteroid surveys---there are only three known NEOs with $H>32$ as of Oct 22, 2011---and could provide a remote laboratory for detailed long-term studies of the physical properties of the smallest asteroids. The long-term future concept of a spacecraft mission to retrieve an entire meteoroid from Earth orbit would provide an unprecedented scientific opportunity.  \\begin{table}[!h] \\centering \\begin{tabular}[c]{ll} \\hline NES & Natural Earth Satellite \\\\ SFD & Size-Frequency Distribution \\\\ NEO & Near-Earth Object \\\\ EMS & Earth-Moon System \\\\ TP  & Test Particle \\\\ TC  & Temporary Capture \\\\ TCO & Temporarily-Captured Orbiter \\\\ TCF & Temporarily-Captured Flyby \\\\ ISP & Intermediate Source Population \\\\ SP  & Source Population \\\\ \\hline \\end{tabular} \\caption{List of acronyms.}\\label{table:acronyms} \\end{table}  Despite a large body of work on satellite capture by the gas giants, mainly Jupiter and Saturn, there has been surprisingly little published about the Earth's natural satellites other than the Moon. The origin and evolution of the population of temporarily-captured irregular natural Earth satellites (NES) is entirely unknown.  To the best of our knowledge, the earliest paper mentioning NESs other than the Moon was \\citeasnoun{1913JRASC...7..145C}. He explained that the great meteor procession witnessed in North America on Feb 9, 1913 ``had been traveling through space, probably in an orbit about the sun, and that on coming near the earth they were promptly captured by it and caused to move about it as a satellite.'' A few years later \\citeasnoun{den1916a} concluded that ``the large meteors'' that passed over Northern America in 1913 must have been temporary Earth satellites because they traveled 2600 miles in the atmosphere suggesting that the orbits were ``concentric, or nearly concentric, with the Earth's surface.''  Since about 1957 a large body of work has been carried out in the field of dynamics of low-altitude satellites with the main application being artificial spacecraft. \\citeasnoun{bak1958a} considered the possibility that artificial satellites ``may be accompanied in their journey through space by certain \"natural\" satellites.'' He hypothesized that the heliocentric orbits of Earth-grazing meteors would become geocentric and elliptical as a result of atmospheric drag. In the next decade, \\citeasnoun{cas1965a} suggested that the line of craters and hexahedrite meteorites associated with the Campo del Cielo craters in Argentina originated in a high-altitude break-up of a temporary Earth satellite in a decaying orbit. They also hypothesize that the North Chilean hexahedrites may be fragments of the same body which made one more revolution before coming to ground.  One line of studies on the dynamics of NESs was motivated by the hypothesis that the Moon would be the origin of the terrestrial tektites but a nontrivial transport mechanism was needed to explain their uneven distribution on the Earth's surface \\citeaffixed{oke1961a}{see, e.g.,}. They are now thought to originate in terrestrial impact events.  \\citeasnoun{cli1979a} estimated the maximum speed of an object that can be captured by the Earth-Moon system (EMS) if the object has a very close encounter with the Moon. He concluded that in the four-body system (Sun-Earth-Moon-asteroid) an asteroid on a heliocentric orbit may be captured through a very close encounter with the Moon.  The first identification of an NES was only in the last decade.  While \\citeasnoun{bag1969a} provided direct and indirect evidence for electrically charged NESs, \\citeasnoun{mee1973a} later revealed elementary misunderstandings and contradictions in the earlier work that completely refute any evidence of NESs. \\citeasnoun{tan1997a} discussed the origin of 1991 VG and suggested that this recurrently temporarily-captured object could be a piece of lunar ejecta formed by a large impact. The estimated absolute magnitude has a wide range ($26.7 < H < 29.0$) even when just considering slope parameters typical for natural objects with high albedos. However, the possibility that 1991 VG could be artificial has not yet been ruled out although this is an unlikely scenario due to the large projected area (G. Tancredi, personal communication). Interestingly, 1991 VG is currently flagged as a Virtual Impactor by the Jet Propulsion Laboratory's impact-monitoring system, SENTRY. \\citeasnoun{kwi2009a} presented photometry of asteroid 2006 RH$_{120}$ ($H=29.9\\pm0.3$) which orbited the Earth within the Earth's Hill sphere from July 2006 until July 2007.  It has been concluded that it can not be a man-made object based on its low area-to-mass ratio (P.\\ Chodas, personal communication) and high circular-polarization ratio and low albedo from radar observations (L.\\ Benner, personal communication). The preliminary estimate for its diameter is $>2.3\\meter$ based on continuous-wave radar measurements.  We have not tested NEO orbits in the Minor Planet Center catalogue for temporary capture by the EMS, but have assumed that if such objects would exist they would have been reported in other publications. E.g., \\citeasnoun{kwi2009a} give an extensive and detailed list of NES candidates, and argue that 2006 RH$_{120}$ is the only object certainly known to be a NES. Testing for past or future captures is not trivial --- assuming that astrometric measurements are only available after the escape in the previous case --- because the orbital uncertainties and long extrapolation intervals combined with the fractal nature of the orbit distribution for capturable objects \\cite{mur1989a} might lead either to missing some captures or to false positive captures. The utility of adding possible NES candidates would thus be questionable for the purpose of estimating the validity of our results on the NES SFD. On the other hand, if there are observations of an object both before and after a capture it is likely that there also are some observations (or observers have at least attempted to get them) during the capture. It is likely that these events or their mere possibility would have been published in the literature.  The Earth's quasi-satellites have some common characteristics with the NESs. (For a description of their dynamics see, e.g., \\citeasnoun{mik2006a} and references therein.) The essential difference between satellites and quasi-satellites is that the NES orbit depends critically on the gravity of the EMS while the orbit of a quasi-satellite would hardly change if the EMS suddenly ceased to exist because it is orbiting the Sun on an Earth-like orbit in the vicinity of the Earth. Quasi-satellites thus form a subgroup of objects with co-orbital motion with respect to planets (objects on tadpole and horseshoe orbits form the other two subgroups). Based on the NEO model by \\citeasnoun{bot2002a}, the steady-state population of Earth's co-orbitals was estimated by \\citeasnoun{mora2002a} to be around 13--19 objects with $H<22$ and the length of the co-orbital motion episodes range from 25~kyr to 1~Myr in their integrations. The origin and evolution of Earth's known co-orbitals has recently been studied by \\citeasnoun{bra2008a} who conclude that these objects should exist on Earth-like orbits for around 10kyr (note the discrepancy with \\citename{mora2002a}) --- three orders of magnitude shorter than the average lifespan of an NEO.  The general approach we have chosen to determine the NES SFD is to first calculate the capture probability as a function of the source population's heliocentric orbital elements and then combine the probabilities with the best available NEO population models \\cite{bot2000a,bot2002a,bro2002a}. The questions we will answer are: \\begin{itemize}  \\item What is the steady-state population of temporary-captured NESs?  \\item What are the pre- and post-capture orbit distribution for NESs?  \\item How long do capture events last?  \\item How many orbits does an NES typically complete before escaping the EMS?  \\item What fraction of NESs impact the Earth and Moon while captured?  \\item What are the orbit characteristics of an NES?  \\item Can an NES become a temporary Moon satellite?  \\item Are capture events equally likely to happen throughout the year?  \\end{itemize} ",
        "conclusions": "\\label{sec:conclusions}  We provide the first estimate of the orbit and size-distribution for temporarily captured natural irregular satellites of the Earth. We predict that there is a one-meter-diameter or larger NEO temporarily orbiting the Earth at any given time. The NEO orbit and SFD model is currently the main factor limiting the accuracy of our predictions. Given the orbit distribution of \\citeasnoun{bot2002a} the \\citeasnoun{bro2002a} NEO SFD is consistent with the only known TCO of natural origin, 2006 RH$_{120}$, while the \\citeasnoun{bot2002a} and \\citeasnoun{rab2000a} SFDs are each off by about an order of magnitude.  Our integrated TCO population will allow us to examine different scenarios for their detection and to estimate, e.g., the average time from discovery to an accurately known orbit. It seems plausible that the long-lived TCOs could have stable enough orbits to allow successful searches to be carried out in specific regions of the sky. Once TCOs can be reliably and frequently identified early enough in a capture event they create an opportunity for a low-cost low-delta-v meteoroid return mission \\citeaffixed{elv2011a}{cf.}. The scientific potential of being able to first remotely characterize a meteoroid and then visit and bring it back to Earth would be unprecedented.       \\newpage"
    },
    "1112/1112.2391_arXiv.txt": {
        "abstract": "We propose a cosmological inflationary scenario based on the supergravity-embedded Standard Model supplemented by the right-handed neutrinos. We show that with an appropriate % K\\\"{a}hler potential the $L$-$H_u$ direction gives rise to successful inflation that is similar to the recently proposed gravitationally coupled Higgs inflation model but is free from the unitarity problem. The mass scale $M_R$ of the right-handed neutrinos is subject to the seesaw relation and the present 2-$\\sigma$ constraint from the WMAP7-BAO-$H_0$ data sets its lower bound $M_R\\gtrsim$ 1 TeV. Generation of the baryon asymmetry is naturally implemented in this model. We expect that within a few years new observational data from the Planck satellite will clearly discriminate this model from other existing inflationary models arising from the same Lagrangian, and possibly yield stringent constraints on $M_R$. ",
        "introduction": "Today observational cosmology is a precision science. Cosmological inflation, which is supported by all observational data, is now an indispensable theoretical ingredient not only in astrophysics but also in particle phenomenology. A remaining mystery of this otherwise extremely successful paradigm is embedding into a particle theory model. By virtue of Occam's razor, a plausible % possibility may be that the fields responsible for cosmological inflation (inflatons) are those already included in the Standard Model (SM), or its (not too large) extension. The recently proposed SM Higgs inflation model \\cite{HiggsInflation} is an interesting idea to test this possibility. This model is attractive due to its minimalistic nature and the remarkable agreement with the present day observational data. It also relates the dynamics of inflation with the electroweak scale physics, making a prediction on the SM Higgs mass from the cosmological microwave background (CMB) data. A rather unfavourable feature of this type of model is that it requires extremely large nonminimal coupling to gravity, which could lead to violation of the unitarity bound \\cite{unitarity}. The model also suffers from the hierarchy problem, which may be cured by supersymmetrisation \\cite{Einhorn:2009bh,FKLMP,Arai:2011nq}. % See \\cite{Pallis:2011ps} for related models.   Certainly, there are more traditional ways of embedding inflation into supersymmetric SMs. It has been known for a while that the flat directions in the minimal supersymmetric Standard Model (MSSM), lifted by soft supersymmetry breaking terms and other effects, can serve as inflatons (reviewed in \\cite{Enqvist:2003gh}; more recent developments include \\cite{ATI}). Another type of embedding is into a supersymmetric SM with right-handed neutrinos \\cite{MSYY}, in which one of the right-handed sneutrinos is identified as the inflaton. These models are phenomenologically well motivated; the hierarchy problem is solved by supersymmetry, and the models with the right-handed neutrinos are furthermore consistent with the small but nonzero neutrino masses indicated by neutrino oscillation.   In this paper we present a new scenario of inflation, inspired by these developments. Our model has the following features: (i) the scenario is based on the simplest supersymmetric extension of the SM that includes the right-handed neutrinos, naturally explaining the small neutrino masses through the seesaw mechanism \\cite{seesaw}; (ii) the problem associated with the large nonminimal coupling that afflicts the SM Higgs inflation is alleviated; (iii) the CMB data gives predictions on the mass scale of the right-handed neutrinos through the seesaw relation; (iv) leptogenesis is naturally implemented; (v) the predicted cosmological parameters fit well in the present day observational constraint, and (vi) the model can be tested by the upcoming observational data from the Planck satellite. We discuss construction of the model and describe these features below. ",
        "conclusions": ""
    },
    "1112/1112.3732_arXiv.txt": {
        "abstract": "{High-resolution echelle spectra taken with the 6-m telescope in 2003--2011 are used to study features of the optical spectrum and the velocity field in the atmosphere of the semiregular variable LN~Hya in detail. The weak symmetric photospheric absorptions indicate radial velocity variations from night to night (by as much as 3\\,km/s), resulting from small pulsations. Peculiarities and profile variations were found for strong lines of Fe\\,I, Fe\\,II, Ba\\,II, Si\\,II, etc. The profiles of these lines were asymmetric: their short-wave wings were extended and their cores were either split or distorted by emission. During the 2010 observing season, the position and depth of the H$\\alpha$ absorption component, the intensities of the shortand long-wave emission components, and the intensity ratio of the latter components varied from spectrum to spectrum. Weak emissions of neutral atoms (VI, MnI, CoI, NiI, FeI) appeared in the spectrum of June 1, 2010. All these spectral peculiarities, recorded for the first time, suggest that we have detected rapid changes in the physical conditions in the upper atmospheric layers of LN\\,Hya in 2010.}  \\authorrunning{\\it Klochkova \\& Panchuk} \\titlerunning{\\it Spectrum of LN Hya in 2010} ",
        "introduction": "The semiregular variable LN\\,Hya (HD\\,112374, IRAS\\,12538$-$2611) is located at a high Galactic latitude (b\\,=\\,36$\\lefteqn{.}$\\degr4), indicating that the star belongs to an old Galactic population. Its brightness variation amplitude is $\\Delta$B\\,=\\,0.32$^{\\rm m}$ ; the variability period is not firmly established, and possible durations range from 44 to 86 days [1]. Semiregular variables are far evolved post-AGB stars. Stars of intermediate masses (with initial masses of $3\\div 8{\\mathcal M}_{\\odot}$) are observed in this short-lived stage; they lose matter via their stellar winds as they evolve of the AGB. The secular variability of the principal parameters of post--AGB stars has stimulated spectroscopic monitoring. Such spectroscopic monitoring of selected AGB and post-AGB stars has been performed with the 6-m telescope of the Special Astrophysical Observatory over the past decade. The main aim of our monitoring is to detect peculiarities and possible variations of the spectrum, and to study the time behavior of the velocity field in the stellar atmospheres and shells. Observations on the 6-m telescope have detected variations in the spectra of V510\\,Pup [2, 3], V2324\\,Cyg [4], and the optical counterpart of the IR--ray source IRAS\\,01005\\,+\\,7910 [5]. New results concerning variations of the optical spectra and atmospheric velocity fields have also been obtained for the variable stars QY\\,Sge [6], V354\\,Lac [7, 8], V448\\,Lac [9], CY\\,CMi [10].  Stellar evolution from the AGB stage to the planetary-nebula stage remains incompletely understood. This concerns the mechanisms and details of the stellar mass loss and the complex morphology of their circumstellar gaseous-dusty shells. The role of pulsations and binarity has also been poorly studied, and the detection of these effects for stars with complex atmospheric velocity fields is laborious and time consuming. These considerations also apply to studies of the variable supergiant LN\\,Hya. An image taken with the Hubble Space Telescope [11] shows LN\\,Hya as a point-like object. The star is surrounded with a gas and dust shell, whose presence in the optical is revealed by the emission profile of the H$\\alpha$ line. After publication of the IRAS results and optical identification of IRAS objects via catalog cross correlation, LN\\,Hya was identified with the infrared source IRAS\\,12538$-$2611. However, the fluxes in the IRAS bands are low [12] compared to the star's radiation in the visible. Note that low IR--fluxes have also been observed for several peculiar UU\\,Her-type supergiants at high Galactic latitudes. LN\\,Hya is usually considered to be a member of this rare Galactic population (e.g. [13, 14]). No infrared excess was detected for the high-latitude supergiant UU\\,Her, whose properties will also be discussed below.  The brightest star in its class, LN\\,Hya has been studied spectroscopically many times. Already in the early 1980s, Luck et al. [15] determined the main parameters (effective temperature Teff\\,=\\,6000\\,K, surface gravity log\\,g\\,=\\,0.4 and 0.8) and detailed chemical composition of LN\\,Hya, which displays a metal deficiency, [Fe/H]\\,=$-1.2$, and large nitrogen excess, [N/Fe]\\,=\\,+0.5. Somewhat later, Klochkova and Panchuk [16] used photographic spectra taken with the main stellar spectrograph of the 6-m telescope to study the spectrum of LN\\,Hya among a sample of supergiants at high Galactic latitudes, and confirmed the main conclusion of Luck et al. [15]: LN\\,Hya is a post-AGB star with a deficiency of heavy elements. Ten years later, Giridhar~et~al.~[17] obtained new estimates of the fundamental parameters and chemical composition of LN\\,Hya based on CCD spectra, reaching similar conclusions. The main focus of the recent study [18] was abundances of elements with low condensation temperatures, which are not subject to selective separation processes (C, N, O, S, Zn).  Thus, several papers devoted to the chemical composition of LN\\,Hya have been published over the last 30 years, but very little information is available about the peculiarities of its spectrum or possible time variability. We have obtained new high-quality optical spectra of LN\\,Hya over the last decade. This paper presents our results on variability and previously unknown peculiarities of the star's spectrum and its detailed radial-velocity pattern.  \\begin{table}[t] \\caption{Log of observations and the heliocentric radial velocities Vr of LN\\,Hya} \\medskip \\begin{tabular}{ c| c | c | c | r c} \\hline Date    &    JD 2450000+ & $\\Delta\\lambda$\\,,\\AA{} &\\multicolumn{3}{c}{Vr, km/s}           \\\\ \\cline{4-6} &      &  &Vr(metals)& Vr(H$\\alpha$$_{\\rm abs}$)& Vr(NaI$_{\\rm emis}$) \\\\ \\hline 21.02.2003 &  2692.52 &     5150--6660 &$-$23.20$\\pm$0.12 (125)&$-$25.5 &$-$21.9  \\\\ 02.04.2010 &  5289.39 &     5160--6680 &$-$23.68$\\pm$0.08 (304)&$-$20.5 &$-$21.1  \\\\ 05.04.2010 &  5292.33 &     5160--6680 &$-$21.93$\\pm$0.07 (312)&$-$20.7 &$-$22.0  \\\\ 01.06.2010 &  5349.27 &     5220--6690 &$-$28.22$\\pm$0.11 (211)&$-$7.8  &$-$21.6  \\\\ 12.01.2011 &  5574.59 &     5210--6680 &$-$23.30$\\pm$0.06 (235)&$-$22.3 &$-$21.8  \\\\ 14.03.2011 &  5635.42 &     5220--6690 &$-$24.88$\\pm$0.07 (153)&$-$23.6 &$-$21.6  \\\\ \\hline \\end{tabular} \\label{date} \\end{table}    Vr(metals) is the mean velocity from symmetric absorption features; the number of lines used to derived the mean velocity is given in brackets. Vr(H$\\alpha$$_{\\rm abs}$) is the velocity from the H$\\alpha$ absorption component, and Vr(Na$_{\\rm emis}$) the velocity from the emission components of the D--lines NaI. ",
        "conclusions": "Our echelle spectra of the low-mass, high-latitude supergiant LN\\,Hya acquired with the 6-m telescope over several years over a wide spectral range and with high spectroscopic resolution have revealed new spectroscopic phenomena for this star, and enabled us to measure radial velocities for spectral features formed at various depths in the stellar atmosphere.  Variations in the radial velocities measured from weak photospheric absorption lines indicate the presence of low-amplitude pulsations in deep (nearly photospheric) layers of the atmosphere.  We have found complex, time-variable profile shapes for strong absorption lines (Fe\\,I, Fe\\,II, Si\\,II, Ba\\,II, etc.) formed in upper layers of the star's extended atmosphere. We have detected variations of the asymmetries of these absorption lines from spectrum to spectrum, including variations in their splitting.  A previously unknown property of LN\\,Hya appeared in the spectrum of 1~June, 2010: weak symmetric emissions identified with lines of neutral atoms (V\\,I, Mn\\,I, Co\\,I, Ni\\,I, Fe\\,I). The time--stationary mean Vr from these emission lines, formed in the circumstellar medium (Vsys\\,=$-21.6$\\,km/s), can be adopted as the systemic velocity.  During the 2010 observing season, the position and depth of the H$\\alpha$ absorption component, intensities of the short-wave and long-wave emission lines, and the intensity ratio of these lines varied from spectrum to spectrum. These properties of the stellar spectrum, detected for the first time in these observations, suggest that we detected a rapid change of the physical conditions in the upper atmospheric layers of LN\\,Hya in 2010, possibly due to the passage of a shock."
    },
    "1112/1112.3959_arXiv.txt": {
        "abstract": "The halos of galaxies preserve unique records of their formation histories. We carry out the first combined observational and theoretical study of phase-space halo substructure in an early-type galaxy: M87, the central galaxy in the Virgo cluster. We analyze an unprecedented wide-field, high-precision photometric and spectroscopic data set for 488 globular clusters (GCs), which includes new, large-radius Subaru/Suprime-Cam and Keck/DEIMOS observations. We find signatures of two substructures in position-velocity phase-space. One is a small, cold stream associated with a known stellar filament in the outer halo; the other is a large shell-like pattern in the inner halo that implies a massive, hitherto unrecognized accretion event. We perform extensive statistical tests and independent metallicity analyses to verify the presence and characterize the properties of these features, and to provide more general methodologies for future extragalactic studies of phase-space substructure. The cold outer stream is consistent with a dwarf galaxy accretion event, while for the inner shell there is tension between a low progenitor mass implied by the cold velocity dispersion, and a high mass from the large number of GCs, which might be resolved by a $\\sim0.5\\,L^*$ E/S0 progenitor. We also carry out proof-of-principle numerical simulations of the accretion of smaller galaxies in an M87-like gravitational potential. These produce analogous features to the observed substructures, which should have observable lifetimes of $\\sim$~1~Gyr. The shell and stream GCs together support a scenario where the extended stellar envelope of M87 has been built up by a steady rain of material that continues until the present day. This phase-space method demonstrates unique potential for detailed tests of galaxy formation beyond the Local Group. ",
        "introduction": "\\label{sec:intro}  The ultimate record of the formational histories of galaxies is found in the six-dimensional phase space of stellar positions and velocities, supplemented by additional information about ages and elemental abundances (primarily iron-based metallicities). The halos of galaxies are natural hunting grounds for assembly clues in phase-space because of the preservative effects of the long dynamical times, and since any material entering a galaxy naturally has to traverse its halo.  Indeed, it is a fundamental prediction of the current cosmological paradigm that the outer regions of galaxies should be constantly bombarded with smaller, infalling systems that persist as significant substructures while gradually becoming disrupted. This idea has been borne out qualitatively in recent years by observations of halo streams and shells in the nearby Universe (e.g., \\citealt{1999Natur.402...53H,2001Natur.412...49I,2006ApJ...642L.137B,2009AJ....138.1417T,2010AJ....140..962M,2011arXiv1111.2864C}).  Massive elliptical galaxies are touchstones for this issue, since their spheroidal nature and their residence in the densest areas of the Universe suggest particularly vigorous accretion histories. Comparisons of these galaxies' stellar luminosity profiles at high and low redshifts now provide indirect evidence that their outer envelopes grow tremendously in size up to the present day (e.g., \\citealt{2008ApJ...687L..61B,2008ApJ...688...48V,2009ApJ...695..101D,2010ApJ...709.1018V,2011arXiv1106.4308C}), which may be naturally explained by accretion and mergers in a cosmological context \\citep{2006ApJ...648L..21K, 2009ApJ...699L.178N,2010MNRAS.401.1099H,2010ApJ...725.2312O,2011arXiv1106.5490O}.  The most extreme examples of ongoing assembly are expected to be the brightest cluster galaxies (BCGs), whose typical locations at the centers of clusters should involve very active merger histories (e.g., \\citealt{2009ApJ...696.1094R}).\\footnote{Although the central galaxy is often not technically the brightest object in a cluster, we will follow the convention of calling this the BCG.} Current cosmologically-based models find that around half of the stellar mass growth of a present-day BCG occurred over the past 5~Gyr via gas-poor mergers (after an early formation epoch for the core regions; \\citealt{2007MNRAS.375....2D}).  This theoretical picture has received very mixed support from comparisons of BCG stellar masses and surface brightness distributions at different redshifts \\citep{2008MNRAS.387.1253W,2009MNRAS.395.1491B,2009Natur.458..603C,2010ApJ...721L..19V,2011ApJ...726...69A,2011MNRAS.414..445S}. Surveys for major mergers in action have found that these contribute significantly to recent BCG growth (\\citealt{2008MNRAS.388.1537M,2009MNRAS.396.2003L}; see also \\citealt{2011MNRAS.414L..80B}), but the contributions from minor mergers are unknown and the overall comparison to theory unclear. Part of the difficulty is that even if accretion and mergers are important growth processes, the frequency of observing an ongoing event can be fairly low. Thus it would be invaluable to go beyond purely photometric observations and delve into the phase-space of position, velocity, and metallicity for longer lived signatures (e.g., \\citealt{2008ApJ...689..936J,2009MNRAS.394..641Z}).  Within the Local Group, detailed phase-space studies are being undertaken using individual stars (e.g., \\citealt{2008ApJ...689..958K,2009ApJ...705.1275G,2009ApJ...698..567S,2011ApJ...738...79X}), but are currently impossible in more distant galaxies. Instead, bright stellar proxies may be used, such as planetary nebulae (PNe; \\citealt{2003ApJ...582..170D,2003MNRAS.346L..62M,2010ApJ...725L..97S,2011MNRAS.414..642C}) and globular clusters (GCs). GCs are thought to form along with field stars, and as collections of $\\sim$~$10^6$ stars each, they are visible at much greater distances, allowing their individual positions, line-of-sight velocities, and chemical properties to be measured as far away as $\\sim$~50~Mpc \\citep{2011A&A...531A.119R}.  The ``chemodynamics'' of GCs provided the watershed evidence for the accretion origin of the Milky Way's outer halo \\citep{1978ApJ...225..357S}, where it is now thought that disrupting satellite galaxies have deposited many accompanying GCs (e.g., \\citealt{2003AJ....125..188B,2007ChJAA...7..111G,2007PASP..119..939G,2007ApJ...661L..49L,2010MNRAS.404.1203F}). A similar scenario is now evident in the nearby spiral M31, where some GCs are associated with halo substructures \\citep{2009MNRAS.396.1619C,2010ApJ...717L..11M}. In BCGs, the presence of GCs in enormous numbers (up to $5\\times10^4$ per system) has long motivated suggestions that accretion is an important factor for these galaxies \\citep{1982AJ.....87.1465F,1998ApJ...501..554C}, although it is not clear how many of these GCs may have formed during in situ processes or mergers rather than being accreted.  Various GC and PN-based studies of the halo kinematics of giant ellipticals and BCGs have turned up evidence for substructures that may imply active accretion events or mergers \\citep{2003ApJ...591..850C,2009AJ....137.4956R,2010A&A...513A..52S,2010A&A...518A..44M,2011AJ....141...27W}. In other cases, accretion events are suggested by photometric substructures or by broad kinematical or chemical halo transitions \\citep{2009AJ....138.1417T,2009MNRAS.398...91P,2009MNRAS.394.1249C,2010MNRAS.407L..26C,2010ApJ...715..972J,2011MNRAS.413.2943F,2011ApJ...736L..26A,2011MNRAS.415..993M}. However, in no case has a halo feature been observed and modeled in enough detail to determine its origin and the implications for galaxy formation.  To go beyond the previous work on nearby BCGs by carrying out detailed phase-space studies of their halos, there are two critical observational requirements beside sheer instrumental throughput: wide-field coverage and spectroscopic resolution. The former permits blind searches for halo substructures and demands a field of view of at least $\\sim$~0.5~degree in order to probe galactocentric radii out to $\\sim$~100~kpc or more in galaxies at $\\sim$~15~Mpc distances. The latter allows for a wide mass range of accreted galaxies to be probed; e.g., if a BCG has a characteristic velocity dispersion of $\\sigma \\sim$~300~\\kms\\ and the precision of velocity measurements for halo tracers is $\\Delta v \\sim$~50--100~\\kms, as in many past surveys, then the observations are sensitive to mass ratios down to $\\sim (\\Delta v/\\sigma)^2 \\sim$~1:40--1:10 in dynamical mass. More minor mergers than these are certainly expected to be more frequent, and may comprise a dominant mode for halo assembly, so a velocity resolution of tens of \\kms\\ ($\\Delta v/\\sigma \\lsim 0.1$) is preferable.  We present here the first wide-field, large-sample spectroscopic survey of a BCG with velocity resolution of $\\Delta v/\\sigma \\sim 0.04$. Our subject is M87, the central elliptical in Virgo, the nearest galaxy cluster at a distance of $\\sim$~16.5~Mpc. The GC system of M87 has seen decades of spectroscopic study (e.g., \\citealt{1987AJ.....93..779H,1997ApJ...486..230C,2001ApJ...559..812H}), but extending to radii of only $\\sim$~40~kpc and with $\\Delta v/\\sigma \\gsim 0.25$.  Until now, M87 could be characterized as a dynamically quiet galaxy, with only mild signs of interactions such as very faint stellar filaments in its far outer halo \\citep{2005ApJ...631L..41M,2010ApJ...715..972J}.  Our new GC study reveals the kinematics of one of these faint streams, and unveils an unsuspected, enormous shell-like substructure in phase-space.  Our paper proceeds as follows.  Section~\\ref{sec:obs} describes the observations and data reduction, and presents a basic overview of the halo substructures. Sections~\\ref{sec:shell} and \\ref{sec:outer} analyze the characteristics and statistical significance of the large inner shell and outer stream, respectively. Theoretical analyses of the substructure origins and dynamics are explored in Section~\\ref{sec:sim}. Section~\\ref{sec:concl} summarizes the findings and outstanding questions. ",
        "conclusions": "\\label{sec:concl}  We have used an unprecedentedly wide-field data set of high-precision globular cluster velocities to search for phase-space substructure around the central cluster galaxy M87. The large majority of the data are from S+11, with a few additional velocities obtained from the literature, and from a new sample of Keck/DEIMOS spectroscopy presented here that targets a known filament of stellar light in the outer halo. Analysis of ($g-i$) colors and spectroscopic metallicities for a subset of our GC sample confirms that the colors provide a good proxy for metallicity, allowing for chemical tagging of substructures.  We identify two candidate features: one a small stream near to the outer-halo stellar filament, and the other a previously unknown colossal ``shell'' in the inner halo. We have carried out extensive statistical tests and characterizations of these substructures.   The outer GC stream appears to be a collection of $15\\pm10$ objects (after extrapolating to fainter magnitudes) residing in a $\\sim$~5~$\\times$~20~kpc clump with an internal velocity dispersion of $\\sim 30\\pm20$~\\kms. We have found the presence of this stream to be significant both in a group-finding analysis $(p=0.09)$, from a maximum likelihood model $(p=0.025)$, and from its distinctive color distribution compared to the surrounding GCs $(p=0.22)$.  The stream dispersion, its GC colors and numbers counts, the color of the stellar filament, and the filament's narrowness all suggest the accretion of a dwarf galaxy. Low-mass kinematical substructures in elliptical galaxies have been inferred in other studies \\citep{2003ApJ...591..850C,2011AJ....141...27W}, but in this case the association with a visible counterpart should allow for a much clearer determination of the accretion dynamics.  The shell of GCs has a chevron-like shape in phase-space that resembles classic expectations for a disrupted infalling system \\citep{1988ApJ...331..682H}, although the interpretation in detail is more challenging. The appearance of the shell is sharpened by a change of coordinates from simple radius to elliptical, circular-equivalent radius. The significance of the shell in phase-space is established using a group-finding algorithm ($p<0.01$), a crude entropy metric ($p=0.31$), and a maximum-likelihood chevron-based model ($p<0.01$ compared to a pure Gaussian model). The color distribution of the shell and non-shell GCs from the same region is different $(p=0.13)$, providing independent evidence for the existence of the shell. The phase-space morphology resembles a simple chevron model more than a wedge or tapered Gaussian, but may very well have a complex, intermediate morphology; however, such uncertainties do not compromise our basic finding of a relatively massive and cold substructure.  We find from maximum-likelihood fitting that the shell contains 500--1100 GCs, which is $\\sim$~5 times the entire Milky Way system, and comprises around 20\\% of M87's GCs in the 50--95~kpc region. This number suggests a stellar debris field of (2--6)$\\times10^{10} L_\\odot$, brought in by either a large group of dwarf galaxies, or by a single giant elliptical or lenticular galaxy.  The dramatic appearance of the shell in phase-space is due to its relatively cold kinematics, with an estimated velocity dispersion of 10--30~\\kms\\ (or up to $\\sim$~120~\\kms\\ if we make generous allowances for uncertainties concerning the shell morphology). The lower velocity dispersion estimates are difficult to reconcile with the number of inferred GCs, as they most naturally imply a dwarf galaxy hosting only a handful of GCs. However, after allowing for systematic uncertainties in our analyses, we suggest a solution involving an E/S0 progenitor with a luminosity of $\\sim 0.5\\,L^*$, which could very well also be the parent galaxy for one of the tidally-stripped galaxies now visible close in to M87.  There is a critical need for additional data from the shell region, including new velocities of both GCs and PNe, in order to confirm the presence of the substructure and to refine our estimates of its characteristics. Along these lines, we have recently obtained another cycle of GC spectroscopy from MMT/Hectospec that includes $\\sim$~35 velocity measurements of independent objects in the shell region. A preliminary analysis yields no obvious sign of substructure, but the statistics from the new data are poor, and at least another $\\sim$~100 velocities are needed.  We have carried out simplified $N$-body simulations of satellite galaxy accretion in a static cluster$+$BCG potential, finding that these readily produce streams and chevrons in phase space. In one simulation, a dwarf galaxy has recently fallen in and is partially disrupted, forming a classic, narrow tidal stream in both positional and phase space. In a second case, the initial infall orbit had much lower angular momentum, so that the satellite galaxy passed close to the center of M87 and was quickly disrupted. Its remnants are strewn more haphazardly across positional space and could be difficult to discern against the main body of M87, while a distinct chevron persists in phase space, with the apex corresponding to a shell turning point.  These simulations compare favorably to the M87 observations because the low internal velocity dispersions of the dwarf progenitors result in cold kinematics for the tidal debris.  Simulations of a more massive progenitor produce hotter kinematics that do not as strongly resemble the data for the ``shell'', reiterating the tension between the different constraints on the shell progenitor mass, which remains unresolved pending further observations and modeling. It would be worth also investigating alternative explanations for the substructure, such as a shell-like boundary between well-mixed GC subpopulations formed or accreted in different epochs.  We have explored the use of the velocity-radius slopes of the substructures in phase-space estimate the underlying gravitational potential and dark matter distribution around M87. However, the interpretation of the data appears to be complicated by non-spherical effects.  Taken at face value, the shell GCs imply a substantial, recent accretion event in the halo of M87. This is the first large substructure identified around a central cluster galaxy, the first large stellar substructure with a clear kinematical detection in {\\it any} type of galaxy beyond the Local Group, and by far the largest shell or stream discovered in the Local Universe.  M87 was a favored target for our survey as well as previous ones because of its proximity, rich GC system, and {\\it general lack of obvious dynamical disturbance}. The large shell we have discovered is an example of the lurking accretion signatures that could be found from detailed phase-space studies of apparently placid galaxies.  There are however more indirect lines of evidence for a recent minor merger in M87, such as an off-center supermassive black hole \\citep{2010ApJ...717L...6B}, a lopsided central GC velocity distribution (S+11), an accretion powered jet, and several surrounding peculiar low-luminosity satellites that could be the nuclei of since-shredded disk galaxies. Further theoretical work is needed to see if these features could all be explained by a unified scenario involving a recent gas-rich minor merger.  The shell and stream in M87 have probable lifetimes of $\\sim$~0.5~Gyr and $\\sim$~1.5~Gyr, respectively.  Given our best guesses for the progenitors of these features, the total luminosity of M87 would be built up by $\\sim$~10 or $\\sim$~1000 events corresponding to the shell or the stream, respectively. If these recent events are representative of the long-term evolution of M87, and by extension of the class of brightest cluster galaxies in general, then they support a picture of late-epoch growth by hierarchical assembly \\citep{2007MNRAS.375....2D}.  More generally, the M87 results are another demonstration of the potency of GCs for providing unique information about the stellar halos of galaxies. We note the concluding prediction of \\citet{2009ApJ...699L.178N} for their landmark two-phase galaxy formation scenario: ``...the outer parts of massive giant ellipticals will tend to be old, blue, metal-poor, and relatively uniform from galaxy to galaxy since they are all composed essentially of the debris from tidally destroyed accreted small systems.'' Globular clusters have in fact long provided evidence for exactly this scenario.  \\vskip 5pt"
    },
    "1112/1112.3990_arXiv.txt": {
        "abstract": "We present results of an analysis of the local ($z \\sim 0$) morphology-environment relation for 911 bright ($M_B < - 19$) galaxies, based on matching classical RC3 morphologies with the SDSS-based group catalog of Yang et al., which includes halo mass estimates. This allows us to study how the relative fractions of spirals, lenticulars, and ellipticals depend on halo mass over a range of $10^{11.7}$--$10^{14.8} \\, h^{-1} \\Msol$, from isolated single-galaxy halos to massive groups and low-mass clusters. We pay particular attention to how morphology relates to central vs.\\ satellite status (where ``central'' galaxies are the most massive within their halo). The fraction of galaxies which are elliptical is a strong function of stellar mass; it is also a strong function of halo mass, but \\textit{only} for central galaxies. We interpret this as evidence for a scenario where elliptical galaxies are always formed, probably via mergers, as central galaxies within their halos, with satellite ellipticals being previously central galaxies accreted onto a larger halo. The overall fraction of galaxies which are S0 increases strongly with halo mass, from $\\sim 10$\\% to $\\sim 70$\\%.  Here, too, we find striking differences between the central and satellite populations. $20 \\pm 2$\\% of central galaxies with stellar masses $\\Mstellar > 10^{10.5}\\Msol$ are S0 regardless of halo mass, but satellite S0 galaxies are only found in massive ($> 10^{13} \\, h^{-1} \\Msol$) halos, where they are $69 \\pm 4$\\% of the $\\Mstellar > 10^{10.5}\\Msol$ satellite population.  This suggests \\textit{two} channels for forming S0 galaxies: one which operates for central galaxies, and another which transforms lower mass ($\\Mstellar \\lesssim 10^{11}\\Msol$) accreted spirals into satellite S0 galaxies in massive halos.  Analysis of finer morphological structure (bars and rings in disk galaxies) shows some trends with stellar mass, but none with halo mass; this is consistent with other recent studies which indicate that bars are not strongly influenced by galaxy environment.  Radio sources in high-mass central galaxies are common, similarly so for elliptical and S0 galaxies, with a frequency that increases with the halo mass. Emission-line AGN (mostly LINERs) are more common in S0s, but show no strong trends with environment. ",
        "introduction": "\\label{sec:intro}  Galaxies come in many different shapes and sizes, but primarily consist of two dynamically stable components, bulges and disks, with additional contributions from quasi-stable features such as bars and rings. \\citet{Hubble26} devised what has become known as the ``tuning fork'' classification to describe this morphological schema: elliptical galaxies ($\\sim$ pure bulge); spiral galaxies (disks containing spiral features, both barred and unbarred, with a sequence Sa--Sc\\footnote{\\citet{RC1} extended this sequence further to Sd and Sm} of decreasing bulge component and increasing spiral arm opening angle); and lenticular ``S0'' galaxies (defined by the presence of a disk with no discernable spiral arms).  The abundance of these galaxy types is now known to correlate strongly with environment: elliptical galaxies live preferentially in regions of very high local density \\citep{Dressler80}, inhabiting the cores of clusters and groups rather than their outskirts \\citep{Melnick77,Whitmore91,Wilman09}. However, the total luminosity-limited fraction of ellipticals is similar in a wide range of environments \\citep{Desai07,Wilman09,Just10} and has evolved very weakly since $z \\sim 1$ \\citep{Dressler97,Fasano00,Smith05,Postman05}.  Lenticular galaxies, on the other hand, are globally no more abundant in clusters than they are in groups \\citep{Wilman09}, with a much weaker dependence on local density than ellipticals \\citep{Dressler80,Dressler97,Postman05,Poggianti08}. A lower fraction at fixed luminosity is found only in the lower density field \\citep{Wilman09}.  The global fraction of lenticular galaxies has grown by a factor $\\gtrsim 2.5$ in groups and clusters since $z \\sim 0.5$ \\citep{Dressler97,Fasano00,Wilman09}, potentially more rapidly in groups and low mass clusters than in high-mass clusters \\citep{Poggianti09,Just10}.  Since lenticulars and ellipticals make up the bulk of the low redshift passive population \\citep[e.g.][]{Bundy10}, this means that the majority of galaxies which have ceased forming stars since $z \\sim 0.5$ have retained their disks as lenticulars.  Beyond $z \\sim 0.5$ (up to $z \\sim 1$), there is no evidence for further evolution in the elliptical or S0 fraction in the high-mass clusters sampled to date \\citep{Smith05,Postman05}.  To build a global picture of galaxy evolution, it is important to understand how the evolution of galaxies and their environment is intertwined, and which physical processes drive the suppression of star formation and the morphological transformations which are required to explain these observations.  It is attractive to explain morphological evolution in the context of bulge growth through galaxy mergers \\citep{Springel05,Hopkins10}, which take place in the center of halos as a natural consequence of hierarchical clustering and dynamical friction \\citep[e.g.][]{deLucia07,deLucia11}.  Elliptical galaxies almost certainly form through a mixture of major and minor mergers.  However, the role of mergers in the formation of bulges in galaxies with disks is less clear.  Cooling flows should form at the center of halos, and will reform a disk around the merger remnant so long as the gas is not too efficiently reheated. SPH simulations suggest that during mergers of galaxies with high gas fractions, much of the gas retains its angular momentum and reforms a disk, and that some component of the stellar disk can survive in all but 1:1 mergers \\citep{Hopkins09}. Meanwhile, secular processes might contribute $\\lesssim 10\\%$ of the galaxy mass in a ``pseudobulge'' \\citep{Kormendy04}, or potentially a larger fraction at high redshift where disks are less dynamically stable \\citep{Genzel08}.  That S0 galaxies exhibit more significant bulges on average than spiral galaxies \\citep[e.g.][]{Dressler80,Wilman09,Laurikainen10} is consistent with a merger origin.  Nonetheless, there is currently no more than circumstantial evidence that bulge growth and the suppression of star formation in most S0s are causally linked.  Our observational picture still misses one vital and surprisingly straightforward ingredient: a full picture of how galaxy morphology depends upon environment at low redshift.  Early studies of environment focussed almost exclusively on galaxy clusters, for which good statistics and strong trends are to be found even without high spectroscopic completeness.  The more recent emphasis has been on extending these trends to high redshift with high resolution Hubble Space Telescope (HST) imaging.  Modern local surveys such as the Sloan Digital Sky Survey (SDSS) are so large that much effort has been invested in the automated classification of morphological properties, such as the simple SDSS ``concentration'' index, and multi-component decomposition with different levels of sophistication \\citep[resolution dependent, e.g.][]{Allen06,Gadotti09}. An interesting alternative solution is the ``Galaxy Zoo'', providing visual classifications for the whole SDSS, by employing an enthusiastic public to classify galaxies \\citep{Lintott11}.  The disadvantage of all these approaches is that new classification schemes tend to be relatively simplistic, with unknown systematics, and produce data which is difficult to assess on the basis of our existing understanding of the Hubble scheme.  Standard Hubble type classifications within SDSS are available only for subsamples \\citep[e.g.][]{Fukugita07}, although these are becoming larger \\citep{Nair10}.  In this paper we have taken the simple step of taking a sample of galaxies from the Third Reference Catalog of Bright Galaxies \\citep[RC3,][]{RC3} which is matched to the SDSS Data Release 4 (DR4) to provide classical morphological information for a large sample of galaxies for which the selection function and environment can be properly characterized.  For our purposes this provides an ideal sample to study environmental trends for galaxies with detailed and well understood classifications.  In section~\\ref{sec:sample} we describe the sample, discussing the selection function of RC3 galaxies and the group catalog which we use to describe galaxy environment.  Section~\\ref{sec:results} presents our results, including Hubble type fraction as a function of stellar and halo mass, a dichotomy between central and satellite galaxies, the nature of activity in the SDSS fiber spectra of galaxies with different morphological types and environments, and the dependence (or lack thereof) of morphological components (bars, rings etc) on environment. This is further discussed in the context of a hierarchically evolving Universe in section~\\ref{sec:disc}, and in section~\\ref{sec:concl} we present our conclusions.  To compute distances, absolute magnitudes, and stellar masses, we assume a $\\Lambda$CDM cosmology with $\\Omega_M$, $\\Lambda$, and $H0$ equal to $0.3$, $0.7$, and 71 km~s$^{-1}$ Mpc$^{-1}$, respectively. Halo masses are presented in $h^{-1} \\Msol$, to retain the native units of \\citet{Yang07}.\\footnote{$h=0.71$ in our adopted cosmology.} ",
        "conclusions": "\\label{sec:concl}  We have created a catalog of 1064 nearby (median $z \\sim 0.02$), bright ($B < 16$ and $\\MB < -19$) galaxies, which combines RC3 morphological classifications and the NYU-VAGC version of the SDSS DR4 photometric and spectroscopic data; magnitudes have been corrected to account for the undersubtracted background, and stellar masses were determined using $g - i$ color-based mass-to-light ratios calibrated by \\citet{Zibetti09} and applied to the $i$-band absolute magnitudes. The main morphological classifications (elliptical vs.\\ S0 vs.\\ spiral) were checked by visual examination of SDSS images; a total of 165 galaxies ended up being re-classified, and 55 more were classified for the first time.  To this dataset we added halo masses and \\textit{central} vs.\\ \\textit{satellite} status (where ``central'' = most massive galaxy in its group) from the group catalog of \\citet[][Y07]{Yang07}, which resulted in a total of 911 galaxies with halo mass assignments; 729 of these are central galaxies and 182 are satellites.  The main advantage of the Y07 catalog is that it spans the full range of halo masses down to $\\Mhalo = 10^{11.63} \\, h^{-1} \\Msol$. This allows us to describe the dependence of morphological fractions on halo mass at $z \\sim 0$ for the first time, from smaller clusters down to single-galaxy halos. (We also determined total number of galaxies per group and group velocity dispersions for the Y07 groups in our sample, in order to check the robustness of the halo-mass-based results.) Using the full SDSS DR4 catalog as a parent sample, we characterized the selection function of our catalog as a function of $B$ magnitude (synthesized using SDSS colors) and galaxy size ($r_{90}$, the radius containing 90\\% of the Petrosian flux in r-band). Our sample is robust to a luminosity limit of $\\MB = -19.0$, which corresponds to a stellar mass of $\\Mstellar \\sim 10^{10.5}\\Msol$ for E and S0 galaxies. Galaxies are weighted to correct for the selection bias, and are also weighted by $V/V_{\\rm max}$ to correct for Malmquist bias. This allows us to examine the fraction of galaxies of various types as a function of stellar and halo mass. We use a weighted logistic regression method to allow us to assess the statistical significance of apparent trends. This method has the advantages of modeling a binomial property without binning of data, and it allows for individual data points to have weights.  We find that the global fraction of elliptical galaxies increases with galaxy luminosity and with stellar mass, but not with halo mass. The fraction of S0s declines to high stellar mass, but increases with halo mass, group velocity dispersion and number of neighbours. These results are consistent with previous work at higher redshift which found the S0 fraction to increase in groups relative to the field population, whilst the elliptical fraction remains roughly constant \\citep{Wilman09}.  The fraction of \\textit{central} galaxies with elliptical morphology increases with stellar and halo mass, consistent with their formation in mergers. In contrast, the fraction of satellite ellipticals is globally low at all halo masses, but increases with stellar mass. We interpret this as evidence that satellite ellipticals were formed as the central galaxies of progenitor halos, which were subsequently accreted onto their present halo.  Limited to $\\Mstellar > 10^{10.5}\\Msol$, a modest fraction of central galaxies are S0s ($20.1 \\pm 1.6$\\%, with little dependence on stellar or halo mass). The remaining S0s are satellites of massive halos only -- we find that the fraction of satellites with S0 morphology rises from $0^{+2.6}_{-0}$\\% in halos with $\\Mhalo < 10^{13} \\, h^{-1} \\Msol$ to $69.0^{+4.3}_{-4.6}$\\% above this threshold.  Presuming S0s to be spirals in which star formation has been suppressed (leading in turn to the suppression of spiral arms), we interpret our result as a strong indication that there are \\textit{two} populations of S0s, in which star formation has been suppressed in different ways.  Central S0s may be formed via suppression of star formation during minor mergers and/or by feedback from AGN, with similarities to the elliptical population. Satellite S0s which became S0s while they were still central galaxies within their progenitor halos constitute a \\textit{pre-processed} population which can account for up to $\\sim 20\\%$ of the satellite S0s, including all of those with $\\Mstellar \\geq 10^{11}\\Msol$.  However, the higher fraction of S0s in higher-mass halos implies that many satellite S0s were accreted as spiral galaxies. These accreted spirals were then {\\it post-processed}, becoming satellite S0s, and are the dominant source of satellite S0s in the range $10^{10.5}\\Msol \\leq \\Mstellar \\leq 10^{11}\\Msol$. Altogether, we estimate that $64 \\pm 11$\\% of our satellite S0s were accreted as spirals.  Central S0 and elliptical galaxies frequently host radio sources, consistent with radio-mode heating of the surrounding hot gas. This heating may offset cooling onto these galaxies, and thus suppress star formation. However, S0s host ionized gas components -- mostly LINERs -- much more frequently than elliptical galaxies of the same mass.  We find no strong dependence of structural subcomponents -- bars, inner rings/spirals, outer rings -- on environment, in contrast to some earlier studies \\citep[e.g.][]{Elmegreen92}, though we do find evidence that the frequency of both bars and inner rings increases, and the frequency of outer rings decreases, with galaxy mass.  By comparing our results with the study of \\citet{Wilman09}, we find tentative evidence that the fraction of bright S0s in intermediate-mass groups (those with velocity dispersions $175 \\kms \\lesssim \\sigma \\lesssim 500 \\kms$) has increased in the last $\\sim 4$ Gyr, rising from $28\\pm4$\\% at $z \\sim 0.4$ to $52.3^{+3.9}_{-4.0}$\\%. This is at least qualitatively consistent with increases in the fraction of S0s in clusters reported by other studies."
    },
    "1112/1112.4040_arXiv.txt": {
        "abstract": "{Gamma-ray bursts (GRBs) are attractive objects for constraining the nature of dark energy in a way complementary to other cosmological probes, especially at high redshifts. However, the apparent magnitude of distant GRBs can be distorted by the gravitational lensing from the density fluctuations along the line of sight.} {We investigate the gravitational lensing effect on the cosmological parameters and dark energy equation of state from GRBs.} {We first calibrated the GRB luminosity relations without assuming any cosmological models. The luminosity distances of low-redshift GRBs were calibrated with the cosmography method using a latest type Ia supernova (SNe Ia) sample. The luminosity distances of high-redshift GRBs were derived by assuming that the luminosity relations do not evolve with redshift. Then we investigated the non-Gaussian nature of the magnification probability distribution functions and the magnification bias of the gravitational lensing.} {We find that the gravitational lensing has non-negligible effects on the determination of cosmological parameters and dark energy. The gravitational lensing shifts the best-fit constraints on cosmological parameters and dark energy. Because high-redshift GRBs are more likely to be reduced, the most probable value of the observed matter density $\\Omega_M$ is slightly lower than its actual value. In the $\\Lambda$CDM model, we find that the matter density parameter $\\Omega_M$ will shift from $0.30$ to $0.26$ after including the gravitational lensing effect. The gravitational lensing also affects the dark energy equation of state by shifting it to a more negative value. We constrain the dark energy equation of state out to redshift $z\\sim 8$ using GRBs for the first time, and find that the equation of state deviates from $\\Lambda$CDM at the $1\\sigma$ confidence level, but agrees with $w=-1$ at the $2\\sigma$ confidence level.} {} ",
        "introduction": "\\label{sec:introduction} Our Universe is currently in accelerating expansion. This is revealed by observations of high-redshift type Ia supernovae (SNe Ia) (Riess et al. 1998; Perlmutter et al. 1999), cosmic microwave background (CMB) fluctuations (Spergel et al. 2003, 2007), and the large-scale structure (LSS) (Tegmark et al. 2006). These observations suggest that the composition of the universe may consist of an extra component such as dark energy, or that the equations governing gravity may need a variation to explain the acceleration of the universe at the present epoch.  In order to measure the expansion history of our Universe, we need the Hubble diagram of standard candles. Type Ia supernovae that have played an important role in constraining cosmological parameters are the well-known standard candles. Unfortunately, it is difficult to observe SNe Ia at $z>1.7$, even with excellent space-based projects such as SNAP (Aldering et al. 2004). They cannot provide any information on the cosmic expansion beyond redshift 1.7. On the other hand, gamma-ray bursts (GRBs) are ideal candidates for this investigation. The high luminosities of GRBs make them detectable out to the edge of the visible universe (Bromm \\& Loeb 2002, 2006). The farthest GRB is GRB 090429B at $z=9.4$ (Cucchiara et al. 2011). Schaefer (2007) compiled 69 GRBs to simultaneously use five luminosity indicators, which are relations of $\\tau_{\\rm lag}-L$ (Norris, Marani \\& Bonnell 2000), $V-L$ (Fenimore \\& Ramirez-Ruiz 2000), $E_{\\rm peak}-L$ (Schaefer et al. 2003), $E_{\\rm peak}-E_{\\gamma}$(Ghirlanda, Ghisellini \\& Lazzati 2004), and $\\tau_{\\rm RT}-L$ (Schaefer 2007). Here the time lag ($\\tau_{\\rm lag}$) is the time shift between the hard and soft light curves, $L$ is the peak luminosity of a GRB, the variability $V$ of a burst denotes whether its light curve is spiky or smooth and it can be obtained by calculating the normalized variance of an observed light curve around a smoothed version of that light curve (Fenimore \\& Ramirez- Ruiz 2000), $E_{\\rm peak}$ is the photon energy at which the $\\nu F_{\\nu}$ spectrum peaks, $E_{\\gamma}=(1-\\cos\\theta_j)E_{\\rm \\gamma, iso}$ is the collimation-corrected energy of a GRB, and the minimum rise time ($\\tau_{\\rm RT}$) in the gamma-ray light curve is the shortest time over which the light curve rises by half of the peak flux of the pulse. Qi \\& Lu (2010) found a new correlation in the X-ray band of GRB afterglows. After calibrating with luminosity relations, GRBs may be used as standard candles to provide information on the cosmic expansion at high redshift and, at the same time, to tighten the constraints on cosmic expansion at low redshift (Dai et al. 2004; Ghirlanda et al. 2004; Friedman \\& Bloom 2005; Liang \\& Zhang 2005, 2006; Wang \\& Dai 2006; Schaefer 2007a; Wright 2007; Wang, Dai \\& Zhu 2007; Wang 2008; Qi, Wang \\& Lu 2008a,b; Liang et al. 2008; Amati et al. 2008; Cardone et al. 2009; Liang et al. 2009; Qi, Lu \\& Wang 2009; Izzo et al. 2009). Gamma-ray bursts can also potentially probe the cosmographic parameters to distinguish between dark energy and modified gravity models (Wang, Dai  \\& Qi 2009a, b; Vitagliano et al. 2010; Capozziello \\& Izzo 2008; Xia, et al. 2011). Samushia \\& Ratra (2010) also derived constraints on the $\\Lambda$CDM and XCDM models using a smaller set of 69 GRBs analyzed in two different ways, following Schaefer (2007) and Wang (2008). They found that GRB data disfavor standard $\\Lambda$CDM at a 1$\\sigma$ confidence level, and is consistent with this model at a 2$\\sigma$ confidence level. They also found that the two techniques give somewhat different cosmological constraints. This means that model-independent calibration method and more GRB data are needed to improve this situation. We will use the cosmographic parameters to calibrate the GRB luminosity correlations. More recently, Wang, Qi \\& Dai (2011) enlarged the GRB sample and put constraints on cosmological parameters and equation of state of dark energy. In this paper, we use this GRB sample which includes 116 GRBs.    However, the gravitational lensing by random fluctuations in the intervening matter distribution induces a dispersion in GRB brightness (Oguri \\& Takahashi 2006; Schaefer 2007), degrading their value as standard candles as well as SNe Ia (Holz 1998). Gamma-ray bursts can be magnified (or reduced) by the gravitational lensing produced by the structure of the Universe. The gravitational lensing has sometimes a great impact on high-redshift GRBs. First, the probability distribution functions (PDFs) of gravitational lensing magnification have much higher dispersions and are markedly farther from the Gaussian distribution more remarkably (Valageas 2000a; Wang et al. 2002; Ogri \\& Takahashi 2006). Second, there is effectively a threshold for the detection in the burst apparent brightness. With gravitational lensing, bursts just below this threshold might be magnified in brightness and detected, whereas bursts just beyond this threshold might be reduced in brightness and excluded. Schaefer (2007) considered the gravitational lensing biases and the Malmquist biases of 69 GRBs. He found that the gravitational lensing and Malmquist biases are much smaller than the intrinsic error bars. Our method differs in two ways from the one used in Schaefer (2007). First, we use the latest GRB sample that includes 116 GRBs. Second, we calculate the distance dispersions from the universal probability distribution function (UPDF) of the gravitational lensing amplification (Wang et al. 2002; Wang 2005). But Schaefer (2007) used the method of Gonzalez \\& Faber (1997), which relies on the poorly known luminosity function and number densities of GRBs.  We explore the gravitational lensing effects on constraints of cosmological parameters and equation of state from GRBs. We focus on the non-Gaussian nature of magnification probability distribution functions and the magnification bias of the gravitational lensing. The structure of this paper is arranged as follows: in section 2 we calibrate the luminosity relations of GRBs in a cosmological model-independent way. The constraints on the cosmological parameters and dark energy including gravitational lensing are presented in section 3. In section 4 we present model-independent constraints on the dark energy equation of state to $z\\sim8$ including the weak gravitational effect. In section 5 we summarize our findings and give a brief discussion. ",
        "conclusions": "We have presented the gravitational lensing effects on constraints of cosmological parameters and dark energy from GRBs. We mainly focussed on the non-Gaussian nature of magnification probability distribution functions and the magnification bias of gravitational lensing. We first used an SNe Ia sample to calibrate the luminosity relations of GRBs. Because the luminosity distances of SNe Ia are completely cosmological-model-independent, the GRB luminosity relations can be calibrated in a cosmology-model-independent way. Then we calculated the PDFs of gravitational lensing. The probability distribution functions at high redshifts have higher variance and a lower height of the maximum. The peaks reduce to a smaller magnification factor $\\tau$. From the probability distribution functions we can see that the gravitational lensing is more important for high-redshift objects. Finally we presented constraints on cosmological parameters and dark energy. We found that the gravitational lensing had non-negligible effects on the determination of cosmological parameters and dark energy. The gravitational lensing shifts the best-fit constraints of cosmological parameters and dark energy. Because high-redshift GRBs are more likely to be reduced, the most probable value of the observed matter density $\\Omega_M$ is slightly lower than its actual value. The gravitational lensing also biases a more negative value of the dark energy equation of state. We also constrained the dark energy equation of state out to redshift $z\\sim 8$ in a model-independent way using GRBs for the first time, and found that the equation of state deviates from $\\Lambda$CDM at the $1\\sigma$ confidence level, but agrees with $w=-1$ at a $2\\sigma$ confidence level.   As shown in Samushia \\& Ratra (2010), the cosmological constraints from the two methods of Schaefer (2007) and Wang (2008) may be different when using 69 GRBs. Therefore we emphasize that we need detailed studies of new correlations with a much greater number of GRBs and an examination of systematic errors to be able to regard GRBs as more accurate standardizable candles. Now ongoing missions like \\emph{Swift}, \\emph{Fermi} and \\emph{Suzaku}, and the collaboration of many observers on ground will promise the progression of GRB cosmology."
    },
    "1112/1112.1981_arXiv.txt": {
        "abstract": "Galileon gravity offers a robust gravitational theory for explaining cosmic acceleration, having a rich phenomenology of testable behaviors. We explore three classes of Galileon models -- standard uncoupled, and linearly or derivatively coupled to matter -- investigating the expansion history with particular attention to early time and late time attractors, as well as the linear perturbations.  From the relativistic and nonrelativistic Poisson equations we calculate the generalizations of the gravitational strength (Newton's constant), deriving its early and late time behavior.  By scanning through the parameters we derive distributions of the gravitational strength at various epochs and trace the paths of gravity in its evolution.  Using ghost-free and stability criteria we restrict the allowed parameter space, finding in particular that the linear and derivative coupled models are severely constrained by classical instabilities in the early universe. ",
        "introduction": "The origin of the current cosmic acceleration is a fundamental mystery, even to whether it arises from a new field or a change in the laws of gravity.  Scalar-tensor gravity is an intriguing possibility, partaking of elements of both classes of explanation.  Those theories that do not have a scalar potential avoid some naturalness problems associated with this, and those which involve a symmetry or geometric origin moreover sidestep difficulties with high energy physics corrections.  A class of theories possessing all these desiderata is Galileon gravity \\cite{arXiv:0811.2197, arXiv:0901.1314,arXiv:0906.1967,arXiv:1007.5278}, involving a shift symmetry in the field, while the field itself can be viewed as a geometric object arising from higher dimensions and entering in invariant combinations that assure second order field equations and protect against various pathologies.  The derivative self couplings of the field also screen deviations from general relativity in high gradient regions (e.g.\\ small scales or high densities), thus satisfying solar system and early universe constraints \\cite{75796,hep-th/0106001,arXiv:1111.5090}. The galileon field, and closely related models \\cite{arXiv:1106.2000,arXiv:1005.0868,arXiv:1006.0281,arXiv:1009.6151,arXiv:1008.0048,Qiu:2011cy,arXiv:1109.1047,arXiv:1103.5360,arXiv:1108.4242}, have been extensively studied in both theoretical \\cite{arXiv:1008.1305,arXiv:1008.4128,arXiv:1008.4580,arXiv:1007.5424,arXiv:1008.3312,arXiv:1008.0745,arXiv:1003.5917,arXiv:1007.0443,arXiv:1104.0155} and observational \\cite{arXiv:0909.4538,arXiv:1009.0677,arXiv:1003.3281,arXiv:1008.4236,arXiv:1011.6132,arXiv:1010.0407,arXiv:1007.2700,arXiv:1008.1588} contexts (and also with respect to inflation; \\cite{Kobayashi:2010cm,arXiv:1009.2497,Kobayashi:2011nu}).    Galileon gravity thus merits investigation in detail, from both the expansion properties (homogeneous background) and structure growth (inhomogeneous perturbations) perspectives.  Regarding the expansion, we will particularly be interested in the early universe corrections to radiation and matter domination, the effective equation of state evolution, and late time attractor behavior, especially to a de Sitter state.  For the behavior of linearized perturbations (concentrating on subhorizon scales), we explore the form of the modified gravitational strengths (Newton constants) in the generalized Poisson equations, looking for characteristic features in their evolution.  The parameter space of the theory will be constrained by general physical considerations, such as  positivity of the effective energy density and freedom from ghosts (Hamiltonian bounded from below).  Observationally, we ask that the radiation and matter eras remain intact and the effective dark energy density is of order 3/4 the critical density today.  These conditions still permit considerable diversity in the gravitational behavior, which is quite interesting to explore.  In Section~\\ref{sec:model} we briefly review the standard Galileon model and include terms with linear and derivative coupling to matter, setting up the field equations and modified Friedmann equations for a homogeneous, isotropic universe.  Section~\\ref{sec:pert} sets up the linear perturbation theory and the generalized Poisson equations, deriving expressions for the gravitational strength modifications.  Solutions for the expansion history for the uncoupled, linear-, and derivative-coupled models are presented in Section~\\ref{sec:expan}, including discussion of early and late time attractors.  Evolution of the gravitational strengths is studied in Section~\\ref{sec:geff}, along with the no-ghost and stability conditions. We summarize the results and conclude in Section~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  Galileon gravity is a well defined theory with a rich phenomenology. We explored the predictions for expansion history giving rise to current acceleration, despite the absence of any cosmological constant or potential, and found early time scaling and tracing solutions and late time de Sitter states.  Equally important is the behavior of the inhomogeneities in the Galileon field, and we discussed requirements for a sound theory in terms of the no-ghost condition, Laplace stability, and positivity of energy density. By creating a large ensemble of expansion histories, we were able to use these theoretical conditions to efficiently rule out large regions of the multi dimensional parameter space. In this work we focussed solely on the dynamics of scalar perturbations; if we include tensor modes then we expect even more stringent bounds on the Galileon parameters (see for example \\cite{Kimura:2011qn}). Confronting the model with data to further constrain the parameters will be the subject of future work.  The modifications to gravity induced by the Galileon terms can be characterized by gravitational strengths (effective Newton's constants) entering the Poisson equations for the various combinations of metric potentials, $\\geff^{(\\phi)}$, $\\geff^{(\\psi)}$, and $\\geff^{(\\phi+\\psi)}$. We derived expressions for these and solutions in the early time and late time limits.  Considering special cases and general scans over the parameter space, we presented the ``paths of gravity'', the evolution of the effective gravitational strengths over cosmic history.  Galileon models do not have simple, monotonic evolution of gravity.  One special feature discovered is that in the de Sitter limit the ``matter gravity'' $\\geff^{(\\psi)}$ entering into growth of structure and the ``light gravity'' $\\geff^{(\\phi+\\psi)}$ entering into light deflection are identical (unlike in many other theories), although generally different from general relativity.  Gravity can suddenly become stronger at some epoch in $1\\lesssim z\\lesssim 1000$, and then be restored to general relativity; it is intriguing to consider the possible effects on structure formation such as enabling early formation of massive clusters. The observed phenomenology of $\\geff^{(\\psi)}$ at intermediate redshifts will significantly increase the (linear) power of the matter perturbations at late times, and this will be used to severely constrain the uncoupled Galileon in an upcoming study.   Models which include a coupling between the Galileon field and the energy-momentum tensor, or equivalently the field and the Ricci scalar in the case of linear coupling or the field kinetic term and the Einstein tensor in the case of derivative coupling, are severely constrained by conditions for soundness of the theory.  In particular, we find that each coupling term must be subdominant to uncoupled Galileon terms in the early universe, and generally small.  Deviations from the uncoupled case therefore only tend to show up at late times.  For example, the linearly coupled case does not have a true de Sitter fixed point.  Galileon gravity offers a rich variety in both the expansion behavior (e.g.\\ equations of state that can track the matter component and give rise to early dark energy) and gravity behavior (spikes of deviations from, and restoration to, general relativity).  It provides a theoretically viable alternative to Einstein gravity, motivated by higher dimension geometric theories and protected by symmetries, without many of the naturalness issues of the cosmological constant or scalar field potentials. Galileon cosmology's implications for observations at high and low redshift make it an excellent theory to explore, and to test as an origin for cosmic acceleration and an extension to Einstein gravity."
    },
    "1112/1112.6273_arXiv.txt": {
        "abstract": "We have studied the star-formation and AGN activity of massive galaxies in the redshift range $z=0.4-2$, which are detected in a deep survey field using the AKARI InfraRed (IR) astronomical satellite and {\\em Subaru} telescope toward the North Ecliptic Pole (NEP). The AKARI/IRC Mid-InfraRed (MIR) multiband photometry is used to trace their star-forming activities with the Polycyclic-Aromatic Hydrocarbon (PAH) emissions, which is also used to distinguish star-forming populations from AGN dominated ones and to estimate the Star Formation Rate (SFR) derived from their total emitting IR (TIR) luminosities. In combination with analyses of their stellar components, we have studied the MIR SED features of star-forming and AGN-harboring galaxies, which we summarize below: 1) The rest-frame 7.7-$\\mu$m and 5-$\\mu$m luminosities are good tracers of star-forming and AGN activities from their PAH and dusty tori emissions, respectively. 2) For dusty star-forming galaxies without AGN, their SFR shows a correlation that is nearly proportional to their stellar mass and their specific SFR (sSFR) per unit stellar mass increases with redshift.  Extinctions estimated from their TIR luminosities are larger than those from their optical SED fittings, which may be caused by geometric variations of dust in them. 3) Even for dusty star-forming galaxies with AGN, SFRs can be derived from their TIR luminosities with subtraction of the obscured AGN contribution, which indicates that their SFRs were possibly quenched around $z \\simeq 0.8$ compared with those without AGN. 4) The AGN activity from their rest-frame 5-$\\mu$m luminosity suggests that their Super Massive Black Holes (SMBH) could already have grown to $\\simeq 3 \\times 10^8 M_{\\odot}$ in most massive galaxies with $10^{11} M_{\\odot}$ at $z>1.2$, and the mass relation between SMBHs and their host galaxies has already become established by $z\\simeq 1-2$. ",
        "introduction": "\\label{sec:intro} A major issue in observational cosmology is to reconstruct the cosmic evolution of star formation and the detailed growth history of the present-day galaxies. Even though it is not easy to reconstruct the whole path of the stellar mass for galaxies as a function of the present-day galaxy stellar mass $M_{\\ast;0}$ and the cosmic time $t$ as $M_{\\ast}(M_{\\ast;0}, t)$, it is essential to obtain knowledge of their Star Formation Rate (SFR) as a function $SFR(M_{\\ast},t)$.  Recent near-infrared (NIR) observations suggest that most of the stellar mass is contained in massive galaxies, and that this situation had already become established at $z \\simeq 1$. This appears to have been followed by a rapid evolution of the stellar mass density between $z = 1-3$ \\citep{papovich_stellar_2001, dickinson_evolution_2003, fontana_assembly_2003, franx_significant_2003, drory_comparing_2004, fontana_k20_2004, glazebrook_high_2004, van_dokkum_stellar_2004, labb_irac_2005, fontana_galaxy_2006, papovich_spitzer_2006, yan_stellar_2006, webb_star_2006, kriek_direct_2006, pozzetti_vimos_2007, arnouts_swire-vvds-cfhtls_2007, vergani_vimos_2008}. These recent studies imply that most of the massive galaxies must already have formed at $z\\simeq 1$, supporting the so-called {\\it downsizing} scenario \\citep{cowie_evolution_1997}. In this scenario, star formation shifts from higher to lower mass galaxies as a function of cosmic time $t$, although the stellar mass in massive galaxies may not directly follow the growth history of their Dark Matter Halos (DMHs) which build up hierarchically from less massive systems to more massive ones at a later time.  However, several studies of massive galaxies at $z < 1$ \\citep{brown_evolving_2007, faber_galaxy_2007, brown_red_2008} report the evolution from predominantly blue to red galaxies, suggesting that the number and stellar mass of galaxies in the red-sequence may continue to grow, even after $z = 1$ \\citep{bell_nearly_2004, zucca_vimos_2006}, in agreement with the hierarchical merging scenario. Recently, a mass-dependent quenching of star formation has been proposed to explain these rapid increase of the number density of the massive galaxies relative to lower-mass galaxies at $z=1-2$ (Peng et al. 2010; Kajisawa 2011; Brammer et al. 2011).  The evolutionary path of the $M_{\\ast}(M_{\\ast;0}, t)$ up to the present mass $M_{\\ast;0}$ can be also reconstructed from knowledge of the $SFR(M_{\\ast},t)$.  Recent studies have confirmed the features for rapid SFR evolution up to $z \\simeq 2$ \\citep{brinchmann_mass_2000, fontana_assembly_2003, prez-gonzlez_spitzer_2005, feulner_specific_2005, feulner_connection_2005, papovich_spitzer_2006, iglesias-pramo_uv_2007, zheng_infrared_2007, cowie_integrated_2008, damen_star_2009, dunne_star_2009, santini_star_2009, pannella_star_2009}. However, we have not found a general consensus concerning the stellar mass dependence for specific SFRs: $sSFR =SFR/M_{\\ast}$ (the SFR per unit stellar mass).  Most studies based on UV derived SFR's have found evidence to support an anti-correlation with stellar mass, not only local galaxies, but also the distant Universe, supporting the {\\it downsizing} scenario. It is not easy, however, to accept that UV estimates are the only way to estimate the SFR that will capture all the emission originally from new born stars in a galaxy, which can be independently estimated from rest-frame luminosities in the UV, IR, Submillimeter (Submm), and radio wavelength regions by introducing appropriate calibration factors.  Although the UV luminosity is used as a proxy of the SFR, its reliability often depends on the use of large, and rather uncertain corrections for dust extinction.  Most of the short wavelength photons from star-forming galaxies are reprocessed through absorption and scattering with dust in their star-forming regions, and then re-emitted in the IR-Submm wavelength regions. Thus, the mid/far-IR and their submillimeter emissions become an important tracer of the total star forming activity in high-$z$ galaxies.  The total IR (TIR) luminosity $L_{IR}$ integrated across the long wavelength spectral region is a good proxy for the SFR, which is relatively independent of the dust obscuration \\citep{kennicutt_star_1998}.  Luminous InfraRed Galaxies (LIRGs) and Ultra Luminous InfraRed Galaxies (ULIRGs) up to $z=1-2$, detected at the Mid-InfraRed (MIR) wavelength with the IR observing satellites as the ISO \\citep{mann_observations_2002, elbaz_bulk_2002, pozzi_mid-infrared_2004, dennefeld_firback._2005, taylor_properties_2005, marcillac_mid_2006, marcillac_star_2006}, the AKARI \\citep{takagi_multi-wavelength_2007, takagi_polycyclic_2010}, and the {\\it Spitzer} \\citep{egami_spitzer_2004, shapley_ultraviolet_2005, caputi_linking_2006, choi_star_2006, marcillac_star_2006, marcillac_mid_2006, papovich_spitzer_2006, papovich_spitzer_2006-1, reddy_star_2006, elbaz_reversal_2007, papovich_spitzer_2007, rocca-volmerange_12_2007, shim_massive_2007, teplitz_measuring_2007, brand_spitzer_2008, magnelli_irac_2008, dasyra_0.9_2009, hernn-caballero_mid-infrared_2009, murphy_balancingenergy_2009, santini_star_2009, strazzullo_deep_2010}, have shown that the SFR rapidly increases with the redshift, as they contribute $\\simeq 70 \\%$ of the cosmic infrared luminosity density \\citep{le_floch_infrared_2005}. The MIR emission from LIRGs/ULIRGs is evidence of dusty starburst or AGN surrounded by a dusty torus. Recently, the LIRGs/ULIRGs are expected to suppress star formation at $z=1-2$ with feedback activity from starbursts or AGNs. Understanding of the LIRGs/ULIRGs at $z=1-2$ can be a key to solve the cosmic history of star formation and AGN activities. The SFR for galaxies detected with the {\\it Spitzer} is frequently estimated with the TIR luminosity estimated from the 24 $\\mu$m flux \\citep{daddi_new_2004, santini_star_2009, pannella_star_2009}, which is nearly proportional to their stellar mass up to $z \\simeq 2$. These recent results with the {\\it Spitzer} are also basically consistent with recent results from stacking analysis for their radio emissions \\citep{dunne_star_2009, pannella_star_2009} as another SFR estimation immune to dust extinctions. These two radio stacking analyses, however, show a discrepancy in their details as \\cite{dunne_star_2009} report sSFR with a negative correlation to stellar mass while \\cite{pannella_star_2009} report sSFR with stellar mass independence. Thus, there may be various unknown factors in the SFR estimations. The TIR luminosity $L_{IR}$ of LIRGs/ULIRGs is generally derived from the fitting of the observed Spectral Energy Distribution (SED) with synthetic templates. In the SED fitting, we could sometimes use only a few bands in the IR wavelength without longer $100-1000~\\mu$m wavelengths, which may cause large uncertainties in the $L_{IR}$ estimation. Obscured AGNs are another notable problem in the $L_{IR}$ estimation. It means that a sample selection without significant AGN contamination is also important to obtain an accurate $SFR_{IR}$. Thus, systematically designed surveys are required to construct samples for distant star-forming galaxies.  Optical surveys have been successful in selecting numerous Lyman Break Galaxies (LBGs) as star forming galaxies beyond $z \\sim 2$ \\citep{steidel_spectroscopy_1996, steidel_spectroscopic_1996, steidel_lyman-break_1999, lehnert_luminous_2003, iwata_lyman_2003, ouchi_subaru_2003, ouchi_subaru_2004, dickinson_color-selected_2004, sawicki_keck_2006-1, sawicki_keck_2006} and even up to z $\\simeq 6$ \\citep{yan_high-redshift_2004, stanway_lyman_2003, bunker_star_2004, bouwens_galaxy_2004, shimasaku_number_2005}. They may, however, fail to detect substantial numbers of massive galaxies at $z < 3$, in which the Lyman break photometric technique does not work well with ground-based telescopes because of telluric absorption.  Furthermore, even if galaxies are undergoing intensive star formation, the internal dust extinction of their UV radiation may reduce the prominence of their Lyman break features, to appear as LIRGs/ULIRGs in the local Universe. However, alternative photometric techniques can be used to select massive high-$z$ galaxies through their Balmer/4000~\\AA\\ break or the 1.6 $\\mu$m bump features in their Spectral Energy Distributions (SEDs), as evidence of the presence of an evolved stellar population.  We will designate the former and the latter as 'Balmer' Break Galaxies (BBGs) and as IR 'Bump' Galaxies (IRBGs), respectively.  It is known that both star-forming and passive BBGs in the redshift desert around $z = 2$ can be selected with a simple and sophisticated two-color selection based on $B,z$, and $K$ band photometry \\citep{daddi_new_2004}. This three filter optical-NIR technique was designed to detect the Balmer/4000~\\AA\\ break, nearly free from the dust extinction which caused some difficulties for previous `red color' techniques relying on a color diagnostic of $(R-K)_{Vega} >5-6$, corresponding to $(R-K)_{AB}> 3.35-4.35$ \\citep{elston_deep_1988, thompson_surface_1999, daddi_eros_2000, firth_las_2002, roche_clustering_2002, roche_nature_2003} or for $(J-K)_{Vega}>2.3$ corresponding to $(J-K)_{AB}>1.3$ \\footnote{$J_{AB}=J_{Vega}+0.89, Ks_{AB} = Ks_{Vega}+1.84$} \\citep{pozzetti_k20_2003, franx_significant_2003}. With the near-IR (NIR) observations from the {\\it Spitzer} and from the AKARI, it has become possible to select IRBGs at $z>1$ with NIR colors at $> 3$ $\\mu$m of the 1.6 $\\mu$m bump \\citep{simpson_detection_1999, sawicki_1.6_2002, berta_contribution_2007}. These optical-NIR observations, detecting the Balmer/4000~\\AA\\ break and the 1.6 $\\mu$m bump, have been also useful to trace how stellar mass is correlated with the absolute NIR luminosity.  Although optical-NIR multi-band photometry has the potential to select a sample for star-forming massive galaxies, it may not be completely free from the AGN contamination. MIR multi-band photometry tracing PAH emissions at 6.2, 7.7 and 11.3 $\\mu$m and Silicate (Si) absorption at 10~$\\mu$m (rest-frame) from dusty star-forming regions, can be used to discriminate between the star-forming activity in LIRGs/ULIRGs and that of AGN dominated galaxies. Thus, systematic multi-wavelength observations that include observations at $> 2~\\mu$m are important for studies the star-forming activity of massive high-$z$ galaxies, and in looking back at the evolution of their SFRs. In order to connect the history of stellar mass growth in massive galaxies traced in optical-NIR wavelength with ground based telescopes and that of star formation traced in MIR wavelength observed outside of the atmosphere, we undertook multi-wavelength surveys of deep and wide areas with $\\simeq 0.38$ and $\\simeq 5.8$ deg$^2$, respectively, in a field close to the North Ecliptic Pole (NEP), which were extensively observed at the IR wavelength with the first Japanese infrared astronomical satellite AKARI.  In this paper, we attempt to reconstruct the cosmic history of dusty massive galaxies up to $z \\simeq 2$, selected estimated from optical-IR images in the NEP survey field. We report on: 1) photometric redshift $z_{phot}$, stellar mass $M_{\\ast}$, and $SFR_{UV;cor}$, which are estimated with their optical-NIR SED fittings, 2) the distinction between dusty starburst and AGN dominated galaxies using the MIR SED feature, 3) TIR luminosities of starburst dominated galaxies with the subtraction of the obscured AGN contribution, from which we can derive $SFR_{IR+UV}$ complementary to the $SFR_{UV:cor}$, and 4) IR luminosities of the dusty tori of AGN dominated galaxies, which show evolutionary trends in the connection between AGN and star forming activities.  In section~\\ref{sec:data}, we summarize the data we used; in section~\\ref{sec:photo}, we present the photometry for $z'$-band detected galaxies in our multi-wavelength dataset; in section~\\ref{sec:irbg}, we also report AKARI/IRC detected IRBGs classified using the redshifted 1.6-$\\mu$m IR bump in the NIR bands, which are divided into three redshifted subgroups; in section~\\ref{sec:pah}, we present AKARI/IRC MIR detected galaxies which correspond to LIRGs/ULIRGs and classify them into dusty starbursts and AGN from their MIR SEDs with/without PAH features in the MIR bands; in section~\\ref{sec:ir_lum}, we derive the TIR luminosities for the MIR detected galaxies; in section~\\ref{sec:evol_sfr}, we estimate their star formation rates and extinctions; in section~\\ref{sec:evol_pop}, we reconstruct their intrinsic stellar populations with extinction correction; in section~\\ref{sec:evol_agn} we also study the AGN activities in the MIR detected galaxies; in section~\\ref{sec:discuss}, we discuss some expectations for on-going and future observations; finally in section~\\ref{sec:concl}, we summarize the conclusions of our study.  In this paper we adopt $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$ throughout.  All magnitudes are represented in the AB system (e.g. Oke 1974). ",
        "conclusions": "\\label{sec:concl}  We have studied BBGs/IRBGs in three redshift ranges of $z=0.4-0.8, 0.8-1.2$, and $>1.2$ with MIR multiband photometry, which revealed their star-forming/AGN activities with/without PAH emissions deriving SFR from their TIR luminosities.  \\begin{itemize} \\item The MIR colors and SED analysis can classify MIR bright N3 Red galaxies (MbN3Rs) into three classes; starburst dominates (sb-MbN3Rs), starburst/AGN co-existence(s/a-MbN3Rs), and AGN dominates (agn-MbN3Rs). The ratio of rest-frame luminosities at 5 and 10 $\\mu$m to that at $7.7~\\mu$m is systematically higher in the s/a-(agn-)MbN3Rs than in the sb-(s/a-)MbN3Rs, which is consistent with the picture of the s/a-MbN3Rs as the starburst/AGN co-existent subpopulation in the dusty star-forming MbN3Rs (s-MbN3Rs). \\item MIR marginally-detected N3Rs (MmN3Rs) can be also classified into two classes; sb-MmN3Rs as mimics of sb-MbN3Rs, and non sb-MmN3Rs. \\item Their rest-frame 7.7-$\\mu$m luminosity is a good tracer of TIR luminosity as the PAH emission dominates for the sb- and s/a-MbN3Rs and sb-MmN3Rs which are classified as dusty star-forming galaxies even up to $z=2$. \\item The $SFR_{IR+UV}$ derived from the total bolometric luminosity $L_{IR+UV}$ of the sb-MbN3Rs/MmN3Rs shows a correlation that is nearly proportional to the stellar mass $M_{\\ast}$, which is consistent with that of the extinction corrected UV star formation ratio $SFR_{UV;corr}$. \\item The $SFR_{IR+UV}$ of the s/a-MbN3Rs tends to be smaller than that of the sb-MbN3Rs at $z<0.8$. \\item The IR-derived specific SFR of the sb-MbN3Rs/MmN3Rs rises with the redshift at all stellar masses. \\item The ratio of the rest-frame luminosity at $10~\\mu$m to that at $7.7~\\mu$m of the sb-MbN3Rs/MmN3Rs may trace the optical depth due to the dust extinction. \\item The AGN activity could be traced with the rest-frame 5-$\\mu$m luminosity for the agn- and the s/a-MbN3Rs. The evolutionary trend $(1+z)^6$ of the AGN luminosity for the agn-MbN3Rs is more rapid than the $(1+z)^3$ of SFR for the sb-MbN3Rs/MmN3Rs. \\item SMBHs in the galaxies have grown to $\\simeq 3 \\times 10^8 M_{\\odot}$ in maximum luminous populations with $10^{12}L_{\\odot}$ and $10^{11} M_{\\odot}$ of the MIR selected AGNs at $z>1.2$, which suggests the mass relation between the SMBH and its host is already established at $z\\simeq 1-2$ and the same as that at the present day. \\end{itemize}  \\bigskip This study is based on data mostly collected with the AKARI IR astronomical satellite and at the Subaru Telescope, which are operated by the ISAS/JAXA and the National Astronomical Observatory of Japan, respectively.  We would like to thank their staff for their invaluable assistance.  We thank B. Weiner and C. Papovich for their help in the spectroscopic observation with the Hectospec/MMT under TSIP program of National Optical Astronomy Observatory. We are grateful to an anonymous referee for his/her comments which have helped to clarify the paper. HH thanks P. Langman, Iwate University, for proofreading the English writing. This research was mostly supported for HH by the Grant-in-Aid for Scientific Research (21340042) from the Japan Society for the Promotion of Science (JSPS), and partly supported for MI by the Korea Science and Engineering Foundation (KOSEF) grant no. 2010-0000712 from the Korea government (MEST).  \\appendix"
    },
    "1112/1112.1386_arXiv.txt": {
        "abstract": "We consider the the effect of the electron cloud about a vorton in the \\CFLKz\\ high-density phase by numerically solving the ultrarelativistic Thomas-Fermi equation about a toroidal charge. Including electrons removes the electric monopole contribution to the energy, and noticeably decreases the equilibrium radius of these stable vortex loops. ",
        "introduction": "% \\label{sec:introduction}  The behavior of matter at very high densities, like those found at the core of neutron stars, remains  largely an open problem. This is due to a lack of reliable methods to extract information directly from QCD. But since the temperatures relevant for the physics of neutron stars is very low (as compared to the Fermi energy) a number of observables of  interest depend only on properties of the low lying excitations of the ground state. In the case of dense matter it is believed that those excitations are pseudo-Nambu-Goldstone (NG) bosons arising from the spontaneous breaking of certain approximate symmetries. The study of these quasiparticles using the methods of effective field theory opens up the possibility of a more reliable understanding of dense matter.  The breaking of symmetries in high-density QCD is related to quark pairing. At high density, evidence now points towards color superconductivity where quarks pair in different combinations depending on the temperature and chemical potential\\cite{Barrois:1977xd,Bailin:1983bm,Alford:1997zt,Rapp:1997zu}.  This pairing uses the BCS mechanism, just like electron pairing in conventional electrical superconductors.  However, color superconductivity is a much richer phenomenon than electrical superconductivity because quarks carry flavor and color while electrons do not---different color superconducting phases arise depending on which quarks pair. The nature of the NG bosons depends crucially on the pattern of symmetry breaking and, consequently, on the kinds of quark pairing found in the ground state.  At asymptotically high density, three-flavor QCD is believed to pair in a pattern known as color-flavor-locking (CFL), as it pairs all quarks\\cite{Alford:1998mk}.  At lower densities, where the QCD coupling is not small, the problem is more complicated and a plethora of  phases have been suggested (for a review see \\cite{Alford:2001dt,Rajagopal:2000wf}).  Perhaps the most likely phase at very high density is the \\CFLKz\\ phase, which seems guaranteed to be the energetically favored phase for some range of large densities\\cite{Bedaque:2001p107,Kaplan:2001qk}.  This range may include densities relevant for compact astrophysical objects and thus may be phenomenologically relevant for extracting predictions about extreme environments, like the inner cores of neutron stars.  The NG modes of the CFL phase form an SU(3) octet, and are named in direct analogy with the \\emph{in vacuo} SU(3) flavor octet, but at high density the ``kaons'' are the lightest particles \\cite{Son:1999cm,Beane:2000ms,Son:2000tu,Casalbuoni:1999wu,Casalbuoni:1996pg}.  As they live in an SU(3) octet, their behavior is described by a theory resembling familiar chiral perturbation theory.  As the chemical potential $\\mu$ is lowered and the quark masses become relevant some of the NG modes become lighter.  When $\\mu$ is low enough, the neutral ``kaons'' would have a negative mass-squared, and instead condense with some vacuum expectation value, turning the CFL phase into \\CFLKz \\cite{Bedaque:2001p107,Kaplan:2001qk}.  This neutral kaon condensate provides a superfluid background that is the stage for interesting topological defects. The reason the neutral, as opposed to the charged, kaons condense is the small differences in their mass induced by isospin violating quark masses and electromagnetic effects. These differences are small so a condensate of charged kaons is almost degenerate with  the favored neutral kaon condensate.  This peculiar situation provides a physical realization of the so-called ``superconducting string\" solutions, suggested in more abstract contexts\\cite{Witten:1984eb} and also appearing in some models describing beyond the Standard Model physics\\cite{Hindmarsh:1994re,vilenkin2000cosmic,Davis1989209,Davis:1988ij} . Regular global vortices appear due to the \\Kminus condensation. At their core, the value of the \\Kzero\\ condensate vanishes, as it is the case in all vortices. Since there is a close competition between neutral and charged kaon condensates, the charged kaons will then condense in the core of the \\Kzero\\ vortex. The condensation of a charged particle leads to \\emph{electrical} superconductivity (not color superconductivity) along the vortex. Loops of superconducting vortices (``vortons\") can be stabilized by the supercurrent running around them. The mechanism, roughly speaking, is that the energy of a vorton has two main contributions. One is the  tension along their length, which creates an energy that scales with its radius $R$ and drives a vorton to be small.  However, the energy of a vorton includes terms that scale like  $J^{2}/Q^{2}R$ where  $J$ is its angular momentum and $Q$ its electric charge.  These terms prevent a vorton from collapsing. This mechanism was invoked in \\cite{Kaplan:2001hh} in the context of quark dense matter. It was subsequently pointed out \\cite{Buckley:2002ur,Buckley:2002p109}  that, just like in the case of cosmological vortons, the currents necessary to stabilize a vorton could be high enough that they would quench the superconductivity of the vortex core, destroying its supercurrent and rendering the stability mechanism mute. The quenching effect can, however, be counteracted by electric charge. A vorton with an overall electric charge, in addition, would be more stable on account of Coulomb repulsion between opposong sides of the loop.   A more quantitative theory of quark matter vortons appeared only recently. In Ref.~\\cite{Bedaque:2011fg}, the equilibrium size of vortons containing charge as well current was estimated by making a variety of approximations and assumptions.  It was found that stabilizing a vorton so that its equilibrium radius was several times the thickness of its core $\\delta$ and thus could really be identified as a torus, required large electric charges ($Z\\sim10^{4}$).  With such a large charge, it is possible that ambient electrons which are in the bulk to ensure charge neutrality could be pulled close to the vorton, shielding its electric fields, and significantly reducing the effectiveness with which the electric charge aided in stabilizing the vorton in the first place.  We call these vortons with neutralizing orbiting electrons vortonium to easily distinguish them from unshielded or bare vortons.  It is the purpose of this paper to account for these electrons and to find the equilibrium radius of a vortonium ``nucleus''.  We use the Thomas-Fermi approximation to find the electrons' semiclassical charge distribution, allowing them to shield the electric field.  Second, having found the electron distribution, we can compute the energy in the electric field numerically, without making geometrical or multipole approximations.  Intuitively, when including electrons we should find that vortons have a smaller equilibrium radius, because there are exactly enough electrons to kill the monopole field.  We find that making these  improvements decreases the equilibrium radius of a vorton of given angular momentum and electric charge significantly. ",
        "conclusions": "% \\label{sec:summary_and_outlook}   In this paper we consider the effect of electron shielding on the equilibrium radius of vortons in the \\CFLKz\\ phase at high density.  For parameter values where unshielded vortons were reasonably large and had a clear toroidal shape with an aspect ratio $R/\\delta$ of a few ($Z\\gtrsim 6000$), the electron shielding of vortonium, as expected, shrinks both the radius and aspect ratio for a given $Z$ and also diminishes the response of $R/\\delta$ to a change in $Z$.  By allowing that portion to be shielded with electrons and applying a relativistic Thomas-Fermi approximation, we have accounted for the electron cloud surrounding such a vorton and have relieved some of uncontrolled approximations concerning the toroidal geometry of vortons in \\CFLKz.  We have shown that the energy in the electric field considered in Ref.~\\cite{Bedaque:2011fg} is unrealistically optimistic insofar as it drastically changes the equilibrium shape of the vorton as compared to the more realistic vortonium.  The reduced size of vortonium may be understood by realizing that the far-fields considered for unshielded vortons is predominantly a monopole piece while for vortonium it is a quadrupole.  Two directions seem obvious for pursuing the physics of \\CFLKz\\ vortons further.  The first is to try to cure all of the issues we have touched on in this work:  attempt to find, in a toroidal geometry, the actual mean-field wavefunctions of the neutral and charged kaon condensates for a vorton, while simultaneously incorporating electron shielding and allowing electrons to penetrate the core and not artificially excluding them, and could account for curvature effects.  Work in this direction will almost certainly require intensive numerical calculations.  The other possible direction is to try to estimate the vorton production rate  in the early times in the neutron star life as it goes from a high temperature disordered phase to a low temperature \\CFLKz\\ phase. This number is important to gauge the vorton influence on neutron star phenomenology.   Along the same lines, one might consider the decay rate of these objects through weak effects, to try and estimate how long any created population might persist in these extreme environments.  This line of inquiry would help prioritize any other work on vortons at high density, and would provide the most direct route to understanding if vortons might contribute to observable astrophysical phenomena."
    },
    "1112/1112.3383_arXiv.txt": {
        "abstract": "We report the discovery and orbital solutions for two new OB binaries in the Cygnus~OB2 Association, MT311 (B2V~$+$~B3V) and MT605 (B0.5V~$+$~B2.5:V).  We also identify the system MT429 as a probable triple system consisting of a tight eclipsing 2.97 day B3V+B6V pair and a B0V at a projected separation of 138 AU.  We further provide the first spectroscopic orbital solutions to the eclipsing, double-lined, O-star binary MT696 (O9.5V~$+$~B1:V), the double-lined, early B binary MT720 (B0-1V~$+$~B1-2V), and the double-lined, O-star binary MT771 (O7V~$+$~O9V).  These systems exhibit orbital periods between 1.5 days and 12.3 days, with the majority having P$<$6 days.  The two new binary discoveries and six spectroscopic solutions bring the total number of known massive binaries in the central region of the Cygnus~OB2 Association to 20, with all but two having full orbital solutions. ",
        "introduction": "The ubiquity of close massive binaries in open clusters is becoming well documented (e.g., \\citealt{Sana09}, \\citealt{Mahy09}, \\citealt{Sana08}, and \\citealt{debeck06}). However, the number of orbital solutions for massive binaries is still relatively small. In our previous works documenting the number and statistics of massive binaries in Cygnus OB2 \\citep{Kiminki07, Kiminki08, Kiminki09}, we discovered nine systems (MT059, MT145, MT252, MT258, MT372, MT720, MT771, Schulte~3, and Schulte~73; notation from \\citealt{MT91} and \\citealt{Schulte56}). We provided spectroscopic solutions to these nine and two additional systems (A36 and A45; notation from \\citealt{Comeron02}) as part of our ongoing radial velocity survey in one of the Galaxy's super OB clusters. The primary objective of the Cyg~OB2 radial velocity survey is to aid our understanding of numerous open questions involving massive stars, including their formation, the predicted rates of supernovae type Ib/c arising through binary channels, and whether there is a correlation between cluster density and the binary fraction therein.  In this work, we report discovery of two new double-lined spectroscopic binaries (SB2s), and we provide full orbital solutions (periods, mass ratios, eccentricities and probable inclinations) for these and four known binaries previously lacking full solutions. This brings the number of massive binaries having full or partial solutions to 20, the most of any single cluster or association. In a companion paper \\citep{Kiminki11b}, we summarize the current state of massive binaries in the literature and use the combined binary orbital parameter information and ephemerides for 113 massive stars in Cyg~OB2 to infer the intrinsic orbital parameter distributions and the binary fraction of the Association.  We present these findings as a complement to analyses of massive stars in other clusters (such as the work with Sco~OB2 by \\citealt{Kouwenhoven07}, NGC~6231 by \\citealt{Sana08}, NGC~2244 by \\citealt{Mahy09}, and NGC~6611 by \\citealt{Sana09}). For a more comprehensive discussion of the scope of this project see \\citet{Kiminki11b}.  Section~2 reviews the observations of this ongoing 12-year radial velocity survey and discusses the method we used to determine radial velocities. Section~3 discusses the method for computing orbital solutions and provides new spectroscopic orbital solutions for MT311, MT429, MT605, MT696, MT720, and MT771. Finally, Section~4 summarizes the survey findings to date. ",
        "conclusions": "We present orbital solutions for two new SB2s, MT311 and MT605. We also provide the first spectroscopic solutions to the eclipsing system, MT696, the early B, eclipsing triple SB1, MT429, the early B SB2, MT720, and the late-O SB2, MT771.  Table~\\ref{Binaries4} lists the 20 known OB binaries in Cyg~OB2. Column~1 gives the system designations.  Column~2 shows the binary types (i.e., eclipsing, SB1, SB2, etc). Column~3 provides the spectral types of the components. The spectral types in square brackets have been determined indirectly from the data (i.e., based on the computed mass ratio rather than the component's spectrum). Columns~4--6 give the orbital periods, eccentricities, and mass-ratios respectively. Finally, column~7 lists the corresponding surveys responsible for the orbital information. Figure~\\ref{binaries4} shows the locations of these systems relative to the other OB stars in this survey.  There are no apparent clusterings or groups of massive stars, nor are there tendencies of the massive binaries to reside in the center of the Association versus the outskirts.  Of the 20, we know reliably the orbital periods of all 20, the orbital eccentricities of 18, and the mass-ratios of all 20 (the mass ratios for the SB1s are actually range estimates based on the range of probable inclinations). There are six SB1s, 14 SB2s, five eclipsing binaries of the Algol type, two eclipsing binaries of the $\\beta$~Lyr type, and one eclipsing binary of the W~UMa type. At least five binaries are both spectroscopic and eclipsing systems. In a companion paper \\citep{Kiminki11b}, we use the orbital parameter information and ephemerides associated with the short period systems ($P<26$~days) listed in Table~\\ref{Binaries4}, along with the ephemerides of 94 additional OB stars in Cyg~OB2, to compute the intrinsic binary orbital parameter distributions (i.e., period, mass ratio, and eccentricity) and intrinsic binary fraction of the cluster."
    },
    "1112/1112.3510_arXiv.txt": {
        "abstract": "We built models for low bulge mass spiral galaxies (late type as defined by the Hubble classification) using a 3-D self-gravitating model for spiral arms, and analyzed the orbital dynamics as a function of pitch angle, going from 10$\\deg$ to 60$\\deg$. Testing undirectly orbital self-consistency, we search for the main periodic orbits and studied the density response. For pitch angles up to approximately $\\sim 20\\deg$, the response supports closely the potential permitting readily the presence of long lasting spiral structures. The density response tends to ``avoid'' larger pitch angles in the potential, by keeping smaller pitch angles in the corresponding response. Spiral arms with pitch angles larger than $\\sim 20\\deg$, would not be long-lasting structures but rather transient. On the other hand, from an extensive orbital study in phase space, we also find that for late type galaxies with pitch angles larger than $\\sim 50\\deg$, chaos becomes pervasive destroying the ordered phase space surrounding the main stable periodic and quasi-periodic orbits and even destroying them. This result is in good agreement with observations of late type galaxies, where the maximum observed pitch angle is $\\sim 50\\deg$. ",
        "introduction": "\\label{sec:intro} It is assumed that the Hubble sequence reveals a strong correlation between the morphology of galaxies and their formation process. The study of this relation is a very active field where the merger hypothesis plays a fundamental role; however, it is very unlikely that all trends observed in this sequence can be explained solely by this hypothesis. In particular, the long-term galactic dynamics of non-interacting galaxies is also determined by their inner structure. These highly non-linear stellar systems are prone to exhibit the coexistence of an astonishing dichotomy, an exquisite order right together with counterintuitive complex regions of chaos, where non-axisymmetric features such as spiral arms, bars, etc, play a key part.  Regarding to ordered motion, smooth and weak long-lasting large scale structures such as spiral arms (in general, low mass and/or low pitch angle), when modeled, present a relatively simple orbital structure made of families of quasi-periodic orbits surrounding the main periodic obits that sculpt spiral arms. The density response is also smooth and coincides nicely with the imposed potential in the case of low pitch angle spiral arms, for example. There are indications that even chaotic orbits could reinforce observed morphological features in this case (Kaufmann \\& Contopoulos 1996; Patsis, Athanassoula \\& Quillen 1997; Harsoula, Kalapotharakos \\& Contopoulos 2011). However, orbital analysis on dynamical models suggests that chaotic motion plays a significant role (Contopoulos 1983,1995; Contopoulos, Varvoglis \\& Barbanis 1987; Grosb\\o l 2003) in spirals.  Indeed, large scale structures are not expected to emerge on systems built out of pure chaos, this is, as long as chaos does not become pervasive, large scale structures of discs are practically unaffected. Recently, there has been an interesting discussion about the possible chaotic nature of the spiral structure (Patsis 2006; Voglis, Stavropoulos \\& Kalapotharakos 2006; Romero-G\\'omez \\et 2007; Voglis, Tsoutsis \\& Efthymiopoulos 2006; Contopoulos \\& Patsis 2008; Patsis \\et 2009).  Although the best known part of the Hubble classification regarding pitch angles, categorizes galaxies assuming that late types possess the most opened spiral arms (largest pitch angles), this is just the envelope of the classification. Late type spirals actually present a large scatter in this parameter (Kennicutt 1981; Ma \\et 2000), going from about 10$\\deg$ to 50$\\deg$. In particular, late type spirals (Sb to Scd), are better fitted with strong spirals (Contopoulos \\& Grosb\\o l 1986; Patsis et al. 1991; Patsis, Grosb\\o l \\& Hiotelis 1997, and references therein), meaning they are far from being a slight perturbation that can be reproduced by the Lin \\& Shu (1967) spiral arm potential with a cosine function, solution of the Tight Winding Approximation (TWA).  We present a first result of a detailed orbital study on models of late type spiral galaxies as defined by Hubble (1926). In particular, this work is devoted to one of the structural parameters of spiral arms: the pitch angle. In this study some restrictions are imposed theoretically on their steady or transient nature, and on their maximum pitch angle.  This letter is organized as follows. In Section \\ref{model}, the 3-D galactic potential used to compute orbits is briefly described. In Section \\ref{results} we present our results with periodic orbit analysis, density response and phase space studies. Finally, in Section \\ref{conclusions} we present our conclusions. ",
        "conclusions": "With the use of an axisymmetric fixed model to simulate a typical late type galaxy as a background, we superposed a bisymmetric steady spiral arm potential ({\\tt PERLAS}) and studied the evolution of orbital behavior in the plane of the disk, as we change the pitch angle going from 10$\\deg$ to 60$\\deg$, in order to set some structural restrictions to the spiral arms, based on orbital dynamics. Observed galaxies classified as late type spirals present a wide scatter in pitch angles, going from $\\sim 10\\deg$ to $50\\deg$.  With these family of models, we have carried out an exhaustive orbital study (order and chaos) with periodic orbits and with phase space diagrams.  In this paper we present the first restriction relative to the pitch angle. In the case of ordered motion, with periodic orbits, a limit in the pitch angle of the density response is found at approximately $20\\deg$, up to which, the density response reinforces the spiral arms potential at all radii, i.e. with a more long-lasting nature. Beyond this limit, the density response ``avoids'' to follow the spiral arm potential, producing pitch angles much smaller than the background spiral potential. Spiral arms beyond this limit might be better explained as transient structures.  A second restriction is obtained out of chaos behavior. With the phase space orbital study, going from pitch angles of 10$\\deg$, where order reigns, to more than 50$\\deg$, where chaos becomes pervasive. With this orbital study we are able to set a limit value for the maximum pitch angle before the system becomes completely chaotic. This limit closely coincides with the observational maximum pitch angle of spirals ($\\sim 50\\deg$), suggesting a possible relation between the structural characteristics of the galaxy and chaos."
    },
    "1112/1112.0513_arXiv.txt": {
        "abstract": "We present model spectra of stellar populations with variable chemical composition. We derived the [$\\alpha$/Fe] abundance ratio of the stars of the most important libraries (ELODIE, CFLIB and MILES) using full spectrum fitting and we generated PEGASE.HR models resolved in [$\\alpha$/Fe]. We used a semi-empirical approach that combines the observed spectra with synthetic stellar spectra. We tested the models using them to derive  [$\\alpha$/Fe] in galaxies and star clusters using full spectrum fitting. The present models are available from \\url{http://ulyss.univ-lyon1.fr} ",
        "introduction": "The comparison between models and observations is the key to derive the metallicity, age and star formation history of galaxies and star clusters. The evolutionary synthesis of galaxies consists in choosing an initial mass function (IMF) and following the evolution of the population over the time according to some theoretical evolutionary tracks. This approach predicts the distribution of the stars in the effective temperature (\\teff) vs. surface gravity (log g) plan at any epoch, and the final step is to draw spectra from a library to assemble the integrated spectrum.  The modelling of stellar populations steadily progressed since the first attempts in the 80s \\citep[e.g.][]{bruzual1983}. Beside the metallicity, age and history, the detailed chemical composition became an important parameter to measure. This  allows one to trace the source of metal enrichment, because of the different yields of type {\\sc i}a and {\\sc ii} supernovae. In turn, since both SN types correspond to different time-scales, this informs about the duration of the enrichment process. A population that produced its metal and formed its stars before the  SN{\\sc i}a onset will be deficient in Fe (or as more commonly said, over-abundant in $\\alpha$-elements, like O and Mg). To vary the yield in the models, both the stellar library and the evolutionary tracks have in principle to be changed. However, \\cite{dotter2007} have shown that the tracks essentially depend on the total metallicity, rather than of the actual mixture. Therefore, the scaled-solar tracks can reasonably be used for any composition, provided they are associated with stellar spectra of the same total metallicity.   The models predict colours, used for SED fitting, spectra, and line-strength indices. Indices are equivalent-widths defined for a number of prominent spectral features and their study brings an approximate handle on the physical parameters (age, metallicity and/or detailed composition).  Several authors modeled indices with variable abundances of $\\alpha$-elements relative to Fe \\citep{weiss1995, trager2000, thomas2003}. The approach consisted in evaluating the response of individual indices to changes of composition using theoretical spectra \\citep{tripicco1995}. Then, assuming that the abundance pattern in the library is matching that of the solar neighborhood (justified by the fact that the observed libraries consist of nearby stars), the indices could be estimated for any \\teff{} log g, \\feh{} and \\afe. In the visible region of the spectrum, the variation of \\afe{} results in a prominent change of the Mg$_b$ index near 5175~\\AA, and therefore a combination of this index with some Fe indices could be used to derive the \\afe{} in galaxies and clusters.  Models of indices were developing faster and further than the spectral predictions, because of the lack of stellar libraries calibrated in relative flux. However, the quality of the libraries also progressed. Libraries of about 1000 stars covering a wide range of the parameter space, with an extended coverage of the optical spectrum at mediun spectral resolution and a fair flux calibration became available. The three most important are ELODIE \\citep[R=$\\lambda/\\Delta\\lambda \\approx$10000, $\\lambda\\lambda = 3900 - 6800$ \\AA]{elodie31}, CFLIB \\citep[R$\\approx$4000, $\\lambda\\lambda = 3460 - 9400$ \\AA]{valdes2004} and MILES \\citep[R$\\approx$2000, $\\lambda\\lambda = 3536 - 7410$ \\AA]{sanchez-blazquez2006}.  In this paper, we present the first spectral models with variable \\afe{} based on these libraries and on precise measurements of \\afe{} in their stars.    \\begin{figure}[b] \\begin{center} \\includegraphics[width=0.95 \\textwidth]{prugniel_fig1.eps} \\caption{Distribution of \\afe{} in the three libraries, compared to the distribution in \\cite{milone2011}.} \\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "1112/1112.3040_arXiv.txt": {
        "abstract": "Current studies on Submillimeter Galaxies (SMGs) mostly focus on bright sources with 850~$\\mu$m flux greater than 2~mJy, and the results have shown that they are likely high redshift mergers with z $>$ 2 and could be a dominant population on star formation in the early Universe. However, bright SMGs only contributes 20-30\\% of the 850~$\\mu$m extragalactic background light (EBL), meaning the bulk of the cosmic star formation still hidden by dust and our current understanding is biased. We have started a program to study an unbiased sample of highly-amplified and intrinsically faint SCUBA detected SMGs in the field of massive lensing clusters. Here we report the newly obtained SMA observations at 850~$\\mu$m on one of our sample source, A2390-5, behind the massive lensing cluster A2390. We successfully detect the source with a flux of 3.95~mJy. Surprisingly, it does not have any counterpart in any other wavelengths even though there are tentative candidates, which implies a very dusty and high-z nature. With less than 1$\\arcsec$ positional accuracy and the adoption of z = 5, we obtain the amplification factor of 12 using current lensing model, which makes A2390-5 a faint SMG with a de-lensed flux of 0.33~mJy. Together with our previous detection on another faint SMG, both of them have no counterpart in other wavelengths and their properties are very different than previously thought from the single-dish data. We emphasize the importance of direct submillimeter high-resolution studies on faint SMGs, which could be the dominant population of the high-z star formation. ",
        "introduction": "Blank-field submillimeter surveys with ground-based single-dish telescopes have resolved a rare but cosmologically important population, submillimeter galaxies (SMGs; \\citealt{Blain:2002p8120}). Detail follow-up studies on both continuum and emission lines have shown that SMGs are at high redshift (2$<$z$<$5), gas rich ($M_{gas} > 10^{10} M_\\odot$), relatively short lived ($<$~100 Myr), highly clustering, likely mergers and the progenitors of massive elliptical galaxies in present-day Universe. Despite the rareness, their extreme star formation (SFR $\\sim$ 100$-$1000 M$_\\odot$ yr$^{-1}$ ) makes them a major contributor to star formation in the early Universe (\\citealt{Barger:2000p2144}).  However, the blank-field SMGs currently under extensive study are usually bright with 850~$\\mu$m fluxes above 2~mJy, and they only contribute 20-30\\% of the 850~$\\mu$m extragalactic background light (EBL; \\citealt{Coppin:2006p9123}), which is the integrated light from sources outside our Milky Way galaxy. Thus bulk of the cosmic star formation is still unresolved and the formation mechanism of the faint SMGs emitting the rest of the 850~$\\mu$m EBL naturally becomes a key issue to address some important questions such as are mergers the major contributor to the cosmic stellar assembly in the early Universe.  Unfortunately, the poor resolution (e.g., $\\sim14\\arcsec$ FWHM on the JCMT at 850~$\\mu$m) prevents us from directly measuring the faint SMGs below the 2~mJy confusion limit. Almost all of our knowledge about faint SMGs comes from observations with single-dish telescopes in the fields of massive lensing clusters. Due to the presence of the intervening cluster mass, the intrinsically faint fluxes of background sources are gravitationally amplified to a detectable level, and the confusion limit is reduced by the expansion of the source plane. Faint SMGs with fluxes between 0.1 and 2~mJy have been detected in this way, and they contribute more than 50\\% of the 850~$\\mu$m EBL (\\citealt{Cowie:2002p2075, Knudsen:2008p3824}). As with the bright SMGs, pinning down their exact location is the first and the most crucial step for any follow-up studies. We have started a program to study an unbiased sample of highly-amplified and intrinsically faint SCUBA detected SMGs in the field of massive lensing clusters. Here we discuss the newly obtained high-resolution Submillimeter Array (SMA) observations at 850~$\\mu$m on one of our sample source, A2390-5, behind the massive lensing cluster A2390. ",
        "conclusions": "The coarse resolution of the 850~$\\mu$m maps from single-dish survey has always been a major issue on the studies of the SMGs. It not only leads to confusion limits, it also makes locating the true counterparts of the SMGs at other wavelengths and deciding the accurate amplification for those strongly lensed SMGs difficult: Using high resolution interferometric imaging with the SMA to resolve blank-field SCUBA detected sources, \\citet{Wang:2011p9293} found that the multiplicity appears to be common and the selection techniques from other wavelengths are usually inadequate. A similar study by \\citet{Chen:2011p11605} on several strongly lensed faint SMGs shows that the large uncertainty on the source amplifications caused by the poor resolution can lead to severe misinterpretations about their nature.  What is more important is that, like A2390-5, the only faint SMG detected in \\citet{Chen:2011p11605} shows no counterpart in other wavelengths either (Figure \\ref{a23903}), which leads to the fact that so far all SMA-observed faint SMGs from the non-preselected sample are not traceable. This happens much more frequently than it is expected from bright SMGs (\\citealt{Younger:2007p6982, Younger:2009p9502}), meaning that we are detecting an unexpectedly large number of high redshift dusty galaxies. This could be of considerable importance given that most studies on faint SMGs use selection techniques from other wavelengths and their results could be misleading if many faint SMGs are not traceable in this way. We therefore emphasize the importance of direct submillimeter high-resolution studies of the faint SMGs, which could be the dominant population of the high-z star formation."
    },
    "1112/1112.0619_arXiv.txt": {
        "abstract": "{We found the first direct evidence that the Cepheid class namesake, $\\delta$~Cephei, is currently losing mass. These observations are based on data obtained with the \\emph{Spitzer} Space Telescope in the infrared, and with the Very Large Array in the radio. We found that $\\delta$~Cephei is associated with a vast circumstellar structure, reminiscent of a bow shock. This structure is created as the wind from the star interacts with the local interstellar medium. We measure an outflow velocity of $\\approx 35.5$~km\\,s$^{-1}$ and a mass loss rate of $\\approx 10^{-7}$ -- $10^{-6} \\ M_\\odot$~yr$^{-1}$. The very low dust content of the outflow suggests that the wind is possibly pulsation-driven, rather than dust-driven as common for other classes of evolved stars.}  \\abstract*{We found the first direct evidence that the Cepheid class namesake, $\\delta$~Cephei, is currently losing mass. These observations are based on data obtained with the \\emph{Spitzer} Space Telescope in the infrared, and with the Very Large Array in the radio. We found that $\\delta$~Cephei is associated with a vast circumstellar structure, reminiscent of a bow shock. This structure is created as the wind from the star interacts with the local interstellar medium. We measure an outflow velocity of $\\approx 35.5$~km\\,s$^{-1}$ and a mass loss rate of $\\approx 10^{-7}$ -- $10^{-6} \\ M_\\odot$~yr$^{-1}$. The very low dust content of the outflow suggests that the wind is possibly pulsation-driven, rather than dust-driven as common for other classes of evolved stars.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.0317.txt": {
        "abstract": "The interstellar medium in star-forming galaxies is a multiphase gas in which turbulent support is at least as important as thermal pressure.  Sustaining this configuration requires continuous radiative cooling,  such that the overall average cooling rate matches the  decay rate  of turbulent energy into  the medium.   Here we carry out a set of numerical simulations of a stratified, turbulently stirred, radiatively cooled medium, which uncover a fundamental transition at a critical one-dimensional turbulent velocity of $\\approx 35$ km/s. At turbulent velocities below  $\\approx 35$ km/s,  corresponding to temperatures below $10^{5.5}$K, the medium is stable, as the time for gas to cool is roughly constant as a function of temperature.    On the other hand, at turbulent velocities above the critical value, the gas is shocked into an unstable regime in which the cooling time increases strongly with temperature, meaning that a substantial fraction of the interstellar medium is unable to cool on a turbulent dissipation timescale.   This naturally leads to runaway heating and ejection of gas from any stratified medium with a  one-dimensional turbulent velocity above $\\approx 35$ km/s, a result that has  implications for galaxy evolution at all redshifts.\\\\ ",
        "introduction": "The cosmic history of dark matter  is relatively simple.  Starting as weak  perturbations in a Gaussian random field, dark matter structures  accrete and merge to form ever-larger objects, leading to a hierarchical evolution that is in excellent agreement with a wide range of observations (\\eg Fu \\etal 2008; Vikhlinin \\etal 2009; Komatsu \\etal 2011).  On the other hand, the formation of galaxies within these potentials is extremely complex.   Although gas initially collapses along with the dark matter,  it soon feels pressure support, cools radiatively,  forms molecules, and experiences a wide range of intricate physical processes.  Perhaps the most important of these is the formation of stars, which liberate vast quantities of energy that cause much of the gas to  reverse its history, parting ways from the dark matter and traveling outwards into the depths of intergalactic space.  Such large-scale outflows have been observed in star-forming galaxies over a wide range of masses and redshifts (\\eg Heckman, Armus \\& Miley 1990;  Lehnert \\& Heckman 1996;  Lehnert, Heckman, \\& Weaver 1999; Martin 1999, 2005; Strickland \\etal 2000, 2004; Pettini \\etal  2001; Veilleux, Cecil  \\& Bland-Hawthorn 2005; Westmoquette \\etal 2007, 2009), and have been shown to be of crucial importance  in reconciling the distributions of galaxies and dark-matter halos (\\eg Scannapieco \\etal 2001; Benson \\etal 2003). Yet, at the same time, galaxy outflows are very poorly-understood, as the interstellar medium (ISM) from which they form can not be modeled as a single-temperature gas.  Rather it is both continuously cooling and being stirred by turbulence, which in star forming regions is likely to be driven primarily by massive stars  (\\eg McKee \\& Ostriker 1977).  In fact, this simultaneous cooling and turbulent energy input is so extreme that it often leads to a supersonic medium in which turbulent velocities exceed the thermal velocities, and turbulent motions simultaneously support the disk and compress a fraction of the gas to drive star formation (\\eg Elmegreen \\& Scalo 2004;  Mac Low \\& Klessen 2004).  The extremely short cooling times within the ISM make it impossible to model stellar feedback by simply adding thermal energy to the gas, but the range of physical scales involved does not allow for the direct simulation of supernova remnants within a galaxy-scale simulation. To deal with this problem, galaxy outflow simulations have suppressed cooling with an empirical heating function (Mac Low \\etal 1989; Mac Low \\& Ferrara 1999), imposed an artificial temperature floor (Tomisaka \\& Bregman 1993; Tenorio-Tagle \\& Mu\\~ noz-Tu\\~n\\'on 1998; Suchkov \\etal 1994; D'Ercole \\& Brighenti 1999;   Strickland \\& Stevens 2000;  Fujita \\etal 2003; 2004),  delayed cooling by an arbitrary time delay  (Gerritsen \\& Icke 1997; Thacker \\& Couchman 2000), or simply added outflowing gas directly on large scales (Navarro \\& White 1993; Mihos \\& Hernquist 1994; Scannapieco \\etal 2001; Springel \\& Hernquist 2003).  While outflows arise from all of these approximations, the approaches are so simplified that  several studies have suggested that supernovae and stellar winds may not be effective at driving outflows at all, but rather the primary driver may instead be radiation pressure on dust (Thompson \\etal 2005) or non-thermal pressure from cosmic rays (Socrates \\etal  2008).   On the other hand, Scannapieco \\& Br\\\" uggen (2010), showed that by using a subgrid model to track the unresolved velocities and length scales of turbulence driven by massive stars, one could produce galaxy-scale outflows while maintaining radiative cooling throughout the simulation.  Furthermore, the material ejected in these simulations was a multiphase distribution that grew out of the simultaneous turbulent heating and radiative cooling of the gas.   These results, although built on an approximate subgrid approach, illustrated that strong, stellar driven turbulence indeed had the potential to produce outflows with the properties of observed starbursting galaxies.  Here we examine this possibility more closely by conducting a series of direct numerical simulations that study the behavior of a stratified, turbulently supported medium, representing a small patch of a larger galaxy.    Previous theoretical work on stratified media at these scales  can be roughly grouped into two categories: (i)  two-and three dimensional studies that deposit the kinetic energy from a large number of supernovae within a single low-density superbubble (\\eg Tomisaka \\& Ikeuchi 1986; Mac Low \\etal 1989; Tenorio-Tagle \\etal 1990; Tomisaka 1998; Stil \\etal 2009); and (ii) three-dimensional studies of the evolution of stratified turbulence driven by the stochastic injection of thermal and/or kinetic energy from individual or small groupings of supernovae (\\eg Korpi \\etal 1999;  de Avillez \\& Breitschwert 2004, 2005, 2007, 2010; Dib \\etal 2006; Joung \\& Mac Low 2006; Cooper \\etal 2008; Shetty \\& Ostriker 2008; Joung \\etal 2009; Ostriker \\& Shetty 2011).  Here we adopt a third approach, abstracting the problem to study the long-term stability of a stratified medium that is continuously stirred at a fixed scale, and radiatively cooled such that the energy input from turbulent decay is matched to the overall cooling rate.   In this way our method is more like the ones adopted in studies of compressible turbulence in molecular clouds (\\eg  Stone \\etal 1998; Porter \\etal 1999; Padoan \\& Nordlund 1999; Cho \\& Lazarian 2003; Cho \\etal 2003; Padoan \\etal 2004; Kritsuk \\etal 2007; Benzi \\etal 2008; Lemaster \\& Stone 2008, 2009; Burkhart 2009), although unlike these studies we maintain vertical stratification of the medium, do not assume isothermal gas,  and include the atomic chemical reactions that are important in modeling the larger scale, higher temperature ISM.  While this more abstract approach limits the physical processes probed by our simulations, it nevertheless has three key advantages.  First, it sidesteps the many uncertain issues surrounding the coupling of massive stars with the surrounding medium.   In particular, the ISM simulations described above not only adopt many approximate techniques to deposit the energy from massive stars, but most of them also include an additional large source of diffuse heating, either taken from a purely numerical model that balances cooling in the initial configuration (\\eg de Avillez \\& Breitschwerdt 2004) or a model that approximates the photoelectric gas heating caused by the high-energy electrons that are removed from dust grains by stellar ultraviolet radiation (\\eg Joung \\& Mac Low 2006).   Our simulations, on the other hand, allow for cooling at all times in all locations, and they do not include any additional sources of energy beyond that of the turbulent stirring itself.  Secondly, our approach allows us to study the evolution of media with a wide variety of velocity dispersions without having to attempt to reproduce them by adjusting several indirect parameters such as diffuse heating rate, supernova rate, and the method used to deposit supernovae energy in the ISM.  Instead, as described below, we are able to increase the velocity dispersion simply by changing the mean density of the medium,  as this increases the overall average cooling rate, which, in equilibrium, is balanced by more vigorous turbulence.   Indeed, recent simulations that attempt to reproduce the turbulent ISM through direct supernova modeling find that the one-dimensional turbulent velocity dispersion changes only weakly with star formation rate, remaining $\\lesssim 15$ km/s over a wide range of parameters (Dib \\etal 2006; Shetty \\& Ostriker 2008; Agertz \\etal 2009; Joung \\etal 2009).  On the other hand, many galaxies are observed to have velocity dispersions exceeding 30 km/s, and in some cases even exceeding 100 km/s (F\\\"orster Schreiber \\etal 2006; 2009; Wright \\etal 2007; van Starkenburg 2008; Epinat \\etal 2008, 2009, 2010; Cresci 2009; Law \\etal 2009; Genzel \\etal 2008; 2011; Green \\etal 2010; Jones \\etal 2010; Lemoine-Busserolle \\& Lamareille 2010; Lemoine-Busserolle \\etal 2010).  Finally, and perhaps most importantly, our approach lets us separate effects arising from the direct deposition of hot gas onto the grid from those caused purely by the presence of widespread turbulence. This in turn allows us to identify for the first time a global instability that occurs in turbulently-supported media driven above a critical velocity dispersion.   The structure of this work is as follows.  In \\S 2, we describe our numerical simulations of stratified, turbulently-stirred, radiatively-cooled media. In \\S3 we present the results of these simulations, focusing on the physics of a transition that occurs at a critical turbulent velocity dispersion.    A  discussion is given in \\S4, which also explores the implications of our results for galaxy outflows across cosmic time. ",
        "conclusions": "Taken together, our simulations imply that while there is more to be learned about the details of gas escape from the overall gravitational potential, we can nevertheless be confident in our identification of a fundamental transition in stratified,  turbulently-stirred, radiatively-cooled media.    The question remains, however, if the transition we have found in our simplified models directly corresponds to a similar transition in the much more complex ISM distributions that occur in real galaxies.   Based on both theory and observations, we argue that there are many reasons to believe that it does. From a theoretical point of view, one can list a number of  ISM processes that we do not include in our simulations, yet none of them is likely to impede the turbulent runaway we see above  $\\sigma_{\\rm 1D, M} = 35$ km/s.  This runaway arises because any medium with sustained turbulence will be continuously heated by turbulent decay, and this thermal energy must be radiated away to maintain equilibrium.   This means that if turbulence is vigorous enough to move a fraction of the gas into a regime in which the cooling time increases strongly with temperature, the continued deposition of thermal energy will heat this gas dramatically, driving it out of the midplane, and in many cases out of the host galaxy itself.   The simplicity of this mechanism suggests that more complete ISM models will only refine the connection between this processes and galaxy outflows.  Thus while our nonrotating simulations do not include the Coreolis force, this is unlikely to play a key role.   In particular, its impact can be estimated by calculating the Rossby number, the ratio of the typical velocity to the shear across the region, which is greater than 1 in cases in which rotation is negligible. In a simulation offset from the center of the disk  by a distance $d_0,$ this is $R \\approx \\sigma_{\\rm 3D,T d_0} / (2H v_0),$ where $v_0$ is the rotational velocity at $d_0$.   For a galaxy with a flat rotation curve $v_0^2 = g_0 d_0$ where $g_0$ is the gravitational acceleration as in eq\\, (\\ref{eq:g}), and for solid body rotation, as in a dwarf galaxy  or near the center of a large galaxy, $v_0^2 = 3 g_0 d_0$.    This means $R \\approx (\\sigma_{\\rm 1D,T}/c_{\\rm s,grav})\u00ca(d_0/H)^{1/2}$ for a flat rotation curve, and $R \\approx (\\sigma_{\\rm 1D,T}/c_{\\rm s,grav}) (3d_0/H)^{1/2}$ for solid body rotation,  where $c_{\\rm s,grav} = (kT_{\\rm grav}/\\mu m_p)^{1/2} = 37 \\, {\\rm km/s} \\, (T_{\\rm grav}/10^5 K)^{1/2}.$   This means that $R \\approx 1$  if we were to consider our lowest velocity simulation to take place at the center of a galaxy, and $R \\gsim 1$ for all other simulations and choices of $d_0$.  Furthermore, the Coreolis force has no effect for velocities perpendicular to the disk, the direction of the gas ejection driven by turbulent runaway.  Rotation may be a mechanism worth exploring, but it is one that is unlikely to drastically effect the transition identified here.  Magnetic fields are also not included in our simulations, although they are known to exist in galaxies, and they can contribute to turbulence via the magentorotational instability (Piontek \\& Ostriker 2004; 2005), modify the  flow dynamics, and affect the decay of turbulent velocities into heat.  On the other hand,  simulations of supersonic magnetohydrodynamic turbulence (Stone \\etal 1998; Mac Low \\etal 1998) have shown that the energy density of the magnetic fields built up in the MHD turbulence are typically $\\approx 2-3$ smaller than the turbulent energy density, a result that is consistent with observations (Ferri\\' ere 2001).  Furthermore, Lemaster \\& Stone (2009) showed that large changes in magnetic field strength only have a minor impact on the dissipation rate of turbulence.   On the other hand, much like gravity, ordered fields have the potential to influence the flow near the plane of the galaxy,  even in the presence of large velocities and high temperatures.  Thus, magnetic fields are likely to lead to interesting changes in the simulations, but unlikely to affect the  dynamics strongly enough to avoid turbulent runaway at high velocity dispersions.  While heat conduction has the potential to affect turbulent heating by moving energy from hot rarified regions into the cooler clumps, the addition of this processes to our simulations is very  unlikely to have any effect on our conclusions. Although the precise value of the conductivity is relatively uncertain and dependent on magnetic field structure (\\eg Cowie \\& McKee 1977; Malyshkin \\& Kulsrud  2001), it is clear that it is a slow, second order process on ISM scales.   Instead,  mixing between hot and cold regions is primarily determined by the cascade of structures down to the very small scales at which conduction can operate (\\eg Pan \\& Scannapieco 2010).   In fact, if anything, as  turbulence is driven at marginally-resolved scales in our simulations, the presence of unavoidable numerical diffusion is likely to somewhat {\\em overestimate} heat conduction and mixing relative to the real ISM, as evidenced in the increased mixing in our fiducial run, S34, relative to the high resolution run, S35HR.  Our simulations also do not include self-gravity effects, which can both serve to enhance shear-driven turbulence and lead to the formation of Jeans unstable molecular clouds.  While the swing amplifier instability that occurs in self-gravitating rotating disks (Goldreich \\& Lynden-Benn 1965; Julian \\& Toomre 1966)  has been shown to be important in low density regions near the edges of disk galaxies (\\eg Wada, Meurer, \\& Norman 2002; Kim \\& Ostriker 2007; Koyama \\& Ostriker 2009), it is unlikely to drive turbulence at the levels considered here.  On the other hand, molecular clouds can drive more vigorous turbulence by gravitational scattering (Fukunaga \\& Tosa 1989; Gammie, Ostriker, \\& Jog 1991), and much more importantly, by the formation of stars.    For any complete model of ISM evolution, tracking the formation of such clouds is essential, which requires not only self-gravity, but a full model of the molecular chemistry and cooling.   On the other hand, as our goal here is not to capture star formation in detail, but rather to explore the evolution of stratified media as a function of turbulent velocity dispersion, this process is only of secondary importance.  Our aim is not to capture how turbulence is driven, but rather study how driven turbulence affects the subsequent evolution of the ISM.  Besides driving turbulence, star formation will also have two major effects  that are not considered in our simulations.   First it will lead to diffuse gas heating through the photoionization of neutral gas (\\eg Wolfire et al 2003), and the photoelectric gas heating caused by the high-energy electrons that are removed from dust grains by stellar ultraviolet radiation (\\eg Joung \\& Mac Low 2006).  To the extent to which such heating varies with space and time (Parravano \\etal 2003) or contributes to gas motions through radiation pressure (Thompson \\etal 2005) it will contribute to the overall level of turbulence.   To the extent to which such heating is uniform, however, it will offset the overall level of cooling, which in turn will mean that less turbulent driving is needed to maintain equilibrium.  In this case, a given turbulent velocity dispersion will correspond to a larger overall column depth, again making the point that $H \\bar n$ in our simulations is likely to be much less than measured in similar systems in nature, and that our results are much better interpreted as a function of $\\sigma_{\\rm 1D,M}.$  Secondly, supernovae and stellar winds arising from star formation will not only drive turbulence, but will also deposit hot gas directly into the interstellar medium.   In fact, the overlapping of such hot regions into a superbubble  is traditionally thought of as the mechanism by which outflows are driven (\\eg Tomisaka \\& Ikeuchi 1986; Mac Low \\etal 1989;  Tenorio-Tagle \\etal 1990; Mac Low \\& Ferrara 1999; Strickland \\& Stevens 2000;  Fujita \\etal 2003).   On the other hand, there is considerable observational evidence supporting the idea that galaxy outflows are  driven instead by the more widespread input from the massive stars (Strickland \\etal 2004).  For example, observed outflow pressure profiles demonstrate that mass and energy are ejected relatively smoothly over large regions  that match up well with region of intense star formation (Heckman \\etal 1990; Lehnert \\& Heckman 1996; Strickland \\& Stevens 2000; Strickland \\etal 2000).  This suggests that turbulent runaway may in fact be the primary mechanism producing the hot gas observed to be pouring out of starbursting galaxies, and that the fraction of hot gas directly added by SNe is not only uncertain theoretically, but of secondary importance physically.  In fact, adding this process to our simulations is only likely to obscure the main results.   The ISM in galaxies with velocity dispersions below $\\sigma_{\\rm 1D,M}$ would contain hot gas, but this gas would cool and condense within the disk relatively quickly. On the other hand, supernova material in $\\sigma_{\\rm 1D,M}$ disks would either be mixed into the ISM if it were deposited within a dense clump, or escape to high latitudes if it were ejected into a low-density, superheated region.   As seen in our more approximate galaxy-scale modeling in Scannapieco \\& Br\\\"uggen (2010), this would lead to outflows that were more metal enriched than the overall ISM, but this difference would be much less than in a ``superbubble\" outflow picture in which the metals from supernovae were injected into a single hot cavity.  However, as our simulations include energy input from stars as random turbulent motions its is not possible for us to associate metals from stars with particular parcels of gas in our present simulations to quantify this effect further.  Furthermore, the best way to implement supernovae into ISM scale simulations is extremely unclear, and to date has lead to turbulent velocity dispersions that vary only weakly with star formation rate and increase only up to $\\approx 15 $ km/s, (Dib \\etal 2006; Shetty \\& Ostriker 2008; Agertz \\etal 2009; Joung \\etal 2009), falling far short of the critical velocity dispersion identified here.  On the other hand, with the advent of modern integral-field unit (IFU) spectroscopy, galaxies with turbulent velocities above 35 km/s have been well observed at a variety of redshifts, with a variety of methods.  At cosmological distances, Lemoine-Busserolle \\etal (2010) and Lemoine-Busserolle \\& Lamareille (2010) conducted  Adaptive Optics (AO) IFU spectroscopy with the VLT/SINFONI instrument, using rest-frame optical nebular emission lines to measure average velocity dispersions of $60-70$ km/s in four $z \\approx 3$ galaxies and velocity dispersions of $35-100$ km/s in eight $z= 1.0- 1.5$ galaxies.  At Keck, Jones \\etal (2010) conducted AO OSIRIS IFU spectroscopic observations of six strongly-lensed $z = 1.7-3.1$ galaxies, at resolutions of up to 100 pc, obtaining H$\\alpha$ velocity dispersions between $\\sigma_{\\rm 1D, H\\alpha} = 50-100$ km/s.  Wright \\etal (2007) and Law \\etal (2009) used AO and the  OSIRIS IFU  to study a sample of unlensed  galaxies between $z=1.5$ and $3.1$ and found $\\sigma_{\\rm 1D, H\\alpha}$  between $50-100$ km/s in most cases, and $139 \\pm 3.2$ in an extreme starbursting region of a single galaxy (Q1623-BX543).  van Starkenburg \\etal (2008)  used SINFONI at the VLT to observe a $z=2.03$ galaxy with  $\\sigma_{\\rm 1D, H\\alpha}$ ranging between $30-100$ km/s, and Epinat \\etal (2009) used this instrument to observe nine  $z = 1.2 - 1.6$ galaxies with $\\sigma_{\\rm 1D, H\\alpha} \\approx 40 - 100$ km/s.  Finally, the large SINS survey  to study  the dynamics of high-redshift galaxies also used SINFONI IFU spectroscopy,  and to date has measured over twenty $z \\approx 2$ disk galaxies and star-forming clumps with $\\sigma_{\\rm 1D, H\\alpha} = 30-100$ km/s (F\\\" orster Schreiber \\etal 2006; 2009; Cresci 2009; Genzel \\etal 2008; 2011).  In the nearby universe, the Gassendi H$\\alpha$ survey of SPirals (GHASP) has used spatially-resolved Fabry-P\\' erot spectroscopy to observe a  sample of 153 isolated disk galaxies with $H\\alpha$ velocity dispersions  ranging from 20-40 km/s (Epinat \\etal 2008; 2010).  Selecting a sample of higher H$\\alpha$ luminosity disks from the Sloan Digital Sky Survey, Green \\etal (2010) used  the SPIRAL and WiFeS IFU spectrographs to measure luminosity-weighted  $H\\alpha$ velocity dispersions on $\\approx 2$ kpc scales in a sample of  65 $z \\approx 0.1$ rapidly star-forming galaxies.   These velocity dispersions ranged from $25-100$ km/s, and they were strongly correlated with star formation rate and uncorrelated with stellar mass, suggesting massive stars as the primary drivers of vigorous turbulence.  In observed star forming disk galaxies, the total pressure in the disk is just sufficient to stabilize all perturbations on scales too small to be stabilized by rotation, meaning that the Toomre parameter $Q = (\\sigma^2_{\\rm 1D}+ c_s^2)^{1/2} \\kappa/ \\pi G  \\Sigma_g \\approx 1-2,$ where $\\kappa$ is the epicyclic frequency and $\\Sigma_g$ is the gas density per unit area (Leroy \\etal 2008).  This means that the turbulent velocity dispersion is closely tied up with the gas surface density, which in turn is closely tied up with the star formation rate per unit area $\\dot \\Sigma_\\star.$  In fact Genzel \\etal (2011), compiled a large sample of the higher-redshift data described above to show that all galaxies with $\\sigma_{\\rm 1D,H\\alpha} \\geq 35$ km/s have high star formation rate densities $\\dot \\Sigma_\\star > 0.1 M_\\odot \\, {\\rm yr}^{-1} \\, {\\rm kpc}^{-2}.$   At the same time, outflows are ubiquitous in all  galaxies in which the star formation rate per unit area exceeds this same value of $0.1 M_\\odot \\, {\\rm yr}^{-1} \\, {\\rm kpc}^{-2}$ (Heckman 2002).  Indeed, the star forming nucleus in M82, the archetypal example of an outflowing starburst, has a star formation surface density of $\\approx 5 M_\\odot \\, {\\rm yr}^{-1} \\, {\\rm kpc}^{-2}$ (Strickland \\& Heckman 2009) and an average velocity dispersion of  $\\approx 90 \\pm 30$ km/s (Westmoquette \\etal 2009).  We may be just beginning to uncover the role of turbulent runaway in the history of galaxy formation."
    },
    "1112/1112.0569_arXiv.txt": {
        "abstract": "We show that the 5-GHz radio flux of transient ballistic jets in black hole binaries correlates with the dimensionless black hole spin parameter $a_*$ estimated via the continuum-fitting method. The data suggest that jet power scales either as the square of $a_*$ or the square of the angular velocity of the horizon $\\Omega_H$.  This is the first direct evidence that jets may be powered by black hole spin energy. The observed correlation validates the continuum-fitting method of measuring spin.  In addition, for those black holes that have well-sampled radio observations of ballistic jets, the correlation may be used to obtain rough estimates of their spins. ",
        "introduction": "Accreting black holes, both supermassive and stellar mass, are commonly observed to produce powerful relativistic jets \\citep{Zensus97, Mirabel_Rodriguez99}.  Although there now exists a wealth of data and many detailed models, the mechanism that powers these jets remains a mystery.  The popular idea that jets are powered by the black hole (BH) goes back to the work of \\citet{Penrose69}, who showed that a spinning BH has free energy.  \\citet{BZ77} proposed a plausible mechanism whereby this free energy could be used to power an astrophysical jet.  They suggested that magnetic fields in the vicinity of an accreting BH would be twisted as a result of the dragging of space-time by the rotating BH. The twisted field lines will carry away energy from the BH, producing an electromagnetic jet. The broad outlines of this model have been confirmed in numerical simulations (e.g. \\citealt{Tchekhovskoy+11}).  While a connection between BH spin energy and relativistic jets is theoretically very appealing, there has been no direct observational evidence for such a link. This is because, until recently, there was no BH with a believable measurement of the dimensionless spin parameter $a_* \\equiv cJ/GM^2$, where $M$ and $J$ are the mass and angular momentum of the BH.  The situation has now changed.  Methods are now available --- one in particular, the continuum-fitting (CF) method\\footnote{A second method, based on modeling the relativistically broadened X-ray iron K$\\alpha$ line, is not considered in this Letter (see Section 4 for a discussion).} \\citep{Zhang+97,Gierlinski+01,Shafee+06,Davis+06,McClintock+06} --- that have enabled us to make plausibly reliable measurements of $a_*$ for several stellar mass BHs. With this sample of spin measurements, we are now in a position to test whether jet power is related to BH spin.  Such a test is the goal of this Letter.  In Section 2, we describe our sample of stellar mass BHs and collect together the relevant observational data on BH spins and jet power. In Section 3, we plot radio power against BH spin and demonstrate that there is a significant correlation between the two quantities. We summarize and discuss in Section 4. ",
        "conclusions": "\\label{discussion}  Since the correlations shown in Figs~2 and 3 are based on only four objects, one wonders whether we are merely seeing chance alignment of intrinsically uncorrelated data.  The chief argument against this hypothesis is that $a_*$ varies over the full allowable range of prograde spins, $\\Omega_H$ varies by more than a factor of 10, and $P_{\\rm jet}$ varies by 2.4 orders of magnitude (or 3 orders of magnitude if one corrects for beaming assuming $\\gamma_{\\rm jet}=2$). Also, the plotted points differ from one another by several standard deviations, which is not statistically likely.  Therefore we conclude that the power of ballistic jets is most likely correlated with the spin of the accreting BH.  At the same time, and as a corollary, the strong apparent correlation validates the CF method of measuring spin.  The CF method is based on a number of assumptions, most of which have been independently validated (see \\citealt{McClintock+11}).  The results presented here provide yet another validation. Caveats to the above conclusions include the small size of the sample, insufficient data on one object (4U 1543-47), and uncertainties in how well jet power and radio luminosity track each other.  The existence of a correlation between jet power and BH spin does not necessarily mean that the energy source for the jet is BH spin.  The power could possibly be supplied by the accretion disc \\citep{Ghosh_Abramowicz97,Livio+99}.  Since the binding energy of a particle at the ISCO increases with increasing $a_*$, the disc power increases with BH spin and this might cause the observed correlation. However, we note that the radiative luminosity of a thin accretion disc varies by only a modest factor with BH spin; the radiative efficiency $\\eta=0.061$ for $a_*=0.12$ (the spin of A0620--00) and $\\eta = 0.23$ for $a_*=0.98$ (the spin of GRS 1915+105). If jet power scales similarly, and if radio luminosity is roughly proportional to jet power, we expect no more than a factor of 4 variation in the radio powers in our sample. Instead, we see a factor $\\sim250-1000$.  Moreover, the observed spread is rather close to what is expected theoretically if jets are powered by BH spin. The evidence thus suggests that ballistic jets are powered directly by the spin energy of the accreting BH.  Based on the above arguments, we view our results as a confirmation of the Penrose-Blandford-Znajek mechanism of powering relativistic jets by BH spin energy.  Theoretically, this mechanism depends on both the BH spin and the magnetic field strength at the horizon.  The latter is believed to depend on the mass accretion rate $\\dot{M}$ (e.g. \\citealt{Tchekhovskoy+11}).  Since ballistic jets are seen during a very specific phase of the evolution of a transient BHB, it is reasonable to assume that $\\dot{M}$ (normalized by the BH mass) is roughly the same in different objects when they exhibit ballistic jets, or indeed in different ballistic jet episodes in the same object.  This allows us to eliminate $\\dot{M}$ from our analysis and to treat ballistic jets as `standard candles', thereby making the comparisons shown in Figs~2 and 3 meaningful.  In addition to the CF method, a second method based on fitting the profile of the relativistically broadened Fe K$\\alpha$ line has been used to estimate BH spins \\citep{Reynolds_Nowak03,Miller07}.  In this Letter, we opt to use only CF spin data for two reasons: (1) The Fe-line models are complex and therefore relatively less reliable. CF spins are obtained rather simply by modeling a dominant thermal disc component, while Fe-line spins require modeling the thermal disc plus a Compton component plus a disc reflection component, which includes the Fe K$\\alpha$ line.  The Fe-line method furthermore requires characterizing a luminous corona of unknown geometry. (2) For several well-studied systems, the Fe-line method has generated widely inconsistent values of the spin parameter or shown to be strongly model dependent (e.g. for Cyg X-1, see Sec. 7.1 in \\citealt{Gou+11}; for MCG--6-30-15, see \\citealt{Miller+09}).  The CF method, on the other hand, gives consistent results for multiple and independent observations of individual sources.  For example, for the BHBs listed in Table 1 (excluding A0620--00), consistent results were obtained for $\\approx50$ {\\it RXTE} spectra (XTE J1550--564); 2 {\\it ASCA} and 31 {\\it RXTE} spectra (GRO J1655--40); 1 {\\it ASCA} and 5 {\\it RXTE} spectra (GRS 1915+105); and 34 {\\it RXTE} spectra (4U 1543--47).  The standout example is LMC X-3 with 411 spectra collected by eight X-ray missions over 26 years \\citep{Steiner+10}.  After considering separately both ballistic and hard state steady jets, \\citet{Fender+10} find no evidence for a correlation between jet power and BH spin.  We have already given in Section 2.2 a plausible reason for the absence of evidence in the case of the steady jets.  We now focus on rationalizing the very different results obtained by Fender et al. (presented in their Section 2.2.1 and fig.~6) and ourselves for the ballistic jets.  Our data sample (Table~1) is identical to their comparable sample (see the right panel of their fig.~6).  The only significant difference in data selection is that for GRO J1550--564 we use the new spin value of \\citet{Steiner+11}, $a_*=0.34 \\pm 0.24$, while they used the earlier \\citet{Davis+06} limit of $a_*<0.8$.  The substantial difference between our results and those of \\citet{Fender+10} is, in the end, determined by the choice of the quantity used to represent jet power.  We simply use the maximum observed flux density at 5 GHz expressed as a luminosity.  Fender et al.  compute jet power from the peak radio luminosity and the rise time of some particular synchrotron event.  The authors clearly state that their approach `is useful to provide lower limits on, and order-of-magnitude estimates of, jet power but is very susceptible to errors resulting from poor sampling of events, uncertainties in Doppler boosting, assumptions about equipartition, etc.'  Their estimates of jet luminosity for three sources are given in Table~1 of \\citet{Fender+04}, but it is not clear how the luminosities of A0620--00 and 4U~1543--47 were estimated.  The authors further adopt a formula relating jet power to X-ray luminosity, log$_{\\rm 10}L_{\\rm jet} = c~+~0.5($log$_{\\rm 10}L_{\\rm x}~-~34)$, and estimate the normalization constant $c$ in the preceding formula, which they treat as their proxy for jet power.  In short, their proxy for jet power is heavily model dependent and ours is model independent.  The correlation shown in Fig.~2 can be used to obtain rough estimates of spin for any transient BH that has undergone a major outburst cycle and that has been closely monitored at radio wavelengths. For instance, radio observations of Nova Muscae 1991 (GRS 1124--68) by \\citet{Ball+95} suggest a maximum 5-GHz radio flux $\\approx0.2$~Jy. Assuming a distance $D\\approx6$~kpc and a typical BH mass $M\\approx8M_\\odot$ \\citep{Ozel+10}, we obtain $\\log[P_{\\rm jet}]\\approx 0.65\\pm0.3$.  Fig.~2 then gives $a_* \\approx0.3-0.6$. In the case of GX 339-4, the brightest X-ray and radio outburst \\citep{Gallo+04} had a maximum 5-GHz flux of 0.055~Jy. Taking $D\\approx 9$~kpc, $M\\approx 8M_\\odot$ \\citep{Ozel+10}, we find $\\log[P_{\\rm jet}]\\approx 0.4\\pm0.3$ and $a_* \\approx 0.2-0.5$.  The latter estimate is consistent with the strict upper limit $a_*<0.9$ derived by \\citet{Kolehmainen_Done10} using the CF method with conservative assumptions.  These examples illustrate the importance of obtaining good radio coverage for all future transient BHBs, including especially the recurrent system 4U 1543--47.  Those systems that have CF-based spin measurements will flesh out the correlations plotted in Figs~2 and 3. For the many other BH transients that lack a sufficiently-bright optical counterpart and are therefore out of reach of the CF method, the radio data can either be used as a check on Fe-line measurements of spin or serve as our only estimate of spin.  Acknowledgements: The authors thank Jack Steiner for useful discussions.  This work was supported in part by NASA grants NNX11AE16G (to RN) and NNX11AD08G (to JEM)."
    },
    "1112/1112.2683_arXiv.txt": {
        "abstract": "We present constraints on the number of relativistic species from a joint analysis of cosmic microwave background (CMB) fluctuations and light element abundances (helium and deuterium) compared to big bang nucleosynthesis (BBN) predictions.  Our BBN calculations include updates of nuclear rates in light of recent experimental and theoretical information, with the most significant change occuring for the $d(p,\\gamma)^3\\mathrm{He}$ cross section.  We calculate a likelihood function for BBN theory and observations that accounts for both observational errors and nuclear rate uncertainties and can be easily embedded in cosmological parameter fitting. We then demonstrate that CMB and BBN are in good agreement, suggesting that the number of relativistic species did not change between the time of BBN and the time of recombination.  The level of agreement between BBN and CMB, as well as the agreement with the standard model of particle physics, depends somewhat on systematic differences among determinations of the primordial helium abundance. We demonstrate that interesting constraints can be derived combining only CMB and D/H observations with BBN theory, suggesting that an improved D/H constraint would be an extremely valuable probe of cosmology. ",
        "introduction": "The sensitivity of early-universe cosmology to the number and properties of light particle species was recognized early in the development of the hot big bang model \\cite{hayashi50,alpher53,peebles66,shvartsman69}.  The expansion rate of the universe at early times increases with number of relativistic particle species in thermal equilibrium, and this in turn sets timescales for big bang nucleosynthesis (BBN).  One can then use the BBN yields of light nuclei to constrain the number of light species quantitatively \\cite{steigman77}.  Constraints on the expansion rate are therefore often expressed in units of an effective number of light fermion species \\neff, so that for the standard model of particle physics $\\neff \\sim 3$, while $\\neff\\neq 3$ corresponds to some altered expansion rate that might be caused by additional neutrino species or by some other effect (like a nonzero neutrino chemical potential or violations of general relativity) \\footnote{Even within the standard model, $\\neff\\neq 3$ after the temperature drops below the electron rest mass.  At that time, all electron-positron pairs annihilate, transferring their energy mainly into the photon bath. A small fraction of annihilations produce neutrino pairs that fail to thermalize, contributing energy equivalent to 0.046 of a relativistic species to the subsequent evolution \\cite{mangano05}. In the CMB literature, \\neff\\ within the standard model is generally taken to be 3.046, but since pair annihilation occurs during BBN, most BBN work refers to the standard-model case as $\\neff=3$.  We report results in terms of \\neff\\ at late times, so that the standard model has $\\neff=3.046$.}.  Recent BBN-derived constraints may be found in Refs. \\cite{steigman10,hamann10,mangano11}.  The cosmic microwave background (CMB) is sensitive to the expansion rate much later: at the times of equal matter and radiation densities and of hydrogen recombination.  Recent measurements of the CMB power spectrum now allow interesting constraints on \\neff\\ at those times \\cite{wmap7,hamann11-1,keisler11,hou11}.  Constraints on light species from BBN and the CMB have been examined many times as both observational data and theoretical calculations have improved.  The consistency or not of these two types of constraints is an important question, particularly since they probe expansion rates at widely separated times in the evolution of the universe.  In the present work, we check for consistency between the two data sets, examine the sensitivity of that consistency to model assumptions, and then attempt to place the strongest constraint possible on \\neff\\ by combining the two types of data.  The BBN yields -- especially of deuterium -- depend not just on \\neff\\ but also on the cosmological baryon density, which in units of the critical density of the universe is $\\Omega_b\\equiv 8 \\pi G \\rho_B/3h^2$.  ($h$ is the Hubble constant in units of 100 $\\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$, $G$ is Newton's constant, and $\\rho_B$ is the baryon density expressed as mass per unit volume.) Prior to the precise inference of $\\obh$ from CMB measurements, the strongest constraint on \\neff\\ came from using the deuterium-to-hydrogen number ratio D/H to restrict $\\obh$ and then exploiting the \\neff\\ dependence of the primordial helium mass fraction \\yhe.  However, D/H has its own dependence on \\neff\\ \\cite{cardall96,steigman10}; a strong external constraint on $\\obh$ allows BBN constraints on \\neff\\ that are independent of \\yhe. We show below that, particularly because precise measurements of \\yhe\\ are difficult, the constraint on \\neff\\ from D/H is competitive with that from \\yhe.  In this paper, we focus on three issues: a full update of the BBN reaction rates including a revision to the $d(p,\\gamma)^3\\mathrm{He}$ rate based on recent theoretical calculations; an implementation of BBN constraints that includes both observational and theoretical uncertainties within CosmoMC \\cite{lewis02} (a popular Markov chain Monte Carlo fitting program used for cosmological data); and the utility of D/H as a probe of new physics that is independent of uncertainties in the primordial helium abundance. ",
        "conclusions": "In this paper, we have presented a comparison of determinations of the effective number of neutrino species, $\\neff$, from recent CMB measurements (WMAP7 and SPT) and BBN calculations, applying observed abundances of helium and deuterium in several ways.  Our results are similar to those of other analyses in the recent literature \\cite{archidiacono11,hamann11-1,gonzalez11,hamann11-2}, though comparison is difficult because of the use of different cosmological data and model spaces and the common practice of setting $Y_p=0.24$, which seems on the basis of Fig. \\ref{fig:bbn-variants} to bias \\neff\\ a bit high.  Like these and previous analyses, ours generally favor $\\neff > 3$.  The most directly comparable result in the literature combines BBN and CMB data with constraints on large-scale structure and the Hubble constant to find 95\\% confidence limits of $\\neff = 3.90^{+0.39}_{-0.56}$ \\cite{hamann11-1}.  This is similar to our results.  The BBN calculation used here includes a new estimate of a key cross section, $d(p,\\gamma)^3\\mathrm{He}$, based on detailed new theoretical calculations, which shifts deuterium yields by $\\sim 6\\%$ and lithium yields by $\\sim 10\\%$. This new cross section shifts \\obh\\ or \\yhe\\ inferred from D/H and slightly aggravates the large discrepancy between standard BBN and observed $^7\\mathrm{Li}$ abundances.  The BBN calculations and observed abundances have been embedded in a module that can be used in a widely used MCMC code (CosmoMC) and combined with CMB measurements or other cosmological data. The BBN likelihood calculation includes both theoretical and observational uncertainties, including the weak correlation between computed D/H and \\yhe.  We found that there is good agreement between BBN-based and CMB-based determinations of \\neff. These tests probe two very different epochs of the universe (a few minutes versus a few hundred thousand years), providing evidence that the same cosmological model describes both epochs.  However, both probes show hints of new physics in that the preferred number of effective neutrino species is above the standard model expectation. This is especially true in the case of the two higher \\yhe\\ determinations used (based on overlapping data), with as much as a $3.4 \\sigma$ excess.  Setting aside the helium constraints and combining the CMB with BBN theory and the observed D/H, the excess is $2\\sigma$, at $\\neff=3.9\\pm 0.44$. Improvements in the measurement of the D/H abundance would be extremely valuable as both a cosmological probe and an external check on \\yhe\\ determinations.  The results here underscore the complementarity of CMB measurements, BBN theory, and light element abundance measurements.  CMB measurements allow both an increased precision in BBN predictions (through a robust measurement of the baryon density) and a strong test for BBN (through comparisons with \\neff\\ and \\yhe\\ determinations). This is a consequence of the strong linkages in the standard cosmology between light element yields in the early universe and observable signatures at recombination, which connect the underlying cosmological model, the acoustic oscillations and Silk damping observed in the CMB, and the abundances observed in low-metallicity systems.  Precise measurements of all of these interlocking observables provide strong tests of the model and its parameters."
    },
    "1112/1112.2967_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "Intercalibration of heliospheric UV and EUV measurements requires a common basis of heliospheric ionization processes, which affect the distribuition of neutral interstellar gas in the heliosphere and thus influence both the spectrum and intensity distribution of the heliospheric backscatter glow. In this chapter we present a review of the solar factors affecting the distribution of neutral interstellar gas in the heliosphere, namely we discuss the radiation pressure, solar EUV ionizing radiation, and solar wind. We review the history of measurements of these factors and develop a model of time and heliolatitude evolution of solar wind speed and density based on data available from in situ measurements of the solar wind parameters, remote sensing interplanetary scintillation observations of the solar wind speed structure, and from remote sensing observations of the heliospheric Lyman-alpha backscatter glow. The results of this model are used as input in the global heliospheric models, discussed by Izmodenov et al. in another chapter of this book. The results of the global heliospheric modeling can in turn be used to fine tune the absolute calibration of heliospheric EUV measurements and, in a new iteration, to further refine the model of evolution of solar wind. The refined model will cover at least the two recent solar cycles and thus will provide a common homogeneous basis for interpretation of the present and past heliospheric experiments. It will also be used to interpret the observations of Energetic Neutral Atoms by the Interstellar Boundary Explorer (IBEX) \\citep{mccomas_etal:09a, mccomas_etal:09c}."
    },
    "1112/1112.0709_arXiv.txt": {
        "abstract": " ",
        "introduction": "The past three-year period has seen steady efforts to collect large numbers of radial-velocity (RV) measurements, as well as important applications of radial velocities to astrophysics.  Improvements in precision continue to be driven largely by exoplanet research. A workshop entitled ``Astronomy of Exoplanets with Precise Radial Velocities'' took place in August of 2010 at Penn State University (USA), and was attended by some 100 researchers from around the world. The meeting included thorough discussions of the current capabilities and future potential of the radial velocity technique, as well as data analysis algorithms to improve precision at visible and near-infrared wavelengths.  \\smallskip  While most of these discussions were focused on the search for and characterization of exoplanets, it is clear that more classical applications of radial velocities are also benefiting from the improvements, as evidenced by recent work on binary stars described herein. Below is a summary of other activity in the field of radial velocities during this triennium. Due to space limitations, we include only a selection of efforts and results in this area. ",
        "conclusions": ""
    },
    "1112/1112.0223_arXiv.txt": {
        "abstract": "We analyze the possibilities of detection of hypothetical exoplanets in coorbital motion from synthetic radial velocity (RV) signals, taking into account different types of stable planar configurations, orbital eccentricities and mass ratios. For each nominal solution corresponding to small-amplitude oscillations around the periodic solution, we generate a series of synthetic RV curves mimicking the stellar motion around the barycenter of the system. We then fit the data sets obtained assuming three possible different orbital architectures: (a) two planets in coorbital motion, (b) two planets in a 2/1 mean-motion resonance (MMR), and (c) a single planet. We compare the resulting residuals and the estimated orbital parameters.  For synthetic data sets covering only a few orbital periods, we find that the discrete radial velocity signal generated by a coorbital configuration could be easily confused with other configurations/systems, and in many cases the best orbital fit corresponds to either a single planet or two bodies in a 2/1 resonance. However, most of the incorrect identifications are associated to dynamically unstable solutions.  We also compare the orbital parameters obtained with two different fitting strategies: a simultaneous fit of two planets and a nested multi-Keplerian model. We find that, even for data sets covering over ten orbital periods, the nested models can yield incorrect orbital configurations (sometimes close to fictitious mean-motion resonances) that are nevertheless dynamically stable and with orbital eccentricities lower than the correct nominal solutions.  Finally, we discuss plausible mechanisms for the formation of coorbital configurations, by the interaction between two giant planets and an inner cavity in the gas disk. For equal mass planets, both Lagrangian and anti-Lagrangian configurations can be obtained from same initial condition depending on final time of integration. ",
        "introduction": "The possible existence of exoplanets in coorbital motion has fascinated planetary scientists for several years. Since the diversity of exoplanetary configurations continue to surprise us, even more than 15 years after the discovery of Peg51b, and it seems almost natural to expect Trojan planets to exist somewhere and the announcement of their discovery to be only a matter of time. Probably the first detailed analysis of hypothetical coorbital planets is due to Laughlin and Chambers (2002). They studied three types of coorbital configurations: tadpole orbits (around $L_4$ and $L_5$ equilateral points), horse-shoe configurations and ``eccentric resonances''. They proposed to distinguish between coorbital and single planet fits from RV data by observing residuals from long-term observations (more than ten orbital periods), because the coorbital configurations have large mutual interactions due to resonant motion.  For systems with more than one planet, it is well known that the existence of resonant motion may be possible evidence of a past large-scale planetary migration due to interactions with the gaseous disk. Although their importance is unquestionable, it is still intriguing why some commensurabilities are very populated (e.g., 2/1 MMR) and others that are currently empty (particularly the 1/1 MMR).  Lagrange (1873) discovered stable solutions for three massive bodies such that at all times their relative positions are located in the vertices of an equilateral triangle of variable size ($L_4$ and $L_5$ solutions). Linear stability analyses traditionally focused on the restricted three-body problem (where one of the masses vanishes, e.g Morais 2001, Namouni et al. 1999 and references therein). Recently, Nauenberg (2002) numerically investigated the dynamical stability of general three body problem as a function of the eccentricity of the orbits and the Rouths' mass parameter. The resonant solutions were found stable when $\\frac{m_1+m_2}{M_\\odot} \\lesssim 0.038$ and up to eccentricities $\\sim 0.6$. The existence of such periodic orbits in the general three-body problem also was derived for example in Siegel \\& Moser (1971) and Laughlin \\& Chambers (2002). As mentioned in Laughlin \\& Chambers (2002), horse-shoe orbits are stable for planetary masses up to $\\sim 0.4 m_{jup}$, and despite first impression, Keplerian fit are good enough to reveal their radial velocity signature with only four orbital periods. We note also that those type of orbits named ``eccentric resonances'', are simply an effect of angular momentum conservation in a coplanar three-body problem, that predicts when the eccentricity of one planet is maximum the other eccentricity reach a minimum and they have high amplitude of oscillation of resonant angles.  More recently Hadjidemetriou et al. (2009) studied the properties of motion close to a periodic orbit by computing the Poincar\\'e map on a surface of section. They constructed the symmetric families of stable and unstable motion describing one previous unknown symmetric configuration in the general problem, the Quasi-Satellite (QS) solution.  In a previous work (Giuppone et al. 2010), we studied the stability regions and families of periodic orbits of two planets in the vicinity of a 1/1 MMR, using numerical integrations and developing a semi-analytical method. We considered different ratios of planetary masses and orbital eccentricities, although we assumed that both planets share the same orbital plane (coplanar motion). As result we identified two separate regions of stability, symmetric and asymmetric types of motion easily described with the behavior of resonant angles $(\\sigma, \\Delta \\varpi) = (\\lambda_2 \\!-\\!\\lambda_1 , \\varpi_2-\\varpi_1)$, summarizing\\footnote{We recommend see the Figs. 2 and 5 from Giuppone et al. (2010) in order to easily identify these kind of coorbital motions in same plane of astrocentric parameters. We center this work on  $L_4$ and $AL_4$ configurations since, due to symmetry consideration, they give same results as $L_5$ and $AL_5$ respectively.}  \\begin{itemize} \\item \\textbf{Quasi-Satellite region:} correspond to oscillations around a fixed point located at $(\\sigma,\\Delta\\varpi) = (0,180^{\\circ})$ independent of mass ratio. Although not present for quasi-circular trajectories, they fill a considerable portion of the phase space in the case of moderate to high eccentricities.  \\item \\textbf{Lagrangian region:} Two distinct types of asymmetric periodic orbits exists in which both $\\sigma$ and $\\Delta\\varpi$ oscillate around values different from $0$ or $180^\\circ$. The first is the classical equilateral Lagrange solution associated to local maxima of the averaged Hamiltonian function. Independently of the mass ratio $m_2/m_1$ and their eccentricities, these solutions are always located at $(\\sigma,\\Delta\\varpi) = (\\pm 60^\\circ,\\pm 60^\\circ)$. However, the size of the stable domain decreases rapidly for increasing eccentricities, being practically negligible for $e_i > 0.7$.  The second type of asymmetric solutions correspond to local minimum of the averaged Hamiltonian function. They were named Anti-Lagrangian points ($AL_4$ and $AL_5$) and, for low eccentricities, are located at $(\\sigma,\\Delta\\varpi) = (\\pm 60^\\circ,\\mp 120^\\circ)$. Each is connected to the classical $L_i$ solution through the $\\sigma$-family of periodic orbits in the averaged system (a family of solutions where the angle $\\sigma$ has zero amplitude of oscillation, e.g. Michtchenko et al. 2008ab). Contrary to the classical equilateral Lagrange solution, their location in the plane $(\\sigma,\\Delta\\varpi)$ varies with the planetary mass ratio and eccentricities. Although their stability domain also shrinks for increasing values of $e_i$, they do so at a slower rate than the classical Lagrangian solutions, and are still appreciable for eccentricities are high as $\\sim 0.7$. An empirical relation exist for eccentricities below 0.6 (e.g. Giuppone et al. 2010, Hadjidemetriou \\& Voyatzis 2011) \\be e_1 \\simeq \\biggl( \\frac{m_2}{m_1} \\biggr) e_2. \\label{e1e2} \\ee \\end{itemize}  In this paper we address the question of extrasolar Trojans, both from its detectability and possible origin. In a recent work, Anglada-Escude et al. (2010) showed that two planets in a 2/1 MMR resonance could mimic a single planet in a more eccentric orbit. We wonder whether the opposite could also occur, and if two planets in a coorbital configuration could give similar RV signals as other dynamical systems including the 2/1 resonance. In Section 2 we review the radial velocity signal generated by such systems and how they can be easily mistaken by the signal generated by a single planet. In the following section we select representative configurations and generate synthetics data-sets for several mass ratio. Sections 4 analyzes the dynamical stability of the different possible configurations, while the dispersion of the best-fit parameters is discussed in Section 5. A comparison between two different strategies for planetary detection (two-planet fit vs. nested Keplerian fits) is presented in Section 6. Finally, a possible formation mechanism of extrasolar coorbitals is discussed in Section 7, where we present a series of numerical simulations (both hydro and N-body) of two planet systems immersed in a gaseous disk with an inner gap. Conclusions close the paper in Section 8. ",
        "conclusions": "In this paper we have analyzed the detectability and possible formation mechanism of hypothetical massive planets in stable coorbital configurations. So far there are no known extrasolar planetary systems containing coorbital bodies, which may imply that either these configurations are extremely difficult to form or to detect from RV surveys.  We have studied the detectability of three different types of coorbital motion (QS, $L_4/L_5$ and $AL_4/AL_5$), trying to evaluate possible bias in detections and identify what kind of compatible configuration could be detected. The analysis of Keplerian contributions to radial velocities allowed us to predict the value for the signal of one single planet that could be confused with coorbital configurations.  For even low values of errors in radial velocities measurements ($\\epsilon$ = 3 m/s) and observation time-spans covering four orbital periods (which include most of the presently detected exoplanets), coorbital configurations appear hard to identify: the results of the fitting process could easily confuse the RV data with that stemming from other configurations/systems (single planet or 2/1 MMR system). For large mass ratios, a correct identification of coorbital configuration is even more complicated and more easy to confuse with the other configurations. In all cases, the residuals of the different systems are comparable, even more so for large mass ratios (see Fig \\ref{rms-stat}).  Observation data sets covering longer time-spans allow to detect mutual perturbations, but once again the best fits are not always associated to coorbital motion. We have found several cases where other resonant configurations (i.e. 3/1 commensurability) actually give smaller residuals and better results.  Coorbital motion is not only sensitive to the data set, but also to the fitting procedure. Nested two-planet strategies may also confuse the real dynamics and yield results widely different from the nominal orbits; in this sense simultaneous two-planet fits seem more robust.  Transit observations should help distinguish coorbitals planets from other solutions.  Finally, we have found that stable coorbital systems with two massive planets may be formed from originally non-resonant orbits through their interaction with a inner cavity in the protoplanetary disk, as long as the surface density of the disk is sufficiently large. Both our N-body and hydro-simulations indicate a preference towards coorbitals with similar masses (i.e. $m_1 \\sim m_2$). In all our simulations with large mass ratios, the smaller planet was either pushed inside the cavity or trapped in a mean-motion commensurability outside the density jump."
    },
    "1112/1112.4770_arXiv.txt": {
        "abstract": "We report on the results of a spectroscopic survey of 130 comets that was conducted at McDonald observatory from 1980 through 2008. Some of the comets were observed on only one night, while others were observed repeatedly.  For 20 of these comets, no molecules were detected.  For the remaining 110 comets, some emission from CN, OH, NH, C$_{3}$, C$_{2}$, CH, and NH$_{2}$ molecules were observed on at least one occasion.  We converted the observed molecular column densities to production rates using a Haser (1957) model. We defined a restricted data set of comets that had at least 3 nights of observations.  The restricted data set consists of 59 comets. We used ratios of production rates to study the trends in the data. We find two classes of comets: typical and carbon-chain depleted comets.  Using a very strict definition of depleted comets, requiring C$_{2}$ \\underline{and} C$_{3}$ to both be depleted, we find 9\\% of our restricted data set comets to be depleted.  Using a more relaxed definition that requires only C$_{2}$ to be below a threshold (similar to other researchers), we find 25\\% of the comets are depleted. Two-thirds of the depleted comets are Jupiter Family comets, while one-third are Long Period comets. 37\\% of the Jupiter Family comets are depleted, while 18.5\\% of the Long Period comets are depleted. We compare our results with other studies and find good agreement. ",
        "introduction": "Since the time that the first spectra of comets were obtained, it has been remarked that all cometary spectra are very similar, with the most marked difference being the amount of continuum present.  This simple observation leads to the question of whether all comets share the same composition or if there are comets with fundamentally different compositions.  If differences are seen, then an additional question of the origin of the differences would be raised: are differences the result of different formational scenarios or different evolutionary scenarios?  These questions are not just curiosities.  We believe that comets are leftovers from the formation of the Solar System, that they have undergone only small changes since their time of formation, and that they represent the building blocks of the jovian planets.  Thus, understanding their chemistry is important for our understanding of the conditions at the start of the formation of the Solar System.  With these questions in mind, we began a program in 1980 to obtain spectra of comets using low resolution spectrographs at McDonald Observatory. This program continued through 2008.  We obtained observations of 132 different comets, some at more than one apparition.  This paper reports on the observations of 130 of those comets. 29P/Schwassmann-Wachmann~1 is left out because its almost-circular, 5\\,AU orbit and random outbursts make it an unusual comet that has been described in several previous papers (Cochran {\\it et al.} 1980, 1982, 1991a, 1991b; Cochran and White 1993). \\nocite{cobaco80,cocoba82,coco91sw1p1,cocobast91,cowh93} D/Shoemaker-Levy 9 broke into many pieces prior to impacting Jupiter in 1994. Observations of the pieces had to be obtained in a different manner from the other comets and these are detailed in Cochran {\\it et al.} (1994). \\nocite{cosl994}  Some of the comets described in this present paper were already discussed in prior papers (e.g. Cochran {\\it et al.} 1987, 1989, 1992, 2009; Cochran and Barker 1987).  They are included again for completeness so that all the observations can be found in the same place.  In addition, all observations were analyzed in a consistent manner in this paper. Thus, this paper supersedes its predecessors. \\nocite{coba87gz,cobarast92,cobacari09,cogrba89,co87corr}  This is not the only survey of its kind to be carried out.  A'Hearn et al. (1995, hereafter AH95) reported on observations of photometry of 85 comets. Fink (2009, hereafter F09) presented spectroscopic observations of 92 comets. Langlund-Shula and Smith (2011, hereafter LS11) observed 26 comets.  We will compare our results with these studies near the end of this paper. \\nocite{ahetal95,fi09,lasm11} ",
        "conclusions": "When we began this survey in 1980, we fully expected that we would find some comets with very different spectra than the vast majority of comets.  Instead, we found a mostly homogeneous group of objects with subtle compositional differences.  In particular, we found:  \\begin{itemize} \\item The vast majority of comets have remarkably similar compositions.  \\item There exists a group of comets which are depleted in the carbon-chain molecules.  Using a very strict definition of carbon-chain depletion requiring depletion in both C$_{2}$ and C$_{3}$, we find that 9\\% of the comets are depleted.  Using a more relaxed definition that requires only the C$_{2}$ be depleted (similar to other authors), we find 25\\% of the comets are depleted.  This is consistent with other authors.  \\item For the most part, we agree with other studies on which individual comets show depletions.  \\item Carbon-chain depleted comets can be of any dynamical type (though we did not observe any HTCs with depletion our HTC sample is very small).  \\item Using the relaxed definition, two thirds of the depleted comets are JFCs and one third are LPCs (with the stricter definition it is 80\\% and 20\\% respectively).  \\item Depleted comets make up 37\\% of the JFCs and 18.5\\% of the LPCs.  \\item On average, there are few differences in composition with increasing current orbit semi-major axis.  However, C$_{2}$ and NH may show significant trends  of lower C$_{2}$ and higher NH for JFCs than for LPCs.  \\end{itemize}  \\vspace*{0.75in} \\begin{center}Acknowledgments\\end{center} This research was supported by NASA Grants NNG04G162G, NNX08A052G and predecessor grants.  McDonald Observatory is operated by The University of Texas at Austin.  We thank David Schleicher for performing a cluster analysis for us that, unfortunately, did not clarify the groupings. We thank Dr. Michael A'Hearn and Dr. Uwe Fink for very helpful suggestions to clarify the paper and for identifying some errors."
    },
    "1112/1112.2589_arXiv.txt": {
        "abstract": "Of the approximately 400 known Galactic classical novae, only ten of them, the recurrent novae, have been seen to erupt more than once.  At least eight of these recurrents are known to harbor evolved secondary stars, rather than the main sequence secondaries typical in classical novae.  In this paper, we propose a new nova classification system, based solely on the evolutionary state of the secondary, and not (like the current schemes) based on the properties of the outbursts.  Using archival optical and near-infrared photometric observations of a sample of thirty eight quiescent Galactic novae we show that the evolutionary state of the secondary star in a quiescent system can predicted and several objects are identified for follow-up observations; CI~Aql, V2487~Oph, DI~Lac and EU~Sct. ",
        "introduction": "\\label{sec1}  A classical nova (CN) outburst occurs in an interacting binary system comprising a white dwarf (WD; the primary) and typically a late-type main-sequence star (the secondary) that fills its Roche Lobe \\citep{1956ApJ...123...44C}.   Matter is transferred, usually via an accretion disk around the WD, from the secondary to the WD surface, becoming degenerate in the process. Once extensive CNO cycle nuclear burning commences under these degenerate conditions, a thermonuclear runaway occurs ejecting $\\sim 10^{-5} - 10^{-4} \\msun$ of matter from the system with velocities from hundreds to thousands of \\kms\\ \\citep[see, e.g.][and references therein]{2008clno.book.....B}. These systems generally have orbital periods of a few hours \\citep{2008clno.conf....1W} and recurrence times of a few $\\times\\;10^{3}-10^{6}$ years are often quoted \\citep{1986ApJ...308..721T}.  CNe are traditionally sub-divided using solely their eruptive properties, typically speed class \\citep{1957gano.book.....G}, or spectral class \\citep{1992AJ....104..725W}.  More recently, novae have also been classified as ``disk'' or ``bulge'' novae, on the basis of their parent stellar population and location within the host galaxy \\citep{1990LNP...369...34D,1992A&A...266..232D,1998ApJ...506..818D}.  Closely related to these systems are the recurrent novae (RNe), which, unlike CNe, are observed to recur on time-scales of order tens to a hundred years. However, it should be noted that the upper limit to this range is likely to be an observational effect based on the paucity of historical records. Such a short recurrence time-scale is usually attributed to a high mass WD, probably close to the Chandrasekhar limit, together with a high accretion rate \\citep{1985ApJ...291..136S,2005ApJ...623..398Y}.  The orbital periods of RNe are less homogeneous than those of CNe, with periods ranging from a few hours to several hundreds of days \\citep{2008ASPC..401...31A}.  Recurrent nova systems have been sub-divided into three general sub-classes \\citep{2008ASPC..401...31A} via a combination of the properties of the eruption and via the properties of the progenitor system whilst at quiescence:  {\\it RS~Oph/T~CrB systems} (hereafter the RS~Oph-class) are observed to contain red giant secondaries and hence have much longer orbital periods ($\\sim$ a year), and typically smaller outburst amplitudes than CNe.  Their outbursts exhibit rapid declines from maximum, with large ejection velocities ($\\ge 4000\\;\\mathrm{km\\;s}^{-1}$).  These systems show evidence of an interaction between the ejecta and material from the pre-existing red giant wind.  The ejected mass from these systems is typically two orders of magnitude less than that observed from CN systems \\citep[see papers within][]{2008ASPC..401.....E}.  {\\it U~Sco systems} are observed to contain evolved main sequence or sub-giant secondaries, with orbital periods closer to those of CNe (hours up to of order a day).  They again exhibit rapid declines, being amongst the fastest declining novae observed, with very high ejection velocities (up to $10000\\;\\mathrm{km\\;s}^{-1}$).  Their post-outburst spectra resemble those of the He/N sub-class of CNe \\citep{1992AJ....104..725W}.  These systems eject a similar mass of material as RS Oph systems.  The {\\it T~Pyx systems} (also CI~Aql and IM~Nor) are much more akin to CNe.  They exhibit short orbital periods and spectroscopically resemble Fe~II CNe \\citep{1992AJ....104..725W}.  Their optical decline rate classifies them as moderately fast or slow novae.  The ejected mass in these systems is consistent with the range observed in CNe ($\\mathrm{M}_{\\mathrm{ej}}\\sim 10^{-5}\\msun$).  These systems are generally only distinguishable from CNe due to their shorter recurrence times, not by the properties of the progenitor system or the outburst.  The short inter-outburst time observed in RNe is likely due to some combination of a higher mass WD and an accretion rate greater than the typical for CN systems.  Indeed, both RS~Oph and U~Sco (amongst others from these classes) appear to have WDs close to the Chandrasekhar limit.  A number of authors \\cite[see e.g.][]{2007ApJ...659L.153H,2011ApJ...727..124O,SumnerConf} have indicated that the WD mass may be increasing over time and these systems have been proposed as a Type Ia supernova (SN) progenitor candidate \\citep[see e.g.][]{2008ASPC..401..150K}.  However, \\citet{2011arXiv1107.4013M} reported that U~Sco could contain an O-Ne WD, rather than a C-O WD, and hence may not evolve in to a Type Ia SN as is widely predicted for this system.  Any CN that has been observed to erupt more than once is re-classified as a recurrent nova.  This approach has given us ten Galactic RNe amongst $\\sim400$ known Galactic CNe from an underlying rate of $34^{+15}_{-12}$~year$^{-1}$ \\citep{2006MNRAS.369..257D}.  Due to observational effects - which typically worsen the further back in time one looks - it is highly likely that a considerable proportion of these Galactic CNe are in fact (short time-scale) recurrents for which only a single outburst has been observed.  Either a previous outburst has been missed or the inter-outburst period is particularly long.  So it would be seem advantageous to better develop a nova classification system that relies more on the fundamental properties of the systems than on observational selection effects.  Over the past few years observations of a number of Galactic ``transients'' have shown properties consistent with being RNe, although they had never been seen in outburst previously.  For example, the nova V2487~Ophiuchi (Nova Oph 1998) was observed to have a rapid optical decline and plateau phase \\citep{2002ASPC..261..629H,2002AIPC..637..381H}, suggestive of a RN.  Consultation of the Harvard College Observatory archival photographic collection revealed a previously unknown outburst on 1900 June 20 \\citep{2009AJ....138.1230P}.  Additionally, the novae V1721~Aql \\citep{2011arXiv1104.3068H}, V2491~Cyg \\citep{2011arXiv1104.3482D}, KT~Eri \\citep{RibKTTEri} and V2672~Oph \\citep{2010MNRAS.tmp.1484M} and also the extragalactic nova M31N~2007-12b \\citep{2009ApJ...705.1056B} have been shown to harbor secondary stars akin to either the U~Sco or RS~Oph systems, albeit with only one recorded outburst.  In this paper we will attempt to simplify the nomenclature somewhat by only classifying each system using the evolutionary state of the secondary; a main sequence star (MS-Nova), a sub-giant star (SG-Nova), or a red-giant branch star (RG-Nova).  We introduce such classifications in an attempt to avoid the accretion rate / WD mass / recurrence time degeneracy.  All the known U~Sco-class RNe would be placed into the SG-Nova group, all the known RS~Oph-class RNe into the RG-Nova group, the T~Pyx-class RNe and Classical Novae will populate the MS-Nova group.  By virtue of the geometry of these systems, the above classifications would effectively be the same as sub-dividing novae by orbital period, if one assumes that the secondaries fill their Roche lobes.  The MS-Novae would have orbital periods of order hours, the SG-Novae of order a day, and the RG-Novae of order a year.  Indeed \\citet{1995CAS....28.....W} indicates that systems with orbital periods longer than eight hours should contain evolved secondaries.  In this paper we present a simple method which will allow differentiation of RG-Novae and SG-Novae from the MS-Novae population based solely on their quiescent (inter-outburst) broad-band optical and near-IR (NIR) properties.  The format of the remainder of this paper/letter is as follows: In Section~\\ref{data} we briefly describe our method and present the data used to illustrate this technique; in Section~\\ref{results} we present our results, in Section~\\ref{discussion} we discuss our findings and present a number of predictions and finally, in Section~\\ref{conc} we summarize our conclusions. ",
        "conclusions": "  \\begin{itemize} \\item We have proposed a new classical nova classification system, based solely on the evolutionary state of the secondary star; a main sequence star (MS-Nova), a sub-giant star (SG-Nova) or a red-giant branch star (RG-Nova). \\item The position of a quiescent novae on optical or NIR color-magnitude diagrams is a strong function of both the evolutionary state of the secondary star and the luminosity of any accretion disk.  The more evolved the secondary star, the weaker the position dependency is on the accretion disk. \\item A standard ($V$ versus $(B-V)$) color-magnitude diagram is useful for differentiating between the MS-Novae (which all lie bluer than the main sequence) and the SG-/RG-Novae (which all lie on, or are redder, than the main sequence).  Although this plot can be strongly sensitive to the accretion rate. \\item Additionally NIR color-magnitude diagrams may be used to differentiate between all three sub-classes of secondaries. \\end{itemize}  We also make a number of predictions and observations about systems included in this paper:  \\begin{itemize} \\item The recurrent nova CI~Aql - often grouped with T~Pyx - is in fact a SG-Nova, a system more akin to U~Sco. \\item The recurrent nova V2487~Oph is a system similar to KT~Eri, containing a low-luminosity red giant branch star.  This contradicts the short orbital period reported in \\citet{2010ApJS..187..275S} and we recommend follow-up observations to directly determine the orbital period.  Additionally, we cannot rule out the possibility that the secondary in this system is a horizontal branch star. \\item Based on the high quiescent luminosity and long orbital period, we predict that the DI~Lac system may contain a sub-giant secondary star. \\item In line with \\citet{1994MNRAS.266..761W} and \\citet{2003MNRAS.340.1011W}, we propose that EU~Sct contains a luminous red-giant secondary. \\end{itemize}  Follow-up (quiescent) spectroscopic observations of novae, such as those initially reported in \\citet{MomayConf}, can greatly aid in the classification of the secondary stars in these systems."
    },
    "1112/1112.5226_arXiv.txt": {
        "abstract": "The magnetic flux that is generated by dynamo processes inside the Sun emerges in the form of bipolar magnetic regions. The properties of these directly observable signatures of the dynamo can be extracted from full-disk solar magnetograms. The most homogeneous, high-quality synoptic data set of solar magnetograms has been obtained with the MDI instrument on the SOHO spacecraft during 1995-2011. We have developed an IDL program which has, when applied to the 73,838 magnetograms of the MDI data set, automatically identified 160,079  bipolar magnetic regions that span a range of scale sizes across nearly four orders of magnitude. The properties of each region have been extracted and statistically analysed, in particular with respect to the polarity orientations of the bipolar regions, including their tilt angle distributions and their violations of Hale's polarity law. The latitude variation of the average tilt angles (with respect to the E-W direction), which is known as Joy's law, is found to closely follow the relation $32.1^\\circ \\times \\sin$(latitude). There is no indication of a dependence on region size that one may expect if the tilts were produced by the Coriolis force during the buoyant rise of flux loops from the tachocline region. A few percent of all regions have orientations that violate Hale's polarity law. We show explicit examples, from different phases of the solar cycle, where well defined medium-size bipolar regions with opposite polarity orientations occur side by side in the same latitude zone in the same magnetogram. Such oppositely oriented large bipolar regions cannot be part of the same toroidal flux system, but different flux systems must coexist at any given time in the same latitude zones. These examples are incompatible with the paradigm of coherent, subsurface toroidal flux ropes as the source of sunspots, and instead show that fluctuations must play a major role at all scales for the turbulent dynamo. To confirm the profound role of fluctuations at large scales we show explicit examples where large bipolar regions differ from the average Joy's law orientation by an amount between $90^\\circ$ and $100^\\circ$. We see no observational support for a separation of scales or a division between a global and a local dynamo, since also the smallest scales in our sample retain a non-random component that significantly contributes to the accumulated emergence of a north-south dipole moment that will lead to the replacement of the old global poloidal field with a new one that has the opposite orientation. ",
        "introduction": "\\label{sec:intro} The origin of solar activity with its 11 yr cycle is generally understood in terms of an oscillatory dynamo inside the Sun \\citep[cf.][]{stenflo-brandsub05,stenflo-brandenburg05}. While a toroidal magnetic field is generated from a poloidal field through differential rotation, the essence of the global dynamo is the regeneration of the poloidal field from the toroidal one through interaction with turbulent convection in the rotating Sun. The systematic twisting due to the Coriolis force generates an N-S dipole moment from the originally E-W oriented toroidal field. This dipole moment spreads by turbulent diffusion to reverse the old global poloidal field and replace it with a new one of opposite sign.  This general scenario provides a framework for the explanation of various observed properties of the cycle: (1) Hale's polarity law \\citep{stenflo-haleetal19}, according to which sunspots appear in pairs approximately aligned with the E-W direction, but with opposite polarity orientations in the two hemispheres, and with reversal of the orientations between two successive 11-yr cycles. (2) The butterfly diagram, according to which the latitude zones, in which the sunspots appear, migrate from high ($\\sim 40^\\circ$) to low (near the equator) latitudes as the cycle progresses. (3) Joy's law \\citep[also reported in the seminal paper by][]{stenflo-haleetal19}, according to which the orientations of the magnetic polarities of the sunspot groups deviate from the E-W direction in a systematic way, such that the preceding (W) part of the region is on average closer to the equator than the following (E) part.  While much of the empirical studies of the solar cycle have been in terms of the properties and distributions of sunspots, we know that sunspots only represent proxies for the underlying fundamental agent, the magnetic field. The solar dynamo operates in the Sun's interior, which is not accessible to direct observations (except indirectly through helioseismology). At the surface of the Sun the main directly observable signatures of the dynamo are the bipolar magnetic regions, which represent magnetic flux that has emerged from the interior, and the large-scale flux patterns that are shaped by the accumulated patterns of flux emergence combined with the action of flux dispersion and other transport processes (like meridional circulation).  Bipolar magnetic regions occur at all scales with a size spectrum that follows a power law \\citep{stenflo-harveyphd93,stenflo-harveyzwaan93,stenflo-schrharv94,stenflo-parnelletal09}, from the largest regions that harbor major sunspots often accompanied by violent flare activity, to ephemeral active regions \\citep{stenflo-harveymartin73,stenflo-harveyetal75,stenflo-martinharvey79} without sunspots, and the still smaller intranetwork fields. The global contribution to the overall flux emergence rate dramatically increases as we go down in scale size \\citep{stenflo-zirin87}, suggesting that the global flux balance is dominated by the smallest scales, in contrast to the general visual impression that one may get from magnetograms.  Also the smaller scales seem to statistically follow Hale's polarity law and the butterfly diagram distribution, although with increasing statistical spread as we go down in scale. It is not yet known at what scales the polarity orientations become so randomized that the accumulated effect of still smaller scales no longer contribute significantly to the operation of the global dynamo. Magnetoconvection is expected to cause magnetic structuring down to scales of order 10-100\\,m, where the magnetic Reynolds number becomes unity so that the magnetic field ceases to be frozen-in and decouples from the turbulent plasma \\citep{stenflo-dewijn09}.  The observational signature of the dynamo mechanism that most directly represents the regeneration of the poloidal field from the toroidal one is the systematic tilt of the bipolar magnetic regions that is statistically described by Joy's law. Combined with Hale's polarity law this tilt describes how the emerging bipolar regions bring to the surface an N-S bipolar moment that is the seed for the regeneration of the new global poloidal field of reversed polarity, from which the subsequent 11 yr cycle of solar activity is generated. The observed properties of the tilt angles therefore give guidance to the dynamo theories and constrain the ways in which the dynamo is allowed to operate.  The tilt angles of sunspot pairs have been studied in many papers since the discovery by \\citet{stenflo-haleetal19} of Joy's law, using Kitt Peak magnetograms \\citep{stenflo-wangsheeley89}, Mount Wilson magnetograms \\citep{stenflo-howard91a}, and sunspot group data \\citep{stenflo-howard91b}, confirming the overall magnitude and crude latitude dependence originally found by \\citet{stenflo-haleetal19}. In addition \\citet{stenflo-wangsheeley89} found that 4\\,\\%\\ of the bipolar magnetic regions had inverse polarity orientations, and that the spread in the distribution of tilt angles increased significantly as one goes to smaller regions. The superb quality of the SOHO/MDI data set will allow us in the present work to explore these properties with much better precision and in much greater detail than has been possible before.  Recently \\citet{stenflo-ks08} used the set of MDI magnetograms to study the tilt changes in emerging bipolar magnetic regions. They found that during the first few days after emergence the tilt angles relaxed towards the value expected from Joy's law and not towards the E-W orientation, in agreement with the findings of \\citet{stenflo-sivaramanetal07} from an analysis of Kodaikanal white-light images. \\citet{stenflo-ks08} further found no dependence of the tilt behavior on the amount of flux or size of the bipolar regions. Both these findings contradict the paradigm that the tilt is caused by the Coriolis force acting on initially untilted flux loops that rise from a toroidal source region near the bottom of the convection zone and emerge at the surface as tilted bipolar regions \\citep{stenflo-dsilva93,stenflo-fisheretal95}. Instead the tilt, which is the source of the N-S dipole moment that leads to the reversal and regeneration of the poloidal field, appears to have been established already in the dynamo region in the Sun's interior. The tilt observed at the surface reflects this property regardless of the size or amount of flux of the observed regions. ",
        "conclusions": "\\label{sec:concl} The bipolar magnetic regions represent the most conspicuous directly observable signatures of the dynamo that operates in the Sun's interior. While their general E-W polarity orientations indicate that they represent amplified largely toroidal flux, their systematic tilt shows that the emerging flux also brings to the surface an N-S dipole moment that represents the seed of the new poloidal field that will replace the old one, with reversed orientation. The tilt angles thus represent explicit signatures of the dynamo process that is responsible for the regeneration of the poloidal field from the toroidal one.  A prevailing paradigm for the origin of the tilt angles has been in terms of the Coriolis force acting on flux loops that buoyantly rise from the tachocline region at the bottom of the convection zone, where the dynamo is assumed to operate \\citep{stenflo-dsilva93,stenflo-fanetal94,stenflo-fisheretal95}. This paradigm requires superstrong (60-100\\,kG) toroidal magnetic fields without appreciable tilt in the tachocline zone, much stronger than the value for equipartition with the convective motions. As the speed of the buoyant rise depends both on the amount of flux and the strength of the field, one would expect that the observed tilt at the surface should depend on the size of the bipolar regions.  Work by \\citet{stenflo-sivaramanetal07} and \\citet{stenflo-ks08} has however shown that this paradigm is untenable, since the tilt angles are found to relax after emergence not towards the E-W orientation but towards the angle prescribed by Joy's law, and this behavior is found to be independent of size or amount of flux of the regions. This is evidence that the tilt (or N-S dipole moment) is already established in the source region inside the Sun and not during the buoyant rise of the flux loops. The present work supports this conclusion by showing, with much enhanced statistics, that the average tilt does not change appreciably as we go to regions with orders of magnitude smaller flux contents.  An even much more ingrained, long-term paradigm that is shown to be untenable by the present work is the phenomenological scenario, according to which differential rotation creates from a poloidal field a coherent toroidal flux system, from which the sunspots arise. Such a scenario is not compatible with the violations of Hale's polarity law that we have presented here. Our illustrated examples where well defined medium-size bipolar magnetic regions occur side by side {\\it in the same magnetogram and the same latitude zone} unambiguously show that these oppositely oriented regions cannot be part of the same flux system, but that there must be a coexistence of oppositely oriented toroidal flux in the same latitude zones at any given time.  Although the occurrencies of violations of Hale's polarity law represent only a few percent of all cases, they rule out the possibility of well defined, coherent toroidal flux systems as a source of all active regions, even the large ones. Our results make clear that {\\it fluctuations} represent an essential inherent physical property of the solar dynamo, as expected to various degrees from turbulent dynamo theory \\citep{stenflo-brandsub05}.  The fundamental role of large fluctuations at all scales is further illustrated in Fig.~\\ref{fig:90deg} for four selected cases of well defined bipolar magnetic regions, including large ones, which have orientations that differ by $90^\\circ$ - $100^\\circ$ from the orientations prescribed by Joy's law. The existence of many such cases shows that we have fluctuations over the whole range of orientation angles, not only between ``Hale and anti-Hale'' (proper and reversed) orientations. Also the largest bipolar regions are subject to these fluctuations.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig9_stenflo.eps}} \\caption{In contrast to Fig.~\\ref{fig:violate}, which showed examples of regions with orientations differing by nearly $180^\\circ$ from the orientations expected from Hale's and Joy's laws, we here show four examples where the orientation differs by an amount between $90^\\circ$ and $100^\\circ$ (marked by the enclosing rectangular boxes). The location of the solar equator is marked by the dashed line. The grey-scale cuts are set at $-50$ (dark) and $+50$\\,G (white). }\\label{fig:90deg} \\end{figure}  Our results showing the profound role of fluctuations together with the finding that the average tilt angles described by Joy's law are independent of the amount of flux of the regions, over several orders of magntiude in flux, cast doubt on the validity of tachocline dynamo theories, and seem to be more compatible with distributed dynamo models \\citep[cf.][]{stenflo-brandenburg05}.  Other concepts that are not supported by our analysis is the separation of scales and the separation between a global and a local dynamo. Previous work has shown that the size spectrum of magnetic regions follows a power law indicating scale invariance over the whole range of resolved scales \\citep[cf.][]{stenflo-parnelletal09}. Our finding that the average tilt angle is independent of region size shows that also the smallest regions in our sample contribute significantly to the N-S dipole moment that leads to the reversal of the global dipole field. There is no observational evidence that below a certain scale size the fields no more contribute to the global dynamo. On the contrary the accumulated global effect of the small regions may even be the dominating one, because it has long been well known that the smallest regions dominate the global flux emergence rate \\citep{stenflo-zirin87}. Since our analysis indicates that even the smallest regions, although increasingly randomized in their orientations, retain a non-random component, their net global effect may be substantial. Although it is beyond the scope of the present paper to quantify the relative global contributions of the various scales, we feel that the artificial separation of scales and division between a global and local dynamo are theoretical idealizations that are not justified by observations of the Sun."
    },
    "1112/1112.5821_arXiv.txt": {
        "abstract": "We observed the first unidentified TeV $\\gamma$-ray source TeV J2032$+$4130 with Suzaku. Owing to Suzaku's high sensitivity for detection of diffuse X-ray emission, we found two small structures in the TeV emitting region. One of them is coincident with a $\\gamma$-ray pulsar PSR~J2032$+$4127, which was discovered by the Fermi Gamma-ray Space Telescope. By subtracting contribution of point sources estimated by Chandra data, we obtained diffuse X-ray spectrum. The X-ray spectrum can be reproduced by a power-law model with a photon index of $\\sim$~2, and an X-ray flux of $2\\times10^{-13}$ erg~s$^{-1}$~cm$^{-2}$. The ratio of the $\\gamma$-ray flux to the X-ray flux is about 10. If the origin of the TeV $\\gamma$-ray is inverse Compton scattering of microwave background by high energy electrons, the ratio corresponds to the magnetic field strength of $\\sim$~1~$\\mu$G. However, the smaller size of the X-ray emission than that of the TeV emission suggests that energy loss of the electrons can explain the large ratio of the $\\gamma$-ray flux with a reasonable magnetic field strength of a few $\\mu$G. ",
        "introduction": "The stereoscopic technique of atmospheric Cerenkov telescopes improved the angular resolution for detecting TeV $\\gamma$-rays and thus increased the number of TeV $\\gamma$-ray sources. Some of new TeV objects have no counterparts at other wavelengths and are called unidentified TeV $\\gamma$-ray objects. These objects provide key information in the investigation of the origin of high-energy cosmic rays. Multi-wavelength observations of these objects are very important in elucidating the emission mechanism and successful identification of them.  TeV J2032+4130 was the first unidentified TeV $\\gamma$-ray source discovered by HEGRA (\\cite{Aha2002}). The TeV emission exhibits a significant extension with a radius of \\timeform{6.2'} and a center of gravity of (R.A., Dec) = (\\timeform{20h31m57.0s}, \\timeform{41D29'56.8''}) \\citep{Aha2005a}. The position is coincident with an OB association, Cyg OB2, and located in north of a microquasar, Cyg X-3. These two sources had been suggested to be possible origins of TeV $\\gamma$-rays, but no firm evidence to determine the counterpart has been found. TeV $\\gamma$-ray emission of this region has also been reported by other telescopes: Whipple, MAGIC, and Milagro (\\cite{Konopelko2007}, \\cite{Alb2008}, \\cite{Abdo2007}).  In the X-ray band, Chandra and XMM-Newton observed the TeV~J2032+4130 region. The first observation by Chandra detected 27 point sources within the observed field ($\\sim$ \\timeform{17'}) with the exposure time of 5 ks (\\cite{MHG2003}, MHG2003 hereafter,). \\citet{But2003} also resolve 19 point sources above the threshold of 2.5 $\\sigma$ by adapting the $\\tt wavdetect$ tool to the same data. Then follow-up deeper 50~ks observation remarkably increased the number of detected sources: $240$ sources in almost the same filed (\\cite{But2006}, But2006 hereafter).  XMM-Newton also detected many point sources in the wider FOV ($\\sim$ \\timeform{30'}). By subtracting the contribution of detected point sources, \\citet{Horns2007} indicates a hint of diffuse emission extending about the size of TeV emission region. However, this result may still include contribution of faint point sources which were resolved by Chandra, due to the moderate angular resolution of XMM-Newton.  The detection of a diffuse X-ray emission for the 50~ks Chandra data has been also reported by \\citet{MGH2007} with the same analysis technique. However, the spectral model can not be constrained because of the low photon statistics observation by Chandra.  Recently, the Large Area Telescope on the Fermi Gamma-ray Space Telescope detected $\\gamma$-rays from this region, with the energy from 20~MeV to 300~GeV (1FGL J2032.2+4127, \\cite{Abdo2010a}). In addition, this $\\gamma$-ray source showed a pulsation with a pulse period of 143~ms (PSR~J2032$+$4127; \\cite{Abdo2009}, 2010b). Subsequent observation of a radio band also detected a pulsation with the consistent position and the pulse frequency of the $\\gamma$-ray pulsar. The position of the pulsar is coincident with the optical point source of the number 213 in \\citet{MT91} (MT91 213; \\cite{Camilo2009}).  These results imply that the origin of TeV $\\gamma$-rays also relates to this $\\gamma$-ray pulsar. Active pulsars are losing a significant part of the energy via relativistic particles, and forms pulsar wind nebulae (PWNe). PWNe emit synchrotron radiation from radio to X-ray bands. In addition, some PWNe are found to be TeV emitter (\\cite{Gaensler2006}, \\cite{Kar2010}). Thus PSR~J2032$+$4127 is a possible candidate for the counterpart of TeV~J2032+4120.  The distance to the PSR~J2032$+$4127 is estimated to be 3.6~kpc by measuring the dispersion measure in the radio band \\citep{Camilo2009}. While, the distance to the Cygnus OB2 is estimated to be 1.7~kpc by a spectroscopic observation of OB stars \\citep{Hanson2003}. We adopt former value as the distance to X-ray emission.  In this paper, we report X-ray observation of TeV J2032+4130 with Suzaku, which has a higher sensitivity for detecting diffuse X-ray emission with the large effective area and the low stable background. We analyze the diffuse X-ray spectrum of the PWN in detail.  Though Suzaku has the advantage in detecting diffuse X-ray sensitivity, the angular resolution is not sufficient to resolve point sources. To properly estimate the contribution of point sources, we also reanalyze the Chandra data. There are many point sources in and near the Cygnus OB2 region, which are one of the candidate for the origin of TeV emission (e.g., \\cite{But2006}). We resolve point sources within strict parameters, and subtract the point source flux from the diffuse emission. Thus we investigate diffuse emission by combination of Suzaku and Chandra. ",
        "conclusions": "We obtained X-ray spectrum of diffuse emission around the $\\gamma$-ray pulsar PSR~J2032$+$4127 (structure~1). The photon index is determined to be $\\sim$~2. This value is coincident with the typical index of X-ray spectrum from PWN: $\\Gamma~\\simeq$~1--2 \\citep{Kar2008}. In this section, we consider that the diffuse X-ray emission is radiated by the pulsar wind nebula, and discuss the radiation mechanism.  \\subsection{Energy injection} First, we discuss the energy injection rate. The spin-down power of the $\\gamma$-ray pulsar, PSR~J2032$+$4127, is calculated to be about $2.63\\times10^{35}$~erg~s$^{-1}$ (\\cite{Abdo2010b}). Meanwhile, the isotropic luminosity of the GeV $\\gamma$-ray pulsar is $1.4~\\times~10^{35}~(d/3.6 {\\rm kpc})^{2}$~erg~s$^{-1}$ (\\cite{Camilo2009}), where $d$ is the distance to the pulsar. At most about a half of the energy is emitted as GeV $\\gamma$-rays with an assumption of isotropic radiation. The ratio can be smaller in the case of collimated radiation. The intensity of off-pulse emission is almost the same as background level, which is estimated from surrounding annulus region of the pulsar (Fig. A-42 in \\cite{Abdo2010b}). Indeed, detailed analysis of off-pulse spectrum cannot constrain the flux level \\citep{Ack2011}. A large portion of GeV $\\gamma$-rays originate from the pulsar's magnetosphere, and the luminosity of PWN is negligible in this energy band. Consequently, we cannot obtain meaningful GeV $\\gamma$-ray flux for PWN. We take no account of the GeV $\\gamma$-ray band for discussing the energy injection rate from the pulsar.  X-ray luminosity and TeV $\\gamma$-ray luminosity, which are considered to be diffuse emission, are $\\sim 3\\times10^{32}$~erg~s$^{-1}$ and $\\sim 2\\times10^{33}$~erg~s$^{-1}$, respectively (1.0--10 TeV; \\cite{Aha2005a}, \\cite{Alb2008}). These are two or three orders of magnitude smaller than the spin-down energy. The spin-down energy is enough to supply X-ray and TeV $\\gamma$-ray emission. Therefore the energy of the diffuse emission in TeV and X-ray could be supplied by the pulsar, as a PWN. The ratio of $\\sim 10^{-3}$ between X-ray luminosity and the spin-down energy of the pulsar is typical for X-ray emission of PWN (e.g., Figure~10 in \\cite{Kar2009}).   \\subsection{Emission mechanism} In Fig.~\\ref{fig:wide_spec}, we plot the fluxes of X-rays and TeV $\\gamma$-rays. The X-ray flux is indicated by best-fit model of the diffuse component spectrum of structure~1. The ratio of flux between TeV and the X-ray bands $F_{\\rm TeV}/F_{\\rm X}$ is $\\sim~10:1$. Although TeV $\\gamma$-ray emission dominates the X-ray flux in some PWNe (\\cite{Funk2007}, \\cite{Kar2009}, \\cite{MGH2009}), it is difficult to explain such a large ratio with a simple energy distribution of high energy electrons.  If the origin of TeV $\\gamma$-ray emission is inverse Compton scattering of cosmic microwave background (CMB) by TeV electrons, the same high energy electrons also emit X-rays by synchrotron radiation. In this case, the ratio of TeV flux to the X-ray flux depends only on the magnetic field strength. The synchrotron emission model of TeV electrons are shown by solid lines in Fig.~\\ref{fig:wide_spec} for some assumed magnetic fields of 1, 3, and 10~$\\mu$G. Our data corresponds to the magnetic field of about 1~$\\mu$G. However, this value is much lower than that expected in this region. The total magnetic field strength in the Galaxy disk is larger than 3~$\\mu$G in the whole area \\citep{Beck2001}. In addition, active star forming region indicates stronger magnetic field.  If a photon energy density of the local radiation field is much higher than CMB, Compton scattered TeV $\\gamma$-rays exhibit larger intensity and might explain our data. The number of OB stars are $\\sim 1000$ in Cyg OB2 region. Although we could not rule out the possibility of the strong radiation background, we only consider CMB in this study.   One possibility for resolving the discrepancy is a hadronic origin of TeV $\\gamma$-ray emission (e.g., \\cite{Bed2003}). In such a case, X-ray emission can be much lower than leptonic case, and a typical strength of magnetic field of PWNe, 3~$\\mu$G, could be acceptable.  Another possibility is an existence of a high energy cut-off of electrons. The energy of electrons, which is responsible for TeV emission, is lower than X-ray emitting electrons by about one order of magnitude. If the energy distribution of TeV electrons exhibit strong cut-off above the TeV region, the flux of X-ray emission is decreased (dashed line in Fig.~\\ref{fig:wide_spec}). \\citet{Mat2009} investigated the evolution of the PWNe by comparing 14 samples, and concluded that the ratio of TeV and X-ray luminosities has a positive correlation with the characteristic age of the pulsar. The X-ray luminosity decreases with the pulsar age because of radiative cooling in the earlier stage (by a factor $\\sim 10^6$ in $10^5$ yr), while the $\\gamma$-ray luminosity is constant. Thus, the ratio becomes larger. PSR~J2032$+$4127 is an old pulsar of 120~kyr \\citep{Abdo2010b}. The large luminosity ratio might be natural for such an evolved PWN.  This cut-off hypothesis also explains the difference in size between TeV and X-ray emissions as discussed in \\citet{Aha2005b}. X-rays are emitted by young electrons, because higher energy electrons lose their energy quickly. On the other hand, older electrons radiate TeV $\\gamma$-rays. Such a difference in age could cause the concentration of X-ray emission near the pulsar and the diffusion of TeV emission. Thus the size of X-ray emitting region could be much smaller than that of TeV.  \\subsection{Other structures} The best-fit parameters of structure~2 are almost the same as structure~1. In addition, the X-ray emission size is also comparable to structure~1. Although no pulsar is found at the location of structure~2, the X-ray emission can be explained by PWN in the same manner as structure~1. If structure~2 is also a PWN, a part of TeV emission could originate from this source.  The remaining diffuse emission also exhibits the same spectral shape as structure~1. However, the X-ray flux and the size is much larger than structure~1, 2. We cannot insist that the remain is related to TeV emission only by the spectral similarities."
    },
    "1112/1112.4827_arXiv.txt": {
        "abstract": "{In this paper we report on the broadband X-ray properties of a complete sample of 33 absorbed Seyfert galaxies hard X-ray selected with \\integral. The high quality broadband spectra obtained with  both \\xmm, and \\integral-IBIS data are well reproduced with an absorbed primary emission with a high energy cutoff and its scattered fraction below 2-3 keV, plus the Compton reflection features. A high energy cut-off is found in 30\\% of the sample, with an average value below 150 keV. The diagnostic plot N$_H$ $vs$ F$_{obs}$(2--10 keV)/F(20--100 keV) allowed the isolation of the Compton thick objects, and may represent a useful tool for future hard X-ray observations of newly discovered AGN. We are unable to associate the reflection components with the absorbing gas as a torus, a more complex  scenario being necessary. In the Compton thin sources, a fraction (but not all) of the Fe K line needs to be produced in a gas possibly associated with the optical Broad Line Region, responsible also for the absorption.  We still need a Compton thick medium (not intercepting the line of sight) likely associated to a torus, which contributes to the Fe line intensity and produces the observed reflection continuum above 10 keV. The so-called Iwasawa-Taniguchi effect can not be confirmed with our data. Finally, the comparison with a sample of unobscured AGN shows that, type 1 and type 2 (once corrected for absorption) Seyfert  are characterized by the same nuclear/accretion properties (luminosity, bolometric luminosity, Eddington ratio), supporting the ''unified'' view. }  \\FullConference{The Extreme and Variable High Energy Sky - extremesky2011,\\\\ September 19-23, 2011\\\\ Chia Laguna (Cagliari), Italy}   \\begin{document} ",
        "introduction": "The widely accepted Unification Model (UM) for active galactic nuclei \\cite{antonucci93} is strongly supported by spectropolarimetric observations of hidden broad-line regions (HBLRs) in several type 2 AGN, and by the X-ray observations, showing that Seyfert 2  have absorption columns exceeding the Galactic ones \\cite{cappi06}. Despite this, and in view of  several pieces of evidence coming from recent observations of different classes of AGN, the UM is still being debated. In particular, the association of the Compton thick (N$_{H}\\gtrsim \\sigma_{T}^{-1} \\backsim$ 10$^{24}$ cm$^{-2}$) obscuring medium with a  uniform torus at intermediate scale between the  BLR and NLR, has been questioned by the presence of N$_{H}$ variability on timescales of hours, requiring an X-ray absorber no larger than the BLR \\cite{risaliti07}.  These absorbing media are also responsible for the reflected features in the X-ray range: a Compton hump above 10 keV and a Fe line at 6.4-6.9 keV \\cite{perola02}. This kind of study has been performed in a statistical way with a variety of different samples over a broadband energy range: BeppoSAX \\cite{risaliti02}, \\integral\\ and \\xmm\\ \\cite{derosa11,derosa08}, \\suzaku\\ and Swift-BAT (both type 1 and type 2 AGN \\cite{eguchi09,winter09,fukazawa11}. The 4th INTEGRAL/IBIS catalogue \\cite{bird10} includes more than 700 hard X-ray sources detected over a large part of the sky above 20 keV  and with positional accuracies ranging between 0.2 and 5.5 arcmin, depending on the source significance. To obtain the best quality broadband high energy data, the \\integral/AGN team has undertaken a programme of  follow-up observations of hard X-ray selected AGN with \\xmm\\ \\cite{molina08,derosa08,derosa11} In this paper we report on the broadband spectral properties of 33 type 2 AGN included in the  hard X-ray selected AGN complete sample as  defined by \\cite{malizia09} optically classified as type 1.8-2. The main goals of this study, widely presented and discussed in \\cite{derosa11}, are (1) to measure the spectral properties of obscured AGN: the intrinsic continuum shape  together with the absorption distribution and (2) to accurately constrain the strength of the reflection components. The line properties tell us about the obscuring/reflecting medium while the Compton continuum is an important yet uncertain parameter (like the high energy cut-off) in population synthesis models of the CXB \\cite{gilli07}. ",
        "conclusions": "\\subsection{The primary absorbed and scattered continuum}   The BLM  reproduces quite well all the spectra analysed since the $\\chi^{2}$/dof ranges between 0.77-1.12. The mean value of the photon index and its standard deviation for the whole sample is $\\Gamma$=1.68, $\\sigma$=0.30. This value is in very good agreement with those found in previous analyses performed in a limited energy range (2--10 keV) in both type 1 and type 2 AGN (e.g, \\cite{singh11,bianchi09}). We measured the high energy cutoff in 10 out of 33 sources (i.e. 30 per cent of the sample), while only an upper limit at 300 keV has been found in IGR J16351-5806. In the remaining sources we obtained lower limits to E$_{c}$. The average value of the ten measured high energy cutoffs is E$_c$=73 keV  ($\\sigma$=25 keV). In Fig. \\ref{fig1} (left panel) we plot E$_{c}$ as a function of $\\Gamma$. All the measured values are below 150 keV, while lower limits (filled triangles) are well grouped below 300 keV. These values are consistent with those observed in the Seyfert 1 by BeppoSAX \\cite{perola02,derosa07}. We do not find any evidence of correlation between $\\Gamma$ and E$_c$,  the correlation coefficient being r=0.1. Our analysis strongly suggests that the high energy rollover is an ubiquitous property of Seyfert galaxies, and that its value  ranges in the same intervals for type 2 and type 1 Seyferts \\cite{molina08}.  We found a value of $f_{sc}$ of a  few per cent in all sources (with the exception of IGR J14515-5542 and IGR J16351-5806  see  discussion in \\cite{derosa11}). This value is in good agreement with that found in different samples of obscured AGN, observed with  \\xmm\\ and \\chandra\\ \\cite{bianchi06}. No evidence of correlation between $f_{sc}$  and relative reflection R has been detected in our sample.  The average value of the absorption column density is 23.15 ($\\sigma$=0.89). Excluding the six Compton thick sources we obtain an average value ($\\sigma$) of 22.85 (0.67). In Fig. \\ref{fig1}, central panel, we plot N$_{H}$ as a function of the ratio of the observed 2--10 keV flux $vs$ the 20--100 keV flux. This diagnostic plot was proposed by \\cite{malizia07} in an attempt to isolate peculiar objects among  a sample of 34 \\integral\\ AGN observed with Swift-XRT. The plot shows a clear anticorrelation between N$_{H}$ and the softness ratio: this ratio becomes lower as the column density increases. The lines define a region obtained by assuming an absorbed power-law with high energy cutoff, where we have used  the average values of $\\Gamma$ and E$_c$, with errors, as found in our sample. In this diagnostic plot, Compton thick AGN are well isolated, all six Compton thick objects of our sample are located above the line.   \\begin{figure} \\centering \\includegraphics[width=0.3\\linewidth]{GEc_4.ps} \\includegraphics[width=0.3\\linewidth]{softness_nh6.ps} \\includegraphics[width=0.3\\linewidth]{NhEW_6.ps} \\caption{ \\textit{Left panel}: High-energy cutoff vs photon index for the whole sample.\\textit{Central panel}: The column density distribution as a function of the softness ratio F$_{oss}$(2--10 keV)/F(20--100 keV). \\textit{Right panel}: Fe K$\\alpha$ EW  as a function of N$_{H}$. See text for details} \\label{fig1} \\end{figure}   \\subsection{The reflection features: Fe K line and Compton hump} \\label{refl_sec} We have  measured the reflection fraction (R) in 22 sources (average value R=1.63 with $\\sigma$=1.15), while for 8 objects we only found upper limits. With the sole exception of IGR J17513-2011, a neutral Fe line is detected in all the sources of the sample (marginally significant for IGR J00040+7020, 95\\%). The average value of the equivalent width (EW), calculated with respect to the total direct+reflected continuum is 430 eV (with a somewhat high standard deviation of 564 eV),  while, excluding the Compton thick objects, it becomes 196 eV ($\\sigma$=176 eV). It is worth noting that the fit of our BLM produces a degeneracy between the reflection fraction R and the photon index $\\Gamma$: a  higher $\\Gamma$  produces a higher value of R. Another point to stress is that the soft and high energy data (below and above 10 keV) are non-simultaneous. This point is discussed in \\cite{derosa11}. In the right panel of Fig. \\ref{fig1} we show the Fe K$\\alpha$ EW  as a function of N$_{H}$. Below N$_{H}$=10$^{23.5}$cm$^{-2}$ the EW trend is almost flat, with a best fit constant value of EW$_{0}$=120$\\pm$16 eV; this value is in a very good agreement with that found in other X-ray sample of Compton thin  type 2 \\cite{fukazawa11}  as well as type 1 AGN \\cite{bianchi09} observed with  \\suzaku, and \\xmm, respectively. However, above N$_{H}$=10$^{23.5}$cm$^{-2}$ the behaviour is different, and a clear correlation between N$_{H}$ and EW is evident, as expected from reflection from a Compton thick torus \\cite{awaki91}. The dashed line in that figure represents the EW values as obtained  from transmission through a uniform slab of medium with solar abundances subtending 4$\\pi$ to a continuum source of photon index $\\Gamma$ = 2.0 \\cite{leahy93}.  Large EW above 10$^{23.5-24}$cm$^{-2}$ could be originated in a Compton thick torus, also responsible for the observed absorption. On the other hand, the large EW  observed in Compton thin AGN (N$_{H}$ below 10$^{23.5}$cm$^{-2}$) are far above the dashed line, suggesting that the line is due to reflection from a material different from the absorber. The evidence that the same Fe K$\\alpha$ EW ($\\sim$120 eV) is found  for Compton thin and type 1 AGN, suggests that the origin of this component is in the same medium, e.g. the BLR. The solid line in  Fig. \\ref{fig1}  represents the function EW(N$_{H}$)=EW$_{0}$exp($\\sigma_{Fe} N_H$), with $\\sigma_{Fe} $ being the photoelectric cross-section at 6.4 keV and EW$_{0}$=120 eV. This trend well reproduces the Fe EW below 10$^{23.5}$cm$^{-2}$, and represents a scenario in which the Fe line is generated by reflection off the BLR seen along an unobscured line-of-sight, while the underlying continuum is obscured by matter with a column density $N_{H}$;  the total effect will be a larger value of EW increasing  N$_{H}$. It is worth noting that in 4 sources of our sample a broad Fe line is required to fit the broadband data. The intrinsic width of these  lines is completely compatible with an origin within the BLR. This scenario is supported by the recent observations of rapid variation of the absorbing column density in both Compton thin and thick AGN. However, we stress that a Compton thin gas, also covering 4$\\pi$ solid angle to the source, cannot create a Fe EW of 50-150 eV \\cite{yaqoob01} as observed in our sample below N$_{H}$=10$^{23.5}$cm$^{-2}$; then, a Compton thick medium like a torus (not intercepting the line of sight in Compton thin or type 1 AGN), should exist.  Our conclusion is that in heavily obscured Seyferts (N$_{H}$ above 10$^{23.5-24}$cm$^{-2}$), the Fe line is produced in a torus also responsible for the absorption;  for Compton thin sources (N$_{H}$ below 10$^{23.5}$cm$^{-2}$), as well as for type 1 AGN, a high fraction of the Fe K line (but not the total line photons) has to be produced in a gas  located closer to the black hole than to the torus. This region can be associated with the BLR that could also be responsible for  the observed moderate absorption. We also investigated the X-ray Baldwin/''Iwasawa-Taniguchi'' effect \\cite{baldwin77} indicating the existence of an  anticorrelation between the 2--10 keV luminosity corrected for absorption and the EW of the Fe K line. A  linear fit  to the data  of the form log(EW)=5.3$\\pm4.8$ - (0.07$\\pm$0.11)log(L$_{corr}$(2--10 keV)), indicates no evidence of a correlation (r=0.109). The X-ray Baldwin effect is not confirmed in our sample; although, we note that using the average EW in each luminosity bin, the correlation increases to r=0.2.  Considering the hard X-ray luminosity in the 20--100 keV energy range, the correlation coefficient increases in this case to r=0.3, still indicating a lack of correlation. It is worth noting that the luminosity range we are exploring is rather narrow (log(L)=41.5-44.5), and this could make it difficult to measure any correlation as observed by \\cite{bianchi09} over a wider luminosity range.    \\subsection{Type 1 $vs$ type 2 accretion properties} \\label{ty1_ty2}   From the samples of hard X-ray selected obscured (this work) and unobscured \\cite{molina08} AGN, we collected  all available published mass measurements obtained with reverberation mapping techniques, maser, X-ray variability or M-$\\sigma$ relation. A sub-sample of 33 (17 type 1+16 type 2) AGN with known masses has been obtained. This sub-sample can be considered as representative of the whole one (see discussion in \\cite{derosa11}). We used this hard X-ray selected sample to draw  a statistical comparison between type 1 and type 2 Seyfert of (1) the absorption corrected 2--10 keV and 20--100 keV luminosity, L$_{corr}$(2--10 keV) and L(20--100 keV); (2) the bolometric luminosity, L$_{bol}$; and (3) the Eddington ratio, L$_{bol}$/L$_{Edd}$, with the main goal of testing the prediction of the UM for AGN. The distribution of the corrected 2--10 keV luminosity (left panel of Fig. \\ref{fig2}),  has averaged value of 43.52 ($\\sigma$=0.71) and 43.06 (0.72)  for type 1 and type 2, respectively. Using the  Kolmogorov-Smirnov test we do not find a significant difference in the distributions, having a probability  P=0.111 to be drawn from the same parent population. This result is also confirmed in the higher, unbiased for absorption, energy range (20--100 keV) (see distribution in the central panel of Fig. \\ref{fig2}), with a KS test probability of 0.137. To be conservative and fully convinced to have properly corrected  the 2--10 keV luminosities for absorption, we excluded the  Compton thick sources from this analysis and, even in this case,  the KS test probability  confirms the similarity of the two populations (P=0.165).  \\begin{figure} \\centering \\includegraphics[width=0.3\\linewidth]{Lun210_withMass.ps} \\includegraphics[width=0.3\\linewidth]{L20100_withMass.ps} \\includegraphics[width=0.3\\linewidth]{G_Lbol_Ledd_ty1_ty2_withMass.ps} \\caption{\\textit{Left and Central panels}: Luminosity distribution (corrected for absorption) in (2--10 keV) and 20--100 keV for the sub-sample of type 1 and type 2 AGN with measured masses. Type 1 are shown as horizontal large spaced  line, type 2 are diagonal small spaced lines. \\textit{Right panel}: $\\Gamma$ as a function of the ratio L$_{bol}$/L$_{edd}$  for the sub-sample of type 1 and type 2 AGN with measured masses. Filled triangles are type 1 while open circles represent type 2 AGN.} \\label{fig2} \\end{figure}   For all sources in the sub-sample we calculated the bolometric luminosity  derived from the 2--10 keV luminosities, adopting the luminosity-dependent bolometric correction given in  \\cite{marconi04}. The L$_{bol}$ average value is 44.91 ($\\sigma$=0.91) and 44.33 (0.90) for type 1 and type 2, respectively. A KS test gives a probability P=0.111 that the distributions belong to the same parent population, suggesting, even in this case, the similarity between the two classes of AGN. In the right panel of Fig. \\ref{fig2} we plot $\\Gamma$ versus the Eddington ratio for our sub-sample. The average value of the ratio L$_{bol}$/L$_{edd}$  is 0.17 ($\\sigma$=0.20)  and 0.32 (0.72) for type 1 and type 2 AGN, respectively. Although the dispersion is huge, we note that, as for the luminosity, no evidence of separation in the accretion rate distribution is evident between absorbed and unabsorbed AGN (KP probability P=0.292). These  results concerning luminosities is also reflected in the spectral properties. In fact a comparison between the primary photon index in type 1 and type 2 class does not show any evidence for a separation being the average value of $\\Gamma$ equal to 1.80 ($\\sigma$=0.18) and 1.75 ($\\sigma$=0.26) for type 1 and type 2 respectively. A KS test  in this case gives a probability P=0.862.  These same conclusions have  recently been reached in \\cite{singh11} through the X-ray analysis of an optical selected sample of type 1 and type 2 Seyfert. The value of the L$_{bol}$/L$_{edd}$ ratio in our sub-sample of type 1 and type 2 AGN  range in 0.001 - 3, extending to higher values with respect to a optical selected sample \\cite{singh11} We are aware that this result could be affected by two issues: mass estimation and bolometric corresction, both are discussed in \\cite{derosa11} Finally, we also stress that, using the whole hard X-ray selected sample of  41 type 1 + 33 type 2 , we find that L$_{corr}$(2--10 keV) and L(20--100 keV) have a  small  probability  (P=0.05 and 0.02, respectively) to be drawn from the same parent population. This effect is likely to be due to a bias in the sample selection through high lunimosity AGN."
    },
    "1112/1112.1319_arXiv.txt": {
        "abstract": "{% 3D MHD global solar simulations coupling the turbulent convective zone and the radiative zone have been carried out. Essential features of the Sun such as differential rotation, meridional circulation and internal waves excitation are recovered. These realistic models are used to test the possibility of having the solar tachocline confined by a primordial inner magnetic field. We find that the initially confined magnetic fields we consider open into the convective envelope. Angular momentum is transported across the two zones by magnetic torques and stresses, establishing the so-called Ferarro's law of isorotation. In the parameter space studied, the confinement of the magnetic field by meridional circulation penetration fails, also implying the failure of the tachocline confinement by the magnetic field. Three-dimensional convective motions are proven responsible for the lack of magnetic field confinement. Those results are robust for the different magnetic field topologies considered, i.e. aligned or oblique dipole.} ",
        "introduction": "\\label{sec:introduction}  Since its discovery \\citep[][]{Brown:1989hq}, the solar tachocline has puzzled the scientific community. In particular, a clear explanation for its extreme thinness \\citep[less than 5\\% of the solar radius $R_\\odot$, see][]{Charbonneau:1999es} is still lacking. The first theoretical work dedicated to the tachocline was carried out by \\citet{Spiegel:1992tr}. They considered an hydrodynamic tachocline and showed that such an interface layer will spread into the radiative interior because of thermal diffusion. As a result, they estimated that the differential rotation should extend down to $0.3\\, R_\\odot$ after $4.5\\, $Gyears, which is in total contradiction with the helioseismic inversions \\citep{Schou:1998bi,Thompson:2003bk}. Consequently, additional physical mechanisms have to be considered in order to properly explain the observed thinness of the tachocline.  \\citet{Spiegel:1992tr} suggested that the stratification in the tachocline would imply an anisotropy in the turbulence, making it predominantly horizontal. It would then erode the latitudinal gradient of angular-velocity \\citep{Elliott:1997wz}. However \\citet{Gough:1998ik} pointed out that such anisotropic turbulence would on the contrary act as an anti-diffusion \\citep[][]{Dritschel:2008em}. In fact, the question is still hardly settled when both radial and latitudinal shears of angular-velocity are taken into account \\citep{Miesch:2003eh,Kim:2005fe,Leprovost:2006ft,Kim:2007cy}. Further, \\citet{Tobias:2007ks} showed that the introduction of (even weak) magnetic fields in the bulk of the tachocline would erase both the diffusive and anti-diffusive behaviors of 2D turbulence.  \\citet{Gough:1998ik} then proposed that a (fossil) dipolar magnetic field confined in the solar interior could oppose the thermal spreading of the tachocline. This solution offered also an explanation for the solid body rotation of the solar radiation zone and such a fossil field was also invoked by \\citet{Rudiger:1997wz} to confine the tachocline. The fossil field confinement scenario is in fact a double confinement problem. First, the imposed magnetic field erodes latitudinal gradients of angular-velocity, thus confining the tachocline. Second, the magnetic field has to remain confined in the radiation zone against its outward ohmic diffusion. In order to confine the magnetic field, it was first argued that a meridional flow coming from the convection zone down to the radiative interior through the tachocline at the high latitudes could eventually prevent the field from diffusing outward in this region \\citep{Gough:1998ik}. At the equator,  \\citet{Wood:2011es} suggested that if the magnetic field were to connect to the convection zone, the magnetic pumping \\citep{Tobias:2001ho,Dorch:2001jx,Ziegler:2003fd} would 'confine' it below the region of intense shear.  Despite the appeal of this scenario, none of the many numerical simulations that have been carried out in the past succeeded in recovering it completely. The first simulations of the radiation zone in 2D \\citep{Garaud:2002fq} and in 3D \\citet{Brun:2006bb} showed that if the confinement of the magnetic field fails, angular momentum is transported along the field lines into the radiation zone, making the radiative interior rotate differentially. The Ferraro's law of iso-rotation is then established \\citep{Ferraro:1937ts}, in complete contradiction with solar observations. However these models did not allow for flows penetrating from the convection zone, such as plumes or meridional circulation. In their simulations, \\citet{Sule:2005bq,Rudiger:2007be,Garaud:2008fi} imposed a meridional circulation at the top of the radiation zone. They recovered partially the \\citet{Gough:1998ik} scenario, but one may object that those results are highly sensitive to the prescribed (profiles and amplitude) meridional circulation. Simulations coupling self-consistently the two zones of the Sun were then carried out by \\citet{Rogers:2011kw} in 2D and by \\citet{Strugarek:2011cx} (hereafter SBZ11) in 3D in order to take into account the motions of the convection zone. Both studies were not able to validate the magnetic confinement scenario. Finally, \\citet{Wood:2011id} conducted an analytical study in a reduced cartesian model and proposed an improved theory for a magnetic confinement scenario \\textit{\\`a la} \\citet{Gough:1998ik}.  We are interested here in the fact that SBZ11 showed that the convective motions were responsible for the lack of confinement of the buried magnetic field, and that in most of the works previously cited, the magnetic field enters the convection zone primarily at the equator. Since different magnetic topologies may lead to significantly different interactions with the convective motions, the universality of the results obtained with our axisymmetric dipolar field can be questioned. We thus conduct here numerical simulations based on the SBZ11 model to investigate the role of the magnetic field topology. In Sect. \\ref{sec:model-both-conv} we summarize the main ingredients of the model used by SBZ11, and in Sect. \\ref{sec:infl-magn-topol} we study the impact of the magnetic topology on the confinement of the field. Conclusions and perspectives are reported in Sect. \\ref{sec:consl-persp}. ",
        "conclusions": "\\label{sec:consl-persp}  In a previous paper (SBZ11), we tested the \\citet{Gough:1998ik} magnetic confinement scenario of the solar tachocline with a 3D non-linear MHD model of $90\\%$ of the Sun which couples together the convective and radiative zones. We showed that a fossil axisymmetric magnetic field does not remain confined in the radiation zone, and thus cannot prevent the spread of the tachocline.  We showed that the motions (both meridional circulation and convection) were not strong enough to confine the field and that upward motions at the base of the convection zone were on the contrary helping the field to expand inside the convection zone. Although our numerical experiment is far from representing the real Sun, we are confident that our parameter regime allows us to model the main ingredients of the tachocline confinement scenario. More precisely, we demonstrated in SBZ11 that even if we consider high magnetic diffusivity, our magnetic Reynolds number is higher than $1$ in the overshooting/penetration region and in the convection zone. The magnetic field evolution is thus dominated by advection and not by diffusion. The overshooting depth of the convective plumes is likely to be overestimated because our Peclet number is lower than the solar one, and the penetration of the meridional circulation may be underestimated \\citep[see][]{Garaud:2008fe,Wood:2011id} in our model. However, penetration and overshooting are clearly observed in our model. If more 'realistic' values of the parameters could have consequences on the polar dynamics, we stress here that it may not have much influence at the equator, where the confinement of the magnetic field fails.  The magnetic topology considered in SBZ11 was very simple (axisymmetric dipole), and the robustness of the results is not obvious when considering the geometry of the field. We demonstrated in Sect. \\ref{sec:opening-field-at} that an axisymmetric dipole preferentially diffuses at the location where its maximum radial gradient is, \\textit{i.e.,} at the equator. This has some importance since SBZ11 showed that the lack of confinement specifically at the equator was responsible for the failure of the \\citet{Gough:1998ik} scenario. In order to test other topologies, we reported in Sect. \\ref{sec:diff-magn-topol} the study of an oblique dipole buried in the radiation zone. A tilted dipole can be decomposed into an axisymmetric and a non-axisymmetric part, and the axisymmetric component was shown to evolve similarly to the purely axisymmetric dipole of SBZ11. As a result, the magnetic confinement also fails for the tilted dipole.  The non-axisymmetric component of the titled dipole exhibits interesting three dimensional dynamics and transports angular momentum at the equator. In order to get rid of the SBZ11 dynamics and to isolate the effect of the non-axisymmetric dipole, simulations with a purely perpendicular dipole (\\textit{i.e.}, with no axisymmetric component, $\\beta=90$ in Eq. \\eqref{eq:dipole_nax_ASH}) will be reported in a future paper."
    },
    "1112/1112.3764_arXiv.txt": {
        "abstract": "In this work I investigate the redshift evolution of pair fraction of a sample of 196 massive galaxies from $z$ = 0 to 3, selected from the COSMOS field. We find that on average a massive galaxy undergoes $\\sim 1.1 \\pm 0.5$ major merger since $z=3$. I will review the current limitations of using the pair fraction as a probe for quantifying the impact of mergers on galaxy evolution. This work is based on the paper \\citet{Man2011}. ",
        "introduction": "Mergers are thought to be one of the primary drivers of galaxy evolution, at least since $z \\sim 3$. Theories and simulations have shown that mergers are able to transform the structure of a galaxy, and are required for the mass assembly of galaxies. However, it remains an open issue how important merging is, compared to other physical mechanisms such as cold gas accretion or secular evolution. Hence there is a motivation to quantify the importance of merger through observations.  Typically merger candidates are identified from photometric catalogues through either pair selection \\citep{Bluck2009, Williams2011, Newman2011} or morphology \\citep{Lotz2008, Bluck2011}. Irregular morphologies at high redshift could be due to cosmological surface-brightness dimming or clumpy star formation, and are not always related to merging activity. Hence we focus on using pair counts as a probe for identifying mergers in this work. \\vspace{-10pt} ",
        "conclusions": "\\begin{itemize} \\item The pair fraction evolves only mildly with redshift since $z=3$. \\vspace{-10pt} \\item A massive galaxy experiences $N_{m} = 1.1 \\pm 0.5$ major mergers on average since $z=3$. This amount of major merging is sufficient to explain the number of newly formed massive galaxies, but insufficient to explain the observed size evolution. This hints that massive galaxies undergo other mechanisms that is efficient in puffing up their sizes, but not so important in growing the mass. Minor merging is compatible with this scenario. \\vspace{-10pt} \\item It remains a challenge to constrain the observability timescale of mergers due to the complexity of the problem. Its importance oughts to be highlighted, as it is a crucial quantity for interpreting the impact of mergers on the growth of galaxy population as a whole. \\end{itemize} \\vspace{-10pt}"
    },
    "1112/1112.4372_arXiv.txt": {
        "abstract": "{ A method for estimating the true meteor rate $\\lambda$ from a small number of observed meteors $n$ is derived. We employ Bayesian inference with a Poissonian likelihood function. We discuss the choice of a suitable prior and propose the adoption of Jeffreys prior, $P(\\lambda)=\\lambda^{-0.5}$, which yields an expectation value $E(\\lambda) = n+0.5$ for any $n\\geq 0$. We update the ZHR meteor activity formula accordingly, and explain how 68\\%- and 95\\%-confidence intervals can be computed. } } ",
        "introduction": "The formula commonly used to estimate the Zenithal Hourly Rate (ZHR) of meteors as given in the Handbook of the International Meteor Organization (IMO; Rendtel \\& Arlt, 2008) and used in many meteor activity graphs (e.g., Arlt \\& Barentsen, 2006), is given by: \\begin{equation} \\E(\\ZHR) = \\frac{ n_{\\rm tot} + 1 } { T } \\quad\\mbox{with}\\quad \\sigma=\\frac{ \\sqrt{n_{\\rm tot} + 1} } { T }, \\label{zhr} \\end{equation} where $n_{\\rm tot}$ is the total number of meteors counted in a number of observing intervals ($n_{\\rm tot} = \\sum_{i} n_i$) and $T$ is the observing time weighted by a given correction factor for each interval ($T = \\sum_{i} T_{{\\rm eff},i} / C_i $).  The use of $n_{\\rm tot}+1$ rather than $n_{\\rm tot}$ often surprises observers, because it yields a ZHR which is larger than zero even when no meteors are observed. Whilst Arlt (1999) already indicated that $n_{\\rm tot}+1$ is used to account for the effects of small-number statistics, this paper will explain the formula in more detail. However, contrary to the earlier publications, we will suggest that $n_{\\rm tot}+0.5$ rather than $n_{\\rm tot}+1$ is the most appropriate formula.  We will first describe the problem of small-number statistics in \\S2. We then describe the solution using Bayesian inference in \\S3, followed by a discussion on the choice of the prior assumptions in \\S4. In \\S5 we explain how confidence intervals may be computed, and in \\S6 we provide examples. Finally in \\S7 we discuss the effect of correction factors and in \\S8 we present the conclusions. ",
        "conclusions": "In conclusion, we recommend to compute ZHR values using the term $n+0.5$: \\begin{equation} \\E(\\ZHR) = \\frac{\\left(n_{\\rm tot}+0.5\\right) \\, r^{6.5-{\\rm lm}}}{T_{\\rm eff}\\sin h_{\\rm R}}, \\end{equation} with error margins: \\begin{equation} \\Delta{\\rm ZHR} = \\frac{\\rm ZHR}{\\sqrt{n_{\\rm tot}+0.5}}. \\label{finalerror} \\end{equation} Note that for $n_{\\rm tot}\\leq 30$, the error margins listed in Table~1 should be used instead of Eqn.~\\ref{finalerror} to obtain a 68\\%-confidence interval.  This method of computing the ZHR adds a small bias taking into account the asymmetry of possible rates. In particular, it is essential to adopt the method whenever rates are based on less than $\\sim10$ meteors. Such situations commonly occur when a major shower is analysed using very short (e.g.\\ 1-minute) intervals. When computing a rate based on a number of observing periods (indexed with $i$), {\\em never ever\\/} compute the rate from $n_i+0.5$ for individual periods and average them afterwards. A more accurate estimate is based on the sum of meteors from these observing periods, and hence on $n_{\\rm tot}+0.5$. Finally, it is, of course, much better to have enough observations and large enough $n_{\\rm tot}$ that the subtleties of choosing a good prior are no longer important, i.e.\\ $n_{\\rm tot} \\gg 0.5$.  Comprehensive information about Bayesian inference in general and the choice of priors can be found in e.g.\\ Kass \\&Wasserman (1996) and Bolstad (2007; Chapter~10 for priors)."
    },
    "1112/1112.1427.txt": {
        "abstract": "During stellar core collapse, which eventually leads to a supernovae explosion, the stalled shock is unstable due to  the standing accretion shock instability (SASI). This instability induces large-scale non spherical oscillations of the shock, which have crucial consequences on the dynamics and the geometry of the explosion. While the existence of this instability has been firmly established, its physical origin remains somewhat uncertain. Two mechanisms have indeed been proposed to explain its linear growth. The first is an advective-acoustic cycle, where the instability results from the interplay between advected perturbations (entropy and vorticity) and an acoustic wave. The second mechanism is purely acoustic and assumes that the shock is able to amplify trapped acoustic waves. Several arguments favouring the advective-acoustic cycle have already been proposed, however none was entirely conclusive for realistic flow parameters. In this article we give two new arguments which unambiguously show that the instability is not purely acoustic, and should be attributed to the advective-acoustic cycle. First, we extract a radial propagation timescale by comparing the frequencies of several unstable harmonics that differ only by their radial structure. The extracted time matches the advective-acoustic time but strongly disagrees with a purely acoustic interpretation. Second, we present a method to compute purely acoustic modes, by artificially removing advected perturbations below the shock. All these purely acoustic modes are found to be stable, showing that the advected wave is essential to the instability mechanism. ",
        "introduction": "Several observational as well as theoretical arguments show that most core collapse supernovae are asymmetric explosions. The observational evidence include polarisation measures \\citep{leonard06}, double peaked oxygen lines \\citep{maeda08}, and the large velocities of neutron stars \\citep{lai01}. On the theoretical side, the failure of realistic spherically symmetric calculations to produce an explosion \\citep{liebendorfer01} suggests that the breaking of the initial spherical symmetry is essential to the explosion mechanism. Indeed most proposed mechanisms are fundamentally asymmetric: the modern version of the neutrino-driven mechanism \\citep{marek09}, the acoustic mechanism \\citep{burrows06}, and the magneto-centrifugal mechanism \\citep{wheeler02}.  The spherical symmetry can be broken by several phenomena. In the magneto-centrifugal mechanism, the asymmetry is due to a fast rotation in conjunction with a strong magnetic field. In the case of moderate rotation, two non radial instabilities can develop during the stalled shock phase: the neutrino-driven convection, and the stationary accretion shock instability usually called SASI which induces non spherical oscillations of the shock \\citep{blondin03}. Of these two instabilities, SASI is the most efficient in creating a global asymmetry \\citep{foglizzo06} and helping the explosion \\citep{marek09,burrows06}. In addition to its crucial role in the explosion mechanism, SASI has a number of potentially important consequences: it affects the kick and the spin of the neutron star \\citep{scheck06,blondin07a}, causes an emission of gravitational waves \\citep{kotake07}, and creates a time dependence of the neutrino signal \\citep{marek09b}.  But what is SASI ? Which physical phenomenon drives these shock oscillations ? Surprisingly, the answer to this basic question remains somewhat uncertain. Indeed, two mechanisms have been proposed -- the advective-acoustic cycle \\citep{foglizzo07} and a purely acoustic mechanism \\citep{blondin06} -- and distinguishing them remains a controversial issue. The aim of this article is to clarify this question.  In Section~\\ref{sec:review}, we review previous works by explaining the two instability mechanisms and the arguments proposed to distinguish them. On the basis of these works, the advective-acoustic cycle is generally favoured but the purely acoustic mechanism cannot be ruled out for realistic flow parameters. In Section~\\ref{sec:radial_time}, we propose a new method to extract a radial propagation time from the frequency of the unstable modes. For this purpose we compare the frequencies of several harmonics that differ only by their radial structure. As the radial advective-acoustic and purely acoustic times are significantly different, this allows us to distinguish between the two cycles and discards the purely acoustic mechanism. In Section~\\ref{sec:acoustic_modes},  we describe a method to compute purely acoustic modes, which are shown to be stable. In Section~\\ref{sec:details}, we give a detailed analysis of the frequency spectra of advective-acoustic and purely acoustic cycles, which confirms the validity of the method used in Section~\\ref{sec:radial_time}. Finally in Section~\\ref{sec:conclusion}, we summarise our work and conclude that the instability mechanism behind SASI is the advective-acoustic cycle. ",
        "conclusions": "\\label{sec:conclusion} Two new arguments have been presented which allow us to conclude that SASI cannot be explained by a purely acoustic mechanism, and should thus be interpreted as an advective-acoustic cycle. Contrary to previous works, these two arguments do not rely on a WKB analysis of the acoustic waves, and are therefore valid for realistic values of the shock radius.  First we proposed a method to extract a radial propagation timescale from the frequency difference between two successive SASI modes. This timescale agrees fairly well with the advective-acoustic time, but strongly disagrees with the radial acoustic time. This discards the purely acoustic cycle as the driving mechanism of SASI, and argues in favour of the advective-acoustic cycle. A detailed analysis of the frequency spectrum provided in Section~\\ref{sec:details} confirmed the validity of this argument.  Second, we described a new method to compute purely acoustic modes by artificially subtracting the advected wave created by the shock. All these acoustic modes were found to be stable, thus showing that the advected entropy-vorticity wave created by the oscillation of the shock plays a central role in the instability mechanism.  As a side product, we provided an analytical description of the frequency spectra of the advective-acoustic and purely acoustic cycles. This description was checked in the planar toy model of \\citet{foglizzo09} and the model of \\citet{blondin06}. The similarity between the frequency spectra in these two models, supports the relevance of the toy model for core collapse supernovae. Finally, we note that unexplained phase shifts during the coupling processes were observed in the second model. This calls for a better understanding of the coupling process in a cooling layer."
    },
    "1112/1112.1885_arXiv.txt": {
        "abstract": "The dense molecular clump P1 in the infrared dark cloud (IRDC) complex G28.34+0.06 harbors a massive protostellar cluster at its extreme youth. Our previous Submillimeter Array (SMA) observations revealed several jet-like CO outflows emanating from the protostars, indicative of intense accretion and potential interaction with ambient natal materials. Here we present the Expanded Very Large Array (EVLA) spectral line observations toward P1 in the \\nh3 (J,K) = \\11, \\22, \\33 lines, as well as \\water\\ and class I \\meth\\ masers. Multiple \\nh3 transitions reveal the heated gas widely spread in the 1 pc clump. The temperature distribution is highly structured; the heated gas is offset from the protostars, and morphologically matches the outflows very well. Hot spots of spatially compact, spectrally broad \\nh3 \\33 emission features are also found coincident with the outflows. A weak \\nh3 \\33 maser is discovered at the interface between an outflow jet and the ambient gas. These findings suggest that protostellar heating may not be effective in suppressing fragmentation during the formation of massive cores. ",
        "introduction": "The majority of stars, from low to high mass ones are born in clusters \\citep{Lada2003}. In a clustered environment, feedback from protostellar radiation and outflows can play an important role in shaping their parental molecular cloud and thereby regulate further cloud fragmentation and star formation. In particular, protostellar feedback may be crucial for the formation of massive stars which typically contain masses at least one order of magnitude higher than the global Jeans mass in the molecular clump ($M_J \\sim 1$ \\msun\\ for $n_{\\rm H} \\sim 10^5$ \\cm3 and $T \\sim 15$ K; \\citealt{pillai06, rath06}). Calculations have shown that the accretion luminosity of a low mass protostar can heat gas to 100 K out to hundreds of AU, thereby raising the local thermal Jeans mass \\citep{Masunaga2000, Krumholz2006}. Subsequent numerical simulations indicate that protostellar heating from low mass stars in the cluster may change the equation of state \\citep{Larson1985}, and suppress the fragmentation \\citep{Krumholz2010b}, thereby aiding the formation of massive cores. On the other hand, protostellar outflows in the cluster can also inject momentum and energy into the local medium, and affect the evolution of the cloud \\citep{Quillen2005,Curtis2010,WangP2010}.   To investigate the protostellar feedback, we observed the dense molecular clump P1 in the IRDC complex G28.34+0.06. Located in the southern part of G28.34+0.06, P1 is one of the dense molecular clumps among the filamentary network in the IRDC complex which stretches over 6 pc at a kinematic distance of 4.8 kpc \\citep{Carey2000, rath06, wy08}. P1 is a northeast-southwest (NE-SW) orientated molecular clump containing a large mass reservoir of $10^3$ \\msun\\ within a size of 1 pc. \\cite{wy08} presented an \\nh3 mosaic of the entire IRDC, and found that the gas temperature decreases from 20 to 14 K toward the inner part of P1, indicative of external heating. A similar trend was found in the \\nh3 line width, suggesting a turbulence dissipation toward smaller scales. Our SMA observations resolved the P1 clump into a string of five cores (named SMA1--5) containing 22--64 \\msun\\ in mass and $\\lsim 0.1$ pc in size; at higher angular resolution and better sensitivity, the cores are further resolved into ten compact condensations (named, e.g. SMA2a--d) with masses of 1.4--10.6 \\msun\\ and sizes of $\\sim$0.01 pc, indicative of a hierarchical fragmentation \\citep{qz09,me11}. In addition, all five cores are associated with highly collimated CO (3--2) outflows. The large mass reservoir and the ongoing accretion indicate that P1 will eventually form a massive star cluster after the accretion is completed. This young protocluster provides an opportunity to investigate how protostellar heating and outflows affect the parent molecular environment, and the formation of massive cores. The new EVLA observations presented in this Letter reveal the temperature profiles at unprecedented high angular resolution, allowing us to evaluate the protostellar feedback in detail. ",
        "conclusions": "\\label{sec:discuss}  \\subsection{Outflow Heating} \\label{sec:heat}  Intensity ratios of the main and hyperfine components from multiple \\nh3 (J,K) transitions constrain rotational temperature which well approximates {kinetic} temperature in the regime of $<$20 K \\citep{Ho1983, Walmsley1983}. In the following analysis, we assume that the rotational temperature is the same as the kinetic temperature when the former is $<$20 K. Following the procedure of \\cite{Ho1983}, we compute a rotational temperature (\\t21) map based on the \\nh3 \\11 and \\22 images. The \\nh3 images were first smoothed to a common beam of $2''.8$ before computing the \\t21 map. The uncertainty of \\t21 is estimated to be 3 K. Figure \\ref{fig:heat}(a) shows \\t21 with the dust condensations and the CO (3--2) outflows superposed. The \\t21 map reveals a highly structured temperature profile ranging from  8 to 30 K, opposite to the smooth temperature profiles observed in other IRDCs \\citep{Ragan11}. This may be due to beam dilution since the data in \\cite{Ragan11} have a larger beam. A number of remarkable enhancements ($T_{21}>26$ K) are present among the entire map, but surprisingly none of the temperature enhancements coincide with any dust condensations. Instead, all the \\t21 enhancements spatially match the footprint of the outflows.  Since the \\nh3 \\33 is more widely spread than the \\11 toward the northwest where the signal-to-noise ratio of \\11 prohibits us from computing opacity, it is difficult to derive a reliable rotational temperature map from the \\nh3 \\11 and \\33 lines. Instead, we compute the intensity ratio of the main hyperfine of \\33/\\11, \\r31, as plotted in Figure \\ref{fig:heat}(b). To compute the intensity ratio, the \\nh3 \\11 image was first smoothed to the beam of the \\33 image. The intensity ratio is positively correlated with the rotational temperature \\citep{Ho1983, qz02}. For reference, assuming optically thin emission and the same abundance for the ortho and para \\nh3 \\citep{qz07}, $R_{31}=0.5$, 2, and 5 correspond to rotational temperatures of $T_{31} = 31$, 55, and 109 K, respectively. Like the \\t21 map, the \\r31 map is also highly structured. \\r31 varies continuously from 0.3 to 12.5 over the map. Again several hot spots represented by \\r31 peaks (corresponding to \\nh3 \\33 emission peaks A--E) match the outflow footprint and are offset from the dust condensations. In addition to the peaks, there is a low \\r31 valley along SMA2 to SMA3. Optically thin thermal \\nh3 \\33 emission should have \\r31 $\\leqslant 6.4$ \\citep{Ho1983}. Higher \\r31 values are seen toward peaks A--C, indicative of non-thermal emission. Toward peak B, the (3,3) emission is weakly amplified (Figure \\ref{fig:spec}), resembling an \\nh3 \\33 maser similar to those found in other massive star formation regions where outflows are present \\citep[e.g.][]{Mangum1994, qz95, Beuther2007, Brogan2011}. However the spatial resolution is insufficient to constrain the brightness of this emission feature.   The morphological match between the heated gas and the outflow footprint suggests that the heating in the P1 region is due to protostellar outflows. Indeed, all the hot spots exclusively coincide with large \\nh3 \\33 line widths compared to relatively small line width toward the dust condensations (Figure \\ref{fig:spec}), likely indicative of gas motion and extra turbulence injected by the outflow shocks. The outflow heating has two important features. First, the heating alters thermal property of the gas throughout the clump, from a distance 0.1 pc up to 0.5 pc away from the driving protostars. Second, the heating is localized and far from well mixed. The temperature at the low \\r31 valley is not high enough to suppress further fragmentation into Jeans mass of 1 \\msun\\ \\citep{me11}. We probably have witnessed through EVLA and SMA the very beginning of an outflow induced gas heating.    \\subsection{Radiative Heating}  It is apparent that protostellar outflows can introduce significant heating to the ambient gas. What is the effect of protostellar radiation? Is radiative heating from low mass protostars sufficient to suppress fragmentation and aid the formation of massive cores? The fact that all dense cores correspond with low \\r31 is intriguing. In addition, the temperature map in Figure \\ref{fig:heat}(a) indicates that dense cores have typical temperatures of 18 K, and are offset from the high temperature peaks. The temperature distributions suggest that the heating of low mass protostars may not play a significant role in the formation of dense cores detected with the SMA. Otherwise, massive cores should be found at the peaks of the temperature distribution. \\cite{Longmore2011} made a similar statement for a hot molecular core G8.68--0.37.  It is also interesting to note that the protostars embedded in the SMA cores have not produced significant heating. SMA1 to SMA4 have a rather uniform \\t21 of 18 K, while SMA5 has a lower \\t21 of 9 K. We caution that the \\t21 map has a $2''.8$ beam, so any temperature structure has been smoothed to a spatial scale of 0.065 pc. How much luminosity is required to heat the gas to 18 or 9 K at a radius of 0.065 pc from the central source? Assuming thermal equilibrium between gas and dust, the luminosity can be estimated following the equation \\citep{scoville76,garay99}: $$ T_\\mathrm{dust} = 65 \\left( \\frac{0.1\\, \\mathrm{pc}}{r} \\right) ^{2/(4+\\beta)} \\left( \\frac{L_\\mathrm{star}}{10^5 L_\\odot} \\right) ^{1/(4+\\beta)} \\left( \\frac{0.1}{f} \\right) ^{1/(4+\\beta)} \\,, $$ where $\\beta = 1.5$ is the dust emissivity at far infrared wavelengths \\citep{rath10}, $f = 0.08 \\, \\rm{cm}^2 \\, \\rm{g}^{-1}$ is the dust emissivity at 50 \\um\\ \\citep{scoville76}, and $r = 0.065$ pc is the radius from the central protostar. We estimate a luminosity of 30 \\lsun\\ for each protostar embedded in SMA1, SMA2a, SMA3, and SMA4a, and 1 \\lsun\\ for the protostar SMA5. These values should be taken with caution because the beam dilution may introduce large uncertainty in both temperature and radius. Nevertheless, we note that the luminosity of SMA2a is in rough agreement with the luminosity inferred from SED for the 24 \\um\\ source in the SMA2 vicinity \\citep[$\\sim$10$^2$ \\lsun,][]{wy08}. A lack of significant heating from the growing protostars is consistent with non detection of line emission from organic molecules \\citep{qz09,me11}.    \\subsection{Origin of the Minor Filament} \\label{sec:mag}  The origin of the low-density minor filament is rather interesting. A low-density streak perpendicular to a high-density oblate region is expected if magnetic field is dynamically important compared to self-gravity and turbulence during the formation of the system \\citep{hb11}. In this case, gravitational contraction along the field lines is more efficient, resulting in a high-density region perpendicular to the magnetic field, while on the other hand turbulence is also channeled by the field and gives rise to streaks aligned with the magnetic field in the lower-density region. This phenomenon has been seen in simulations \\citep{Price2008,Nakamura2008} and in observations of NGC 2024 in Orion \\citep{hb11}. In G28.34-P1, the $X$-shaped filamentary system projected on the sky is almost perpendicular. If the system were formed from the same mechanism, one should expect a mean magnetic field along the minor filament and perpendicular to the main filament. Based on the JCMT dust polarization data cataloged by \\cite{Matthews2009}, we derived a mean B-field direction of $\\mathrm{P.A.} = -25^{\\circ}.6 \\pm 45^{\\circ}$ as labeled on Figure \\ref{fig:mnt}. Because the main filament has a P.A. of $48^{\\circ}$ \\citep{me11}, the separation angle between the mean B-field and the main filament is $73^{\\circ}.6$. The mean magnetic field at the clump scale is indeed along with the minor filament and roughly perpendicular to the main filament. The large deviation in the B-field direction may indicate that the strong outflows have started to disturb the initially ordered field lines in the inner part of P1 \\citep{hb09}. Therefore we speculate that the minor filament originates from turbulence anisotropy during the formation of the $X$-shaped filamentary system, and that the formation process has been governed by the interplay between strong magnetic field, self-gravity, and turbulence."
    },
    "1112/1112.1627_arXiv.txt": {
        "abstract": "\\vspace*{5mm}  We derive constraints on the annual modulation signal in Dark Matter (DM) direct detection experiments in terms of the unmodulated event rate. A general bound independent of the details of the DM distribution follows from the assumption that the motion of the earth around the sun is the only source of time variation. The bound is valid for a very general class of particle physics models and also holds in the presence of an unknown unmodulated background. More stringent bounds are obtained, if modest assumptions on symmetry properties of the DM halo are adopted. We illustrate the bounds by applying them to the annual modulation signals reported by the DAMA and CoGeNT experiments in the framework of spin-independent elastic scattering. While the DAMA signal satisfies our bounds, severe restrictions on the DM mass can be set for CoGeNT. ",
        "introduction": "\\label{sec:intro}  A smoking-gun signature of the dark matter (DM) signal in DM direct detection experiments is the annual modulation of the event rate. This arises because the earth rotates around the sun, while at the same time the sun moves relative to the DM halo~\\cite{Drukier:1986tm, Freese:1987wu}. Currently, two experiments report annual modulation of the signals, DAMA with 8.9$\\sigma$ significance~\\cite{Bernabei:2008yi}, and CoGeNT with a significance of $2.8\\sigma$~\\cite{Aalseth:2011wp}.  An important question is whether the observed modulation rates are consistent with the interpretation that the signal is due to DM.  Ideally, one would like to answer this question without requiring a detailed knowledge of the DM halo. As we will show below, we come very close to this ideal.  The DM scattering rate in the detector is determined by the particle physics properties of DM and by the astrophysical properties of the DM halo. The latter carry large uncertainties. When presenting results of DM direct detection experiments, it has become customary to use a ``standard\" Maxwellian DM velocity distribution in order to show the constraints on DM mass and scattering cross section. This is very likely an oversimplification, with N-body simulations indicating a more complicated structure of the DM halo, see e.g., \\cite{Kuhlen:2009vh}.  It is well known that the annual modulation signal is sensitive to such deviations from the Maxwellian halo \\cite{Fornengo:2003fm, Green:2003yh,Fairbairn:2008gz}. DM streams are particularly relevant for the annual modulation signal and can lead to large effects. These can, for instance, be comparable in size to the modulation effect due to the whole Maxwellian halo \\cite{Gelmini:2000dm, Savage:2006qr, Chang:2008xa, Natarajan:2011gz}.  In the following we will derive constraints on the amplitude of the annual modulation using measured unmodulated rates.  The derivation of these bounds is made possible by the fact that the velocity of the earth rotating around the sun, $v_e\\sim 30$ km/s, is much smaller than the typical velocity of DM particles in the halo, $\\langle v \\rangle \\sim 200$~km/s, and we can expand the velocity integral to first order in $v_e$. The derived constraints on the annual modulation then serve as a consistency check for the hypothesis that the annual modulation is due to the DM signal. The check that we provide is (almost) independent of DM halo properties.  Previously halo independent methods to interpret DM direct detection data have been developed and applied in \\cite{Drees:2007hr, Drees:2008bv, Fox:2010bz, Fox:2010bu, McCabe:2011sr, Frandsen:2011gi}, with different goals than ours. In ref.~\\cite{Drees:2008bv} a method for determining the DM mass independently of the halo was developed, while refs.~\\cite{Drees:2007hr, Frandsen:2011gi} focused on extracting halo properties from the data. Compatibility studies of different experiments without adopting assumptions on the halo have been performed in~\\cite{Fox:2010bz, Fox:2010bu, McCabe:2011sr}. Our goal is to obtain a consistency check between the modulation amplitude and the unmodulated signal within a given experiment.  The structure of the paper is as follows: in sec.~\\ref{sec:notation} we set up the notation, followed in sec.~\\ref{sec:general} by a derivation of the general bound on the modulation signal under the very mild assumption that the properties of the DM halo in our vicinity do not change on time scales of months. In sec.~\\ref{sec:spec} we impose further symmetry requirements on the DM velocity distribution and obtain successively more stringent bounds. In sec.~\\ref{sec:data} we then demonstrate the usefulness of our bounds by applying them to the modulation signals reported in CoGeNT and DAMA within elastic spin-independent scattering. We find that our method leads to non-trivial restrictions on DM interpretations of the modulation signal in CoGeNT whereas the DAMA signal satisfies our bounds for the relevant DM masses. We summarize in sec.~\\ref{sec:conclusion}. In appendix~\\ref{sec:maxw} we consider the Maxwellian DM halo and show that the expansion in $v_e$ is rather accurate in this case. Furthermore, we illustrate the various bounds compared to the modulation signal expected for the Maxwellian halo. Appendix~\\ref{sec:translation} provides technical details necessary to translate the bounds derived for the halo velocity integral into observable event rates. Supplementary material related to the treatment of an unknown background is given in appendix~\\ref{app:proc2}. ",
        "conclusions": "\\label{sec:conclusion}  The annual modulation is probably the most distinctive signature of dark matter and is playing (and will certainly play) a central role in revealing its existence and nature. However, it is not always clear that the modulation detected by a DM direct detection experiment is caused by the dark matter particles and not by any other unknown background source. In this work we have presented a consistency check for the amplitude of the modulation and the unmodulated count rate, which is a necessary (but not sufficient) condition for DM being the origin of an observed annual modulation.  We have derived upper bounds on the energy integrated modulation amplitude in terms of the unmodulated rate by expanding the DM velocity integral to first order in the earth's velocity $v_e \\approx 30$~km/s. It holds for a wide class of particle physics models where the differential scattering cross section $d\\sigma/dE_{nr}$ is proportional to $1/v^2$. Although we have only focused on elastic scattering the method can also be generalized to the inelastic case. The important aspect of our work is that our bounds hold for very general assumptions about the DM velocity distribution $f(\\vect{v})$ in the sun's vicinity. We have presented bounds under the hypothesis of a single DM species and the following assumptions on the DM halo:  \\begin{itemize} \\item {\\bf Assumption 1, bound in eq.~\\eqref{eq:bound_gen}:} We assume that the only time dependence is induced by the rotation of the earth around the sun. The halo itself is static on the time scale of months to years and spatially constant at the scale of the sun--earth distance. Otherwise the DM velocity distribution can have an arbitrary structure, including, for example, several streams coming from various directions.  In order to saturate this bound the halo has to be very peculiar, with rather unrealistic properties.  \\item {\\bf Assumption 2, bound in eq.~\\eqref{eq:bound_spec}:} In addition to Assumption~1 we assume that there is just one preferred direction of the DM velocity distribution in the rest frame of the sun (independent of the minimal velocity $v_m$ in the halo integral). This requires certain symmetries of the velocity distribution, which are specified in eqs.~\\eqref{eq:cond2} and \\eqref{eq:cond3}. It covers typically situations where the DM distribution is dominated by one single component, which may come from an arbitrary (but constant) direction. Assumption~2 requires that the phase of the modulation is independent of the recoil energy (up to a phase flip of half a year).  \\item {\\bf Assumption 2a, bound in eq.~\\eqref{eq:bound_spec_alpha}:} We require that the preferred direction from Assumption~2 is aligned with the motion of the sun. This is fulfilled for any isotropic halo, and also for tri-axial halos up to corrections due to the peculiar velocity of the sun.  Furthermore, it includes the possibility of streams parallel to the motion of the sun, such as a possible dark disk. Assumption~2a requires that the maximum (or minimum) of the event rate is around June~2nd. \\end{itemize}  The theoretical bounds are obtained in terms of the minimal velocity $v_m$ and have to be related to observable quantities by translating into recoil energy. We have outlined two possible procedures for this task in section~\\ref{sec:data} for the case of elastic scattering, taking into account statistical errors and the possibility that an unknown background may contribute to the unmodulated rate, but that it has no time-dependence. As an example we have applied the proposed consistency checks to the annual modulation signals reported by the CoGeNT and DAMA experiments.  While DAMA data are compatible with a dark matter origin for its modulation, severe restrictions on the dark matter mass can be set for the case of CoGeNT. Applying our bounds we find that the CoGeNT modulation amplitude can be consistent with the unmodulated rate at the $68\\%$~CL only for DM masses $m_{\\chi} \\gtrsim 7.3, 18, 16$~GeV for the Assumptions~1, 2, 2a, respectively. If preliminary results on a possible surface events contamination of the unmodulated rate at low energies in CoGeNT are confirmed, those bounds would become even more restrictive. In this case, CoGeNT modulation data would be inconsistent with the DM hypothesis under assumptions 2 and 2a at about 97\\% and 90\\% CL, respectively. DAMA has a relatively large unmodulated count rate of about 1~cnt/kg/day/keV compared to the modulation amplitude of about 0.02~cnt/kg/day/keV. Therefore, our bounds are not very stringent and the modulation amplitude is consistent with the unmodulated rate. In contrast, CoGeNT reports a modulation amplitude in the range of 0.5 to 1~cnt/kg/day/keV, compared to an unmodulated rate of 3 to 4~cnt/kg/day/keV. Because of this relatively large ratio between modulated and unmodulated rate our method provides stringent constraints in the case of CoGeNT. Several comments are in order:  \\bigskip   $(i)$ In deriving the bounds we assume a certain smoothness of the halo, since we are expanding in $v_e$. Spikes in the velocity distribution much narrower than 30~km/s are not covered by our procedure.  $(ii)$ Our bounds assume that a modulation signal is present in the data. If the significance of the modulation is weak, the bounds are more easily satisfied. This effect is explicitly illustrated in sec.~\\ref{sec:data} by using CoGeNT data under different assumptions regarding the modulation signal. The method discussed here cannot be used to establish the presence of a modulation, it can only test whether a given modulation signal is consistent with the unmodulated rate.  $(iii)$ In this work we have used the relation between the minimal velocity $v_m$ and nuclear recoil energy $E_{nr}$ implied by elastic scattering, see eq.~\\eqref{eq:vm} and appendix~\\ref{sec:translation}. The bounds can be generalised in a straightforward way also to inelastic scattering.  However, in that case $v_m = (m_A E_{nr} / \\mu_{\\chi A} + \\delta)/\\sqrt{2m_A E_{nr}}$ has a minimum in $E_{nr}$, and therefore there is no one-to-one correspondence between $v_m$ and $E_{nr}$. When translating from $v_m$ to $E_{nr}$ one has to take into account that different disconnected regions in $E_{nr}$ can contribute to a given interval in $v_m$.  $(iv)$ The type of DM--nucleus interaction (i.e., the particle physics) has to be specified before applying our bounds. Apart from the $1/v^2$ dependence of the differential cross section $d\\sigma/dE_{nr}$, the $E_{nr}$ dependence of the interaction is encoded in the form factor $F(E_{nr})$. It can describe conventional spin-independent or spin-dependent interactions, but also a possible non-trivial $q^2$-dependence of the DM--nucleus interaction, all of which can be absorbed into $F(E_{nr})$.  \\bigskip  To conclude, we have presented a consistency check for the amplitude of a DM induced annual modulation compared to the unmodulated event rate in a given DM direct detection experiment. Our bounds rely only on very mild and realistic assumptions about the DM halo. We believe that the proposed method is a useful check which any annual modulation signal should pass. A violation of our bounds suggests a non-DM origin of the annual modulation, or requires rather exotic properties of the DM distribution, for example very sharp spikes and edges. Such extreme features of the halo should have other observable consequences, such as surprising spectral features or strong energy dependence of the modulation phase. Such features could be used as additional diagnostics, beyond the signatures explored here, which are restricted to the energy integrated modulation amplitude, irrespective of the phase."
    },
    "1112/1112.1411_arXiv.txt": {
        "abstract": "We use $\\sim$12,000 spectra of $\\sim$3,500 magnetically active M0-M9 dwarfs from the Sloan Digital Sky Survey taken at 10-15 minute intervals, together with $\\sim$300 spectra of $\\sim$60 M0-M8 stars obtained hourly with the Hydra multi-object spectrometer, to probe H$\\alpha$ variability on timescales of minutes to weeks.  With multiple observations for every star examined, we are able to characterize fluctuations in H$\\alpha$ emission as a function of activity strength and spectral type. Stars with greater magnetic activity (as quantified by L$_{\\rm{H}\\alpha}$/L$_{bol}$) are found to be less variable at all spectral types. We attribute this result to the stronger level of persistent emission in the high activity stars, requiring a larger heating event  in order to produce measurable variability. We also construct H$\\alpha$ structure functions to constrain the timescale of variability. The more active objects with lower variability exhibit a characteristic timescale longer than an hour, likely due to larger, longer lasting heating events, while the less active objects with higher variability have a characteristic timescale shorter than 15 minutes. ",
        "introduction": "\\label{sec:intro}  M dwarfs, the least massive and most numerous stars in our Galaxy, provide remote laboratories for observing high-energy magnetic phenomena in stellar atmospheres.  The red and dim nature of M dwarfs makes the H$\\alpha$ spectral feature the most accessible tracer of magnetic activity in their convective outer layers \\citep{Bopp1978,Walkowicz2009}.  M dwarfs with H$\\alpha$ in emission are classified as magnetically active, with a greater fraction of stars showing H$\\alpha$ emission at later spectral types \\citep{Hawley1996,Gizis2000,West2011}.  Magnetic activity is correlated with energetic flaring events that last seconds to hours \\citep{moffett74,kowalski2010} and photometric brightness increases as large as 8.2 magnitudes in the \\emph{B}-band \\citep{Almeida2011}. The incidence of H$\\alpha$ emission and large flares decreases for stars farther from the Galactic Plane \\citep{West2006,West2008,kowalski2009,Hilton2010, 2011AJ....141...50W}.  This is interpreted as an age effect, leading to activity timescales of 1--8 billion years for early--late M dwarfs respectively.  Variability in quiescent (non-flaring) H$\\alpha$ emission has been attributed to both stellar rotation and M dwarf activity cycles. \\citet{Bopp1978} related sequential H$\\alpha$ measurements to orbital phase and suggested that the observed changes were caused by the rotation of localized active regions in and out of view on timescales of days.  Activity cycles with periods of a few years have been identified for M dwarfs by \\citet{Cincunegui2007b} and \\citet{Buccino2011}, and may affect the observed strength of H$\\alpha$ emission.  Although large flares are relatively infrequent events \\citep{Hilton2011}, sporadic low-level magnetic field reconnection events (``microflares'') may also contribute to inherent H$\\alpha$ emission variations in active M dwarfs.  There have been some recent investigations of H$\\alpha$ variability on timescales of minutes to days.  \\citet{Gizis2002} measured the H$\\alpha$ equivalent widths (EWHA) for a sample of more than 600 active M0--M6 dwarfs from the Palomar/Michigan State University (PMSU) survey of nearby stars \\citep{reid1995}, and found that the variation in repeated measurements was near 15\\% for stars with mean EWHA, $\\langle$EWHA$\\rangle$ $<$ 5\\AA, while the variability for stars with $\\langle$EWHA$\\rangle$ $>$ 5\\AA\\ was in excess of 30\\%. In their study of EWHA variations on later type stars, \\citet{Lee2010} identified significant variability in $\\sim$80\\% of their sample of 43 M4--M8 stars.  They found that the majority of variability events had timescales of 30 minutes or longer, and that the later-type stars were more variable, as measured by $R$(EW) = max(EWHA)/min(EWHA). Similarly, \\citet{Kruse2010} found an increase in the same metric (which they refer to as $R_{EW}$) at later spectral types for timescales from 15 minutes to an hour in the Sloan Digital Sky Survey \\citep[SDSS;][]{york2000} Data Release 7 \\citep[DR7;][]{dr7} component spectra.  They also noted that $R_{EW}$ was larger for intermittent sources.  Our analysis uses a set of spectra with exposure times of $\\sim$1 hour from the multi-object Hydra spectrograph on the WIYN telescope at Kitt Peak National Observatory, as well as the component spectra from SDSS, to extend the evaluation of H$\\alpha$ variability using new metrics and a broader range of timescales.  We employ a construct known as the structure function, which is commonly used in quasar variability analyses \\citep[e.g.][]{Simonetti1985,MacLeod2011}, to further constrain the timescales of H$\\alpha$ emission variability.  The spectroscopic M dwarf samples we used are detailed in Section 2, along with our methods for determining spectral types and measuring H$\\alpha$ emission strength. Analysis of H$\\alpha$ emission variability with spectral type, activity strength and timescale is presented in Section 3, and we summarize our findings in Section 4. ",
        "conclusions": "\\label{sec:discussion}  \\begin{figure}[t] \\centering \\scalebox{0.5}{\\includegraphics*[0.05in,0.31in][6.94in,4.67in]{fig8}} \\caption{The structure functions for a) the high-variability (less active) stars; and b) the low-variability (more active) stars, for the SDSS (black asterisks) and Hydra (blue filled circles) samples.  The low-variability stars show an increasing structure function amplitude between 0.4-1 hour, indicating a timescale of more than an hour for the variability.  The high-variability stars apparently vary on timescales shorter than our smallest time separation ($\\sim$15 minutes). The Hydra data are discussed separately in the text.} \\label{strfunc} \\end{figure}  We used two spectroscopic samples of active M dwarfs to examine H$\\alpha$ emission variability on timescales from minutes to weeks.  We found that the H$\\alpha$ variability measured using the $R_{EW}$ metric remains relatively constant as a function of spectral type, in contrast to previous results. However, using our fractional variability metric ($\\sigma_{EWHA}$/$\\langle$EWHA$\\rangle$), we showed that the higher-variability stars have relatively low activity strength  (L$_{\\rm{H}\\alpha}$/L$_{bol}$), and conversely the low-variability stars have high activity strength. This naturally leads to an apparent dependence on spectral type because the later type stars typically have lower activity strength. We speculated that the physical reason for the higher variability in the low activity stars is that small changes in H$\\alpha$ emission, e.g. due to small-scale flaring or the appearance of a new active region on the stellar surface, would have a relatively larger effect on the less active stars.  We investigated the timescales of H$\\alpha$ emission variability using structure functions.  The observed H$\\alpha$ variations for low-variability stars ($\\sigma_{EWHA}$/$\\langle$EWHA$\\rangle<$ 0.16) occur on a timescale longer than an hour, while the high-variability M dwarfs exhibit a timescale that is shorter than our sampling time of $\\sim$15 minutes. Neither the low nor high variability samples show significant changes in the structure function on timescales longer than a day. We speculated that the low-variability sample, which is mainly composed of the higher activity strength stars, requires more energetic and hence longer-lived emission events in order to be detected above the strong persistent surface activity, leading to a longer characteristic variability timescale.  Better time resolution in spectroscopic M dwarf monitoring, as well as a wider and better-sampled range of time separations, will allow for a more accurate determination of the timescales of H$\\alpha$ variability in low-mass dwarfs and may lead to a better understanding of the physical mechanisms that cause these changes.  We acknowledge NSF grant AST 08-07205, and thank David Schlegel for help obtaining the individual component SDSS spectra.  Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England.  The SDSS Web site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck Institute for Astronomy (MPIA), the Max-Planck Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "1112/1112.0417_arXiv.txt": {
        "abstract": "We derive the Layzer-Irvine equation in the presence of a homogeneous (or quasi-homogeneous) dark energy component with an arbitrary equation of state. We extend the Layzer-Irvine equation to homogeneous and isotropic universes with an arbitrary number of dimensions and obtain the corresponding virial relation for sufficiently relaxed objects. We find analogous equations describing the dynamics of cosmic string loops and other p-branes of arbitrary dimensionality, discussing the corresponding relativistic and non-relativistic limits. Finally, we generalize the Layzer-Irvine equation to account for a non-minimal interaction between dark matter and dark energy, discussing its practical use as a signature of such an interaction. ",
        "introduction": "The Newtonian energy conservation equation when generalized to an expanding cosmological background becomes \\be {\\dot E} + H(2K+U)=0,\\label{eq:LI} \\ee where a dot represents a total derivative with respect to physical time, $H$ is the Hubble parameter, $E=K+U$, with $K$ and $U$ being the peculiar kinetic and gravitational potential energies, respectively, of a system of non-relativistic particles interacting through gravity. This equation was derived independentely in \\cite{1961PhDT.........2I,1963ApJ...138..174L} and it is known as the cosmic energy or Layzer-Irvine equation (see also \\cite{citeulike:2083812}). Eq. (\\ref{eq:LI}) is valid throughout the entire process of structure formation and in the ${\\dot E}=0$ limit one recovers the usual virial relation $K=-U/2$ that holds for collapsed objects that have reached the state of hydrostatic equilibirum.  The Layzer-Irvine equation has been establishing itself as one of the most renowned equations of modern cosmology with its many applications including the determination of the matter density, cluster mass and size, and the galaxy peculiar velocity field \\cite{Davis:1997vg,citeulike:1284470,Fukugita:2004ee,Zaroubi:2004vu}. More recently, some authors \\cite{Bertolami:2007zm,Bertolami:2011yp,Abdalla:2007rd,Abdalla:2009mt} have been using the Layzer-Irvine equation as a tool to detect a possible non-minimal interaction between the dark matter (DM) and the dark energy (DE) which, together, account for approximately 96$\\%$ of the energy content of the Universe today \\cite{Amanullah:2010vv,Komatsu:2010fb} and whose fundamental nature is still largely unkown. The existence of such an interaction would in general invalidate the energy balance dictated by Eq. (\\ref{eq:LI}). Consequently, by measuring the properties of sufficiently relaxed structures, such as galaxy clusters, one may expect to be able to detect a signature of an interaction between dark matter and dark energy through deviations from the usual virial relation \\cite{Bertolami:2007zm,Abdalla:2007rd,Abdalla:2009mt,Bertolami:2011yp}.  In this paper, our main goal is to present a broad discussion of the Layzer-Irvine equation in a generalized framework (see also  \\cite{Shtanov:2010iy}). In Sec. \\ref{sec:LIa} we start by deriving the Layzer-Irvine equation in the presence of a homogeneous dark energy component, extending it to Friedmann-Robertson-Walker (FRW) cosmologies with more than three spatial dimensions. In Sec. \\ref{sec:PB} we show that the dynamics of cosmic string loops and other p-branes of arbitrary dimensionality are described by analogous equations, discussing the corresponding relativistic and non-relativistic limits. In Sec. \\ref{sec:INT} the Layzer-Irvine equation is generalized to the case where the dark matter is non-minimally coupled to the dark energy background and the implications of such coupling are discussed. Finnally we conclude in Sec. \\ref{conc}. ",
        "conclusions": "}  In this paper we studied the Layzer-Irvine equation and discussed some of its generalizations. In particular, we derived the Layzer-Irvine equation in the presence of a general homogeneous dark energy background showing that the final form of the equation is not affected explicitly by the dark energy component. We further generalized the equation and the virial relation to FRW  cosmologies with $N+1$ dimensions (with $N>2$). We have also demonstrated that the macroscopic dynamical energy equations of cosmic string loops and other p-branes of arbitrary dimensionality are, in the non-relativistic limit, analogous to the Layzer-Irvine equation. Finally, we generalized the Layzer-Irvine equation to account for a non-minimal interaction between dark matter and a homogeneous dark energy form. We have shown that, in general, gravitationally bound systems are not expected to reach hydrostatic equilibrium in the presence of a coupling between these two components. This constrasts with the usual assumption in the literature where the equilibrium relation ${\\dot E}=0$ is assumed a priori. Hence, a non-minimal coupling between dark matter and dark energy will generally lead to the breakdown of the usual virial relation $K = -U/2$, providing a crucial signature of such an interaction."
    },
    "1112/1112.3658_arXiv.txt": {
        "abstract": "\\indent A code, \\Epic, has been developed which computes, in the two-dimensional case, the initially circular orbits of guiding centra in an arbitrary axisymmetric potential with an arbitrary, weak perturbing potential in solid body rotation. This perturbing potential is given by its Fourier expansion. The analytic solution solves the linear epicyclic approximation of the equations of motion. To simulate the motion of interstellar matter and to damp the Lindblad resonances, we have in these equations introduced a friction which is proportional to the deviation from circular velocity. The corotation resonance is also damped by a special parameter. The program produces, in just a few seconds, orbital and density maps, as well as line of sight velocity maps for a chosen orientation of the galaxy.  We test \\Epic~ by comparing its results with previous simulations and observations from the literature, which gives satisfactory agreement. The aim is that this program should be a useful complement to elaborate numerical simulations. Particularly so are its abilities to quickly explore the parameter space, to construct artificial galaxies, and to isolate various single agents important for developing structure of interstellar matter in disc galaxies. ",
        "introduction": "Since the emergence of high spatial resolution two-dimensional velocity field for galaxies, the interpretation of observed velocity fields has become one of the key concerns of the extragalactic scientific community (e.g., \\citealt{Boulesteix1987}; \\citealt{Adam1989}; \\citealt{Bacon1992}; \\citealt{Bacon2001}; \\citealt{Allington1997}; \\citealt{Hernandez2003}). Of particular importance is the class of galaxies which host bars and/or spiral arms, as they are estimated to be the most abundant types of galaxies in the local Universe \\citep{Loveday1996}.  The theoretical foundation for understanding the kinematics and dynamics of spiral galaxies has been worked out over the last century. % Bertil \\citet{Lindblad1963,Lindblad1964} conceived a picture of circulation of stars between the spiral arms of a quasi-steady rotating spiral potential. Simultaneously, \\citet{LinShu1964} presented their density wave theory approaching the problem from a different theoretical point of view. Subsequently, \\citet{Shu1973} and \\citet{Roberts1979} derived solutions for the circulation of the interstellar medium through such a spiral potential and predicted large scale galactic shocks along the spiral arms. These processes make bars and/or spiral arms potential actors to redistribute angular momentum which will lead to dramatic effects such as intense star formation, build up of bulges, or onset of nuclear activity. In turn, the host galaxies may undergo dramatic secular evolution on time-scales compared to the galaxy dynamical time (see \\citealt{Kormendy2004} for a comprehensive discussion).  Beginning with \\citet{Holmberg1941} in the pre-computer era and in the 1950s with electronic computers (e.g. P.O. \\citealt{Lindblad1960}), a large number of simulations of development and evolution of structure in galaxies have been made with a variety of computer codes. These studies have often aimed at establishing the details of the flow of gas in and around bars (eg. \\citealt{Athanassoula1992b}), as well as the mechanisms which trigger starbursts and nuclear activity in the centers of galaxies (e.g., \\citealt{Shlosman1989}). However, they often involve extensive computational power leading to relatively long computing times. %   Several successful attempts have been carried out in order to deliver analytic solutions for observable parameters for some of the effects involved in the above questions (Sakhibov \\& Smirnov 1989; \\citealt{Canzian1993}; \\citealt{Wada1994}; \\citealt{Lindblad1994}; \\citealt{Schoenmakers1997}; \\citealt{Wong2004}; \\citealt{Fathi2005}; \\citealt{Byrd2006}; \\citealt{vandeVenfathi2010}), though they have all exclusively treated one specific morphological feature in each analytic model. While a linear analysis can be a guide for a physical understanding of complicated numerical results from simulations, in these analytic studies, bars and spiral arms have been accounted for separately, and interpretations have been based on marginalizing the effect of the other features. A notable difference between these results and numerical simulations is that the stepwise integration in a detailed simulation requires appreciable computing times, while for analytical calculations the computing time is negligible.   In this paper, we present an analytic solution within the epicyclic approximation, in which we introduce an arbitrary gravitational potential % and a damping coefficient for adequate appearance of the corotation resonance. Our solution has been coded as a development of the program \\textsc{Epic} \\citep{Lindblad1994} which was based on the first order epicyclic approximation. The current code, \\Epic, derives the response of interstellar matter, originally in circular orbits, to the gravitational potential % in a matter of few seconds, which makes it an efficient code for surveying a large parameter space. In section \\ref{sec:method}, we present the analytical solution for the epicyclic theory with a generic potential, and we describe our code \\Epic. We compare results generated by \\Epic~ with \\citet{Wada1994} in section \\ref{sec:models}, and in section \\ref{sec:new}, we demonstrate a standard case of a barred galaxy. In section \\ref{sec:1365} we compare results obtained by \\Epic~with observations and simulations made for the galaxy NGC~1365. Finally, we conclude in section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  \\Epic~cannot fully replace elaborate hydrodynamical simulations. \\Epic~ solves the equations of motion in a solid body rotating, time independent galactic potential. This is done in the first order epicyclic approximation. \\Epic~ computes the deformation of the initially circular orbits of the guiding center, the velocity fields, and the structure created. \\Epic~ pictures a steady state created by a weak perturbing potential, it does not give the time evolution of galactic structure. Due to the degeneracy of the strength of the gravitational potential and damping coefficients, we can only derive relative values of the density.  Our comparison with a more elaborate simulation, as well as observations, show that it still can reproduce structural and dynamic properties. One of the advantages with \\Epic, involving an analytic solution, is the speed by which it can survey the parameter space. Further, the possibilities, by varying parameters, to isolate artificially various agents important for forming structure of galaxies will be of importance for understanding the formation of galactic structure. \\Epic~will be a useful complement to time consuming, elaborate simulations.        \\vspace{5mm}  We would like to thank to J. A. Sellwood and R. Z\\'anmar S\\'anchez for letting us use their H$\\alpha$ data of NGC~1365. We also want to thank to the reviewer, Gene Byrd, for his useful comments. NP-F acknowledges financial support from NOTSA, and KF is supported by the Swedish Research Council (Vetenskapsr\\aa det)."
    },
    "1112/1112.1377_arXiv.txt": {
        "abstract": "We investigate the levels of small scale structure in surface brightness images of the core of the X-ray bright cool-core galaxy cluster AWM\\,7. After subtraction of a model of the smooth cluster emission, we find a number of approximately radial surface brightness depressions which are not present in simulated images and are seen in both the \\emph{Chandra} and \\emph{XMM-Newton} data.  The depressions are most strongly seen in the south of the cluster and have a magnitude of around 4 per cent in surface brightness. We see these features in both an energy band sensitive to the density (0.6 to 5~keV) and a band more sensitive to the pressure (3.5 to 7.5~keV).  Histograms of surface brightness in the data, when compared to realisations of a smooth model, reveal stronger surface brightness variations. We use the $\\Delta$-variance technique to characterise the magnitude of the fluctuations as a function of length scale. We find that the spectrum in the 0.6 to 5 keV band is flatter than expected for Kolmogorov index fluctuations. If characterised by a power spectrum, on large scales it would have an index around $-1.7$, rather than $-3.7$. The implied 3D density fluctuations have a standard deviation of around 4 per cent. The implied 3D pressure variations are at most 4 per cent.  Most of the longer-scale power in the density spectrum is contributed by the southern half of the cluster, where the depressions are seen. The density variations implied by the spectrum of the northern sector have a standard deviation of about 2 per cent. ",
        "introduction": "The velocity structure and degree of turbulence in the intracluster medium (ICM) of galaxy clusters is a matter of much debate as there are few direct observational constraints on its dynamical properties. The nonthermal contributions to gas pressure, if significant, are important to include for the determination of galaxy cluster masses from X-ray data.  The viscosity and turbulence of the ICM have a bearing on how energy can be dissipated within the ICM, which connects to how the AGN can feed energy back into its environment \\citep[e.g.][]{Kunz11}.  Theoretical calculations predict that the fraction of pressure support in the ICM from gas motions, generated by cluster growth and mergers, ranges from 5 to 20 per cent in the centres, increasing at larger radii (e.g. \\citealt{Lau09, Vazza09, Vazza11}). Relaxed clusters have significantly lower amounts of turbulence than those merging or have recently merged, with turbulent pressure fractions of less than 10 per cent. However, the amount of turbulence in simulations depends crucially on transport processes such as viscosity or magnetohydrodynamical effects. The active galactic nuclei (AGN) in the centres of clusters are also often seen to be affecting the surrounding ICM, injecting jets and bubbles into the cluster \\citep[e.g.][]{McNamaraNulsen07} and likely causing gas motions. These generated motions have been predicted to range from 500 to $1000\\kmps$ \\citep{Bruggen05,Heinz10}.  One direct method of measuring or limiting gas motions in galaxy clusters is to examine the width of X-ray emission lines using the X-ray reflection grating spectrometer (RGS) on \\emph{XMM-Newton}. We recently placed a limit on ICM random gas motions in the core of Abell 1835 of $274 \\kmps$, at the 90 per cent level \\citep{Sanders10_A1835}. The implied limit on the turbulent to thermal energy density is less than 13 per cent. For a sample of objects we were able to find at least 15 sources with less than 20 per cent of the thermal energy density in turbulence \\citep{Sanders10_Broaden}. This technique is, however, limited to sources with compact X-ray emitting cores and we can only examine the innermost regions. Other techniques to obtain limits or measure the velocity structure in clusters or elliptical galaxies includes examining the degree of resonance scattering in high resolution X-ray spectra \\citep{Werner09}, which imply strong upper limits on the gas motions.  Each of the above probes for gas motions indicates rather small amounts of gas motion in the cluster cores. There is additional indirect evidence that gas motions in the ICM are not strong in relaxed objects. These include the smooth changes in velocity of the filaments in the core of the Perseus cluster \\citep{HatchPer06}, which are apparently tracing gas motion \\citep{FabianPerFilament03} and their linear morphology.  Comparisons of X-ray and optical gravitational profiles in elliptical galaxies indicate additional non-thermal pressure contributions, some of which may be due to gas motions. The non-thermal pressure is around 10 per cent of the thermal gas pressure in M\\,87 and NGC\\,1399 \\citep{Churazov08} or 20 to 30 per cent in a sample \\citep{Churazov10}. In the Perseus cluster, \\cite{Churazov04} found evidence from the existence of resonance line emissions for gas motions at around half the sound speed.  \\cite{Schuecker04} noted that for isotropic turbulence, like the velocity distribution, the pressure energy spectrum should have a powerlaw dependence with wave number. For the pressure, $E_P(k) \\propto k^{-7/3}$ \\citep[see for example, section 6.9 in][]{LesieurTurbFluids08}, rather than the standard 5/3 Kolmogorov relation for velocity. In three-dimensions (3D) this corresponds to a pressure power spectrum with a spectral index of $-13/3$. \\cite{Schuecker04} examined the spectrum of the fluctuations in a pressure map of the Coma cluster. They found that the spectrum between length scales of 40 and 90 kpc was well described by a projected Kolmogorov type spectrum. Integrating the spectrum after deprojection implied a lower limit of 10 per cent of the total ICM pressure in turbulence.  \\cite{ChurazovComa11} recently reexamined \\emph{Chandra} and \\emph{XMM-Newton} data of the Coma cluster, finding that density variations in this object are rather small. The amplitude of the density fluctuations varies from 7 to 10 per cent on 500~kpc scales down to 5 per cent on scales of around 30~kpc.  In the interstellar medium, the electron density power spectrum as measured with interstellar scintillation is close to the value for a Kolmogorov process ($-11/3$) over at least five orders of magnitude \\citep[e.g.][]{Armstrong95}. The spectral slope when observing H\\textsc{i} 21cm line emission can vary, however. The value for the Small Magellanic Cloud is $-3.04$ \\citep{Stanimirovic99}, but measurements of the interstellar medium of dwarf galaxies gives slopes from $-1.5$ to $-2.6$ \\citep{Dutta09}.  Simulations of density fluctuations using hydrodynamics \\citep{Kim05} found a Kolmogorov slope for low Mach numbers, flattening as the Mach number increased. For 3D compressible magnetohydrodynamical simulations of turbulence \\cite{Kowal07} found that the energy spectrum of density fluctuations with subsonic turbulence is consistent with $-7/3$ for strong magnetic fields and the Kolmogorov index of $-5/3$ for weak magnetic fields. For supersonic turbulence the density spectrum flattens, like the hydrodynamic simulations.  In this paper we use deep X-ray observations of a nearby ($z=0.0172$) galaxy cluster AWM\\,7 to look for fluctuations in surface brightness which can indicate the degree of gas motions or turbulence. AWM\\,7 is an X-ray bright ($F_{2-10\\keV} \\sim 9 \\times 10^{-11} \\ergpcmsqps$; \\citealt{Edge92}) poor cluster of galaxies located close to the edge of the Perseus galaxy cluster. The X-ray surface brightness is elongated along the east-west direction in the same direction as the Perseus-Pisces chain of galaxies \\citep{Neumann95}. This is claimed to be caused by the infall of material along that direction rather than tidal effects. There is evidence for a clump of galaxies to the east of the cluster \\citep{Beers95}. The central dominant galaxy and central X-ray peak is offset 30~kpc to the west of the centre of the larger scale emission and the galaxy is aligned close to perpendicularly to the X-ray emission \\citep{Neumann95}. The redshift of the central galaxy agrees with the cluster redshift and the cluster has a velocity dispersion of $\\sim 730\\kmps$ \\citep{Koranyi00}. No significant velocity substructure is found \\citep{Koranyi02}.  \\cite{Sato08} found no evidence for bulk motions greater than the sound speed $\\sim 1000\\kmps$, using the \\emph{Suzaku} observatory. They found that the ICM temperature declines from a peak of 3.85~keV to 3.35~keV at 500~kpc. The cluster has a steeply peaked surface brightness profile, from which would be inferred a cooling rate of 20--35$\\Msunpyr$ in the absence of heating \\citep{Peres98,Neumann95}, and a cool 2 keV core \\citep{Furusho03}. The central mean radiative cooling time is short at $0.4$~Gyr \\citep{Cavagnolo09}.  In this paper we restrict ourselves to an analysis of the surface brightness of the galaxy cluster AWM\\,7, leaving a spectral analysis to a future work.  We assume, unless stated otherwise, that $H_0 = 70\\kmpspMpc$. At a redshift of 0.0172, 1 arcsec on the sky corresponds to 0.351~kpc. ",
        "conclusions": "We examine deep \\emph{Chandra} and \\emph{XMM-Newton} images of the bright AWM\\,7 galaxy cluster. After representing the large scale surface brightness distribution with a model made by interpolating ellipses fitted to surface brightness contours, we find residual features which are approximately radial, have a projected magnitude of 4 per cent in the 0.6 to 5 keV band and are strongest in the south of the cluster. Examining the 3.5 to 7.5 keV band, which is more sensitive to pressure variations, we also find evidence for such features at a similar magnitude.  A histogram of the surface brightness distribution after subtraction of a smooth model also shows a wider distribution than expected from the Poisson statistics.  If, instead of the ELFIT model for the surface brightness, we fit a $\\beta$ or double-$\\beta$ model to the cluster, the residuals are considerably larger because of the offset peak in X-ray emission in this cluster and a number of edges in surface brightness.  We use $\\Delta$-variance spectra to characterise the scales on which the fluctuations occur. The magnitude of the fluctuations is roughly consistent with a model with powerlaw power spectra with an index and a standard deviation of 8 per cent pressure fluctuations in 3D emissivity the 0.6 to 5 keV band. The spectrum appears flatter than expected, however. The implied density fluctuations have a standard deviation of around 4 per cent. In the pressure-squared-sensitive 3.5 to 7.5 keV band, the 3D pressure fluctuations appear to be 4 per cent or less.  We find that there is a large difference between the signal in the $\\Delta$-variance spectra between the north and south of the cluster. On the longest scales there is roughly a factor of two less $\\Delta$-variance signal in the north relative to the south."
    },
    "1112/1112.1188_arXiv.txt": {
        "abstract": "{ In recent years, with new ground-based and HST measurements of proper motions of the Magellanic Clouds being published, a need of a reanalysis of possible orbital history has arisen. As complementary to other studies, we present a partial examination of the parameter space  -- aimed at exploring the uncertainties in the proper motions of both Clouds, taking into account the updated values of Galactic constants and Solar motion, which kinematically and dynamically influence the orbits of the satellites. In the chosen setup of the study, none of the binding scenarios of this pair could be neglected. } % ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.4849_arXiv.txt": {
        "abstract": "\\PRE{\\vspace*{.3in}} We derive bounds on the dark matter annihilation cross-section for low-mass ($5-20~\\gev$) dark matter annihilating primarily to up or down quarks, using the Fermi-LAT bound on gamma-rays from Milky Way satellites.  For models in which dark matter-Standard Model interactions are mediated by particular contact operators, we show that these bounds can be directly translated into bounds on the dark matter-proton scattering cross-section.  For isospin-violating dark matter, these constraints are tight enough to begin to constrain the parameter-space consistent with experimental signals of low-mass dark matter.  We discuss possible models that can evade these bounds. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1112/1112.1836_arXiv.txt": {
        "abstract": "The Major Atmospheric Gamma-ray Imaging Cherenkov (MAGIC) experiment is an array of two 17-meter telescopes located in the Canary Island of La Palma that observes the very-high energy (VHE) gamma-ray sky in stereoscopic mode since 2009. MAGIC is distinguished by its low-energy threshold of approximately 50 GeV, which grants the system a unique potential in the study of distant extragalactic sources whose gamma-ray emission is significantly attenuated due to absorption by the extragalactic background light (EBL). The observation of non-thermal gamma rays in the GeV-TeV range from extragalactic sources is a characteristic signature of their relativistic nature and therefore fundamentally important for our understanding of the physics of these objects. Since the beginning of its stereo operation, MAGIC has observed a large number of active galactic nuclei (AGN) of different classes, including several blazars and distant quasars. In this paper we will review some of the most important results of these observations. ",
        "introduction": "Extragalactic relativistic sources such as active galactic nuclei (AGN) are among the strongest emitters of VHE gamma-ray radiation ($E > 100$ GeV), an unavoidable signature of intense particle acceleration up to tens of TeVs. The study of such energetic processes is key to understanding the nature and evolution of these objects, as it has been revealed that gamma-ray emission accounts for a large fraction of the radiative output from many of them. Efficient particle acceleration in the rarefied and strongly magnetised plasmas characteristic of relativistic jets and other astrophysical environments is \\textit{per se} a question of great interest, with direct impact to the issue of the origin of the VHE cosmic rays\\cite{Felix_rev}$^,$\\cite{Jim_rev}.  In this brief review we are going to present a selection of observational results by MAGIC which we consider to be the most relevant achievements in the extragalactic science programme of the collaboration in the past two years. Ground-based gamma-ray astronomy has gone through a continuous development in the past years and in the extragalactic realm this coincides with a doubling of the number of detected sources (which now total over 45 objects) as well as the discovery of new classes of objects.  \\subsection{The MAGIC Telescopes}  Being the largest imaging atmospheric Cherenkov (IAC) instruments currently in operation, with 17-meter diameter dishes, the MAGIC telescopes enjoy the lowest observational threshold of all ground-based gamma-ray experiments ($E_{\\rm{Thresh}} \\simeq 50$ GeV). Given the opacity effect imposed by the EBL on the propagation of VHE photons\\cite{Dominguez11}$^,$\\cite{Mazin07}, this characteristic gives the experiment a natural advantage to the study of the extragalactic universe. In fact, of the circa 20 new extragalactic objects discovered since 2009, over a third were due to observations by MAGIC, including one new source class and the farthest VHE AGN known\\cite{Atels}.  MAGIC has passed through a major upgrade in the beginning of 2009, whereby a second telescope was added to the then monoscopic experiment, allowing for 3D stereoscopy imaging and reconstruction of the air showers, with a consequent improvement of the sensitivity of the new array across its entire observational range. The experiment currently counts with an integral sensitivity above 250 GeV of 0.8\\% of the Crab flux (after 50 hours observation). Its angular resolution is of 0.07$^\\circ$ and the improved energy resolution now falls at the 15\\% level\\cite{Emiliano11}. Although most of the data presented here was taken in the new instrumental configuration, some of it contains observations with the stand-alone MAGIC-I telescope, which will be indicated whenever it is the case. For a view of the technical specifications of the MAGIC mono-system, please refer to the following publication\\cite{MAGIC08}. ",
        "conclusions": "We gave a short overview of the main extragalactic results of the MAGIC collaboration. Emphasis was given to those which are most interesting regarding their physical implications. Among topics left out due to space constraints is the status of the MAGIC searches for GRBs, for which a recent exposition can be found here\\cite{GRB}."
    },
    "1112/1112.0001_arXiv.txt": {
        "abstract": "{Analytical radially self-similar models are the best available solutions describing disk-winds but need several improvements. In a previous article, we introduced models of jets from truncated disks, i.e. evolved in time numerical simulations based on a radially self-similar MHD solution but including the effects of a finite radius of the jet-emitting disk and thus the outflow. In paper I of this series, we compared these models with available observational data varying the jet density and velocity, the mass of the protostar and the radius of the aforementioned truncation.} {In paper I, we assumed that the jet lies in the plane of the sky. In this paper, we investigate the effect of different inclinations of the jet.} {In order to compare our models with observed jet widths inferred from recent optical images taken with HST and AO, we create again emission maps in different forbidden lines and from such emission maps, we determine the jet width as the full-width half-maximum of the emission.} {We can reproduce the jet width of DG Tau and its variations very well and the derived inclination of 40$^\\circ$ is in excellent agreement with literature values of 32--52$^\\circ$. In CW Tau we overestimate the inclination in our best-fit model. In the other objects, we cannot find appropriate models which reproduce the variations of the observed jet widths, only the average jet width itself is well modeled as in paper I.} {We conclude that truncation -- i.e. taking into account the finite radius of the jet launching region -- is necessary to reproduce the observed jet widths and our simulations limit the possible range of truncation radii. The effects of inclination are important for modeling the intrinsic variations seen in observed jet widths. Our models can be used to infer independently the inclinations in the observed sample, however, a parameter study with a finer grid of parameters is needed.} ",
        "introduction": "Astrophysical jets and disks \\citep{Liv09} seem to be inter-related, notably in the case of young stellar objects (YSOs), wherein jet signatures are well correlated with the infrared excess and accretion rate of the circumstellar disk \\citep{CES90, HEP04}. Disks provide the plasma which is outflowing in the jets, while jets in turn provide the disk with the needed angular momentum removal in order that accretion in the protostellar object takes place \\citep{Har09}. On the theoretical front, the most widely accepted description of this accretion-ejection phenomenon \\citep{Fer07} is based on the interaction of a large scale magnetic field with an accretion disk around the central object. Then, plasma is channeled and magneto-centrifugally accelerated along the open magnetic field lines threading the accretion disk, as firstly described in \\citet{BlP82}. Several works have extended this study either by semi-analytic models using radially self-similar solutions of the full magnetohydrodynamics (MHD) equations with the disk treated as a boundary condition \\citep{VlT98}, or, by selfconsistently treating numerically the disk-jet system \\citep[e.g.][]{ZFR07}.  The original \\citet{BlP82} model, however, has serious limitations for a needed meaningful comparison of its predictions with observations. Singularities exist at the jet axis, the outflow is not asymptotically superfast, and most importantly, the disk has no intrinsic scale with the result that the jet formally extends to radial infinity, to mention just a few. The singularity at the axis can be easily taken care of by numerical simulations extending the analytical solutions close to this symmetry axis \\citep[][GVT06 hereafter]{GVT06}. The outflow speed at large distances may be tuned to cross the corresponding limiting characteristic, with the result that the terminal wind solution is causally disconnected from the disk and hence perturbations downstream of the superfast transition (as modified by self-similarity) cannot affect the whole structure of the steady disk-wind outflow \\citep[][V00 hereafter]{VTS00} which has also been shown to be structurally stable \\citep[][M08 hereafter]{MTV08}. The next step of introducing a scale in the disk has been undertaken in a previous paper \\citep{STV08}, wherein we presented numerical simulations of truncated flows whose initial conditions are based on analytical self-similar models.  In order to test our truncated models, we then applied our simulations to observations \\citep[][paper I hereafter]{SGT10}. In recent years, many NIR and optical data have become available exploring the morphology and kinematics of the jet launching region \\citep[e.g.][and references therein]{DCL00,RDB07,Dou08}. HST and adaptive optics observations give access to the innermost regions of the wind, where the acceleration and collimation occurs \\citep{RMD96,DCL00,WRB02,HEP04}. Since YSO jets emit in a number of atomic (and molecular) lines, we use a synthetic emission code to create emission maps in different forbidden lines which were used by other authors to extract the jet width from images. In paper I, we compared the observed jet widths with those extracted from our synthetic images. A similar study has been done by \\citet{Fer97,CaF00,GCF01,DCF04}, however, using untruncated disk wind models.  In paper I, however, we always assumed the inclination of the jet to be 90$^\\circ$. The influence of varying inclinations on the derived synthetic jet widths will be examined now in this paper.  The remainder of the paper is organized as follows: we briefly review the initial set-up of the numerical simulations and describe our procedure for the comparison with observations in Sec. \\ref{sec_models}. The results of our studies are presented in Sec. \\ref{sec_results}. In Sec. \\ref{sec_bestfit} we describe our best-fit models for each object in the sample and compare derived inclinations with literature values. Finally, we conclude with the implications of the results in terms of the structure of the disk and the respective launching radii of the jets in YSOs. In appendix \\ref{sec_all_lw}, we present the extracted line widths for all models, runs and inclinations for the sake of completeness. Jet velocities derived from synthetic position- velocity diagrams are given in Appendix \\ref{sec_pv_vel}. ",
        "conclusions": "In paper I, we showed as a proof of concept that jet widths derived from numerical simulations extending analytical MHD jet formation models can be very helpful for understanding recently observed jet widths from observations with adaptive optics and space telescopes. Here we investigated qualitatively the influence of different inclinations.  We examined the origin of variations in the extracted jet widths and identified physical and geometrical effects. We showed in which range our results are reliable and not affected by numerical artefacts.  We found that the distance of the first maximum from the jet source increases for increasing mass of the central object and decreasing velocity of the jet. It also increases for increasing degree of truncation. In several models, a second maximum appears whose distance decreases for increasing degree of truncation. The most noticeable effect of a decreasing inclination is a change in the position of the maxima in the synthetic jet widths. In all cases, the distance of the first maximum from the jet source decreases for decreasing inclination. The lower the inclination, the higher is also the number of maxima in the jet width.  We compared our synthetic jet widths with observations of our sample including the position of maxima in the jet width variations as diagnostics for the inclination. Only for DG Tau and CW Tau, we could unambiguously find models which fitted the observed jet widths. For the former, the derived inclination is perfectly consistent with values found in the literature; for the latter our derived inclination is higher than commonly measured (60$^\\circ$ compared to 41$^\\circ$).  In a future work, we have to refine our grid of models for reproducing the observed jet widths and their variations even more accurate. These models have to be done with higher numerical resolution in order to give significant results in terms of synthetic emission maps even within a few AU to the jet source. Furthermore we have to extend the dimensions of our computational domain in order to avoid the described geometrical artefacts."
    },
    "1112/1112.2835_arXiv.txt": {
        "abstract": "Models for the internal composition of Dense Compact Stars are reviewed as well as macroscopic properties derived by observations of relativistic processes. Modeling of pure neutron matter Neutron Stars is presented and crust properties are studied by means of a two fluid model. ",
        "introduction": "Neutron stars (NS) are stellar objects that represent one possible end of stellar evolution. They are in hydrostatic equilibrium in which the gravitational force is balanced by the internal one resulting from the Pauli exclusion principle. Theorized first by Landau (1932) and Zwicky (1934), Neutron Stars have been first observed as pulsars in 1967 by Jocellyn Bell. The latest observations of NS point out that the most extreme physical conditions are present in them.  They are the densest compact objects ($\\rho \\sim 10^{14}-10^{16}$ $\\textrm{g}/\\textrm{cm}^{3}$), fastest rotating stars (as fast as 716 MHz as measured)  and fastest moving objects in the galaxy ($v\\sim 1083$ $\\textrm{km}/\\textrm{s}$). They possess the highest magnetic fields ($B=10^{15}$ G and largest surface gravity ($10^{14}$ $\\textrm{cm}/\\textrm{s}^{2}$). Their interior contains superconducting material with the highest expected temperature value ($T_{c}=10^{9}$ K),  and even neutrinos can be trapped in proto-neutron stars with temperatures at birth of about 700,000 million K \\cite{2004Sci...304..536L}. \\vspace{-.5cm} ",
        "conclusions": "Neutron Stars are dense compact objects where the most extreme physical conditions exist. The macroscopic description being in terms of the General Theory of Relativity. For the interior of the star different nuclear models are being used based on terrestrial laboratory experiments and extrapolated to the conditions found in NS. Theoretical models should be able to reproduce observations, and with the forthcoming data such models should start to converge  since some of them could be tweaked in their parameters to reach the desired values.  In particular the latest 2 $\\textrm{M}_\\odot$ measurement imposes a strong constrain in the EoS.  Finally other relativistic processes like the cooling of the proto-neutron star could also provide evidence on the interior of NS, since the cooling rate is modified by the presence of superfluid matter. All this is reflected in the $M$ vs $R$ relation."
    },
    "1112/1112.2973_arXiv.txt": {
        "abstract": "{} {Weak G-band (WGB) stars are a rare class of cool luminous stars that present a strong depletion in carbon, but also lithium abundance anomalies that have been little explored in the literature since the first discovery of these peculiar objects in the early 50's. Here we focus on the Li-rich WGB stars and report on their evolutionary status. We explore different paths to propose a tentative explanation for the lithium anomaly. } {Using archive data, we derive the fundamental parameters of WGB (\\teff, \\logg, \\lL) using Hipparcos parallaxes and recent temperature scales. From the equivalent widths of Li resonance line at 6707 \\AA, we uniformly derive the lithium abundances and apply when possible NLTE corrections following the procedure described by Lind et al. (2009). We also compute dedicated stellar evolution models in the mass range 3.0 to 4.5 \\msun, exploring the effects of rotation-induced and thermohaline mixing. These models are used to locate the WGB stars in the H-R diagram and to explore the origin of the abundance anomalies.} { The location of WGB stars in the H-R diagram shows that these are intermediate mass stars of masses ranging from 3.0 to 4.5 M$_{\\odot}$ located at the clump, which implies a degeneracy of their evolutionary status between subgiant/red giant branch and core helium burning phases. The atmospheres of a large proportion of WGB stars (more than 50\\%) exhibit lithium abundances $A(Li) \\geq 1.4$ dex similar to Li-rich K giants. The position of WGB stars along with the Li-rich K giants in the H-R diagram however indicates that both are well separated groups. The combined and tentatively consistent analysis of the abundance pattern for lithium, carbon and nitrogen of WGB stars seems to indicate that carbon underabundance could be decorrelated from the lithium and nitrogen overabundances.} {} ",
        "introduction": "\\label{sec:intro} \\subsection*{ Defining WGB stars} The weak G-band (hereafter WGB) stars are G and K giants whose spectra show very weak or absent G-bands of the CH A$^{2}\\Delta$ - X$^{2}\\Pi$ system at 4300\\AA. These stars were first identified as stellar class by \\citet{bidelman1951}, and have been mainly studied in the late seventies, early eighties with a total number of dedicated papers not exceeding 20 \\citep[see e.g. for instance][]{sneden1978,rao1978,partha1980,day1980}. They are rare with less than 30 known to date among the population of G-K giants in the Galaxy.\\\\  Chemical composition studies \\citep{sneden1978,cottrell1978} demonstrated that they are very much underabundant in carbon (typical [C/Fe] $\\approx -1.4$) and present small overabundances of nitrogen and normal oxygen. \\citet{sneden1978} analyzed the $^{13}$CN red-system features in the high resolution spectra of weak G-band stars and found \\ciso = 4.  \\citet{hartoog1977} from the photometry of the 2.3$\\mu$m CO vibration-rotation bands confirmed the underabundances of carbon in a large sample of weak G-band stars. The CH band strengths observed by \\citet{rao1978} also indicate that the weakening of the G-band is due to underabundance of carbon. \\citet{sneden1984} obtained high resolution spectra of the 2$\\mu$m first-overtone CO bands in the weak G-band giants and found excellent agreement between the carbon abundances derived from CO data and those determined using features of the CH G-band. They have found \\ciso $\\le$ 4  which is in agreement with the predicted ratio for the CN-cycle in equilibrium. The normal abundance of oxygen and sodium shows that the atmospheres of WGB stars are probably not mixed with ON-cycle processed material \\citep{sneden1978,drake1994}.\\\\  Additional chemical constraints are given by lithium and and beryllium abundance determinations.\\\\ Li abundances have been derived for several WGB stars and several stars were found to be lithium rich, e.g.  with A(Li)\\footnote{With the classical notation A(Li) = log(${\\rm N}_{Li}/{\\rm N}_H) + 12$} $\\ge 1.4$ \\citep[see][]{brown1989} and reaching up to  A(Li) = 3 \\citep{lambert1984,partha1980}. \\citet{lambert1984} also report on the possible presence of $^6{\\rm Li}$ in HR 1299, but the profile of the 6707 \\AA $^7{\\rm Li}$ line that is better reproduced with $^6{\\rm Li}/^7{\\rm Li} \\ne 0$, could also be reproduced without having to invoke the presence of $^6{\\rm Li}$. \\\\Be abundances were derived from IUE spectra for 3 Li-rich weak G-band stars by \\citet{partha1984}. Providing a differential analysis between the WGB stars and K giants in the Hyades, Be is found to be similar in both groups, and compatible within presumably large errorbars\\footnote{The authors do not indicate any errorbar in the paper and give abundances of Be ``{\\em rounded to the nearest  0.5 dex}''.}, with the expected post dredge-up Be abundances according to standard stellar evolution models.  \\subsection*{Tentative evolutionary status and initial mass estimate}  Due to their relative low temperature and gravity, WGB stars are classified as giants and fall amidst the very crowded area of the HR diagram populated with stars on the subgiant branch, the RGB, the core helium burning phase and the AGB. This has made the definition of their evolutionary status and initial mass difficult. They seem however to have been unanimously qualified as core helium burning star. In particular \\citet{sneden1978} tentatively attribute a mass of $\\approx 1 \\msun$ to HR 6766 (e.g. HD 165634) and exclude it from being more massive than 3 \\msun because of the magnitude they derive. With such a low-mass, that star should then have undergone the He flash. Let us note that their conjecture is based on a very crude estimation for the luminosity due to the large error bars impairing the parallax of HR 6766. \\citet{cottrell1978} also identify WGB stars in their sample as being in the mass range 1.5 \\msun $\\leq$ M $\\leq$ 3 \\msun past the He flash. Later on, \\citet{lambert1984} proposed that WGB stars could be the progeny of magnetic Ap stars, with initial masses of 2-3 \\msun that have undergone a dredge-up, but they do not make any clear statement of their evolutionary status. Much recently \\citet{partha2000,partha2002} suggest that WGB stars should be the descendants of intermediate-mass stars with 2 \\msun $\\leq$ M $\\leq$ 5 \\msun. In that case, WGB should not go through the He flash, and will undergo a different evolution than that of the lower mass stars. This in turn may have an impact on the possible mechanisms that can be invoked to explain the abundance anomalies.  \\subsection*{Comparison with other carbon depleted and Li-rich giants}  WGB stars are thus possibly low to intermediate mass evolved stars exhibiting strong depletion of carbon that can be accompanied in several cases by a Li enrichment. One may wonder how these objects compare with the other stars that also present signatures of CN-cycle processed material in their atmospheres and exhibit unexpected Li enrichment.\\\\ Giant stars with depleted carbon and enhanced nitrogen abundances are found in almost all environments among giants. These are essentially low-mass stars on the RGB. \\citet{charbon1998} show that 96\\% of field RGB stars with metallicities from solar to 1/1000 solar and M$_V \\leq -0.5$ have a carbon isotopic ratio lower than what is predicted by standard theory. In their survey of metal-poor field stars, \\citet{gratton2000} show that the unexpected decrease of the carbon isotopic ratio occurs at the RGB bump, when the outgoing hydrogen burning shell (hereafter HBS) erases the mean molecular weight discontinuity left behind by the retreating convective envelope as a result of the first dredge-up. It is furthermore associated with a decrease of the carbon abundance [C/Fe], an increase of the nitrogen abundance [N/Fe] and a strong decrease of the lithium abundance, all characteristic of CN-cycle processed material. Oxygen does not seem to be affected, as in the case of WGB stars\\\\% Typically, \\citet{charbon1998} report that around solar metallicity the lower envelope for the \\ciso ratio in low-mass RGB stars is of about 12. \\citet{gratton2000} report  [C/Fe] = -0.58 $\\pm$ 0.03 as a mean carbon abundance for upper RGB low-metallicity field stars (e.g. \\lL $>$ 2 and -2 $\\le$ [Fe/H] $\\leq$ -1), and \\ciso between 6-10. These values are higher than the mean obtained for WGB stars.\\\\ In addition to that Li-enrichment also exists among K giants. It is a very rare phenomenon found in less than 2\\% of the K giant stars \\citep{kumar2011}. \\citet{charbon2000} have shown that around solar metallicity, giants exhibiting a genuine overabundance of Li (e.g. excluding those giants that are still undergoing the first dredge-up) appear to cluster around the RGB bump and the clump. Very recently however \\citet{monaco2011} have studied Li-rich RGB stars in the Galactic thick disk, and show that at low metallicity Li-rich giants do not cluster at the bump and are found scattered between the bump and the tip of the RGB \\citep[see also][]{lebzelter2011}. \\citet{kumar2011} confirmed the finding by \\citet{charbon2000}, with Li-rich stars clustered at the RGB bump, and some new objects that seem to be more evolved, may be past the RGB tip and undergoing central He burning. Let us note that Li-rich K giant stars around the bump region do not present ``anomalous'' carbon isotopic ratios, e.g., the \\ciso determined in these stars well agrees with the predictions of post dredge-up dilution \\citep{charbon2000,kumar2011}.\\\\ In these K giants, Li enrichment is thought to have an internal origin \\citep[so called][mechanism]{cf1971} and to be a short-lived phase {\\em preceding} the phase of transport CN-cycle processed material from the HBS to the surface \\citep{charbon2000,palacios2001}. The extra-mixing mechanism capable of transporting Li in regions where it is preserved from proton captures and connecting the convective envelope to these regions is not clearly identified \\citep{palacios2001,palmerini2011}.  As a summary, in low-mass RGB stars, extra-mixing processes are clearly responsible for bringing up to the stellar surface the products of the CN-cycle occurring in the outer layers of the HBS deeper inside the star, and may be also responsible for strong and short-lived episodes of Li enrichment. The specific location of Li-rich K giants \\citep[see, for example][]{jasniewicz1999} points towards an extra-mixing process that becomes efficient only when no strong molecular weight barrier shields the CN-cycle processed regions. This occurs only after the RGB bump in low-mass stars ($M \\leq 2.3 M_\\odot$) or on the early AGB for the stars that do not go through the He flash ($M > 2.3 M_\\odot$). Carbon depletion and nitrogen enhancement at the stellar surface that would result from internal mixing are thus expected to happen during these specific evolutionary stages according to the proposed self-consistent scenarios \\citep{charbon2000,charbon2010}. \\subsection*{Paper outline} In this paper, we propose a new re-analysis of the data available for WGB stars and investigate the possible origin for their abundance anomalies, giving particular attention to the lithium abundance anomalies. In particular we investigate the possible physical mechanisms that could account for the {\\em simultaneous} lithium overabundances and strong carbon underabundances. In the following, we describe the homogeneous redetermination of atmospheric parameters and lithium abundances for the WGB stars known to date in \\S~\\ref{sec:param} and \\ref{sec:li}. We then place the WGB stars in the HR diagram and make a tentative classification in terms of evolutionary status and initial masses in \\S~\\ref{sec:evol}. In \\S~\\ref{sec:abund} and \\ref{sec:discussion} we finally explore possible scenarios to consistently account for the light elements (Li, C, N and O) abundances found in WGB stars based on the results of dedicated stellar evolution models. ",
        "conclusions": "In this study we have made a tentative re-analysis of the data existing for the WGB stars, in particular concerning fundamental parameters and lithium abundances. We have confirmed that WGB stars are within the mass range 3-4.5\\msun. However, the too few abundance determinations of Li, N and C and the fact that the WGB stars are located at the clump where stellar evolution tracks corresponding to the subgiant, red giant and core helium burning phases are intertwined prevent us to assess any sound evolutionary status to the WGB stars. Concerning the chemical peculiarities, we have used dedicated stellar evolution models to test their endogenous origin, and shown that for those WGB stars that are not Li-rich, lithium and nitrogen abundances can be reproduced by rotating models . On the other hand the carbon deficiency of all WGB stars is difficult to reconcile with an internal origin and we propose that it is decorrelated from the other chemical peculiarities. The case of lithium is still under debate as it is not clear, because of the degeneracy of the evolutionary status, whether the WGB stars with $A(Li) \\geq 1.4$ dex are actually Li-rich (core helium burning stars) or normal (subgiant or RGB stars undergoing the first dredge-up). \\\\ Carbon depletion could be associated to an early mass transfer from a more massive (5-6 \\msun) companion under specific hypothesis, but yields of intermediate mass stars in that mass range are very uncertain and no firm conclusion should be drawn. As the rate of binarity among WGB stars is far from clear, and is at least similar to that in normal giants, a new search for duplicity would be helpful.\\\\ It is quite clear from this study that more data are needed to shed light on the weak G-band stars puzzle. Considering the work by \\citet{miglio2011} on the use of asteroseismic data to lift the status degeneracy between giants and clump stars, applying such an analysis to the oscillation spectra of WGB stars could be a very important step to clarify the evolutionary status of WGB stars.\\\\ Actually deriving proper abundances of light elements for these stars is also needed to start and better understand these overlooked puzzling stars."
    },
    "1112/1112.3090_arXiv.txt": {
        "abstract": "The cross-correlation search for gravitational wave, which is known as 'radiometry', has been previously applied to map of the gravitational wave stochastic background in the sky and also to target on gravitational wave from rotating neutron stars/pulsars. We consider the Virgo cluster where may be appear as `hot spot' spanning few pixels in the sky in radiometry analysis. Our results show that sufficient signal to noise ratio can be accumulated with integration times of the order of a year. We also construct numerical simulation of radiometry analysis, assuming current constructing/upgrading ground-based detectors. Point spread function of the injected sources are confirmed by numerical test. Typical resolution of radiometry analysis is a few square degree which corresponds to several thousand pixels of sky mapping. ",
        "introduction": "Current constructing and upgrading ground-based gravitational wave detectors may reach the strain sensitivity in order of $10^{-24} {\\rm [1/\\sqrt{Hz}]}$ in near future. These detector sites are dislocated, and make possible to have a independent observation of gravitational waves. Since the dislocated detectors may not have a correlation in noise component but have a correlation in gravitational waveforms, cross-correlation product by two or more detectors make possible to search for unknown gravitational waveforms. This is useful property in burst gravitational wave search for short duration, or in stochastic gravitational wave search for long duration. Distance between detectors also make possible triangulation for signal measurement. Time delay between detectors suggest the incident direction of the gravitational wave. It is used in short duration signal analysis. On the other hand, an employment of the time delay term in long duration signal search or stochastic gravitational wave which come from any directions is not simple.  However, we can extract the signal strength if we resolve Earth's rotation and convoluted signal from anywhere in the sky. Using multiple detectors and long time integration, the deconvolution will be possible. The directional cross-correlation search for gravitational wave is proposed as 'radiometry'.  Radiometry and any other triangulation analysis require long distance among detectors, and variation of arrangement (zenith direction, and azimuthal rotation of interferometers) make possible to cover the sky and resolve the polarization. Recent starting of the construction of LCGT\\cite{LCGT} in addition to previously constructed advanced detector sites of LIGO\\cite{LIGO} and Virgo\\cite{VIRGO} encourage the development of radiometry analysis.  \\subsection{Radiometry Filter}  Gravitational wave radiometry is a kind of stochastic gravitational wave search, but it is a directional analysis. Radiometry analysis assumes that gravitational wave come from particular direction as a point like source or a hot spot of sky. The radiometry filter uses at least two detector observation data, and do not require the waveform. Gravitational wave coming from particular direction will have a time delay of arrival at two detectors shown in Fig.\\ref{fig:IdeaOverview}. Since gravitational wave signals on both detectors are coherent in same polarization component, we can extract gravitational wave signals in the cross-correlation product from two detectors observation data with appropriate time delay. Moreover, signal-to-noise ratio for gravitational wave will be emphasized with long integration time, because the two detector noises are independent. Changing the assumption of source direction step by step, we can search gravitational wave `sky map' by radiometry analysis.  \\begin{figure}[htb] \\begin{center} \\includegraphics[width=9cm]{Idea_overview.eps} \\caption{Ideal setup of a radiometry} \\label{fig:IdeaOverview} \\end{center} \\end{figure}  Radiometry analysis is given and studied in previous works\\cite{mitra}\\cite{bose}. We explain the radiometry filter briefly here. The source direction or particular direction of search can be represent as unit vector $\\hat{\\Omega}$. Two detectors are noted as 1 and 2, and a base-line from detector 1 to detector 2 can be given as vector $\\overrightarrow{\\Delta x}$. An arrival time delay for gravitational wave between two detectors is $\\displaystyle \\frac{\\hat{\\Omega}\\cdot\\overrightarrow{\\Delta x}}{c}$, where $c$ is a speed of light. Detector output from each detector $\\displaystyle s_{1}(t), s_{2}(t\\ -\\ \\frac{\\hat{\\Omega}\\cdot\\overrightarrow{\\Delta x}}{c})$ in time domain have same gravitational wave timing. In Fourier domain, this time delay represent as a phase difference of two detectors output $\\tilde{s}_{1}(f), \\tilde{s}_{2}(f)$ in Fourier transform of $s_1(t)$ and $s_2(t)$. Therefore, we can define the radiometry filter as following. A directed overlap reduction function $\\gamma(f,\\hat{\\Omega})$ can be defined to correct phase difference as \\begin{equation} \\gamma(f,\\hat{\\Omega})=\\sum_{A=+,\\times}\\!\\!\\!F_1^A F_2^A e^{i 2\\pi f \\frac{\\hat{\\Omega}\\cdot\\overrightarrow{\\Delta x}}{c}}\\ , \\label{eq:gamma} \\end{equation} where $F_{1,2}^{A}$ are antenna pattern\\cite{antennaPat} of each detector 1 or 2 for each two polarization $A=+$ or $\\times$ of gravitational wave. To make optimal filtering, we employ two detectors noise power spectrum $P_{1,2}(f)$ respectively, and assume that gravitational wave spectrum as $H(f)$. Radiometry filter can defined as \\begin{equation} Q(f,\\hat{\\Omega})=\\lambda\\frac{\\gamma^*(f,\\hat{\\Omega})H(f)}{P_1(f)P_2(f)}\\ , \\label{eq:Q} \\end{equation} and filter output is \\begin{equation} \\Delta S(\\hat{\\Omega})=\\int_{-\\infty}^{\\infty}df\\ \\tilde{s}_1^*(f) \\tilde{s}_2(f)Q(f,\\hat{\\Omega}). \\end{equation}  The base-line $\\overrightarrow{\\Delta x}$ will change relatively in celestial coordinate according to Earth's rotation in the case of ground-based detectors. Therefore, we calculate $\\Delta S$ at short time slice (`chunk') during $t - \\Delta t/2$ and $t + \\Delta t/2$. Fourier transforms of detector signals can be represented using short Fourier transform of $\\tilde{s}_{1,2}^*(t;\\!f)$, and $Q(f)$ can be denoted as $Q(t,f)$. \\begin{equation} \\Delta S(t,\\hat{\\Omega})=\\int_{-\\infty}^{\\infty} df\\ \\tilde{s}_1^*(t;\\!f) \\tilde{s}_2(t;\\!f) Q(t,\\!f,\\!\\hat{\\Omega}). \\label{eq:radiometryFilter} \\end{equation}  Processing many chunks for long duration observation data, we can get the statistics of $\\Delta S$. Since the gravitational wave component will appear in real part of $\\Delta S$, we employ a mean $\\mu(\\hat{\\Omega})$ \\begin{equation} \\mu(\\hat{\\Omega})=\\langle\\Re[\\Delta S]\\rangle\\ , \\label{eq:mu} \\end{equation} and a standard deviation $\\sigma(\\hat{\\Omega})$ for a real part of $\\Delta S$. Signal-to-noise ratio of particular direction can be given as \\begin{equation} \\rho(\\hat{\\Omega})=\\mu(\\hat{\\Omega})/\\sigma(\\hat{\\Omega}). \\label{eq:sn} \\end{equation} ",
        "conclusions": "We studied radiometry analysis for stochastic gravitational waves from astronomical origin, which might have an anisotropy or local excess. Virgo cluster has a possibility to appear as a `hot spot' on the radiometry map of the sky with advanced and third generation gravitational wave detectors. We also studied numerical simulation of radiometry filter, and displayed the spreading image for point like gravitational wave sources, with LCGT and advanced LIGO combination. Also other combinations in global network detectors will be useful to emphasize the signal behavior. We are planning to proceed the injection test of our numerical simulation studies for multiple sources and area spreading sources.   \\ack S. Dhurandhar acknowledges the DST and JSPS Indo-Japan international cooperative programme for scientists and engineers for supporting visits to Osaka City University, Japan and Osaka University, Japan. H. Tagoshi, N. Kanda and H. Takahashi thank JSPS and DST under the same Indo-Japan programme for their visit to IUCAA, Pune, India. N.Kanda's work was supported in part by a Monbu Kagakusho Grant-in-aid for Scientific Research of Japan (No. 23540346). H.Tagoshi's work was also supported in part by a Monbu Kagakusho Grant-in-aid for Scientific Research of Japan (Nos. 20540271 and 23540309). H.Takahashi's work was also supported in part by a Monbu Kagakusho Grant-in-aid for Scientific Research of Japan (No. 23740207)."
    },
    "1112/1112.4165_arXiv.txt": {
        "abstract": "{ The ESA space astrometry mission Gaia, planned to be launched in 2013, has been designed to make angular measurements on a global scale with micro-arcsecond accuracy. A key component of the data processing for Gaia is the astrometric core solution, which must implement an efficient and accurate numerical algorithm to solve the resulting, extremely large least-squares problem. The Astrometric Global Iterative Solution (AGIS) is a framework that allows to implement a range of different iterative solution schemes suitable for a scanning astrometric satellite. } { Our aim is to find a computationally efficient and numerically accurate iteration scheme for the astrometric solution, compatible with the AGIS framework, and a convergence criterion for deciding when to stop the iterations. } { We study an adaptation of the classical conjugate gradient (CG) algorithm, and compare it to the so-called simple iteration (SI) scheme that was previously known to converge for this problem, although very slowly. The different schemes are implemented within a software test bed for AGIS known as AGISLab. This allows to define, simulate and study scaled astrometric core solutions with a much smaller number of unknowns than in AGIS, and therefore to perform a large number of numerical experiments in a reasonable time. After successful testing in AGISLab, the CG scheme has been implemented also in AGIS. } { The two algorithms CG and SI eventually converge to identical solutions, to within the numerical noise (of the order of 0.00001~micro-arcsec). These solutions are moreover independent of the starting values (initial star catalogue), and we conclude that they are equivalent to a rigorous least-squares estimation of the astrometric parameters. The CG scheme converges up to a factor four faster than SI in the tested cases, and in particular spatially correlated truncation errors are much more efficiently damped out with the CG scheme. While it appears to be difficult to define a strict and robust convergence criterion, we have found that the sizes of the updates, and possibly the correlations between the updates in successive iterations, provide useful clues. } {} ",
        "introduction": "\\label{sec-intro}  The European Space Agency's Gaia mission \\citep{gaia2001,lind+08,Lindegren10} is designed to measure the astrometric parameters (positions, proper motions and parallaxes) of around one billion objects, mainly stars belonging to the Milky Way Galaxy and the local group. The scientific processing of the Gaia observations is a complex task that requires the collaboration of many scientists and engineers with a broad range of expertise from software development to CCDs. A consortium of European research centres and universities, the Gaia Data Processing and Analysis Consortium (DPAC), has been set up in 2005 with the goal to design, implement and operate this process \\citep{mign+08}. In this paper we focus on a central component of the scheme, namely the astrometric core solution, which solves the corresponding least-squares problem within a software framework known as the Astrometric Global Iterative Solution, or AGIS \\citep{Lammers+09,agis11,WOM2010}.  In a single solution, the AGIS software will simultaneously calibrate the instrument, determine the three-dimensional orientation (attitude) of the instrument as a function of time, produce the catalogue of astrometric parameters of the stars, and link it to an adopted celestial reference frame. This computation is based on the results of a preceding treatment of the raw satellite data, basically giving the measured transit times of the stars in the instrument focal plane \\citep{Lindegren10}. The astrometric core solution can be considered as a least-squares problem with negligible non-linearities except for the outlier treatment. Indeed, it should only take into account so-called primary sources, that is stars and other point-like objects (such as quasars) that can astrometrically be treated as single stars to the required accuracy. The selection of the primary sources is a key component of the astrometric solution, since the more that are used the better the instrument can be calibrated, the more accurate the attitude can be determined, and the better the final catalogue will be. This selection, and the identification of outliers among the individual observations, will be made recursively after reviewing the residuals of previous solutions \\citep{agis11}. What remains is then, ideally, a `clean' set of data referring to the observations of primary sources, from which the astrometric core solution will be computed by means of AGIS.  From current estimates, based on the known instrument capabilities and star counts from a Galaxy model, it is expected that at least 100~million primary sources will be used in AGIS. Nonetheless, the solution would be strengthened if even more primary sources could be used. Moreover, it should be remembered that AGIS will be run many times as part of a cyclic data reduction scheme, where the (provisional) output of AGIS is used to improve the raw data treatment \\citep[the Intermediate Data Update; see][]{OMull+09}. Hence, it is important to ensure that AGIS can be run both very efficiently from a computational viewpoint, and that the end results are numerically accurate, i.e., very close to the true solution of the given least-squares problem.  Based on the generic principle of self-calibration, the attitude and calibration parameters are derived from the same set of observational data as the astrometric parameters. The resulting strong coupling between the different kinds of parameters makes a direct solution of the resulting equations extremely difficult, or even unfeasible by several orders of magnitude with current computing resources \\citep{bom-lin-hol-09}. On the other hand, this coupling is well suited for a block-wise organization of the equations, where, for example, all the equations for a given source are grouped together and solved, assuming that the relevant attitude and calibration parameters are already known. The problem then is of course that, in order to compute the astrometric parameters of the sources to a given accuracy, one needs to know first the attitude and calibration parameters to corresponding accuracies; these in turn can only be computed once the source parameters have been obtained to sufficient accuracy; and so on. This organization of the computations therefore naturally leads to an iterative solution process. Indeed, in AGIS the astrometric solution is broken down into (at least) three distinct blocks, corresponding to the source, attitude and calibration parameter updates, and the software is designed to optimize data throughput within this general processing framework \\citep{Lammers+09}. Cyclically computing and applying the updates in these blocks corresponds to the so-called simple iteration (SI) scheme (Sect.~\\ref{sec-simple}), which is known to converge, although very slowly.  However, it is possible to implement many other iterative algorithms within this same processing framework, and some of them may exhibit better convergence properties than the SI scheme. For example, it is possible to speed up the convergence if the updates indicated by the simple iterations are extrapolated by a certain factor. More sophisticated algorithms could be derived from various iterative solution methods described in the literature.  The purpose of this paper is to describe one specific such algorithm, namely the conjugate gradient (CG) algorithm with a Gauss--Seidel preconditioner, and to show how it can be implemented within the AGIS processing framework. We want to make it plausible that it indeed provides a rigorous solution to the given least-squares problem. Also, we will study its convergence properties in comparison to the SI scheme and, if possible, derive a convergence criterion for stopping the iterations.  Our focus is on the high-level adaptation of the CG algorithm to the present problem, i.e., how the results from the different updating blocks in AGIS can be combined to provide the desired speed-up of the convergence. To test this, and to verify that the algorithm provides the correct results, we need to conduct many numerical experiments, including the simulation of input data with well-defined statistical properties, and iterate the solutions to the full precision allowed by the computer arithmetic. On the other hand, since it is not our purpose to validate the detailed source, instrument and attitude models employed by the updating blocks, we can accept a number of simplifications in the modelling of the data, such that the experiments can be completed in a reasonable time. The main simplifications used in the present study are as follows: \\begin{enumerate} \\item For conciseness we limit the present study to the source and attitude parameters, whose mutual disentanglement is by far the most critical for a successful astrometric solution \\citep[cf.][]{bom-lin-hol-09}. For the final data reduction many calibration parameters must also be included, as well as global parameters \\citep[such as the PPN parameter $\\gamma$;][]{Hobbs+09}, and possibly correction terms to the barycentric velocity of Gaia derived from stellar aberration \\citep{Butkevich+08}. These extensions, within the CG scheme, have been implemented in AGIS but are not considered here. \\item We use a scaled-down version of AGIS, known as AGISLab (Sect.~\\ref{sec-agislab}), which makes it possible to generate input data and perform solutions with a much smaller number of primary sources than would be required for the (full-scale) AGIS system. This reduces computing time by a large factor, while retaining the strong mutual entanglement of the source and attitude parameters, which is the main reason why the astrometric solution is so difficult to compute. \\item The rotation of the satellite is assumed to follow the so-called nominal scanning law, which is an analytical prescription for the pointing of the Gaia telescopes as a function of time. That is, we ignore the small ($< 1$~arcmin) pointing errors that the real mission will have, as well as attitude irregularities, data gaps, etc. The advantage is that the attitude modelling becomes comparatively simple and can use a smaller set of attitude parameters, compatible with the scaled-down version of the solution. \\item The input data are `clean' in the sense that there are no outliers, and the observation noise is unbiased with known standard deviation. This highly idealised condition is important in order to test that the solution itself does not introduce unwanted biases and other distortions of the results. \\end{enumerate} An iterative scheme should in each iteration compute a better approximation to the exact solution of the least-squares problem. In this paper we aim to demonstrate that the SI and CG schemes are converging in the sense that the errors, relative to an exact solution, vanish for a sufficient number of iterations. Since we work with simulated data, we have a reference point in the true values of the source parameters (positions, proper motions and parallaxes) used to generate the observations. We also aim to demonstrate that the CG method is an efficient scheme to solve the astrometric least-squares problem, i.e., that it leads, in a reasonable number of iterations, to an approximation that is sufficiently close to the exact solution. An important problem when using iterative solution methods is how to know when to stop, and we study some possible convergence criteria with the aim to reach the maximum possible numerical accuracy.  The paper provides both a detailed presentation of the SI and CG algorithms at work in AGIS and a study of their numerical behaviour through the use of the AGISLab software \\citep{Holl+09}. The paper is organized as follows: Section~\\ref{sec-iter} gives a brief overview of iterative methods to solve a linear least-squares problem. Section~\\ref{sec-algo} describes in detail the algorithms considered here, viz., the SI and CG with different preconditioners. In Sect.~\\ref{sec-conv} we analyze the convergence of these algorithms and some properties of the solution itself. Then, Sect.~\\ref{sec-agis} presents the implementation status of the CG scheme in AGIS before the main findings of the paper are summarized in the concluding Sect.~\\ref{sec-conc}. ",
        "conclusions": "}  We have shown how the conjugate gradient algorithm, with a Gauss--Seidel type preconditioner, can be efficiently implemented within the already existing processing framework for Gaia's Astrometric Global Iterative Solution (AGIS). This framework was originally designed to solve the astrometric least-squares problem for Gaia using the so-called Simple Iteration (SI) scheme, which is intuitively straightforward but computationally inefficient. The conjugate gradient (CG) scheme, by using the same kernel operations as SI, takes about the same processing time per iteration but requires a factor 3--4 fewer iterations. Both schemes have been extensively tested using the AGISLab test bed, which allows to perform scaled-down and simplified simulations of Gaia's astrometric observations and the associated least-squares solution, using (in our case) $10^5$ sources and a total of about 2~million source and attitude unknowns.  To within the numerical noise of the double-precision arithmetic, corresponding to $<10^{-5}~\\mu$as in parallax, the SI and CG schemes converge to identical solutions. In the case when no observational noise was added to the simulated observations, the solutions agree with the true values of the astrometric parameters to within the numerical noise. Thus we conclude that the iterative method provides the correct solution to the least-squares problem, provided that a sufficient number of iterations is used (full numerical convergence reached). As theoretically expected, the resulting solution is completely insensitive to the initial values of the astrometric parameters, although the rate of convergence may depend on the initial errors.  Although the SI and CG schemes eventually reach the same solution, to within the numerical precision of the computations, the truncation errors obtained by prematurely stopping the iteration have quite different character in the two schemes. In the SI scheme the truncation error maps are often strongly correlated over large parts of the celestial sphere, while in the CG scheme there is less spatial correlation. Thus, the CG scheme not only converges faster than SI, but the truncation errors at a given level of the updates are considerably more `benign' in terms of large-scale systematics.  Most of the numerical experiments use a highly idealised, uniform distribution of sources, and a uniform level of the standard error of the observations. However, it has been demonstrated that the solution works flawlessly also in the case when the observations have a much larger weight in a small part of the sky (representing, for example, a bright stellar cluster). Although such a situation needs more iterations, the converged solution correctly reflects the weight distribution of the observations -- i.e., the accuracy of the astrometric parameters of the cluster stars is increased roughly in proportion to their increased observational accuracy, without any noticeable negative impact on the results for other stars.  We have stressed the need to drive the iterations to full numerical convergence, at least in the final astrometric solution for the Gaia catalogue. This is important in order to avoid that the final catalogue contains truncation errors that are unrelated to the mission and satellite itself, but merely caused by inadequate processing. Such truncation errors could mimic `systematic errors' in the catalogue. The Gaia catalogue will certainly not be free of systematic errors, but at least we should insist that they are not produced by prematurely stopping the astrometric iterations. However, achieving full numerical convergence may require many iterations beyond the point where simple metrics indicate convergence. It is in fact quite difficult to define a numerically strict convergence criterion, although we have found that the correlation between updates in successive iterations may provide a useful clue. At the point where the updates have not yet settled at their final level, the magnitude of the updates gives a good indication of the remaining truncation errors. If one does not insist on full numerical convergence, it is therefore relatively safe to stop iterating when the updates have reached a sufficiently low level, say below 0.01~$\\mu$as.  The CG algorithm described in this paper only considers the two major processing blocks in AGIS, namely the determination of the source and attitude parameters. This restriction was intentional, in order to simplify the description and numerical testing. At the same time, the successful disentangling of the source and attitude parameters is the key to a successful solution, as shown by numerous experiments. Nevertheless, the proposed algorithm is readily extended to the fully realistic problem that includes calibration and global parameters, and has in fact been realised in this form and successfully demonstrated in the current AGIS implementation at the Gaia data processing centre in ESAC (Madrid).  Our aim has been to investigate the applicability of the CG algorithm for solving Gaia's astrometric least-squares problem efficiently within the AGIS framework. To this end we have considered a scaled-down and highly idealized version of the problem where many detailed complications in the real Gaia data are ignored. Nevertheless, within the given assumptions, we have successfully demonstrated how the CG algorithm can be adapted to the astrometric core solution. Moreover, by means of numerical simulations we have shown that the numerical accuracy achieved with this method is high enough that it will not be a limiting factor for the quality of the astrometric results."
    },
    "1112/1112.4679_arXiv.txt": {
        "abstract": "\\textbf{Small-scale extreme ultraviolet (EUV) dimming often surrounds sites of energy release in the quiet Sun. This paper describes a method for the automatic detection of these small-scale EUV dimmings} using a feature based classifier. The method is demonstrated using sequences of 171\\AA\\ images taken by STEREO/EUVI on 13 June 2007 and by SDO/AIA on 27 August 2010. The feature identification relies on recognizing structure in sequences of space-time 171\\AA\\ images using the Zernike moments of the images. The Zernike moments space-time slices with events and non-events are distinctive enough to be separated using a Support Vector Machine (SVM) classifier. The SVM is trained using 150 event and 700 non-event space-time slices. We find a total of 1217 events in the EUVI images and 2064 events in the AIA images on the days studied. Most of the events are found between latitudes -35\\degree\\ and +35\\degree. The sizes and expansion speeds of central dimming regions are extracted using a region grow algorithm. The histograms of the sizes in both EUVI and AIA follow a steep power law with slope about -5. The AIA slope extends to smaller sizes before turning over. The mean velocity of 1325 dimming regions seen by AIA is found to be about 14~\\kms. ",
        "introduction": "{Extreme ultraviolet (EUV) dimming surrounding small bursts of EUV brightening in the quiet Sun were first noticed by Innes et al. (2009) in 171\\AA\\ images taken by the Extreme UltraViolet Imager (EUVI) on STEREO. The dimmings are usually more extended and last longer than the brightening (Podladchikova et al. 2010; Innes et al 2010). They seem to be related chromospheric eruptions at the junctions of supergranular cells (Innes et al. 2009). In 171\\AA\\ solar disk images small-scale eruptions have similar characteristics to face-on CMEs but on a smaller scale. Apart from microflare brightening and dimming of the surrounding corona, several events have wave-like features propagating from the eruption site.}  Podladchikova et al. (2010) focused on the small-scale diffuse coronal fronts and their associated dimmings for a couple of events. The coronal wave and dimming properties were extracted using a semi-automatic algorithm adapted from their method developed for CMEs (Podladchikova \\& Berghmans 2005). They show that waves develop from micro-flaring sites and propagate up to a distance of 40,000~km in 20~min. The dimming region is two orders of magnitude smaller than for large-scale events. The waves associated with the micro-eruptions had a velocity of 17~\\kms, and were diagnosed as slow mode, propagating almost perpendicular to the background magnetic field.  Innes et al. (2009) estimated a rate of 1400 events per day over the whole Sun by picking out events by eye in a period of 24 hours. Now the Solar Dynamics Observatory/Atmospheric Imaging Assembly (SDO/AIA) is providing full Sun images through ten UV and EUV filters with a cadence of about one image every second. To optimize statistical analyses of various kinds of solar phenomena, it is necessary to develop automatic detection techniques. Aschwanden (2010) gives an extensive overview of solar image processing techniques that are used in automated feature detection algorithms. One of the modules in the Computer Vision Center for SDO is designed to automatically detect and extract coronal dimmings from AIA images (Attrill \\& Wills-Davey 2010).   Here, we give an automatic method to detect {small-scale dimmings} in the 171\\AA\\ STEREO/EUVI and SDO/AIA images using Zernike moments and the Support Vector Machine (SVM) classifier. The method exploits the characteristics of the Zernike moments in order to represent event/non-event patterns. The Zernike moments are invariant to rotation, scaling, and translation, and this allows the proposed method to detect events even if they are translated, scaled or rotated. First the Zernike moments are extracted for a selection of events and non-events, and fed into the SVM as training data sets so that it can successfully identify events/non-events in subsequent 171\\AA\\ image sequences.   The paper layout is as follows: data reduction is discussed in Section \\ref{data}; Zernike moments are explained in Section 3; the Support Vector Machine is discussed in Section \\ref{svm}; and the results and conclusions are presented in Section \\ref{con}. ",
        "conclusions": "\\label{con} \\subsection{Results} The automatic detection scheme  {for small-scale EUV dimmings} has been applied to full Sun STEREO/EUVI and SDO/AIA 171\\AA\\ data. On 13 June 2007, 1217 events were detected from the EUVI images. This is approximately twice the rate found by Innes et al. (2009) by eye for the same dataset. In the AIA images, 2064 events were detected on 27 August 2010. AIA is therefore giving another factor two in event rate. Figure~\\ref{fig7} shows the positions of events found during 20~min of EUVI and AIA images.   \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{fig6.eps} \\caption{Positions of events detected in 20 min by EUVI (left) and AIA.} \\label{fig7} \\end{figure}   The event numbers from EUVI are plotted versus time, \\sunx, and \\suny\\ in Figure~\\ref{fig8}. We see that approximately 60 events were detected every hour. On average, there are 45 events per 50\\arcsec\\ in both \\sunx\\ and \\suny. The numbers for AIA were double the EUVI numbers: 115 events per hour and 100 events for every 50\\arcsec\\ in space. For both EUVI and AIA, the events are mostly found in the central part of the disk (between +35 degrees and -35 degrees). They are most easily identified against the quiet Sun because in active regions they disappear into the general background activity and off-limb there is too much line-of-sight confusion (Innes et al. 2009).  {To make an estimate of how reliable the algorithm is we checked 400 of the detected events. We found that 30 of them are not clear events, like those shown in Figure~\\ref{fig6}. We also checked for events that the algorithm missed and could not find any additional events. Therefore we may be over-estimating small events numbers by about 10\\%.}  \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{fig7.eps} \\caption{Number of events per hour and per 50\\arcsec\\ detected in STEREO/EUVI images on 13 June 2007 plotted versus time (left) \\sunx\\ (middle), and \\suny\\ (right). } \\label{fig8} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{fig8.eps} \\caption{Three less clear events detected by SVM from AIA data. } \\label{fig6} \\end{figure}  The size and expansion rates of the dimmings have been computed with a region grow algorithm applied to the dimming region in the space-time slices (Figure~\\ref{fig9}). {First a region (2 by 2) including the the original $x_{min}$ and $t_{min}$ is selected (left panel of Figure \\ref{fig9}). Then the region is expanded to include all connected pixels which fall below the specified threshold. The width, $\\Delta x$, is the maximum width of the dimming region. The event duration time interval, $\\Delta t$, is the maximum along the time direction.} {Because the same event is often seen in several neighbouring $y$-slices, we first group events together and then apply the region grow algorithm to all events in the $x$ and $y$ directions. The longest time intervals and biggest widths are taken as the event sizes.}  {The size of the region depends on the threshold used. If it is too low then external dark areas will be included and if it is too high the region is too small. Our automatic method uses the full width at half maximum of the histograms of intensities of each space-time slice as the threshold. We tested this automatic level against manually chosen ones for a random 10\\%\\ of the slices. Here we took special care to ensure no external areas were included. The difference between the width and duration of the two cases (automatic thresholds and manual thresholds) are calculated. Only two out of 140 events differed by more than 2 spatial pixels (1000~km). The duration has a slightly large uncertainty: 15 of the 140 events differed by for than 2 time steps (180~s).}   The histograms of the AIA and EUVI sizes are shown in Figure~\\ref{fig10}. For larger events the slopes are the same (about -5) for both instruments. The difference between the two curves is that AIA curve turns over at a smaller size. The size refers to the central core dimmings not the event size so comparison with the large-scale CME distribution (Schrijver 2010) is not possible.  {We compute the average velocity ($\\Delta x/\\Delta t$) for 1325 AIA dimming regions. } The histogram of average velocities are shown in Figure~\\ref{fig11}. This histogram is similar to the distribution of CME's speeds obtained from solar cycle 23 (Mittal \\& Narian 2009). The average velocities concentrate in the range of 3-80~\\kms\\ with mean of 14~\\kms. This is same as the average mini-filament velocity measured by Wang et al. (2000). {Based on the uncertainly in space and duration the velocity uncertainties are small, typically 1~\\kms. Only 5\\%\\ have an uncertainly greater than 5~\\kms.}    \\subsection{Conclusion} An automatic method for the detection of  {small-scale EUV dimmings} has been developed and tested on SDO/AIA and STEREO/EUVI 171\\AA\\ images. The method exploits the fact that the Zernike moments of events that move have a specific pattern in space-time slices . This structure can be recognized by a Support Vector Machine classifier (SVM). Due to the invariant characteristics of the Zernike moments, the proposed method can detect events even if they are small or very faint.  {At the moment we detect motion of the dimming front but the algorithm is not able to distinguish between events due to coronal field evolution with and without eruption. Since eruptive events almost always have sudden brightening at onset, we are working on an algorithm that uses event groups to find the position and intensity of the brightest pixel at the earliest time in a group. Then the event characteristics relative to this brightening can be used to classify events as eruptive or not.}   \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{fig9.eps} \\caption{Demonstration of the mini-dimming extraction from a space-time slice, The frames are 7.8 arcsec by 10.9 min from SDO/AIA.} \\label{fig9} \\end{figure}    \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{fig10.eps} \\caption{Histogram of  {small dimming events} detected by AIA and EUVI.} \\label{fig10} \\end{figure}   \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{fig11.eps} \\caption{ The histogram of average velocities of 1325 events are plotted versus velocities. } \\label{fig11} \\end{figure}"
    },
    "1112/1112.1430_arXiv.txt": {
        "abstract": "The black hole mass function of supermassive black holes describes the evolution of the distribution of black hole mass. It is one of the primary empirical tools available for mapping the growth of supermassive black holes and for constraining theoretical models of their evolution. In this review we discuss methods for estimating the black hole mass function, including their advantages and disadvantages. We also review the results of using these methods for estimating the mass function of both active and inactive black holes. In addition, we review current theoretical models for the growth of supermassive black holes that predict the black hole mass function. We conclude with a discussion of directions for future research which will lead to improvement in both empirical and theoretical determinations of the mass function of supermassive black holes. ",
        "introduction": "\\label{s-intro}  Understanding how and when supermassive black holes (SMBHs) grow is currently of central importance in extragalactic astronomy. A significant amount of empirical work has established correlations between SMBH mass and host galaxy spheroidal properties, such as luminosity \\cite{korm95,mclure01,mclure02}, stellar velocity dispersion (the $M_{BH}$--$\\sigma_*$ relationship, e.g., \\cite{gebh00a, merr01, trem02}), concentration or Sersic index \\cite{graham01,graham07}, bulge mass \\cite{mag98,marc03,haring04}, and binding energy \\cite{aller07,fundplane}. These scaling relationships imply that the evolution of spheroidal galaxies and the growth of SMBHs are intricately tied together. The currently favored mechanism for linking the growth of SMBHs and their hosts is black hole feedback, whereby black holes grow by accreting gas in so-called ``active'' phases, possibly fueled by a major merger of two gas-rich galaxies, until feedback energy from the SMBH expels gas and shuts off the accretion process \\cite{silk98,fabian99,begel05,murray05}. Alternatively, it has been suggested that the origin of the scaling relationships does not necessarily require SMBH feedback, but emerges from the stochastic nature of the hierarchical assembly of black hole and stellar mass through galaxy mergers. \\cite{peng07,jahnke11}.  Feedback-driven `self-regulated' growth of black holes has been able to reproduce the local $M_{BH}$--$\\sigma_*$ relationship in smoothed particle hydrodynamics simulations \\cite{dimatteo05,spring05,johan09}. Moreover, AGN feedback has also been invoked as a means of quenching the growth of the most massive galaxies \\cite{bower06,croton06}. There have been numerous models linking SMBH growth, the quasar phase, and galaxy evolution \\cite{haehn98,kauff00,haehn00,wyithe03,vol03,catt05,dimatteo08, somer08,booth09}. While feedback is likely important for regulating the growth of SMBHs and galaxies, the fueling mechanisms that contribute to growing the SMBH are likely diverse. Major-mergers of gas-rich galaxies may fuel quasars at high redshift, and grow the most massive SMBHs. However, major-mergers alone do not appear to be sufficient to reproduce the number of X-ray faint AGN \\cite{marulli07}, and accretion of ambient gas via internal galactic processes \\cite{ciotti01,stochacc}, may fuel these fainter, lower $M_{BH}$ AGN at lower $z$. This is supported by the fact that many AGN are observed to live in late-type galaxies out to $z \\approx 1$ \\cite{guyon06,gabor09}, and the X-ray luminosity function of AGN hosted by late-type galaxies suggests that fueling by minor interactions or internal instabilities represents a non-negligible contribution to the accretion history of the Universe \\cite{georg09}.  The black hole mass function (BHMF) provides a complete census of the mass of SMBHs and their evolution. Because of this, the BHMF is one of the primary empirical tools available for investigating the growth of SMBHs, and for constraining theoretical models for the growth of the SMBH population. Because SMBHs and galaxies are thought to be linked in their evolution, the BHMF provides insight into the fueling mechanisms that dominate black hole growth, and therefore into the role of feedback in the evolution of the host galaxy. The BHMF is also an important tool in planning future surveys, as it provides an estimate of the distribution of SMBH mass expected for the survey. This in turn is important because mass is a fundamental quantity of the black hole, and therefore is an important observational quantity for empirical studies of black hole accretion physics \\cite{merloni03,falcke04,mchardy06,davis11,sob11,kelly11}. Of course, further improvement to our understanding of black hole accretion physics will further improve our modeling and understanding of black hole accretion and feedback, which in turn will improve our understanding of black hole-galaxy coevolution. Therefore, the BHMF is an important empirical quantity for SMBH studies.  In this review we discuss the current status of BHMF estimation and theoretical modeling. In \\S~\\ref{s-est} we discuss the non-trivial task of estimating the BHMF. In \\S~\\ref{s-local} we discuss current estimates of the local SMBH BHMF. In \\S~\\ref{s-continuity} we discuss BHMF estimates derived by combining the local BHMF with the AGN luminosity function via a continuity equation. In \\S~\\ref{s-agn} we discuss BHMFs estimated for AGN only. In \\S~\\ref{s-theory} we review theoretical models for SMBH growth that predict the SMBH BHMF. Finally, in \\S~\\ref{s-future} we discuss directions for future improvements to the empirical and theoretical studies of the BHMF. We note that unlike, say, the luminosity function, the division between `observational' and `theoretical' studies is not as clear for the BHMF, as some amount of modeling is necessary in order to estimate the BHMF from strictly observational quantities. We have attempted to divide the studies according to whether the BHMF is constrained empirically, as in, say, a formal statistical fitting procedure, or if it is predicted from a theoretical model for SMBH growth. In reality, the line between theoretical and empirical studies is blurry, and some procedures which we have considered to be empirical may be thought of as theoretical. ",
        "conclusions": ""
    },
    "1112/1112.4215_arXiv.txt": {
        "abstract": "%  A technological milestone for experiments employing Transition Edge Sensor (TES) bolometers operating at sub-kelvin temperature is the deployment of detector arrays with 100s--1000s of bolometers. One key technology for such arrays is readout multiplexing: the ability to read out many sensors simultaneously on the same set of wires. This paper describes a frequency-domain multiplexed readout system which has been developed for and deployed on the APEX-SZ and South Pole Telescope millimeter wavelength receivers. In this system, the detector array is divided into modules of seven detectors, and each bolometer within the module is biased with a unique $\\sim$MHz sinusoidal carrier such that the individual bolometer signals are well separated in frequency space. The currents from all bolometers in a module are summed together and pre-amplified with Superconducting Quantum Interference Devices (SQUIDs) operating at 4\\,K. Room-temperature electronics demodulate the carriers to recover the bolometer signals, which are digitized separately and stored to disk. This readout system contributes little noise relative to the detectors themselves, is remarkably insensitive to unwanted microphonic excitations, and provides a technology pathway to multiplexing larger numbers of sensors. ",
        "introduction": "A new generation of mm-wavelength receivers instrumented with hundreds or thousands of transition edge sensor (TES) bolometric detectors, e.g., \\cite{2011ApJS..194...41S,apexSZ_instrument,2010ApJ...711.1141T,2009arXiv0907.4445C,2008AIPC.1040...66L}, is allowing for unprecedented precision in measurements of the Cosmic Microwave Background (CMB). Large TES arrays are expected to bring forth similar advances in other observational bands such as the sub-mm (e.g., \\cite{2006SPIE.6275E..45H}), X-ray~(e.g., \\cite{2008JLTP..151..363S}), and gamma-ray~(e.g., \\cite{2007ApPhL..90s3508D,2004ApPhL..85.4762D,dreyer_llnl2007}) as well as phonon detection in dark matter direct detection experiments (e.g., \\cite{Wang:1990qk}).  TES detectors in the millimeter to far-IR range are approaching the photon statistics limit, meaning their noise performance is dominated by fluctuations in the arrival rate of photons at the detector, and not by the intrinsic noise performance of the detectors themselves. Consequently, higher mapping speed can best be realized by increasing the number of detectors, rather than by improving the detectors themselves. A key challenge in building focal planes with large sensor counts has been multiplexing the signals between the cold sub-kelvin operating temperature of the detectors and room temperature electronics. Multiplexing minimizes the heat-load on the detector cold stage and reduces the complexity of the cold wiring.  Two complementary multiplexing strategies are commonly used: (1) time domain multiplexing (\\tMUX)~\\cite{chervenak99,2000NIMPA.444..107C, 2002AIPC..605..301I,2008JLTP..151..908B}, where the detectors are read out sequentially one at a time with a revisit rate that is faster than the detector time constant, and (2) frequency-domain multiplexing (\\fMUX) where each detector is biased at a unique location in frequency space and read out continuously. While the \\tMUX\\ system was developed earlier, both technologies are now quite mature. \\mynewtext{ An eight-channel \\tMUX\\ proof-of-concept system was first to make sky observations~\\cite{2003SPIE.4855..100S}. } A new TES readout scheme that multiplexes signals from GHz frequency RF \\squid s~\\cite{2004ApPhL..85.2107I,2006NIMPA.559..802I} is in the early stages of development, as is a new code-domain multiplexing readout~\\cite{2010SuScT..23c4004I}.  In this paper, we describe the \\fMUX\\ readout system \\cite{yoon01,2002AIPC..605..305Y,Spieler02,2003SPIE.4855..172L,2004NIMPA.520..548L,2005ApPhL..86k2511L,2006NIMPA.559..793L,2006PhDT........32L} as it has been deployed for mm-wavelength observations with the APEX-SZ~\\cite{apexSZ_instrument,2006NewAR..50..960D,2003NewAR..47..933S} and South Pole Telescope (SPT) SZ~\\cite{2009arXiv0907.4445C,2004astro.ph.11122R} receivers. \\mynewtext{ These experiments are designed for high angular resolution observations of the CMB.  For example, the unique combination of sensitivity and resolution of these instruments makes it possible to study the structure and evolution of galaxy clusters using the Sunyaev-Zel'dovich (SZ) effect~\\cite{1972CoASP...4..173S}, a small distortion imprinted on the CMB when primordial photons inverse-Compton scatter off hot electrons in the intra-cluster medium. }  The resolved SZ image of the Bullet Cluster (1E 0657-56) obtained with APEX-SZ~\\cite{2008arXiv0807.4208H} was the first published scientific result with a large array of multiplexed TES detectors. SPT achieved the first SZ detections of previously unknown galaxy clusters~\\cite{2008arXiv0810.1578S} with its array of multiplexed TES detectors. \\mynewtext{ A wealth of new measurements have been presented by both SPT~\\cite{ 2011arXiv1111.0932R, 2011ApJ...743...28K,2011ApJ...735L..36S,2011ApJ...738..139W, 2011ApJ...736...61S,2010ApJ...721...90B,2011ApJ...738...48A, 2010ApJ...722.1180V,2010ApJ...723.1736H,2010ApJ...719.1045L, 2010ApJ...718..632H,2010ApJ...719..763V,2010ApJ...716.1118P} and APEX-SZ~\\cite{2010A&A...519A..29B,2009ApJ...701.1958R,2009A&A...506..623N}. }  Other implementations of frequency multiplexed TES readout systems are described in Refs.~\\cite{2009AIPC.1185..245V,2008ITNS...55...21D}.   This paper is outlined as follows. In \\S\\ref{s_general_overview}, a brief overview of the system, its advantages, and the performance requirements is presented. A detailed description of the sub-kelvin components of the system is provided in \\S\\ref{s_cold_mux}. This description is accompanied by a summary of the design process for specifying the basic multiplexer parameters, such as channel spacing. \\S\\ref{s_squid_and_electronics} details the \\squid\\ pre-amplifier and back-end electronics. The performance for the system operating in the SPT and APEX-SZ experiments is described in \\S\\ref{s_performance}, and conclusions are presented in \\S\\ref{s_conclusions}. ",
        "conclusions": "\\label{s_conclusions}  We have presented the design and performance of the SQUID-based frequency-domain multiplexed (\\fMUX) readout system that has been deployed for TES detectors on the APEX-SZ instrument and the South Pole Telescope SZ-camera. The readout reduces the heat-load on the sub-kelvin cryogenic stage by sharing wiring among many detectors using a simple architecture that minimizes the sub-kelvin system complexity. It has near-zero power dissipation on the sub-kelvin stage, is extensible to much higher multiplexing factors, allows bias levels to be set individually for each detector, is modular with no shared components between readout modules, and exhibits strong reduction of microphonic response and low-frequency EMI because the sky signals are modulated up to high frequency.  The system has been characterized in terms of noise performance, sensitivity, detector yield, and cross-talk in both the laboratory and with on-sky observations. The measured noise equivalent power of the system is consistent with theoretical expectations, and contributions from the readout system are small in comparison to detector and photon noise above 1~Hz. At low frequencies, the system noise is dominated by residual atmospheric fluctuations.  These implementations have shown that frequency-domain multiplexing is an effective technology for scaling to very large arrays of TES detectors. An upgraded digital version of the analog \\fMUX\\ back-end electronics has recently been developed~\\cite{2008ITNS...55...21D} and has been deployed for the EBEX~\\cite{2010SPIE.7741E..52A} balloon-borne CMB polarimeter and the ground-based POLARBEAR~\\cite{2010SPIE.7741E..39A} CMB polarimeter. \\mynewtext{ A polarimeter camera~\\cite{2009AIPC.1185..511M} for the South Pole Telescope is being commissioned during the austral 2011/2012 summer and uses the digital \\fMUX. } The digital \\fMUX\\ system achieves an order of magnitude lower power consumption, higher multiplexing factors, faster detector and \\squid\\ setup, better low frequency noise performance, and lower cost. Development is underway to develop \\fMUX\\ electronics for satellite applications~\\cite{2011_FMUX_FOR_FLIGHT}. \\revisedtext{ Present prototypes use Digital Active Nulling~\\cite{2012_DAN_DAF} in place of the flux-locked loop. This extends the \\squid\\ bandwidth to allow for 64 bolometers per \\squid\\ module using a single pair of sub-Kelvin wires. }"
    },
    "1112/1112.5163_arXiv.txt": {
        "abstract": "This paper reports follow-up photometric and spectroscopic observations, including warm {\\em Spitzer} IRAC photometry of seven white dwarfs from the SDSS with apparent excess flux in UKIDSS $K$-band observations.  Six of the science targets were selected from 16\\,785 DA star candidates identified either spectroscopically or photometrically within SDSS DR7, spatially cross-correlated with $HK$ detections in UKIDSS DR8.  Thus the selection criteria are completely independent of stellar mass, effective temperature above 8000\\,K, and the presence (or absence) of atmospheric metals.  The infrared fluxes of one target are compatible with a spatially-unresolved late M or early L-type companion, while three stars exhibit excess emissions consistent with warm circumstellar dust.  These latter targets have spectral energy distributions similar to known dusty white dwarfs with high fractional infrared luminosities (thus the $K$-band excesses). Optical spectroscopy reveals the stars with disk-like excesses are polluted with heavy elements, denoting the ongoing accretion of circumstellar material.  One of the disks exhibits a gaseous component -- the fourth reported to date -- and orbits a relatively cool star, indicating the gas is produced via collisions as opposed to sublimation, supporting the picture of a recent event.  The resulting statistics yield a lower limit of 0.8\\% for the fraction dust disks at DA-type white dwarfs with cooling ages less than 1\\,Gyr.  Two overall results are noteworthy: all stars whose excess infrared emission is consistent with dust are metal-rich; and no stars warmer than 25\\,000\\,K are found to have this type of excess, despite sufficient sensitivity. ",
        "introduction": "Near-infrared photometry of white dwarfs examines a scientific phase space that is difficult or impossible to probe for main-sequence stars.  The first to recognize this potential was \\citet{pro81}, who realized the compact nature of white dwarfs allowed a straightforward photometric detection of very cool stellar and substellar companions.  Brown dwarfs, which often require sophisticated detection techniques when searched for at main-sequence stars, are up to 10 times larger than a typical white dwarf and can be readily detected in {\\em spatially-unresolved} observations as excess near-infrared flux \\citep{pro83}. Thus, white dwarfs are an excellent tool for studies of the low-mass stellar and substellar mass function via companions \\citep{far05,zuc92,pro82}.  Photometry with {\\em Spitzer} has extended this potential to include closely orbiting brown dwarfs and massive planets although no candidates are yet apparent \\citep{far08b,mul07}.  \\begin{table*} \\begin{center} \\caption{Multi-Wavelength Photometry and Derived Parameters for SDSS White Dwarfs\\label{tbl1}} \\begin{tabular}{@{}rccccccc@{}} \\hline\\hline  SDSS Prefix\t\t\t\t\t&0959\t\t\t&1159$^a$\t\t&1221\t\t\t&1247\t\t\t&1320\t\t\t&1506\t\t\t&1557\\\\  \\hline  Spectral Type\t\t\t\t\t&DAZ\t\t\t&DQp\t\t\t&DAZ\t\t\t&DA\t\t\t\t&DA\t\t\t\t&DA\t\t\t\t&DAZ\\\\ Spectroscopic Data\t\t\t\t&WHT\t\t\t&SDSS\t\t\t&WHT\t\t\t&SDSS\t\t\t&SDSS\t\t\t&WHT\t\t\t&SDSS+WHT\\\\ $\\alpha$ (J2000; h\\,m\\,s)\t\t\t&09 59 04.69\t\t&11 59 33.10\t\t&12 21 50.81\t\t&12 47 40.93\t\t&13 20 44.68 \t\t&15 06 26.18\t\t&15 57 20.77\\\\ $\\delta$ (J2000; $\\degr \\ ' \\ ''$)\t\t&$-$02 00 47.6\t\t&$+$13 00 31.6\t&$+$12 45 13.3\t&$+$10 35 56.1\t&$+$00 18 54.9 \t&$+$06 38 45.9\t&$+$09 16 24.6\\\\ FUV (0.15\\,$\\mu$m; AB\\,mag)\t\t&...\t\t\t\t&$>20$\t\t\t&$19.55\\pm0.06$\t&$18.23\\pm0.02$\t&$16.52\\pm0.03$\t&$18.40\\pm0.02$\t&$18.03\\pm0.02$\\\\ NUV (0.23\\,$\\mu$m; AB\\,mag)\t\t&...\t\t\t\t&$19.62\\pm0.09$\t&$18.78\\pm0.02$\t&$18.33\\pm0.01$\t&$16.76\\pm0.01$\t&$17.31\\pm0.01$\t&$18.16\\pm0.01$\\\\ $u$ (0.36\\,$\\mu$m; AB\\,mag)\t\t&$18.52\\pm0.02$\t&$18.22\\pm0.03$\t&$18.54\\pm0.02$\t&$18.63\\pm0.02$\t&$17.08\\pm0.02$\t&$16.96\\pm0.01$\t&$18.50\\pm0.02$\\\\ $g$ (0.47\\,$\\mu$m; AB\\,mag)\t\t&$18.09\\pm0.03$\t&$18.14\\pm0.02$\t&$18.19\\pm0.02$\t&$18.50\\pm0.01$\t&$17.13\\pm0.01$\t&$16.57\\pm0.01$\t&$18.45\\pm0.01$\\\\ $r$ (0.62\\,$\\mu$m; AB\\,mag)\t\t&$18.36\\pm0.02$\t&$17.75\\pm0.01$\t&$18.40\\pm0.02$\t&$18.82\\pm0.02$\t&$17.47\\pm0.02$\t&$16.63\\pm0.01$\t&$18.81\\pm0.01$\\\\ $i$ (0.75\\,$\\mu$m; AB\\,mag)\t\t&$18.54\\pm0.02$\t&$17.70\\pm0.02$\t&$18.56\\pm0.02$\t&$19.13\\pm0.02$\t&$17.72\\pm0.02$\t&$16.79\\pm0.01$\t&$19.13\\pm0.01$\\\\ $z$ (0.89\\,$\\mu$m; AB\\,mag)\t\t&$18.88\\pm0.06$\t&$17.82\\pm0.02$\t&$18.86\\pm0.05$\t&$19.33\\pm0.06$\t&$18.05\\pm0.03$\t&$16.93\\pm0.01$\t&$19.34\\pm0.04$\\\\ $J$ (1.25\\,$\\mu$m; mag)\t\t\t&$18.32\\pm0.06$\t&$17.40\\pm0.03$\t&$18.43\\pm0.07$\t&$18.95\\pm0.08$\t&$17.62\\pm0.04$\t&$16.41\\pm0.01$\t&$18.82\\pm0.06$\\\\ $H$ (1.64\\,$\\mu$m; mag)\t\t\t&$18.17\\pm0.09$\t&$17.32\\pm0.06$\t&$18.39\\pm0.10$\t&$18.90\\pm0.11$\t&$17.71\\pm0.08$\t&$16.37\\pm0.02$\t&$19.03\\pm0.14$\\\\ $K$ (2.20\\,$\\mu$m; mag)\t\t\t&$17.62\\pm0.11$\t&$17.08\\pm0.08$\t&$18.01\\pm0.15$\t&$18.47\\pm0.20$\t&$17.53\\pm0.11$\t&$16.35\\pm0.03$\t&$18.35\\pm0.15$\\\\ IRAC1 (3.55\\,$\\mu$m; $\\mu$Jy)\t&$80\\pm4$\t\t&$42\\pm2$\t\t&$40\\pm5$\t\t&$25\\pm2$\t\t&$35\\pm2$\t\t&$75\\pm4$\t \t&$35\\pm2$\\\\ IRAC2 (4.49\\,$\\mu$m; $\\mu$Jy)\t&$76\\pm4$\t\t&$28\\pm2$\t\t&$40\\pm6$\t\t&$17\\pm2$\t\t&$25\\pm2$\t\t&$50\\pm3$\t \t&$40\\pm2$\\\\  $T_{\\rm eff}$ (K)\t\t\t\t&$13280\\pm20$\t&$9500\\pm500$\t&$12250\\pm20$\t&$18540\\pm150$\t&$20220\\pm180$\t&$10670\\pm10$\t&$22810\\pm40$\\\\ $\\log\\,g$ (cm\\,s$^{-2}$)\t\t\t&$8.06\\pm0.03$\t&8.0\t\t\t\t&$8.20\\pm0.03$\t&$7.86\\pm0.03$\t&$8.37\\pm0.04$\t&$8.21\\pm0.03$\t&$7.54\\pm0.01$\\\\ $M$ ($M_{\\odot}$)\t\t\t\t&$0.64\\pm0.02$\t&0.59\t\t\t&$0.73\\pm0.02$\t&$0.54\\pm0.02$\t&$0.86\\pm0.02$\t&$0.73\\pm0.02$\t&$0.42\\pm0.01$\\\\ $d$ (pc)\t\t\t\t\t\t&$203\\pm4$\t\t&$110\\pm00$\t\t&$180\\pm4$\t\t&$396\\pm9$\t\t&$150\\pm4$\t\t&$71\\pm2$\t\t&$566\\pm2$\\\\ IR Excess$^b$\t\t\t\t\t&Dust\t\t\t&Atm\t\t\t&Dust\t\t\t&dM/L\t\t\t&Bkgd\t\t\t&None\t\t\t&Dust\\\\  \\hline  \\end{tabular} \\end{center}  \\flushleft {\\em Note}.  Ultraviolet, optical, and near-infrared photometry are from {\\em GALEX} \\citep{mar05}, SDSS \\citep{aba09}, and UKIDSS \\citep{law07} respectively.   For stars with multiple spectroscopic or photometric datasets, table values are the weighted average of available measurements.  The SDSS photometry is given in PSF magnitudes; the only catalog entries appropriate for point sources.  Stellar parameters are derived by fitting the Balmer line profiles as described in \\S3. \\smallskip  $^a$This star was selected independently from the six primary science targets (see \\S3.2). \\smallskip  $^b$See \\S3 for the various infrared excess descriptions.  \\end{table*}  Another way in which near-infrared photometry of white dwarfs is advantageous is that circumstellar dust orbiting within the Roche limit of the stellar remnant becomes heated sufficiently to emit in this wavelength range \\citep{kil06}.  This potential is precluded for main-sequence stars because the analogous spatial region, where km-size or larger solid bodies are tidally destroyed, does not extend significantly above their surfaces \\citep{dav99}.  Furthermore, any material generated in that narrow region above the star will rapidly dissipate due to radiation pressure and drag forces.  Thus, while asteroids and comets commonly pass sufficiently close to the Sun, and presumably other main-sequence stars that host planetary systems, the tidal disruption and subsequent accretion of this debris can only be witnessed at white dwarfs (and possibly neutron stars; \\citealt{wan06}).  This paper reports warm {\\em Spitzer} IRAC measurements and optical spectroscopy of six white dwarf candidates selected for apparent $K$-band excess fluxes based on the ground-based photometric surveys the Sloan Digital Sky Survey (SDSS) and the UKIRT Infrared Deep Sky Survey (UKIDSS).  Target selection criteria and survey data are given in \\S2, along with a description of the {\\em Spitzer} and ground-based observations.  The spectroscopic and photometric data analysis is presented in \\S3, with a summary of the overall results in \\S4. ",
        "conclusions": "\\subsection{An Unbiased Dust Disk Frequency}  What makes this study so far unique is that previous IRAC programs have either targeted known metal-rich white dwarfs, or near-infrared bright white dwarfs.  While the first approach has been highly successful and accounts for the bulk of white dwarf disk discoveries to date \\citep{far10b,far09b,jur07}, two drawbacks exist. First, it requires 8\\,m class telescope time using high-resolution spectroscopy to efficiently identify (weak) metal lines in white dwarfs \\citep{koe05,zuc03}, and second, only about 20\\% of metal-contaminated white dwarfs have disks \\citep{far09b}.  The second method, employed by the \\citet{mul07} survey of more than 130 white dwarfs with $K_s<15$\\,mag, was unbiased by the presence or absence of atmospheric metals, but only two stars with dust were found \\citep{von07}, and is hence only 1.5\\% efficient by number of targets. Here, five candidate DA white dwarfs with $K$-band excess have been selected with a success rate of 60\\% for circumstellar dust and 80\\% for confirmed infrared excesses physically associated with the system.  One of five targets has an infrared excess probably due to a background galaxy.  Based on the candidate selection criteria and numbers given in Table 6 of \\citet{gir11}, of 1884 candidate and confirmed DA white dwarfs with the necessary UKIDSS detections, there are 12 stars whose data are consistent with disk-like, $K$-band excesses.  Several years of Spitzer studies have shown that only 50\\% of dusty white dwarfs have such $K$-band excesses \\citep{far10b}, implying that a decent estimate of the disk fraction among DA white dwarfs is at least 1.2\\%. However, as seen in this study, in some cases the $K$-band data can be the result of neighboring background objects, as well as real companions such as low mass stars or brown dwarfs.  Thus a more realistic estimate is $3/5$ of this number, implying at least 0.8\\% of DA (and presumably all) white dwarfs have dust disks (and atmospheric metals), consistent with previous estimates \\citep{far09b}.  These results only apply to stars with cooling ages less than 1\\,Gyr, as this corresponds to the approximate cutoff in effective temperature (8000\\,K; \\citealt{fon01}) for the DA color selection.  However, {\\em Spitzer} studies have shown a distinct lack of dust disks at older and cooler metal-contaminated stars; only G166-58 has an (anomalous) infrared excess and a cooling age beyond 1\\,Gyr \\citep{far08b}.  \\subsection{A Lack of $K$-Band Emitting Dust Disks at Warm to Hot White Dwarfs}  Interestingly, the DA selection criteria is quite sensitive to warm to hot, hydrogen-rich white dwarfs, as the Balmer decrement gives these stars unique colors in that wavelength region \\citep{gir11}. However, there are no candidates for dust among thousands of stars(!).  Either dust\tdisks\t at white dwarfs with\t$T_{\\rm eff} >25\\,000$\\,K exist yet rarely exhibit $K$-band excess or dust at such stars is itself rare.  The latter possibility is somewhat contradictory based on the paradigm of planetary systems that are dynamically rejuvenated at the onset of the post-main sequence \\citep{bon11,bon10,deb02}.  If this apparent dearth of warmer systems at $K$-band is real, it would likely require a physical mechanism to preclude such emission (such as rapid removal of warm grains, or long timescales for dust disk spreading toward the star).  \\citet{far11} plot the fractional infrared luminosity of all known white dwarfs with dust circa mid-2010 as a function of cooling age, revealing a potential trend of increasing dust emission towards cooler stars.  This could be a natural consequence of the potential for cooler white dwarfs to host wider disks: as a star cools and its luminosity decreases, dust grains may persist closer to the stellar surface prior to sublimation.  Thus, any disk extending inward from the Roche limit has the potential to be wider at cooler white dwarfs.  Conversely, at increasingly warmer stars, the inner dust disk edge will be physically further from the star as the radius at which grains are rapidly sublimated increases (see Table \\ref{tbl4}), eventually exceeding the Roche limit for large solid bodies \\citep{von07}.  While it is uncertain at which precise effective temperature an evolving dust disk might be completely sublimated, the results of this work clearly show that disk emission analogous to G29-38 is rare among warm and hot white dwarfs.  More sensitive surveys with {\\em Spitzer} and {\\em WISE} are needed to address this matter more comprehensively."
    },
    "1112/1112.4868_arXiv.txt": {
        "abstract": "{  We introduce a new method to determine the redshift of unknown-redshift BL Lac Objects. The method relies on simultaneous multi-wavelength (MWL) observations of BL Lac objects in optical, X-ray, HE (E>100 MeV) gamma-rays and VHE (E>100 GeV) gamma-rays. It involves best-fitting spectral energy distribution (SED) from optical through HE gamma-rays with a Synchrotron-Self-Compton (SSC) model. We extrapolate such best fitting model into VHE regime, and assume that it represents the intrinsic emission of the object. We then compare the observed VHE flux which has been affected by the interaction with Extragalactic Background Light (EBL). Constraining the measured vs intrinsic emission leads to the determination of gamma-gamma opacity. Comparing the obtained opacity with the predicted opacity based on EBL model, we obtain the redshift of the photon source.  } ",
        "introduction": "In a recent paper \\cite{bib:man010}, some of us (NM, MP and VT) have shown how it is possible to use simultaneous multiwavelength (i.e. optical, X-ray, HE\\,$\\gamma$, and VHE\\,$\\gamma$) data to obtain a measure the Extragalactic Background Light (EBL). EBL is the integrated light produced and re-processed in all the history of the universe. Knowledge of the EBL is particularly important in studying the VHE\\,$\\gamma$-ray emission from extragalactic sources (typically, blazars), because interaction of TeV photons with the EBL attenuates their observed intensity while the resulting cascade will enhance the observed GeV emission. These effects are clearly more marked for more distant sources, because of the larger optical depth, $\\tau_{\\gamma \\gamma}$, for the $\\gamma_{\\rm VHE}-\\gamma_{\\rm EBL}$ interaction: for a given EBL model, the quantitative details of the modification to the VHE\\,$\\gamma$-ray spectrum emitted by a source will only depend on the redshift of the source.  By reversing the line of reasoning in \\cite{bib:man010}, if we assume to have a reliable EBL model, then simultaneous multi-$\\lambda$ observations of a source may allow for an independent derivation of the source redshift, $z$. For instance: (1) low- and high-energy (eV, keV, GeV) spectral-energy-distribution (SED) data (that are unaffected by EBL) are fitted with an emission model; (2) the model SED is extrapolated into the TeV domain, and the extrapolation is assumed to represent the insinsic (emitted) TeV emission; (3) the 'emitted' VHE spectrum is compared with the measured VHE spectrum, and the discrepancy between the two is taken to be a measure of $\\tau_{\\gamma \\gamma}$; (4) given a specific EBL model, $\\tau_{\\gamma \\gamma}$ can be solved as a function of the source distance (i.e., for a given cosmology, of $z$).  In what follows we present preliminary results of an approach to the problem of distance determination for a blazar, which is slightly different from the one above and in which we employ a dynamical EBL correction during a SED fit in which $z$ also appears as a parameter. Just below, we describe the source that we are using for our analysis. Then we provide some details about the numerical algorithm that, elaborating what is done in \\cite{bib:man011}, is used to perform the numerical fit to a chosen SED model which will return, {\\rm inter alia}, an estimation of the source redshift. ",
        "conclusions": ""
    },
    "1112/1112.3165_arXiv.txt": {
        "abstract": "The \\textit{Fermi} Large Area Telescope (LAT) has revolutionized high-energy (HE) astronomy, and is making enormous contributions particularly to $\\gamma$-ray pulsar science. As a result of the many new pulsar discoveries, the $\\gamma$-ray pulsar population is now approaching~100. Some very famous millisecond pulsars (MSPs) have also been detected: J1939+2134 (B1937+21), the first MSP ever discovered, as well as J1959+2048 (B1957+20), the first black widow pulsar system. These, along with other MSPs such as PSR~J0034$-$0534 and J2214+3000, are rare among the pulsar population in that they exhibit nearly phase-aligned radio and $\\gamma$-ray light curves (LCs). Traditionally, pulsar LCs have been modelled using standard HE models in conjunction with low-altitude conal beam radio models. However, a different approach is needed to account for phase-aligned LCs. We explored two scenarios: one where both the radio and $\\gamma$-ray emission originate in the outer magnetosphere, and one where the emission comes from near the polar caps (PCs) on the stellar surface. We find best-fit LCs using a Markov chain Monte Carlo (MCMC) technique for the first class of models. The first scenario seems to be somewhat preferred, as is also hinted at by the radio polarization data. This implies that the phase-aligned LCs are possibly of caustic origin produced in the outer magnetosphere, in contrast to the usual lower-altitude conal beam radio models. Lastly, we constrain the emission altitudes with typical uncertainties of $\\sim10\\%$ of the light cylinder radius. The modelled pulsars are members of a third $\\gamma$-ray MSP subclass, in addition to two others with non-aligned radio and $\\gamma$-ray LCs. ",
        "introduction": "For quite some time, there were only a handful of pulsars detected in the $\\gamma$-ray waveband~\\cite{Thompson04}. This small sample already provided a starting point for spectral and light curve (LC) modelling as well as preliminary population studies. However, the launch of the \\textit{Fermi} Large Area Telescope (LAT)~\\cite{Atwood09} heralded an exciting new era, particularly for high-energy (HE) pulsar physics.    \\subsection{$\\gamma$-ray pulsars galore} Following the first \\textit{Fermi} Catalog which included about 1~500 $\\gamma$-ray sources~\\cite{Ferim_1st_Cat}, the second \\textit{Fermi} Catalog is in now production, containing nearly 1~900 $\\gamma$-ray sources~\\cite{Fermi_2nd_CatA,Fermi_2nd_CatB} including almost~100 $\\gamma$-ray pulsars~\\cite{PSR_cat10,Smith11}. About a third of these are millisecond pulsars (MSPs), some having been found by observing non-variable, unassociated \\textit{Fermi} sources at high latitudes with radio telescopes and searching for pulsed radio signals~\\cite{Cognard11,Keith11,Ransom11}. Some of these MSPs exhibit the unusual phenomenon that the peaks of their $\\gamma$-ray and radio LCs occur at the same observer phase (`longitude' $\\phi$), i.e., they are phase-aligned.  \\subsection{Traditional models} Two classes of models have been invoked to explain HE pulsar radiation. The first class, low-altitude polar cap (PC) models~\\cite{Daugherty82,DH96}, assume that primary electrons are accelerated above the neutron star surface, and that magnetic pair production of curvature radiation or inverse-Compton-scattered $\\gamma$-rays occurs in the intense B-fields close to the stellar surface. In the case that pair creation is suppressed along the last open field lines, which mark the boundary between the corotating closed field line zone and the active open one, a slot gap (SG)~\\cite{Arons83} may  form, corresponding to a two-pole caustic (TPC) geometry~\\cite{Dyks03} which extends from the stellar surface up to near the light cylinder. The second class of models are the outer gap (OG) models~\\cite{CHR86a,Romani96}, which assume the production of HE radiation along the last open field lines above the so-called null charge surface, where the Goldreich-Julian charge density becomes zero. In both classes of outer-magnetospheric models (SG and OG), the HE pulse profiles are the result of the formation of \\textit{caustics} -- the accumulation of photons in narrow phase bands due to a combination of special relativistic effects and magnetic field line curvature. Also, the narrow gaps in these models require abundant production of electron-positron pairs which will be able to screen the accelerating E-field.  For the case of MSPs, LCs and spectra have been modelled~\\cite{Frackowiak05,HUM05,VdeJ05} using a pair-starved polar cap (PSPC) model \\cite{MH04_PS}. The latter is similar to the traditional PC model, but in this case the number of pairs is not sufficient to screen the accelerating E-field, such that the accelerating region above the star is pair-starved. MSP spectra and energetics have also been modelled using an OG model~\\cite{Zhang03,Zhang07}, while LC modelling using an annular gap model can furthermore reproduce the salient features of the $\\gamma$-ray LCs~\\cite{Du10}.  The outer magnetospheric models usually invoke relatively low-altitude `core' and `conal' radio beams centred on the magnetic axis, in addition to the HE radiation being produced in extended regions reaching up to the light cylinder. The difference in the location of the $\\gamma$-ray and radio emission regions thus implies that the corresponding pulse profiles will have non-zero lags (phase offsets) between the $\\gamma$-ray and radio LCs.  \\subsection{Phase-aligned radio and $\\gamma$-ray light curves: a new MSP subclass} The LCs of the first~8 \\textit{Fermi}-detected $\\gamma$-ray MSPs~\\cite{Abdo09_MSP} have been modelled~\\cite{Venter_MSP09}, yielding two distinct MSP subclasses: those whose LCs are well fit by a standard OG or TPC model, and those whose LCs are well fit by a PSPC model. These fits are mutually exclusive. Importantly, this unexpectedly implied that there is copious pair production taking place even in MSP magnetospheres with their characteristic low B-fields.  A third class of MSPs emerged with the discovery of PSR~J0034$-$0534~\\cite{Abdo10_J0034}, which exhibited (nearly) phase-aligned radio and $\\gamma$-ray LCs. Prior to this detection, such behaviour has only been observed for the Crab pulsar~\\cite{Kniffen74}. This phenomenon has now also been seen for \\JfirstC \\Jblack~\\cite{Guillemot11_2msps}, and \\JRansom~\\cite{Ransom11}. Some other MSPs, including PSR~B1821$-$24 and PSR~J0737$-$3039A, may also have phase-aligned radio and $\\gamma$-ray LCs, on the basis that their radio and X-ray profiles are phase-aligned~\\cite{Backer97,Chatterjee07}. ",
        "conclusions": "Future studies should include development of full radiation models which will be able to reproduce both the multiwavelength LC shapes, polarization properties, as well as the energy-dependent behaviour of the spectra of the $\\gamma$-ray pulsars. Ways in which to induce abundant pair creation in low-B MSP magnetospheres, such as using offset-dipole~\\cite{HM11} or multipole~\\cite{Zhang03} B-fields, will need to be investigated.  \\ack CV is supported by the South African National Research Foundation. AKH acknowledges support from the NASA Astrophysics Theory Program. CV, TJJ, and AKH acknowledge support from the \\textit{Fermi} Guest Investigator Program as well as fruitful discussions with Dick Manchester and Matthew Kerr."
    },
    "1112/1112.1356_arXiv.txt": {
        "abstract": "A completely new mechanism to generate the observed amount of large-scale cosmological magnetic fields is introduced in the context of three-form inflation.  The amplification of the fields occurs via fourth order dynamics of the vector perturbations and avoids the backreaction problem that plagues most previously introduced mechanisms. ",
        "introduction": "\\label{I&C}  Large scale magnetic fields exist in most observed structures in our Universe such as galaxies, clusters, filaments, and beyond, stretching to regions in the sky where little to next to no gas is present.  Observational evidence is growing by the day, unveiling an unsettling uniform and ubiquitous presence; yet, their nature and origin remains draped by a hazy veil of mystery, see the reviews~\\cite{Han:2002ns,Subramanian:2008tt,Kandus:2010nw,Widrow:2011hs}, and~\\cite{Govoni:2004as,Beck:2008ty,Neronov:1900zz,Tavecchio:2010mk,Dolag:2010ni}.  The property which is perhaps most intricate to frame within a consistent theoretical model for cosmic magnetic fields is their pervasively wide coherence length, which can stretch well beyond the Mpc mark; cosmological inflation might be the only viable candidate for its automatic transformation of short wavelength fluctuations into beyond-the-horizon correlations~\\cite{Martin:2007ue}.   Many simple options for inflationary magnetogenesis suffer from a backreaction issue, see~\\cite{Demozzi:2009fu,Kanno:2009ei,Urban:2011bu,Byrnes:2011aa}.  In short, as inflation proceeds at amplifying electromagnetic waves, the overall energy density of the latter catches up with that of the former, leading to instability in the de Sitter background.  When this occurs inflation stops, and if this happens too early then we are not able to generate an almost scale invariant spectrum of curvature perturbations, as well as expand the causal region of the Universe beyond the observed one today -- a possible viable model is perhaps the simplest~\\cite{Martin:2007ue}, but it appears to be successful only in the strong coupling (hence non perturbatively trustable) regime~\\cite{Demozzi:2009fu,Kanno:2009ei} (but see the proposal of~\\cite{Caldwell:2011ra} and criticism in~\\cite{Barnaby:2012tk}).  The reasons behind this unhappy ending are twofold.  In the first place, what is generally found is that the electric field, while strongly suppressed in a highly conductive plasma (such as the primordial soup during and after reheating), plays the part of the giant in vacuum, see~\\cite{Urban:2011bu}.  Secondly, the electric field tends to be redder than the magnetic field~\\cite{Urban:2011bu}, and since inflation generally blindly processes an extremely wide range of modes, the magnetic spectrum needs to be close to being red or all the allowed energy density will be snatched by small scales; but if this is so, the electric field will greatly dominate in the IR, leaving only crumbs of power to the magnetic field.  In this work we show how these general issues can be circumvented in a particular model of inflation, three-form driven inflation~\\cite{Germani:2009iq,Koivisto:2009sd}, thanks to the peculiar re-shaping of the magnetic spectrum the coupling between the three-form and electromagnetism produces.  Three-form inflation is a generally viable alternative to scalar field models.  Details of the rich dynamics that three-form inflation allows can be found in various cases in the references~\\cite{Germani:2009iq,Koivisto:2009sd,Koivisto:2009ew,Boehmer:2011tp,Ngampitipan:2011se,Koivisto:2009fb}; for instance, the simplest models cannot produce phantom inflation without instabilities or blue-tilted spectrum, but there are viable models generating slight red tilt with or without slow roll.  Once coupled to EM, three-form inflation provides us with the two key ingredients for a promising magnetisation of the large scale Universe.  First, it entails the possibility of exponential growth in the EM sector, driven by an instability in the three-form rotational sector.  Moreover, and more importantly, provides an effective UV cutoff at which the instability dies off; it then automatically splits it from the scale of inflation, allowing for a careful selection of the modes that are amplified, thereby upsetting the democraticity of inflation.  This can be understood in the dual description of three-form inflation as a scalar field inflation, since the duality maps kinetic into potential and vice versa.  However, the duality can break down in the presence of nontrivial self-interactions~\\cite{Koivisto:2009fb} of the three-form or nonminimal couplings to geometry~\\cite{Germani:2009iq} or to matter~\\cite{Ngampitipan:2011se}.  Thus three-form inflation is equivalently described by a scalar field model only in the simplest minimally coupled cases.  The novel aspect here is the coupling to the electromagnetic field: this both promotes the model to a physically distinct alternative to scalar-field driven constructions by breaking the duality (as it is apparent when one sees that the three-form has propagating vector degrees of freeedom), and resolves a generic shortcoming of theirs by enabling sufficient magnetic field production without the backreaction issue.   Our result are shown in Figs.~\\ref{fig2} and~\\ref{fig2s}: despite the electric field still being the one to watch out for, magnetic fields which today have strengths in the $10^{-15}$ Gauss range can be easily produced without excessive fine-tuning of the parameters, and without jeopardising the development and completion of a long epoch of inflation. ",
        "conclusions": "\\label{I&C}  Large scale magnetic fields exist in most observed structures in our Universe such as galaxies, clusters, filaments, and beyond, stretching to regions in the sky where little to next to no gas is present.  Observational evidence is growing by the day, unveiling an unsettling uniform and ubiquitous presence; yet, their nature and origin remains draped by a hazy veil of mystery, see the reviews~\\cite{Han:2002ns,Subramanian:2008tt,Kandus:2010nw,Widrow:2011hs}, and~\\cite{Govoni:2004as,Beck:2008ty,Neronov:1900zz,Tavecchio:2010mk,Dolag:2010ni}.  The property which is perhaps most intricate to frame within a consistent theoretical model for cosmic magnetic fields is their pervasively wide coherence length, which can stretch well beyond the Mpc mark; cosmological inflation might be the only viable candidate for its automatic transformation of short wavelength fluctuations into beyond-the-horizon correlations~\\cite{Martin:2007ue}.   Many simple options for inflationary magnetogenesis suffer from a backreaction issue, see~\\cite{Demozzi:2009fu,Kanno:2009ei,Urban:2011bu,Byrnes:2011aa}.  In short, as inflation proceeds at amplifying electromagnetic waves, the overall energy density of the latter catches up with that of the former, leading to instability in the de Sitter background.  When this occurs inflation stops, and if this happens too early then we are not able to generate an almost scale invariant spectrum of curvature perturbations, as well as expand the causal region of the Universe beyond the observed one today -- a possible viable model is perhaps the simplest~\\cite{Martin:2007ue}, but it appears to be successful only in the strong coupling (hence non perturbatively trustable) regime~\\cite{Demozzi:2009fu,Kanno:2009ei} (but see the proposal of~\\cite{Caldwell:2011ra} and criticism in~\\cite{Barnaby:2012tk}).  The reasons behind this unhappy ending are twofold.  In the first place, what is generally found is that the electric field, while strongly suppressed in a highly conductive plasma (such as the primordial soup during and after reheating), plays the part of the giant in vacuum, see~\\cite{Urban:2011bu}.  Secondly, the electric field tends to be redder than the magnetic field~\\cite{Urban:2011bu}, and since inflation generally blindly processes an extremely wide range of modes, the magnetic spectrum needs to be close to being red or all the allowed energy density will be snatched by small scales; but if this is so, the electric field will greatly dominate in the IR, leaving only crumbs of power to the magnetic field.  In this work we show how these general issues can be circumvented in a particular model of inflation, three-form driven inflation~\\cite{Germani:2009iq,Koivisto:2009sd}, thanks to the peculiar re-shaping of the magnetic spectrum the coupling between the three-form and electromagnetism produces.  Three-form inflation is a generally viable alternative to scalar field models.  Details of the rich dynamics that three-form inflation allows can be found in various cases in the references~\\cite{Germani:2009iq,Koivisto:2009sd,Koivisto:2009ew,Boehmer:2011tp,Ngampitipan:2011se,Koivisto:2009fb}; for instance, the simplest models cannot produce phantom inflation without instabilities or blue-tilted spectrum, but there are viable models generating slight red tilt with or without slow roll.  Once coupled to EM, three-form inflation provides us with the two key ingredients for a promising magnetisation of the large scale Universe.  First, it entails the possibility of exponential growth in the EM sector, driven by an instability in the three-form rotational sector.  Moreover, and more importantly, provides an effective UV cutoff at which the instability dies off; it then automatically splits it from the scale of inflation, allowing for a careful selection of the modes that are amplified, thereby upsetting the democraticity of inflation.  This can be understood in the dual description of three-form inflation as a scalar field inflation, since the duality maps kinetic into potential and vice versa.  However, the duality can break down in the presence of nontrivial self-interactions~\\cite{Koivisto:2009fb} of the three-form or nonminimal couplings to geometry~\\cite{Germani:2009iq} or to matter~\\cite{Ngampitipan:2011se}.  Thus three-form inflation is equivalently described by a scalar field model only in the simplest minimally coupled cases.  The novel aspect here is the coupling to the electromagnetic field: this both promotes the model to a physically distinct alternative to scalar-field driven constructions by breaking the duality (as it is apparent when one sees that the three-form has propagating vector degrees of freeedom), and resolves a generic shortcoming of theirs by enabling sufficient magnetic field production without the backreaction issue.   Our result are shown in Figs.~\\ref{fig2} and~\\ref{fig2s}: despite the electric field still being the one to watch out for, magnetic fields which today have strengths in the $10^{-15}$ Gauss range can be easily produced without excessive fine-tuning of the parameters, and without jeopardising the development and completion of a long epoch of inflation."
    },
    "1112/1112.4435_arXiv.txt": {
        "abstract": "The spectral and spatial characteristics of two special categories of O stars found in the VFTS dataset are presented.  One of them comprises very rapid rotators, including several more extreme than any previously known.  These objects are distributed around the peripheries of the main 30~Doradus clusters, suggesting a runaway nature for which their radial velocities already provide preliminary supporting evidence.  The other category consists of a large number of Vz stars, previously hypothesized on spectroscopic grounds to be on or very near the ZAMS.  Their distribution is the inverse of that of the rapid rotators: the Vz are strongly concentrated to the ionizing clusters, plus a newly recognized band of recent and current star formation to the north, which provides strong circumstantial evidence for their extreme youth. ",
        "introduction": "The VLT-FLAMES Tarantula Survey (VFTS; Evans et~al.\\ 2011) has produced an unprecedented spectroscopic dataset for the stellar content of the 30~Doradus Nebula in the Large Magellanic Cloud, the largest in the Local Group.  About 800 OB stars have been observed with a resolving power of $\\sim$10$^4$ at the Very Large Telescope of the European Southern Observatory (program 182.D-0222).  It is reasonable to expect that such an advance over previous material will entail both observational and theoretical progress on the formation and evolution of massive stars.  Indeed, some early, unexpected empirical developments, resulting from the spectral classification of 167 O stars with no detected radial-velocity variations, are briefly described here; they will be amplified shortly by Walborn et~al.\\ (2012, in prep.), while the optimized classification atlas generated for this work will be presented by Sana et~al.\\ (2012, in prep.). ",
        "conclusions": ""
    },
    "1112/1112.6430_arXiv.txt": {
        "abstract": "The measurement of the mass $1.97\\pm 0.04~M_\\odot$ for PSR J1614-2230 provides a new constraint on the equation of state and composition of matter at high densities. In this contribution we investigate the possibility that the dense cores of neutron stars could contain strange quarks either in a confined state (hyperonic matter) or in a deconfined one (strange quark matter) while fulfilling a set of constraints including the new maximum mass constraint. We account for the possible appearance of hyperons within an extended version of the density-dependent relativistic mean-field model, including the $\\phi$ meson interaction channel. Deconfined quark matter is described by the color superconducting three-flavor NJL model. ",
        "introduction": "The measurement of the mass $1.97\\pm 0.04~M_\\odot$ for PSR J1614-2230 by Demorest et al. \\cite{Demorest:2010bx} has renewed the interest in the question of the internal structure of a neutron star (NS). This concerns in particular the question of its composition and the possibility of exotic forms of matter in the cores of NSs which  has been tested extensively against this new measurement \\cite{Ozel:2010bz,Lattimer:2010uk,Weissenborn:2011qu,Bonanno:2011ch}, thus reviving an old controversy \\cite{Ozel:2006bv,Alford:2006vz} in the course of which it had been shown that hybrid stars with quark matter cores could not be ruled out by the observation of a $2~M_\\odot$ compact star \\cite{Klahn:2006iw}.  This was demonstrated not only within extended bag models \\cite{Ozel:2010bz,Weissenborn:2011qu} but also within the field theoretical Nambu--Jona-Lasinio (NJL) model description of quark matter \\cite{Klahn:2006iw,Klahn:2011fb}. In the latter approach it was found that the hybrid star becomes unstable as soon as strange quark matter (in its color superconducting CFL phase) appears in the very core of the hybrid star. If one combines the NJL-model description of color superconducting quark matter with an additional (possibly density dependent) bag pressure, then it is possible to achieve stability of hybrid stars even with strange quark matter interior \\cite{Pagliara:2007ph}, eventually fulfilling the $2~M_\\odot$ constraint from PSR J1614-2230 \\cite{Ippolito:2007hn,Bonanno:2011ch,Blaschke:2010vj}. There appears a new question which we would discuss here: Given a high-mass hybrid star contains quark matter, % could this be strange quark matter?  A separate issue is the possible appearance of hyperons at supersaturation densities, before a phase transition to quark matter occurs. It was found that the appearance of hyperons softens the equation of state (EoS) to the extent that sufficiently massive hybrid stars can not be described. In the Brueckner-Bethe-Goldstone (BBG) theory even the typical NS mass of $\\sim 1.4~M_\\odot$, well measured for binary radiopulsars, could not be reached \\cite{Baldo:1999rq}. An early deconfinement transition to bag model quark matter with a density-dependent bag constant was suggested as a possible solution \\cite{Baldo:2003vx}. Within the density-dependent relativistic mean field (RMF) theory of neutron star matter \\cite{Hofmann:2000mc} higher maximum masses than in BBG theory could be reached and the inclusion of repulsive $\\phi$ meson interations provided a stiff enough EoS to support a  $2~M_{\\odot}$ star with hyperonic interior \\cite{Weissenborn:2011ut,Bonanno:2011ch}.  In the present contribution we discuss hybrid stars with both, hyperonic and quark matter interior in the light of the new maximum mass constraint. ",
        "conclusions": "We conclude that one can obtain hyperon stars fulfilling the new $2~M_{\\odot}$ mass constraint. This requires the repulsive $\\phi$ meson interaction channel. Straightforward attempts to construct three-phase hybrid stars fail to meet the new mass constraint. With a phenomenological $\\mu$-dependent contribution to the quark pressure motivated by backreaction from the gluon sector a deconfinement phase transition could be obtained with hybrid star sequences fulfilling the new maximum mass constraint. It is plain that traditional phase transition constructions introduce inconsistencies like a ``second crossing'' of pressure curves and should be replaced by a microscopic description of hadron dissociation at high densities, to be developed."
    },
    "1112/1112.2213_arXiv.txt": {
        "abstract": "Exploiting a fundamental characteristic of galaxy assembly in the $\\Lambda$CDM paradigm, the observed spatial biasing and kinematics of metal-poor globular star clusters are used to constrain the {\\it local} reionization epoch around individual galaxies.  Selecting three galaxies located in different environments, the first attempt at constraining the environmental propagation of reionization in the local Universe is carried out.  The joint constraint from the three galaxies (\\allz) agrees remarkably well with the latest WMAP constraint on $z_{reion}$ for a simple instantaneous reionization model.  More importantly, the range of $z_{reion}$ values found here are consistent with the global range of $z_{reion}$ estimates from other observations.  We furthermore find a $1.7\\sigma$ indication that reionization completed in low-density environments before the intergalactic medium in high-density environments was reionized. This is consistent with certain theoretical models that predict that reionization was globally prolonged in duration, with neutral hydrogen pockets surviving in high-density environments, even after the surrounding regions were reionized.  More generally, this work provides a useful constraint on the formation history of galaxy stellar halos. ",
        "introduction": "Reionization marks the phase change of the Universe's intergalactic hydrogen gas from a neutral to an ionized state.  Though the general reionization picture follows from the fundamentals of Big Bang cosmology \\citep{barkana_beginning:_2001}, a detailed timeline of how it propagated throughout the Universe is difficult to constrain because most of the direct observational signatures of this event are $\\ga12$ billion light years distant from Earth.  Despite this challenge, observations show that the intergalactic medium was reionized no later than a redshift of $z\\sim5$, not before $z\\sim12$ and completed on a timescale of at least $\\Delta z > 0.06$ \\citep{fan_constraining_2006,ouchi_statistics_2010,robertson_early_2010,bolton_first_2010,bowman_lower_2010,larson_seven-year_2011}.  Though the dominant ionizing source contributing to reionization is still an open question, there is general agreement that the ionizing sources first appeared in high-density environments \\citep{dijkstra_limit_2004,fan_constraining_2006,lidz_quasar_2007,wise_how_2008,power_primordial_2009,baek_reionization_2010,bunker_contribution_2010,robertson_early_2010,srbinovsky_fraction_2010,willott_canada-france_2010,yan_galaxy_2010,bouwens_lower-luminosity_2011,dopita_re-ionizing_2011,lorenzoni_star-forming_2011,mitra_reionization_2011,mitra_joint_2012}.  A key piece of information that is not yet constrained is the environmental propagation of reionization:  did reionization complete in dense environments first or were low-density voids the first locations to ionize?  While high-density environments of the Universe probably had a ``head-start'' and thus contained more ionizing sources, the same locations were also more likely to be shielded by dense clouds of dust and neutral hydrogen (HI) gas.  Furthermore, if the sources were predominately small galaxies, the UV background would have been more uniformly distributed than if the ionizing sources were typically active galaxy nuclei of rare, massive proto-galaxies.  Also, unlike ultraviolet radiation, X-rays produced by gas accretion onto black holes can escape local gas and dust absorption and ``pre-heat'' gas in lower-density regions \\citep[e.g. ][]{baek_reionization_2010}.  Indeed, large-scale radiative transfer models show that the ionizing front can travel vast distances, from massive galaxies in dense regions to low-density locales \\citep[e.g. ][]{weinmann_dependence_2007,finlator_late_2009,iliev_reionization_2011}. Understanding how reionization propagated through various environments will help constrain the properties of the ionizing sources and the state of the intergalactic medium during the reionization epoch.\\\\   One possible surviving relic of the reionization epoch may be the bimodal distribution of globular star clusters within large galaxies \\citep{forbes_origin_1997,van_den_bergh_how_2001,beasley_formation_2002,santos_high-redshift_2003,rhode_metal-poor_2005,rhode_global_2007,rhode_wiyn_2010,moore_globular_2006,bekki_galactic_2005,bekki_origin_2008,griffen_globular_2010}:  systems of globular clusters show two distinct metallicity subpopulations \\citep{zinn_globular_1985,gebhardt_globular_1999,neilsen_color_1999,larsen_properties_2001,kundu_new_2001,peng_acs_2006,strader_globular_2006,strader_globular_2007,kundu_bimodal_2007,spitler_extendingbaseline:_2008}.  The subpopulations include the ``metal-poor'' globular clusters (MPGCs) with typical metallicities a few percent of Solar metallicity and the ``metal-rich'' globular clusters (MRGCs) with metallicities $\\sim30\\%$ Solar.  Various theoretical works have invoked reionization as a way to produce the bimodal metallicity distribution in GC systems \\citep[e.g. ][]{beasley_formation_2002,santos_high-redshift_2003,bekki_galactic_2005,moore_globular_2006,bekki_origin_2008,griffen_globular_2010}.  More specifically, these studies explored models where MPGCs form within small ($M_{\\rm vir}\\sim10^{8}M_{\\odot}$) dark matter halos at $z\\ga6$ \\citep{bromm_formation_2002,boley_globular_2009}.  In this model, the MPGCs formed until cold molecular clouds within these halos was heated up during the reionization epoch.  This suppressed globular cluster formation for a period of time and allowed gases within the host galaxy to enrich in metals.  After some time had passed, a second generation of globular clusters, the metal-rich GCs, started to form during galaxy merger events \\citep[e.g. ][]{ashman_formation_1992,beasley_formation_2002,griffen_globular_2010} and/or the growth of galaxy disks \\citep{shapiro_star-forming_2010} at redshifts $z\\sim2-4$ \\citep{shapiro_star-forming_2010,spitler_first_2010}.  Although the global nature of the reionization epoch may provide a convenient way to explain the ubiquitous GC metallicity bimodality observed in nearby galaxies, a number of interesting alternative models have been proposed to explain GC metallicity distributions without reionization \\citep[e.g. ][]{ct_formation_1998,scannapieco_triggering_2004,pipino_formation_2007,hasegawa_formation_2009,shapiro_star-forming_2010,muratov_modeling_2010,gray_formation_2010}.   If reionization truncated globular cluster formation, the typical ages of metal-poor globular clusters should correlate with the {\\it local} reionization epoch.  This means metal-poor globular clusters can be used to measure $z_{reion}$ for individual galaxies.  Current MPGC age estimates overlap with the expected reionization epoch.  Galactic MPGCs are constrained to be older than $\\sim11$ Gyrs or $z\\sim2$ \\citep{krauss_age_2003}, with more stringent constraints on individual MPGCs suggestive of a pre-reionization epoch formation \\citep[e.g. ][though exceptions do exist: \\citealt{hansen_white_2007}]{hansen_hubble_2004}.  Constraints on the ages of individual extragalactic MPGCs are poorer \\citep[$\\ga 10$ Gyrs, see refs. in ][]{brodie_extragalactic_2006}, but are consistent with a pre-reionization formation.  Also, within the 1-sigma uncertainties the mean MPGC metallicities of various early-type galaxies are statistically consistent with preliminary measurements for galaxies at redshifts $z \\sim 6-8$ \\citep{finkelstein_stellar_2010,bouwens_very_2010,labbe_star_2010}.  This is shown in Fig.~\\ref{figmpgc} and is consistent with a scenario where MPGCs were in place at redshifts greater than $z \\sim 7$, well into the expected reionization epoch.   Unfortunately, the limited observational information about extragalactic globular clusters and large modeling uncertainties mean that direct methods to age-date extragalactic globular clusters do not provide a useful constraint on the local reionization epoch around individual galaxies.  A novel way to determine when the metal-poor GC formation epoch finished was recently proposed by \\citet{diemand_distribution_2005} and \\citet{moore_globular_2006}. The technique is based upon a fundamental property of a $\\Lambda$CDM universe, where the most massive or rarest dark matter halos (and anything that formed within them) will tend to be centrally concentrated within their final host halos and show unique kinematic signatures at the present epoch.  Thus by constraining the spatial and kinematic properties of metal-poor globular clusters today, information about their progenitor halos can be recovered.  Since the rarity of a halo for a given virial mass depends on redshift, this information can be used to age-date metal-poor GCs with enough accuracy to constrain the reionization epoch around the galaxy hosting the GCs \\citep[][see also \\citealt{bekki_galactic_2005}]{diemand_distribution_2005,moore_globular_2006}.  This technique was demonstrated in \\citet{moore_globular_2006} using the Milky Way's globular cluster system.  They found evidence for a reionization redshift of $z_{reion}=10\\pm2$ \\citep[updated for a WMAP7 cosmology,][]{komatsu_seven-year_2011}.  In the present work, this technique is applied to two additional galaxies located in denser environments than the Milky Way to help constrain the temporal propagation of reionization through the local Universe.\\\\  Section~\\ref{theoryintro} provides a detailed background, describing in detail both the theoretical framework of the ``Diemand-Moore'' methodology and relevant globular cluster observations. In Section~\\ref{observations}, the observational data are presented. The main mass modeling and $z_{reion}$ measurements are presented in Section~\\ref{method}.  It is followed by a discussion in Section~\\ref{discussion}. ",
        "conclusions": "The spatial and kinematic properties of metal-poor globular clusters (MPGCs) provide an exciting hint that reionization was inhomogeneous.  In the local Universe, reionization completed first in low-density environments around $z_{reion}\\sim11-12$ (for the Milky Way and large central group elliptical NGC~1407) and finished last in the high-density environment around the Virgo cluster at $z_{reion}\\sim8$.  While uncertainties on the galaxy halo mass models limit the strength of this conclusion, another property, the relative number of metal-poor globular clusters, supports this interpretation.  The apparent environmental dependence of $z_{reion}$ also aligns with theoretical work that finds that the ionization front propagated from high- to low-density environments, with the high densities finishing reionization last.  Furthermore, the apparent agreement between the range of $z_{reion}$ constraints from MPGC observations and those from the literature may provide additional support to models where reionization was globally prolonged in duration.  More observational work, especially to improve the individual galaxy halo mass models, is required to confirm the results found here.  It will also be important to model MPGC formation in large-scale, radiative transfer simulations with realistic models of star cluster dynamical evolution.  This will help refine the technique utilized here and clarify the relationship between reionization and the truncation of MPGC formation.  For example, given that high star formation rates and gas densities were needed to produce such dense star clusters, it is possible that MPGCs trace {\\it the end} of reionization, where ionizing radiation was able to finally penetrate the densest gas clouds.  Expanding this initial study to other galaxies will yield a reionization map of the local Universe, further aiding our understanding of galaxy formation and evolution over cosmic time.  If reionization is not responsible for the truncation of MPGC formation, the results presented here provide unique constraints on the formation histories of MPGCs.   Furthermore, since MPGCs and stellar halos of the Milky Way share spatial and chemical properties \\citep[][ see also M31, \\citealt{huxor_exploring_2011}]{helmi_stellar_2008,martell_building_2011} they might share a common origin.  If this is the case for other galaxies, then MPGC observations can be used to constrain the star formation history of their stellar halos.  Finally, the power of the Diemand-Moore technique is that it does not care about the detailed star formation and enrichment physics -- it only requires an observable tracer object that accreted onto the galaxy at early times \\citep{diemand_distribution_2005}.  \\vspace{-0.7cm}"
    },
    "1112/1112.2355_arXiv.txt": {
        "abstract": "If dark matter self-annihilates then it may produce an observable signal when its density is high. The details depend on the intrinsic properties of dark matter and how it clusters in space. For example, the density profile of some dark matter candidates may rise steeply enough toward the Galactic Center that self-annihilation may produce detectable $\\gamma$-ray emission.  Here, we discuss the possibility that an annihilation signal may arise near a compact object (e.g., neutron star or black hole) even when the density of dark matter in the neighborhood of the object is uniform. Gravitational focusing produces a local enhancement of density, with a profile that falls off approximately as the inverse square-root of distance from the compact star. While geometric dilution may overwhelm the annihilation signal from this local enhancement, magnetic fields tied to the compact object can increase the signal's contrast relative to the background. ",
        "introduction": "Dark matter accounts for the majority of the mass in the Universe, yet its identity remains elusive. Candidates include weakly interacting massive particles (WIMPs) like the neutralino ($\\chi$), the supersymmetric partner of the neutrino~\\citep{PagPri82}, although their properties are only loosely constrained by theory and experiment.  In some cases, plausible values of the mass and cross section suggest that self-annihilation signatures may be detectable in regions where the density of dark matter is high~\\citep{BerGurZyb92,BerGon96,BerHooSil04}.  For example, \\citet{BerUllBuc98} calculate the gamma ray flux from neutralino self-annihilation in the \\GC\\ \\citep[see also][]{ZahHoo06}, while \\citet{Tyl02} and \\citet{BerHoo06} provide estimates of the annihilation signal from the nearby Draco dwarf galaxy. Intermediate-mass black holes may yield a WIMP annihilation signal \\citep{BerZenSil05}, as may remnant dark matter minihalos distributed throughout the Galaxy \\citep[e.g.][]{BerDokEro03,SanDieFre11}.  The observability of an annihilation signal critically depends on the number density of particles, since the local event rate is proportional to the density squared. Estimates of the strength of this signal typically derive from the assumption that dark matter density profiles follow a power law, $\\rho \\sim r^{-\\gamma}$ in many astrophysical contexts. The power-law index $\\gamma$ is between 1 and 2 in the central regions of galaxies according to cosmological simulations \\citep{NavFreWhi96, Mooetal99,Powetal03}, corresponding to a density ``cusp.''  If the simulations realistically describe the distribution of dark matter, then self-annihilation of WIMPS may indeed be observable in the centers of galaxies.  The presence of massive black holes in the centers of galaxies may further enhance the steep rise of a dark matter density profile. \\citet{GonSil99} model the adiabatic growth of a central black hole to show that a density ``spike,'' with a profile steeper than $r^{-2}$, forms around the black hole.  Dynamical processes, such as scattering by stars and capture the black hole, may erode a spike over time, although it may remain largely intact \\citep{BerMer05}. More problematic is the inspiral of smaller massive black holes captured from accreted galaxies that may disrupt a density spike \\citep{Meretal02}. Binary black hole coalescence can destroy density structures in the vicinity of the central black hole. However, even if merger event disrupts a density spike, relaxation processes may regenerate a dark matter ``crest,'' with a density profile that falls off as $r^{-1.5}$ \\citep{MerHarBer07}.  The time scale for the growth of a dark matter crest can be long, $\\sim$10~Gyr in the case of the Milky Way, so that its structure could reflect the merger history of the Galaxy.  From an observational perspective, strong density cusps or spikes are not obviously common. For instance, \\citet{Kraetal98} find that dwarf and low surface brightness galaxies have shallow core profiles, with $\\gamma < 0.5$, consistent with results from \\citet{vanSwa01} in a study of dwarf galaxies.  \\citet{Miletal02} identify an anticorrelation between profile steepness and mass of the central black hole, lending support for a scenario in which the black hole grows through a sequence of merger events which tend to reduce or destroy any density cusp.  Even if no density cusp or spike exists in a galactic nucleus, the presence of a massive black hole, such as \\sgr\\ in the center of the Milky Way \\citep{MelFal01}, can facilitate annihilation radiation. The reason is that gravitational focusing inevitably leads to a dark matter density enhancement near the black hole.  This phenomenon will occur for any star in a field of dark matter~\\citep{DanCam57, Gri88, SikWic02,AleGon06}.  The purpose of this paper is to review the gravitational focusing effect and discuss its observational consequences in terms of dark matter annihilation radiation. ",
        "conclusions": "Gravitational focusing will generally enhance the annihilation signal around any massive compact object.  However, because of geometrical dilution, this signal may not be distinguishable from the background annihilation radiation. Fortunately, local conditions around a compact object may cause the annihilation signal to appear in wavebands which are different from the background. For example, magnetic fields tied to compact objects can generate a distinctive synchrotron radiation. Of the sources considered here---neutron stars, stellar-mass black holes, and massive black holes---only the latter seems capable of producing a detectable signal by gravitational focusing alone.  Neutron stars, specifically magnetars, may be the weakest sources, but they still promise interesting physics involving WIMP decay products.  Theoretical models of WIMP annihilation in the Galactic Center, bolstered by recent observations with the Fermi Gamma-ray Space Telescope \\citep[e.g.,][]{Abd-etal10}, have the dark matter concentrated about \\sgr\\ as a result of uncertain galactic evolution and black hole growth dynamics. The main point of this work is that gravitational focusing will generate {\\em some} signal, no matter what the formation history. While distinctive gamma ray signal could be lost in a uniform bath of WIMPs, synchrotron radiation may retain some contrast with the background. The trade-off with synchrotron signal is that the accretion flows that presumably generate the localized magnetic field may contaminate the signal. Our best hope may be to find isolated intermediate mass black holes \\citep{BerZenSil05}. Alternatively we might seek an underluminous active galactic nucleus. Fortunately, our own Galaxy has one."
    },
    "1112/1112.0172_arXiv.txt": {
        "abstract": "Here we present an overview of some of the most significant observational and theoretical studies of the broad Fe K$\\alpha$ spectral line, which is believed to originate from the innermost regions of relativistic accretion disks around central supermassive black holes of galaxies. The most important results of our investigations in this field are also listed. All these investigations indicate that the broad Fe K$\\alpha$ line is a powerful tool for studying the properties of the supermassive black holes (such as their masses and spins), space-time geometry (metric) in their vicinity, their accretion physics, probing the effects of their strong gravitational fields, and for testing the certain predictions of General Relativity. ",
        "introduction": "According to General Relativity \\citep{ein16}, a black hole \\citep[named by][]{whe68} is a region of space-time around a collapsed mass with a gravitational field that has became so strong that nothing (including electromagnetic radiation) can escape from its attraction, after crossing its event horizon \\citep{wald84}. The problem of such catastrophic gravitational collapse was first adressed by \\citet{chandra31} when he discovered the upper mass limit for ideal white dwarfs, composed of a degenerate electron-gas. Subsequent studies by \\citet{lan32,opp39} and \\citet{pen65} established the modern theory of gravitational collapse. The theory of black holes, as well as their fundamental properties are presented in numerous monographs and papers \\citep[see e.g.][]{car73,chandra83}, and therefore in this review we will focus only on those studies which are necessary for the discussion of the relation between black holes and the Fe K$\\alpha$ spectral line.  All black holes in nature are commonly classified according to their masses as: supermassive black holes (with masses $M_{BH}\\sim 10^5-10^{10}\\ M_\\odot$), intermediate-mass black holes ($M_{BH}\\sim 10^2-10^5\\ M_\\odot$), stellar-mass black holes ($M_{BH} < 10^2\\ M_\\odot$), mini and micro black holes ($M_{BH} \\ll M_\\odot$). A crucial event for the acceptance of black holes was the discovery of pulsars by \\citet{hb68}, because it was the clear evidence of the existence of neutron stars, and therefore, confirmation of Chandrasekhar limit. The first detection of a solar mass black hole came in 1972, when the mass of the rapidly variable X-ray source Cygnus X-1 was proven to be above the maximum allowed for a neutron star \\citep[see e.g.][and references therein]{ff05}.  Nowadays, it is widely accepted that supermassive black holes are located in the centers of most of galaxies, and thus have a fundamental influence on galactic formation and evolution. Some of the first indirect arguments that they exist in galactic nuclei are given by \\citet{lyn69} and \\citet{lyn71}. Also, according to the unification model of active galactic nuclei -- AGN \\citep{ant93,pet97,krol99}, they are most likely powered by the accretion of gas onto their central supermassive black holes with mass ranging from $10^5$ to $10^9\\ M_\\odot$ \\citep{Kaspi00,pet04}.   \\subsection{Space-time geometry in vicinity of supermassive black holes}  In general, black holes have three measurable parameters (not including the Hawking temperature): charge, mass (and hence gravitational field) and angular momentum (or spin). In the case of supermassive black holes, only the latter two are of sufficient importance because they are responsible for several effects which can be observationally detected \\citep[for more details see e.g.][]{jov08}.  A central supermassive black hole of a galaxy can be stationary or rotating. In the first case, space-time geometry in the black hole vicinity depends only on its mass $M$ and is described by the Schwarzschild metric \\citep{schw16}, while in the latter case, space-time geometry depends also on spin $a$ (i.e. angular momentum $J=a M$) of the black hole and is described by Kerr metric \\citep{kerr63}. Therefore, Schwarzschild metric describes the spherically symmetric gravitational field around a time-steady non-rotating black hole in vacuum, while the Kerr metric describes the gravitational field outside an uncharged rotating black hole and, contrary to the Schwarzschild metric, is no longer spherically symmetric. \\citet{bl67} discovered a generalized coordinate system in which Kerr metric is most commonly used today \\citep[see e.g.][]{car73}.  Several characteristic radii can be defined around black holes, and the most important are \\citep[see e.g.][]{jov09}: 1. Schwarzschild radius $R_S=2GM/c^2$ (representing the limiting radius below which a collapsed mass form a spherically symmetric non-rotating black hole), where $G$ is the gravitational constant and $c$ is the speed of light; 2. gravitational radius $R_g$ (being a half of $R_S$ and usually used as a unit for distance around a black hole); 3. radius of event horizon $R_h$ (representing space-time boundary below which events cannot affect an outside observer); and 4. radius of marginally stable orbit $R_{ms}$ (representing the minimum allowed radius of a stable circular equatorial orbit around a black hole).  The equations governing photon trajectories in the Kerr metric are derived by \\citet{car68}. \\citet{mtw73} showed that certain relativistic effects in vicinity of a black hole, such as light banding, gravitational and Doppler shifts as seen by distant observer could be computed by solving the equation of the photon orbit. \\citet{cun72,cun73} calculated radiation of a point source in a circular orbit in the equatorial plane around an extreme Kerr metric black hole (see Fig. \\ref{fig01}).  \\begin{figure}[ht!] \\centering \\includegraphics[width=\\columnwidth]{fig01.eps} \\caption{A circular orbit (solid line) of a point source in the equatorial plane around an extreme Kerr metric black hole. Image Credit: Cunningham \\& Bardeen, Copyright: ApJ, 183, 237 (1973).} \\label{fig01} \\end{figure}  \\citet{bard70} studied the effects of accretion from a disk of gas orbiting a black hole on its properties, and showed that the Kerr metric is more appropriate for describing supermassive black holes than the Schwarzschild metric. This conclusion is strengthened even more by \\citet{th74}, who showed that a black hole at the center of an accretion disk would be spun up to the maximum possible value of $a \\approx 0.998$ in a relatively short time. In such a case the disk would be extended down to about 1.23 $R_g$, while in the case of a non-rotating black hole it could extend only down to 6 $R_g$.   \\begin{figure*}[ht!] \\centering \\includegraphics[width=\\textwidth]{fig02.eps} \\caption{The distribution of the temperature as a function of the disk radius $R$, given for two different values of angular momentum $a$. Negative values of $R$ correspond to the approaching and positive values to the receding side of the disk. Image Credit: Popovi\\'{c} et al., Copyright: ApJ, 637, 620 (2006).} \\label{fig02} \\end{figure*}  \\begin{figure*}[ht!] \\centering \\includegraphics[width=0.9\\textwidth]{fig03.eps} \\caption{Magnetohydrodynamic simulations of an accretion disk as seen by a distant observer at the following inclination angles \\citep{Armitage03}: 5$^\\circ$ (upper left), 30$^\\circ$ (upper right), 55$^\\circ$ (lower left) and 80$^\\circ$ (lower right). The disk rotation is in counterclockwise direction. Image Credit: Armitage \\& Reynolds, Copyright: MNRAS, 341, 1041 (2003).} \\label{fig03} \\end{figure*}  \\subsection{Accretion onto the supermassive black holes}  Some observed quasars have luminosities of up to $L_{bol} \\sim 10^{47}$ erg/s \\citep{Schneider06}. The corresponding total energy emitted during the lifetime of such quasars can be estimated to $E \\geq 3\\times 10^{61}$ erg, assuming that their minimum age is about $10^7$ years and that their luminosities  do not change significantly over the lifetime \\citep{Schneider06}. Accretion is the only mechanism which can yield sufficient efficiency $\\epsilon$ (defined as the mass fraction of fuel that is converted into energy, according to $E=\\epsilon m c^2$) to explain such high energies. The maximum efficiency of accretion is about $\\epsilon \\sim 6\\%$ for a non-rotating black hole and $\\epsilon \\sim 29\\%$ for a maximally rotating one \\citep{Schneider06}.  On the other hand, if the cold matter which was initially at rest and without magnetic field, was subjected to free radial infall it would accrete to the central black hole without any energy release or observational effects \\citep{Shakura73}. However, in the case of AGN the accreting matter has a significant angular momentum which does not allow its free infall. At the marginally stable orbit of central black hole with mass of $10^8M_\\odot$, the specific angular momentum is $\\approx 1\\times 10^{24}\\ \\mathrm{cm}^2/\\mathrm{s}$, which is much less even in comparison to a typical galaxy where the speciffic angular momentum of orbiting material is $\\approx 6\\times 10^{28}\\ \\mathrm{cm}^2/\\mathrm{s}$ \\citep{krol99}. It means that approach of accreting material toward a black hole requires a loss of the greatest fraction of its initial angular momentum. Mechanisms which can contribute to such loss of angular momentum are viscosity, nonaxisymmetric gravitational forces, magnetic forces \\citep{krol99}, outflowing winds \\citep{beg83,elv00}, or the Magnetorotational Instability (MRI) which was discovered by \\citet{vel59} and \\citet{chan60}, but did not come to the attention of the astrophysical community until rediscovered by \\citet{balb91,balb98}.  Since the orbit of minimum energy for fixed angular momentum in any spherically symmetric potential is a circle, the infall of accreting material due to loss of its angular momentum will be in form of successively smaller and smaller concentric circles. Matter traveling along orbits inclined to each other will eventually collide in the plane of intersection, and as a result, the angular momenta of different gas steams will be mixed and equalized. Consequently, all accreting matter will orbit in a single plane and will have the same specific angular momentum at any given radius, meaning that accretion is most likely performed through an accretion disk \\citep{krol99}. The assumption of a disk geometry for the distribution of the X-ray emitters in the central parts of AGN is also supported by the spectral shape of the observed Fe K$\\alpha$ line \\citep{Nandra97,Nandra99,Nandra07}.  \\citet{Shakura73} developed what is now called ''standard model of accretion disk'', in which accretion occurs via an optically thick and geometrically thin disk and where the spectrum of thermal radiation emitted from the disk surface depends on its structure and temperature, and therefore on the distance to the central black hole. The distribution of the temperature along the radius of the disk in this model is presented in Fig. \\ref{fig02} \\citep[see][for more details]{pop06}. This model was originally developed to describe the accretion disks around stellar sized black holes in the binary systems, but with certain modifications it could be also applied on accretion disks around central supermassive black holes of AGN.  The total energy release of emitted radiation is mainly determined by the rate of matter inflow into the disk on its outer boundary, i.e. by its accretion rate $\\dot{M}$. If the accretion converts matter to radiation with fixed radiative efficiency $\\eta$ (in rest-mass units), then a characteristic scale for the accretion rate is the Eddington accretion rate for which the total release of energy in the disk $L=\\eta\\dot{M}c^2$ is equal to the Eddington luminosity, a critical luminosity beyond which the radiation force overpowers gravity for any given mass \\citep{krol99,Shakura73}. Observed AGN have luminosities from $\\sim 10^{42}$ to $\\sim 10^{48}$ erg/s, which means that their central black holes must have masses from $10^5$ to $10^9\\ M_\\odot$ \\citep{krol99}.  Using the Faint Object Spectrograph (FOS) on the Hubble Space Telescope (HST), \\citet{harm94} found strong evidence for a disk of ionized gas around a supermassive nuclear black hole in M87.  \\citet{Armitage03} performed magnetohydrodynamic simulations of an accretion disk as seen by a distant observer at four different inclinations (see Fig. \\ref{fig03}). These simulations showed the existence of especially bright regions in form of arcs within the turbulent flow, which trace out the photon trajectories close to the radius of marginally stable orbit. These arcs are much brighter on approaching side of the disk due to higher surface temperature, as shown in Fig. \\ref{fig03}, and also due to Doppler boosting and relativistic beaming which both enhance the flux observed in the direction of motion and diminish the flux in the opposite direction. The influence of these relativistic effects on the shape of the Fe K$\\alpha$ line emitted from the accretion disk will be explained in more details in \\S2.3. ",
        "conclusions": "The results of previously mentioned observational and theoretical studies indicate that the broad Fe K$\\alpha$ line which originates in vicinity of the supermassive black holes is a powerful tool for studying their properties, space-time geometry (metric) in their vicinity, their accretion physics, probing the effects of their strong gravitational fields, and for testing the certain predictions of General Relativity. However, some important scientific issues still need to be addressed, such as decoupling the narrow and broad Fe K$\\alpha$ line components and explaining their different variabilities \\citep[see e.g][]{sul98}. This will be one of the crucial steps in determining the real importance of relativistic effects on the observed Fe K$\\alpha$ line profiles.  \\bigskip\\noindent\\textit{Acknowledgements.} This work is part of the project 176003 ''Gravitation and the Large Scale Structure of the Universe`` supported by Ministry of Education and Science of the Republic of Serbia. It was presented as an invited talk at special workshop ''Spectral lines  and super-massive black holes'' held on June 10, 2011 as a part of activity in the frame of COST action 0905 ''Black holes in a violent universe'', during the ''8th Serbian Conference on Spectral Line Shapes in Astrophysics''."
    },
    "1112/1112.0491_arXiv.txt": {
        "abstract": "{The Small Magellanic Cloud (SMC) is ideally suited to investigating the recent star formation history from X-ray source population studies. It harbours a large number of Be/X-ray binaries (Be stars with an accreting neutron star as companion), and the supernova remnants can be easily resolved with imaging X-ray instruments.} {We search for new supernova remnants in the SMC and in particular for composite remnants with a central X-ray source.} {We study the morphology of newly found candidate supernova remnants using radio, optical and X-ray images and investigate their X-ray spectra.} {Here we report on the discovery of the new supernova remnant around the recently discovered Be/X-ray binary pulsar CXO\\,J012745.97$-$733256.5 = \\sxp in radio and X-ray images. The Be/X-ray binary system is found near the centre of the supernova remnant, which is located at the outer edge of the eastern wing of the SMC. The remnant is oxygen-rich, indicating that it developed from a type Ib event. From XMM-Newton observations we find that the neutron star with a spin period of 1062 s (the second longest known in the SMC) shows a very high average spin-down rate of 0.26 s per day over the observing period of 18 days.} {From the currently accepted models, our estimated age of around 10000$-$25000 years for the supernova remnant is not long enough to spin down the neutron star from a few 10 ms to its current value. Assuming an upper limit of 25000 years for the age of the neutron star and the extreme case that the neutron star was spun down by the accretion torque that we have measured during the XMM-Newton observations since its birth, a lower limit of 0.5~s for the birth spin period is inferred. For more realistic, smaller long-term average accretion torques our results suggest that the neutron star was born with a correspondingly longer spin period. This implies that neutron stars in Be/X-ray binaries with long spin periods can be much younger than currently anticipated.} ",
        "introduction": "One of the most cataclysmic events in the universe is the explosion of a massive star as a supernova, which can create a neutron star (NS). NSs are thought to be born rapidly spinning with rotation periods of a few 10 ms. Their rotation is first slowed down by magnetic dipole braking and then by the propeller effect when the NS is in a binary star system and mass loss from the companion star begins \\citep[for a review see e.g.][]{1991PhR...203....1B}. % When rotating slowly enough, accretion onto the NS sets in \\citep{1978ApJ...223L..83G} and the system can be detected as an X-ray binary. The spin evolution of the NS in a high-mass X-ray binary (HMXB) system depends on the initial magnetic field strength of the NS and on the mass accretion rate \\citep{1998MNRAS.299...73U}.  Be/X-ray binary (BeXRB) systems are a subgroup of HMXBs with an NS accreting matter from the circumstellar disk of a Be star \\citep{2011Ap&SS.332....1R}. The Small Magellanic Cloud (SMC) harbours an extraordinarily large number of BeXRBs \\citep{2005MNRAS.356..502C,2005MNRAS.362..879S,2008A&A...489..327H}. These are believed to have been created during a burst of star formation about 42 million years ago \\citep{2010ApJ...716L.140A}. For more than fifty BeXRBs in the SMC the NS spin period is known from the detection of coherent pulsations in their X-ray flux. The spin periods range from 2.16~s \\citep[XTE\\,J0119-731 = SXP\\,2.16; ][]{2003IAUC.8064....4C} to 1320 s \\citep[RX\\,J0103.6-7201 = SXP\\,1323; ][]{2005A&A...438..211H} with a bimodal structure indicated in their distribution \\citep{2011arXiv1111.2051K}.  From ROSAT HRI data, \\citet{1994AJ....107.1363H} proposed two BeXRBs within supernova remnants (SNRs) in the SMC. For one of them, located in the direction of the SNR IKT\\,21 \\citep{2004A&A...421.1031V}, pulsations of 345~s were discovered by \\citet[][SAX J0103.2-7209]{2000ApJ...531L.131I}. The other - still unconfirmed - candidate BeXRB in the direction of IKT\\,25 was not detected by XMM-Newton or Chandra. In both cases it is not clear whether the BeXRB (or candidate) is physically associated with the SNR because they are not located near the centre of the remnant and the high number of BeXRBs in the bar of the SMC make chance coincidences likely. The SNR, which was created by the supernova explosion, is expected to have faded beyond detectability before accretion onto the NS starts. Finding a BeXRB associated to a SNR would therefore suggest a much faster spin-down than predicted by the standard models, or else the NS star was born with slow rotation. Here, we present the discovery of an SNR in radio and X-rays around the BeXRB pulsar \\cxoj\\ = \\sxp. The discovery of the new SMC pulsar and its possible association with a shell nebula around it has recently been reported by \\defcitealias{Henault}{HB11}\\citet[][ hereafter \\citetalias{Henault}]{Henault}.  \\begin{figure*} \\begin{center} \\resizebox{0.92\\hsize}{!}{\\includegraphics[clip=]{ima2.ps}} \\caption{Images of the region around the BeXRB SXP\\,1062. {\\it Left:} Continuum-subtracted MCELS images. Red, green, and blue correspond to H$\\alpha$, \\OIII\\ and \\SII. The cross marks the centre of the white circle (R.A. = 01\\rahour27\\ramin44\\fs15, Dec. = $-$73\\degr33\\arcmin01\\farcs6, J2000), which indicates the estimated size of the SNR in all images. The Be star counterpart of \\sxp is visible close to the cross. {\\it Middle:} MOST 36\\,cm radio image, logarithmically scaled from 0.5 to 3 mJy. {\\it Right:} Combined, instrument background-subtracted EPIC-pn and MOS colour image obtained from the XMM-Newton observations. Red, green, and blue denote X-ray intensities in the 0.2$-$1.0, 1.0$-$2.0, and 2.0$-$4.5 keV bands. The images from the individual bands were adaptively smoothed with a Gaussian $\\sigma$ of 4\\arcsec\\ (at high intensities) to 20\\arcsec\\ (low intensities). Circles indicate the extraction regions for the X-ray spectra: Source region for SXP\\,1062 (A, 29\\arcsec\\ radius) and associated background (B, 25\\arcsec). For the SNR spectrum (C, 84.5\\arcsec), regions D (40\\arcsec), E (11\\arcsec), and F (14\\arcsec) were excluded and region G (60 \\arcsec) was used to estimate the background. Regions E and F were excluded because they contain possible point sources in the direction of the SNR as found in source detection runs on the EPIC images. } \\label{fig-ima} \\end{center} \\end{figure*} ",
        "conclusions": "Our multi-wavelength morphological studies of the field around the new BeXRB pulsar \\sxp confirm the previously unknown SNR proposed by \\citetalias{Henault}. Its SNR nature is further supported by its X-ray spectrum obtained from both the XMM-Newton EPIC data and its steep radio spectrum. The strong \\OIII\\ line emission suggests the SNR as oxygen-rich type \\citep{2005MNRAS.360...76A}, which usually comes from a type~Ib event, the explosion of a massive O, B, or WR star.  \\sxp is detected close to the centre of the SNR. At the outer edge of the eastern wing of the SMC, only very few BeXRBs are detected, and all known SNRs are located in the bar of the SMC \\citep{2007MNRAS.376.1793P,2008A&A...485...63F}. This makes a chance coincidence of \\sxp with the SNR very unlikely. However, the ages of SNRs and BeXRBs are expected to be rather different. While SNRs can usually be seen for at most a few \\oexpo{5} years in X-rays and radio \\citep[for the methods used to estimate SNR ages see e.g.][]{2005ChJAA...5..165X},  % BeXRBs are expected to become X-ray active much later (see below). Using the relation t$_{\\rm y}$\\,=\\,3.8\\expo{2}\\,R$_{\\rm pc}$(kT)$_{\\rm keV}^{-1/2}$ from \\citet{2005ChJAA...5..165X} of temperature derived from X-ray spectral modelling (0.23 keV) and the size of the SNR ($\\sim$2.5\\arcmin\\ in diameter, corresponding to $\\sim$40 pc at a distance of 60\\,kpc), we estimate the age of the SNR to $\\sim$16 ky. If one compares the ages estimated for SNRs in the SMC by \\citet{2004A&A...421.1031V} with their sizes as listed in \\citet{2010MNRAS.407.1301B}, one finds ages of at most 25 ky for remnants with diameters around 40-50 pc.  We estimate the distance between the BeXRB and the centre of the SNR to about 9\\arcsec, which corresponds to $\\sim$8\\expo{13} km at SMC distance. From evolutionary investigations, kick velocities around 100$-$200 \\kms\\ were invoked for neutron stars in HMXB systems \\citep{1995A&A...296..691P}. Using this as maximum space velocities for the binary systems and neglecting a velocity component in radial direction, it would take 13$-$26 ky for \\sxp to move from the SNR centre to its current position, compatible with the SNR age estimates from above, which are also consistent with those of \\citetalias{Henault}.  \\begin{figure} \\begin{center} \\resizebox{0.99\\hsize}{!}{\\includegraphics[clip=,angle=-90]{spin.ps}} \\caption{Evolution of the spin period of \\sxp as obtained from the combined EPIC-pn and EPIC-MOS2 data of the four XMM-Newton observations. Spin periods and 1-$\\sigma$ errors were derived with the Bayesian odds ratio method.} \\label{fig-spin} \\end{center} \\end{figure}  Statistical measurements of BeXRBs in the SMC result in lower space velocities: \\citet{2005MNRAS.358.1379C} inferred a lower limit of 30 \\kms\\ from measuring the separation of BeXRBs to nearby young stellar clusters. Similarly \\citet{2010ApJ...716L.140A} argue from the correlation between the number of BeXRBs and the local star formation rate 42 My ago at their current position for a maximum velocity of 15$-$20 \\kms. Both cases should actually be regarded as minimum velocities as the BeXRB might not have been born in the nearest star cluster. The work of \\citet{2011arXiv1111.2051K} might also suggest that BeXRB pulsars with long spin period received stronger kicks. In either case, using lower space velocities (and also components in radial direction) would correspondingly increase the SNR age estimate.  According to the model calculations of \\citet{1998MNRAS.299...73U}, it seems very unlikely that an NS can be spun down to a period around 1000 s within a few \\oexpo{4} years if it was born with a spin period around 10~ms. Of course, such model calculations assume simplified approximations of the conditions in a BeXRB (in particular concerning the highly variable accretion usually observed in such systems), but only for models with extreme conditions (high magnetic field strength for the NS and low accretion rates), spin periods in excess of a few 100 s are reached after more than 10$^6$ years. If the NS in \\sxp is indeed the compact remnant from the supernova explosion, which also created the SNR around it, then either the NS was born with a much longer spin period than 10 ms, or some other effect plays an important role to efficiently decelerate the rotation of the NS.  \\sxp was discovered as a transient BeXRB pulsar during XMM-Newton and Chandra observations in 2010 when its X-ray luminosity was at the level of 6\\ergs{35}. The BeXRB was not active during the ROSAT PSPC observations of the SMC \\citep{2000A&AS..142...41H}. Based on the longest PSPC observation from 1991-10-07 and assuming the spectral shape measured by EPIC, we derive an upper limit of 1.5\\ergcm{-13} (0.2$-$10 keV), assuming an upper limit of 10 cts and a vignetting-corrected exposure of 9.7 ks at the position of \\sxp. This is about a factor of 10 lower than the flux during the XMM-Newton observations. While BeXRB transients generally exhibit spin-up during intervals of increased accretion \\citep[e.g.][]{1997ApJS..113..367B}, \\sxp is peculiar as it showed a strong spin-down consistent with a constant rate of 0.26 s day$^{-1}$ over 18 days. \\citet{2010MNRAS.401..252C} investigated RXTE monitoring data covering $\\sim$10 years to study the spin period changes of BeXRB pulsars in the SMC over their active periods, which typically last 50$-$500 days. They find four out of 15 systems with measured short-term spin period changes, which showed spin-down during outburst. In Table~\\ref{tab-fdot}, we summarise the pulse period changes for some X-ray binary pulsars, which show times of spin-down. We list the frequency derivative \\fdot\\ = -\\Pdot/P$^2$, which is directly related to the accretion torque exerted on the NS. We include the three SMC BeXRB pulsars with well measured values of \\Pdot\\ and for comparison Cen~X-3 (a super-giant HMXB with Roche-lobe overflow) and GX~1+4, a low-mass system with M5 III donor star. \\sxp exhibits spin-down about five times larger than the other listed SMC BeXRBs (which have remarkably similar values). It should be noted that the values derived by \\citet{2010MNRAS.401..252C} are average values over longer periods of time and that intermittent intervals of spin-up \\citep[as usually seen from accreting pulsars,][]{1997ApJS..113..367B} reduce the longer-term average.  \\sxp shows a remarkably large accretion torque, similar in magnitude to that of Cen~X-3. In this HMXB intervals of steady spin-up and spin-down alternate, which last typically 10$-$100 days \\citep{1997ApJS..113..367B}, but at much higher X-ray luminosities of $\\sim$\\oergs{38} \\citep[e.g.][]{1992ApJ...396..147N}. Also GX~1+4 was still observed at ~\\oergs{36} during its extended lowstate \\citep{1988Natur.333..746M}, about a factor of 2 brighter than \\sxp. This demonstrates that the generally accepted model for accretion, where the torque is directly proportional to the mass accretion rate (hence luminosity) is probably too simple, as similar torques can be exerted, although the luminosity is more than a factor of 100 different. The work of \\citet{1997ApJS..113..367B} suggests that disk-accreting pulsars are subject to instantaneous torques with similar magnitude but opposite sign. Different average long-term spin-up or spin-down values of individual pulsars would then be the result of different time scales for reversals between spin-up and -down.  \\sxp most likely also shows intervals of alternating spin-up and spin-down, which would reduce the magnitude of the long-term spin change. Assuming the extreme case that the NS was spun down with -2.6\\expo{-12} Hz s$^{-1}$ (-8.4\\expo{-5} Hz y$^{-1}$) over its whole life of 25 ky (the maximum age) an upper limit of 2 Hz for the total spin-down is derived, which corresponds to a lower limit of 0.5~s for the spin period at birth. For a more realistic, smaller long-term average accretion torque a correspondingly longer birth spin period is expected. Therefore, if the NS in \\sxp is indeed the compact remnant from the supernova explosion that created the SNR, our results show that the NS was most likely born with a spin period much longer than a few tens of ms as generally adopted for an NS at birth.  \\begin{table} \\caption[]{X-ray binary pulsars with periods of spin-down.} \\begin{center} \\begin{tabular}{lccl} \\hline\\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{Object} & \\multicolumn{1}{c}{Pulse} & \\multicolumn{1}{c}{\\fdot} & \\multicolumn{1}{c}{Reference} \\\\ \\multicolumn{1}{c}{name} & \\multicolumn{1}{c}{period [s]} & \\multicolumn{1}{c}{[Hz s$^{-1}$]} & \\multicolumn{1}{c}{} \\\\  \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} Cen~X-3          &    4.82   &    -(1$-$5)\\expo{-12}     &  \\citet{1997ApJS..113..367B} \\\\ SXP\\,8.80        &    8.90   &    -5.1\\expo{-13}         &  \\citet{2010MNRAS.401..252C} \\\\ GX~1+4           &    110.2  &    -2.1\\expo{-13}         &  \\citet{1988Natur.333..746M} \\\\ SXP\\,144         &    144.5  &    -5.9\\expo{-13}         &  \\citet{2010MNRAS.401..252C} \\\\ SXP\\,1062        &    1062   &    -2.6\\expo{-12}         &  this work \\\\ SXP\\,1323        &    1325   &    -5.0\\expo{-13}         &  \\citet{2010MNRAS.401..252C} \\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\end{center} \\label{tab-fdot} \\end{table}   We confirm the existence of a new SNR around the BeXRB \\sxp in the SMC, independently suggested by \\citet{Henault}. The SNR is detected in optical, radio, and X-ray wavelengths and shows a shell-like structure with the BeXRB close to its projected centre. We estimate an age of around 10$-$25 ky for the remnant. Such a time scale is too short to spin down the neutron star in a BeXRB to $\\sim$1000\\,s if it was born with a canonical spin period of a few 10 ms. \\sxp is remarkable for showing a strong average spin-down rate \\fdot\\ of -2.6\\expo{-12} Hz s$^{-1}$ during the XMM-Newton observations distributed over 18 days, while it was detected during high X-ray activity at 6.3\\ergs{35}. Assuming the extreme case that such a high reverse accretion torque was exerted on the neutron star for its whole life, a lower limit of 0.5~s for the birth spin period is inferred. If the neutron star in \\sxp is indeed the compact remnant from the supernova explosion that created the SNR, we conclude that the neutron star in \\sxp was born with an even longer spin period for more realistic, lower long-term average accretion torques."
    },
    "1112/1112.0034_arXiv.txt": {
        "abstract": "The outskirts of galaxy clusters, near the virial radius, remain relatively unexplored territory and yet are vital to our understanding of cluster growth, structure, and mass.  In this presentation, we show the first results from a program to constrain the state of the outer intra-cluster medium (ICM) in a large sample of galaxy clusters, exploiting the strengths of three complementary X-ray observatories: \\suzaku\\ (low, stable background), \\xmmnewton\\ (high sensitivity), and \\chandra\\ (good spatial resolution).  By carefully combining observations from the cluster core to beyond \\rth, we are able to identify and reduce systematic uncertainties that would impede our spatial and spectral analysis using a single telescope.  Our sample comprises nine clusters at $z \\sim 0.1$--0.2 fully covered in azimuth to beyond \\rth, and our analysis indicates that the ICM is not in hydrostatic equilibrium in the cluster outskirts, where we see clear azimuthal variations in temperature and surface brightness. In one of the clusters, we are able to measure the diffuse X-ray emission well beyond \\rth, and we find that the entropy profile and the gas fraction are consistent with expectations from theory and numerical simulations. These results stands in contrast to recent studies which point to gas clumping in the outskirts; the extent to which differences of cluster environment or instrumental effects factor in this difference remains unclear.  From a broader perspective, this project will produce a sizeable fiducial data set for detailed comparison with high-resolution numerical simulations. ",
        "introduction": "Exploring the outskirts of galaxy clusters is vital to our understanding of cluster growth, structure, and mass.  The region near the boundary of the dynamically relaxed cluster volume ($\\sim \\rth$ \\cite{Evrardetal1996}), offers a direct view of cluster growth and ICM accretion \\cite{Voit2005}. Simulations predict deviations from hydrostatic equilibrium in the outskirts of clusters at scales on the order of $\\rtfh<r<\\rth$, well beyond the influence of cooling and AGN feedback processes in the  cluster core \\cite{Roncarellietal2006,Burnsetal2010}.  Moreover, it may be that signatures of spatially inhomogeneous accretion from cosmic filaments are observable at these radii. X-ray observations of clusters at $r>\\rfh$ also probe the distribution of mass to large radii, and in so doing provide new constraints on structure formation models \\cite{NFW1997,Borganietal2004,Roncarellietal2006}.  Finally, accurate knowledge of the mass and baryonic fraction of  clusters is essential for cluster-based cosmological tests \\cite{Vikhlininetal2009a,Allenetal2004}. Accurate cluster masses within \\rfh, for example, required for many cosmological tests, rely on knowledge of both the temperature and its radial gradient at \\rfh, and observations to larger radius are essential to determine these quantities reliably.  \\begin{figure}[t] \\begin{minipage}{.38\\textwidth} \\includegraphics[height=\\textwidth,angle=270]{a1795_sb.eps} \\end{minipage} \\begin{minipage}{.24\\textwidth} \\includegraphics[height=\\textwidth,angle=0]{a1795_pretty_wedges.eps} \\end{minipage} \\begin{minipage}{.38\\textwidth} \\includegraphics[height=\\textwidth,angle=270]{a1795_tprof.eps} \\end{minipage} \\caption{Abell 1795 surface brightness (left) and temperature profiles (right).  The image (center) shows color-coding for the radial directions probed; for clarity, only the North and South directions are shown in the surface brightness profile.  A 4-$\\sigma$ difference is seen in the surface brightness between the North and South \\citep{Bautzetal2009}, and a rapid decline in temperature is seen toward all directions.} \\label{fig:a1795tprof} \\end{figure}  Spatially resolved X-ray spectroscopy is one of the prime methods for mapping the distribution of mass and intracluster plasma, and it is the only way to measure heavy elements in the gas in galaxy clusters.  Although \\chandra\\ and \\xmm\\ have successfully probed the central regions of clusters \\cite{Vikhlininetal2009b,Prattetal2007,Snowdenetal2008}, these instruments have difficulty observing the faint emission near \\rth\\ due to their relatively high and time-variable background, and in fact conditions in the ICM beyond $\\sim$0.5\\rth\\ have not been well studied by these observatories.  \\suzaku's low and stable background and substantial effective area, however, have already enabled observations of a handful of clusters that have successfully detected ICM emission out to and even beyond \\rth\\ \\cite{Georgeetal2009,Reiprichetal2009,Bautzetal2009,Hoshinoetal2010,Kawaharadaetal2010,Simionescuetal2011,Akamatsuetal2011}. Our own \\suzaku\\ observations of Abell 1795 \\cite{Bautzetal2009} find significant deviations from hydrostatic equilibrium within \\rtfh--\\rth, in a cluster which appears to be relaxed at smaller radii. These results point to a rapidly declining temperature profile (see Figure \\ref{fig:a1795tprof}), an entropy profile that is flatter than predicted for simple hierarchical structure formation \\cite{Voit2005}, and clearly asymmetric temperature and density profiles at large radius. The steep profile may in turn result from cool, low-entropy infalling substructure, as suggested by the simulations of \\citet{Roncarellietal2006}.  Indeed, \\citet{Burnsetal2010} successfully reproduce with hydrodynamic simulations the temperature and entropy profiles seen in Abell 1795, suggesting that bulk flows and turbulence from accretion greatly affect the energetics of the outer ICM. ",
        "conclusions": ""
    },
    "1112/1112.3037_arXiv.txt": {
        "abstract": "\\textit{Preliminary} constraints are placed on a cluster of YSOs  (J2000 02:54:31.4 +69:20:32.5) discovered in the field of the classical Cepheid SU Cas.  WISE 3.4, 4.6, 12, and 22 $\\mu m$ images reveal that the cluster deviates from spherical symmetry and exhibits an apparent diameter of $3 \\times 6 \\arcmin$.  SEDs constructed using 2MASS $K_s$ ($2.2 \\mu m$) and WISE photometry indicate that 19 (36\\%) class I, 21 (40\\%) class f, and 13 (25\\%) class II objects lie $r<3 \\arcmin$ from the cluster center.   Conversely, 11 (18\\%) class I, 13 (21\\%) class f, and 37 (61\\%) class II objects were detected for $r>3\\arcmin$.   Approximately 50\\% of the class I sources within $r<3 \\arcmin$ were classified solely using WISE photometry owing to the absence of detections by 2MASS. ",
        "introduction": "} The latest generation of near and mid infrared surveys such as the VVV \\citep[\\textit{VISTA Variables in the Via Lactea},][]{mi10}, UKIDSS \\citep[\\textit{UKIRT Infrared Deep Sky Survey},][]{lu08}, GLIMPSE (\\textit{Galactic Legacy Infrared Midplane Survey Extraordinaire}), and WISE \\citep[\\textit{Wide-field Infrared Survey Explorer},][]{wr10} surveys have fostered the discovery of countless stellar nurseries.  \\citet{me05} discovered 92 star clusters using GLIMPSE data (3.6, 5.8, 8.0, $24 \\mu m$), of which $\\sim67$\\% likely host a young stellar demographic. Similarly, a sizable fraction of the 96 clusters discovered by \\citet{bo11} in the VVV survey ($YZJHK_s$) house heavily obscured young populations. Infrared photometry is particularly efficient at revealing embedded sources which are otherwise obscured from optical surveys since $A_{\\lambda} \\propto \\lambda^{-\\beta}$ \\citep[e.g.,][]{ni09}.  The total extinction tied to photometry employed here is $A_{K_s}/ A_V\\sim0.1$ for 2MASS $K_s$ ($\\sim 2.2 \\mu m$), and $A_{[3.4]}$:$A_{[4.5]}$:$A_{[12]}$:$A_{[22]} \\sim0.05 A_{V}$ for WISE photometry \\citep[e.g.,][]{fl07}.   \\begin{figure*} \\begin{center} \\includegraphics[width=6.3cm]{fig1-field.eps} \\includegraphics[width=6.2cm]{fig1-cl.eps} \\end{center} \\caption{\\small{Left/right, optical DSS (\\textit{Digitized Sky Survey}) and WISE 3.4 $\\mu m$ images centered at J2000 02:54:31.4 +69:20:32.5.  The images span $\\sim 15 \\times 15 \\arcmin$ and were aligned using the Aladin software environment \\citep{bo00}.  The YSOs are absent from the optical image owing to significant interstellar and circumstellar extinction.}} \\label{fig-cl} \\end{figure*}  In this study, initial constraints are placed on young stellar objects (YSOs) discovered in the field of the $1.95^{\\rm d}$ classical Cepheid SU Cas.  In \\S \\ref{s-cl}, optical (DSS) and $3.4 \\mu m$ (WISE) images of the cluster are compared, and its apparent extent is determined from the latter; \\S \\ref{s-sed}, SEDs are constructed to classify the YSOs according to the prescription outlined by \\citet{la87}, and to assess whether the YSO class varies as a function of radial distance from the cluster center; \\S \\ref{s-ccd}, $JHK_s$ color-color diagrams are constructed for cluster stars and an adjacent comparison field. ",
        "conclusions": "} Countless YSOs were discovered at J2000 coordinates 02:54:31.4 +69:20:32.5 (Fig.~\\ref{fig-cl}), which is $30 \\arcmin$ north east of the classical Cepheid SU Cas.  The cluster appears non-symmetric and displays an apparent diameter of $6 \\arcmin$ (Fig.~\\ref{fig-cl}).  SEDs constructed using 2MASS and WISE photometry indicate that the cluster hosts 19 (36\\%) class I, 21 (40\\%) class f, and 13 (25\\%) class II objects (Fig.~\\ref{fig-sed}).  At a distance beyond 3$\\arcmin$ from the cluster center the sampling changes markedly to 11 (18\\%) class I, 13 (20\\%) class f, and 37 (60\\%) class II objects.  The statistics reaffirm that the YSO class exhibits a radial dependence.  10 class I sources lacking 2MASS $K_s$ photometry were classified via SEDs tied to WISE observations.  The slope $\\alpha$ (equation \\ref{eqn-fd}) inferred from that approach appears less biased by the wavelength dependence of extinction since WISE passbands exhibit nearly equivalent extinction ratios (as compared to $K_s$).  The extreme extinction obscures many of the cluster's YSOs ($A_V\\sim 35$) beyond optical detection (Fig.~\\ref{fig-ccd}).  The latest generation of infrared surveys (e.g., WISE, GLIMPSE, etc.) shall continue to foster the detection and characterization of new stellar clusters, as reaffirmed here.  Additional research, which includes spectroscopic follow-up of the cluster and broader region, is required.   \\subsection*{{\\rm \\scriptsize ACKNOWLEDGEMENTS}} \\scriptsize{DM is grateful to the following individuals and consortia whose efforts lie at the foundation of the research: 2MASS, WISE, J. Gizis, E. Paunzen (WebDA), W. Dias (DAML), CDS, arXiv, and NASA ADS.  WG is grateful for support from the Chilean Center for Astrophysics FONDAP 15010003 and the BASAL Centro de Astrofisica y Tecnologias Afines (CATA) PFB-06/2007.  This publication makes use of data products from the Wide-field Infrared Survey Explorer, which is a joint project of the University of California, Los Angeles, and the Jet Propulsion Laboratory/California Institute of Technology, funded by NASA.}"
    },
    "1112/1112.4421_arXiv.txt": {
        "abstract": "{ There are extremely luminous quasi stellar objects (QSOs) at high redshift which are absent at low redshift. The lower luminosities at low redshifts can be understood as the external manifestation of either a lower Eddington ratio or a lower mass. To distinguish between both effects, we determine the possible dependence of masses and Eddington ratios of QSOs with a fixed luminosity as a function of redshifts; this avoids the Malmquist bias or any other selection effect. For the masses and Eddington ratios derived for a sample of QSOs in the Sloan Digital Sky Survey, we model their evolution by a double linear fit separating the dependence on redshifts and luminosities. The validity of the fits and possible systematic effects were tested by the use of different estimators of masses or bolometric luminosities, and possible intergalactic extinction effects. The results do not show any significant evolution of black hole masses or Eddington ratios for equal luminosity QSOs. The black hole mass only depends on the bolometric luminosity without significant dependence on the redshift as \\protect{$\\left(\\frac{M_{BH}}{10^9\\ M_\\odot}\\right)\\approx 3.4\\left(\\frac{L_{bol}}{10^{47}\\ {\\rm erg/s}}\\right)^{0.65}$} on average for \\protect{$z\\le 5$}. This must not be confused with the possible evolution in the formation of black holes in QSOs. The variations of environment might influence the formation of the black holes but not its subsequent accretion. It also leaves a question to be solved: Why are there not QSOs with very high mass at low redshift? A brief discussion of the possible reasons for this is tentatively pointed out. ",
        "introduction": "\\label{.introd}  Quasi stellar objects (QSOs) are extremely bright at high redshift, but at low redshift they are much less luminous; in fact from the analysis of the bolometric luminosity function of QSOs at different redshift (Hopkins et al. 2007), it is clear that the relative abundance of high luminosity QSOs decreases quickly at low redshift.  In the Sloan Digital Sky Survey (SDSS), all QSOs at $z<0.4$ are fainter than $M_B=-26$  (with K-corrections) while there are numerous QSOs tens of times brighter than this limit at higher redshifts. The same effect is also observed in other spectral ranges; for instance in X-rays  (Shen et al. 2006) or in radio  (Bridle \\& Perley 1984; Bell 2006, Figs. 9, 10). All these observations require a strong density and luminosity evolution.  Something very different must have happened at high redshift with respect to the low redshift Universe to obtain this different level of luminosity. However, no visible signs of this evolution are observed. There is no indication of any significant evolution in the X-ray properties of  quasars between redshifts 0 and 6, apart from the intrinsic luminosity,  suggesting that the physical processes of accretion onto massive black  holes have not changed over the bulk of cosmic  time (Vignali et al. 2005). Also, the spectral features of low and high  redshift QSOs are very similar (Segal \\& Nicoll 1998). This strong change in luminosity without any additional external sign of evolution is one of the most relevant pending problems in QSOs (L\\'opez-Corredoira 2011).  There are two possible reasons for the lower luminosity of low $z$ QSOs with respect to those ones with high $z$: 1) either their black holes are less massive, and this would explain the lower power of their accretion disks, 2) or they have obtained approximately the same mass but they are less efficient, due to a lower accretion rate. There are several works that estimated masses and Eddington ratios of QSOs' black holes (McLure \\& Dunlop 2004, Kollmeier et al. 2006, Vestergaard \\& Peterson 2006, La Mura et al. 2007, Vestergaard et al. 2008, Shen et al. 2008 [herein S08], Labita et al. 2009a,b). The result in general is that masses and Eddington ratios are larger at higher $z$. Nonetheless, it cannot be directly interpreted as a sign of evolution of QSOs because the samples used are strongly affected by the Malmquist bias, that is, we are comparing low luminous QSOs at low $z$ with high luminous QSOs at high $z$. S08 made a separation in bins of fixed luminosity showing qualitatively that much of this apparent evolution is due to luminosity variation. We think that this analysis requires further attention in order to properly quantify it and the possible systematics involved. This is the main aim of this paper in which we develop a simple method to study the dependence of black hole masses and Eddington ratios on their redshifts and luminosities. ",
        "conclusions": "The main conclusion of our analysis is that both the mass and the Eddington ratio of the black holes for a QSO with a given luminosity do not evolve with redshift. Or in other words, the luminosity of a QSO does not evolve with redshift for a given mass. More precissely and considering  systematic uncertainties  $\\sim 0.2-0.3$ dex in the estimation of masses and luminosities, we conclude that the evolution in redshift, if any, is very small compared  to the change in mean luminosity of the population of QSOs at low redshift with respect to such a population at high redshifts. This implies an important result on the nature of QSOs, i.e. local QSOs are intrinsically less massive than QSOs at high redshift. Labita et al. (2009a) derived that the maximum mass of a black hole in a QSO is a function of the redshift: $\\log _{10}M_{BH}=(0.34z+8.99)$ M$_\\odot$ up to redshift 1.9, or proportional to $(1+z)^{1.64}$ if extended up to a redshift of 4 (Labita et al. 2009b). This lack of the signature of active massive AGN black holes in the local Universe cannot be related with a possible decline in the rate of formation of QSOs (this would affect the density of QSOs but not their average mass; and indeed there is evidence for the change in the comoving density of QSOs of a given mass; Steinhardt \\& Elvis 2011, fig. 3), but because of some mechanism for the formation of huge black holes which took place in the past in the Universe, which is absent in the present Universe.  NOTE: Do not confuse the non-evolution of the black hole mass-luminosity ratio (the result of this paper) with the non-evolution of mechanisms which produce such black holes. Evidently, as said in the introduction, some evolution in the birth of new QSOs must take place in order to explain the absence of very bright QSOs at low redshift.  The fact that the Eddington ratio of the QSO does not change with $z$ for a given luminosity/mass counters scenarios which explain the evolution of the luminosity of QSOs in terms of the change in the environment of the AGNs, like variations in the accretion rate. Our result is at odds with the common assumption that relates the influence  of companion galaxies with the mechanism of feeding the  black hole of the QSO (Stockton 1982; Canalizo \\& Stockton 2001). Horst \\& Duschl (2008) presented the results of a simple cosmological model combined with an evolutionary scenario in which both the formation of the black hole as well as the gas accretion  onto it are triggered by major mergers of gas-rich galaxies. Despite the very  generous number of approximations, their model reproduces the quasar density evolution in  remarkable agreement with some observations. However, we do not see this decrease of gas accretion here, so the environment does not seem to be the major factor responsible for the change in luminosity of QSOs. We do not deny, however, the effect of the environment on the mechanism of formation/turn-off of its own massive black hole, although the synchronization of all very-massive QSOs in an epoch turning off nearly at the same time is not understood, because galaxies continue to merge and virialize at some rate at later epochs (Steinhardt \\& Elvis 2011).  Other results from other papers are also apparently at odds with the idea of powerful AGNs in a rich highly interactive environment. Coldwell \\& Lambas (2006) showed that quasars at $z<0.2$ systematically avoid high density regions, living in regions less dense than cluster environments. At $0.5\\le z\\le 0.8$, only 10\\% of QSOs live in relatively rich clusters, and 45\\% of them in field-like environments (Wold et al. 2001). We might also consider that galaxies in rich clusters are stripped of their interstellar medium by harassment, so it would be reasonable that the QSO activity is less than in the field galaxies, but the ratio of spiral galaxies with non-stripped gas is still high enough to consider that there should be activity being triggered. More recently, Cisternas et al. (2011) showed directly that there is not an enhanced frequency of merger signatures for the AGN hosts with respect to other galaxies, so this points out that mergers should not be an important element for the triggering of activity.  There are other kind of models on the origins and the early evolution of QSOs and supermassive black holes (review at Djorgovski et al. 2008). Many works have related the evolution of QSOs with their star formation ratios; e. g.  Haiman et al. (2007) assume that star formation in spheroids (elliptical galaxies and bulges of late-type galaxies) and black hole fueling are proportional to one another at all times, and fitting conveniently some parameters get a model of luminosity evolution of QSOs in agreement with observed data with quasar lifetimes $\\approx 6-8\\times 10^7$ yr, without a compelling need for any of the model parameters to evolve with redshift between $0<z<6$. This result supports the direct connection between the build up of spheroids and their nuclear supermassive-black-holes.  The existence of very massive black holes is only at high $z$. Perhaps it could be related with the higher ratio of mergers and then star formation in the past. We wonder whether it might have something to do with the excess of very massive galaxies at high-$z$, which is still not completely understood within semianalytical hierarchical $\\Lambda $CDM models (e.g., Fontana et al. 2009).  Indeed, the mass of the black hole has remained proportional to the stellar mass of their host galaxies for at least the last 9 Gyr (Jahnke et al. 2009). Or perhaps it has something to do with the larger  average density of the Universe, or the angular momentum of some components of the galaxies (at high $z$, the black holes rotate faster; Netzer 2010). However, where are the black holes with masses larger than $10^{10}$ M$_\\odot $ which were frequent in the past? We do not know any mechanism by which black holes can reduce its mass. Also, one should not lose sight of some solutions in which, for some reason (e.g., radiation emitted in beams rather than isotropically over $4\\pi$ stereoradians, non-cosmological redshifts, wrong cosmological model, etc.), the luminosity of the QSOs does not correspond to $4\\pi d_L(z)^2Flux$, with $d_L(z)$ the luminosity distance given by standard cosmology, and consequently both the luminosities and the masses\\footnote{The calculation of the mass depends on the luminosity of the continuum, like in the application by S08, or of some line, like in Eq. (\\ref{mbh_beta_green}). Note, moreover, that there is not a  significant variation of the width of broad lines with redshift (S08, fig. 3) so this  reinforces the idea that huge masses are obtained because huge luminosities are measured, and an overestimation of the luminosity directly produces an overestimation of the masses.}  would be erroneously determined. Clearly the question remains open and it is beyond the scope of the present paper.  \\normalem"
    },
    "1112/1112.4798_arXiv.txt": {
        "abstract": "{The origin of the neutrino mass, dark matter (DM), and dark energy (DE) are among the most challenging problems of fundamental physics. We address these questions from analyses of the models where the neutrino mass is generated via Yukawa coupling to the quintessence field which represents the DE. It has been shown in a recent work on the toy model with a single Dirac fermion coupled to the quintessence that by choosing parameters of the DE potential to match the present DE density, the model allows to lock the neutrino mass at $m \\sim 10^{-2}~eV$ and yields consistent estimates for other parameters of the Universe. To include the DM component into this framework we propose here to add the right-handed Majorana term(s) into Lagrangian with the quintessence-generated mass. As a results, the DE field is responsible for generation of the masses of the light active Majorana neutrinos as well as of the heavy sterile neutrinos. The latter are natural DM candidates. }  \\FullConference{The XXth International Workshop  High Energy Physics and Quantum Field Theory     \\\\ September 24 - October 1,  2011 \\\\ Sochi, Russia}  \\begin{document} ",
        "introduction": "It appears now that consensus has being reached on our understanding of the composition of the Universe and fundamental issues demanding extensions of the standard model of particle physics. Roughly 75\\% of the Universe is comprised of dark energy (DE), about 22\\% of it is comprised of dark matter (DM), and only approximately 3\\% corresponds to the contribution of the conventional (baryonic) matter. The present stage evolution of the Universe is governed by the dominant DE contribution, and the Universe experiences an accelerating expansion. The nature of the both DE and DM is still unknown, and this is probably one of the greatest mysteries of modern physics \\cite{DE}. The origin of neutrino mass is also a fundamental problem for particle physics as well as for cosmology \\cite{Dolgov08,Les06}. Current upper limits on the sum of neutrino masses from cosmological observations are of the order of 1eV, while neutrino oscillations give a lower bound of roughly 0.01 eV \\cite{Les06}. Even if  at least \\textit{ad hoc}, minimalistically and intellectually unsatisfactory, the DE can be accounted for by adding a cosmological term to fit current observations, the existence of neutrino masses and DM particles definitely requires revisions of the standard model. There is a vast literature and a great interest in these issues. Due to space constraints we address the reader to ``a short list'' \\cite{DE,Dolgov08,Les06,ShaposhnikovNuMSM,Shaposhnikov,Kayser03} where more references, especially on earlier work could be found.  In this paper we summarize and advance further our recent work \\cite{Us11}. It stems from the earlier proposal by Fardon \\textit{et al} \\cite{Fardon04}. In the approach \\cite{Us11} the origin of the neutrino mass is due to the dark energy in the form of quintessence. Both the DE density and the neutrino mass are determined by a single mass scale of the DE potential, and consistent estimates for other parameters of the Universe are obtained.  Here we propose how to incorporate  the ``third element of the puzzle'', i.e, the DM component of the Universe within a common framework. For this end we analyze a quite straightforward extension of the ``toy-model'' of a single Dirac fermion considered in \\cite{Us11} by adding a right-handed Majorana fermion whose mass is due to the Yukawa interaction with the quintessence field. The Dirac-Majorana action for a single generation of neutrinos is known to be equivalent to two massive Majorana fermions. One of these, the (light) neutrino with mass $m_a$ decreasing in time can be identified with the conventional (active) neutrino, while the (heavy) sterile neutrino with growing mass $m_s$ is a natural candidate for the DM particle. The results and estimates we report below demonstrate that the theory is viable and consistent in unifying neutrino masses and the dark sector of the Universe. ",
        "conclusions": "It appears that the above model of Majorana neutrinos interacting with the quintessence and getting their masses from DE is consistent, and it encourages for further work on it. But we have to stress that below the phase transition the model is probed so far only by consistency checks between various parameters ($\\rho_{DE},~m_a,~m_s$) and observables. Analysis of the equation of motion (\\ref{EqMotion}) for the quintessence field and/or by using more general methods of kinetics of the first-order phase transition is warranted.  However, the model itself needs to be more realistic. It seems to us that the so-called $\\nu$MSM  \\cite{ShaposhnikovNuMSM,Shaposhnikov} is the right model which has already properties needed. This is an extension of the standard model where three additional right-handed Majorana gauge singlets for each of the neutrino flavors are added. With the model's available fitting parameters it is possible \\cite{ShaposhnikovNuMSM} to get along with the three light active neutrinos also the three sterile neutrinos with consistent abundances, masses, and life times.\\footnote{In addition this model also provides a mechanism for baryon asymmetry of the Universe.} One of the sterile neutrinos with the life time comparable to the age of the Universe has a mass in the KeV range and is an excellent ``warm'' DM candidate, while the other short-lived heavy neutrinos have their masses in the GeV range and are decaying DM particles. The results presented here allow us to anticipate that using the quintessence field for generating the right-handed Majorana massive terms (\\ref{LMaj}) in this extended standard model would also allow to obtain the observable DE density and the slow-rolling quintessence dynamics with $w_\\varphi \\equiv P_\\varphi  / \\rho_\\varphi \\approx -1$. This study is currently work in progress."
    },
    "1112/1112.1342_arXiv.txt": {
        "abstract": "{The phase-integral approximation devised by Fr\\\"oman and Fr\\\"oman, is used for computing cosmological  perturbations in the quartic chaotic inflationary model. The phase-integral formulas for the scalar power spectrum are explicitly obtained up to fifth order of the phase-integral approximation. As in  previous reports \\cite{rojas:2007b,rojas:2007c,rojas:2009}, we point out that the accuracy of the phase-integral approximation compares favorably with the numerical results and those obtained using the slow-roll  and uniform approximation methods.} ",
        "introduction": "Chaotic inflationary models like $V \\propto \\phi^\\mathrm{p}$ have been studied and analized with the WMAP data in recent years. The recent observations reported by WMAP favors inflation,  and the quartic chaotic model, with $\\mathrm{p}=4$ \\cite{linde:1983a,linde:1983b}  still remains ruled out by this data  \\cite{komatsu:2011,komatsu:2009}.  Ramirez {\\it et al.} \\cite{ramirez:2009} point out that  if  the scalar power spectrum is considered not scale invariant,  then quartic models would not be excluded by the recent cosmological data.  Ramirez {\\it et al.} \\cite{ramirez:2009}  consider a new scenario where the observable modes cross the horizon near the beginning of inflation.  In order to compare with observations,  we should be able to obtain very accurate results for the predicted power spectrum of primordial perturbations for a variety of inflationary scenarios. In general, most of the inflationary models are not exactly solvable and approximate or numerical methods are mandatory in the computation of the scalar and power spectra.  Traditionally, the method of approximation applied in inflationary cosmology is the slow-roll approximation \\cite{stewart:1993}, which produces reliable results in inflationary models with smooth potentials,  but cannot be improved  on a simple way beyond the leading order. Recently, some authors have applied alternative approximations, such as the WKB method with the Langer modification \\cite{langer:1937,martin:2003a,casadio:2005a}, the Green function method \\cite{stewart:2001}, and the improved WKB method \\cite{casadio:2005c,casadio:2005b}.  Habib \\textit{et al} \\cite{habib:2002,habib:2004,habib:2005b}  have successfully applied the uniform approximation method in the calculation of the scalar and tensor power spectra and the corresponding spectral indices for the quadratic and quartic chaotic inflationary models, showing that the the uniform approximation gives more accurate results than the slow-roll approximation. Casadio \\textit{et al} \\cite{casadio:2006} have applied  the method of comparison equation to study cosmological perturbations during inflation.  The comparison method is based on  the uniform approximation proposed by Dingle \\cite{dingle:1956} and Miller \\cite{miller:1953} and thoroughly discussed by Berry and Mount \\cite{berry:1990}.  Recently \\cite{rojas:2007b,rojas:2007c,rojas:2009}, the phase integral method has been used in the calculation of the scalar and tensor power spectra and the spectral indexes for the power-law and quadratic chaotic model, obtaining better results than  those derived using slow-roll \\cite{stewart:1993}, WKB and the improved uniform approximation methods \\cite{habib:2002,habib:2004,habib:2005b}.  Since the 7-year WMAP data does not show evidence for tensor modes \\cite{larson:2011,alabidi:2010}, in this article we use the phase-integral approximation up to fifth order to calculate the scalar power spectrum for the quartic chaotic inflationary model. We have calculated the tensor power spectrum and  since the results  are very similar to those obtained in the scalar case, we decided to skip their presentation in this article.f  The article is structured as follows: In Sec. \\ref{perturbations} we numerically solve the equation governing the scalar perturbations. In Sec. \\ref{methods} we apply the phase-integral approximation to the quartic chaotic inflationary model.  We present the improved uniform approximation method,  the slow-roll approximation and the numerical solution to the perturbation equation. The comparison among different  methods is presented in Sec. \\ref{results}. Finally in Sec. \\ref{conclusion} we present our conclusions. ",
        "conclusions": "\\label{conclusion}  As in our previous work \\cite{rojas:2007b,rojas:2007c,rojas:2009},  we show that the phase-integral approximation is an efficient method to calculate the scalar (or tensor) power spectrum and the spectral index in several models of inflation. Our results indicate that the scalar power spectrum is almost scale invariant since we have considered the usual scenario where the observable modes cross the horizon before the inflation ends,  unlike the proposal presented by Ramirez {\\it et. al} \\cite{ramirez:2009}. Therefore, the good agreement between the numerical results and those obtained with the phase-integral approximation shows that the phase integral method is a very useful approximation tool for computing the scalar and tensor the power spectra in  a wide range of inflationar scenarios."
    },
    "1112/1112.5341_arXiv.txt": {
        "abstract": "{Many X--ray accreting pulsars have a soft excess below 10 keV. This feature has been detected also in faint sources and at low luminosity levels, suggesting that it is an ubiquitous phenomenon. In the case of the high luminosity pulsars ($L_{\\rm X} > 10^{36}$ erg s$^{-1}$), the fit of this component with thermal emission models usually provides low temperatures ($kT <$ 0.5 keV) and large emission regions ($R \\ge$ a few hundred km); for this reason, it is referred to as a `soft' excess. On the other hand, we recently found that in persistent, low--luminosity ($L_{\\rm X} \\sim 10^{34}$ erg s$^{-1}$) and long--period ($P >$ 100 s) Be accreting pulsars the observed excess can be modeled with a rather hot ($kT_{\\rm BB} >$ 1 keV) blackbody component of small area ($R_{\\rm BB} <$ 0.5 km), which can be interpreted as emission from the NS polar caps. In this paper we present the results of a recent \\XMM\\ observation of the Galactic Be pulsar \\RX, which is a poorly studied member of this class of sources. We have found a best--fit period $P$ = 204.96 $\\pm$ 0.02 s, which implies an average pulsar spin--down during the last 13 years, with $\\dot P \\simeq 6\\times10^{-9}$ s s$^{-1}$. The estimated source luminosity is $L_{\\rm X} \\sim 8\\times10^{34}$ erg s$^{-1}$: this value is higher by a factor $<$ 10 compared to those obtained in the first source observations, but almost two orders of magnitude lower than those measured during a few outbursts detected in the latest years. The source spectrum can be described with a power law plus blackbody model, with $kT_{\\rm BB} = 1.34\\pm0.04$ keV and $R_{\\rm BB} = 273\\pm16$ m, suggesting a polar--cap origin of this component. Our results support the classification of \\RX\\ as a persistent Be/NS pulsar, and confirm that the hot blackbody spectral component is a common property of this class of sources. ",
        "introduction": "Most of the X--ray binary pulsars (XBPs) are High Mass X--Ray Binaries (HMXRBs) in which a Neutron Star (NS) with magnetic field B $\\sim 10^{12}$ G is accreting matter from a high--mass early--type star, either an OB supergiant or a Be star. They can be persistently bright, with luminosities in excess of 10$^{34}$ erg s$^{-1}$, or transient sources characterized by quiescent phases, with emission around 10$^{34}$ erg s$^{-1}$ or less, interrupted by bright outbursts reaching $L_{\\rm X} \\sim 10^{36-38}$ erg s$^{-1}$ \\citep{Negueruela98,Reig07,Sidoli10}.  In these sources the X--ray spectra between 0.1 and 10 keV are usually described by a rather flat power--law, with photon index $\\sim$ 1, but several XBPs have shown a marked \\textit{`soft}' X--ray excess above the main power--law component \\citep[see][ for a review]{LaPalombara&Mereghetti06}; it is well described by a thermal emission model (either blackbody, bremsstrahlung or mekal) with low temperature ($kT_{\\rm SE} <$ 0.5 keV) and large emission area ($R_{\\rm SE} \\ge$ a few hundred km). This feature has been detected not only in the high--luminosity sources (with $L_{\\rm X} \\sim 10^{37-38}$ erg s$^{-1}$) but also in several low--luminosity ($L_{\\rm X} \\sim 10^{35-36}$ erg s$^{-1}$) XBPs observed in the Small Magellanic Cloud (SMC), where its detection is favoured by the low interstellar absorption \\citep{Sasaki+03,Ueno+04,Majid+04,Haberl&Pietsch05,Haberl+08}. Only in a few cases this low--energy component showed coherent pulses and the debate over its origin remains open. \\citet{Hickox+04} have shown that a soft spectral component could be a very common, if not ubiquitous, feature intrinsic to X--ray pulsars: it is visible in all sources with a sufficiently high flux and small absorption, and its origin is related to the source total luminosity.  Recently, based on \\XMM\\ data, we have observed a clear thermal excess also in three of the four \\textit{persistent} Be pulsars originally identified by \\citet{Reig&Roche99}, i.e. \\mbox{RX J0146.9+6121/LS I +61$^{\\circ}$ 235} \\citep{LaPalombara&Mereghetti06}, \\mbox{4U 0352+309/X Persei} \\citep{LaPalombara&Mereghetti07}, and \\mbox{RX J1037.5-5647/LS 1698} \\citep{LaPalombara+09}. These three sources are characterized by a persistently low luminosity ($L_{\\rm X} \\sim 10^{34-35}$ erg s$^{-1}$) and a long pulse period ($P >$ 100 s). These properties suggest that the NS orbits the Be star in a wide and nearly circular orbit, continuously accreting material from the low--density outer regions of the circumstellar envelope; in the case of 4U 0352+309 this picture is supported by the long orbital period of 250.3 days \\citep{Delgado-Marti+01}. For these sources the detection of the thermal component was favoured by the small distance ($d \\le$ 5 kpc) and interstellar absorption ($N_{\\rm H} \\sim 10^{21}$ cm$^{-2}$). We found that their soft excess can be fitted only with a blackbody (other simple models are rejected), which contributes for 30-40 \\% of the total flux; interestingly, in comparison with the other, more luminous sources, their blackbody component is characterized by a higher temperature ($kT_{\\rm BB} >$ 1 keV) and a much smaller emission radius ($R_{\\rm BB} <$ 0.5 km). This \\textit{hot BB} spectral component sets these low--luminosity and long--period sources apart from all the other pulsars, strongly suggesting that they form a distinct class. In their case the thermal component could be due to a different emission mechanism than in the high--luminosity pulsars. Based on the work of \\citet{Hickox+04}, it can be attributed to emission from the neutron--star polar caps. This is supported by the emission area of the blackbody component, which is consistent with the estimated polar cap size, and by the fact that the low energy part of the spectrum is clearly pulsed.  In this paper we present the results of a recent \\XMM\\ observation of \\RX, the remaining member of this class of Be/NS pulsars. This system was discovered during the \\ROSAT\\ Galactic plane survey \\citep{Motch+97} and identified with LS V +44 17, a moderately reddened ($E(B-V)=0.65\\pm0.0$5) B0.2 Ve star at $\\sim$ 3.3 kpc \\citep{Reig11}. Thanks to observations with the PCA instrument on board \\XTE, \\citet{Reig&Roche99} performed the first detailed timing and spectral analysis, and discovered a pulsation with period $P = 202.5 \\pm 0.5$ s. Its spectrum was well fitted with different models (power--law, power--law plus blackbody, two black--bodies, cut--off power--law) and the measured flux implied a source luminosity of 3$\\times$10$^{34}$ erg s$^{-1}$ between 3 and 30 keV. \\RX\\ has been detected also in the hard X--ray range: it is reported (as source PBC J0440.9+4432) in the Palermo \\textit{Swift}-BAT hard X--ray catalogue \\citep{Cusumano+10}, with a 15-150 keV flux of $(2.0\\pm1.1)\\times10^{-11}$ erg cm$^{-2}$ s$^{-1}$, and in the catalogue obtained with the \\textit{INTEGRAL/IBIS} 7--year All--Sky Hard X--ray Survey \\citep{Krivonos+10a}, with a 17-60 keV flux of $(1.36\\pm0.22)\\times10^{-11}$ erg cm$^{-2}$ s$^{-1}$.  In the optical/IR waveband, the H$\\alpha$ line shows a double--peak profile, varying from symmetric shape to completely distorted on one side (V/R phases), with a correlation between the equivalent width of the H$\\alpha$ line and the infrared magnitudes: as the EW(H$\\alpha$) decreases the IR magnitudes become fainter. This long--term optical/IR variability is attributed to structural changes in the Be star's circumstellar disc, which alternates decline and recovery phases on typical time scales $T_{\\rm disc} >$ 10 yrs \\citep{Reig+05}.  During the \\XTE\\ observation this pulsar showed little X--ray variability (by a factor $<$ 10), therefore it was included in the class of persistent, low--luminosity and long--period Be pulsars. However, recently this source has shown three consecutive flux increases, spaced by $\\sim$ 5 months from each other: the first one was detected by \\textit{MAXI/GSC} on 31 March 2010 \\citep{Morii+10} and by \\XTE\\ on 6 April 2010 \\citep{Finger&Camero-Arranz10}; the second one was observed by \\textit{INTEGRAL} on 1 September 2010 \\citep{Krivonos+10b}; the third one was seen by \\textit{Swift} on 29 January 2011 \\citep{Tsygankov+11}. The mean time between the starts of the outbursts was $\\sim$ 155 days, and the peak luminosities of these flares were lower than 10$^{37}$ erg s$^{-1}$ (assuming a source distance of 3.3 kpc). Assuming that these events occured at the periastron passage of the NS, where it approaches the decretion disc of the Be star, their separation is a good estimate of the orbital period of the system. It is interesting to note that the value of 155 days is also in agreement with the period of about 150 days derived from the Corbet diagram of $P_{spin}$ vs $P_{orbit}$ \\citep{Corbet86}. Therefore, after 4U 0352+308/X Persei, \\RX\\ could be the second persistent Be/X-ray pulsar with a known orbital period. ",
        "conclusions": "We have analysed a $\\sim$ 17 ks \\XMM\\ observation of the Be/NS X--ray pulsar \\RX. The source was detected at a luminosity level $L_{\\rm X}\\simeq 8.3\\times10^{34}$ erg s$^{-1}$ in the 0.3--12 keV energy range: this value is less than one order of magnitude higher than the average level of the previous \\ROSAT\\ and \\XTE\\ observations, therefore it confirms the persistent nature of this source. This classification is not denied by the large flux increases recently observed, which imply luminosity levels of a few $10^{36}$ erg s$^{-1}$, since they can be explained by structural changes in the circumstellar disc of the Be mass donor.  The \\XMM\\ observation has provided a refined pulse period $P$ = 204.96 $\\pm$ 0.02 s, which, compared to previous measurements, implies an average pulsar spin--down $\\dot P = (6.4\\pm1.3)\\times10^{-9}$ s s$^{-1}$ between 1998 and 2011. It implies a low average accretion rate, in agreement with the persistent nature of the pulsar. The pulse profile shows a complex structure, which is not sinusoidal, and at all energies the PF is $\\sim$ 55 \\%. However the HR is characterized by significant variations, expecially around the pulse minimum, which indicates a spectral variability of the source.  The source spectrum shows a count excess above the main power--law component, described by a blackbody with $kT_{\\rm BB} \\simeq$ 1.3 keV and $R_{\\rm BB} \\simeq$ 270 m, which contributes $\\sim$ 35 \\% to the source luminosity between 0.3 and 12 keV. We found no evidence of a narrow iron K$_{\\alpha}$ line between 6 and 7 keV, with an upper limit of $\\sim$ 70 eV on its equivalent width. The blackbody radius is comparable to the estimated size of the NS polar--cap, which suggests that the origin of this component is on the NS surface. The phase--resolved spectroscopy neither confirms nor disproves this scenario, since the observed spectral variability along the pulse period can be attributed to the power--law component and is consistent with a constant thermal component.  The spectral properties of \\RX\\ are in full agreement with those observed in the other three persistent Be binary pulsars 4U 0352+309, RX J0146.9+6121 and RX J1037.5--5647, thus confirming that the \\textit{hot BB} spectral component is a common property of this type of sources. Moreover, we have shown that the same type of feature has been detected also in other low--luminosity and long--period pulsars; therefore it is an ubiquitous phenomenon which requires further investigations."
    },
    "1112/1112.2757_arXiv.txt": {
        "abstract": "We present models for the $slow$ neutron-capture process ($s$ process) in asymptotic giant branch (AGB) stars of metallicity [Fe/H]=$-$2.3 and masses 0.9 \\msun\\ to 6 \\msun. We encountered different regimes of neutron-capture nucleosynthesis increasing in importance as the stellar mass decreases: the $^{22}$Ne($\\alpha$,n)$^{25}$Mg reaction activated during the thermal pulses, the $^{13}$C($\\alpha$,n)$^{16}$O reaction activated in radiative conditions during the interpulse periods, and the $^{13}$C($\\alpha$,n)$^{16}$O reaction activated during the thermal pulses, also as a result of mild proton ingestion episodes. The models where the $^{13}$C burns radiatively (masses $\\simeq$ 2 \\msun) produce an overall good match to carbon-enhanced metal-poor (CEMP) stars showing $s$-process enhancements (CEMP-$s$), except they produce too much Na and F. On the other hand, none of our models can provide a match to the composition of CEMP stars also showing $rapid$-process enhancements (CEMP-$s/r$). The models fail to reproduce the observed Eu abundances, and they also fail to reproduce the correlation between the Eu and Ba abundances. They also cannot match the ratio of heavy to light s-process elements observed in many CEMP-$s/r$ stars, which can be more than ten times higher than in the solar system. To explain the composition of CEMP-$s/r$ stars we need to invoke the existence of a different ``$s/r$'' neutron-capture process either with features in-between the $s$ and the $r$ processes, or generated by superpositions of different neutron-capture processes in the same astrophysical site or in sites linked to each other - for example, in multiple stellar systems. ",
        "introduction": "\\label{sec:intro}  Following on from the discovery of a large number of halo stars with peculiar chemical compositions the last decade has seen much interest sparked in detailed spectroscopic observations and nucleosynthetic models of stars of low metallicity \\citep{beers05,sneden08}. These peculiar compositions provide important constraints on the evolution and nucleosynthesis of the low-metallicity stars that existed in the early Galaxy. By interpreting the observations in the light of stellar nucleosynthesis models we can derive clues as to the sites and nature of the first nucleosynthetic processes in the Galaxy. This may ultimately provide important insights into the form of the initial mass function of the first stars. The composition of metal-poor stars may also provide constraints on cold dark matter cosmological models of the formation of galaxies, which predict that galactic halos may be composed of different components of possibly different composition \\citep[e.g.,][]{carollo07}.  Roughly 10-20\\% \\citep{cohen05,lucatello06,carollo11} of old halo stars are enriched in carbon ([C/Fe]\\footnote{The standard spectroscopic notation [X/Y]=log$_{10}$(X/Y)$-$log$_{10}$(X/Y)$_{\\odot}$ is used throughout this paper.}$>$ 1) and are referred to as carbon-enhanced metal-poor (CEMP) stars.  Most CEMP stars are also enriched in nitrogen, with [N/Fe]$>$1. In most cases [N/Fe]$<$[C/Fe], but also [N/Fe]$\\sim$[C/Fe] and [N/Fe]$>$[C/Fe] are found in roughly 10\\% and 5\\%, respectively, of the total number of CEMP stars. Roughly 2/3 of CEMP stars also show enrichments in the elements heavier than iron, which are produced by neutron-capture processes \\citep{aoki07,sneden08}. These enhancements range from pure $s$ enhancements, of elements made primarily via $slow$-neutron captures (the $s$ process) in the case of ``CEMP-$s$'' stars ([Ba/Fe]$>$1), to $s$ and $r$ enhancements, of elements made via the $s$ process and the $r$ process, i.e., $rapid$ neutron captures (the $r$ process) in the case of ``CEMP-$s/r$'' stars with [Eu/Fe]$>$1 and [Ba/Eu]$>$0 \\citep{jonsell06}. There are also halo stars with pure $r$ enhancements, so-called $r$I (usually classified by setting 0.5$<$[Eu/Fe]$<$1) and $r$II ([Eu/Fe]$>$1), one of which, the famous CS\\,22892-052 \\citep{sneden03}, belongs to the CEMP star family. CEMP stars without enhancements in the elements heavier that Fe are commonly referred to as ``CEMP-no'' stars.  To deal with this wide variety of compositions different scenarios for the origin of halo stars with peculiar compositions have been proposed. Halo stars showing pure $r$ enhancements are believed to have been born from material heavily polluted by a primordial core-collapse supernova (SN of Type II, SNII) where the $r$-process elements are supposed to have formed. Though we note that there are still no SNII models able to reproduce the conditions needed for the $r$ process to occur. One of the most favored sites to date, the neutrino wind from the nascent neutron star, has recently been shown to have major difficulties in producing $r$-process nuclei \\citep{roberts10}. Qualitatively, the C and $s$ enhancements in CEMP-$s$ and CEMP-$s/r$ stars can be explained by the presence of an originally more massive binary companion to the observed star, which produced C and $s$-process elements during its asymptotic giant branch (AGB) phase and transferred them to its companion via mass transfer, by either Roche-Lobe overflow or wind accretion \\citep{lucatello05b}. About 50\\% of CEMP stars showing an $s$-process signature also show an $r$-process signature and are thus classified as CEMP-$s/r$ stars. It is puzzling how CEMP-$s/r$ stars could have formed in such a large number given that we believe the $s$ and the $r$ processes to be independent events occurring in different astrophysical sites, AGB stars and SNII, respectively. Many scenarios have been proposed \\citep[as summarised by][]{jonsell06,lugaro09} and most of them seem unlikely. Predictions from a scenario where the $r$-process enhancements are supposed to be primordial to the binary system have been investigated in detail by \\citet{bisterzo11a} and will be discussed in \\S~\\ref{sec:baeu}.  The N enhancements in CEMP stars are also mysterious. The highest enhancements, [N/C]$>$0, may be due to the effect of proton captures at the base of the convective envelope in AGB stars with masses higher than $\\simeq$ 3 \\msun\\ \\citep[hot bottom burning, HBB][]{lattanzio91,boothroyd93}, which very efficiently convert C into N. The milder enhancements can be qualitatively explained by H and He burning occurring together when protons are ingested in the He-burning convective regions of AGB stars \\citep{campbell08,iwamoto04}, or by the operation of some form of deep mixing of the envelope material of the AGB star companion down to regions where H burning occurs. No physical mechanism for this deep mixing has been shown to be capable of producing the required nitrogen enhancement. \\citet{stancliffe10} has shown that it cannot be caused by thermohaline mixing. Another issue that has received much attention is the question of whether material transferred from the AGB companion onto the star now observed as a CEMP star was diluted on the secondary star via thermohaline mixing (or any non-convective process), or kept at the surface until the star reached the giant branch and developed an extended convective envelope \\citep{stancliffe07a,stancliffeglebbeek08}.  In this context, it is obvious that detailed models of the $s$ process in AGB stars of low metallicity are crucial to improving our understanding of CEMP stars. Some models have been presented so far by \\citet{vaneck03}, \\citet{cristallo09a}, and \\citet{bisterzo10} and have been compared to stellar observations by \\citet{bisterzo11a,bisterzo11b}. With this paper we greatly extend the range of published models of the $s$ process in AGB stars by presenting the evolution and nucleosynthesis in a large set of stellar models of masses ranging from 0.9 \\msun\\ to 6 \\msun\\ at the typical metallicity of CEMP stars [Fe/H]=$-$2.3. The $s$-process predictions for low-metallicity AGB stars of masses lower than 1.3 \\msun\\ and higher than 2 \\msun\\ are provided here for the first time in the literature. To test different physical and numerical approaches we computed stellar evolutionary sequences using two different evolutionary codes. We then input these different evolutionary sequences into the same post-processing code, which employs a nuclear network from H to Bi and calculates nucleosynthesis by coupling mixing and burning in convective regions. We then compare our model predictions to the abundances of CEMP stars as a population using the recent compilation of CEMP star compositions from \\citet{masseron10}.  \\subsection{The $s$ process in AGB stars of low metallicity} \\label{sprocessintro}  From centre to surface an AGB star consists of: a C-O degenerate core, a He-burning shell, a He-rich intershell, a H-burning shell, and an extended convective envelope, which is constantly eroded by strong stellar winds. When the whole envelope is lost the hot stellar core is left to evolve into the post-AGB phase, then it becomes the central star of a planetary nebula and eventually a cooling white dwarf. For a detailed review on AGB evolution we refer the reader to \\citet{herwig05}. A peculiarity of the AGB evolution is that the H- and the He-burning shells are activated alternately. Most of the time the H-burning shell is active and freshly made He increases the mass of the intershell for periods of time of around 10$^3$-10$^5$ yr, depending on the stellar mass, with shorter periods associated with more massive AGB stars. The increasing mass of the intershell leads to an increase in the temperature at its base, until He burning is suddenly triggered resulting in a {\\it thermal pulse} (TP). The energy released in this first phase of He-shell burning cannot be carried out radiatively and a convective region develops in the intershell, which lasts a few hundreds years (the exact duration depends on mass) and normally extends up to just below the H-burning shell. In low-mass and low-metallicity AGB stars this convective region can extend into the H-rich layer and peculiar nucleosynthesis is expected to occur there \\citep{campbell08}, including $s$-process nucleosynthesis \\citep{cristallo09b}. Mild proton-ingestion events of this type are found in our low-mass models. They have a strong impact on the $s$-process distributions and will be discussed in detail in \\S~\\ref{sec:results}.  A TP causes the star to expand, which in turn causes the H-burning shell to be quenched. At this stage, the convective envelope can penetrate into the underlying layers and carry to the stellar surface material processed by H and He burning (a process referred to as the {\\it third dredge-up}, or TDU for short), while He burning carries on radiatively. Eventually, He burning is also quenched and the whole cycle of H burning, TP, and TDU starts again. This cycle is repeated a few to hundreds of times depending on the time it takes for the mass loss to erode the envelope, which in turn depends on the initial stellar mass, and possibly the metallicity. The best evidence that TPs and the TDU occur in AGB stars is the formation of carbon stars, which show C$>$O at their surface as well as the presence of the radioactive element Tc \\citep{wallerstein98}. This observed composition is explained by the combined effect of partial He burning during TPs (which produces more C than O) and the $s$ process (which produces Tc), along with the TDU, which carries a fraction of the intershell to the stellar envelope. As mentioned above, HBB is present in AGB stars with masses higher than $\\simeq$ 3 \\msun\\ at the metallicity considered here, converting C into N.  The $s$ process occurs in the intershell of AGB stars where He is abundant and ($\\alpha$,n) reactions can be efficiently activated. For a detailed review on the $s$ process in AGB stars we refer the reader to \\citet{busso99}. There are two possible neutron sources: the \\iso{13}C($\\alpha$,n)\\iso{16}O and the \\iso{22}Ne($\\alpha$,n)\\iso{25}Mg reactions. The \\iso{22}Ne neutron source is activated at temperatures higher than roughly 300$\\times 10^{6}$\\,K (hereafter MK). These high temperatures can only be reached inside the convective TPs, and mainly in AGB stars more massive than $\\simeq$ 3 \\msun\\ \\citep{iben77}. The \\iso{22}Ne is produced via double $\\alpha$ captures on \\iso{14}N when the H-burning ashes are engulfed in the TP. The resulting neutron flux occurs over a timescale of the order of a few years with relatively high neutron densities, up to 10$^{14}$ n/cm$^{3}$ \\citep{gallino98,lugaro03b,vanraai11}. The \\iso{13}C neutron source is instead activated at temperatures as low as 90MK \\citep{cameron55} but there is not enough \\iso{13}C in the H-burning ashes of AGB stars to make it an efficient neutron source. Without this low-temperature neutron source it is not possible to explain the $s$-process enhancements which have been observed in AGB stars since the 1950s \\citep{merrill52,smith86b}. This is because most of the observed $s$-process-enriched AGB stars are of relatively low mass ($<$ 4 \\msun) where the temperature is not high enough to activate the \\iso{22}Ne neutron source. Thus, it has become standard practice to artificially include the \\iso{13}C neutron source in AGB models of the $s$ process.  The most likely way to form the required \\iso{13}C is to allow some extra mixing of protons to occur following each TDU episode at the deepest extent of the convective envelope, where a sharp composition discontinuity forms at the interface between the H-rich envelope and the He-rich intershell. This extra mixing may be physically explained by the effect of semi-convection \\citep{hollowell88}, hydrodynamical overshoot \\citep{herwig00,cristallo09a}, or gravity waves \\citep{denissenkov03a}. However, there is no agreement yet on which of these proposed mechanisms is responsible for the extra mixing and thus we do not know what its detailed features are. The protons are captured by the abundant \\iso{12}C, producing a layer rich in \\iso{13}C and \\iso{14}N \\citep[the $^{13}$C $pocket$,][]{hollowell88,gallino98,goriely00,lugaro03b,cristallo09a}. The resulting \\iso{13}C usually burns under radiative conditions via \\iso{13}C($\\alpha$,n)\\iso{16}O before the onset of the following TP \\citep{straniero95,gallino98}. This burning releases a large number of free neutrons on a long timescale (around $10^4$ yr) and with low neutron densities $< 10^8$ n/cm$^3$. In the region of the pocket where \\iso{14}N is more abundant than \\iso{13}C, no $s$-process nucleosynthesis occurs owing to the dominance of the \\iso{14}N(n,p)\\iso{14}C reaction over neutron captures by Fe seed nuclei and their progeny.  In the intermediate-mass AGB stars that experience HBB, the formation of a \\iso{13}C pocket may be inhibited by proton captures occurring at the hot base of the convective envelope during the TDU, which produce \\iso{14}N and not \\iso{13}C \\citep{goriely04}. Extremely deep TDU may also inhibit the activation of the \\iso{13}C($\\alpha$,n)\\iso{16}O reaction by penetrating into regions of the stellar core with a low abundance of \\iso{4}He \\citep{herwig04a}. While these processes do need further investigation, on first sight a dichotomy appears in models of the $s$ process in AGB stars where the \\iso{13}C neutron source is activated in the low-mass models, while the \\iso{22}Ne neutron source is activated in the higher-mass models. At the metallicity we consider here the mass at which the importance of the two neutron sources switches is $\\simeq$ 3 \\msun\\ \\citep{goriely04}.  An interesting feature of the \\iso{13}C neutron source is that it is a primary neutron source because it is produced from the H and He originally present in the stars. Since the total time-integrated neutron flux $\\tau$ from the \\iso{13}C source is proportional to \\iso{13}C/$Z$, the $s$-process predictions have a strong dependence on the metallicity of the star \\citep{clayton88,busso01}. Specifically, with $\\tau$ increasing when decreasing $Z$, heavier elements are produced at lower metallicities. This property allowed \\citet{gallino98} to predict the existence of low-metallicity Pb-rich stars, which was confirmed by observations of Pb-rich CEMP stars \\citep[e.g.,][]{vaneck03}. On the other hand, the \\iso{22}Ne neutron source is traditionally considered to be a secondary neutron source, since it relies on the original CNO nuclei present in the star. However, at the low metallicities considered here \\iso{22}Ne also has an important primary component because the N abundance in the H-burning ashes increases significantly from its initial value. This is due to the effect of TDU carrying primary \\iso{12}C to the envelope, which is then converted into N by H burning.  The \\iso{13}C and the \\iso{22}Ne neutron sources produce neutron fluxes and neutron-capture nucleosynthesis with different features. In summary, it has been shown that the \\iso{13}C source is responsible for the production of the bulk of the $s$-process elements in low-mass AGB stars reaching, as mentioned above, up to Pb at low metallicities \\citep{gallino98}. Branching points on the $s$-process path are mostly closed during this neutron flux. The \\iso{22}Ne source instead produces smaller abundances of the $s$-process elements but it efficiently activates branching points along the $s$-process path \\citep{abia01,vanraai11}. Note that at low metallicity \\iso{22}Ne is both a neutron source and a neutron poison via (n,$\\gamma$) reactions, leading to the production of light elements such as Na and Mg \\citep{herwig04b,karakas10a,bisterzo10}. Based on our models in \\S~\\ref{sec:results} we revise the different neutron-capture scenarios for AGB stars of low metallicity. ",
        "conclusions": "\\label{sec:conclusions}  We have presented a large set of stellar models at metallicity [Fe/H]=$-$2.3 from two different evolutionary codes for which we have computed detailed $s$-process nucleosynthesis. For the first time we have presented abundance predictions for the elements heavier than iron from stars with masses lower than 1.3 \\msun\\ and above 2 \\msun\\ at this metallicity. The most up-to-date models including the effect of molecular opacities that vary with the envelope composition during the AGB evolution experience less TPs, less TDU episodes, and produce overall lower enhancements of the $s$-process elements with respect to Fe.  We have found that there are four regimes of neutron-capture processes and the dominant regime depends on the initial stellar mass: for the highest mass models the \\iso{22}Ne neutron source is dominant and the inclusion of a \\iso{13}C pocket where \\iso{13}C burns radiatively has an effect only on the production of Pb; for models of mass $<$ 3.5 \\msun, \\iso{13}C is the main neutron source. For masses $\\simeq$ 2 \\msun, \\iso{13}C burns radiatively in the \\iso{13}C pocket during the interpulse periods. For lower masses, \\iso{13}C also burns convectively as the first few \\iso{13}C pocket are engulfed in the following TPs before \\iso{13}C is completely burnt in the interpulse period. For the lowest masses, proton-ingestion episodes associated with the first few TPs also produce \\iso{13}C, which burns convectively inside the TPs. Neutrons released during convective TPs produce lower [Pb/hs] ratios than neutrons released in the \\iso{13}C pocket under radiative conditions, and higher [ls/hs] ratios if a smaller number of protons are ingested. It should be kept in mind that proton ingestion episodes are extremely sensitive to the convective mixing scheme adopted. Our models provide some insight into the effect of these episodes on the $s$-process distribution, and are in agreement with the results presented by \\citet{cristallo09b}. However, more detailed 3D models of the kind presented by \\citet{stancliffe11} are needed to accurately ascertain the effect of proton-ingestion episodes on the $s$-process in AGB stars. The production of Mg and Na increases with the stellar mass as there is more primary \\iso{22}Ne seed available.  From comparison with the observational data-set of CEMP stars compiled by \\citet{masseron10} we have derived that our AGB models, especially those with radiative \\iso{13}C burning, present a very good match to the composition of CEMP-$s$ stars. On the other hand, they cannot match the compositions of CEMP-$s/r$ stars, since they cannot produce the observed Eu enhancements up to one order of magnitude higher than in CEMP-$s$ stars. If the $s$ process and $r$ process components, making Ba and Eu, respectively, in CEMP-$s/r$ stars are completely independent, as in the scenario where the $r$ enrichment is primordial, it is difficult to explain why the observed Ba and Eu are correlated. This is because for any given initial [$r$/Fe] abundance pattern imposed, the AGB models produce the same final [Ba/Fe] abundance. Since this conclusion relies on the observed correlation between Ba and Eu, it is  urgent to confirm and establish this correlation using a larger number of CEMP-$s/r$ star abundances derived from high-resolution spectroscopy and for a larger number of hs elements.  Not only do CEMP-$s/r$ and CEMP-$s$ stars differ in that CEMP-$s/r$ stars have more Eu than CEMP-$s$ stars in absolute terms, but they also differ in that CEMP-$s/r$ stars have on average more Ba (and other hs elements) than CEMP-$s$ stars. This behaviour is not readily explained if the $s$-process component in CEMP-$s$ and CEMP-$s/r$ stars come from the same AGB star sources. Furthermore, the higher Ba (hs) abundances are not accompanied by higher abundances in Mg and in the ls elements, as it is expected in our models. In fact, our AGB models always produce [ls/hs]$>-1$, which is too high to match the distribution of the $s$-process elements in CEMP-$s/r$ with [ls/hs] down to $\\simeq -2$. On the other hand the higher Ba abundances are accompanied by higher Pb abundances, as expected in the $s$ process. In summary, the composition of CEMP-$s/r$ stars is characterised by large enhancements of the hs elements and of Eu (correlated with Ba: $\\simeq 1/200$ Ba), as well as of Pb, but not of the ls elements. It remains to be seen if these compositions could be achieved by a ``$s/r$'' neutron-capture process, which may be a single neutron-capture process with features in-between or a superposition of the $s$ and $r$ processes. Parametric models of neutron-capture processes will be a first step in establishing if this is a possibility.  If the Ba versus Eu correlation in CEMP-$s/r$ stars is confirmed and if it is not due to a new ``$s/r$'' neutron-capture process, but to the $s$ process and the $r$ processes occurring in AGB stars and SNIIe, respectively, then these sites must be somehow correlated to each other. A probably unlikely scenario, which we provide here as an illustrative example, is that of a stable triple stellar system comprising the CEMP-$s/r$ star observed today, a former AGB star, and a former SNII. In this case the orbital parameters such as the distance between the stars, which affect the mass transfer, are correlated, probably resulting in correlated Ba and Eu abundances.  While the AGB models produce a good match to the observed [Mg/hs] ratios in CEMP-$s$ stars, they typically produce too high [Na/hs] and [F/hs] ratios. This may depend on the choice of the proton profile leading to the formation of the \\iso{13}C pocket, the proton-capture reaction rates involved, and the occurrence of extra-mixing processes. Investigation of these issues may result in important constraints on these uncertain stellar and nuclear inputs.  Finally, CS\\,30322-023 may carry the signature of nucleosynthesis in intermediate-mass AGB stars. This would be difficult to reconcile with the suggestion of \\citet{masseron06} that this star is an intrinsic AGB object based on the low log\\,$g=-0.3$. This would make this star an extremely peculiar object born less than a half a Gyr ago, but with an extremely low metallicity typical of the early Universe. Our conclusions are more in line with the higher value of log\\,$g=1$ reported by \\citet{aoki07}, not indicative of an intrinsic AGB star. From the preliminary comparison made here the composition of CS\\,30322-023 indicates that both HBB and the \\iso{13}C pocket were operating in low-metallicity massive AGB stars. However, before drawing conclusions we need to compare the composition of this star with predictions from a model of the appropriate metallicity of [Fe/H]=$-$3.3. Observations of Eu in CS\\,22960-053 and CS\\,29528-041 may reveal that these two stars are also CEMP-low-$s$ with [N/C]$>$1 creating a new class together with CS\\,30322-023. Future studies will focus on comparison of single stellar objects with our AGB model predictions."
    },
    "1112/1112.0108_arXiv.txt": {
        "abstract": "We present CO(3--2) interferometric observations of the central region of the Seyfert 2 galaxy NGC~1068 using the Submillimeter Array, together with CO(1--0) data taken with the Owens Valley Radio Observatory Millimeter Array. Both the CO(3--2) and CO(1--0) emission lines are mainly distributed within $\\sim5$ arcsec of the nucleus and along the spiral arms, but the intensity distributions show differences; the CO(3--2) map peaks in the nucleus, while the CO(1--0) emission is mainly located along the spiral arms. The CO(3--2)/CO(1--0) ratio is about 3.1 in the nucleus, which is four times as large as the average line ratio in the spiral arms, suggesting that the molecular gas there must be affected by the radiation arising from the AGN. On the other hand, the line ratios in the spiral arms vary over a wide range from 0.24 to 2.34 with a average value around 0.75, which is similar to the line ratios of star-formation regions, indicating that the molecular gas is affected by star formation. Besides, we see a tight correlation between CO(3--2)/(1--0) ratios in the spiral arms and star formation rate surface densities derived from \\emph{Spitzer} $8~\\micron$ dust flux densities. We also compare the CO(3--2)/(1--0) ratio and the star formation rate at different positions within the spiral arms; both are found to decrease as the radius from the nucleus increases. ",
        "introduction": "\\label{sect-intro}  The properties of circumnuclear molecular gas (CMG) near an active galactic nucleus (AGN) can be influenced by the AGN itself. Some CMG observations of Seyfert galaxies with a spatial resolution around 100 to 300 pc have shown that the intensities of CO(2--1) line emission are about two times higher in temperature units than those of CO(1--0) emission \\citep[e.g.,][]{mat04,hsi08}. This is different from the properties of molecular clouds in normal or star-forming galaxies, which usually show brightness temperatures in higher-J transitions similar to or lower than those in lower-J lines \\citep[e.g.,][]{dev94,mau99,dum01,oka07,mat09}. From the LVG model, the high values of CO(2--1)/CO(1--0) arise when the emitting molecular gas is in a low opacity environment and the resulting molecular densities are larger than the critical densities of CO(1--0) and CO(2--1) transitions. A possible origin for such unusual behavior is that CMG near AGNs traces X-ray dominated regions due to the strong radiation \\citep[e.g.,][]{mal96,koh01,use04,mei05,mei07}. Unlike far-UV photons, X-ray photons have greater penetration lengths and are more efficient in gas heating \\citep{use04}. On the other hand, CMG can also be affected by mechanical processes, such as gas entrainment by jets \\citep{mat07}.  Star formation in galaxies is well correlated with gas surface density as embodied in the famous Schmidt-Kennicutt law \\citep{sch59,ken98}. In the circumnuclear regions (CNRs) and inner structures of galaxies, the molecular gas usually has relatively higher densities and temperatures than the molecular gas in galaxy disks; in such environments, star formation is expected to be much more vigorous \\citep{ken98,tan00}. \\citet{kt07} suggest that the Schmidt-Kennicutt law might change slope when the averaged gas density is close to the line critical density. We note that the critical densities of CO(1--0) ($\\sim10^3$~cm$^{-3}$) and CO(3--2) $\\sim 3\\times10^4$~cm$^{-3}$ are comparable with the densities of the CNRs of some nearby active galaxies, which are are around $300$ to $10^7$~cm$^{-3}$ \\citep[e.g.,][]{mat04,sak07,per09}. It is unclear whether or how the Schmidt-Kennicutt law would vary in the CNRs of these galaxies when probed with CO(1--0) and CO(3--2) emission.  Among nearby galaxies, NGC~1068 is the best-studied prototypical Seyfert. It has a distance of 14.4~Mpc \\citep[][; $1\\arcsec\\sim70$~pc]{tul88} and is classified as (R)SA(rs)b in the RC3 catalog \\citep{vau91}, with an inner bar in the central kpc and two tightly wound spiral arms starting from the tip of the bar \\citep{sco88,thr89}. The inner 2 kpc is rich in star formation \\citep{tel84}, and the active star forming regions are concentrated along the spiral arms \\citep[e.g.,][]{tel88,dav98,ems06}. Molecular gas is also abundant along the spiral arms \\citep[e.g.,][]{mye87,pla91,kan92,hel95}, and a weak offset ridge of emission along the leading side of the bar is also seen \\citep{hel95,sch00}, which is a typical molecular gas distribution in barred galaxies.  The nucleus of NGC~1068 shows radio jets \\citep[e.g.,][]{wil83,gal96,gal04} and an ionization cone \\citep[e.g.,][]{pog88,mac94}. The nuclear optical spectrum has type 2 characteristics, but the polarized spectrum shows type 1 features, reflecting the existence of an optically thick torus or disk around the central massive black hole \\citep{ant85}. Indeed, a sub-parsec scale ionized gas disk perpendicular to the radio jets has been observed \\citep{gal97,gal04}, with a molecular gas (maser) torus \\citep{gre96} or disk \\citep{gal01} located outside of the ionized gas disk. This maser torus/disk may be surrounded by a pc-scale warm dust torus \\citep{jaf04}. Outside this structure, the existence of a warped molecular gas disk is suggested by interferometric CO(2--1) observations \\citep{bake98,sch00}, but warm molecular gas kinematics dominated by irregular (infalling) motions have been observed in $2.12~\\micron$ 1--0 S(1) molecular hydrogen emission \\citep{mul09}. Interferometric observations of molecular gas emission with angular resolutions at $0.\\arcsec$5--$2\\arcsec$ also indicated non-circular motions in the central $\\sim$100 pc \\citep{krip11}. The intensity ratios of various molecular species exhibit peculiar values, including very high HCN/CO and HCN/HCO$^{+}$ ratios \\citep{jac93,tac94,koh01,use04,krip08,per09}, suggesting that the molecular gas very close to the Seyfert 2 nucleus is irradiated by strong X-ray emission \\citep{use04,koh05,koh08,gar10}.  The central region of NGC~1068 has been mapped in different CO tranistion lines by both single-dish and interferometric observations. Among the former, it has been observed by \\citet[CO(1--0), FCRAO)]{sco83}, \\citet[CO(1--0), Nobeyama]{kan89}, \\citet[CO(1--0) and CO(2--1), IRAM]{pla89}, \\citet[CO(1--0), FCRAO]{you95}, \\citet[CO(2--1) and CO(3--2), JCMT]{pap99}, and \\citet[CO(1--0), CO(2--1), CO(3--2), and CO(4--3), IRAM and JCMT]{israel09}. With interferometers, it has been observed by \\citet[CO(1--0), OVRO]{pla91}, \\citet[CO(1--0), NMA]{kan92}, \\citet[CO(1--0), BIMA]{hel95}, \\citet[CO(1--0) and CO(2--1), IRAM]{sch00}, \\citet[CO(1--0), BIMA]{hel03}, and \\citet[CO(2-1) and CO(3-2), SMA/PdBI]{krip11}. In this paper, we present interferometric observations of the CO(3--2) line in the central 2 kpc of NGC~1068 using the Submillimeter Array (SMA). We also show interferometric observations of the CO(1--0) transition, with $uv$ sampling similar to that of the CO(3--2) data, from the Owens Valley Radio Observatory (OVRO) Millimeter Array. We describe our observations in Sect.~\\ref{sect-obs} and show the overall molecular gas distributions in Sect.~\\ref{sect-res-dist}. Line ratios, molecular gas masses, and kinematics are presented and discussed in Sect.~\\ref{sect-res-ratio}, \\ref{sect-res-mass}, and \\ref{sect-res-rot}, respectively. The relation between line ratios and star formation is discussed in Sect.~\\ref{sect-res-sfr}, and a summary is in Sect.~\\ref{sect-sum}. ",
        "conclusions": "\\label{sect-dis}  \\subsection{CO Line Ratios} \\label{sect-res-ratio}  Before deriving line ratios at various regions, we first matched the $uv$ range between our CO(3--2) and CO(1--0) datasets. In Figure~\\ref{fig3}, we overlay the CO(3--2) (solid contours) and CO(1--0) (dashed contours) intensity distributions. In this image, the shortest $uv$ length is set to 16~k$\\lambda$, and the image resolutions are convolved into the same resolution ($3\\farcs46\\times2\\farcs56$ with a P.A.\\ of $67\\fdg5$). As mentioned above, the overall distributions of both CO(3--2) and CO(1--0) are very similar, but some of the intensity peaks have shifted positions with respect to each other. The most obvious example is in the southern spiral arm, about $12\\arcsec-14\\arcsec$ south of the nucleus, where the CO(3--2) peak is located in the inner part of the spiral arm, but the CO(1--0) peak is shifted toward the outer edge of the spiral arm.  \\begin{figure} \\includegraphics[angle=-90,width=.9\\textwidth]{fig3.eps} \\caption{ Integrated intensity maps of CO(3--2) (solid contours) and CO(1--0) (dashed contours), overlaid on the continuum-subtracted \\emph{HST} F658N image (greyscale) of the central region of NGC~1068. Solid contour levels for CO(3--2) are 10, 30, 50, 70, 90, 110, and 130 $\\times$ 5.3~Jy~beam$^{-1}$~km~s$^{-1}$, and dashed contour levels for CO(1--0) are 2, 4, 6, 8 $\\times$ 5~Jy~beam$^{-1}$~km~s$^{-1}$. The CO(3--2) and CO(1--0) data are matched to the same $uv$ range and have the same synthesized beam size of $3\\farcs46\\times2\\farcs56$ ($\\sim242\\times179\\rm~pc^2$) with a P.A.\\ of $67\\fdg5$, which is shown in the bottom-left corner of the image. Cross is the same as in Figure~\\ref{fig1}. We also plot 25 boxes (C1 and R1 -- R24) that are used to calculate the line ratios. \\label{fig3}} \\end{figure}  We divide the central region of NGC~1068 into 25 regions covering the nucleus and spiral arms, as shown in Figure~\\ref{fig3} (C1 and R1 -- R24). The size of each region is $4\\arcsec\\times5\\arcsec$ except C1, which is $6\\arcsec\\times5\\arcsec$. The intensity scale is then converted to a brightness temperature scale, and the CO(3--2)/CO(1--0) brightness temperature ratios, $R_{31}$, are derived for each region using the MIRIAD task MATHS. Area-averaged spectral peak and integrated brightness temperatures and $R_{31}$ for each region are shown in Table~\\ref{tab-ratio}.  The central core region C1 exhibits a very high integrated intensity ratio with $R_{31}=3.12\\pm0.03$ and a spectra peak ratio of $2.83\\pm0.10$. Figure~\\ref{fig4} shows the spectra of C1. The CO(3--2)/CO(1--0) line ratio is slightly smaller than the results of \\citet{krip11}, which are about 4--6. However, the observations of \\citet{krip11} have a much higher angular resolution and might contain relatively more contribution from the circumnuclear region of the AGN, while our result is an average value over a much larger area and are likely to include emission from outside regions. In fact, if we only consider the line ratio within the brightest beam at the center, we would obtain a line ratio of $\\sim4.6$, which is similar to the results of \\citet{krip11}.  On the other hand, the spiral arm regions R1 -- R24 have a wide range of integrated intensity ratios of 0.24 -- 2.34 and an average value of 0.75 with standard deviation 0.47. We note that the wide range of the intensity ratios might be caused by the different spatial distributions of the CO(1--0) and CO(3--2) emission due to varying physical conditions of the molecular gas. There is no obvious ratio difference between the northern and southern arms; the average ratio of the northern arm regions (R1 -- R4, R18 -- R24) is 0.8 and that of the southern arm regions (R5 -- R17) is 0.7. However, there is a difference between inner arm regions and outer arm regions; the average ratio of the inner arm regions (R1 -- R4, R9 -- R17, and R19 -- R22) is $0.88\\pm0.14$, while the outer arm regions (R5 -- R8, R18, R23, and R24) show $0.42\\pm0.05$. This difference indicates that there is a large-scale gradual decrease of the line intensity ratio from the inner radii to the outer radii. The radial change of molecular line ratios, and therefore of the physical conditions of the molecular gas, has also been detected in the Milky Way and other galaxies \\citep{tur93,aal94,sak94,sak97,pet00,pag01,mat10}, suggesting that this trend is common.  \\begin{figure} \\includegraphics[angle=0,width=\\textwidth]{fig4.eps} \\caption{Area-averaged spectra within the C1 region. The solid line shows the spectrum of CO(3--2), and the dash line shows the spectrum of CO(1--0). The uncertainties per channel are about 0.01~K for the CO(3--2) spectrum and 0.02~K for the CO(1--0) one. \\label{fig4}} \\end{figure}  The line ratios $R_{31}$ in the spiral arm regions are similar to those observed in the centers of nearby normal and starburst galaxies. \\citet{dev94} observed seven nearby starburst galaxies and found that the ratios are in the range of 0.5 -- 1.4 with an average ratio of $0.64\\pm0.06$. A survey toward the centers of 28 nearby star-forming galaxies showed that most of their $R_{31}$ are within the range of 0.2 -- 0.7 \\citep{mau99}. On the other hand, the value of $R_{31}$ in the nucleus is much larger than those observed in the nuclei of the nearby galaxies mentioned above, albeit similar to that observed in the nucleus of the Seyfert 2 galaxy M51. The inner $uv$-truncated $R_{31}$ of the central molecular core of M51 also has a very high value of $\\sim5.2\\pm1.7$ within a beam size of $3\\farcs9\\times2\\farcs6$ or 160~pc $\\times$ 110~pc \\citep[missing flux corrected $R_{31}$ is $\\geq 1.9\\pm0.2$;][]{mat04}. If we only correct the 20\\% missing flux of the CO(1--0) emission in the core region of NGC~1068 and ignore the missing flux of the CO(3--2) following the method of \\citet{mat04}, the $R_{31}$ ratio of the core region will become 2.50 instead of 3.12. This suggests that the physical conditions of molecular gas around the Seyfert 2 AGN in NGC~1068 are very different from those in the centers of nearby star-forming and starburst galaxies, but might be similar to those around the Seyfert 2 nucleus of M51.   \\begin{deluxetable}{cccccccccc} \\rotate \\tabletypesize{\\scriptsize} \\tablecaption{Peak and integrated brightness temperature of CO(3--2) and CO(1--0), and the CO(3--2)/CO(1--0) line ratios, in different areas. \\label{tab-ratio}} \\tablehead{ \\colhead{Area} & \\colhead{Peak CO(3--2)} & \\colhead{Peak CO(1--0)} & \\colhead{Peak $R_{31}$} & \\colhead{$\\int T_{\\rm B}(\\rm{CO}$~$J = 3 - 2)~dv$} & \\colhead{$\\int T_{\\rm B}(\\rm{CO}$~$J = 1 - 0)~dv$} & \\colhead{$R_{31}$} & \\colhead{X offset} & \\colhead{Y offset} \\\\ & \\colhead{[K]} & \\colhead{[K]} & & \\colhead{[K~km~s$^{-1}$]} & \\colhead{[K~km~s$^{-1}$]} & & \\colhead{[arcsec]} & \\colhead{[arcsec]} } \\startdata C1 & $1.73\\pm0.01$ & $0.61\\pm0.03$ & $2.83\\pm0.10$ & $ 438.5\\pm 0.7$ & $ 140.6\\pm  1.5$ & $ 3.12\\pm0.03$ & -1 & 0\\\\ R1 & $1.80\\pm0.03$ & $1.55\\pm0.06$ & $1.16\\pm0.05$ & $ 103.76\\pm 0.97$ & $  44.3\\pm  1.8$ & $ 2.34\\pm0.10$ & 0 & 14\\\\ R2 & $1.53\\pm0.05$ & $1.90\\pm0.08$ & $0.80\\pm0.04$ & $ 104.97\\pm 1.54$ & $ 133.4\\pm  2.8$ & $ 0.79\\pm0.02$ & 4 & 14\\\\ R3 & $1.29\\pm0.06$ & $1.98\\pm0.05$ & $0.65\\pm0.03$ & $ 100.06\\pm 2.07$ & $ 149.8\\pm  1.8$ & $ 0.67\\pm0.02$ & 8 & 12\\\\ R4 & $0.43\\pm0.03$ & $0.73\\pm0.04$ & $0.58\\pm0.05$ & $   6.26\\pm 0.81$ & $  26.6\\pm  1.2$ & $ 0.24\\pm0.03$ & 12 & 9\\\\ R5 & $0.78\\pm0.04$ & $1.92\\pm0.11$ & $0.41\\pm0.03$ & $  56.74\\pm 1.39$ & $  90.0\\pm  3.8$ & $ 0.63\\pm0.03$ & 16 & 6\\\\ R6 & $1.54\\pm0.05$ & $3.65\\pm0.15$ & $0.42\\pm0.02$ & $  78.95\\pm 1.45$ & $ 131.1\\pm  4.4$ & $ 0.60\\pm0.02$ & 18 & 9\\\\ R7 & $0.71\\pm0.04$ & $1.97\\pm0.14$ & $0.36\\pm0.03$ & $  25.72\\pm 1.23$ & $  93.0\\pm  4.4$ & $ 0.28\\pm0.02$ & 18 & 14\\\\ R8 & $0.84\\pm0.03$ & $1.80\\pm0.08$ & $0.47\\pm0.03$ & $  28.85\\pm 1.09$ & $  94.4\\pm  2.5$ & $ 0.31\\pm0.01$ & 20 & 3\\\\ R9 & $1.35\\pm0.03$ & $3.63\\pm0.08$ & $0.37\\pm0.01$ & $ 108.83\\pm 1.11$ & $ 199.9\\pm  2.9$ & $ 0.54\\pm0.01$ & 16 & -1\\\\ R10 & $1.49\\pm0.07$ & $2.20\\pm0.08$ & $0.68\\pm0.04$ & $  62.03\\pm 1.92$ & $ 111.0\\pm  2.2$ & $ 0.56\\pm0.02$ & 12 & -5\\\\ R11 & $1.49\\pm0.07$ & $1.94\\pm0.03$ & $0.77\\pm0.04$ & $  94.26\\pm 2.63$ & $ 172.2\\pm  1.1$ & $ 0.55\\pm0.02$ & 8 & -8\\\\ R12 & $1.74\\pm0.06$ & $3.02\\pm0.04$ & $0.57\\pm0.02$ & $ 112.76\\pm 2.01$ & $ 176.3\\pm  1.4$ & $ 0.64\\pm0.01$ & 4 & -10\\\\ R13 & $1.93\\pm0.09$ & $1.66\\pm0.07$ & $1.16\\pm0.07$ & $ 194.32\\pm 3.64$ & $ 139.4\\pm  2.8$ & $ 1.39\\pm0.04$ & 0 & -12\\\\ R14 & $2.46\\pm0.09$ & $3.40\\pm0.03$ & $0.72\\pm0.03$ & $ 252.08\\pm 3.58$ & $ 192.5\\pm  1.3$ & $ 1.31\\pm0.02$ & -4 & -13\\\\ R15 & $2.68\\pm0.08$ & $3.18\\pm0.09$ & $0.84\\pm0.04$ & $ 239.52\\pm 3.24$ & $ 374.4\\pm  3.4$ & $ 0.64\\pm0.01$ & -8 & -11\\\\ R16 & $3.46\\pm0.11$ & $4.80\\pm0.18$ & $0.72\\pm0.04$ & $ 211.39\\pm 3.69$ & $ 301.1\\pm  6.0$ & $ 0.70\\pm0.02$ & -12 & -8\\\\ R17 & $1.67\\pm0.01$ & $1.68\\pm0.06$ & $0.99\\pm0.04$ & $ 106.98\\pm 0.22$ & $ 141.2\\pm  2.3$ & $ 0.76\\pm0.01$ & -16 & -5\\\\ R18 & $0.79\\pm0.01$ & $2.47\\pm0.07$ & $0.32\\pm0.01$ & $  56.68\\pm 0.62$ & $ 159.4\\pm  2.8$ & $ 0.36\\pm0.01$ & -20 & 1\\\\ R19 & $1.28\\pm0.05$ & $1.72\\pm0.08$ & $0.75\\pm0.04$ & $  49.41\\pm 1.41$ & $  75.5\\pm  2.4$ & $ 0.65\\pm0.03$ & -16 & 6\\\\ R20 & $1.94\\pm0.07$ & $3.19\\pm0.06$ & $0.61\\pm0.02$ & $ 115.44\\pm 2.06$ & $ 168.9\\pm  1.8$ & $ 0.68\\pm0.01$ & -12 & 10\\\\ R21 & $1.83\\pm0.05$ & $1.66\\pm0.06$ & $1.10\\pm0.05$ & $  97.04\\pm 1.48$ & $  81.3\\pm  1.8$ & $ 1.19\\pm0.03$ & -8 & 13\\\\ R22 & $1.10\\pm0.04$ & $0.47\\pm0.04$ & $2.32\\pm0.22$ & $  62.59\\pm 1.26$ & $  46.3\\pm  1.3$ & $ 1.35\\pm0.05$ & -4 & 14\\\\ R23 & $1.02\\pm0.05$ & $2.72\\pm0.07$ & $0.38\\pm0.02$ & $  45.12\\pm 1.70$ & $  94.9\\pm  2.1$ & $ 0.48\\pm0.02$ & -24 & -4\\\\ R24 & $0.96\\pm0.03$ & $2.95\\pm0.16$ & $0.33\\pm0.02$ & $  39.06\\pm 0.86$ & $ 126.5\\pm  4.5$ & $ 0.31\\pm0.01$ & -24 & -9\\\\ \\enddata \\tablecomments{Column (1): Region name. Column (2): Peak brightness temperature of CO(3--2). Column (3): Peak brightness temperature of CO(1--0). Column (4): Line ratio of peak brightness temperature of CO(3--2)/CO(1--0). Column (5): Integrated brightness temperature of CO(3--2). Column (6): Integrated brightness temperature of CO(1--0). Column (7): Line ratio of integrated brightness temperature of CO(3--2)/CO(1--0). The errors Column (8,9): Position offset of region center from the phase center.  Errors were estimated from the statistical errors of the spectra; uncertainties of our flux calibration (about 10\\% for both CO(1--0) and CO(3--2) lines) and missing fluxes are not included.}  \\end{deluxetable}  \\subsection{Molecular Gas Column Density and Mass} \\label{sect-res-mass}  Because CO(1--0) traces the bulk of the molecular gas, we can calculate the molecular gas column density and mass from the CO(1--0) integrated intensities. The column density of the molecular hydrogen can be estimated with a conversion factor, $X_{\\rm CO}$, \\begin{equation} N_{\\rm H_2} = X_{\\rm CO} \\int T_{\\rm B}({\\rm CO}~J = 1 - 0)~dv ~{\\rm [cm^{-2}]}, \\end{equation} where $X_{\\rm CO}\\cong0.4\\times10^{20}$~cm$^{-2}$~(K km s$^{-1}$)$^{-1}$ for circumnuclear molecular gas \\citep{wil95,mau96,weis01,esp09} and $X_{\\rm CO}\\cong3\\times10^{20}$~cm$^{-2}$~(K~km~s$^{-1}$)$^{-1}$ for spiral arm regions \\citep{sol87,you91}. The column density $N_{\\rm H_2}$ for each region is calculated using $\\int T_{\\rm B}({\\rm CO}~J = 1 - 0)~dv$ shown in Table~\\ref{tab-ratio}, and the derived values are shown in Table~\\ref{tab-sfr}. We also derive the molecular gas mass, $M_{\\rm H_2}$, which is displayed in Table~\\ref{tab-sfr}. The total missing flux of CO(1--0) within the central circular area of $30\\arcsec$ radius is about 35\\% \\citep{pap99}. However, the missing flux is dominated by extended structures, so the true missing fluxes in the core and the compact spiral structures should be smaller. For example, the missing flux of CO(1--0) in the core region is about 20\\% as discussed in Sect.~\\ref{sect-res-dist}. We thus expect that the errors of the column densities caused by the missing flux are less than 35\\%.  \\begin{deluxetable}{ccccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{Molecular gas column density and mass, $8~\\mu$m dust and H$\\alpha$ + [N II] emission, and star formation rate surface density. \\label{tab-sfr}} \\tablehead{ \\colhead{Area} & \\colhead{$N_{\\rm H_2}$} & \\colhead{$M_{\\rm H_2}$} & \\colhead{$f_{\\rm 8\\micron, dust}$} & \\colhead{$f_{\\rm H\\alpha+[N II]}$} & \\colhead{$\\Sigma_{\\rm SFR}$} \\\\ & \\colhead{$\\times10^{21}$ [cm$^{-2}$]} & \\colhead{$\\times10^{6}$ [$M_\\odot$]} & \\colhead{[MJy~sr$^{-1}$]} & \\colhead{$\\times10^{4}$ [eps~pixel$^{-1}$]\\tablenotemark{a}} & \\colhead{[$M_\\odot$~yr$^{-1}$~kpc$^{-2}$]} } \\startdata C1 & $  5.6\\pm0.3$ & $ 12.3\\pm0.8$ & $1.05\\times10^{4}\\pm6.4$ & 73.5 & $62.4\\pm0.04$\\\\ R1 & $  13.3\\pm0.5$ & $ 19.5\\pm0.8$ & $238\\pm6.4$ & 1.51 & $1.41\\pm0.04$\\\\ R2 & $  40.0\\pm0.8$ & $ 58.6\\pm1.2$ & $260\\pm6.4$ & $>1.87$\\tablenotemark{b} & $1.54\\pm0.04$\\\\ R3 & $  44.9\\pm0.5$ & $ 65.8\\pm0.8$ & $337\\pm6.4$ & $>2.57$\\tablenotemark{b} & $2.00\\pm0.04$\\\\ R4 & $   8.0\\pm0.4$ & $ 11.7\\pm0.5$ & $241\\pm6.4$ & 2.35 & $1.43\\pm0.04$\\\\ R5 & $  27.0\\pm1.1$ & $ 39.5\\pm1.7$ & $180\\pm6.4$ & 1.08 & $1.07\\pm0.04$\\\\ R6 & $  39.3\\pm1.3$ & $ 57.6\\pm1.9$ & $166\\pm6.4$ & $>0.81$\\tablenotemark{b} & $0.98\\pm0.04$\\\\ R7 & $  27.9\\pm1.3$ & $ 40.9\\pm1.9$ & $107\\pm6.4$ &   ---\\tablenotemark{b} & $0.63\\pm0.04$\\\\ R8 & $  28.3\\pm0.8$ & $ 41.5\\pm1.1$ & $ 89\\pm6.4$ & $>0.72$\\tablenotemark{b} & $0.52\\pm0.04$\\\\ R9 & $  60.0\\pm0.9$ & $ 87.9\\pm1.3$ & $163\\pm6.4$ & $>0.65$\\tablenotemark{b} & $0.97\\pm0.04$\\\\ R10 & $  33.3\\pm0.7$ & $ 48.8\\pm1.0$ & $158\\pm6.4$ & 0.77 & $0.94\\pm0.04$\\\\ R11 & $  51.7\\pm0.3$ & $ 75.7\\pm0.5$ & $174\\pm6.4$ & 0.89 & $1.03\\pm0.04$\\\\ R12 & $  52.9\\pm0.4$ & $ 77.5\\pm0.6$ & $226\\pm6.4$ & 0.97 & $1.34\\pm0.04$\\\\ R13 & $  41.8\\pm0.8$ & $ 61.3\\pm1.2$ & $212\\pm6.4$ & 0.76 & $1.26\\pm0.04$\\\\ R14 & $  57.8\\pm0.4$ & $ 84.6\\pm0.6$ & $236\\pm6.4$ & 1.21 & $1.40\\pm0.04$\\\\ R15 & $ 112.3\\pm1.0$ & $164.5\\pm1.5$ & $311\\pm6.4$ & 2.17 & $1.84\\pm0.04$\\\\ R16 & $  90.3\\pm1.8$ & $132.3\\pm2.6$ & $381\\pm6.4$ & 3.49 & $2.26\\pm0.04$\\\\ R17 & $  42.4\\pm0.7$ & $ 62.1\\pm1.0$ & $345\\pm6.4$ & 2.81 & $2.05\\pm0.04$\\\\ R18 & $  47.8\\pm0.8$ & $ 70.0\\pm1.2$ & $172\\pm6.4$ & 0.62 & $1.02\\pm0.04$\\\\ R19 & $  22.7\\pm0.7$ & $ 33.2\\pm1.0$ & $185\\pm6.4$ & 1.15 & $1.10\\pm0.04$\\\\ R20 & $  50.7\\pm0.6$ & $ 74.2\\pm0.8$ & $207\\pm6.4$ & 1.45 & $1.23\\pm0.04$\\\\ R21 & $  24.4\\pm0.5$ & $ 35.7\\pm0.8$ & $260\\pm6.4$ & 1.33 & $1.54\\pm0.04$\\\\ R22 & $  13.9\\pm0.4$ & $ 20.4\\pm0.6$ & $231\\pm6.4$ & 1.37 & $1.37\\pm0.04$\\\\ R23 & $  28.5\\pm0.6$ & $ 41.7\\pm0.9$ & $123\\pm6.4$ & 0.27 & $0.73\\pm0.04$\\\\ R24 & $  38.0\\pm1.4$ & $ 55.6\\pm2.0$ & $105\\pm6.4$ & 0.15 & $0.62\\pm0.04$\\\\ \\enddata \\tablecomments{Column (1): Region names. Column (2): Column density. Column (3): Molecular gas mass. Column (4): Average intensity of dust emission. Column (5): Average intensity of $\\rm H\\alpha+[N II]$ emission. Column (6): Star formation rate surface density.} \\tablenotetext{a}{Image units in electrons per second per pixels (eps~pixel$^{-1}$)} \\tablenotetext{b}{These areas do not have complete H$\\alpha$ emission information, since these regions are either located at the edge of or outside the \\emph{HST} image.} \\end{deluxetable}  In the spiral arm regions, the box sizes correspond to $280\\times350~$pc$^2$. The average values of $N_{\\rm H_2}$ and $M_{\\rm H_2}$ per region are $39.8\\times10^{21}$~cm$^{-2}$ and $58.9\\times10^6M_{\\odot}$. The standard deviations of $N_{\\rm H_2}$ and $M_{\\rm H_2}$ are large inside the spiral arms, with a value of $24.0\\times10^{21}$~cm$^{-2}$ and $34.8\\times10^6M_{\\odot}$, respectively, indicating that the molecular gas content within the spiral arms varies substantially from region to region. On the other hand, $N_{\\rm H_2}$ and $M_{\\rm H_2}$ for the central core C1 are estimated to be $5.6\\times10^{21}$~cm$^{-2}$ and $12.3\\times10^6$~M$_{\\odot}$ (note that the area of C1 is $\\sim$1.5 times larger than those in the spiral arms). This indicates that the column density and the mass of the central region are much smaller than the average values in the spiral arms. However, we note that the $R_{31}$ of C1 is very different from those of the spiral arm regions (see Sect.~\\ref{sect-res-ratio}), and we have used different conversion factors in estimating the mass and column density \\citep{weis01}.  \\citet{tac94} and \\citet{stern94} showed that the $N_{\\rm H_2}$ column density in the core region is about $4\\times10^{22}$~ cm$^{-2}$ using the Galactic CO to $N_{\\rm H_2}$ conversion factor. Our conversion factor for the core region is about $\\frac{1}{6}$ of the Galactic conversion factor and we obtained a $N_{\\rm H_2}$ column density of $5.6\\times10^{21}$~cm$^{-2}$; in other words, the different results are mainly caused by the different conversion factors adopted. However, we note that a lower conversion factor is typically found in galaxy centers \\citep{wil95,mau96,weis01,esp09}. In particular, the conversion factor in the center of NGC 1068 could be six times lower than the Galactic value \\citep{use04,gar10}.  In Figure~\\ref{fig5}, we compare our CO(3--2) image with the X-ray image obtained with the \\emph{Chandra X-ray Observatory} \\citep{you01}. The X-ray image displays clear emission from the ionization cone emanating from the nucleus of NGC~1068; however, we note that there is an obvious dimmed area at the center of the image. This dimmed area matches well with the central core of our CO(3--2) image, suggesting that the obscuring material of the X-rays is closely related to the molecular gas.  \\begin{figure} \\includegraphics[angle=-90,width=.9\\textwidth]{fig5.eps} \\caption{Comparison between our CO(3--2) image and the X-ray image taken with the \\emph{Chandra X-ray Observatory} \\citep{you01}. The contours show CO(3--2) emission, with levels 10, 20, 30, 40, 50, 70, 90, 110, and $130\\times4.41$~Jy~beam$^{-1}$~km~s$^{-1}$. The gray scale shows the X-ray emission from 0 to 1000 counts pixel$^{-1}$($0.4-5.0$~keV). \\label{fig5}} \\end{figure}   \\subsection{CO(3--2) Rotation Curve} \\label{sect-res-rot}  The rotation curve in the nucleus of NGC~1068 is plotted in Figure~\\ref{fig6}. The data points and the error bars were derived using the MIRIAD task VELFIT. We ignore points within $30\\degr$ of the minor axis (P.A.~$=60\\degr$), and we assume the inclination angle to be $45\\degr$ following \\citet{sch00}. Within the central $2\\arcsec$ ($\\sim140$~pc), the rotational velocity increases with radius (rigid rotation). The enclosed mass within the central $2\\arcsec$ can be estimated using $M(r<2\\arcsec) \\sim rv^{2}/G$, and we find that the total mass within this area is $3.7\\times10^{8}$~M$_{\\odot}$, which is consistent with \\citet{sch00}. The molecular gas mass in C1 is about $12.3\\times10^6$~M$_{\\odot}$, so that the gas-to-dynamical mass ratio is 3\\%. This value is a factor of a few lower than those in star forming galaxies ($\\sim10\\%$) \\citep{sak99,kod05}.  For radii larger than $2\\arcsec$, the rotation curve becomes slowly decreasing. This rotation curve is well fitted with the Brandt rotation curve \\citep{bra60}: \\begin{equation} V = \\frac{V_{\\rm max}\\frac{R}{R_{\\rm max}}} {\\left(\\frac{1}{3}+\\frac{2}{3} \\left(\\frac{R}{R_{\\rm max}}\\right)^{n}\\right)^{\\frac{3}{2n}}} \\end{equation} with $V_{\\rm max}=116$~km~s$^{-1}$, $R_{\\rm max}=2\\farcs7$, and $n=6$ (Figure~\\ref{fig6}). The mass from the Brandt model is: \\begin{equation} M_{\\rm tot} = \\left(\\frac{3}{2}\\right)^{\\frac{3}{n}} \\frac{V^2_{\\rm max}R_{\\rm max}}{G}. \\end{equation} The total mass inside 190~pc ($\\sim2.7\\arcsec$) radius is $7.2\\times10^{8}$~M$_{\\odot}$.  \\begin{figure} \\includegraphics[angle=0,width=.9\\textwidth]{fig6.eps} \\caption{Rotation curve of CO(3--2) in the central region of NGC~1068. Crosses with error bars are the data points, and the solid curve is the fitted Brandt rotation curve (see the main text for details). \\label{fig6}} \\end{figure}    \\subsection{Relation between CO Line Ratio and Star Formation} \\label{sect-res-sfr}  As shown in Sect.~\\ref{sect-res-ratio}, $R_{31}$ varies along the spiral arms. Since the CO(3--2) line is more closely related to star formation \\citep{kom07}, and $R_{31}$ increases with increasing star formation efficiency \\citep{mur07}, the variation in $R_{31}$ along the spiral arms of NGC~1068 may also be related to star formation. To study the cause of the variation, we compare $R_{31}$ with star formation surface density along the spiral arms.  We use both infrared and optical observations as star formation tracers. The infrared data are obtained from the \\emph{Spitzer Space Telescope} IRAC $3.6~\\micron$ and $8~\\micron$ images\\footnote{\\url{http://ssc.spitzer.caltech.edu/archanaly/archive.html}}. \\citet{wu05} showed that dust emission at $8~\\micron$ can be used as a star formation indicator. We estimate the dust emission of in the spiral arms of NGC~1068 using the observed \\emph{Spitzer} IRAC $3.6~\\micron$ and $8~\\micron$ fluxes: \\begin{equation} f_{8\\micron}(\\mathrm{dust}) = f_{8\\micron} - \\eta_{8} f_{3.6\\micron}, \\end{equation} where $\\eta_{8}=0.232$ \\citep{lei99,hel04}. The derived values are shown in Table~\\ref{tab-sfr}.  H$\\alpha$ emission is also considered as a tracer of star formation. For NGC~1068, the wavelength of H$\\alpha$ is shifted to 6587.8~\\AA~ and those of [\\ion{N}{2}] to 6573, 6603\\AA. To obtain the H$\\alpha$ emission, we retrieved the \\emph{Hubble Space Telescope (HST)} F658N and F791W images from the STScI archive\\footnote{\\url{http://archive.stsci.edu}}. The F658N narrow-band filter has a central wavelength of 6590.8~\\AA~ with a bandwidth of $\\sim$30~\\AA; therefore, the F658N image includes the H$\\alpha$ line as well as [\\ion{N}{2}] lines, which are shifted to 6573 and 6603\\AA for NGC~1068. Since this image also includes continuum emission, we correct it using the nearby broad-band filter F791W. The F791W filter has a central wavelength of 7881~\\AA~ with a bandwidth of $\\sim$1231~\\AA. The continuum subtraction can be performed using the following equation: \\begin{equation} f_{\\rm line} = \\frac{f_{\\rm NB}\\Delta\\lambda_{\\rm BB} - f_{\\rm BB}\\Delta\\lambda_{\\rm NB}} {\\Delta\\lambda_{\\rm BB} - \\Delta\\lambda_{\\rm NB}}, \\end{equation} where $f_{\\rm line}$ is the continuum-subtracted line flux, $f_{\\rm NB}$ and $f_{\\rm BB}$ are the observed total fluxes in the narrow-band and broad-band filters, respectively, and $\\Delta\\lambda_{\\rm NB}$ and $\\Delta\\lambda_{\\rm BB}$ are the bandwidths of the narrow-band and broad-band filters, respectively. The resulting continuum-subtracted F658N image can be considered as the H$\\alpha$ + [\\ion{N}{2}] image. The derived values are shown in Table~\\ref{tab-sfr}. We cannot separate the H$\\alpha$ from the [\\ion{N}{2}] lines and we cannot correct the H$\\alpha$ + [\\ion{N}{2}] image for internal extinction with available data. Therefore, we can not obtain any quantitive information from this image; we only use the continuum-subtracted emission as a possible indicator of relative star formation rates.  We compare the \\emph{HST} H$\\alpha$ + [\\ion{N}{2}] line intensity with the \\emph{Spitzer} $8~\\micron$ dust intensity.  As can be seen in Figure~\\ref{fig7}, the H$\\alpha$ + [\\ion{N}{2}] line intensity and the $8~\\micron$ dust intensity are linearly correlated, indicating that the former image is effectively also a good star formation tracer. We note that the CO distribution also matches very well with the dust lanes along the spiral arms in the \\emph{HST} H$\\alpha$ + [\\ion{N}{2}] line image (Figure~\\ref{fig3}). The star forming regions are also located along the spiral arms but slightly shifted toward the outside of the spiral arms, especially in the southern spiral arm. This is similar to the results of previous observations for the spiral arms of the nearby galaxy M51 \\citep{vog88,aal99,kod09}.  \\begin{figure} \\includegraphics[angle=0,width=.9\\textwidth]{fig7.eps} \\caption{Comparison of the \\emph{Spitzer} $8~\\micron$ dust intensity with the \\emph{HST} continuum-subtracted H$\\alpha$ + [\\ion{N}{2}] line intensity. The \\emph{HST} line intensity is not flux-calibrated and is expressed in instrument units. The data points with upper arrows do not have complete H$\\alpha$ emission information because they are located either at the edge of or outside the \\emph{HST} image. \\label{fig7}} \\end{figure}   \\begin{figure} \\includegraphics[angle=0,width=.9\\textwidth]{fig8.eps} \\caption{Comparison between the CO(3--2)/CO(1--0) integrated intensity ratio, $R_{31}$, and the star formation rate surface density ($\\Sigma_{\\rm SFR}$). The error bars have included the $\\sim~10\\%$ uncertainty of our flux calibrations for the CO(3--2) and CO(1--0) emission. \\label{fig8}} \\end{figure}  \\begin{figure} \\includegraphics[angle=90,width=.9\\textwidth]{fig9.eps} \\caption{Relation between the surface density of molecular gas $\\Sigma_{\\rm gas}$ and $\\Sigma_{\\rm SFR}$ for normal galaxies (open squares), starburst galaxies (open circles), and the spiral arm regions of NGC~1068 (filled circles). The gas surface density is derived from CO(1--0) emission. The normal and starburst galaxy samples are obtained from \\citet{ken98}. The solid triangle is the nuclear region of NGC 1068 (C1). The solid line represents the power-law fit to all data points. \\label{fig9}} \\end{figure}   \\begin{figure} \\includegraphics[angle=0,width=.9\\textwidth]{fig10.eps} \\caption{Relation between integrated CO(3--2) and FIR luminosities. The filled circles represent the data of the spiral arm regions of NGC~1068, and the solid triangle is the nuclear region of NGC 1068 (C1). The open circles are the data obtained from \\citet{mao10}. The solid line represents the linear fit to all data points. The outlier point with the lowest CO(3--2) emission is R4. \\label{fig10}} \\end{figure}   We also compare molecular gas properties with the star formation rate surface density ($\\Sigma_{\\rm SFR}$) derived from the \\emph{Spitzer} IRAC $8~\\micron$ and $3.6~\\micron$ data. The $\\Sigma_{\\rm SFR}$ is derived from the $8~\\micron$ dust luminosity: \\begin{equation} \\frac{\\Sigma_{\\rm SFR}}{(M_\\odot~{\\rm yr^{-1}})} = \\frac{\\nu L_{\\nu}[8\\micron(dust)]}{1.57\\times10^9~{L_{\\odot}}} \\end{equation} \\citep{wu05}, where $L_{\\nu}$ is the $8~\\micron$ dust luminosity. The derived values are shown in Table~\\ref{tab-sfr}.  Figure~\\ref{fig8} shows the comparison between $R_{31}$ and $\\Sigma_{\\rm SFR}$ of the spiral arm regions. The correlation coefficient $r$ between the $R_{31}$ and $\\Sigma_{\\rm SFR}$ is 0.33 with a probability $p=89\\%$ of mutual correlation. We note that most of the $R_{31}$ of nearby star-forming galaxies are within 0.2--0.7 \\citep{mau99}. For examples, \\citet{mur07} found that the CO emission of the starburst galaxy M83 has $R_{31} < 1$ and show a good correlation between the $R_{31}$ and SFE of their data. If we consider the data with $R_{31} < 1$ in our results, we find that the correlation between the $R_{31}$ and $\\Sigma_{\\rm SFR}$ becomes highly significant with $r=0.63$ and $p=99\\%$. This suggests that in normal star-forming regions the physical conditions of molecular gas indicated by $R_{31}$ are well correlated with dust emission. The data points where $R_{31} > 1.0$ are obviously out of the correlation, suggesting that the large $R_{31}$ might be caused by other reasons, such as a different heating mechanism for the molecular gas or different distributions of warm and cool molecular gas \\citep[e.g.,][]{ho87}.  Figure~\\ref{fig9} shows the relation between the surface density of molecular gas $\\Sigma_{\\rm gas}$ and $\\Sigma_{\\rm SFR}$. The star formation rate surface densities of the spiral arm regions of NGC 1068 are much higher than those of normal galaxies and similar to those of starburst galaxies; however, they all seem to follow the same star formation law \\citep{ken98}. This result strongly supports the idea that the spiral arms in the inner $\\sim2$ kpc region are experiencing a starburst.  On the other hand, the molecular gas at C1 is obviously offset from the Schmidt-Kennicutt law, suggesting that C1 is mainly affected by AGN activities instead of star formation.  In Figure~\\ref{fig10}, we present the Schmidt-Kennicutt law for the CO(3--2) emission. We used the FIR luminosity instead of the SFR in Figure~\\ref{fig9} so that it is easier to compare with the results of \\citet{mao10}. We first derive the star formation rate from the observed $8\\micron$ dust luminosity \\citep{wu05} and then use the SFR--$L_{FIR}$ relation, SFR$(M_\\odot yr^{-1})\\sim1.7\\times10^{-10}(L_{FIR}/L_\\odot)$ \\citep{ken98}, to derive $L_{FIR}$. We note that this is effectively a Schmidt-Kennicutt law for CO(3--2) emission since $L_{FIR}$ is proportional to SFR, and the integrated CO(3--2) luminosity represents the molecular mass in relatively warm and/or dense regions. We find that most of our results follow the same relation obtained by \\citet{mao10} for nearby galaxies. The only outlier point, which has the lowest CO(3--2) emission, is the interarm region R4. When we combine our results with \\citet{mao10}, we find that the derived power-law index of the $L_{CO(3-2)}$--$L_{FIR}$ relation is $\\sim0.9$. This value is smaller than the index of the traditional Schmidt-Kennicutt law, which is around 1.0 to 2.0, but is reasonable for the excitation conditions of warm and/or dense gas \\citep{kt07,na08}.  In Figure~\\ref{fig9} and Figure~\\ref{fig10}, the physical scale of the NGC~1068 data is about $280\\times350~$pc$^2$, which is much smaller than those of \\citet{ken98} and \\citet{mao10}. We note that there is no correlation between $\\Sigma_{\\rm gas}$ and $\\Sigma_{\\rm SFR}$ in Figure~\\ref{fig9} when only considering the NGC~1068 data. On the other hand, the CO(3--2) emission from the same NGC~1068 regions show a very good correlation with the infrared as shown in Figure~\\ref{fig10}; the correlation coefficient $r$ is 0.789 with a probability $p = 99.9\\%$. The power-law index of the $L_{CO(3-2)}$--$L_{FIR}$ relation for the NGC~1068 data alone (excluding R4) is about 0.49. This value is significantly lower than a typical power-law index of the Schmidt-Kennicutt law.  In Figure~\\ref{fig10}, the CO(3-2) to FIR luminosity distribution is generally consistent with previous studies, which have a power law index of the Schmidt-Kennicutt law $\\sim$1.0. On the other hand, the Schmidt-Kennicutt law derived from the CO(1-0) data has a steeper slope than that from the CO(3-2) data (Figure~\\ref{fig9}). Furtheremore, if we only consider our NGC 1068 data, we find a very flat power law index, which cannot be explained by the model of \\citet{kt07} with a different critical density. One possibility is that the gas is under sub-thermal conditions, which would produce a flatter KS law as shown by \\citet{na08}. This interpretation is also consistent with the relatively large line ratios of CO(3--2)/CO(1--0). We also note that most of our data are under the average value of the KS law, indicating that the SFE of the inner spiral regions of NGC 1068 is smaller than that in most of the sources in Figure~\\ref{fig10}.  There is also a radial variation of star formation activity in the spiral arm regions. The average star formation rate surface density of the inner arm regions (R1 -- R4, R9 -- R17, and R19 -- R22) is 1.45 $M_{\\odot}$ yr$^{-1}$ kpc$^{-2}$, and the outer arm regions (R5 -- R8, R18, R23, and R24) is 0.80 $M_{\\odot}$ yr$^{-1}$ kpc$^{-2}$. In other words, the radial variation of the physical conditions of the molecular gas mentioned in Sect.~\\ref{sect-res-ratio} is correlated with the radial variation of the galaxy's star formation."
    },
    "1112/1112.5207_arXiv.txt": {
        "abstract": "Theoretical models predict that some of the first stars ended their lives as extremely energetic pair-instability supernovae (PISNe). With energies approaching $10^{53}$ ergs, these supernovae are expected to be within the detection limits of the upcoming {\\it James Webb Space Telescope} (JWST), allowing observational constraints to be placed on the properties of the first stars.  We estimate the source density of PISNe using a semi-analytic halo mass function based approach, accounting for the effects of feedback from star formation on the PISN rate using cosmological simulations. We estimate an upper limit of $\\sim$0.2 PISNe per JWST field of view at any given time.  Feedback can reduce this rate significantly, e.g., lowering it to as little as one PISN per 4000 JWST fields of view for the most pessimistic explosion models.  We also find that the main obstacle to observing PISNe from the first stars is their scarcity, not their faintness; exposures longer than a few times $10^4\\,$s will do little to increase the number of PISNe found. Given this we suggest a mosaic style search strategy for detecting PISNe from the first stars. Even rather high redshift PISNe are unlikely to be missed by moderate exposures, and a large number of pointings will be required to ensure a detection. ",
        "introduction": "Understanding the formation of the first stars and galaxies is one of the central challenges of modern cosmology, as they mark a significant increase in complexity from the simple initial conditions of the primordial universe during the `dark ages' (e.g., \\citealt{BarkanaLoeb2001, Miralda-Escude2003, Brommetal2009, Loeb2010}). The basic properties of these so-called Population III (Pop~III) stars have been reasonably well established, with the consensus that the first stars formed in dark matter `minihalos' on the order of $10^5$--$10^6$\\msun at high redshifts \\citep{CouchmanRees1986,HaimanThoulLoeb1996,Tegmarketal1997}. Numerical simulations of the collapse of primordial metal-free gas into these halos, where molecular hydrogen is the only available coolant, had suggested that the first stars were predominantly very massive, with $M_* \\gtrsim100$\\msun and a top-heavy initial mass function (IMF) (e.g., \\citealt{BrommCoppiLarson1999, BrommCoppiLarson2002, AbelBryanNorman2002, BrommLarson2004, Yoshidaetal2006, O'SheaNorman2007}).  More recent work has found that the gas from which the first stars formed underwent significant fragmentation \\citep{StacyGreifBromm2010, Clarketal2011b, Greifetal2011, Greifetal2012}, and experienced strong protostellar feedback \\citep{Hosokawaetal2011, StacyGreifBromm2012}.  These results have revised our picture of Pop~III star formation, with lower characteristic masses (on the order of 50 rather than $100$\\msun) and a much broader IMF now expected.  The light from these stars ended the dark ages and fundamentally transformed the universe, beginning both the reionization (e.g, \\citealt{Meiksin2009}) and the chemical enrichment of the universe (e.g, \\citealt{Karlsson2011}).  Given the top-heavy nature of Pop~III star formation and the fact that low-metallicity stars are unlikely to undergo significant radiatively driven mass loss \\citep{Kudritzki2002}, one interesting possibility is that some of the first stars died as pair-instability supernovae (PISNe), a scenario that has significant consequences for the chemical enrichment history of the universe \\citep{HegerWoosley2002, TumlinsonVenkatesanShull2004, KarlssonJohnsonBromm2008}.  Basic one-dimensional models predict that stars with masses in the range 140--260\\msun will undergo a pair-production instability and explode completely \\citep{BarkatRakavySack1967, Fraley1968}.  During core oxygen burning, a combination of high temperatures and relatively low densities results in the formation of $e^{\\pm}$ pairs, removing pressure support from the core. Following the subsequent contraction, the ignition of explosive oxygen burning completely disrupts the progenitor, resulting in a significant contribution to the metal enrichment of the surrounding medium. More recently, \\citet{ChatzopoulosWheeler2012} and \\citet{YoonDierksLanger2012} have found that stars with initial masses as low as 65\\msun can encounter the pair-production instability if they are rapidly rotating.  In this scenario, strong rotationally induced mixing causes nearly homogeneous evolution, such that the star is converted almost entirely to helium before the next phase in its evolution begins.  These extremely energetic explosions|approaching $10^{53}$ ergs for the most massive models|are very luminous, in part due to the large amount of $^{56}$Ni produced, and are also very temporally extended as a result of the large mass ejected \\citep{FryerWoosleyHeger2001, HegerWoosley2002, Hegeretal2003, JoggerstWhalen2011, KasenWoosleyHeger2011}.  A second possiblility for reaching such extreme explosion energies in slightly lower mass stars is the hypernova scenario for rapidly rotating stars that undergo core collapse \\citep{UmedaNomoto2003, TominagaUmedaNomoto2007}.  During the collapse, accretion onto a central black hole powers a jet which induces a highly energetic explosion. While recent work has decreased the expected mass of the first stars, they have also been found to rotate more rapidly than previously thought \\citep{StacyBrommLoeb2011}, increasing the plausibility of this scenario. Evidence supporting this hypothesis has recently been presented by \\citet{Chiappinietal2011}.  While not an example of a Pop~III star, the recent discovery of the extremely luminous supernova (SN) 2007bi, identified as a possible PISN, in a metal-poor dwarf galaxy at a redshift of $z\\simeq0.1$ \\citep{Gal-Yametal2009} suggests that PISNe may be possible in the local universe under rare circumstances. Further supporting this picture, \\citet{WoosleyBlinnikovHeger2007} have shown that SN 2006gy \\citep{Smithetal2007} is well modeled by a pulsational pair-instability model. For stars with initial masses in the range $\\sim$100--140\\msun, the star encounters the $e^{\\pm}$ production instability, but the resulting explosive ignition of oxygen is insufficient to unbind the star.  Instead it ejects a shell of material before settling back into a stable configuration.  The star encounters this instability several times until the mass of the helium core drops below $\\sim$40\\msun, after which the star can proceed to silicon burning and eventually undergo core-collapse.  With the upcoming launch of the {\\it James Webb Space Telescope} (JWST) we will be able to probe the epoch of first light in unprecedented detail.  While the first stars themselves are unlikely to be visible (e.g, \\citealt{BrommKudritzkiLoeb2001, PawlikMilosavljevicBromm2011}), some of the SNe that end their lives should be within the detection limits of the JWST (e.g., \\citealt{MackeyBrommHernquist2003, Scannapiecoetal2005, Gardneretal2006}). While the basic properties of PISNe and the effect they have on their environment has been well studied \\citep{MoriFerraraMadau2002, BrommYoshidaHernquist2003, FurlanettoLoeb2003, KitayamaYoshida2005, Whalenetal2008, WiseAbel2008, Greifetal2010}, the source density of these events has yet to be well constrained.  The first attempt at estimating the number and observability of SNe at high redshift was made by \\citet{Miralda-EscudeRees1997}, who calculated the all-sky SN rate based on estimates of the total metals produced by a typical SN and the observed metallicity of the intergalactic medium (IGM) at high redshifts.  This yields $\\sim$1 SN yr$^{-1}$ arcmin$^{-2}$ at $z \\sim 5$. Other early work attempted to model the SN rate based on the empirically determined star formation rate out to high redshifts ($z \\sim 5$; e.g., \\citealt{MadauDellaVallePanagia1998, DahlenFransson1999}).  It should be noted that these attempts focused on Type II SNe, not PISNe, which are the focus of the present work. \\citet{MackeyBrommHernquist2003} estimated the PISN rate based on their calculations of the Pop~III star formation rate, predicting $\\sim$$2\\times10^6$ PISNe yr$^{-1}$ over the whole sky above $z=15$.  \\citet{WeinmannLilly2005} performed a similar analysis with more conservative estimates for the star formation rate, finding a PISN rate of $\\sim$4 yr$^{-1}$ deg$^{-2}$ above $z=15$ and $\\sim$0.2 yr$^{-1}$ deg$^{-2}$ above $z=25$, as well as concluding that PISNe should be observable out to $z=50$ with the JWST.  Subsequent work by \\citet{WiseAbel2005} determined the PISN rate based on the collapse of gas into dark matter minihalos. Accounting for radiative feedback, they concluded that $\\sim$0.34 PISNe yr$^{-1}$ deg$^{-2}$ were to be expected above $z=10$, as well as briefly considering the detectability of PISNe at high redshifts based on models from \\citet{HegerWoosley2002}.  \\citet{Scannapiecoetal2005} presented a more thorough analysis of the visibility of PISNe based on a suite of numerical simulations spanning the range of theoretical PISN models using the implicit hydrodynamics code KEPLER. \\citet{MesingerJohnsonHaiman2006} presented a similar but more general halo mass function based analysis, considering the rates and detectability of all SNe; they briefly consider primordial PISNe, but focus on core-collapse SNe and their detectability.  More recently, \\citet{TrentiStiavelliShull2009} explored the observable PISN rate within the context of the metal-free gas supply during the epoch of reionization.  Our work improves upon these investigations by incorporating updated PISN models from \\citet{KasenWoosleyHeger2011} and determining their observability using the published specifications of the Near Infrared Camera (NIRcam) on the JWST\\footnote{http://www.stsci.edu/jwst/instruments/nircam}. In the final stages of the work on this paper we have become aware of the study by \\citet{PanKasenLoeb2012}, who have performed a similar analysis. This paper addresses the question of PISN observability in a nicely complementary way by employing a different normalization strategy. Different from our assumption that viable Pop~III progenitors can only form in unenriched minihalos at $z\\gtrsim6$, \\citet{PanKasenLoeb2012} derive the star formation rates required to produce sufficient photons for reionization.  They then infer the PISN rate corresponding to different choices for the IMF.  Their analysis is thus able to probe the charachter of star formation in the dwarf galaxies that are the main drivers of reionization, whereas we focus on the minihalos where Pop~III stars first begin to form.  This paper is organized as follows.  In Section 2 we describe our semi-analytic model for the PISN rate.  We consider the ability of the JWST to detect PISNe at high redshift in Section 3, and our conclusions are gathered in Section 4.  Throughout this paper we adopt a $\\Lambda$CDM model of hierarchical structure formation, using cosmological parameters consistent with the WMAP 5-year results \\citep{Komatsu2009}: $\\Omega_{\\rm m} = 0.258$; $\\Omega_{\\Lambda}=0.742$; $\\Omega_b=0.0441$; $h=0.719$; $n_s=0.96$; $\\sigma_8=0.796$. ",
        "conclusions": "\\label{conclusions} \\begin{figure*}[t!] \\begin{center} \\includegraphics[width=15cm]{observability} \\caption{\\footnotesize The observability of the R250, He100, B200, and R175 models (clockwise from upper left) from \\citet{KasenWoosleyHeger2011} using the JWST's NIRcam F444W filter. Shown is the possible range for the number of JWST FoVs required to detect 10 sources as a function of exposure time. The blue range is for all PISNe, and the red for PISNe from $z>15$.  The lower boundaries correspond to the no-feedback upper limit to the PISN rate and the upper boundaries to the conservative feedback rate.  From left to right, the black lines represent the number of pointings possible in a total of $10^6$, $10^7$, and $10^8\\,$s for a given exposure time.} \\label{observability} \\end{center} \\end{figure*}  In this work, we have examined the source density of PISNe from Pop~III stars and considered their detectability with the JWST. We conclude that the limiting factor in detecting PISNe will be the scarcity of sources rather than their faintness, in agreement with the conclusions of \\citet{WeinmannLilly2005}. The brightest PISNe should be readily detectable with the longest wavelength NIRcam filters out to $z\\sim25$; the problem is the overall scarcity of sources.  We have derived an estimate for the observable PISN rate, finding an upper limit of just over 0.01 PISNe per year per JWST FoV in the case of negligible chemical and radiative feedback. We also find that the inclusion of feedback can reduce the PISN rate by an order of magnitude, to \\about0.001 per FoV. Accounting for the possibility of enhanced massive star formation in halos affected by LW radiation improves this rate slightly, to \\about0.003 per FoV.  The derived PISN rate then allows us to place an upper limit on the observable number of PISNe in a $10^6\\,$s exposure of \\about0.2 PISNe per JWST FoV in the no-feedback case. The most pessimistic case of a B200-type PISN with strong negative feedback reduces this number to \\about2.5$\\,\\times10^{-4}$ per FoV, or one PISN per 4000 JWST fields of view.  The long duration of PISN lightcurves imply that spectroscopic follow-up of PISNe will likely be of great importance. PISN lightcurves can last for decades when combined with the cosmological time dilation factors at high redshifts, making the detection of PISNe by looking for transients untenable.  However, a multi-year campaign could identify candidates photometrically, and the lack of metal lines in the spectrum at early times could provide a spectroscopic signature for identification.  We find that the main obstacle to observing PISNe is the paucity of sources.  Beyond a moderate exposure time of a few times $10^4\\,$s, the observability of bright PISNe is not a strong function of exposure time and is instead controlled by the source density; this is evident from \\RefFig{observability}, where we have shown the number of JWST FoVs required to detect 10 PISNe (blue) as a function of exposure time for each of our PISN models.  The upper boundaries correspond to our conservative estimate for the PISN rate in the presence of feedback, and the lower boundaries to the estimated rate without feedback. We can see that the decrease in the required number of JWST pointings slows considerably beyond \\about10$^5\\,$s, hence a deep pencil-beam survey would not be advisable in searching for PISNe. Even for only high-redshift sources ($z>15$; red) the dependence on exposure time is still minimal, being controlled by the lack of sources once the required imaging depth is reached.  \\begin{figure*} \\begin{center} \\includegraphics[width=\\textwidth]{area_observability_R250} \\caption{\\footnotesize The total number of PISNe observable with a campaign of $10^6$, $10^7$ and $10^8\\,$s (from left to right) as a function of survey area for the R250 PISN model.  In each case, the total campaign time is apportioned equally over the total survey area to determine the exposure time for individual pointings.  The blue region represents all PISNe, the red only PISNe from $z>15$.  Upper boundaries correspond to the no-feedback upper limit to the PISN rate and lower boundaries to the conservative feedback case. For reference we mark the case of only one PISN visible (dashed line).} \\label{area_obsR250} \\end{center} \\end{figure*} Of particular interest in \\RefFig{observability} are the black lines representing the total number of pointings possible in $10^6$, $10^7$, and $10^8\\,$s for a given exposure time and their location relative to the observability ranges in blue and red.  $10^6\\,$s is approximately the limit of what would be possible with a dedicated deep-field campaign; $10^7\\,$s is the limit of the observations the JWST could make in a year assuming NIRCam is in use one third of the time; $10^8\\,$s (\\about10 years) is the projected mission lifetime.  While the detection of a PISN from a `first' star at very high redshifts would be exciting and is in fact possible given the detection limits of the JWST, the scarcity of sources at these redshifts means that such a detection would be highly contingent on serendipity. Even in the most optimistic case, with all available minihalos producing an R250-type PISN, the observability range for such events lies well above what is possible even in a full year of observations, though a few may be detected over the lifetime of the telescope.  The detection of a PISN at lower redshifts appears to be more realistic.  As the faintest PISNe (R175 and B200) are effectively unobservable, PISN searches should focus on looking for PISNe similar to the R250 and He100 models.  In this case, the strategy with the highest likelihood of detection will be a mosaic survey of many moderately deep exposures.  This is clear from \\RefFig{area_obsR250}, where we show the number of PISNe that will be observable with the JWST in observing campaigns totalling $10^6$, $10^7$ and $10^8\\,$s for the R250 PISN model. The exposure time for each pointing varies with the total area covered by the survey in order to keep the total observing time constant.  Upper boundaries correspond to the number visible in the no-feedback case, lower boundaries to the conservative feedback case. As in \\RefFig{observability}, the blue region shows the observable number from all redshifts, the red region only those from $z>15$.  We see that the observable number increases until the resulting exposure time is no longer sufficient to detect PISNe.  The optimal search strategy then will be to cover as large an area as possible, going only as deep as necessary, possibly in a similar manner to the ongoing Brightest of Reionizing Galaxies survey with the {\\it Hubble Space Telescope} \\citep{Trentietal2011, Bradleyetal2012}.  \\phantom{.}"
    },
    "1112/1112.1612_arXiv.txt": {
        "abstract": "Wide binaries are tracers of the gravity field of the Galaxy, but their study requires some caution. A large list of common proper motion stars selected from the AGK3 were monitored with the {\\sc coravel} spectrovelocimeter, in order to prepare a sample of physical binaries with very wide separations. Sixty-six stars received special attention, since their RV seemed to be variable. These stars were monitored over several years in order to derive the elements of their spectroscopic orbits. In addition, 10 of them received accurate RV measurements from the {\\sc sophie} spectrograph of the T193 telescope at the Observatory of Haute-Provence.  For deriving the orbital elements of double-lined spectroscopic binaries (SB2), a new method was applied, which assumed that the RV of blended measurements are linear combinations of the RV of the components. Thirteen SB2 orbits were thus calculated.  The orbital elements were eventually obtained for 52 spectroscopic binaries (SB), two of them making a triple system. Forty SB received their first orbit, and the orbital elements were improved for 10 others. In addition, 11 spectroscopic binaries were discovered with very long periods for which the orbital parameters were not found. The median period of the 40 first orbits is 1 year, and several SB should be resolved or should receive an astrometric orbit in the future, providing the masses of the components. In addition, it appeared that HD 153252 has a close companion which is a candidate brown dwarf with a minimum mass of 50 Jupiter masses.  The final selection of the wide binaries and the derivation of their statistical properties will be presented in a second paper. ",
        "introduction": "Binarity is a very common stellar property, which covers a very wide range of separations, from two stellar radii to thousands of astronomical units. Close binaries have short orbital periods (from a few hours to a few years), and they are easily detected from variations in their radial velocities (RV). For that reason, they are the most commonly studied of this class of stars (see, eg, \\citealt{Halbwachs03}). Binaries with periods around a few centuries are also rather well known: the components are sufficiently separated to be detected visually (in the past) or on images, and they are still close enough to avoid the risk of confusing a field star and a binary component: they may be discarded by applying a statistical criterion almost as old as the discovery of double stars \\citep{Struve52}.  Very wide binaries are particularly interesting, since the distribution of their separations is a clue to the gravitational perturbers which made the gravity field of the Galaxy (see \\citealt{Jiang10} and references therein). However, the selection of binaries with separations of thousands of astronomical units is difficult: the components have also wide apparent separations, and it is necessary to use additional criteria in order to discard the optical companions. Trigonometric parallax, or another estimation of the distance, is sometimes used, but the most efficient, and the most employed criterion is proper motion: when the semi-major axis of the orbit is large, the orbital motion of the stars around the barycentre of the system generates a difference in proper motion which is negligible, and the components must have similar apparent displacements. Such binaries are called common proper motion (CPM) stars.  In the past, this criterion was used by Luyten to search for wide systems by visual inspection of photographic plates; his prospect was performed over more than 40 years, from \\citet{Luyten40} to \\citet{Luyten87} and led to the discovery of 6121 systems. However, despite their large number, these systems were not used to derive the properties of wide binaries; the main reason is that the selection was based on subjective criteria, and, as a consequence, it would be hard to estimate its incompleteness, and also the contamination by optical pairs. Moreover, since the components of Luyten's double stars are rather faint (around 16th magnitude), it would be difficult to improve the selection with complementary data.  Another selection of CPM pairs was performed by \\citet{Halbwachs86}, on the basis of the AGK2/3 catalogue \\citep{Lacroute74}. Four hundred and thirty-nine CPM double stars were selected, and the number of optical pairs was estimated to be around 40. Since the AGK3 stars are brighter than 12th magnitude, it was decided to select the physical binaries on the basis of the RV of the components. As is the case for the proper motions, the RV are marginally affected by the orbital motion, and the components of wide binaries are therefore also common RV stars. Two hundred and sixty-six stars were selected, and they were measured with the spectrovelocimeter {\\sc coravel}. This programme was initiated in 1986, and it was initially supposed to end after a few years. However, several stars had variable velocities, since these wide pair components were themselves spectroscopic binaries (SB). It was then necessary to extend the observations over about 20 years, in order to derive the SB orbital elements, and thus the systemic velocity of these stars. The present paper is devoted to these variable velocity stars, and to the derivation of their SB elements. The selection of the wide binaries and the derivation of their statistical properties will be treated in a second paper.  The paper is organized as follows: the RV observations are presented in section~\\ref{sec:RV}, and the catalogue of the RV measurements is in section~\\ref{sec:RVcat}. The calculation of the SB orbits is in section~\\ref{sec:SB}; for the double--lined SB (SB2), we present a method to take into account the measurements related to blends of the two components. Some interesting points relating to the new SB that we have discovered are discussed in the conclusion, section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  An observational program initiated for searching common RV stars among CPM stars resulted in the selection of 66 stars suspected to be variable. Thanks to observing runs distributed over a very long time (more than 20 years for 11 stars), we finally derived a first SB orbit for 40 stars. One of these orbits (2:92A=HD~153252) corresponds to a possible brown dwarf companion with a minimum mass equal to 50 Jupiter masses. We derive the elements of 13 SB2 orbits, assuming that the RV of the blended measurements are linear combinations of the RV of the components.  The periods of the 40 new orbits are on average rather long: more than 1 year for 20 stars, including 5 SB2, and 3 stars (1 SB2) even have periods between 10 years and around 15 years. Some of these binaries could be resolved, and the masses of the components could be derived from the combination of the visual and the spectroscopic observations. Among the SB2 brighter than 7~mag, two are ideal targets for speckle interferometry: 2:40A and 1:280A should have a semi-major axis of 58~mas and 17~mas, respectively. Accurate masses should also be obtained in the future, thanks to the astrometric orbits expected from the forthcoming {\\it Gaia} mission \\citep{Halbwachs09}. For that purpose, seven new SB presented in that paper are now measured with {\\sc sophie} in order to improve their orbits.  In addition to the 40 new SB, we improved the orbital elements of ten others. We also found 11 long-period SB, including 2 SB2, for which it was not possible to derive the orbital elements.  We are now able to select a sample of physical wide binaries, including the components which are themselves close binaries, and to investigate the statistical properties of these stars and their relations with the other components of the Galaxy. This will be the topic of the forthcoming second paper."
    },
    "1112/1112.1424_arXiv.txt": {
        "abstract": "The Cosmic Infrared Background Experiment (CIBER) is a suite of four instruments designed to study the near infrared (IR) background light from above the Earth's atmosphere.  The instrument package comprises two imaging telescopes designed to characterize spatial anisotropy in the extragalactic IR background caused by cosmological structure during the epoch of reionization, a low resolution spectrometer to measure the absolute spectrum of the extragalactic IR background, and a narrow band spectrometer optimized to measure the absolute brightness of the Zodiacal light foreground.  In this paper we describe the design and characterization of the CIBER payload.  The detailed mechanical, cryogenic, and electrical design of the system are presented, including all system components common to the four instruments.  We present the methods and equipment used to characterize the instruments before and after flight, and give a detailed description of CIBER's flight profile and configurations. CIBER is designed to be recoverable and has flown twice, with modifications to the payload having been informed by analysis of the first flight data.  All four instruments performed to specifications during the second flight, and the scientific data from this flight are currently being analyzed. \\vspace{0.5cm} ",
        "introduction": "\\label{S:intro}  \\setcounter{footnote}{0}  The extragalactic background light (EBL) is the integrated intensity of all of the photons emitted along a line of sight through the Universe.  Since at near infrared (IR) wavelengths the EBL is thought to be dominated by direct emission from stars, a measurement of the EBL at these wavelengths is an integral constraint on the total energy released via the process of nucleosynthesis through the history of the Universe (see \\citealt{Hauser2001} for a comprehensive review).  Such measurements are a key constraint on all models of galaxy formation and evolution, connecting the global radiation energy density to star formation, metal production, and gas consumption from the present back to the epoch of reionization.  The Cosmic Infrared Background Experiment (CIBER) is a sounding-rocket borne payload comprising four instruments designed to measure the spectrum of the EBL, test models of the Zodiacal light intensity, and constrain the spatial fluctuations in the EBL imprinted during the epoch of reionization (\\citealt{Bock2006}, \\citealt{Zemcov2011}).  The CIBER payload comprises two imaging telescopes (Imagers) which are designed to study spatial fluctuations in the near IR EBL, a low resolution spectrometer (LRS) designed to measure the spectrophotometric properties of the EBL in the range $0.7 \\leq \\lambda \\leq 2.1 \\, \\mu$m, and a narrow band spectrometer (NBS) designed to characterize the absolute brightness of the Zodiacal Light (ZL) in a number of astronomical fields.  CIBER has flown twice, first on Feb 25, 2009 and second on Jul 10, 2010, and two additional flights of essentially the same payload configuration are planned.  A preliminary measurement of the spectrum of the ZL as measured by CIBER is presented in \\citet{Tsumura2010}.  This paper presents the overall mechanical, cryogenic and electronic design and implementation of the CIBER hardware.  In addition, the customized laboratory test equipment and a description of the typical CIBER flight profile and performance are described.  The design, scientific motivation and implementation of the individual instruments are presented in separate papers: \\citealt{Sullivan2011} discusses the design, implementation and performance of the Imagers, while \\citealt{Tsumura2011} describes the LRS and \\citealt{Renbarger2011} the describes NBS.  Together, these three papers are hereafter referred to as the ``CIBER Instrument Companion'' papers.  Since the CIBER payload was modified between its two flights to improve its performance, throughout this paper we describe the final configuration from the second flight, and specify the design changes informed by the first flight. ",
        "conclusions": ""
    },
    "1112/1112.1754_arXiv.txt": {
        "abstract": "{We present two large cosmological $N$-body simulations, called Horizon Run 2 (HR2) and Horizon Run 3 (HR3), made using $6000^3$ = 216 billions and  $7210^3$ = 374 billion particles, spanning a volume of (7.200 $h^{-1}\\textrm{Gpc}$)$^3$ and (10.815 $h^{-1}\\textrm{Gpc}$)$^3$, respectively. These simulations improve on our previous Horizon Run 1 (HR1) up to a factor of 4.4 in volume, and range from 2600 to over 8800 times the volume of the Millennium Run. In addition, they achieve a considerably finer mass resolution, down to $1.25 \\times 10^{11} h^{-1} $M$_{\\odot}$, allowing to resolve galaxy-size halos with mean particle separations of 1.2$h^{-1}\\textrm{Mpc}$ and 1.5$h^{-1}$Mpc, respectively. We have measured the power spectrum, correlation function, mass function and basic halo properties with percent level accuracy, and verified that they correctly reproduce the $\\Lambda\\textrm{CDM}$ theoretical expectations, in excellent agreement with linear perturbation theory. Our unprecedentedly large-volume $N$-body simulations can be used for a variety of studies in cosmology and astrophysics, ranging from large-scale structure topology, baryon acoustic oscillations, dark energy and the characterization of the expansion history of the Universe, till galaxy formation science -- in connection with the new SDSS-III. To this end, we made a total of 35 all-sky mock surveys along the past light cone out to $z=0.7$ (8 from the HR2 and 27 from the HR3), to simulate the BOSS geometry. The simulations and mock surveys are already publicly available at \\textrm{http://astro.kias.re.kr/Horizon-Run23/.}} ",
        "introduction": "\\label{sec_intro}   State-of-the-art observations from the Cosmic Microwave Background (CMB), such as data provided by the Wilkinson Microwave Anisotropy Probe (WMAP; Spergel et al. 2003, Komatsu et al. 2011), and from the Large Scale Structure (LSS) as in the Sloan Digital Sky Survey (SDSS; York et al. 2000), in the 2dF galaxy redshift survey (Colless et al. 2001) and in the WiggleZ survey (Blake et al. 2008), support a model of the Universe dominated by Cold Dark Matter (CDM) and Dark Energy (DE), with baryons constituting only a percent of the total matter-energy content. Structures form and grow hierarchically, from the smallest objects to the largest ones. In this framework, dark matter collapses first into small haloes which accrete matter and merge to form progressively larger halos over time. Galaxies form subsequently, out of gas which cools and collapses to the centers of dark matter halos (White \\& Rees 1978; Peebles 1982; Davis et al. 1985); therefore, halo and galaxy properties are strongly related. In synergy with independent data from Type I supernovae (SNI$a$; Riess et al. 1998, Perlmutter et al. 1999, Kowalski et al. 2008), we also have increasing confirmations that the Universe is geometrically flat, and currently undergoing a phase of accelerated expansion.   What is driving this expansion, and in particular the nature of dark energy, still remains to be explained. To date, it is perhaps one of the most important open questions in cosmology, which would deeply impact our understanding of fundamental physics (see for example Albrecht et al. 2006). In addition, investigating the nature of the (nearly) Gaussian primordial fluctuations -- which eventually led to the formation of halos and galaxies -- is also of utmost importance for shedding light into the structure formation process. In fact, constraining possible deviations from Gaussianity will impact for example the LSS topology (Gott et al. 1986; Park et al. 1998, 2005; Gott et al. 2009; Choi et. al 2010), as any primordial non-Gaussianity might modify the clustering properties of massive cosmic structures forming out of rare density fluctuations (LoVerde et al. 2008; Desjacques et al. 2009; Jeong \\& Komatsu 2009), and generate a scale-dependent large-scale bias in the clustering properties of massive dark matter halos (Dalal et al. 2008; Verde \\& Matarrese 2009; Desjacques \\& Seljak 2010). Moreover, studying global halo properties such as halo formation, density profiles, concentrations, shapes, kinematics, assembly times, spin, velocity distributions, substructures and the effects of baryons is particularly important in order to gain insights into the formation and evolution of galaxies, their large-scale environment in the cosmic web (Bett et al. 2007;  Gao \\& White 2007; Jing Suto \\& Mo 2007; Li et al. 2008; Gott et al. 2009; Shandarin et al. 2010), and to test the validity of cosmological models.   Therefore, at the moment there is considerable interest in constraining cosmological models (i.e. the main cosmological parameters) from real datasets with percent level accuracy, and especially the dark energy equation of state (DE EoS). In particular, most of the efforts towards characterizing DE and its properties are focused on measurements of the baryon acoustic oscillation (BAO) scale. The BAO distinct signature, imprinted in the large-scale galaxy distribution by acoustic fluctuations in the baryon-photon fluid prior to recombination, appears as a quasi-harmonic series of oscillations of decreasing amplitude in the galaxy power spectrum at wavenumbers $0.01h$Mpc$^{-1} \\le k \\le 0.4 h $Mpc$^{-1}$ (Sugiyama 1995; Eisenstein \\& Hu 1998, 1999), or as a broad and quasi-Gaussian peak in the corresponding two-point correlation function (Matsubara 2004). The measurement of its scale is often achieved by using a well-controlled sample of luminous red galaxies (LRGs), as observed for example in the SDSS or in the 2dFGRS surveys (Colless et al. 2001; Eisenstein et al. 2005; Cole et al. 2005; Sanchez et al 2006, 2009; Percival et al. 2007, 2010; Gaztanaga et al. 2009; Kazin et al. 2010; Reid et al. 2010; Carnero et al. 2011). Since the BAO scale depends solely on the plasma physics after the Big Bang and  can be calibrated using CMB data, it can be used as a standard ruler to measure the redshift dependence of the Hubble parameter $H(z)$ and the angular diameter distance (Montesano et al. 2010), and thus to constrain the DE EoS parameter $w(z)$ as a function of redshift -- along with the other main cosmological parameters. Complementary to the BAO technique, the LSS topology can also be used as a standard ruler, by measuring and comparing the amplitude of the genus curve at different redshifts and smoothing lengths (see Park \\& Kim 2010). Cosmological parameters can also be obtained through the abundance of high-mass halos, identified as galaxy clusters, and via several other methods that require and utilize precise knowledge of galaxy clustering (i.e. Zheng \\& Weinberg 2007; Yoo et al. 2009).   To this end, planned next generation photometric redshift galaxy surveys will span volumes considerably larger than the current data sets, providing a dramatic improvement in the accuracy of the constraints on cosmological parameters. Examples are the Panoramic Survey Telescope \\& Rapid Response System (Pan-STARRS; Kaiser et al. 2002), the Dark Energy Survey (DES; Abbott et al. 2005), the Baryonic Oscillation Spectroscopic Survey (BOSS; Schlegel  et al. 2009) as well as BigBOSS (Schlegel et al. 2011), the Large Synoptic Survey Telescope (LSST; Tyson et al. 2004), the Hobby Eberly Telescope Dark Energy Experiment (HETDEX; Hill et al. 2004), and the space-based Wide-Field Infrared Survey Telescope (WFIRST; Green et al. 2011) and Euclid (Cimatti et al. 2008). The lack of precision in the redshift determination is compensated with the larger volume of the survey, the larger density of galaxies, and the possibility of analyzing different galaxy populations.   With large-volume surveys becoming the norm, it is timely to devise large-volume numerical simulations able to mimic, reproduce and control the various observational surveys. This paper aims at presenting two large cosmological $N$-body simulations, called Horizon Run 2 (HR2) and Horizon Run 3 (HR3), made using innovative computational facilities and forefront capabilities. Cosmological $N$-body simulations provide in fact a powerful tool to test the validity of cosmological models and galaxy formation mechanisms (Bertschinger 1998), as well as they represent a complementary and necessary benchmark to design and support the observational surveys. In particular, the advent of realistic cosmological $N$-body simulations, capable of reproducing successfully the global properties of the observed structure at large scales, has enabled significant progress in understanding the structure of dark matter halos, the DM clustering over a huge range of scales, the distribution of galaxies in space and time, the effects of the environment, the mass assembly history, and the nature of DM itself. For example, numerical studies have indicated that the  hierarchical assembly of CDM halos yields approximately universal mass profiles (Navarro et al. 1996, 1997), strongly triaxial shapes with a slight preference for nearly prolate systems (i.e. Jing \\& Suto 2002), presence of abundant but non-dominant substructure within the virialized region of a halo, and cuspy inner mass profiles (for instance see Springel et al. 2005). Numerical simulations are also crucial to DE studies, especially for characterizing BAO systematics (Crotts et al. 2005; Crocce \\& Scoccimarro 2008; Kim et al. 2009). These are well controlled only with a detailed knowledge of all nonlinear effects associated, and can only be addressed by studying mock catalogs constructed from large $N$-body simulations. In this respect, comparison with $N$-body simulations is critical, as it allows one to calibrate the DE experiments with high accuracy.   While the current and future surveys of large-scale structure in the Universe demand larger and larger simulations, simulating large cosmological volumes with good mass resolution is clearly a major computational challenge. However, it is crucial to be able to reach volumes big enough to assess the statistical significance of extremely massive halos, almost always underrepresented in simulations that survey a small fraction of the Hubble volume. For example, till recently our understanding of the mass profile of massive haloes has been rather limited, derived largely from small numbers of individual realizations or from extrapolation of modes calibrated on different mass scales (Navarro et al. 1996, 1997; Moore et al 1998; Klypin et al. 2001; Diemand et al. 2004; Reed et al. 2005). This is because, at high masses, enormous simulation volumes are required in order to collect statistically significant samples of these rare DM halos. Instead, individual halo simulations may result in biased concentration estimates in a variety of ways (Gao et al. 2008). At high masses we also know only approximately the redshift and mass dependence of halo concentration, even for the current concordance cosmology. Another example of how important is to simulate big volumes is related to DE science: larger volumes allow one to model more accurately the true power at large scales and the corresponding power spectrum, particularly in the case of the LRG distribution, which is critical to the BAO test (Kim et al. 2009). A large box size will guarantee small statistical errors in the power spectrum estimates: we need in fact to measure the acoustic peak scale down to an accuracy better than 1\\%.  \\begin{figure*}[!t] \\centering \\epsfxsize=12cm \\epsfbox{fig1.eps} \\caption{Evolution of the number of particles in $N$-body simulations versus time (years). Filled hexagons (pale green) are PM simulations made by our group; stars (pink) are PM-Tree simulations carried out by our group (the HR2 and HR3 are located in the upper-right corner); open circles (blue) are those performed by other authors -- see the detailed legend in the figure. Note in particular the position of the Millennium Run (Springel et al. 2005; 10 billion particles, box size 500$h^{-1}$Mpc), and the French collaboration (Teyssier et al. 2009; 68.7 billion particles, box size 2000$h^{-1}$Mpc). The solid line is the mean evolution of the simulation size.} \\label{fig_simulation_particles} \\end{figure*}  As explained in Neto et al. (2007) and in Gao et al. (2008), trying to combine different simulations of varying mass resolution and box size is instead potentially very risky, along with extrapolation techniques. To this end, Neto et al. (2007) brings the example of Macci\\`{o} et al. (2007), who combined several simulations of varying mass resolution and box sizes in order to reach scales with $M \\ll M_{*}$. Unfortunately, this approach comes with pitfalls because in order to resolve a statistically significant sample of haloes with masses of the order of $10^{10}h^{-1} $M$_{\\odot}$ one must use a considerably smaller simulation box (i.e. 14.2$h^{-1}$Mpc on a side); this implies a substantial lack of large-scale power, due to such small periodic realization, which may influence the result. The obvious way out is to increase the dynamic range of the simulation, so as to encompass a volume large enough to be representative, while at the same time having enough mass resolution to extend the analysis well above (or below) $M_{*}$.  In this view, the enormous volume of our HR2 and HR3, together with the large number of particles used, make our simulations ideal to characterize with minimal statistical uncertainty the dependence of the structural parameters of CDM halos on mass, formation time, departures from equilibrium, etc., and will be of great use for DE science and for all the upcoming photometric redshifts surveys.   The paper is organized as follows. In Section \\ref{sec_num_met} we briefly discuss the standard numerical techniques in LSS analyses, and compare the number of particles and volumes of our simulations with other recent numerical studies. In Section \\ref{sec_our_sims}, after a synthetic description of our previous Horizon Run 1 (HR1), we present our new simulations HR2 and HR3. In particular, we provide some technical details on the code used, the mass and force resolution, and the halo selection procedure. In Section \\ref{sec_mocks} we introduce the 35 all-sky mock surveys along the past light cone (8 from the HR2 and 27 from the HR3), made in order to simulate the BOSS geometry. These mock catalogs are particularly suitable for DE studies (i.e. LRGs), and are already publicly available. We then present in Section \\ref{sec_first_results} the first results from testing our new simulations; namely, the computation of the starting redshifts, mass functions, power spectra and two-point correlation functions. We conclude in Section \\ref{sec_conclusions}, with a brief description on the importance and use of our simulations and catalogs. We are making the simulation data and mocks available to the community at \\textrm{http://astro.kias.re.kr/Horizon-Run23/}. ",
        "conclusions": "\\label{sec_conclusions}   The main goal of this paper was to present in detail two new large cosmological $N$-body simulations, HR2 and HR3, made using $6000^3$ = 216 billions and  $7210^3$ = 374 billion particles, and spanning a volume of (7.200 $h^{-1}$Gpc)$^3$ and (10.815 $h^{-1}$Gpc)$^3$, respectively. The volumes of these two simulations are considerably bigger than our previous HR1 volume, up to a factor of 4.4, and the mass resolution has been improved (down to $1.25 \\times 10^{11} h^{-1} $M$_{\\odot}$); this allows to resolve large-galaxy (i.e. LRG galaxy) sized halos, with mean particle separations of 1.2$h^{-1}$Mpc (HR2) and 1.5$h^{-1}$Mpc (HR3). The main characteristics of the simulations have been described in detail in Section \\ref{sec_our_sims}, including the code developed and used to perform the runs, the mass and force resolution, and the halo selection procedure. The basic simulation parameters are listed in Table \\ref{table_sim_par}.  Planned next generation photometric redshift galaxy surveys such as BOSS (Schlegel  et al. 2009), BigBOSS (Schlegel et al. 2011), DES (Abbott et al. 2005), LSST (Tyson et al. 2004), WFIRST (Green et al. 2011) and Euclid (Cimatti et al. 2008) will span volumes considerably larger than the current data sets, providing a dramatic improvement in the accuracy of the constraints on cosmological parameters; to this end, our simulations will be of great help for these surveys.  In particular, in support of the SDSS-III, we made a total of 35 all-sky mock catalogs along the past light cone up to $z=0.7$, 8 from the HR2 and 27 from the HR3, to simulate the BOSS geometry (Section \\ref{sec_mocks}). The simulations and mock surveys are already publicly available to the community, and they should be useful for a variety of studies in cosmology and astrophysics, ranging from large-scale structure topology, baryon acoustic oscillations, dark energy and the characterization of the expansion history of the Universe, till galaxy formation science.  In Section \\ref{sec_first_results} we also presented the first results from the analysis of the HR2 and HR3 simulations. In particular, we have measured power spectra, correlation functions, mass functions and basic halo properties with percent level accuracy, and verified that they correctly reproduce the $\\Lambda\\textrm{CDM}$ theoretical expectations, in excellent agreement with linear perturbation theory. For example, by computing the halo multiplicity functions we proved that our large-volume simulations are able to resolve correctly about six decades in mass. Much more scientific analyses can be carried out with these simulations, and in the upcoming papers we will focus in particular on topology, BAO, and galaxy formation science.   Clearly, one may question the utility of performing such large-volume, expensive and computationally challenging simulations. Why not running instead many smaller simulations, with different initial conditions? Several reasons were already pointed out in the introductory part; essentially, the main goal in doing so is to cover large scale modes accurately, which is crucial for dark energy and BAO science (i.e. systematic effects), for the statistical study of rare massive halos, and for the formation of unusually large scale structures. In fact, as we argued at the beginning, extremely massive halos are almost always underrepresented in simulations that survey a small fraction of the Hubble volume, and in order to reach a statistically significant sample of these objects one needs simulations with enormous dynamic range. By contrast, combining different simulations of varying mass resolution is instead potentially very risky (i.e. Neto et al. 2007; Gao et al. 2008).  While we strive to carry out simulations in as large volumes as possible, we also need high resolutions to sample accurately the underlying gravitational potential wells. The HR2 and HR3 achieve both criteria: their volumes are unprecedentedly large, but their mass resolutions allow us to resolve large-galaxy-sized (LRG) halos; therefore, they are ideal to characterize with minimal statistical uncertainty the dependence of the key structural parameters of CDM halos on mass, spin, formation time, etc., and will be of great use for DE science (i.e. BOSS) and for all the upcoming photometric redshifts surveys.  In particular, a topical issue which can be addressed with our simulations is the nature of DE and the determination of the DE EoS parameter $w(z)$ as a function of redshift. A standard technique, which has received much attention lately, is the measurement of the BAO scale using LRGs (Eisenstein et al. 2005). In this case, being able to control the systematics is imperative, and can only be achieved via large volume $N$-body simulations such as the HR2 and HR3. Note in fact that the typical box sizes usually adopted for these studies in previous literature range from 500$h^{-1}$Mpc to 2000$h^{-1}$Mpc, but those boxes are inadequate because they are too close to the baryon oscillation scale -- which is of the order of 108 $h^{-1}$Mpc. Instead, it is essential to be able to model the power spectrum accurately at large scales, and to have a large box size so that the statistical errors in the power spectrum are small.  The HR2 and HR3 also permit us to assess the statistical significance of unusually large structures such as the Sloan Great Wall which, with a length of 1.37 billion light years (Gott et at. 2005), is considerably more massive than predicted by previous $N$-body simulations (i.e. Springel et al. 2005). The large number of halos in our simulations allows us to study in detail deviations from the mean trends and the possible presence of such systems with unusually high concentrations of galaxy haloes, or unusually low densities. We will present a detailed study focused on this issue in a separate publication. We will also investigate several other applications, which involve the large scale topology in relation to dark energy and distance measurements, the detailed modeling of the integrated Sachs-Wolf effect (Yoo, Fitzpatrick, \\& Zaldarriaga 2009) and the comparison with counts of LRGs in the sky, and even applications of first order corrections to our simulations in order to capture general relativity effects.   Our simulations and mock surveys are already publicly available at \\textrm{http://astro.kias.re.kr/Horizon-Run23/}. Please cite this paper whenever you use the simulation data."
    },
    "1112/1112.4347_arXiv.txt": {
        "abstract": "We consider the distribution of many samples of Gamma-Ray Bursts (GRBs) when plotted in a diagram with their bolometric fluence ($S_{bolo}$) versus the observed photon energy of peak spectral flux ($E_{peak,obs}$).  In this diagram, all bursts that obey the Amati relation (a luminosity relation where the total burst energy has a power law relation to $E_{peak,obs}$) must lie above some limiting line, although observational scatter is expected to be substantial.  We confirm that early bursts with spectroscopic redshifts are consistent with this Amati limit.  But, we find that the bursts from BATSE, {\\it Swift}, Suzaku, and Konus are all greatly in violation of the Amati limit, and this is true whether or not the bursts have measured spectroscopic redshifts.  That is, the Amati relation has definitely failed.  In the $S_{bolo}-E_{peak,obs}$ diagram, we find that every satellite has a greatly different distribution. This requires that selection effects are dominating these distributions, which we quantitatively identify.  For detector selections, the trigger threshold and the threshold for the burst to get a measured $E_{peak,obs}$ combine to make a diagonal cutoff with the position of this cutoff varying greatly detector to detector.  For selection effects due to the intrinsic properties of the burst population, the distribution of $E_{peak,obs}$ makes for bursts with low and high values to be rare, while the fluence distribution makes for bright bursts being relatively uncommon.  For a detector with a high threshold, the combination of these selection effects serves to allow only bursts within a region along the Amati limit line to be measured, and these bursts will then appear to follow an Amati relation.  Therefore, the Amati relation is an artifact of selection effects within the burst population and the detector.  As such, the Amati relation should not be used for cosmological tasks.  This failure of the Amati relation is in no way prejudicial against the other luminosity relations. ",
        "introduction": "Luminosity relations are a tool for which we can connect a measurable quantity (called a luminosity indicator) from a long-duration Gamma-Ray Burst (GRB) to the burst's peak luminosity or total energy.  The inverse square law and the observed burst brightness can then be used to determine the distance to the burst.  In doing this, GRBs become a `standard-candle' (in the same sense as Cepheid variables and Type Ia supernovae) where observed properties can be used to determine luminosity and then distance.  GRBs are seen out to very high distances, at least to redshift z=8.2 (Tanvir et al. 2009), and so the GRB luminosity relations offer a unique and powerful opportunity for new cosmology.  Currently, GRBs have seven well-established luminosity relations (Fenimore \\& Ramirez-Ruiz 2000; Norris et al. 2000; Amati et al.  2002; Schaefer 2002; Schaefer 2003; Ghirlanda et al.  2004).  The luminosity indicators are the spectral lag time ($\\tau_{lag}$, Norris et al. 2000), `variability' (V, Fenimore \\& Ramirez-Ruiz 2000), the peak of the $\\nu F_\\nu$ power spectrum ($E_{peak}$; Schaefer 2003), the minimum rise time in the light curve ($\\tau_{RT}$; Schaefer 2002), and the number of peaks in the light curve ($N_{peak}$, Schaefer 2002).  Five of the luminosity relations connect the luminosity indicators to the burst's peak luminosity.  Two luminosity relations instead connect to the total energy of the burst, with the physics relating to the burst energetics instead of the conditions at the time of peak luminosity.  These are: the (A) Amati relation (Amati et al.  2002;2006), which connects $E_{peak}$ to the total energy of the burst assuming for isotropic emission ($E_{\\gamma,iso}$), and (B) the Ghirlanda relation (Ghirlanda et al.  2004) which connects $E_{peak}$ to the total energy of the burst, but with a correction for the beaming factor ($F_{beam}$), resulting in the energy emission ($E_{\\gamma}$).  In addition, there are a variety of newer luminosity relations have been proposed, but not yet fully tested and confirmed by the community (e.g.  Dainotti et al.  2008; 2010; 2011) and while others have proven to not be viable.  An example of this is the Firmani relation (Firmani et al.  2006), which was found to be no significant improvement over existing relations (Collazzi \\& Schaefer 2008).  Physical explanations for a variety of the established relations have been made (e.g.  Kobayashi et al.  2002; M\\'{e}sz\\'{a}ros et al.  2002; Schaefer 2003; Eichler \\& Levinson 2004; Liang et al.  2004; Schaefer 2004;  Rees \\& M\\'{e}sz\\'{a}ros 2005; Giannios \\& Spruit 2007; Thompson et al.  2007).  Recently, Collazzi et al.  (2011) performed an exhaustive study on the sources of error on $E_{peak}$.  This study showed a considerable amount of scatter was hidden in how $E_{peak}$ is measured, which was much larger than that from the reported Poisson errors alone.  The sources of this scatter included: (A) the choices of different analysts, (B) which $E_{peak}$ is measured (e.g. the time-integrated or time-resolve $E_{peak}$), and (C) the detector response matrix, all in addition to the regular Poisson statistical error.  This scatter can be as large as 0.43 in log space and has a typical value of 0.24.  This scatter can explain the scatter seen in the luminosity relations that use $E_{peak}$.  Still, even the currently accepted luminosity relations have their drawbacks.  The best (i.e.  the tightest) of these relations, the Ghirlanda relation, can only be applied if there is an observed jet break.  Jet breaks are a well understood phenomena (Rhoads 1997; Sari et al. 1999).  Measuring a jet break is fairly difficult for a variety of reasons, and has only been observed in a small percentage of bursts.  Melandri et al.  (2008) and Kocevski \\& Butler (2008) have pointed out problems in identifying these jet breaks with the X-ray data.  Most notably, the Amati relation has been criticized for several reasons.  Li (2006) originated a test that demonstrated an ambiguity when the measured properties are used to determine the redshift.  This result was confirmed by other groups later (e.g.  Schaefer \\& Collazzi 2007).  This ambiguity also exists for the Ghirlanda relation, but at a substantially higher redshift.  This ambiguity does not exist for any of the other confirmed luminosity relations, and so the ambiguity problem can go away if multiple relations are used to determine redshift.  This problem does \\textit{not} affect current work on the GRB Hubble Diagram because those luminosity indicators were derived from a known redshift obtained spectroscopically.  Another major criticism towards the Amati relation came from Nakar and Piran (2005).  In this work, Nakar and Piran developed a test specifically for the Amati relation, the beauty of the test being that a redshift was not needed.  This test has since been generalized in several independent investigations (e.g.  Band \\& Preece 2005; Schaefer \\& Collazzi 2007; Goldstein et al. 2010).  They combined the Amati relation (equation \\ref{eq:AmatiRel}) with the inverse square law for fluences (equation \\ref{eq:ISLF}) to eliminate $E_{\\gamma,iso}$ (equation \\ref{eq:NaP}). \\begin{equation} E_{\\gamma,iso} = (9.2 \\times 10^{47}\\unit{erg \\; keV^{-2.04}}) \\left[E_{p}\\left(1+\\textit{z}\\right)\\right]^{2.04} \\label{eq:AmatiRel} \\end{equation} \\begin{equation} E_{\\gamma,iso} = \\frac{4 \\pi d_{L}^2 S_{bolo}}{1+\\textit{z}} \\label{eq:ISLF} \\end{equation} \\begin{equation} \\frac{E^{2.04}_{peak}}{S_{bolo}} = \\frac{4 \\pi d_{L}^2}{(1+\\textit{z})^{3.04} (9.2\\times10^{47} \\unit{erg\\; keV^{-2.04}})} \\label{eq:NaP} \\end{equation} Here, $E_{\\gamma,iso}$ is the isotropic gamma ray energy, $d_L$ is the luminosity distance as derived with the concordance cosmology ($\\Omega_\\Lambda = 0.7$, $\\Omega_m = 0.3$, H$_0$ = 74 km/s/Mpc), $S_{bolo}$ is the bolometric fluence (the fluence over the burst rest frame 1-10,000 keV range), and \\textit{z} is the redshift of the burst.  The quantity $\\frac{E^{2.04}_{peak}}{S_{bolo}}$ has been called the `energy ratio' for the Amati relation (e.g. Band \\& Preece 2005).  The left side of equation 3 uses only directly observable quantities (albeit, they are model dependent), while the right side is only a function of distance.  As the distance rises, $d_L^2$ gets larger and $(1+\\textit{z})^{-3.04}$ gets smaller, which gives a maximum value for the right side.  When the concordance cosmology is used, the function peaks at $\\textit{z}\\sim3.6$.  Specifically, the right side of the equation cannot exceed $1.13 \\times 10^9 \\unit{keV^{2.04} \\; erg^{-1} \\; cm^{2}}$ and, therefore, \\begin{equation} \\frac{E^{2.04}_{peak,obs}}{S_{bolo} }  \\leq 1.13 \\times 10^{9} \\unit{keV^{2} \\; erg^{-1} \\; cm^{2}}. \\label{eq:AmatiLimit} \\end{equation} This becomes a simple way to test the Amati relation even for bursts without redshifts.  Similarly, for the Ghirlanda relation, \\begin{equation} E_{\\gamma} = (1.35 \\times 10^{47}\\unit{erg \\; keV^{-1.43}}) \\left[E_{p}\\left(1+\\textit{z}\\right)\\right]^{1.43} \\label{eq:GhirRel} \\end{equation} \\begin{equation} E_{\\gamma} = \\frac{4 \\pi d_{L}^2 S_{bolo}F_{beam}}{1+\\textit{z}} \\label{eq:GISLF} \\end{equation} \\begin{equation} \\frac{E^{1.43}_{peak}}{S_{bolo}} = \\frac{4 \\pi d_{L}^2 F_{beam}}{(1+\\textit{z})^{2.43} (1.35\\times10^{47} \\unit{erg\\; keV^{-1.43}})} \\label{eq:NaPG} \\end{equation} The beaming factor, $F_{beam}$, is defined as $(1-\\cos \\theta_{jet})$, where $\\theta_{jet}$ is the opening angle of the jet of the burst.   The right hand side has a maximum value at $\\textit{z}_{max} = 12.6$ with a value of $2.7 \\times 10^{10} \\unit{keV}^{1.43} \\unit{erg}^{-1} \\unit{cm}^{2}$ for $F_{beam}=1$.  Thus, the Ghirlanda relation forces the limit, \\begin{equation} \\frac{E^{1.43}_{peak,obs}}{S_{bolo} }  \\leq 2.7 \\times 10^{10} \\unit{keV^{1.43}  \\; erg^{-1} \\; cm^{2}}. \\label{eq:GhirLimit} \\end{equation} So we have a simple observational test for compliance with the Ghirlanda relation. We also have reproduced the result that the `energy ratio' for the Ghirlanda differs from the Amati relation (e.g. Band \\& Preece 2005).  At first glance, it might appear that this it is being overly generous to apply a beaming factor of $F_{beam}=1$ to calculate the limit of the Ghirlanda relation.  The whole point of applying such a beaming factor is to give an illustration of the lowest value the limit can have.  If we were apply some sort of average beaming factor, the value of the Ghirlanda limit would increase, resulting in many more rejected bursts.  As an example, let us use a typical jet angle, $\\theta_{jet}=8.5^{o}$, which corresponds to a beaming factor,  $F_{beam}=0.01$.  In this case, the Ghirlanda limit would increase by a power of two to $\\approx 2.7 \\times 10^{8} \\unit{keV^{1.43}  \\; erg^{-1} \\; cm^{2}}$.  There are two reasons, however, for choosing to keep the Ghirlanda limit with a $F_{beam}=1.0$.  The first is the simple mathematical statement that we are looking for the maximum value for $\\frac{E^{1.43}_{peak}}{S_{bolo}}$.  The second, and more physical explanation, is that the Ghirlanda relation accounts for all beaming factors, and thus the limit \\textit{should} given as when $F_{beam}=1.0$.  Indeed, as we will see, the Amati relation can be thought of as a sort of Ghirlanda relation that was constructed with some average beaming factor of bursts.  Nakar \\& Piran (2005) analyzed 751 \\textit{BATSE} bursts, finding 48\\% of the bursts to violate the Amati limit (equation \\ref{eq:AmatiLimit}).  They used this result to declare the Amati relation as not reliable.  Schaefer \\& Collazzi (2007) later showed that this fraction was what is expected for this relation due to simple scattering effects.  In the ideal case with no measurement uncertainties, there are no violators, but as soon as noise is introduced, some of the points (particularly those near the limit) would become violators.  Of course, there would be an equal number of bursts that would go up (and thus above the limit) to those who go down (and below the limit).  With the function on the right side of equation 3 being near its maximum value (when compared to the uncertainties in E$_{peak}$ and the scatter in the luminosity relations) for the redshift of most bursts (1 $<$ \\textit{z} $<$ 6), this means that about half of the GRBs will be apparent violators even if the Amati relation is correct.  So the real finding is that the differences between the limit and the equation are very small given the known scatter (Schaefer \\& Collazzi 2007).  Thus, the original test by Nakar and Piran actually \\textit{confirmed} the Amati relation.  Schaefer \\& Collazzi (2007) extended this test to 69 bursts with known redshifts from many satellites.  The result was that the Amati relation had an expected amount of violators, 44\\%.  The paper goes on to show that this test could be generalized for use to test \\textit{all} of the luminosity relations.  In most cases, this resulted in no maximum, or at the very least no maximum out to very high redshift ($\\sim 20$).  All luminosity relations were shown to have either no bursts failing the test (like the lag-luminosity relation), or a number of failures that is acceptable given measurement errors).  The end conclusion is that all accepted luminosity relations of the time had passed the Nakar and Piran test.  A complication arises in Band \\& Preece (2005), which examined (largely) the same BATSE bursts as Nakar and Piran without known redshifts, finding a $>$80\\% failure rate for the Amati relation.  This result was later confirmed by Goldstein et al. (2010).  In particular, the spread of points was greatly over the limit, and this cannot be attributable to measurement error or scatter in the luminosity relation.  With the differences in the violator fraction and the interpretation, we have a core dilemma for the Amati relation, and the understanding of these differences is the core of this paper.  Another criticism of the Amati relation is that the Amati relation is both dependent on the satellite and that it arises from selection effects (Butler et al.  2007).  Butler et al.  (2007) pointed out an apparent shift in the Amati relation between \\textit{Swift} and pre-\\textit{Swift} data sets.  Their claim that selection effects will produce the Amati relation were never substantiated by any analysis, examples, or derivations, and the cause of the selection effects was never identified.  These claims have scared off some workers from using any luminosity relations (e.g.  Bromm \\& Loeb 2007); however, it has been widely rejected for a variety of strong causes (e.g.  Cabrera et al.  2007; Amati et al. 2009; Ghirlanda et al. 2009; Krimm et al.  2009; Nava et al. 2009; Xiao \\& Schaefer 2009) and most recently the authors themselves recanted their previous findings (Butler et al.  2009).  Nevertheless, we are here reconsidering the cause and effects for the basic claims of satellite-to-satellite differences.  In this paper, we first start by presenting and explaining the Nakar and Piran test, which following Band \\& Preece (2005) we extend by considering bursts in a plot of their $S_{bolo}$ versus $E_{peak,obs}$.  In addition, we explain why a certain amount of violators are expected, and what the observed distributions of bursts tell us to expect.  We follow this by showing gathered data from various detectors and providing a comprehensive examination of how each detector's data performs under the Nakar and Piran test.  Following this, we provide an explanation for why the vast majority of the data sets have too many violators of the Amati limit, and therefore the Amati relation is not good as a luminosity relation.  Finally, we examine several sources of systematic offsets that are actually the cause of the Amati relation in the first place, which only further condemns the Amati relation's usefulness. ",
        "conclusions": "The $S_{bolo} - E_{peak,obs}$ diagram has two limit lines, where bursts cannot be below that line if the Amati or Ghirlanda relation holds.  Actually, with the fairly large total uncertainties, substantially larger than the simple measurement errors quoted in the literature, we can expect nearly half of the bursts to be scattered below the Amati limit line.  So a simple test of the Amati relation is whether the {\\it average} burst falls below the Amati limit.  (This is similar to the original test proposed by Nakar \\& Piran, except that agreement with the Amati relation corresponds to about 40\\% violators.)  We apply this test to many burst samples.  The samples of early bursts with spectroscopic redshifts (as originally used to calibrate the Amati relation) pass our test, as does the sample of HETE bursts (even though the scatter about the Amati relation is unusably large).  All other satellites have a large fraction of violators far below the Amati limit line.  This is true whether we look at bursts with or without measured spectroscopic redshifts.  This constitutes a proof that the Amati relation could possibly apply, at best, to only a small and unrecognizable fraction of GRBs.  Indeed, the wide variations in distribution from detector to detector constitute a proof that selection effects must dominate the Amati relation.  We find four selection effects restrict the distribution on all sides.  The best known detector selection effect is the trigger threshold, which produces a roughly horizontal and fuzzy cutoff.  A more subtle and more restrictive selection effect is that for an $E_{peak,obs}$ value to be reported, the burst must be brighter than some threshold, with this threshold rising fast with increasing $E_{peak,obs}$.  These two detector selection effects will cut out bursts that are some combination of faint and hard, with these effects changing greatly from detector to detector.  The third and fourth selection effects operate to restrict the burst population as it appears in the sky.  The third selection effect is that bursts have a log-normal distribution of $E_{peak,obs}$ with the mean value shifting to lower values for faint bursts.  This effect will also reduce the number of detectable bursts that are faint and hard.  The fourth selection effect is that bright bursts are much rarer than faint bursts, as quantified by the usual power-law $\\log [N(>P)]-\\log[P]$ curve.  The combination of the third and fourth effects means that the bright and soft bursts are  doubly-rare, so that the upper-left side of the $S_{bolo} - E_{peak,obs}$ diagram will be empty.  For a detector with a range of spectral sensitivity and a low detection threshold, the distribution in the $S_{bolo} - E_{peak,obs}$ diagram will extend relatively low, with a large fraction of violators below the Amati limit (like for BATSE).  For a detector with a low energy range of sensitivity and a low detection threshold, the cutoff will be a diagonal line just below the Amati limit.  When combined with the paucity of bright-soft bursts in the GRB population (i.e., those above the Amati limit line), we have a combined selection effect that picks out bursts near the Amati limit.  Such a burst sample would then appear to follow the Amati relation.  Thus, the very strong selection effects for the early bursts with spectroscopic redshifts will create the Amati relation without any need for a physical connection between the $E_{peak,obs}$ and $S_{bolo}$.  That is, the Amati relation is not real, but its appearance in some data sets is simply a result of various selection effects by the detectors and within the GRB population.  With these strong results, the Amati relation should clearly not be used for purposes of cosmology, as has been previously done by many groups.  In particular, BATSE provides the strongest evidence for not using the Amati relation for cosmological purposes.  We note that our groups have never used the Amati relation for any cosmological purpose (e.g., Schaefer 2007; Xiao \\& Schaefer 2011).  Our community has expected a connection between $E_{peak}$ and energy for some time (e.g. Lloyd et al. 2000).  With the Amati relation clearly unsuitable, the other possible interpretation is non-physical correlation between $E_{peak}$ and $E_{iso}$.  As demonstrated above, this relation is extremely weak (e.g., see Table 1).  That is, for any given detector, some $\\langle\\log\\frac{E^{2.04}_{peak}}{S_{bolo}}\\rangle$ of observed bursts can be used as a sort of intercept for such a correlation with the given dispersion.  For some detectors, this empirical correlation can serve as a useful description of the data, and indeed this is what gave rise to the proposed Amati relation.  However, for many detectors, this interpretation is not useful, however, because the given dispersions are as wide as the data itself, so any such statements regarding a correlation between the two would be largely meaningless.  We emphasize that the failure of the Amati relation in no way carries any implications for any other GRB luminosity relation.  The fault of the Amati relation can be viewed as if it is merely a version of the Ghirlanda relation except that the beaming correction is unknown, so isotropic emission was assumed. The result, however, is that the Amati relation is biasing itself towards some average of whatever the beaming factors of the calibrating bursts are.  All the other GRB luminosity relations do {\\it not} involve beaming corrections, and the known physics of the beaming is already accounted for in the physics derivations of these laws.  The Ghirlanda relation is in essence just a conservation of energy statement, while the other luminosity relations (all involving the peak flux, not the fluence) just involve relativistic effects in the visible region of colliding jets.  Indeed, most of the other GRB luminosity relations were {\\it predicted} from the physics and then later confirmed.  In all, the failure of the Amati relation is zero evidence for the validity of the other relations (many of which were confirmed predictions) and there are good physical reasons to know that they are valid physical laws for GRBs."
    },
    "1112/1112.6174_arXiv.txt": {
        "abstract": "We examine the hypothesis that plasma associated with ``Type II'' spicules % is heated to coronal temperatures, and that the upward moving hot plasma constitutes a significant mass supply to the solar corona. % 1D hydrodynamical models including time-dependent ionization are brought to bear on the problem. These calculations indicate that heating of field-aligned spicule flows should produce significant differential Doppler shifts between emission lines formed in the chromosphere, transition region, and corona. % At present, observational evidence for the computed 60-90 \\velu{} differential shifts is weak, but the data are limited by difficulties in comparing the proper motion of Type-II spicules, with spectral and kinematic properties of  associated transition region and coronal emission lines. % Future observations with the upcoming IRIS instrument should clarify if Doppler shifts are consistent with the dynamics modeled here.  % ",
        "introduction": "\\label{sec:introduction}  Traditionally, research on the ``coronal heating problem'' has focused on the transport and dissipation of non-radiative energy directly into plasma at heights typical of the solar corona \\citep[e.g.][]{Kuperus+Ionson+Spicer1981,Parker1994,Walsh+Ireland2003}, 2Mm and more above the visible solar surface. At a first glance, this seems odd, since the corona is a tenuous structure fed by mass, momentum and energy from below. Indeed, processes operating solely in coronal plasma cannot actually produce a corona unless one already exists. Yet  five decades of detailed spectral analysis from space experiments have implied that there is a significant net downflow of energy in the form of heat conduction and enthalpy from the corona towards the underlying chromosphere \\citep[e.g.][]{Mariska1992}. The evidence in support of this picture includes agreement between emission measures, above $10^5$ K,  from models dominated by (downward-directed) heat fluxes, and the observation that the profiles  of transition region lines, while highly variable, are nevertheless mostly red-shifted. In addition, much magnetic (free) energy can be stored in the coronal volume in the form of current systems and/or waves. Thus the community has become accustomed to working largely within the picture of {\\em in situ} dissipation of mechanical energy directly within the coronal plasma.  However, observations from the Solar Optical Telescope (SOT, \\citealp{Tsuneta+others2008}), mounted on the stable, seeing-free platform of the Hinode spacecraft \\citep{Kosugi+others2007} have revived some earlier ideas concerning the coupling of the corona to the underlying chromosphere. Data obtained and analyzed by \\citet{dePontieu+others2007} have produced quantitative information about two populations of ``spicules'', jets of material moving supersonically above the limb and into the corona, something that was not achievable with the previous generation of ground-based observations \\citep[e.g.][]{Roberts1945,Beckers1968, Beckers1972}. De Pontieu and colleagues identified a class of finely structured spicules with apparent motions far faster ($\\sim 100$ \\velu) than previously measured ($\\sim 20-30$ \\velu{}). These ``Type-II'' spicules, unlike their longer lived brethren, disappear in a fashion suggesting that some of the cool material is heated during the spicule's lifetime. Statistical correlations of these new spicules with coronal counterparts have since been obtained using data from the Hinode, TRACE, STEREO, and SDO spacecraft \\citep{dePontieu+others2009,McIntosh+DePontieu2009a,McIntosh+others2010,dePontieu+others2011}.  This work suggests that significant energy deposition occurs in the plasma associated with spicular events.  Recent MHD simulations have provided one scenario in which Lorentz forces and Joule heating on a wide range of different field lines can explain some of the observed phenomena \\citep{Martinez-Sykora+others2011}. % Here we test the possibility that the heating and acceleration of plasma occurs in a field-aligned flow, the freshly heated spicular material filling the overlying atmosphere with material heated to higher, perhaps coronal, temperatures. In other words, we test whether a simple hydrodynamic approach can reproduce the observed phenomena. These studies are of interest because  the association of hot plasma with spicules suggested that ``energy deposition at coronal heights cannot be the only source of coronal heating'' \\citep{dePontieu+others2009}.  Ideas along these lines have been around for some time. \\citet{Thomas1948a}, interpreting Roberts' spicules as hydrodynamic jets, noted that ``the directed mechanical energy of the jet becomes converted to random thermal energy, in part, as the jet moves though the atmosphere''. \\citet{Miyamoto1949a} studied the viscous dissipation of kinetic energy of spicular motions. But with the observations then available, the 30 \\velu{} speeds of spicules was insufficient to produce coronal temperatures. \\citet{Pneuman+Kopp1978} estimated mass fluxes from available data, which with mass conservation, prompted them to suggest that ``the observed [transition region] downflow represents spicular material returning to the chromosphere after being heated to coronal temperatures''. \\citet{Athay+Holzer1982} noted that ``spicular material is raised well above the height that would be achieved by a projectile of the same initial velocity, thereby obtaining gravitational potential energy much in excess of its initial kinetic energy''. Athay and Holzer concluded that ``if sufficient heat is added to spicules, in conjunction with their acceleration, the spicule phenomenon may also play a major role in the production and maintenance of much of the solar corona.\" In a paper entitled ``the coronal heating paradox'', \\citet{Aschwanden+others2007}, proposed that the the phrase ``coronal heating problem'' is a misnomer. Instead, they argue that coronal heating should be considered as two processes, a ``chromospheric heating problem'' and ``coronal loop filling process'', moving the problem of energy dissipation into chromospheric plasma. \\citet{Hansteen+others2010} present sophisticated 3D radiation MHD calculations from sub-photosphere to corona and conclude that, for calculations driven by the work done on photospheric fields by convection, and dissipated by Ohmic heating, ``the heating per unit mass [\\ldots] necessarily is concentrated toward the transition region and low corona''.   Before proceeding, we clarify some terminology, following \\citet{Athay+Holzer1982}. Phrases such as ``heating in the chromosphere'', ``heating in the transition region'', as well as ``heating in the corona'' abound.  We will avoid this usage, as these phrases seem to imply that there is something about the thermal properties of these plasmas (e.g. ionization state, conductivities) that delineate different mechanisms affecting the three regimes.  This may well be the case, but few studies deal with such differences (one example that does is the work of \\citealp{Goodman2004}). Also, this language suggests, perhaps subliminally, that energy transport mechanisms must be found that fit in a pre-existing thermal structure. Instead we see the chromosphere, transition region, and corona, including the associated zoo of observed phenomena we call spicules, fibrils, explosive events, sprays, blinkers, etc, as {\\em observational manifestations} of the transport and dissipation of mass, momentum and non-radiative energy throughout the Sun's atmosphere. When viewed in a more physical sense, one avoids difficulties arising from observationally-based language use, such has, how can so much heating appear to be confined to the thin solar transition region? How can an observational phenomenon like a spicule can be considered as a source of mechanical heating? Why need one consider coronal heating to be paradoxical at all? ",
        "conclusions": "Our 1D hydrodynamic calculations attempt to capture the evolution of the emitted spectrum of a plug of accelerated and heated plasma flowing along magnetic field lines extending upwards from the upper chromosphere.  The dynamics and emergent spectra are  computed using full non-equilibrium in the atomic rate equations. The calculations reveal that the coronal emission which accompanies the heating needed to produce coronal plasma from the cool spicular plasma is blueshifted by $\\sim 50-90$ \\velu{} relative to the chromospheric emission (compare lines of He~I, C II Si II with the coronal lines in figs.~\\pref{fig:profiles} \\pref{fig:profiles1}). This is because, in the hydrodynamic regime, the heating produces an over-pressure and a subsequent expansion of the plasma into the corona.  Thus, if indeed the chromospheric plasma in Type II spicules is injected into the corona, at the same time being heated in this fashion, then we must expect the velocity distributions of chromospheric and coronal lines to differ. Accompanying the prediction of Doppler shift differences, changes in the intensities of line ratios that are traditionally viewed as being sensitive to electron temperature, are also predicted in our rapidly heated outflow models. These intensity changes result from the slowness of ionization/recombination processes relative to dynamical timescales. It would seem worthwhile to investigate atomic systems of the Li and Na-like isoelectronic sequences observationally, following earlier work \\citep{Heroux+others1972} as well as other atomic systems with temperature  sensitive line ratios \\citep[e.g.][]{Pinfield+others1999}.    Such a systematic difference between cool and hot flows associated with Type II spicules has not been reported in the observational literature, but we note that the observations have significant uncertainties. The cool components of Type II spicules at the limb show apparent velocities of order 50-100 \\velu{} \\citep{dePontieu+others2007}, with rapid blueshifted events, thought be their disk counterparts, showing line-of-sight velocities of order 50 \\velu{} \\citep{Rouppe+others2009} and apparent velocities of order 75 \\velu{} \\citep{dePontieu+others2011}. The velocities of coronal counterparts are also difficult to determine precisely. They are derived from weak blueshifted emission components, seen at the base of much stronger emission cores, near the limit at which EUV spectrographs can operate (e.g., instrumental broadening, signal-to-noise, ...). These measurements suggest line-of-sight velocities of order 50-150 \\velu{} for the coronal counterparts of spicules. Because both the chromospheric and coronal measurements are significantly impacted by viewing geometry, radiative transfer and instrumental limitations, it is not yet clear that the observed velocity distributions are sufficiently well determined to reveal a velocity difference of 50 \\velu{} between the chromosphere and coronal counterparts of spicules.  If the large relative shift in the Doppler motions predicted here are not confirmed by future high resolution measurements (such as those from the IRIS mission\\footnote{http://iris.lmsal.com/index.htm}), we would conclude that the suggestion put forward by \\citet{dePontieu+others2007,dePontieu+others2009,dePontieu+others2011}, namely the close correspondence of velocities from spicules with those of Doppler shifted lines in the transition region and corona, cannot credibly be accounted for in the type of 1D hydrodynamic flow modeled here. One possible resolution might simply be that the magnetic fields within the tube are significantly twisted such that the spicular plasma has an azimuthal as well as axial component as the Lorentz force directs the plasma along field lines.  In this case the Doppler shifts computed should be multiplied by the cosine of the magnetic pitch angle.  %  The recent MHD calculations of \\citet{Martinez-Sykora+others2011} might help resolve the problems identified here, in that the system modeled is far more complex than our simple calculations and does not seem to show such large velocity differences between chromospheric and coronal lines.  In their work, they find that chromospheric material is injected into the corona via, first, compression due to horizontal Lorentz forces associated with emerging flux, and second, via the field-aligned gas pressure gradient.  As the field-aligned flow progresses it is heated via Joule dissipation, in a fashion not dissimilar from the ad-hoc calculations we present.  However, in the modeled MHD system, the acceleration and heating acts, with varying strength, on plasma that occurs on an ensemble of neighboring field lines. As a result, the coronal emission % has contributions from many magnetic field lines with varying mixes of chromospheric and coronal plasma/flows.  Thus, the line profiles of coronal plasma must be carefully synthesized to find the relative contributions of these different components.  Notably though, the cool plasma that actually flows into the corona in these MHD calculations is indeed a field-aligned flow.  Thus, if our parametrized heating and acceleration profiles are representative of what occurs in the complex 3D configuration, the heated plasma in that part of the calculated flow would be expected to behave like the calculations presented in the present paper. Current EUV observations lack the spatio-temporal resolution to resolve the finely structured dynamics and energetics predicted by the model of \\citet{Martinez-Sykora+others2011}.    The potential problem of Doppler shift differences may be avoided in an entirely different scenario \\citep{Judge+Tritschler+Low2011},   where it was suggested that Type II spicules, in particular those with apparent upward velocities in excess of 50 \\velu{}, correspond not to jets but to the line-of-sight superposition of warped current sheets modulated by Alfv\\'enic\\footnote{Simply meaning motions dominated by the balance between magnetic tension and inertial terms.} fluctuations driven from below. Low-frequency Alfv\\'enic-type wave motions have been reported before in movies of Type II spicules \\citep[][]{DePontieu+others2007b}. Should Type II spicules be proven to be sheets, the calculations here are largely irrelevant, and the physical connection between the chromospheric, transition region and coronal emission proposed earlier would be called into question. But, as suggested by \\citet{Judge+Tritschler+Low2011} the 1D flow picture might  apply to a subset of observed spicules with plasma genuinely flowing at lower speeds. % However, several questions regarding the overall interpretation of spicule properties at the limb, their relation to RBEs,  and statistical analyses would then arise.   Whatever the outcome, we conclude that the role of Type II spicules in supplying mass and energy to the corona \\citep{dePontieu+others2007,dePontieu+others2009,dePontieu+others2011}, is a subject ripe for future study. Further observational work on spicule properties and their relationship to the properties of the associated hotter plasma seems warranted. The upcoming IRIS mission seems well positioned to address this question."
    },
    "1112/1112.4940_arXiv.txt": {
        "abstract": "The absolute neutrino mass scale is currently unknown, but can be constrained by cosmology. The WiggleZ high redshift, star-forming, and blue galaxy sample offers a complementary dataset to previous surveys for performing these measurements, with potentially different systematics from non-linear structure formation, redshift-space distortions, and galaxy bias. We obtain a limit of $\\sum m_\\nu < 0.60\\, \\eV$ ($95\\%$ confidence) for WiggleZ+Wilkinson Microwave Anisotropy Probe. Combining with priors on the Hubble parameter and the baryon acoustic oscillation scale gives $\\sum m_\\nu < 0.29\\, \\eV$, which is the strongest neutrino mass constraint derived from spectroscopic galaxy redshift surveys. ",
        "introduction": "Neutrinos are the lightest massive known particles, yet they are treated as exactly massless by the Standard Model of particle physics. {We know they have non-zero masses because neutrino oscillation experiments using solar, atmospheric, and reactor neutrinos have measured mass differences between the three species to be $\\Delta m_{32}^2 = 2.43\\times10^{-3}\\, \\eV^2$ and $\\Delta m_{21}^2 = 7.59\\times10^{-5}\\, \\eV^2$ \\cite{Fukuda:1998,Amsler:2008}. The Heidelberg-Moscow experiment has limited the mass of the electron neutrino to be less than $0.35\\, \\eV$ using $\\beta-$spectroscopy \\cite{Klapdor:2006}, but no current experiment has sufficient sensitivity to measure the absolute neutrino mass.  Massive neutrinos affect the way large-scale cosmological structures form by suppressing the gravitational collapse of halos on scales smaller than the free-streaming length at the time the neutrinos become non-relativistic. This leads to a suppression of the small scales in the galaxy power spectrum that we observe today, and consequently we can infer an upper limit on the sum of neutrino masses from the matter distribution of the Universe \\cite{Lesgourgues:2006}. Combining with the lower limit of $\\sum m_\\nu > ~0.05 \\eV$ provided by the mass differences from oscillation experiments, allows us to narrow the range of possible neutrino masses.  The cosmic microwave background (CMB) provides an upper limit of $\\sum m_\\nu < 1.3\\, \\eV$ \\cite[][all limits are 95\\% confidence]{Komatsu:2010}. Combining with large-scale structure measurements such as the galaxy power spectrum \\cite{Reid:2010,Thomas:2010,Gonzalez-Garcia:2010,dePutter:2012}, galaxy luminosity function \\cite{Jose:2011}, cluster mass function \\cite{Mantz:2010,Benson:2011}, or the scale of baryon acoustic oscillations \\cite[BAO,][]{Reid:2010,Komatsu:2010} tightens the constraints to $\\sum m_\\nu \\lesssim 0.3\\, \\eV$ by breaking degeneracies with other parameters. Consequently neutrino mass constraints are important goals of current and future galaxy surveys such as e.g.~Baryon Oscillation Spectroscopic Survey \\cite{Eisenstein:2011}, Dark Energy Survey \\cite{Lahav:2010} and Euclid \\cite{Euclid}. In this letter we use the galaxy power spectrum from the WiggleZ Dark Energy Survey to constrain the sum of neutrino masses.  The WiggleZ galaxy survey has several complementary aspects and potential advantages over previous surveys: \\\\1) The neutrino suppression of the galaxy power spectrum is degenerate with effects from non-linear structure formation. Non-linearities increase with time so for the distant galaxies probed by WiggleZ, the contamination from non-linearities is smaller than for previous surveys. This is illustrated in \\figref{Pkratio} where we show the ratio between a simulated WiggleZ power spectrum and the linear power spectrum for $z=0.2$ (dashed blue) and $z=0.6$ (solid black). For comparison we also show the ratio for simulated highly-biased massive haloes at $z=0.2$ (dotted red). \\\\2) The relationship (bias) between the observed galaxy distribution and the dark matter distribution, which is influenced by massive neutrinos, depends on the observed galaxy type. Previous studies \\cite[e.g.][]{Percival:2009,Reid:2010} measured red galaxies, which tend to cluster in the centers of dark matter halos, whereas the star-forming blue WiggleZ galaxies avoid the densest regions. This leads to a lower overall bias, which makes WiggleZ less susceptible to any possible systematics that could arise from a scale-dependence of the bias.  \\begin{figure} \\includegraphics[width=0.99\\columnwidth]{./1_pkratio.eps} \\caption{The ratio between the simulated WiggleZ halo power spectrum (from GiggleZ) and the corresponding linear power spectrum (normalised to the first bin) at $z=0.2$ (blue dashed) and $z=0.6$ (black solid) and for luminous red galaxies (dotted red). The vertical lines indicate our fitting range of  $k=0.02-0.3\\,h\\,\\Mpc^{-1}$. Non-linear corrections are clearly less significant for the high-redshift, low-bias WiggleZ halos than at lower redshifts.} \\label{fig:Pkratio} \\end{figure}  Galaxy redshifts are not entirely determined by the Hubble flow, but also by their peculiar motions (redshift-space distortions) providing a challenge when comparing observations (redshift-space) with theory (real-space). Through exhaustive tests using numerical dark matter simulations of the WiggleZ survey, we demonstrate the breakdown of common models at small scales, and calibrate a new non-linear fitting formula. ",
        "conclusions": "\\label{sec:results} When fitting the observed WiggleZ power spectra we obtain the results presented in \\figref{kmaxmnuchi2}. The upper panel shows the $\\chi^2$ as a function of $\\kmax$ for the best fitting parameter values for each of the six approaches, and the lower panel shows the corresponding neutrino mass constraints. Although all models produce similar $\\chi^2$ values, our comparison with the full $N$-body simulation catalogue (\\figref{GiggleZchi2}) revealed that systematic errors arise when models A) to E) are fit across the range of scales $\\kmax < 0.3\\,h\\,\\Mpc^{-1}$. Using the fully-calibrated model F), we obtain $\\sum m_\\nu < 0.60\\, \\eV$ for WMAP+WiggleZ with $\\kmax=0.3\\,h\\,\\Mpc^{-1}$. Combining with BAO+$\\Hub$ reduces the uncertainty in $\\Omega_m$ and $H_0$, leading to stronger neutrino mass constraints. Without WiggleZ, the WMAP+$\\Hub$+BAO dataset gives $\\sum m_\\nu < 0.55\\, \\eV$ whereas combining with WiggleZ adds information about the power spectrum tilt ($n_s$). The resulting neutrino mass constraint is  $\\sum m_\\nu < 0.29\\, \\eV$ for model F) and $\\kmax=0.3\\,h\\,\\Mpc^{-1}$.  \\begin{figure} \\includegraphics[width=0.99\\columnwidth]{./4_mnu.eps} \\caption{{\\it Upper:} Reduced $\\chi^2$ as a function of $\\kmax$ for each of the six approaches. {\\it Lower:} Upper limits on $\\sum m_\\nu$ as a function of $\\kmax$. The models are: A) blue dotted, B) green dashed, C) magenta dot-dashed, D) cyan triple dot-dashed, E) red long dashed, F) thick orange solid.. The dashed grey line is the lower limit from oscillation experiments, and the black lines are upper limits from WMAP+BAO+$\\Hub$ (dotted) and WiggleZ+WMAP+BAO+$\\Hub$ (solid).} \\label{fig:kmaxmnuchi2} \\end{figure}  The relative probability distributions of $\\sum m_\\nu$ for model F) with $\\kmax=0.3\\,h\\,\\Mpc^{-1}$ are shown in \\figref{prob}. It is clear how adding WiggleZ data to the fit narrows the distributions (dotted to solid) both with (orange) and without (black) the inclusion of BAO+$\\Hub$. This is the strongest neutrino mass limit so far derived from spectroscopic redshift galaxy surveys. The advantages of WiggleZ are a higher redshift for which the structure formation is linear to smaller scales, and a simple galaxy bias for the strongly star-forming blue emission line galaxies.   \\begin{figure} \\includegraphics[width=0.99\\columnwidth]{./5_prob.eps} \\caption{The relative probability distribution of $\\sum m_\\nu$ from fitting model F) with $\\kmax=0.3\\,h\\,\\Mpc^{-1}$ for WMAP (dotted orange), WiggleZ+WMAP (solid orange), WMAP+BAO+$\\Hub$ (dotted black) and WiggleZ+WMAP+BAO+$\\Hub$ (solid black). The dashed grey line is the lower limit from oscillation experiments, and the vertical lines are 95\\% confidence upper limits.} \\label{fig:prob} \\end{figure}  Our result is comparable to that obtained using photometric redshift galaxy surveys \\cite[$\\sum m_\\nu < 0.28\\, \\eV$, ][]{Thomas:2010,dePutter:2012}, but the systematics in the two data set are completely different. For example, imaging surveys are potentially susceptible to systematic errors from the imprint of stars on the selection function \\cite{Ross:2011} and the shape of the redshift distribution. WiggleZ contains negligible star contamination and much higher redshift resolution compared to photometric surveys. The high redshift and blue galaxies of WiggleZ allow us to fit the power spectrum to smaller scales than previous surveys (both spectroscopic and photometric), where the effect of the neutrinos is larger, and get a similar neutrino mass constraint from a smaller, but well understood galaxy sample. Also the result from galaxy clusters \\cite[$\\sum m_\\nu < 0.33\\, \\eV$, ][]{Mantz:2010,Benson:2011} is similar, but with different systematics.  Since the data sets are all independent, they can potentially be combined in the future to provide even stronger constraints. This is particularly interesting in light of recent results  \\cite{Komatsu:2010,Kopp:2011} that hint at the existence of additional neutrino species ($N_{\\mathrm{eff}}>3.04$). Allowing for additional neutrino species degrades the constraining power of large scale clustering alone, and the combination of $N_\\mathrm{eff}$ and $\\sum m_\\nu$ is therefore poorly constrained with current data.  In the future, galaxy surveys such as the Baryon Acoustic Oscillation Survey, Dark Energy Survey and Euclid will be far more sensitive, giving cosmological neutrino mass constraints of order $\\sum m_\\nu < 0.05-0.1\\, \\eV$ \\cite{Carbone:2011,Carbone:2011b}. This will be small enough to distinguish between the ordering of the neutrino masses (normal hierarchy where $m_1<m_2<<m_3$ or inverted where $m_3<<m_1<m_2$). However, as demonstrated in this paper, the small details of the modelling of non-linear effects become very important, so robust modelling either theoretically or calibrated to simulations with massive neutrinos will be necessary."
    },
    "1112/1112.1081_arXiv.txt": {
        "abstract": "Numerical relativity simulations predict that coalescence of supermassive black hole (SMBH) binaries not only leads to a spin flip but also to a recoiling of the merger remnant SMBHs. In the literature, X-shaped radio sources are popularly suggested to be candidates for SMBH mergers with spin flip of jet-ejecting SMBHs. Here we investigate the spectral and spatial observational signatures of the recoiling SMBHs in radio sources undergoing black hole spin flip. Our results show that SMBHs in most spin-flip radio sources have mass ratio $q\\ga 0.3$ with a minimum possible value $q_{\\rm min} \\simeq 0.05$. For major mergers, the remnant SMBHs can get a kick velocity as high as $2100\\, {\\rm km~s^{-1}}$ in the direction within an angle $\\la 40^\\circ$ relative to the spin axes of remnant SMBHs, implying that recoiling quasars are biased to be with high Doppler-shifted broad emission lines while recoiling radio galaxies are biased to large apparent spatial off-center displacements. We also calculate the distribution functions of line-of-sight velocity and apparent spatial off-center for spin-flip radio sources with different apparent jet reorientation angles. Our results show that the larger the apparent jet reorientation angle is, the larger the Doppler-shifting recoiling velocity and apparent spatial off-center displacement will be. We investigate the effects of recoiling velocity on the dust torus in spin-flip radio sources and suggest that recoiling of SMBHs would lead to ``dust poor'' AGNs. Finally, we collect a sample of 19 X-shaped radio objects and for each object give the probability of detecting the predicted signatures of recoiling SMBH. ",
        "introduction": "In the hierarchical galaxy formation model of the $\\Lambda$CDM cosmology, massive galaxies are produced by successive galaxy mergers \\citep{KH00,spr05}. During the merging of galaxies, the supermassive black holes (SMBHs) at centers quickly form hard supermassive black hole binaries (SMBHBs) at a separation of {\\it pc}-scale, and then may stall for a timescale longer than Hubble time, if { the gravitational potential at galactic nucleus is spherical and stellar relaxation is dominated by two-body scattering} \\citep{beg80,yu02}. However, the hardening rate of SMBHBs can be boosted and SMBHBs become merged within Hubble time because of triaxial or rotating structures, massive perturbers, or massive gas disks at galactic centers \\citep[][and references therein]{mm05,col11}. The strong gravitational wave (GW) radiations from the coalescence of SMBHBs are the main targets of the GW detector Laser Interferometer Space Antenna (LISA) and the GW detecting program Pulsar Timing Array (PTA).  Spatially resolved dual or binary active galactic nuclei (AGNs) at {\\it kpc}-scale separation were discovered in the past decade first in a few merging galaxy systems \\citep[e.g.][]{kom03,kom06,gre10}, and recently in a large sample of AGNs showing double-peak narrow emission lines \\citep{zho04,XK09,com09,wan09,liu10a,smith10,liu10b, fu10,shen11,rosario11}. Meanwhile, more compact or hard SMBHBs are extremely difficult to observe with present telescopes \\citep[see][for the only pc-scale SMBHB candidate so far discovered by VLBA]{rod06}, and for sub-pc SMBHBs only indirect observational evidences are suggested in a few sub-classes of AGNs with peculiar radio morphologies \\citep{beg80,liu04,liu07}, periodic variabilities \\citep{sil88,val00,liu95,liu97,liu06,liu02}, or special broad-emission-line features \\citep{SF91,gas96,bor09}. Although it is still unclear how SMBHBs can pass through the final pc-scale and become merged, some observations of AGNs and QSOs are suggested to be due to the coalescence of SMBHBs, e.g. the interruption and restarting of jet activities in double-double radio galaxies \\citep{liu03}, and the rapid jet reorientation in X-shaped radio sources \\citep{mer02,zie05}. For SMBHBs in quiescent galaxies some observational signatures have been predicted \\citep{che08,che09,che11,KM08b,liu09} but need to be confirmed observationally. Among all the observational evidences for the coalescence of SMBHBs in AGNs, X-shaped radio source is one of the most popular. X-shaped radio sources are a subclass of double-lobed radio sources, consisting of a pair of long and diffusive inactive ``wings'' oriented at a large angle relative to the active lobes \\citep{lw84,den02}. The ``spin flip model'' suggests that the wings are the fossil plasma of old lobes and jets are misaligned with the wings because of swift jet reorientation due to the spin-flip of jet-ejecting SMBHs during the coalescence of SMBHBs \\citep{mer02}. Meanwhile, many other formation scenarios, which do not require binary black hole coalescence, are also proposed in the literature, including merger of two jet-emitting AGNs \\citep{lal07}, hydrodynamic interaction between jet and intra-/inter-galactic medium \\citep{wir82,lw84,wor95,cap02,zie05,SS09,HK10}, and precession of jet  axis \\citep{eke78,beg80,liu04}.  Numerical simulations show that the coalescence of two black holes not only leads to a spin-flip { black} hole but also to a kick of the post-merger black holes due to the anisotropic emissions of gravitational waves. Therefore, it is expected that the coalescence of SMBHBs in spin-flip radio sources should lead to the recoiling of post-merger SMBHs surrounded by accretion disk. The recoiling velocity depends on the angular momentum and mass ratio of the pre-merger SMBHs, and can be as large as $175\\, {\\rm km~s^{-1}}$ for Schwarzschild black holes \\citep{gon07},  $3800\\, {\\rm km~s^{-1}}$ for extremely spinning equal-mass black hole binaries with anti-aligned spin angular momenta in the orbital plane \\citep{pre05,her07,bak06,bak08,cam06,cam07a,cam07b}, or even $10^4 \\, {\\rm km~s^{-1}}$ in hyperbolic encounters \\citep{hea09}. The gravitational recoiling SMBHs in AGNs may lead to signatures such as kinematic and spatial offsets \\citep{bon07,kom08,shi09,civ10,rob10,tsa11,era11,MQ04,mag05, loe07,bat10,jon10,BL08,gue09,civ10,gue11,ble11}, along with some other electromagnetic signatures {\\citep[see][for a review]{kom10}}.  In the following, we refer to the radio sources with spin-flip signatures, such as those in X-shaped radio sources, as ``spin-flip radio sources''. In this paper, we adopt the spin-flip model and investigate the spectral and spatial signatures of the recoiling SMBHs in such objects, using the empirical formula obtained in numerical relativity to calculate the spin angular momentum and recoiling velocity of the post-merger black hole. We neglect the slow precession of the spin axes of SMBHs during the in-spiraling of SMBHBs, We assume that the old and active jets, respectively, form along the spin axes of the primary and post-merger SMBHs and any jet-ejecting SMBH has spin parameter $a=c S / G M^2 \\geq 0.9$, where $S$ is the black hole spin angular momentum, $M$ is the black hole mass, $G$ is the gravitational constant, and $c$ is the speed of light. For the secondary SMBH, no constraint on the spin parameter $a_2$ can be given. We assume that $a_2$ can be any value between 0 and 1 of flat probability distribution. Our results show that the SMBHBs producing detectable spin-flip angles should form in major mergers with mass ratio $\\ga 0.2 - 0.3$ and that the recoiling SMBHs preferentially move in a direction with an angle $\\la 40^\\circ$ relative to jet orientation. We calculate the Doppler kinematic shift of broad emission lines, the AGN off-center displacement in host galaxy, and the effects of recoiling SMBHs on dust torus in spin-flip radio galaxies and quasars as a function of apparent jet reorientation angle and jet size. We finally collect a sample of 19 X-shaped radio objects and calculate the detection probabilities for the proposed observational signatures.  The paper is organized as follows. In Section~\\ref{sec:method}, we introduce the empirical formula for the spin angular momenta and the recoiling velocity of post-merger SMBHs obtained in numerical relativity. Section~\\ref{sec:results} describes our numerical simulations and { general results for spin-flip radio sources}. We collect a sample of 19 X-shaped radio objects and calculated the observational signatures of recoiling SMBHs in Section~\\ref{sec:sample}. Our discussions and conclusions are given in Section~\\ref{sec:dc}. Throughout the paper, we assume a $\\Lambda$CDM cosmology with parameters $H_0 = 73\\, {\\rm km~s^{-1}~Mpc^{-1}}$, $\\Omega_{\\rm \\Lambda}  = 0.73$, and $\\Omega_{\\rm m}  = 0.27$. ",
        "conclusions": "\\label{sec:dc}  A popular model in the literature suggested that the X-shaped radio sources form because of jet reorientation following the spin flip of SMBHs at SMBHB mergers. Numerical relativity suggested that the black hole merger not only leads to a spin flip but also to a gravitational kick of the post-merger black holes. In this paper, we investigated with Monte Carlo simulations the distribution function of the spin-flip angles and recoiling velocities of the post-merger SMBHs, and the detectability of recoiling SMBHs, in the forms of Doppler-shifted spectral lines and off-center AGNs, in radio sources showing spin-flip signatures  {of jet reorientation}. We also collected a sample of 19 X-shaped radio objects and calculated for all the sample objects the spectral kinematic offsets $V_{\\parallel}$, the apparent spatial off-center displacements $L_\\bot$, and the degree of dust-torus truncation because of the gravitational kick.  Because the spin flip and recoiling velocity strongly depend on the parameters $q$, ${\\bf a}_1$, ${\\bf a}_2$, and $\\bf{e_s}$ of SMBHB system, we first investigated the distribution of mass ratio $q$ in spin-flip radio sources with given apparent jet reorientation angles, assuming that the spin parameter $a$ of jet-ejecting SMBHs has $a\\ge 0.9$. Our results indicate that to cause significant apparent jet {reorientation} the mass ratio cannot be smaller than a minimum value $q_{\\rm min} \\sim 0.1$. The larger the apparent jet reorientation angle is, the larger the minimum mass ratio will be. For spin-flip radio sources with apparent jet reorientation angle $15^{\\circ}<\\delta<40^{\\circ}$, $q_{\\rm min} \\sim 0.05$ and the most probable mass ratio is $q \\simeq 0.3$, while for spin-flip radio sources with $65^{\\circ} < \\delta < 90^{\\circ}$, $q_{\\rm min} \\simeq 0.15$ and the most probable mass ratio is $q \\ga 0.6$. Our results suggest that if  {an observed jet reorientation} is indeed due to the spin flip following a SMBHB merger, the merger must be major merger with $q\\ga 0.2-0.3$. If the spin-flip model for radio sources  {of jet reorientation} could be confirmed, our results suggest that we can measure the distribution of SMBHB mass ratio in luminous radio sources and give statistic constraints on the hierarchical galaxy formation model and the growth history of SMBHs.  {Having} the mass ratio distribution, we investigated the distributions of spin flip angle and recoiling velocity in spin-flip radio sources. Our results showed that the larger the SMBHB mass ratio is, the larger the the most possible spin flip angle and the recoiling velocity will be. For typical major merger mass ratio $q=0.3$, $q=0.7$ and $q=1.0$, the most possible spin flip angles are moderate and increase with mass ratio from $\\Delta = 19^{\\circ}$ for $q=0.3$, and $25^{\\circ}$ for $q=0.7$ to $30^{\\circ}$ for $q=1$. The spin flip angle is hardly ever larger than $\\Delta \\sim 55^{\\circ}$. Our results imply that any apparent jet reorientation angle $\\delta \\ga 60^{\\circ}$ is due to projection effect. For $q = 0.3$, the recoiling velocity ranges from about zero to larger than $1200\\, {\\rm km~s^{-1}}$ with a most possible velocity $V \\sim  600\\, {\\rm km s^{-1}}$. While for $q \\ga 0.7$, the distribution of recoiling velocity is similar and ranges from about $0$ to larger than $2100\\, {\\rm km~s^{-1}}$ with most probable velocity $\\hat{V}\\sim 900\\, {\\rm km~s^{-1}}$. For major mergers with $q\\ga 0.3$, almost all recoiling SMBHs move nearly along the final jet orientation within a angle of $40^\\circ$. Alignment  {of} spin axes with orbital angular momentum of SMBHB prior to  {the black hole} coalescence would significantly reduce the recoiling velocity and spin-flip angle, but tighten the correlation between the orientations of final spin and recoiling velocity. Our results suggest that to efficiently detect recoiling SMHBs with Doppler shifted spectral broad emission lines, one should preferentially observe  {nearly face-on AGNs, e.g.} radio loud quasars or blazars among which the broad emission lines are nearly maximumly Doppler-shifted for a given recoiling velocity. This conclusion is also consistent with present observations in the literature that all candidates for recoiling SMBHs are bright quasars and the measured Doppler-shifting recoiling velocity is preferentially very larger \\citep[e.g.][]{kom08}. However, to detect recoiling SMBHs in radio galaxies, one should try to measure the apparent spatial off-center displacements of AGNs relative to the centers of the host galaxies because the recoiling SMBHs are moving nearly along the jet axes, which are at large angles with respect to LOS. Even in some quasars, recoiling SMBHs may be discovered by measuring the apparent spatial off-center displacements because some radio galaxies and Type II quasars may  {transit} to Type I quasars due to the recoiling oscillations as suggested by \\citet{KM08a}. For minor mergers with $q\\sim 0.1$, both the spin-flip angle and the recoiling velocity narrowly distribute with peaks at the most possible angle $\\Delta \\sim 12^\\circ$ and velocity $V\\sim 150\\, {\\rm km~s^{-1}}$.  In order to use the observed jet reorientation angles in spin-flip radio sources to constrain the observable spectral kinematic offsets $V_\\parallel$ and apparent spatial off-center displacements $L_\\bot$, we projected the spin-flip angles and the recoiling velocities along and vertical to LOS and investigated their correlations. An analytical 1-D model is used to calculate the deceleration of a recoiling SMBH due to dynamical friction, and the elapsed time since the black hole gets kicked is estimated from the size of the radio morphology. In general, $V_\\parallel$ in spin-flip radio sources is smaller than $2000~\\kms$ and $L_\\bot$  {hardly ever} exceeds 2 kpc. However, we should note that the conclusion is valid only for quasars with SMBHBs of circular orbits. For two SMBHs merging along hyperbolic orbits, it is possible to have a recoiling velocity significantly larger than $2000~\\kms$ \\citep{hea09}. Our results indicate that in radio sources showing spin-flip signatures  {of jet reorientation}, detectable Doppler-shifted broad emission lines are mainly contributed by GPS and CSS sources with large apparent spin-flip angles. The detection rate of Doppler-shifted broad emission lines in a sample of GPSs/CSSs with $\\delta>40^\\circ$ could be as high $50\\%$. For GRSs, deceleration due to dynamical friction has significantly reduced $V_\\parallel$, so the probability of detecting Doppler-shifted broad emission lines is typically $\\la10\\%$. Detecting AGN off-center displacement requires astrometric observations with high angular resolution.  {For current telescopes with limit resolution up to} $10$ mas, detection of off-center displacement would be limited to the GPS/CSS sources within a distance of 100 Mpc.  {However,} future interferometer missions with sub-mas resolution  {or higher}, such as Gaia, would allow detection of off-center AGNs  {in GPS, CSS, and GRSs at} a much larger distance. Our results suggest that to efficiently detect the recoiling SMBHs in spin-flip radio sources, one should select a sample of {\\rm GPS/CSS} sources with apparent jet reorientation angles $\\delta \\ga 40^\\circ$.  When the SMBH in an AGNs gets kicked, the outer part of the dust torus in AGN unification model may become unbound and the dust torus is truncated, leading to the transition of the AGN from dust rich to ``dust poor'' or even to ``dust free''. When a FRII radio galaxy or an obscured quasar becomes ``dust free'', it should transform to an unobscured broad emission line quasar. We calculated the minimum recoiling velocities for partial or full truncation of dust torus and estimated their detectability in typical spin-flip radio sources with different apparent jet reorientation angles $\\delta$. Our results suggest that about $11\\%$ spin-flip radio galaxies with $\\delta \\ga 40^\\circ$ and about $1\\%$ with $\\delta \\la 40^\\circ$ will become ``dust free'' and broad emission line AGNs, if broad line regions are present. Our results imply that quasars with larger bolometric luminosity and smaller black hole mass should have higher probability to become ``dust free''. The velocity for dust torus to be partially truncated is typically lower than the most-probable recoiling velocity, therefore most X-shaped radio sources should be ``dust poor''. Quasars may become weak emission line objects, if the broad emission regions are partially truncated due to gravitational recoiling. This is possible for a quasar with small black hole mass ($\\sim10^6~\\msun$) radiating at the Eddington limit. We note that without gas replenishing the accretion disks, the dust-free and weak-line quasars formed by recoiling SMBHs are relatively short-living ($\\la1$ Myr), because the combination of high recoiling velocity ($\\ga10^3~\\kms$), small black hole mass ($\\la10^7~\\msun$), and large bolometric luminosity ($\\ga10^{45}~{\\rm erg~s^{-1}}$) will significantly reduce the lifetimes of the accretion disks around the recoiling black holes \\citep[e.g.][]{loe07}. This argument implies that among X-shaped radio sources whose host  galaxies are typically gas-poor elliptical, the ones with long active jets ($\\ga 100$ kpc or $t_a \\ga 3\\times 10^6$ yr for a typical advancing speed $0.1c$ of hot spot in FRII radio galaxies) and small black hole masses ($\\la 10^7~\\msun$) are unlikely to host dust-free and weak-line quasars.  We then selected a sample of 19 X-shaped radio objects, consisting of 18 X-shaped radio galaxies and 1 quasar, and for each sample object we calculated the distribution functions of the LOS recoiling velocity $V_\\parallel$ and the apparent spatial off-center displacement $L_\\bot$, as well as the critical recoiling velocities $V_{\\rm cr}$ and $V_{\\rm dp}$ for partial and complete truncation of dust torus. To do the calculations, we measured the apparent jet reorientation angles from the radio images and collected the central SMBH masses, the optical luminosities of AGNs, the viewing angles and dynamic ages of the active lobes from the literature. Our results showed that most X-shaped radio objects in the sample have negligible $V_\\parallel$ ($P_{v,\\parallel}<10\\%$), except for J0941-0143, 4C+01.30, J1348+4411, and J1614+2817. All the 15 sample objects with valid Monte Carlo solutions have significant probabilities ($>20\\%$) to show AGN off-center displacements when a spatial resolution of $0.1$ mas is considered. In particular, the radio galaxy J1348+4411 have appreciable probability to show both Doppler-shifted broad emission lines ($P_{v,\\parallel}\\simeq40\\%$) and off-center spatial displacement ($P_{\\theta,\\bot}\\simeq55\\%$), therefore would be intriguing target for follow-up observations. Our results also showed that the sample X-shaped radio objects have high probabilities to be ``dust-poor''. The probabilities for X-shaped radio galaxy J1348+4411 and quasar 4C+01.30 are always higher than $80\\%$, despite the uncertainties in torus size and bolometric correction factor. Because of the high probability of being both dust-poor and offset from galaxy center, these two objects would be interesting  targets for IR follow-up observations to test the dust-poorness. Besides, J1348+4411 and 4C+01.30 also have the highest probability in our sample to be ``dust-free'' and non-obscured. The fact that 4C+01.30 is already a quasar implies that this object could be an obscured AGN before binary black hole merger, but may have transformed into a quasar after recoil takes place. Detecting the predicted off-center displacement and Doppler shifted broad emission lines would strongly support that 4C+01.30 is a recoiling SMBH. For J1348+4411, the lack of broad line in this object  suggests that the recoiling velocity is toward the lower end of the predicted distribution function, but the probability of being dust-poor is still higher than $70\\%$.  The detection of the predicted observational signatures of recoiling SMBHs in X-shaped radio sources would give final confirmation of the spin-flip model for X-shaped radio sources. While, if none of the predicted signatures is detected in X-shaped radio sources, the SMBHB spin-flip model for the formation of X-shaped radio feature would be most likely excluded, unless our assumptions for the initial conditions of SMBHB mergers are severely unrealistic. However, this does not imply that X-shaped radio sources do not harbor SMBHBs at center, because jet reorientation may form due to the interaction of SMBHB and accretion disk \\citep{liu04}. According to the disk-binary interaction model, the final recoiling velocity of a SMBH is small due to the alignment between the black hole spin axes and the orbital angular momentum before the merger, therefore neither the spectral kinematic offset nor the spatial AGN off-center displacement are as large as those in the spin flip model."
    },
    "1112/1112.4171_arXiv.txt": {
        "abstract": "The spectrum of an isotropic extragalactic $\\gamma$-ray background (EGB) has been measured by the Fermi-LAT telescope at high latitudes. Two new models for the EGB are derived from the subtraction of unresolved point sources and extragalactic diffuse processes, which could explain from 30\\% to 70\\% of the Fermi-LAT EGB. Within the hypothesis that the two residual EGBs are entirely due to the annihilation of dark matter (DM) particles in the Galactic halo, we obtain $conservative$  upper limits on their annihilation cross section \\sigmav. Severe bounds on a possible Sommerfeld enhancement of the  annihilation cross section are set as well. Finally, would {\\sigmav} be inversely proportional to the WIMP velocity, very severe limits are derived for the velocity-independent part of the annihilation cross section. ",
        "introduction": "\\label{EGB} The high latitude  ($|b|>10^{\\rm o}$) $\\gamma$-ray emission measured by Fermi-LAT \\cite{LATIsotropicSpectrum}, given its reduced contamination by galactic sources, is a powerful tool to set limits on the contribution of DM to the measured flux. The spectrum has been obtained after the subtraction from the data of the sources resolved by the telescope, the (indeed model dependent) diffuse galactic emission, the cosmic ray (CR) background in the detector and the solar $\\gamma$-ray emission. For each low--flux source there may be a large number of \\textsl{unresolved} point sources which have not been detected because of selection effects, or too low emission.  Most of the unassociated high latitude sources are blazars,  a class of Active Galactic Nuclei (AGNs), and they pile to the EGB with the largest flux \\cite{2010ApJ...720..435A}.  Galactic resolved pulsars and Milli-Second Pulsars (MSPs) represent the second largest population  in the Fermi-LAT catalog \\cite{1FGL,2010ApJS..187..460A} and they are expected to contribute significantly to the putative EGB. A non-negligible $\\gamma$-ray flux seems to be guaranteed by unresolved normal star-forming galaxies \\cite{EGBStarformgal2010}.  Ultra-high energy CRs (UHECRs) may induce secondary electromagnetic cascades, originating neutrinos and $\\gamma$-rays at Fermi-LAT energies \\cite{EGBCMB1}. Unresolved blazars and MSPs are believed to contribute at least few percent to the Fermi-LAT  EGB, while predictions for star-forming galaxies and UHECRs are highly model dependent. In the following, we describe few classes of $\\gamma$-ray emitters whose unresolved flux is firmly estimated in a non-negligible Fermi-LAT EGB percentage. In a conservative scenario (Model I), we will subtract AGN and MSPs to the Fermi-LAT EGB as derived in Ref. \\cite{LATIsotropicSpectrum}. A more relaxed model (Model II) will be drawn by the further subtraction of a minimal flux from star-forming galaxies and CRs at the highest energies. The derivation of each contribution is described in details in \\cite{CDD}.   \\subsection{Models for the EGB} \\label{sec_residui} The aim of this Section is to subtract from the Fermi-LAT EGB \\cite{LATIsotropicSpectrum} additional contributions from unresolved sources at latitudes $|b|>10^{\\rm o}$. The contributions removed  from the Fermi-LAT spectrum are minimal. In fact, the predictions that we will take into account for the relevant unresolved sources are the lowest ones according to the literature. For the Model I, we subtract from the Fermi-LAT EGB \\cite{LATIsotropicSpectrum} the unresolved contributions for both BL Lacs and FSRQs, and  MSPs. Model II for the EGB refers to the scenario where the additional contributions from star-forming galaxies and UHECRs add to explaining the Fermi-LAT EGB. For all the details, the interested reader is addressed to Ref. \\cite{CDD}. The ensuing fluxes are displayed in Fig. \\ref{fig:EGB_Model_II}. At 100 MeV,  Model II explains about $70\\%$ of the Fermi-LAT EGB, while above  1-2 GeV they count about $50\\%$ of the total. \\begin{figure}[h] \\begin{minipage}{20pc} \\includegraphics[width=20pc]{fig2.eps} \\caption{\\label{fig:EGB_Model_II}$\\gamma$-ray spectrum for $|b|>10^{\\rm o}$ latitudes. Fermi-LAT data points are displayed along with their power--law fit (solid black curve) \\cite{LATIsotropicSpectrum}. Dots and long dashed-curve (yellow) correspond to the UHECRs  and star-forming galaxies $\\gamma$-ray fluxes, respectively. Dot-dashed (blue) curve: sum of BL Lacs, FSRQs and MSPs fluxes. Short-dashed (red): sum of all the unresolved components. Solid (green) curve is derived by subtracting all the contributions to the Fermi-LAT result (Model II). } \\end{minipage}\\hspace{2pc}% \\begin{minipage}{20pc} \\includegraphics[width=20pc]{fig3.eps} \\caption{\\label{fig:upper_bounds} Upper bounds on {\\sigmav} from $\\gamma$-ray in the high latitude galactic halo, as a function of the DM mass. From top to bottom, solid lines refer to 90\\% C.L. limits  from the comparison with  Fermi-LAT EGB   (black lines), Model I (red lines), Model II (blue lines) (see text for details). Dotted, solid and dashed lines correspond DM annihilation into $\\mu^+\\mu^-$, $b\\bar{b}$, $\\tau^+\\tau^-$, respectively.} \\end{minipage} \\end{figure} ",
        "conclusions": "The $\\gamma$-ray EGB measured by Fermi-LAT \\cite{LATIsotropicSpectrum} likely includes contributions from galactic and extragalactic $unresolved$ sources. We have discussed two residual EGB fluxes derived by the subtraction of unresolved BL Lacs, FSRQs and galactic MSPs, star-forming galaxies and UHECRs to the $\\gamma$-ray EGB measured by Fermi-LAT. From our new residual EGB fluxes, we have set upper limits on the DM annihilation cross section into $\\gamma$-rays. A conservative upper bound on {\\sigmav} is derived by assuming that our new residual fluxes  are entirely due to WIMPs pair-annihilating in the halo of our Galaxy. Furthermore, we have shown these EGB residuals bound the Sommerfeld enhancement of {\\sigmav} to a factor of 3-10-50-200 for $m_\\chi$=10-100-1000-5000 GeV, respectively. In case of a Yukawa-like potential, a force carrier heavier than 1 GeV is required. Finally, within the hypothesis that {\\sigmav} is inversely proportional to the WIMP velocity, very severe limits are derived for the velocity-independent part of the annihilation cross section."
    }
}